{
    "1004/1004.3907_arXiv.txt": {
        "abstract": "In recent years a large number of Hot Jupiters orbiting in a very close orbit around the parent stars have been explored with the transit and doppler effect methods. Here in this work we study the gravitational microlensing effect of a binary lens on a parent star with a Hot Jupiter revolving around it. Caustic crossing of the planet makes enhancements on the light curve of the parent star in which the signature of the planet can be detected by high precision photometric observations. We use the inverse ray shooting method with tree code algorithm to generate the combined light curve of the parent star and the planet. In order to investigate the probability of observing the planet signal, we do a Monte-Carlo simulation and obtain the observational optical depth of $\\tau \\sim 10^{-8}$. We show that about ten years observations of Galactic Bulge with a network of telescopes will enable us detecting about ten Hot Jupiter with this method. Finally we show that the observation of the microlensing event in infra-red band will increase the probability for detection of the exo-planets. ",
        "introduction": "Gravitational microlensing as one of the applications of general relativity is proposed by Paczy\\'nski (1986) for detecting the dark compact halo objects so-called MACHOs in the Galactic halo. While not enough MACHOs have been detected in the halo \\cite{mil}, however the microlensing technic has been used as an astrophysical tool for studying the atmosphere of the stars and exploring the exo-planets. In the standard method for exploring planets with the microlensing, a star with the companion planets can play the role of lens and produce caustic lines where crossing the caustics by the source star produces a high magnification on the light curve \\cite{pac91}. In this case in addition to the standard microlensing light curve we can detect a short duration spark due to the caustic crossing, formed by the planet. A precise photometry of the event is essential to find out this short duration signature of the planet. The advantage of this technic compare to the other methods of the exo-planet detection is that it is sensitive to the observation of earth mass planets \\cite{bal05} and also those planets located beyond the snow line \\cite{gould01}. There is also other methods in gravitational microlensing such as planetary microlensing signals from the orbital motion of the source star around the common barycenter of source star--planet system \\cite{rahvar}.  In addition to the mentioned methods, Graff \\& Gaudi (2000) proposed caustic crossing of a close-in Jupiter size planet, produced by a binary lens. In this case the planet's light is magnified so much that it can be detected by a 10-m class telescope. Here we extend this work looking to the details of the light curves and study the most favorite pass band for this observation. Since in a close-in Jupiter, the thermal emission due to the high temperature of the planet is more significant than the reflected light from the parent star, the observations in the Infra-red pass band is more favorable than the visual pass band. We also do a Monte-Carlo simulation with a given observational strategy to obtain the number of observable events in terms of the parameters of the planet and the parent star. We emphasize that while the observations of the hot Jupiters is simpler in nearby stars via the eclipsing and doppler methods, the microlensing method can detect distant systems and enable us to compare the statistics of the hot Jupiters with the nearby observations. One of the interesting features of the light curve for the planet caustic crossing is that the planet can cross the caustic more than that of parent star, as it traces effectively a longer path due to the revolving motion around the parent star.  The paper is organized as follows. In Section \\ref{lightcurve} we will introduce the caustic crossing of the parent star and planet system and generate light curve with inverse ray shooting technic, introducing a new development in tree-code algorithm. In section \\ref{char} we study the characteristics of the light curve in terms of the orbital parameters of the planet and the parent star. In section \\ref{mc} we explain our Monte-Carlo simulation for estimating the probability of illuminating Hot Jupiters with this method. In section \\ref{conc} we give the conclusions. ",
        "conclusions": "\\label{conc} In this work we examined the possibility of hot Jupiter detection through caustic crossing of a binary lens by a planet as the source object. The effect of this caustic crossing due to the small size of the planet is like an illumination on the microlensing light curve of the parent star. Taking the flux of the parent star as the background light and the illumination of the planet as the signal, we studied the physical characteristics of the planet as the orbital size, the atmospheric property and the size of the planet from one hand and the characteristics of the parent star and lens from the other hand on the observability of planet.  In the next step, we did a Monte-Carlo simulation to obtain the detection efficiency of the planet with this method. We showed that takeing just geometrical caustic crossing of the planets, $36\\%$ of population can be illuminated. However in reality due to the photometric error the peak generated by the planet may not be detected. We used three different photometric precision in the observation and obtained the detection efficiency, assuming that we use a network of the telescopes to have enough sampling of data much less than one hour, a typical time of the caustic crossing of the planet. We showed that for longer pass-bands the detection efficiency is increased due to the more relative emissivity of planet compare to the parent star.   Finally we estimate the number of the Hot Jupiters that can be observed with this method towards the Galactic bulge. With a ten years monitoring of $10^7$ stars towards the Galactic Bulge, we can detect in the order of $10$ Hot Jupiter with this method. This observation may be done by the next generation of the microlensing surveys towards the Galactic Bulge."
    },
    "1004/1004.1902_arXiv.txt": {
        "abstract": "{Because of new telescopes that will dramatically improve our knowledge of the interstellar medium, chemical models will have to be used to simulate the chemistry of many regions with diverse properties.  To make these models more robust, it is important to understand their sensitivity to a variety of parameters.   } {In this article, we report a study of the sensitivity of a chemical model of a cold dense core, with homogeneous and time-independent physical conditions, to variations in the following parameters: initial chemical inventory, gas temperature and density, cosmic-ray ionization rate,  chemical reaction rate coefficients, and elemental abundances.} {We used a Monte Carlo method to randomly vary individual parameters and groups of parameters within realistic ranges. From the results of the parameter variations, we can quantify the sensitivity of the model to each parameter as a function of time.    Our results can  be used in principle with observations to constrain some parameters for different cold clouds.  We also attempted to use the Monte Carlo approach with all parameters varied collectively.} {Within the parameter ranges studied, the most critical parameters turn out to be the reaction rate coefficients at times up to $4 \\times 10^{5}$ yr and elemental abundances at later times.  At typical times of best agreement with observation, models are sensitive to both of these parameters.  The models are less sensitive to other parameters such as the gas density and temperature.       } {The improvement of  models will require that the uncertainties in rate coefficients of important reactions  be reduced.  As the chemistry becomes better understood and more robust, it should be possible to use model sensitivities concerning other parameters, such as the elemental abundances and the cosmic ray ionization rate, to yield detailed information on cloud properties and history. Nevertheless, at the current stage, we  cannot determine the best values of all the parameters simultaneously based on purely observational constraints.} ",
        "introduction": "Molecules are powerful observational tools for the study of the physical conditions and dynamics of star-forming regions. Each step of the formation of a stellar and planetary system is characterized by a chemical composition, which directly reflects the physical conditions of the medium and its evolutionary stage. The far-infrared space telescope Herschel and (sub)-millimeter interferometer ALMA promise to open new windows on the wealth of information found in the interstellar medium (ISM).  With the improvement of observational instrument sensitivity and resolution,  and the opening of new wavelength ranges, more molecules will be detected in the ISM and many molecules will be detected under a wider variety of physical conditions.  The more molecules detected and the more diverse physical conditions  discovered, the more complex chemical models will have to be to reproduce observations.  The results of chemical models, or simulations,  depend on a number of parameters or groups of parameters often poorly constrained. In this paper, we will be concerned with the sensitivity of simple gas-phase chemical simulations of cold dense interstellar cores to these parameters.  The models are based on the  pseudo-time-dependent approximation, in which the chemistry evolves under fixed and homogeneous density and temperature.  Homogeneous models are also referred to as 0D (zero-dimensional).   Although more physically reasonable models should include the formation of cold cores from more diffuse material as the chemistry progresses, such models are rare given the lack of unanimity of how the collapse proceeds.  Among the parameters in pseudo-time-dependent/0D models determined only to a limited extent by observations or their comparisons with model results are (1) the gas kinetic temperature and density, (2) the cosmic-ray ionization rate,   (3) the elemental gas-phase abundances, and (4) the initial chemical inventory.    We add to these parameters the rate coefficients (5) of the many chemical reactions despite the fact that they are best determined by laboratory experiments rather than by observational constraints. We discuss the parameters in turn.  \\begin{enumerate}  \\item The gas kinetic temperature and density are usually determined by the excitation conditions of observed molecular emission lines.  In addition to the uncertainties in the radiative transfer analysis of the lines, inhomogeneities of the gas may exist along the same line of sight.  \\item There is no direct way to measure the cosmic-ray ionization rate $\\zeta$ in dense clouds.  There are two different approaches to constrain $\\zeta$.  The first one, usually used, is to determine it by comparison of modeled and  observed abundances for specific ionic species or their ratios.  Currently, there appears to be a dichotomy between diffuse clouds, in which the high abundance of H$_{3}^{+}$ leads to a large value of $\\zeta$, and dense clouds, in which the ionization conditions often lead to a value of $\\zeta$ 1-2 orders of magnitude lower  \\citep{1982ApJ...255..160W,1998ApJ...499..234C,2003Natur.422..500M, 2004A&A...417..993L,2009ApJ...694..257I}.  There is some indication that the difference is simply due to the inability of low energy cosmic rays to penetrate into the center of dense clouds \\citep{2009A&A...501..619P}.  A second approach is to use the value measured in the solar system \\citep{1968ApJ...152..971S,1998ApJ...506..329W} and consider that $\\zeta$ is constant in the Galaxy. However, in the solar system, the low energy end of the cosmic ray spectrum cannot be determined directly because of the solar wind.  \\item The computed abundances of atomic and molecular species depend on the choice of elemental abundances in the gas phase. Elements are measured in absorption in the diffuse medium \\citep{2004oee..symp..336J}. In order to reproduce the observed gas-phase composition of dense clouds, it is often assumed that heavy atoms, including S, Si, Fe, Na and Mg, are depleted further from the gas between the diffuse (or translucent) and dense phases of a cloud life. Specifically,  the initial abundances of these elements in pseudo-time-dependent models are taken to be 2 to 3 orders of magnitude lower than in diffuse clouds but the efficiency of the depletions onto the grains is quite uncertain \\citep{1982ApJS...48..321G}.  Moreover, the depletions of the elements may depend on the depth into the molecular cloud.  \\item Since the $t=0$ time of pseudo-time-dependent models is artificial, the initial chemical inventory chosen can be artificial as well.  Chemical models usually start with all the elements in the atomic form, except for H$_{2}$, assuming that the density of the gas had suddenly increased from a diffuse to a dense medium.   Shock models of the conversion of diffuse to dense gas show, on the other hand, that the process is slow and that a significant amount of CO may be synthesized before a sizable dense cloud core is produced \\citep{2004ApJ...612..921B}.   Preliminary results show that the formation of complex molecules is far slower if the starting form of carbon is CO (Hassel, Herbst, \\& Bergin, in preparation). Thus, the initial chemical abundances and the cloud age, as determined by comparison of observational and model abundances, are correlated.  \\item Finally, the chemistry in the dense ISM involves thousands of gas-phase reactions, especially when one is dealing with complex molecules. Only a small fraction of these reactions have been studied at the low temperatures  present in cold dense cores.  Thus, even for systems studied in the laboratory at higher temperatures, there can be a difficulty in extrapolating results to significantly lower temperatures.   The most important reactions at assorted cloud ages have been discussed in previous papers with our sensitivity approach and those of others \\citep{2004AstL...30..566V,2005A&A...444..883W,2006A&A...451..551W,2008ApJ...672..629V}.  \\end{enumerate}  In order to quantify the reliability of the chemical models and to determine how best to improve them, we need to know how sensitive they are to these parameters so that we can reduce the uncertainties in the more important ones.  For example, if it is found that model results are more sensitive to current uncertainties in chemical rates than to elemental abundances, then future measurements of such rates should rank at high priority.  If the opposite is true, priority might be given to developing better dynamic models of clouds to follow the depletion of elements from the gas more appropriately.   Moreover, if the models are compared with well-known cold clouds such as TMC-1 and L134N, the observations can be used to further constrain some of the model parameters within their computed uncertainties.  In this article, we present sensitivity analyses of gas-phase cold dense cloud chemistry to all of the parameters discussed above.   Details on the model are given in Section 2, while the sensitivity method and the ranges of values over which the parameters are varied are described in Section 3. The results of these analyses are discussed in Section 4 while the constraints that observations of specific cold cores can yield concerning the model parameters are presented in Section 5.  Finally, Section 6 contains a discussion.  In our sensitivity calculations, one can proceed via the independent variation of parameters or groups of parameters, or one can attempt to vary all parameters collectively.  Although the latter approach is attempted, as discussed below, most of our useful results stem from the former treatment. ",
        "conclusions": "We have studied the sensitivity of the chemistry of cold dense cores to an assortment of parameters using the 0D gas-phase model Nahoon \\citep{2005A&A...444..883W} in which chemical abundances evolve from given initial abundances. As variable parameters, we considered the temperature and density of the gas, the cosmic-ray ionization rate, the elemental abundances, the reaction rate coefficients, and initial concentrations.  These parameters {or groups of parameters} were all varied within realistic ranges of values either individually or simultaneously using Monte Carlo methods \\citep{2004AstL...30..566V,2005A&A...444..883W,2006A&A...451..551W,2008ApJ...672..629V} used previously mainly to study the sensitivity of molecular abundances to chemical reaction rates.  For the case of the elemental abundances, our variations were carried out randomly from a base of abundances based on a new set first defined in \\citet{2008ApJ...680..371W} as appropriate for cold dense clouds.  Our results depend to some extent upon whether individual parameters are varied while others are held fixed at standard values or whether all parameters are varied simultaneously. Nevertheless, certain trends can be found.   Among the most important is the finding that, when averaged over all molecules,  the dominant source of uncertainty in predicted molecular abundances for times less than $4\\times 10^5$~yr is the uncertainty in rate coefficients, whereas the dominant uncertainty at later times is caused by uncertainties in elemental abundances.  The former source of uncertainty can be reduced by improved laboratory and theoretical determinations of rate coefficients at appropriate temperatures, especially when sensitivity analyses point out the specific reactions of greatest importance, as done in \\citet{2009A&A...495..513W}.  The latter source of uncertainty is astrophysical in origin, and suggests that more attention be paid to a physical understanding of how gas-phase elemental abundances evolve as dense clouds are formed from more diffuse gas.     Within the ranges shown in Table \\ref{param}, the sensitivity to the other parameters is generally less. As long as the temperature and density of the gas uncertainties lie within their listed ranges, the molecular abundances predicted by the gas-phase model are quite robust; i.e., they do not depend much on $T$ and $n_{\\rm H}$. The only exceptions are the N-bearing species containing only nitrogen and hydrogen, such as NH$_3$, which are produced much less efficiently at temperatures lower than the standard value of 10 K.  Many molecular abundances depend on the value of the cosmic-ray ionization rate, either at all times or only  between cloud ages of  $10^4$ and $10^7$~yr.  An example of the former is OH and one of the latter is  HC$_{3}$N, as shown in Fig.~\\ref{corr_zeta}.  Of all the parameters, the initial concentrations with the chosen elemental abundances seem to be the parameter to which model results at all reasonable times are least sensitive.   Nevertheless, we found that cloud chemistry has to start with a significant fraction of the carbon in the atomic form and a very small amount of the hydrogen in atomic form to synthesize complex molecules.   Our standard model meets these requirements, but it may be in conflict with with some recent 21 cm studies by \\cite{2009AAS...21348509K}, which show a larger amount of atomic hydrogen.  The range of atomic hydrogen abundances used when the concentrations are varied differs strongly from the standard model in which all hydrogen starts in its molecular form.  It is thus critical to understand both the physics and chemistry of the stages of collapse leading to the formation of cold cores for this problem as well as the elemental abundance problem. A new treatment of the gas-grain chemistry of cold cores in the process of formation is being prepared by Hassel, Herbst, \\& Bergin.   In addition to the study of model sensitivity to assorted parameters, we have utilized the simultaneous variations of these parameters to attempt to determine their optimum values by comparison of calculated and observational results for the cold cloud cores TMC-1 and L134N.  In this paper, we used an indicator of agreement that minimizes the sum of the absolute values of the difference of the logarithms of the calculated and observed abundances. Constraints on individual model parameters are highly sensitive to the values of the other parameters (especially reaction rate coefficients) when variations are run simultaneously. It is currently not possible to constrain all the parameters by comparison with observed abundances by  rigorous methods so that efforts have to be made to reduce the uncertainties for some of the parameters. Uncertainties in rate coefficients for instance can be reduced by laboratory measurements or theoretical calculations on reactions clearly identified as quantitatively important for the model predictions.   Finally, although the strength of sensitivity methods for gas-phase chemical models has been clearly demonstrated in this and earlier studies for both cold clouds and more complex objects such as protoplanetary disks \\citep{2008ApJ...672..629V},  the methods will need to be generalized to estimate the sensitivity to both chemical processes on grain surfaces and adsorption and desorption processes. The use of sensitivity methods for surface chemistry may become quite feasible within the next decade given the rapid pace of advance of both theory and experiments in this complex field."
    },
    "1004/1004.3070_arXiv.txt": {
        "abstract": "We present the first extensive photometric results of CL Aur from our $BVRI$ CCD photometry made on 22 nights from 2003 November through 2005 February. Fifteen new timings of minimum light were obtained. During the past 104 years, the orbital period has varied due to a periodic oscillation superposed on a continuous period increase. The period and semi-amplitude of the oscillation are about 21.6 yrs and 0.0133 d, respectively.  This detail is interpreted as a light-travel-time effect due to a low-luminosity K-type star gravitationally bound to the CL Aur close system. Our photometric study indicates that CL Aur is a relatively short-period Algol-type binary with values of $q$=0.602 and $i$=88$^\\circ$.2. Mass transfer from the secondary to the primary eclipsing component is at least partly responsible for the observed secular period change with a rate of $dP$/$dt$ = +1.4$\\times$10$^{-7}$ d yr$^{-1}$. A cool spot model has been calculated but we think that an alternative hot-spot model resulting from a gas stream impact on the hot star is more reasonable despite two difficulties with the explanation. Absolute dimensions of the eclipsing system are deduced and its present state is compared with tracks for single star and conservative close binary evolution. Finally, we examine the possible reconciliation of two different calculations of the luminosity of the hot spot and a re-interpretation of the secular term of the period variability. ",
        "introduction": "Algol-type close binaries are semi-detached interacting systems in which one type of interaction is mass transfer between the component stars by means of a gas stream. They have been known as good astrophysical laboratories for studying accretion processes because a number of them are bright.  They are in the slow phase of mass transfer with $dM/dt \\simeq 10^{-11}-10^{-7}$ M$_\\odot$ yr$^{-1}$ and do not undergo violent eruptions that interfere with the accretion process. The circumstellar structures produced by the mass-transfer process in these systems have been sorted according to orbital period by Richards \\& Albright (1999) but do not depend upon it significantly. Rather, their natures can be easily understood from the position of the mass-gaining component in the so-called $r$-$q$ diagram in which the fractional radius $r$ = ($R/a$) of a gainer is plotted {\\it versus} the mass ratio $q$ and compared with the semianalytical computations of the gas stream hydrodynamics of Lubow \\& Shu (1975). In the short-period Algols located above the $\\omega_d$ curve of the diagram (cf. Figure 2 of Richards \\& Albright), the hot, detached primary star is large relative to the orbital radius and the two components are too close to each other to form an accretion disk or even a stable accretion annulus. Instead, it is possible that an impact region, and hence a hot spot, can be formed on the surface of the primary star somewhat displaced from the line of centers due to the Coriolis acceleration imposed on the flowing gas. If the secondary stars are sufficiently cool, they likely display enhanced magnetic activity due to deep outer convective layers and rapid rotation. This magnetic mechanism may contribute to the period and light variations for systems with spectra later than F-type (Hall 1989).  CL Aur (GSC 2393-1455, HV 6886, TYC 2393-1455-1) was discovered to be a variable star by Hoffleit (1935) based on photographic plate estimates. Kurochkin (1951) presented the first (partial) photographic light curve of the star and the original light elements, Min. I = HJD 2,432,967.262 + 1.2443666$E$. The value of the period positions this object toward the short-period limit for Algols. The spectral type of the primary star was classified to be A0 by G\\\"otz \\& Wenzel (1968). Since then, times of minimum light have been published assiduously by numerous workers but, to our knowledge, a complete light curve and the fundamental parameters for the binary system have not been made so far. Changes of the orbital period have been considered by Heged\\\"us (1988) and Wolf et al. (1999). Heged\\\"us selected this system as a possible candidate for the study of apsidal motion. However, the later authors ruled out this possibility from CCD timings for primary and secondary eclipses. They suggested the cause of period variation to be a light-travel-time (LTT) effect due to the presence of a third body in the binary system.  Most recently, Wolf et al. (2007, hereafter W07) reported that a long-term period increase is superimposed on an LTT orbit with a period of $P_3 $=21.7 yrs, a semi-amplitude of $K$=0.014 d, and an eccentricity of $e$=0.32.  In the Simbad data base\\footnote {http://simbad.u-strasbg.fr/simbad/}, the system is described as an eclipsing binary of $\\beta$ Lyr type. $BVJHK$ magnitues are listed for the star but these are from heterogeneous sources and are not mutually consistent. Part of this inconsistency arises because the magnitudes refer to different phases in the Keplerian cycle. With the well-known transformations (ESA 1997), standard photometric values for CL Aur in the Johnson system were calculated to be $V=+$11.62 and $(B-V)=+$0.33 from the Tycho results. These refer to some unknown Keplerian phase. These are not consistent with those in Simbad presumably because of the large phase-locked variations in magnitude and color index of the binary.  At present, CL Aur is known only as a neglected eclipsing system composed of an A-type primary and a cooler companion. In order to derive photometric solutions and to examine whether the W07 suggestion is appropriate for the orbital period change, we decided to obtain light curves with multiband photometry. In this paper, we present the first mutual analyses of the $O$-$C$ diagram and the light curves. ",
        "conclusions": "We next probed the credibility of the hot (inferentially an impact) spot model motivated largely by two recognitions: (a) the two arrays of light curve residuals in Figure 4 do not appear very different one from another although formally the light curve fitting is improved by the assumption of a spot and (b) the center of the spot is sensibly distant from the stellar equator and orbital plane.  The first matter is the consequence of the envelope of the residuals responding weakly to the spot modeling although the residuals of small absolute value did respond well with considerably reduced values on average.  This can only be construed as a situation in which there exist residuals of noise much greater than most of the rest of them. In principle, these residuals could be produced by a combination of physical causes such as magnetic activity from the cool secondary and a change in the mass transfer rate. However, the actual noise level is not significantly larger than that ($\\pm$0.01 mag) of SOAO data made during the last few observing seasons and probably can be traced to marginal observing conditions on some nights.  The location of the spot has more physical interest.  First, an extensive set of restricted 3-body calculations examined the longitude coordinate of the spot on the hot primary for a range (0.0000003 km s$^{-1}$ -- 450 km s$^{-1}$) of initial velocities streaming from the L$_1$ point as shown in the first column of Table 9. This range obviously includes the thermal velocity in the cool star envelope with the lower limit of the range essentially that of free-fall. For the greatest extent of this velocity range the impact spot is close (within 21$^\\circ$) to the line of centers. In the hydrodynamic computations of Lubow \\& Shu (1975), just such an effect appears with the streaming gas deflected about 20$^\\circ$ from the line of centers, not nearly enough to avoid contact with the large, detached primary. From this point of view, the credibility of the concept of impact is therefore high, and conservative mass transfer of the streaming material must be considered likely.  These calculations also make it difficult to imagine any stable accretion structures around the hot star but they invariably led to an orbital-plane spot rather than the modeled one of co-latitude very different from 90$^\\circ$. Another suite of calculations, also given in Table 9, showed the expected result: inital non-zero-velocity $z$-components in the streaming gas led to impacts away from the hot stellar equator and these did not have to be large in order to fall 15--20$^\\circ$ away from the equator. Is this result to be taken as evidence that the spot really exists at its modeled location or is it just a dynamical truism and doesn't verify the spot existence at all?  The resolution of this quandary requires a mechanism to move the gas appropriately and three possible ones come to mind.  (a) Turbulence in the gas moving from the L$_1$-point caused a concentration of it to fall at the modeled position during the two observing seasons. (b) A weak magnetic field seated in the cool star had at least one component channeling the ionized fraction of the streaming gas to the modeled impact spot. (c) The solution for the spot is possibly not so unique as W-D indicates. For want of independent evidence, each of these is an unprovable hypothesis.  More information does exist, however.  We calculated the impact luminosity from the stream at the rate of mass transfer given by the interpretation of the $O$--$C$ diagram.  For the same range (and in the same sense as above) of stream velocity, the impact luminosity varied between 3.2 $L_\\odot$ and 0.7 $L_\\odot$. If the impact energy is partitioned between virialization into the gainer star and spot luminosity as Hilditch (1989) proposes, the least luminous output from a spot should therefore be about 0.35 $L_\\odot$. This is to be compared to the spot luminosity from its black body W-D model of about 4$\\times10^{-5}$ $L_\\odot$.  The magnitude of the discrepancy means either (a) that conversion from kinetic to luminous energy is very inefficient or (b) that the mass-motions rate indicated by the $O$--$C$ diagram is not restricted to material leaving the L$_1$ point or (c) some combination of these two ideas or (d) that the concept of a spot is itself incorrect.  For (a) to be a realistic interpretation, deep penetration of the impacting gas into the hot star is required so that most of the impacting gas becomes thermalized in the hot star. It is also necessary to postulate an envelope circulation pattern that re-surfaces at the modeled spot position for the residual gas that has not been thermalized.  If (b) is to be entertained seriously, there must be substantial mass lost by evaporation from the entire Roche lobe that is the photosphere of the cool star.  This demands only a small ($\\sim 10^{-4}$) contribution to the $O$--$C$ diagram from the gas leaving L$_1$ with the majority of the mass and angular momentum loss contributed by that departing the rest of the cool star's surface. On a small scale the transfer would be conservative as usual but large scale mass loss would be the dominant cause of the $A$-term in the ephemeris. Hypothesis (d) would be the most economical interpretation but conflicts with the impersonal syntheses of the light curves. At this time, we favor a combination of interpretations (a) and (b).  In summary, our study of the orbital period and the light curves reveals that CL Aur is a classical Algol-type interacting system with the less massive and cool secondary star filling its inner Roche lobe. The possibility of a hot-spot model due to impact of streaming gas onto the hot star has led to confusing difficulties that we have resolved only tentatively. High-resolution, near-IR spectroscopy should reveal the lines of the secondary star and lead to accurate absolute parameters to replace our estimates. Moreover, there is also the possibility of obtaining direct evidence of mass-transfer activity, such as complex and variable line profiles of H$\\alpha$ or variations in the strength of the O I 7774 absorption line. The evolutionary status of the system will then be more convincingly in hand.  The more general reality of things is that there are many short-period Algols which have not been studied so deeply as this present work has examined CL Aur. It would make a significant advance if more of these - having different algebraic signs for the secular term of the period variability - could be brought to the same level of knowledge as the present binary. For instance, should all such binaries with $A >$ 0 require a spot on the hot star near the line of centers, there would be major support for the reasoning concerning mass movements.  Should it also happen that the spot luminosity was consistently found to be lower than the kinetic impact required for 50\\% energy conversion, there would be good reason to believe that the majority of the mass lost from the cool secondaries is, in fact, lost to the systems and not transferred conservatively. The short-period binaries found to require mode 4 (semi-detached systems with the primary stars filling its inner Roche lobes) for their representations would be expected, then, to show small-scale mass transfer to the secondaries and systemic mass loss from the  primaries. Much valuable observational and modeling work remains."
    },
    "1004/1004.1243_arXiv.txt": {
        "abstract": "We discuss  outer gap closure mechanism in the trans-field direction with the magnetic pair-creation process near the stellar surface. The gap closure by the magnetic pair-creation is possible if some fraction of the pairs are produced with an outgoing  momentum. By assuming that  multiple magnetic  field will affect the local field near the stellar surface,  we show a specific magnetic field geometry near the stellar surface resulting in the outflow of the pairs. Together with the fact that the electric field is weak below null charge surface, the characteristic curvature photon energy emitted by incoming particles, which were accelerated in the outer gap, decreases drastically to $\\sim 100$MeV near the stellar surface. We estimate the height measured from the last-open field line, above which 100~MeV photons is converted into pairs by the magnetic pair-creation. We also show the resultant multiplicity due to the magnetic pair-creation process  could acquire  $M_{e^{\\pm}}\\sim 10^4-10^5$. In this model the fractional outer gap size is proportional to $P^{-1/2}$.  The predicted gamma-ray luminosity ($L_{\\gamma}$) and the characteristic curvature photon energy ($E_c$) emitted from the outer gap are proportional to $B^2P^{-5/2}$ and $B^{3/4}P^{-1}$ respectively.  This model also predicts that  $L_{\\gamma}$ and $E_c$ are related to the spin down power ($L_{sd}$) or the  spin down age of pulsars ($\\tau$) as $L_{\\gamma} \\propto L_{sd}^{5/8}$ or  $L_{\\gamma} \\propto \\tau^{-5/4}$, and  $E_c \\propto L_{sd}^{1/4}$ or $E_c \\propto \\tau^{-1/2}$ respectively. ",
        "introduction": "The mechanism of  particle acceleration and high-energy emission processes in the pulsar magnetospheres are one of the unresolved physics of the pulsar activities. The particle acceleration process and resultant high-energy $\\gamma$-ray  emission process have been discussed with the polar cap model (Ruderman \\& Sutherland 1975; Daugherty \\& Harding 1982), the slot gap model (Arons 1981; Harding, Usov \\& Muslimov 2005; Harding et al. 2008) and the outer gap  model (Cheng, Ho \\& Ruderman 1986a,b; Hirotani 2008; Takata \\& Chang 2009). The polar cap model assumes the emission site  close to the stellar surface above the polar cap, and the slot gap and the outer gap models assume the emission site in the outer magnetosphere. The different acceleration  models have predicted the different properties of the $\\gamma$-ray emissions from the pulsar magnetospheres.  The observations of the pulsar  emitting electromagnetic radiation in  the high-energy $\\gamma$-ray bands have been facilitated by recent space and ground based telescopes.    In  particular,  the $Fermi$ $\\gamma$-ray telescope has measured the  $\\gamma$-ray emissions from  $\\sim 46$  pulsars (Abdo et al. 2010, 2009a,b), including 21 radio-loud,  17 radio-quiet and 8 millisecond  pulsars. In addition, $AGILE$ (Astro-rivelatore Gamma a Immagini LEggero) has also reported the detection of $\\gamma$-ray emissions from 4 new pulsars with 4 candidates (Pellizzoni et al. 2009). In more higher energy regime, $MAGIC$ (Major Atmospheric Gamma Imaging Cherenkov) telescope has detected for the first time pulsed gamma-ray radiation above 25~GeV from the Crab pulsar  (Aliu et al. 2008). These observations will be useful to discriminate between the emission models.  For example, the $Fermi$ telescope  has measured the  spectral properties  above 10~GeV  with a  better sensitivity than $EGRET$. It was   found that the spectral shape of $\\gamma$-ray emissions from the Vela pulsar is well fitted with a power low (photon index $\\Gamma\\sim$1.5) plus exponential cut-off ($E_{cut}\\sim 3$~GeV) model. The discovered exponential cut-off feature predicts that  the emissions from the outer magnetosphere  (Abdo et al. 2009c) is more favored   than the polar cap region (Daugherty \\& Harding 1996), which predicts a super exponential cut-off with the magnetic pair-creation. Furthermore, the detection of the radiation above 25~GeV bands associated with the Crab pulsar has also predicted the high-energy emission in  the outer magnetosphere (Aliu et al. 2008).   The pulse profiles observed by the $Fermi$ telescope allow us to study the site of the $\\gamma$-ray emissions in the pulsar magnetosphere. Venter et al. (2009) fitted the pulse profiles of the 8 millisecond pulsars with the geometries predicted by the different emission models. They showed that most of the pulse profiles  can be best fit with the outer gap   (Takata et al. 2007; Takata \\& Chang 2007; Tang et al. 2008)  or the two pole caustic (Dyks \\& Rudak 2003; Dyks et al.  2004) geometries, which have a slab like  geometry along the last-open field lines. However, they also found that the pulse profiles  of  two out of eight   millisecond pulsars cannot be fitted by either the  geometries with the outer gap or  the caustic models.  They proposed a pair-starved polar cap model, in which the multiplicity of the pairs is not high enough to completely screen the electric field above the polar cap,  and the particles are continuously accelerated up to high altitude over full open field line region.  The increase of the $\\gamma$-ray pulsars allows us to perform a detail  statistical study of the $\\gamma$-ray pulsars. In particular,  the $Fermi$ $\\gamma$-ray pulsars including millisecond pulsars will reveal the relation between the $\\gamma$-ray luminosity ($L_{\\gamma}$) and the spin down power ($L_{sd}$), for which $L_{\\gamma}\\propto L_{sd}^{\\beta}$ with $\\beta\\sim 0.5-0.6$ was predicted by $EGRET$ measurements (Thompson 2004). Also, the $Fermi$ $\\gamma$-ray pulsars will enable us to discuss the general trend of the relation among the spectral properties of the $\\gamma$-ray emissions (e.g. the cut-off energy and photon index)  and the pulsar parameters (e.g. rotation period and surface magnetic field). Together with the observed pulse profiles and the spectra, these  general properties of the $\\gamma$-ray emissions will discriminate between the $\\gamma$-ray emission models in the pulsar magnetospheres.  In this paper, we discuss the $\\gamma$-ray emissions from the outer gap accelerator. We propose a new outer gap closure mechanism by the magnetic pair-creation process near the stellar surface. The pairs produced by the magnetic pair creation will be able to close the gap  if the sufficiently strong surface multiple field exists and  affects  the dipole field near the surface. In section~\\ref{closure},  we first summarize results of gap closure process  by   photon-photon pair-creation process, and  then we discuss our  new gap closure mechanism by  the magnetic pair-creation process. In section~\\ref{outergap}, we describe the model predictions of the properties of the $\\gamma$-ray emissions.  In section~\\ref{discussion}, we will compare the model predictions with the results of the  $Fermi$ observations. We also discuss applicability of  our model. In section~\\ref{conclusion},  we will summarize our gap  closure model and predictions for the outer gap accelerator. ",
        "conclusions": "\\label{conclusion} In this paper, we have studied  the outer gap accelerator model closed  by magnetic pair-creation process. We argued  that  below null charge surface, the curvature loss is not compensated by the acceleration due to the electric field in the gap. In such a case, the incoming particles, which were produced in the outer magnetosphere,   will emit curvature photons with about $E_{min}\\sim m_ec^2/\\alpha_{f}\\sim 100$~MeV. The 100~MeV curvature photos propagating toward the stellar surface will be converted into pairs by the pair-creation process with the strong local magnetic field near the stellar surface, where the multiple magnetic field affects to the global field lines.  For the canonical pulsar, the synchrotron radiation of the created pairs produce  $\\sim$5~MeV photon, which will be furthermore converted into pairs.  As a result, multiplicity of an incoming particle could acquire $M_{e^{\\pm}}\\sim 10^4-10^5$.  With the local field lines bending away from the last-open field line (such as illustrated in Figure~\\ref{Pulsar}),  the created pairs via the magnetic pair-creation process can have the outgoing  momentum and migrate into the outer magnetosphere. If $\\sim 10$ pairs out of  $M_{e^{\\pm}}\\sim 10^4-10^5$ migrate  outward to the outer magnetosphere, those pairs could close  the outer gap accelerator. According to this scenario, the main results of this paper are as follows. The fractional thickness of the outer gap becomes $f_m\\sim 0.25K P^{1/2}_{-1}$, which has a less dependency on the rotational period compared with the outer gap model proposed by Zhang \\& Cheng (1997).  With the present model, the spectral properties of the $\\gamma$-ray emissions depend on the local parameter $K\\sim \\chi_{-1}^2B^{-2}_{m,12}s_7(R_s/R_i)^{3/2}$, which is determined by the local magnetic structure near the star. We expect that the local parameter $K$ takes a vale of $K\\sim 2$  for  the canonical pulsars and $K\\sim 15$ for the millisecond pulsars (Figures~\\ref{cutoff} and~\\ref{lumin}). The present model predicts that the canonical pulsars and the millisecond pulsars are connected in the plots of $L_{\\gamma}/E_{c}$ versus $B_{s}^{1/2}P^{-1/2}$, in which the effect of the fractional gap thickness is carried away (Figure~\\ref{basic}).  The present model predicts that the cut-off energy ($E_c$) and the $\\gamma$-ray luminosity ($L_{\\gamma}$) depend on the spin down age or the spin down power  as $E_{c}\\propto \\tau^{-1/2}$ and $L_{\\gamma}\\propto \\tau^{-5/4}$ or  $E_{c}\\propto L_{sd}^{1/4}$ and $L_{\\gamma}\\propto L_{sd}^{5/8}$ (Figures~\\ref{age-cutoff}-\\ref{spin-lumi}).   In addition to the cut-off energy and $\\gamma$-ray luminosity, which have been discussed in this paper,  the $Fermi$ $\\gamma$-ray telescope provides  the photon index of $\\gamma$-ray spectrum  and the pulse profiles for each pulsar (Abdo et al. 2010). It must be  important to discuss the photon index and the pulse profile with the acceleration model, because they will contain information of the electric structure (e.g. the distribution of the electric field) in the acceleration region and the three-dimensional geometry of the emission region (e.g. Romani \\& Yadigaroglu 1995; Cheng, Ruderman \\& Zhang 2000; Spitkovsky 2006). However, a more  detail  model, which has to consider the electrodynamics  in the gap and three-dimensional structure, is required to study  the shape of the $\\gamma$-ray spectra and the pulse profiles.  Studying the emission properties with the electrodynamics in the present outer gap closure model will be done in the subsequent papers."
    },
    "1004/1004.2595_arXiv.txt": {
        "abstract": "{ The High Frequency Instrument of Planck will map the entire sky in the millimeter and sub-millimeter domain from 100 to 857~GHz with unprecedented sensitivity to polarization ($\\Delta P/T_{\\mbox{\\tiny cmb}} \\sim 4\\cdot 10^{-6}$ for $P$ either $Q$ or $U$ and $T_{\\mbox{\\tiny cmb}}\\simeq 2.7\\,$K) at 100, 143, 217 and 353 GHz.  It will lead to major improvements in our understanding of the Cosmic Microwave Background anisotropies and polarized foreground signals. Planck will make high resolution measurements of the $E$-mode spectrum (up to $\\ell \\sim 1500$) and will also play a prominent role in the search for the faint imprint of primordial gravitational waves on the CMB polarization.\\\\ This paper addresses the effects of calibration of both temperature (gain) and polarization (polarization efficiency and detector orientation) on polarization measurements. The specific requirements on the polarization parameters of the instrument are set and we report on their pre-flight measurement on HFI bolometers.\\\\ We present a semi-analytical method that exactly accounts for the scanning strategy of the instrument as well as the combination of different detectors.  We use this method to propagate errors through to the CMB angular power spectra in the particular case of Planck-HFI, and to derive constraints on polarization parameters.\\\\ We show that in order to limit the systematic error to 10\\% of the cosmic variance of the $E$-mode power spectrum, uncertainties in gain, polarization efficiency and detector orientation must be below 0.15\\%, 0.3\\% and 1\\deg\\ respectively. Pre-launch ground measurements reported in this paper already fulfill these requirements.}  \\authorrunning{C. Rosset, M. Tristram, N. Ponthieu {\\it et al}} \\titlerunning{Planck-HFI: Polarization Calibration} ",
        "introduction": "The Planck satellite, launched on May 14th, 2009, will map the whole sky in the range 30--857~GHz. One of the most exciting challenges for Planck is to measure the polarization anisotropies of the Cosmic Microwave Background (CMB), which offers a unique way to constrain the energy scale of inflation.  CMB polarization can be decomposed into modes of even-parity ($E$-mode) and odd-parity ($B$-mode). Gravitational waves generated during inflation (hereafter ``primordial'' gravitational waves) create $B$-modes with a specific angular power spectrum, whose amplitude is related to the energy scale of inflation. A detection of these ``primordial'' $B$-modes would therefore provide the first measure of the energy scale of inflation.  $E$-modes were first detected by \\textsc{Dasi} in 2002, followed by other ground and balloon-borne experiments \\citep{Kovac:2002lt, Readhead:lk, Wu:ni, Montroy:2006vk, QUaD-collaboration:-C.Pryke:2008wo} covering a few percent of the sky. These detections are complemented by the \\textsc{Wmap} satellite observations of the whole sky \\citep{Page:2007rw}. All these measurements have confirmed the existence of an $E$-mode polarization compatible with the $\\Lambda CDM$ model, and are compatible with a $B$-mode polarization of zero. The tensor-to-scalar ratio $r$ parametrizes the amplitude of $B$-mode polarization. The most stringent upper limit on $r$ is obtained by \\cite{Komatsu:2009zs}, combining \\textsc{Wmap} measurements of TT, TE and EE power spectra with Baryon Acoustic Oscillations and supernovae data. They obtain $r<0.22$ if the scalar spectral index $n_S$ is constant, or $r<0.55$ if a running spectral index is allowed.    Planck has been designed to map the $E$-mode of polarization with high precision and good control over the polarization foreground contamination up to multipoles as large as $\\ell \\sim 1500$. Planck may also detect the $B$-mode polarization anisotropies, if tensor modes contribute at a level of a few percent or more of the amplitude of the scalar modes \\citep{Efstathiou:2009kx}. However, various instrumental systematic effects, induced by error on the knowledge of detector characteristics, may alter these measurements.  Most of the properties of the detectors, such as the gain, time constant, bandpass and beam, are independent of the sensitivity to linear polarization. These properties are described in detail in companion papers \\citep{Pajot:2010fr,Lamarre:2010fr,Tauber:2010ay,Maffei:2010ly}. In this paper, we study the systematic effects induced by uncertainties in temperature and polarization calibration (gains, polarization efficiencies and orientations) on Stokes parameters and $E$ and $B$-mode power spectra.  We also report on the ground calibration of the polarization efficiencies and orientations of High Frequency Instrument (HFI) detectors. A study of polarization systematics for the Low Frequency Instrument (LFI) of Planck is presented in \\cite{leahy:2010}.  The paper is organized as follows. In Sect.~\\ref{sec:detfp}, we present the Polarization Sensitive Bolometers (PSBs) used by the Planck HFI and the layout of the focal plane. Section~\\ref{sec:photo} gives the generic expression of the polarized photometric equation and introduces the polarization-related systematic effects discussed in Sect.~\\ref{sec:syste}. In Sect.~\\ref{sec:method}, we describe a semi-analytical method to propagate uncertainties on temperature and polarization calibration of detectors up to angular power spectra while exactly accounting for the scanning strategy and the combination of multiple detectors. We apply this method to the Planck HFI in Sect.~\\ref{sec:hfi} and derive requirements on the knowledge of these parameters. Finally, Sect.~\\ref{sec:ground_cal} describes the procedure used to measure polarization parameters of the detectors on ground and compares them to the requirements derived in the previous section. ",
        "conclusions": "This paper focuses on the impact of polarized calibration parameters (gain, polarization efficiency and detector orientation) on power spectra in the context of Planck-HFI. We have developed a semi-analytical method that allows us to compute quickly and easily the impact of uncertainties on gain, polarization efficiency and orientation on the $E$ and $B$-mode power spectra, while exactly accounting for the scanning strategy and the combination of different detectors. We used this method in the particular case of Planck-HFI and derived constraints on the gain, polarization efficiency and detector orientation needed to achieve Planck-HFI's scientific goals.  Planck will use the orbital dipole to calibrate the total power for each detector.  We find that the relative uncertainty on the gain must be lower than 0.15\\% to keep systematic error on $E$-mode power spectrum below 10\\% of the cosmic variance in the multipole range $\\ell = 2-1000$. Given the 0.2\\% accuracy on relative gain obtained by WMAP \\citep{Hinshaw:2009os}, we expect that HFI can achieve the 0.15\\% requirements, thanks to the higher gain stability expected for HFI.  We show that the polarization efficiency uncertainty must be below 0.3\\% in order to achieve the required sensitivity for the $E$-mode. The error on the primordial $B$-mode power spectrum will be kept below 10\\% of the signal expected from a tensor-to-scalar ratio $r=0.05$ in the multipole range $\\ell=2-100$ if the polarization efficiency is known to better than 10.3\\%. In this paper, we have presented the results of the ground measurements on HFI PSBs polarization efficiency, which show an accuracy of 0.3\\% that fulfills the requirements for both $E$ and $B$-modes.  For the polarization orientation, we have distinguished a global orientation error of the focal plane (which affects identically all detectors) from a relative error (different for each detector). For $E$-modes, we show that the requirement is 2\\pdeg1 on the global orientation knowledge and 1\\deg\\ on the relative orientation to keep the error below 10\\% of the cosmic variance in the range $\\ell=2-1000$. Both these requirements are already fulfilled by the ground measurements, in which we found 0\\pdeg3 and 0\\pdeg9 respectively. In order to measure a $B$-mode signal with a systematic error lower than 10\\% for a tensor-to-scalar ratio $r=0.05$, the global orientation must be known to better than 1\\pdeg2 and the relative orientation at better than 0\\pdeg75. While the ground measurements fulfill the requirement on global orientation, the relative orientation knowledge will need to be improved in flight. For Planck, we plan to use the Crab nebula as the primary polarization calibrator \\citep{Aumont:2009rc}, which will also allow the results presented in this paper to be cross-checked. The accuracy of the ground measurements of polarization efficiencies and orientations will allow the $E$-mode power spectrum to be measured, with systematic errors lower than 10\\% of the cosmic variance, provided that the other sources of systematic effects are controlled.  \\acknowledgement{Planck \\emph{(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 in a collaboration between ESA and a scientific Consortium led and funded by Denmark.}"
    },
    "1004/1004.0073_arXiv.txt": {
        "abstract": "{ Recent observations of the Crab pulsar show no evidence for a spectral break in the infrared regime. It is argued that the observations are consistent with a power-law spectrum in the whole observable infrared - optical range. This is taken as the starting point for an evaluation of how self-consistent incoherent synchrotron models fare in a comparison with observations. Inclusion of synchrotron self-absorption proves important as does the restriction on the observed size of the emission region imposed by the relativistic beaming thought to define the pulse profile. It is shown that the observations can be used to derive two independent constraints on the distance from the neutron star to the emission region; in addition to a direct lower limit, an indirect measure is obtained from an upper limit to the magnetic field strength. Both of these limits indicate that the emission region is located at a distance considerably larger than the light cylinder radius. The implications of this result are discussed and it is emphasized that, in order for standard incoherent synchrotron models to fit inside the light cylinder, rather special physical conditions need to be invoked. } ",
        "introduction": "Incoherent synchrotron radiation was recognized early on as a likely emission mechanism for the infrared-optical pulses in the Crab pulsar \\citep[e.g.,][]{S70, P71}. However, as discussed by \\citet{OS70} and \\citet{EP73}, some of the model constraints imposed by the pulsed nature of the emission were initially not explicitly included; e.g., the assumption that the pulses are due to a combination of rotation and relativistic streaming/small pitch angles implies that the frequency of the observed emission should be at least as large as the Doppler boosted cyclotron frequency. This puts an upper limit to the magnetic field in the emission region. With a dipolar magnetic field structure, this results in a minimum distance to the emission region corresponding roughly to the light cylinder radius or even somewhat larger \\citep[e.g.,][]{golden00b}.  The spectral characteristics in the infrared-optical range have long been a matter of debate. The discussion has centered on two issues \\citep[see, e.g.,][ hereafter SS09]{ss09}; namely, the value of the spectral index $\\alpha_\\nu$, defined such that the flux $F(\\nu) \\propto \\nu^{\\alpha_\\nu}$, in the optical and the existence of a possible break and/or bump in the infrared. The main question regarding $\\alpha_\\nu$ is whether or not it is consistent with a value of $1/3$. Since this value is the largest possible for optically thin incoherent synchrotron radiation, such a spectrum would indicate a distribution of electron energies with a sharp low energy cut-off and that the typical synchrotron frequencies for these electrons lie above the optical frequency band. In addition to its implication for the acceleration process, it has consequences for the deduced upper limit of the magnetic field in the emission region. The value discussed above assumes that the observed frequencies correspond to the typical synchrotron frequencies, i.e., that the observed spectral range is determined by the range of electron energies. Hence, the lower frequency limit corresponds to electrons being non-relativistic in the frame where there is no streaming (i.e., the average pitch-angle is $\\pi/2$). This is in contrast to the case of a low energy cut-off for which the lowest energy electrons can be highly relativistic also in the no-streaming frame. Since the inertial mass of the electron increases with the Lorentz factor, the synchrotron spectrum extends below the cyclotron frequency. Hence, the arguments leading to the upper limit of the magnetic field and the associated lower limit of the distance to the emission region are no longer valid and, {\\it a priori}, no restrictions regarding the location of the emission region can be set.  The synchrotron self-absorption frequency is an important parameter for restricting the properties of the emission region. This was used already by \\citet[][ see also \\citealt{PS83}]{S70}  together with the expected brightness temperature to estimate the source size, i.e., its lateral extent. Simple scaling relations have also been derived for the expected emission outside the infrared-optical range for the Crab pulsar itself \\citep[e.g.,][]{PS87} as well as for the infrared-optical emission in other pulsars. In the latter case, the synchrotron self-absorption frequency plays a central role \\citep{oconnor05}.  Another important constraint in the synchrotron scenario is the relation, imposed by the geometry of the emission region, between the observed lateral extent of the source and its distance from the neutron star. Although some aspects of this have been considered, for example by \\citet{PS87} in their scaling relations, a more extensive discussion including synchrotron self-absorption still seems to be lacking. Motivated by recent observational progress regarding the spectral properties of the Crab pulsar, it is the aim of the present paper to provide such an analysis.  In Sect.~\\ref{Obs} the observational situation is summarized and evaluated. It is concluded that $\\alpha_\\nu = 1/3$ is consistent with, although not required by, the observations. Furthermore, it is emphasized that recent observations show no indications of either a break or a bump in the infrared. The temporal structure of the Crab pulses is also briefly discussed. Standard synchrotron theory is applied to a few different settings for the emission region in Sect.~\\ref{Synch}. It is shown that the observations can be used to derive two independent constraints on the properties of the emission region; the first is an upper limit to the strength of the magnetic field, even in the case of a low energy cut-off in the electron energy distribution, while the second one is a direct lower limit of the distance to the emission region. Although the various settings give somewhat different limits, they all suggest that the distance to the emission region is considerably larger than the light cylinder radius. The implications of these results are discussed in Sect.~\\ref{Disc}. It is concluded that an incoherent synchrotron radiation scenario is still tenable although severely restricted. A few alternative origins for the infrared-optical emission are also suggested. Throughout this paper, $cgs$-units are used. ",
        "conclusions": "\\label{Disc} Several explicit scenarios for the emission site of the Crab pulsar have been developed in the past, including the outer gap model \\citep{cheng86} and the polar cap model \\cite[e.g.,][]{harding81}. In this paper we have instead focused on the fundamental constraints inherent to any synchrotron emission model. The main result from Sect.~\\ref{Synch} is that in the incoherent synchrotron scenario, the implied distance to the emission region is considerably larger than the light cylinder radius. This conclusion is supported by two independent pieces of evidence; namely, a direct lower limit of the distance and an indirect one through the upper limit to the magnetic field strength. Most of the uncertain physics is contained in one parameter, which affects these limits in the same way. Hence, by changing its value, one of the limits will indicate a smaller distance to the emission site, while, at the same time, the other limit would suggest an even larger distance. It is therefore not possible to make both limits compatible with a distance comparable to or smaller than the light cylinder radius by invoking a particular value for this parameter (cf. eq. [\\ref{eq:1.14a}]). Furthermore, these limits are quite robust; in particular, both of them are valid using just one of two independent observations: (1) The upper limit on the self-absorption frequency ($\\nu_{\\rm amax}$) or (2) the upper limit on the spectral index (i.e., a $\\nu^{1/3}$-spectrum).  Although the distance to the emission region may be much larger than the light cylinder radius (i.e., $R_{\\rm norm} \\gg 1$), it can still be situated within the light cylinder. This requires a small inclination angle ($i$) between the rotation and magnetic axes. The allowed values for $i$ are bounded from below ($i > 1/ \\Gamma_{\\rm min}$) to assure pulsed emission and from above ($i < 1/R_{\\rm norm}$) by the requirement that the emission region lies within the light cylinder. These inequalities are satisfied in the streaming scenario, since $\\Gamma_{\\rm min} > R_{\\rm norm}$ is needed in order for the lateral extent of the emission region to be smaller than the light cylinder radius. It is seen from equation (\\ref{eq:1.11}) that $\\Gamma_{\\rm min} > R_{\\rm norm}$ is possible for $\\Gamma_{\\rm min} \\approx 10^2$. Hence, a streaming scenario together with a small inclination angle are compatible with observations but certainly for inclinations much smaller than typically envisioned in most models for an oblique rotator \\cite[e.g.,][]{cheng86}.  The constraints imposed by observations on an emission region located close to the light cylinder are qualitatively similar to those corresponding to streaming along the magnetic field lines within the pulsar magnetosphere. In models with the emission site in the vicinity of the last closed magnetic field lines, $R_{\\rm norm} \\approx 1$ is expected; hence, small inclination angles cannot be invoked to make such scenarios compatible with observations. The limits discussed above for the distance to the emission region hinge on the size of the emitting surface and it was argued that an upper bound to this size could be found in the streaming scenario. A similar upper bound is likely to apply also to emission sites close to the light cylinder, in which case such models would be untenable. However, the magnetospheric properties in the vicinity of the light cylinder remain rather uncertain, which leaves open the possibility that this upper bound could be exceeded.  The restrictions on potential incoherent synchrotron models are quite severe. The main problem afflicting them is that their maximum brightness temperature is too low to easily fit the emission region inside the light cylinder. The brightness temperature can be increased by invoking coherent /amplified radiation. An alternative that would preserve many of the attractive features of the standard synchrotron scenario is, therefore, coherent/amplified synchrotron radiation. The conditions needed for synchrotron radiation to be amplified within a pulsar magnetosphere have been discussed  by \\citet{Sto82}. In this case, the emission would likely come from well within the light cylinder and a connection to the radio emission is possible. Another possibility is emission regions located outside the light cylinder.  \\citet{ler70} has considered a scenario, in which the dipole field of the rotating neutron star induces oscillations at the interface with an external plasma. Although the properties of the synchrotron-like emission expected from such models have not been worked out in any detail, their applicability should not be constrained by the size of the emission region."
    },
    "1004/1004.0629_arXiv.txt": {
        "abstract": "We examine whether the accretion of dark matter onto neutron stars could ever have any visible external effects.  Captured dark matter which subsequently annihilates will heat the neutron stars, although it seems the effect will be too small to heat close neutron stars at an observable rate whilst those at the galactic centre are obscured by dust.  Non-annihilating dark matter would accumulate at the centre of the neutron star.  In a very dense region of dark matter such as that which may be found at the centre of the galaxy, a neutron star might accrete enough to cause it to collapse within a period of time less than the age of the Universe.  We calculate what value of the stable dark matter-nucleon cross section would cause this to occur for a large range of masses. ",
        "introduction": "Observations of the kinematics of self gravitating objects such as galaxies and clusters of galaxies consistently send us the same message - if we are to believe in Einstein's theory of gravity on these scales, then there appears to be an invisible quantity of dark matter in each of these objects which weighs more than the matter we can observe. Cosmological observations add weight to this hypothesis and tells us that this invisible matter cannot consist of baryons, rather it must be a new kind of matter which interacts with the rest of the standard model rather feebly - dark matter \\cite{bergstrom}.  The exact coupling and mass of this dark matter is not known but has been constrained. One hypothesis is that the dark matter annihilates with itself and interacts with the rest of the standard model via the weak interaction. This weakly interacting massive particle (WIMP) scenario has gained favour because such particles would fall out of equilibrium with the rest of the plasma at such a temperature that their relic abundance today would be approximately correct to explain the astronomical observations.  Such a scenario also predicts a direct detection signal due to the recoil of atoms which are hit by dark matter particles, recoils which are being searched for at several purpose built experiments (e.g. \\cite{xenon,cdms,zeplin}). We also expect to see signals from the self annihilation of WIMP dark matter in regions of the galaxy where the density is large, although there are many uncertainties with regards to the magnitude of this signal.  Neither of these signals has yet been detected although international efforts to find such signals are intensifying to coincide with the opening of the LHC which also may create WIMP dark matter particles.  Since we only understand the thermal history of the Universe back to the start of nucleosynthesis, we cannot say with any surety whether or not the WIMP scenario makes sense. Furthermore there are many other scenarios of dark matter which involve much more massive particles or particles which cannot annihilate with themselves \\cite{wimpzillas,baryoniccharge}.  There is roughly 5-7 times the amount of dark matter in the Universe by mass relative to baryonic matter.  This ratio is rather close to one, a mystery which is only solved within the WIMP framework by a happy coincidence.  The closeness of these numbers has led some researchers to suggest that, like baryons, dark matter also possesses a conserved charge and there is an asymmetry in this charge in the Universe.  If the two asymmetries are related then one would require the dark matter mass to be approximately 5-7 times the mass of a nucleon. This intriguing possibility would be consistent with the controversial DAMA experiment \\cite{dama} and the slight hint of anomalous noise in the cogent experiment\\cite{cogent}.  Such a dark matter candidate could also have interesting implications for solar physics \\cite{frandsen}.   Since any constraints on the nature of the mass and cross section of dark matter particles are interesting, in this paper we will consider both of these paradigms and see whether or not it is possible to obtain any new constraints from a new angle - namely by considering the capture of dark matter by neutron stars.  The accretion of dark matter onto stellar objects has been considered by various groups looking at both stars \\cite{bouquet,freese,malcolmstars,iocco,taoso} and compact objects \\cite{moskalenkowai,mccullough,Kouvaris2007}. In particular, the ultimate fate of neutron stars which accrete non-annihilating dark matter has been discussed before \\cite{fairbairnbertone,nussinov}.  Our aim is to consider the accretion of dark matter onto neutron stars in greater detail in order to examine whether or not it would ever be possible to either observe the heating of a neutron star due to dark matter annihilation within the object, or the collapse of a neutron star which accretes non annihilating dark matter.  In the next section we will outline our estimate for the accretion rate of dark matter onto a neutron star.  Then we will explain which densities we will be assuming for dark matter in the Milky Way.  We will then go on to work out how hot we can expect a neutron star to get simply due to the accretion of dark matter and compare this with observations.  Finally we will look at whether it is at all sensible to imagine a situation where the accretion of non-annihilating dark matter onto a neutron star would give rise to its subsequent collapse before concluding. ",
        "conclusions": "In this work we have investigated whether or not it would ever be possible to use the accretion of dark matter onto neutron stars in order to understand the properties of dark mmatter better.  We first looked at the effects of annihilating dark matter on the temperature of the neutron stars.  As can be seen in Fig. \\ref{fig:Final-surface-temperatures}, the highest final surface temperatures which could be caused by the heating of neutron stars with dark matter lie around 10$^{6}$ K even in the most optimistic circumstances. Given the surface area of the neutron stars, these values would produce luminosities in the vicinity of 10$^{-2}$ $L_{\\odot}$ with a peak wavelength at about 3 nm (corresponding frequency $\\sim$ 100 PHz). These sources would thus radiate mainly in the range between extreme ultraviolet (UV) and soft X-rays. Given the important absorption due to dust between us and the centre of our galaxy and the presence of other luminous X-ray sources in this region, we believe that the objects in question would prove rather tricky to detect.  Perhaps more interesting are the constraints on non-annihilating dark matter which come from the fact that if enough of such dark matter were to accumulate onto a neutron star it would form a degenerate star at the centre.  If this internal star were to get too large, it might reach its own Chandrasekhar mass, which is smaller than that of the neutron star since the dark matter mass is greater than the nucleon mass in most models.  In the event of the mass of dark matter in the star reaching the Chandrasekhar mass of the star, the dark matter would lead to the collapse of the neutron star which collected it.  Such an event might happen for the kind of values of dark matter mass and cross section currently being probed by direct detection experiments but only in regions of extremely high density.  On the other hand, for higher mass dark matter particles, required cross sections are much smaller, since a much smaller mass of dark matter would need to be accumulated in order for collapse to occur.  The idea that the accretion of stable dark matter could be responsible for the collapse of Neutron stars is very exciting, in this paper we have quantified how likely that is.  For low mass dark matter particles, it seems extremely unlikely."
    },
    "1004/1004.4136_arXiv.txt": {
        "abstract": "A method based on Monte Carlo techniques is presented for evaluating thermonuclear reaction rates. We begin by reviewing commonly applied procedures and point out that reaction rates that have been reported up to now in the literature have no rigorous statistical meaning. Subsequently, we associate each nuclear physics quantity entering in the calculation of reaction rates with a specific probability density function, including Gaussian, lognormal and chi-squared distributions. Based on these probability density functions the total reaction rate is randomly sampled many times until the required statistical precision is achieved. This procedure results in a median (Monte Carlo) rate which agrees under certain conditions with the commonly reported recommended ``classical\" rate. In addition, we present at each temperature a low rate and a high rate, corresponding to the 0.16 and 0.84 quantiles of the cumulative reaction rate distribution. These quantities are in general different from the statistically meaningless ``minimum\" (or ``lower limit\") and ``maximum\" (or ``upper limit\") reaction rates which are commonly reported. Furthermore, we approximate the output reaction rate probability density function by a lognormal distribution and present, at each temperature, the lognormal parameters $\\mu$ and $\\sigma$. The values of these quantities will be crucial for future Monte Carlo nucleosynthesis studies. Our new reaction rates, appropriate for {\\it bare nuclei in the laboratory}, are tabulated in the second paper of this series (Paper II). The nuclear physics input used to derive our reaction rates is presented in the third paper of this series (Paper III). In the fourth paper of this series (Paper IV) we compare our new reaction rates to previous results. ",
        "introduction": "The most influential charged-particle thermonuclear reaction rate evaluations of the 20th century were published by Fowler and collaborators in a series of several papers, with the latest being published in 1988 \\cite{CF88}. The latter work provided compiled rates in tabular and in analytical format for 128 proton- and $\\alpha$-particle induced reactions on A=1 to 30 nuclei. About a decade later, a new reaction rate evaluation by the NACRE collaboration \\cite{Ang99} updated many of the previously published results. The NACRE evaluation contains the rates of 86 reactions on A=1 to 28 nuclei in tabular and analytical format. It represented a major improvement, not only by including newly available nuclear physics input, but it provided for the first time: (i) estimates of reaction rate uncertainties at each temperature in tabular format, and (ii) most of the nuclear data and the associated references used to derive the reaction rates. Another evaluation was published in 2001 by Iliadis and collaborators \\cite{Ili01}. These authors  provided 55 reaction rates involving A=20 to 40 target nuclei in tabular format. They presented reaction rate uncertainties in graphical format and most of the nuclear physics input used to compute the rates. The two major innovations of the latter work were: (i) an extension of the rate evaluation effort to reactions involving radioactive target nuclei, and (ii) the normalization of many resonance strengths to a ``backbone\" of selected and carefully measured standard strengths.  The fast progress seen in the field of nuclear astrophysics over the past few years warrants a new reaction rate evaluation. The original aim was to publish in a short paper the reaction rates that were recently updated by one of us (CI) while working on a textbook \\cite{Ili07} and thus to make them available to the community of stellar modelers. However, it became quickly obvious that there are significant problems in all previously published reaction rate evaluations when the results are confronted with some basic ideas of statistics: what is the statistical meaning of published reaction rates and their uncertainties? Do the published rate uncertainties represent standard deviations of Gaussian distributions or do they perhaps correspond to some other coverage probability? What is the precise meaning of published ``upper\" and ``lower\" limits? And, finally, how can published reaction rate uncertainties be used in the calculations they are mainly intended for, that is, in stellar models?  We argue here that reaction rates from previously published evaluations have no precise statistical meaning. The present work is part of a series of four papers on a new evaluation of charged-particle thermonuclear reaction rates on A=14 to 40 target nuclei. In the first paper, referred to as Paper I, we present a method based on Monte Carlo techniques of estimating statistically meaningful reaction rates and their associated uncertainties\\footnote{It is regrettable that the terms {\\it uncertainty} and {\\it error} are used interchangeably in the nuclear astrophysics literature. According to the {\\it ISO Guide to the Expression of Uncertainty and Measurement (GUM)} \\cite{GUM,NIST} these expressions ``...are not synonyms, but represent completely different concepts; they should not be confused with one another or misused...\". The  {\\it uncertainty} is defined as a ``parameter, associated with the result of a measurement, that characterizes the dispersion of the values that could reasonably be attributed to the measurand\". Uncertainty of measurement comprises, in general, many components; some of these may be evaluated from the statistical distribution of the results of series of measurements and can be characterized by experimental standard deviations; other components, which also can be characterized by standard deviations, are evaluated from assumed probability distributions based on experience or other information. On the other hand, if we use the term {\\it error} in connection with a reaction rate, it means that we think the rate is wrong since perhaps a correction for some systematic effect was disregarded.}. Paper II contains our numerical results in tabular format, while in Paper III we provide the complete nuclear physics data input used to derive our new reaction rates. In Paper IV we compare our new reaction rates to previous results.  The aim of the present work is to evaluate and compile charged-particle thermonuclear reaction rates for A=14 to 40 nuclei on a grid of temperatures ranging from T=0.01 GK to 10 GK. These reaction rates are assumed to involve {\\it bare nuclei in the laboratory}. For use in stellar model calculations, the results presented here must be corrected, if appropriate, for (i) electron screening at elevated densities, and (ii) thermal excitations of the target nucleus at elevated temperatures. Although we occasionally used results from nuclear theory, the present reaction rates are overwhelmingly based on {\\it experimental nuclear physics information}. Only in exceptional situations, for example, when a nuclear property had not been measured yet, did we resort to nuclear theory.  Paper I is organized as follows. In Sec. \\ref{formalism} we present the formalism and the expressions used for computing reaction rates. The commonly employed and accepted procedure of estimating reaction rates and their associated uncertainties is briefly presented in Sec. \\ref{classicalrates}. We refer to all results derived from this method, including those presented in Refs. \\cite{CF88,Ang99,Ili01}, as ``classical reaction rates\". It will become obvious that there are major problems from the statistics point of view with this method. Statistical distributions are briefly reviewed in Sec. \\ref{statdistr} in order to provide a basis for the following discussion. Our method of estimating reaction rates, which is based on Monte Carlo techniques, is presented in Sec. \\ref{MonteCarlo}. We will refer to the new results as ``Monte Carlo reaction rates\". A summary and suggestions for future work are given in Sec. \\ref{summary}. ",
        "conclusions": "\\label{summary} The present work describes a method, based on Monte Carlo techniques, of evaluating  thermonuclear reaction rates. The point is made that reaction rates reported up to now in the literature have no rigorous statistical meaning. As a first step toward a new method, we associate each nuclear physics quantity entering in the calculation of reaction rates with a specific probability density function, including Gaussian, lognormal and chi-squared distributions. Based on these (input) probability density functions the total reaction rate is randomly sampled many times until the required statistical precision is achieved. This procedure results in a median Monte Carlo rate that agrees under certain conditions with the commonly reported recommended ``classical\" rate. For each temperature a low rate and a high Monte Carlo rate is computed, corresponding to the 0.16 and 0.84 quantiles of the cumulative reaction rate distribution. These quantities differ in general from the statistically meaningless ``minimum\" (or ``lower limit\") and ``maximum\" (or ``upper limit\") reaction rates which are commonly reported in the literature. In addition, we approximate the (output) reaction rate probability density function by a lognormal distribution and present, at each temperature, the lognormal parameters $\\mu$ and $\\sigma$. The values of these quantities will be important in future Monte Carlo nucleosynthesis studies. Our new reaction rates, appropriate for {\\it bare nuclei in the laboratory}, are tabulated in the second paper of this series (Paper II). The nuclear physics input used to compute the reaction rates, together with a description of the Monte Carlo code \\texttt{RatesMC}, is presented in the third paper (Paper III). In the fourth paper (Paper IV) we compare our new reaction rates to previous results.  We summarize below certain aspects of our work that call for future efforts from the nuclear astrophysics community: \\\\ (1) We can hardly overemphasize that incomplete information is usually published when  the results of a measurement are reported. It is not sufficient to provide a value and its standard  deviation, but the probability density function on which these values are based should also be reported. This is especially important for null-results: to report simply an ``upper limit\" together with a confidence level is insufficient, unless the most important piece of information, that is, the corresponding probability density function, is reported as well.\\\\ (2) Null-results are incorporated in a consistent way into the present Monte Carlo method within the framework of Porter-Thomas distributions. In order to draw a random sample of a reduced width (or a spectroscopic factor) from a Porter-Thomas distribution, the {\\it local mean value} must be known. It is reasonable to assume that this mean depends on the nuclear mass number $A$ and the orbital angular momentum $\\ell$. In the present work we use values for the mean that have been obtained from our preliminary analysis, that is, by binning together all values of a large data sample and fitting them to a Porter-Thomas distribution, regardless of their $A$ or $\\ell$ values. What is required are systematic studies of nuclear statistical properties that provide improved local mean values for proton and $\\alpha$-particle reduced widths. Similar studies should be performed for reduced $\\gamma$-ray transition probabilities. Theoretical investigations, perhaps employing the shell-model, could be helpful in this regard.\\\\ (3) We approximate the output reaction rate probability density function by a lognormal distribution and provide reasonable arguments for justifying this assumption. However, in some cases, especially when the uncertainty on the resonance energy is large or when undetected low-energy resonances become important, the output reaction rate distributions deviate strongly from lognormality. In such cases we obtain results which can only be approximated by a statistical distribution that depends on more than two parameters. Further studies are required to decide if more complicated expressions of reaction rate probability density functions are needed for future nucleosynthesis studies.\\\\"
    },
    "1004/1004.4907_arXiv.txt": {
        "abstract": " ",
        "introduction": "In this review I will discuss early problems which confused our understanding of NGC 5128 and provide an overview of what we have learned about the properties of its old stars, with an emphasis on the globular cluster and field star populations. Because NGC 5128 is $<$ 4Mpc away ($(m=M)_0 = 27.9$) we can study its stellar component in greater depth and detail than is possible for any other large elliptical galaxy.  Consequently we have data for hundreds of globular clusters and planetary nebulae which are telling a rich history of how and when its stars formed. In addition (and unknown to many readers) we can now resolve individual halo stars, long period variables and Cepheids in NGC 5128, telling us its distance and providing additional clues to its history. For recent results on these and other stellar components see: distance \\citep{har10}, ages of field stars and long period variables \\citep{rej03, rej05}, globular clusters (\\citep{peng04b, woodoz, wood10}, and planetary nebulae \\citep{wal99, peng04a, wal10}. ",
        "conclusions": "$\\bullet$ There is now agreement, within a few $\\%$ as to the distance to NGC 5128, making it possible to interpret observed properties with new confidence. \\\\ $\\bullet$ Although there is still debate as to whether this is an S0p or Ep galaxy, the properties of its halo and individual stars are consistent with classifying NGC 5128 as an Ep. \\\\ $\\bullet$ We have a sample of $\\sim$600 globular clusters confirmed by radial velocity and/or angular resolution, approximately 50$\\%$ of the estimated total population.  Roughly half of these are metal-rich and the majority are old.  \\\\ $\\bullet$ In spite of these numbers, global properties of the system are not well known because of major biases in the spatial coverage. \\\\ $\\bullet$ Observations of the spatial structure of $>$200 GCs shows that they follow the same fundamental plane relation as do GCs in the Milky Way and M31. The NGC 5128 system allows us to trace this relation to higher GC luminosity and mass than we can in the Milky Way or M31. \\\\ $\\bullet$ Over a wide range in galactocentric distance, the halo stars are remarkably similar and predominantly metal-rich; and it appears that these metal-rich stars are old, possibly with ages similar to those of the metal-rich GCs. We may need to observe the halo to galactocentric distances twice that of current datasets in order to uncover its metal-poor component."
    },
    "1004/1004.1525_arXiv.txt": {
        "abstract": "{We present  the preliminary results of the study of an interesting target in the first CoRoT  exo-planet field (IRa1): CoRoT~102918586. Its light curve presents additional variability on the top of the eclipses, whose pattern suggests multi-frequency pulsations. The high accuracy CoRoT light curve was analyzed by applying an iterative scheme, devised to disentangle the effect of eclipses from the oscillatory pattern.  In addition to the CoRoT photometry we obtained low resolution spectroscopy  with the AAOmega multi-fiber facility at the Anglo Australian Observatory, which yielded a spectral classification as F0~V and allowed us to infer a value of the primary star effective temperature. The Fourier analysis of the residuals,  after subtraction of the binary light curve, gave 35 clear frequencies. The highest amplitude frequency, of 1.22 c/d, is in the expected range for both $\\gamma$ Dor and SPB pulsators, but the spectral classification favors the first hypothesis. Apart from a few multiples of the orbital period, most frequencies can be interpreted as   rotational splitting of the main frequency (an $\\ell =2$ mode) and of its overtones. } ",
        "introduction": " ",
        "conclusions": "The preliminary results on CoRoT 102918586 are based on a single pass-band (though excellent) photometry and low resolution spectra, nevertheless the combination of pulsation and binarity yields already a rich harvest of information. Obviously we need high resolution and time resolved spectroscopy to definitely solve the remaining ambiguities of this interesting system. This target is included with high priority in our spectroscopic follow-up  programs."
    },
    "1004/1004.1180_arXiv.txt": {
        "abstract": "There is a strong decrease in scatter in the $M_\\bullet$ -- $L_{bulge}$ relationship with increasing luminosity and very little scatter for the most luminous galaxies.  It is shown that this is a natural consequence of the substantial initial dispersion in the ratio of black hole mass to total stellar mass and of subsequent galaxy growth through hierarchical mergers.  ``Fine-tuning'' through feedback between black hole growth and bulge growth is neither necessary nor desirable. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.4841_arXiv.txt": {
        "abstract": "{}{We attempt to increase the number of Trans-Neptunian objects (TNOs) whose short-term variability has been studied and compile a high quality database with the least possible biases, which may be used to perform statistical analyses.} {We performed broadband CCD photometric observations using several telescopes.} {We present results of 6 years of observations, reduced and analyzed with the same tools in a systematic way. We report completely new data for 15 objects (1998SG$_{35}$, 2002GB$_{10}$, 2003EL$_{61}$, 2003FY$_{128}$, 2003MW$_{12}$, 2003OP$_{32}$, 2003WL$_{7}$, 2004SB$_{60}$, 2004UX$_{10}$, 2005CB$_{79}$, 2005RM$_{43}$, 2005RN$_{43}$, 2005RR$_{43}$, 2005UJ$_{438}$, 2007UL$_{126}$ (or 2002KY$_{14}$)), for 5 objects we present a new analysis of previously published results plus additional data (2000WR$_{106}$, 2002CR$_{46}$, 2002TX$_{300}$, 2002VE$_{95}$, 2005FY$_{9}$) and for 9 objects we present a new analysis of data already published (1996TL$_{66}$, 1999TZ$_{1}$, 2001YH$_{140}$, 2002AW$_{197}$, 2002LM$_{60}$, 2003AZ$_{84}$, 2003CO$_{1}$, 2003VS$_{2}$, 2004DW). Lightcurves, possible rotation periods and photometric amplitudes are reported for all of them. The photometric variability is smaller than previously thought: the mean amplitude of our sample is 0.1mag and only around 15\\% of our sample has a larger variability than 0.15mag. The smaller variability than previously thought seems to be a bias of previous observations. We find a very weak trend of faster spinning objects towards smaller sizes, which appears to be consistent with the fact that the smaller objects are more collisionally evolved, but could also be a specific feature of the Centaurs, the smallest objects in our sample. We also find that the smaller the objects, the larger their amplitude, which is also consistent with the idea that small objects are more collisionally evolved and thus more deformed. Average rotation rates from our work are 7.5~h for the whole sample, 7.6~h for the TNOs alone and 7.3~h for the Centaurs. All of them appear to be somewhat faster than what one can derive from a compilation of the scientific literature and our own results. Maxwellian fits to the rotation rate distribution give mean values of 7.5~h (for the whole sample) and 7.3~h (for the TNOs only). Assuming hydrostatic equilibrium we can determine densities from our sample under the additional assumption that the lightcurves are dominated by shape effects, which is likely not realistic. The resulting average density is 0.92~g/cm$^3$ which is not far from the density constraint that one can derive from the apparent spin barrier that we observe.} {} ",
        "introduction": "The rotational properties of the asteroids provide plenty of information about important physical properties,  such as density, internal structure, cohesion and shape (e.g., \\cite{Pravec-Harris2000,Holsapple2001, Holsapple2004}). As for asteroids, the rotational properties of the Kuiper Belt Objects (KBOs) provide a wealth of knowledge about the basic physical properties of these icy bodies (e.g. \\cite{Sheppard-Jewitt2002, Ortiz2006,Trilling-Bernstein2006,Sheppard2008}). In addition, rotational properties provide valuable clues about the primordial distribution of angular momentum, as well as the degree of collisional evolution of the different dynamical groups in the Kuiper Belt. Rotational properties can also provide empirical tests of predictions based on models of the collisional evolution of the Kuiper Belt \\citep{Davies1997, Benavidez2009}.  Studies of short-term photometric variability of Kuiper Belt Objects allow us to retrieve rotation periods from the photometric periodicities  and these studies also provide constraints on shape (or surface heterogeneity) by means of the amplitude of the lightcurves. Therefore, observational programs on time series CCD photometry are a good tool for studing the Kuiper Belt. Unfortunately, most KBOs are faint and CCD photometry programs are time expensive and require medium to large telescopes. For this reason and in contrast to the case for asteroids,  the sample of KBOs for which short-term variability has been studied is not large and more importantly, the sample of objects with known rotation periods is severely biased toward large photometric amplitudes, the reason being that the large amplitude objects produce lightcurves that are much easier to observe and from which a rotation period can be unequivocally obtained. Very long rotation periods are also difficult to determine and scientists rarely publish null results or failed attempts to derive lightcurves, which causes a bias in the literature. In other words, the scientific literature in the Kuiper Belt field consists mostly of high amplitude lightcurves, which are not truly representative of the rotation properties of the whole Kuiper Belt population.  Other biases are present in the literature, such as an overabundance of large objects because larger bodies are usually brighter and easier to observe. One can try to avoid this bias by studying Centaur objects, which are not TNOs because they are not farther away than Neptune but are widely accepted to originate in the Kuiper Belt; they are thus KBOs that recently came to the inner solar system vicinity, with ``recently\" meaning time frames of several mega-years, the typical lifetime of Centaurs \\citep{Tiscareno2003}. The currently known Centaurs are smaller than TNOs with the same brightness. Therefore, targeting Centaurs provides information about small size objects of the Kuiper Belt that would otherwise be too faint for telescopic studies.  Based on the aforementionned ideas we initiated a programme to photometrically monitor as many KBOs as possible, trying to build a reasonably good sample in terms of number of objects observed and also in terms of the amplitude biases. Thus, we published even dubious rotation periods of low-amplitude lightcurve objects rather than omitting them. In their review of asteroid rotation rates, \\cite{Binzel1989} emphasized that excluding poor reliability objects results providing more weight to asteroids with large amplitudes and short periods, introducing a significant bias.  We present new results from our survey and a reanalysis of several bodies for which we had already published results. We reanalyzed the data for these bodies either because we had acquired more data or because we had developed superior analysis tools that resulted in what we consider an improvement. All the objects presented here were analyzed with the same tools and software and therefore represent a homogeneous data set. The work presented here summarizes a considerable effort in which more than 5000 images were reduced and analyzed.  This paper is divided into 6 sections. Section 2 describes the observations and the data sets. Section 3 describes our software reduction tools and the methods used to derive e.g., periodicities, rotation periods and photometric range. Section 4 deals with the main results obtained for each object and Section 5 discusses the results altogether. Finally, our findings are summarized in Section 6. ",
        "conclusions": "We have presented a large sample of Kuiper Belt Objects whose short-term variability has been studied in detail to increase the number of objects studied so far and try to avoid observational biases. Amplitudes and rotation periods have been derived for all of them with different degrees of reliability, but we have compiled an ensemble of all of them to study the whole population. We present therefore a homogeneous data set from which some conclusions can be drawn. We found that the percentage of low amplitude rotators is higher than previously thought and that in our sample the rotation rates appear to be slightly higher (faster objects) than previously suggested.  A simple idea investigated in detail in \\cite{Duffard2009} to explain the large abundance of small amplitude objects might be that hydrostatic equilibrium is applicable to the overwhelming majority of the bodies and that the usual KBO shapes are MacLaurin spheroids which therefore do not cause any shape induced variations (and whose variability is caused by albedo variegations exclusively). We estimate that 0.1~mag seems to be a good measure of the typical variability caused by albedo features.  The plots of both amplitude versus size and rotation rate versus size seem to be compatible with the typical collisional evolution scenario in which larger objects have been only slightly affected by collisions, whereas the small fragments are highly collisionally evolved bodies with usually more rapid spins of larger amplitudes.  Based on the assumption of hydrostatic equilibrium, one can derive densities for all the bodies and we found a possible trend of higher densities toward higher sizes, which is a physically plausible scenario. There appears to be a spin barrier that allows us to obtain a density limit that is also compatible with the average density derived based on hydrostatic equilibrium assumptions. Nevertheless, a more appropriate derivation of mean densities is presented in \\cite{Duffard2009}."
    },
    "1004/1004.0379_arXiv.txt": {
        "abstract": "We show that a possible astrophysical experiment, detection of lensed images of stars orbiting close to Sgr A*, can provide insight into the form of the metric around a black hole. We model Sgr A* as a black hole and add in a $\\varpropto \\frac{1}{r^2}$ term to the Schwarzschild metric near the black hole. We then attempt to determine the effect of this extra term on the properties of the secondary images of the S stars in the Galactic Center. When the $\\frac{1}{r^2}$ term is positive, this represents a Reissner-Nordstrom (RN) metric, and we show that the there is little observational difference between a Schwarzschild black hole and a RN black hole, leading to the conclusion that secondary images may not be a useful probe of electrical charge in black holes. A negative value for the $\\frac{1}{r^2}$ term can enter through modified gravity scenarios. Although physically unlikely to apply in the case of a large black hole, the Randall-Sundrum II braneworld scenario admits a metric of this form, known as tidal Reissner-Nordstrom (TRN) metric. We use values of tidal charge ($Q$ in $\\frac{Q}{r^2}$) ranging from $-1.6 M^2$ to $0.4 M^2$. A negative value of $Q$ enhances the brightness of images at all times and creates an increase in brightness of up to 0.4 magnitudes for the secondary image of the star S2 at periapse. We show that for other stars with brighter secondary images and positions  more aligned with the optic axis, using the Tidal Reissner-Nordstrom metric with negative $Q$ enhances the images as well, but the effect is less pronounced. This effect is related to the increase in the size of the photon sphere, and therefore, should be noticeable in other metrics with a similar effect on the photon sphere. With the next generation of instruments and increased knowledge of radiation from Sgr A*, using properties of secondary images to place constraints on the size of the $\\frac{1}{r^2}$ term. This knowledge will be useful in constraining any modified gravity theory that adds a similar term into the strong field near a black hole. ",
        "introduction": "Gravitational lensing provided the first experimental verification of general relativity (GR) through observations of starlight bending around the Sun during an eclipse in 1919 \\cite{schneider, 1919eclipse} and continues to be a major source of insight into gravitation \\cite{LensingReview, schneider, jainreview}. Excitingly, increasingly precise observations of the compact radio source Sgr A* at the Galactic center and its surrounding stars have given us very high confidence that only a very gross deviation from GR could allow for the absence of a black hole there \\cite{NatureHorizon,smbhdetails, ghez2008, propSgrA, propSgrA2}.  According to \\cite{propSgrA}, the black hole is estimated to have a mass of about $4.31 \\times 10^6 M_{\\odot}$ and a distance of about 8.33 kpc from Earth. Black holes are unique laboratories for gravitational lensing because their compactness allows light to closely approach the photon sphere and its path will bend significantly there due to gravity. Most studies of gravitational lensing \\cite{petters, petterskeeton} are in the weak deflection limit, when the point of closest approach of any lensed photons is far from the lensing mass. This allows for the simple expression of the bending angle as $\\alpha = \\frac{4M}{r_0}$, where $r_0$ is the point of closest approach of the null geodesic (this paper uses geometric units $G=c=1$). This expression is valid in the limit $\\frac{M}{r_0} \\ll 1$. When a photon closely approaches a  black hole's photon sphere, the weak deflection limit does not hold and using this approximation for the bending angle leads to inaccurate results. For a spherically symmetric, static metric with line element  \\begin{equation} ds^2= -A(r)dt^2 + B(r) dr^2 +C(r) r^2 d\\Omega^2 \\label{metric} \\end{equation} the bending angle is an elliptic integral based on the functions of the metric \\cite{VE2002} and is \\begin{eqnarray} \\nonumber \\alpha (r_0) &=& 2 {\\int_{r_0}}^{\\infty}\\left(\\frac{B(r)}{C(r)}\\right)^{1/2} \\left[(\\frac{r}{r_0})^2\\frac{C(r)}{C(r_0)}\\frac{A(r_0)}{A(r)}-1\\right]^{-1/2} \\\\ & \\times & \\frac{dr}{r}- \\pi \\label{bending} \\end{eqnarray}  If the point of closest approach is very close to the photon sphere, the strong deflection limit approximation of this integral \\cite{bozza2002, bozza2007} can be used. In all cases, a full numerical treatment of the bending angle can be used \\cite{VE2000}, and it should be used for studies, such as this one, where neither the strong nor weak deflection limits are satisfied. This is further explained in Sec. \\ref{sec:lensing}. Large bending angles yield interesting results: when photons approach close enough to the photon sphere, they can loop around the lens before reaching the observer and produce an infinite sequence of images on both sides of the optic axis \\cite{darwin, VE2000, bozzaetal, bozza2002, V2009}. The properties of these images are sensitive to the form of the black hole metric \\cite{eiroad, whisker2005, me, mmreview, whiskerthesis} with \\cite{eiroarn} studying the effect of the RN metric. There has been further study of an astrophysical scenario which utilizes lensing with a large bending angle. S2 was the first star that was examined for a ``retrolensing\" \\cite{depaolis} effect (lensing with a bending angle of $ \\alpha \\approx \\pi$).  By using orbital parameters of the S stars in the Galactic center provided by \\cite{propSgrA}, \\cite{bozzas2, bozzas2revised,bozzamancini2009} study the properties of secondary images of those S stars. This study showed 9 stars with secondary images with a maximum brightness in the K-band brighter than $m_K=30$, a cutoff that maximizes observational prospects.  In this paper, we use a black hole metric with an additional term that is $\\varpropto \\frac{1}{r^2}$. The metric is \\begin{equation} ds^2 = -(1-\\frac{2M}{r} + \\frac{Q}{r^2})dt^2 + (1-\\frac{2M}{r}+  \\frac{Q}{r^2})^{-1} dr^2 +r^2 d\\Omega^2 \\label{eq:trnmetric} \\end{equation} with $ Q$ a free parameters often expressed by $q \\equiv \\frac{Q}{4M^2}$. When $q$ is positive, this represents the Reissner-Nordstrom solution for a charged black hole. Static black holes with a large amount of electric charge are not expected to exist in nature, and the existence of rotating, charged black holes is controversial \\cite{punsly}. In addition, the amount of charge is limited to $Q < M^2$ or $q < 0.25$ because the saturation of this bound would lead to a naked singularity and the violation of cosmic censorship \\cite{waldcensorship}. However, since the calculation of the properties of secondary images of S stars has only been done with a Schwarzschild metric, it is useful to examine such a fundamental case. We have also found it useful to explore negative values of $q$. There are only weak constraints on a lower bound for $q$ that come from studies of neutron star binary systems \\cite{trnconstraints}. Using a $\\frac{1}{r^2}$ term is interesting because it is motivated by alternative gravity frameworks, particularly braneworld theories that construct gravity as a higher-dimensional theory. As these theories usually predict a correction that strengthens gravity, this would correspond to a negative value for $q$. For any non-trivial value of $q$, observational constrains from Solar System observations \\cite{dmptidalrn} disallow the $\\frac{1}{r^2}$  term from applying anywhere but near the black hole. Strengthening gravity using a negative value for $q$ yields a brighter secondary image; therefore, observations of these secondary images can place constraints on the size of the $\\frac{1}{r^2}$ term near the black hole.  An extra $\\frac{1}{r^2}$  term in the metric comes directly from a proposed black hole metric in the Randall-Sundrum II theory. The Randall-Sundrum I model \\cite{rs1} is a braneworld scenario inspired by heterotic M-theory \\cite{ovrut1, ovrut2, ovrut3}, a 5 dimensional effective framework that arises from the dimensional reduction of 11-dimensional Harova-Witten theory on a Calabi-Yau manifold with the imposition of $S^1/Z_2$  symmetry \\cite{witten1}. Six dimensions are compactified, making gravity effectively 5 dimensional. The Randall-Sundrum II scenario is this model with the second brane taken to infinity \\cite{rs2, MaartensReview}. It is not clear whether a static black hole solution exists for a RS II braneworld. Although \\cite{emparan} has shown the existence  of a static black hole in a 2+1 brane setup, it is unclear what the solution is and whether there even is a static black hole solution in the  3 + 1 brane scenario \\cite{kanti, creek, shiromizu2, wiseman, wiseman2}. Attempts to study the problem have yielded several possible black hole metrics \\cite{dmptidalrn, whisker2005, blackstring, casadio, gtmetric}. The effect of these metrics on lensing has been examined in several studies \\cite{eiroad, whisker2005, me, petterskeeton, mmreview}. For a supermassive black hole such as Sgr A*, the Garriga-Tanaka \\cite{gtmetric} and ``black string\"  \\cite{blackstring} metrics will not show any results \\cite{me} and this paper aims to show that a modification of the metric near the horizon can affect lensing observables. One particular metric that has been studied in connection with lensing by the black hole at Sgr A* is the ``tidal\" Reissner-Nordstrom metric \\cite{dmptidalrn, eiroad, whiskerthesis, me} which is of the form of Eq. (\\ref{eq:trnmetric}). We therefore use it as an example of a metric with a $\\frac{1}{r^2}$ potential term. There are several possible objections to using the TRN metric for a supermassive black hole: Studies suggest that supermassive black holes in the braneworld should have induced metrics that are no more than negligibly different than the Schwarzschild metric \\cite{wiseman}. In addition, the TRN metric has no known completion in the bulk \\cite{whiskerthesis} and there are likely naked singularities in the bulk for the TRN metric \\cite{trnpath}. However, we will study lensing properties of the TRN metric to gain understanding of the effects of adding a $\\varpropto \\frac{1}{r^2}$ term in the potential, whether it comes from the braneworld scenario or any other gravitational framework.  In Sec. II, we discuss S stars near the Galactic center. Now that their orbital parameters are well known from decades of observation \\cite{propSgrA}, we show how to reconstruct the orbits of these stars and characterize the variables in the Ohanian lens equation \\cite{bozzalens, ohanian} in terms of the star's orbital parameters. Although current uncertainties in lensing parameters can interfere with the test we are proposing, future instruments and observations will further constrain these orbital parameters. In Sec. III, we use the orbital parameters to construct a light curve for the secondary images of S2, S6, and S14. For these stars, we show the light curve of the secondary image when assuming a Schwarzschild spacetime, a TRN spacetime with $q= -1.6$, and, for the case of S2, an extremal RN spacetime with $q= 0.25$. We show that for high enough values of the tidal charge parameter $q$, an appreciable difference appears for these light curves in the TRN spacetime. We also show the relevant properties, such as image magnitude and image position at peak brightness for several values of $q$. In Sec. IV, we discuss the observational prospects for these images as well as the observing the difference in brightness for these images. We conclude that observation of these images is possible, as is observing the difference in image properties due to a modified gravity theory. ",
        "conclusions": "Discussion}  Of the three stars studied in this paper, S14 and S6 have brighter secondary images at periapse compared to S2's image because of the more edge-on nature of their orbits relative to our line of sight with Sgr A*. While the orbit of S2 is not aligned close to edge-on, its periapse is the closest amongst known stars, causing its secondary image to be very bright. In addition, S2 will next be at periapse in 2018, allowing for a more immediate study of its image's properties.  As mentioned above, the brightness of secondary images is dependent on $q$ because the images are either pushed closer or further from the optic axis when there is a $\\frac{1}{r^2}$ term in the metric. This is directly related to the difference in the size of the event horizon and photon sphere due to the value of $q$. Another, and perhaps easier, way of determining the metric around the black hole would be to measure the shape and size of the black hole's horizon or photon sphere. A preliminary attempt at this has been made \\cite{NatureHorizon}, but they were unable to identify the observed structure with the black hole itself. At present, there are no constraints which would prevent the size of the event horizon from being significantly larger or smaller than predicted in a Schwarzschild spacetime (but not by an order of magnitude \\cite{flaring}). In addition, since the photon sphere is so small, even a relatively large percentage change in its size corresponds to only a few $\\mu$as. Resolving the difference between two proposed photon sphere sizes may be beyond the capabilities of projected future instruments. In this case, observing secondary images may provide more information.  There has been some discussion about observing these secondary images with the upcoming generation of telescopes \\cite{bozzamancini2009}. A very promising project is the MCAO Imaging Camera for Deep Observations (MICADO) telescope \\cite{micado} at E-ELT. It is expected to be able resolve images as faint as $m_K=30.1$ and will have a resolution of up to 6 mas in imaging mode (10 mas in the K-band) and an astrometric accuracy of up 10 $\\mu$as. It also will have a photometric accuracy of 0.03 magnitudes. From the data shown above, trying to resolve the difference between an image's position in the Schwarzschild metric and an image's position in the TRN metric is next to impossible given the small separation between these positions.  Therefore, it is important to explore not only image (and photon sphere) positions, but to use image magnfications as a complementary avenue for exploring the metric. If the secondary images in this paper were isolated point sources, MICADO would not have any problems detecting them and even differentiating between two predicted values for image brightness (assuming the difference is larger than the photometric accuracy). However, these images will be very close to and essentially unresolvable from Sgr A* and its crowded environment. This is not a fatal flaw, because Sgr A* is very faint in the Near Infared K-band ($\\lambda = 2.2 \\mu$m), and it may be feasible to subtract out the quiescient state of Sgr A* in the K-band. Some studies \\cite{flares, flares2, ghezconstant, constant2} claim that Sgr A* has a highly variable ``quiescient\" state of $m_K \\approx 17$. In addition, there are occasional flares that can be brighter than $m_K = 16$ and last on a scale of hours. While the flares are thought to originate very close to the black hole, it is not clear whether the ``quiescient\" radiation in the K-band comes from Sgr A* itself or whether it comes from the fact that the lower resolution (65 mas) in the survey includes one or more sources in the area near Sgr A*. The currently known star with the closest approach is S2, which appears to be 11 mas from Sgr A* at the point of closest approach- at all points, MICADO should be able to resolve S2 and many other closely moving stars from Sgr A*. There may be other, unidentified stars that are currently conflated with radiation from Sgr A* but will be separately with MICADO's increased resolution. Better resolution, combined with further understanding of the K Luminosity Function near the Galactic center means that observations of secondary images is a real possibility. MICADO's photometric accuracy should be able to distinguish many of the image brightness differences discussed in this paper. This assumes that when viewed with a fine enough resolution, Sgr A* does not emit too brightly in the K-band and source crowding is not insurmountable. Even if flares are observed, they persist for a time scale much shorter than secondary images (which remain bright for months) and should be easily distinguished from the quiescient state. Additionally, it is expected that an additional $10^0$-$10^2$ S stars will be found in the central milli-parsec of the galaxy \\cite{milliparsec}, yielding additional and perhaps even better candidates to observe secondary images with the right properties to give us insight into the metric near the black hole. It may not be possible to accurately treat stars that close to Sgr A* with the thin-lens approximation, but in that case, a more exact numerical treatment can be utilized \\cite{bozzas2, cunningham}.  Observing secondary images of S stars will be challenging, but we may very well find that it is possible to observe faint secondary images and use their properties to give us insight into the true nature of gravity. This exciting prospect should be an additional motivation for the next generation of observational instruments aimed at the Galactic center."
    },
    "1004/1004.1910.txt": {
        "abstract": "This talk discusses the formation of primordial intermediate-mass black holes, in a double-inflationary theory, of sufficient abundance possibly to provide all of the cosmological dark matter. There follows my, hopefully convincing, explanation of the dark energy problem, based on the observation that the visible universe is well approximated by a black hole. Finally, I discuss that Gell-Mann is among the five greatest theoreticians of the twentieth century. name{Abstract} \\newcommand\\bibname{References}% % \\newcommand\\today{\\ifcase\\month\\or January\\or February\\or March\\or April\\or May\\or June\\or July\\or August\\or September\\or October\\or November\\or December\\fi \\space\\number\\day, \\number\\year} \\newcount\\minute \\newcount\\hour \\def\\currenttime{% \\minute\\time \\hour\\minute \\divide\\hour60 \\the\\hour:\\multiply\\hour60\\advance\\minute-\\hour\\the\\minute} % \\newdimen\\trimheight \\newdimen\\trimwidth \\newdimen\\typeheight \\newdimen\\typewidth \\newdimen\\tempdimen \\newdimen\\tablewidth \\newdimen\\normaltextheight \\newbox\\tempbox \\newdimen\\tablewd % %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Fonts %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % \\def\\foliofont{\\fontsize{8}{10}\\selectfont} \\def\\bibfont{\\fontsize{9}{11}\\selectfont} \\def\\rhfont{\\footnotesize\\itshape{}} \\def\\catchlinefont{\\footnotesize} 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        "introduction": "\\noindent It is an honor to talk at a festschrift for Murray Gell-Mann, who dominated research in particle phenomenology for at least twenty years.  \\bigskip  \\noindent At the beginning of my talk, I shall discuss a recent paper \\cite{FKTY} on the production of primordial intermediate-mass black holes of mass $M_{BH}=M_{\\odot}^p$ with $-8 \\leq p \\leq +5$, providing a sufficient abundance, that the primordial IMBHs can possibly act as all the cosmological dark matter. \\bigskip  \\noindent I then discuss my solution\\cite{Frampton} for the difficult dark energy problem which was first identified from observations of supernovae, twelve years ago. Although I knew all the correct theoretical ingredients back then, the solution hit me only on February 6, 2010. Because this was an overwhelming human experience, I self-indulgently discuss it.  \\bigskip  \\noindent Finally, I discuss why Gell-Mann, who must himself have experienced a similar personally fulfilling moment, for the $\\Omega^{-}$ particle\\cite{MGMomega,EXPTomega}, is to be correctly, regarded as among the five greatest theoreticians, of the twentieth century.  \\bigskip ",
        "conclusions": ""
    },
    "1004/1004.0009_arXiv.txt": {
        "abstract": "We present continent-scale very-long-baseline interferometry (VLBI) -- obtained with the European VLBI Network (EVN) at a wavelength of 18~cm -- of six distant, luminous submillimetre-selected galaxies (SMGs). Our images have a synthesized beam width of $\\approx$30\\,milliarcsec {\\sc fwhm} -- three orders of magnitude smaller in area than the highest resolution Very Large Array (VLA) imaging at this frequency -- and are capable of separating radio emission from ultra-compact radio cores (associated with active super-massive black holes -- SMBHs) from that due to starburst activity. Despite targeting compact sources -- as judged by earlier observations with the VLA and MERLIN -- we identify ultra-compact cores in only two of our targets. This suggests that the radio emission from SMGs is produced primarily on larger scales than those probed by the EVN, and therefore is generated by star formation rather than an AGN -- a result consistent with other methods used to identify the presence of SMBHs in these systems. ",
        "introduction": "\\label{sec:intro}  With bolometric luminosities rivalling quasars, submillimetre galaxies (SMGs) are some of the most extreme objects in the Universe \\citep[for a review, see][]{blain2002}, yet they outnumber similarly luminous quasars by many orders of magnitude \\citep{chapman2005}. Selected via their dust-reprocessed thermal emission, they are thought to be powered by intense bursts of star formation triggered by major mergers of gas-rich galaxies \\citep{tacconi2006, tacconi2008, younger2008highres,younger2010highres} at $z\\sim 2-3$ \\citep{chapman2005} with a significant tail extending out to higher redshifts \\citep{eales2003, younger2007, younger2009.aztecsma, wang2007, wang2008, greve2008, capak2008, schinnerer2008, daddi2009, coppin2009}, and are the likely progenitors of massive galaxies in the local Universe \\citep{scott2002, blain2004, smail2004, swinbank2006}. In addition, it is thought that infrared- (IR-)luminous objects ($L_{\\rm IR} \\gsim 10^{12-13} \\mathrm{L}_{\\sun}$) come to dominate the cosmic star-formation rate density at high redshift \\citep[$z\\sim 2-4$;][]{blain1999, hopkins2009.ulirg, hopkins2009.ulirg2}, making the SMG population an important contributor to the build-up of stellar mass during the epoch of galaxy formation.  At the same time, over the past twenty years it has become clear that super-massive black holes (SMBHs) in the nuclei of galaxies are common, even at high redshift \\citep[e.g.][]{ivison2008} and that their masses are strongly correlated with the stellar component of their host galaxies \\citep[e.g.,][]{kormendy1995, magorrian1998, Gebhardt2000, Ferrarese2000, Tremaine2002, novak2006, Hopkins2007obs}. These correlations are indicative of a close and fundamental link between SMBHs and stellar populations, and the enormous difference ($\\sim 10^9$) in linear scales suggests that this may be accomplished via accretion-related radio jets or radiatively-driven winds.  Recent theoretical models cast the co-evolution of SMBHs and galaxies in the context of a cosmic cycle driven by major mergers \\citep[][and references within]{hopkins2007a, hopkins2007b}, wherein gas-rich, spiral galaxies collide and trigger a gas inflow, thereby fuelling a nuclear starburst \\citep{hernquist1989a, mihos1994, mihos1996} and feedback-regulated SMBH growth \\citep{silk1998, Page2004, diMatteo2005, hopkins2007theory, younger2008.smbh}, eventually revealing a bright quasar, after which gas exhaustion and violent relaxation transform the remnant into a red elliptical galaxy \\citep{barnes1992b}.  These models of galaxy evolution, which match constraints from both galaxy \\citep{hopkins2006} and SMBH populations \\citep{robertson2006a, hopkins2007theory, younger2008.smbh}, predict that hyperluminous starbursts at high redshift will be associated with periods of active SMBH growth. SMGs are thus likely candidates for the transition objects that theory predicts should be powered by a mixture of star formation and active black hole growth \\citep[see also][]{younger2009.warmcold}.  To date, searches for actively growing SMBHs in SMGs have focused on X-ray observations \\citep{alexander2005, alexander2008} and mid-IR photometry and spectroscopy \\citep{ivison2004, lutz2005, menendez2007, menendez2009, valiante2007, pope2008b}. Whilst promising, these approaches are often compromised: even hard X-ray searches will miss Compton-thick SMBHs, and the interpretation of mid-IR spectra can be very sensitive to modelling assumptions, especially regarding the origin of power-law spectra \\citep{yun2001, younger2009.warmcold}.  Very-high-resolution radio imaging, provided by continent-scale very-long-baseline interferometry (VLBI), avoids many of these issues; even in the most extreme environments, the interstellar medium is optically thin to radio emission and the dominant radio emission mechanisms are well understood. In fact, there is a well-known upper limit to the brightness temperature ($T_{\\rm b}$) a starburst can achieve \\citep[$T_{\\rm b} \\lsim 10^5$ {\\sc k} --][]{condon1991b}, and any compact object detected at VLBI resolution will exceed (or be close to) this limit.  A number of authors have used this technique in order to distinguish between AGN and starburst activity in both high-redshift QSOs \\citep{beelen2004,momjian2005,momjian2007} and in ultra-luminous infrared galaxies (ULIRGs), the low-redshift analogues of the SMGs. The latter are particularly important as they present a far more detailed view of the relationship between the star formation and AGN emission in an intense starburst than will ever be possible in a distant SMG, and aid in the interpretation of these less spatially resolved sources; the synthesised beam of a typical VLBI array is high enough, for example, to resolve individual supernova remnants in local ULIRGs e.g.\\ Arp~220 \\citep{lonsdale2006} and IRAS~17208-0014 \\cite{momjian2006}. Examples of ULIRGs that are believed to contain an AGN include Arp~220 \\citep{downes2007} and Mrk~273 \\citep{carilli2000}.  Only one VLBI observation of a classical, high-redshift, optically-faint SMG \\citep{momjian2010} has been published to date, continuum emission from GOODS~850-3 having been detected in a tapered High Sensitivity Array (HSA) image. In this paper we present the first VLBI survey of SMGs, using the European VLBI Network (EVN) to search for ultra-compact cores in a sample of objects in the Lockman Hole. These data, with a synthesized beam size of $\\approx$30\\,milliarcsec (mas) {\\sc fwhm}, are the highest-resolution radio detections of SMGs ever achieved and have allowed us to put the tightest constraints to date on the brightness temperatures of a significant sample of high-redshift starbursts.  This paper is organised as follows: in \\S\\ref{sec:targets} we describe our target selection and in \\S\\ref{sec:reduce} give details of the observing strategy, data reduction and imaging. In \\S\\ref{sec:results} we present our high-resolution images of each SMG and describe each source in detail. In \\S\\ref{sec:discuss} we discuss our findings before presenting our conclusions in \\S\\ref{sec:conclude}.  Where necessary we have assumed a flat $\\Lambda$CDM cosmology of $\\Omega_{\\Lambda} = 0.73$, $\\Omega_{m} = 0.27$ and $H_0 = 71$~km\\,s$^{-1}$\\,Mpc$^{-1}$ \\citep{hinshaw2009}. ",
        "conclusions": "\\label{sec:conclude}  We have undertaken a program designed to uncover AGN radio cores in a sample of SMGs in the Lockman Hole, using very-long-baseline interferometry with the EVN. Our sensitive, high-resolution images have detected only two of our six targets, despite a strong selection bias in favour of compact emission. From their brightness temperatures, there can be little doubt that we are seeing radio emission from active nuclei. For the other four SMGs we place upper limits on the radio flux density of any compact component associated with an AGN and conclude that star formation is probably the dominant source of their radio emission."
    },
    "1004/1004.2532_arXiv.txt": {
        "abstract": "{}{}{}{}{}  \\abstract { } {The neutral oxygen resonance $\\lambda$1302\\AA\\ line can, if the optical depth is sufficiently high, de-excite by an intercombination transition at $\\lambda$1641\\AA\\ to a metastable state. This has been noted in a number of previous studies but never systematically investigated as a diagnostic of the neutral red giant wind in symbiotic stars and symbiotic-like recurrent novae.} {We used archival $IUE$ high resolution, and GHRS and STIS medium and high resolution, spectra to study a sample of symbiotic stars. The integrated fluxes were measured, where possible, for the O I $\\lambda$1302\\AA\\ and O I] $\\lambda$1641\\AA\\ lines. } { The intercombination 1641\\AA\\ line is detected in a substantial number of symbiotic stars with optical depths that give column densities comparable with direct eclipse measures (EG And) and the evolution of the recurrent nova RS Oph 1985 in outburst. In four systems (EG And, Z And, V1016 Cyg, and RR Tel), we find that the O I] variations are strongly correlated with the optical light curve and outburst activity. This transition can also be important for the study of a wide variety of sources in which an ionization-bounded H II region is imbedded in an extensive neutral medium, including active galactic nuclei, and not only for evaluations of extinction. } {} ",
        "introduction": "Symbiotic stars present the unusual situation of a nearly neutral, stable environment centered on a cool giant star, in which a hot source, along with its surrounding ionized region, is imbedded. The radius of the H II region is determined only by the mass gainer's effective temperature and luminosity, which in turn depend only on the accretion rate from the wind (or in the cases where a disk forms, from the flux distribution of the surrounding disk along with that of the underlying star). Since these can be separated using multiwavelength observations, and the incident spectra are simple at ultraviolet (UV) wavelengths (a hot white dwarf and/or an accretion disk continuum and emission line continuum), it is possible to model the formation of the spectrum comparatively easily. This is mainly because the wind of the companion red giant is at nearly its terminal velocity (see Vogel 1991, Pereira et al 1999)  and, even if structured by the orbital motion and hydrodynamic processes related to the accretion (e.g. Dumm et al. 2000; Walder et al. 2008) this happens on a length scale far larger than the gainer and its ionized zone.  In such an environment several radiative processes, not usually encountered under nebular conditions, are observable. Principal among these are fluorescence due to various scattering mechanisms. Accidental resonances account for much of the down-conversion of UV emission to emission in the optical and near infrared. Perhaps the best known are those Fe II and related ions that can be excited by UV resonance transitions of highly-ionized species, e.g. C IV and its coincidence with ground state multiplets of Fe II that de-excite through optical forbidden transitions (Johansson 1983, 1988). Raman scattering (e.g. Schmid 1989), a nearly coincident resonance process that produces broad, down-converted emission lines, is particularly spectacular in the symbiotics, the most notable lines being those of O VI $\\lambda\\lambda$6825, 7082\\AA\\ that are produced by the near coincidence of the resonant O VI doublet $\\lambda\\lambda$1031, 1037\\AA\\ and H Ly$\\beta$. There is, in addition, a process whereby the UV resonance line of a {\\it neutral} species can, by virtue of absorption in a surrounding neutral gas, produce both optical and UV emission lines through otherwise inaccessible forbidden transitions. This happens because the ionization potential of several neutral atoms, in particular oxygen, is slightly higher than that of hydrogen and can therefore form in the H II region along with those formed by recombination. Furthermore, resonant absorption by the neutral gas at energies significantly below the ionization limit can, if the optical depth is sufficiently large, lead to emission in alternate channels even in the resonant scattering case.  The O I spectrum is a case in point. The $\\lambda$1302 resonance line is one of the strongest emission features in the spectrum of late-type symbiotics. It forms in the H II region around the degenerate gainer since O I has a slightly higher ionization potential than neutral hydrogen. In addition to the ground state, the O I $^3$S$\\rm^o$ $-$ $^3$P ($\\lambda\\lambda$1302, 1304, 1305\\AA) multiplet 5 is connected to two long-lived states through emission at $\\lambda$1641\\AA\\ and $\\lambda$2324\\AA, both spin forbidden (intercombination) transitions (see Fig. 1), and their associated decay channels to the ground state. One such decay channel, $\\lambda$6300 \\AA, is well known from terrestrial auroral spectra. Also, the $\\lambda$1641 line has been used as a proxy measure of solar activity variability and its effect on the atmosphere (e.g. Bowers et al. 1987, see below). These lines are also well known from planetary nebulae (e.g. Feibelman 1997) and have been discussed in the literature for studies of interstellar extinction in Seyfert galaxies (Grandi 1983) and the determination of the oxygen abundance in cool stars ([O I] $\\lambda$6300, Nissen et al. 2002).  A difficulty presented by any neutral or singly-ionized resonance transitions is that the interstellar medium, possessing the same resonance transitions, is opaque along many lines of sight, especially for distances of several kiloparsecs that are typical of symbiotic and planetary nebular targets. This is exacerbated for cosmological distances where the intervening Ly$\\alpha$ forest potentially contaminates the whole redshift range from that of the host galaxy to nearly the local standard of rest. These systems should, therefore, present sufficient line of sight optical depths to produce detectable O I] emission. \\begin{figure} \\centering \\includegraphics[width=7.5cm, angle=270]{fig1.ps} \\caption{Grotrian diagram for the principal transitions of O I involving all levels up to 10$^5$ cm$^{-1}$. Each multiplet in the figure is labeled with the shortest wavelength (in \\AA) in the multiplet, the number of lines of the multiplet (in parentheses) and the total Einstein $A$ transition probability (in italics, units of s$^{-1}$).   The transition coincident with H Ly$\\alpha$ is shown as a long dash, while lines of the prominent decay path and the O I] $\\lambda$1641\\AA\\ line are presented in bold. For clarity, the transition $^5$P (99092 cm$^{-1}$) - $^3$S$^o$ (96225 cm$^{-1}$) ($\\lambda$34860 \\AA, $A$ = 1.1e01 s$^{-1}$) has not been included. Atomic data are taken from Ralchenko et al. (2008).  } \\label{spectra}% \\end{figure}  In a study of the UV spectra of the recurrent nova RS Oph during its 1985 outburst, Shore \\& Aufdenberg (1991) noted the presence of a transient emission line on the red wing of He II $\\lambda$1640 relatively early in the outburst and identified this as O I] $\\lambda$1641. This line was also identified by Aufdenberg (1993) in the $IUE$ spectrum of RR Tel. In a recent study of the 2006-2009 outburst of the S-type symbiotic star AG Dra, we discussed the variations of the optical spectra, concentrating on the optical Raman features (Shore, Wahlgren, Genovali, et al. 2010). This survey included an examination of archival material as well as optical high-resolution spectra. The absence of the [O I] $\\lambda$6300 line was noted but it was suggested that it would be worthwhile checking the existing archive of high resolution UV data for O I] $\\lambda$1641. This symbiotic is a special case, having a radial velocity of $-$144 km s$^{-1}$; any resonance line originating from the star is well shifted in wavelength with respect to its ISM components. In this paper we report on our search of the archives for the presence of O I] $\\lambda$1641, as well as other O I lines, in the spectra of symbiotic stars. As mentioned above, the O I spectrum originates in the vicinity of the red-giant star. Observations of the O I] line may prove to be useful diagnostics of the red-giant wind and its sources of excitation. Correlation of the UV lines with optical and near-infrared (near-IR) O I lines would therefore enable studies of symbiotic star properties and behavior in the absence of UV spectra. ",
        "conclusions": "Several fluorescence processes can lead to the population of levels that will ultimately lead to emission of O I] 1641 \\AA. Coincidence of H Ly$\\beta$ $\\lambda$1025.722 with O I $\\lambda$1025.762 will populate the O I 3d $^3$D$^o$ (97488 cm$^{-1}$) level from the ground level. The dominant decay chain (according to their Einstein transition probabilities, see Fig. 1) from the 3d $^3$D$^o$ level ($\\lambda\\lambda$11285 and 8446) will populate the 3s $^3$S$^o$ (76794 cm$^{-1}$) level, which subsequently decays through three channels ($\\lambda\\lambda$1302, 1641, 2324), two of which are commonly detected in symbiotic star spectra. In addition to the Ly$\\beta$ O I pumping, McMurry \\& Jordan (2000) identified CO emission fluorescently-excited by O I UV 2 resonance line emission near $\\lambda$1302 in the UV spectrum of $\\alpha$ Tau.  A second pumping mechanism for the O I 76794 cm$^{-1}$ level is He II $\\lambda$1640 for those stars which have a broad He II line. This is evident for RR Tel and EG And, and less so for RW Hya and AG Peg, in Fig. 2. Other possible pumping mechanisms that might lead to population enhancement of this O I level are: 1) H Ly$\\epsilon$ $\\lambda$937.803 \\AA\\ coincident with O I $\\lambda$937.841 to pump the O I 106765 cm$^{-1}$ level, which can decay to the 76794 cm $^{-1}$ level through the chain $\\lambda\\lambda$14110, 4368 \\AA, among others. 2) H Ly$_6$ $\\lambda$930.748 and He II $\\lambda$930.342 can pump the O I 8s $^3$S$^o$ 107497 cm$^{-1}$ level, which decays to the 76794 cm$^{-1}$ level through the chains $\\lambda\\lambda$12790, 4368 or $\\lambda\\lambda$5298, 8446. 3) C II $\\lambda$2324.69 emission is coincident with O I $\\lambda$2324.738 and can pump the O I 76794 cm$^{-1}$ level from the metastable O I level 33792 cm$^{-1}$. The five lines of C II multiplet 2 are seen in emission in a number of symbiotics and symbiotic novae (RR Tel). Direct excitation by the resonance line should, however, be more effective in symbiotics, as in the terrestrial case since the Ly lines are so optically thick and the illumination is from the companion, not {\\it in situ} from the chromosphere (there will, of course, be a contribution from the spectrum of the late-type component but this is small compared to that from white dwarf environment).  Population of the O I 76794 cm$^{-1}$ level via electron recombination is possible through additional decay chains. Spectral observations at infrared wavelengths may offer a means of determining the dominant excitation mechanisms by detecting other emission lines. The number of lines from the O I spectrum that have been observed in astronomical targets, in particular symbiotic stars and novae, are few. Common UV lines detected include transitions at wavelengths $\\lambda\\lambda$1302, 1304, 1305, 1355, 1358, 1641. At optical wavelengths $\\lambda\\lambda$6300, 6363 are found in planetary nebula spectra  with $\\lambda$6300 commonly used for abundance analysis in cool stars. For near-IR wavelengths, detections, or suspicions of detections, have been mentioned for 8446 \\AA\\ in AG Dra (Iijima et al. 1987), and $\\lambda$11289 (Evans et al. 2007). Absorption lines at $\\lambda\\lambda$7771, 7773, 7774 are commonly used in abundance analysis in a variety of stars. There is also the curious appearance of an undiscussed weak emission line near 2.9 $\\mu$m (Schild, Boyle \\& Schmid 1992) in spectra of several symbiotics. Conspicuous by their absence from discussion and published spectra are lines of large transition probability ($\\lambda\\lambda$4368, 9204, 9260 for example) and small transition probability ($\\lambda\\lambda$1727, 2324, 2958, 2972). A full accounting of O I lines for any target would be useful for determining the excitation conditions and better enable the physical modeling.  The importance of the O I] $\\lambda$1641 line for the symbiotics is as a possible tool for as long as ultraviolet spectroscopy is available. Oddly, although this line has been included in a number of identification lists at high resolution, it has never been exploited as a diagnostic for symbiotics or related systems. It has, however, been noted as a contributor to the energetics of AGN when the O I resonance line is sufficiently optically thick. Grandi (1983), in discussing reddening determinations for AGN using the resonance and Bowen fluorescence O I UV2 lines ($\\lambda\\lambda$1302, 1304 vs $\\lambda$8446) noted that the line ratio O I] $\\lambda$1641 (UV146) to the resonance multiplet is often unusually large, given the branching ratio. This can be accounted for by a large enough optical depth to strongly self-absorb the ground state lines. With a Ly$\\alpha$ optical depth as large as $10^6$ the reduction in $\\lambda$1302 is sufficient to produce an integrated flux of only a factor of 2 larger than the forbidden transition. The inhomogeneous regions around the central engine often show such large opacities while still permitting observation of the nucleus along a given sight line. More recently, the chromospheric O I spectrum has been rediscussed for a few main-sequence and evolved F, G, and K stars by Koncewicz (2005) and Koncewicz \\& Jordan (2007). There is another mode to produce the O I] $\\lambda$1641 emission, the coincidence of Ly$\\beta$ $\\lambda$1025.72 and the O I resonance line (UV4) at $\\lambda$1025.77 that pumps the $^3$D$^o$ 97488 cm$^{-1}$ level , which then decays through the $\\lambda\\lambda$11286, 8446, 1641, 6300 chain.  To date, however, most of the literature deals with the O I lines in the context of planetary -- specifically, terrestrial -- atmospheric structure and composition. Atomic oxygen forms in excited states by dissociative collisions between O$_2$ and electrons. These transitions have been used for studying the oxygen abundance and temperature structure of the troposphere in a number of papers, e.g. Meier \\& Conway (1985), and Conway et al. (1988). Doering \\& Gulcicek (1989) include the $\\lambda\\lambda$1355, 1358 lines. Since the O I] transition is always optically thin and absent in the reflected solar spectrum this transition probes almost the entire terrestrial stratosphere and ionosphere. The branching ratio (Garstang 1961, Erdman \\& Zipf 1986) is O I] $\\lambda$1641/O I $\\lambda$1302 = 5.1$\\times$10$^{-6}$ with an uncertainty of 30\\% . The most recent compilation, Wiese et al. (1996), gives $A$($\\lambda$1641)/$A$($\\lambda$1302) = 5.4$\\times$10$^{-6}$. Following Grandi (1983), based on escape probability formalism (Kwan \\& Krolik 1981), we can estimate the required column density in the resonance line. The observed branching ratio is $\\approx$ 1 for all systems in which the O I] 1641\\AA\\ line is detected in our survey, the implied optical depth for $\\lambda$1302 is $\\tau_{\\lambda1302} \\approx 7\\times 10^{-14}f v_{50}^{-1} N_O \\approx 2\\times10^5$, where $f$ is the oscillator strength and $v_{50}$ is the wind velocity in km s$^{-1}$, that  for a solar O/H ratio ($5\\times 10^{-4}$, see Asplund et al. 2009) gives a column density N$_H$ $\\approx$ 10$^{23}$ cm$^{-2}$ for the neutral absorption region.  This is the same order of magnitude as the column density in absorption required to explain the narrow UV emission line variations during the early RS Oph outburst (Shore et al. 1996) and similar to that derived by Crowley et al. (2008) from eclipse spectra of EG And. Using the length scale from the photoionization modeling in Crowley et al. (2008), who obtain a standoff distance for the neutral region from the white dwarf in EG And of about 10$^{13}$ cm, gives a characteristic number density of about 10$^{10}$ cm$^{-3}$. For a wind velocity of 50 km s$^{-1}$ and using R$_{13}$ =(R/10$^{13}$ cm) gives an estimate of the mass loss rate for the red giant of $\\approx$ 10$^{-6}$ R$^2$$_{13}$ M$_{\\sun}$ yr$^{-1}$. This estimate is different between systems (there are several with lower branching ratios, others with higher, and in many cases the stellar radial velocity is not sufficient to displace the $\\lambda$1302 line from within the interstellar absorption. The optical depth expected for the red-giant wind in the FUV O I doublet, combined with the opacity of the Ly$\\beta$ transition, suggest that this is not a dominant mechanism in producing the $\\lambda$1641 line and that the O I opacity suffices."
    },
    "1004/1004.1657.txt": {
        "abstract": "Solar activity and helioseismology show the limitation of the standard solar model and call for the inclusion of  dynamical processes in both convective and radiative zones. In this paper, we concentrate on the radiative zone. We first recall  the sensitivity of boron neutrinos to the microscopic physics included in solar standard and seismic models. We confront the neutrino predictions of the seismic model to all the detected neutrino fluxes.  Then we compute new models of the Sun including a detailed transport of angular momentum and chemicals due to internal rotation that includes meridional circulation and shear induced turbulence. We use two stellar evolution codes: CESAM and STAREVOL to estimate the different terms.  We follow three temporal evolutions of the internal rotation which differ by their initial conditions:  very slow, moderate and fast rotation, with magnetic braking at the arrival on the main sequence for the last two. {bf We find} that the meridional velocities in the present solar radiative zone  is extremely small in comparison with those of the convective zone (smaller than 10$^{-6}$ cm/s instead of m/s). All models lead to a radial differential rotation profile in the radiative zone but with a significantly different contrast. % We compare the %result to the observed solar internal rotation and show that this work %favors a slow rotation of the young Sun. We compare these profiles to the presumed solar internal rotation and show that  if meridional circulation and shear turbulence  were the only mechanisms transporting angular momentum within the Sun,  a rather slow rotation in the young Sun  is favored. We confirm the small influence of the transport by rotation on the sound speed profile but its potential impact on the chemicals in the transition region between radiation and convective zones. These models are physically more representative of the real Sun than the standard or seismic solar models but a high initial rotation, as it has been considered previously, increases the disagreement with neutrinos and the sound speed in the radiative zone. This present work pushes us to pursue the inclusion of the other dynamical processes  to better reproduce the observed solar profile in the whole radiative zone and  to better describe the young active Sun. We also need to get a better knowledge of solar gravity mode splittings to use their constraints. ",
        "introduction": "Evidence for the presence of dynamical processes exist in stars all over the Hertzsprung-Russell diagram. In the case of the Sun, rotation and magnetic activity have been confirmed for more than 4 centuries. Rapid rotation is found mainly in young stars and intermediate-mass to massive stars while mild to slow rotation is found in low-mass stars and in giant stars. \\cite{Kraft} was the first to study in details the projected rotational velocity $\\rm v sin_i$ of young stars in clusters and shows the transition between rapid rotation for early-type stars and slow rotation for late-type stars. Then, Weber explained that the angular momentum loss is due to a magnetized wind in solar-like stars.   In studying Pleaides, Ursa  Major and Hyades,  \\cite{Skumanich} used solar and cluster data to establish the law of variation of the rotation with age, this simple law is in fact more complex  and has been studied in great details for different masses \\citep{Stauffer,Bouvier94,MaederMeynet00a,MaederMeynet04,Bouvier09}.  Rotation affects the internal stellar structure and evolution both directly via the modification of the gravitational potential, and by means of associated transport processes.  It has a direct impact on the shape of stars, that can be directly probed with interferometric data for the strongly deformed fast rotating stars \\citep{Domiciano03,McAlister05,vanBelle06,Kervella2006,Domiciano2008}. In the case of the Sun, the effect is small  as shown in \\cite{Piau} but of considerable importance (Zahn 2009; Rozelot 2009) to check the influence of the deep core rotation and of the deep magnetic field \\citep{Duez09, Duez10} on the tachocline and on the solar surface.  The standard solar and stellar evolution models do not take into account the effects of such dynamical processes as rotation or magnetic fields. If helioseismology required the improvement of the solar model and motivated the introduction of microscopic diffusion, going beyond the zero-order model, it is the far too rough agreement between observational data and theoretical predictions for other stars that has led to progressively introduce the dynamical physical processes likely to affect the transport of momentum and chemicals in the stellar evolution codes.  Using either a purely diffusive approach (\\cite{Endal, Pinsonneault, Chaboyer, Langer, Heger00}) or the more complex formalism developed by \\cite{Zahn92} and \\cite{MaederZahn98}, the introduction of the rotation-induced transport of angular momentum and chemicals to model both massive and low-mass stars improves the comparison with observations \\citep[see e.g.][]{MeynetMaeder00,TalonCharbonnel98,Palacios03}. These results encourage us to pursue this effort for the Sun to get a proper interpretation of the existing helioseismic observations and the coming asteroseismic ones.   Acoustic and gravity mode detections provide a unique insight on the internal solar sound speed, density and rotation profiles \\citep{Kosovichev,Turck2001a,Thompson,Couvidat,Turck04a,Mathur,Mathur2}. This gives us a unique opportunity to validate the complex formalism introduced in rotating stellar evolution models.  This information justifies the development of a Dynamical Solar Model to be confronted with the helioseismic and neutrino probes. The solar status must improve our knowledge of solar-like stars where only external stellar rotation rates or abundance anomalies are used to confront theoretical assumptions to observational facts. Moreover, building a consistent dynamical evolution model for the Sun will largely contribute to establish a complete and consistent MHD representation of the Sun and a good connection between internal and external magnetism.  The inner solar rotation is the first evidence of the internal dynamics that needs to be understood. This is not an easy task because the present-day profile results from the interplay between several distinct processes \\citep{ZahnTalonMatias97,Chaboyer,Eggenberger05,CharbonnelTalon05} and their influence depends on the adopted theoretical prescriptions.  The most recent of these studies include the transport of angular momentum by magnetic field  or internal gravity waves, in addition to the ``purely'' rotational transport by meridional circulation and shear-induced turbulence, to obtain a more efficient extraction of angular momentum and a relative flat rotation profile in the radiative interior in better agreement with helioseismic data.  Actually, both magnetic field and internal gravity waves probably contribute to the transport of angular momentum in stellar interiors, yet no models have been published up to now including all the processes. Moreover, the way to account for the effect of magnetic fields in 1--D stellar evolution codes is still puzzling and a matter of debate, and the introduction of gravity waves produces some irregularities in the radiative rotation profile \\citep[see e.g.][]{CharbonnelTalon05} in disagreement with the helioseismic results.   In the present paper, we limit ourselves to the sole effect of rotational transport and carefully study the hypotheses and the order of magnitude of the terms that we introduce for the inner rotation. This approach pushes further the previous works by \\citet{ChaboyerZahn92,Chaboyer,CharbonnelTalon05,Eggenberger05,Yang1,Yang2}. We focus our analysis on the solar core and we compare our results with all the recent existing seismic and neutrino indicators.  This paper is a first of a series where we will discuss in details the influence of rotation, magnetic field and internal waves on the solar dynamical model.  We do not  present and discuss  the abundances of lithium and beryllium as it has been done extensively by \\cite{Pinsonneault}. Indeed we have shown in our previous works on the tachocline \\citep{Brun2} and on young stars that  the abundance of these elements are sensitive to the early evolution as well as to the presence of the tachocline.  These studies lead us to the conclusion that we need to treat properly the magnetic field in the early phase to better estimate the lithium evolution during this phase \\citep{Piau}.  In Section 2, we present the status of the solar classical models and compare the seismic model predictions to all the present detected neutrinos. Then in Section 3 we recall the formalism proposed by \\citet{Zahn92} and \\citet{MaederZahn98}, and slightly improved by \\citet{Mathis1} that we use to introduce the rotation effects in the stellar evolution codes. In that section, we also give a brief description of the CESAM and STAREVOL stellar evolution codes used in the paper, and present the three different types of solar models which differ by their initial rotation and the presence of magnetic braking. In Section 4 we compare the results of these three rotating solar models obtained with both codes. These models are compared with seismic observations and neutrino detections in Section 5.  In the last section, we summarize the important points in a more general context. ",
        "conclusions": "In this paper we have shown the following facts:  - We have examined  three initial rotation rates (models A, B and C), choosing initial angular momentum contents corresponding to 2.5 km/s, 20 km/s and 50 km/s at the ZAMS.  The last two  values  have also been adopted in other studies of the rotating Sun and solar-type stars \\citep{Yang1,Eggenberger05,TalonCharbonnel2005,Chaboyer}.   At 4.6 Gyrs, the three cases show a radial gradient of rotation in the core. %and we compare the transport of angular momentum and %chemicals due to the solar rotation history in two evolutionary %codes. Its amplitude  depends strongly on the initial rotation rate. If this one is small, the radial gradient is established during the contraction phase and is amplified during the subsequent evolution up to the present Sun.  For models rotating faster, one notes much steeper gradients at the age of the Sun, the angular momentum losses associated with magnetic braking on the ZAMS are responsible for the rather small decrease on the ZAMS.   - The transport of angular momentum in the solar radiative zone during the main sequence appears extremely small  and the meridional circulation in the radiative zone is smaller by 10 orders of magnitude in comparison with the observed convective meridional circulation velocity at 99 \\% R$_\\odot$. This process is even slower than the microscopic diffusion (gravitational settling) that we use in the radiative zone. As a first consequence such implementation in a stellar evolution code is delicate mainly because one needs to solve four coupled equations with derivation of quantities that practically do not vary.  In order to settle our conclusions concerning the form and order of magnitude of the different quantities associated to the rotation-induced mixing, we have  confronted models obtained with two different codes, CESAM and STAREVOL, using distinct numerical approaches and the old composition of Grevesse and Noels. We have shown, for models A and B, that they lead to rather similar results and the same kind of profile for the present Sun, thus validating the results and numerical approaches. Such a very large difference between the meridional velocity in the radiative zone ($10^{-7}- 10^{-6}$ cm/s) and in the convective zone ( m/s) would naturally produce some turbulent hydrodynamical layer called and this is an interesting result of the present calculation.  - Although the combined effect of meridional circulation and shear-induced turbulent associated to rotation is small, this study allows us to present radial rotation profiles that can be directly compared to the seismically observed one.  This study sustains the idea that the Sun was not  a rapid rotator after the contraction phase. The angular velocity profile we get for models A is not far from the presumed solar one in the core. Of course  this model should not be considered at face value since it does not take into account any magnetic braking generally observed in young clusters, but it can be useful to estimate the interplay between processes. The second model (moderate rotation) is more realistic but shows a greater steep gradient in disagreement with the published detection of gravity modes. So one can imagine that the Sun has been in an intermediate case arriving at 5-10km/s at ZAMS and that its rotation profile would have being eroded by some other process.  - Let us stress however that other dynamical processes known to generate efficient transport of angular momentum such as magnetic fields and internal gravity waves were not yet included in our models. The inclusion of these processes is in the scope of a further study, and is expected to significantly affect the choice of the preferred model. For instance, internal fossil magnetic field produces certainly a very small effect on the solar structure \\citep{Duez09} but may lead to efficient transport of angular momentum that would help flattening the angular velocity profile in the radiative zone.  But before we would like to study the activity of the very young Sun. Previous works  show how magnetic field can modify the rotation profile \\citep{Eggenberger05, Denissenkov07} but the action could be improved by the introduction of a more sophisticated field topology which preserves the stability of such a field \\citep{Duez09b}.  The understanding of the magnetic field during the contraction phase has probably a crucial impact on the radial rotation gradient and deserves premature estimate of lithium and beryllium destruction at this stage.  - Other processes may modify the present conclusions. For example in the present treatment we neglect the fundamental role of the tachocline which must be established at least since the arrival on the main sequence. The hydro (or magnetohydrodynamical) nature of such a region may alter the angular momentum and/or chemicals tranport due to the internal rotation but we have already noticed that a crude treatment of this region might slightly amplify the present tendancy on chemicals gradient and structure effect \\citep{Brun2}.  In the present study, we note that if the initial rotation is accompanied by an efficient magnetic braking, it generates some turbulent flow at the base of the convective zone which smoothes the helium profile and increases its abundance at the surface.  - We note that the impact of the rotation on the solar structure is rather small. As the transport of angular momentum and chemicals goes from the radiative zone to the convective zone, it implies a slight reduction of the central temperature and of the helium content in the radiative zone. Consequently it increases slightly the present discrepancy between model and the observed sound speed. In fact, one cannot exclude other momentum transport which may come from the convective zone and play the inverse role \\citep{Garaud08,Gough09}. All these other processes must be included in stellar evolution codes before getting a proper DSM.  We see in this study that the description of the dynamics of the PMS phase (especially the related magnetic field evolution) will be a crucial issue too. In the comparison between CESAM and STAREVOL we have noticed some non negligible difference during the contraction phase for models with a moderate (the same for high) initial rotation rate, this phase illustrates the sensitivity of the rotation gradient to the numerical details and to the magnetic braking.  Moreover we show in this study that independently of the initial rate, the central rotation value does not change by more than 50\\% during the main sequence. So the final comparison of the DSM to the observations will partly  depend on the way we shall treat the contraction phase, the inner corresponding magnetic field evolution and the magnetic braking phase. The same conclusion was already reached discussing the problem of lithium in young stars and in the Sun \\citep{Brun2, Piau}.  - The sound speed profile is practically unchanged when the slow initial rotation is assumed.  So even if  this model is probably more representative of the dynamics of the solar interior than the standard model, it is still an incomplete model and its predictive character remains limited. For this reason, and considering the consequent discrepancy on the sound speed  predicted by the present standard model and the observed one, which will be  amplified by the recent CNO composition, we continue to recommend the seismic model for any global predictions (gravity modes or neutrinos).  We have shown in this paper that the predictions of the seismic model are in very good agreement with all the neutrino detections including BOREXINO."
    },
    "1004/1004.2704_arXiv.txt": {
        "abstract": "The decay of non-topological electroweak strings formed during the electroweak phase transition in the early universe may leave an observable imprint in the universe today. Such strings can naturally seed primordial magnetic fields. Protogalaxies then tend to form with their axis of rotation parallel to the external magnetic field, and moreover, the external magnetic field produces torque which forces the galaxy axis to align with the magnetic field, even if the two axis were not aligned initially. This can explain an (observed, but as of yet unexplained) alignment of the quasars' polarization vectors. We demonstrate that the shape of a magnetic field left over from two looped electroweak strings can explain the non-trivial alignment of quasar polarization vectors and make predictions for future observations. ",
        "introduction": "Recently, Hutsem\\'{e}kers \\cite{Hutsemekers} made two interesting observations on a sample of 355 quasars\\footnote{In this paper we use the term 'quasar' to describe both optically and radio selected quasi-stellar objects.}. They observed that the polarization vector of quasars appear to be i.) somewhat aligned over large (cosmologically interesting) volumes of space and ii.) the angle of these vectors seem to rotate coherently with increasing redshift. As discussed in their paper, these two observations seem unlikely to be attributable to either natural contamination such as intervening dust particles or unaccounted instrumental bias. Instead, the effect appears to be cosmological.  The direction of the optical polarization vector can be attributed to the physical orientation of the quasar itself \\cite{Elvis,BorguetQuasarPolarizationAxisDirection,BorguetSmallPaper}. We propose that the quasars themselves are somewhat aligned on cosmological scales. Any model that explains the coherent alignment of quasars on such scales should also address the rotation through $\\sim 240^{o}$ that is observed in the sample. This feature cannot be easily accommodated in generic models.  We propose that the orientation of these quasars is caused by a magnetic field left over from two linked loops of electroweak strings. From the time of the electroweak phase transition to today, magnetic field lines seeded by these strings are stretched by the expansion of the universe and act as a background magnetic field at the time of quasar formation. We fit the alignment data and find that the electroweak string loops can explain this alignment very well. We emphasize that our explanation is based on known and pretty well understood physics of the standard model and its embedded defects like electroweak strings. ",
        "conclusions": "Electroweak strings are predicted to exist in the early universe. Although searching for stable strings in the universe today is a difficult endeavor (as many models predict $\\sim 1$ will exist in a given horizon volume), the potential to observe the imprint left over from an electroweak string remains an intriguing possibility. Linked strings may leave behind lines of magnetic flux imprinted in the universe which could be stretched to cosmological scales by both the expansion of the universe and by the fact that parallel lines of magnetic flux repel. Quasars that form in the vicinity of these magnetic fields are essentially forming in a background magnetic field. The quasars would therefore preferentially form with their axes aligned parallel to the magnetic field. The other effect that synergetically works with this is that the external magnetic field tends to align the quasar axis with itself even if initially the two axis were not aligned. Other nearby quasars will also be forming in essentially the same average background magnetic field which could explain the observed alignment of quasar polarization vectors. On large enough scales, however, the effects of the two looped magnetic fields would be observable as a rotation of the average direction of quasar polarization vectors.  The agreement between our theoretical model and the observational data is very good. In particular we were able to explain the rotation of the polarization angle with the redshift, a feature which is not easily accommodated in simple adhoc models. Our model gives clear predictions that can be tested once a greater sample size of quasar polarization data is available, since we predict an overall trend of quasar polarization vector behavior based on the model given by Eq.~(\\ref{TwoLoops}). The A1-A3 axis seems to lie somewhat along the line connecting the NGP and SGP. We would expect other quasar polarization angles in this region to follow the same pattern as observed in Fig. (\\ref{A1A3ErrorBarsAndBestFitLines}). Alternately, we may also look for quasar polarization angles away from the A1-A3 axis that still follow the pattern as predicted by Eq.~(\\ref{TwoLoops}). This observation would likely be somewhat more difficult, as this would require observing a large sample of objects through the galactic disk. Another interesting test of our model would be to look for the systematic effects such a magnetic field configuration would have on CMB photons such as Faraday rotation.  We emphasis that our explanation of the observed large scale alignment of quasars' polarization angles is based on conventional cosmology and minimal standard model, without invoking any exotic physics or non-standard cosmology. In particular, the formation and subsequent decay of the electroweak strings, and their seeding of the primordial magnetic fields can not be avoided. In this paper we just inked this fact with the large scale alignment of quasars' polarization angles."
    },
    "1004/1004.1131.txt": {
        "abstract": "{} {Analyse the distribution of matter around the progenitor star of gamma-ray burst GRB 021004 as well as the properties of its host galaxy with high-resolution echelle as well as near-infrared spectroscopy. } {Observations were taken by the 8.2m Very Large Telescope with the Ultraviolet and Visual Echelle spectrograph (UVES) and the Infrared Spectrometer And Array Camera (ISAAC) between 10 and 14 hours after the onset of the event.} {We report the first detection of emission lines from a GRB host galaxy in the near-infrared, detecting H$\\alpha$ and the [O {\\small  III}] doublet. These allow an independent measurement of the systemic redshift ($z=2.3304\\pm0.0005$) which is not contaminated by absorption as the Ly$\\alpha$ line is, and the deduction of properties of the host galaxy. From the visual echelle spectroscopy, we find several absorption line groups spanning a range of about 3,000 km s$^{-1}$ in velocity relative to the redshift of the host galaxy. The absorption profiles are very complex with both velocity-broadened components extending over several 100 km s$^{-1}$ and narrow lines with velocity widths of only $\\sim$ 20 km s$^{-1}$. By analogy with QSO absorption line studies, the relative velocities, widths, and degrees of ionization of the lines (``line-locking'', ``ionization--velocity correlation'') show that the progenitor had both an extremely strong radiation field and several distinct mass loss phases (winds).} {These results are consistent with GRB progenitors being massive stars, such as Luminous Blue Variables (LBVs) or Wolf--Rayet stars, providing a detailed picture of the spatial and velocity structure of the GRB progenitor star at the time of explosion. The host galaxy is a prolific star-forming galaxy with a SFR of $\\sim$40 M$_{\\odot}$yr$^{-1}$.} ",
        "introduction": "The afterglows of long-duration Gamma-Ray Bursts (GRBs), which are linked with the explosions of massive stars \\citep[see][for a recent review]{woo06}, are the most luminous optical sources in the universe for short periods of time \\citep{kan07,blo09}. Low-resolution optical spectroscopy was initially only usable to determine the redshift and thus place them at cosmological distances \\citep{met97}. Deeper insight came with the first medium-resolution spectrum, obtained with Keck ESI of the afterglow of GRB 000926 \\citep{cas03}. The first true high-resolution echelle spectra were obtained for GRB 020813 \\citep{fio05}, but they were of low signal-to-noise ratio. The first echelle spectra with good S/N were finally obtained for GRB 021004, the focus of this work \\citep[see also ][]{fio05}.  Such spectroscopy allows deep insight into the environments of GRBs \\citep{pro06}. Some highlights include the possible detection of a Galactic superwind in the host galaxy of GRB 030329 \\citep{tho07}, and variable absorption lines which result from direct UV pumping by the luminous GRB afterglow \\citep{des06, vre07, del09}. These detections have recently been possible for more afterglows due to the rapid localization capabilities of the \\emph{Swift} satellite \\citep{geh04} in combination with the rapid-response mode (RRM) which is now available for the Ultraviolet-Visual Echelle Spectrograph (UVES) at the Very Large Telescope (VLT) \\citep[e.g.,][]{vre07,del09}. Recently, the covered wavelength region has been expanded all the way from the ultraviolet into the $K$ band near-infrared (nIR) by the second-generation instrument X-Shooter at VLT \\citep[][see \\citealt{deu09} for a first result]{dod06}.  GRB 021004   was detected at   12:06:14 universal time  (UT)  on 2002, October 4 with the gamma-ray  instrument FREGATE, the wide-field X-ray Monitor (WXM) and  the soft X-ray  camera (SXC) aboard the High-Energy Transient Explorer (\\emph{HETE-2}) \\citep{shi02}. \\object{GRB~021004}, was  a moderately  bright, long-duration  ($T_{90}=100$ s) event  with fluences  of   $6.4\\times10^{-7}$   erg   cm$^{-2}$ (7-30     keV) and $2.3\\times10^{-6}$ erg cm$^{-2}$ (30-400 keV) \\citep{bar02}.  The GRB was rapidly localized in flight and  the position was reported in  less than a minute,  allowing   rapid  ground-based follow-up  which revealed the  presence of the fading optical   afterglow of GRB 021004 \\citep{fox03}.   This    prompted  follow-up  observations   at   many observatories, which led to  an extensive long-term coverage  at X-ray \\cite{sak02a, sak02b}, radio \\citep{fra02, ber02, poo02a, poo02b}\\footnote{See http://www.aoc.nrao.edu/$\\sim$dfrail/grb021004.dat for the complete VLA data set.}, millimeter \\citep{deu05}, near-IR  \\citep{deu05,fyn05} and      optical    wavelengths,          both            ground-based \\citep{fox03,ber03,uem03,hol03,pan03,mir03,kaw04,deu05}  as  well   as space-based \\citep{fyn05}.  The isotropic  energy release during  the prompt emission of  this GRB was    modest,    with   log    $E_{iso}=52.65^{+0.12}_{-0.17}$ \\citep{kan10}, and the dust extinction  for this somehow reddish afterglow (R-K $\\sim$ 3) was also low as is typical for many well-observed  afterglows \\citep{kan06}.  Despite  the low energy promptly released,  this is  among the  most luminous  afterglows  ever detected \\citep{kan06}, even in comparison  with a much larger \\emph{Swift}-era sample  \\citep{kan10}. The multiwavelength  temporal evolution  of the GRB 021004  afterglow can be  explained by multiple  energy injections \\citep{bjo04, deu05},  nevertheless  other  scenarios  can not  be  discarded \\citep{laz02}.   GRB  021004 remains  one  of  the most  well-observed afterglows ever.  Early-time  low  and   medium  resolution  spectroscopic  observations allowed a  redshift  of  $z=2.33$ to  be  determined on  the  basis of Ly-$\\alpha$ absorption   and  emission  lines \\citep{cho02}.   Several absorption systems with outflows velocities  of  few 1000 km  s$^{-1}$ were   also  reported  \\citep{sal02,sav02} and  studied  in  detail by \\citet{mol02b}, \\cite{wan03}, \\citet{mat03}, \\citet{sch03}, \\citet{mir03}, \\citet{sta05} and \\cite{laz06}. \\cite{jak05} also report the detection of the host galaxy Ly$\\alpha$ line in narrow-band imaging.  \\citet{fio05} have constrained the ionization parameters of the various absorption  components detected above Ly$\\alpha$ (at an observer  restframe of 4,050  \\rm \\AA) in GRB 021004,  interpreting them  as density fluctuations, like \\citet{laz02} did. Within the context of larger samples, the UVES spectra of this GRB have also been studied by \\cite{che07} (the general lack of wind signatures in high-resolution spectra of GRB afterglows), \\cite{pro08} (N {\\small V} absorption lines toward GRB afterglows), \\cite{fox08} (high-ionisation line systems toward GRB afterglows) and \\cite{tej07, tej09, ver09} (study of Mg {\\small II} foreground absorption systems).   The original expectation, as detailed high quality spectroscopy of GRB afterglows became possible, was that we would rapidly learn much about the GRB progenitors via the study of the complex absorption line systems they were expected to exhibit due to ejection events leading up to the final collapse. This expectation has far from proven true.  Despite the large number of long GRB afterglow spectra obtained since, GRB 021004 still stands out as the one with the most complex set of intrinsic (0-3,000 km/s ejection velocity) absorption systems. At the time it also held the place as the lowest detected HI column density, and it still ranks between the lowest seven found. The complexity originally suggested that it indeed represented a display of ejecta \\citep{mol02a, mir03, fio05} but this interpretation was later disputed by \\citet{che07}. A final interpretation of the complex systems has not yet been agreed on, and this object still stands our as the best candidate of a seemingly rare case where the signatures of events prior to the collapse are displayed. The rarity alone would warrant a more detailed discussion of the spectral features. In addition we are adding UVES data below 4050 ang for a more detailed discussion of the absorption systems (para 3.1) and near-IR VLT/ISAAC spectroscopy allowing a more accurate determination of host properties and redshift (para 3.2). We discuss our results in the light of progenitor models for GRBs in \\kref{Discussion}, and summarize the work in \\kref{Conclusions}.     \\begin{figure*}[t!] \\begin{center} \\resizebox{15.2cm}{!}{\\includegraphics[clip]{021004_FIG1.eps}}  \\caption{Overall   view of    GRB  021004    optical   afterglow spectrum. These VLT/UVES data were obtained $\\sim$0.6 days after the GRB  and show the Ly$\\alpha$  emission line arising from the host galaxy redshift plus some  of the most prominent absorption lines systems at (or very  close  to) the host galaxy  redshift: low ionization  lines (like Fe  {\\small II}, Si {\\small II}) and high ionization lines  (like Si {\\small  IV},  C {\\small IV},  N {\\small    V}), all  labelled   in red   colour. Also shown  for completeness are  the   absorption  systems for  the  foreground systems at $z  = 1.6020$  (labelled in green)  and $z  = 1.3820$ (labelled in  blue). For clarity,  the  original data have  been smoothed. The telluric   lines in the  A-band (7600--7630 \\AA{}) are noticeable. The dotted line is the error spectrum.} \\label{uves spectrum} \\label{SpectrumFull} \\end{center} \\end{figure*}  For a Hubble constant of H$_{0}$ = 72 km s$^{-1}$ Mpc$^{-1}$, a matter density $\\Omega_{m}$ = 0.3, and a cosmological constant $\\Omega_{\\Lambda}$ = 0.7, the luminosity distance to the host is d$_{L}$ = 18.22 Gpc, and the look-back time is 10.42 Gyr. All errors are given at a $1\\sigma$ level of confidence for a parameter of interest unless stated otherwise. ",
        "conclusions": "\\label{Conclusions}  We find several absorption line groups spanning a range of about 3,000 km s$^{-1}$  in velocity relative to the  redshift of the host galaxy. The absorption profiles are  very complex with both velocity-broadened components extending over  several 100  km  s$^{-1}$ and  narrow lines with velocity widths  of only $\\sim$ 20 km  s$^{-1}$.  By analogy with QSO absorption line studies, the    relative velocities, widths,   and degrees      of    ionization  of     the   lines   (``line-locking'', ``ionization--velocity correlation'')  show  that the  progenitor  had both an extremely  strong  radiation field  and  several distinct mass loss phases (winds). These results are consistent with GRB progenitors being massive  stars, such as  LBVs or  Wolf--Rayet stars, and provide further   insight into the   nature of  these   progenitors  and their immediate environments.  The host galaxy is a prolifically star-forming galaxy at a systemic redshift  $z=2.3304$,  with a SFR of $\\sim$40  M$_{\\odot}$yr$^{-1}$ as also found by \\citet{fyn05} and \\citet{deu05}, reinforcing the potential association of some GRB with starburst galaxies (\\citet{Chri04,gor05} and references there in).  The   {\\it Swift}  mission with a predicted lifetime of ten  years \\citep{geh04} will certainly  bring  us the opportunity to   carry out high-resolution spectroscopy for dozens of future GRBs, and to set physical/chemical properties common to all GRB outflows."
    },
    "1004/1004.2474_arXiv.txt": {
        "abstract": "We investigate cosmological scenarios of generalized Chaplygin gas in a universe governed by Ho\\v{r}ava-Lifshitz gravity. We consider both the detailed and non-detailed balance versions of the gravitational background, and we include the baryonic matter and radiation sectors. We use observational data from  Type Ia Supernovae (SNIa), Baryon Acoustic Oscillations (BAO), and Cosmic Microwave Background (CMB), along with requirements of Big Bang Nucleosynthesis (BBN), to constrain the parameters of the model, and we provide the corresponding likelihood contours. We deduce that the present scenario is compatible with observations. Additionally, examining the evolution of the total equation-of-state parameter, we find in a unified way the succession of the radiation, matter, and dark energy epochs, consistently with the thermal history of the universe. ",
        "introduction": "Recently  Ho\\v{r}ava proposed a power-counting renormalizable theory with consistent ultra-violet (UV) behavior \\cite{hor2,hor1,hor3,hor4}. Although presenting an infrared (IR) fixed point, namely General Relativity, in the  UV the theory exhibits an anisotropic, Lifshitz scaling between time and space. Due to these novel features, there has been a large amount of effort in examining and extending the properties of the theory itself \\cite{Cai:2009ar,Cai:2009dx,Nishioka:2009iq,Charmousis:2009tc,Li:2009bg,Visser:2009fg, Sotiriou:2009bx,Germani:2009yt,Chen:2009bu,Bogdanos:2009uj,Kluson:2009rk,Afshordi:2009tt,Myung:2009ur,Alexandre:2009sy, Blas:2009qj,Capasso:2009fh,Chen:2009vu,Kluson:2009xx,Kiritsis:2009vz,Garattini:2009ns,Kluson:2010aw,Son:2010qh,Carloni:2010nx,Eune:2010qk,Wang:2010mw,Gullu:2010wb}. Additionally, application of Ho\\v{r}ava-Lifshitz gravity as a cosmological framework gives rise to Ho\\v{r}ava-Lifshitz cosmology, which proves to lead to interesting behavior \\cite{Calcagni:2009ar,Kiritsis:2009sh}. In particular, one can examine specific solution subclasses \\cite{Lu:2009em,Minamitsuji:2009ii,Wu:2009ah,Cho:2009fc,Boehmer:2009yz,Momeni:2009au,Cai:2010ud,Huang:2010rq}, the phase-space behavior \\cite{Carloni:2009jc,Leon:2009rc,Myung:2009if,Bakas:2009ku,Myung:2010qg}, the gravitational wave production \\cite{Mukohyama:2009zs,Park:2009gf,Park:2009hg,Myung:2009ug}, the perturbation spectrum \\cite{Mukohyama:2009gg,Piao:2009ax,Chen:2009jr,Gao:2009ht,Cai:2009hc,Wang:2009yz,Kobayashi:2009hh,Wang:2009azb,Kobayashi:2010eh}, the matter bounce \\cite{Brandenberger:2009yt,Brandenberger:2009ic,Cai:2009in,Gao:2009wn, Saridakis:2011pk}, the black hole properties \\cite{Danielsson:2009gi,Cai:2009pe,Kehagias:2009is,Myung:2009va,Park:2009zra,BottaCantcheff:2009mp,Lee:2009rm,Kiritsis:2009rx,Greenwald:2009kp}, the dark energy phenomenology \\cite{Saridakis:2009bv,Park:2009zr,Chaichian:2010yi,Jamil:2010vr}, the observational constraints on the parameters of the theory \\cite{Dutta:2009jn,Dutta:2010jh}, the astrophysical phenomenology \\cite{Kim:2009dq,Harko:2009qr,Iorio:2009qx,Iorio:2009ek,Izumi:2009ry, Lobo:2010hv,Cardone:2010tb,Saridakis:2011eq}, the thermodynamic properties \\cite{Wang:2009rw,Cai:2009qs,Cai:2009ph,Wei:2010yw,Jamil:2010di} etc. However, despite this extended research, there are still many ambiguities if Ho\\v{r}ava-Lifshitz gravity is reliable and capable of a successful description of the gravitational background of our world, as well as of the cosmological behavior of the universe \\cite{Charmousis:2009tc,Li:2009bg,Sotiriou:2009bx,Bogdanos:2009uj,Koyama:2009hc,Papazoglou:2009fj,Kimpton:2010xi,Bellorin:2010je}.  Although the foundations and the possible conceptual and phenomenological problems of Ho\\v{r}ava-Lifshitz gravity and its associated cosmology are still an open issue, it is worth investigating different cosmological scenarios in this gravitational background. In this regard, it would be interesting to examine cosmological scenarios where apart from the gravitational sector there exists a generalized Chaplygin gas \\cite{kamenshchik,Bilic:2001cg,Bento:2002ps,Gorini:2002kf,Bento:2002yx,Bento:2003we,Bento:2003dj,Bertolami:2004ic}.  The generalized Chaplygin gas (GCG) is an alternative to the Quintessence model which had attracted great interest in recent times. The scenario can explain the acceleration of the universe via an exotic equation of state, which mimics  a pressureless fluid at the early stages of evolution of the Universe, and a cosmological constant at late times. It is therefore interesting to consider the GCG scenario as a unified description for dark matter and dark energy \\cite{Bento:2002ps}. The background evolution fits well the observational data \\cite{Bento:2002yx, Bento:2003we}, however the cosmological behavior is indistinguishable from that of the $\\Lambda$CDM scenario while fitting with the structure formation data as well as with data from Cosmic Microwave Background radiation \\cite{Sn,bean}. Additionally, the scenario is plagued by the presence of instabilities as well as oscillations which are not observed in the matter power spectrum. Although this is a serious drawback of GCG scenario it is not the final verdict for its fate as a model for the unified description of dark matter and dark energy. As it was shown by Reis {\\it{et al.}} in  \\cite {reis:2003}, allowing for small entropy perturbations can eliminate instabilities and oscillations in the matter power spectrum, even in the linear regime, for a region of the parameter space where the GCG model behaves quite differently from the $\\Lambda$CDM one. Furthermore, as it was shown by Avelino {\\it{et al.}} in \\cite{Large},  in the GCG scenario the transition from dark matter to dark energy behavior is never smooth, and hence the linear theory which was used by \\cite{Sn, bean} to rule out GCG as a unified model, may break down late in the matter-dominated era, even on large cosmological scales. Therefore, nonlinear effects should be necessarily taken into account when confronting cosmological observations. Moreover, the addition of baryons in the GCG scenario can also improve the behavior of the matter power spectrum \\cite{Large2}. In summary, the GCG as a unified model of dark matter and dark energy is not completely ruled out and it deserves further investigations. Finally, we mention that GCG in the presence of cold dark matter as well as baryonic matter (where GCG acts as a normal dark energy candidate) is one of the most well-fit scenarios with cosmological observations, amongst all the exotic models that have been considered so far \\cite{davis}.   In the present work we construct the cosmology of a generalized Chaplygin Gas in a universe governed by Ho\\v{r}ava-Lifshitz gravity. Additionally, we use observational data from Type Ia Supernovae (SNIa) \\cite{SNIadata}, Baryon Acoustic Oscillations (BAO) \\cite{BAOdata} and Cosmic Microwave Background (CMB) \\cite{CMBdata}, together with the Big Bang Nucleosynthesis (BBN) conditions, to constrain the various parameters of the model. Furthermore, in order to be general we perform the analysis both with and without the detailed-balance condition of the gravitational sector.  The manuscript is organized as follows: In section \\ref{model} we present Ho\\v{r}ava-Lifshitz cosmology in both its detailed-balance and beyond-detailed-balance version. In section \\ref{GCGHL} we construct the scenario of a generalized Chaplygin Gas in Ho\\v{r}ava-Lifshitz gravitational background and we extract the cosmological equations. In section \\ref{observations} we use observational data in order to constrain the various parameters of the scenario. In section \\ref{cosmimpl} we discuss the cosmological implications, focusing on the evolution of the total equation-of-state parameter of the universe and of the expansion rate. Finally, section \\ref{conclusions} is devoted to the summary of our results. ",
        "conclusions": "\\label{conclusions}  In this work we investigated the cosmological scenario of generalized Chaplygin Gas in a universe governed by Ho\\v{r}ava-Lifshitz gravity, both in the detailed-balance as well as in the beyond-detailed balance version of the theory. Furthermore, in order to obtain a realistic cosmology we have included the baryonic matter and standard-model radiation sectors.  After extracting the cosmological equations, we used data from SNIa, BAO and CMB observations, as well as arguments from Big Bang Nucleosynthesis, in order to impose constraints on the various parameters of the scenario. For the  detailed-balance case we have three free parameters, namely the exponent $\\beta$ of the generalized Chaplygin Gas, its present equation-of-state parameter value $A_{s}$, and the effective neutrino parameter $\\Delta N_{\\nu}$ which quantifies the total amount of Ho\\v{r}ava-Lifshitz dark-radiation and kination-like components allowed during BBN. The best fit values of the parameters together with the corresponding $1\\sigma$ confidence intervals, arisen from the likelihood analysis for open and closed universe, were shown in Table \\ref{db_bestfits}, while in Fig. \\ref{dbcontours} we presented the corresponding likelihood contours. We deduced that the scenario at hand is compatible with observations, however the data lead to strong bounds on $A_{s}$ and $\\Delta N_{\\nu}$ and the spatial curvature $\\Omega_{K0}$. Finally, the scenario at hand depends only slightly on $\\beta$.  In the beyond-detailed-balance scenario the free parameters are five, namely  $\\beta$, $A_{s}$, $\\Delta N_{\\nu}$, the curvature density $\\Omega_{K0}$, and the parameter $\\alpha$, which is the ratio of the dark-radiation energy density to the sum of dark-radiation and kination-like energy densities at the time of BBN. The best fit values of the parameters together with the corresponding $1\\sigma$ confidence intervals, arisen from the likelihood analysis, were shown in Table \\ref{ndb_bestfits}. Additionally, in Figures \\ref{cont_varying_alpha}-\\ref{cont_varying_beta} we presented the likelihood contours of the model parameters. From this analysis we deduced that the entire (consistently with BBN) range of $0\\leq\\Delta N_{\\nu}\\leq2.0$ is allowed, for suitable choices of the other parameters. We also found tight constraints on the curvature and on $A_s$.   After the observational elaboration, we investigated the cosmological implications of the examined scenarios. In particular, we focused on the evolution of the total equation-of-state parameter of the universe, and of the expansion rate. In the detailed-balance case, we saw that in general at very early-times is radiation that dominates, at intermediate redshifts is the Chaplygin gas that dominates, behaving like matter for a long time, and finally at late times the dominant Chaplygin gas behaves like dark energy, triggering the accelerating expansion (see Fig \\ref{db_w_bf}). This evolution is consistent with the thermal history of the universe, and this is an advantage of the present scenario. Note moreover that the present accelerated era is described in a unified way, without the need of any additional mechanism. Concerning the qualitative dependence on the various parameters, we found that the expansion rate has a strong dependence on both $A_{s}$ and $\\Delta N_{\\nu}$ at low redshifts, while at high redshifts, where the Chaplygin gas is not dominant, this dependence weakens and disappears (see Fig \\ref{db_variations}). Lastly, the dependence on $\\beta$ is very weak at all redshifts.  In the beyond-detailed-balance scenario we found a similar behavior for the total equation-of-state parameter of the universe, that is we obtained successively a radiation, a matter, and a dark energy era (see Fig \\ref{w_bf}). Qualitatively, $A_{s}$ has a weak impact on the expansion rate at lower redshifts, $\\Delta N_{\\nu}$ has a significant effect at large redshifts, while $\\Omega_{K0}$ and especially $\\alpha$ and $\\beta$ have a negligible impact on the expansion rate at all redshifts (see Figures \\ref{H_varying_As}-\\ref{H_varying_beta}).  In summary, the generalized Chaplygin gas in Ho\\v{r}ava-Lifshitz gravitational background is compatible with observations, and can successfully reproduce the expansion history of the universe. However, we should mention that the present analysis does not enlighten the discussion about the possible conceptual problems and instabilities of Ho\\v{r}ava-Lifshitz gravity,  nor it can address the questions concerning the validity of its theoretical background, which is the subject of interest of other studies. It just analyzes the phenomenological consequences and the cosmological implications of the generalized Chaplygin Gas in such a gravitational background, and thus its results can be taken into account only if Ho\\v{r}ava-Lifshitz gravity passes successfully the necessary theoretical tests.  In the same lines, the present work does not address the problem of undesirable instabilities and oscillations in the matter power spectra, that GCG scenario faces in standard Einstein gravity. Although, as discussed in the Introduction, there are several approaches which try to solve this problem in conventional gravity, it may be quite interesting to see whether it can be addressed in the context of Ho\\v{r}ava-Lifshitz gravity. Let us make some comments in this perspective. As discussed in \\cite{Bento:2002ps, Bertolami:2004ic}, GCG can be modelled in terms of scalar fields having both canonical as well as non-canonical kinetic energy terms, and thus perturbing the GCG one has to essentially perturb these scalar fields. In the case of Ho\\v{r}ava-Lifshitz gravity one expects in general the linear perturbations to be quite different due to higher-order curvature-terms present in the action. These higher-order gradient terms can in principle generate non-adiabatic pressure perturbation, which can help to cure the instabilities or oscillations in the power spectra even in the large scales. The effect can be prominent with scalar fields having non-canonical kinetic energy terms, for instance of Dirac-Born-Infeld form, which is typical for GCG-like equation of state.  In conclusion, the generalized Chaplygin gas scenario in Ho\\v{r}ava-Lifshitz cosmology, can have rich cosmological consequences. The present study is a first step in this direction, where we have considered the background evolution and have confronted the model with a variety of currently available observational data. The next step is to consider the inhomogeneous GCG model in Ho\\v{r}ava-Lifshitz gravity, which will be our goal in the near future."
    },
    "1004/1004.2368_arXiv.txt": {
        "abstract": "We have done a statistical analysis of Very Long Baseline Array (VLBA) data of water masers in the star-forming regions (SFRs) Cepheus A and W75 N, using correlation functions to study the spatial clustering and Doppler-velocity distribution of these masers. Two-point spatial correlation functions show a characteristic scale size for clusters of water maser spots $\\la$1~AU, similar to the values found in other SFRs. This suggests that the scale for water maser excitation tends to be $\\la$1 AU. Velocity correlation functions show power-law dependences with indices that can be explained by regular velocity fields, such as expansion and/or rotation. These velocity fields are similar to those indicated by the water maser proper-motion measurements; therefore, the velocity correlation functions appear to reveal the organized motion of water maser spots on scales larger than 1 AU. ",
        "introduction": "Water maser emission is commonly found in star-forming regions (SFRs) within clumps of molecular gas with high densities ($\\sim10^8-10^9$~cm$^{-3}$) and warm temperatures ($\\sim$400~K) \\citep{Eli89}. These physical conditions can be reached in both the inner parts of accretion disks around young stellar objects (YSOs) and the gas compressed by shock waves generated by outflows associated with these objects.  Water masers are very compact and intense ($\\sim10^{13}$~cm, $T_b\\sim10^{12}$~K), allowing studies through Very Long Baseline Interferometry (VLBI) observations to analyze in detail the spatio-kinematical distribution of the gas around YSOs with very high angular and spectral resolutions. In particular, sources with a large number of masers allow us to obtain through a statistical analysis the clustering and velocity distribution of masers from scales of a few hundred of AUs down to less than 1~AU. Statistical studies give important information about the typical size for clusters of masers that may be related to a fundamental scale of its excitation, as well as the statistical properties of both the clustering and the velocity field.  In this context, \\citet{Gwi92} carried out VLBI water maser observations of the SFR W49. They identified 271 maser features with 625 maser spots, which were used for a statistical study. A maser spot is an individual component of the maser emission occurring at a given velocity channel and position. Maser spots tend to cluster in position and Doppler velocity, typically 1~AU and 0.5~km~s$^{-1}$ \\citep{Gwi94a}. Such clusters or groups of maser spots are called maser features. Physically, maser features should be small clouds supporting population inversion by a pumping mechanism \\citep{Gwi94b}. The statistical analysis of spatial and velocity distributions for both spots and features reveals that turbulent motions dominate on a spatial scale of $\\sim$1--300~AU \\citep{Gwi94a}.  \\citet{Ima02} carried out a similar study using Very Long Baseline Array (VLBA) maser data of the SFR W3 IRS 5. For the statistics, \\citet{Ima02} used 905 maser spots grouped in 152 maser features. They found that spots form features with a typical size of $\\sim$0.5~AU. The statistical analysis of the Doppler velocities, specifically, the velocity correlation functions for maser spots follow a power-law dependence in the range of 0.04$-$300~AU with an index of $\\sim$0.29, that is consistent with the Kolmogorov value of 1/3, the value expected for incompressible fluids with a turbulent velocity field \\citep{Kol41,Str07}. Similar results have been also obtained by \\citet{Str02} in other five SFRs.  VLBA water maser observations toward the SFRs Cepheus A and W75 N revealed remarkable microstructures \\citep{Tor01b,Tor03}. These microstructures exhibit a coherent and well-ordered spatio-kinematical behavior at AU scales \\citep{Usc05}. Proper-motion measurements of water masers suggest the presence of  organized motions of structures with sizes from tens to a few hundreds of AUs \\citep{Tor01b,Tor03}. Here we study the spatial and velocity distribution of the water masers in these two SFRs using a statistical analysis to investigate whether organized or turbulent motions dominate over spatial scales from a few hundred of AUs down to less than 1~AU.  This paper is organized as follows. In Section 2, we define the correlation functions used in this statistical study. In Sections 3 and 4, we give a brief introduction to the SFRs, Cepheus A and W75~N, and present the results of the correlation functions, respectively. In Section 5, we discuss the clustering of water maser spots and give a possible interpretation for the velocity field traced by the water masers, based on the statistical analysis. Finally, we summarize our conclusions in Section 6. ",
        "conclusions": "In this paper, we present a statistical analysis of the spatial and velocity distribution of water masers in the SFRs Cepheus A and W75~N, using the data of the VLBA maser observations carried out by \\citet{Tor01b,Tor03}. Our conclusions are as follows:  \\begin{enumerate} \\item We have found a characteristic scale size for clusters of water maser spots $\\la$1~AU, indicated by the break in the slope of the two-point spatial correlation function. Specifically, $\\sim$0.31$\\pm$0.07~AU in Cep A R5, $\\sim$0.85$\\pm$0.10 in W75~N VLA 1, and $\\sim$0.9$\\pm$0.3~AU in W75~N VLA2. These values are close to the typical sizes of the water maser features found in other SFRs, e.g., $\\sim$0.5~AU in W3~IRS~5 \\citep{Ima02} and $\\sim$1~AU in W49~N \\citep{Gwi94a}. Probably, these results may indicate that the scale for water maser excitation tends to be $\\la$1 AU, as it was pointed out by models of water masers excited by shocks \\citep{Eli89,Kau96}.  \\item Two-point spatial correlation functions follow power-law dependences, indicating self-similar spatial distributions of water masers. The power-law indices found for separations smaller than the characteristic scale size of clusters ($\\alpha=-$4.3 to $-$2.4) are steeper in Cepheus A and W75~N than the value found in W3 IRS~5 \\citep[$\\alpha=-2.09$,][]{Ima02}. This could be due to a different spatial distribution of spots inside the clusters of each region. The power-law indices found for separations larger than the characteristic scale size ($\\alpha=-1.56$ to $-$1.27) are similar to the value found in W49~N \\citep[$\\alpha=-1.33$,][] {Gwi94a}. This suggests similar statistical properties for the spatial distribution of masers at scales larger than 1 AU in the SFRs Cepheus A, W75~N, and W49~N.  \\item Velocity correlation functions follow power-law dependences. In Cep~A R5, the power-law index is about zero, consistent with pure expansion of dots randomly distributed in a thin spherical shell. In W75~N VLA~2, the value for the power-law index can be explained by the presence of expansion and rotation within an arc of a circle observed with a small inclination angle. These results are in agreement with proper-motion observations of water masers \\citep{Tor01a,Tor01b,Tor03}. In W75~N VLA~1, the value for the power-law index seems to be consistent with organized motions but further studies are needed. In the SFRs Cepheus A and W75~N, the values estimated for the power-law indices suggest that water masers are tracing an organized component of the velocity field instead of a turbulent component, probably due to a prevalence of organized motions over turbulent motions at spatial scales from a few hundred of AUs down to less than 1~AU.  \\item Based on the value of the power-law index followed by the velocity differences of maser spot pairs as a function of their separation, it may be possible to find evidence for expansion and/or rotation motions by analyzing water maser data from a single observed epoch. \\end{enumerate}"
    },
    "1004/1004.0647_arXiv.txt": {
        "abstract": "Using simulations of geosynchrotron radiation from extensive air showers, we present a relation between the shape of the geosynchrotron radiation front and the distance of the observer to the maximum of the air shower. By analyzing the relative arrival times of radio pulses at several radio antennas in an air shower array, this relation may be employed to estimate the depth of maximum of an extensive air shower if its impact position is known, allowing an estimate for the primary particle's species. Vice versa, the relation provides an estimate for the impact position of the shower's core if an external estimate of the depth of maximum is available. In realistic circumstances, the method delivers reconstruction uncertainties down to $30$\\unit\\gcm\\ when the distance to the shower core does not exceed $7$\\unit{km}. The method requires that the arrival direction is known with high precision. ",
        "introduction": "One of the most important open questions in astroparticle physics is the nature of cosmic-ray particles at the highest energies. At energies exceeding $10^{15}$\\unit{eV}, at present, the only practical way to investigate cosmic-ray particles is to register extensive air showers induced by cosmic rays in the atmosphere. In such experiments it is only possible to make statements on the composition of primary cosmic rays based on statistical evaluations. Abundances of primary particle types of an ensemble of air showers are frequently derived by looking at the depth of the shower maximum, i.e.\\ the depth at which the number of particles in a shower reaches its maximum.  In recent years, there has been a surge of interest in the detection of extensive air showers by means of the radio emission produced by the shower particles~\\citep{2005:Falcke,2005:Ardouin}. This observational technique allows one to look all the way up to the shower maximum, and it has the advantage over detecting the particles themselves at ground level that there is no \\change{attenuation}{absorption} of the signal. Several theories explaining the emission mechanism have been proposed \\citep{2003:Falcke,2008:Scholten,2008:MeyerVernet}. The former of these explains the observed radio emission from the principle of geosynchrotron radiation, and using a sophisticated model of geosynchrotron emission it was shown that the position of the maximum of inclined showers can be derived from the lateral slope of the electric field strength at ground level~\\citep{2008:Huege}.  In this work, we use simulations of air showers and their geosynchrotron radiation to estimate the value of the depth of maximum and the impact position of the shower core. The method developed exploits delays in the arrival time of the signal at different positions on the ground. ",
        "conclusions": "Through detailed simulations of air showers and their geosynchrotron radio emission, we have derived an empirical relation between the relative delay of the radio pulse emitted by the air shower front and the atmospheric depth of the shower maximum. By analysis of the radio pulse arrival delays in radio antennas in an array of low-frequency radio antennas, this relation can be used to estimate the depth-of-maximum if the impact position is known or vice versa.  We have confirmed that both methods work in principle, with no information other than radio signal delays used in the reconstruction. When the algorithm is tested under realistic conditions, however, the accuracy of the method is reduced. In the case of determining the shower maximum, reconstruction down to a useful confidence level is possible only for shower maxima up to~$\\sim7$\\unit{km} away, and only if the shower core impact position is known down to a few meters. When the parameterization is used to derive this position, the critical quantity is the accuracy in the zenith angle of the shower, which needs to be significantly less than a degree to reconstruct the shower impact location to an accuracy of $10$\\unit{m} at high inclinations up to~$60$\u00ba.  \\begin{ack} This work is part of the research programme of the `Stich\\-ting voor Fun\\-da\\-men\\-teel Onder\\-zoek der Ma\\-terie (\\textsc{fom})', which is financially supported by the `Neder\\-landse Orga\\-ni\\-sa\\-tie voor Weten\\-schap\\-pe\\-lijk Onder\\-zoek (\\textsc{nwo})'. T.~Huege was supported by grant number VH-NG-413 of the Helmholtz Association. \\end{ack}"
    },
    "1004/1004.4252_arXiv.txt": {
        "abstract": "Recently \\citet{Knutson.10} have demonstrated a correlation between the presence of temperature inversions in the atmospheres of hot Jupiters, and the chromospheric activity levels of the host stars. Here we show that there is also a correlation, with greater than 99.5\\% confidence, between the surface gravity of hot Jupiters and the activity levels of the host stars, such that high surface gravity planets tend be found around high activity stars. We also find a less significant positive correlation between planet mass and chromospheric activity, but no significant correlation is seen between planet radius and chromospheric activity. We consider the possibility that this may be due to an observational bias against detecting lower mass planets around higher activity stars, but conclude that this bias is only likely to affect the detection of planets much smaller than those considered here. Finally, we speculate on physical origins for the correlation, including the possibility that the effect of stellar insolation on planetary radii has been significantly underestimated, that strong UV flux evaporates planetary atmospheres, or that high mass hot Jupiters induce activity in their host stars, but do not find any of these hypotheses to be particularly compelling. ",
        "introduction": "\\label{sec:intro}  With more than 70 transiting exoplanets (TEPs) now known\\footnote{e.g. http://exoplanets.org}, it has been become possible to detect statistically robust correlations between the parameters of TEPs and their host stars, which in turn yields insights into the processes that are important for determining the physical properties of exoplanet systems. Several correlations have already been noted, including correlations between the masses and orbital periods of TEPs \\citep{Gaudi.05, Mazeh.05,2008ApJ...677.1324T}, between their surface gravities and orbital periods \\citep{Southworth.07,2008ApJ...677.1324T}, between the inferred core mass of planets and the metallicity of their host stars \\citep{Guillot.06, Burrows.07}, between Safronov number and the host star metallicity \\citep{2008ApJ...677.1324T}, and between the radii of planets and their average equilibrium temperature and host metallicity \\citep{Enoch.10}.  Very recently \\citet[][hereafter KHI10]{Knutson.10} have demonstrated a correlation between the emission spectra of TEPs and the chromospheric activity levels of their host stars, as measured from the strength of the emission lines at the Ca~II H and K line cores. Planets with spectra consistent with noninverted temperature models appear to be found around high activity stars, while planets with spectra consistent with temperature inversions are found around low activity stars. In demonstrating this correlation KHI10 also published a catalogue of $\\log R^{\\prime}_{HK}$ values for 39 TEPs. This new, homogeneous sample enables statistical studies of the relationships between stellar activity and the physical properties of TEPs.  In this paper we use the sample of $\\log R^{\\prime}_{HK}$ values from KHI10 to investigate correlations between stellar activity and other TEP parameters. We find that there is a significant correlation between $\\log R^{\\prime}_{HK}$ and the planet surface gravity $\\log g_{\\rm P}$. A similar correlation between stellar activity (as traced by the temporal variation in an index related to $\\log R^{\\prime}_{HK}$) and the minimum planetary mass $M_{\\rm P}\\sin i$ was previously noted by \\citet{Shkolnik.05} for a sample of 10 RV planets, though the authors deemed the correlation to be only suggestive. Here we demonstrate that the $\\log R^{\\prime}_{HK}$-$\\log g_{\\rm P}$ correlation is robust with greater than $99.5\\%$ confidence.  The structure of this paper is as follows: in Section~\\ref{sec:dataandanal} we describe the data and conduct the statistical analysis to establish the correlation, in Section~\\ref{sec:selecteffect} we discuss a potential observational bias which might lead to this correlation, and in Section~\\ref{sec:discussion} we speculate on the physical origins of this correlation. ",
        "conclusions": "\\label{sec:discussion}  Assuming that the observed correlation between $\\log R^{\\prime}_{\\rm HK}$ and $\\log g_{\\rm P}$ is not due to an observational bias, it is not obvious what physical processes might give rise to it. It is well-known that $\\log R^{\\prime}_{\\rm HK}$ is a decreasing function of age for FGK stars \\citep[e.g.][]{Soderblom.91}, but the $\\log g_{\\rm P}$-age relation implied from the $\\log g_{\\rm P}$-$\\log R^{\\prime}_{\\rm HK}$ relation is opposite of what is expected--that planets should contract with age, and not expand with age. For example, by interpolating the \\citet{Fortney.07} models while accounting for the increase in stellar luminosity over time, we find that a coreless $1.0~M_{\\rm J}$ planet orbiting a $1.0~M_{\\odot}$ star on a $2.5~{\\rm day}$ period should decrease in radius from $1.22~R_{\\rm J}$ to $1.13~R_{\\rm J}$ between 300~Myr and 4.5~Gyr. Models such as these, however, are known to underpredict the radii of many hot Jupiters. If the effect of insolation on planetary radii is substantially larger than anticipated, so that the inflation due to the increase in stellar luminosity with time is greater than the gravitational contraction of the planet over time, the result would be a positive $\\log g_{\\rm P}$-$\\log R^{\\prime}_{\\rm HK}$ correlation.  Another possibility is that strong stellar UV flux increases the evaporation of hydrogen from the atmospheres of hot Jupiters, leading to higher metallicity, more compact planets \\citep[e.g.][and references therein]{LecavelierdesEtangs.10}. If this were the case, we might expect to see a correlation between planet radius and stellar activity, and no correlation between planet mass and activity. The fact that the opposite effect is observed casts doubt on this hypothesis. Moreover, since the stellar activity should decrease with age, this hypothesis does not explain how planets could re-inflate when the activity is lowered.  Alternatively, the presence of hot Jupiters may induce activity on the host star, either by tidally spinning-up the star's convection zone, or via a magnetic star-planet interaction \\citep[see the review by][]{Shkolnik.09}. In this case stars with high $\\log R^{\\prime}_{\\rm HK}$ may not necessarily be younger than stars with lower $\\log R^{\\prime}_{\\rm HK}$. Evidence that the presence of a hot Jupiter is correlated with increased stellar X-ray activity has been presented by \\citet{Kashyap.08}, while \\citet{Pont.09} found that hot Jupiter host stars may exhibit excess rotation. Other investigations have found evidence of magnetic activity variations correlated with planet properties \\citep[e.g.][]{Shkolnik.05,Lanza.09}. For both tidal and magnetic star-planet interactions, the strength of the interaction increases with planet mass. If the $\\log R^{\\prime}_{\\rm HK}$-$\\log g_{\\rm P}$ correlation is a by-product of a more fundamental $\\log R^{\\prime}_{\\rm HK}$-$M_{\\rm P}$ correlation, one might wonder why the former is detected with higher significance than the latter. A possible explanation is that $\\log g_{\\rm P}$ is determined directly from measurable parameters while $M_{\\rm P}$ is directly proportional to the stellar mass $M_{\\rm S}$, which in turn is dependent on stellar models. As a result $\\log g_{\\rm P}$ is generally determined with better precision, and presumably with better accuracy, than $M_{\\rm P}$ for TEPs. However, by simulating data sets with the observed $M_{\\rm P}$-$\\log R^{\\prime}_{\\rm HK}$ correlation and $M_{\\rm P}-R_{\\rm P}$ correlations, assuming the scatter about these relations is intrinsic, and assuming the observational errors for $M_{\\rm P}$ and $\\log g_{\\rm P}$ are realistic, we find that there is only a $\\sim 1\\%$ probability of the FAP of $\\log g_{\\rm P}$-$\\log R^{\\prime}_{\\rm HK}$ being less than 0.1\\% while the FAP of $M_{\\rm P}$-$\\log R^{\\prime}_{\\rm HK}$ is greater than 1\\%. Even if we assume the true observational error on $M_{\\rm P}$ is $\\sim 0.5M_{\\rm P}$, the probability is only $\\sim 6\\%$. It is therefore unlikely that the $\\log R^{\\prime}_{\\rm HK}$-$M_{\\rm P}$ relation is driving the $\\log R^{\\prime}_{\\rm HK}$-$\\log g_{\\rm P}$ relation.  In summary, we have identified a significant positive correlation between stellar activity and planetary surface gravity. As far as we are aware this correlation is unanticipated, and its cause is unclear."
    },
    "1004/1004.1870.txt": {
        "abstract": "We develop a theory of nonlinear cosmological perturbations on superhorizon scales for a single scalar field with a general kinetic term and a general form of the potential. We employ the ADM formalism and the spatial gradient expansion approach, characterised by $O(\\epsilon^m)$, where $\\epsilon=1/(HL)$ is a small parameter representing the ratio of the Hubble radius to the characteristic length scale $L$ of perturbations. We obtain the general solution for a full nonlinear version of the curvature perturbation valid up through second-order in $\\epsilon$ ($m=2$). We find the solution satisfies a nonlinear second-order differential equation as an extension of the equation for the linear curvature perturbation on the comoving hypersurface. Then we formulate a general method to match a perturbative solution accurate to $n$-th-order in perturbation inside the horizon to our nonlinear solution accurate to second-order ($m=2$) in the gradient expansion on scales slightly greater than the Hubble radius. The formalism developed in this paper allows us to calculate the superhorizon evolution of a primordial non-Gaussianity beyond the so-called $\\delta N$ formalism or separate universe approach which is equivalent to leading order ($m=0$) in the gradient expansion. In particular, it can deal with the case when there is a temporary violation of slow-roll conditions. As an application of our formalism, we consider Starobinsky's model, which is a single field model having a temporary non-slow-roll stage due to a sharp change in the potential slope. We find that a large non-Gaussianity can be generated even on superhorizon scales due to this temporary suspension of slow-roll inflation. ",
        "introduction": "\\label{sec:intro}  Recent observations of the cosmic microwave background anisotropy show very good agreement of the observational data with the predictions of conventional, single-field slow-roll models of inflation, that is, adiabatic Gaussian random primordial fluctuations with an almost scale-invariant spectrum~\\cite{Spergel:2006hy,Komatsu:2010fb}. Nevertheless, as the observational accuracy improves, it has become observationally feasible to detect a small non-Gaussianity in the data~\\cite{Bartolo:2004if,Komatsu:2001rj,Komatsu:2010fb}. In particular, the PLANCK satellite \\cite{Planck:2006uk} launched last year is expected to bring us much finer data and it is hoped that non-Gaussianity may actually be detected. As a consequence, non-Gaussianity from inflation has been a focus of much attention in recent years~\\cite{Seery:2005wm,Seery:2005gb, Rigopoulos:2003ak,Lyth:2005fi}.  To study possible origins of non-Gaussianity, one must go beyond the linear perturbation theory~\\cite{Sasaki:1998ug,Lyth:2004gb,Langlois:2006vv}. The conventional models of inflation cannot explain an observationally detectable level of non-Gaussianity, since the magnitude of it is extremely small, suppressed by slow-roll parameters~\\cite{Maldacena:2002vr}. Then a variety of ways to generate a large non-Gaussianity have been proposed. They may be roughly classified into two; multi-field models that produce non-Gaussianity classically on superhorizon scales~\\cite{Suyama:2007bg,Sasaki:2008uc,Malik:2006pm,Sasaki:2006kq,Yokoyama:2007uu, Byrnes:2008wi}, and non-canonical kinetic term models that produce non-Gaussianity quantum mechanically on subhorizon scales~\\cite{Alishahiha:2004eh,Chen:2006nt,Chen:2006xjb}. In particular, in the former case, the $\\delta N$ formalism~\\cite{Starobinsky:1986fxa,Sasaki:1995aw,Wands:2000dp} turned out to be a powerful tool for the estimation of non-Gaussianity~\\cite{Sasaki:1998ug,Rigopoulos:2003ak,Lyth:2004gb}.  In order to parameterize the amount of non-Gaussianity of primordial perturbations, the nonlinear parameter $f_{NL}$ is commonly used \\cite{Komatsu:2001rj,Spergel:2006hy}. This is related to the bispectrum of the curvature perturbation on $\\zeta$~\\cite{Wands:2000dp}, and is generally defined as \\begin{eqnarray} f_{NL}={5\\over 6}{\\prod_{i=1}^3 k_i^3 \\over \\sum_{i=1}^3 k_i^3} {B_{\\zeta}({\\bm k}_1,{\\bm k}_2,{\\bm k}_3)\\over 4\\pi^4 {\\cal P}_{\\zeta}^2 }\\,. \\end{eqnarray} Here ${\\cal P}_\\zeta$ and $B_{\\zeta}$ are the power spectrum and bispectrum of $\\zeta$, respectively, and they are defined in Fourier space by \\begin{eqnarray} \\langle {\\zeta}_{{\\bm k}_1} {\\zeta}_{{\\bm k}_2} \\rangle&=& (2\\pi)^3 \\delta^3({\\bm k}_1+{\\bm k}_2){2\\pi^2 \\over k_1^3} {\\cal P}_{\\zeta}({\\bm k}_1)\\,, \\nonumber\\\\ \\langle {\\zeta}_{{\\bm k}_1} {\\zeta}_{{\\bm k}_2}{\\zeta}_{{\\bm k}_3} \\rangle&=&(2\\pi)^3 \\delta^3({\\bm k}_1+{\\bm k}_2+{\\bm k}_3) B_{\\zeta}({\\bm k}_1,{\\bm k}_2,{\\bm k}_3)\\,, \\end{eqnarray} Corresponding to the two different origins of non-Gaussianity mentioned above, the nonlinear parameter $f_{NL}$ can be mainly classified into two types; the {\\it local\\/} type, $f_{NL}^{\\rm local}$, which may arise from multi-scalar models on superhorizon scales, and {\\it equilateral\\/} type, $f_{NL}^{\\rm equil}$, which arises from non-canonical kinetic term models on subhorizon scales.\\footnote{A new type of $f_{NL}$ has been studied recently~\\cite{Senatore:2009gt}, called the {\\it orthogonal\\/} type. This may be generated from higher derivative terms in the action.} %  %%%%%%%%%%%%%%%%%%% The local type is called so because it represents a local, point-wise non-Gaussianity given by \\begin{eqnarray} {\\zeta}({\\bm x})={\\zeta}_G({\\bm x})+{3\\over 5}f_{NL}^{\\rm local} {\\zeta}^2_G({\\bm x})\\,, \\end{eqnarray} where ${\\zeta}_G$ is the Gaussian random field. On the other hand, the equilateral form of the bispectrum is given by \\begin{eqnarray} f_{NL}^{\\rm equil}({\\bm k}_1, {\\bm k}_2, {\\bm k}_3) ={10\\over 3} {{\\cal A}_{NL}\\over \\sum_i k_i^3}, \\end{eqnarray} with the shape function ${\\cal A}_{NL}$ typically in the form~\\cite{Chen:2006nt}, \\begin{eqnarray} {\\cal A}_{NL}\\propto {1\\over 8}\\sum_i k_i^3 -{1\\over K}\\sum_{i<j} k_i^2 k_j^2+{1\\over 2K^2} \\sum_{i\\neq j} k_i^2 k_j^3,~~~~{\\rm with\\ \\ } K\\equiv k_1+k_2+k_3. \\end{eqnarray} Note that the sign convention of $f_{NL}$ here follows WMAP's sign convention and it is opposite to Maldacena~\\cite{Maldacena:2002vr}. On the observational side, the current bounds on the parameter $f_{NL}$ by WMAP seven years~\\cite{Komatsu:2010fb} are $-10<f_{NL}^{\\rm local}<74$ (95\\% C.L.) for the local form of the bispectrum and $ -214<f_{NL}^{\\rm equil}<266$ (95\\% C.L.) for the equilateral form. By PLANCK~\\cite{Planck:2006uk}, it is expected that non-Gaussianity of the level $|f_{NL}|\\gsim 5$ can be detected~\\cite{Komatsu:2001rj}.  In this paper, we investigate another possible origin of non-Gaussianity, namely, non-Gaussianity due to a temporary non-slow roll stage on superhorizon scales. In order to investigate such a case, however, the $\\delta N$ formalism is not sufficient. The reason is as follows. On superhorizon scales, one can employ the spatial gradient expansion~\\cite{Starobinsky:1986fxa,Nambu:1994hu,Salopek:1990jq, Kodama:1997qw}, which is characterised by the expansion parameter $\\epsilon=1/(HL)$ representing the ratio of the Hubble horizon radius to the characteristic length scale $L$ of the perturbation. The $\\delta N$ formalism or the separate universe approach is equivalent to the leading order approximation, i.e., $O(\\epsilon^0)$ in the gradient expansion. It is valid at a slow-roll stage when local values of the inflaton field at each local point (averaged over each horizon-size region) determine the evolution of the universe at each point. In the context of perturbation theory, this implies one can ignore the decaying mode of the curvature perturbation. However, when the slow-roll condition is violated, the decaying mode cannot be neglected any more and the gradient expansion to $O(\\epsilon^2)$ is known to play a crucial role already at the level of linear perturbation theory~\\cite{Seto:1999jc,Leach:2001zf,Jain:2007au}. Thus, to evaluate non-Gaussianity from a non-slow-roll stage of inflation, it is necessary to develop a nonlinear theory of cosmological perturbations valid up through $O(\\epsilon^2)$ in the gradient expansion~\\cite{Tanaka:2006zp,Tanaka:2007gh,Takamizu:2008ra}.  This paper is organized as follows. In Sec.~\\ref{sec:gradexp}, we briefly review the spatial gradient expansion to $O(\\epsilon^2)$ in the uniform Hubble slicing~\\cite{Tanaka:2006zp,Tanaka:2007gh,Takamizu:2008ra}. In Sec.~\\ref{sec:nlcurvpert}, we develop a theory of full nonlinear curvature perturbations on superhorizon scales valid up through $O(\\epsilon^2)$ in the gradient expansion. In doing so, we introduce a variable that represents the nonlinear curvature perturbation as an extension of the linear comoving curvature perturbation. Then we derive an explicit expression for the nonlinear solution. We find a nonlinear second-order differential equation which the general solution satisfies. In Sec.~\\ref{sec:matching}, we develop a general formalism for matching a $n$-th-order perturbative solution to the general nonlinear solution we found on superhorizon scales. Then in Sec.~\\ref{sec:linearmatch} we consider a special case when a linear perturbative solution is matched to the nonlinear solution. This applies to the case when the inflaton is still slow-rolling when the scale of interest crosses the Hubble horizon scale. As an application of our formalism, we study Starobinsky's model in Sec.~\\ref{sec:starobinsky}. Sec.~\\ref{sec:summary} is devoted to summary and discussion. ",
        "conclusions": "\\label{sec:summary}  We have developed a theory of nonlinear cosmological perturbations on superhorizon scales for a single scalar field with a general kinetic term and a general form of the potential to second-order in the spatial gradient expansion. The solution to this order is necessary to evaluate correctly the final amplitude of the curvature perturbation for models of inflation with a temporary violation of the slow-roll condition. We have employed the ADM formalism and obtained the general solution for full nonlinear curvature perturbations valid up through second-order in the gradient expansion. We have introduced a reasonable variable that represents the nonlinear curvature perturbation on comoving slices ${\\cal R}_ c^{\\rm NL}$, which reduces to the comoving curvature perturbation ${\\cal R}_c^{\\rm Lin}$ in the linear limit. Then we have found that ${\\cal R}_c^{\\rm NL}$ satisfies a nonlinear second-order differential equation, (\\ref{eq: basic eq for NL}), as a natural extension of the linear second-order differential equation.  Then we have formulated the matching of the nonlinear solution to a perturbative solution at $n$-th-order, ${\\cal R}_c^{\\rm pert}$, on superhorizon scales, and obtained a formula for the final value of the nonlinear curvature perturbation expressed in terms of ${\\cal R}_c^{\\rm pert}$ and its time derivative at the time of matching. Since the evolution of superhorizon curvature perturbations is genuinely due to the $O(\\epsilon^2)$ effect, our formulation can be used to calculate the primordial non-Gaussianity beyond the $\\delta N$ formalism which is equivalent to leading order in the gradient expansion.  Then we have considered the case when the linear approximation is valid up to the time of horizon crossing for wavenumbers of physical interest. In this case, we have developed a method to determine quantities corresponding to ${\\cal R}_c^{\\rm pert}$ and its time derivative at the matching time in terms of the linear solution.  As an example of such a case, we have investigated Starobinsky's model~\\cite{Starobinsky:1992ts} in which there is a temporary non-slow-roll stage during inflation due to a sudden change of the potential slope. We have found that non-Gaussianity can become large if the parameter $T\\approx {A_+/A_-}$, which characterises the ratio of the slope before and after the transition, is large. For $T\\gg1$, we have found that the non-Gaussianity parameter for the bispectrum $f_{NL}({ k}_1,{k}_2,{k}_3)$ is peaked at the wavenumbers forming an equilateral triangle, $k=k_1=k_2=k_3$, denoted by $f_{NL}^{eql}(k)$. It is found to be positive and takes the maximum value $f_{NL}^{eql}(k)\\simeq 2 T$ at $y=\\sqrt{T}k/k_0\\simeq 1.5$ where $k_0$ is the comoving wavenumber that crosses the horizon at the time when the potential slope changes. This implies that, even for a relatively small $T$, say for $T=10$, it is possible to generate a fairly large non-Gaussianity $f_{NL}\\sim 20$ at wavenumber $k\\simeq 0.5 k_0$.   Our formalism can be applied to many other interesting circumstances in which the slow-roll condition is temporarily violated. To mention a couple of examples, a non-slow-roll stage appears in a double inflation model~\\cite{Saito:2008em} or in a specific case of DBI inflation~\\cite{Jain:2007au}. It is of interest to investigate the non-Gaussianity in these models by applying our formalism.  It is also of interest to investigate the case when there is a step in the inflaton potential instead of a change in the slope, which was proposed to explain the `features' in the cosmic microwave background anisotropy~\\cite{Chen:2006xjb,Joy:2008qd}. The case of time-varying sound speed for models with non-canonical kinetic terms~\\cite{Khoury:2008wj} may also deserve future study, since a rapid temporal variation of the sound velocity violates a certain type of the slow-roll condition.  Finally, here we have focused on the case of a single scalar field. An immediate issue is to extend the present formalism to the case of a multi-component scalar field. We plan to work on this and hope to report the result in the near future.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1004/1004.4122_arXiv.txt": {
        "abstract": "{The mechanisms giving rise to diffuse radio emission in galaxy clusters, and in particular their connection with cluster mergers, are still debated.}{We aim to obtain new insights into the internal dynamics of the cluster Abell 2294, recently shown to host a radio halo.}  {Our analysis is mainly based on redshift data for 88 galaxies acquired at the Telescopio Nazionale Galileo. We combine galaxy velocities and positions to select 78 cluster galaxies and analyze its internal dynamics. We also use new photometric data acquired at the Isaac Newton Telescope and X--ray data from the Chandra archive.}{We re--estimate the redshift of the large, brightest cluster galaxy (BCG) obtaining $\\left<z\\right>=0.1690$, well at rest within the cluster. We estimate a quite large line--of--sight (LOS) velocity dispersion $\\sigma_{\\rm V}\\sim 1400$ \\ks and X--ray temperature $T_{\\rm X}\\sim \\,$10 keV. Our optical and X--ray analysis detects evidence for substructure. Our results are consistent with the presence of two massive subclusters separated by a LOS rest frame velocity difference $V_{\\rm rf}\\sim 2000$ \\kss, very closely projected in the plane of sky along the SE--NW direction.  The observational picture, interpreted through the analytical two--body model, suggests that A2294 is a cluster merger elongated mainly in the LOS direction and catched during the bound outgoing phase, a few fractions of Gyr after the core crossing.  We find Abell 2294 is a very massive cluster with a range of $M=2-4$ \\mquii, depending on the adopted model.  Moreover, contradicting previous findings, our new data do exclude the presence of the H$\\alpha$ emission in the spectrum of the BCG galaxy.}{The outcoming picture of Abell 2294 is that of a massive, quite ``normal'' merging cluster, as found for many clusters showing diffuse radio sources. However, maybe due to the particular geometry, more data are needed for a definitive, more quantitative conclusion.} ",
        "introduction": "\\label{intr}  Merging processes constitute an essential ingredient of the evolution of galaxy clusters (see Feretti et al. \\cite{fer02b} for a review). An interesting aspect of these phenomena is the possible connection of cluster mergers with the presence of extended, diffuse radio sources: halos and relics. The synchrotron radio emission of these sources demonstrates the existence of large--scale cluster magnetic fields and of widespread relativistic particles. Cluster mergers have been suggested to provide the large amount of energy necessary for electron reacceleration up to relativistic energies and for magnetic field amplification (Tribble \\cite{tri93}; Feretti \\cite{fer99}; Feretti \\cite{fer02a}; Sarazin \\cite{sar02}). Radio relics (``radio gischts'' as referred by Kempner et al. \\cite{kem04}), which are polarized and elongated radio sources located in the cluster peripheral regions, seem to be directly associated with merger shocks (e.g., Ensslin et al. \\cite{ens98}; Roettiger et al. \\cite{roe99}; Ensslin \\& Gopal--Krishna \\cite{ens01}; Hoeft et al. \\cite{hoe04}).  Radio halos, unpolarized sources which permeate the cluster volume similarly to the X--ray emitting gas (intracluster medium, hereafter ICM), are more likely to be associated with the turbulence following a cluster merger (Cassano \\& Brunetti \\cite{cas05}; Brunetti et al. \\cite{bru09}). However, the precise radio halos/relics formation scenario is still debated since the diffuse radio sources are quite uncommon and only recently one can study these phenomena on the basis of a sufficient statistics (few dozen clusters up to $z\\sim 0.3$, e.g., Giovannini et al. \\cite{gio99}; see also Giovannini \\& Feretti \\cite{gio02}; Feretti \\cite{fer05}; Giovannini et al. \\cite{gio09}) and attempt a classification (e.g., Kempner et al. \\cite{kem04}; Ferrari et al. \\cite{ferr08}). It is expected that new telescopes will largely increase the statistics of diffuse sources (e.g. LOFAR, Cassano et al. \\cite{cas09}).  From the observational point of view, there is growing evidence of the connection between diffuse radio emission and cluster merging, since up to now diffuse radio sources have been detected only in merging systems. In several cases the cluster dynamical state has been derived from X--ray observations (see Buote \\cite{buo02}; Feretti \\cite{fer06} and \\cite{fer08} and refs. therein).  Optical data are a powerful way to investigate the presence and the dynamics of cluster mergers (e.g., Girardi \\& Biviano \\cite{gir02}), too. The spatial and kinematical analysis of member galaxies allow us to detect and measure the amount of substructure, to identify and analyze possible pre--merging clumps or merger remnants.  This optical information is really complementary to X--ray information since galaxies and intra--cluster medium react on different time scales during a merger (see, e.g., numerical simulations by Roettiger et al. \\cite{roe97}).  In this context we are conducting an intensive observational and data analysis program to study the internal dynamics of clusters with diffuse radio emission by using member galaxies (Girardi et al. \\cite{gir07}\\footnote{see also the web site of the DARC (Dynamical Analysis of Radio Clusters) project: http://adlibitum.oat.ts.astro.it/girardi/darc.}). Most clusters showing diffuse radio emission have a large gravitational mass (larger than $0.7\\times 10^{15}$ within $2$ \\hh; see Giovannini \\& Feretti \\cite{gio02}) and, indeed, most clusters we analyzed are very massive clusters with few exceptions (Boschin et al. \\cite{bos08}).  During our observational program we have conducted an intensive study of the cluster \\object{Abell 2294} (hereafter A2294).  A2294 is a very rich, X--ray luminous, and hot Abell cluster: Abell richness class $=2$ (Abell et al. \\cite{abe89}); $L_\\mathrm{X}$(0.1--2.4 keV)=6.6$\\times 10^{44} \\ h_{50}^{-2}$ erg\\ s$^{-1}$; $T_{\\rm X}=\\,$8-9 keV recovered from ROSAT and Chandra data (Ebeling et al. \\cite{ebe98}; Rizza et al. \\cite{riz98}; Maughan et al. \\cite{mau08}). Optically, the cluster is classified as Bautz--Morgan class II (Abell et al. \\cite{abe89}) and is dominated by a central, large brightest cluster galaxy (BCG, see Fig.~\\ref{figimage1}).  From both ROSAT and Chandra data A2294 is known for having no cool core (Rizza et al. \\cite{riz98}; Bauer et al. \\cite{bau05}). As for the presence of possible substructure, using ROSAT data, Rizza et al. (\\cite{riz98}) found evidence of a centroid shift and detect a Southern excess in the X--ray emission. Moreover, using Chandra data, Hashimoto et al. (\\cite{has07}) classified A2294 as a ``distorted'' cluster due to its large value of the asymmetry parameter. Indeed, A2294 is a very peculiar cluster since, in contrast with the absence of a cooling core, it is very compact in its X-ray appearance (see Fig.3 of Bauer et al. \\cite{bau05}). Out of a sample of 115 clusters recently analyzed using Chandra data, A2294 is one with the smallest ellipticity, while shows a not large, but highly significant centroid shift (Maughan et al. \\cite{mau08}).  A2294 is also peculiar for another aspect.  Out of a sample of 13 clusters at $z\\sim 0.15-0.4$ showing evidence for H$\\alpha$ emission in the BCG spectrum, it is the only one not showing a cool core (see Fig.5 of Bauer et al. \\cite{bau05}). The correlation between BCG H$\\alpha$ emission and the presence of a cool core is also true for nearby clusters where `` H$\\alpha$ luminous galaxies lie at the center of large cool cores, although this special cluster environment does not guarantee the emission-line nebulosity in its BCG'' (Peres et al. \\cite{per98}).  More recent observations also agree that H$\\alpha$ emission is more typical of cool core clusters than of non--cool core clusters ($\\sim 70\\%$ against $\\sim 10\\% $, Edwards et al. \\cite{edw07}).  As for the diffuse radio emission, Owen et al. (\\cite{owe99}) first reported the existence of a detectable diffuse radio source in this cluster. Despite of the presence of some disturbing pointlike sources in the central region of the cluster, Giovannini et al. (\\cite{gio09}) could detect a radio--halo 3\\arcm in size. In particular, the position of A2294 in the $P_{1.4\\,\\rm{GHz}}$ (radio power at 1.4 GHz) - $L_{\\rm X}$ plane is consistent with that of all other radio--halo clusters (see Fig.~17 of Giovannini et al. \\cite{gio09}).  To date poor optical data are available. The cluster redshift reported in the literature ($z=0.178$) is only based on the BCG H$\\alpha$ emission line (Crawford et al. \\cite{cra95}). Instead, the real cluster redshift, as estimated in this paper, is rather $\\left<z\\right>=0.169$ fully consistent with that measured on the BCG on the base of our data which, indeed, do not show any evidence of H$\\alpha$ emission (see \\S~\\ref{data}).  Our new spectroscopic and photometric data come from the Telescopio Nazionale Galileo (TNG) and the Isaac Newton Telescope (INT), respectively. Our present analysis is based on these optical data and X--ray Chandra archival data.  This paper is organized as follows. We present our new optical data and the cluster catalog in Sect.~2. We present our results about the cluster structure based on optical and X--ray data in Sects.~3 and 4, respectively.  We briefly discuss our results and give our conclusions in Sect.~5.  \\begin{figure*} \\centering \\includegraphics[width=18cm]{figimage1.ps} \\caption{INT $R$--band image of the cluster A2294 (North at the top and East to the left) with, superimposed, the contour levels of the Chandra archival image ID~3246 (thick contours; photons in the energy range 0.5--2 keV) and the contour levels of a VLA radio image at 1.4 GHz (thin contours, see Giovannini et al. \\cite{gio09}). Labels and arrows highlight the positions of radio sources listed by Rizza et al. (\\cite{riz03}).} \\label{figimage1} \\end{figure*}  Unless otherwise stated, we give errors at the 68\\% confidence level (hereafter c.l.).  Throughout this paper, we use $H_0=70$ km s$^{-1}$ Mpc$^{-1}$ in a flat cosmology with $\\Omega_0=0.3$ and $\\Omega_{\\Lambda}=0.7$. In the adopted cosmology, 1\\arcm corresponds to $\\sim 173$ \\kpc at the cluster redshift.  \\begin{figure*} \\centering \\includegraphics[width=18cm]{figottico3.ps} \\caption{INT $R$--band image of the cluster A2294 (North at the top and East to the left). Circles and squares indicate cluster members and non--members, respectively (see Table~\\ref{catalogA2294}). Solid circle in the center highlights the position of the BCG galaxy. Annuli and box annuli show member and non--member emission line galaxies, respectively. Labels indicate the IDs of cluster galaxies cited in the text. A diamond at the right border of the image highlights a QSO at $z\\sim$2.1.} \\label{figottico2} \\end{figure*} ",
        "conclusions": ""
    },
    "1004/1004.2124_arXiv.txt": {
        "abstract": "{} {Recent theoretical predictions of the lowest very high energy (VHE) luminosity of SN~1006 are only a factor 5 below the previously published H.E.S.S. upper limit, thus motivating further in-depth observations of this source.} {Deep observations at VHE energies (above 100~GeV) were carried out with the High Energy Stereoscopic System (H.E.S.S.) of Cherenkov Telescopes from 2003 to 2008. More than 100 hours of data have been collected and subjected to an improved analysis procedure. } {Observations resulted in the detection of VHE $\\gamma$-rays from SN~1006. % The measured $\\gamma$-ray spectrum is compatible with a power-law, the flux is of the order of 1$\\%$ of that detected from the Crab Nebula, and is thus consistent with the previously established H.E.S.S.\\ upper limit. The source exhibits a bipolar morphology, which is strongly correlated with non-thermal X-rays. } {Because the thickness of the VHE-shell is compatible with emission from a thin rim, particle acceleration in shock waves is likely to be the origin of the $\\gamma$-ray signal.  The measured flux level can be accounted for by inverse Compton emission, but a mixed scenario that includes leptonic and hadronic components and takes into account the ambient matter density inferred from observations also leads to a satisfactory description of the multi-wavelength spectrum.} ",
        "introduction": "The source SN~1006 is the remnant of one of the few historical supernovae. It appeared in the southern sky on 1006 May 1 and was recorded by Chinese and Arab astronomers \\cite{Stephenson02}. The remnant of this explosion was first identified at radio wavelengths on the basis of historical evidence \\cite{Gardner}. The evolution of its luminosity indicates that it is the result of a Type Ia supernova \\cite{Schaefer}, probably the brightest supernova in recorded history. A distance of 2.2~kpc was derived by Winkler et~al.~(2003) based on comparing the optical proper motion with an estimate of the shock velocity derived from optical thermal line broadening assuming a high Mach number single-fluid shock.  Contemporary interest in the very high energy (VHE) emission from supernova remnants (SNRs) has arisen due to their association as prime candidates for Galactic cosmic-ray acceleration. Firstly, Galactic SNRs have sufficient kinetic energy to explain the estimated Galactic luminosity in cosmic rays of $10^{40}$~erg/s. Secondly, and more importantly, it has been shown that diffusive shock acceleration provides a viable mechanism which can efficiently accelerate charged particles in the blast waves of SNRs (e.g. Drury 1983; Blandford \\& Eichler 1987; Jones \\& Ellison 1991; Berezhko et al. 1996). Indeed, most shell-type SNRs are non-thermal radio emitters, which confirms that electrons are accelerated up to at least GeV energies. Moreover, the limb-brightened non-thermal radio emission traces the site of effective particle acceleration.  The source SN~1006 was also the first SNR in which a non-thermal component of hard X-rays was detected in the rims of the remnant by ASCA \\cite{Koyama} and ROSAT \\cite{Willingale}, whereas the interior of the remnant exhibits a thermal spectrum with line emission. The hard featureless power-law spectrum strongly implies a synchrotron origin of the radiation, which in turn suggests that electrons can be accelerated up to energies of $\\sim\\mathrm{100}$~TeV. % Subsequent arcsecond resolution images by Chandra revealed a small-scale structure in the nonthermal X-ray filaments of the NE rim of SN~1006 \\cite{Bamba,Long}, supporting the idea of high B-fields in the bright limbs of the remnant \\cite{Berezhko2002}. % An analysis of the X-ray observations from XMM-Newton by Rothenflug et.~al (2004) leads to the conclusion that the magnetic field in the remnant is oriented in the NE-SW direction. The synchrotron emission would then be concentrated in regions where the shock is quasi-parallel \\cite{Voelk2003}. %  Also, $\\gamma$-ray observations of SN~1006 were carried out by ground-based $\\gamma$-ray telescopes. A TeV $\\gamma$-ray signal at the level of the Crab flux was claimed by the CANGAROO-I \\cite{Cangaroo-I} and CANGAROO-II \\cite{Cangaroo-II} telescopes, but subsequent stereoscopic observations of the source with the H.E.S.S. telescopes in 2003 and 2004 found no evidence of VHE $\\gamma$-ray emission and derived an upper limit of $\\Phi(>0.26$~TeV$)<2.4\\times 10^{-12}$~ph~cm$^{-2}$~s$^{-1}$ at 99.9\\% confidence level \\cite{HESS_limit}. The CANGAROO-III telescope array found only an upper limit which is consistent with the H.E.S.S. result \\cite{Cangaroo-III}.  The initial non-detection of SN~1006 in VHE $\\gamma$-rays does not invalidate the hypothesis of nuclear particle acceleration in the shock. Indeed, the hadronic $\\gamma$-ray flux is very sensitive to the ambient gas density $\\mathrm{n_H}$ and hence the H.E.S.S. upper limit implies a constraint on $\\mathrm{n_H}<0.1$~cm$^{-3}$ \\cite{Ksenofontov}. Indeed, being 500~pc above the Galactic plane, the remnant is relatively isolated, and the gas density around SN~1006 was recently estimated to be around 0.085~cm$^{-3}$ \\cite{Katsuda}. Ksenofontov et al.~(2005) furthermore showed that the lower limit for the VHE $\\gamma$-ray flux, which is given by the inverse Compton (IC) component derived from the integrated synchrotron flux and field amplification alone, was only a factor 5 below the H.E.S.S. upper limit. These predictions promoted deep observations with the H.E.S.S. telescopes. ",
        "conclusions": "Very high energy $\\gamma$-rays from SN 1006 have been detected by H.E.S.S. The measured flux above 1~TeV is of the order of 1$\\%$ of that detected from the Crab Nebula and therefore compatible with the previously published upper limit \\cite{HESS_limit}. The bipolar morphology apparent in $\\gamma$-rays is consistent with the non-thermal emission regions also visible in X-rays. As the VHE-shell is compatible with a scenario of thin rim emission, particle acceleration in the very narrow X-ray filaments, which are signatures of shocks, is also likely to be at the origin of the $\\gamma$-ray signal.  The measured flux level can be accounted for by inverse Compton emission assuming a magnetic field of about 30~$\\mu$G.  A mixed scenario including leptonic and hadronic processes and taking into account the ambient matter density estimated from observation also leads to a satisfactory description of the multi-wavelength spectrum, assuming a high proton-acceleration efficiency.  None of the models can be excluded at the level of modelling presented here."
    },
    "1004/1004.2915_arXiv.txt": {
        "abstract": "In this paper we present trispectrum estimation methods which can be applied to general non-separable primordial and CMB trispectra.  We review the relationship between the reduced CMB trispectrum and  the reduced primordial trispectrum. We present a general optimal  estimator for the connected part of the trispectrum, for which we derive a quadratic term to incorporate the effects of inhomogeneous noise and  masking.  We describe a general algorithm for creating simulated maps with given arbitrary (and independent) power spectra, bispectra and trispectra. We propose a universal definition of the trispectrum parameter $T_{NL}$, so that the integrated trispectrum on the observational domain can be consistently compared between theoretical models.   We define a shape function for the primordial trispectrum, together with a shape correlator and a useful parametrisation for visualizing the trispectrum; these methods might also be applied to the late-time trispectrum for large scale structure. We derive separable analytic CMB solutions in the large-angle limit for constant and local models. We present separable mode decompositions which can be used to describe any primordial or CMB trispectra on their respective wavenumber or multipole domains.   By extracting coefficients of these separable basis functions from an observational map, we are able to present an efficient estimator for any given theoretical model with a nonseparable trispectrum.   The estimator  has two manifestations,  comparing the theoretical and observed coefficients at either primordial or late times, thus encompassing a wider range of models, such as secondary anisotropies, lensing and cosmic strings. We show that these mode decomposition methods are numerically tractable with order $l^5$ operations for the CMB estimator and approximately order $l^6$ for the general primordial estimator (reducing to order $l^3$ in both cases for a special class of models). We also demonstrate how the trispectrum can be reconstructed from observational maps using these methods. ",
        "introduction": "Single field slow-roll inflationary fluctuations in the standard picture of cosmology predict a nearly scale invariant spectrum of adiabatic perturbations with a nearly Gaussian distribution. Hence it can be described very accurately by its angular power spectrum. These predictions agree well with measurements of the cosmic microwave background (CMB) and large scale structure, such as those provided by WMAP and SDSS. However, it remains possible that there exists a mechanism for generating large non-Gaussianities in the early Universe. Measurements of such non-Gaussianities open up the opportunity of investigating the physics of the early universe including different inflationary models and competing alternative scenarios. In order to study such observations, higher order correlators, beyond the two-point function, offer possibly the best prospects. General methods for comparing the three point correlator, dubbed the bispectrum, were developed in \\cite{Ferg1,Ferg2,Ferg3}. In those papers an integrated measure of the bispectrum was defined, as well as a set of formalisms for comparing, evolving and constraining the bispectrum in the case of both the primordial and CMB three-point correlators. In this paper we will generalise many of these methods to the four-point correlator which is denoted the trispectrum. We will emphasise the application of these methods to the primordial and CMB trispectra. The primary motivation for this paper is to develop formalisms to bring observations to bear on this broader class of cosmological models. We will demonstrate that despite the complexity of trispectrum estimation, these methods are numerically tractable given present resources, even at Planck satellite resolution. \\par In order to get large non-Gaussianity we must move away from the standard single field slow-roll inflation~\\cite{Chen3}. Multifield inflation allows the possibility for superhorizon evolution. Non-Gaussianities are generated when this evolution is nonlinear. We can consider superhorizon behaviour as occurring in patches separated by horizons which evolve independently of each other. This locality in position space translates to non-locality in momentum space and indicates that for such models we expect the signal to peak for $k_4\\ll k_1,k_2,k_3$. This forms the so-called local model. Such models have been investigated in the context of the trispectrum in \\cite{Seery1,Seery2,Seery3,Adshead,Bartolo,Valen,Rodriguez,Huang1,Lehners}. Since subhorizon modes oscillate and so average out, the only chance to have large non-Gaussianity in single field inflationary models is when all modes have similar wavelengths and exit at the same time. A non-standard kinetic term allows for such a possibility. Since the signal peaks when the modes have similar wavelengths this class of forms are known as equilateral models and have been investigated using the trispectrum in \\cite{aChen,bArroja,cChen4,dArroja2,eArroja3,fSenatZal,gGao}. It should be noted that this amplification of nonlinear effects around the time the modes exit the horizon is not possible for slow-roll single field inflation. It has also been shown in \\cite{Huang2,Izumi} that a large trispectrum may be generated in the ghost inflation model. These models are so-called as they are based on the idea of a ghost condensate, i.e. a kind of fluid with equation of state $p=-\\rho$, that can fill the universe, and which provides an alternative method of realising de Sitter phases in the early universe. Of course there are other methods to generate non-Gaussianity such as having sharp features in the potential or a non-Bunch-Davies vacuum. Also there are models which have features that resemble the aforementioned forms in different regimes, e.g. quasi-single field inflation~\\cite{Chen2}, or have mixed contributions, e.g. in multifield DBI inflation~\\cite{RenauxDBI}. \\par One of the motivations for studying the four-point correlator is that it may be possible that the bispectrum is suppressed but still have a large trispectrum. In particular, this behaviour may be realised in quasi-single field inflation~\\cite{Chen2} or in the curvaton model~\\cite{Sasaki}. It also occurs in the case of cosmic strings where the bispectrum is suppressed by symmetry considerations \\cite{Hindmarsh2009,Regan}. The effects of non-Gaussianity  could also be detectable in a wide range of astrophysical measurements, such as cluster abundances and the large scale clustering of highly biased tracers. In \\cite{Sefusatti} the possibility of using the galaxy bispectrum to constrain the local form of the trispectrum has been reviewed. \\par The trispectrum, $T(k_1,k_2,k_3,k_4)$, is generally parametrised using the variable $\\tau_{NL}$ which schematically is given by the ratio $\\tau_{NL}\\approx T(k,k,k,k)/P(k)^3$. Standard slow-roll inflation predicts $\\tau_{NL}\\lesssim r/50$ where $r<1$ is the tensor to scalar ratio~\\cite{Seery1}. Such a low signal would be undetectable since it is below the level of non-Gaussian contamination that would be expected from secondary anisotropies $\\tau_{NL}\\approx \\mathcal{O}(1)$. Using the analysis of N-point probability distribution of the CMB anisotropies~\\cite{vielva}, where a local non-linear perturbative model $\\Phi=\\Phi_L+f_{NL}(\\Phi^2_L-\\langle \\Phi_L^2\\rangle)+g_{NL}\\Phi_L^3+\\mathcal{O}(\\Phi_L^4)$ is used to characterise the large scale anistropies, the constraint $-5.6\\times 10^5<g_{NL}<6.4\\times 10^5$ was obtained\\footnote{It should be noted that for single field local inflation $\\tau_{NL}^{\\rm{loc}}=\\left(\\frac{5}{6}f_{NL}\\right)^2$. Since $f_{NL}$ is constrained by the bispectrum, $g_{NL}$ is the quantity that is constrained by the trispectrum directly in this case.}. For the more general case, there is only a weak experimental bound imposed on non-Gaussianity by the trispectrum, which is roughly $|\\tau_{NL}|\\lesssim 10^8$~\\cite{Lyth}. In~\\cite{Cooray,Cooray2} an improved constraint on $\\tau_{NL}$ was presented using estimators to allow a joint fit of $f_{NL}$ and $g_{NL}$ using the trispectrum of WMAP5 data. However, the analysis therein included an incomplete formula for the CMB trispectrum due to local non-Gaussianity\\footnote{The formula for the reduced local CMB trispectrum has been used in place of the full local CMB trispectrum, which appears to simplify the analysis.}. Nonetheless, the approach indicates that vast improvements to trispectrum constraints should be achievable in the near future. In fact, it is expected that the Planck satellite will be sensitive to a value of $|\\tau_{NL}|\\sim 560$~\\cite{Kogo}. \\par The analysis of the trispectrum is a computationally intensive operation. In fact only the trispectrum induced by the local shape has been constrained so far by CMB data. The local form is an example of a separable shape - a notion which we will define more concretely in this paper. Essentially, since the primordial trispectrum is a six dimensional quantity, separability means the trispectrum is the product of one dimensional functions of each of these variables. Exploiting this separability reduces the problem from one of $\\mathcal{O}(l_{\\rm{max}}^7)$ operations to a more manageable $\\mathcal{O}(l_{\\rm{max}}^5)$. In special cases we get a further reduction to $\\mathcal{O}(l_{\\rm{max}}^3)$. \\par In the next section we shall describe the CMB trispectrum and its relation to the primordial equivalent. We will make use of a particular parametrisation of the reduced primordial trispectrum and exploit a Legendre series expansion in terms of one of these parameters to write an expression for the reduced CMB trispectrum which is valid in general. We will also outline a general correlation method for comparing different trispectra. In this section we will also give a formula for the kurtosis in terms of the multipoles. In section III we define a shape function which is a scale invariant form of the trispectrum. Using this function we define a shape correlator that is expected to predict closely the correlation between the respective trispectra. Next, we show how to decompose this shape in order to provide a method for visualising trispectra. We apply this visualisation to the case of the local and equilateral models which we describe in section IV. We also present the Sachs Wolfe limit ($l<100$) for the local and constant models. In section V we describe how to form a mode expansion for general non-separable shapes. This provides a rigorous method to find a separable approximation to any shape and therefore makes analysis of the trispectrum far more tractable. This expansion can be performed for both the primordial and CMB trispectra. Of immediate relevance in terms of Planck is to find a general measure for the size of the trisectrum. This is addressed in section VI in both the primordial and CMB cases. It is clearly desirable to be able to reconstruct the underlying trispectrum given the data. As we shall describe in section VII this is a computationally intensive task, but it is tractable. We will observe here that there is a degeneracy in reconstruction of the primordial trispectrum, implying that only the zeroth Legendre mode is recoverable. Finally in section VIII we outline a method for performing CMB map simulations for given general bispectra and trispectra. ",
        "conclusions": "We have described in this paper two comprehensive pipelines for the analysis of general primordial or CMB trispectra. The methods are based on mode expansions, exploiting a complete orthonormal eigenmode basis to efficiently decompose arbitrary trispectra into a separable polynomial expansion. These separable mode expansions allow for a reduction of the computational overhead to tractable levels, regardless of whether the reduced trispectrum is being computed at Planck resolution, or we are directly finding an estimator for the size of the trispectrum from a real data set. A shape decomposition has been described allowing for a visualisation of a scale invariant reduced trispectrum on particular slices. \\par We have presented a correlator for comparing trispectra. We have also defined a correlator for comparing the shape functions that is expected to closely approximate the former. However, the main purpose of this paper was to present a detailed theoretical framework for finding an estimator for the size of the trispectrum using separable eigenmode expansions. Using this efficient method for finding an estimator for the trispectrum we have defined a universal measure $T_{NL}$ which will allow for consistent comparison between theoretical models. This measure can be calculated for both primordial models and late-time models, e.g. due to active models such as cosmic strings. The completeness of the orthogonal eigenmodes should allow for a reconstruction of the full CMB trispectrum from the data, assuming the presence of a sufficiently significant non-Gaussian signal. We have also detailed an algorithm for producing non-Gaussian simulations with a given power spectrum, bispectrum and trispectrum. The implementation of these methods will be discussed in a future publication~\\cite{Regan2}. Clearly, the full implementation of the primordial and late-time pipelines represents a significant challenge. However the generality and robustness of the methodology described here indicates that there is an intriguing possibility of exploring and constraining a wide class of non-Gaussian models using the trispectrum."
    },
    "1004/1004.2853_arXiv.txt": {
        "abstract": "{% As a result of the variability survey in $\\chi$~Persei and NGC\\,6910, the number of $\\beta$~Cep stars that are members of these two open clusters is increased to twenty stars, nine in NGC\\,6910 and eleven in $\\chi$~Persei. We compare pulsational properties, in particular the frequency spectra, of $\\beta$~Cep stars in both clusters and explain the differences in terms of the global parameters of the clusters. We also indicate that the more complicated pattern of the variability among B type stars in $\\chi$~Persei is very likely caused by higher rotational velocities of stars in this cluster. We conclude that the sample of pulsating stars in the two open clusters constitutes a very good starting point for the ensemble asteroseismology of $\\beta$~Cep-type stars and maybe also for other B-type pulsators.} ",
        "introduction": "From the point of view of seismic modelling, the most important stellar parameters are masses, radii, luminosities, and effective temperatures. For single stars these parameters are usually poorly determined, especially for distant stars. On the other hand, if pulsators are members of an open cluster, the individual stellar parameters can be constrained much better. First, it can be safely assumed that the members have the same age, distance and chemical composition. In consequence, one can adopt that they are located on the same isochrone.  This helps to pinpoint stellar masses, especially if there are eclipsing binaries among cluster members that can be used to yield their masses and radii.  These considerations form a background for an {\\it ensemble} asteroseismology of stars in open clusters, where individual pulsators need no longer be modelled independently. Through the location on the isochrone, the masses of stars and consequently their frequency spectra are closely related and can be matched simultaneously. In general, we would like to point out that this kind of modelling has great potential in terms of (i) understanding the pulsations of different groups of stars observed in a cluster, (ii) estimating the cluster parameters, and (iii) testing the physics, e.g., the opacities.  The possibility of ensemble asteroseismology motivated us to carry out multisite campaigns on three open clusters that were known to contain $\\beta$~Cep-type pulsators, and for which we expected to discover several more. The first cluster, NGC\\,3293, will not be discussed here. Preliminary results of the campaign on this cluster were presented by Handler et al.~(2008). The other two clusters we selected are NGC\\,884 ($\\chi$~Persei) and NGC\\,6910. Before the campaign, two $\\beta$~Cep stars and some candidates were known in $\\chi$~Persei (Krzesi\\'nski \\& Pigulski 1997, 2000), four in NGC\\,6910 (Ko{\\l}aczkowski et al. 2004).  The two open clusters, $\\chi$~Persei and NGC\\,6910, are different in many ways. The mass of $\\chi$~Persei is estimated to be 3700~M$_\\odot$, the age is about 14~Myr (Currie et al.~2010). NGC\\,6910 is about an order of magnitude less massive and much youn\\-ger; the age of this cluster was estimated to be 6\\,$\\pm$\\,2~Myr by Ko\\-{\\l}acz\\-kow\\-ski et al.~(2004). NGC\\,6910 is also more reddened than $\\chi$~Persei with considerable differential reddening: the range of $E(B-V)$ colour excess amounts to 0.52--0.56~mag for $\\chi$~Persei and 1.0--1.4~mag for NGC\\,6910. The colour-mag\\-ni\\-tu\\-de diag\\-rams (CMDs) for $\\chi$~Persei and NGC\\,6910 are shown in Figs.~\\ref{chiper-cmd} and \\ref{6910-cmd}, respectively. While the main sequence of member stars can be clearly seen in the CMD of $\\chi$~Persei with only a small contamination of field stars, the CMD for NGC\\,6910 is severely contaminated. Due to the large (and differential) reddening, the cluster main sequence in Fig.~\\ref{6910-cmd} is smeared and located right of the bluest field stars. ",
        "conclusions": "Since we discovered many $\\beta$~Cep stars, we will focus on the differences in the frequency spectra of these stars in both clusters in our short discussion. However, there are other main-sequence pulsators in both clusters, in particular slowly pulsating B (SPB) and $\\delta$~Sct stars, that can also be used in seismic modelling.  In general, the separation in frequency between p and g modes in massive stars is relatively good for non-evolved stars. In the course of main sequence evolution, however, the frequencies of g modes increase eventually replacing the frequencies of p modes through avoided crossings. These modes are called `mixed' due to their mixed character, g in the interior and p in the envelope. Although the occurrence of a p mode (or at least a mixed mode) is a prerequisite for the classification of a star as a $\\beta$~Cep-type variable, for lack of mode identification we usually use a more practical definition of a $\\beta$~Cep-type star. For example, early B-type stars showing periodic variations with periods shorter than 0.3~d can be termed $\\beta$~Cep stars (see, e.g., Sterken \\& Jerzykiewicz 1993). Following this definition, we classify eleven stars in $\\chi$~Persei  and nine in NGC\\,6910 as $\\beta$~Cep stars (see Figs.~\\ref{chiper-cmd} and \\ref{6910-cmd}). This makes the two clusters comparable to four other young open clusters known to be rich in $\\beta$~Cep stars: NGC\\,3293, NGC\\,4755, NGC\\,6231 and h~Persei (NGC\\,869).  \\begin{figure}[!t] \\includegraphics[width=3.2in]{Camp_Lan_Fig3.eps} \\caption{Schematic frequency spectra for nine $\\beta$~Cep stars in NGC\\,6910. The vertical line denotes a period equal to 0.3~d. The two light gray bars denote a combination mode and an harmonic. See text for the explanation of the dark gray bars.} \\label{6910-frsp} \\end{figure}  \\begin{figure}[!t] \\includegraphics[width=3.2in]{Camp_Lan_Fig4.eps} \\caption{The same as in Fig.~\\ref{6910-frsp} but for eleven $\\beta$~Cep stars in $\\chi$~Persei.} \\label{chi-frsp} \\end{figure}  Accounting for differential extinction, we have plotted in Fig.~\\ref{6910-frsp} the frequency spectra of the nine $\\beta$~Cep stars found in NGC\\,6910 going from the brightest (top) to the faintest (bottom). This is also a sequence of decreasing mass. In addition, because of the same age and different masses, the sequence goes from more evolved to less evolved stars. A striking feature that can be seen in Fig.~\\ref{6910-frsp} is a strip of modes with increasing frequency when we go from more massive (bigger) to less massive (smaller) stars. This is exactly what is expected for p modes. Additionally, three $\\beta$~Cep stars in this cluster, NGC\\,6910-14, 27, and 36\\footnote{The numbers we use follow the Oosterhoff (1937) for $\\chi$~Persei and the WEBDA (http://www.univie.ac.at/webda) designation for NGC\\,6910.}, show the presence of low-frequency modes, which can be interpreted as g mo\\-des. These stars are therefore good candidate hybrid $\\beta$~Cep/SPB stars. Moreover, we see that, in the four faintest stars, another group of modes occurs (shown with dark gray bars). All they have frequencies in the range between 4 and 5.5~d$^{-1}$. Without detailed modelling it is difficult to say if they are p or g modes; f mode(s) are also a possibility.  \\begin{figure}[!t] \\includegraphics[width=3.2in]{Camp_Lan_Fig5.eps} \\caption{The same as in Fig.~\\ref{6910-frsp} but for five other stars from $\\chi$~Persei showing variability with periods shorter than 0.3~d.} \\label{chi-midb} \\end{figure}  The frequency spectra for $\\beta$~Cep stars in $\\chi$~Persei are presented in Fig.~\\ref{chi-frsp}. When we go from the most to the least massive $\\beta$~Cep stars in this cluster (from top to bottom in Fig.~\\ref{chi-frsp}), the frequencies of the detected modes do not change as monotonically as in NGC\\,6910. The most likely explanation of this behaviour is the fast rotation of stars in $\\chi$~Persei. The projected rotational velocities, $V\\sin i$, were measured for many stars in $\\chi$~Persei and for at least seven $\\beta$~Cep stars from our sample the values of $V\\sin i$ exceed 100~km\\,s$^{-1}$ (Strom et al.~2005, Huang \\& Gies 2006). Fast rotation results in considerable shifts in frequency for non-axi\\-sym\\-met\\-ric modes. Stars in NGC\\,6910 have no measurements of their rotational velocities, but there is only one Be star known in this cluster, whereas in $\\chi$~Persei 20 Be stars are known (Keller et al.~2001). This might mean that, on average, $\\beta$~Cep-type stars in NGC\\,6910 rotate much slower than stars of this type in $\\chi$~Persei. We also note two hybrid stars, Oo\\,2114 and 2185, with modes that have frequencies below 1~d$^{-1}$, presumably g modes.  Fast rotation in $\\chi$~Persei might also be responsible for the occurrence of several stars with periods shorter than 0.3~d which we do not classify as $\\beta$~Cep stars because, as can be judged from Fig.~\\ref{ubbv}, they are mid- to late-B type stars. We divided them into two groups which are shown with different symbols in Figs.~\\ref{chiper-cmd} and \\ref{ubbv}. The first group (shown with filled squares) consists of five stars. Their frequency spectra are shown in Fig.~\\ref{chi-midb}. These stars show bi-periodic variability; the period of one term is shorter than 0.3~d, the period of the other, longer. Four of the five stars have measured projected rotational velocities. They range from 150 to 250~km\\,s$^{-1}$. In addition, Oo\\,2091 is known as Be star.  According to many different observations, including ground-based data (Uytterhoeven et al.~2007), satellite data (Walker et al.~2005, Huat et al.~2009, Guti\\'errez-Soto et al.~2010) and the surveys in the Magellanic Clouds (Ko\\-{\\l}acz\\-kow\\-ski et al.~2006), this type of behaviour seems to be very common in fast-rotating stars. Their frequency spectra contain usually a large number of terms with frequencies clustering around two values; the ratio of these values is usually close to 2.  Although not all five stars in $\\chi$~Persei we discussed above are known as Be stars, following Walker et al.~(2005), we tentatively classify them as SPB(e) stars.  \\begin{figure}[!t] \\includegraphics[width=3.2in]{Camp_Lan_Fig6.eps} \\caption{The $(U-B)$ vs.~$(B-V)$ diagram for the discussed stars in $\\chi$~Persei and NGC\\,6910. The thick line stands for the relation for unreddened luminosity class V stars (Caldwell et al.~1993). The two thin continuous lines denote the same relation shifted by the reddening vectors corresponding to $E(B-V)$ = 0.46 and 0.68~mag, a minimum and maximum reddening in $\\chi$~Persei. Similar relations for NGC\\,6910 are shown with dotted lines for $E(B-V)$ = 0.9 and 1.3~mag. Reddening lines, delimiting the range of spectral types for $\\beta$~Cep stars, are shown with dashed lines.} \\label{ubbv} \\end{figure}  The second group consists of five other variable mid-to late-B type stars, Oo 2752, 2611, 2019, 2406, and 2228, shown with open squares in Figs.~\\ref{chiper-cmd} and \\ref{ubbv}. However, compared to the first group, they show only terms with a period shorter than 0.3~d. If due to pulsations, the short periods can be explained by fast rotation: rotation can easily shift the frequencies of non-axisymmetric prograde g modes to the observed values. It is therefore quite possible that they are also SPB(e) stars in which long-period (LP), i.e., low-frequency term(s) are not detected due to low amplitudes.  The results of the variability surveys in $\\chi$~Persei and NGC\\,6910 which we partially presented here form a very good basis for the ensemble asteroseismology in both clusters.  The fast rotation for $\\chi$~Persei stars is a complication, but, on the other hand, this cluster may help us to understand better the interaction between pulsations and rotation. There are two other open clusters, h Persei and NGC\\,3293, with a similar amount of data obtained during other campaigns. They are of similar age as $\\chi$~Persei and also contain many fast-rotating stars."
    },
    "1004/1004.5197_arXiv.txt": {
        "abstract": "The slope of the nuclear symmetry energy at saturation density $L$ is pointed out as a crucial quantity to determine the mass and width of neutron-star crusts. This letter  clarifies the relation between $L$ and the core-crust transition. We confirm that the transition density is soundly correlated with $L$ despite differences between models, and we propose a clear understanding of this correlation based on a generalised liquid drop model. Using a large number of nuclear models, we evaluate the dispersion affecting the correlation between the transition pressure $P_t$ and $L$. From a detailed analysis it is shown that this correlation is weak due to a cancellation between different terms. The correlation between the isovector coefficients $K_{\\rm sym}$ and $L$ plays a crucial role in this discussion. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.2254_arXiv.txt": {
        "abstract": "We present a multiscale approach to measurements of galaxy density, applied to a volume-limited sample constructed from the Sloan Digital Sky Survey Data Release 5 (SDSS DR5). We populate a rich parameter space by obtaining independent measurements of density on different scales for each galaxy. This parameterization purely relies on observations and avoids the implicit assumptions involved, e.g., in the construction of group catalogues. As the first application of this method, we study how the bimodality in galaxy colour distribution (specifically $u-r$) depends on multiscale density.  The $u-r$ galaxy colour distribution is described as the sum of two gaussians (a red and a blue one) with five parameters: the fraction of red galaxies ($\\fr$) and the position and width of the red and blue peaks ($\\mur$, $\\mub$, $\\sigmar$ and $\\sigmab$). Galaxies mostly react to their smallest scale ($<0.5\\Mpc$) environments: in denser environments red galaxies are more common (larger $\\fr$), redder (larger $\\mur$) and with a narrower distribution (smaller $\\sigmar$), while blue galaxies are redder (larger $\\mub$) but with a broader distribution (larger $\\sigmab$).  There are residual correlations of $\\fr$ and $\\mub$ with $0.5 - 1\\Mpc$ scale density, which imply that total or partial truncation of star formation can relate to a galaxy's environment on these scales.  Beyond $1\\Mpc$ ($0.5\\Mpc$ for $\\mur$) there are no positive correlations with density.  However $\\fr$ {\\it anti-correlates} with density on $>2\\Mpc$ scales at fixed density on smaller scales, and $\\mur$ anti-correlates with density on $>1\\Mpc$ scales.  We examine these trends qualitatively in the context of the halo model, utilizing the properties of haloes within which the galaxies are embedded, derived by \\citet{Yang07} and applied to a group catalogue.  This yields an excellent description of the trends with multiscale density, including the anti-correlations on large scales, which map the region of accretion onto massive haloes.  Thus we conclude that galaxies become red only once they have been accreted onto haloes of a certain mass.  The mean colour of red galaxies $\\mur$ depends positively only on $<0.5\\Mpc$ scale density, which can most easily be explained if correlations of $\\mur$ with environment are driven by metallicity via the enrichment history of a galaxy within its subhalo, during its epoch of star formation. ",
        "introduction": "\\label{sec:intro}  \\subsection{Motivation}  The evolution of galaxies is heavily intertwined with the growth of the dark matter dominated structure in which they live, as baryons react to their local gravitational potential.  It is therefore one of the most challenging goals of observational astronomy to explain the continuously evolving galaxy population in the context of a hierarchical universe.  Most fundamentally perhaps, galaxies exhibit a distinct bimodality of properties which correlates strongly with the properties of the embedding potential.  This is true whether the potential is measured on galaxy scales via properties such as luminosity, stellar or dynamical mass, or on larger scales via measurement of the galaxies' local environment.  In particular, galaxies living in deeper potential wells are more likely to have formed stars at earlier times, both in terms of the average stellar age \\citep[e.g.][]{Thomas05,Smith06} and the last episode of star formation \\citep[and thus a lower fraction are continuing to form stars, e.g.][]{Lewis02,Baldry04,Balogh04Ha,Balogh04,Kauffmann04}. Such galaxies are apparently unable to restart star formation, since they have no stable access to a suitable reservoir of gas which can cool to form a star-forming disk \\citep[e.g.][]{Gallagher75,Caon2000}. Simultaneously the morphology of these galaxies must evolve from the typical rotation-dominated spiral or irregular morphologies of star-forming galaxies to form a pressure dominated elliptical, or a bulge plus smooth disk lenticular \\citep[e.g.][]{Dressler80,Dressler97,Wilman09}.  The precise nature of these transformations remains controversial, despite evidence from simulations and observations for various mechanisms which can explain some or all of these observations.  Galaxy mergers are especially attractive in the context of a hierarchical Universe, and can explain both morphological changes and the rapid exhaustion of cold gas, although additional physics may be required to keep a galaxy from reforming a gas disk \\citep[e.g.][]{Springel05,Hopkins09,Johansson09}. Stripping or evaporation of cold and/or hot gas associated to an infalling galaxy will take place within a dense hot intra-cluster medium \\citep[e.g.][]{GunnGott72,Chung09,LTC}, although it is unclear whether these processes can be important within lower density environments \\citep{Kawata08,McCarthy08,Jeltema08}. Tidal interactions between galaxies can drive gravitational instabilities within galaxies, with gas loss or triggering of star formation as likely consequences. Numerous fast interactions (harrassment) can also drive morphological evolution \\citep[e.g.][]{Moore99}.  To statistically describe the evolutionary path of galaxies, it is necessary not only to take a firm grasp of the relevant physical processes but also to place robust observational constraints on precisely how the galaxy population reacts to its environment.  It has been customary in this field to use neighbouring galaxies as test particles in order to `measure' the local environment of each galaxy. One computes a `local density' by simply counting neighbours within some fixed radius and velocity range or, almost equivalently, by computing the distance to the N$^{th}$ nearest neighbour, where N is typically chosen to be 5 or 10 (\\citet{Dressler80,Balogh04Ha}, see \\citet{Cooper05,Kovac09} for a discussion of these and more complex methods).  As with so many traditional measures, this measurement is fairly arbitrary in nature \\footnote{N is typically selected to obtain sufficient S/N whilst retaining a ``local'' measurement.}, and it is not easy to see precisely how it relates to the underlying density field of galaxies and dark matter.  The two point correlation function or power spectrum of galaxies on the other hand, measure the excess probability over random of finding two galaxies with a given separation. The comparison of galaxy subsets selected by e.g. colour, provides a measure of the relative bias of these galaxy types as a function of the scale of pair separation, diagnosing the scale at which one class dominates relative to another. \\citep[e.g.][]{Norberg01,Budavari03,Zehavi05}.  Marked correlation functions weight the galaxies in each pair by a given property, normalizing by the unweighted correlation function \\citep[e.g.][]{Skibba06}. This examines how galaxies with large or small values of the chosen property are clustered. The bias or marked correlation as a function of scale are intimitely related to the details of how galaxy properties depend on their environments.  These statistics simply provide the mean behaviour of galaxies by tracing their overdensity on different scales which, however, are themselves closely correlated with one another. A better approach would be to examine how galaxy properties correlate with overdensity on one scale for a fixed overdensity on another, independently computed scale.  Now that such a large, homogeneous sample of galaxies is available with the Sloan Digital Sky Survey (SDSS), there have been a few attempts at isolating such behaviour on different scales. \\citet{Kauffmann04}, \\citet{Blanton06} and \\citet{Blanton07} conclude that galaxies only react to their environments on $<1\\Mpc$ scales, with no significant residual dependence on larger scale overdensity. This contradicts the apparent dependence on large scales found by \\citet{Balogh04Ha}, which may have resulted from low signal to noise measurement of density on small scales \\citep{Blanton06}. These works, however, contained correlated measurements of density on both scales.  \\subsection{A New Method}  In this paper we approach this problem in a different way.  Firstly, we parameterize the environment using annular, non-overlapping measurements of density on different scales to ensure that they are (formally) independent. Secondly, we consider the $u-r$ colour of galaxies and utilize the remarkable fact that its distribution is bimodal and can be modelled as the sum of two gaussians.  This relates to the bimodality of galaxy properties, since galaxies dominated by old stellar populations have red $u-r$ colours and actively star-forming galaxies are blue. \\citet{Baldry04} initially demonstrated the power of this technique, illustrating the low chi-squared values produced with such a fit to $u-r$ SDSS colours, and demonstrating that the luminosity dependence of galaxy colours was strong and easily parameterized using this model.  \\citet{Balogh04} and \\citet{Baldry06} have since extended this method, with a powerful demonstration of the strong and fundamental dependence of the fraction of red galaxies on environment (measured using the distance to the $4^{th}$ and $5^{th}$ nearest neighbours) and stellar mass.  The parameterization of the double gaussian fit (the fraction of red galaxies $\\fr$, and the position and width of each peak) provides five quantities which can be interpreted in terms of changes in the spectral energy distribution and the potential physical trigger, and which can be correlated with environment measured on different scales.  \\subsection{Expectations from the Halo Model}  Finally, to ease the physical interpretation of our results, we would like to know how measurements of density map onto the dark matter dominated `cosmic web'.  To this goal we rely on the models that describe the clumpy distribution of dark matter in the Universe in terms of haloes, i.e. the potential wells in which galaxies form, infall, virialize and interact with one another and with the intra-cluster/group medium (ICM/IGM) \\citep[e.g.][]{White78, Cooray02, Diaferio01}.  The halo model (which applies on both galaxy and group/cluster scales) is a simplification of reality, but has been hugely successful in describing the dynamics and large scale structure statistics of galaxies \\citep[e.g.][]{Rubin83,Berlind03}.  In particular, the two point correlation function appears to demonstrate an inflection at $\\sim2\\Mpc$ which is attributed to the switch from sub-halo scales (the one halo term) to halo-halo scales (the two halo term) \\citep[e.g.][]{Zehavi04,Zehavi05,Cooray06}.  Models of galaxy formation and evolution pay homage to the halo model by assuming galaxies to care about their environment only in terms of their embedding dark matter halo \\citep[e.g.][]{Kauffmann93,Cole00}.  This puts a natural scale on the expected dependence of galaxy properties on their environments: they should not care about scales beyond their halo.  However simulations also show that the formation (collapse) time of a halo depends upon the overdensity on larger scales \\citep[e.g.][]{Sheth04,Gao05}. In this case, one might expect galaxies to have evolved further in a region of the Universe which is more overdense on large scales.  \\subsection{Structure of the paper}  Section~\\ref{sec:sample} introduces the sample used in this paper, selected from the SDSS Data release 5 (DR5).  In Section~\\ref{sec:density} we describe in detail the method used to compute the density of galaxies on different scales, paying close attention to the completeness corrections required to ensure a robust measurement in the presence of incompleteness in SDSS DR5.  Multiscale density information for the 73662 galaxies of our photometrically complete sample are available for public use on request to the authors.  The method of fitting a double gaussian model is described in Section~\\ref{sec:fitting}.  Our results are presented in Section~\\ref{sec:results}.  Firstly we examine the simple dependence of the galaxy colour distribution on luminosity and density on a single scale in Section~\\ref{sec:deplumdens}.  Then, in Section~\\ref{resultsscale}, we present our main results in which the importance of different scales are tested independently.  This forms the purely observational part of the paper.  In Section~\\ref{sec:interpretation} we utilize a catalogue of embedding haloes (groups) which has been assigned to SDSS galaxies within the context of a halo model and with a few implicit assumptions \\citep{Yang07}.  We examine the dependence of halo-based properties such as mass and halo-centric radius on the density computed on different scales and compare this to the way in which the galaxy colour distribution depends upon the same measurements of density. What emerges is a remarkably consistent picture, in which the halo model is extremely successful at explaining the colours of galaxies: smaller scale density can drive some aspects of colour evolution, but star forming galaxies do not seem to feel the halo environment prior to infall.  This exercise should provide the first, qualitative step towards a fully quantative, model-independent description of the way in which galaxies trace their environment.  Models can then be constrained independently of the assumptions that are required in the construction of group catalogues or modeling of the correlation function. The rich parameter space of multiscale densities will help overcoming degeneracies and difficulties in interpreting single aperture measurements of local density.  In this paper we assume the cosmological parameters: $(\\Omega_M,\\Lambda,h) = (0.3,0.7,0.75)$. ",
        "conclusions": "\\label{sec:conclusions}  We have developed a new multiscale method to characterize the dependence of galaxy properties on environment.  This provides a model-independent, rich parameter space in which to evaluate galaxy properties.  In this paper we have examined the dependence of the $u-r$ colour distribution for galaxies in the luminosity range $-21.5\\leq\\Mr\\leq-20$ on multiscale density.  This is parameterized using the double gaussian model, with parameters $\\fr$ (the fraction of red galaxies), $\\mur$ and $\\sigmar$ (the position and width of the red peak) and $\\mub$ and $\\sigmab$ (the position and width of the blue peak).  We confirm and extend known trends with small scale ($\\lesssim 0.5 - 1 \\Mpc$) density: \\begin{itemize} \\item{Galaxies in denser environments are more likely to be red ($\\fr$), implying an almost complete truncation of star formation.} \\item{If they are still blue (star forming), then they are also likely to be relatively redder in denser environments ($\\mub$) although with greater scatter ($\\sigmab$), implying partial but varied truncation of gas supplies, possibly relating to truncated gas disks such as those observed in the Virgo cluster.} \\item{Red Galaxies are redder ($\\mur$) and with less scatter ($\\sigmar$) in denser environments.  This implies that they either form earlier and/or are more metal-rich, and that this formation path is also less varied than galaxies in less dense environments. This depends upon the environment of a galaxy during its epoch of star formation.} \\end{itemize}  On larger scales: \\begin{itemize} \\item{No parameter correlates positively and significantly with density on scales $>1\\Mpc$, excluding trends seen for relatively isolated galaxies where correlated noise in the measurement of density might be important.  This confirms the results of \\citet{Kauffmann04}, \\citet{Blanton06} and \\citet{Blanton07}.} \\item{Although all parameters correlate most strongly with density on $<0.5\\Mpc$ scales, a residual positive correlation of $\\fr$ and $\\mub$ with $0.5-1\\Mpc$ scale density at fixed smaller scale density implies that total or partial truncation of star formation can relate to a galaxy's environment on these scales. A $1\\Mpc$ diameter halo corresponds to a mass of $\\sim2\\times10^{13}\\Msol$ and a typical crossing time on the order of $4\\Gyr$, and one or both of these parameters must relate to a truncation mechanism which is active on these scales.} \\item{On $>2\\Mpc$ scales $\\fr$ {\\it anti-correlates} with density at fixed smaller scale density, and $\\mur$ anti-correlates with density on $>1\\Mpc$ scales.  In other words galaxies are more likely to have ongoing star formation, or if not then to be relatively blue, when the large scale density is relatively high for the density on smaller scales (they live at projected distances $\\sim1-3\\Mpc$ from a local overdensity).} \\end{itemize}  We interpret these trends qualitatively using the halo model.  To make this comparison we utilize the halo catalogue of Y07 which is constructed using a friends of friends group finder, and under the assumptions that galaxies live in haloes, and that the mass of these haloes has a one to one relationship with the total galaxy light.  The properties of embedding haloes has been cross-correlated with our galaxy catalogue to examine how halo properties trace multiscale density.  By comparing this dependence to that of galaxy colours, we find:  \\begin{itemize} \\item{Qualitatively, the multiscale dependence of $\\fr$ and $\\mub$ is remarkably comparable to that seen in the fraction of galaxies living in haloes above a given threshold mass in the range $\\sim10^{12.5-13.5}\\Msol$.  This is consistent with a scenario in which galaxies can experience truncation of their star formation at some time after they are accreted onto a halo.} \\item{The halo model offers a simple and yet profound explanation for anti-correlations with larger (particularly $>2\\Mpc$) scale density.  For galaxies at large radial distances from the centre of massive haloes the halo core is included in measurements of density on large scales.  Thus this region of parameter space traces the accretion region of haloes (effectively marking the crossover from the one to the two halo term in the correlation function).  Anti-correlations on these scales therefore trace the accretion region of haloes.  That many such galaxies are still blue implies that the influence of environment is felt at or after accretion onto a massive halo.} \\item{$\\mur$ correlates positively with density only on $<0.5\\Mpc$ scales, implying that the red colour is set by physics within subhaloes of this size. In this case we suggest metallicity effects must drive these trends, since we know that age must correlate with larger scale densities: i.e. galaxies become red due to the influence of $0.5-1\\Mpc$ scales ($\\fr$ correlates positively with $\\dscc$). However we cannot rule out age effects if the $0.5-1\\Mpc$ scale dependence of $\\mur$ is negated due to a combination of halo mass and radial trends.} \\end{itemize}  Whilst we have restricted ourselves in this paper to qualitative comparisons within the context of the halo model, we advocate the use of multiscale density to quantitatively examine physically motivated models (from simple halo occupation models to semi-analytically populated haloes).  This provides a uniquely rich and observationally motivated parameter space to examine the dependence of galaxy properties on environment."
    },
    "1004/1004.4644_arXiv.txt": {
        "abstract": "As part of our search for young M dwarfs within 25 pc, we acquired high-resolution spectra of 185 low-mass stars compiled by the NStars project that have strong X-ray emission. By cross-correlating these spectra with radial velocity standard stars, we are sensitive to finding multi-lined spectroscopic binaries. We find a low-mass spectroscopic binary fraction of 16\\% consisting of 27 SB2s, 2 SB3s and 1 SB4, increasing the number of known low-mass SBs by 50\\% and proving that strong X-ray emission is an extremely efficient way to find M-dwarf SBs. WASP photometry of 23 of these systems revealed two low-mass EBs, bringing the count of known M dwarf EBs to 15. BD -22 5866, the SB4, is fully described in \\cite{shko08} and CCDM~J04404+3127~B consists of a two mid-M stars orbiting each other every 2.048 days. WASP also provided rotation periods for 12 systems, and in the cases where the synchronization time scales are short, we used $P_{rot}$ to determine the true orbital parameters. For those with no $P_{rot}$, we use differential radial velocities to set upper limits on orbital periods and semi-major axes. More than half of our sample has near-equal-mass components ($q >$ 0.8). This is expected since our sample is biased towards tight orbits where saturated X-ray emission is due to tidal spin-up rather than stellar youth. Increasing the samples of M dwarf SBs and EBs is extremely valuable in setting constraints on current theories of stellar multiplicity and evolution scenarios for low-mass multiple systems. ",
        "introduction": "The multiplicity of stars is an important constraint of star formation theories as most stars form as part of a binary or higher-order multiple system (e.g.~\\citealt{halb03,bate09}). Moreover, double- (or multi-) lined spectroscopic binaries (SBs) allow precise determination of dynamical properties including the mass ratio. The analysis of photometric and spectroscopic data of eclipsing binaries (EBs) provides radii, temperatures, luminosities and masses, arguably the most important stellar parameter, for two stars with same age and metallicity \\citep{Lastennet} making them vital in calibrating stellar evolutionary models.  M~dwarf EBs are particularly important since the radii of active M~dwarfs are known to be 10-15\\% larger than existing models predict \\citep{Mercedesa,Mercedesb,Torres09}. This discrepancy is thought to be caused by magnetic fields on active M~dwarfs which inhibit convection \\citep{Mercedesc,Chabrier} and/or missing opacity sources in the models \\citep{berg06_d}.   Further study of the problem requires the identification and detailed study of more M~dwarf EBs with a range of properties (e.g.~masses, activity levels, metallicities, ages).  Though M dwarfs are ubiquitous in the Galaxy, composing 75\\% of known stars \\citep{boch08}, their intrinsic faintness makes them difficult and costly to observe, and thus the multiplicity of M dwarfs has been a challenge to measure. A binary fraction of 57\\% has been well established for G dwarfs \\citep{duqu91}, and there is clear consensus that the binary fraction of M dwarfs is significantly less than that. Published values range from 25\\% \\citep{lein97} to 42\\% \\citep{fisc92} with the largest uncertainties likely due to incompleteness corrections. In addition to supplying key constraints for star formation theories, completing the M dwarf binary census of the solar neighborhood, which to date, is (near)-complete out to only 9 pc \\citep{delf04}, is yet another goal in finding the nearby low-mass SBs.   \\cite{delf04} presented what may be the most complete statistical study of M dwarfs, including both spectroscopic and visual binaries.  They conclude a binary fraction of 26 $\\pm$ 3\\% and that for M dwarfs, as for G dwarfs, the mass ratio ($q = M_B/M_A$) distribution is a function of orbital period with most shorter period binaries having near equal component masses, while wide binaries ($P_{orb}>$50 days) have a flat \\emph{q} distribution. The different distribution in the two samples may point to two distinct formation mechanisms, one for long-period and one for short-period orbits.   Here we report on 30 low-mass SBs detected from 185 X-ray-selected M dwarfs in the solar neighborhood. Six of the targets were included in the Gliese catalog of which only one was previously detected to be a SB. Prior to this work, 46 M dwarf SBs were published in the literature \\citep{duqu91,delf99,fisc92} with an additional 13 low-mass EBs (\\citealt{shko08} and references therein and \\citealt{blak08}). Thus our work increases the known sample of SBs by 50\\%.  We also searched the WASP photometric database in which we found 2 EBs and measured rotation periods of 12 binaries.  Though rotation periods of single stars are important age indicators (e.g.~\\citealt{barn07}), for short-period binaries, where tidal locking is almost certain, the rotation periods offer true orbital periods instead. ",
        "conclusions": "\\label{summary}   Of our sample of 185 X-ray bright M dwarfs, we find a low-mass, multi-lined spectroscopic binary fraction of 16\\%. These 30 SBs are composed of 27 SB2s, 2 SB3s and 1 SB4, increasing the number of known low-mass SBs by 50\\% and proving that strong X- ray emission is an extremely efficient way to find M-dwarf SBs.  To search for single-lined SBs (SB1), we observed two epochs (separated typically by 2--3 months) of 65 of the 185 targets, none of which showed a significant RV variation between visits to the level of 1 km~s$^{-1}$. This implies that the single-lined binary fraction of stars with orbital periods of less than about 1.5 years in our sample is very low, less than 1.5\\%, and that M dwarf binaries with low mass ratios ($q \\ll 1$) are rare. It is possible that up to 4\\%\\footnote{This 4\\% limit is based on the time a close-in low-mass binary with an orbital period of 5 days would spend near conjunction such that the RVs of the two components would not produce resolved peaks in the CCF.} of the stars with a single observation are indeed double-lined SBs if the systems were in conjunction at the times of the observation. Combining this with the $\\leq$1.5\\% chance of observing an SB1, there are at most a handful of undiscovered SBs in the original 185 ROSAT-selected targets, setting an upper limit of 19\\% to the SB fraction (with $P_{orb} \\lesssim 1.5$ years) of our sample.  WASP photometry of 23 of these systems revealed two low-mass EBs, bringing the count of known M dwarf EBs to 15. The WASP data also provided rotation periods for 12 systems, and in the cases where the synchronization time scales are short, orbital periods and semi-major axes.  This X-ray bright sample of 30 SBs is strongly biased towards high-$q$, tidally-synchronized binaries. In addition to being in short-period orbits, they are also relatively bright, making them excellent targets for a spectroscopic monitoring program to measure the component velocities necessary to determine the Keplerian orbital parameters for more precise mass ratios, and in the case of the eclipsing systems, the individual masses needed to test evolutionary models."
    },
    "1004/1004.2039_arXiv.txt": {
        "abstract": "We investigate the origin of the shape of the globular cluster (GC) system scaling parameters as a function of galaxy mass, i.e. specific frequency ($S_N$), specific luminosity ($S_L$), specific mass ($S_M$), and specific number ($\\hat{T}$) of GCs.  In the low-mass galaxy regime ($M_V\\!\\ga\\!-16$ mag) our analysis is based on {\\it HST/ACS} observations of GC populations of faint, mainly late-type dwarf galaxies in low-density environments. In order to sample the entire range in galaxy mass ($M_V\\!=\\!-11$ to $-23$\\,mag $=10^6\\!-\\!10^{11}L_\\odot$), environment, and morphology we augment our sample with data of spiral and elliptical galaxies from the literature, in which old GCs are reliably detected. This large dataset confirms (irrespective of galaxy type) the increase of the specific frequencies of GCs above and below a galaxy magnitude of $M_V\\simeq-20$\\,mag. Over the full mass range, the $S_L-$value of early-type galaxies is, on average, twice that of late-types. To investigate the observed trends we derive theoretical predictions of GC system scaling parameters as a function of host galaxy mass based on the models of \\cite{Dekel&Birnboim06} in which star-formation processes are regulated by stellar and supernova feedback below a stellar mass of $3\\!\\times\\!10^{10}{\\cal M}_\\odot$, and by virial shocks above it. We find that the analytical model describes remarkably well the shape of the GC system scaling parameter distributions with a  universal {\\it specific GC formation efficiency}, $\\eta$, which relates the total mass in GCs to the total galaxy halo mass. Early-type and late-type galaxies show a similar mean value of $\\eta\\!=\\!5.5\\times10^{-5}$, with an increasing scatter towards lower galaxy masses. This can be due to the enhanced stochastic nature of the star and star-cluster formation processes for such systems. Some massive galaxies have excess $\\eta$ values compared to what is expected from the mean model prediction for galaxies more luminous than $M_V\\simeq-20$\\,mag ($L_V\\gtrsim10^{10}L_\\odot$). This may be attributed to a very efficient early GC formation, less efficient production of field stars or accretion of predominantly low-mass/luminosity high$-\\eta$ galaxies, or a mixture of all these effects. ",
        "introduction": "\\vspace{4mm}One of the very first stellar systems to form in the early Universe are old globular clusters (GCs), which are observed in galaxies of all morphological types. Globular clusters are massive agglomerations of gravitationally bound stars, the majority of which formed almost simultaneously out of gas with similar chemical composition. As fossil records of the first star formation episodes of their host galaxies, the distribution of their integrated properties (age, mass, metallicity, structural parameters) as well as the general properties of all GCs in a galaxy (total numbers, spatial and dynamical distributions) hold important clues to the initial physical conditions at which they have formed and evolved. For that reason, the global properties of globular cluster systems (GCSs) have long been recognized as promising tools to study the major galaxy star formation episodes and to serve as observational constraints to differentiate between various models of galaxy formation \\protect{\\cite[][and references therein]{Kissler-Patig00,vdBergh00, Harris91,Harris03,Brodie06}}.  One of the most commonly used parameters to describe GCSs is the \\emph{specific frequency} ($S_N$), i.e. the number of GCs per unit galaxy luminosity. In essence, $S_N$ measures the formation efficiency of GCs relative to field stars. GCs and field stars are linked through the dissolution of star clusters due to various mechanisms (e.g. tidal shocks, cluster relaxation), which shapes the initial power-law to the observed, present-day Gaussian globular cluster luminosity function \\cite[GCLF, see e.g.][and references therein]{Fall&Zhang01, Goudfrooij04, Goudfrooij07, Gieles06, McLaughlin&Fall08, Kruijssen&Zwart09, Elmegreen10}. The $S_N$ parameter was initially introduced by \\cite{Harris&vdBergh81} as a measure of the richness of GCSs in elliptical galaxies. Since then $S_N$ has been applied by numerous studies to galaxies of different morphological types from early-type, quiescent ellipticals to actively star-forming, late-type spirals, irregulars, and interacting/merger galaxies, covering the entire range in galaxy mass (from giants to dwarfs) and environments (from galaxies in dense clusters to such in loose groups and in the field). Wide-field ground-based and deep Hubble Space Telescope (HST) studies show that the $S_N$ value varies \\emph{greatly} among galaxies, particularly among the most luminous ellipticals and low-mass dwarf galaxies \\cite[e.g.][]{Harris91,Harris01,Wehner08,Peng08}. Spiral galaxies tend to show much less scatter in their $S_N$, with values in the range $0.5-2$ \\citep{Ashman&Zepf98,Goudfrooij03,Chandar04,Rhode07}. The general trend observed for early-type dwarf elliptical (dE) galaxies is that their average $S_N$ value increases with decreasing galaxy luminosity from a few to a few tens \\citep{Durrell96, Miller98,Miller&Lotz07,Peng08}. It was shown recently that a similar behavior of $S_N$ holds for late-type dwarfs \\cite[dIrrs,][]{Seth04, Olsen04, Sharina05, Georgiev08, Puzia&Sharina08}. Conversely, in the high-mass galaxy regime, a significant increase of the $S_N$ values is observed for the most luminous galaxies \\cite[gEs and cDs, e.g.][]{Rhode05, Peng08, Harris09}. As the stellar populations (colours and integrated light spectral indices) of the latter galaxies are not significantly different from galaxies 1-2 mag fainter, such an upturn in the $S_N$ distribution implies that the assembly history of the most massive galaxies must have been different from those of less massive galaxies.  The overall trend for $S_N$ is that with increasing galaxy luminosity from $M_V\\simeq-11$\\,mag to $M_V\\simeq-20$\\,mag, $S_N$ decreases from the range $\\sim0\\!-\\!100$ to $S_N\\sim0.5\\!-\\!3$, respectively. For galaxies more luminous than $M_V\\simeq-20$\\,mag the $S_N$ increases again (to $\\sim 10$).  To explain the increasing fraction of GCs over field stars in the most massive galaxies several studies suggested that the mass in GCs is proportional to the total gas mass supply \\cite[][]{West95,Blakeslee97, Blakeslee99}. In particular, \\cite{McLaughlin99} investigated the $S_N$ behavior for the entire early-type galaxy mass range, from giants to dwarfs. He showed that the enhanced $S_N$ for the most massive galaxies considered in his sample (NGC\\,1399, M\\,87, M\\,49) can be accounted for if the mass in GCs is normalized to the total baryonic mass of the host (stellar plus the hot X-ray emitting gas), implying a constant baryonic GC formation efficiency $\\epsilon=0.26\\%$. The parameter $\\epsilon$ represents the efficiency of converting baryonic matter into GCs, $\\epsilon\\equiv{\\cal M}_{\\rm GCS}/{\\cal M}_b$. However, as suggested by \\cite{Blakeslee97} and \\cite{Blakeslee99}, perhaps more fundamental is the ratio between the mass in GCs and the {\\it total\\/} galaxy halo mass ${\\cal M}_h$, (i.e., $\\eta_h\\equiv{\\cal M}_{\\rm GCS}/{\\cal M}_h$) for which they obtain $\\eta_h\\simeq1.71\\times10^{-4}$. More recently, \\cite{Kravtsov&Gnedin05} in a high-resolution cosmological simulation for galaxies with ${\\cal M}_h  >10^9{\\cal M_\\odot}$ find ${\\cal M}_{\\rm GCS}= 3.2\\times10^6{\\cal M}_\\odot ({\\cal M}_h/10^{11}{\\cal M}_\\odot)^{0.13}$, which predicts $\\epsilon_h=(2-5)\\times10^{-5}$. Assuming a cosmological baryon fraction $f_b=0.17$ \\citep{Hinshaw09} brings the observed and theoretically predicted values of the GC formation efficiency to a common value of $\\sim10^{-5}$ (see Sect.\\,\\ref{Sect:GCefficiency}). Following up on the idea of the proportionality between the mass in GCs and the host halo mass using a statistical stellar-halo mass relation from $\\Lambda$CDM simulations, recent studies with much larger galaxy samples find a good approximation to the data \\citep{Peng08} and a similar result for the GC formation efficiency \\cite[$\\eta_h\\!=\\!7\\times10^{-5}$,][]{Spitler09}. In the low-mass galaxy regime, \\cite{Peng08} find evidence for an environmental bias: the majority of dwarf galaxies with relatively high $S_N$ at a given total luminosity lie within a projected radius of $\\sim1$\\,Mpc from M\\,87, the central giant elliptical in the Virgo galaxy cluster. The study of \\cite{Miller&Lotz07} finds a consistent trend of increasing GC mass fraction with decreasing host galaxy mass for Virgo dwarf elliptical galaxies. Using the \\cite{McLaughlin99} correction for ${\\cal M}^{\\rm init}_{\\rm gas}/{\\cal M}^{\\rm init}_{\\rm star}$, \\citeauthor{Miller&Lotz07} were able to match the trend seen in their observations.  Investigating the increasing $S_N$ with decreasing galaxy mass below ${\\cal M}_\\star\\!<\\!3\\times10^{10}{\\cal M}_\\odot$, \\cite{Forbes05} used the feedback models of \\cite{Dekel&Woo03} which predict ${\\cal M}/L\\propto {\\cal M}^{-2/3}$, and found qualitative agreement with the observations. However, the normalization of this relation remained unconstrained, so that actual GC formation efficiency values $\\epsilon_h$ could not be determined. This quantitative derivation of GC system scaling relations as a function of galaxy mass is one of the goals of this work by including predictions from the latest \\cite{Dekel&Birnboim06} models which include shock-heating regulated star formation for ${\\cal M}_\\star>3\\times10^{10}{\\cal M}_\\odot$. In Section\\,\\ref{Sect:Data} we briefly introduce the galaxy sample used in this study and access contamination. The analysis of the fractions of GCs in late-type dwarf galaxies along with complementing data from the literature is presented in Section\\,\\ref{Sect:Analysis}. In Section\\,\\ref{Sect:Discussion} we discuss the observed trends of the various GCS scaling relations, the specific GC formation efficiency $\\eta$, and implications for galaxy formation scenarios. Combined with dynamical mass measurements of massive and dwarf galaxies from the literature we normalize the models (Sect.\\,\\ref{Sect:GCscalingNorm}) and find a good description of the observed behavior of the frequencies of GCs as a function of galaxy luminosity with a common value of the GC formation efficiency parameter. Using the derived analytical expressions describing the behavior of the specific GC system scaling parameters as a function of galaxy luminosity we discuss (in Sect.\\,\\ref{Sect:impl.4_h.g.f.}) the predictions of a simplistic satellite galaxy accretion model as a function of galaxy luminosity. Our final conclusions are summarized in Sect.\\,\\ref{Sect.Conclusions}. ",
        "conclusions": "We investigate the much debated behavior of the observed GCS scaling parameters as a function of galaxy luminosity, such as the GC specific frequency ($S_N$), specific luminosity ($S_L$), specific mass ($S_M$), and specific number ($\\hat{T}$). Those are the integrated number, luminosity mass and specific number of all globular clusters in a galaxy normalized to the total galaxy luminosity and/or mass, respectively (see Sect.\\,\\ref{Sect:Analysis}). We derive these quantities from {\\it HST/ACS} data of low-mass, faint ($M_V\\!>\\!-16$ mag) dwarf galaxies, mainly late-type irregulars, located in the halo regions of nearby ($\\!<\\!10$\\,Mpc) galaxy groups and in the field \\citep{Georgiev08, Georgiev09}. In order to investigate the scaling relations of their GCSs as a function of galaxy luminosity (mass) we also compiled data from the literature for massive cluster ellipticals \\citep{Peng08} and spiral galaxies \\citep{Spitler08}. To complement our data in the low-mass regime we included 69 cluster dEs  \\citep{Miller&Lotz07} and 12 dIrrs from \\cite{Seth04} and \\cite{Georgiev06}, as well as 24 late-type dwarfs in the local low-density environment from \\cite{Sharina05}. Thus, we cover virtually the entire range in galaxy luminosity from $M_V\\!=\\!-11$ to $-23$ mag ($10^6\\!-\\!10^{11}L_\\odot$) where old GCs ($t\\!\\ga\\!4$ Gyr) are reliably detected. Our main results can be summarized as follows:\\\\  \\noindent$\\bullet$ {\\it Trends in GCS scaling relations hold irrespective of galaxy morphology.} --- The significantly increased number of galaxies in our analysis in comparison with earlier studies over the entire mass and galaxy morphology range allowed us to firmly corroborate the previous observational findings that the GCS scaling parameters vary as a function of galaxy luminosity (Fig.~\\ref{Fig:SNLMT_compare}). We find that this relation holds irrespective of galaxy morphological type, suggesting a universal mode of GC formation. Galaxies show increasing GCS scaling parameters toward low and high-luminosity systems with a minimum at around $M_V\\!\\approx\\!-20.5$\\,mag ($L_V\\approx2\\times10^{10}L_\\odot$).\\\\  \\noindent$\\bullet$ {\\it $S_L$ values of early-type galaxies are twice that of late-types at a given galaxy luminosity.} --- The specific luminosity $S_L$ is the most robust scaling parameter and shows least scatter due to its independence of distance and weak sensitivity to completeness corrections at the faint end slope of the GC luminosity function (Sect.\\,\\ref{Sect:SL}). For late-type galaxies, spirals and irregulars the $S_L$ value is on average two times smaller at a given galaxy luminosity than that for early-type systems (cf. Fig.\\,\\ref{Fig.SNcompare2}). This difference can be partially accounted for by the passive evolutionary fading of the integrated galaxy light ($\\Delta V\\!\\gtrsim\\!1$\\,mag), provided no more GCs are formed or destroyed at the same time (cf. arrow in Fig.\\,\\ref{Fig.SNcompare2}). As cluster dSphs are on average fainter and exhibit higher scaling parameters than field dIrrs, the above analysis supports the idea that dIrr galaxies could be the progenitors of dSphs and fainter dEs \\cite[e.g.][]{Miller98}, but see also \\cite{Grebel03}. However, our sample contains mostly dIrrs  in field and group environments. Their GC scaling parameters might differ in cluster environments due to varying GC formation efficiencies and/or cluster stripping. However, such cluster dIrrs have not been sampled well enough yet.\\\\  \\noindent$\\bullet$ {\\it The \"U-shaped\" behavior of the GC scaling parameters is remarkably well described by a two component model at a transitional halo mass critical for the thermal properties of the inflowing gas.} --- In order to explain the trends in the GCS scaling parameters we have assumed that the total mass in GCs is proportional to the total halo mass of the host galaxy (${\\cal M}_{\\rm GCS}=\\eta{\\cal M}_{\\rm h}$), i.e. the formation of GCs scales with the total galaxy mass. The coefficient, $\\eta$, is the {\\it specific GC formation efficiency} parameter, which measures the present-day mass/luminosity of the GCS drawn from the initial star cluster population mass/luminosity function shaped by various mechanisms over a Hubble time (cluster relaxation, tidal and dynamical friction, dissolution). We have invoked theoretical models of \\cite{Dekel&Birnboim06} which predict the dependence of the galaxy ${\\cal M}/L$ as a function of galaxy mass determined by the thermal properties of the inflowing gas for two different halo mass regimes in which star formation processes are regulated by supernova feedback and virial shocks below and above the critical stellar mass of $3\\!\\times\\!10^{10}{\\cal M}_\\odot$. For low-mass halos these models predict that the mass-to-light ratio varies with host galaxy mass as ${\\cal M}_{h}/L\\propto{\\cal M}_{h}^{-2/3}$ \\citep[see also][]{Dekel&Silk86}. To draw quantitative conclusions, the absolute normalization of this relation is very important which we derive from recent observations of the mass in the inner 0.3\\,kpc (${\\cal M}_{0.3}=10^7{\\cal M}_\\odot$) of nearby low-mass dwarf galaxies \\citep{Strigari08} using its relation to the total halo mass from the models of \\cite{Bullock01}. Above the critical mass, at large halo masses the models predict ${\\cal M}_{h}/L\\propto{\\cal M}_{h}^{1/2}$ which we calibrate with dynamical mass measurements based on galactic group dynamics of massive early-type galaxies \\citep{Eke06}. This results in an analytical model which describes remarkably well the observed trends in the GC scaling relations as a function of galaxy luminosity (Sect\\,\\ref{Sect.ContinuousRelations}). In the \\cite{Dekel&Birnboim06} model, compact and dense molecular clouds (the sites of GC formation) are efficiently shielded against shock heating and ionizing stellar feedback above and below the critical galaxy mass, respectively. This implies that in this picture, the shape of the GC scaling relations can be described as a universal GC formation efficiency for the entire galaxy mass range but with an evolving field star-formation efficiency. The latter is equally low for low- and high-mass systems (compared to the formation of massive/dense star clusters) due to being efficiently suppressed by the stellar and SNe ionizing feedback and shock heating of the inflowing gas below and above the critical galaxy mass ($\\sim\\!10^{10.5}{\\cal M}_\\odot$), respectively. This causes an evolution of the mass in clusters to mass in field stars ratio as a function of galaxy mass. The evolution of this ratio is at the roots of the shape of the GC scaling relations. The best fit of this model to the observed distributions yielded a value of the observed GC formation efficiency parameter ($\\eta$) for the luminosity and mass-normalized GC formation efficiency of $\\langle\\eta_L\\rangle \\simeq5.5\\times10^{-5}$ and $\\langle\\eta_M\\rangle \\simeq6.47\\times10^{-5}$ (see Sect.\\,\\ref{Sect:GCefficiency}). The $\\eta$ distributions are very similar for all galaxy types, if passive evolutionary fading of the late-type galaxy sample of $\\sim\\!1$\\,mag is applied (cf. Fig.\\,\\ref{Fig:eta-hist}).\\\\  \\noindent$\\bullet$ {\\it The differences between model predictions and observations can be attributed to a mixture of effects of varying GC formation efficiencies, galaxy merging histories, and a variation in cluster destruction mechanisms as a function of galaxy mass lacking in the current model.} --- The most massive galaxies, whose $S_M$ and $S_L$ values increase more rapidly than expected from the theoretical predictions for a fixed $\\eta$ can be understood as systems which have either undergone an early episode of extremely efficient star-cluster formation or less efficient formation of field stars or which have preferentially accreted high-$\\eta$ dwarf galaxies. In a simple merging picture of satellite galaxies of a given luminosity we showed in Sect.\\,\\ref{Sect:impl.4_h.g.f.} that it is possible to boost the $S_L$ values of galaxies by accretion of intermediate luminosity, high-$\\eta$ dwarfs, thus offering an efficient mechanism for explaining the high $\\eta$ values of cD galaxies. In addition, the observed spread in $\\eta$ may be caused by the stochastic nature of star and star-cluster formation at low galaxy mass or by differences in SFH and SFR intensity for galaxies at the same luminosity. \\\\  To understand whether the difference in the GCS scaling-parameter ($S_L$ or $S_M$) distributions between late- and early-type galaxies is statistically significant and, if so, to understand the nature of this difference (e.g.~age, chemical enrichment, environment, etc.) it is crucial to sample with more observations (e.g. near-UV/IR photometry or spectroscopy) the GCSs of spiral galaxies at intermediate to high masses, late-type dwarfs in dense cluster environments, both of which are significantly under represented in the current study."
    },
    "1004/1004.1383_arXiv.txt": {
        "abstract": "The published Mount Wilson Doppler-shift measurements of the solar velocity field taken in 1967--1982 are revisited with a more accurate model, which includes two terms representing the meridional flow and three terms corresponding to the convective limb shift. Integration of the recomputed data over the visible hemisphere reveals significant variability of the net radial velocity at characteristic time scales of $0.1$--$10$ years, with a standard deviation of 1.4 \\ms. This result is supported by independent published observations. The implications for exoplanet detection include reduced sensitivity of the Doppler method to Earth-like planets in the habitable zone, and an elevated probability of false detections at periods of a few to several years. ",
        "introduction": "The spectacular success of the exoplanet search program, resulting in the discovery of over 300 planets and planetary system to date, has been achieved mostly through the Doppler-shift technique. Indeed, spectroscopic detections constitute the majority of known systems \\citep{but, udr}, with the bias toward short-period, massive planets probably due to the selection effects of this method. Undoubtedly, there is much room for further progress with the Doppler technique \\citep{egg}, with the accuracy of spectroscopic instruments steadily improving and now reaching $\\sim 1$ \\ms\\ and the sensitivity of telescopes extending toward fainter stars. Combining precision photometric observations with radial velocity measurements provides a range of important physical characteristics of transiting stars invaluable for our understanding the physics and the origin of exoplanets. The strategic goal of detecting rocky, habitable planets outside the Solar system now seems to be coming within reach. As the instrumental precision steadily improved, a growing attention has been paid to the intrinsic perturbations in the observable parameters used in exoplanet detection. Stochastic, uncorrelated physical perturbations increase the level of noise, making it difficult to achieve the threshold signal-to-noise ratio, while possible cyclic processes in the host stars can mimic exoplanet signatures. In particular, the rotating pattern of photospheric spots  \\citep{saa} and irregularities in the convective structure on the surface \\citep{meu} can lead to an intrinsic scatter of radial velocities of up to a few \\ms\\ for solar type G dwarfs. In the younger Hyades, which are more magnetically active and rotate faster than the Sun, this effect is magnified to $\\sim 16$ \\ms\\ in standard deviation, as observed by \\citet{pau}. As the level of activity is not constant in solar-type stars, the intrinsic radial velocity scatter is correlated with the magnetic cycle, opening possibilities of more sophisticated spectroscopic analysis in order to mitigate these difficulties. For example, the index of chromospheric activity is correlated with the area of star spots (and hence, with the observed scatter) and can serve as an indicator of magnetically induced cycles \\citep{san00}. Careful selection of target stars can further improve the prospect of detection of smaller planets with the Doppler technique. Recent simultaneous measurements of chromospheric activity and radial velocity imply that K-type dwarfs may be significantly less variable than the Sun, and the intrinsic RV jitter is less than 1 \\ms\\ \\citep{san10}. Some stars older than 6 Gyr and evolved off the main sequence have sharply reduced levels of magnetic activity \\citep{wri} compared to the Sun. In this paper, we consider another important source of intrinsic radial velocity variation, related to the physical motion of the surface layers of stars, which remained largely outside the scope of previous papers.  The surface layers of the Sun, where the spectroscopic lines are formed, are known to be involved in a complex pattern of radial and tangential motion. Furthermore, it is now an observational fact that this velocity field is not static. In this paper, the series of Doppler measurements taken at the Mount Wilson Observatory is revisited, and the published fitting model coefficients are transformed to a different model of differential rotation, convective blueshift and meridional flow, which may more adequately represent the reality (\\S~\\ref{data.sec}). The transformed model is integrated to produce the net radial velocity of the Sun as a star (\\S~\\ref{rv.sec}). The results are compared to other data on the variability of the main components of the velocity field in \\S~\\ref{comp.sec}, and a good agreement is found. The impact of the intrinsic variation of solar RV on the detectability of exoplanets, especially within the habitable zone, are investigated in \\S~\\ref{impact.sec} by means of $\\chi^2$ and spectral density periodograms, as well as by a planet detection experiment, which results in two bogus planets. The relative importance of surface flows compared with other sources of intrinsic RV perturbations are discussed in \\S~\\ref{dis.sec}. ",
        "conclusions": "\\label{dis.sec} Variable surface flows are only one of the known physical processes on the Sun that can change the integrated radial velocity. They appear to be long-term in nature, and therefore, may be the main obstacle to detecting Earth-like habitable planets with ultra-precise Doppler measurements. The main pulsation modes of solar-type stars ($p$-modes) have periods of 10--20 min, and can be successfully suppressed in observations by taking longer exposures. The supergranulation pattern, discussed in \\S~\\ref{comp.sec}, should produce variations of up to a few \\ms\\, but their time scale is several hours. Rotation of stars and the non-uniform distribution of surface brightness due to photospheric spots and plages is probably the main source of RV perturbations on the time scales up to 50 d. The impact of spots has been investigated in numerous papers, and the latest estimates suggest relatively small, but non-negligible, dispersion for the Sun of $\\sigma_{\\rm RV} \\simeq 0.4$ \\ms\\ \\citep{mak}. As was shown by \\citet{pau} for the Hyades, the photometric and the RV jitter of more active stars than the Sun are strongly correlated. Extrapolating the empirical relation found in that paper, one would expect a jitter of $\\sigma_{\\rm RV} = 1.7$ \\ms\\ for the Sun. The large discrepancy between these estimates is probably related to the different morphology of photospheric inhomogeneities in the Hyades stars. Dark spots are the dominating magnetic features on active and young stars, whereas the contribution of bright plages on older, solar-type stars tends to match the impact of spots, or even to prevail \\citep{loc}. Besides, active stars often have only one or two giant long-lived spots on the surface, resulting in a strong rotational modulation of the light curve, whereas the surface of an active sun is usually marked with a few spot groups at a time, fairly uniformly distributed in longitude. \\citet{meu} estimated the combined impact of sunspots and plages at $\\simeq 1$ \\ms\\ in RV amplitude, and concluded that it should be confined to periods less than 100 d. Our study suggests that on the time scales $0.6$--$1.4$ yr, characteristic of orbits within the habitable zones, slowly evolving surface flows and the distribution of the convective blueshift is the major source of confusion and error.  Our conclusions are in agreement with the observations of solar radial velocity by \\citet{mkm}. These authors observed the solar light reflected from the Moon for 5 years nightly in violet absorption lines and determined an upper limit of 4 \\ms\\ for the overall dispersion of radial velocities. The periodogram of RV variations (their Fig. 4) shows four peaks rising above the $3\\sigma$ threshold, all in the long-period part of the spectrum. Three of them may be of instrumental origin, but the unresolved peak at period longer than 3 yr appears to be genuine and is called \"intriguing\" in the paper. The two bogus planets found in the reconstructed RV data in this paper (\\S~\\ref{impact.sec}) have periods $6.35$ and $9.35$ yr. Therefore, it is possible that \\citet{mkm} already detected long-term variations in the RV of the Sun as a star. \\citet{meu} pointed out that the effects of convective blueshift, which is included in our RV model (\\S~\\ref{data.sec}), may be smaller in the deep violet lines than in the redder lines that are normally used for planet search.  Our final note is that the model of meridional flow depicted in Fig.~\\ref{sunV.fig} is a simplified representation of the actual fine structure of the velocity field. By necessity, the model for the Mount Wilson data is coarse, where the contribution of convective blueshift and meridional flow can not be decoupled. The former is especially problematic, because it is wavelength-dependent, and uncalibrated changes in the spectrograph setup can bring about long-term systematic errors \\citep[cf.][]{ulr}. The low-order polynomials used to fit the surface velocity field (\\S~\\ref{data.sec}) only reveal the underlying, very much smoothed structure of the actual distribution of tangential and radial motion, characterized by smaller spatial scales and larger amplitudes. For example, the local flows are known to converge toward active regions, reaching 50--100 \\ms\\ \\citep{giz4}, and these local inflows may be responsible for the temporal variation of the overall velocity field on the time scales of supergranulation (8 hours) and the solar cycle (11 years). Thus, the temporal behavior of the velocity field may be in part the average of many stochastic components, and in part the evolution of the global magnetic field \\citep{ulb}, both still poorly understood for the Sun. Therefore, it is doubtful that a good diagnostics can be devised for other stars to separate the physical motion of the surface from the signatures of small planets."
    },
    "1004/1004.3386_arXiv.txt": {
        "abstract": "{ We present FEROS high-resolution (R$\\sim$ 45000) optical spectroscopy of 34 Herbig Ae/Be star candidates with previously unknown or poorly constrained spectral types. A total of 32 sources are from the Th\\'{e} et al. (1994) catalog and two are new nearby Herbig Ae/Be star candidates from Vieira et al. (2003). Within the sample, 16 sources are positionally coincident with nearby (d$<$250 pc) star-forming regions  (SFRs). All the candidates have reported infrared excess. We determine the spectral type and luminosity class of the sources, derive their radial and projected rotational velocities, and constrain their distances employing spectroscopic parallaxes and photometry from the literature. We confirm 13 sources as Herbig Ae/Be stars and find one classical T Tauri star. Three sources are emission line early-type giants (B, A, and F stars with luminosity class III) and may be Herbig Ae/Be stars. One source is a main-sequence A-type star. Fourteen sources are post-main-sequence giant and supergiant stars (7 with H$\\alpha$ emission and 7 without). Two sources are extreme emission-line stars and no accurate spectral classification was possible because of strong veiling. Most of the sources appear to be background stars at distances over 700 pc. We show that high-resolution optical spectroscopy is a crucial tool for distinguishing young stars (in particular Herbig Be stars) from post-main sequence stars in samples taken from emission-line star catalogs based on low-resolution spectroscopy. Within the sample, three young stars (CD-38 4380, Hen~3-1145, and HD~145718) and one early-type luminosity class III giant with emission lines (Hen 3-416) are at distances closer than 300~pc and are positionally coincident with a nearby SFR. These 4 sources are likely to be nearby young stars and are interesting for follow-up observations at high-angular resolution. Furthermore, seven confirmed Herbig Ae/Be stars at $d>700$~pc (Hen~2-80, Hen~3--1121~N\\&S, HD 313571, MWC~953,  WRAY~15-1435, and Th~17-35) are inside or close ($<5'$) to regions with extended 8 $\\mu$m continuum emission and in their 20' vicinity have astronomical sources characteristic of SFRs (e.g., HII regions, molecular clouds, dark nebulae, masers, young stellar-objects). These 7 sources are likely to be members of SFRs. These regions are attractive for future studies of their stellar content. ",
        "introduction": "Herbig Ae/Be stars are intermediate-mass (2-8 M$_\\odot$) pre-main sequence (PMS) stars. They exhibit emission lines (e.g., H$\\alpha$, H$\\beta$, \\ion{Ca}{ii}) in their optical spectra and infrared (IR) excess in their spectral energy distributions. These observational characteristics provide indirect evidence that Herbig Ae/Be stars have an accreting circumstellar disk. The infrared excess is interpreted as emission from small dust grains present in the hot surface layer of the disk. By analogy  with the lower-mass T Tauri stars (e.g., Hartmann 1999, Muzerolle et al. 2004), the \\ion{H}{i} emission lines can be interpreted as originating in the magnetospheric accretion shock\\footnote{The origin of the \\ion{H}{i} emission lines in Herbig Ae/Be stars is in fact controversial. Several authors have alternative scenarios for magnetospheric accretion (e.g., strong stellar winds, outflows, direct disk accretion) to explain the origin of the \\ion{H}{i} lines (e.g., B\\\"ohm \\& Catala 1993,  Mottram et al. 2007). The magnetospheric accretion model can explain the H$\\alpha$ line and its spectropolarimetry signal in Herbig Ae stars (e.g., Pontefract et al. 2000, Vink et al. 2002, 2005). However, this is less clear in the case of Herbig Be stars, since wind or outflows contributions to the H$\\alpha$ line are likely, and the spectropolarimetry signal does not unambiguously support the magnetospheric accretion scenario (Mottram et al. 2007). As several stars in our sample are Herbig Be stars, it could well be that different emission mechanisms work for different stars in our sample.} when the gas of the disk reaches the surface of the star with free-fall velocities of a few hundred km/s. More recently, spatially resolved dust and molecular line observations in the millimeter and sub-millimeter domain (e.g., Mannings and Sargent 1997, Semenov et al. 2005), together with scattered light coronographic imaging (e.g., Fukagawa et al. 2004, Grady et al. 2005) provided direct evidence that Herbig Ae/Be stars are effectively surrounded by a disk (for a detailed review about Herbig Ae/Be stars see Waters and Waelkens 1998).  From the observational point of view, bright nearby (d$<$250 pc)  Herbig Ae/Be stars  (and CTTS) are particularly relevant, because they permit detailed studies of the structure of their disks. Disks are interesting because they play a key role in early stellar evolution and are the sites of planet formation. In nearby sources, the disk can be spatially resolved with 8 -- 10 m class telescopes and infrared and mm interferometers. In bright sources, high--resolution spectroscopy in the near and mid-IR can be obtained to study the gas in the disk (see reviews by Najita et al. 2007 and Carmona 2010). Since the amount of identified nearby PMS with spatially resolved disks is still relatively small, the identification of bright nearby Herbig Ae/Be stars is an important step for future observational studies of protoplanetary disks.  \\begin{table*} \\begin{minipage}[t]{\\textwidth} \\caption{Studied stars, SFRs positionally coincident and summary of the observations.} \\label{table:1} \\centering \\scriptsize \\renewcommand{\\footnoterule}{} \\begin{tabular}{l c c c c c c c c c c } \\hline\\hline &  $\\alpha$ (J2000.0)  & $\\delta$ (J2000.0)  & l & b    &  & d$_{\\rm SFR}$ &  t$_{exp}$ & &  Date(s)\\\\ Star\\footnote{All sources from Table IVb of Th\\'e et al. (1994), except for Hen 3-1145 (Table II), Hen 2-80 (Table IVa) and CD-38 4380 and HD145718 (Vieira et al. 2003).} & [~h~~m~~~ s~ ] & [~$^\\circ$ ~~'~~~ ''~ ]& [deg] & [deg] & SFR & [pc] & [s] & S/N\\footnote{The S/N is the average S/N in the continuum close to the H$\\alpha$ line.} &  [yyyy-mm-dd]  \\\\ \\hline CD-38 4380       & 08 23 11.86 & -39 07 01.5\t  &257.32\t& -1.06 &Gum Nebula                             & 200 \\-- 240 & 1200 & 25 & 2004-04-03 \\\\ WRAY 15-488    & 10 01 48.11 & -59 12 12.5    & 282.70 & -3.17  & ScoOB2\\--5                             & 145            & ~800 & 25 & 2004-04-04 \\\\ WRAY 15-522    & 10 12 12.42 & -62 32 33.1    & 285.69 & -5.13  & ScoOB2\\--5                             & 145            & 1500 & 25 & 2004-04-05 \\\\ Th 35-41             & 10 25 40.07 & -58 22 17.4    & 284.78 & -0.73  & ScoOB2\\--5                             & 145            & 1800 & 15 & 2004-04-08 \\\\ Hen 3-416          & 10 25 44.51 & -58 33 52.2    & 284.89 & -0.89  & ScoOB2\\--5                             & 145             & 1200 & 30 & 2004-04-07 \\\\ WRAY 15-566    & 10 25 51.36 & -60 53 13.2    & 286.13 & -2.86  & Low.Cen.Crux \\-- ScoOB2\\--5 & 118 \\-- 145  & 1800  & 10 & 2004-04-07 \\\\ HD 305773       & 10 56 03.88 & -60 29 37.6    & 289.18 & -0.74  & Low.Cen.Crux \\-- ScoOB2\\--5 & 118 \\-- 145  & ~700 & 100 &2008-05-21 \\\\ WRAY 15-770    & 11 11 28.49 & -63 00 23.7    & 291.87 & - 2.30 & Low.Cen.Crux \\-- ScoOB2\\--5  & 118 \\-- 145  & 1200 & 10 & 2004-04-05 \\\\ Hen 2-80            & 12 22 23.18 & -63 17 16.8    & 299.67 &  -0.60 & Low.Cen.Crux \\-- ScoOB2\\--4  & 118 \\-- 145  & 1800 & 25 & 2004-04-08 \\\\ Hen 3-823          & 12 48 42.39 & -59 54 35.0    & 302.59 &  2.96 & Low.Cen.Crux \\-- ScoOB2\\--4  & 118 \\-- 145  & ~900 & 60 & 2008-04-21 \\\\ Th 17-35             & 13 20 03.59 & -62 23 54.0    & 306.24 &  0.29 & Low.Cen.Crux \\-- ScoOB2\\--4  & 118 \\-- 145 & 1800 & 25 & 2004-04-08 \\\\ WRAY 15-1104   & 13 29 51.02 & -56 06 53.7   & 308.30 &  6.36 & Low.Cen.Crux \\-- ScoOB2\\--4  & 118 \\-- 145 & 1500 & 40 & 2004-04-06 \\\\ WRAY 15-1372 % & 15 53 50.59 & -51 43 05.1    & 329.21   &1.59  & Up.Cen.Lup                            & 140              & 1500 & 30 & 2004-04-08 \\\\ Hen 3-1121N    & 15 58 09.62 & -53 51 18.3 & 328.345 & -0.463 & ... & - & 1500 & 70 & 2008-04-21 \\\\ Hen 3-1121S    & 15 58 09.67 & -53 51 34.9 & 328.342 & -0.466 & ... & - & 1500 & 85 & 2008-04-24 \\\\ Hen 3-1145     & 16 08 54.69 & -39 37 43.1    & 339.21 & 8.95    & Up.Cen.Lup   & 140 & 2000 & 16& 2004-04-08 \\\\ WRAY 15-1435   & 16 13 06.68 & -50 23 20.0 & 332.36  &  0.58  & ... & - & 1800 & 55 & 2008-04-24 \\\\ HD 145718      & 16 13 11.59 & -22 29 06.6    & 352.43 & 20.44    & $\\rho$Oph \\-- Sco OB2 & 110 \\-- 160& 1000 & 20 & 2004-04-08 \\\\ HD 152291       & 16 54 24.20 & -40 39 09.0 & 344.40  &  1.88  & Up.Cen.Lup   & 140 & ~600 &  100& 2008-04-21 \\\\ Hen 3-1347     & 17 10 24.15 & -18 49 00.7    & ~~~4.10 & 12.26   & ... & - & 1800 & 70 & 2007-08-04 \\\\ WRAY 15-1651   & 17 14 45.03 & -36 18 38.4  & 350.27\t&  1.35 & ... & - & 2400 & 17 & 2008-05-01 \\\\ WRAY 15-1650   & 17 15 32.79 & -55 54 22.7    & 334.28  &\t-10.07  & ... & - & 900  & 8& 2007-08-17 \\\\ HD 323154        & 17 23 02.36 &\t-39 03 52.5    & 348.96  & -1.57   & ... & - & 600  & 120 &  2007-08-04 \\\\ WRAY 15-1702   & 17 24 30.88 & -37 34 27.7  & 350.35\t& -0.97 & ... & - & 1700 &8 &  2008-04-30 \\\\ MWC 878        & 17 24 44.70 &\t-38 43 51.4\t    & 349.42  &\t-1.65   & ... & - & 900 & 70 & 2007-08-04 \\\\ AS 231% & 17 30 21.66 &\t-33 45 29.6\t    & 354.18  & 0.17\t& ... & - & 1400 & 140 & 2004-07-31 \\\\ Hen 3-1428      & 17 35 02.49 & -49 26 26.4    & 341.41 &~-9.03& ... & - & 1200 & 50 & 2007-08-07 \\\\ HD 320156     & 17 37 58.51 &\t-35 23 04.3     & 353.66  & -2.03 & ... & - & 600 & 90 & 2007-08-07 \\\\ MWC 593        & 17 49 10.16 &\t-24 14 21.2\t    & ~~~4.43 & ~~1.75 & ... & - & 720 & 85 & 2007-07-31 \\\\ HD 313571        & 18 01 07.18 & -22 15 04.0\t\t& ~~~7.53 &\t~~0.39 & ... & - & 900 & 70 & 2007-08-02 \\\\ MWC 930        & 18 26 25.24 & -07 13 17.8\t\t& ~23.65  &\t~~2.23 & ... & - & 2100 & 50 & 2007-08-04 \\\\ MWC 953        & 18 43 28.43 & -03 46 16.9    \t& ~28.69  & ~~0.05 & ... & - & 1200 & 60 & 2007-07-31 \\\\ AS 321         & 18 47 04.80 &\t-11 41 02.3\t    & ~22.02  & ~-4.35 & ... & - & 1500 & 75 & 2004--07-31\\\\ MWC 314        & 19 21 33.97 & +14 52 57.0\t    & ~49.57  & ~~0.25 & ... & - & 900 & 65 & 2007-08-04 \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*}  Herbig Ae/Be stars were initially identified based on the presence of emission lines (i.e. H$\\alpha$) in their optical spectra and their physical association with a dark cloud or nebulosity (e.g., Herbig, 1960). Herbig (1960) used the last condition to exclude the post-main sequence B[e] stars (i.e. giant or supergiant B-type stars with emission lines). Thanks to the advent of IR space observatories such as IRAS, ISO, and Spitzer this last criterion has been relaxed and replaced by the presence of  near- or far-IR excess, in addition to the emission lines, as membership criteria to the Herbig Ae/Be stellar group (e.g., Finkenzeller \\& Mundt 1984; Th\\'e et al. 1994; Vieira et al. 2003).  However, published samples of Herbig Ae/Be stars may well be contaminated with other classes of objects. One should bear in mind three aspects of the identification: ({\\it i\\,}) in general, Herbig Ae/Be star candidates have been identified in surveys for emission-line stars based on low-resolution data, in particular slit-less spectra (this makes no difference on the detection of the emission lines, but it matters for the determination of the luminosity class, see below); ({\\it ii\\,})  post-main sequence B[e] supergiants can also have IR excess (e.g., Miroshnichenko et al. 2005 and references therein); and ({\\it iii\\,}) as beam sizes for infrared observations employed in previous studies have typically been large (e.g., 30\" in the case of IRAS), confusion with other infrared sources may have occurred.  Since hydrogen lines are observed in emission in Herbig Ae/Be stars, the hydrogen lines width, the usual means for determining the luminosity class, cannot be used. Thus the observation of gravity sensitive lines (e.g., \\ion{N}{ii} at 3995~\\AA, \\ion{Si}{ii} at 4128 and 4131~\\AA,  \\ion{C}{ii} at 4267~\\AA, \\ion{Si}{iii}  at 4553 and 4561~\\AA, and \\ion{O}{ii} lines at 4070 and 4976~\\AA) is required. However, these lines are relatively weak and are barely visible in low-resolution spectra. Therefore, low spectral resolution studies that have identified Herbig Ae/Be candidates have the important limitation that background B[e] supergiants can be mistakenly classified as Herbig Ae/Be stars. Consequently, to confirm that a Herbig Ae/Be candidate is indeed a young star - and not a post-main sequence object - observations at high spectral resolution are necessary.  In this paper, we present the results of a high-resolution optical spectroscopy campaign aimed at identifying and characterizing new nearby Herbig Ae/Be stars. We obtained FEROS high-resolution (R$\\sim$45000) optical spectra of 34 candidates to Herbig Ae/Be  stars. We studied sources positionally coincident with nearby (d$<$250 pc) star-forming regions (SFRs) and ``isolated\" sources. Our goal was to determine whether the candidates belong to the Herbig Ae/Be stellar group by searching emission lines in their spectra and by determining their spectral type and luminosity class. We then constrained their distances employing spectroscopic parallaxes and derived their radial and projected rotational velocities. In the case of sources positionally coincident with SFRs, we used the estimated distance to determine whether the sources are members of nearby SFRs. Finally, for the confirmed Herbig Ae/Be stars that are not members of nearby SFRs, we searched the Spitzer archive for 8 $\\mu$m imaging and the SIMBAD database to find evidence for extended near-infrared emission and astronomical objects characteristic of SFRs (e.g., HII regions, molecular clouds, dark nebulae, masers, young stellar-objects). Our aim was to find evidence of whether these distant Herbig Ae/Be stars might be members of distant SFRs.  \\begin{figure*} \\centering \\includegraphics[angle=0,width=\\textwidth]{R_K_band_charts.eps} \\caption{Optical (R band) and near-IR (K band) images of the 0.5' $\\times$ 1' field centered on the target stars. The target is indicated by an arrow.} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[angle=0,width=\\textwidth]{R_K_band_chartsB.eps} \\caption{Optical (R band) and near-IR (K band) images of the 0.5' $\\times$ 1' field centered on the target stars. The target is indicated by an arrow.} \\end{figure*} ",
        "conclusions": "We obtained high-resolution optical spectroscopy of 34 candidates to Herbig Ae/Be stars with unknown or poorly constrained spectral types from the Th\\'e et al. (1994) catalog and two candidates from Vieira et al. (2003). We observed 16 candidates positionally coincident with nearby (d$<$250 pc) SFRs and 18 relatively bright ($V < 14$) ``isolated\" candidates. All our candidates have reported IR-excess from IRAS. Our aim was to determine whether the candidates are Herbig Ae/Be stars or background giants, and in the specific case of the candidates positionally coincident with SFRs, we wanted to further find out whether they are members of the SFR. We determined the spectral types of the sources by careful comparison with spectral templates, we measured their radial and projected rotational velocities, finally, we constrained their distances employing spectroscopic parallaxes based on the intrinsic colors of the established spectral type and luminosity class and photometry from the literature.  From the 34 Herbig Ae/Be candidates studied, 26 objects exhibit H$\\alpha$ in emission ($\\sim$80\\%, see Figure 3 and Table 3). From these 26 objects, 14 are dwarfs and subgiants (luminosity classes V and IV), 10 are giants and super giants (luminosity classes III, II, and I), and 2 are unclassified extreme emission line objects. From the 8 objects {\\it without} H$\\alpha$ emission, 7 are giants and one (AS 321) is a main-sequence A-type star.  Among the 14 emission line dwarfs and subgiants, 13 objects are confirmed Herbig Ae/Be stars and one is a CTTS. In addition to these 13 confirmed Herbig Ae/Be stars, 5 additional sources might be Herbig Ae/Be stars: 3 emission-line early type luminosity class III giants, and 2 extreme emission line objects. However, our data did not allowed us to firmly establish whether these 5 sources are truly Herbig Ae/Be stars.  Two confirmed Herbig Ae/Be stars (CD-38 4380, HD~145718) and the CTTS (Hen 3-1145) are at distances closer than 250 pc. These sources are likely members of nearby SFRs. One emission line giant star (Hen 3-416) is at closer than 300 pc. If this source is a young star, it may be associated with Sco OB 2-5. These 4 sources are likely to be nearby young stars and are interesting for  follow-up observations at high-angular resolution. The rest of our confirmed Herbig Ae/Be stars (11 sources) are at distances greater than 700 pc. From this subsample, 7 stars (Hen~2-80, Hen~3--1121~N\\&S, HD 313571, MWC~953,  WRAY~15-1435, and Th~17-35) are inside or close (separation $<5'$) to regions with extended infrared (IR) continuum emission at 8$\\mu$m and have astronomical sources characteristic of SFRs in their 20' vicinity. These 7 sources are likely to be members of distant SFRs. Such regions are attractive for future studies of their stellar content. Two confirmed Herbig Ae/Be stars at $d>700$ pc, MWC 878 and WRAY 15-1372, may be truly ``isolated\" sources.  From our 34 Herbig Ae/Be candidates we found that $\\sim$50\\% (15 of 34) turned out to be background giant stars and not young stars. They show us that high-resolution optical spectroscopy is an important tool for distinguishing young stars (in particular Herbig Be stars) from post-main sequence stars in samples taken from catalogs based on low-resolution spectroscopy. A systematic study of large samples of candidates to Herbig Ae/Be stars employing high-spectral resolution spectroscopy is fundamental for firmly establishing their genuinly young nature."
    },
    "1004/1004.3665_arXiv.txt": {
        "abstract": "{Gamma-ray bursts are cosmological sources emitting radiation from the gamma-rays to the radio band. Substantial observational efforts have been devoted to the study of gamma-ray bursts during the prompt phase, i.e. the initial burst of high-energy radiation, and during the long-lasting afterglows. In spite of many successes in interpreting these phenomena, there are still several open key questions about the fundamental emission processes, their energetics and the environment.} {Independently of specific gamma-ray burst theoretical recipes, spectra in the GeV/TeV range are predicted to be remarkably simple, being satisfactorily modeled with power-laws, and therefore offer a very valuable tool to probe the extragalactic background light distribution. Furthermore, the simple detection of a component at very-high energies, i.e. at $\\sim 100$\\,GeV, would solve the ambiguity about the importance of various possible emission processes, which provide barely distinguishable scenarios at lower energies.} {We used the results of the MAGIC telescope observation of the moderate resdhift ($z\\sim0.76$) \\object{GRB\\,080430} at energies above about 80\\,GeV, to evaluate the perspective for  late-afterglow observations with ground based GeV/TeV telescopes.} {We obtained an upper limit of $F_{\\rm 95\\%\\,CL} = 5.5 \\times 10^{-11}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ for the very-high energy emission of \\object{GRB\\,080430}, which cannot set further constraints on the theoretical scenarios proposed for this object also due to the difficulties in modeling the low-energy afterglow. Nonetheless, our observations show that Cherenkov telescopes have already reached the required sensitivity to detect the GeV/TeV emission of GRBs at moderate redshift ($z \\lesssim 0.8$), provided the observations are carried out at early times, close to the onset of their afterglow phase.} {} ",
        "introduction": "\\object{GRB\\,080430} was detected by the \\textit{Swift} satellite \\citep{Geh04} on April 30, 2008 at 19:53:02\\,UT \\citep{Gui08}. The prompt emission lasted $\\sim 16$\\,s \\citep{Stam08} allowing to assign this event to the long duration class \\citep{Kouv93}. X-ray and optical counterparts were discovered and followed-up by many groups. Optical spectroscopy was rapidly carried out allowing to derive a redshift of $z = 0.758$. The redshift estimate has been revised recently with a more accurate wavelength calibration \\citep[][and de Ugarte Postigo et al. in preparation, hereinafter DEUG10]{deUg08,CuFo08}. The relatively modest redshift made it an interesting target for the Major Atmospheric Gamma-ray Imaging Cherenkov (MAGIC) telescope\\footnote{http://wwwmagic.mpp.mpg.de/} observations. In the past, upper limits for several Gamma-Ray Bursts (GRBs) at energies greater than about 100\\,GeV were reported both for single event observations and for a sample of events \\citep[e.g.][]{Alb06,Tam06,Alb07,Aha09}. In this paper we try to predict the Very-High Energy (VHE) flux for \\object{GRB\\,080430} by modeling the detected X-ray and optical afterglow and adopting as a reference the cosmological fireball model \\citep{Pir99,Zha07}. In Sect.\\,\\ref{sec:magicobs} we report the results of the MAGIC observation, in Sect.\\,\\ref{sec:aft} we discuss the lower energy afterglow, in Sect.\\,\\ref{sec:ssc} we introduce the adopted modeling scenario for the VHE flux, in Sect.\\,\\ref{sec:ebl} we discuss the effect of Extragalactic Background Light (EBL) attenuation and finally, in Sect.\\,\\ref{sec:concl}, conclusions and considerations about future perspectives are drawn. Throughout the paper we assume a $\\Lambda{\\rm CDM}$ cosmology with $\\Omega_{\\rm m} = 0.27$, $\\Omega_\\Lambda = 0.73$ and $h_0 = 0.71$. At the redshift of the GRB the luminosity distance is $\\sim 4.8$\\,Gpc ($\\sim 1.5 \\times 10^{28}$\\,cm, corresponding to a distance modulus $\\mu = 43.4$\\,mag). All errors are $1\\sigma$ unless stated otherwise. Throughout this paper the convention $Q_x = Q/10^x$ has been adopted in CGS units. Results presented in this paper supersede those reported in \\citet{Cov09b}. ",
        "conclusions": "\\label{sec:concl}  The prediction of the expected SSC flux for an afterglow is not straightforward since it is required to know, or at least to reliably estimate, the parameters of the underlying afterglow (see Fig.\\,\\ref{fig:ul}). In the case of \\object{GRB\\,080430} the sampling of the X-ray and optical afterglow allowed us to estimate the various afterglow parameters to derive meaningful predictions for the expected SSC flux. However, a satisfactory modeling of the \\object{GRB\\,080430} can not be obtained within the standard fireball scenario. At least two different components are required for the early-time afterglow, as discussed in detail in DEUG10. Our present discussion is based on the assumption that one of these components is the regular afterglow \\citep[i.e. the forward-shock][]{Pir99,Zha07} which is the main responsible for the late-afterglow emission although other components are likely playing a role.  The results appear to be well below the reported upper limits. Furthermore, our assumed low opacity for the EBL is in agreement with current observations \\citep[see also][]{Gilm09EBL}. At any rate, this pilot case shows fairly interesting perspectives for a late-afterglow detection at high energies.  In general, to increase the flux expected from a GRB afterglow (for SSC) it is mandatory to try to decrease the observation energy (due to the $\\nu^{-p/2}$ dependence above the cooling SSC frequency), which is also very important for the minimization of the EBL attenuation. If the telescope sum trigger hardware upgrade had already been implemented before the observations, a limit above an energy of 45\\,GeV would have been obtained \\citep[see also][]{GRB090102}. At these energies, the strong effect of the EBL could probably be neglected and the low energy threshold together with the expected performances of MAGIC\\,II would undoubtedly increase the chances of positive detections.  As a matter of fact, \\object{GRB\\,080430} was an average event in terms of energetics. More energetic GRBs are indeed relatively common, and due to the positive dependence on the isotropic energy of a GRB, much higher fluxes than in the present case can be foreseen. This is also true if we consider the uncertainty in the present total energy determination, which is based on an average value for the prompt emission efficiency.  The time delay of the observation from the GRB has a clear impact, essentially because the observed SSC component is strictly related to the underlying synchrotron component which rapidly decays in intensity with time, depending on the specific environment and micro-physical parameters. Eq.~\\ref{eq:extrap} goes roughly with $t^{-1.1}$ which means that had MAGIC been able to start observations right at the start of the late afterglow phase (e.g. at $T_0 + 1$ks), the flux predictions would have increased by more than an order of magnitude. The time delay of about two hours, coupled with the poor observing conditions, were more than enough to depress the observed flux and raise the reported upper limits.  Given the uncertainties in the modeling of the afterglow, many possible modifications to the standard afterglow model \\citep{Pir99,Zha07} can be applied. In some scenarios, substantially higher VHE flux can be predicted \\citep[e.g.][]{Pan08,Mur10}, which makes observations at VHE energies powerful potential diagnostic tools.  The case of \\object{GRB\\,080430} in this pilot study demonstrates that if three conditions are met: 1) a moderate redshift ($z \\lesssim 0.8$), 2) start of observations right at the beginning of the afterglow phase or even during the prompt emission and 3) the use of the MAGIC sum trigger enabling reaching energy thresholds below 50\\,GeV, detection is within reach. The recent detection of $\\sim 30$\\,GeV photons during the prompt or afterglow phases of \\object{GRB\\,090510} \\citep{Abd09} and \\object{GRB\\,090902B} \\citep{deP09a,deP09b} by the Fermi satellite \\citep{Band09} indeed shows that, with a threshold energy of a few tens of GeV and with the collecting area of a ground-based Cherenkov telescope, GRB VHE astrophysics is becoming a promising observational field."
    },
    "1004/1004.3515_arXiv.txt": {
        "abstract": "{Recent X-ray observations have proved to be very effective in detecting previously unknown supernova remnant shells around pulsar wind nebulae (PWNe), and in these cases the characteristics of the shell provide further clues on the evolutionary stage of the embedded PWN. However, it is not clear why some PWNe are still ``naked''. } {We carried out an X-ray observational campaign targeted at the PWN \\src, the ``close cousin'' of the Crab, with the aim to detect the associated SNR shell.} {We analyzed an XMM-Newton and Suzaku observations of \\src\\  and we model out the contribution of dust scattering halo.} {We detected an intrinsic faint diffuse X-ray emission surrounding the PWN up to $\\sim 6\\arcmin$ ($\\sim 10$ pc) from the pulsar, characterized by a hard spectrum, which can be modeled either with a power-law ($\\gamma= 2.9$) or with a thermal plasma model ($kT=2.0$ keV.)} {If the shell is thermal, we derive an explosion energy $E=0.5-1.6\\times 10^{51}$ erg, a pre-shock ISM density of 0.2 cm$^{-3}$ and an age of $\\sim 2000$ yr. Using these results in the MHD model of PWN-SNR evolution, we obtain an excellent agreement between the predicted and observed location of the shell and PWN shock.} ",
        "introduction": "One of the most intriguing problems in the field of the Pulsar Wind Nebulae (PWNe) study is the lack of a shell around some of these objects. This is somehow disturbing for the consolidated picture of a remnant of a core-collapse supernova, which indicates that the PWN is expanding inside the host supernova remnant, giving rise to a variety of complex phenomena, like reverberation, Rayleigh-Taylor instability at the interface between the PWN and ejecta, rejuvenating the shell, etc. (e.g. \\citealt{vag01}; \\citealt{bcf01}; \\citealt{rac05}; \\citealt{gsz09} and references therein). One of the reasons could be the lack of deep observations aimed at the PWN surroundings. Indeed, recently a shell-like component has been observed in many objects, such as G21.5--0.9 (\\citealt{bb04}; \\citealt{bvc05}), G0.9+0.1 (\\citealt{pdw03}), 3C58 (\\citealt{bwm01}; \\citealt{ghn07}).  Therefore, X-ray observations are very effective for the discovery of associated shell components, even in the presence of high absorption column densities (G21.5--0.9 has $N_H\\sim2\\E{22}\\U{cm^{-2}}$; G0.9+0.1 even $\\sim10^{23}\\U{cm^{-2}}$).  The objects in which pulsar, plerion, and shell are all detected (collectively known as composite SNRs) are extremely important to set the physical conditions for their modeling.  The properties and the evolution of a PWN are determined by the interaction of the pulsar wind with the ambient medium.  The effectiveness by which the surrounding matter confines the PWN is very important to determine the the level of the synchrotron emission from the nebula. Therefore, measuring density and pressure in the shell component is needed for better constraining the models of the PWN. Moreover, the pulsar, the PWN and the shell would allow us to estimate in independent ways some quantities, such as the actual age of the object, or the internal pressure of the nebula.  This redundancy would allow us also to verify our assumptions, like that on the level of equipartition in the PWN and that on how reliable is the age estimated from the pulsar spin-down properties.  \\src\\  is the Galactic PWN that most closely resembles the Crab Nebula: this is the reason why \\citet{lwa02} have dubbed it ``a close cousin of the Crab Nebula''.  Using Chandra data, \\citet{lwa02} have shown the presence of a well defined torus of $\\sim10''$ in diameter, together with elongations, toward E and W directions, which could be ascribed to X-ray jets.  From those data, the size of the X-ray nebula appears $\\sim1'$, but the outer part of the nebula is very faint, and its edge is poorly defined.  At radio wavelengths, instead, the nebular size is $\\sim1.5'$ (\\citealt{vb88}), corresponding to $\\sim2.7\\,d_{6.2}$~pc where $d_{6.2}$ is the distance of G54.1+0.3 in units of that estimated by \\citet{ltw08}, namely $d_{6.2}=d/6.2^{+1.0}_{-0.6}$ kpc.  It is important to understand to which extent this difference in size is real (i.e.\\ due to synchrotron losses of the emitting electrons), or it is an artifact of the limited X-ray sensitivity.  Radio maps show a rather amorphous structure, but the radio emission from G54.1+0.3 is highly polarized, up to 20--30\\% (\\citealt{vb88}), and this indicates (similarly to the case of the Crab Nebula) that the nebular field is highly ordered.  The X-ray spectrum is a power law with a photon index $-1.9$, an absorption column density $N_H\\sim  1.6\\E{22}\\U{cm^{-2}}$ and an X-ray luminosity $L_X\\sim  1.3\\E{33}d_{6.2}^2\\U{erg\\,s^{-1}}$. \\citet{clb02} have detected the pulsar PSR J1930+1852 at the center of the nebula, which has a period of 136 ms, a characteristic age of 2900 yr and a spin-down luminosity of $1.2\\times 10^{37}$ erg s$^{-1}$.  Recently, \\src\\  has attracted some interest because \\citet{kml08} have identified an IR shell surrounding the PWN at a distance of $\\sim 1.5'$ from the pulsar. The shell contains a dozen of IR compact sources. \\citet{kml08} suggests that the sources are young stellar objects, whose formation has  been triggered by the wind of the progenitor of the SN. This intriguing possibility has been questioned by \\citet{tsr10}, who pointed out that the IR shell may be ejecta dust, rather than a pre-existing ISM dense cloud.  \\citet{ltw08} reported the presence of a molecular cloud partially interacting with the PWN, on the basis of the CO emission around the nebula.  Therefore, even if there is no hint of a radio shell around this PWN, there are a number of evidences for interaction between the PWN and the surroundings, so it is worth searching for an X-ray shell.  In Sect. 2 we present a deep X-ray campaign aimed to this PWN, which led to the detection of such a shell. In Sect. 3 we estimate the contribution of the dust scattering halo, showing that it is negligible at the shell location, and in Sect. 4 we discuss our findings comparing them to a PWN-SNR evolution model. ",
        "conclusions": "We have analyzed an \\xmm\\  and a \\suz\\  observation of the PWN \\src, in the framework of a program aimed to survey the region around this isolated nebula in search for the X-ray shell of the associated supernova remnant. We detected very faint X-ray emission around the PWN, extending form the outskirts of the PWN (at $\\sim 1.5^\\prime$ from the central pulsar) until a radius $\\sim 3.8$ times the PWN radius (i.e. $\\sim 5.7^\\prime$, around 10.3 pc at the distance of the nebula). This extended diffuse emission is more evident toward south and it has an irregular morphology on a angular scale of $\\sim 1^\\prime$. We modeled the X-ray dust scattering halo around \\src, and we have found that the detected faint diffuse emission cannot be due to this effect, but it must be intrinsic to the source.  We modeled the X-ray spectrum of the diffuse emission with a thermal model, finding a best-fit temperature of $\\sim 2$ keV, { which may imply electron heating by the shocked ions}. This value, together with the apparent size and the emission measure of the X-ray emitting plasma, is consistent with a SNR shell expanding into a $\\sim 0.2$ cm$^{-3}$ ISM, whose explosion energy is $\\sim 10^{51}$ erg, { and whose most probable age is 1800-2400 yr, a bit less than the characteristic age of the pulsar PSR J1930+1852}, located at the center of the PWN. However, due to limited counting statistics, the X-ray spectrum of the diffuse emission can be { alternatively well fitted} with a non-thermal power-law model, whose photon index ($\\gamma = 2.9$) is roughly consistent with an interpretation in terms of synchrotron emission from accelerated particles.  The morphology of the large diffuse emission neither seems to be directly linked to the IR shell observed around the PWN by \\citet{kml08}, nor to the molecular cloud detected by \\citet{ltw08} and reported as contours in Fig. \\ref{images}, but the fact that the X-ray shell is incomplete is probably related to the interaction between the PWN and these inhomogeneities of the ISM.  We have compared the PWN and SNR sizes with the prediction of the evolutionary model of \\citet{gsz09} for composite SNRs, and we find an excellent agreement. We conclude that the faint diffuse emission around the PWN \\src\\  may indeed be the shell of the associated remnant. However, { deeper X-ray and radio observations are required to definitely distinguish between thermal and non-thermal interpretation. } Given the recent detections of X-ray shells around other PWNe, our results suggest that the lack of shell around remaining isolated PWNe may be simply the result of high absorption and/or lack of long observations, and that the X-ray band may be very effective in discovering them."
    },
    "1004/1004.2333_arXiv.txt": {
        "abstract": "In a landscape with metastable minima, the bubbles will inevitably nucleate. We show that during the bubbles collide, due to the dramatically oscillating of the field at the collision region, the energy deposited in the bubble walls can be efficiently released by the explosive production of the particles. In this sense, the collision of bubbles is actually highly inelastic. The cosmological implications of this result are discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.0100_arXiv.txt": {
        "abstract": "{In the last years, space missions such as COROT, Kepler or MOST have provided very accurate photometric observational data. In the particular case of $\\delta$ Scuti stars, the observed frequency spectra have hundreds (if not thousands) of modes and a clear amplitude distribution. In this work we present new techniques for modelling these observations and the results obtained. We searched for regular patterns in the observational data, which yields something resembling the large separation. This allows to reduce the possible positions of the star in the HR diagram, yielding a value of the mean density with an accuracy never reached before for isolated stars of this type. Finally, we answer whether a $\\delta$ Scuti star is stable despite all of the observed frequencies are simultaneously excited} ",
        "introduction": "}   The $\\delta$ Scuti stars are intermediate-mass pulsating variables with spectral types ranging from A2 to F0. They are located on and just off the main sequence in the lower part of the Cepheid instability strip (luminosity classes V \\& IV). Nowadays, the $\\delta$ Scuti stars are considered as particularly suitable for asteroseismic studies of poorly known hydrodynamical processes occurring in stellar interiors such as convective overshoot, mixing of chemical elements and redistribution of angular momentum (Zahn 1992), etc. Due to the complexity of the oscillation spectra, their pulsating behaviour is not fully understood, in particular in what regards the rotation-pulsation interaction (see a complete rewiew on such effects in Goupil et al. 2005). In the last decade, numerous interpretation works have taken the effects of rotation into account (Foz Machado et al. 2006; Su\\'arez et al. 2007a; Su\\'arez et al. 2007b; Bruntt et al. 2007)   The very precise space photometry supplied by the CoRoT mission give us the possibility to deal with a range and an amount of frequencies not reached by usual ground based observations. In this work we focus on the star HD\\-174936. This is a field $\\delta$ Scuti star observed during the first short run SRc01 for which we have a frequency resolution corresponding to 27 days of observation, namely 0.45~$\\mu$Hz.  The evolutionary code CESAM (Morel 1997; Morel \\& Lebreton 2008) and the pulsation codes GraCo (Moya et al. 2004; Moya \\& Garrido 2008) have been used as numerical codes to calculate frequencies, growth rates and other physical quantities. GraCo provides, in particular, non-adiabatic variables and growth rates. Comparing the theoretical predictions given by these numerical codes with the observed range of excited frequencies and studying the periodic spacing in the oscillation frequency distribution, we have performed a seismic study in order to constrain the physical parameters of the star.  Finally, we have used a representative model of HD\\-174936 to study if the star can excite, at the same time, the huge number of pulsational modes observed in $\\delta$ Scuti stars from space (Poretti et al. 2009; Moya \\& Rodr\\'iguez-L\\'opez 2010). ",
        "conclusions": "}  We have performed an asteroseismic analysis of the $\\delta$ Scuti star HD\\,174936, observed by CoRoT during its first short run, SRc01. The very precise space photometry provides the possibility of dealing with a significantly large number of frequencies (around 400), for which we demonstrated that the star has enough energy to excite. However, Kallinger \\& Matthews (2010) have recently suggested that most of the lowest amplitude frequencies can be a consequence of surface granulation.  We have combined the classical seismic analysis with the use of statistical properties of the modes. In particular, we have searched for periodic patterns in the frequency spectrum of HD\\,174936 in order to find new observational constraints. We have found a peak distribution in the frequency spectrum that seems to correspond with a large separation about 50~$\\mu$Hz.  We have then performed a theoretical analysis in which models have been constrained to fit the observed large separation and to predict unstable the observed modes. These restrictions yield a range of models with [7801, 8192]~K in $T_{eff}$ and [4.07, 4.19] in $\\log g$, which corresponds to a range of [1.47, 1.82]~$M_{\\odot}$ in mass, [788.5, 1705.9]~Myr in age and [1.61, 2,05]~$R_{\\odot}$ in radius."
    },
    "1004/1004.5349_arXiv.txt": {
        "abstract": "We calculate the most massive object in the Universe, finding it to be a cluster of galaxies with total mass $M_{200}=3.8\\times10^{15}\\,M_{\\odot}$ at $z=0.22$, with the $1\\sigma$ marginalized regions being $3.3\\times10^{15}\\,M_{\\odot}<M_{200}<4.4\\times10^{15}\\,M_{\\odot}$ and $0.12<z<0.36$.  We restrict ourselves to self-gravitating bound objects, and base our results on halo mass functions derived from N-body simulations.  Since we consider the very highest mass objects, the number of candidates is expected to be small, and therefore each candidate can be extensively observed and characterized.  If objects are found with excessively large masses, or insufficient objects are found near the maximum expected mass, this would be a strong indication of the failure of $\\Lambda$CDM. The expected range of the highest masses is very sensitive to redshift, providing an additional evolutionary probe of $\\Lambda$CDM. We find that the three most massive clusters in the recent SPT $178\\,\\mbox{deg}^2$ catalog match predictions, while XMMU J2235.3--2557 is roughly $3\\sigma$ inconsistent with $\\Lambda$CDM. We discuss Abell 2163 and Abell 370 as candidates for the most massive cluster in the Universe, although uncertainties in their masses preclude definitive comparisons with theory. Our findings motivate further observations of the highest mass end of the mass function. Future surveys will explore larger volumes, and the most massive object in the Universe may be identified within the next decade. The mass distribution of the largest objects in the Universe is a potentially powerful test of $\\Lambda$CDM, probing non-Gaussianity and the behavior of gravity on large scales. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.2043_arXiv.txt": {
        "abstract": "Several neutral species ({\\MgI}, {\\SiI}, {\\CaI}, {\\FeI}) have been detected in a weak {\\MgII} absorption line system ($W_r(2796) \\sim 0.15~${\\AA}) at $z \\sim 0.45$ along the sightline toward HE0001-2340.  These observations require extreme physical conditions, as noted in \\citet{D'Odorico07}.  We place further constraints on the properties of this system by running a wide grid of photoionization models, determining that the absorbing cloud that produces the neutral absorption is extremely dense ($\\sim 100-1000$ {\\cc}), cold ($<100$K), and has significant molecular content ($\\sim 72-94 \\%$).  Structures of this size and temperature have been detected in Milky Way CO surveys, and have been predicted in hydrodynamic simulations of turbulent gas.  In order to explain the observed line profiles in all neutral and singly ionized chemical transitions, the lines must suffer from unresolved saturation and/or the absorber must partially cover the broad emission line region of the background quasar.  In addition to this highly unusual cloud, three other ordinary weak {\\MgII} clouds (within densities of $\\sim 0.005$~{\\cc} and temperatures of $\\sim10000$~K) lie within 500~{\\kms} along the same sightline.  We suggest that the ``bare molecular cloud'', which appears to reside outside of a galaxy disk, may have had in situ star formation and may evolve into an ordinary weak {\\MgII} absorbing cloud. ",
        "introduction": "\\label{sec:1}  Observations of {\\MgIIdblt} in intervening quasar absorption line systems are vital to the study of the interstellar medium of low-redshift galaxies and their environments, as {\\MgIIdblt} lies in the optical from $z \\sim 0.3-2.4$, and serves as a probe of low-ionization gas.  Through photoionization models, it is possible to derive many properties of the absorbing gas.  This includes the line-of-sight extent, density, temperature, and molecular content, which are constrained by absorption in the different ionization states of various chemical elements.  \\citet{CWC99b} and \\citet{Rigby} found that roughly one-third of weak {\\MgII} absorbers at redshifts $0.4 < z < 1.4$, with $W_r(2796) < 0.3${\\AA}, are in multiple-cloud systems.  They propose that many of the weak absorbers contain multi-phase structures with a range of levels of ionization.  Such absorbers are generally thought to arise in sub-Lyman limit system environments, with $N({\\HI}) < 10^{17.2}~{\\cmsq}$, and metallicities log$[Z/Z_{\\odot}]> -1$ \\citep{Churchill00}.  Strong {\\MgII} absorbers ($W_r(2796) > 0.3$~{\\AA}), in contrast, are almost always associated with Lyman limit systems ($N({\\HI}) > 10^{17.2}~{\\cmsq}$), and many with damped {\\Lya} absorbers (DLAs; with $N({\\HI}) > 2$x$10^{20}~{\\cmsq}$). \\citet{Rao06} found that $36\\%$ of MgII absorbers with $W_r(2796) > 0.5$~{\\AA} and {\\FeII} $W_r(2600) > 0.5$~{\\AA} were DLAs in an HST survey for $z<1.65$ systems.  In that sample, the average $N({\\HI})$ was $9.7 \\pm 2.2$ x $10^{18}~{\\cmsq}$ for $0.3 < W_r(2796) < 0.6$~{\\AA}, and $3.5 \\pm 0.7$ x $10^{20}~{\\cmsq} $ for $W_r(2796) > 0.6$~{\\AA}.  Most DLAs at low $z$ are thought to be associated with galaxies with a variety of morphological types, from $0.1 L^*$ galaxies to low surface brightness galaxies \\citep{LeBrun97, RT98, Bowen01, RT03}. \\citet{Ledoux03} find molecular hydrogen in 13-20\\% of DLA systems at high redshift, but note that there is no correlation between the detection of molecular hydrogen and {\\HI} column density. Despite this lack of correlation, \\cite{Petitjean06} find molecular hydrogen in 9 out of 18 high metallicity systems ([X/H]$>-1.3$) at high redshift.  In this paper, we study a multiple-cloud weak {\\MgII} system toward HE0001-2340 ($z_{em}=2.28$, \\citealt{Reimers98}) at $z= 0.4524$.  The four weak {\\MgII} clouds are spread over a velocity range of $\\sim 600$~{\\kms}. {\\MnII} and {\\CaII}, which are generally only detected in the very strongest absorbers, are found in one of the four clouds comprising this system. {\\FeI}, {\\SiI}, and {\\CaI} are also detected in this cloud; these neutral states have not been reported to exist in any other extragalactic environment, even in most DLAs, and have only been found in several dense galactic molecular clouds \\citep{Welty03}. \\citet{D'Odorico07} notes that the ratios of ${\\MgI}/{\\MgII}$ and ${\\CaI}/{\\CaII}$ in this system are orders of magnitude higher than in other absorbers, implying a very low ionization state.  She also observes that there is an extreme underabundance of Mg with respect to Fe, which cannot be explained by nucleosynthesis or dust depletion, and cannot be reproduced by photoionization models.  The metallicity of the system cannot be directly determined from {\\Lya} due to the absorption from a Lyman limit system at $z=2.18$ \\citep{Reimers98}.  Due to a lack of metallicity constraints, \\citet{D'Odorico07} assumes DLA-regime column densities, noting that the observed $N({\\MgI})$ of the system is comparable to the sample of 11 DLA by \\citet{DZ06}.  She further constrains her parameters by comparing to the local cold interstellar clouds of \\citet{Welty03}, finding that systems with similar amounts of ${\\FeI}$ have metallicities $-3.78 < [\\rm{Fe/H}]]<-2.78$.  With such a low assumed metallicity and high ${\\HI}$ column density, \\citet{D'Odorico07} is unable to reproduce the observed ratios of $\\frac{N({\\MgI})}{N({\\MgII})}, \\frac{N({\\CaI})}{N({\\CaII})}$ and $\\frac{N({\\FeI})}{N({\\FeII})}$ in photoionization models, and is thus unable to draw concrete conclusions about the properties of this cloud. We propose that these noted differences suggest that this unusual cloud is part of class of systems unrelated to previously observed DLAs, which \\citet{D'Odorico07} notes have $N({\\MgII})$ two orders of magnitude higher than this system, as well as no previous {\\FeI}, {\\CaI}, and {\\SiI} detections.  We therefore explore a range of metallicities and $N({HI})$ in our modeling process.  As it is difficult to reproduce the line ratios in this unusual cloud, we also expand the consideration of parameter space to explore the possibility of unresolved saturation.  The effect of partial covering of the background quasar is also explored, since photoionization models of the system suggest the cloud is compact enough to partially cover the broad emission line region of the quasar, with high densities ($n_H=1-34$ {\\cc}), cold temperatures ($<200$K), and a molecular hydrogen fraction larger than $20\\%$.  Partial covering has only been observed once before in an intervening quasar absorption line system. In the lensed quasar APM 08279+5255 at $z=3.911$ \\citep{Kobayashi02}, one-third of the components of a strong {\\MgII} absorber were not detected toward the second image of the QSO, while fits to the components suggested $C_f=0.5$ in the other image, constraining the absorber size to be as small as $200$~pc.  We begin, in \\S\\ref{sec:2}, with a description of the VLT/UVES spectrum of HE~$0001-2340$, and display and quantify the observed properties of the $z=0.4524$ multiple cloud weak {\\MgII} system. \\S\\ref{sec:3} details the Voigt profile fit performed on this four-cloud system, including covering factor analysis of the first cloud.  It also describes the photoionization modeling method used to constrain the ionization parameters/densities of the four clouds. \\S\\ref{sec:4} gives modeling results for each cloud, while \\S\\ref{sec:5} discusses the implications of the cloud models on the origin of the gas, while \\S\\ref{sec:6} summarizes the findings and considers the properties of the absorption system in the context of broader questions relating to galaxy environments. ",
        "conclusions": "\\label{sec:6} Although {\\FeI} detections are extremely rare in extragalactic absorbers, photoionization models of Cloud 1 suggest that its small size ($<1$pc) may cause many such absorbers to go undetected; we estimate that two percent of the areas of $z\\sim0.45$ halos (24-49~kpc in radius) should be covered by such objects. With cold temperatures ($<100$K), high densities ($30<n_H<1100$), and a large molecular hydrogen fraction ($72-94\\%$), the properties of Cloud 1 are similar to the dense, small Milky Way clouds at high galactic latitudes observed in CO by \\cite{Heithausen02, Heithausen04} and \\citet{Penprase93}, suggesting that pockets of gas like Cloud 1 may be analogs of these Milky Way clouds in the halos of other galaxies.  Small scale clumps, with sizes of hundreds of AU, are predicted to exist both as part of the fractalized ISM in the halos of galaxies \\citep{Pfenniger94a, Pfenniger94b} and as condensates in turbulent gas \\citep{Glover07, Fujita09}.  A Type-Ia supernova may have been responsible for both the turbulence and observed iron abundance in Cloud 1.  In contrast, the ${\\alpha}$-enhancement observed in Cloud 2, with its low-Fe content, is suggestive of an origin in a Type II supernova-driven superwind. Multiple-cloud weak {\\MgII} absorption line systems are generally thought to originate in lines of sight that pass through more than one dwarf galaxies, through sparse regions of luminous galaxies at high impact parameters, or in gas-poor galaxies.  In this case, one of the four components happens to be a relatively rare type of cold and dense weak {\\MgII} absorber.  Via imaging of the HE0001-2340 field it may be possible to further constrain the absorption origin, perhaps by identifying a host galaxy or galaxies, or by finding a nearby galaxy whose environment is being sampled by these absorbers.  It is also of interest to find another similar {\\FeI} cloud, by searching large numbers of high-$S/N$, high-resolution quasar spectra.  Perhaps in another case, the sightline would be clear enough to allow access to the Lyman series lines, so that its metallicity could be directly measured, or to higher ionization absorption lines so that any surrounding, lower density gas phases could be identified."
    },
    "1004/1004.3874_arXiv.txt": {
        "abstract": "Cosmological $N$-body simulations indicate that the dark matter haloes of galaxies should be generally triaxial. Yet, the presence of a baryonic disc is believed to alter the shape of the haloes. Here we aim to study how bar formation is affected by halo triaxiality and how, in turn, the presence of the bar influences the shape of the halo. We perform a set of collisionless $N$-body simulations of disc galaxies with triaxial dark matter haloes, using elliptical discs as initial conditions. Such discs are much closer to equilibrium with their haloes than circular ones, and the ellipticity of the initial disc depends on the ellipticity of the halo gravitational potential. For comparison, we also consider models with initially circular discs, and find that the differences are very important. In all cases, the mass of the disc is grown quasi-adiabatically within the haloes, but the time-scale of growth is not very important. We study models of different halo triaxialities and, to investigate the behaviour of the halo shape in the absence of bar formation, we run models with different disc masses, halo concentrations, disc velocity dispersions and also models where the disc shape is kept artificially axisymmetric. We find that the introduction of a massive disc, even if this is not circular, causes the halo triaxiality to be partially diluted. Once the disc is fully grown, a strong stellar bar develops within the halo that is still non-axisymmetric, causing it to lose its remaining non-axisymmetry. In triaxial haloes in which the parameters of the initial conditions are such that a bar does not form, the halo is able to remain triaxial and the circularisation of its shape on the plane of the disc is limited to the period of disc growth. We conclude that part of the circularisation of the halo is due to disc growth, but part must be attributed to the formation of a bar. Bars in the halo component, which have been already found in axisymmetric haloes, are also found in triaxial ones. We find that initially circular discs respond excessively to the triaxial potential and become highly elongated. They also lose more angular momentum than the initially elliptical discs and thus form stronger bars. Because of that, the circularisation that their bars induce on their haloes is also more rapid. We also analyse halo vertical shapes and observe that their vertical flattenings remain considerable, meaning that the haloes become approximately oblate by the end of the simulations. Finally, we also analyse the kinematics of a subset of halo particles that rotate in disc-like manner. These particles occupy a layer around the plane of the disc and their rotation is more important in the spherical halo than in triaxial ones. We also find that, even though the final shape of the halo is roughly independent of the initial shape, the initially triaxial ones are able to retain the anisotropy of their velocity dispersions. ",
        "introduction": "\\label{sec:intro}  The vast majority of papers studying bar formation and evolution use idealised galaxy models with an exponential, or near-exponential, disc and a spherical halo. Yet, already in the nineties, cosmological $N$-body simulations had shown that dark matter haloes are generally non-spherical, with a tendency to be more prolate than oblate \\citep{Frenk1988, Dubinski1991, Warren1992, Cole1996}. Later simulations \\citep{Jing2002, BailinSteinmetz2005, Allgood2006, Novak+2006} have a sufficiently large number of particles to allow an adequate statistical analysis of the halo properties. Typically it is found that such haloes have isodensity axis ratios of $b/a \\sim 0.8$ and $c/a \\sim 0.6$, depending on the mass of the halo and on the details of the cosmological simulations. In particular, more massive haloes tend to be more triaxial (or rather, prolate). This is presumably due to the fact that massive haloes undergo a larger number of merging events, which take place non-isotropically, along preferred directions linked to the filaments of the global large-scale structure. Furthermore, the axis ratios show some radial dependence and the triaxiality is usually found to increase towards the centre, as seen by \\citet*{Hayashi2007} who measured the shapes of the isopotential surfaces of cosmological haloes.  Simulations of large scale structure used to be restricted to dark matter, ignoring the baryonic component, due to the computational cost. More recently, cosmological simulations that include gas and the treatment of several physical processes (such as star formation, gas cooling, chemical enrichment and supernova feedback) have been able to form discs. And it is found that haloes tend to become axisymmetric due to the disc \\citep{Kazantzidis2004, Tissera+2009}. Thus a halo which contains a baryonic disc is expected to be less triaxial than a pure dark matter halo. Indeed, observational constraints on halo shape find that present-day haloes are very mildly elongated on the plane of the disc, or even consistent with an axisymmetric potential \\citep{Trachternach2008}. This is also reflected on the statistics of disc galaxy shapes (\\citealt*{Lambas1992}; \\citealt{Fasano1993}; \\citealt{RixZaritsky1995}; \\citealt{Ryden2004, Ryden2006}), whose ellipticities are small compared to the ellipticities of haloes from cosmological simulations.  The shapes of discs are deformed by the aspheric halo potentials, such that in the equilibrium configuration the discs should be generally elliptical. At the same time, the presence of the disc acts to oppose the halo ellipticity on the plane \\citep{Jog2000, Bailin2007}, since the disc elongation is perpendicular to the halo elongation. The interplay of the baryonic disc with its halo should somehow reconcile the highly triaxial shapes of pure dark matter haloes from cosmological simulations with the near circularity of present-day observed galaxies.  Such disc galaxies may develop bars, whose dynamics must be influenced both by the non-circular disc and the non-spherical halo. In order to investigate the effect of a triaxial halo on bar formation, \\citet{GadottideSouza2003} performed $N$-body simulations of a spheroid embedded in a rigid triaxial halo potential. In their simulations, the spheroid was distorted into a bar-like structure, due to the halo triaxiality.  The effects of a cosmological setting were studied by \\citet*{Curir+2006}, who embedded circular discs in the haloes of cosmological simulations. They argue that the large scale anisotropies of the mass distribution influence the bar strengths and details of bar evolution, due to the continuous matter infall and substructures in the halo.  With simulations of considerably higher mass resolution, \\citet*{BerentzenShlosmanJogee2006} examined the effect of mildly triaxial haloes on bar evolution. In two of their three simulations with a live triaxial halo, they found that the bar quickly dissolved. The third case had a very massive and very mildly triaxial halo. In particular, its isopotential axis ratio in the equatorial plane is about 0.9. It this last case the bar does not dissolve.  \\citet{BerentzenShlosman2006} performed $N$-body simulations in which they grew circular seed discs in assembling triaxial haloes, within a quasi-cosmological setting. The final shape of the halo depends on the mass of the disc, but not on the timescale of its growth. They show that massive discs completely wash out the halo prolateness and then develop long-lived bars, whereas discs that contribute less to the rotation curve are less efficient in axisymmetrising their haloes. They claim that in less massive discs, the bar instability is damped by the halo triaxiality.  \\citet*{HellerShlosmanAthanassoula2007a} investigated the formation of discs by following the collapse of an isolated cosmological density perturbation. They include star formation and stellar feedback in their simulations, so that a baryonic disc forms inside the assembling dark matter halo. The halo triaxiality is decreased in their models during disc growth. Bars that are formed early decay within a few Gyr, but such bars are driven by the prolateness of the halo and do not follow the usual bar instability evolution. Also, they find that the tumbling of the triaxial halo figure is insignificant.  \\citet{Widrow2008} used the adiabatic squeezing method of \\citet{HolleyBockelmann2001} to produce triaxial halo models, which he applied to a study of F568-3. This work was focused on a study of the rotation curve of this galaxy and did not discuss the reasons for the changes of the halo and disc shapes. \\citet{Debattista2008} carried out control simulations in which a rigid disc was adiabatically first grown and then evaporated and found that the haloes were substantially rounder when the disc was near full mass, but that they returned to their initial shape after the disc was evaporated. The main goal of this paper and of its sequel \\citep{Valluri2009arXiv} was to search for the changes of orbital structure that could account for the changes of halo shape due to the baryonic component.  In the present work we continue along the lines of the above mentioned work and particularly investigate bar formation and evolution in triaxial haloes and the corresponding effects on the halo properties. Does the triaxiality of the halo inhibit bar formation, or change drastically the bar properties? Or, alternatively, are strong bars able to form inside triaxial haloes and then cause them to lose their remaining triaxiality? We investigate different models and different types of initial conditions. As initial conditions, we use elliptical discs which are designed to be in equilibrium with the elliptical potential of the halo and compare the results with those obtained with the more straightforward, but out of equilibrium, initially circular discs. We also focus on evaluating the separate effects of two factors that contribute to changing the halo shape: the growth of the disc mass, and the formation of a bar.  In Sect.~\\ref{sec:initialconditions} we present our initial conditions and in Sect.~\\ref{sec:ECdisks} we discuss bar growth and compare the results in initially axisymmetric and in initially elongated discs. In Sects.~\\ref{sec:coresizes} and \\ref{sec:timescales} we consider different halo core sizes and different time-scales for disc growth, respectively. In Sect.~\\ref{sec:nobars} we present a series of simulations, made in order to distinguish how much of the evolution of the halo shape is due to the introduction of the disc and how much to the growth of the bar. Vertical shapes are the subject of Sect.~\\ref{sec:covera}. In Sect.~\\ref{sec:disklike} we consider the effect of triaxiality on halo kinematics. Finally we summarise and conclude in Sect.~\\ref{sec:conclusions}.  \\begin{table} \\caption{Initial shapes of the haloes as given by the intermediate-to-major ($b/a$) and minor-to-major ($c/a$) axis ratios.} \\label{tb:halos} \\begin{center} \\begin{tabular}{l c c} \\hline halo & $b/a$ & $c/a$\\\\ \\hline 1  & 1.0 & 1.0 \\\\ 2  & 0.8 & 0.6 \\\\ 3-- & 0.7 & 0.5 \\\\ 3  & 0.6 & 0.4 \\\\ 3+ & 0.5 & 0.3 \\\\ \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "\\label{sec:conclusions}  Cosmological $N$-body simulations have shown that dark matter haloes of galaxies should be triaxial, at least in cases where there are no baryons \\citep{Allgood2006}. This can be the result of asymmetric mergings, or of a radial orbit instability \\citep{Bellovary2008}, coupled to tidal effects from other galaxies or from groups and clusters. The resulting prolate halo has very little or no figure rotation.  The aim of this paper was to investigate how such a halo will influence bar formation and, more generally, how it will influence the secular evolution of disc galaxies. Our combined disc and halo initial conditions were built so as to be as near equilibrium as possible, with discs which are initially elliptical. We have, nevertheless, also considered initially circular discs, to test what the effect of such less realistic initial conditions would be.  The growth and evolution of such discs drive the haloes rapidly towards axisymmetry, except for the innermost parts, where the final shape of the halo is elongated. This latter effect is independent of the initial halo triaxiality and is found also in initially circular haloes \\citep[e.g.][]{Athanassoula2005a, Athanassoula2007, Colin+2006}. It is linked to the angular momentum exchange within the galaxy and the formation of a `halo bar', which is shorter and less elongated than the disc bar, but rotates with the same pattern speed.  \\begin{figure} \\includegraphics[width=\\columnwidth]{figs/fig25.eps} \\caption{Left: velocity dispersions $\\sigma_{\\varphi}$ (solid lines), $\\sigma_{R}$ (dashed lines) and $\\sigma_{z}$ (dotted lines) for the $|z|<0.5$ region. For comparison, the average velocity dispersion of the spherical case is show in all panels (dot-dashed line). Right: velocity dispersions for the permanently disc-like halo particles, by definitions 1 (thick lines) and 2 (thin lines), as discussed in the text. The line types are as for the left panels. All panels, both left and right, correspond to $t = 800$.} \\label{fig_sigs_keep_800} \\end{figure}  This innermost prolate elongation put aside, the remaining halo tends towards axisymmetry even for models with considerable initial triaxiality, and this from the moment the disc starts growing. One can distinguish two different axisymmetrisation phases. Initially, while the disc grows, this trend towards halo axisymmetry is quite rapid. The second phase depends on whether or not a bar is formed: in the presence of a bar, the circularisation continues.  The disc shape also changes with time. Initially it is elongated, its ellipticity depending on that of the halo. Thus the $A_{2}$ is non-zero at $t = 0$, although there is no bar. It decreases with time and reaches a minimum around the time that the disc has reached its maximum mass. After that, the bar starts forming and induces a further increase of halo axisymmetry. In our models, bar formation takes place in the presence of haloes that are still considerably triaxial, even though the triaxiality is later erased. In the more triaxial cases, bar formation is somewhat delayed, but the subsequent evolution of the bar proceeds in a manner similar to the spherical halo case. We have also presented triaxial models that do not develop bars at all; but in such cases, the equivalent spherical ones do not either. The general agreement of our results with previous studies is in the sense that a truly triaxial halo cannot coexist with a strong bar for very long: one of these non-axisymmetries must give way. We have presented simulations in which the bar prevails and the halo triaxiality yields. This argues that in situations where the parameters are such that a bar is known to form in the spherical case, it would also form in the (initially) triaxial cases, further erasing the triaxiality as it does.  \\begin{figure} \\centering \\includegraphics[scale=1]{figs/fig26.eps} \\caption{Top: anisotropy parameter for the entire halo (filled squares) and for the $|z|<0.5$ region (open squares). Bottom: anisotropy parameter with definitions 1 (solid symbols) and 2 (open symbols), measured at $R=1$ (circles) and at $R=3$ (triangles). Note that the two panels have very different scales.} \\label{fig_betas} \\end{figure}  Using circular discs, instead of very near-equilibrium elliptical ones, may give rise to quite different $A_{2}$ evolutions, particularly if the triaxiality of the halo is important. In the case of a very triaxial halo containing a circular disc, the $A_{2}$ initially increases very abruptly to reach a strong peak and then decays. That was the only simulation in which halo triaxiality damped bar formation, but we stress that this is only when a circular disc was used. In that very triaxial halo, a circular disc is even further from equilibrium than in the less triaxial ones. There are other situations in which $A_{2}$ grows very abruptly in the very beginning and then decays without forming a bar, and this only happens if we use circular discs as initial conditions. For instance, in the case of low mass discs, the initial $A_{2}$ increase is due to the circular disc becoming excessively distorted. After the peak, the remaining non-zero $A_{2}$ is due to a mild oval distortion in the centre of the disc (and also some vague spiral structure), but none of this amounts to actual bar formation (furthermore, there is no exchange of angular momentum between disc and halo).  Such behaviour of $A_{2}$ might be analogous to those of \\citet{BerentzenShlosmanJogee2006}, who also have initially circular discs and find that the bars in their live triaxial haloes dissolve after a few Gyr (except in the case of their more cuspy halo, which does not retain its triaxiality). Similarly, \\citet{BerentzenShlosman2006}, also with initially circular discs, find that bar formation is damped by the triaxiality of the halo. In their models, strong bars form in discs that have erased the halo prolateness almost completely.  In all our models, as in previously run models with initially spherical haloes, bar formation is always necessarily accompanied by angular momentum exchange; more specifically, angular momentum is transferred from the disc to the halo. In the beginning of the simulations, practically all the angular momentum is in the disc. In all models where there is bar formation, the disc-to-total ratio of angular momentum decreases steadily. Further properties of the angular momentum exchange are similar to those found in spherical haloes; e.g. it is related to the bar strength; strong bars causing more transfer from disc to halo. Also, in haloes with large core, the bars are weaker and the angular momentum lost by the disc is correspondingly small. There are simulations where the $A_{2}$ is non-zero due to oval distortions of the disc which, however, do not qualify as real bars. In such cases, there is very little, if any, angular momentum exchange between disc and halo. We also find differences in angular momentum transfer depending on whether we use circular or elliptical discs as initial conditions. The initially circular discs become excessively elongated during the very short phase of disc growth and they thus develop strong bars faster then the equivalent model with an elliptical disc. These discs, that host stronger bars, lose angular momentum faster. The other quantity that is also closely related to angular momentum transfer and bar strength is the shape of the halo: a stronger bar causes the halo to become axisymmetric faster.  Our standard models have haloes with a small core, of the type called MH in \\citet{Athanassoula.Misiriotis02}. Since the core radius is an essential feature in determining the evolution of the galaxy \\citep{Athanassoula.Misiriotis02, Athanassoula2003}, we also ran simulations with large cores, i.e. of the MD type. Here the effect of the disc is even stronger. In all models, by the time the disc was fully grown, the halo has already become nearly axisymmetric, so that any further evolution follows closely that of the MD models of \\citet{Athanassoula.Misiriotis02}, namely, compared to MH models, the bar is less strong and there is less angular momentum exchange between the disc and the halo.  We also presented models in which the growth time of the disc was either much smaller, or much bigger than that of the standard models. We found that this does not affect much the results, limiting itself mostly to a temporal shift of the evolution, backwards or forwards in time, respectively. This is in good agreement with similar tests made by \\citet{BerentzenShlosman2006}, albeit with different models.  A crucial question in this context is how much of the loss of the halo triaxiality is due to the introduction of the disc and how much to the formation of the halo. In the case of MD discs, i.e. of discs which are heavy relative to the halo in the inner parts of the galaxy, we saw that by the time the disc has reached its full mass, the halo has become nearly spherical. It is thus reasonable to expect that low mass discs would have a smaller effect on the halo shape, as already found in the simulations of \\citet{BerentzenShlosman2006}. We repeated such simulations here, but with initially elliptical discs. In these cases, our disc has sufficiently low mass so that very little angular momentum can be emitted from its inner regions. Thus although the halo was ready to receive angular momentum, not much exchange was possible for lack of sufficient emitters. As a result, the bar was exceedingly weak, as in the very low mass discs of \\citet{Athanassoula2003}, and so brought no effect on the halo triaxiality. Thus in these cases the slight decrease of the halo ellipticity when the disc is introduced, is not followed by further decrease caused by the bar. We point out that our corresponding spherical halo model with low mass disc does not develop a bar either, showing that it was not the triaxiality of the other models that was responsible for inhibiting bar formation. In fact, in none of our sets of models with initially near-equilibrium discs, do we find cases where a bar develops in the spherical halo, while failing to do so in the equivalent triaxial ones. Such models, with discs of low mass relative to the halo, give us useful information on galaxies with low surface brightness discs, and argue that such galaxies should not have suffered much loss of halo triaxiality, and that their halo shape should be near what it was from galaxy formation and as due to effects of interactions and mergings.  We made further simulations to find what fraction of the triaxiality loss can be attributed to the disc introduction and what to the formation and evolution of the bar. For this we ran models where the disc was kept artificially axisymmetric, models where the disc was very hot (so that bar grew very slowly and was more of an oval than a strong bar) and models where the disc was rigid. In these three cases (as well as in the low mass case), the simulations were designed so as not to form bars. In all cases where there is no bar, the halo is able to remain triaxial until the end of the simulation. Any circularisation that takes place is restricted to the period of disc growth. This circularisation is particularly small if the disc mass is small or if the disc is rigid. The rigid discs -- where the potential of a (permanently circular) disc is represented by an analytic potential instead of live particles -- are the ones that cause the smallest effects on their haloes, suggesting that it is not merely the presence of disc mass that alters the halo shape, but also the active response of the disc orbits that oppose the halo elongation.  We find that the vertical shape of the triaxial haloes is also affected both by disc growth and by bar formation, in the sense that haloes become less flattened with time. These changes are, however, always small compared to the changes in the equatorial shape. Nevertheless, in cases where the $b/a$ increase is more pronounced, that of $c/a$ is important also. The general evolution of vertical shapes is thus analogous to that of the halo shape on the plane of the disc: disc introduction causes some degree of change and the subsequent formation of a bar determines whether or not there will be further loss of flattening. But even in models where a strong bar forms (and the halo is thus completely axisymmetrised), the vertical flattening remains considerable, meaning that the haloes have become approximately oblate by the end of the simulation.  As the bar forms, the vertical structure of the disc is also affected. As with previous simulations of axisymmetric galaxies, the discs in the triaxial models also undergo buckling episodes and peanut formation. In our initially triaxial models that form strong bars, the first buckling of the disc occurs later than in the spherical case. Consequently, peanut formation is delayed in the triaxial models. The models with initially circular discs develop stronger peanuts before the corresponding models with elliptical discs. All this is in agreement with the fact that peanut strength is related to bar strength \\citep{Athanassoula.Martinez-Valpuesta.2010}, and the fact that bars grow faster in initially circular discs.  The spherical halo particles in a layer near the equatorial plane acquire during the evolution some rotation in the same sense as the disc rotation, as in \\citet{Athanassoula2007}. We extended that analysis to triaxial haloes. We select halo particles that are permanently in the vicinity of the $z=0$ plane; by analysing their radial profiles of tangential velocities, we find that these `disc-like' halo particles show considerable rotation, with velocities of the order of half of that of the disc particles. Such rotation is present as well in the triaxial models and the peak tangential velocities depend systematically on the initial triaxiality of the model. There is more rotation in the spherical halo than in the triaxial ones, even though their shapes are roughly the same by the end of the simulation. Apart from the rotation, there is another kinematic feature which is dependent on the initial triaxiality: the anisotropy of the velocity dispersions. If we consider the entire haloes, we see that the in spherical one the velocities are isotropic, but the initially triaxial models retain their anisotropy even after their shapes have become approximately spherical. We find that the velocity anisotropy at the end of the simulation depends systematically on the \\textit{initial} triaxiality.  To summarise, we have presented simulations of bar formation in triaxial haloes and showed that the haloes become more axisymmetric due to two separate factors: disc growth and bar formation. Typically half the circularisation can be attributed to the introduction of the disc and the other half to the formation of a strong bar. Halo vertical flattening is also affected, but to a much lesser degree, meaning that haloes become roughly oblate by the end of the simulation. This is the first study of bar formation and disc bar evolution in which live elliptical discs within live triaxial haloes were employed as the initial conditions for $N$-body simulations. Elliptical discs have the advantage of being closer to equilibrium and of not responding to the presence of the aspherical halo potential by becoming excessively and unphysically distorted, as circular discs tend to do. We have also analysed the kinematics of the halo and pointed out that even after the haloes lose their triaxiality, they are able to retain the anisotropy of their velocity dispersions. We also note that the disc-like particles of the halo rotate less in the haloes that were initially triaxial.  Interesting issues remain to be explored. Simulations such as these, in which all components are live, can be used to study in detail the orbital structure, not only during the period of disc growth, but also during the period of bar formation. Another interesting issue is the effect of gas on the dynamics of these systems. Simulations analogous to the ones we have presented, but including also a gas component are currently under way and will be the subject of a forthcoming paper."
    },
    "1004/1004.1330_arXiv.txt": {
        "abstract": "The 3D magnetic topology of a solar active region (NOAA 10956) was reconstructed using a linear force-free field extrapolation constrained using the twin perspectives of \\emph{STEREO}. A set of coronal field configurations was initially generated from extrapolations of the photospheric magnetic field observed by the Michelson Doppler Imager (MDI) on \\emph{SOHO}. Using an EUV intensity-based cost function, the extrapolated field lines that were most consistent with 171~\\AA\\ passband images from the Extreme UltraViolet Imager (EUVI) on \\emph{STEREO} were identified. This facilitated quantitative constraints to be placed on the twist ($\\alpha$) of the extrapolated field lines, where $\\nabla \\times {\\bf B} = \\alpha {\\bf B}$. Using the constrained values of $\\alpha$, the evolution in time of twist, connectivity, and magnetic energy were then studied. A flux emergence event was found to result in significant changes in the magnetic topology and total magnetic energy of the region. ",
        "introduction": "Extreme solar events such as solar flares and coronal mass ejections (CMEs) originate in the corona of active regions. As the corona is a low-$\\beta$ plasma, the magnetic field rather than the gas pressure governs its structure and evolution \\citep{gary:2001}. This enables force-free field extrapolations of the photospheric field to be used to study the detailed topology of the solar corona \\citep{longcope:2009}. Of particular interest to space weather applications is how the three-dimensional (3D) field topology of active regions governs the production of extreme solar events. Unfortunately, previous attempts to reconstruct the 3D topology of active region fields have been limited by observational issues, such as incomplete measurements of the photospheric vector field, or theoretical issues such as ill-constrained extrapolation algorithms (e.g., \\citeauthor{derosa:2009} \\citeyear{derosa:2009} and references within). With the launch in December 2006 of the \\emph{Solar TErrestrial RElations Observatory (STEREO)} satellites, the observational limitations on reconstructing the 3D topology of active regions has largely been overcome, thus opening a new era in the study of the ever-changing solar corona.  Observations of the coronal magnetic field are difficult to obtain due to the tenuous nature of the corona and the absence of strong magnetically sensitive emission lines. Estimates of the strength of the coronal magnetic field are possible using radio observations, however, this remains a difficult problem requiring simultaneous knowledge of both the temperature and density of the solar atmosphere~\\citep{brosius:2006}. In order to overcome these issues, modellers have developed a variety of theoretical techniques to examine the topology of the coronal magnetic field~\\citep{gary:1989,mikic:1990}. These methods require observations of the photospheric or chromospheric magnetic field as boundary conditions for the extrapolations. Well constrained extrapolations are normally produced using photospheric vector magnetograms, although \\citet{metcalf:1995,metcalf:2002} have developed methods that make use of chromospheric vector magnetograms. As the chromosphere is in a force-free state, chromospheric vector magnetograms are ideally suited as boundary conditions in force-free extrapolation methods. However, these observations have a limited field-of-view and spatial and temporal resolution. As such, photospheric magnetograms are commonly used as boundary conditions for magnetic field extrapolations.  The extrapolation of the photospheric field into the corona requires that a number of assumptions be made. Firstly, it is assumed that the coronal magnetic field remains stable for the duration of the observation. Therefore, a static magnetic field model is normally used. In addition, it is reasonable to assume that the field is in a force-free configuration. This results in a vanishing Lorentz force, $\\mathbf{j}\\times\\mathbf{B} =  0$, where $\\mathbf{j}$ is the current density, $\\mathbf{B}$ is the magnetic field. The force-free condition requires that the current density remains parallel to the magnetic field, $\\mu_0\\mathbf{j} = \\mathbf{\\alpha} \\mathbf{B}$, and therefore $\\mathbf{\\nabla} \\times \\mathbf{B} =\\mathbf{\\alpha} \\mathbf{B}$, where $\\mathbf{\\alpha}$ is the amount of twist in the field and is some times referred to as the torsion function. In the case of $\\mathbf{\\alpha} = \\alpha(\\mathbf{r})$, the governing equations becomes $\\mathbf{\\nabla} \\times \\mathbf{B} =  \\alpha(\\mathbf{r}) \\mathbf{B}$ and $\\mathbf{B}\\cdot\\nabla\\alpha(\\mathbf{r})=0$. This non-linear force free (NLFF) field is the generic solution to the force-free field condition. Complete knowledge of the photospheric magnetic field allows for the inversion of the preceding equations and the calculation of $\\alpha$ at each point on the photosphere. As such, vector magnetograms are required to calculate the NLFF field in the corona~(\\citeauthor{Schrijver:2006} \\citeyear{Schrijver:2006}). While NLFF methods offer the possibility of increased accuracy, they are computationally intensive. Additionally, NLFF extrapolation suffer from several other sources of error, such as the limited field-of-view, non-force-free nature of the photospheric field, the intrinsic 180-degree ambiguity, and low signal to noise of the transverse component of vector magnetograms~\\citep{derosa:2009}. Forcing $\\mathbf{\\alpha} = 0$ results in the potential or current-free solution, which is a first order approximation to the coronal magnetic topology. The potential field solution is the first of two subclasses of solutions to the force-free field. Given the reduced complexity of the potential solution, the only boundary condition necessary to form a solvable set of equation are measurements of the line-of-sight (LOS) magnetic field~\\citep{gary:1989}. The second and final subclass of solution to the force-free field is the linear force-free (LFF) field, where $\\mathbf{\\alpha}$ = constant and $\\mathbf{\\nabla} \\times \\mathbf{B}=\\alpha \\mathbf{B}$. Assuming that \\al{} is constant for all field lines allows for the formation of a closed set of equations with only knowledge of the LOS magnetic field, provided \\al{} is predefined~\\citep{alissandrakis:1981}.  A major difficulty in studying the extrapolated field line distribution is deciding on what is the most appropriate way to constrain $\\alpha$. \\citet{leka:1999} and more recently \\citet{burnette:2004} compared three methods and applied them to a single active region. The first method considered was based on the best fit value of $\\alpha$ from a LFF extrapolation of a longitudinal magnetogram that best matched the horizontal field of the corresponding vector magnetogram~\\citep{pevtsov:1995}. Secondly, the value of $\\alpha$ was chosen from the mean of the distribution of local $\\alpha_z = \\alpha(x,y)$~\\citep{pevtsov:1994}. In the final method, the value of $\\alpha= J_z/B_z$ was inferred from the slope of a least-squares linear fit to the distribution of local $J_z(x, y)$ versus $B_z(x, y)$.  \\citet{burnette:2004} conducted a detailed study of these three methods, concluding that they give statistically consistent values of $\\alpha$. As the methods described above require the availability of high quality vector magnetograms, they are not always an available option for constraining the values of $\\alpha$.  EUV and X-ray images of emitting coronal structures, mainly loops, can also be used to constrain the choice of \\al{} in magnetic field extrapolations. Initially, loop tracing methods were used to provide a first guess at the structure of coronal loops and through the application of numerical methods the 3D structure found~\\citep{aschwanden:1999}.  Other authors have developed methods that minimize the separation between the field lines calculated from observations and extrapolations as a constraint on the choice of \\al{}~\\citep{lim:2007}. \\citet{wiegelmann:2005},~\\citet{carcedo:2003}, and ~\\citet{aschwanden:2008a} developed related techniques that examine the EUV emission of coronal loop structures and retrieve their two dimensional structure using a cost function method. \\citet{wiegelmann:2005} developed a cost function grading method to evaluate the agreement of calculated field lines and the observed emission of coronal loops. This enabled them to gauge the optimum value of $\\alpha$ for a small number of LFF extrapolations. The cost function method was found to allow for the selection of the optimum value of $\\alpha$ used in the LFF extrapolations, however, the number of possible values of $\\alpha$ examined was insufficient to allow for a detailed study of the active region. Most recently, \\citet{aschwanden:2008,aschwanden:2009} expanded on the methods discussed above with the use of the \\emph{STEREO} spacecraft to study the structure and oscillations of coronal loops.  Once the value of the magnetic twist ($\\alpha$) in a region has been well constrained, an accurate estimate of the magnetic energy or more importantly, the free magnetic energy, can be estimated. The magnetic energy budget of active regions can be calculated using volume-integral and surface-integral methods such as the magnetic Virial Theorem~\\citep{emslie:2004}. The free-energy is calculated as the difference between the magnetic energy of the NLFF field and that of a lower bound such as the potential (current-free) field. These methods have shown the free energy in active region's to be in the range of $10^{32}-10^{33}$~ergs depending on the size of the region~\\citep{gary:1989,regnier:2007,metcalf:1995, metcalf:2005}. The magnetic Virial Theorem assumes the measured field is a force-free field and is only possible if the measured data is force-free consistent. Additionally, knowing the topology of the active region at instances during its evolution, changes in the regions connectivity can be examined. Studies of the changing connectivity of active regions using a combination of active region segmentation techniques and magnetic field extrapolations have been achieved~\\citep{longcope:2001,longcope:2007,longcope:2009}. \\citet{longcope:2009} concluded that comparison of the changing connectivity in active region over time could provide insight into energetics and reconnection in the coronal fields.   The methods described in this paper allow for the 3D structure of the magnetic field of an active region to be recovered and a detailed analysis of its evolution to be performed using \\emph{SOHO} Michelson Doppler Imager (MDI,~\\citeauthor{scherrer:1995} \\citeyear{scherrer:1995}) LOS magnetograms and \\emph{STEREO} Extreme UltraViolet Imager (EUVI, \\citeauthor{howard:2000}~\\citeyear{howard:2000}) images. The observation and the transformation of co-ordinates between \\emph{STEREO} and \\emph{SOHO} reference frames are outlined in Section~\\ref{observation}, while details of the extrapolation method and the cost function are then given in Section~\\ref{method}. Our results, relating to twist, connectivity and magnetic energy are presented in Section~\\ref{results}, and our conclusions are given in Section~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  NOAA 10956 rotated onto the solar disk on 2007 May 15 and on 2007 May 19 the region produced a B9 class flare, starting at 12:34~UT, peaking at 13:02~UT, and ending at 13:19~UT. The flare location suggests that the source of the flare was the loop structure connecting regions 1$^-$ to 2$^+$ in our analysis. According to the \\emph{SOHO}/LASCO CME Catalogue\\footnote{\\url{http://cdaw.gsfc.nasa.gov/CME\\_list}} , the region had an associated CME which first appeared at 13:24~UT in the C2 field of view. In addition, a disappearing filament was observed erupting from the lower right section of the regions beginning at 12:31~UT~\\citep{long:2008}.  Our analysis has highlighted a number of physical changes within the magnetic structure of the active region around this time. These are: \\begin{itemize} \\item Magnetic flux cancellation and emergence. \\item Increase in the amount of $\\alpha$ or twist along the observed loop. \\item Changes in the connectivity between sub-regions. \\item Significant changes in the magnetic energy stored in the region. \\end{itemize}   The increase in the amount of twist or \\al{} along the observer loop from 09:35~UT to 11:11~UT, was accompanied by an increase in the magnetic flux within sub-region 1$^+$, and decreases in the magnetic flux within sub-regions 1$^-$ and 2$^+$ and an increase in the amount of magnetic energy stored within the volume. The increase in magnetic energy can be explained by the increase in \\al{} used to calculate it. However, as \\al{} is obtained from a comparison to EUV observations this increase is taken as real. Following the increase in \\al{} at 11:11~UT, the amount of twist along the observed loop decreases to 09:35~UT levels by 14:27~UT. During this time period the magnetic flux in sub-regions 1$^+$, 1$^-$, 2$^+$, and 3$^+$ increases by $0.4-1.5\\times10^{20}$~Mx and there is a $1.1\\times10^{20}$~Mx decrease in the magnetic flux in region 4$^+$. Additionally, the magnetic energy decreases to 09:35~UT levels over the same time period. Changes in the connectivity from region 1$^-$ can be associated with the changes in \\al{} prior to 11:11~UT. However, by 12:47~UT the amount of flux connecting to region 1$^-$ to regions 1$^+$ and 3$^+$ rose significantly.  From 09:35~UT to 11:11~UT the difference in magnetic energies increased from $1.3\\times10^{30}$ to $6.1\\times10^{30}$~ergs in the sub-volume centered on the observed loop and from $2.1\\times10^{30}$ to $10.9\\times10^{30}$~ergs in the full computational domain. This result is particularly important as the energy calculation at 11:11~UT is significantly different from the potential case, as indicated by the error estimate. It should be noted that smaller error bounds may be possible provided that a more detailed analysis of the cost function (smaller $\\Delta_{\\alpha}$) is performed and that a clear minimum is the cost function exists. By 12:47~UT this difference in magnetic energy in the system had dropped by a factor of three and by 14:27~UT had decreased below the difference at 09:35~UT. The build-up and subsequent decrease in magnetic energy in the system is correlated with the mechanisms thought to be responsible for the release of the B9 flare at 12:30~UT. The method is unable to calculated the free energy of the system due to the limitations of LFF methods~\\citep{seehafer:1978}. However, the reported changes in the difference in magnetic energy between the potential and LFF fields provide a estimate for changes in the free energy contained in the region during this time.   The above investigation of magnetic properties of NOAA 10956 has shown that from 09:35~UT there was a significant amount of flux cancellation within the region, accompanied by an increase in the amount of twist and free energy along the observed loop. Following the onset of the flare at 12:34~UT, the amount of twist and magnetic energy in the loops decreased. This was accompanied by the decrease in open flux leaving region 1$^-$ and increase in the magnetic flux in the region. It is suggested that the sudden decrease and subsequent increase in magnetic flux within the active region and the changes in twist and magnetic energy caused small changes in the observed connectivity between sub-regions, Table~\\ref{connectivity}. The emergence of magnetic flux within the region and resulting increase in magnetic energy could have resulted in the region passing a critical threshold, passed which it was in an unstable configuration. Under the theory of self-organized criticality this would have resulted in a series of events returning the region to a stable configuration~\\citep{vlahos:2004}. Whether the mechanism responsible for passing this critical threshold is the unbalanced emergence of flux within the region or the resulting non-potentiality of the flux ropes remains unknown. The decrease in both \\al{} and magnetic energy and increase in connectivity of loop connecting region 1$^-$ to region 3$^+$ suggest that the loop had room to expand and relax into the neighboring sub-region following the flare and eruption in the neighboring sub-region.  The method described in this paper allows for the detailed analysis of the amount of current or $\\alpha$ in coronal loop structures and with the added use of the twin perspectives of \\emph{STEREO} EUVI images, are an expansion of the methods developed by \\citet{wiegelmann:2005} and \\citet{carcedo:2003}. These previous studies involved the use of cost functions as a mechanism for constraining the free parameters in extrapolations of the solar corona and were restricted to the investigation of one or two properties and singular instances in time. More advanced NLFF method have been used to investigate the free magnetic energy of active regions~\\citep{regnier:2007,Schrijver:2008} and in some instance its evolution~\\citep{thalmann:2008,thalmann:2008a}. However these methods have been restricted by the limited time cadence of vector magnetogram data and provide little detail into the physical changes in the structure of active regions surrounding the buildup and release of magnetic energy. While the method developed here are unable to accurately calculate the free energy of the active region, changes in the energy difference can be used as a proxy to changes in the free energy.  The method has been shown to detect changes in the amount of twist (and hence current) within coronal loop structures using the twin perspective of \\emph{STEREO}. Coupling the temporal analysis of the structure of NOAA 10956 with an investigation of the magnetic flux within each sub-region has provided great insight into the evolution and driving forces within the active region. Future work is needed improve the method and remove the need for user input. Expanding the method for the analysis of multiple loop structures in active regions would greatly increase the diagnostic power of the algorithm."
    },
    "1004/1004.3784_arXiv.txt": {
        "abstract": "On  large  angular scales  ($\\gtrsim$  60$^{\\circ}$), the  two-point angular correlation  function of  the  temperature of  the cosmic  microwave background (CMB), as  measured (outside of the plane  of the Galaxy) by   the   {\\it  Wilkinson   Microwave   Anisotropy  Probe},   shows significantly lower large-angle  correlations than expected from the standard inflationary cosmological  model. Furthermore, when derived from  the full CMB  sky, the  two lowest  cosmologically interesting multipoles, the  quadrupole ($\\ell=2$) and  the octopole ($\\ell=3$), are unexpectedly  aligned with each other.  Using randomly generated full-sky and  cut-sky maps,  we investigate whether  these anomalies are correlated at a statistically significant level. We conclusively demonstrate  that,   assuming  Gaussian  random   and  statistically isotropic  CMB anisotropies, there  is no  statistically significant correlation between the missing power on large angular scales in the CMB and the alignment of  the $\\ell=2$ and $\\ell=3$ multipoles.  The chance to  measure the sky  with both such a  lack of large-angle correlation  and such an alignment  of the  low multipoles is  thus quantified  to be below $10^{-6}$. ",
        "introduction": "Several anomalies in the cosmic microwave background maps observed by the WMAP satellite \\citep{Bennett2003,Jarosik2010} have been recently observed and much discussed in the literature.  These include alignments of the lowest modes of CMB anisotropy with each other, and with geometry and direction of motion of the Solar System, as well as unusually low power at these largest scales. Attempts to explain these anomalies include astrophysical, instrumental, cosmological causes, as well as arguments that faulty data analysis or {\\it a posteriori} statistics are at work (for a review, see \\citet{CHSS_AA_review}). Two particularly puzzling features at large angular scales are the alignments of large-angle anisotropies \\citep{oliveira04,Schwarz2004,Land2005a,lowl2}, and the low power at large scales \\citep{Spergel2003,wmap123}; each of these anomalies are $\\ll 1\\%$ likely. For reviews with differing points of view see \\citet{wmap7_anomalies,CHSS_AA_review}.  A natural question that arises is whether the alignments and low power at large scales are correlated. Naively, the answer is negative, since the alignments are defined by orientation of the multipoles, and are independent of the total power. In the language of multipole vectors, the normalization at each multipole defines power, and is independent of the multipole vectors that define orientations of multipoles. Such a conclusion has been verified, for full sky maps, by \\citet{Rakic2007}.  However, the answer for the {\\it cut-sky} maps is less clear {\\it a priori}. Consider the case where the large-scale multipoles are largely planar (as is the case for our own quadrupole and octopole), thus creating alignments. If the sky cut happens to be parallel to this plane, then the cut-sky power may be unusually low, and this may well be happening more often than when the cut is applied to unaligned (statistically isotropic) maps. Clearly, the only safe way to investigate the correlation between the low power and alignments in the presence of the realistic sky cut is to perform Monte Carlo comparisons with Gaussian random, statistically isotropic maps.  In this {\\it Letter} we investigate the correlation between the alignments and low power for the cut-sky as well as full-sky maps assuming a Gaussian random, statistically isotropic, cosmological model. ",
        "conclusions": "\\label{sec:conclude}  Applying the multipole vector formalism to Gaussian random and statistically isotropic realizations of CMB maps (both full-sky and cut-sky) lacking correlation on large angular scales, we find no increased probability for alignment of the quadrupole and the octopole than expected in the case of unconstrained random maps. These low-power maps also do not show higher angular momentum dispersion. On the flip side, we also find that realizations of CMB maps (both full-sky and cut-sky), constrained to have aligned quadrupole and octopole at a level greater than that exhibited by WMAP, do not have lower power on large scales than expected in the case of unconstrained random realizations.  Our results conclusively demonstrate that, under the standard Gaussian and isotropic model, there is no statistically significant correlation between the missing power on large angular scales in the CMB and the alignment of the $\\ell=2$ and $\\ell=3$ multipoles. Therefore, in the context of the standard model, their combined statistical significance is equal to the product of their individual significances. For example, simultaneous observation of the missing large-angle correlations with probability $P(S_{1/2}^{\\rm cut\\, sky})\\lesssim 0.1\\%$ {\\it and} alignments with the probability $P(S^{(3,3)})\\simeq 0.1\\%$ is likely at the $\\lesssim 0.0001\\%$ level.  Given that both anomalies occur at the largest observable scales and are correlated with special directions in the sky (ecliptic and/or dipole), they clearly require a causal explanation."
    },
    "1004/1004.1054_arXiv.txt": {
        "abstract": "{The non-linear Hall term present in the induction equation in the electron-magneto-hydrodynamics limit is responsible for the Hall drift of the magnetic field and, in some cases, for the so-called Hall instability. } {We investigate whether or not the growth rates and eigenfunctions found in the linear analysis are consistent with the results of non-linear numerical simulations.} {Following the linear analysis of Rheinhardt \\& Geppert, we study the same cases for which the Hall instability was predicted by solving the non-linear Hall induction equation using a two-dimensional conservative and divergence-free finite difference scheme that overcomes intrinsic difficulties of pseudo-spectral methods and can describe situations with arbitrarily high magnetic Reynolds numbers. } {We show that unstable modes can grow to the level of the background field without being overwhelmed by the Hall cascade, and cause a complete rearrangement of the field geometry. We confirm both the growth rates and eigenfunctions found in the linearized analysis and hence the instability. In the non-linear regime, after the unstable modes grow to the background level, the naturally selected modes become stable and oscillatory. Later on, the evolution tends to select the modes with the longest possible wavelengths, but this process occurs on the magnetic diffusion timescale. } {We confirm the existence of the Hall instability. We argue against using the misleading terminology  that associates the non-linear Hall term with a turbulent Hall cascade, since small-scale structures are not created everywhere. The field evolves instead in a Burgers-like manner, forms local structures with strong gradients which become shocks in the zero resistivity limit, and Hall waves are launched and propagated through the entire domain.} ",
        "introduction": "The influence of the Hall term  on the magnetic field evolution of a neutron star (NS) has been debated for many years. By analogy with the vorticity equation in ordinary hydrodynamics, \\cite{GR1992} conjectured that the transfer of magnetic energy from large to small scales proceeds in a similar way to ordinary turbulence. However, the analogy of the Hall induction equation with the vorticity equation is not complete, and the conjecture remained to be confirmed by multidimensional numerical simulations. \\cite{Vai2000} proposed a mechanism for the fast dissipation of magnetic field based on the Hall drift in stratified media. They correctly pointed out that Hall currents are able to create current sheets (which are sites for efficient dissipation) and that the evolution of the toroidal field resembles the Burgers equation. The same Burgers-like equation is applicable even to non-stratified media, but in a spherical shell \\citep{PonsGeppert2007}, in which the Hall term in the induction equation  tends to create current sheets instead of ordinary turbulence. \\cite{RG2002} (hencefort RG) showed by a linear analysis that, in a one-component (electron) plasma, a large-scale background magnetic field may become unstable to smaller scale perturbations. This Hall-drift induced instability (HI) occurs when the magnetization parameter is high and the background field has enough curvature. Since these conditions may be realized in the crust of a NS, the problem of the HI became interesting not only from the MHD point of view but also for the astrophysics community. Although the Hall-drift is a non-dissipative process, the growth of small-scale magnetic field components modifies the overall magnetic field structure and opens the possibility of more rapid field decay than pure ohmic dissipation would predict. This effect may be relevant, at least, in relatively young neutron stars.  Determining its effect on the long-term evolution, and which configurations are stable on long time-scales are different issues. A review by \\cite{Reis2009} discusses the main properties of stable magnetic configurations emphasizing the importance of a non-barotropic stratification in neutron stars, and provides insight into the expected long-lived magnetic field structures. Some observational implications of rapid magnetic field decay were also discussed by \\cite{GR2002,PonsLink2007,Aguilera2008a,Aguilera2008b}. Other attempts to understand the long-term evolution of neutron star magnetic fields by numerical simulations were performed by \\cite{Hoyos2008}, but were restricted to 1-dimensional geometries.  The appearance of the HI was disputed by \\cite{HR2002}: they did not find any evidence of the instability in their numerical study of field decay under the influence of the Hall drift in a constant density spherical shell. The study was later extended to include density gradients \\citep{HR2004}. However, it was clear that they could not find the HI because of numerical limitations of their code, which is unable to handle situations with sufficiently high magnetic Reynolds number. The relevance of the HI to astrophysical scenarios had already been questioned by RG. Later \\cite{CAZ2004} conjectured that the gap in the energy spectrum created by the HI could be rapidly filled in by the Hall cascade. They also computed Hall-wave eigenfunctions, including the effect of density gradients across the crust, and found agreement with the results of RG. An important component of their study was their estimate of the response of the crust to the magnetic stresses induced by Hall waves. \\cite{PonsGeppert2007} performed 2D simulations by solving the non-linear Hall-induction equation in a realistic NS crust and found some evidence for a rapid reorganization of the field during a relatively short period, but no clear confirmation of the HI could be concluded. The code again had numerical limitations that avoided additional exploration of extreme limits to confirm or rule out the existence of the HI. \\cite{WH2009}(hereafter WH) used their 3D pseudo-spectral code and concluded that the HI proposed by RG is  a consequence of particular initial conditions and irrelevant compared to the Hall cascade. They considered a similar, but not identical, geometry as RG and used a code limited to employing periodic boundary conditions in all directions.  In this Letter, we solve the non-linear induction equation using a new code {(Pons et al., 2010, in preparation)} based on a divergence-free, conservative, finite-difference scheme that resembles numerical methods widely employed in astrophysical MHD problems to handle situations where shocks are present. It solves problems caused by the transition from a parabolic equation to the hyperbolic regime as the magnetic resistivity vanishes. Since the code is designed by following methods developed to solve hyperbolic equations, which are similar to standard techniques used in perfect fluid hydrodynamics and ideal MHD \\citep{Anton2006}, it  is able to deal with zero physical resistivity, therefore extending all previous studies to the limit of arbitrary high magnetic Reynolds numbers. Our numerical scheme leads to a fourth order numerical hyper-resistivity, in a similar way to other studies on the basis of which we developed our own algorithm \\citep{OD2006,Toth2008}. We performed a number of tests including the purely diffusive case and the exactly non-resistive case (e.g., propagation of whistler waves on a homogeneous background). In all cases  where analytical solutions are available, the correct modes, decay times, and wave propagation speeds were reproduced.  This allows us to study the same model originally analyzed by RG and answer open questions about the occurrence of the HI and its interplay with the so-called  {\\it Hall cascade}. ",
        "conclusions": "We have proven that there are cases in which the HI dominates the early evolution of the magnetic field and completely reorganizes the field structure in a few Hall times. It is not always subsumed by the Hall cascade, as argued by other authors. We have found that the finite size of the computational domain, and the use of different boundary conditions significantly affect  the outcome. Since the HI is not entirely a scale-free problem and, when a slab geometry is considered, the relative sizes of the length and thickness of the slab make a difference, our findings raise new and interesting questions. In simple terms, a thin slab is  easier subject to the HI than a thicker one but, what happens in a real NS crust ?  The crust of a NS is not an infinite slab, but a finite-size,  curved domain. To gain insight into realistic scenarios, one can mimic the axially symmetric case with a plane parallel model of a longitudinal size $\\approx70$ km and a thickness of 1 km (but with reflecting boundary conditions instead of periodic boundary conditions at the poles as used in this work). This resembles one of our models that has the largest grid size and therefore for which the HI occurs. By analogy with this toy model, one may be tempted to claim that the HI is truly relevant and will affect the evolution of the magnetic field in NSs. When we compare the structures and evolutions of the models shown in this paper with those of \\cite{PonsGeppert2007}, both the qualitative shape of the field and the typical evolution are very similar, despite important differences in the physical model. Perhaps the octupolar structures that survive to late times in the realistic NS crust simulations of \\cite{PonsGeppert2007} originate from the HI. However, caution is required in accepting this interpretation. Given the sensitivity of the results to details such as the size of the grid (which allows only a finite number of unstable modes) and the existence of non-linear couplings between modes, it becomes necessary to perform multidimensional simulations in a spherical shell with a realistic crust model before we can make any claim about the relevance of the HI in NSs. In addition, toy models that have constant density and cubic boxes with periodic boundary conditions, such as those studied in WH, cannot be used to argue against the relevance of the HI in NSs, which is a global, more complex problem, that also depends on the true dimensions and the nature of the physical boundaries. As we have proven, the numerical truncation of some modes due to the computational domains being too small can suppress the instability. An extension of our present code to spherical coordinates and to contain a realistic NS model is under development. Very preliminary results seem to indicate that an active Hall stage exists, but only a more detailed analysis will provide a definitive answer."
    },
    "1004/1004.3098_arXiv.txt": {
        "abstract": "We study the spectral evolution of PWNe taking into account the energy injected when they are young. We model the evolution of the magnetic field inside a uniformly expanding PWN. Considering time dependent injection from the pulsar and coolings by radiative and adiabatic losses, we solve the evolution of the particle distribution function. The model is calibrated by fitting the calculated spectrum to the observations of the Crab Nebula at an age of a thousand years. The spectral evolution of the Crab Nebula in our model shows that the flux ratio of TeV $\\gamma$-rays to X-rays increases with time, which implies that old PWNe are faint in X-rays, but not in TeV $\\gamma$-rays. The increase of this ratio is because the magnetic field decreases with time and is not because the X-ray emitting particles are cooled more rapidly than the TeV $\\gamma$-ray emitting particles. Our spectral evolution model matches the observed rate of the radio flux decrease of the Crab Nebula. This result implies that our magnetic field evolution model is close to the reality. Finally, from the viewpoint of the spectral evolution, only a small fraction of the injected energy from the Crab Pulsar needs to go to the magnetic field, which is consistent with previous studies. ",
        "introduction": "A pulsar releases its rotational energy as a relativistic magnetized outflow called a pulsar wind. The pulsar wind collides with the surrounding supernova ejecta, forms the termination shock, and creates a PWN \\citep{kc84a}. The acceleration of the pulsar wind particles occurs at the termination shock and the PWN consists of the magnetic field and the ultrarelativistic particles \\citep{rg74, kc84b}. Created PWN emits photons ranging from radio to TeV $\\gamma$-rays via the synchrotron radiation and the inverse Compton scattering. Current status of the theoretical models as well as the observational confrontations is reviewed by \\citet{gs06}.  The Crab Nebula is one of the best studied PWN at almost all observable wavelengths including its central pulsar, called the Crab Pulsar. Many studies have been made to explain the observed properties of the Crab Nebula as a typical PWN. \\citet{kc84a} studied the spatial structure of the Crab Nebula, assuming that it is a steady state object (KC model). They found that the magnetization parameter $\\sigma$, the ratio of the electromagnetic energy flux to the particle energy flux just upstream the termination shock, must be as small as 0.003 to explain the observed dynamical properties of the Crab Nebula. \\citet{aa96} succeeded to reconstruct the current observed broadband spectrum of the Crab Nebula by the use of the KC model.  The KC model is not fit to consider the evolution because it is a steady state model. However, it is important to consider the spectral evolution of the Crab Nebula. To explain the flux decrease rates of the Crab Nebula in radio and optical wavelengths \\citep[e.g.][]{ar85, v07, s03}, we need to consider the spectral evolution. Moreover, we need to understand the spectral evolution of PWNe in general. Recent observations have found many PWNe which have a variety of characteristics in terms of age, expansion velocity, morphology, radiation spectrum and others. Some of these characteristics may be understood by the spectral evolution of PWNe.  For example, several old PWNe which are faint in X-rays have been detected in TeV $\\gamma$-rays and \\citet{dd08} discussed the possibility that some of the TeV $\\gamma$-ray sources without an X-ray counterpart may be old PWNe. When we study the broadband spectrum of old PWNe, we need to take into account the evolution of the injected energy from the pulsar, such as the magnetic field and the relativistic particles. For old PWNe, most of the energy inside the PWNe is the energy which was injected when they were young because the spin-down power of the pulsar decreases fast with time.  Several spectral evolution models of PWNe have been studied \\citep[e.g.][]{zet08, det08, get09}. Although they included the evolution of the energy injection from the pulsar, several issues still remain to be further clarified. These studies assume different evolution model of the particle distribution inside the PWNe. \\citet{zet08} considered the escape of the particles from the PWN which has a fixed volume. Because a PWN is expanding while it is confined in the expanding supernova ejecta, it is more realistic to regard that particles rather suffer from an adiabatic loss than they escape from the nebula. \\citet{zet08} and \\citet{det08} used the conversion efficiency of the spin-down power to the particle energy $\\sim 10-50 \\%$ to normalize the observed flux. In their model, the magnetic field inside the PWN does not relate to the spin-down power and the remaining $50-90 \\%$ of the spin-down power is not explicitly considered. \\citet{aa96} has a similar problem of the flux normalization. \\citet{get09} did not compare their calculated spectrum with the observations because they were interested mainly in the dynamical evolution of the PWN.  In this paper, we revisit a spectral evolution model of PWNe, paying attention to the issues mentioned above. We do not take into account effects of the spatial structure of the PWN because it is somewhat costly and too detailed to discuss the whole radiation spectrum. In our model, the PWN is simply treated as an expanding uniform sphere. The energy inside the PWN is injected from the pulsar spin-down energy, which is divided between the magnetic field and the relativistic particles with a constant ratio. Our simple model can describe the observed basic features. Of course, many details are not treated in this simplified model, such as the filamentary structures and the spatial variations of photon indices.  We study the spectral evolution of the Crab Nebula as the first application of our model. The Crab Nebula can be used as a calibrator of our model when it will be applied to other PWNe in future. In section \\ref{model}, we describe our model of the PWN evolution. In section \\ref{crab}, we apply this model to the Crab Nebula. In section \\ref{discussion_conclusion}, discussions and conclusions are made. ",
        "conclusions": "The evolutions of the magnetic field and the particle distribution determine the spectral evolution of a PWN. The evolution of the particle distribution is affected by the assumptions of the magnetic field evolution model, uniformity of the PWN, particle injection spectrum, and expansion evolution of the PWN. We discuss about the effects of these assumptions which are made in our model. Finally, we summarize the conclusions of this paper.  \\subsection{Discussion}\\label{discussion} Our model of the magnetic field evolution is somewhat ad hoc. As discussed in section \\ref{mag_evol}, however, the time dependence of the magnetic field strength $B \\propto t^{-a}$ for $t > \\tau_0$ seems to be in a range of $1.5 \\le a \\le 2.0$ from other theoretical considerations and we adopt $a = 1.5$. Moreover, because our result of the radio flux decrease of the Crab Nebula is almost consistent with the observation, our model of the magnetic field evolution can be near the truth.  For the assumption of the uniform PWN, many non-uniform structures have been observed, such as the filamentary structures and the spatial variations of photon indices. However, for the calculation of the total spectrum of the PWN, the energetics of the PWN is important in the lowest order. We consider that the assumption of the uniform PWN is reasonable for the calculation of the total spectral evolution of the PWN.  For the injection spectrum of the particle distribution, the acceleration of the particles is an unsolved problem and we adopt the broken power-law injection. It should be noted that one of the important conclusions in our study that old PWNe can be observed as $\\gamma$-ray sources with no or weak X-ray counterpart is not affected by the broken power-law assumption. This is because low energy particles do not contribute to X-ray and high energy $\\gamma$-ray emissions.  The use of the time independent parameters $\\gamma_{\\rm min}$, $\\gamma_{\\rm b}$ and $\\gamma_{\\rm max}$ can be improved as time dependent parameters, because the physical condition of pulsar wind termination shock may be change with the decrease of the spin-down power of the pulsar. Considering the time dependence of $\\gamma_{\\rm max}(t)$, in \\citet{vd06}, they use the condition $r_{\\rm L}(t) < 0.5 r_{\\rm s}$ with time dependent magnetic field, where $r_{\\rm L}$ is the electrons' Larmor radius and $r_{\\rm s}$ is the radius of the termination shock. As discussed in section \\ref{other}, our model satisfy this condition. Both $\\gamma_{\\rm min}$ and $\\gamma_{\\rm b}$ are important parameters because these may include the information about the pulsar magnetosphere and the pulsar wind. However, the time dependences of them are uncertain. For simplicity, we used all of them as the time independent parameters in the present paper.  Constant velocity expansion is a good assumption for young PWNe, although the expansion of the PWN should be calculated by taking account for the environment of the PWN \\citep[e.g.][]{get09}. In our model, the magnetic field evolution explicitly depends on the expansion velocity $v_{\\rm PWN}$ (see equation (\\ref{eq8})). To understand a little more about the effects of the expansion evolution, we study how the Crab Nebula would be observed in the context of the constant velocity expansion. The Crab Nebula is one of the sources without observable SNR shell. It may be because the surrounding interstellar medium is less dense than other PWNe with observable SNR shell. If the Crab Nebula were in the different surroundings, the expansion velocity would also be affected. That is, if it were in a less or more dense surroundings, the expansion velocity would be more rapid or slower.  An example of a twice rapid expansion is shown in Figure \\ref{f6}. All the parameters except the expansion velocity are the same as in Figure \\ref{f1}. An example of a half velocity expansion is shown in Figure \\ref{f7} and all parameters except for the expansion velocity are the same as in Figure \\ref{f1}. Note that both of them are hypothetical PWNe, not the Crab Nebula itself. For the spectrum of the rapid expansion case, the absolute value of the flux becomes smaller and the flux ratio of the inverse Compton scattering to the synchrotron radiation is larger than the real Crab Nebula shown in Figure \\ref{f1} and for the spectrum of the slow expansion case vice versa. Comparing the particle distribution in Figure \\ref{f6} with that in Figure \\ref{f7}, the low energy particles take the same distribution, but the high energy particle distribution in Figure \\ref{f7} is steeper than that in Figure \\ref{f6}. These spectral behaviors against the expansion velocity can be understood from equations (\\ref{eq8}), (\\ref{eq14}) and (\\ref{eq16}). In our model, the magnetic field becomes smaller when the expansion velocity becomes larger. This leads to the difference in the absolute flux and the synchrotron cooling which changes the high energy particle distribution. On the other hand, because the adiabatic cooling does not depend on the absolute value of the expansion velocity, the low energy particle distribution does not change.  Lastly, in \\citet{aa96}, they included the infrared photons and the starlight for the target photon fields of the inverse Compton scattering. Although these soft photons can significantly contribute to the $\\gamma$-ray flux of other PWNe \\citep[e.g.][]{pet06}, this is not the case of the Crab Nebula. Because the Crab Nebula is located far away from the galactic center ($\\sim 10 \\rm kpc$) and galactic plane ($\\sim 200 \\rm pc$), the energy density of these soft photon fields is less than the solar neighborhood ($\\sim 8 \\rm kpc$ from the galactic center). Even if we assume that the energy density of the infrared photons and the starlight is the same as the solar neighborhood, the inverse Compton scattering off these photon fields contributes less than 30 \\% of the current total $\\gamma$-ray flux. Note that it also does not much affect the $\\gamma$-ray spectrum, when the SSC flux decreases if the energy density of these soft photon fields is less than the half that in the solar neighborhood since inverse Compton scattering off the CMB dominates there.   \\subsection{Conclusions}\\label{conclusion} In this paper, we built a model of the spectral evolution of PWNe and applied to the Crab Nebula as a calibrator of our model. We solved the equation for the particle distribution function considering adiabatic and radiative losses with a simple model of magnetic field evolution.  The flux decrease of the $\\gamma$-rays is more moderate than radio to X-rays, because the magnetic field decreases rapidly, which implies that old PWNe can be observed as $\\gamma$-ray sources with no or weak X-ray counterpart. Although \\citet{dd08} obtained the same result but for a different reason that the X-ray emitting particles are cooled more rapidly than TeV $\\gamma$-ray emitting particles.  The current observed spectrum of the Crab Nebula is reconstructed when the fraction parameter has a small value $\\eta = 0.005$. This is consistent with the prediction of the magnetization parameter $\\sigma \\ll 1$ obtained by \\citet{kc84a}. They obtained $\\sigma \\ll 1$ from the viewpoint of the current dynamical structure of the Crab Nebula, while we determine $\\eta \\ll 1$ from the viewpoint of the spectral evolution.  The smaller value of the current magnetic field $B_{\\rm now} = 85 \\mu \\rm G$ is needed to reconstruct the observed spectrum of the Crab Nebula. This is consistent with that of \\citet{vet08} for the spatially averaged magnetic field strength $\\sim 100 \\mu \\rm G$ from their relativistic MHD simulation, but smaller than $\\sim 300 \\mu \\rm G$ in most of other papers \\citep[e.g.][]{aa96}.  Our model can predict the spectral evolution of the Crab Nebula, and the observed flux decrease of the Crab Nebula at radio wavelengths can be explained by our model. This conclusion does not depend on the assumption of the broken power-law injection, and gives the validity of the our magnetic field evolution model. The observed flux decrease of the Crab Nebula in optical wavelengths is somewhat larger than our model, but the trend that the decreasing rate increases with frequency matches observations.  The minimum energy $\\gamma_{\\rm min}$ is related to the pair production multiplicity in the pulsar magnetosphere, since low energy particles are assumed to be injected in the same way as high energy particles in our model. Our result of the minimum energy $\\gamma_{\\rm min} = 1.0 \\times 10^2$ and the low energy power-law index $p_1 = 1.5$ means that the multiplicity $\\kappa \\sim 10^6$ is necessarily larger than other models which adopt a separate origin of low energy particles."
    },
    "1004/1004.5578.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) {We performed a detailed study of maser and radio continuum emission toward the high-mass star-forming region \\G23. This study aims at improving our knowledge of the high-mass star-forming process by comparing the gas kinematics near a newly born young stellar object (YSO), analyzed through  high spatial resolution maser data, with the large-scale environment of its native hot molecular core (HMC), identified in previous interferometric observations of thermal continuum and molecular lines.} % methods heading (mandatory) {\\textbf{Using the VLBA and the EVN arrays, we conducted phase-referenced observations of the three most powerful maser species in \\G23: H$_2$O at 22.2~GHz (4 epochs), CH$_3$OH at 6.7~GHz (3 epochs), and OH at 1.665~GHz (1 epoch). In addition, we performed high-resolution ($\\geq0\\farcs1$), high-sensitivity ($<0.1$~mJy) VLA observations  of the radio continuum emission from the HMC at 1.3 and 3.6~cm.}} % results heading (mandatory) {We have detected H$_2$O, CH$_3$OH, and OH maser emission clustered within 2000 AU from the center of a flattened HMC, oriented SE-NW, from which emerges a massive $^{12}$CO outflow, elongated NE-SW, extended up to the pc-scale. Although the three maser species show a clearly different spatial and velocity distribution and sample distinct environments around the massive YSO, the spatial symmetry and velocity field of each maser specie can be explained in terms of expansion from a common center, which possibly denotes the position of the YSO driving the maser motion. Water masers trace both a fast shock (up to 50~\\kms) closer to the YSO, powered by a wide-angle wind, and a slower (20~\\kms) bipolar jet, at the base of the large-scale outflow. Since the compact free-free emission is found offset from the putative location of the YSO along a direction consistent with that of the maser jet axis, we interpret the radio continuum in terms of a thermal jet. The velocity field  of methanol masers can be explained in terms of a composition of slow (4~\\kms\\ in amplitude) motions of radial expansion and rotation about an axis approximately parallel to the maser jet. Finally, the distribution of line of sight velocities of the hydroxyl masers suggests that they can trace gas less dense ($\\rm n_{H_{2}} \\le$~10$^6$~cm$^{-3}$) and more distant from the YSO than that traced by the water and methanol masers, which is expanding toward the observer. A few pairs of OH masers, with different circular polarization, are well aligned in position on the sky and we interpret them as Zeeman pairs. From Zeeman splitting, the derived typical values of the magnetic field are of a few mG. } % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "Hot, dense, molecular cores (HMCs) of dust and gas located in giant molecular clouds (GMCs) are the birth sites of high-mass young stellar objects (YSOs). Ultraviolet light from the newly formed stars ionizes the surrounding gas, creating a compact H~\\textsc{ii} region (e.g., \\citealt{Hoare2007}), excites several molecular transitions (e.g., \\citealt{Sridharan2002}), and heats the dust, which reradiates the energy in the far infrared (FIR) band (e.g., \\citealt{Molinari2008}). At the radio wavelengths, maser emissions of several molecular transitions are observed during the earliest evolutionary phases of  high-mass (proto)stars (e.g., \\citealt{Szymczak2005}), even before the appearance of an ultra compact H~\\textsc{ii} region (UCH~\\textsc{ii}). From an observational point of view, however, the study of high-mass star-forming regions (HMSFRs) is challenging because of three main limitations: the dusty environment is obscured at optical/NIR frequencies; time scales of massive star formation are short, and hence, the chance of observing massive YSOs is small; massive YSOs are rare and statistically found far away from the observer, at distances of a few kpc, clustered in tight associations.  Since a few years, we have started an observational campaign to study the high-mass star-forming process by comparing interferometric thermal data, tracing the large-scale environment (e.g., \\citealt{Codella1997,Furuya2008}), with Very Long Baseline Interferometry (VLBI) measurements of maser transitions, tracing the inner kinematics of the (proto)stellar cocoon. Details about our VLBI observational program to measure molecular masers in HMSFRs are extensively presented in \\citet[][hereafter Paper~\\textrm{I}]{Sanna2010}. The present paper focuses on our observations and analysis of the HMSFR \\object{G23.01$-$0.41}.  In Sect.~2, we provide an up-to-date review of previous observations toward the HMSFR \\G23. Section~3 describes our VLBI observations of the 22.2~GHz water (H$_2$O), 6.7~GHz methanol (CH$_3$OH) and 1.665~GHz hydroxyl (OH) maser transitions, together with the new Very Large Array (VLA) observations of the radio continuum emission at 1.3 and 3.6~cm. In Sect.~4, we illustrate the spatial morphology, kinematics, and time-variability of individual maser species, and present results from our VLA observations, constraining the properties of the radio continuum observed associated with the masers. Section~5 discusses the spatial association of the maser species and their overall kinematics, and draws a comprehensive picture of the phenomena observed in the HMSFR \\G23 on angular scales from a few mas to tens of arcsec.  Main conclusions are summarized in Sect.~6. ",
        "conclusions": "Using the VLBI technique, we observed the HMSFR \\G23 in the three most powerful maser transitions: 22.2~GHz H$_2$O, 6.7~GHz CH$_3$OH, and 1.665~GHz OH. The source \\G23 was also observed with the VLA, detecting faint ($\\approx$~mJy) radio continuum emission with 0\\farcs1 resolution towards the center of the HMC. Our main conclusions can be summarized as follows:  \\begin{enumerate}  \\item H$_2$O, CH$_3$OH, and OH maser emissions are distributed within  $\\approx2000$~AU from the center of a HMC powering a massive bipolar outflow. The three maser species, although associated with the same YSO, present different, although complementary, spatial distributions and kinematical properties.  \\item Water masers trace fast (20--50~\\kms) outflows emitted from the YSO. An arc-like structure of water masers superposed on the radio continuum marks a fast shock propagating through dense gas, and is probably driven by a YSO wind. The 1.3~cm continuum source and  the two clusters of water masers, aligned along the NE--SW direction detected at larger distance from the YSO, are likely tracing the jet driving the massive CO outflow observed at larger scales.  \\item Methanol masers present a N--S oriented, funnel-like spatial distribution with red- and blueshifted features located to the south and the north, respectively. Observing 3 different EVN epochs spanning 2~yr, we have measured accurate (relative errors $<30\\%$) proper motions of 6.7~GHz methanol features. The maser transverse velocities range from a few to about 15~\\kms. The pattern of 6.7~GHz maser proper motions can be interpreted in term of a composition of expansion and rotation around a YSO of about 20~M$_{\\odot)}$, with the rotation axis oriented on the sky at similar P.A. to the axis of the jet/outflow system traced by the water masers. It is then plausible that the  CH$_3$OH masers trace the internal portions of the toroid, elongated along the SE--NW direction, observed in the CH$_3$CN and NH$_3$ lines on arcsec scale.  \\item Hydroxyl masers superposed on the radio continuum are probably seen in the foreground and expand outward from the central source tracing a lower density environment than that harboring the methanol and water masers.   \\end{enumerate}  This study demonstrates that multi-epoch VLBI observations of different maser species provide information useful to explore the complex phenomena occurring at distances of tens up to thousands of AU around massive YSOs. The best scientific return from VLBI maser observations can be expected when such observations will be compared to data of new generation (sub)millimeter interferometers (such as ALMA). Clearly, thermal line observations of both outflow(s) and core(s) with an angular resolution closer to that of VLBI maser data will help to solve ambiguities in the interpretation of kinematic structures, thanks to a better sampling of the (proto)stellar environment and its physical properties."
    },
    "1004/1004.0863_arXiv.txt": {
        "abstract": "{The plateau in the abundance of \\element[][7]{Li} in metal-poor stars was initially interpreted as an observational indicator of the primordial lithium abundance. However, this observational value is in disagreement with that deduced from calculations of  Big Bang nucleosynthesis (BBN), when using the Wilkinson microwave anisotropy probe (WMAP) baryon density measurements. One of the most important factors in determining the stellar lithium abundance is the effective temperature. In a previous study by the authors, new effective temperatures ($T_{\\rm eff}$) for sixteen metal-poor halo dwarfs were derived using a local thermodynamic equilibrium (LTE) description of the formation of Fe lines. This new $T_{\\rm eff}$ scale reinforced the discrepancy.} {For six of the stars from our previous study we calculate revised temperatures using a non-local thermodynamic equilibrium (NLTE) approach. These are then used to derive a new mean primordial lithium abundance in an attempt to solve the lithium discrepancy.} {Using the code ${\\rm\\sc MULTI}$ we calculate NLTE corrections to the LTE abundances for the \\ion{Fe}{i} lines measured in the six stars, and determine new $T_{\\rm eff}$'s. We keep other physical parameters, i.e. log $g$, [Fe/H] and $\\xi$, constant at the values calculated in Paper I. With the revised $T_{\\rm eff}$ scale we derive new Li abundances. We compare the NLTE values of $T_{\\rm eff}$ with the photometric temperatures of Ryan et al. (1999, ApJ, 523, 654), the infrared flux method (IRFM) temperatures of Mel\\'{e}ndez {\\&} Ram\\'{i}rez (2004, ApJ, 615, 33), and the Balmer line wing temperatures of Asplund et al. (2006, ApJ, 644, 229).} {We find that our temperatures are hotter than both the Ryan et al. and Asplund et al. temperatures by typically $\\sim$ 110 K - 160 K, but are still cooler than the temperatures of Mel\\'{e}ndez {\\&} Ram\\'{i}rez by typically $\\sim$ 190 K. The temperatures imply a primordial Li abundance of 2.19 dex or 2.21 dex, depending on the magnitude of collisions with hydrogen in the calculations, still well below the value of 2.72 dex inferred from WMAP + BBN. We discuss the effects of collisions on trends of \\element[][7]{Li} abundances with [Fe/H] and $T_{\\rm eff}$, as well as the NLTE effects on the determination of log $g$ through ionization equilibrium, which imply a collisional scaling factor $\\rm S_{H} >$ 1 for collisions between \\ion{Fe}{} and H atoms.} {} ",
        "introduction": "Since its discovery by \\citet{SpiteSpite1982}, many studies of the plateau in lithium in metal-poor dwarfs have been undertaken, e.g. \\citet{Spiteetal1996}, \\citet{Ryanetal2000}, \\citet{MelendezRamirez2004}, \\citet{Bonifacioetal2007} and \\citet{Aokietal2009}, confirming its existence. Most studies find a comparable Li abundance ($A$(Li)\\footnote{$A$(Li)$\\equiv log_{10} \\left(\\frac{N(\\rm Li)}{N(\\rm H)}\\right)+12.00$} $\\approx$ 2.0 - 2.1 dex)  yet discrepancies still exist, in particular the high value found by \\citet{MelendezRamirez2004} ($A$(Li) = 2.37 dex). However, the biggest discrepancy comes from a comparison of the primordial abundances inferred from observations and that derived from Big Bang Nucleosynthesis (BBN) with the WMAP constraint on the baryon density fraction, $\\Omega_{\\rm B}h^{2}$, which leads to $A$(Li) = 2.72 dex \\citep{Cyburt2008}. This is what has become known as the ``lithium problem''.  Several possibilities have been proposed to explain this discrepancy. Broadly these are: systematic errors in the derived stellar Li abundances; errors in the BBN calculations due to uncertainties in some of the relevant nuclear reaction rates; the destruction of some of the BBN-produced Li prior to the formation of the stars we have observed; the introduction of new physics that may affect BBN \\citep{JedamzikPospelov2009} ; or the removal of Li from the photospheres of the stars through their lifetimes \\citep[see introduction to][Paper I, for more details]{Hosfordetal2009}. The possible explanation under study in this work is that of systematic errors in the effective temperature ($\\textit{T}_{\\rm eff}$) scale for metal-poor stars. The effective temperature is the most important atmospheric parameter affecting the determination of Li abundances. This is due to the high sensitivity of $A$(Li) to $\\textit{T}_{\\rm eff}$, with $\\partial{A}/\\partial{T_{\\rm eff}} {\\sim}$ 0.065 dex per 100 K. One reason for the spread in the observed $A$(Li) is the differences in the $\\textit{T}_{\\rm eff}$ scales used by different authors. For instance,  \\citet{Spiteetal1996} and \\citet{Asplundetal2006} derive a $\\textit{T}_{\\rm eff}$ of 5540 K and 5753 K for the star HD140283, respectively. The scale of \\citet{MelendezRamirez2004} is on average $\\sim$ 200 K hotter than other works. This goes some way to explaining their higher $A$(Li); other factors, such as the model atmospheres with convective overshooting used in their work, may also contribute to the discrepancy. It is important to confirm, or rule out, whether systematic errors in $\\textit{T}_{\\rm eff}$ are the cause of the Li problem, and in doing so address the need for other possible explanations.  In previous work \\citep[- Paper I]{Hosfordetal2009}, we utilised the exponential sensitivity in the Boltzmann distribution to $\\chi/T$, where $\\chi$ is the excitation energy of the lower level of a transition. Using this, we determined $\\textit{T}_{\\rm eff}$'s for eighteen metal-poor stars close to the main-sequence turnoff. This was done by nulling the dependence of $A$(Fe) on $\\chi$ for approx 80 -- 150 \\ion{Fe}{i} lines. Two $\\textit{T}_{\\rm eff}$ scales were generated due to uncertainty in the evolutionary state of some of the stars under study. It was found that our temperatures were in good agreement with those derived by a Balmer line wing method by \\citet{Asplundetal2006} and those derived by photometric techniques by \\citet{Ryanetal1999}. However, our $\\textit{T}_{\\rm eff}$ scale was on average $\\sim$ 250 K cooler than temperatures from the infrared flux method (IRFM) as implemented by \\citet{MelendezRamirez2004}. This is not the case for all work done using the IRFM, the IRFM effective temperatures of \\citet{Alonsoetal1996} are similar to ours, for stars we have in common.  The derived mean abundances in Paper I were $A$(Li) = 2.16 dex assuming main-sequence (MS) membership and $A$(Li) = 2.10 dex assuming sub-giant branch (SGB) membership. For the five stars that have a known evolutionary state, we calculated a mean $A$(Li) = 2.18 dex. It is clear that these values are not high enough to solve the lithium problem. However, the analysis of \\citet{Hosfordetal2009} assumed that the spectrum was formed in local thermodynamic equilibrium (LTE). This is a standard way of calculating spectra, but oversimplifies the radiative transfer problem, and it was acknowledged in \\citet{Hosfordetal2009} that LTE simplification affect those results. Consequently, although it was shown that, within the LTE framework, systematic errors in the $\\textit{T}_{\\rm eff}$ scale are not the cause of the disparity between spectroscopic and BBN+WMAP values for the primordial Li, we also need to assess the impact of non-local thermodynamic equilibrium (NLTE) on the determination of stars effective temperatures. That is the aim of the current work.  This work is not intended to be a full dissection of the methods of NLTE, but rather an application of those more complex (and possibly more accurate) methods to derive a new $\\textit{T}_{\\rm eff}$ scale and to assess their impact on the lithium problem. However, to do this we need to delve, with some depth, into the processes of NLTE line formation, which we do in Sect. 2. This will give some understanding of the complexities and uncertainties that are involved and give the opportunity to make some generalisations on the important aspects that need to be addressed. In Sect. 3 -- 5 we detail our calculations and results, and discuss these further in Sect. 6. ",
        "conclusions": "We have discussed the processes of NLTE line formation of Fe lines. Here, we have shown the challenges posed by such calculations and the uncertainties that still arise, in particular due to the unknown magnitude of \\ion{H}{} collisions. As there is at present no better theoretical or experimental description of the role of H collisions, one obvious next step would be to tie down the value of $\\rm S_{H}$ for metal-poor stars, for example by forcing the equality of {\\sc hipparcos} gravities and those determined by ionization equilibrium by changing $\\rm S_{H}$ \\citep{Kornetal2003}. For this reason we have discussed the effect of NLTE corrections on the ionization equilibrium and the magnitude of the effect on log $g$.  Six of the original program stars from Paper I have been analysed to calculate the effects of NLTE on the $T_{\\rm eff}$ scale derived from \\ion{Fe}{i} lines via excitation equilibrium. We have found that the effect of the correction is to cause an increase in $T_{\\rm eff}$ ranging from 2 K to 150 K for $\\rm S_{H}$ = 0 and 41 K to 122 K for $\\rm S_{H}$ = 1. There is one exception; the star CD$-$33$^{\\circ}$1173 has a negative correction ($-93$ K) for the $\\rm S_{H}$ = 0 case. This may be due to the limited number of Fe lines available for this star, but also emphasises the intricacies of NLTE work which make it difficult to make reliable generalisations.  Our new temperatures have been compared to the photometric temperatures of \\citet{Ryanetal1999}, the IRFM of \\citet{MelendezRamirez2004}, and the Balmer line wing method of \\citet{Asplundetal2006}. We find that the NLTE temperatures are hotter than \\citet{Ryanetal1999} by an average of 132 K for $\\rm S_{H}$ = 0 and 162 K for $\\rm S_{H}$ = 1. Similar results are found when comparing against \\citet{Asplundetal2006} with average differences of 76 K and 110 K for $\\rm S_{H}$ = 0 and 1 respectively. The difference between our temperatures and the \\citet{Asplundetal2006} temperatures may be removed if the Balmer line wing method suffers from NLTE effects \\citep{Barklem2007}, or the effects of granulation are properly described. We find that even with NLTE corrections we are unable to match the high $T_{\\rm eff}$'s of \\citet{MelendezRamirez2004}. However, it has been acknowledged that their temperatures suffer from systematic errors (Mel$\\rm \\acute{e}$ndez 2009 - private communication) and a revision of their temperature scale is under way.  With our new $T_{\\rm eff}$ scale, new Li abundances have been calculated. This has led to an increase of the mean Li abundance from \\citet{Hosfordetal2009} to values of 2.19 dex with a scatter of 0.07 dex and 2.21 dex with a scatter of 0.06 dex for $\\rm S_{H}$ = 0 and 1 respectively, both of which lie well below the value of 2.72 dex inferred from WMAP+BBN \\citep{Cyburt2008}. This has shown that systematic errors in the $T_{\\rm eff}$ scale of metal-poor stars are not the cause for the discrepency."
    },
    "1004/1004.5488_arXiv.txt": {
        "abstract": "{Among multi-planet planetary systems there are a large fraction of resonant systems. Studying the dynamics and formation of these systems can provide valuable informations on processes taking place in protoplanetary disks where the planets are thought have been formed. The recently discovered resonant system HD~60532 is the only confirmed case, in which the central star hosts a pair of giant planets in 3:1 mean motion resonance.} {We intend to provide a physical scenario for the formation of HD~60532, which is consistent with the orbital solutions derived from the radial velocity measurements. Observations indicate that the system is in an antisymmetric configuration, while previous theroretical investigations indicate an asymmetric equilibrium state. The paper aims at answering this discrepancy as well.} {We performed two-dimensional hydrodynamical simulations of thin disks with an embedded pair of massive planets. Additionally, migration and resonant capture are studied by gravitational N-body simulations that apply properly parametrized non-conservative forces.} {Our simulations suggest that the capture into the 3:1 mean motion resonance takes place only for higher planetary masses, thus favouring orbital solutions having relatively smaller inclination ($i=20^{\\circ}$). The system formed by numerical simulations qualitatively show the same behaviour as HD~60532. We also find that the presence of an inner disk (between the inner planet and the star) plays a very important role in determining the final configurations of resonant planetary systems. Its damping effect on the inner planet's eccentricity is responsible for the observed antisymmetric state of HD~60532.} {} ",
        "introduction": "A significant fraction of multi-planet planetary systems contain a pair of giant planets engaged in a mean motion resonance (MMR). These planets are mainly in the low order 2:1 MMR (as GL~876, HD~128311, and HD~73526), but in a few systems higher order resonances have been suggested as well; two planets around 55~Cancri may be in 3:1 MMR, or HD~202206 may host a pair of planets in 5:1 MMR. In the recently discovered system, HD~45364, the giant planets are revolving in 3:2 MMR \\citet{Correiaetal2009A&A}.  The observed orbital solutions and the formation of the majority of resonant systems have been studied thoroughly by many authors. It has been shown, for instance, by \\citet{Kleyetal2004A&A} that a sufficiently slow migration process of two giant planets embedded in an ambient protoplanetary accretion disk ends with either a 3:1 or 2:1 resonant configuration depending on the speed of migration. The formation of resonant systems in 2:1 MMR has been modelled exhaustively by hydrodynamical and N-body simulations as well. The system GJ~876 has been investigated by \\citet{LeeandPeale2002ApJ}, Kley et al (2005), \\citet{Cridaetal2008A&A}, and the systems HD~128311 and HD~73526 by \\citet{SandorandKley2006A&A} and \\citet{Sandoretal2007A&A}, respectively. Most recently, the formation of the system HD~45364 with planets in the 3:2 MMR has been investigated by \\citet{Reinetal2010A&A}. Analytical studies related to the stationary solutions of the 2:1 and 3:1 MMR have been done by \\citet{Beaugeetal2003ApJ} and \\citet{Beaugeetal2006MNRAS}, for instance.  Regarding the 3:1 MMR case, the presence of the planet 55~Cancri-c has already been questioned \\citep{Naefetal2004A&A}. Additionally, recent orbital fits indicate that the planets 55~Cancri-c and 55~Cancri-d are not in a resonant configuration \\citep{Fischeretal2008ApJ}. Thus, until the recent discovery of the two planets in a 3:1 MMR around the F-type star HD~60532 by \\citet{Desortetal2008A&A}, there has been a serious lack of knowledge about the observed behaviour of 3:1 resonant systems. The results of the dynamical study performed by \\citet{Desortetal2008A&A} did not prove without any doubt that the giant planets are in fact in a 3:1 MMR.  The final confirmation of the 3:1 MMR between the giant planets in HD~60532 is given in a recent paper by \\citet{LaskarandCorreia2009A&A}, in which two new orbital fits are provided, slightly improving on the previous fit of \\citet{Desortetal2008A&A} with $i=90^{\\circ}$. Through a detailed stability analysis of different orbit integrations, \\citet{LaskarandCorreia2009A&A} suggest that an inclination of $i=20^{\\circ}$ is the most likely configuration for a co-planar model. By this assumption, the planetary masses are increased by a factor of $1/\\sin(i)\\approx 3$ in comparison to \\citet{Desortetal2008A&A}. In a recent paper, \\citet{2009MNRAS.400.1373L} analyse the excitation of {\\it mutual} inclination for planetary systems driven into resonance by planet-disk interaction. They use a damped $N$-body evolution and find that for stronger eccentricity damping, the mutual inclination is less excited. However, they could not place any constraints on the observed inclination of the system HD~60532.  The very large difference in planetary masses between the two orbital solutions having  $i=20^{\\circ}$ or $90^{\\circ}$ also raises the important question of which orbital solutions are preferred by formation based on the planetary migration scenario. In the present paper we aim at answering this question by performing fully hydrodynamical evolution of planets embedded in the disk. Through this procedure, we obtain realistic migration and eccentricity damping rates that will allow us to determine the most probable final state of the system.  The paper is organized as follows. First, we numerically integrate the orbits of giant planets using the two sets of orbital solutions given by \\citet{LaskarandCorreia2009A&A} as initial conditions. Then by considering different planetary masses, we investigate the possible capture into the 3:1 MMR between the giant planets. After having formed a resonant system, we compare the results of our simulations to the orbital behaviour of giant planets obtained from numerical integrations of Sect. 2. Finally, by performing gravitational three-body numerical integration with dissipative forces for migration, we study how an inner disk influences the behaviour of the system toward its stationary solutions. We demonstrate that the presence of the inner disk determines the final resonant configuration of the system. ",
        "conclusions": "After the discovery of the system HD 60532 \\citep{Desortetal2008A&A}, the thorough dynamical analysis of \\citet{LaskarandCorreia2009A&A} has reliably established the first resonant system that contains two giant planets in a higher order mean motion resonance, here 3:1. Until now, only one system, 55 Cancri, has been proposed as a candidate for a system where two giant planets might be in 3:1 MMR. On the other hand, for 55 Cancri, the existence of the resonant solution could not be confirmed by \\citet{Naefetal2004A&A}, and a new non-resonant solution was later found by \\citet{Fischeretal2008ApJ}. Thus according to our present knowledge, HD~60532 is the only known system containing giant planets in the 3:1 MMR.  Recently, \\citet{LaskarandCorreia2009A&A} has improved the orbital solution found by \\citet{Desortetal2008A&A}. In their new fit that assumes a relatively small inclination between the common orbital plane of the planets and the tangent plane of the sky, $i=20^{\\circ}$, the giant planets have quite high masses: $m_1=3.15 M_J$ and $m_2=7.46 M_J$.  The main path to forming a resonant planetary system is thought to be by convergent migration of planets. For an alternative scenario, formation through a scattering process has been proposed by \\citet{2008ApJ...687L.107R}. Migration is the the result of dissipative forces originating in an ambient accretion disk, which act along with the mutual gravitational forces between the planets and the central star. Being a dissipative process, a sufficiently long-lasting migration process brings the system close to a stationary solution, which corresponds to a periodic solution of the system. The stationary (minimum energy) solutions in the systems yielding 3:1 MMR have been studied by a Hamiltonian approach by \\citet{Beaugeetal2003ApJ}. It has been found that, depending on the ratio of the eccentricities, the resonant angles can exhibit antisymmetric (symmetric anti-aligned) or asymmetric libration. These results have also been confirmed by numerical simulations of \\citet{Kleyetal2004A&A}.  In a first step, we integrated the system HD~60532 numerically using the (observed) initial conditions calculated from Fits I and II of \\citet{LaskarandCorreia2009A&A}. We found that all resonant angles librate and that the system is in apsidal corotation. Since $\\Delta \\omega$ oscillates around $180^{\\circ}$, observationally the system lies in an antisymmetric configuration. However, for the observed mean eccentricities, the Hamiltonian approach would suggest an asymmetric configuration.  To model the formation of the system HD~60532 and find an explanation for this special antisymmetric configuration, we assumed that the system formed through a planet-disk interaction process. We performed a series of fully hydrodynamical simulations with low, intermediate, and high planetary masses $m_i/\\sin i$, assuming $i=90^\\circ$, $30^\\circ$, and $20^\\circ$, respectively. Since the eccentricity of the inner planet oscillates around a moderate mean value $e_1\\sim 0.3$, we assumed that during the migration process an inner disk (between the inner planet and the central star) was present, providing an efficient damping mechanism on the inner planet's eccentricity. We found that the convergent migration through planet-disk interaction, which takes the planets into the 3:1 MMR only occurred for the highest planetary masses. The dynamical behaviour of the resulting resonant planetary system is then indeed very similar to the one obtained from the radial velocity observations. Through our full hydrodynamical simulations, we support the small inclination $i=20^{\\circ}$ of the system as suggested by \\citet{LaskarandCorreia2009A&A}.  To understand why the system does not appear to be in a miminum energy configuration to the 3:1 MMR, we performed a series of dedicated 3-body simulations with additonal forces emulating the effects of the protoplanetary disk. In particular, we studied situations with and without the influence of an inner disk.  We found that it is exactly the effect of the inner disk that distiguishes between an antisymmetric or asymmetric final configuration. Here, its presence is responsible for the system not reaching the minimum energy asymmetric configuration. Nevertheless, we point out that the solutions obtained in the antisymmetric state are indeed stationary solutions, which are stable when the effects of the disk are reduced.  In related earlier works, we had already established that the observed smallness of the mean eccentricities in resonant planetary systems can also be attributed to the effect of an inner disk \\citep{Sandoretal2007A&A,Cridaetal2008A&A}. We conclude that an inner disk during the migration process has directly observable consequences on the post-formation dynamical behaviour of resonant planetary systems in mean motion resonances."
    },
    "1004/1004.4896.txt": {
        "abstract": "We analyze the color gradients (CGs) of $\\sim 50\\,000$ nearby Sloan Digital Sky Survey (SDSS) galaxies estimated by their photometrical parameters (S$\\rm \\acute{e}$rsic index, total magnitude, and effective radius). From synthetic spectral models based on a simplified star formation recipe, we derive the mean spectral properties, and explain the observed radial trends of the color as gradients of the stellar population age and metallicity. Color gradients have been correlated with color, luminosity, size, velocity dispersion, and stellar mass. Distinct behaviours are found for early- and late-type galaxies (ETGs and LTGs), pointing to slightly different physical processes at work in different morphological types and at different mass scales.  In particular, the most massive ETGs ($\\mst \\gsim 10^{11} \\, \\Msun$) have shallow (even flat) CGs in correspondence of shallow (negative) metallicity gradients. In the stellar mass range $10^{10.3-10.5} \\lsim \\mst \\lsim 10^{11} \\, \\Msun$, the metallicity gradients reach their minimum of $\\sim -0.5 \\, \\rm dex^{-1}$. At $\\mst \\sim 10^{10.3-10.5} \\, \\Msun$, color and metallicity gradient slopes suddenly change. They turn out to anti-correlate with the mass, becoming highly positive at the very low masses, the transition from negative to positive occurring at $\\mst \\sim 10^{9-9.5} \\, \\Msun$. These correlations are mirrored by similar trends of CGs with the effective radius and the velocity dispersion. We have also found that age gradients anti-correlate with metallicity gradients, as predicted by hierarchical cosmological simulations for ETGs. On the other side, LTGs have gradients which systematically decrease with mass (and are always more negative than in ETGs), consistently with the expectation from gas infall and supernovae feedback scenarios.  Metallicity is found to be the main driver of the trend of color gradients, especially for LTGs, but age gradients are not negligible and seem to play a significant role too. Owing to the large dataset, we have been able to highlight that older galaxies have systematically shallower age and metallicity gradients than younger ones.  The emerging picture is qualitatively consistent with the predictions from hydrodynamical and chemo-dynamical simulations. In particular, our results for high-mass galaxies are in perfect agreement with predictions based on the merging scenario, while the evolution of LTGs and younger and less massive ETGs seems to be mainly driven by infall and SN feedback. ",
        "introduction": "\\label{sec:intro}    Galaxy formation is a complicated matter which has not yet come to a complete and coherent understanding. The standard cosmological paradigm, the so-called $\\Lambda$CDM, predicts that dark matter (DM) haloes evolve hierarchically since early epochs, with smaller units merging into massive structures (\\citealt{Kauff+93}, \\citealt{deLucia06}, \\citealt{Ruszkowski+09}), which find positive evidences in the strong size evolution in massive galaxies since $z \\sim 3$ (\\citealt{Glazebrook+04}, \\citealt{Daddi+05}, \\citealt{Trujillo+06}). This scenario is at variance with the evidences that today's high mass galaxies formed most of their stars in earlier epochs and over a shorter time interval than low-mass ones. This ``downsizing'' scheme (\\citealt{Cowie96,Gallazzi05,Jimenez+05,Nelan+05,Thomas2005,Treu2005b, Bundy+06,Cimatti+06,Pannella+06,Panter+07,Fontanot+09}) seems at odd with the $\\Lambda$CDM hierarchical growth and seems rather to support a ``monolithic--like'' formation scheme (see e.g. Chiosi \\& Carraro 2002) which broadly predicts an inside-out formation of stars after an early dissipative collapse.  However, it is increasingly clearer that the downsizing is compatible with the hierarchical model if the feedback processes in the mass accretion history of the $\\Lambda$CDM dark haloes are taken into account (see e.g., \\citealt{Neistein+06}, \\citealt{Cattaneo+08}, Conroy \\& Wechsler 2008, \\citealt{Cattaneo+10}). On the one hand, the gas cooling and the shock heating (\\citealt{dek_birn06}, \\citealt{Cattaneo+08}) are the main drivers of the star formation (SF) activity during the hierarchical growth of the DM haloes; on the other side supernovae (SNe), active galactic nuclei (AGN), merging, harassment, strangulation etc., generally inhibit the stellar formation (\\citealt{Dekel86,Recchi01,Pipino02,Scannapieco06,DiMatteo05, deLucia06,Cattaneo+06,Kaviraj07,Schawinski07,Cattaneo+08,AS08, Khalatyan08,Romeo+08,Tortora2009AGN,weinmann09}).    The different processes might rule the star formation at the global galaxy scale, or act at sub-galactic scales (e.g. the nuclear regions vs outskirts) such that they are expected to introduce a gradient of the main stellar properties with the radius that shall leave observational signatures in galaxy colours.  This paper is motivated by the fact that color gradients (CGs) are efficient markers of the stellar properties variations within galaxies, in particular as they mirror the gradients of star ages and metallicities (\\citealt{1972ApJ...176...21S}), although it is not yet fully clear whether metallicity is the main driver of the CGs in normal spheroidal (\\citealt{Saglia+00, TO2000, LaBarbera2002, LaBarbera2003, Ko04, Spolaor09, Rawle+09}) and late-type systems (\\citealt{MacArthur+04}, \\citealt{Taylor+05}), or age plays also an significant role (\\citealt{Saglia+00, LaBarbera2003, MacArthur+04, Spolaor09, Rawle+09}).  CGs are primarily a tool to discriminate the two broad formation scenarios (monolithic vs hierarchical), but more importantly they provide a deeper insight on the different mechanisms ruling the galaxy evolution. As these mechanisms depend on the galaxy mass scale, the widest mass (and luminosity) observational baseline is needed to remark the relative effectiveness of the different physical processes and their correlation with the observed population gradients.   The observational picture is so complex to justify a schematic review of the evidences accumulated so far. The facts are as follows.  1) On average, nearby elliptical and spiral galaxies are bluer outwards (\\citealt{FIH89}, \\citealt{Peletier+90a,Peletier+90b}, \\citealt{BP94}, \\citealt{TO2003}), while dwarfs show mainly redder outskirts (\\citealt{Vader+88}, \\citealt{KD89}, \\citealt{Chaboyer94}, \\citealt{Tully+96}, \\citealt{Jansen+00}). Recently, deeper investigations of later-type galaxies have been done and have highlighted the presence of non-monotonic colour and stellar population gradients in these systems (\\citealt{MacArthur+04}, \\citealt{BTP08}, \\citealt{MacArthur+09}, \\citealt{MS+09}, \\citealt{SB+09}).  2) Spectral line indices (mainly measured in early-type galaxies, ETGs) are somehow a more efficient tracer of the stellar population properties and generally found to change with the radius like CGs (\\citealt{KoAr99}, \\citealt{Kuntschner06}, \\citealt{MacArthur+09}, \\citealt{Rawle+09}).   3) Interpreting CGs in terms of metallicity gradients, in high-mass galaxies typically $d \\log Z / d \\log R \\sim -0.3$, which is shallower than predicted by simulations of dissipative ``monolithic'' collapse (where $-0.5$ or a steeper value is expected, see e.g. \\citealt{Larson74, Larson75}, \\citealt{Carlberg84}, \\citealt{AY87}, \\citealt{Ko04}, see also below).  4) There are contradictory evidences of a correlations of CGs with galaxy mass and luminosity. Earlier studies have claimed a weak correlation with the physical properties of galaxies (e.g., mass, luminosity, etc., \\citealt{Peletier+90a}, \\citealt{Davies+93}, \\citealt{KoAr99}, \\citealt{TO2003}), at variance with the typical monolithic collapse predictions. Recently, a stronger correlation with mass has started to emerge (e.g. \\citealt{Forbes+05}), pointing to a metallicity gradient decreasing with the mass for low-mass galaxies (\\citealt{Spolaor09}, \\citealt{Rawle+09}), accordingly with the monolithic scenario. Instead, high-mass galaxies show too shallow gradients to match the predictions of the monolithic collapse, but compatible with galaxy merging.  This wealth of observational evidences must be confronted with the model predictions coming from the different galaxy formation scenarios.  1) Steep gradients are expected when stars form during strong dissipative (monolithic) collapses in deep potential wells of galaxy cores where the gas is more efficiently retained, with a consequent longer star formation activity and a longer chemical enrichment in the inner than in outer regions (negative metallicity gradients). On top of that, the delayed onset of winds from supernovae, causing a further metal supply in the central regions, would contribute to reinforce the steepness of these gradients (see, e.g. \\citealt{Pipino08}). As these processes are regulated by the galaxy potential depth, which is somehow related to the galaxy mass (and luminosity) \\footnote{It remains to see whether there might be a correlation with the the DM fraction of the systems which is also a function of the galaxy stellar mass (see e.g., \\citealt{Nap05}, \\citealt{Cappellari06}, \\citealt{Conroy08}, \\citealt{Tortora2009}).}, high-mass galaxies are expected to have metallicity gradients which are steeper than lower-mass ones. The latter seem to have almost no gradients (\\citealt{Gibson97}, \\citealt{CC02}, \\citealt{KG03}), although the results here are based on limited galaxy samples.  2) Within the hierarchical picture, merging induce a meshing of stellar populations, and lead to a more uniform metallicity distribution and to shallower CGs (\\citealt{White80}, \\citealt{Ko04}). Strong jets from powerful AGNs can also quench the star formation on the global galaxy scale and flatten colors gradients in the host systems (\\citealt{Tortora2009AGN}).  3) Metallicity and color gradients seem also to depend on the efficiency of the dissipative processes in dark haloes (\\citealt{Hopkins+09a}), with just a weak dependence on the remnant mass (\\citealt{BS99}, \\citealt{Ko04}, \\citealt{Hopkins+09b}), in a way more similar of the observed gradients. However, while merging is crucial to smear out the color radial variation in high-mass galaxies (\\citealt{Ko04}), it seems unimportant for low-mass systems (\\citealt{deLucia06}, \\citealt{Cattaneo+08}). Instead, the energy from stellar winds and supernovae and the effect of dissipative collapse might induce even positive steep gradients (\\citealt{Mori+97}).  ~\\\\  CGs are the most direct observables to investigate the effect of the physical processes (such as merging, AGN, SNe, stellar feedback), which drive the galaxy evolution as a function of the main galaxy parameters: luminosity, mass and central velocity dispersion. In this paper we analyze optical CGs in about $50\\,000$ SDSS local galaxies, spanning a wide range of luminosities and masses. CGs are obtained in an indirect way from the photometrical parameters of the individual galaxies (effective radius, S$\\rm \\acute{e}$rsic index, total magnitude). Using synthetic spectral models, we compute age and metallicity gradients for early- and late-type systems, and analyze the trends with mass. The sample allows us to investigate the distribution of the CGs over an unprecedented baseline of galaxy sizes, luminosities, velocity dispersions, and stellar masses. The observed trends are interpreted by means of quite different physical phenomena at the various mass scales. This approach, while exposed to the uncertainties on the structural parameters (e.g. S$\\rm \\acute{e}$rsic index, $n$, and effective radius, \\Re) has the advantage of dealing with large statistics; furthermore it is exportable to higher redshift galaxy surveys which are generally limited to rest-frame visual bands. In this respect we pay particular care in the check of the consistency of our results with independent analyses. We will mainly concentrate on ETGs, for which a wide collection of results (having an homogenized gradient definition) is available from literature. On the other hand, due to the uncertainties of the structural parameters of late-type galaxies (LTGs), we will take the results on this galaxy sample with the right caution, being well aware that our findings might be only a benchmark for more accurate analyses.  In \\S \\ref{sec:data} we present the data sample and spectral models. In \\S \\ref{sec:results} we show the main results of our analysis, discussing CGs as a function of structural parameters, stellar masses, and stellar properties derived from the fitting procedure. In the same Section, age and metallicity gradients are shown, and discussed within the galaxy formation scenarios in \\S \\ref{sec:disc}. We finally draw some conclusions in \\S \\ref{sec:conclusions}.  In the following, we use a cosmological model with $(\\Omega_{m}, \\, \\Omega_{\\Lambda}, \\, h) = (0.3, \\, 0.7, \\, 0.7)$, where $h = H_{0}/100 \\, \\textrm{km} \\, \\textrm{s}^{-1} \\, \\textrm{Mpc}^{-1}$ (\\citealt{WMAP, WMAP2}), corresponding to a Universe age of $t_{\\rm univ}=13.5 \\, \\rm Gyr$.    \\begin{figure*} \\psfig{file= fig1a.eps, width=0.43\\textwidth} \\psfig{file= fig1b.eps, width=0.43\\textwidth}\\\\ \\psfig{file= fig1c.eps, width=0.43\\textwidth} \\psfig{file= fig1d.eps, width=0.43\\textwidth} \\psfig{file= fig1e.eps, width=0.43\\textwidth} \\psfig{file= fig1f.eps, width=0.43\\textwidth}\\\\ \\caption{Gradients of the $g-i$ colour, as a function of various observed and derived quantities. Medians of binned CGs are plotted together with the $25-75\\%$ quantiles, shown as error bars. {\\it Top-left.} CG as a function of color, black open and gray squares are for colors at $R_{1}$ and $R_{2}$, respectively. {\\it Top-right.} CG as a function of $r$-band magnitude. {\\it Middle-left.} CG as a function of logarithm of $\\sigc$, which is defined as the velocity dispersion within a circular aperture of radius $0.1 \\Re$ (using $r$-band $\\Re$) from the SDSS velocity dispersion $\\sigma_{ap}$, using the relation in \\citet{Jorgensen+95,Jorgensen+96}. {\\it Middle-right.} Gradient as a function of total stellar mass $\\mst$ (assuming a Chabrier IMF), which is the output of our stellar population analysis fitting the total colors. {\\it Bottom-left.} CG as a function of logarithmic of $r$-band \\Re\\ (the results does not depend on the band). {\\it Bottom-right.} CG as a function of i-band S$\\rm \\acute{e}$rsic index $n$; the lines are the the median trends for the S$\\rm \\acute{e}$rsic index in the other bands.}\\label{fig: fig0} \\end{figure*} ",
        "conclusions": "\\label{sec:conclusions}   We have investigated the colour gradients in a sample of $\\sim 50\\,000$ local galaxies from the SDSS as a function of structural parameters, luminosity, and stellar mass. CGs have been found to correlate mainly with luminosity and stellar mass. They have a negative minimum ($\\ggi \\sim -0.2$) at $\\mst \\sim 10^{10.3} \\, \\rm \\Msun$, and increase with the mass for $\\mst \\gsim 10^{10.3} \\, \\rm \\Msun$, the very massive galaxies, $\\mst \\gsim 10^{11} \\, \\rm \\Msun$, having the shallower values ($\\sim -0.1$). On the other mass side, the gradients decrease with mass and become positive for $\\mst \\lsim 10^{9}  \\, \\rm \\Msun$. These trends are mirrored by similar behaviours with luminosity and galaxy size, e.g. the gradients have a minimum at $r \\sim -20$ mag and $\\log \\Re \\sim 0.5$, then increase toward the small and the large end of the parameter distribution, turning to positive values for $\\log \\Re\\lsim 0$ and $r\\gsim -18$ mag. The dependence on velocity dispersion is very loose. A clear dichotomy is also found when looking at the dependence on the S$\\rm \\acute{e}$rsic index $n$ which suggests a distinct behaviour between the ETG and LTG. In fact, these two families mark clear differences in their trends with structural parameters (e.g. mass and \\sigc).  For LTGs, \\ggi\\ monotonically decreases with the mass, with more massive systems having the lowest CGs ($\\sim -0.4$) and less massive systems ($\\mst \\lsim 10^{8-8.5}\\, \\rm \\Msun$) showing even positive gradients. ETGs have negative gradients mildly increasing with mass for $\\log\\mst \\gsim 10^{10.3-10.5}\\, \\rm \\Msun$, which marks the mass scale for the gradient slope inversion (see Table \\ref{tab:slopes_grad_Z_age}). This result is consistent with \\cite{Spolaor09} and reminiscent of the typical mass scale where the star formation and structural parameters in galaxies drastically change (\\citealt{Capaccioli92a, Capaccioli92b}, \\citealt{Kauffmann2003}, \\citealt{Croton06}, \\citealt{Cattaneo+08}). A similar trend is observed when plotting the gradients as a function of effective radius, while a tighter trend with velocity dispersion is evident for ETGs only (Fig. \\ref{fig: fig4}). This was masked when plotting LTGs and ETGs together.  We have used galaxy colors at $0.1\\Re$ and $1\\Re$ in our synthetic spectral models to determine the variation of age and metallicity at these radii. The observed trends of the CGs with mass and \\sigc\\ are correlated with a similar trends for metallicity and age gradients. Despite the large scatter of the data, the strong correlation of metallicity gradients with the central velocity dispersion of ETGs is clear, with a turnoff point at $\\log \\sigc \\sim 2.2 \\, \\rm km/s$  (see Table \\ref{tab:slopes_grad_Z_age}). These results are in very good agreement with a collection of results from literature (\\citealt{Mehlert+03}, \\citealt{Proctor03}, \\citealt{Ogando+05}, \\citealt{Reda+07}, \\citealt{SB+07}, \\citealt{Koleva+09a, Koleva+09}, \\citealt{Spolaor09}, \\citealt{Rawle+09}), in particular for the old galaxies in our sample.  A remarkable result of our analysis is the confirmation that the galaxy (central) age is one of the main drivers of the scatter of the age and metallicity gradients, with the older systems showing generally the shallower gradients with respect the young ones. In Fig. \\ref{fig: fig6}, for instance, we show that in the low mass regime, at fixed \\sigc, older galaxies have on average shallower gradients, consistently with the results in \\cite{Spolaor09}, while systems with late formation or with a recent SF episode have a larger spread (and metallicity gradients down to very low values, $\\sim -0.6$), consistently with findings in \\cite{Koleva+09a,Koleva+09}. The measured scatter might be the consequence of a variety of phenomena affecting the dwarf galaxy evolution, such as (soft) SN feedback, interaction/merging in the high density environment, and star-formation by shell expansion (Fig. 8). On the other mass side, a further factor of the spread of the massive systems gradients might be the randomness of the initial conditions of the mergings (\\citealt{Rawle+09}).   Previous attempts to quantify the correlations of gradients with luminosity, mass, or velocity dispersion have often failed, mainly because of the exiguity of the galaxy sample (e.g. \\citealt{Peletier+90a}, \\citealt{KoAr99}, \\citealt{TO2003}). Only recently a clear correlation of the color and metallicity gradients with mass has been ascertained (see e.g., \\citealt{Forbes+05}), pointing to different trends for high and low mass galaxies (\\citealt{Spolaor09}, \\citealt{Rawle+09}). Our analysis has confirmed and reinforced these results, as it relies on one of the largest local galaxy samples including dwarf, normal, and giant galaxies. Using CGs derived by the structural parameters in B05, we obtained statistically meaningful trends, which allow us to fix the link with physical phenomena as well as to make direct comparisons with predictions from simulations (\\citealt{Mori+97}, \\citealt{BS99}, \\citealt{Kawata01}, \\citealt{KG03}, \\citealt{Ko04}, \\citealt{Hopkins+09a}).  We can finally draw the physical scenario sketched in Fig. \\ref{fig: fin}, resting on the results discussed so far. The formation of LTGs and less massive ETGs is mainly driven by a monolithic-like collapse as they have almost null or positive age gradients and negative metallicity gradients decreasing with the mass. These are in fact well explained by simple gas inflow and feedback from SN and evolved stars, and also reproduced in simulations of dissipative collapse and SN feedback models in Fig. \\ref{fig: fig9} (\\citealt{Larson74, Larson75}, \\citealt{Carlberg84}, \\citealt{AY87}, \\citealt{KG03}). The difference in the magnitude of the gradients between LTGs and lower mass ETGs mainly resides in the effect of the environment as the former live in very low density environments, while the latter stay in massive haloes, together with more massive ETGs. Here they experience strong gravitational interactions producing shallower color and metallicity gradients. The efficiency of such phenomena is strong enough for lower mass ETGs, thus shaping the observed decreasing trend with mass (see e.g. \\citealt{dek_birn06}).  \\begin{figure} \\psfig{file= fig9.eps, width=0.45\\textwidth} \\caption{Color-mass diagram. The contours show the density of data-points with gray scale going from darker (low density of galaxies) to brighter regions (high density). The thin white line sets the separation between RS and BC; the thick vertical one gives the mass scale, $\\mst \\sim 10^{10.3}$, which separates galaxies belonging to the RS in normal and dwarf ETGs, and sets a qualitative upper mass for LTGs. The arrows give information about the efficiency of the phenomena which drive the two-fold trend we discuss in the paper. AGN feedback and merging are important at high mass with an efficiency increasing with mass, while SN feedback and gas inflow drive the galaxy evolution in the less massive side, with an efficiency that is larger at the lowest masses. Galaxies in RS and BC lie in environments with a different density, which is manifested as a difference of the gradients in the two samples.} \\label{fig: fin} \\end{figure}  The most massive ETG systems have shallower gradients (\\citealt{BS99}, \\citealt{Ko04}, \\citealt{Hopkins+09a}), flattening out with the mass due to the increasing intervention of (gas-rich and -poor) merging, tidal interactions, and quasar/radio mode AGN feedback (\\citealt{Capaccioli92a}, \\citealt{FM2000}, \\citealt{deLucia06}, \\citealt{dek_birn06}, \\citealt{Liu+06}, \\citealt{Sijacki+07}, \\citealt{Cattaneo+08}). All the simulations collected in literature fail to reproduce the fine details of the trends we find, suggesting that a full understanding of physical processes involved in the galaxy evolution is still missing.  In future analysis we plan to enlarge the wavelength baseline and also to use line-strength measurements for the stellar model (when available). We will investigate the systematics induced when other synthetic prescriptions are assumed and the dependence of the CGs on the galaxy star formation history. We will derive the $M/L$ gradients, if any, and discuss their impact on the dark matter content of ETGs and correlate the CGs with the DM fraction of these systems as probe of the galaxy potential wells. Finally, the realization of hydrodynamical simulations of jets emitted by AGN would be useful to shape, together with galaxy merging, the color and metallicity gradients for massive galaxies."
    },
    "1004/1004.4926_arXiv.txt": {
        "abstract": "The flux and line shape of the fine-structure transitions of \\NeII\\ and \\NeIII\\ at 12.8 and 15.55\\,$\\mu$m and of the forbidden transitions of \\OI\\ $\\lambda6300$ are calculated for young stellar objects with a range of mass-loss rates and X-ray luminosities using the X-wind model of jets and the associated wide-angle winds. For moderate and high accretion rates, the calculated \\NeII\\ line luminosity is comparable to or much larger than produced in X-ray irradiated disk models. All of the line luminosities correlate well with the main parameter in the X-wind model, the mass-loss rate, and also with the assumed X-ray luminosity --- and with one another. The line shapes of an approaching jet are broad and have strong blue-shifted peaks near the effective terminal velocity of the jet. They serve as a characteristic and testable aspect of jet production of the neon fine-structure lines and the \\OI\\ forbidden transitions. ",
        "introduction": "More than 50 detections of the \\NeII\\ 12.8\\,$\\mu$m line have been made in young stellar objects (YSOs) by the {\\it Spitzer Space Telescope} (Espaillat et al.~2007; Lahuis et al.~2007; Pascucci et al.~2007; Ratzka et al.~2007; Carr \\& Najita~2008) and by ground based telescopes (Herczeg et al.~2007; Najita et al.~2009; van Boekel et al.~2009; Pascucci \\& Sterzik 2009). The neon fine-structure lines were predicted by Glassgold et al. (2007; henceforth GNI07) to arise in disk atmospheres irradiated by X-rays which ionize neon and generate the warm and high-ionization conditions needed to excite the lines. However, the formation of even a single low-mass star occurs in a multi-component system consisting of a collapsing cloud core, a disk, an accretion funnel and outflows arising close to the star. Thus various authors have suggested that the \\NeII\\ 12.8\\,$\\mu$m line might be generated elsewhere in the system and not just by the disk (e.g., Meijerink et al. 2008, henceforth MGN08; Alexander 2008; van Boekel et al.~2009; Najita et al.~2009; Flaccomio et al. 2009; Guedel et al. 2010).  Information on the origin of the \\NeII\\ emission can be obtained from the velocity resolved profiles of the 12.8\\,$\\mu$m line. Najita et al.~(2009) discussed four such measurements: TW Hya (Herczeg et al.~2007); T Tau (van Boekel et al.~2009); AA Tau, and GM Aur (Najita et al.~2009). TW Hya and GM Tau are transitional disks, the former seen almost face-on and the latter at an inclination of $54 \\arcdeg$, whereas AA Tau is a classical T Tauri star (TTS) with a large inclination angle ($75 \\arcdeg$). Pascucci \\& Sterzik (2009) have recently reported four more line shape measurements: TW Hya, T Cha, Sz 73, and CS Cha; T Cha, and CS Cha are also transitional disks. Najita et al.~(2009) concluded that the line profiles for TW Hya, AA Tau, and GM Aur are consistent with the emission arising mainly from a gravitationally bound disk atmosphere. Pascucci \\& Sterzik (2009) measured small blue shifts in three transition disks which they interpreted as arising from low-velocity ($\\sim 10\\, \\kmps$) photo-evaporative outflows, following Alexander (2008). Observations of T~Tau by van Boekel et al.~(2009) show that the \\NeII\\ 12.8\\,$\\mu$m emission in this case arises from more than one location within a complicated and only partially revealed triplet stellar system. T Tau is one of several accreting sources with very strong \\NeII\\ emission (Guedel et al.~2008). Its total \\NeII\\ luminosity is 20 times larger than measured for many revealed TTSs observed by {\\it Spitzer}. The line shape of Sz 73 shows mainly blue-shifted emission extending well beyond $-100\\, \\kmps$, indicative of a strong stellar outflow. Neufeld et al.~(2006) detected the \\NeII\\ 12.8\\,$\\mu$m line in HH objects in the HH 7-11 outflow.  We show here that jets from YSOs, long known to emit the forbidden optical lines of heavy atomic ions, also generate the mid-infrared transitions of \\NeII\\ and \\NeIII. Using the X-wind model, we calculate the luminosities and shapes of these lines and also the forbidden \\OI\\ $\\lambda6300$ transitions for comparison. We obtain results for classical TTSs (Class I and II YSOs) with varying mass-loss rates, and we show that active jets can dominate the observed flux of the \\NeII\\ line. We also show that the line profiles have a distinctive shape that offers observational opportunities to verify the origin of the emission in strong outflow sources and to test the predictions of the X-wind model of jets. ",
        "conclusions": "We have calculated the emission of the \\NeII\\ 12.8 $\\micron$, \\NeIII\\ 15.55 $\\micron$ and \\OI\\ $\\lambda$6300 lines for the X-wind model for the outflows from YSOs. The oxygen lines are well-established diagnostics of jets, whereas the \\NeII\\ 12.8 $\\micron$ line has only recently become a probe of the circumstellar gas of young stars, as described in the Introduction. We have calculated the dependence of the line emission on the two main parameters of the model, the wind mass-loss rate $\\Mw$ and the X-ray luminosity $\\LX$. The mass-loss rate is an intrinsic parameter of X-wind theory, whereas the X-ray luminosity was added by SGSL in modeling the physical properties of the wind. Figure 2(a) expresses our first basic result. It shows that the line luminosities correlate well with the X-wind parameter $\\Mw$. Figure~2(a) also shows that, for mass-loss rates in excess of those characteristic of Class II TTSs, the luminosity of the \\NeII\\ 12.8 $\\micron$ is well above the observed range for this line given by Guedel et al.~(2010). It is interesting that the combination parameter, $\\LX\\Mw$, gives the best correlation of the three shown in Figure~2. Figure 3 gives another example of a good correlation, in this case between the \\OI\\ forbidden transitions and the \\NeII\\ fine-structure line. It suggests that the \\NeII\\ 12.8 $\\micron$ line is an excellent probe of the jets from YSOs, as are of course the well-established \\OI\\ $\\lambda$6300 line.  The theoretical line profiles displayed in Figures~4 and 5 have a distinctive shape. For an approaching one-sided jet, they have a strong blue-shifted peak near the wind terminal velocity that is associated mainly with the inner part of the jet within $\\varpi < 5$\\,AU. There is also a broad shoulder centered around the stellar velocity that arises near the source of the jet and the acceleration region of the outflow. These features reflect some of the unique properties of the X-wind model of outflows from YSOs. Observations to test these predictions for high luminosity sources are feasible for both the \\NeII\\ 12.8 $\\mu$m and the \\OI\\ $\\lambda6300$ lines. High-spectral resolution measurements would be of great interest for understanding the source of the \\NeII\\ and \\OI\\ line emission, the role of X-rays in determining the physical properties of outflows, and the dynamics of jet formation. The blue-shifted peak of the \\NeII\\ line is probably the most important signature of jet emission, because the low-velocity shoulder may be contaminated in some cases by emission from the X-ray irradiated disk, either from the disk proper (MGN08) or from a photo-evaporated wind (Owen et al.~2010; Ercolano \\& Owen 2010).  Line profile measurements have already been carried out for seven low-mass YSOs, as discussed in the Introduction. Four of the seven are transition objects which have evolved disks with substantial gaps in the inner dust distribution (TW Hya, GM Aur, CS Cha, and T Cha), typically on the AU scale. These four objects all show clear signatures of accretion onto the stellar surface, which implies that there is a significant amount of gas present within the inner radius of the dust gap. These are not the most appropriate systems for applying the X-wind model, originally conceived to describe YSOs with strong outflows. Two of the seven YSOs (T Tau and Sz 73) are outflow sources. The simple X-wind model used in this paper is inadequate to deal with the T Tau triple stellar system, which has multiple sources of X-rays and \\NeII\\ line emission. Of all the line shapes reported by Pascucci \\& Sterzik (2009), the broad, blue-shifted line of Sz 73 is closest to the profiles we show in Figures~4 and 5. However, little is known about this YSO, and it would be premature to conclude that it agrees (or disagrees) with our model until more information is obtained, e.g, its X-ray properties determined, and detailed modeling is carried out. This leaves AA Tau, whose line shape has been measured by Najita et al.~(2009) to be broad (FWHM $\\sim 70\\, \\kmps$) and with a red-shifted peak at $+15\\, \\kmps$. There is also emission corresponding to the blue-shifted velocity, but the spectrum is noisy and difficult to characterize. With its large inclination angle, it is possible that higher signal-to-noise observations may show the line shape expected on the basis of our X-wind calculations, especially since AA Tau manifests some evidence for outflow (Bouvier et al.~2003, 2007). On the other hand, the flux of its \\NeII\\ 12.8 $\\micron$ line is at the level obtained for X-ray irradiated disk models (MGN08). Thus any serious model calculation would have to include all the sources of emission, the accretion funnel, the disk and the wind. We would like to recommend that searches be made for \\NeII\\ 12.8 $\\micron$ line emission from strong jets already detected in forbidden transitions from \\OI\\ and other heavy ions using high-resolution MIR spectrometers on large telescopes.  Hollenbach \\& Gorti (2009) have provided an alternative view of the \\NeII\\ 12.8 $\\micron$ line emission for sources with high accretion rates. In their theory, the main source of ionization of neon and the excitation of the 12.8 $\\micron$ line is high-velocity ($\\sim 100\\, \\kmps$) shocks that affect the entire jet. While high-velocity shocks may generate large enough temperatures to ionize neon in localized regions, there is little observational evidence for the strong and continuous shocking of jets. The bulk of the magnetized flow is at most mildly shocked, despite the high terminal speeds reached close to the source. It is these mild shocks that are responsible for heating the outflow in the present model. Hollenbach \\& Gorti (2009) do not calculate line shapes for their high-velocity shock model. Although they are likely to be broad ($\\sim 100\\, \\kmps$), the line shapes can be expected to be different from the ones shown here because the broadening arises downstream along the jet where the strong shocks occur. In contrast, most of the \\NeII\\ and \\NeIII\\ emission in our work is generated within or close to the acceleration region that surrounds the source, so that the line shape takes its final form close to the source, as shown in Figure~4.  Guedel et al.~(2010) have made extensive searches for observational correlations of the \\NeII\\ 12.8 $\\micron$ line luminosity with X-ray luminosity and other properties of YSOs, as have Flaccomio et al.~(2009). They do not find any particular trend for low-accretion systems, where the measured \\NeII\\ luminosities span a range of a factor of 10 or more about a median of approximately $L($\\NeII$) \\approxeq 5 \\times 10^{-6}\\, L_{\\sun}$ (or $ \\log L($\\NeII$) = 28.25$ in cgs units). This value is close to our reference theoretical value, the large open triangle in Figure~2. It is also similar to the value predicted by MGN08 for an X-ray irradiated disk of a typical TTS, although the \\NeII\\ emission from disks has not been calculated for any other case. Calculations for photo-evaporated winds give similar values (Alexander 2008; Ercolano \\& Owen 2010). One explanation for the large number of measurements for both optically thick and transition disks with this level of \\NeII\\ luminosity is that other variables are important, and not just the X-ray luminosity. Schisano et al.~(2010) illustrate this possibility by showing that variations in disk flaring and X-ray spectrum give rise to a scatter in the predicted \\NeII\\ luminosity similar to what is observed. According to our calculations, outflows may also contribute in some cases, but they are not the dominant source of emission for low or even moderate accretion rates and X-ray luminosities. However, as shown in Figure~2, they may dominate for high mass-loss rates and X-ray luminosities. Thus our calculations support the ``bi-modal'' picture of Guedel et al.~(2010) and van Boekel et al.~(2009), especially their suggestion that the \\NeII\\ luminosity of high-accretion sources arises in outflows.  On the basis of a regression analysis, Guedel et al.~(2010) also concluded that the \\NeII\\ luminosities of jet sources correlate, not only with X-ray luminosity, but with accretion rate and mass-loss rate. The significance of these correlations is diminished by the relatively small sample of sources for which the parameters have been measured and by scatter in the data. This is especially true in the case of the mass-loss rate, where only two sources (T Tau N and DG Tau) have exceptionally large \\NeII\\ luminosities. Without relying on these two ``outliers'',\\footnote{T Tau N is an outlier in that it is part of a compact triple system; DG Tau has an unusually small {\\it observed} X-ray luminosity.} the correlation of the \\NeII\\ luminosity with mass-loss rate is not yet definitively established. Such a correlation would be directly relevant to the present calculations because, in X-wind theory, it is the mass-loss rate that directly determines the properties of the outflow, rather than the accretion rate, although the two are related. Mass-loss rates of YSO jets are often determined from bright optical emission lines such as \\OI\\ $\\lambda$6300 (e.g., Hartigan et al.~1994, 1995), but the results for YSOs of the same mass can range over several dex (White \\& Hillenbrand 2004). In light of such uncertainties, it would be better to compare the luminosity of the \\OI\\ $\\lambda$6300 with that of the \\NeII\\ line, as we do in Figure~3. Guedel et al.~(2010) test this correlation with data for sources they identify as optically thick disks without jets, transition disks, and jet sources. A preliminary inspection of the available data for the seven jet sources in their Figure~4 with $L($\\NeII$) > 8 \\times 10^{-6}\\, L_{\\sun}$ suggests a good correlation with our theory. Further observations of the \\NeII\\ and \\OI\\ emission from strong jets including line shape measurements would be extremely useful in this connection.  The correlation predictions in Figure 2 suggest another possible test of our calculations using the \\NeIII\\ 15.55\\,$\\mu$m line, although only a few detections of this line have been reported so far. In Figure~2(a) the predicted but uncertain ratio of the \\NeII\\ 12.8\\,$\\mu$m to the \\NeIII\\ 15.55\\,$\\mu$m line flux is in the range $\\sim 1-10$; the \\NeIII\\ flux is $\\sim 10^{-14}-10^{-13}\\, {\\rm erg}\\,{\\rm cm}^{-2}\\,{\\rm s}^{-1}$. Lahuis et al.~(2007) reported a tentative detection of the \\NeIII\\ line with {\\it Spitzer} for Sz 102, a poorly studied source with a prominent jet, known to emit primarily soft X-rays (Guedel et al.~2010). Lahuis et al.~(2007) give the neon line fluxes as $L($\\NeII$) = 3.6 \\times 10^{-14}\\, {\\rm erg}\\,{\\rm cm}^{-2}\\,{\\rm s}^{-1}$ and $L($\\NeIII$) = 2.3 \\times 10^{-15}\\, {\\rm erg}\\,{\\rm cm}^{-2}\\,{\\rm s}^{-1}$, or a \\NeIII/\\NeII\\ luminosity ratio of $\\sim 1/16$. This small value is consistent with a disk origin for the neon lines, but the measured fluxes are smaller than predicted by MGN08 for the generic TTS disk and the measured X-ray luminosity. Flaccomio et al.~(2009) detected both neon lines from the F7 Class III YSO WL5/GY246 in the $\\rho$ Oph cluster. Although the luminosities are in rough accord with a disk origin, the \\NeIII/\\NeII\\ luminosity ratio of $\\sim 1/4$ is somewhat higher than given by MGN08 and more like the values obtained here for jets. However, we would not expect to find a strong jet in a Class III YSO, so the disk model is favored in this case. A careful determination of upper limits if not detections in {\\it Spitzer} spectra of high luminosity sources of the \\NeIII\\ 15.55\\,$\\mu$m line could provide further checks of our theory. Since {\\it Spitzer} operations have entered the warm phase, further direct detections of this line from space will not be available for some time. However the optical lines of \\NeIII\\ $\\lambda$3869/3967 are accessible from the ground, and they have been observed from H\\,{\\sc ii} regions and other highly-ionized sources. Preliminary calculations indicate that these lines have about the same luminosity as the \\NeIII\\ 15.55\\,$\\mu$m line. We suggest that these lines be used to further probe the ionization state and emission characteristics of jets and disks around low-mass YSOs.  In conclusion, we have calculated the emission of the fine-structure lines of neon and of the forbidden \\OI\\ transitions near 6300\\,\\AA, according to the X-wind model of jets. These lines trace similar aspects of the outflow. They correlate well with the X-wind model parameters and with the combination variable $\\LX\\Mw$, as well as with one another. We have suggested several tests of the calculations, including the measurement of the unique blue-shifted peak in the line profile near the wind terminal velocity. We also suggest that jets may account for the luminous mode in the bimodal picture of \\NeII\\ correlations proposed by Guedel et al.~(2010) and van Boekel et al.~(2009)."
    },
    "1004/1004.3995_arXiv.txt": {
        "abstract": "Recent claims in the literature have suggested that the {\\it WMAP\\/} quadrupole is not primordial in origin, and arises from an aliasing of the much larger dipole field because of incorrect satellite pointing. We attempt to reproduce this result and delineate the key physics leading to the effect. We find that, even if real, the induced quadrupole would be smaller than claimed.  We discuss reasons why the {\\it WMAP\\/} data are unlikely to suffer from this particular systematic effect, including the implications for observations of point sources.  Given this evidence against the reality of the effect, the similarity between the pointing-offset-induced signal and the actual quadrupole then appears to be quite puzzling. However, we find that the effect arises from a convolution between the gradient of the dipole field and anisotropic coverage of the scan direction at each pixel. There is something of a directional conspiracy here -- the dipole signal lies close to the Ecliptic Plane, and its direction, together with the {\\it WMAP\\/} scan strategy, results in a strong coupling to the $Y_{2,\\,-1}$ component in Ecliptic co-ordinates. The dominant strength of this component in the measured quadrupole suggests that one should exercise increased caution in interpreting its estimated amplitude.  The {\\it Planck\\/} satellite has a different scan strategy which does not so directly couple the dipole and quadrupole in this way and will soon provide an independent measurement. ",
        "introduction": "Measurements of Cosmic Microwave Background (CMB) anisotropies from space have laid the foundations of the standard model of cosmology. These observations provide prima facie evidence that the Universe is close to spatially flat with nearly scale-invariant initial density fluctuations. It is remarkable that only a handful of other parameters, specifying the fraction of baryonic and dark matter components, together with the local expansion rate and the redshift of reionization, can fit the data over such a wide range of scales.  Despite the success of the standard model (see e.g~\\cite{scott}), cosmologists have searched for evidence of new physics. Unfortunately this practice is fraught with uncertainty because of the prevalence of a posteriori statistics. One particular avenue of investigation has focused on the CMB anisotropies at large angular scales. A low quadrupole, for example,  was first noted by the Cosmic Background Explorer ({\\it COBE})~\\cite{Smoot:1992td} and subsequently confirmed by the Wilkinson Microwave Anisotropy Probe ({\\it WMAP})~\\cite{Bennett:2003ba}. Several other features have been discovered in {\\it WMAP\\/} data (see~\\cite{Bennett:2010jb} and references therein), with subsequent debate about their statistical significance. Hence it is crucial to investigate all potential sources of systematic error which could be important at large angular scales.  Recently, it has been claimed that an important effect has been overlooked in the {\\it WMAP\\/} analysis~\\cite{Liu}. This particular issue relates to a 25.6\\,ms offset between recording the pointing and differential temperature data from the satellite, which translates into an angular error of about $7^\\prime$. As part of the {\\it WMAP\\/} processing pipeline a dipole signal is first removed from the time-ordered data (TOD), and this has a much higher amplitude than the primordial anisotropies. If the pointing is incorrect a residual signal will remain in the TOD, which has been found to induce a quadrupole pattern in the final temperature sky maps. Interestingly, this quadrupole has  similar $a_{2m}$ spherical harmonic coefficients to those of the {\\em primordial} signal measured by WMAP.  In this short article we show why the claimed result cannot be correct. In doing so we uncover the physics behind the coupling of such systematic effects to certain harmonic modes. We also investigate the implications for the {\\em Planck\\/} experiment, which measures absolute rather than differential data, and scans the sky very differently to {\\it WMAP}. ",
        "conclusions": "We have investigated recent claims in the literature that the {\\it WMAP\\/} quadrupole is systematic in origin, arising from an offset between the recording of pointing and temperature data from the satellite. Due to the {\\it WMAP\\/} scan pattern coupling with the direction of the dipole field, this effect results in a strong $Y_{2\\,,-1}$ mode in Ecliptic co-ordinates, which happens to be similar in phase to the actual quadrupole.  We find that the size of the effect for an offset of $25.6\\,$ms cannot be large enough to match the observed quadrupole.  We have also described reasons why the claimed effect is unlikely to be in the {\\it WMAP\\/} data, in addition to the {\\it WMAP\\/} team insisting that any error in their timing  could not be nearly as large as claimed.  Nevertheless, because of the similarity between the induced and primordial quadrupole signals, one should exercise increased caution in interpreting the amplitude of the primordial component.  {\\it Planck\\/} has a significantly different scan strategy and will soon provide an independent measurement of the quadrupole."
    },
    "1004/1004.2488_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro}  The formation of large-scale structures in the universe is characterized by two fundamental scales:\\footnote{In a universe filled with a combination of different matter components there are additional scales associated with the horizon at the time of transition between the different eras; {\\it e.g.}~in our universe the horizon at matter-radiation equality defines the scale $k_{\\rm eq}^{-1} \\equiv {\\cal H}^{-1}(\\eta_{\\rm eq})$. This complication is of little consequence for the arguments made in this section, but will of course be taken into account in the remainder of the paper.} \\begin{enumerate} \\item[i)] the (comoving) {\\it Hubble scale} ${\\cal H}^{-1}(\\eta)$ defines the extent of the observable universe at any given time $\\eta$ and limits the range over which interactions can influence the evolution of large-scale perturbations; \\item[ii)] the {\\it non-linear scale} $k_{\\rm NL}^{-1}(\\eta)$ describes the size of structures whose density contrast $\\delta(\\eta, \\x) \\equiv [\\rho(\\eta, \\x)/\\bar \\rho(\\eta) - 1]$ exceeds unity. \\end{enumerate} The large hierarchy between these two scales, $\\varepsilon \\equiv k^{-1}_{\\rm NL}/{{\\cal H}^{-1}} \\ll 1$, is of fundamental importance for cosmology. It is this hierarchy of scales that is responsible for the success of linear perturbation theory: the most important features of the anisotropies of the cosmic microwave background (CMB) and the large-scale structure (LSS) observed in galaxy surveys are accurately described by linear perturbations around a homogeneous Friedmann-Robertson-Walker (FRW) background. However, with the advance of observations the study of small non-linear corrections to the long-wavelength dynamics is becoming more and more relevant.  \\vskip 6pt \\noindent {\\sl UV-IR coupling in cosmology.} \\hskip 8pt While for linearized perturbations different Fourier modes evolve independently, at the non-linear level two short-wavelength (UV) perturbations can couple to produce a long-wavelength (IR) perturbation, {\\it i.e.}~starting at quadratic order Fourier modes don't evolve independently. Beyond linear perturbation theory short-wavelength perturbations can therefore, in principle, affect the evolution of the long-distance universe. In particular, small-scale, non-linear fluctuations can give subtle backreaction effects both on the evolution of the background spacetime and the dynamics of long-wavelength perturbations. In fact, it has been claimed that the renormalization of the background from small-scale structures can be large enough to explain the acceleration of the universe without the need for dark energy~\\cite{Kolb:2005da}.  In addition, the effect of small-scale non-linearities on the evolution of superhorizon perturbations from inflation provides an interesting case study~\\cite{Boubekeur:2008kn}. These effects and others will be explored in this paper.  We expect corrections to the linear evolution of scales with wavelengths comparable to (or larger than) the Hubble scale to be suppressed by the aforementioned hierarchy between the non-linear scale and the Hubble scale. This decoupling of short-wavelength (high-energy) fluctuations from the long-wavelength (low-energy) theory is of course a common feature of effective field theories~\\cite{Appelquist:1974tg} (see {\\it e.g.}~\\cite{Weinberg:1996kr,Goldberger:2007hy, Burgess:2007pt, Pich:1998xt} for recent reviews). It should therefore be possible to derive a long-wavelength effective theory in which corrections to the linear evolution appear systematically suppressed by powers of $\\varepsilon^2$ (corrections linear in $\\varepsilon$ are forbidden by the isotropy of the background). In this paper we aim to formalize the effective field theory approach as applied to cosmological perturbations. Specifically, we are interested in an effective description of the long-wavelength universe obtained by `integrating out' short-wavelength modes.\\footnote{We define this procedure in detail in \\S\\ref{sec:PerfectFluid}. There we explain that {\\it integrating out short-wavelength fluctuations} amounts to smoothing the equations of motion and taking expectation values of the short-wavelength modes in the presence of long-wavelength perturbations, so that one is left with equations in terms only of the long-wavelength modes.} We will derive the small corrections to the effective theory of long-wavelength perturbations arising from the UV-IR coupling imposed by the non-linearities of the Einstein equations and the matter sources.  \\vskip 6pt \\noindent {\\sl Matter fluctuations and perturbation theory.} \\hskip 8pt When following this logic, one may worry that the density contrast $\\delta$ becomes large below the non-linear scale and that small scales therefore lead to large backreaction effects on the long-wavelength modes. However, while the density contrast indeed becomes large, the spacetime perturbations and the particle velocities remain small (at least outside of black holes). The system is therefore still amenable to perturbation theory if organized in terms of the gravitational potential $\\Phi$ and the average particle velocity $v$ rather than the density perturbation $\\delta$. Such an analysis reveals that very small scales in fact decouple from the large-scale evolution. We believe that this decoupling of short-wavelength non-linearities should even apply in the extreme case that the universe is filled with a gas of black holes. In this case, our perturbative scheme breaks down, but the effective theory can be matched continuously  to the effective theory for the dynamics of black holes by Goldberger and Rothstein~\\cite{Goldberger:2004jt}, making a large backreaction of gravitational non-linearities even in this case unlikely.  Below the virial scale $k_{\\rm vir}^{-1}$ there exist definitive relations between the gravitational potentials $\\Phi$ and the velocities $v$, with large cancellations between the potential and kinetic energies. One of the main results of our work will be a proof that virial scales indeed decouple completely -- {\\it i.e.}~they don't even lead to $\\varepsilon^2$ suppressed contributions to the effective pressure, although they of course lead to a small renormalization of the background density. In other words, the non-linear source terms for the evolution of large-scale modes can be expressed in a form that vanishes identically in the virial limit. We stress that this is more than standard effective field theory decoupling: indeed, according to the latter, the leading long-distance effect of short-distance physics is just a renormalization of the parameters of the long-distance effective theory \\cite{Appelquist:1974tg}. Here instead, we claim that one such parameter -- the effective pressure -- does not even get renormalized in the virial limit. Therefore, our result goes more properly under the name of a `non-renormalization theorem'. This significantly constrains our expectations for backreaction effects of small-scale non-linearities on the long-distance cosmological dynamics. We stress that the decoupling of virialized structures holds at all orders in the post-Newtonian expansion. Therefore it applies equally well to {\\em relativistic} virialized systems, like for instance those containing black holes.  \\vskip 6pt \\noindent {\\sl The effective theory.} \\hskip 8pt Ultimately, our theory will be formulated as an FRW universe with small (quasi-linear) long-wavelength perturbations evolving in the presence of an effective fluid whose properties are determined by non-linear short-wavelength modes. We present the details of the effective theory in \\S\\ref{sec:ImperfectFluid}. The key element of the theory is the effective stress-energy tensor $\\tau_{\\mu \\nu}$ induced by the short-wavelength modes.  Given its importance\\footnote{The non-linear scalar source terms captured by $\\tau_{\\mu \\nu}$ have the following effects: \\begin{enumerate} \\item the generation of vector perturbations $\\omega_i^{(2)}$ \\cite{vectors1, vectors2, matta}; \\item the generation of tensor perturbations $\\chi_{ij}^{(2)}$ \\cite{matta, wands, baum, wands2}; \\item the superhorizon evolution of scalar perturbations $\\dot \\Phi^{(2)} \\ne 0$\\ \\cite{Boubekeur:2008kn} (see \\S\\ref{sec:evolution}); \\item the viscous damping of density perturbations (see \\S\\ref{sec:BAO}). \\end{enumerate} All of these effects are absent in linear perturbation theory. }, we will derive $\\tau_{\\mu \\nu}$ in two different ways:  \\vskip 6pt \\noindent {\\sl Effective stress-energy via (post-)Newton -- constructive approach.} \\hskip 8pt We aim at understanding the effects of the short-scale non-linearities on the background Hubble expansion and on the long-wavelength fluctuations. In \\S\\ref{sec:PT} we will study non-linear cosmological perturbations in a general-relativistic framework. However, we will encounter a number of technical complications that hide to some extent the physical relevance and intuitive nature of our findings. Fortunately, there are two important points emerging from the former discussion that simplify the problem considerably, allowing us to give a quicker yet rigorous derivation. First, the scale at which non-linearities in the perturbation equations become relevant is much smaller than the horizon scale. This allows us to concentrate on subhorizon distances and neglect the general-relativistic effects associated with the {\\em background} FRW expansion. Second, for non-relativistic structures like clusters or galaxies, these non-linearities involve the matter sector only. That is, the short-scale gravitational dynamics is Newtonian to a very good approximation. Of course, the effect that these non-linear  structures then have on the long-scale perturbations is {\\em post}-Newtonian in nature -- since it involves the coupling of gravity to itself -- but given the above considerations the following simplified approach suggests itself: short-scale perturbations evolve according to flat-space Newtonian equations, where all gravitational fields (short-scale as well as long-scale) are encoded in the Newtonian potential $\\Phi$. The total stress-energy tensor $\\tau^{\\mu\\nu}$ of this system is conserved, in the ordinary sense, {\\it i.e.}~$\\partial_\\mu \\tau^{\\mu \\nu} =0$. Indeed, the existence of a conserved stress-energy tensor follows from locality and from invariance under space-time translations, regardless of Lorentz-invariance. The tensor $\\tau^{\\mu\\nu}$ has a gravitational contribution, of order $\\rho \\Phi$, because the gravitational potential energy obviously participates in the stress-energy conservation already in Newtonian physics. We then smooth this stress-energy tensor over some scale $\\Lambda^{-1}$ larger than the typical inhomogeneity scale, and define an effective long-scale stress-energy tensor. The post-Newtonian leap is now to {\\em declare} that the effective stress-energy tensor thus defined is what couples to long-wavelength gravitational fields. However, there is not much freedom in this assumption. Long-wavelength gravitational fields must couple to anything (including themselves) through a locally conserved symmetric tensor~\\cite{Weinberg:1964ew, Weinberg:1965rz}. For any given system, the only tensor with such properties is the (symmetrized) stress-energy tensor, which is unique up to total derivative terms of the form \\cite{Weinberg:1965rz} \\beq \\label{Tmn_ambiguity} \\partial _\\alpha \\partial _\\beta \\Sigma^{[\\alpha \\mu] [\\beta \\nu]} \\; . \\eeq Here, the tensor $\\Sigma$ is symmetric under the exchange of the two index pairs, and antisymmetric within each pair. The addition of a term of the form of (\\ref{Tmn_ambiguity}) to a system's stress-energy tensor is identically conserved, does not affect the associated global charges ({\\it i.e.}, the total four-momentum), and most importantly for our purposes vanishes like two powers of momentum in the small-momentum limit. This ambiguity thus belongs in the class of `higher-derivative' corrections (which we will discuss below), and can be neglected for long wavelengths. We work out the details of this approach in \\S \\ref{sect:Newton}.   \\vskip 6pt \\noindent {\\sl Effective stress-energy via Einstein -- deductive approach.} \\hskip 8pt A second equivalent description of the UV-IR coupling of cosmological fluctuations arises from a simple reorganization of the Einstein equations. First, we decompose the Einstein tensor into a homogeneous background (denoted by overbars) and terms that are linear (L) and non-linear (NL) in the metric perturbations, collectively denoted by $\\delta X(\\eta,  \\x)  \\equiv X(\\eta, \\x) - \\bar X(\\eta)$. The Einstein equations can then be written as \\beq \\bar G_{\\mu \\nu}[\\bar X] + (G_{\\mu \\nu})^{\\rm L}[\\delta X] + (G_{\\mu \\nu})^{\\rm NL}[\\delta X^2] = 8 \\pi G \\, T_{\\mu \\nu}\\, . \\eeq The background equations, $\\bar G_{\\mu \\nu} = 8\\pi G \\, \\bar T_{\\mu \\nu}$, and the linearized Einstein equations, $(G_{\\mu \\nu})^{\\rm L} = 8\\pi G\\, (T_{\\mu \\nu})^{\\rm L}$, are then defined in the standard way. The non-linear Einstein equations can be written in a form that is very similar to the linear equations, \\beq \\label{equ:nonlinear} (G_{\\mu \\nu})^{\\rm L} = 8\\pi G\\, (\\tau_{\\mu \\nu} - \\bar T_{\\mu \\nu})\\, , \\eeq where we defined the {\\it effective stress-energy pseudo-tensor} \\beq \\tau_{\\mu \\nu} \\ \\equiv \\  T_{\\mu \\nu} - \\frac{(G_{\\mu \\nu})^{\\rm NL}}{8\\pi G} \\, . \\eeq The stress-energy pseudo-tensor has a long history in studies of General Relativity ({\\it e.g.}~as an approach to studying gravitational raditation; see the books by Weinberg~\\cite{Weinberg} or Landau and Lifshitz~\\cite{Landau}) and will play a key role in this paper.  \\vskip 6pt \\noindent {\\sl Properties of the effective fluid.} \\hskip 8pt Given the form of $\\tau_{\\mu \\nu}$ in terms of short-scale (high-momentum) fields -- $\\Phi_s$, $v_s$ -- we can analyze its effects on the large-scale modes -- $\\delta_\\ell$, $\\Phi_\\ell$, $v_\\ell$ -- and on the homogeneous background. For the benefit of the impatient (curious) reader, we will now state some of the highlights of that analysis, leaving detailed derivations and explanations to the main text and the appendices: \\begin{enumerate} \\item In the absence of long-wavelength perturbations or on very large (superhorizon) scales, the gravitational small-scale (subhorizon) non-linearities mimic an isotropic fluid whose effective density and pressure simply renormalize the background by terms of order of the velocity dispersion, $\\langle v^2_s \\rangle$. The effective pressure of the fluid is always positive and much too small to significantly affect the background evolution. Moreover, the spatial part of the stress-energy tensor is equal to the second time-derivative of the moment of inertia tensor \\beq [\\tau_{ij} ]_\\Lambda = \\frac{1}{2} \\frac{d^2 I_{ij}}{d \\eta^2}\\, , \\eeq where $[\\dots]_\\Lambda$ denotes spatial averaging over a region of size $\\Lambda^{-1}$, and $I_{ij}$ is the moment of inertia associated with the same region. This shows that virialized structures decouple completely from the effective theory at large scales. The backreaction effects that we capture in our effective treatment therefore {\\it cannot} explain the acceleration of the universe. Finally, the small induced pressure and the associated renormalization of the background explain the apparent superhorizon evolution of primordial curvature perturbation $\\zeta$ \\cite{Boubekeur:2008kn}. After the renormalization of the background is taken into account, $\\zeta$ is indeed constant on superhorizon scales.   \\item The fluid is an imperfect fluid in the sense that the small-scale non-linearities induce dissipative terms and non-negligible anisotropic stress into the evolution of long-wavelength perturbations. At long wavelengths the fluid is characterized by only a few parameters like an equation of state, a sound speed and a viscosity parameter.\\footnote{We point out that the effective viscosity of the fluid leads to a damping in the non-linear evolution of density fluctuations $\\delta_\\ell$ and suggest that this intuitively explains the non-linear broadening of the peak of baryon acoustic oscillations (see \\S\\ref{sec:Application}).} For instance, to leading order, the source term in the Euler equation may be written as \\beq \\frac{k_i k_j}{k^2} \\frac{\\langle [\\tau_{ij}]_\\Lambda \\rangle}{\\bar \\rho} = c_s^2 \\delta_\\ell - c_{\\rm vis}^2 \\frac{k_i v^i_\\ell}{{\\cal H}}\\, , \\eeq where $c_s$ and $c_{\\rm vis}$ are time-dependent coefficients. These parameters can be calibrated by computing the small-scale dynamics with numerical $N$-body simulations. Alternatively, the fluid parameters may simply be retained as free parameters to be measured by fitting predictions of the effective theory to observations. This kind of matching calculation is of course common in effective field theory.  \\item The short-wavelength fluctuations provide a source of noise to the dynamics of the long-wavelength perturbations. Although this statistical noise has a negligible effect on the evolution of the background cosmology, it is not irrelevant in all contexts. For example, the stochastic contribution to the pressure fluctuations can be comparable to the pressure fluctuations induced by long-wavelength density fluctuations $\\delta_\\ell$ for a wide range of scales. However, these fluctuations are uncorrelated at leading order with the amplitude of the long-wavelength modes, which results in a suppression of their importance in averaged quantities such as the power spectrum.  \\end{enumerate}  \\vskip 6pt \\noindent {\\sl An alternative to conventional perturbation theory.} \\hskip 8pt After the fluid parameters are determined from $N$-body simulations of scales with high momenta, $k > \\Lambda$, the effective fluid has small expansion parameters -- $\\{ \\delta_\\ell , \\Phi_\\ell, v_\\ell\\} \\ll 1$ -- allowing for a controlled perturbative expansion at low momenta, $k < \\Lambda$ (see Fig.~\\ref{fig:scales0}). Conceptually, our approach therefore offers a well-defined and controlled treatment of the effects of short-distance non-linearities on the long-wavelength universe. This is to be contrasted with the failure of many cosmological perturbation theory techniques to include the effects of gravitational non-linearities~\\cite{Carlson:2009it}.  \\begin{figure}[h!] \\centering \\includegraphics[width=0.55\\textwidth]{wavenumbers.pdf} \\caption{\\sl Hierarchy of scales and perturbative expansions. In the effective theory loop integrals only contain modes with $k < \\Lambda$, while conventional perturbation theory contains modes with $k \\sim k_{\\rm NL}$ where the perturbative expansion is known to break down (see \\S\\ref{sec:Application}).} \\label{fig:scales0} \\end{figure}  \\vskip 8pt \\noindent \\newpage {\\sl Outline}: \\hskip 8pt The outline of the paper is as follows: In Section~\\ref{sect:Newton} give a bottom-up construction of the effective stress-energy tensor using the Newtonian approximation on subhorizon scales. Alternatively, in Section~\\ref{sec:PT} we explain our basic perturbative approach within General Relativity and present second-order results for metric and matter perturbations in Poisson gauge. In Section~\\ref{sec:PerfectFluid} we derive the gravitational stress-energy pseudo-tensor and define an effective fluid by taking its long-wavelength limit. We show that on very large scales the fluid behaves as an isotropic fluid and gravitational non-linearities only renormalize the background density and pressure. Furthermore, we prove that virial scales decouple completely from the long-wavelength theory. In Section~\\ref{sec:ImperfectFluid} we show that on scales comparable to and smaller than the horizon scale the fluid anisotropic stress is non-negligible and important for the evolution of perturbations. We characterize the properties of this imperfect fluid in detail. In Section~\\ref{sec:Application} we suggest possible applications of our effective theory. Finally, we present our conclusions in Section~\\ref{sec:conclusions}.  A number of appendices contain technical details: Appendix~\\ref{sec:Euler} gives an alternative derivation of the effective fluid properties starting from the Newtonian conservation equations. This derivation provides considerable intuition for the physical origin of the effective fluid. Appendix~\\ref{sec:Einstein} presents more details of second-order cosmological perturbation theory. We collect results in Poisson gauge, cite second-order gauge transformations and discuss the long-wavelength limit of the spacetime. In Appendix~\\ref{sec:FluidReview} we review key elements of dissipative fluid dynamics that are used in Section~\\ref{sec:ImperfectFluid}. Finally, in Appendix~\\ref{sec:estimates} we illustrate some of our ideas by presenting example calculations in perturbation theory.  \\vskip 6pt \\small \\hrule \\vskip 1pt \\hrule \\vskip 4pt \\small Short {\\bf Proofs} and {\\bf Examples} are separated from the main text by horizontal lines. These parts can be omitted without loss of continuity, but they often illuminate the underlying physics. \\vskip 4pt \\hrule \\vskip 1pt \\hrule \\vskip 6pt \\normalsize  Our work builds on a long history of studies in non-linear perturbation theory \\cite{Peebles, Bernardeau:2001qr, Matarrese:1997ay,  Crocce:2005xy, Jeong:2006xd, Shoji:2009gg, Pietroni:2008jx, Matarrese:2007wc, Matsubara:2008wx, Boubekeur:2008kn}, numerical simulations \\cite{Siegel:2005xu, Heitmann:2008eq, Heitmann:2009cu, Widrow:2009ru}, and effective field theory \\cite{Goldberger:2007hy, Burgess:2007pt, Goldberger:2004jt, Cheung:2007st}. Related ideas have appeared recently in Refs.~\\cite{Dominguez:2000dt, Ishibashi:2005sj, Gruzinov:2006nk, Pueblas:2008uv, McDonald:2009hs, Peebles:2009hw, Shoji:2010hm}.  \\vskip 8pt \\noindent {\\sl Notation and conventions}: \\hskip 8pt Except for the Newtonian analysis of \\S \\ref{sect:Newton}, we will work exclusively with conformal time $\\eta$ and in units where the speed of light is set to unity, $c \\equiv 1$. Greek indices $\\mu, \\nu = 0,1,2,3$ are used for four-dimensional spacetime coordinates, while Latin indices $i, j = 1,2,3$ are reserved for spatial coordinates. We will use overdots to indicate derivatives with respect to conformal time. We use commas to denote partial derivaties and semi-colons for covariant derivatives, {\\it i.e.} \\beq (\\dots)_{,\\mu} =\\frac{\\partial}{\\partial x^\\mu} (\\dots) \\equiv \\partial_\\mu (\\dots) \\qquad {\\rm and} \\qquad  (\\dots)_{;\\mu} = \\nabla_\\mu (\\dots) \\, . \\eeq Our Fourier convention will be \\beq \\phi_{\\k} =  \\int_{\\x} e^{-i \\k \\cdot \\x } \\phi(\\x)\\, ,  \\qquad \\phi(\\x) = \\int_{\\k}  e^{i \\k \\cdot \\x } \\phi_{\\k}\\, , \\eeq where we used the notation \\beq \\int_\\x \\equiv \\int d^3 \\x \\, , \\qquad \\int_\\k \\equiv  \\int \\frac{d^3 \\k}{(2\\pi)^3}\\, . \\eeq Reality of $\\phi(\\x)$ demands that $\\phi^*_\\k = \\phi_{-\\k}$. The product of two functions in real space is a convolution in Fourier space \\beq \\int_\\x e^{-i \\k \\cdot \\x} \\phi(\\x) \\psi(\\x) = \\int_\\q \\phi_{\\q} \\psi_{\\k - \\q} \\, . \\eeq We define two forms of the power spectrum, \\beq \\langle \\phi_{\\k}(\\eta) \\phi_{{\\k}'}(\\eta) \\rangle = (2\\pi)^3 \\delta_{\\rm D}({\\k} + {\\k}') P_\\phi(k, \\eta)\\, , \\eeq and \\beq \\Delta^2_\\phi(k, \\eta) \\equiv \\frac{k^3}{2\\pi^2} P_\\phi(k, \\eta)\\, , \\eeq such that $\\langle \\phi^2 \\rangle = \\int d \\ln k\\, \\Delta_{\\phi}^2(k)$. The time-evolution of the gravitational potential on subhorizon scales during the radiation era leads to a momentum-dependent transfer function $T_\\phi(k, \\eta)$, where \\beq \\phi_{\\k}(\\eta) = T_\\phi(k, \\eta) \\phi_\\k(0)\\, . \\eeq We take the initial conditions to be scale-invariant, with an amplitude fixed by cosmic microwave background observations~\\cite{Komatsu:2010fb}, $\\Delta_\\phi^2(k, 0) \\approx 10^{-9}$. ",
        "conclusions": "\\label{sec:conclusions}  Cosmology has made significant progress by studying {\\it linear} perturbations around a homogeneous Friedmann-Robertson-Walker background \\cite{Dodelson:2003ft}. A crucial simplification of the linear treatment is that large scales (IR) are decoupled from small scales (UV), {\\it i.e.}~at linear order different Fourier modes evolve independently. Furthermore, in linear theory there exists a useful classification of the perturbations into independent scalar, vector and tensor modes. Beyond linear theory, the Einstein equations couple the UV and the IR, with small-scale fluctuations providing sources for the formation and evolution of large-scale perturbations. In this paper, we studied this UV-IR coupling of cosmological fluctuations. By integrating out short-wavelength perturbations we derived an effective theory for the long-wavelength universe. We observed that on sufficiently large scales the universe can be described by quasi-linear perturbations evolving in the presence  of an effective fluid whose properties are determined by the small-scale structures. The fluid is somewhat unconventional in the following sense: higher moments in the Boltzmann hierarchy are small not because of sizable interactions like in a conventional fluid, but because they haven't had sufficient time to develop during one Hubble time ${\\cal H}^{-1}$. In the absence of gravity ${\\cal H}^{-1}$ would go to infinity and a macroscopic effective  fluid description would not be applicable at non-zero momentum. It is in this sense that we refer to the fluid as a `gravitational fluid', highlighting the fundamental importance of gravity. On superhorizon scales, the only effect of small-scale structures is to renormalize the background density and pressure by terms of order the velocity dispersion. Moreover, we proved that the virial theorem naturally filters out the contributions from very small scales. On subhorizon scales, dissipative effects like viscosity are induced by the small-scale non-linearities. The imperfect fluid is described by a few parameters like the equation of state, the sound speed and a viscosity parameter. We proposed that this effective description of the long-wavelength universe be used to formulate a well-defined alternative to conventional perturbation theory. At scales larger than some smoothing scale $\\Lambda^{-1}$ the theory is defined as an expansion in small perturbation variables -- $\\delta(k \\ll \\Lambda)$ and $\\theta(k \\ll \\Lambda)$ -- which evolve in a fluid whose physical parameters are determined by a numerical matching calculation. It remains to be quantified if a simple version of this procedure ({\\it e.g.}~a one-loop calculation in the effective theory) will be sufficient to reach the required level of precision in a practical application like baryon acoustic oscillations \\cite{Future}.   \\subsubsection*"
    },
    "1004/1004.0819_arXiv.txt": {
        "abstract": "We present SHARC-2 350\\um data on 20 luminous $z \\sim 2$ starbursts with $S_{1.2{\\rm mm}}$$>$2\\,mJy from the {\\em Spitzer}-selected samples of Lonsdale et al.\\ and Fiolet et al. All the sources were detected, with $S_{350\\mu{\\rm m}}$$>$25\\,mJy for 18 of them. With the data, we determine precise dust temperatures and luminosities for these galaxies using both single-temperature fits and models with power-law mass--temperature distributions. We derive appropriate formulae to use when optical depths are non-negligible. Our models provide an excellent fit to the 6$\\mathum$--2\\,mm measurements of local starbursts. We find characteristic single-component temperatures $T_1$$\\simeq$35.5$\\pm$2.2\\,K and integrated infrared (IR) luminosities around 10$^{12.9\\pm0.1}$\\,L$_{\\odot}$ for the SWIRE-selected sources. Molecular gas masses are estimated at $\\simeq$4$\\times$10$^{10}$\\,$M_{\\odot}$, assuming $\\kappa_{850\\mu{\\rm m}}$=0.15\\,m$^2$\\,kg$^{-1}$ and a submillimeter-selected galaxy (SMG)-like gas-to-dust mass ratio. The best-fit models imply $\\gtrsim$2\\,kpc emission scales. We also note a tight correlation between rest-frame 1.4\\,GHz radio and IR luminosities confirming star formation as the predominant power source. The far-IR properties of our sample are indistinguishable from the purely submillimeter-selected populations from current surveys. We therefore conclude that our original selection criteria, based on mid-IR colors and 24\\um flux densities, provides an effective means for the study of SMGs at $z$$\\sim$1.5--2.5. ",
        "introduction": "Star-forming galaxies release a significant fraction of their energy output at infrared (IR) and submillimeter wavelengths ($\\lambda$$\\sim$5--1000\\um in the rest frame). The luminosities of starbursts, with star formation rates over 100\\,M$_{\\odot}$\\,yr$^{-1}$, are almost exclusively carried at IR wavelengths. Whereas luminous and ultra-luminous infrared galaxies (LIRGs and ULIRGs) are extremely rare in the local universe, they are prevalent at higher redshifts \\citep[e.g.][]{LeFloch2005, Caputi2007}. By studying this population we can learn about the star formation history of the universe, and because the starbursting is triggered by merger events, we can also probe models of structure formation and halo dynamics through the ages.  Around 1\\,mm wavelengths it is possible to find similar starbursts across much of the volume of the universe, because the dimming of radiation from increasing distances is countered by the steeply rising energy spectrum between 2\\,mm and 200$\\mathum$ in the rest frame. Thus, an IR galaxy would produce nearly the same (sub)millimeter flux regardless of its distance in the range of $z \\sim 0.5$--10.  The nearly bias-free selection, together with the fundamental desire for studying starbursts, make submillimeter surveys \\citep{Coppin2006, Weiss2009, Austermann2009, Ivison2009} especially relevant. However, while finding submillimeter-selected galaxies (SMGs) is relatively straightforward, studying them in any detail has proved difficult. This is due to two factors: they are faint at other wavelengths, and there are often multiple possible optical/NIR counterparts \\citep[see][]{Pope2006, Younger2009b} due to the poor spatial resolution of most (sub)millimeter telescopes (typically 10''--30'' FWHM).  As a result, most spectroscopic redshifts have been obtained for SMGs with radio-detected counterparts providing the required positional accuracy. Thus, \\citet{Chapman2003, Chapman2005} and \\citet{Kovacs2006} discovered that most SMGs lie around a median redshift of 2.3, are extraordinarily luminous (10$^{12}$--10$^{13}$\\,L$_\\odot$) and have large molecular gas reservoirs (10$^{10}$--10$^{11}$\\,M$_\\odot$). The close correlation between IR and radio luminosities of SMGs resembles that of local star-forming galaxies \\citep{Helou1985, CondonBroderick1991, Condon1992, Yun2001}, implying that star-formation is the principal power source.  Despite these successes however, the number of SMGs with spectroscopic redshifts is only around a hundred, and less than half of these are characterized in the far-infrared (FIR). Furthermore, the radio-undetected SMG population is almost completely unexplored. A truly unbiased understanding of SMGs requires improved selection methods.   \\subsection{SMGs and {\\em Spitzer}}   A major step forward in understanding SMGs was provided by {\\em Spitzer}. With its sensitive mid-IR imaging and spectroscopic capabilities, {\\em Spitzer} detected most SMGs in multiple mid-IR bands, which can provide accurate photometric redshifts $dz/(1+z)$$\\simeq$0.1 \\citep[][Lonsdale et al.~2009]{Pope2006}. Furthermore, many {\\em Spitzer}-selected objects are predicted to be submillimeter bright (Lonsdale et al.~2009; hereafter \\citet{Lonsdale2009}). Unfortunately, the identification of mid-IR counterparts to SMGs is plagued by problems similar to those of optical association \\citep{Pope2006}, and so the reliance on radio associations remains.  Selecting strongly luminous starbursts at $z$$\\sim$2 from {\\em Spitzer} data can overcome these problems, and the resulting samples are expected to contain a large proportion of SMGs within them. Thus, a two-step process has been suggested. First, selecting sources with peak {\\em Spitzer}--Infrared Array Camera (IRAC) flux densities in the 5.8\\um channel should remove galaxies with a mid-IR luminous active galactic nucleus \\citep[AGN;][]{Weedman2006, Farrah2008}, leaving objects with a clear rest-frame 1.6\\um opacity minimum at $z$$\\approx$2. The required IRAC detections ensure that stellar masses are large, especially when applied to limited sensitivity samples such as SWIRE \\citep{Lonsdale2003}. Second, a bright 24\\um flux density cut ($S_{24\\mathum}$$>$400\\,$\\mu$Jy) favors starbursts at the same redshift as the strong rest-frame 7.7\\um polycyclic aromatic hydrocarbon (PAH) emission feature is redshifted into the band. Redshifts from InfraRed Spectrograph (IRS) spectra or from IRAC photometry confirm that most sources selected in this way lie at $z$$\\simeq$2 \\citep[][N.\\ Fiolet et al., in preparation]{Weedman2006}. The effectiveness of such a two-pronged selection was confirmed by \\citet{Lonsdale2009} and Fiolet et al.~(2009, hereafter \\citet{Fiolet2009}), who detected a significant fraction of these sources at 1.2\\,mm using the MAMBO camera at the IRAM 30\\,m telescope (i.e.,~$S_{1.2_{\\rm mm}}$$\\gtrsim$2\\,mJy). Clearly, the proposed {\\em Spitzer} selection (see Table~\\ref{tab:selection}) yields distant luminous starburst galaxies without radio preselection.  \\begin{deluxetable}{l c}[!bth] \\tablewidth{\\columnwidth} \\tablecolumns{2} \\tablecaption{Sample Selection Criteria\\label{tab:selection}} \\tablehead{ Band/Instrument & Criterion } \\startdata IRAC & $S_{3.6\\mathum} < S_{4.5\\mathum} < S_{5.8\\mathum} > S_{8.0\\mathum}$ \\\\[1pt] MIPS 24 & $S_{24\\mathum} > 400$\\,$\\mu$Jy \\\\[1pt] MAMBO & $S_{1.2{\\rm mm}} \\gtrsim 2$\\,mJy \\enddata \\tablecomments{Summary of the selection criteria for the MAMBO-detected 5.8$\\mathum$ peaker samples of \\citet{Lonsdale2009} and \\citet{Fiolet2009}. } \\end{deluxetable}  A key question is to what extent the SWIRE-selected samples are in fact representative of the purely submillimeter-selected population. A straightforward way to test this is to constrain where the peak of the IR emission lies, thus bridging the gap between the cold dust detected by MAMBO, and the hot dust and PAHs detected by {\\em Spitzer}. Such constraints can be used to derive effective dust temperatures and accurate IR luminosities, for comparison with the overall SMG population. Because the sample consists of galaxies at z$\\sim$2, and dust temperatures are expected in the range of 30--40\\,K, the 350\\um band offers an ideal opportunity for providing the constraints we seek.  For our study, we targeted 12 galaxies from \\citet{Fiolet2009} and another 8 from \\citet{Lonsdale2009}. The former is an almost complete sample of sources from 0.5\\,deg$^2$ satisfying the criteria of Table~\\ref{tab:selection}, whereas the \\citet{Lonsdale2009} sources are drawn arbitrarily from $\\sim$30\\,deg$^2$. To enhance our chances of detection at 350$\\mathum$, we selected objects with clear MAMBO detections, with an additional preference toward the highest 1.2\\,mm flux density measurements.  The observations are described in Section~\\ref{sec:observations} and the results are summarized in Section~\\ref{sec:results}, where we also discuss one of the submillimeter sources, which appears to be a close association of several IR galaxies. In Section~\\ref{sec:sed}, we develop appropriate spectral energy distribution (SED) models to interpret the measurements. With these models, we characterize the sample in Section~\\ref{sec:discussion}, deriving precise dust temperatures, dust/gas masses and luminosities, and with them we verify the (far-)infrared to radio correlation.  We assume a cosmology with $H_0$=71\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_M$=0.27 and $\\Omega_{\\Lambda}$=0.73, and calculate distances as prescribed by \\citet{Hogg1999}. ",
        "conclusions": "Our main conclusion is that the mid-IR selection criteria of the \\citet{Lonsdale2009} and \\citet{Fiolet2009}, picking the bright 24\\um sources that are also 5.8\\um peakers, is effective in identifying a significant fraction of the SMGs around $z$$\\simeq$2, thus providing a means of enhancing our understanding of the elusive SMG population through studies at mid-IR and shorter wavelengths. The FIR characteristics (dust masses, temperature, and IR luminosities) of the SWIRE-selected sources are essentially identical to those of the classical SMGs at the same redshift. Additionally, the SWIRE-selected sample exhibits a correlation between radio and IR luminosities that is indistinguishable from that seen for SMGs. Our other conclusions are as follows.  \\begin{enumerate}  \\item{To provide realistic FIR characterizations for the sample, we developed new models for the SEDs of galaxies with power-law mass distributions of temperature components. We demonstrated that such models describe local starbursts extremely well, under a tightly constrained dust emissivity index $\\beta$$\\simeq$1.5 and mass-temperature index $\\gamma$$\\simeq$7.2. The $\\gamma$ values are consistent with what we expect for heating dominated by star formation.}  \\item{We find that the most likely IR spectral continuum shapes of the distant SWIRE-selected starbursts differ only slightly from the local examples, with best-fit $\\gamma$$\\simeq$6.7$\\pm$0.1. The lower $\\gamma$ values are consistent with star formation under the high optical depths ($\\tau_{\\rm pk}$$\\simeq$1.2) in our sample.} \\item{Relying on the stacked IRS spectra of \\citet{Lonsdale2009} for estimating the 24\\um continuum contribution, we place the typical diameter of starburst activity in these SMGs above 1.2\\,kpc (with 95\\% confidence), and find a most likely value of 2\\,kpc.} \\item{The emission scales, the typical $\\gamma$, and the consistently low $q$ values all confirm that star formation is the main power source in the galaxies of the sample.} \\item{One of our targets shows a close clustering of three or four 350\\um components, each of which contributes comparably to the observable 350\\um flux density under the 15''--30'' resolution of most current (sub)millimeter surveys. A commonality of unresolved SMG multiplets could bias our understanding of this population. We need further studies to explore the importance of close multiplets in current and future submillimeter surveys.}  \\end{enumerate}"
    },
    "1004/1004.2417_arXiv.txt": {
        "abstract": "Numerical simulations of the magnetorotational instability (MRI) with zero initial net flux in a non-stratified isothermal cubic domain are used to demonstrate the importance of magnetic boundary conditions.In fully periodic systems the level of turbulence generated by the MRI strongly decreases as the magnetic Prandtl number ($\\Pm$), which is the ratio of kinematic viscosity and magnetic diffusion, is decreased. No MRI or dynamo action below $\\Pm=1$ is found, agreeing with earlier investigations.  Using vertical field conditions, which allow the generation of a net toroidal flux and magnetic helicity fluxes out of the system, the MRI is found to be excited in the range $0.1\\le\\Pm\\le10$, and that the saturation level is independent of $\\Pm$. In the vertical field runs strong mean-field dynamo develops and helps to sustain the MRI. ",
        "introduction": "The realization of the astrophysical signifigance of the magnetorotational instability \\citep{BH91}, first discovered in the context of Couette flow \\citep{V59,C60}, seemed to resolve the long-standing problem of the mechanism driving turbulence in accretion disks. Early numerical simulations produced sustained turbulence, large-scale magnetic fields and outward angular momentum transport \\citep[e.g.][]{BNST95,HGB95}.  These results also showed that a significant qualitative difference exists between models where an imposed uniform magnetic field is present as opposed to the situations where such field is absent: the saturation level of turbulence and angular momentum transport are substantially higher when a non-zero vertical net flux is present \\citep[e.g.][]{BNST95,Sea96}.  Also the presence of an imposed net toroidal field seemed to enhance the transport \\citep{Sea96}.  In the meantime, a lot of numerical work has been done with zero net flux setups that omit stratification and adopt fully periodic or perfectly conducting boundaries in order to study the saturation behaviour of the MRI in the simplest possible setting \\citep[e.g.][]{FP07,FPLH07,LKKBL09,KKV10}. Due to the boundary conditions, the initial net flux in conserved and no magnetic helicity fluxes out of the system are allowed.  The results of these investigations have shown that as the numerical resolution of the simulations increases, or equivalently as the explicit diffusion decreases, the level of turbulence and angular momentum transport transport decrease, constituting a convergence problem for zero net flux MRI \\citep{FPLH07}.  Runs with explicit diffusion show that sustaining turbulence becomes increasingly difficult as the magnetic Prandtl number, $\\Pm=\\nu/\\eta$, where $\\nu$ is the viscosity and $\\eta$ the magnetic diffusivity, is decreased \\citep{FPLH07}. Currently the convergence problem is without a definite solution. It has been suggested that this issue could be related to the $\\Pm$-dependence of the fluctuation dynamo \\citep[e.g.][]{Schekea07}.  It has even been argued that the MRI in periodic zero net flux systems would vanish in the limit of large Reynolds numbers and that a large-scale dynamo would be needed to sustain the MRI and turbulence \\citep{V09}.  Notably, large-scale dynamos have no problems operating at low magnetic Prandtl numbers as long as the relevant Reynolds and dynamo numbers exceed critical values \\citep{B09}.  From the point of view of mean-field dynamo theory \\citep{BS05}, systems with fully periodic or perfectly conducting boundaries are rather special. In such closed setups magnetic helicity, defined as a volume integral of $\\bm{A}\\cdot\\bm{B}$, where $\\bm{A}$ is the vector potential and $\\bm{B}=\\bm\\nabla\\times\\bm{A}$ is the magnetic field, is a conserved quantity in ideal MHD. In the presence of magnetic diffusion, magnetic helicity can change only on a timescale based on microscopic diffusivity, which is usually a very long in any astrophysical setting. Such a behaviour, which has been captured in numerical simulations \\citep{B01}, is well described by simple mean-field models taking into account magnetic helicity conservation \\citep[e.g.][]{BB02}. This would mean that generating appreciable large-scale magnetic fields, which are possibly vital for sustaining the MRI, can take a very long time. Furthermore, the saturation value of the mean magnetic field decreases inversely proportional to the magnetic Reynolds number \\citep[e.g.][]{CH96,B01}. In dynamo theory this detrimental effect to the large-scale dynamo is known as the catastrophic quenching \\citep{VC92}.  The situation, however, changes dramatically if magnetic helicity flux out of the system is allowed. In particular, the \\cite{VC01} flux, which requires large-scale velocity shear to be present and flows along the isocontours of shear, is a potential mechanism that can drive a magnetic helicity flux out of the system and alleviate catastrophic quenching.  Indirect evidence for its importance exists from convection simulations in a shearing box setup \\citep{KKB08,KKB10b}, where dynamo excitation is easier in systems with boundaries that allow a net magnetic helicity flux. However, these results can be explained by a somewhat higher critical dynamo number in the perfect conductor case \\citep{KKB10b}, which is a purely kinematic effect. More dramatic differences between different boundary conditions are seen in the nonlinear saturation regime, with strong quenching of large-scale magnetic fields in the perfect conductor case \\citep{KKB10b}. The reason for this behaviour is not yet clear, especially in light of recent results of \\cite{HB10} who failed to find evidence of the Vishniac--Cho flux in a numerical setup similar to ours.  In the present paper we demonstrate that the boundary conditions play a crucial role for the excitation of the MRI and the associated large-scale dynamo. Following previous work that has shown that open boundary conditions allow more efficient dynamo action \\citep{KKB08,KKB10b}, we model a system that is isothermal, non-stratified, and the magnetic field has a zero net flux initially. We then apply vertical field boundary conditions which allow a magnetic helicity flux through the vertical boundaries by letting the magnetic field cross them. We show that if the MRI is excited, a large-scale dynamo is also excited and that the saturation level of the turbulence, large-scale magnetic field, and angular momentum transport are essentially independent of $\\Pm$.  This is contrasted by periodic simulations where we find a strong $\\Pm$-dependence in accordance with earlier studies. Our results also suggest that for a given $\\Pm$ the results (level of turbulence and angular momentun transport) are independent of the magnetic Reynolds number \\citep[see also][]{F10}.  The remainder of the paper is organised as follows: in Sect.~\\ref{sec:model} we describe our model, and in Sect.~\\ref{sec:results} and \\ref{sec:conclusions}, we present our results and conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We present three-dimensional numerical simulations of the magnetorotational instability in an isothermal non-stratified setup with zero net flux initially. Using fully periodic boundaries, that do not allow the generation of a mean toroidal flux or magnetic helicity fluxes out of the system, we encounter the convergence problem \\citep{FPLH07} of the MRI: turbulent kinetic and magnetic energies, and the angular momentum transport increase approximately proportional to the magnetic Prandtl number. Intermittent large-scale magnetic fields are observed in the periodic runs. Increasing the Reynolds numbers moderately at a given $\\Pm$ does not appear to markedly change the results in the saturated state.  When vertical field boundary conditions, allowing the generation of a mean flux and a magnetic helicity flux, are used, the MRI is excited at least in the range $0.1\\le\\Pm\\le10$ for our standard value of $\\Cem=1.5\\cdot10^4$. We find that the saturation level of the turbulence and the angular momentum transport are only weakly dependent on the Prandtl number and that strong large-scale fields are generated in all cases. The Shakura--Sunyaev viscosity parameter has consistently a value of $\\aSS\\approx6\\cdot10^{-3}$ in the vertical field case. Exploring even lower values of $\\Pm$ is infeasible at the moment due to prohibitive computational requirements but there are no compelling arguments against a large-scale dynamo operating at low $\\Pm$ \\citep{B09}. We conjecture that the operation of the MRI at low $\\Pm$ is due to the efficient large-scale dynamo in the system. It is conceivable that the dynamo only works if magnetic helicity is allowed to escape \\citep[see also][]{V09} or annihilate at the disk midplane due to an internal diffusive flux \\citep{MCCTB10}.  However, measuring the magnetic helicity fluxes in the presence of boundaries is difficult due to the fact that they are in general gauge dependent \\citep[e.g.][]{BDS02,HB10}.  The current results highlight the close connection between dynamo theory and the theory of magnetised accretion disks \\citep[see also][]{Bl10} and the importance of studying the results in a common framework \\citep[e.g.][]{G10}. Clearly, a more thorough study is needed in order to substantiate the possible role of magnetic helicity fluxes for the excitation and saturation of the MRI. We plan to address these issues in future publications."
    },
    "1004/1004.2883_arXiv.txt": {
        "abstract": "We use a highly homogeneous set of data from 132 early-type galaxies in the Virgo and Fornax clusters in order to study the properties of the globular cluster luminosity function (GCLF). The globular cluster system of each galaxy was studied using a maximum likelihood approach to model the intrinsic GCLF after accounting for contamination and completeness effects. The results presented here update our Virgo measurements and confirm our previous results showing a tight correlation between the dispersion of the GCLF and the absolute magnitude of the parent galaxy. Regarding the use of the GCLF as a standard candle, we have found that the relative distance modulus between the Virgo and Fornax clusters is systematically lower than the one derived by other distance estimators, and in particular it is 0.22\\textsl{mag} lower than the value derived from surface brightness fluctuation measurements performed on the same data. From numerical simulations aimed at reproducing the observed dispersion of the value of the turnover magnitude in each galaxy cluster we estimate an intrinsic dispersion on this parameter of 0.21\\textsl{mag} and 0.15\\textsl{mag} for Virgo and Fornax respectively. All in all, our study shows that the GCLF properties vary systematically with galaxy mass showing no evidence for a dichotomy between giant and dwarf early-type galaxies. These properties may be influenced by the cluster environment as suggested by cosmological simulations. ",
        "introduction": "\\label{sec:INTRO}  The distribution of globular cluster (GC) magnitudes has the remarkable property that it is observed to peak at a value of $M_V \\approx -7.5$ \\textsl{mag} in a near universal fashion (e.g., Jacoby et~al.~1992, Harris 2001, Brodie \\& Strader 2006). This distribution, usually referred to as the GC luminosity function (GCLF), has been historically described by a Gaussian.  By virtue of its near universality, the derived mean or ``turnover'' magnitude $\\mu$ has seen widespread use as a distance indicator (e.g.~Secker 1992, Sandage \\& Tammann 1995), even though some dispersion and discrepant results have been reported in the literature (see discussion in Ferrarese et~al.~2000a).  There is nevertheless no solid theoretical explanation for the observed universality of the turnover magnitude. The luminosity function is a reflection of the more fundamental mass spectrum of the GCs, and as such the ``universal'' turnover magnitude corresponds to a cluster mass of $\\sim 2 \\times 10^5 \\mathcal{M_{\\odot}}$. Vast efforts have been undertaken from the theoretical point of view in order to explain the underlying universal mass function. The many publications on this topic can be separated into those trying to identify some particular initial condition that selects a certain mass scale for star formation (e.g., Peebles \\& Dicke 1968, Fall \\& Rees 1985, West 1993), and those looking for a destruction mechanism that selects clusters in a particular mass range starting from an initially wide mass spectrum (e.g., Fall \\& Rees 1977, Gnedin \\& Ostriker 1997, Prieto \\& Gnedin 2008)  At the high-mass end (i.e.~$m_{gc} < \\mu$) the mass function of globular clusters resembles very closely the mass function of young clusters and molecular clouds in the Milky Way and other nearby galaxies (see e.g.~Harris \\& Pudritz 1994, Elmegreen \\& Efremov 1997, Gieles et~al.~2006). On the other hand, neither young clusters nor molecular clouds show a turnover on their mass distributions, but they keep rising monotonically following a power-law to lower masses. Fall \\& Zhang (2001) used simple analytical models (including evaporation by two-body relaxation, gravitational shocks and mass loss by stellar evolution) to study the evolution of the GC mass function. They showed that, for a wide variety of initial conditions, an initial power-law mass function develops a turnover that, after 12 Gyr, is remarkably close to the observed turnover of the GCLF. Vesperini (2000, 2001) reaches a similar conclusion, but finds that a log-normal mass function provides a better fit to the data. Fainter than the turnover, the evolution would be dominated by two-body relaxation, and the mass function would end up having a constant number of GCs per unit mass, reflecting the fact that the masses of tidally limited clusters are assumed to decrease linearly with time until they are destroyed (other authors propose different mass-loss rates, see e.g., Lamers et~al.~2006). Brighter than the turnover, the evolution is dominated by stellar evolution at early times and by gravitational shocks at late times. Recently, McLaughlin \\& Fall (2008) have shown that the GC mass function in the Milky Way depends on cluster half-mass density (i.e.~the mean density within a radius containing half the total mass of the GC), in the sense that the turnover mass increases with half-mass density, while the width of the GC mass function decreases.  But while there is currently a fairly good understanding of the dynamical processes that shape the GCLF, many details are still missing. In particular none of the theories proposed has been entirely successful on addressing the question of how the turnover magnitude can remain constant regardless of environmental properties and the mass of the host galaxy.  The use of deep HST data during the last years has resulted in high quality GCLF data, reaching 2 magnitudes beyond the turnover at the distance of the Virgo cluster ($\\sim$16.5 Mpc, Mei et~al.~2007). The use of these deeper observations has recently uncovered a strong correlation between the GCLF dispersion and the absolute magnitude of the parent galaxy (Jord\\'an et~al.~2006,~2007b), demonstrating the non-universality of this parameter and, as a consequence, of the GCLF as a whole. Here we present a study of the GCLF of 132 early type galaxies aimed to perform a precise test of the GCLF as a distance indicator by comparing the relative distance between the Virgo and Fornax clusters derived using the GCLF to the one derived using an analysis of surface brightness fluctuations (SBF, Tonry \\& Schneider 1988) based on the same data (Blakeslee et~al.~2009). Previous papers in the this series have presented an introduction to the survey (Jord\\'an et~al.~2007a), the properties of the central surface brightness profiles of early-type galaxies (C\\^ot\\'e et~al.~2007) and a catalog of SBF distances and a precise measurement of the Virgo-Fornax distance (Blakeslee et~al.~2009).  The organization of this paper is as follows. In \\S\\ref{sec:OBS&DATA} we present a description of the observations and data reduction procedures. In \\S\\ref{sec:GCLF} we describe the GCLF model fitting, and in \\S\\ref{sec:sigma_M} we compare the properties of the fits to previous results regarding the dispersion of the GCLF. Section \\S\\ref{sec:dist} is focused on determining how universal the value of the turnover magnitude is, while in \\S\\ref{sec:StCand} we look for a better understanding of the external parameters that might affect this value. Finally, in \\S\\ref{sec:CONC} we summarize our results and the main conclusions of this paper. ",
        "conclusions": "\\label{sec:CONC} We used ACS/HST data in order to study the GCLF of 89 early-type galaxies in the Virgo cluster and 43 galaxies in the Fornax cluster, which constitute the most homogeneous set of data used to date for this purpose. The GCLF of these galaxies was fitted by using a maximum likelihood approach to model the intrinsic Gaussian distribution after accounting for contamination and completeness effects. From the derived values of the turnover magnitude and the dispersion of the Gaussian fits we conclude that:  \\begin{enumerate}  \\item The analysis of 43 early-type galaxies belonging to the Fornax cluster shows that the dispersion of the GCLF decreases as the luminosity of the host galaxy decreases, confirming our previous results obtained with Virgo galaxies (Jord\\'an et~al. 2006, 2007b).  \\item By using the GCLF turnover magnitude as a distance indicator on our homogeneous data set we derive a relative distance modulus between the Virgo and the Fornax clusters of $\\Delta(m-M)_{\\rm GCLF}=0.20\\pm 0.04$ \\textsl{mag}, which is lower than the one derived using SBF measurements on the same data, $\\Delta (m-M)_{\\rm SBF}=0.42\\pm0.03$ \\textsl{mag}.  \\item Setting the relative Virgo-Fornax distance as that given by SBF implies a difference in the value of $\\langle \\mu_{TO} \\rangle$ in the two closest clusters of galaxies, suggesting that this quantity is influenced by the environment in which a GC system is formed and evolves. These results support a previous study by Blakeslee \\& Tonry (1996), who found a correlation between GCLF turnover magnitude and velocity dispersion of the host cluster, in the sense that galaxy clusters with higher velocity dispersions (higher masses) host galaxies with fainter turnovers in their GC systems.  \\item The discrepancy in the absolute magnitude of the GCLF turnovers in Virgo and Fornax can be accounted for if GC systems in the Fornax clusters were on average $\\sim$3 Gyrs younger than those in Virgo (thus making them brighter). Recent results from high-resolution numerical simulations (e.g.~Springel et~al.~2005, De Lucia et~al.~2006) suggest that stellar populations of Virgo-like galaxy clusters (high mass and high velocity dispersion) were formed mostly at higher redshift compared to less massive and lower-dispersion clusters like Fornax. This trend could therefore be at least partially responsible for the observed discrepancy in the absolute GCLF turnover magnitudes between both clusters.  \\item We have measured a total dispersion on the value of the turnover magnitude of 0.31 and 0.28 \\textsl{mag} for Virgo and Fornax respectively. We show using simulations that these values can be only partially accounted by the dispersion produced by cluster depth and observational uncertainties. The additional dispersion can be modeled by an intrinsic dispersion on the value of $\\mu_{0}$ of 0.21 \\textsl{mag} for the Virgo cluster and 0.15 \\textsl{mag} for Fornax.  \\item The measured GCLF turnover is found to be systematically fainter for low luminosity galaxies, showing a $\\sim$0.3 \\textsl{mag} decrease on dwarf systems, although we suffer from large uncertainties in that galaxy luminosity regime. The luminosity (i.e.~$\\sim$ mass) of the parent galaxy seems to play an important role on shaping the final form of the luminosity distribution. This might be at least partly accounted for by the effects of dynamical friction if all other processes that contribute on shaping the mass function (two-body relaxation, tidal shocks, etc.)  were to lead to a roughly constant M$_{TO}$ (Jord\\'an et~al. 2007b).  \\item Overall we find that GCLF parameters vary continuously and systematically as a function of galaxy luminosity (i.e.~mass). The correlations we present here show no evidence for a dichotomy between giant and dwarf early-type galaxies at $M_z \\sim -19.5$ ($M_B \\sim -18$) in terms of their GC systems. This is consistent with results presented in several recent studies (e.g.~Graham \\& Guzm\\'an, 2003; Gavazzi et~al.~2005, C\\^ot\\'e et~al.,~2006), and is at odds with earlier claims by Kormendy (1985).  \\end{enumerate}"
    },
    "1004/1004.3627_arXiv.txt": {
        "abstract": "Giant elliptical galaxies, believed to be built from the  merger of lesser galaxies, are known to house a massive black hole at their center rather than a compact star cluster.  If low- and intermediate-mass galaxies do indeed partake in the hierarchical merger scenario, then one needs to explain why their dense nuclear star clusters are not preserved in merger events.  A valuable clue may the recent revelation that nuclear star clusters and massive black holes frequently {\\it co-exist} in intermediate mass bulges and elliptical galaxies.  In an effort to understand the physical mechanism responsible for the disappearance of nuclear star clusters, we have numerically investigated the evolution of merging star clusters with seed black holes.  Using black holes that are  1-5\\% of their host nuclear cluster mass, we reveal how their binary coalescence during a merger dynamically heats the newly wed star cluster, expanding it, significantly lowering its central stellar density, and thus making it susceptible to tidal destruction during galaxy merging.  Moreover, this mechanism provides a pathway to explain the observed reduction in the nucleus-to-galaxy stellar mass ratio as one proceeds from dwarf to giant elliptical galaxies. ",
        "introduction": "The central regions of inactive galaxies are receiving increasing attention as astronomers begin to accurately quantify their inner-most features.  These range from partially evacuated cores housing a massive black hole (MBH) in giant galaxies to excess light in the form of a dense nuclear star cluster (NC)  in less massive spheroids\\footnote{The term spheroid is used to denote either an elliptical galaxy or  the bulge of a disk galaxy.}.  Curiously, an unexpected connection between MBHs and NCs  is starting to emerge (Ferrarese et al.\\ 2006a; Wehner \\& Harris 2006; Balcells et al.\\ 2007; Graham \\& Spitler 2009).  At the low mass end, dwarf elliptical (dE) galaxies are frequently observed to contain a dense cluster of stars near their centre (e.g.\\ Sandage \\& Binggeli 1984; Binggeli et al.\\ 1985; Ferguson \\& Binggeli 1994).  The stellar mass of these NCs relative to their host spheroid's stellar mass is known to systematically decrease as the dE mass increases (e.g.\\ Graham \\& Guzm\\'an 2003; Grant et al.\\ 2005).\\footnote{The bulges of disc galaxies also commonly contain a NC (e.g., Philips et al.\\ 1996; Carollo, Stiavelli \\& Mack 1998; B\\\"oker et al.\\ 2002), and the same general trend in the nuclear-to-spheroid stellar flux ratio is observed (e.g., Balcells et al.\\ 2003, 2007).}  Therefore, if nucleated elliptical galaxies are players in an hierarchical Universe (White \\& Rees 1978), one cannot simply merge such galaxies and double the mass of the new host galaxy and its NC.  At the high mass end are massive elliptical galaxies --- the end product of major mergers.  However, such galaxies are observed not to contain NCs, instead they possess central stellar deficits relative to the inward extrapolation of their outer S\\'ersic light profile (e.g.\\ Graham et al.\\ 2003; Trujillo et al.\\ 2004; Graham 2004; Ferrarese et al. 2006b). While it has been advocated that the dry merging of elliptical galaxies will result in the partial evacuation of the new galaxy's core due to the binary coalescence of pre-existing MBHs (Begelman et al.\\ 1980; Ebisuzaki et al.\\ 1991; Milosavljevi{\\'c} \\&  Merritt 2001; Merritt \\& Milosavljevi{\\'c}  2005; Merritt et al.\\ 2007), there must be more going on.  There must be a phase which erases the NCs ---  not included in the above mentioned studies --- that are prevalent in the less massive progenitor galaxies. While Berczik et al.\\ (2005) used Plummer models to represent galaxies in their detailed analysis of the impact that binary MBHs can have, here we dramatically rescale the problem, using Plummer models to represent NCs that are $\\sim$$10^3$ times smaller and, globally, $\\sim$$10^5$ times denser than galaxies.  Our motivation arises because it has recently been recognised that intermediate mass elliptical galaxies, and the similarly massive bulges of disc galaxies, regularly contain both a NC and massive black hole (e.g. Graham \\& Driver 2007; \\ Gonz{\\'a}lez Delgado et al.\\ 2008; Seth et al.\\ 2008).  Graham \\& Spitler (2009) have quantified how the $M_{\\rm BH}/M_{\\rm NC}$ mass ratio ($F_{\\rm BH}$) increases with increasing host spheroid stellar mass, $M_{\\rm sph}$, % until only a MBH is present at the centre.  While the runaway merger of NC stars during a merger event may lead to their conversion into a MBH (e.g., Zel'dovich \\& Podurets 1965; Frank \\& Rees 1976; Quinlan \\& Shapiro 1987; Lee 1993), and feedback processes may also impact $F_{\\rm BH}$ (McLaughlin et al.\\ 2006; Nayakshin et al.\\ 2009), this Letter explores whether dense NCs with seed MBHs might evaporate during a collision due to dynamical heating by the MBHs. ",
        "conclusions": "So far we  have  focused on the internal density profiles of NCs and have not  discussed their  {\\it projected} (two-dimensional, 2D) radial density profiles,  which can be directly compared with recent observational studies for (i) the origin of the apparent MBH-NC connection  (e.g., C\\^ot\\'e et al. 2006) and (ii) the observed $F_{\\rm BH}-M_{\\rm sph}$ relation (Graham \\& Spitler 2009). Fig.\\ 5 reveals that the {\\it projected} radial density profiles of NC merger remnants, ${\\Sigma}_{\\rm NC}(R)$, are significantly different between our four models with different $F_{\\rm BH}$. The rather low central ${\\Sigma}_{\\rm NC}$ value at $R=0.05R_{\\rm NC}$ and shallow inner density profile, in the model with  $F_{\\rm BH}=0.05$ suggests that if this merger remnant is located in the central region of a galaxy, it is less likely to be observed as a distinct NC. While the order of magnitude drop in surface density may seem like overkill, especially given the apparently small levels of excess nuclear light seen in most resolution-limited images, we note that well-resolved galaxies can have NC light up to 5 mag arcsec$^{-2}$ (100$\\times$) brighter than the underlying galaxy (e.g., Graham \\& Spitler 2009).  The present study confirms that more evolved  NCs --- by which we mean NCs, with MBHs, that are further along the merger tree --- can have lower inner  densities (${\\rho}_{\\rm NC}$ and ${\\Sigma}_{\\rm NC}$) and shallower inner 2D density  profiles than their progenitors.  This  suggests   that boundaries between distinct stellar nuclei and background field stars in galaxies are less clear for more evolved  systems as the NCs are effectively washed-out and dissolve into the host galaxy. Such diffuse NCs are also more susceptible to tidal destruction during galaxy merging. The present study therefore suggests that the observed $f_{\\rm BH}$-$M_{\\rm sph}$ relation can be understood in terms of the structural evolution of merging NCs with MBHs.   Measurements of partially-depleted galaxy cores, relative to a galaxy's outer light-profile,  have revealed a correlation between the central stellar mass deficit and the luminosity of the host spheroid and its MBH mass (e.g., Graham 2004; Ferrarese et al.\\ 2006b). As detailed in Graham \\& Guzm\\'an (2003) and C\\^ot\\'e et al.\\ (2007), the transition between massive galaxies with  partially-depleted cores and those without --- which frequently have excess nuclear light instead --- occurs around $M_B = -20.5$ mag.  Previous numerical simulations proposed that the origin of these central stellar deficits can be understood in the context of core formation through dynamical heating of stars by inspiralling MBHs in galaxy merging (e.g., Ebisuzaki et al. 1991). The present study has, for the first time, addressed one of the over-looked problems related to the nuclear structures  of galaxies: why and how can dense NCs  disappear during galaxy growth through galaxy merging ?   We advocate here that core-depletion due to the gravitational slingshot of host galaxy stars by inspiralling MBHs will not occur in earnest until the NCs surrounding the MBHs have first been eroded away by this same mechanism: once  the NCs are effectively gone, the binary MBHs, perhaps from additional merger events, can then commence to eat into the inner light profile of the host galaxy to produce the observed partially-depleted cores. This important step can explain why NCs disappear along the spheroid mass sequence and it also offers a process through which to understand the observed $F_{\\rm BH}$-$M_{\\rm bulge}$ relationship in terms of galaxy formation within the hierarchical merging scenario.   The present study suggests that if  NCs in low-mass galaxies  have seed MBHs, then their inner densities should progressively decline as galaxies grow through merging. This is at odds with the simple superposition of the NC density field for merging NCs without MBHs (Fig.\\ 3, lower panel). Although many previous theoretical studies investigated how NCs are formed, either by merging of SCs (e.g.,  Tremaine et al. 1975;  Capuzzo-Dolcetta \\& Miocchi 2008) or by dissipative gas dynamics in galaxies (e.g., Bekki et al. 2006; Bekki 2007b), they did not predict  (i) how BHs can be formed in NCs and (ii) what a reasonable value is for $F_{\\rm BH}$. The formation of seed MBHs {\\it within  NCs} may be different from that of intermediate-mass BHs in {\\it isolated} globular clusters through merging of stellar-mass black holes (e.g., O'Leary et al. 2006), because the deeper gravitational potential wells of the NC host galaxies  would  play a role in retaining interstellar gas more efficiently. It is thus our future study to investigate how seed MBHs can be formed in NCs at the epoch of NC formation in low-mass galaxies based on more sophisticated numerical simulations."
    },
    "1004/1004.1308_arXiv.txt": {
        "abstract": "The idea of a magnetic axion helioscope was first proposed by Pierre Sikivie in 1983. Tokyo axion helioscope was built exploiting its detection principle with a dedicated cryogen-free superconducting magnet and PIN photodiodes for x-ray detectors. Solar axions, if exist, would be converted into x-ray photons in the magnetic field. Conversion is coherently enhanced even for massive axions by filling the conversion region with helium gas. Its start up, search results so far and prospects are presented. ",
        "introduction": "Existence of the axion is implied to solve the strong CP problem of Quantum chromodynamics (QCD)\\cite{RDP,SW,FW,JEK}. The axion would be produced in the solar core through the Primakoff effect. It can be converted back to an x-ray in a strong magnetic field in the laboratory by the inverse process. This is the principle of the solar axion detection by the axion helioscope and appeared for the first time in  the paper\\cite{PS} entitled \"Experimental Tests of the 'invisible' axion\" written by Pierre Sikivie in 1983, 27 years ago. Sikivie indeed mentioned the idea of the axion helioscope in this paper.  The conversion process is coherent when the axion and photon remain in phase over the length of the magnetic field\\cite{R-S} if the axion is almost massless. Coherence can be maintained even for higher mass axions if the conversion region is filled with a low-Z buffer gas like helium\\cite{KVB,R-S}. ",
        "conclusions": ""
    },
    "1004/1004.4007_arXiv.txt": {
        "abstract": "Dynamical Chern-Simons gravity is an extension of General Relativity in which the gravitational field is coupled to a scalar field through a parity-violating Chern-Simons term. In this framework, we study perturbations of spherically symmetric black hole spacetimes, assuming that the background scalar field vanishes. Our results suggest that these spacetimes are stable, and small perturbations die away as a ringdown. However, in contrast to standard General Relativity, the gravitational waveforms are also driven by the scalar field. Thus, the gravitational oscillation modes of black holes carry imprints of the coupling to the scalar field. This is a smoking gun for Chern-Simons theory and could be tested with gravitational-wave detectors, such as LIGO or LISA. For negative values of the coupling constant, ghosts are known to arise, and we explicitly verify their appearance numerically. Our results are validated using both time evolution and frequency domain methods. ",
        "introduction": "In Chern-Simons gravity \\cite{Deser:1982vy,Lue:1998mq,Jackiw:2003pm} the Einstein-Hilbert action is modified by adding a parity-violating Chern-Simons term, which couples to gravity via a scalar field. This correction could explain several problems of cosmology \\cite{Weinberg:2008mc,GarciaBellido:2003wd,Alexander:2004xd,Alexander:2004us,Konno:2008np}. Furthermore, a Chern-Simons term arises in many versions of string theory \\cite{Polchinski:1998rr} and of loop quantum gravity \\cite{Ashtekar:1988sw,Taveras:2008yf,Mercuri:2009zt}, and Chern-Simons gravity can be recovered by truncation of low energy effective string models \\cite{Smith:2007jm,AD08}.  When Chern-Simons gravity was first formulated, the scalar field was considered as a prescribed function. Later on, it was understood that this {\\it a priori} choice is not really motivated (see the discussion in Ref.~\\cite{Yunes:2009hc}). Then, dynamical Chern-Simons (DCS) gravity has been formulated \\cite{Smith:2007jm}, in which the scalar field is treated as a dynamical field.  Since DCS gravity has a characteristic signature (the Chern-Simons term violates parity), there is the exciting prospect of testing its predictions against astrophysical observations. This has motivated a large body of work on the subject (for a review on DCS gravity and its astrophysical consequences see Ref.~\\cite{Alexander:2009tp}).  In this context, the study of black hole (BH) perturbations is very promising, since astrophysical black holes are probably the most appropriate objects to probe the strong field regime of General Relativity \\cite{Sopuerta:2009iy}.  The first study of BH perturbations in the context of DCS gravity has been carried out in Ref.~\\cite{Yunes:2007ss}, where it was found that, if the background solution contains a (spherically symmetric) scalar field, polar and axial perturbations of DCS BHs are coupled, and the equations describing them are extremely involved. Recently, in Ref.~\\cite{Cardoso:2009pk} (hereafter, Paper I), some of us found that, when the background scalar field vanishes, polar and axial gravitational perturbations of a Schwarzschild BH decouple, and only axial parity perturbations are affected by the Chern-Simons scalar field. We also found that under this assumption the gravitational and scalar perturbations are described by a coupled system of two second order ordinary differential equations (ODEs). The numerical integration of this system to find the quasi-normal modes (QNMs) of Schwarzschild DCS BHs is challenging, due to the same asymptotic divergence which prevented for many years the numerical computation of QNMs for Schwarzschild BHs \\cite{Chandrasekhar:1975zz,Nollert:1999ji,Berti:2009kk,Ferrari:2007dd}. Therefore, in Paper I the QNMs of Schwarzschild DCS BHs were not investigated thoroughly. It is remarkable that there are very few studies of this kind of system, i.e., QNMs described by coupled ODEs (one interesting work is presented in Ref. \\cite{Seahra:2005}). In Paper I we also claimed that Schwarzschild DCS BHs are unstable for a specific range of the parameters of the theory. This result was the consequence of a sign error in the derivation of the perturbation equations; on the contrary, as we discuss in this paper, there is strong evidence that these spacetimes are stable.  In this paper we complete the study of Schwarzschild DCS perturbations, performing a thorough numerical analysis of the perturbation equations. We employ two different -- and completely independent -- numerical approaches: time evolution and a formulation of the frequency domain approach \\cite{Watanabe:1980} which has never been applied before to the study of instability in black hole spacetimes. The results of the two independent methods agree very well, typically within an accuracy of $0.1\\%$, validating each other.  The main result we find is that any perturbation decays at late-time as a damped sinusoid. This is known as the ringdown phase, where the black hole radiates all excess hairs in its lowest QNMs \\cite{Berti:2009kk,Ferrari:2007dd}. What is new here, and with important implications for tests of DCS gravity, is that the gravitational sector has two distinct sets of QNMs. For large values of the constant $\\beta$ (associated to the dynamical coupling of the scalar field), these two sets coincide with the usual gravitational QNMs and scalar field QNMs of General Relativity. This result enables simple, yet fundamental tests on DCS gravity. By measuring (or not) these two different modes, one could effectively constrain DCS gravity through gravitational-wave observations. For instance, detection of ringdown modes with a signal-to-noise ratio $\\gtrsim 6$ (feasible with both the Earth-based LIGO and the space-based detector LISA), could allow one to test DCS gravity if the mass of the BH is known, for instance through observations of the inspiral phase of black hole binaries. For signal-to-noise ratios $\\gtrsim 150$ one could be able to discriminate between DCS gravity and standard General Relativity without any further knowledge of the BH parameters. \\subsection*{A summary of our results}  For the reader wishing to skip the technical details of the rest of the paper, the following is a brief summary of our results.  \\begin{itemize}  \\item[(i)] Two complementary numerical methods were developed and employed. They are completely independent and their concordance is very good.  \\item[(ii)] For small values of the coupling constant ($M^4 \\beta\\lesssim0.5$), the perturbative dynamics is characterized by a stable exponentially decaying phase. The intermediate late time evolution is dominated by \\begin{equation} \\Phi(t,r_{\\rm fixed}) = e^{\\omega_{\\rm no} \\, t} \\left(\\begin{array}{c} a\\\\ b \\end{array} \\right) \\end{equation} with $\\textrm{Re}[\\omega_{\\rm no}] = 0$ and $\\textrm{Im}[\\omega_{\\rm no}] < 0$ (with our sign conventions, a QNM is stable if $\\textrm{Im}[\\omega]<0$). Our results for the non-oscillatory frequency values are compatible with the expression: \\begin{equation} \\omega_{\\rm no} = - 0.04024 (M^4 \\beta)^{0.44} \\ell \\left(1 + \\frac{2.0953}{\\ell} - \\frac{3.4460}{\\ell^2} \\right)\\,. \\end{equation}  \\item[(iii)] For intermediate values of $M^4 \\beta$, field evolution is dominated by a stable oscillatory phase. We have detected two oscillatory modes, named here ``gravitational'' and ``scalar'' modes. Although the time profiles of the gravitational perturbation $\\Psi$ and of the scalar field $\\Theta$ are usually different, they consist on different superpositions of the {\\it same} modes.  \\item[(iv)] In the $\\beta\\to\\infty$ limit, these ``gravitational'' and ``scalar'' branches coincide with actual gravitational and scalar modes of Schwarzschild BHs in General Relativity. In this regime, we report that for $\\ell=2$, we find that the gravitational perturbation oscillates with a combination of the two modes \\begin{eqnarray} M \\omega_{\\rm grav} (\\Psi)&=&0.3736 - {\\rm i} \\,\\, 0.08899\\,,\\\\ M \\omega_{\\rm sc} (\\Psi)&=&0.4837 - {\\rm i} \\,\\, 0.09671\\,. \\end{eqnarray} These numbers correspond to the lowest mode of pure gravitational and scalar quasi-normal frequencies in Einstein's theory \\cite{Berti:2009kk}. The scalar field perturbation, instead, oscillates with the mode $\\omega_{\\rm sc}$ only. This behavior can be easily understood by looking at the form of the equations in this limit.  \\item[(v)] At late times, the field decays with a power-law tail, as $t^{- (2\\ell + 3)}$. The tails do not depend on $\\beta$ or $M$. Note that the same behavior characterizes Schwarzschild BHs \\cite{CPM78}, implying that a gravitational-wave observation of the tail would not be able to discriminate DCS gravity from General Relativity.  \\item[(vi)] An extensive investigation of BH oscillations, performed using two different numerical approaches, only yields stable modes, either oscillating or non-oscillating. This gives strong indications that Schwarzschild BHs in DCS modified gravity are stable against axial and polar perturbations.  \\item[(vii)] We also discuss how the inclusion of a non-vanishing scalar potential in the Lagrangian affects the QNM spectrum.  We focus on potentials of the form \\begin{equation} V(\\vartheta)=m^2\\vartheta^2+{\\cal O}(\\vartheta^3) \\label{potential} \\end{equation} and find that in the $\\beta\\to\\infty$ limit this inclusion only affects the scalar branch of QNMs, while the gravitational branch is unaltered. When $M^4\\beta\\lesssim100$, also the gravitational sector is affected by the scalar potential. \\end{itemize}  The paper is organized as follows. In Section \\ref{eqns} we briefly review the derivation of the perturbation equations in DCS gravity. In Section \\ref{num1} we describe the time domain and frequency domain numerical approaches that we have employed to solve the perturbation equations. In Section \\ref{results} we present our results in the time and frequency domains.  In Section \\ref{tests}, a possible observational signature of DCS gravity is discussed. Implications and final remarks are presented in Section \\ref{concl}.  In Appendix \\ref{sec:negative_beta} we discuss ghost-like instabilities arising when the wrong sign of the kinetic term in the action is chosen, i.e. when $\\beta<0$ in Eq.~\\ref{action} below. ",
        "conclusions": "We have found that Schwarzschild BHs in DCS modified gravity are stable against axial and polar perturbations. Indeed, an extensive investigation of BH oscillations, performed using two different numerical approaches, only yields stable modes, either oscillating or non-oscillating.  Polar perturbations obey exactly the same master equation as in General Relativity, and therefore BHs in DCS gravity oscillate at the same polar frequencies. Axial perturbations, instead, couple to a scalar field, enlarging the spectrum of resonances in the gravitational sector. In particular, the ringdown of a BH in DCS gravity is a superposition of two different QNM sectors. For large values of the constant $\\beta$, which is associated to the dynamical coupling of the scalar field, one of these sectors corresponds to the gravitational and the other sector to scalar-field QNMs of Schwarzschild BHs in General Relativity.  Thus, a golden opportunity to test these theories is by detection of BH ringdowns. As shown in Section \\ref{sec:nohair}, a modest SNR ($\\gtrsim 6$) could be sufficient to discriminate between General Relativity and DCS modified gravity. These estimates assume very special relative amplitudes between the modes. Accurate estimates, as well as constraints on the coupling parameters, require the calculation of accurate waveforms for physically interesting processes exciting these ringdown modes.  The problem dealt with here is also interesting for a number of other reasons, in particular because we expect such kind of problems, i.e. QNMs described by a system of coupled second order ODEs, to be a general feature of alternative and more intricate theories; surprisingly there are very few studies of this kind of system in General Relativity.  Finally, we detail in Appendix \\ref{sec:negative_beta} how ghost-instabilities develop in this theory when $\\beta<0$, by a careful analysis of the instability timescale and other features.   Generalization of our results to rotating black holes is of utmost importance, given that many astrophysical black holes are rapidly rotating. Rotating solutions in DCS gravity are only partially understood \\cite{Yunes:2009hc,Konno:2009kg}, we hope to come back to this issue in the near future."
    },
    "1004/1004.0697_arXiv.txt": {
        "abstract": "{Recently, the CoGeNT experiment has reported events in excess of expected background. We analyze dark matter scenarios which can potentially explain this signal.  Under the standard case of spin independent scattering with equal couplings to protons and neutrons, we find significant tensions with existing constraints. Consistency with these limits is possible if a large fraction of the putative signal events is coming from an additional source of experimental background. In this case, dark matter recoils cannot be said to explain the excess, but are consistent with it.  We also investigate modifications to dark matter scattering that can evade the null experiments.  In particular, we explore generalized spin independent couplings to protons and neutrons, spin dependent couplings, momentum dependent scattering, and inelastic interactions.  We find that some of these generalizations can explain most of the CoGeNT events without violation of other constraints. Generalized couplings with some momentum dependence, allows further consistency with the DAMA modulation signal, realizing a scenario where both CoGeNT and DAMA signals are coming from dark matter. A model with dark matter interacting and annihilating into a new light boson can realize most of the scenarios considered. }   \\begin{document} ",
        "introduction": "Numerous experiments have been designed to search for the dark matter (DM), which constitutes the vast majority of all matter.  Most searches have only placed limits.  How do these null results influence the interpretation of a putative signal?  This depends upon  the properties of Weakly Interacting Particles (WIMPs)  themselves. Variations in the couplings of WIMPs to protons, neutrons and spin \\cite{Jungman:1995df}, as well as possible inelastic \\cite{iDM,iDMUpdate} or momentum dependent \\cite{MDDM,FFDM}  scatterings can all change expectations for signal rates.  They also impact comparisons between different experiments.   In addition, differences in target compositions and energy thresholds cause astrophysical unknowns, such as the local velocity distribution of WIMPs to impact experiments in different ways.  When confronting a new signal, a broad exploration of these ideas is important to discern whether it can be of DM origin, or simply an unexpected background.  Recently, the CoGeNT collaboration reported on a low energy ionization spectrum not immediately identifiable with background~\\cite{Aalseth:2010vx}. One possible explanation for these low energy events is a genuine signal from DM nuclear recoils against the germanium target. However, since CoGeNT does not discriminate between nuclear and electron recoils, one must be wary of unexpected backgrounds.   It is thus important to determine the viability of a DM origin of the signal for various scenarios, given other experimental limits.  Initial work in this direction has already appeared in \\cite{Fitzpatrick:2010em} and the updated appendix of \\cite{Kopp:2009qt}.  In finding a DM fit,  \\cite{Fitzpatrick:2010em} allows the possibility of a substantial contribution to the CoGeNT data from both DM and background, while \\cite{Kopp:2009qt} allows only a DM contribution to CoGeNT.  They also differ in their treatment of possible systematic errors associated with the XENON experiment.   Reference \\cite{Fitzpatrick:2010em} entertains the possibility that systematic errors at the XENON experiment might degrade their published limits.  In this work, we further explore the viability of DM recoils as an explanation to the CoGeNT data.  We place an emphasis on the effect of the inclusion of background.  In the absence of an exponential background which makes a significant contribution to the data, an interpretation of the CoGeNT data as spin-independent scattering is strongly constrained.  With equal couplings to protons and neutrons,  the CoGeNT $99\\%$ confidence region is entirely excluded by recent results from CDMS silicon (CDMS-Si) run~\\cite{Filippini}. We demonstrate that parameter space can open up for other relations between these couplings.  We also examine other scenarios, such as light inelastic DM (iDM), spin-dependent couplings, and momentum dependent interactions. Aside from elastic spin-dependent interactions, we find that these generalizations are viable interpretations for CoGeNT not excluded by other experiments.  These analyses can be sensitive to the inclusion of backgrounds. Therefore we consider the effect of an additional exponential background component at low energies. Depending on the amount of background allowed, one lessens the tension with CDMS-Si.   However, completely avoiding the Si constraints requires a large amount of background, with at least one bin below the known background sources containing more than 50\\% of its events from a new background source. In this case,  the CoGeNT data is not really ``explained\" by DM recoil events, rather it would be fair to say the data is not {\\em inconsistent} with a DM signal.  We also investigate whether the preferred regions for the DAMA modulation signal can be made consistent with CoGeNT.  Typically, the DAMA region requires a significantly lower cross section than CoGeNT, however this assumes the level of channeling that the DAMA collaboration has claimed from Monte Carlo simulations \\cite{Bernabei:2007hw}.  In some models, it is possible to get consistency between them.  This typically requires lowering the CoGeNT DM contribution by introducing additional background and/or reducing the amount of channeling in NaI crystals to raise the required cross section for DAMA.  For certain specialized cases, it is even possible to get consistency by only moving DAMA up by a reduction in channeling.  This raises the possibility of consistency with all constraints, while simultaneously explaining CoGeNT and DAMA as true DM signals.   The outline for the rest of the paper is as follows.  In Section \\ref{sec:methodology}, we discuss how we interpret the potential signals of CoGeNT and DAMA and how we derive constraints from the other direct detection experiments that impact light mass WIMPs.  Those who are interested in our main results can skip to Section \\ref{sec:SI}, where we begin our discussion on whether WIMPs which scatter with spin independent interactions can explain the data.  We will find that there is tension between the signals and constraints for the standard assumption of equal proton and neutron couplings. The inclusion of a background opens regions of parameter space, but the points in these regions generally fail to explain at least half of the events in some bin below 1 keVee.  In Section \\ref{sec:alternatives}, we discuss alternative scenarios, starting with generalized couplings for protons and neutrons, spin dependent scattering, moving on to momentum dependent interactions and finally, inelastic DM.  For these scenarios, we find that generalized couplings along with  momentum dependent interactions are the most promising in terms of explaining all experiments.  We also comment on the types of models which might explain the necessary couplings and interactions. Models where the DM annihilates into and interacts via a new light boson provide simple explanations for most scenarios we consider.  Finally, in Section \\ref{sec:conclusions}, we conclude. ",
        "conclusions": "} The recent CoGeNT results are intriguing, both because of the lack of an obvious background, and the possible connection to the DAMA result. In this paper, we have taken the CoGeNT results and asked the question, in what cases can DM interactions consistently explain a substantial number of their lower energy events?  The limits of XENON10, SIMPLE, and especially CDMS-Si give significant constraints. Taken together, they make it difficult to answer this question in the affirmative without considering somewhat unconventional scenarios.  To summarize our results:  \\begin{itemize} \\item An interpretation of the low energy excess events as coming entirely from elastic scattering, either spin-independent or spin-dependent, of WIMPs against nucleons is possible and results in a good fit to the data. However, the implied parameter space is strongly disfavored by other searches.  \\item Postulating an additional background component at low energies helps to avoid existing constraints by allowing a reasonable fit to the data with lower WIMP mass and/or lower WIMP-nucleon cross-section. However, the WIMP signal in these regions of parameter space cannot be said to explain the data as it contributes less (and generally substantially less) than 50\\% to at least one of the first 5 bins. At  best one can say that a WIMP signal is not inconsistent with the data if we allow additional background component at low energies.  \\item  Modifications beyond standard scattering can improve the situation. \\begin{itemize} \\item Light inelastic DM with SI or SD interactions can explain the low energy events in the CoGeNT data, but have difficulty explaining both DAMA and CoGeNT with DM alone. \\item Generalized SI couplings to protons and neutrons are the most promising, in particular when $f_p\\approx -f_n$. This is especially true when considered in conjunction with momentum dependence. This allows situations where DM is simultaneously responsible for the CoGeNT and DAMA signals. \\end{itemize} \\end{itemize}  We leave it for future work to construct models which explain the necessary couplings of these nonstandard scattering mechanisms.  In the near future, further results from CDMS, from either additional silicon data or a low threshold germanium analysis, could impact whether these scenarios are still viable. More optimistically,  a confirmation of these light mass WIMPs could be coming soon from COUPP, XENON100, XMASS, and LUX, depending on their low energy threshold.  \\vskip 0.15in \\textbf{Note Added:} We would like to thank the authors of Refs.~\\cite{slac-paper,Graham:2010ca} for bringing to our attention their related work. Our analysis of the elastic spin-independent constraints are in agreement with what was found in \\cite{Graham:2010ca} except in the limits coming from the CDMS-Si runs. Ref.~\\cite{Graham:2010ca} used in addition the results of CDMS-SUF~\\cite{Akerib:2003px} which we have omitted for two reasons. First, CDMS-SUF was a fairly different experiment with larger backgrounds (not shielded like in CDMS-Soudan). Second, it had a much smaller exposure and most of its power is derived from the low threshold ($5\\keV$) where the efficiency is uncertain. To remain conservative in our limits we opted to include only the 2-tower data \\cite{Akerib:2005kh} and the 5-tower data shown in \\cite{Filippini}."
    },
    "1004/1004.3694_arXiv.txt": {
        "abstract": "{The fidelity of Smoothed Particle Hydrodynamics (SPH) simulations of accretion disks depends on how Artificial Viscosity (AV) is formulated.}{We investigate whether standard methodology is reliable in this regard.}{We test the operation of two methods for selective application of AV in SPH simulations of Keplerian Accretion Disks, using a ring spreading test to quantify effective viscosity, and a correlation coefficient technique to measure the formation of unwanted prograde alignments of particles.} {Neither the Balsara Switch ($B$) nor Time Dependent Viscosity (TDV) work effectively, as they leave AV active in areas of smooth shearing flow, and do not eliminate the accumulation of alignments of particles in the prograde direction. The effect of both switches is periodic, the periodicity dependent on radius and unaffected by the density of particles. We demonstrate that a very simple algorithm activates AV only when truly convergent flow is detected and reduces the unwanted formation of prograde alignments. The new switch works by testing whether all the neighbours of a particle are in Keplerian orbit around the same point, rather than calculating the divergence of the velocity field, which is very strongly affected by Poisson noise in the positions of the SPH particles.} ",
        "introduction": "In Smoothed Particle Hydrodynamics (SPH) simulations which involve strongly convergent flows, it is essential to include Artificial Viscosity (AV) (Monaghan $\\&$ Gingold 1983; Monaghan $\\&$ Lattanzio 1985; Monaghan 1992). It causes close, approaching SPH particles to repel each other, with a force which increases with approach velocity. The result is that particles in colliding streams are rapidly decelerated and the streams do not pass through each other. Such particle interpenetration would be an unphysical situation and cannot be allowed in a simulation. It is important to note that AV was never intended to replicate the behaviour of real viscosity (Monaghan 2005).  SPH is now used widely to simulate the evolution of disks (Bate et al 2003; Lodato $\\&$ Rice 2004; Rice et al 2003a, 2003b, 2004; Mayer et al 2004; Schaefer et al 2004; Boffin et al 1998; Watkins et al 1998a and 1998b, Stamatellos et al 2008 and 2009). The behaviour of viscosity in real circular shear flow in fluids is understood analytically and experimentally (eg Feynman 1964), and the evolution of accretion disks is predicted to be dominated by effects such as turbulence which behave like viscosity, causing angular momentum to be transferred outwards, enabling matter to fall in and accrete onto a central object (Shakura $\\&$ Sunyaev 1973; Lynden Bell $\\&$ Pringle 1974). It is often assumed that turbulence is generated and moderated by the magneto-rotational instability (Balbus $\\&$ Hawley 2002).  However, AV, if allowed to operate unchecked in a simulated Keplerian Disk, does not operate only when convergence occurs, but also acts to slow down neighbouring particles as they overtake because of the differential rotation. The magnitude of AV is therefore, systematically, too high. Clarke (2009) has shown that it is a key requirement for correct modelling of the evolution of disks that numerical viscosity is kept very low, otherwise it dominates the pseudo viscous forces, such as self-gravity, in the simulation. It is therefore crucial to disable AV in normal, equilibrium Keplerian shear flow, and to activate it only when it is needed to capture a shock.  In this paper we set out to test the Balsara Switch ($B$), and Time Dependent Viscosity (TDV), for their effectiveness in disabling AV in non-convergent Keplerian shear flow. In Sect. 2 we explain the theoretical operation of $B$ and in Sect. 3 we examine the actual values of $B$ calculated for a snapshot of a smoothly rotating accretion disk. In Sect. 4 we investigate the variation of $B$ with time for individual SPH particles. In Sect. 5 we turn to TDV and look at actual values calculated for a simulated Keplerian disk in which there is no convergence. In Sect. 6 we look at the spreading of a ring of particles, as a quantitative measure of the effectiveness of $B$ and TDV in reducing viscous shear. We also quantify, for the first time, the previously reported tendency for AV to cause SPH particles to form prograde alignments. A new approach to controlling AV in disks, by simply checking whether neighbours are in Keplerian orbit around the same object, is explained and evaluated in Sect. 7 and these results are discussed in Sect. 8. ",
        "conclusions": "Both $B$ (Balsara 1989) and the TDV approach (Morris and Monaghan 1997) to controlling AV in circular shear flow fail for the same reasons: SPH estimates of $\\nabla \\cdot {\\bf v}$ appear to detect convergence in steady shear flow. AV controlled by these methods is therefore applied in all areas of a simulated Keplerian accretion disk, with results which will be unprepresentative of real viscous evolution, and compromise the results (Clarke, 2009). AV results in the selective formation of prograde alignments of SPH particles, which can be measured using a correlation coefficient method. We find that a disk populated with randomly placed SPH particles initially has 18$\\%$ particles in prograde alignments and 18$\\%$ retrograde. After being allowed to evolve using pressure forces and AV, these figures change to 28$\\%$ and 12$\\%$. Clearly the use of AV changes the arrangement of particles in the experiment, even when very low levels of AV are applied.  An alternative method of controlling AV, based on identifying Keplerian velocity characteristics, shows promise, being more effective both at switching on in regions of convergence and off in pure Keplerian flow, and significantly reducing the accumulation of excess prograde alignments of particles."
    },
    "1004/1004.3377_arXiv.txt": {
        "abstract": "In this paper we consider a stable particle with flavor mixing. We demonstrate that incoherent conversion of heavy mass eigenstates into light ones and vice versa can occur, as a result of elastic scattering. This effect is nontrivial for non-relativistic particles, for which the standard flavor oscillation ceases rapidly due to incoherence. We also prove that if a heavy state is bound in a gravitational potential and a light state is unbound, the mass-state conversion can lead to gradual ``evaporation'' of the mixed particle from the potential. A number of implications, ranging from the cosmic neutrino background distortions to scenarios of cold dark matter evaporation from halos, are addressed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.3141_arXiv.txt": {
        "abstract": "We consider a model of dark energy/matter unification based on a k-essence type of theory similar to tachyon condensate models. Using an extension of the general relativistic spherical model which incorporates the effects of both pressure and the acoustic horizon we show that an initially perturbative k-essence fluid evolves into a mixed system containing cold dark matter like gravitational condensate in significant quantities. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.4557_arXiv.txt": {
        "abstract": "We report radial velocities (RVs), projected equatorial velocities ($v\\sin i$) and Ca\\,{\\sc ii}~triplet (CaT) chromospheric activity indices for 237 late-K to mid-M candidate members of the young open cluster NGC~2516.  These stars have published rotation periods between 0.1 and 15\\,days. Intermediate resolution spectra were obtained using the Giraffe spectrograph at the Very Large Telescope.  Membership was confirmed on the basis of their RVs for 210 targets. For these stars we see a marked increase in the fraction of rapidly rotators as we move to cooler spectral types. About 20 per cent of M0--M1 stars have $v \\sin i >15$\\,km\\,s$^{-1}$, increasing to 90 per cent of M4 stars. Activity indices derived from the first two lines of the CaT (8498\\AA\\ and 8542\\AA) show differing dependencies on rotation period and mass for stars lying above and below the fully convective boundary. Higher mass stars, of spectral type K3--M2.5, show chromospheric activity which increases with decreasing Rossby number (the ratio of period to convective turnover time), saturating for Rossby numbers  $<0.1$.  For cooler stars, which are probably fully convective and almost all of which have Rossby numbers $<0.1$, there is a clear decrease in chromospheric activity as $(V-I)_0$ increases, amounting to a fall of about a factor of 2--3 between spectral types M2.5 and M4. This decrease in activity levels at low Rossby numbers is not seen in X-ray observations of the coronae of cluster M-dwarfs or of active field M-dwarfs. There is no evidence for supersaturation of chromospheric activity for stars of any spectral type at Rossby numbers $<0.01$. We suggest that the fall in the limiting level of chromospheric emission beyond spectral type M3 in NGC~2516 is, like the simultaneous increase in rotation rates in field stars, associated with a change in the global magnetic topology as stars approach the fully  convective boundary and not due to any decrease in dynamo-generated magnetic flux. ",
        "introduction": "Measurements of the masses and radii of M-dwarfs are significantly discrepant from the predictions of evolutionary models (Ribas et al. 2008). Initial evidence for this comes from eclipsing binaries, where radii are 10--15 per cent higher at a given mass than predicted (L\\'opez-Morales 2007; Morales et al. 2009). It has been suggested that the presence of dynamo-generated magnetic fields in what are relatively fast rotating stars, can suppress convection, produce cool star spots and hence reduce the stellar effective temperature (D'Antona, Ventura \\& Mazzitelli 2000; Mullan \\& MacDonald 2001; Chabrier, Gallardo \\& Baraffe 2007). Jackson, Jeffries \\& Maxted (2009) measured the radii of single, rapidly rotating M-dwarfs in the young open cluster NGC~2516. They found that their radii, at a given luminosity, are also larger than predicted by evolutionary models. The discrepancy increases from a few per cent for early (M0) M-dwarfs, to some 50 per cent for mid-M dwarfs ($\\simeq $M4). The same evolutionary models correctly predict the radii of magnetically inactive M-dwarfs, thus implicating rotationally induced magnetic activity as the source of the discrepancy.  Whilst this appears credible in qualitative terms, further data are required to correlate measurements of mass and radii with measurements of rotation, magnetic field strength and indicators of chromospheric and coronal activity.  In low-mass F-, G- and K-type stars, the ratio of coronal X-ray to bolometric flux, $L_{\\rm x}/L_{\\rm bol}$, or a variety of similarly defined chromospheric flux indicators, are used as proxies for magnetic activity. Magnetic flux and X-ray/chromospheric activity both appear to depend primarily on rotation rate, but also on the convective turnover time, as expected from simply dynamo models (e.g. Mangeney \\& Praderie 1984).  Magnetic flux and magnetically induced emissions increase with rotation speed and with decreasing Rossby number -- defined as the ratio of rotation period to convective turnover time ($N_R = P/\\tau_c$). However, for $N_R < 0.1$, magnetic activity reaches a saturation plateau where $L_{\\rm x}/L_{bol} \\simeq 10^{-3}$ (Stauffer et al. 1994). Similar saturation plateaus are also found in chromospheric emission lines (Soderblom et al. 1993; James \\& Jeffries 1997). More limited observational evidence shows that for extremely fast rotating G- and K-stars with $v \\sin i$ in the range 100 to 200\\,km\\,s$^{-1}$ (and hence $N_R\\simeq 0.01$) there is a downward trend in $L_{\\rm x}/L_{bol}$ from the saturation plateau (Pizzolato et al. 2003). This effect has been dubbed supersaturation by Prosser et al. (1996).  Less is known about the behaviour of magnetic activity in fast rotating M-dwarfs, but these may harbour crucial clues to the explanation of saturation and supersaturation. They have much deeper convection zones than hotter stars and as a result have longer convective turnover times and hence lower Rossby numbers at the same rotation period. Furthermore, main-sequence M-dwarfs with spectral types of M3 and cooler are probably fully convective (Siess, Dufour \\& Forestini 2000), so a dynamo operating at the interface between radiative core and convective envelope can no longer explain their magnetic activity.  Rotation and magnetic activity in low-mass M-dwarfs have been the focus of much recent work. The key points are that in field M-dwarfs there appears to be an abrupt change in rotational properties at spectral type M3. Hotter than this there are few rapid rotators, but for cooler stars there are many fast rotators and almost no slowly rotating stars (Delfosse el al. 1998; Reiners \\& Basri 2008; Jenkins et al. 2009). This has been interpreted as a rapid lengthening of the spin-down timescale, roughly coinciding with the transition to fully convective stars. Field M-dwarfs earlier than M3 show clear evidence of a rotation-activity relation similar to F-K stars for coronal indicators, including saturation at small Rossby numbers (Pizzolato et al. 2003; Kiraga \\& Stepien 2007). For fully convective stars the data are sparse. Most fully convective field M-dwarfs are likely to have very small Rossby numbers. Their coronal activity appears to saturate at $L_{\\rm x}/L_{\\rm bol} \\simeq 10^{-3}$ out to spectral types of at least M6 (Delfosse et al. 1998; Reiners, Basri \\& Browning 2009), but peak levels of chromospheric emission, as measured by $L_{{\\rm H}\\alpha}/L_{\\rm bol}$, are lower in stars cooler than M6 and may begin to decline at spectral type M4 (Mohanty \\& Basri 2003).  Supersaturation either in coronal or chromospheric indicators is largely uninvestigated for M-dwarfs, though tentative evidence for the effect has been claimed at X-ray wavelengths by James et al. (2000) and by Reiners \\& Basri (2010) for the chromospheric H$\\alpha$ emission of rapidly rotating M7-M9 dwarfs.   A difficulty in most of these studies is that they (understandably) concentrate on nearby field M-dwarfs.  However, this inevitably leads to samples with a range of ages and metallicities, a wide spread of rotation velocities and often no information about rotation periods. A complementary approach is to target M-dwarfs in open clusters which are presumably coeval and chemically homogeneous and where ages and chemical compositions have already been determined from higher mass stars. The disadvantage here is the distance, but this can be mitigated using multiplexing instruments which operate over a significant area within a cluster -- for example, simultaneous time-series montitoring of many M-dwarfs to find rotation periods or fibre spectroscopy of many targets in one exposure.  In this paper we report the results of intermediate resolution spectroscopy, over the wavelength range 8060\\AA~to 8600\\AA, for 237 late-K to mid-M dwarfs.  These are all photometric candidate members of the open cluster NGC~2516 with published rotation periods and an age of $\\sim 150$\\,Myr (Irwin et al. 2007). The spectra were analyzed to identify cluster members and measure radial velocities (RVs), projected equatorial velocities ($v \\sin i$) and chromospheric activity using the CaT lines.  In section~2 we review the properties of NGC~2516. In sections 3 and 4 we report on our target selection and the observations and data analysis of fibre spectroscopy taken with the Very Large Telescope (VLT). Section 5 discusses the selection of cluster members from our data and presents their rotational properties.  In section 6 the strengths of two of the CaT lines are used to determine levels of chromospheric activity. In section 7 the results are investigated with respect to spectral type, period and Rossby number to look for evidence of chromospheric saturation or supersaturation.  In section 8 the results are compared with other observations and discussed in the context of current theories for the generation of magnetic fields and the magnetic topology in stars with masses above and below the fully convective boundary. ",
        "conclusions": "We have measured rotation rates and RVs, which confirm 210 late-K to mid-M dwarfs as members of the open cluster NGC~2516.  These were previously identified as photometric members by Irwin et al. (2007) and have measured rotation periods. The RVs of cluster members are tightly bunched about the mean showing an intrinsic dispersion of $0.65\\pm0.17$~km\\,s$^{-1}$.  The projected equatorial velocities show an increase in the proportion of fast rotators for later spectral types, such that 90 per cent of M4 stars have $v\\sin i > 15$~km\\,s$^{-1}$. Fewer stars are rapid rotators at earlier spectral types, but still a much greater proportion than are found in a field star sample, which is expected because the cluster is younger than the probable spindown timescales for all M-dwarfs.  We have gauged chromospheric activity using intermediate resolution spectroscopy of the first two of the near infrared calcium triplet lines (8498\\AA\\ and 8542\\AA) and a spectral subtraction technique to remove the photospheric contribution. The chromospheric activity of stars that are hotter or cooler than spectral type M2.5 ($(V-I)_0 \\simeq 2.3$) show a differing dependence on rotation period and colour/mass. Our main findings are: (i) Stars with spectral type earlier than M2.5 behave like other samples of young G- and K-type stars. Their chromospheric activity increases with decreasing period, or decreasing Rossby number, and reaches a saturated plateau for Rossby numbers smaller than about 0.1. (ii) Cooler stars almost all rotate fast enough that they have Rossby numbers less than 0.1. For these stars we find almost no dependence on rotation period or Rossby number, but a rather steep decline in chromospheric activity with colour, amounting to factors of 2--3 between spectral types M2.5 and M4 ($2.3 < (V-I)_0 < 2.9$). (iii) There is no evidence in our data for any systematic fall in activity at very low Rossby numbers ($\\la 0.01$), a phenomenon that has been seen in the coronal emission from fast-rotating G- and K-stars and dubbed ``supersaturation''.  It is tempting to identify changes in the properties of chromospheric activity with the disappearance of the radiative core of a star at $\\simeq 0.35\\,M_{\\odot}$ and presumably the emergence of a new, distributed or turbulent dynamo that operates in fully convective M-dwarfs. We do not see any corresponding decline in peak levels of X-ray emission across the fully convective boundary in NGC~2516 and none has been reported in field stars. This, combined with literature suggesting that magnetic flux continues to be generated strongly in fully convective stars leads us to favour a changing magnetic topology as the cause of both the decline in chromospheric emission and the rapid increase in angular momentum loss timescales as stars approach or cross the fully convective boundary."
    },
    "1004/1004.3231_arXiv.txt": {
        "abstract": "A constraint on the viable $f(R)$ model is investigated by confronting theoretical predictions with the multipole power spectrum of the luminous red galaxy sample of the Sloan Digital Sky survey data release 7. We obtain a constraint on the Compton wavelength parameter of the $f(R)$ model on the scales of cosmological large-scale structure. A prospect of constraining the Compton wavelength parameter with a future redshift survey is also investigated. The usefulness of the redshift-space distortion for testing the gravity theory on cosmological scales is demonstrated. ",
        "introduction": "Experimental tests of gravity on the scale of the Solar System show good agreement with predictions of general relativity (e.g., \\cite{Adelberger}). The nature of the Newtonian gravity is the attractive force, which naturally predicts a decelerated expansion of the universe. Contrary to this expectation, it has been discovered that our universe is undergoing an accelerated expansion epoch \\cite{Riess98,Perlmutter99,WMAP03}. Though the accelerated expansion is explained by introducing a cosmological constant, its small but nonzero value cannot be explained naturally \\cite{Weinberg}. The problem might be deeply rooted in the nature of fundamental physics.  This problem has attracted many researchers, and many works have been done, both theoretically and observationally. As a generalisation of the cosmological constant, a dynamical field, called the dark energy model and its variants, are proposed to explain the accelerated expansion of the universe (see \\cite{Peebles} and references therein). As an alternative to the dark energy model, modification of gravity may explain the accelerated expansion. General relativity is not considered to be the complete theory, because its quantum theory cannot be formulated in a well defined manner. The theory of gravity might need to be reformulated within a more general framework.  From the observational point of view, the constraint on the gravity theory on cosmological scales has not been well investigated, compared with the constraint on the scales of the Solar System. Many future projects to produce large galaxy surveys are in progress or planned \\cite{DETF06,sdss3,sumire,lsst,ska,Robberto}, which aim to explore the nature of the dark energy. These surveys are useful for testing the theory of gravity at cosmological scales (e.g., \\cite{Yamamoto06}). The dynamical dark energy models may have similar expansion rates as models of modified gravity, but predict different histories for the growth of structures. The key to testing the gravity theory is the measurement of the evolution of cosmological perturbations, as many authors have concluded recently \\cite{Uzan,Shirata,Sealfon, Linder05,Ishak06,UzanII, CahnLinder,Yamamoto07,Huterer07,Kunz07,Spergel08,Song09}.  The cosmic microwave background anisotropies are useful for investigating the cosmological perturbations through the measurements of the integrated Sachs-Wolfe effect or the lensing effect on the angular power spectrum \\cite{Smoot}. Imaging surveys of galaxies are also useful through the weak lensing statistics or cluster number counts \\cite{JainTakada,Bean}. Similarly, redshift surveys of galaxies are helpful for testing gravity \\cite{Linder,Guzzo,Yamamoto08,WhitePercival,Peacock09,Reyes}. In the present paper, we revisit the problem of testing the gravity theory through a measurement of the multipole power spectra in the sloan digital sky survey (SDSS) luminous red galaxy (LRG) sample \\cite{Yamamoto08}. Measuring the multipole power spectra  is a way to estimate the redshift-space distortions, which reflects the linear growth rate of the matter density perturbations \\cite{Cole,Hamilton,YBN}.  Many authors have investigated the clustering nature of the SDSS LRG sample \\cite{Eisenstein,Hutsi,PercivalI,PercivalII, Tegmark,PercivalIII,Okumura,HutsiII,Cabre,Sanchez}. In the references \\cite{Percival2009,Reid2009}, recent results on LRGs from the SDSS data release (DR) 7 are reported. In the reference \\cite{Cabre}, a test of gravity is considered using the observed anisotropic correlation function. Three of the authors of the present paper have shown that the SDSS LRG sample is useful to test the gravity theory by measuring the quadrupole power spectrum of galaxy distribution, which represents the redshift-space distortions \\cite{Yamamoto08}. In the present paper, we revisit the issue of testing the gravity theories on the cosmological scales using the SDSS LRG sample of the DR 7, especially focusing on the $f(R)$ gravity model.   The $f(R)$ models proposed in \\cite{HuSawicki,Starobinsky,Tsujikawa,Nojiri} are viable models of modified gravity, which include some function of the Ricci scalar, $f(R)$, added to the Einstein Hilbert action. As the modification of gravity involves the introduction of extra degree of freedom in general, one must be careful with the resulting behaviour. Furthermore, any theory must reduce to the general relativity on the scales of the Solar System. In the $f(R)$ model, the general relativity is supposed to be recovered by the {\\it chameleon mechanism} \\cite{KhouryWeltman,Mota}, which hides the field of the extra degree of freedom because the mass of the field becomes large for a dense region. The cosmological bounds on the $f(R)$ model have been investigated with the cosmic microwave background anisotropies \\cite{SPH} and also using the abundance of galaxy clusters \\cite{Schmidt}. However, our approach is based on the redshift-space distortion \\footnote{After we have completed this manuscript, we noticed of the paper by Girones et al.~\\cite{Girones}, in which the similar cosmological constraint on the $f(R)$ model is investigated.}.   This paper is organised as follows: In section 2, we briefly review the $f(R)$ model and the characteristic evolution of the matter density perturbation. In section 3, we present our results for the multipole power spectrum of the SDSS  LRG sample of the DR 7. In section 4, cosmological constraint is discussed by confronting the observed multipole spectra with the theoretical predictions. In section 5, a prospect of constraining the $f(R)$ model is discussed on the basis of the Fisher matrix analysis, assuming a future large redshift survey. Section 6 is devoted to summary and conclusions. Throughout this paper, we use units in which the velocity of light equals 1, and adopt the Hubble parameter $H_0=100h {\\rm km/s/Mpc}$ with $h=0.7$. ",
        "conclusions": "In this paper, we determined a cosmological constraint on the viable $f(R)$ model based on the redshift-space distortion by measuring the monopole and quadrupole spectra of the SDSS LRG sample of DR7. The monopole and the quadrupole spectra are used to fit the bias parameters and to constrain the growth factor and the growth rate of the density perturbations, which depend on the Compton scale of the $f(R)$ model.  Our results show that short Compton wavelength model fits the data better, while the long Compton wavelength model is excluded, though the constraint depends on the evolution parameter $n$. For the case $n=1/2$, our constraint is similar to that from the cluster number counts reported in \\cite{Schmidt}. When we adopt more accurate model for the redshift-space power spectrum \\cite{Jennings}, the constraint becomes slightly weaker. However, the long Compton wavelength case of the $f(R)$ model with the smaller value of $n$ is excluded. Our results exemplify that the redshift-space distortion is quite useful in testing gravity theory. We also demonstrated that a future redshift survey like the WFMOS/SUMIRE is potentially useful in obtaining a constraint on the Compton wavelength scale.  We acknowledge that the widely used theoretical model of the anisotropic power spectrum adopted in the present paper might need careful improvements. We adopted the Peacock and Dodds formula and the Smith formula for the nonlinear modelling of the mass power spectrum. Our results do not significantly depend on the choice. However, there might be a need to adopt a more sophisticated formula for the precise nonlinear modelling within the framework of the modified gravity, as demonstrated by Koyama, Taruya, Hiramatsu \\cite{KTH}. The treatment of the Finger of God effect in our paper was simple, which assumed the exponential distribution function for the pairwise velocity and introduced one free parameter -- the pairwise velocity dispersion. In reality it might not be an adequate model to describe the nonlinear region of the redshift-space power spectrum~\\cite{Scoccimarro}. We checked the reliability of our results by adopting the other possible model proposed in Ref.\\cite{Jennings}, extensively applying the fitting formula to the $f(R)$ model, whose accuracy in this case, however, has not been demonstrated. We found that there is a non-negligible effect on the constraint on the $f(R)$ model. Therefore, a more precise modelling of the redshift-space power spectrum should arguably be needed in the future. Concerning the modelling of the clustering bias, we adopted a simple scale-dependent bias. Here too there is potentially a lot of room for improvement. These issues are out of scope for the present paper, but need to be elaborated for a precise test of gravity with the future redshift surveys.    \\vspace{0.3cm} {\\it Acknowledgement} We thank T.~Kobayashi, J.~Yokoyama, H.~Motohashi, T.~Nishimichi, S.~Saito, A.~Taruya and M.~Takada, and Y.~Suto for useful comments and discussions. This work was supported by Japan Society for Promotion of Science (JSPS) Grants-in-Aid for Scientific Research (Nos.~21540270,~21244033). This work is also supported by JSPS Core-to-Core Program ``International Research Network for Dark Energy''. T.N. acknowledges support by a research assistant program of Hiroshima University.    \\vspace{0.3cm}"
    },
    "1004/1004.1002_arXiv.txt": {
        "abstract": "We present the correlations between the spectroscopic metallicities and ninety-three different intrinsic colors of M31 globular clusters, including seventy-eight BATC colors and fifteen SDSS and near infrared $ugriz$K colors. The BATC colors were derived from the archival images of thirteen filters (from $c$ to $p$), which were taken by Beijing-Arizona-Taiwan-Connecticut (BATC) Multicolor Sky Survey with a 60/90 cm f/3 Schmidt telescope. The spectroscopic metallicities adopted in our work were from literature. We fitted the correlations of seventy-eight different BATC colors and the metallicities for 123 old confirmed globular clusters, and the result implies that correlation coefficients of twenty-three colors $r>0.7$. Especially, for the colors $(f-k)_0$, $(f-o)_0$, and $(h-k)_0$, the correlation coefficients are $r>0.8$. Meanwhile, we also note that the correlation coefficients ($r$) approach zero for $(g-h)_0$, $(k-m)_0$, $(k-n)_0$, and $(m-n)_0$, which are likely to be independent of metallicity. Similarity, we fitted the correlations of metallicity and $ugriz$K colors for 127 old confirmed GCs. The result indicates that all these colors are metal-sensitive ($r>0.7$), of which $(u-$K$)_0$ is the most metal-sensitive color. Our work provides an easy way to simply estimate the metallicity from colors. ",
        "introduction": "\\label{sect:intr}  Globular clusters (hereafter GCs) are the oldest bound stellar systems in the Milky Way and other galaxies. They provide a fossil record of the earliest stages of galaxy formation and evolution. Located at a distance of $\\sim780$ kpc \\citep{sg98, mac01} away, Andromeda (M31) is the nearest large spiral galaxy in our local group. According to the latest RBC v.4.0 \\citep{gall04,gall06,gall07}, M31 contains more than one thousand GCs and candidates. Therefore M31 GCs provide us an ideal laboratory for us to study the nature of extragalactic GCs and to better understand the formation and evolution of M31. Furthermore, since M31 is Sb-type, which is similar to our Galaxy, it may offer us the clues to formation and evolution history of our Galaxy.  The correlation of metallicity and colors of globular clusters has been previously studied by many authors: \\citet{bh90} fitted the correlations of metallicities and ($V-$K), ($J-$K) for Galactic GCs and applied the calibrated relation to M31 GCs to derive the metallicities of M31 GCs. \\citet{barmby00} analyzed the correlations of ten intrinsic colors ($(B-V)_0$, $(B-R)_0$, $(B-I)_0$, $(U-B)_0$, $(U-V)_0$, $(U-R)_0$, $(V-R)_0$, $(V-I)_0$, $(J-$K$)_0$, $(U-$K$)_0$) and metallicity of 88 Galactic GCs. They find that the correlation coefficients \\emph{r} for all these ten colors range from 0.91 to 0.77, of which $(U-R)_0$ is the best metallicity indicator while $(V-R)_0$ is not so good. The authors also studied the $(J-$K$)_0$ and $(V-$K$)_0$ color-metallicity  correlations for M31 and Milky Way globular clusters. In this work, we attempt to find the most metal-sensitive colors, which could be used for the derivation of the metallicities of GCs in the future work. It is well known that, spectroscopy is relatively expensive in terms of observational time and is rarely available for many extragalactic globular clusters. Hence it is useful to determine metallicities from photometry.  The studies on identification, classification and analysis of the M31 GCs have been started since the pioneering work of \\citet{hub32}, \\citep[see e.g.][]{vet62,sar77,batt80,batt87,batt93,cra85,barmby00}. These studies provided a large amount of photometric data in different photometric systems, i.e. CCD photometry, photoelectric photometry, and photographic plates, even visual photometry. In order to obtain homogeneous photometric catalogue of M31 GCs, \\citet{gall04, gall06, gall07} took the photometry of \\citet{barmby00} as reference and transformed others to this reference and compiled the famous The Revised Bologna Catalogue of M31 GCs and GC candidates (The latest version is RBC v.4.0), where 654 confirmed GCs and 619 GC candidates of M31 have been included, and 772 former GC candidates have been proved to be stars, asterism, galaxy, region HII or extended cluster. The catalogue also includes the newly discovered star clusters from \\citet{mac06}, \\citet{kim07} and \\citet{hux08}. \\citet{mac06} reported the discovery of eight remote GCs in the outer halo of M31 by using the deep ACS images; \\citet{kim07} found 1164 GCs and GC candidates in M31 with KPN 0.9 m telescope and the WIYN 3.5 m telescope, of which 559 are previously known GCs and 605 are newly found GC candidates; later, \\citet{hux08} detected 40 new GCs in the halo of M31 with Isaac Newton Telescope and Canada-France-Hawaii Telescope data. Recently, \\citet{cald09} presented a new catalogue of 670 likely star clusters, stars, possible stars and galaxies in the field of M31, all with updated high-quality coordinates being accurate to $0.2''$ based on the images from the Local Group Survey \\citep{massey06} or Digitized Sky Survey (DSS). Very recently, \\citet{peac10} identified M31 GCs with images from Wide Field CAMera (WFCAM) on the UK Infrared Telescope and SDSS archives, and performed the photometry for them in SDSS $ugriz$ and infrared $K$ bands. Furthermore, the authors combined all the identifications and photometry of M31 GCs from references and those of their new work and updated the M31 star cluster catalog, including 416 old confirmed clusters, 156 young clusters and 373 candidate clusters. Our work is based on the 416 old confirmed M31 GCs, from which we will select our sample GCs.  Studies of M31 GCs based on BATC observations are also numerous: \\citet{jiang03} presented the BATC photometry of 172 GCs in the central $\\sim 1$ deg$^2$ region of M31 and estimated the ages for them with the Simple Stellar Population (SSP) models. After that, \\citet{fan09} added six more new BATC observations fields surrounding the central region and performed the photometry for thirty GCs, which did not have the broadband photometry before. With the photometry, the authors, for the first time, suggested the blue tilt of M31 GCs. \\citet{ma09} fitted the ages of thirty-five GCs of the central M31 field which were not included in \\citet{jiang03} with BATC, 2MASS and GALEX data and the SSP models. Later, \\citet{wang10} performed the photometry for another 104 GCs of M31 and estimated the ages by fitting the SEDs with SSP models, revealing the existence of young, the intermediate-age and the old populations in M31.  In this paper, first we used the BATC data to analyze the correlations of spectroscopic metallicities and the colors of M31 GCs. Moreover, we did the similar work with the $ugriz$K band data. The paper is organized as follows: Sect. \\ref{sect:data} describes the data utilized in our work, including the spectroscopic metallicity from literature, BATC archival images and the data reduction, the $ugriz$K photometry from literature, along with the reddening of M31 GCs. Sect. \\ref{sect:ana} present the analysis and the results on the correlations of the metallicities and intrinsic colors for M31 GCs. The summary and remarkable conclusions of our work are in Sect. \\ref{sect:sum}. ",
        "conclusions": "\\label{sect:sum}  Our work provides the most metal-sensitive colors of M31 GCs which could be utilized for roughly estimating the metallicity of GCs from their colors. The observational spectroscopic metallicities applied in our analysis were from \\citet{per02} and the colors are derived from BATC photometry and \\citet{peac10}.  First, we performed the photometry with BATC archival images from $c$ to $p$ band, covering the wavelength from $\\sim$ 4,000 to $\\sim$ 10,000 {\\AA} and obtained the BATC colors for 123 confirmed old GCs. The reddening values for the color corrections are from \\citet{fan08} and \\citet{barmby00}. For the thorough study, all the seventy-eight different BATC colors are used for the correlation analysis. Our fitting result shows that correlation coefficients of 23 colors $r>0.7$. Especially, for the colors $(f-k)_0$, $(f-o)_0$, $(h-k)_0$, the correlation coefficients $r>0.8$. Meanwhile, the correlation coefficients approach zero for $(g-h)_0$, $(k-m)_0$, $(k-n)_0$, $(m-n)_0$, which can be explained by lacking absorption lines.  Further, we also fitted the correlations of metallicity and $ugriz$K colors for the 127 old confirmed GCs with the photometry from \\citet{peac10}. Similarly, all the fifteen colors are utilized for analysis. The fitting result indicates that all these colors are metal-sensitive with correlation coefficients $r>0.7$. In particular, $(u-$K$)_0$ is the most metal-sensitive color while $(i-z)_0$ is the metal-nonsensitive color."
    },
    "1004/1004.1528_arXiv.txt": {
        "abstract": "Galaxy harassment is an important mechanism for the morphological evolution of galaxies  in  clusters. The spiral galaxy NGC~4254 in the Virgo cluster is believed to be a harassed galaxy. We have analyzed the  power spectrum of HI emission fluctuations from NGC~4254 to investigate  whether it carries any imprint of galaxy harassment. The power  spectrum, as determined using the $16$ central  channels which contain  most of the HI emission,  is found to be well fitted by a  power law $P(U)=AU^{\\alpha}$  with  $\\alpha\\ =-\\ 1.7\\pm 0.2$ at length-scales $1.7 \\, {\\rm k pc}$ to  $ 8.4 \\, {\\rm kpc}$. This is similar to other normal spiral galaxies which  have a  slope of $\\sim -1.5$ and  is interpreted as arising from two dimensional turbulence at length-scales larger than the galaxy's scale-height. NGC~4254 is hence yet another example of a spiral galaxy that exhibits  scale-invariant density fluctuations out to length-scales  comparable to the diameter of the HI disk. While a large variety of possible energy sources like proto-stellar winds, supernovae, shocks, etc.  have been proposed to produce turbulence, it is still to be seen whether  these are effective  on length-scales comparable to that of the entire HI disk. On separately analyzing the HI power spectrum in different  parts of NGC~4254, we find that the outer parts have a different slope  ($ \\alpha = -2.0\\pm0.3$)  compared to the central part of the galaxy ($\\alpha = -1.5\\pm0.2$). Such a change in slope is not seen in other, undisturbed galaxies. We suggest that, in addition to changing the overall morphology, galaxy harassment also effects the fine scale structure of the ISM, causing the power spectrum to have a steeper  slope in the outer parts. ",
        "introduction": "\\label{sec:intro} Galaxy harassment (frequent high speed galaxy encounters, \\citealt{MKL96}) is believed to be an important process in driving the morphological transformation of spiral galaxies to ellipticals inside clusters. Typically, the first encounters convert a normal spiral galaxy to a disturbed spiral with dramatic features drawn out from the dynamically cold gas.  The spiral galaxy NGC~4254, located in the nearby Virgo cluster, is found to have a tail \\citep{M05} with neutral hydrogen (HI) mass $2.2 \\times 10^8 \\ M_{\\odot}$ within a distance of $120 \\ {\\rm kpc}$ from the galaxy. This gaseous tail, without any  stellar counterpart, is believed to have been produced by an act of galaxy harassment \\citep{HGK07}. Each act of harassment  has the potential to induce a burst of star formation and to  change the internal properties of the galaxy, including the properties of the inter stellar medium (ISM).  Here we study the effect of the harrasment on the fine scale structure of the ISM. We use the power spectrum of HI intensity fluctuations to quantify the fine scale structure.  Power spectrum analysis of the HI $21$-cm intensity fluctuation has been  widely used to probe the ISM \\citep{CD83,GR93,SBH98,SSD99,DD00,EK01,ES04I, ES04II}. These studies, mainly of our Galaxy and its satellites,   find a  power law  power spectrum $P(U)=A \\ U^{\\alpha}$ with index $\\alpha$ in the range $-2.5$ to $-3.0$.  This scale invariant power spectrum is interpreted as arising from  turbulence in the ISM.  Recently \\citet{AJS06} have presented a visibility based formalism for determining the power spectrum of  galaxies with extremely weak HI emission. This has been applied to a sample of  dwarf and   spiral galaxies (\\citealt{DBBC08}; \\citealt{DBBC09a}; \\citealt{DBBC09b}). These studies indicate a dichotomy in $\\alpha$ with values $\\sim -1.5$ in some galaxies and $\\sim -2.5$ in others.  The two values $\\sim -1.5$ and $\\sim -2.5$ are interpreted as arising respectively from two (2D) and three (3D) dimensional turbulence.  Some support for this interpretation has been provided by analysis of the fluctuation power spectrum for NGC~1058 where a transition from 2D to 3D turbulence is observed at an angular scale corresponding to the scale height of the galaxy's disk.  NGC~4254 the galaxy whose power spectrum we estimate in this paper, is a  lopsided spiral galaxy (morphological type SA(s)c), with  an inclination of $\\sim 42^{\\circ}$ \\citep{PH93}. The galaxy is located  at a distance of 1 Mpc from the core of Virgo cluster and is  believed to be falling into the cluster with a relative velocity of $1300 \\ {\\rm km\\ s}^{-1}$ \\citep{VH05}. The distance to this galaxy is estimated to be 16.7 Mpc \\citep{M07}; at this distance 1$^{\\prime\\prime}$ corresponds to 81 pc. \\begin{figure} \\begin{center} \\includegraphics[angle=0,width=2.8in]{fig1.eps} \\end{center} \\caption{ The $6.5^{\\prime\\prime} \\times 8.0^{\\prime \\prime}$ resolution integrated HI column  density map of NGC~4254 (contours) is overlayed with the optical image of the galaxy (density).  The contours levels are  3., 5., 10., 20., 25., 35., 40. and 45. $\\times 10^{20}$ atoms cm$^{-2}$. } \\label{fig:mom0} \\end{figure} ",
        "conclusions": "\\label{sec:discuss} Figure~\\ref{fig:powsp24} shows the real and imaginary parts of $\\hat{\\rm P}_{\\rm  HI}(\\vU)$ evaluated using $16$ channels from $27$ to $42$ which have relatively high HI emission. Here the estimator  $\\hat{\\rm P}_{\\rm  HI}(\\vU)$ was separately evaluated for each channel and then averaged over all $16$ channels to increase the signal to noise ratio. As expected  from the theoretical considerations (\\citealt{AJS06}), the imaginary part is well suppressed compared to  the real part.  We also estimate $\\hat{\\rm P}_{\\rm  HI}(\\vU)$ using  the line free  channels to test  for  any  contribution from the residual continuum. This is found to be much smaller than the signal (Figure~\\ref{fig:powsp24}) indicating that the continuum has been adequately subtracted out.  The power law $P_{\\rm HI}(U)=AU^{\\alpha}$, with  $\\alpha\\ =-\\ 1.7\\pm 0.2$  is found to give a good fit to the  HI power  spectrum for the $U$ range $2.0 \\, {\\rm k}  \\lambda$ to $10.0 \\, {\\rm k} \\lambda$ (Table~\\ref{table:t1}).  The best-fit power-law and the $1-\\sigma$ error-bar on $\\alpha$ were determined using   $\\chi^2$ minimization  \\citep{AJS06,DBBC09b}. We use $D=16.7 \\ {\\rm Mpc}$ (distance to the   galaxy)  to convert a baseline $U$ to a length-scale $D/U$ in the plane of the galaxy's image.  The $U$ range  $2.0 \\, {\\rm k}  \\lambda$ to $10.0 \\, {\\rm k} \\lambda$ corresponds to  the range of length-scales  $8.4$ to $1.7 \\ {\\rm kpc}$.  \\citet{DBBC09a} have broadly classified the observed turbulence   in their galaxy sample   as 2D turbulence at the large scales in the plane of the galaxies disk and 3D turbulence at scales smaller than the scale height of the disk. They also show that the slope of the power spectrum changes from $\\sim -1.5$ to $\\sim -2.5$ as one goes from 2D to 3D turbulence. For the present case the largest  length-scale   ($8.4 \\, {\\rm kpc}$) is  definitely larger than the  typical HI scale heights within the Milky-Way \\citep{LH84,WB90} and in  external spiral galaxies (e.g. \\citealt{narayan}). It is thus quite reasonable to conclude that   the slope  $\\alpha = -1.7 \\pm 0.2$ is  of 2D turbulence  in the plane of the galaxy's disk. The fact that we do not observe the transition to 3D turbulence, which is expected   to occur at a baseline $U=D/\\pi  z_h$ \\citep{DBBC09a},   allows us to place an upper limit on the galaxy's scale height $z_h \\le 2.6\\ {\\rm  k pc}$.  Our present observation is another confirmation (Dutta et al 2008, 2009a 2009b) of the fact that spiral galaxies exhibit scale-invariant density fluctuations that extend to length-scales of  $\\sim 10 \\, {\\rm kpc}$  (eg. NGC~628, NGC~1058) which is comparable to the diameter of the HI disk. While a large variety of possible energy sources like proto-stellar winds, supernovae, shocks, etc. have been proposed to drive turbulence  \\citep{ES04I},  it is still to be seen whether these are effective  on length-scales as  large as $10 \\, {\\rm kpc}$.   \\begin{figure} \\begin{center} \\includegraphics[angle=0,width=6in]{fig3.eps} \\caption{Integrated HI column density maps of the galaxy NGC~4254 using data cubes A, B and C. Note the diagonal movement of the centroid of emission from North east (A) to South west (C). The contours levels are  5., 8. and 12. $\\times 10^{20}$ atoms cm$^{-2}$.} \\label{fig:abc} \\end{center} \\end{figure}   \\begin{figure} \\begin{center} \\includegraphics[angle=-90,width=2.8in]{fig4.eps} \\caption{HI power spectrum for data cubes A, B and C plotted with arbitrary offsets to prevent them from overlapping. } \\label{fig:aps} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[angle=0,width=6in]{fig5.eps} \\caption{Integrated HI column density maps of the galaxy NGC~628 using data cubes D, E and F. Note the diagonal movement of the centroid of emission from North east (D) to South west (F). The contours levels are  6.,  12. and 18. $\\times 10^{20}$ atoms cm$^{-2}$. } \\label{fig:DEFM} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\includegraphics[angle=-90,width=2.8in]{fig6.eps} \\caption{Power spectra of the HI emission for the galaxy NGC~628 is shown for D (channels 108-119), E (channels 120-131) and F (channels 132-143) with  arbitrary offsets to prevent them from overlapping. Details of the data used for this estimation can be found in \\citet{DBBC08}} \\label{fig:628PS} \\end{center} \\end{figure}    Galaxy harassment is expected  to have  different effects on the inner and outer parts of the galaxy. While gas is stripped from the outer parts, the inner part looses angular momentum and gradually collapses to the center through repeated galaxy encounters. We can selectively study different parts of NGC~4254, whose rotation axis is tilted at  $42^{\\circ}$  to the line of sight, by considering different velocity channels. Our analysis till now has used only the central $16$ channels, we now use the central $24$ channels ($23-46$) for the subsequent analysis.  We  construct $3$  different data cubes namely A, B and C containing channels $23-30$, $31-38$ and $39-46$ respectively. We can now separately probe the North east,  central and South west parts of the galaxy (Figure~\\ref{fig:abc}) using these three data cubes. For each data cube,  we  evaluate the HI power spectrum using individual channels and then average over the channels in the respective cubes. We estimated $U_m$ seperately for A, B and C and the best fit power law is obtained for $U \\ge U_m$ only. We find that for each  data cubes the HI power spectrum is well fit by a power law (Figure~\\ref{fig:aps}), the details being shown in Table~\\ref{table:t1}. We find that the slope $\\alpha$ is $-2.0\\pm 0.3$ for A and C which probe the outer parts of the disk while it is  $-1.5\\pm 0.2$ in B which probes the central region. To verify  that this change in slope is due to harrassment, we also consider the HI power spectrum of the  spiral galaies NGC~628 and NGC~1058 \\citep{DBBC08,DBBC09b} which are not undergoing galaxy harassment.  Figure \\ref{fig:DEFM} shows  the regions  of the spiral galaxy    NGC~628  corresponding to  each of the channel range 108-119 (D),  120-131 (E) and  132-143 (F).  We find that the power spectra of these three data cubes (Figure~\\ref{fig:628PS}) all  have the same slope $\\sim -1.6$,  which is also similar to the slope of the central part of NGC~4254. The results are similar for NGC~1058 and hence we do not explicitly show these here. Based on this we conclude that the difference in slope between the inner and outer parts of NGC~4254 is a consequence of galaxy harassment.  Not only does galaxy harassment affect the global morphology of the galaxy, it also affects the fine scale structure in the ISM as reflected by  HI power spectrum.   We note that NGC~4254 have an inclination of $42^{\\circ}$, on the other hand NGC 628 and NGC~1058 are more face-on galaxies (inclination $\\sim 10^{\\circ}$). In our analysis we selectively study different parts of a galaxy by considering different velocity channel ranges. However, since face-on galaxies do not have much range in radial velocity from the rotation curve, the spatial extent of HI in NGC~628 and NGC~1058 will not be very different for the three velocity channel ranges, which is unlike the case for NGC~4254. Hence it is likely that we failed to find a change of slope in NGC~628 and NGC~1058 because of this effect of inclination. However, as seen in Figure~\\ref{fig:DEFM} for NGC~628, all three data cubes D, E and F have a significant contribution of HI from the outer region of NGC~628, leading us to believe that the power spectrum has a slope of $\\sim\u22121.6$ in the outer parts of NGC~628. This is similar to the value of slope seen in the central parts of NGC~4254 and significantly different from the slope in the outer parts of NGC~4254. Hence it indicates an impact of harassment in NGC~4254. Further analysis of spiral galaxies with large inclination angles would possibly be able to resolve this issue.   We currently do not have  an understanding of how galaxy harassment caused a steepening of the HI power spectrum in the outer parts of the galaxy.  Theoretical modeling and the analysis of other Virgo cluster spiral galaxies are  needed for further progress in this direction.  \\begin{table} \\centering \\begin{tabular}{|c|c|c|c|c|c|} \\hline Data & Channels & $U_{m}$ & $U$ range & $\\alpha$   \\\\ & & (k $\\lambda$) & (k $\\lambda$) & \\\\ \\hline \\hline 16 central  & $27-42$ & $1.8$ & $2.5 -10.0$ & $-1.7\\pm 0.2$ \\\\ channels &&&& \\\\ \\hline A & $23-30$ & $1.3$ & $2.0 -10.0$ & $-2.0\\pm 0.3$ \\\\ B & $31-38$ & $2.2$ & $2.5 -10.0$ & $-1.5\\pm 0.2$ \\\\ C & $39-46$ & $1.3$ & $2.0 -10.0$ & $-2.0\\pm 0.3$ \\\\ \\hline \\hline \\end{tabular} \\caption{Results of the power spectrum analysis of NGC~4254 for different data cubes spanning over different velocity ranges.} \\label{table:t1} \\end{table}"
    },
    "1004/1004.3862.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatoryms_Dec20.tex  \\abstract % context heading (optional) % {} leave it empty if necessary {The second parameter (the first being metallicity) defining the distribution of stars on the horizontal branch (HB) of globular clusters (GCs) has long been one of the major open issues in our understanding of the evolution of normal stars. Large photometric and spectroscopic databases are now available: they include large and homogeneous sets of colour-magnitude diagrams, cluster ages, and homogeneous data about chemical compositions from our FLAMES survey.} {We use these databases to re-examine this issue.} {We use the photometric data to derive median and extreme (i.e., the values including 90\\% of the distribution) colours and magnitudes of stars along the HB for about a hundred GCs. We transform these into median and extreme masses of stars on the HB, using the models developed by the Pisa group, and taking into account evolutionary effects. We compare these masses with those expected at the tip of the red giant branch (RGB) to derive the total mass lost by the stars.} {We find that a simple linear dependence on metallicity of this total mass lost describes quite well the median colours of HB stars. Assuming this mass loss law to be universal, we find that age is the main second parameter, determining many of the most relevant features related to HBs. In particular, it allows us to explain the Oosterhoff dichotomy as a consequence of the peculiar age-metallicity distribution of GCs in our Galaxy, although both Oosterhoff groups have GCs spanning a rather large range in ages. However, at least an additional - third - parameter is clearly required. The most likely candidate is the He abundance, which might be different in GC stars belonging to the different stellar generations whose presence was previously derived from the Na-O and Mg-Al anticorrelations. Variations in the median He abundance allow us to explain the extremely blue HB of GCs like NGC~6254 (=M~10) and NGC~1904 (=M~79); such variations are found to be (weakly) correlated with the values of the R-parameter (that is the ratio of the number of stars on the HB and on the RGB). We also show that suitable He abundances allow deriving ages from the HB which are consistent with those obtained from the Main Sequence. Small corrections to these latter ages are then proposed. We find that a very tight age-metallicity relation (with a scatter below 4\\%) can be obtained for GCs kinematically related to the disk and bulge, once these corrections are applied. Furthermore, star-to-star variations in the He content, combined with a small random term, explain very well the extension of the HB. There is a strong correlation between this extension and the interquartile of the Na-O anticorrelation, strongly supporting the hypothesis that the third parameter for GC HBs is He. Finally, there are strong indications that the main driver for these variations in the He-content within GCs is the total cluster mass. There are a few GCs exhibiting exceptional behaviours (including NGC~104=47 Tuc and in less measure NGC~5272=M~3); however, they can be perhaps accommodated in a scenario for the formation of GCs that relates their origin to cooling flows generated after very large episodes of star formation, as proposed by Carretta et al. (2009d). } {}  % ",
        "introduction": "Sandage \\& Wallerstein (1960) noticed that the distribution with colour/temperature of stars on the horizontal branch (HB) of globular clusters (GCs) is roughly correlated with their metal content. A few years later, this observation was explained by the first successful models of HB stars describing the effect of metal content on the efficiency of H-shell burning in low mass stars where He is burning in the core (Faulkner 1966). However, soon after this important theoretical achievement, van den Bergh (1967) and Sandage \\& Wildey (1967) pointed out that the correlation between colour/temperature and metallicity had several exceptions, a difficulty that has become known as {\\it the second parameter problem}. In the following forty years, a large number of tentative explanations for this discrepancy appeared in the literature, but no overall satisfactory scenario has yet been found. A proof of the large interest raised by this issue is that entering \"globular cluster\" and \"second parameter\" in the ADS data base\\footnote{http://adsabs.harvard.edu/} resulted in 231 abstracts (24 since 2006) with 6031 citations on a query made on April 15th, 2009. Of course, this search is probably incomplete, because there are many related issues, e.g., the Oosterhoff dichotomy in the mean periods of RR Lyrae in galactic GCs (Oosterhoff 1944; Sandage 1982), the UV upturn in the spectra of bulges and elliptical galaxies (Code 1969), the ages and the He abundances of GCs (Iben 1968), the mass loss law for low mass stars, which use different keywords. The second parameter problem is certainly one of the major open issues in our understanding of the evolution of normal stars. For reviews of this topic, we refer to Moehler (2001) and Catelan (2009).  There are various reasons why the second parameter issue has been insofar so difficult to solve. The most important is that there is most likely more than a single second parameter. The colour of HB stars is very sensitive to several physical stellar quantities in the age and metallicity regime typical of GCs (see e.g., Rood 1973; Renzini 1977; Freeman \\& Norris 1981; Fusi Pecci et al. 1993). Zinn (1980) and many authors after him convincingly showed that younger ages might explain the red colours of the HB of several of the outer halo GCs (see Dotter et al. 2008 for a similar line of thought). However, progress in the determination of the (relative) ages of GCs (e.g., Stetson et al. 1996; Rosenberg et al. 1999; De Angeli et al. 2005; Mar\\'in-Franch et al. 2009) demonstrated that this cannot be the only second parameter. The same result had previously been obtained even more directly from the broad distribution in colours of HB stars within some individual GCs (see e.g., Ferraro et al. 1990).  Since a change in the mass of stars on the HB may well cause even large variations in their colours (Rood 1973), a special mass-loss law has become a popular explanation (see Dotter 2008 for an example of this approach). Unfortunately, the physics of mass loss is very poorly known at present (see e.g., Willson 2000; Meszeros et al. 2009; Dupree et al. 2009). Many different mechanisms may affect mass loss (see e.g., Rich et al. 1997; Green et al. 1997; Soker \\& Harpaz 2000) and empirical evidence is inadequate for fully constraining them (see e.g., Peterson 1982; Origlia et al. 2007, 2008). Given these limitations, it is not yet possible to build an ab-initio model for the estimation of mass loss from GC stars. Hence we prefer to carefully restrict our assumptions: we looked for solutions with a mass loss law based on as few simple parameters as possible, for which we may obtain constraints from independent observations.  Additional second parameters considered included He abundances, the ratio of CNO to Fe abundances, stellar rotation or binarity (see e.g., Freeman \\& Norris 1981), and cluster concentration (Fusi Pecci et al. 1993). However, all of these explanations were found to be unsatisfactory overall. In the most successful cases, they might explain some groups of stars with anomalous colours on the HB (e.g., most field O-B subdwarfs are binaries, see e.g., Maxted et al. 2001; Napiwotzki et al. 2004; Han et al. 2003; however most of these in GCs seem to be single stars, see e.g., Moni Bidin et al. 2006). In the least successful cases, they are inconsistent with the data (CNO abundances: see the case of the second parameter pair NGC~362 and NCC~288: Shetrone \\& Keane 2000 and Catelan et al. 2001). The situations for rotation and He abundances are more complicated. For rotation, after the initial enthusiasm triggered by the pioneering observations of Peterson (1982), the problem was found to be less straightforward. More extensive data sets (Deliyannis et al. 1989, Behr et al. 2000a, 2000b; Behr 2003; Recio-Blanco et al. 2002) revealed an intricate pattern, so that it seems difficult to use direct observations of rotation along the HB to confirm its r\\^ole in the second parameter issue. While these observations do not rule out the possibility that rotation is indeed important, we cannot avoid concluding that the evidence for and against remains poor (see e.g., Sweigart 2002). We discuss the case of He abundances in subsequent sections.  Over the years, an enormous wealth of observational data has been collected, not only in terms of the distribution of stars along the HB with colour, but also the chemical composition, the period distribution of RR Lyrae, and other properties of GCs. Several authors (see e.g., Fusi Pecci et al. 1993; Catelan et al. 2001; Recio-Blanco et al. 2006; Carretta et al. 2007) pointed out the existence of correlations between global cluster parameters (such as luminosity, concentration, or Galactic orbit) and phenomena related to the second parameter. However, the mechanism connecting these global properties to the evolution of individual stars remained elusive until a few years ago.  A revised approach to the problem of the second parameter can now be developed. It is based on what was initially considered to be a separate characteristic of GCs, that is the abundance anomalies observed for GCs in the light elements CNO, Na, Mg, and Al (see Kraft 1994 and Gratton et al. 2004 for reviews of this topic). Early suggestions that there might be a correlation between these two sets of observations were made by Norris (1981), and more recently by Kraft (1994) and Catelan \\& de Freitas Pacheco (1995) and several other authors after them. However, the exact mechanism linking the two phenomena remained unclear. The two innovative steps taken subsequently were: \\begin{itemize} \\item the recognition that typical GCs host at least two generations of stars: this is required to explain the abundance anomalies observed for main sequence (MS) stars by Gratton et al. (2001) and Cohen et al. (2002). These observations contradicted the paradigm of GCs as single stellar populations, and opened a new view on GC formation and evolution that we are now only beginning to explore (see Gratton et al. 2004 for early results). \\item the understanding by D'Antona et al. (2002) that (large) variations in the abundance of He, which are expected to be correlated with the variations in Na and O and other elements, might result in large differences in the turn-off (TO) masses of stars of similar age: this is because He-rich stars evolve faster than He-poor ones and thus, at a given age, He-rich stars at TO are less massive. \\footnote{As pointed out by the referee, the knowledge that He abundance variations lead to variations in turn-off masses is much older (see Iben \\& Rood 1970). However, D'Antona et al. were the first to relate variations in Na, O, and other light elements to variations in He abundances, and then TO masses. This created the link between {\\it abundance anomalies} on the RGB and the second parameter issue. } Therefore, similar mass losses along the red giant branch (RGB) would lead to HB stars of very different masses, and hence colours. Not all authors estimated variations in the He abundances as large as those proposed by D'Antona et al. (2002; see for instance Marcolini et al. 2009). We note however that these small spreads in He abundance (which are usually justified on the basis of chemical evolution models) have difficulties in reproducing the observed spread in masses along the HB and the splitting of the MS of NGC~2808 (Lee et al. 2005, Piotto et al. 2007). \\end{itemize}  That a combination of age and He differences may explain the second parameter is very attractive for several reasons: (i) The large variations in He and the Na-O anticorrelation are explained by the presence of different generations of stars in GCs, and it is then easy to link them to general cluster properties, such as their mass or location in the Galaxy, which seem to play a r\\^ole in the second parameter issue; (ii) These different stellar generations may well be used to explain discontinuous and often discrete distributions of stars along the HB, such as that observed e.g., in NGC~2808 (D'Antona et al.  2005); (iii) Very accurate photometric data detect multiple main sequences in some GCs (Bedin et al. 2004; Piotto et al. 2007) that can only be explained by assuming large variations in the He content (Norris et al. 2004; Piotto et al. 2005; Milone et al. 2010).  While these observations are extremely interesting, they are rather limited in number: the data discussed insofar only concern a handful of massive GCs. A more comprehensive study of a large set of GCs, analysed in a homogeneous way, was lacking until recently. Such an analysis is now possible, thanks to the large databases of colour magnitude diagrams (CMD) and accurate ages provided by ground (Rosenberg et al. 2000a,b) and space observations (e.g., Piotto et al. 2002), and the extensive data on the Na-O anticorrelation from our FLAMES survey (Carretta et al. 2009a,b and references therein), complemented by literature data. %(Shetrone \\& Keane 2000; Sneden et al. 2004; Cohen \\& Melendez 2005). In this paper, we present an exploration of these unprecedented databases. In the first part of the paper, we consider the evidence provided by extensive photometric datasets, from both ground-based and HST observations for a sample of almost a hundred GCs, deriving the properties of HBs as defined by their median values and extension, and examining their correlations with metallicity and age. This analysis produces a simple mass-loss law, that explains the median colours of HB stars. However, as found by several authors before, an additional parameter is needed to explain the HB colours of GCs with an extreme blue HB (BHB) and the spread of colours in many other cases. In the second part of the paper, we consider variations in the He content as a possible explanation of these discrepancies, we derive the He abundance variations required explaining the observed properties of the HBs, and discuss the implications for MS photometry. In the third part of the paper, we search for additional evidence that He is indeed the additional parameter required to explain the HB morphology. In particular, we explore the correlations with other chemical anomalies, namely the Na-O anticorrelation. For this purpose, we consider a smaller but still quite large sample of 24 GCs, including classical second parameter cases (NGC~288 and NGC~362; NGC~5272 and NGC~6205), a list of blue HB clusters (such as NGC~1904 or NGC~6752), and the very extended HB cases (such as NGC~2808, NGC~6388 and NGC~6441). The aim of our discussion is to convince the reader that a combination of age and He abundance variations, the latter being related to multiple generations of stars within each GC, is a promising scenario to clarify most of the so far unexplained characteristics of the second parameter issue. We emphasise that we do consider that additional effects (e.g., binarity) may affect the colour of stars along the HB, but the r\\^oles played by the age and He abundance variations are probably dominant. ",
        "conclusions": "As we have seen, our reanalysis of public extensive photometric databases of GCs demonstrates that {\\it age is the main second parameter} affecting the HB morphology. This hypothesis is able to explain quite well most of the observables related to median HB stars of GCs, when coupled with a simple mass-loss law that is a linear function of [Fe/H]. Among the many observables that are successfully explained, we note the Oosterhoff dichotomy that we attribute to the peculiar age-metallicity distribution of Galactic GCs. Oo~{\\sc ii} clusters are mostly old, while Oo~{\\sc i} are predominantly young, although young Oo~{\\sc ii} and old Oo~{\\sc i} GCs exist. However, {\\it at least a third parameter is required} (and possibly even others) to fully explain the median colours of HBs (in particular those with very blue HBs) as well as their extension. There are various reasons to identify this third parameter with {\\it variations in the He content}. These include the variation in the scatter with metallicity, some correlation with the R-parameter, and the clear links with chemical anomalies observed in GCs. This result is strongly indicative of a possible link between the colours of the stars on the HB and their original composition, in a multiple generation scenario for the formation and early evolution of GCs. Self-pollution in GCs is possibly responsible for a large variety of the second parameter features, and may be in part described using the Na-O anticorrelation, although some modulation according to cluster luminosity is required.\\footnote{The scenario we propose should of course not only explain the Milky Way GCs, but e.g., those of Fornax (Buonanno et al. 1998). This case is quite puzzling, since clusters 1, 3, and 5 are nearly coeval, and have similar metallicity, and still have very different HB's. We did not quantify these variations in terms of mean colours and magnitudes as done for the Milky Way GCs considered in the present paper, hence we cannot provide any quantitative analysis. We only note that the ranking of HBR ratios for the three coeval clusters of similar metallicity (1, 3, and 5: -0.2, 0.50, and 0.44) is the same as the ranking in absolute magnitude $M_V$ (-5.32, -7.66, -6.82). A quantitative analysis is required to settle this issue.}  A combination of age and He variation therefore appears to be an explanation of the long-standing problem of the second parameter, although we do not exclude additional parameters such as the CNO abundances or the presence of binaries (e.g., of the progeny of blue stragglers) possibly playing some r\\^ole. However, this issue is still far from being completely settled. We need to make some progress in developing models, and a number of observational tests. A short list includes: \\begin{itemize} \\item Understanding the nature of the polluters. This requires advances in the modeling of AGB stars and rotating massive stars. Furthermore, detailed spectroscopic data for stars in massive and very young clusters, such as RSGC1 and RSGC2 (Davies et al. 2007), or intermediate age clusters in the LMC, where multiple MS TO's have been observed (Mackey et al. 2008; Milone et al. 2009), may provide a crucial test of this scenario. \\item The present discussion has focused on He, Na and O, but the abundances of other elements may also play an important r\\^ole. For instance, Al might be a proxy for He that is more reliable than Na. Unfortunately, our data for Al are not as extensive as those for Na and O, but the relation between the production of He (which most likely occurred within MS stars), and the proton capture processes (which might have occurred in the same main sequence stars, if massive and fast rotating, or later during the AGB phase, if the stars were of intermediate mass) must be clarified. \\item A number of confirmations of this scenario are required. These include (i) direct determinations of He, Na, and O in HB stars, which were shown to be possible in at least some cases by Villanova et al. (2009); and (ii) a discussion of the luminosity of the RGB bump that we defer to another paper currently in preparation. \\item In addition, we ask: do properties of RR Lyrae variables agree with expectations? Are anti-correlations found where expected (important clusters such as M~54 and NGC~1851 do not yet have adequate data)? Are multiple sequences observed where they are expected? We note here that while the connections between multiple MSs and variations in the He abundance is quite clear, the case of multiple sub-giant branches (SGB) is more ambiguous. SGB splitting measured using visual-red-near infrared colours (see e.g., Milone et al. 2008) might be due to a variation either in age or most likely total CNO content (see e.g., D'Antona et al. 2009 and Cassisi et al. 2008), or even [Fe/H] (in this case however some spread in the MS and RGB is also expected). Interpretation of splitting is even more ambiguous when considering ultraviolet colours, which are sensitive to N excesses. Variations in He abundances are only marginally effective in these cases, because sequences differing only in Y are very close each other on the SGB (D'Antona et al. 2002). Variations in total CNO content can be most likely attributed to the contribution of thermally pulsing AGB stars (Cassisi et al. 2008), which have a rather low mass and probably do not contribute much to He abundance variations. It is then unclear that there should be any correlation between the SGB splitting and large spreads on the HBs. In fact, NGC~2808 has a quite narrow SGB (Piotto et al. 2007). SGB splitting has been detected using visual-red-near infrared colours in 47 Tuc (Anderson 2009), NGC~1851 (Milone et al. 2008), and NGC~6388 (Moretti et al. 2009). These clusters have very different HB morphologies, ranging from very short (47 Tuc), to bimodal (NGC~1851), to very extended (NGC~6388). We obtain very different estimates of the He spread (0 for 47 Tuc; 0.048 for NGC~1851, a possibly too large value compared with those determined by Salaris et al. 2008, and Catelan et al. 2009b; 0.037 for NGC~6388). This lack of correlation suggests that the two phenomena are somewhat different, as expected if the mass range of the polluters changes from cluster-to-cluster. \\item The scenario requires a number of refinements. All analytic dependencies adopted throughout this discussion should be reviewed, and possibly replaced by a comparison with synthetic HBs. This may allow us to detect additional effects, e.g., a variation in total CNO abundances, not included in the present analysis. \\item Finally, hydrodynamical simulations of the formation and early evolution of massive star clusters are urgently needed. There are aspects of the present scenario that are necessary to explain observations, but should be understood more clearly. The most intriguing is the existence of a pool of gas from which second generation stars formed, which is composed of material processed through H-burning at high temperature diluted with pristine gas. How this pool of gas is generated, and how the stars form from it within the potential well of the young GC, while other stars of the original population evolve, remains unclear. Some explorative results were obtained by D'Ercole et al. (2008) which while very promising should be placed on a sounder basis. \\end{itemize}"
    },
    "1004/1004.1199_arXiv.txt": {
        "abstract": "We present a method that we developed to discern where the optical microvariability (OM) in quasars originates: in the accretion disk (related to thermal processes) or in the jet (related to non-thermal processes). Analyzing nearly simultaneous observations in three different optical bands of continuum emission, we are able to determine the origin of several isolated OM events. In particular, our method indicates that from nine events reported by \\cite{Ram09}, three of them are consistent with a thermal origin, three to non-thermal, and three cannot be discerned. The implications for the emission models of OM are briefly discussed. ",
        "introduction": "Optical microvariability (OM; variations with small amplitude in time scales from minutes to hours) in quasars can be a powerful tool to constrain the energy emission models of active galactic nuclei. Some studies indicate that OM does not depend on the radio properties of quasars (see \\citealt{de98}, hereafter PI; \\citealt{Stalin04}; \\citealt{Gupta05}; \\citealt{Ram09}, hereafter PII). However, the mechanism responsible for OM has not been characterized. Concerning this topic, variability studies in the optical bands acquired a particular relevance, because the optical/UV excess is identified with the emission from the accretion disk (e.g., \\citealt{Krishan94}; \\citealt{Lawrence05}; \\citealt{Pereyra06}; \\citealt{Li08}; \\citealt{Wilh08}). Unfortunately, all the proposed scenarios are very complex when the emission from additional spectral components is considered.  In photometric studies, the properties of the spectral energy distribution (SED) of quasars era characterized by the spectral index, which is defined as the slope of a curve in a plane $log(f_\\nu)~versus~log(\\nu)$, and calculated as magnitude differences at different bands (e.g., \\citealt{Masa98}; \\citealt{Treve02}). During variability events, changes in the shape of this SED can give us important information about the emission processes. For instance, evidence from variability studies strongly indicates that OM has non-thermal nature in both BL Lac objects and flat spectrum radio quasars (FSRQs; e.g., \\citealt{Damicis02}; \\citealt{Vagne03}; \\citealt{Villata04b}; \\citealt{Villata04a}; \\citealt{Hu06a}, \\citealt{Hu06b}, \\citealt{Hu07}), but in the case of FSRQs the contribution of a thermal component must be considered in the spectrum of these objects (e.g., \\citealt{Gu06}; \\citealt{Hu06b}; \\citealt{Ram04}). Concerning quasars, short-term variability (variations with amplitudes of tenths of a magnitude and time scales of weeks to months) might have thermal origin (e.g., \\citealt{Treve02}). Nevertheless, all these studies usually employ two bands, or only the average of color values is used, losing thus information about {\\it texture} of shortest variations (see \\citealt{Giveon99}; \\citealt{Webb00}; \\citealt{Treve02};  \\citealt{Vagne03}).  In order to discern where the OM in quasars originates, we assume that emission from the accretion disk must be related to thermal processes, while the emission from the jet must be related to non-thermal processes. During OM events, the spectral component responsible for the variation may display characteristic color changes. Then, we developed a method with quasi-simultaneous observations at three optical bands to analyze these color changes that accompanied microvariability events. This paper is organized as follows: in Section~\\ref{datos}, we shortly comment on the data treatment; in Section~\\ref{sed}, we define a spectral variability index and how to determine the OM origin; in Section~\\ref{results}, we analyze the data; finally, we discuss the results in Section~\\ref{discusion}; while a summary and conclusions are given in Section~\\ref{summary}. ",
        "conclusions": "\\label{summary}     The optical microvariations in RQQ and RLQ reported by \\citet{Ram09} were analyzed to determine whether their origin is thermal or non-thermal. We suppose that while thermal variations are related to the accretion disk, non-thermal variations are related to the relativistic jet.  With such a purpose, we developed a method to discriminate between color changes by thermal and non-thermal processes in quasars during an OM event. This method consists in modeling the optical broadband continuum of quasars considering thermal and non-thermal components, and comparing the observed color changes with spectral variations derived from this model. The main result of this work is the possibility to distinguish between thermal and non-thermal origin of an optical microvariation event in quasars.  We identified that some events are consistent with a thermal origin, while some others with a non-thermal one. In other cases either component can give an explanation to the detected variation. Thus, another important contribution of this research is that, analyzing the {\\it spectral microvariability}, we have found that the OM might be generated in either the disk or the jet, regardless of the radio classification of the quasars, CRLQ or RQQ. Additionally, some variations in RLQs are consistent with thermal OM and vice versa. Thus, our results indicate that microvariability in both quasar type might be originated under similar conditions.  In addition, although the continuum emission can be fitted in all cases by a unique non-thermal component, the broad band spectral variation may require to assume the presence of a second component, of thermal origin. The relative contribution of each component is an important parameter that should be taken into account for the description of spectral variability. Additional studies must be carried on to investigate the effects of these contributions on the type of quasar that we observe, and if this can be expended to the blazar class.  This result agrees with the discovery of small relativistic jets in RQQs (\\citealt{Blundell01}) and explains previous results (\\citealt{de98};  \\citealt{Ram09}). Although this method yields interesting results applied to our data, it is desirable to have a monitoring  with a better temporal resolution,  as well as to observe in more bands. More realistic models are also necessary."
    },
    "1004/1004.1685_arXiv.txt": {
        "abstract": "{\\it Swift} observations suggest that the central compact objects of some gamma-ray bursts (GRBs) could be newly born millisecond magnetars. Therefore, by considering the spin evolution of the magnetars against {\\it r}-mode instability, we investigate the role of the magnetars in GRB X-ray afterglow emission. Besides modifying the conventional energy injection model, we pay particular attention to the internal X-ray afterglow emission, whose luminosity is assumed to track the magnetic dipole luminosity of the magentars with a certain fraction. Following a comparison between the model and some selected observational samples, we suggest that some so-called ``canonical\" X-ray afterglows including the shallow decay, normal decay, and steeper-than-normal decay phases could be internally produced by the magnetars (possibly through some internal dissipations of the magnetar winds), while the (energized) external shocks are associated with another type of X-ray afterglows. If this is true, from those internal X-ray afterglows, we can further determine the magnetic field strengths and the initial spin periods of the corresponding magnetars. ",
        "introduction": "Gamma-ray bursts (GRBs) are short, intense flashes of soft gamma-rays ($\\sim0.01-1$ MeV), which are always followed by long-lasting low-frequency afterglow emission. Usually, the afterglow emission is attributed to an external forward shock arising from the interaction of the GRB outflow with the circumburst medium, whereas the mechanisms responsible for the bursts are under more debate. Since the launch of the {\\it Swift} spacecraft (Gehrels 2004), many detailed features of the GRB X-ray afterglows have been revealed by the X-Ray Telescope (XRT) aboard. Then Nousek et al. (2006) and Zhang et al. (2006) phenomenologically summarized a ``canonical\" X-ray light curve (LC) with four smooth segments (sometimes superposed by some sharp flares), although the LCs owning all components are actually in the tiny minority. Specifically, the four different emission phases defined by them include: (1) {\\it Initial steep decay phase} that is widely accepted to be the tail of the prompt emission (i.e., curvature effect; Fenimore et al. 1996; Kumar \\& Panaitescu 2000). (2) {\\it Shallow decay (even a plateau) phase} that is usually ascribed to a continuous energy injection into the external shock (Rees \\& {\\mes} 1998; Dai \\& Lu 1998a,b; Zhang \\& {\\mes} 2001). (3) {\\it Normal decay phase} that does not contradict with the standard external shock model. (4) {\\it Steeper-than-normal decay phase} that is often connected to the jet break (Rhoads 1997). To summarize, in the conventional picture as described above, the emission during the shallow decay, normal decay, and steeper-than-normal decay phases (of interest in this paper) are usually considered to be associated with the external shock, while the initial steep decay emission as well as the flares are probably of internal origin.  The X-ray shallow-decay and flare emission strongly suggest that GRB central objects should have long activities after the bursts. Therefore, highly magnetized, rapidly spinning pulsars (i.e., millisecond magnetars) gradually become a popular candidate of the central compact objects (e.g., De Pasquale et al. 2007; Metzger et al. 2007; Troja et al. 2007; Zhang \\& Dai 2008, 2009; Bucciantini et al. 2009; Corsi \\& {\\mes} 2009; Lyons et al. 2009), although the black hole model is still an attractive choice. It is a difficult task to observationally distinguish between the magnetar and black hole models. However, some magnetohydrodynamic simulations showed that the magnetar model could be more mature in the sense that it provides quantitative explanations for the durations, energies, Lorentz factors, and collimation of long GRB outflows (Metzger 2010). Moreover, with a spinning-down magnetar, some {\\it Swift}-XRT features can be explained well (e.g., Dai et al. 2006; Fan \\& Xu 2006; Yu \\& Dai 2007) . Especially, the unusual X-ray afterglow LC of GRB 070110, where a nearly constant X-ray emission is followed by a very steep decline of $\\alpha\\sim 9$ ($\\alpha$ is the decline index of $t^{-\\alpha}$), can be understood by ascribing the plateau emission to magnetar-driven internal emission (Troja et al. 2007).  The internal plateau emission of GRB 070110 also tells us that, in addition to the initial steep decay and flare phases, the X-ray emission during any other afterglow phases could also arise from some internal dissipation mechanisms. Therefore, it is fair to consider that the long active central objects of some other GRBs could also play an essential, relatively more direct role in their afterglow emission, at least in the X-ray band. In some extreme situations, we suspect that such an internal-origin emission component could even dominate the total X-ray afterglow emission of a GRB, whereas the external shock emission is outshined in the X-ray band. In other words, the conventional external shock model is only one choice among various afterglow origin models, as also proposed by some authors before (Ghisellini et al. 2007; Kumar et al. 2008; Cannizzo \\& Gehrels 2009; Lindner et al. 2009; Lyutikov 2009). A wide investigation on the GRB X-ray afterglows given by Willingale el al. (2007) indeed showed that more than one hundred X-ray afterglows can be divided into two emission components, one of which is probably of internal origin.  Following the observational results and the above two theoretical considerations, in this paper, we investigate in more detail the role of millisecond magnetars in the X-ray afterglow emission of some GRBs, based on a careful analysis on the spin evolution of the magnetars. Besides modifying the conventional energy injection model (Dai \\& Lu 1998a, b; Zhang \\& {\\mes} 2001), we pay particular attention to the internal X-ray afterglow emission that is produced by the magnetars, possibly through some internal dissipations in the magnetar winds. The luminosity of this internal emission is assumed to simply track the magnetic dipole luminosity of the magnetars with a certain fraction. To summarize, the observed X-ray afterglows could be emitted from two different regions (i.e., an internally-dissipated magnetar wind and an energized external shock) at very different radii. The competition between these two emission components leads to a diversity of the X-ray afterglow LCs.  In Section 2, we briefly review the spin-down of magnetars against {\\it r}-mode instability. We analyze the temporal behaviors of X-ray afterglows by combining the contributions from a magnetar wind and an external shock in Section 3, where the energy injection from the wind to the shock is also taken into account. In Section 4, some observational samples are selected and fitted in order to confront the model with observations. Meanwhile, some implications to the magnetars from the afterglow data are discussed. Finally, a summary and discussion are given in Section 5. ",
        "conclusions": "Based on two assumptions of that (i) some GRB central objects are millisecond magnetars and (ii) the magnetar winds can continuously produce X-ray emission whose luminosity tracks the magnetic dipole luminosity, we investigate the temporal behaviors of the GRB X-ray afterglows arising from an emitting magnetar wind and an energized external shock together. The competition between the internal- and external-origin emission components determines the diversity of the observed X-ray afterglow LCs. A comparison between the model and observations shows that the model-predicted shock- and wind-dominated emission is qualitatively consistent with the observed one-break and two-break afterglows, respectively.  In the conventional shock model, the second break of the two-break afterglows is always connected to the jet break, which is however seriously challenged by the usually observed chromatic breaks (Liang et al. 2008). In contrast, the chromatic breaks could be acceptable for the internal-origin emission. On one hand, such an argument is supported by the lack of the optical counterparts of X-ray flares, which are of internal origin. On the other hand, in view of the possible small radii where the internal dissipations occur, the low-frequency emission of the wind is quite likely to be suppressed, for example, by some self absorption effects. Therefore, for some GRBs, while their X-ray afterglows are contributed by the magnetar winds, the optical emission could be still dominated by the external shocks. In this case, chromatic breaks would be detected naturally.  In this paper, we mainly concern the ordinary GRB X-ray afterglows that may be associated with a magnetar with a relatively low magnetic field ($B<B_c$). For simplicity we do not pay much attention to some unusual X-ray afterglows such as the afterglows behaving as a plateau followed by a steep decay (e.g., GRBs 060607A and 070110). As discussed in Section 3, such X-ray afterglows could be internal afterglows produced by the magnetars with relatively high magnetic fields ($B>B_c$). In this case, however, if a low value of $B$ is found, the related magnetar could be a candidate of strange quark stars rather than neutron stars, since only strange stars can suppress the {\\it r}-mode instability effectively (Yu et al. 2009b)."
    },
    "1004/1004.2365_arXiv.txt": {
        "abstract": "s{ Directional detection of galactic Dark Matter is a promising  search strategy for discriminating genuine WIMP events from  background ones. However,  to take full advantage of this powerful detection method, one need to be able to extract information from an observed recoil map to identify a WIMP signal. We present a comprehensive formalism, using a map-based likelihood method allowing to recover the main incoming direction of the signal, thus proving its galactic origin, and the corresponding significance. Constraints are then deduced in the ($\\sigma_n, m_\\chi$) plane. } ",
        "introduction": "Taking advantage of the astrophysical framework, directional detection of Dark Matter is an interesting strategy in order to distinguish WIMP events from background ones. Indeed, like most spiral galaxies, the Milky Way is supposed to be immersed in a halo of WIMPs which outweighs the luminous component by at least an order of magnitude. As the Solar System rotates around the galactic center through this Dark Matter halo, WIMPs should mainly come from the direction to which points the Sun velocity vector and which happens to be roughly in the direction of the Cygnus constellation. Then, a directional WIMP flux is expected to enter any terrestrial detectors (see fig.\\ref{fig:DistribRecul} left) infering a directional WIMP induced recoil distribution which should be pointing toward the Cygnus Constellation, {\\it i.e.} in the ($\\ell_\\odot = 90^\\circ,  b_\\odot =  0^\\circ$) direction (see fig.\\ref{fig:DistribRecul} middle). The latter corresponds to the expected WIMP signal probed by directional detectors and as it is shown on the fig.\\ref{fig:DistribRecul} (middle), a strong anisotropy is expected \\cite{spergel} while the background should be isotropic.\\\\  Several project of directional detectors are being developed \\cite{white} and in this paper, we present a map-based likelihood analysis \\cite{billard} in order to extract from an observed recoil map the main incoming direction of the events and its significance. This way, the galactic origin of the signal, thus the identification of a genuine WIMP signal, can be proved by showing its correlation with the direction of the solar motion. This blind analysis is intended to be applied to directional data of any detector and as an example we will apply this method to a realistic simulated data.  \\begin{figure}[t] \\begin{center} \\includegraphics[scale=0.2,angle=90]{fig1a.eps} \\includegraphics[scale=0.2,angle=90]{fig1b.eps} \\includegraphics[scale=0.2,angle=90]{fig1c.eps} \\caption{From left to right : WIMP flux  in the case of an isothermal spherical halo,   WIMP-induced recoil distribution and a typical simulated measurement :  100 WIMP-induced recoils and 100 background events with a low angular resolution. Recoils maps are produced for a Fluorine target, a  100 GeV.c$^{-2}$ WIMP and considering recoil energies in the range  5 keV $\\leq E_R \\leq$ 50 keV. Maps are Mollweide equal area projections.} \\label{fig:DistribRecul} \\end{center} \\end{figure} ",
        "conclusions": "\\label{sec:conclusion} We have presented a statistical analysis tool to extract information from  a  data sample of a directional detector in order to identify a galactic WIMP signal. As a proof of principle, it has been tested within the framework of an isothermal spherical halo model. We have shown the feasibility to extract from an observed map the main incoming direction of the signal and its significance, thus  proving  its galactic origin. Systematical studies have been done\\cite{billard} in order to show that this analysis tool gives satisfactory results on a large range of exposure and background contamination."
    },
    "1004/1004.0156_arXiv.txt": {
        "abstract": " ",
        "introduction": "Astronomers study the physical phenomena outside the Earth's atmosphere by observing cosmic particles and electromagnetic waves impinging on the Earth. Each type of observation provides another perspective on the universe thereby unraveling some mysteries while raising new questions. Over the years, astronomy has become a true multi-wavelength science. A nice demonstration is provided in Fig.\\ \\ref{fig:NGC5055}. In this image, the neutral hydrogen gas observed with the Westerbork Synthesis Radio Telescope (WSRT) exhibits an intricate extended structure that is completely invisible in the optical image from the Sloan digital sky survey \\cite{SDSS}. The radio observations therefore provide a radically different view on the dynamics of this galaxy.  \\begin{figure} \\centering \\includegraphics[width=\\columnwidth]{fig1.eps} \\caption{Image of the spiral galaxy NGC 5055, showing the structure of the neutral hydrogen gas observed with the Westerbork Synthesis Radio Telescope (blue) superimposed on an optical image of the same galaxy from the Sloan digital sky survey (white) \\cite{Battaglia2006-1}. \\label{fig:NGC5055}} \\end{figure}  Images like Fig.\\ \\ref{fig:NGC5055} are only possible if the instruments used to observe different parts of the electromagnetic spectrum provide a similar resolution. This poses quite a challenge since the resolution of any telescope is determined by the ratio of the wavelength and the telescope diameter. Consequently, the aperture of radio telescopes has to be 5 to 6 orders of magnitude larger than that of an optical telescope to provide comparable resolution, i.e.\\ radio telescopes should have an aperture of several hundreds of kilometers. Although it is not feasible to make a dish of this size, it is possible to synthesize an aperture by building an interferometer, i.e., an array.  Radio astronomy started with the discovery by Karl Jansky, at Bell Telephone Laboratories in 1928, that the source of unwanted interference in his short-wave radio transmissions actually came from the Milky Way. For this, he used the large antenna mounted on a turntable shown in Fig.\\ \\ref{fig:instruments}$(a)$. Subsequent single-antenna instruments were based on larger and larger dishes, culminating in the Arecibo telescope (Puerto Rico 1960, 305 m non-steerable dish) and the Effelsberg telescope (Bonn, Germany, 1972, 100 m steerable dish, Fig.\\ \\ref{fig:instruments}$(b)$). Making larger steerable dishes is not practical.  \\begin{figure*} \\parbox[b]{.658\\columnwidth}{ \\includegraphics[width=.658\\columnwidth]{fig2a.eps} \\includegraphics[width=.658\\columnwidth]{fig2b.eps} } \\parbox[b]{.73\\columnwidth}{ \\includegraphics[width=.73\\columnwidth]{fig2c}\\\\[-4.7mm] } \\parbox[b]{.7\\columnwidth}{ \\includegraphics[width=.7\\columnwidth]{fig2d.eps} \\includegraphics[width=.7\\columnwidth]{fig2e.eps}\\\\[-4.3mm] } \\caption{The radio telescopes of $(a)$ Jansky, $(b)$ Effelsberg, $(c)$ WSRT, $(d)$ VLA, $(e)$ concept for ALMA.} \\label{fig:instruments} \\end{figure*}  An interferometer measures the correlation between two antennas spaced at a certain distance. Initially used to study a single source passing over the sky, the principle was used in optical astronomy in the Michelson stellar interferometer (1890, 1920); the first radio observations using two dipoles were done by Ryle and Vonberg in 1946 \\cite{ryle52}.  Examples of subsequent instruments are: the Cambridge One Mile Telescope in Cambridge, UK (1964, 2 fixed and one movable 18.3 m dishes); the 3~km WSRT in Westerbork, The Netherlands (1970, 12 fixed and 2 movable 25 m dishes, Fig.\\ \\ref{fig:instruments}$(c)$); the 36~km Very Large Array (VLA) in Socorro, New Mexico, USA (1980, 27 movable 25 m dishes, Fig.\\ \\ref{fig:instruments}$(d)$); the 25~km Giant Meter-Wave Radio Telescope (GMRT) in Pune, India (1998, 30 dishes with 45 m diameter). These telescopes use the Earth rotation to obtain a sequence of correlations for varying antenna baseline orientations relative to the desired sky image field, resulting in high-resolution images via {\\em synthesis mapping}. Even larger baselines (up to a few thousand km) were obtained by combining these instruments into a single instrument using a technique called VLBI (very long baseline interferometry), where the telescope outputs are time-stamped and post-processed by correlation at a central location. An extensive historical overview is presented in \\cite{Thompson2001-1}. In the near future, astronomers are building even larger arrays, such as the Atacama Large Millimeter Array (ALMA, Chile, 2011, 50 movable 12 m dishes with possible extension to 64 dishes, Fig.\\ \\ref{fig:instruments}$(e)$), the Low Frequency Array (LOFAR, The Netherlands (2009, about 30,000 dipole antennas grouped in 36 stations, Fig.\\ \\ref{fig:hierarchy}), and the Square Kilometer Array (SKA, 2020+, Fig.\\ \\ref{fig:hierarchy}). A recent issue of the Proceedings of the IEEE (Vol.97, No.\\ 8, Aug.\\ 2009) provides overview articles discussing many of the recent and future telescopes.  High-resolution synthesis imaging would not be possible without accurate calibration. Initially, the complex antenna gains and phases are unknown; they have to be estimated. Moreover, propagation through the atmosphere and ionosphere causes additional phase delays that may create severe distortions. Finally, image reconstruction or {\\em map making} is governed by finite sample effects: we can only measure correlations on a small set of baselines. Solving for these three effects is intertwined and creates very interesting signal processing problems.  In this overview paper, we focus on the calibration aspects, whereas imaging is covered in a companion paper \\cite{leshem09spm}. The examples provided in this paper are generally borrowed from low frequency ($<$ 1.5 GHz) instruments, but the framework presented is applicable to high frequency instruments like ALMA as well. ",
        "conclusions": ""
    },
    "1004/1004.2709_arXiv.txt": {
        "abstract": "The integrated luminosity of the TP-AGB phase is a major uncertainty in stellar population synthesis models. We use the white dwarf initial final mass relation and stellar interiors models to demonstrate that a significant fraction of the core mass growth for intermediate ($1.5<\\msol<6$) mass stars takes place during the TP-AGB phase. We find evidence that the peak fractional core mass contribution for TP-AGB stars is $\\sim20\\%$ and occurs for stars between $2\\ \\msol$ and $3.5\\ \\msol$. Using a simple fuel consumption argument we couple this core mass increase to a lower limit on the TP-AGB phase energy output. Roughly half of the energy released in models of TP-AGB stars can be directly accounted for by this core growth; while the remainder is predominantly the stellar yield of He. A robust measurement of the emitted light in this phase will therefore set strong constraints on helium enrichment from TP-AGB stars, and we estimate the yields predicted by current models as a function of initial mass. Implications for stellar population studies and prospects for improvements are discussed. ",
        "introduction": "\\label{sec:intro}  Stars can experience a core-collapse supernova only if they are born with a mass much higher than the Sun, even though a chemically evolved core of order only one solar mass is required to ignite the advanced burning stages.  The culprit is mass loss severe enough to strip off the stellar envelope before a sufficiently massive processed core develops.  Although there is some mass loss even in shell hydrogen burning giants, the vast majority in intermediate mass stars occurs in the presence of thermal pulsations involving interactions between hydrogen and helium burning shells; we refer to this as the thermally pulsing AGB (TP-AGB) phase.  Our understanding of even the basic properties of the TP-AGB phase, such as the lifetime or light emitted, is limited because their severe mass loss is difficult to constrain observationally or predict theoretically.  The uncertainties in TP-AGB evolution have profound consequences for stellar population studies.  Stellar interiors models can accurately predict evolutionary properties up to the onset of the TP-AGB phase, and such models have been extensively used to reconstruct star formation histories in both resolved and unresolved populations.    Models of the TP-AGB phase, however, are extremely dependent on the assumptions related to mass loss and their predictive power is thus limited.  There has been intriguing work suggesting that the fraction of light emitted by TP-AGB stars in LMC star clusters is very high, of order 40 \\% \\citep{Persson83}. Such a substantial flux, largely redistributed to the far-IR by dust, can be especially important for interpreting intermediate redshift galaxy properties \\citep{Conroy09}. There has been a strong emphasis on updating population synthesis to include the TP-AGB phase \\citep[\\eg][]{Bruzual03, Maraston05}. \\citet{Maraston06} demonstrated that modeling of the TP-AGB phase has become a defining characteristic of different stellar population synthesis (SPS) codes while \\citet{Conroy09} identified the substantial uncertainties in TP-AGB properties as a major component of the error budget for galaxy evolution models.  In this paper we employ a fuel consumption argument to set a firm lower bound on the fraction of light emitted during the TP-AGB phase.  We use stellar interiors models to set the core mass at the onset of the TP-AGB phase, and demonstrate that it is surprisingly insensitive to the choice of input physics.  The nuclear processed core then grows until the envelope is expelled, at which point the final white dwarf mass is set.  The white dwarf initial-final mass relationship (IFMR) can in turn be inferred from open clusters.  We use the difference between the final and starting masses in the TP-AGB phase as a bound on the fuel consumed, and thus the emitted light.  Although it is common for investigators to compare their models to initial-final mass relationships, and fuel consumption arguments have been used to check population synthesis models, the quantitative bounds from the IFMR are typically not used in population synthesis calculations.  We argue that the additional empirical information encoded there permits a narrower set of possibilities than considered by \\citet{Conroy09}. Processed fuel (especially helium) can be ejected in winds, so the white dwarf data formally sets only a lower bound on the emitted light.  This raises interesting links between the light emitted in this phase and chemical evolution studies, which we discuss in our conclusions. The plan of our paper is straightforward.  Our sample and methods are discussed in Section 2, our results are presented in Section 3, and we discuss their broader implications in Section 4. ",
        "conclusions": "\\label{sec:disc}  Using the WD IFMR and robust stellar evolutionary theory, we have shown that a significant mass fraction of the final, stellar core is generated during the TP-AGB phase. Through simple fuel consumption arguments, the TP-AGB phase must therefore emit a substantial amount of light- proportional to the aforementioned core growth.  There is measurable core growth during the TP-AGB phase in all nine clusters in our sample. The representative progenitor masses range from $\\sim 1.7 < \\msol\\ < 6$ and they span $\\sim 0.30$ dex in \\feh. The least constrained theoretical assumption in stellar evolutionary models up to the onset of the TP-AGB phase is the depth of convective overshooting and whether or not it takes place. Our main result is valid over a broad swath of theoretical input parameter space, including the uncertain nature of core overshooting. Therefore, we have constructed a strict lower bound on the light output during the TP-AGB phase as a function of initial mass that is predominantly dependent on observational constraints from white dwarfs and their host cluster properties.  The fractional contribution of TP-AGB stars to the total light emitted over their lifetime peaks when $\\mi\\ \\sim 2-3.5 \\msol$. In general, $\\sim20\\%$ of the core is built during the TP-AGB phase in this mass range, double that of lower and higher mass stars. The rise and fall of the TP-AGB's importance with initial mass generally agrees with the relative lifetimes of TP-AGB stars predicted by current models \\citep[\\eg][]{Girardi07, Marigo07, Bertelli08, Bertelli09}. The maximum of the TP-AGB fractional energy output at $\\sim 2-3.5 \\msol$ has strong implications for interpreting the integrated light of elliptical galaxies. The Infrared light from TP-AGB stars will be a one of stronger principle components of the spectral energy distributions of galaxies harboring stellar populations $2$ to $3\\ Gyr$ old.  Constraining the importance of the TP-AGB phase via the IFMR becomes increasingly difficult for progenitor masses above $\\sim5 \\msol$. High mass stars have short lifetimes, implying that $\\tclus\\ - \\tcool$ is a small number. In that case, nominal errors in \\tclus\\ and \\tcool\\ have a fractionally larger impact on the uncertainty in progenitor age at high mass than at low mass. Additionally, different convective overshooting efficiencies produce more discrepant core masses at various stages of evolution for higher initial masses. Still, we find compelling evidence that $\\sim10\\%$ of the final core is built up during the TP-AGB stage at high mass regardless of convective overshoot efficiency. The upper limit in progenitor mass for stars experiencing the TP-AGB phase is still controversial and empirical data is scare \\citep{Williams09b}. While this work suggests $5.5\\ \\msol$ stars experience a TP-AGB phase, the IFMR will provide a strong means to determine the highest initial masses to become TP-AGB stars if progenitor masses in this regime can be better constrained in the future.  Our results have potentially important consequences for chemical evolution models. We determine the total energy output in the TP-AGB phase required by the nuclear reactions that directly lead to core mass growth. Assuming nucleosynthesis is the only source of energy generation in the TP-AGB phase, the only other tracer of energy output would be nucleosynthetic byproducts expelled to the ISM. Any measurement of the total light output of a TP-AGB stellar population, combined with this work, places direct constraints on the yield of that population (principally helium). If the helium yield is significant, it would impact the $\\Delta Y/\\Delta Z$ relationship. Additional helium in the ISM from TP-AGB stars would alter the luminance and lifetimes of subsequent generations of stars. Assuming the models of \\citet{Bertelli08, Bertelli09} correctly predict the total light in the TP-AGB phase, we find that $0.08\\pm0.01\\ \\msol$ of helium are released into the ISM by TP-AGB stars. TP-AGB stars are short lived; as such, observations of TP-AGB populations and limits on their luminosity have proven difficult in galaxies more distant than the Magellanic Clouds. However, the above argument can be reversed: given limiting case assumptions of the helium yield of TP-AGB stars and IFMR data, one can establish both upper and lower bounds on the integrated light from TP-AGB stars.  Early evidence suggests that modern models of TP-AGB stars make realistic predictions for their total light output. Mass addition to the core accounts for $\\sim 50\\%$ of the integrated luminosity from the TP-AGB phase predicted by \\citet{Bertelli08, Bertelli09}. If \\fecore\\ had exceeded $100\\%$, we would conclude that the chosen models underestimated the amount of light emitted by TP-AGB stars. Alternatively, if \\fecore was small ($<10\\%$), we would calculate a relatively large stellar yield. In the future, empirical evidence describing the behavior of the $\\Delta Y/ \\Delta Z$ relationship or integrated light measurements of TP-AGB populations will constrain TP-AGB stellar yields and impose tighter restrictions on the upper bound of total energy release. Currently, the \\citet{Bertelli08, Bertelli09} models predict an amount of luminous energy release by TP-AGB stars that is in agreement with observations. \\citet{Bertelli08, Bertelli09} predict that, at its peak, the TP-AGB phase is responsible for $\\sim33\\%$ of the total light output of stars with $2.2\\ \\msol\\leq\\ \\mi\\ \\leq 2.5 \\msol$; this fractional contribution falls to $20\\%$ when \\mi$=3.5\\ \\msol$ and is $\\sim 10\\%$ at \\mi$=4.0\\ \\msol$. Though our current results and errors cannot rule it out, it is unlikely that TP-AGB stars are responsible for as much as $40\\%$ of cluster light unless the stellar yields of TP-AGB stars are significantly higher than those predicted in Section~\\ref{sec:Hefuel}.  This work suggests additional tests of population synthesis models. The predicted light emission from stellar populations in SPS models provides a upper bound on the expected remnant mass as a function of initial mass. In the future, population synthesis codes should quantify their agreement with the IFMR. Various theoretical question marks in SPS modeling can be answered, or at least constrained, by the IFMR. The degree of convective overshooting in stellar evolutionary models has an appreciable impact on implied progenitor masses and the \\mtp$-$\\mi\\ relationship; hence, for a predicted light output in the TP-AGB phase, there will be a statistically significant difference between the remnant mass functions predicted with different overshooting parameters. If we demand that \\mtp\\ is such that the core grows significantly during the TP-AGB phase, we can place instructive bounds on convective overshoot prescriptions for some clusters. Additionally, the luminosity of TP-AGB populations represent a major uncertainty in SPS models \\citep{Conroy09}. Our results showing core mass growth in the TP-AGB phase as a function of initial mass represent a floor to TP-AGB star luminosity and can already exclude some limiting case models. Obviously, models assuming that TP-AGB stars do not contribute to a population's integrated light can no longer be physically motivated. As mentioned above, this work and assumptions as to the chemical yields of TP-AGB stars is sufficient to constrain TP-AGB light output. These new constraints are valuable to future SPS models and will greatly reduce the uncertainties associated with the one of the most theoretically dubious stellar evolutionary phases..  The most prominent sources of error in this analysis are the WD and cluster observations. Measurements of WD surface gravities and temperatures provide constraints on remnant masses and cooling times while cluster observations yield the cluster age and composition. Uncertainties in theory linking these observations and the parameters of interest are typically small compared to the observational errors (S09). By assuming either perfect WD or cluster measurements and repeating our analysis, we find that uncertainties in these two types of observables contribute similarly to the resulting error in core mass growth during the TP-AGB phase. More precise WD or cluster measurements will provide stricter bounds on the energy output of TP-AGB stars. Another potential method to improve the precision of the IFMR (and our results) would be a survey specifically designed to find young white dwarfs in open clusters. In any given cluster, the youngest white dwarfs would correspond to the oldest progenitor lifetimes, reducing the impact of uncertainties in the WD and cluster observations on the implied initial mass and \\mtp. More precise measurements of the IFMR would directly lead to more precise calculations in our procedure.  Determining the metallicity dependence of TP-AGB star energy output is an intriguing next step. However, there are not enough open clusters at extreme metallicities to empirically calibrate such a dependence. We would need to increase the number of clusters in our sample, even if the overall spread in metallicity does not change. With many more clusters, one could reproduce our results in bins of metallicity. After empirically characterizing how the energy output of the TP-AGB phase scales with composition (albeit over a limited range), theory may be able to provide a physical model for TP-AGB star luminosity as a function of initial mass and stellar density. If said model could be extrapolated to low metallicity, it would provide a theoretical constraint on the contribution of TP-AGB stars to the spectral energy distributions of high redshift galaxies. At high metallicity, the lowest mass stars may go directly to the white dwarf stage, missing the core He burning stages and beyond \\citep{Kilic07}.  This effect is not currently included in population synthesis models and might be important for the low redshift properties of giant elliptical galaxies. Physically motivated priors on TP-AGB star light output would be a substantial advancement of semi-analytic galaxy evolution models as SPS codes still have widely varied characterizations of these stars.  We have shown that the WD IFMR places powerful constraints on the energy released during the TP-AGB stage of stellar evolution. Future observations of TP-AGB populations will constrain the helium yields of these stars- an important factor for chemical evolution models to consider. Alternatively, nominal assumptions regarding the yield of these stars, in conjunction with this work, result in a narrower range of possible TP-AGB population luminosity than previously considered."
    },
    "1004/1004.5403_arXiv.txt": {
        "abstract": "% {% The profiles of emission lines formed in the corona contain information on the dynamics and the heating of the hot plasma. Only recently has data with sufficiently high spectral resolution become available for investigating the details of the profiles of emission lines formed well above $10^{6}$K. These show enhanced emission in the line wings, which has not been understood yet. } {% We study the underlying processes leading to asymmetric line profiles, in particular the responsible plasma flows and line broadening mechanisms in a highly filamentary and dynamic atmosphere. } {% Line profiles of \\ion{Fe}{15} formed at 2.5\\,MK acquired by the Extreme ultraviolet Imaging \\NNN{Spectrometer} (EIS) onboard the Hinode solar space observatory are studied using multi Gaussian fits, with emphasis on the resulting line widths and Doppler shifts.} {% In the major part of the active region, the spectra are best fit by a narrow line core and a broad minor component.  The latter contributes some 10\\% to 20\\% to the total emission, is about a factor of 2 broader than the core, and shows strong blueshifts of up to 50\\,km/s, especially in the footpoint regions of the loops. On average, the line width increases from the footpoints to the loop top for both components. A component with high upflow speeds can be found also in small restricted areas. } {% The coronal structures consist of at least two classes that are not resolved spatially but only spectroscopically and that are associated with the line core and the minor component. Because of their huge line width and strong upflows, it is proposed that the major part of the heating and the mass supply to the corona is actually located in source regions of the minor component. It might be that these are identical to type II spicules. The siphon flows and draining loops seen in the line core component are consistent with structures found in a three-dimensional magneto-hydrodynamic (3D MHD) coronal model. Despite the quite different appearance of the large active region corona and small network elements seen in transition region lines, both show similar line profile characteristics. This indicates that the same processes govern the heating and dynamics of the transition region and the corona. } ",
        "introduction": "Understanding the dynamics of the plasma in the corona of the Sun is pivotal for unveiling the processes that govern the heating of the outer atmosphere of the Sun to temperatures of millions of degrees. Imaging instruments can give vital information on the evolving structures. However, they only provide information on the apparent motions, which do not necessarily have to be real plasma motions, and they are broad band in wavelength, thereby mixing several spectral lines formed at different temperatures. Spectroscopy provides details on the emission line profiles and thus supplies crucial information for investigating the thermal and dynamic structure of the coronal plasma.   Various scenarios and models make detailed predictions on the spectral profiles to be observed. Only selected aspects should be mentioned here. The nanoflare-heating models of \\cite{Patsourakos+Klimchuk:2006} based on 1D hydrodynamic loop models predict high-temperature upflows of hot plasma following transient heating events (introduced adhoc in the model). These would show up as a weak additional component in the wing of emission lines formed well above 10$^6$K with a blueshift of some 100\\,km/s (see their Fig.\\,3). Based on their observations, \\cite{DePontieu+al:2009:roots.of.heating} propose a scenario for the mass cycle of the active region corona where hot plasma related to type II spicules is propelled upwards to supply mass to the corona (which subsequently rains down after cooling). This scenario would also imply a strongly blueshifted component in the coronal emission lines. Similarly, \\cite{Tu+al:2005} or, more recently, \\cite{Tian+al:2010} have found evidence of the plasma in open field regions near the limb above temperatures of several 100.000\\,K being accelerated to form the open wind, with blueshifts of some 30 km/s seen in lines formed at 2\\,MK. If (locally) open regions on the disk, e.g. adjacent to an active region, were to produce a similar outflow, this could show up in the spectra as excess emission in the blue wing of the emission line, too.    Line profiles originating in the transition region from the chromosphere to the corona have been intensively investigated using vacuum ultraviolet spectrometers such as HRTS \\citep[High Resolution Telescope and Spectrograph;][]{Brueckner+Bartoe:1983} or SUMER \\citep[Solar Ultraviolet Measurements of Emitted Radiation;][]{Wilhelm+al:1995}. These provided a high spectral resolution of about 20\\,000 or more allowing them not only to derive the shifts of the line centroid and width, but also to detect peculiarities in the line profiles. The most noticeable features in the transition region line profiles are found during explosive events \\citep{Dere+al:1989:expl.events}: two distinct satellites to the blue and red indicating a bi-directional flow, most likely induced by a small reconnection event. The spatial and temporal evolution of these were first described by \\cite{Innes+al:1997}, confirming the nature of a bi-directional flow. \\citet{Kjeldseth-Moe+Nicolas:1977} were the first to show that transition region spectra show enhanced emission in the line wings. \\cite{Peter:2000:sec:err} presented evidence that these enhanced wings are predominantly found above the chromospheric network and are best fitted by a double Gaussian with a narrow line core and a broad second (or wing) component. A double Gaussian fit with one broad component is even more significant than a triple Gaussian fit with two narrow components accounting for the wing emission \\citep{Peter+Brkovic:2003}. This questions the previous interpretation of the wing excess as caused by small counterparts of explosive events \\citep{Dere+Mason:1993}. The width of the broad wing component seems to increase monotonically with line formation temperature \\citep{Peter:2001:sec}, which is consistent with line broadening due to magneto-acoustic waves.   Our knowledge of the detailed spectral profiles of lines formed in the hotter corona is more limited by instrumentation in the era of SoHO \\citep[Solar and Heliospheric Observatory;][]{Domingo+al:1995} and before. The Coronal Diagnostic Spectrometer  \\citep[CDS;][]{Harrison+al:1995} did not provide sufficient spectral resolution to reveal details in the line profiles (except for extreme cases, e.g.\\ flares), and the SUMER spectra of lines formed at hot temperatures are problematic in terms of signal-to-noise ratio. This situation changed with the the Extreme ultraviolet Imaging Spectrometer \\citep[EIS;][]{Korendyke+al:2006,Culhane+al:2007} onboard the Hinode solar space observatory \\citep{Kosugi+al:2007}. While the spectral resolution of about 4000 is worse than SUMER or HRTS, the good signal-to noise ratio provides the possibility of studying the details of the line profile.  \\begin{figure*} \\includegraphics[bb=56 283 566 722]{fig02} \\caption{% Spatial maps of line shift and width for \\ion{Fe}{15} (284\\,{\\AA}). These are shown for single Gaussian fits (left column), as well as for the core component and the minor (wing) component for a free double Gaussian fit. The panels are co-spatial and co-temporal with respect to \\fig{F:context}. The respective color scales for the single Gaussian and the core of the double Gaussian are the same, while the scales for the minor component of the double Gaussians span far wider ranges. The Doppler shifts have been corrected for the tilt of the spectrograph slit and the orbital motion of the spacecraft, while the line widths (exponential widths) have not been corrected for instrumental broadening. The crosses, rectangles, and contours have the same meaning as in \\fig{F:context}. See \\sects{S:single} and \\ref{S:free.double}. \\label{F:shift.width.maps} } \\end{figure*}   The first study of the asymmetries of coronal lines in an active region was conducted by \\citet{Hara+al:2008} using EIS data. As a first approach they performed single Gaussian fits to two emission lines and found deviations from a single Gaussian, showing excess emission in the blue wing near loop foot points. Consequently they interpret their findings by favoring the nanoflare heating model by \\cite{Patsourakos+Klimchuk:2006}.  However, as demonstrated in this paper, a closer inspection of the data by a more advanced data processing shows a more detailed picture. Indeed, evidence is found for \\NNN{upflows near the loop footpoints, but these are probably not consistent with the nanoflare scenario}.   The paper is organized as follows. First the data and their reduction are discussed in \\sect{S:obs} before the multi Gaussian line profile fits are analyzed in \\sect{S:multi.Gauss}. Subsequently, the results are discussed with respect to the high-velocity outflows in the active region periphery (\\sect{S:high.velo.outflow}), to the heating and mass supply in the active region (\\sect{S:broad.component}), and to the self-similarity of processes in the active region corona and the quiet Sun network transition region (\\sect{S:tr}). Finally, the results are put into the context of recent three-dimensional magneto-hydrodynamic (3D MHD) coronal models in \\sect{S:loop.model} and the conclusions summarized in \\sect{S:conclusions}. ",
        "conclusions": "\\label{S:conclusions}   This study emphasizes the importance of information hidden in details of the spectral line profiles, which show severe deviations from a single Gaussian in coronal lines formed in active regions. In this manuscript the interpretation is adopted that (two) different spatial structures contribute to the line profile and, through this, lead to a double Gaussian profile. The line fits show that the line profile is in general composed of a narrow Gaussian accounting for the line core emission and a minor broad Gaussian component for the excess emission in the line wings. The results for the line shift for the line core can differ significantly from a single Gaussian fit. In some areas even opposite results can be found when comparing the Doppler shifts for the line core and the single Gaussian (as in the region with fast upflow components, cf.\\ \\fig{F:shift.width.maps} and \\sect{S:single}).   In two small areas at the periphery of the sunspot some evidence of a high upflow velocity component in the line profiles is found. This is located in regions of weak coronal emission, supposedly along low-density structures, which might be interpreted as the footpoint regions of the (slow) solar wind (\\sect{S:wind}). Evidence is found that these structures are highly filamentary with small fingers of high-velocity upflows feeding the wind sticking up into the background funnel (cf.\\ \\fig{F:cartoon}).   In the active region, basically all spectral profiles show a minor broad component that is about two times broader than the line core component and shows systematic blueshifts up to 50\\,km/s, while the line core remains more or less unshifted (cf.\\ \\fig{F:shift.width.histo}). The data indicate that the line core \\emph{and} the minor component are narrower near the loop footpoints than at the apex, indicating some asymmetry in the broadening mechanism. This could be caused by (Alfv\\'en) waves or a heating mechanism that leads to preferential perpendicular heating of the ions (\\sect{S:wave.asymm}).   Comparing the present observations with previous work leads to a picture where individual strands of a loop are feeding the loop with heated material, leaving the signature of a broad blueshifted component in the line profile. After a while, these strands will no longer be heated, will slowly cool down and become part of the background loop that produces the narrow more or less unshifted emission of the line core. Now other strands are heated and fill the loop, and so forth. These heated strands might be identified with the type II spicules proposed by \\cite{DePontieu+al:2009:roots.of.heating} and \\cite{McIntosh+DePontieu:2009:upflows}. This is depicted in \\fig{F:cartoon}. In contrast to previous studies, this analysis also provides information on the heating and dynamics of the EUV counterparts of the type II spicules, which might be the source region of the minor components, because in the present study information on line widths and shifts of the minor component is extracted (\\sect{S:mass supply}).  Depending on the magnetic structure and the distribution of the heat input between the footpoints, the background loops might host siphon flows or they might be draining in response to a lower heating rate (cf.\\ \\fig{F:cartoon}). Both cases can be found in the present data set and have also been found in 3D MHD simulations (\\sect{S:loop.model}).  No clear evidence could be found to support the nano-flare model of \\cite{Patsourakos+Klimchuk:2006}. This model predicts that there is a narrow high-velocity upflow component in the blue wing (at some 100\\,km/s, which is not present in the footpoint regions of the active region loops. However, further (more realistic) modeling efforts would be needed for any conclusive statement (\\sect{S:nanoflares}).   The relation of line width to line shift for the line core, as well as for the minor component, is very similar in the coronal line investigated here and the transition region emission lines originating in bright network patches (\\sects{S:stat}  and \\ref{S:tr}). This shows that the basic processes driving the heating and dynamics on smaller scales in the network and on larger scales in the active region are similar in nature. (At least they have to produce the same spectral signatures.)  This is supported by the match of the minor component line width of the coronal line from the active region to the extrapolation from the transition region lines from the network patch.   In summary, these observations reveal the complex flow systems in the active region corona leading to a mass cycle with small structures (probably type II spicules) feeding mass into the corona, which could host different types of flows as depicted in \\fig{F:cartoon}. It has to be noted, however, that this is only one possible interpretation. Other appealing ideas, such as an asymmetric ion velocity distribution function  causing asymmetric line profiles (\\sect{S:velo.distribution}), will have to be followed and checked more closely against the new observations of coronal EUV emission line profiles.  New studies that test the ideas proposed here further will have to include the analysis of both emission lines with different mass-to-charge ratios, as well as of emission lines covering a range of line formation temperatures from the upper transition region to the hot corona.    {"
    },
    "1004/1004.1400_arXiv.txt": {
        "abstract": " ",
        "introduction": "The wealth of new data provided by the Cassini mission (and in particular, the collection of independent radial wave optical depth profiles of Saturn's rings) opens the possibility to use wave dynamics to establish physical diagnostics of ring physics to a level of accuracy substantially higher than what was possible with Voyager. However, such a research program requires an accurate determination of the wave kinematic parameters throughout the wave propagation region. This has prompted us to reinvestigate this last problem.  Many density waves in Saturn's rings are strong and nonlinear -- that is they cannot be modeled by the linear theory of density waves fully described e.g. by \\cite{S84} -- but this linear theory is nevertheless used due to the absence of a well-established nonlinear inversion procedure ({\\sl e.g.}, \\citealt{NCP90, Rosen91a, Rosen91b, S04, Tiscareno06, Tiscareno07}). This paper presents a new approach to analyze linear and nonlinear density waves. The procedure is established through the use of a number of simulated optical depth profiles, as only simulated data allow us to compare the reconstructed kinematic parameters with the original ones. It is applied next to the Mimas 5:3 density wave as observed by the Cassini Radio Science occultation experiment, as an example. This procedure will serve as the basis of several subsequent studies of real data. The approach adopted here considerably sharpens the method described by \\cite{LB86}, which extracted the maximum information from one Voyager photopolarimeter stellar occultation profile. However, one profile is insufficient to accurately constrain all the parameters of the model to the level of accuracy required to produce new and detailed diagnostics of the ring physics.  There are several motivations for analyzing density waves, and especially nonlinear density waves. The first objective of determining the kinematic behavior of ring density waves is to test the background kinematic model described in section \\ref{theory}, as possible systematic deviations from this model might provide information on physics that are yet unknown or not well constrained. Cassini provides many different radio occultation profiles of the rings, which can be used to verify that a single set of wave parameters can be determined for various observations of a wave at various longitudes with respect to that of the associated satellite (at least when modulations due to satellite orbital variations or other physical effects are weak enough).  For each wave, the set of kinematic parameters determined by the method described in this paper, coupled to the analysis of the evanescent part of the wave (which may require extraneous dynamical constraints to be accomplished), form the basis of the next objective, which is to measure as precisely as possible the torques exerted by the satellites on the rings. Such a direct measurement was previously attempted in \\cite{LB86}, but with limited success due to insufficient constraints in the first wavelength of the studied wave, and to failure of the WKBJ approximation in the vicinity of the resonance. This type of torque measurement based on a determination of the wave kinematics can only be performed in Saturn's rings, and therefore constitutes a unique way to verify the dynamical theoretical predictions for such torques. The measurement and verification of the torque are directly relevant to the dynamics of disk-satellite interactions in general. The extent to which dynamical constraints are required on top of a purely kinematic description of the wave remains to be seen; this bears on the model dependence of such a torque measurement.  Another important study that will follow from this paper is the detailed analysis of the ring stress tensor, with the hope that this may in turn provide useful constraints on the ring particles collisional properties. We wish to identify statistically the various stress behaviors that might be found, depending, for instance on the ring region or ring background optical depth. Several models of damping exist. For instance, \\cite{BGT83} and \\cite{SDLYC85} predict a bimodal behavior as a function of the ring mean optical depth in dilute rings. \\cite{BGT85} generated models for dense rings. This research program requires combining the kinematic reconstruction method developed here with the nonlinear (pressure-corrected) dispersion relation and at least one generic dynamical equation describing wave propagation and damping (such as the ones derived in \\citealt{SDLYC85} or {\\citealt{BGT86}).  In the bulk of this paper, we simulate a density wave that has some similarities with the Mimas 5:3 density wave, as observed in eight diffraction-corrected Cassini radio occultation profiles. We did not try to simulate this wave precisely, because the purpose of the simulation was to establish the procedure. A preliminary and simplified analysis of real data pertaining to this wave is presented in section \\ref{mimas}. The reader might wonder why we started with simulated data. The reason is that such an analysis is in itself already very complex and it is important to separate this complexity from that inherent to the analysis of real data. There are several reasons for which the analysis of simulated data is difficult.  The first reason is the nonlinearity in itself. The equation describing the shape of the wave, with its large troughs and narrow peaks, does not provide a model that one can readily fit to the data. The fit turns out to be extremely sensitive to the dependence of the wave phase on semi-major axis, and so the procedure must allow for a very precise determination of this quantity. By using simulated data, we can compare how the reconstructed wave kinematic parameters such as the wave phase compare to the ones used in constructing the data; we can therefore quantify the error in the reconstruction method developed here.  The next difficulty is associated with the fact that the absolute radial scale of Saturn's rings is presently known to only $\\pm 2$ km (\\citealt{NCP90, French93, Jacobson06}). Long period variations in the satellite's orbit may produce similar radial shifts in the resonance location, and are expected to produce variations in the wave propagation which are not yet modeled. The uncertainty in the radial scale, at least, will be reduced by future kinematic models to the level of about 100 m.  Another major problem comes from the fact that one must distinguish between radius $r$ and semi-major axis $a$; the relation $r=a\\left[1-e\\cos(m\\phi+m\\Delta)\\right]$ (see section \\ref{theory} for definitions of the variables and discussion of this relation) is simple only superficially.  Other effects add their own complexity to the data. Interestingly, the noise present in real data (which we took into account in our simulations) is not a real problem, although the need for several profiles to constrain a unique self-consistent solution results mostly from the presence of noise. However, and as indicated by the work of \\cite{LS00, Lewis05}, gravitational clustering can cause significant effects on the structure of local density maxima. Also, the gravitational wakes observed in the rings (see for instance \\citealt{Hedman07, Colwell06, Colwell07a}) make it necessary to normalize the background optical depth among the various profiles, and may affect the optical depth profile of the wave in an as yet uncontrolled way. Focusing on simulated data allows us to ignore this added complexity in the first step .  Before describing our procedure in section \\ref{procedure}, we begin in section \\ref{theory} by gathering the basic equations of the theoretical streamline kinematic model. Section \\ref{simulation} describes and illustrates the method to obtain the simulated data. Section \\ref{mimas} applies the procedure to the Mimas 5:3 density wave. Section \\ref{conclusion} summarizes and concludes the paper with a discussion of future work. ",
        "conclusions": "We have established a robust procedure to analyze nonlinear density waves and we have applied it to the Mimas 5:3 density wave. In principle, the number of independent profiles provided by the Cassini mission would seem to make it possible to devise an inversion procedure relying only on the kinematic characteristics of the wave, but we found that in practice, and relying only on the radio data, this fails for three major reasons: the differences in the absolute radial scales of different profiles, the complex way in which the $a(r)$ relation comes into play, and the sensitivity to errors in the dependence of the wave phase on semi-major axis.  Instead, we have extended and considerably improved the method introduced in \\cite{LB86}, which relies not only on the wave kinematics, but on the WKBJ ordering as well. The intrinsic nonlinearity of the problem makes the procedure complex. Nevertheless, it is possible to reconstruct the wave parameters as well as the data, in the presence of uncertainties in the absolute scale of about $\\pm 2$ km and of noise in the data; the data resolution adopted was 250 m but the results should not be sensitive to this parameter.  As such, this inversion procedure performs very well in the far wave region. This opens new possibilities of physical diagnostics of wave regions. In particular, with the help of the nonlinear dispersion relation and the wave damping equation, one may possibly constrain the surface density and ring's stress tensor in a joint way, which would provide us with an indirect window on the ring' collisional physics. We plan to make use of this inversion procedure to undertake such systematic studies in the future.  The major inaccuracy of this reconstruction procedure is its limited efficacy within the first wavelength of propagation of the wave, where the WKBJ ordering is expected to fail on theoretical grounds, an expectation confirmed by the behavior observed in actual wave profiles. This limitation is most drastic in the reconstruction of the eccentricity $e(a)$ and lag angle $\\Delta(a)$. This is unfortunate for the purpose of the determination of the exchange of angular momentum with the satellite, as the generic expression of the torque results from a radial integral involving these two quantities [see, {\\sl e.g.}, Eq.~(35) of \\citealt{LB86}]. The results of the present inversion method (which yields a fairly educated guess of the form of the kinematic parameters even inside the first wavelength, and has allowed us to remove the problem of inconsistencies of the absolute radial scale between the various profiles), may ensure the convergence of a refined, direct inversion method of the type we initially tried, provided enough high quality profiles are available. This approach may require to a different method of inversion of Eqs.~(\\ref{qcosgamma}) and (\\ref{qsingamma}), if needed; a possible way to do this is to re-express these equations as a set of nonlinear algebraic equations, and solve them by an algebraic iteration method. Additionally, further dynamical constraints may be needed to ensure that such a direct inversion method is well-defined. Three levels of such constraints are available in the literature. In their most useful form, they are all related to the dynamical equation governing the behavior of $Z\\equiv e\\exp(i m\\Delta)$ (see \\citealt{SYL85}): they are the dynamical equation itself, its solution in the linear limit, and the asymptotic behavior of this solution in the evanescent region. Defining the most appropriate strategy on this issue requires substantial additional work, which will be reported elsewhere.  The application to the actual data of the Mimas 5:3 density wave gives good results, although the determination of $\\tau_0$ and $q$ is impaired by the noise in the radio data; better results are expected from the use of, e.g., the UVIS data.  \\appendix"
    },
    "1004/1004.5156.txt": {
        "abstract": "% context heading (optional), leave it empty if necessary CONTEXT {} % aims heading (mandatory) AIMS {The knowledge of accurate stellar parameters is a keystone in several fields of stellar astrophysics, such as asteroseismology and stellar evolution. Although the fundamental parameters can be derived both from spectroscopy and multicolour photometry, the results obtained are sometimes affected by systematic uncertainties. In this paper, we present a self-consistent spectral analysis of the pulsating star RR~Lyr, which is the primary target for our study of the Blazhko effect.} % methods heading (mandatory) METHODS {We used high-resolution and high signal-to-noise ratio spectra to carry out a consistent parameter determination and abundance analysis for RR~Lyr. The \\llm\\ code was employed for model atmosphere calculations, while \\synth\\ and \\width\\ codes were used for line profile calculation and LTE abundance analysis.  We provide a detailed description of the methodology adopted to derive the fundamental parameters and the abundances. Stellar pulsation reaches high amplitudes in RR~Lyrae stars, and as a consequence the stellar parameters vary significantly over the pulsation cycle.  The abundances of the star, however, are not expected to change. From a set of available high-resolution spectra of RR~Lyr we selected the phase of maximum radius, at which the spectra are least disturbed by the pulsation. Using the abundances determined at this phase as a starting point, we expect to obtain a higher accuracy in the fundamental parameters determined at other phases. } % results heading (mandatory) RESULTS {The set of fundamental parameters obtained in this work fits the observed spectrum accurately. %In particular, the surface gravity was adjusted to fit pressure-sensitive spectral features. Through the abundance analysis, we find clear indications for a depth-dependent microturbulent velocity, that we quantified. } % conclusions heading (optional), leave it empty if necessary {We confirm the importance of a consistent analysis of relevant spectroscopic features, application of advanced model atmospheres, and the use of up-to-date atomic line data for the determination of stellar parameters. These results are crucial for further studies, e.g., detailed theoretical modelling of the observed pulsations. } ",
        "introduction": "% The modelling of pulsational signals requires the knowledge of stellar parameters and primarily accurate values of the effective temperature (\\Teff) and metallicity ($Z$). The determination of fundamental parameters can be performed by different methods (some examples for stars from B- to G- type are Fuhrmann et al. 1997, Przybilla et al. 2006, and Fossati et al. 2009) that do not always lead to consistent results. Thus, it is important to choose a methodology that allows us to constrain the parameters of the star from the available observables (usually photometry and spectroscopy) in the most robust and reliable way.  RR~Lyr is the prototype and eponym of its class of pulsating stars. RR Lyrae stars play a crucial role as distance indicators.  Their evolutionary stage (He burning in core, H burning in shell) makes them useful tracers of galactic history.  These classical pulsators display radial oscillations (the simplest type of pulsation) with large amplitudes which makes them useful touchstones for theoretical modelling. RR~Lyr is one of the best studied stars of its class.  Almost a century ago Shapley (1916) discovered that it shows a strong Blazhko effect, i.e. a (quasi-)periodic modulation of its light curve shape in amplitude and phase. The Blazhko effect in RR~Lyr has been closely followed over the past century, and changes have been reported both in the strength and the duration of its Blazhko cycle (Szeidl 1988, Kolenberg et al. 2006). Some well-studied stars even show multiple (variable) modulation periods (see, e.g., LaCluyz\\'e et al. 2004). Despite numerous attempts to model the phenomenon, the Blazhko effect has eluded a satisfactory explanation so far. Recently obtained high-precision photometry from ground-based or space-borne precise instruments also indicate that Blazhko modulation may be a much more common phenomenon than initially thought: as many as half of galactic RRab stars may be modulated  (Jurcsik et al. 2009;  Szabo et al. 2009; Kolenberg et al. 2010).  In order to constrain the viable models for the Blazhko effect, it is vital to obtain accurate values for the fundamental parameters (and their variations) for modulated and non-modulated RR Lyrae stars.  This has been the main motivation for the study presented in this article.  RR~Lyr is the only star of its class to have a directly determined parallax, recently measured with the HST/FGS, by Benedict et al. (2002), to be $\\pi$(FGS) = 3.82 $\\pm$ 0.2 mas (d = 262 $\\pm$ 13 pc). When a small ISM correction of $A_v$ = 0.07 is applied, this new distance results in an $M_{v}^{\\rm RR} \\simeq +0.61_{-0.11}^{+0.10}$ mag which corresponds to $\\simeq 49 \\pm 5 L_{\\odot}$.  Fundamental parameters of RR Lyr have been obtained by several authors with a variety of methods (e.g., Lambert et al. 1996; Manduca et al. 1981; Siegel 1982; for a summary see Kolenberg 2002). The published fundamental parameters of RR~Lyr display a considerable range both in \\Teff\\ and \\logg\\  due to the large pulsation amplitudes. According to these analyses, RR~Lyr's  \\Teff\\ varies over its 13h36min pulsation cycle between approximately 6250 and 8000 K and its \\logg\\ between 2.5 and 3.8 (extreme values). Superposed on the large variation, the Blazhko cycle leads to an additional variation of the fundamental  parameters. Jurcsik et al. (2008) recently showed that also the {\\it mean} properties of  modulated RR~Lyrae stars change over the Blazhko cycle. Element abundances of RR~Lyr were obtained previously by, e.g., Clementini et al. (1996), Lambert et al. (1996), and Takeda et al. (2006).  The main goal of the present work is to perform a self-consistent atmospheric and abundance analysis of RR~Lyr that reproduces all of its photometric and spectroscopic data. Furthermore, we want to investigate the degree to which the derived fundamental parameters depend on the applied methods. Considering the structure of the available models, especially the position of the convective zones and the zones of nuclear fusion, the measured abundances of the star are not expected to change over the pulsation (and the Blazhko) cycle.  Hence, if the abundances are accurately determined at one phase in the pulsation cycle, they can be of help to determine (or at least constrain) the fundamental parameters at other phases.   In this paper we also selected the optimal phase for determining the abundances of the star. This is the first of a series of planned papers devoted to a detailed spectroscopic study of RR~Lyr. In forthcoming papers we will discuss the spectral variations over the pulsation and Blazhko cycle of the star.  % ",
        "conclusions": ""
    },
    "1004/1004.3319_arXiv.txt": {
        "abstract": "With the Blue Channel Spectrograph (BCS) on the MMT telescope, we have obtained spectra to the atmospheric cutoff of quasars previously known to show at least one absorption system at $z>1.6$ with very strong metal lines.   We refer to these absorbers as candidate metal-strong damped \\lya\\ systems (cMSDLAs), the majority of which were culled from the Sloan Digital Sky Survey.  The BCS/MMT spectra yield precise estimates of the \\ion{H}{1} column densities (\\nhi) of the systems through Voigt profile analysis of their \\lya\\ transitions.  Nearly all of the cMSDLAs (41/43) satisfy the \\nhi\\ criterion of DLAs, $10^{20.3}$ \\atomspercm. As a population, these systems have systematically higher \\nhi\\ values than DLAs chosen randomly from quasar sightlines.  Combining our \\nhi\\ measurements with previously measured metal column densities, we estimate metallicities for the MSDLAs.  These systems have significantly higher values than randomly selected DLAs; at $z \\approx 2$, the MSDLAs show a median metallicity [M/H]~$\\approx -0.67$ that is $0.6$\\,dex higher than a corresponding control sample. This establishes MSDLAs as having amongst the most metal-rich gas in the high $z$ universe. Our measurements extend the observed correlation between \\ion{Si}{2}~1526 equivalent width and the gas metallicity to higher values.  If interpreted as a mass-metallicity relation, this implies the MSDLAs are the high mass subset of the DLA population.  We demonstrate that dust in the MSDLAs reddens their background quasars, with a median shift in the spectral slope of $\\delta\\alpha = 0.29$. Assuming an SMC extinction law, this implies a median reddening $E_{B-V} \\approx 0.025$\\,mag and visual extinction $A_V \\approx 0.076$\\,mag.  The latter quantity yields a dust-to-gas ratio of $log(A_V / N_{\\rm HI}) \\approx -22.0$, very similar to estimation for the SMC. Future studies of MSDLAs offer the opportunity to study the extinction, nucleosynthesis, and kinematics of the most chemically evolved, gas-rich galaxies at high $z$. ",
        "introduction": "A Damped \\lya\\ Absorption System (DLA) is defined as a neutral hydrogen absorber along a Quasi-Stellar Object (QSO) sight line with an \\ion{H}{1} column density of $\\mnhi \\ge 10^{20.3}$ \\atomspercm\\ \\citep[see][for a review]{wgp05}. DLAs provide an important way to study the gas and dust content of the early universe \\citep[e.g.][]{pshk94,lu96,pw01}. Their heavy elements were synthesized either in nuclear reactions occurring in stellar cores or when massive stars go supernova.  Studies of DLAs, therefore, probe the processes of stellar evolution in young galaxies. The large surface density of neutral hydrogen that defines DLAs means these systems contain most of the star forming gas in the early universe \\citep{phw05,opb+07}. Since the presence of heavy metals in DLAs implies the existence of stars, DLAs likely trace high redshift galaxies along the sight line between the QSO and earth.  These high redshift galaxies are thought to be precursors of galaxies like the Milky Way \\citep[e.g.][]{nsh04,pgp+08}.  A subset of DLAs exhibit especially strong metal absorption lines. The metal-strong damped \\lya\\ systems (MSDLAs) enable analysis of tens of elements in individual galaxies \\citep{phw03}. The metal-strong criteria were defined by \\cite[][; hereafter, H06]{shf06} to be absorbers showing ionic column densities that satisfy $\\log$ N(Zn$^+$)$\\ge$13.15 and/or $\\log$ N(Si$^+$)$\\ge$15.95. The MSDLAs, therefore, are absorbers that satisfy the metal-strong criteria and also the DLA criterion on \\nhi. H06 identified a number of MSDLA candidates from the Sloan Digital Sky Survey (SDSS) and presented higher resolution follow-up spectra of the metal-line transitions.   The strong metal absorption lines of these candidates suggests a high metallicity. For many of these MSDLA candidates, however, the \\lya\\ absorption line was bluer than the wavelength coverage in the SDSS spectra, precluding measurements of their H I column densities. Therefore, these absorption systems could not be classified as damped \\lya\\ systems nor could the authors provide metallicity estimates for the gas.  The wavelengths covered by SDSS spectra range from 3800 to 9200 \\AA\\ \\citep{sdssdr6}.  This range allows for the identification of DLAs with redshifts corresponding to $z_{\\rm abs} \\gtrsim 2.2$. For absorbers with 1.5 $\\lesssim$ \\zabs\\ $\\lesssim$ 2.2, the \\lya\\ absorption line lies blueward of the SDSS spectral coverage, but these systems can be observed from ground-based observatories equipped with blue-sensitive spectrometers. With this observational goal in mind, we obtained new spectra from the MMT telescope of various QSO sight lines containing MSDLA candidates using the Blue Channel Spectrograph (BCS).    Here we summarize our classification scheme.  Systems with $\\mnhi \\ge 10^{20.3}$ \\atomspercm\\ are classified as DLAs.  Systems with  H I column densities meeting the DLA criteria  stated above and with measured metal column densities of $\\log$ N(Zn$^+$)$\\ge$13.15 and/or $\\log$ N(Si$^+$)$\\ge$15.95 are classified as MSDLAs.  Candidate MSDLAs are systems identified with strong metal lines in the SDSS spectra (or elsewhere for a few systems) for which there was previously no \\lya\\ coverage and/or no precise measurement of the metal column density. We refer to these candidate MSDLAs as cMSDLAs. Each cMSDLA has had its spectrum taken with the MMT/BCS in order to obtain \\lya\\ coverage.   The gas identified with MSDLAs is analogous to the HI regions found in the Milky Way \\citep[e.g.][]{ss96}. Several heavy elements detected in the gas of MSDLAs are also known to make up interstellar dust in the Milky Way and other nearby galaxies. Dust in the Milky Way is comprised of heavy elements including C, O, Mg, Ni, S, K, Mn, Si, Fe, Al, Ti, and Ca \\citep[e.g.][]{savage+mathis79,jenkins09}. \\cite{shf06} presents the abundance ratios of previously observed MSDLAs showing their ratios approach solar. Dust particles preferentially absorb short wavelengths of light, reddening the color of background objects. \\cite{ehl05} have performed the only extinction estimate for a complete survey of DLAs and set an upper limit to their average reddening of $E_{B-V} < 0.04$\\, mag. Reddening has been studied in the low redshift universe where nearby galaxies can be well resolved. In contrast, the properties of dust in high redshift galaxies is very poorly constrained. QSO absorption systems provide a probe for studying reddening in high redshift galaxies due to absorption of the QSO's light by the dust in the absorber. A series of studies have been performed on the extinction properties of DLAs beginning several decades ago \\citep[e.g.][]{oh84,pfb91}. More recently, \\cite{ml04} leveraged the large dataset afforded by SDSS to examine reddening from $\\approx 100$ DLAs in the SDSS-DR2 using the spectrophotometric observations. Their results show no conclusive evidence for reddening ($E_{B-V} < 0.02$ mag). Reddening in DLAs has also been investigated using the SDSS photometry by \\cite{vcl+06,vpw08}, who report evidence for dust. \\cite{vpw08} found evidence for reddening when comparing the colors of quasars background to a large sample of DLAs to a control sample of QSOs without DLAs. These authors estitimate an average reddening $E_{B-V} \\approx 0.006$\\,mag. Given their large metal column densities, MSDLAs may be expected to show significant reddening.   In this paper we first determine if our candidate MSDLAs meet the DLA and/or the metal-strong criteria and then we search for reddening to see if they contain dust.  We cover the observations and data reduction of our candidate MSDLA sample in Section \\ref{sect:obs}.   The measurements for redshift and H I column density are given in  \\textsection\\ \\ref{sect:nhi}.  We determine metallicities for the cMSDLAs and explore the relationship between [M/H] and the EW of the Si II 1256 line in \\textsection\\ \\ref{sect:metal}.  Section \\ref{sect:dust} covers the search for dust reddening and extinction in our data sample and presents the results.  We conclude in $\\S$~\\ref{sec:summary} with a brief summary. ",
        "conclusions": "\\label{sec:summary}  We obtained MMT/BCS spectra of 41 QSO sightlines where previous observations found high metal column densities.  We classify these absorption systems as candidate (c)MSDLAs. 49 absorbers with strong \\lya\\ lines ($W > 5$\\AA) were identified, 43 of which correspond to the targeted cMSDLAs. The redshift for all absorption systems was determined visually by finding the centroid of low ion metal transitions.  The H I column density (\\ncol) for all absorption systems was measured by visually fitting a Voigt profile to the absorber's \\lya\\ line.  Out of 49 absorbers, 42 met the DLA criterion (\\ncol$\\ge 20.3$ \\atomspercm) including all but two of the targeted cMSDLAs. The MSDLAs and cMSDLA show higher systematic H I column densities than comparable systems, in line with the expectation that selecting systems with high metal column density will also select systems with high H I column density.  Using the measured H I column densities and previously measured metal column densities, the metallicity [M/H] for each system was calculated.  MSDLAs and cMSDLAs exhibit systematically higher metallicities than DLAs that are not metal strong.  This suggests that MSDLA sightlines might probe towards the centers of galaxies where metallicity is greater, assuming a metallicity gradient.  Perhaps MSDLAs represent a population of high redshift galaxies that became chemically enriched earlier in their evolution than others.  Using these new calculations, we examined and extended the metallicity vs.\\ \\ion{Si}{2} 1526 EW relation discussed by \\cite{pcw+08}.  If this correlation is set by a underlying mass/metallicity relation, then the MSDLAs may represent the highest mass, gas-rich galaxies at high $z$.  We searched for signatures of dust in the cMSDLAs by studying the reddening of the background quasars.  A power-law fit to the continuum of the SDSS spectra yielded individual power-laws (parameterized by $\\alpha$) that model the continuum of each QSO. The quasars behind cMSDLAs have lower $\\alpha$'s then a control sample ($\\delta\\alpha = 0.29$) indicating significant extinction by dust in the cMSDLAs. Assuming an SMC-like extinction law, we estimate $A_V \\approx 0.076$\\,mag and $E_{B-V} \\approx 0.025$ for these systems. The dust-to-gas ratio for the cMSDLAs was estimated to be $\\log  ( A_V / N_{HI} ) \\approx -22.0$, similar to the ratio for the SMC and roughly following the expected trend of the dust-to-gas ratio scaling with metallicity. The ratio of extinction to dust phase iron column density was estimated to be  $\\left<s_V^{Fe} \\right> \\approx 1 \\times 10^{-17}$.  The MSDLAs represent a unique subset of DLAs which are more chemically enriched, massive, dusty, and possibly evolved then ordinary DLAs. As such, they are excellent targets for future studies of extinction, nucleosynthesis, and searches for stellar light and molecular emission lines at high $z$."
    },
    "1004/1004.3543_arXiv.txt": {
        "abstract": "An analysis of data from the {\\it Spitzer} Space Telescope, {\\it Hubble} Space Telescope, {\\it Chandra} X-ray Observatory, and {\\it AKARI} Infrared Astronomy Satellite is presented for the $z=0.036$ merging galaxy system II~Zw~096 (CGCG~448-020). Because II~Zw~096 has an infrared luminosity of log$(L_{IR}/L_\\sun)=11.94$, it is classified as a Luminous Infrared Galaxy (LIRG), and was observed as part of the Great Observatories All-sky LIRG Survey (GOALS\\footnote{http://goals.ipac.caltech.edu}). The {\\it Spitzer} data suggest that 80\\% of the total infrared luminosity comes from an extremely compact, red source not associated with the nuclei of the merging galaxies. The {\\it Spitzer} mid-infrared spectra indicate no high-ionization lines from a buried active galactic nucleus in this source. The strong detection of the $3.3 \\, \\micron$ and $6.2 \\, \\micron$ PAH emission features in the {\\it AKARI} and {\\it Spitzer} spectra also implies that the energy source of II~Zw~096 is a starburst. Based on {\\it Spitzer} infrared imaging and {\\it AKARI} near-infrared spectroscopy, the star formation rate is estimated to be $120 \\, M_\\sun \\, \\mathrm{yr^{-1}}$ and $> 45 \\, M_\\sun \\, \\mathrm{yr^{-1}}$, respectively. Finally, the high-resolution {\\it B}, {\\it I}, and {\\it H}-band images show many star clusters in the interacting system. The colors of these clusters suggest at least two populations -- one with an age of $1-5$~Myr and one with an age of $20-500$~Myr, reddened by $0-2$ magnitudes of visual extinction. The masses of these clusters span a range between $10^{6}-10^{8} \\, M_\\sun$. This starburst source is reminiscent of the extra-nuclear starburst seen in NGC~4038/9 (the Antennae Galaxies) and Arp~299 but approximately an order of magnitude more luminous than the Antennae. The source is remarkable in that the off-nuclear infrared luminosity dominates the enitre system. ",
        "introduction": "\\object{II~Zw~096} (also known as \\object{CGCG~448-020} or \\object{IRAS~20550+1655}) has an infrared luminosity of log$(L_{IR}/L_{\\sun})=11.94$ and a luminosity distance of 161~Mpc \\citep{Armu09}. Since its luminosity is above $L_{IR} \\geq 10^{11} L_\\sun$, it is classified as a Luminous Infrared Galaxy, or LIRG. From its optical morphology, \\object{II~Zw~096} appears to be a merger of at least two gas-rich spirals. From the optical imaging, \\cite{Arri04} classify \\object{II~Zw~096} as a {\\it Class~III} interacting galaxy, based on the system of \\cite{Surace}, meaning the galaxy has two identifiable nuclei, with well developed tidal tails.  With near-infrared imaging and spectroscopy, \\cite{Gold97} uncovered two extremely red sources to the east of the merging disks that appeared to mark the location of a highly obscured, young starburst.  These authors suggest that \\object{II~Zw~096} is one of a handful of LIRGs (such as \\object{Arp~299} and \\object{VV~114}) that are experiencing enhanced star-formation before the final dissipative collapse of the system. While \\object{II~Zw~096} has previously been observed spectroscopically in the mid-infrared \\citep{Dudl99}, the ground based data were not of high enough signal-to-noise to permit detection of PAH or the fine structure lines. It is also known that \\object{II~Zw~096} hosts an OH megamaser at the B1950 position of R.A.=20$\\mathrm{^h}$55$\\mathrm{^m}$05$\\mathrm{^s}$.3, Dec=+16$\\arcdeg$56$\\arcmin$03$\\arcsec$ \\citep{Baan89}. \\cite{Baan98} and \\cite{Baan06} classified this megamaser as a starburst using the optical line ratios and the radio data.  As part of the Great Observatories All-Sky LIRG Survey (GOALS), we have obtained images and spectra with the {\\it Spitzer} Space Telescope, the {\\it Hubble} Space Telescope ({\\it HST}), the {\\it Chandra} X-ray Observatory, and the {\\it Galaxy Evolution Explorer} ({\\it GALEX}) of a complete sample of LIRGs in the local universe. The GOALS targets consist of all LIRGs found in the {\\it IRAS} Revised Bright Galaxy Sample (RBGS; \\cite{Sand03}), which covers galactic latitudes greater than five degrees and includes 629 extragalactic objects with $60\\,\\micron$ flux densities greater than 5.24~Jy. The median and the maximum redshift of the {\\it IRAS} RBGS are z=0.008 and 0.088, respectively. The GOALS sample includes 179 LIRGs and 23 Ultra Luminous Infrared Galaxies (ULIRGs, $L_{IR} \\geq 10^{12} L_\\sun$), covering the full range of galaxy interaction stages from isolated spirals to late stage mergers. GOALS provides an excellent dataset with which to explore the effect of mergers on infrared activity at low redshift. A critical part of the GOALS survey has been to use the {\\it Spitzer} Space Telescope to identify the location of the infrared emission within LIRGs, and to characterize the source of the power. The GOALS project is fully described in \\citet{Armu09}.  Here, we present the first mid and far-infrared imaging and spectroscopy of \\object{II~Zw~096} from {\\it Spitzer}, together with an analysis of the far-ultraviolet and optical imaging ({\\it HST}), X-ray imaging ({\\it Chandra}), and near-infrared spectra ({\\it AKARI}). In \\S\\ref{sec:obs}, we describe the observations and data reduction, in \\S\\ref{sec:result}, we present our results, identifying the location of the bulk of the far-infrared emission, and in \\S\\ref{sec:dis} we discuss the nature of this buried power source. Cosmological parameters $H_0 = 70 \\, \\mathrm{km \\, s^{-1} \\,Mpc^{-1}}$, $\\Omega_m = 0.28$, and $\\Omega_\\Lambda = 0.72$ are used throughout this paper. ",
        "conclusions": "  \\begin{enumerate}  \\item The {\\it Spitzer} imaging reveals that approximately 80\\% of the far-infrared emission of the \\object{II~Zw~096} comes from a compact, off-nuclear starburst with R.A.=20$\\mathrm{^h}$57$\\mathrm{^m}$24$\\mathrm{^s}$.34, Dec=+17$\\arcdeg$07$\\arcmin$39~1$\\arcsec$ (J2000). The estimated $8-1000 \\, \\micron$ luminosity of this source is $ L_{IR} = 6.87 \\times 10^{11} \\, L_\\sun$. The implied star formation rate of this object is about $120 \\, M_\\sun \\, \\mathrm{yr^{-1}}$.  \\item {\\it HST} NICMOS observations show that the off-nuclear starburst is composed of two prominent, red knots, with a number of smaller peaks spread over approximately 10 square arcseconds ($\\sim 6.6 \\, \\mathrm{kpc^2}$). Most of this emission is not seen in the optical imaging data, either from the ground or with the {\\it HST} ACS. The knot to the southwest of this region, source D, is consistent with the peak in the {\\it Spitzer} MIPS 24 and 70 $\\mu$m emission which dominates the far-infrared luminosity of the entire merging galaxy. Assuming a standard mass-to-light ratio for galaxies, the H-band luminosity of source D corresponds to a mass of $1-4 \\times 10^{9} \\, M_\\sun$.  \\item The {\\it Chandra} X-ray imaging shows that the hardest source in the system is source C+D, and it lies slightly below the correlation of X-ray and infrared luminosities seen for starburst galaxies, although within the 2-$\\sigma$ scatter of the relation.  \\item The {\\it Spitzer} IRS and {\\it AKARI} spectra suggest a starburst-dominated system with strong $3.3\\mu$m and $6.2\\mu$m PAH emission, and no evidence of emission from [\\ion{Ne}{5}] $14.3\\mu$m or [\\ion{O}{4}] $25.9\\mu$m in the IRS data.  While the equivalent width of the $3.3\\mu$m PAH feature as measured in the {\\it AKARI} spectrum is typical for starburst galaxies, the equivalent width of the $6.2\\mu$m PAH feature as measured in the IRS spectrum is about 1/2 the value found in pure starbursts.  This may reflect an excess of hot dust surrounding source C+D, which is easier to detect in the narrow IRS slit.  This is consistent with the [\\ion{Ne}{3}]/[\\ion{Ne}{2}] line flux ratio of $\\sim1$ indicating a relatively hard radiation field. The $9.7\\, \\micron$ silicate optical depth, measured from the IRS spectra suggests an $A_V \\geq 19$~mag toward source C+D, indicating a highly buried source.  \\item The ACS $B$ and $I$-band imaging reveals a large number of (super) star clusters in the \\object{II~Zw~096} system, consisting of at least two populations: one $1-5$~Myr and one $20-500$~Myr, which may have formed at different times during the merger. We find no clear association of cluster age with position in the merger. Most of the cluster masses are in the range of $10^6 - 10^8 \\, M_\\sun$.  \\end{enumerate}  Spitzer imaging and spectroscopy of \\object{II~Zw~096} has revealed a powerful, young extranuclear starburst.  This starburst is reminiscent of those seen in \\object{NGC~4038/9} (\\object{the Antennae Galaxies}) and \\object{Arp~299}, but it is more luminous (more than an order of magnitude more luminous than the Antennae starburst), and it is responsible for nearly all the infrared luminosity in \\object{II~Zw~096} (compared to only $10-15\\%$ in \\object{the Antennae} or \\object{Arp~299}). Source D in \\object{II~Zw~096} is one of the most extreme buried extra-nuclear starbursts yet discovered in the local Universe."
    },
    "1004/1004.3891_arXiv.txt": {
        "abstract": "The K-shell emission line of neutral irons from the Galactic center (GC) region is one of the key for the structure and activity of the GC. The origin is still open question, but possibly due either to X-ray radiation or to electron bombarding to neutral atoms. To address this issue, we analyzed the Suzaku X-ray spectrum from the GC region of intense neutral iron line emission, and report on the discovery of K$\\alpha$ lines of neutral argon, calcium, chrome, and manganese atoms. The equivalent widths of these K$\\alpha$ lines indicate that the metal abundances in the GC region should be $\\sim$1.6 and $\\sim$4 of solar value, depending on the X-ray and the electron origins, respectively. On the other hand, the metal abundances in the hot plasma in the GC region are found to be $\\sim$1--2 solar. These results favor that the origin of the neutral K$\\alpha$ lines are due to X-ray irradiation. ",
        "introduction": "ASCA found clumpy structures of the 6.4~keV line in the Sagittarius (Sgr) B2 and the Radio Arc regions (Koyama et al. 1996). Since the clumps correspond to giant molecular clouds, we refer to them as a \"neutral clump\" hereafter.  The X-ray spectrum of the Sgr B2 neutral clump has a prominent 6.4~keV line on the continuum emission with deep absorption of $N_{\\rm H}\\sim10^{24}$~cm$^{-2}$ (Murakami et al. 2000). Recently, many new neutral clumps M\\,0.74$-$0.09, M\\,0.51$-$0.10 (Sgr~B1), G\\,0.174$-$0.233, M\\,359.47$-$0.15, M\\,359.43$-$0.07, M\\,359.43$-$0.12, M\\,359.38$-$0.00 have been discovered \\citep{Ko07c, Yu07, No08, Fu09, Na09}.  The origin of a neutral clump, particularly the 6.4~keV line, is one of the open questions among high energy phenomena in the Galactic center (GC) region. Koyama et al. (1996) and Murakami et al. (2000) suggested that the 6.4~keV emission from the neutral clumps is due to the K-shell ionization of iron atoms by external X-rays, possibly from the super-massive black hole, Sgr~A* (an X-ray reflection nebula; XRN). Although the present X-ray luminosity of Sgr~A* is $\\sim10^{33-35}$~erg~s$^{-1}$~cm$^{-2}$ (Baganoff et al. 2001, 2003), the luminosity required for the 6.4~keV emission is $10^{38-39}$~erg~s$^{-1}$~cm$^{-2}$. Therefore, Sgr~A* in a few hundred years ago would have been $10^{3-6}$ times brighter than now (e. g. Koyama et al. 1996). On the other hand, Yusef-Zadeh et al. (2007) proposed that the origin of the 6.4~keV emission would be low energy cosmic-ray electrons ($E_{\\rm e}=$10--100~keV), because they found that the X-rays correlated with non-thermal radio filaments.  Neutral atoms of lighter elements, such as Ar and Ca, should also exist in the neutral clumps, and hence would emit K-shell emission lines. These would provide new information to constrain the origin of the neutral clumps in the GC region. However only the K-shell lines of neutral Fe and Ni atoms have been discovered so far (but see Fukuoka et al. 2009).  The Sgr~A region is the best place for the search of neutral K-shell lines of various elements because many bright neutral clumps have been found at ($l$, $b$)$\\sim$(\\timeform{0D.1}, \\timeform{-0D.1}) \\citep{Tsu99} and many observations have been performed with Suzaku \\citep{Ko07b, Hyo09}.  We re-analyzed the X-ray data obtained with the X-ray Imaging Spectrometers (XIS; Koyama et al. 2007a) aboard Suzaku \\citep{Mi07} and found emission lines from neutral Ar, Ca, Cr, and Mn atoms. This paper reports on a detailed analysis, results, and discussion on the origin. ",
        "conclusions": "\\subsection{The GC hot plasma} The absorption column density toward the GC region ($Abs1$) is $\\sim6.7\\times10^{22}$ H~cm$^{-2}$, a typical value to the GC region (e.g., \\cite{Mu04}). This is the first constrain that the low temperature $kT=1.01^{+0.01}_{-0.02}$~keV plasma is also in the GC region.  \\citet{Ko07b}, using the same data set of this paper, reported the temperature of the GC hot plasma is $kT=6.5\\pm0.1$~keV. The 1~keV plasma also emits Fe\\emissiontype{XXV}~K$\\alpha$, but no significant Fe\\emissiontype{XXVI}~K$\\alpha$. Since \\citet{Ko07b} ignored the Fe lines from the 1.0~keV plasma, they under-estimated the flux ratio of Fe\\emissiontype{XXVI}~K$\\alpha$ to Fe\\emissiontype{XXV}~K$\\alpha$, i.e., the plasma temperature. Indeed, about 20\\% of the Fe\\emissiontype{XXV}~K$\\alpha$ line may come from the $1.0$~keV plasma. The flux ratio of the Fe lines is estimated to be $0.35$ and $0.42$ about the 6.5~keV and 7.0~keV plasmas by the APEC model, respectively. Taking the contribution of the low temperature plasma into account, the result in \\citet{Ko07b} is consistent with our work.  The Fe and Ni abundances of 1.1--1.2 and 1.3--2.0 solar are more precise than the previous work \\citep{Ko07b}. The abundances of lighter elements,  Si and S, (Ar, Ca), were respectively determined to be 2.3--2.6 and 1.8--2.0 solar, for the first time. The highly ionized Cr and Mn lines in the GC region were discovered for the first time, but details are beyond the scope of this paper.  \\subsection{Origin of the Neutral Clump}  We discovered the K-shell lines of neutral Ar, Ca, Cr, and Mn from the bright neutral clump toward the Sgr~A region at the significance levels of 6.8, 9.6, 6.1, and 5.4~$\\sigma$, respectively. The absorption column densities, $N_{\\rm H}$($Abs2$) in the Source and Reference regions are $12.0(\\pm1.1)\\times10^{22}$ and $12.7^{+1.2}_{-1.8}\\times10^{22}$ cm$^{-1}$, respectively.  The photon index, $\\Gamma$ of the continuum (power-law) is $1.87\\pm0.04$, which is consistent with the result of \\citet{Ko09} in the same region. The equivalent widths of the neutral K$\\alpha$ lines to the power-law continuum are $\\sim$140, 83, 24, 22, 1150, 83~eV in Ar, Ca, Cr, Mn, Fe, and Ni, respectively.  \\begin{figure}[t] % \\begin{center} \\FigureFile(80mm,50mm){figure5.eps} \\end{center} \\caption{ Equivalent widths of K$\\alpha$ line of various neutral atoms. Black and red lines are the calculated values for the X-ray (XRN) and electron (LECRe) scenarios, respectively. The data points marked with the open circles are the observed value in our work. Errors were estimated at the 90\\% confidence level. The black and red dashed lines are to guide eyes, which are XRN and LECRe scenarios in 1.6 solar and 4.0 solar abundances, respectively. } \\label{fig:EW} \\end{figure}  The major possibility for the origin of the neutral clumps is the ionization of neutral atoms by either low energy cosmic-ray electrons (LECRe: \\cite{Yu07}) with an energy of 10--100~keV, or X-rays of external sources (XRN: \\cite{Ko96, Mu00}). The ionization cross-sections for these processes are very different. They also produce a continuum emission: the bremsstrahlung in the LECRe scenario and Thomson scattering in the XRN scenario. As a result, the two scenarios make different X-ray spectra. In particular, sharp contrasts are the photon index of the continuum and the equivalent widths of the neutral lines.  The continuum emission in the X-ray spectrum produced by the LECRe scenario is an integration of the bremsstrahlung by electrons with various energies. According to \\citet{Ta03}, we calculated the X-ray spectrum with various indexes of their energy distributions into a molecular cloud. As a result, the photon index of $\\Gamma=1.9$ corresponds to an index of the LECRe source spectrum of $\\alpha\\sim3.0$. On the other hand, the Thomson scattering do not change the photon index from the incident spectrum of the external source.  The equivalent width of the line is a good indicator for constraint of the origin. According to the calculation in \\citet{Ta03}, the equivalent widths produced by the LECRe with the index $\\alpha=3.0$ were estimated as shown with the red solid line in figure~5. Here, the elemental abundances in the neutral clump were assumed to be solar. On the other hand, \\citet{Mu00} estimated the XRN spectrum by a numerical simulation. We improved the simulation \\citep{Mu00} to the other relevant elements. These are listed in table~3 and plotted in figure~5 together with the best-fit results.  From figure 5, the equivalent width of each element in the solar abundance is larger than that expected in the both of the LCREe and XRN scenarios. For the LCREe scenario, the abundances of the molecular cloud must be $\\sim$4-times larger than the solar value, while the XRN scenario requires $\\sim$1.6 solar abundance.  Since the molecular cloud may be formed by condensation of the ambient materials, the abundances should be similar to those of 1--2 solar in the GC hot plasma. Accordingly, the neutral lines from the GC region likely come from the fluorescence by external X-rays. The photon index of $\\Gamma\\sim1.9$ is similar to the other neutral clumps in the Sgr~B and C regions (Koyama et al.2007c; Nobukawa et al. 2008; Nakajima et al. 2009), which suggests that the irradiating source is a single object, possibly a super-massive black hole, Sgr~A*(e.g., \\cite{Mu00})."
    },
    "1004/1004.3296_arXiv.txt": {
        "abstract": "In a recent paper, four of the present authors proposed a class of dark matter models where generalized parity symmetry leads to equality of dark matter abundance with baryon asymmetry of the Universe and predicts dark matter mass to be around 5 GeV. In this note we explore how this model can be tested in direct search experiments. In particular, we point out that if the dark matter happens to be the mirror neutron, the direct detection cross section has the unique feature that it increases at low recoil energy unlike the case of conventional WIMPs. It is also interesting to note that the predicted spin-dependent scattering could make significant contribution to the total direct detection rate, especially for light nucleus. With this scenario, one could explain recent DAMA and CoGeNT results. ",
        "introduction": "It is now widely accepted that almost a quarter of the mass-energy in the Universe is dark matter and one of the major challenges of particle physics and cosmology is to discover the nature of the dark matter. Since the standard model of particle physics does not contain any stable particle that can play the role of dark matter, this provides evidence for physics beyond standard model (BSM) and many BSM scenarios have been proposed that include stable or very long-lived Weakly Interacting Massive Particles (WIMPs) which can play this role~\\cite{Bergstrom:2009ib}. Dark matter being pervasive in our galaxy with an energy density of $\\rho_{DM}\\simeq 0.3$\\,GeV/cm$^3$, it could be observable by detection of nuclear recoils produced when it scatters off nuclei in a very low-background detector~\\cite{Goodman:1984dc,Freese:1987wu}. The recoil energy distribution which is in the keV range could provide clues to the nature of the WIMP.  Among the direct detection experiments, {\\it e.g.}, CDMS~\\cite{Ahmed:2008eu} and XENON10~\\cite{Angle:2007uj} have not found any signal from WIMPs and set the most stringent constraints on the WIMP-nucleon elastic scattering cross section. On the other hand, DAMA collaboration has reported an annual modulation signal in the scintillation light from their DAMA/NaI and DAMA/LIBRA experiments, which is interpreted as evidence of dark matter~\\cite{Bernabei:2003za,Bernabei:2008yi}. The CDMS II collaboration has observed two possible dark matter signal events for an expected background of $0.8\\pm 0.2$ events~\\cite{Ahmed:2009zw}. More recently the CoGeNT collaboration has published their results from the ultra low noise germanium detector with a very low energy threshold of $0.4$ keVee in the Soudan Underground Laboratory~\\cite{Aalseth:2010vx}. Although the observed excess is consistent with an exponential background, it also could be explained by a WIMP in the mass range $5\\sim 10$ GeV, with a rather large WIMP-necleon spin-independent (SI) elastic scattering cross section $\\sim 10^{-40}$ cm$^2$~\\cite{Kopp:2009qt,Fitzpatrick:2010em}.   It is well known that the null experiments have already ruled out the case of canonical WIMP masses $\\sim 100$ GeV to be capable of producing the DAMA results. Yet for a low mass (${\\mathcal O}(10)$ GeV) WIMP, the compatibility is possible.   The light dark matter fits of DAMA and CoGeNT motivate many light DM models; among them, a class of very attractive ones are the asymmetric dark matter (ADM) models. The ADM models are different from the usual WIMP models in that whereas the latter have a relic thermal abundance determined by the thermal `freeze-out', ADM abundance is related to the baryon asymmetry in the universe ~\\cite{Nussinov:1985xr,Kaplan:1991ah,Barr:1991qn,Barr:1990ca,Gudnason:2006ug,Dodelson:1991iv,Fujii:2002aj,Kitano:2004sv, Kitano:2008tk,Farrar:2005zd,Berezhiani:2008gi,Kaplan:2009ag,Kribs:2009fy,Cai:2009ia,An:2009vq}.   Recently we proposed an ADM model~\\cite{An:2009vq} in which the standard model is accompanied by a dark standard (or mirror) model which is a complete duplication of the matter and forces in the visible SM. A mirror symmetry guarantees that prior to symmetry breaking there are no free coupling parameters in the dark sector. This is therefore distinct from models where an arbitrary dark sector is appended to the standard model. Symmetry breaking is assumed to be different in the mirror sector compared to the familiar SM sector so that the model is consistent with cosmology.  There are several ways that the two sectors are connected: the first, of course, is via gravity as every matter would couple to gravity. To understand small neutrino masses in our sector, we invoke the seesaw mechanism and add three right-handed neutrinos. A novel aspect of our model~\\cite{An:2009vq} is that instead of adding RH neutrinos separately to two sectors, we add a common set of three RH neutrinos that provides a second link between the two sectors~\\cite{bere}. Finally, we add a kinetic mixing between the $U(1)$ bosons of the two sectors. Other details of the model are reviewed in Sec.~\\ref{sec3}.  The right-handed neutrinos not only help in understanding of the small neutrino masses by a variation of the usual seesaw mechanism~\\cite{MV}, they also play a crucial role in our understanding of dark matter abundance: in the early universe, the RH neutrinos decay out of equilibrium and generate equal leptonic asymmetry in both sectors. These asymmetries are then transferred into baryonic and mirror-baryonic asymmetries through the sphaleron processes in both sectors. Thus the full weak $SU(2)_L$ group in both sectors are essential to our scenario. The lightest mirror baryon is considered as the dark matter particle. Thus baryogenesis via leptogenesis explains both the origin of matter as well as dark matter, making their number densities equal to each other due to mirror symmetry. This allows us to predict the dark matter mass to be $m_N\\Omega_{DM}/\\Omega_B  \\sim 5$ GeV. The $U(1)-U^\\prime(1)$ kinetic mixing  along with a massive mirror photon helps us to maintain consistency of the model with Big Bang Nucleosynthesis (with a mirror photon mass in the $10-100$ MeV range). The mirror photon, therefore, provides a portal linking the two sectors and makes the direct detection of the dark matter possible. Furthermore, the dark matter in our model has self interaction and as pointed out in~\\cite{An:2009vq}, the self interaction cross section is safely below the bullet cluster constraint.   In this work, we investigate the direct detection of the dark baryons that arise in the class of asymmetric mirror models proposed in \\cite{An:2009vq}. We write down the general operators for neutral dark baryon interaction with the visible sector through a light massive mirror photon portal. We find that the interactions are energy/momentum dependent and the differential cross section has non-uniform angular distribution. These new features are absent in the conventional WIMP case for both spin-independent (SI) and spin-dependent (SD) interactions. This provides a way to distinguish between this type of DM from many familiar DM candidates. We also consider the scenarios when the charged dark baryon $p^\\prime$ or $\\Delta^\\prime$ is the dark matter, in which case there is no such momentum dependence.  The paper is organized as follows: in Sec.~\\ref{sec2} we give a general operator analysis of dark matter and nuclear interaction that applies to the asymmetric dark matter and similar models. In Sec.~\\ref{sec3}, we discuss the implications of the general operator analysis and the energy dependent direct detection cross-section that results for this general case. In Sec.~\\ref{sec4} we present our conclusions. ",
        "conclusions": "\\label{sec4} To summarize, we have presented a general operator analysis of an asymmetric dark matter interacting with nucleons via a mirror photon and applied it to an asymmetric mirror dark matter model suggested by four of us in a previous paper. We note that when the dark matter is neutral under dark electromagnetic forces (zero mirror electric charge), {\\it e.g.}, mirror neutron, it interacts with nucleons via the mirror magnetic dipole moment and electric charge radius. In this case, there is an energy dependence in the direct detection cross section as well as an angular dependence different from the usual massive symmetric WIMP case ({\\em e.g.}, SUSY case). As the sensitivities of dark matter searches improve, one can use these results to pinpoint the detailed nature of dark matter interaction with matter."
    },
    "1004/1004.4649_arXiv.txt": {
        "abstract": "{ The relative distribution of  abundances of refractory, intermediate, and volatile elements in stars with planets can be an important tool for investigating the internal migration of a giant planet. This migration can lead to the  accretion of planetesimals and the selective  enrichment of the star with these elements. We report on a spectroscopic determination of the atmospheric parameters and chemical abundances of the parent stars in transiting planets CoRoT-2b and CoRoT-4b. Adding data for CoRoT-3 and CoRoT-5 from the literature, we find a  flat distribution of  the relative abundances as a function of their condensation temperatures. For CoRoT-2, the   relatively  high lithium abundance and intensity of its Li\\,{\\sc i} resonance line permit us to propose an age of 120 Myr, making this stars one of the youngest stars  with planets to date. We introduce a new methodology to investigate a relation between the  abundances of these stars and the internal migration of their planets. By simulating  the internal migration of a planet in a disk formed only by planetesimals, we are able, for the first time, to separate the stellar fractions of refractory, intermediate, and volatile rich planetesimals accreting onto the central star. Intermediate and volatile element fractions enriching the star are similar and much larger than those of pure refractory ones. This result is opposite to what has been considered in the literature for the accreting self-enrichment processes of stars with planets. We also show that these results are highly dependent on the model adopted for the disk distribution regions in terms of refractory, intermediate, and also volatile elements and other parameters considered. We note however, that this self-enrichment mechanism is only efficient during the first 20 -- 30 Myr or later in the lifetime of the disk when the surface convection layers of the central star  for the first time attain  its minimum size configuration. } ",
        "introduction": "\\cite{gonzalez97} and \\cite{santos01} first identified a metallicity excess in stars with giant planets (SWP) relative  to stars without planets. The tendency has been confirmed using  large samples of stars. Nevertheless, planets have also been observed around low metallicity main-sequence stars \\citep{cochran07}. We note, however, that such a large statistical metal excess found among dwarf stars is not present in  giant stars with planets \\citep{pasquini07}.  Using iron as a primary reference element, this metal excess, has been the source of a number of studies trying to explain this property. Two scenarios have been invoked: a self-enrichment mechanism of normal metal stars and a primordial scenario, in which SWP are the result of the formation of entire metal-rich stars in equally metal-rich natal clouds. Which scenario is  correct  remains an open question to this day.   What about  CoRoT exoplanet host stars? Although there are few cases with which  to perform a statistical study, these stars exhibit a different distributions: the majority  have  solar abundances, one has a mild metal excess, and one  has a  low metal abundance. Even at this early stage,  we are tempted to speculate that CoRoT is exploring a different backyard of the Galaxy with different properties from those of the Solar vicinity. Where  does the bulk of the metal excess of the SWP mentioned above  originate?   The relative distribution of the stellar abundances of refractory elements, with a high condensation temperature ($T_{\\rm C}$), with respect to the volatile elements (low $T_{\\rm C}$), has been widely used in the literature as a tool to investigate the nature of  metal enrichment in SWP (see review papers of \\cite{gonzalez03,gonzalez06b}).  This enrichment mechanism should be triggered  by the action of a hot Jupiter. The observed pile up of hot Jupiters is believed to be the result of migration. These planets formed further away in the protoplanetary nebula and migrated afterwards to the small orbital distances at which they are observed. During inward migration, material (planetesimals) from the disk, depleted in H and He, could be accreted onto the star. The presence of a positive slope of [X/H] versus $T_{\\rm C}$, in which the  abundances of volatile elements  are lower than that of the refractory abundances, has been considered as a possible signature of a self-enrichment mechanism \\citep{smith01}. This is based on  the never proven assumption that stellar accretion only favors near star disk regions rich in refractory elements.    In \\cite{winter07}, we explored by means of numerical simulations  stellar accretion caused by an inward forced planetary migration in a disk formed only by planetesimals that cause  a metallic self-enrichment of the central star photosphere. At present, we attempt to identify this stellar accretion by determining the origin in the disk of these particles (planetesimals). In this way, we can separate the contributions of rocky material enriched differentially with refractory, intermediate, and volatile elements. In this work, we apply this methodology to three different systems, namely: CoRoT-2 \\citep{alonso08}, 3 \\citep{deleuil08}, and 4 \\citep{aigrain08}. The main purpose of this work is to analyze the mentioned slopes of [Fe/H] vs. $T_{\\rm C}$ corresponding to CoRoT-2, CoRoT-3, CoRoT-4, and CoRoT-5 and their interpretations. For cases 2 and 4, we obtained new high resolution spectra. Cases 3 and 5 are taken from the  literature. ",
        "conclusions": "In Fig. \\ref{f:temp}, we present one of our most important result: the relative distribution of the abundances of refractory, intermediate, and volatile elements as a function of their condensation temperatures for four CoRoT stars. A flat distribution is found for all of them, even if CoRoT-3 presents a slight tendency to display a positive slope. Clearly, a more reliable determination of this slope would be obtained if a complete collection of abundances of the highly volatile CNO elements was at hand. In any case, we also note the importance to dispose of measurements of the volatile S and Zn elements.  When considering the metallic abundances of CoRoT-2 and CoRoT-4,  derived in this work, we found that the metallicity of CoRoT-4 is somewhat higher than that proposed by \\cite{bouchy08}, hence confirming  a mild metallic excess. We compare the stellar parameters with those in the literature for CoRoT-2 and CoRoT-4  in Table  \\ref{t:parameters}.  CoRoT-2 star deserves  particular attention because of its significant stellar activity and its youth indicators represented, respectively, by its Ca II H and K lines (Fig. \\ref{f:ca}) and its Li resonance line (Fig. \\ref{f:li-c2}), contrasting sharply with those of the older 1 Gyr star CoRoT-4 \\citep{moutou08}, which does not exhibit  these features. Its  age can be determined from the equivalent width (EW) and abundance of  Li by using diagrams of the Li values as a function of the $T_{\\rm eff}$ values. By considering its measured EW of 124 m\\AA\\, we found, using the diagrams presented by \\cite{torres06,torres08}, that this star falls in the lower dispersion limit of the Li abundances found for the members of the Pleiades cluster, with an age of 119 $\\pm$ 20 Myr \\citep{ortega07}. By using its Li abundance  found here of 2.6, we arrived at the same position in the diagrams of \\cite{torres06} and \\cite{dasilva09}. From all these matches, we infer an age of 120 Myr for CoRoT-2, from which \\cite{bouchy08} proposed an  age of $<$ 500 Myr. This new age implies that of CoRoT-2 is  one of the youngest known two stars with planets, the other being HD 70573, a debris-disk type star of a similar age and Li abundance \\citep{setiawan07}. The $T_{\\rm eff}$ values of these two solar-type stars are in a short temperature interval, where old solar-type stars with planets exhibit a peculiar Li depletion with respect to similar stars without planets \\citep{israelian09}. CoRoT-2 and HD 70573 are unaffected by this Li depletion due to their youth (see also \\cite{pinsonneault09}.    We performed simulations of  planetary migration into an internal disk composed of differentiated particles (planetesimals) with refractory, intermediate, and volatile properties, and applied  these simulations to three CoRoT systems. The accreted material contains large and similar contributions of  {\\it I} and  {\\it V} particles and a very small contribution of pure refractory elements, as shown in Fig. \\ref{f:4}. In other words, accretion is mainly ``cool'' and ``warm'' and not  ``hot'' as largely mentioned in the literature. Because zone {\\it I} exhibits a mixture of chemical properties, mainly containing  common elements such as Fe, both intermediate and volatile contributions produce a flat distribution of [X/H] versus $T_{\\rm C}$, as for the observed three CoRoT systems considered here.   When there is a complete primeval mixing of elements in the disk, a single intermediate zone will be maintained  producing a flat distribution of abundances as a function of $T_{\\rm C}$. However, if the mixing is incomplete or partial, some R and V rocky material will be maintained near and far from the star, respectively. An extended I zone (see top of Fig. \\ref{f:temp}) will control a large part of the slope. In that case, the agreement of the model with observations will  depend mainly on the extended distribution of the element abundances  with $T_{\\rm C}$ $<$ 304 K. Nevertheless, this region, containing volatile and highly volatile elements such as CNO and noble gases, in general,  requires spectra of higher S/N than our own to measure  the abundances.  By considering our O abundances for CoRoT 2 and 4 and the C abundances for CoRoT 3 by \\cite{deleuil08} and that for CoRoT 5 \\citep{rauer09}, we  obtained one collection of quite flat gradients   as presented in Fig. \\ref{f:temp}. We note that CoRoT 3, for which the gradient is the least flat, is a peculiar case because  its ``planet'' is  probably a brown dwarf star. We conclude that to obtain a reliable collection of highly volatile element abundances of CNO and noble gases,  data of even higher S/N  are required.    In any case, independently of whether accretion is important or not, depending on the accreting mass, the main result of this work, is that, for the CoRoT systems considered, flat distributions that appear to be observationally the rule, do not represent the absence of a self-enrichment mechanism as sometimes mentioned in the literature. In contrast, we show here that flat distributions of the elements abundances as function of  $T_{\\rm C}$ could be a natural result of accretion.  Does the discussed type of migration represent true situations for the  CoRoT systems 2, 3, and 4? This is not the case for, at least systems 2 and 3, where unrealistically high disk masses are necessary to bring the planets to their observed final distances with respect to the star \\citep{adams03}. System 4 can be considered an exception. A different  migration of type II in a disk containing gas  \\citep{papaloizou06,armitage07} has to be invoked for CoRoT-2. The CoRoT-3 system is particularly difficult and could represent a challenge to migration theories. If this planetary body, with an exceptionally high mass of 21.66 M$_J$ is a real planet and not a brown dwarf, a type II migration study must first avoid the collision of the planet with the star due to the large planetary eccentric orbits developed, which is the case, at least, for a $\\sim$ 10 M$_J$ \\citep{rice08}. These authors mention that if the gas disk dissipates quickly, the eccentric orbits of these massive planets could eventually be  tidally circularized \\citep{ford06}. Is this the case of CoRoT-3 system?  CoRoT-4 system  with a planet of 0.72 M$_J$, also requires a high primeval disk mass of $\\sim$ 0.2 -- 0.3 M$_{\\odot}$. Even if it were an extreme case, it could be acceptable for a star of 1.2 M$_{\\odot}$. Nevertheless, our simulations indicate that in this case only a very rapid migration of 1000 yr can produce a metal enhancement as observed of [I/H] $\\sim$ 0.1., if occurred in the first 20--30 Myr, when the stellar convective layer  for the first time attains its minimum configuration \\citep{ford99}, This could also be possible in principle, even for the smaller planet of mass 0.467 M$_J$ related to CoRoT-5 \\citep{rauer09}. We also show that these results are highly dependent on the model adopted for the disk distribution regions in terms of refractory, intermediate, and volatile elements and the other parameters considered."
    },
    "1004/1004.1201_arXiv.txt": {
        "abstract": " ",
        "introduction": "In the last few years, it has become more and more clear that single-clock inflation, {\\it i.e.} inflationary models where there is only one degree of freedom driving inflation, can produce large and detectable non-Gaussianities. While standard slow roll inflation cannot produce large non-Gaussianities~\\cite{Maldacena:2002vr}, several models in single field inflation have been proposed that result in a detectable level of non-Gaussianities~\\cite{Alishahiha:2004eh,ArkaniHamed:2003uz,Senatore:2004rj,Chen:2006nt,Cheung:2007st,Senatore:2009gt,Flauger:2009ab}. Furthermore thanks to the development of an Effective Field Theory for the Inflationary perturbations \\cite{Cheung:2007st,Senatore:2009gt}, it has become clear that the Lagrangian for the fluctuations can incorporate large interactions without spoiling the naturalness  of the quasi de Sitter background.  Non-Gaussianities represent a powerful signal about inflation for two reasons. First contrary to what happens for the two-point function, non-Gaussianities are sensitive to the interacting part of the inflaton Lagrangian which is clearly more interesting than the free-field part. Second, the symmetries of the inflationary spacetime force the two-point function of the cosmological fluctuations to be quasi scale invariant. This leaves freedom only to vary its amplitude and to add a small scale-dependence as we vary the inflationary models. This means that measurement of the inflationary two-point function reduces in practice to measuring only a couple of numbers. On the other hand, the same symmetries of the inflationary spacetime are not able to constrain as much correlation functions of higher order.  The possible signals in this case are continuous functions of several variables, very similar to the angular dependence of a scattering amplitude~\\cite{Senatore:2009gt,Babich:2004gb}. This is clearly a much more powerful signal.  So far, the observational search for non-Gaussianity as originated from the period of inflation has concentrated on the three-point function. Several models predicted such a signal as the leading source of non-Gaussianity. The purpose of this paper is to show using the Effective Field Theory of Inflation~\\cite{Cheung:2007st} that there are technically natural models of inflation where the level of the (connected) four-point function is much larger than the three-point function, and where the leading source for detection is indeed in the four-point function.  Using the Effective Field Theory of Inflation we show that it is possible to impose on the inflaton fluctuations both an approximate continuous shift symmetry and an approximate parity symmetry that, for models not close to the de Sitter limit, allow for a unique large quartic operator, $\\dot\\pi^4$, where $\\pi$ represent the inflationary fluctuations. These symmetries approximately forbid all  cubic terms.  The resulting shape of the four-point function is unique. A similar construction is carried out for models near the de Sitter limit where more that one shape becomes possible. ",
        "conclusions": "We have shown that in single-clock inflation it is possible to generate a four-point function that could be detected even in the absence of a detectable three-point function. This is possible by imposing an approximate continuos shift symmetry and an approximate parity symmetry in the Effective Lagrangian for the inflationary fluctuations. For models away from the near de Sitter limit, there is a {\\it unique} interaction operator that gives rise to the large four-point function: $\\dot\\pi^4$. This results in only {\\it one} possible shapes for the four-point function. Also in the models close to de Sitter  the four-point function can be the leading signal, and in this case there is more that one possible shape. This motivates us to undertake an analysis of the WMAP data in search for such a signal~\\cite{kensen1}.  \\subsubsection*"
    },
    "1004/1004.3204_arXiv.txt": {
        "abstract": "Using the {\\sc torus} radiative transfer code we produce synthetic observations of the 21\\,cm neutral hydrogen line from an SPH simulation of a spiral galaxy. The SPH representation of the galaxy is mapped onto an AMR grid, and a ray tracing method is used to calculate 21\\,cm line emission for lines of sight through the AMR grid in different velocity channels and spatial pixels. The result is a synthetic spectral cube which can be directly compared to real observations. We compare our synthetic spectral cubes to observations of M31 and M33 and find good agreement, whereby increasing velocity channels trace the main disc of the galaxy. The synthetic data also show kinks in the velocity across the spiral arms, evidence of non-circular velocities. These are still present even when we blur our data to a similar resolution as the observations, but largely absent in M31 and M33, indicating those galaxies do not contain significant spiral shocks. Thus the detailed velocity structure of our maps better represent previous observations of the grand design spiral M81. ",
        "introduction": "The last decade has seen huge advances in computational modelling of galaxies. Allied with observational developments this provides increased opportunities to study and understand the interstellar medium (ISM). Comparing models and observations provides valuable insight into the dynamics and evolution of the ISM, in particular translating the physics incorporated in the simulations into observable features.  High resolution hydrodynamical and magneto-hydrodynamical models can now begin to examine how the interstellar medium is swept into giant molecular clouds, and ultimately forms stars \\citep{Wada99,Shetty06,Wada08,Dobbs08,Tasker09,DGCK08}.  The ability of supernovae to trigger molecular cloud formation can also be studied using numerical models. \\citep{deAvillez_2001,dib_2006}. At the same time, high resolution surveys now provide data on the distribution and properties of molecular clouds \\citep{Heyer98,Eng03,Ros07,Brunt09}.  Advances in observational capability also provide increased insight into the characteristics of the ISM in external galaxies e.g. recent H{\\sc{i}} observations of the nearby spiral galaxies M31 \\citep{chemin_2009} and M33 \\citep{Putman09} and the irregular galaxies Holmberg~{\\sc{II}} \\citep{rhode_1999} and the LMC \\citep{kim_2007}.  Additionally, statistical properties of H{\\sc{i}} in galaxies can be studied using larger samples obtained through surveys \\citep{walter_2008, tamburro_2009}. Thus comparing simulations and observations is particularly timely; however in order to provide a direct comparison the output of the simulations needs to be processed to produce the same data products, such as spectral line datacubes and maps, which are obtained from observations.  An obvious application of such synthetic observations is to test whether numerical models can reproduce the structure of the interstellar medium seen in observations, thus synthetic observations provide an important validation test of numerical models. Synthetic observations may also be generated with very high spatial resolution compared to observed data, hence simulations can show what we will view in nearby galaxies with future facilities. Moreover, synthetic observations are derived from simulated objects with known properties (e.g density, temperature, velocity) whereas when dealing with real observations these properties are calculated from the observations. Within the Milky Way we commonly rely on radial velocities to determine distances, however for simulations, features can readily be located both in velocity and cartesian space. Hence working with synthetic observations makes it easier for features in the observations to be related to the physical processes from which they result.  An important effect in this paper is the impact of velocity structure on H{\\sc{i}} observations.  Previous comparisons of velocity structure in simulated spiral galaxies with observed data range from simply identifying clouds and plotting their position in velocity space \\citep{Dobbs06,Baba09}, to sophisticated radiative transfer models \\citep{Chak09,Naray09} and synthetic galactic plane surveys \\citep{gomez_2004,douglas_2010}. \\citet{Dobbs06} plotted the distribution of molecular clouds in velocity space to show that they reproduce the distribution in the Outer Galaxy reasonably well. \\cite{Baba09} plotted the location of cold gas using the velocities of the gas in their models and an assumed rotation curve, and highlighted the disparity with the actual location of the gas, illustrating the so called `finger of God' effect.  On larger scales, \\citet{Naray09} and \\citet{Chak09} perform analysis on SPH (smoothed particle hydrodynamics) simulations of interacting galaxies, using radiative transfer codes to determine CO emission and spectral energy distributions respectively. \\cite{gomez_2004} and \\cite{douglas_2010} generate synthetic galactic plane surveys and compare their velocity structure to real observations from the Leiden/Dwingeloo H{\\sc{i}} survey and the Canadian Galactic Plane Survey respectively.  On extragalactic scales H\\,{\\sc{i}} is a powerful tracer used to analyse galactic structure, e.g. tracing spiral arms and tidal tails, and showing holes and cavities in the ISM. Synthetic H\\,{\\sc{i}} maps have been constructed previously \\citep{Wada00,Dib05}, though for irregular rather than spiral galaxies. These studies compared synthetic H\\,{\\sc{i}} maps with data from Holmberg II and the LMC, concluding that the structure in these galaxies is primarily due to a combination of turbulence and thermal and gravitational instabilities, with supernovae contributing to a lesser degree. In this paper, we produce synthetic H\\,{\\sc{i}} maps of a simulated grand design spiral galaxy \\citep{DGCK08} in which the structure is instead dominated by spiral density waves, and the gas velocities reveal information about the spiral structure. Velocity gradients reveal streaming motions, due to the presence of spiral density waves (or at least stellar spiral arms which rotate at lower $\\Omega(r)$ than the stars, see \\citet{Dobbs09,wada_2004}) which produce spiral shocks in the gas.  As our model galaxy is a perfectly isolated, grand design system, which has no effects from interactions with companions or high velocity clouds, we expect to see only effects due to spiral shocks.  \\citet{Visser1980} first showed kinks in the velocities along the spiral arms in M81, due to non-circular motions, typically with a magnitude of a few 10's of km s$^{-1}$ \\citep{Adler96}. We expect to see similar effects in our synthetic observations and our idealised model system allows us to study these effects without the additional complexity introduced by environmental interactions.  We used an SPH code to model the galaxy (unlike the grid based calculations by \\citet{Dib05} and \\citet{Wada00}) which we then combine with {\\sc torus}, a grid-based radiative transfer code.  Thus we have the further complication of converting between SPH and a grid code. We begin in Section~\\ref{sec:method} by describing the method used to generate synthetic observations, including the SPH to grid conversion. The technique is used to produce spectral cubes for galaxies with orientations like M31 and M33, which are assessed and compared to real observations of M31 and M33, in Section~\\ref{sec:results}, as a means of validating the method. Although these galaxies do not display grand design structure to the extent of M81 (indeed M33 is a flocculent spiral), we have access to high resolution H{\\sc{i}} data for these nearby systems. ",
        "conclusions": "We have presented a method for generating synthetic spectral cubes of the 21~cm hydrogen line from SPH simulations of spiral galaxies. The method successfully maps the density, temperature and velocity from the SPH particles onto an AMR grid, while preserving important structures (e.g. spiral arms and spurs) and accurately representing the total mass. Synthetic data cubes are generated using a ray tracing method. The synthetic data show good agreement with observations of M31 and M33 whereby increasing velocity channels trace out the main disc of the galaxy. Velocity contours of the synthetic data show perturbations due to non-circular motions, similar to those already observed in M81 \\citep{Adler96,Visser1980,rots_1975}, which are not seen in the observations of M31 and M33.  Our model galaxy is a grand design spiral galaxy and shows velocity structure associated with the spiral perturbation.  The method of generating synthetic observations can also be applied to simulations in which velocity structure is generated by other mechanisms. Internal mechanisms, such as self-gravity and stellar feedback, can generate velocity structure, and indeed are dominant in flocculent spiral galaxies.  External influences can also affect the {H\\sc{i}} morphology of a spiral galaxy e.g. interactions with a companion galaxy, high velocity clouds or the intracluster medium.  As the majority of galaxies reside in groups or clusters, it is likely that external environmental effects will influence the {H\\sc{i}} structure of most real galaxies; our model galaxy is perfectly isolated and our results show idealised behaviour in the absence of external influences.  A large parameter space could potentially be studied, involving galaxies with different internally and externally generated velocity structure. Some of these aspects can already be modelled (e.g. a galaxy undergoing an interaction, \\cite{Dobbs09}) and some require further model development (e.g inclusion of stellar feedback, which is ongoing). A large number of models must be be run in order to generate the SPH input for the radiative transfer calculation. Once such a library of simulations was available, one could compare these with real observations in order to understand the velocity structure of spiral galaxies. However given the potentially complex nature of the velocity structures generated there would need to be a reliable way of matching an observed galaxy with its counterpart in the synthetic observation library. In our present models, we can relate the strength of the spiral shock, and the inclination of the galaxy, to the size of the perturbations in the renzograms. When we include additional physics, e.g. self gravity and feedback, we can investigate whether these perturbations are still observable (for a given shock strength) or whether they are overwhelmed by other motions in the gas.  Our method has also recently been applied to an observer placed inside the galaxy \\citep{douglas_2010} to generate a synthetic galactic plane survey.  H\\,{\\sc{i}} self absorption features and the conversion of cold H\\,{\\sc{i}} to molecular clouds, as seen in a real galactic plane survey (e.g. \\cite{Taylor03,Stil06}), are seen in the synthetic data. A powerful future application of using a simulation for generating a survey is that given a series of time frames we will be able to trace the evolution of molecular clouds, and related observed features, to the physical properties of the material in the cloud. The method can also be extended to generate synthetic observations of other tracer species, such as CO, or dust, for simulated surveys or simulated external galaxies."
    },
    "1004/1004.3032_arXiv.txt": {
        "abstract": "To date, mid-infrared properties of Galactic black hole binaries have barely been investigated in the framework of multi-wavelength campaigns. Yet, studies in this spectral domain are crucial to get complementary information on the presence of dust and/or on the physical processes such as dust heating and thermal \\textit{bremsstrahlung}. Here, we report a long-term multi-wavelength study of the microquasar \\grs1915. On the one hand, we aimed at understanding the origins of the mid-infrared emission, and on the other hand, at searching for correlation with the high-energy and/or radio activities. We observed the source at several epochs between 2004 and 2006 with the photometer IRAC and spectrometer IRS, both mounted on the \\spitzer\\ \\textit{Space Telescope}. When available, we completed our set of data with quasi-simultaneous \\rxte/\\intl\\ high-energy and/or Ryle radio observations from public archives. We then studied the mid-infrared environment and activities of \\grs1915\\ through spectral analysis and broad band fitting of its radio to X-ray spectral energy distributions. We detected polycyclic aromatic hydrocarbon molecules in all but one IRS spectra of \\grs1915\\ which unambiguously proves the presence of a dust component, likely photoionised by the high-energy emission. We also argue that this dust is distributed in a disc-like structure heated by the companion star, as observed in some Herbig Ae/Be and isolated cool giant stars. Moreover, we show that some of the soft X-ray emission emanating from the inner regions of the accretion disc is reprocessed and thermalised in the outer part. This leads to a mid-infrared excess that is very likely correlated to the soft X-ray emission. We exclude thermal \\textit{bremsstrahlung} as contributing significantly in this spectral domain. ",
        "introduction": "In a multi-wavelength study of microquasars, the infrared presents a particular interest as the accretion disc, the jets, or the companion star may all be detected. Nevertheless, most of the previous studies focused on the accretion$-$ejection phenomena as seen is X-ray/near-infrared/radio correlations, and did not  take the mid-infrared (MIR) emission into account \\citep[see \\textit{e.g.}][]{1998Mirabel, 2002Ueda, 2002Corbel, 2003Corbel, 2003Chaty, 2005Homanb, 2006Chatya, 2006Russell}.  Yet, getting both MIR photometric and spectroscopic information is crucial to investigate the presence of dust, the disc illumination, thermal \\textit{bremsstrahlung} from the accretion disc's wind, or the contribution of relativistic ejecta.  \\subsection{\\grs1915}  Discovered by the WATCH all-sky X-ray monitor on board the \\textit{GRANAT} satellite, on 1992 August 15 \\citep{1992Castro, 1994Castro}, \\grs1915\\ is the first microquasar in which apparent superluminal radio ejecta were detected \\citep{1994Mirabelb}. The nature of its companion star was the subject of debate until \\citet{2001Greinera} unambiguously showed that it was a K/M red giant by detecting CO absorption features in its near-infrared (NIR) spectrum. Moreover, the orbital period of the system and the mass of the compact object were found to be $33.5\\pm1.5$~days \\citep[recently refined to $30.8\\pm0.2$~day,][]{2007Neil} and $14\\pm4$~\\msun, respectively \\citep{2001Greinerb}. The inclination is $66^{\\circ}\\pm2^{\\circ}$, and estimates of the distance fall in the range 6$-$12~kpc \\citep{1996Chaty, 1999Fender, 2004Chapuis}.  \\grs1915\\ is strongly variable, on time scales from seconds to days. Using extensive \\rxte\\ timing observations, \\citet{2000Belloni} showed that its X-ray behaviour could be divided in 12 distinct luminosity classes \\citep[up to 14 today,][]{2002Klein, 2003Hannikainen}, and that the source was carrying out transitions between three canonical spectral states, labelled A, B (strong disc domination), and C (corona-dominated, no disc). Previous multi-wavelength studies showed the existence of a strong connection between the accretion disc instabilities and plasma outflows. In particular, discrete ejecta, emitting through optically thin synchrotron, are believed to be triggered during the transition between the C and A states, expanding adiabatically in the environment and detectable gradually from the NIR to the radio domains \\citep[see \\textit{e.g.}][]{1996Mirabel, 1997Fender, 1998Mirabel, 1998Eikenberry, 2000Eikenberry, 2008Rodrigueza, 2008Rodriguezb}. Moreover, the other known class of radio ejecta $-$ the compact jets, emitting simultaneously from the radio to the NIR through optically thick synchrotron $-$ are only detected in the $\\chi$ luminosity class, which is only seen in the C state. The presence of such a jet is characterised by a flat spectrum with a roughly constant flux density between 50 and 100~mJy.  Long periods of the $\\chi$ luminosity class during which compact jets are present are called \\textit{plateau}, and often precede or follow a giant ejection \\citep[see \\textit{e.g.}][]{1996Foster, 1997Pooley, 1999Fender, 2000Dhawan, 2002Klein, 2003Fuchsb}.  \\subsection{Previous MIR studies of microquasars}  \\citet{2002Koch} presented an \\textit{ISO} spectrophotometric study of Cygnus~X$-$3 in quiescence, and found the MIR continuum to be due to free-free emission from the winds of the Wolf-Rayet companion star with perhaps a contribution from a cold dust component \\citep[see][for similar conclusions on SS~433]{2006Fuchs}. Moreover, \\citet{2003Fuchsa} reported ISOCAM photometric data of the \\grs1915 obtained at two different epochs, during a flaring activity and a \\textit{plateau} state. On the one hand, they showed that, despite strong uncertainties, the MIR flux of \\grs1915\\ had likely increased between the two observations, and they argued that during the \\textit{plateau}, the MIR emission of the source was likely due to a compact jet, without excluding \\textit{bremsstrahlung}. They, on the other hand, excluded the contribution of dust, which is the opposite of the conclusion reached by \\citet{2006Muno} to explain the MIR excess they detected in the \\spitzer/IRAC SEDs of A0620$-$00 and XTE~J1118+480 in quiescence. Finally, \\citet{2007Migliari} argued that X-ray/UV irradiation of the disc in the thermal state, and compact jet in the hard state might be responsible for the excess they detected in the MIR emission of GRO~J1655$-$40; the same conclusion was reached  by \\citet{2007Gallo} concerning the quiescence of A0620$-$00, V404~Cyg, and XTE~J1118+480, claiming that dust component is not statistically necessary. \\newline  In this paper, we report a long-term multi-wavelength study of \\grs1915\\, focusing on its spectroscopic and photometric MIR emission.  It aimed at understanding its origins, as well as its possible connection with the high-energy and radio domains. The observations and the data analysis are presented in Sect.~2, while Sect.~3 and Sect.~4 are devoted to the analysis of the broad band X-ray to MIR SEDs of the source built with high-energy and \\spitzer\\ data. We discuss the outcomes in Sect.~5 and we give our conclusions in Sect.~6. ",
        "conclusions": "We presented a multi-wavelength study of \\grs1915\\ whose outcomes suggest that, in the absence of discrete or continuous ejecta, the MIR continuum of the source is mainly due to the X-ray irradiation of the accretion disc and to a photoionised dust component. This might have consequences on the interpretation of the MIR emission of microquasars in presence of compact jets. Indeed, dust might be ubiquitous around isolated compact objects and X-ray binaries, because of mass transfer during the common envelope phase or material from supernova fallback. If so, compact jets could contribute less than expected at infrared wavelengths, with perhaps a cutoff frequency in the millimeter domain. To confirm our results, it is therefore crucial to increase the sample of microquasars and systematically observe them through MIR spectroscopy, as this is the only way to obtain firm information on both their environment and their continuum. In particular, studying microquasars whose variation time scales are longer than the \\grs1915\\ ones, and which do not exhibit such rapid transitions between spectral states would strongly facilitate the multi-wavelength observations and would allow to reach definitive conclusions."
    },
    "1004/1004.4562_arXiv.txt": {
        "abstract": "In this work Gamma Ray Burst (GRB) data is used to place constraints on a putative coupling between dark energy and dark matter. Type Ia supernovae (SNe Ia) constraints from the Sloan Digital Sky Survey II (SDSS-II) first-year results, the cosmic microwave background radiation (CMBR) shift parameter from WMAP seven year results and the baryon acoustic oscillation (BAO) peak from the Sloan Digital Sky Survey (SDSS) are also discussed. The prospects for the field are assessed, as more GRB events become available. ",
        "introduction": "The nature of dark energy  and dark matter  remains an outstanding open problem in cosmology. In spite of the success of the $\\Lambda$CDM parameterization, one must consider more complex models in order to cast some light on the substance of the dark components of the universe. In this work one considers models with interacting \\dedm\\ components. There are several useful tools to probe the phenomenology of these models, such as CMBR data \\citep{wmap7}, BAO \\citep{BAOprim, BAOsec}, SNe data \\citep{SneSDSS} and the deviation from the virial equilibrium of galaxy clusters \\citep{intmodel1,intmodel2,abdalla2009}. It has been suggested \\citep{schaefer2002,bertolamisilva2006} that GRB may be used to extend the Hubble diagram to high redshifts, greater than $z=5$. At these epochs the Universe was dominated by dark matter, from which follows that this tool is less sensitive to dark energy. However, for models where dark energy and matter are coupled \\citep{amendola,intmodel1} or unified \\citep{kamenshcik,GCGprimer,bento2003,GCGwmap5}, GRBs might be a particularly usefull tool \\citep{bertolamisilva2006}.  In the late 1960s, the Vela array of military satellites detected flashes of radiation originating in apparently random directions in space. The observed bursts lasted between tens of milisecond and thousands of seconds, and were composed of soft ($0.01$ to $1~MeV$) gamma rays. Subsequently space missions such as the US Apollo program and the Soviet Venera probes confirmed the existence of the GRBs, even though its rate of occurrence was virtually unknown until the deployment of the Compton Gamma Ray Observatory, in 1991. This observatory was equipped with a sensitive gamma-ray detector, the Burst and Transient Source Explorer (BATSE) instrument which was able to detect one or two events per day. The collected data allowed to divide GRBs into two categories: short duration bursts ({\\it short} bursts) and long duration bursts ({\\it long} bursts). The former usually last for less than two seconds and are dominated by high energy photons; the latter last longer than two seconds and are dominated by lower energy photons. However, this distinction is not always clear.  The physical origin of GRBs has been debated for a long time, before their exact position and a reliable estimate of their distance was lacking (see e.g. \\citep{bertolami1999} and references therein). In 1997, several GRBs were detected by the BeppoSAX sattelite. A GRB {\\it prompt} emission is followed by an {\\it afterglow} emission composed by all wavelengths. Depending on its brightness, an afterglow can last from days to months after the burst itself, the {\\it transient} phase. The detection of the afterglow did manifold the information on GRBs. Through their afterglow, GRB's X-ray, optical and radio counterparts were observed, as well as their redshifts \\citep{review1}, confirming the cosmological origin of most, if not all, the GRBs. When, in 2003, the long GRB 030329 was discovered and linked with the supernova SN2003dh \\citep{GRB&SN}, it became clear that GRBs are linked with the release of gravitational energy during the collapse of stellar mass objects. GRBs are most likely collimated, considering they reach integrated luminosities up to $L \\sim 10^{53}~erg~s^{-1}$, making it hard to associate them with an astrophysical object otherwise. This high energy release creates an outflow that expands relativistically. Two forms of shocks are ensued by the burst, the forward shock and the reverse one, which one separated by a contact discontinuity \\citep{GRBprimer}. If the ejected plasma is too strongly magnetized, only the forward shock is formed. One suggested possibility is that the prompt emission is generated in a baryon dominated ejecta through internal shocks, while the forward and reverse shocks yield the long lasting broadband emission, the afterglow \\citep{GRBprimer}. Actually, the full understanding of the prompt emission mechanism, a basic GRB property, is still lacking. One possibility is that the prompt emission consists of synchrotron radiation \\citep{GRBprimer}, by the relativistic charged particles moving on the magnetized ejected plasma. Currently, the GLAST/FERMI mission, in operation, is continuously increasing the available GRB data and making it worth, as will be discussed and pursued in this work, considering future prospects for the subject. For an overview of most recent missions see e.g. \\citep{mcbreen} and references therein.  GRBs can be used as distance indicators \\citep{ghirlanda2004,bertolamisilva2006,liang&zhang,amati2008}. Its main attractiveness is that the redshift range extends much higher than that of SNe Ia. The main observables that can be measured when studying GRBs are its spherical equivalent energy, its peak isotropic luminosity, the peak energy of its spectrum, the photon fluence, the energy fluence, the pulse duration and the redshift of its host galaxy. Several empirical correlations among these variables can be established. However, there are still large uncertainties in their calibration. Furthermore, there is still no satisfying physical mechanism accounting for them, so that assuming that they hold true can introduce systematic uncertainties in our distance indicator. % From the existing correlations, the very discussed Ghirlanda relation uses the peak energy of the spectrum, \\ep and the collimation corrected energy, $E_{\\gamma}$ \\citep{ghirlanda2004}. % On the other hand, the Liang and Zang relation correlates the isotropic equivalent energy, \\eiso~with \\ep~and the jet break time of the afterglow of the burst \\citep{liang&zhang}. Finally, the Amati relation, correlates the isotropic energy, \\eiso, with \\ep~\\citep{amati2008}. This relation is particularly interesting since the \\ep~- \\eiso~correlation requires only two parameters that can be directly inferred from the obervations. This correlation further emphasizes the relevance of the GRB data. Notice that the aforementioned synchrotron process reproduces the Amati correlation, a quite interesting feature. %  In this work, GRB data and the Amati relation, in particular, are used to probe a generic dark energy - dark matter interacting model. In section \\ref{DEDMint}, the interacting model is presented. The Amati \\ep~- \\eiso~correlation is introduced and discussed in section \\ref{dataanalysis}. The set of real GRB data is then extended to a mock sample of 500 GRBs using a method detailed in subsection \\ref{genmock}. In section \\ref{method}, one discusses the constrains obtained from SNe data in subsection \\ref{SNs}, BAO in subsection \\ref{BAOs} and CMBR shift parameter in subsection \\ref{CMBs}. In section \\ref{results} one presents the obtained results. In section \\ref{conclusion}, conclusions are presented. ",
        "conclusions": "In this work the use of GRBs as cosmological tools is considered. It has been shown \\citep{bertolamisilva2006,amati2008} that GRBs have a great potential to measure the value of $\\Omega_{DM}$ independently from the CMBR constraints. In the current work, it is shown that the present GRB data already gives a better constraint on the \\dedm\\ coupling parameter $\\zeta_0$ than the BAO results. The SNe results provide good bounds on $w$, but are more degenerate in $\\zeta_0$. Despite of that, the experimental GRB sample is still too small to provide significant constraints on $\\zeta_0$. The combined result for SNe and real GRB data further illustrates this, yielding $\\zeta_0\\in[-1.15, 1.66]$, with a width of $\\Delta \\zeta_0=2.81$, as opposed to the combined SNe and CMBR data limit $\\zeta_0\\in[-0.01, 0.13]$ ($\\Delta\\zeta_0=0.14$), both at 68\\% confidence level. These results are obtained without marginalization, for fixed $\\Omega_{DM_0}=0.3$.  A mock population of GRBs was then generated showing that as the number of available events increases, GRB data becomes a more and more valuable tool in constraining the parameter $\\zeta_0$. The GRBs complement the SNe constraint in a similar way the CMBR does. For a fixed $\\Omega_{M_0}=0.3$, the SNe and mock GRB data yields $\\Delta\\zeta_0=0.82$.  For a deeper insight on the future cosmological implications of the GRB observations, the analysis was extended to a marginalization over $\\Omega_{M_0}\\in[0.2, 0.4]$. The combined SNe and CMBR results are $\\Delta\\zeta_0=0.47$ (68\\% CL), while for SNe and the mock GRB yields $\\Delta\\zeta_0=1.03$. Granting that the obtained bounds are not as accurate as the CMBR ones, they still provide a valuable independent measurement of these cosmological parameters.  With the same marginalization in $\\Omega_{M_0}$, the combined result for SNe, CMBR and BAO is $\\zeta_0\\in[-0.27, 0.13]$ (95\\% CL). The updated observations improve the previous result $\\zeta_0\\in[-0.4, 0.1]$ \\citep{Guo}. Note, however, that tighter priors are used in this work and that this model is slightly different, with the inclusion of non-interacting baryonic matter.  It is interesting to compare the present results with the ones arising from estimates of the departure from the virial equilibrium of the Abell cluster A586. The bounds for $\\eta$ from the Abell Cluster A586 yield $\\eta\\in[3.65, 4.00]$ ($\\Delta\\eta=0.35$), with $w=-1$ and $z=0.1708$ and $\\Omega_{M_0}=0.28$, \\citep{intmodel1}. Using SNe and mock GRB data from the present work and fixing $\\Omega_{M_0}=0.28$, one encounters $\\eta\\in[0.93, 2.48]$ ($\\Delta\\eta=1.55$) at 68\\% confidence level. If one considers the particular case of the GCG, the Abell Cluster A586 limits the $\\alpha$ parameter to $\\alpha\\in[0.21, 0.33]$ (68\\% CL). This compares with the combined SNe and mock GRB results, $\\alpha\\in[0.25, 0.83]$ (68\\% CL). One should bear in mind that the GRB results were obtained from a mock population and are, thus, only indicative of what can be expected when the number of observed GRBs increases. The limits for $\\alpha$ with the current observational SNe and CMBR data yield $\\alpha\\in[0.03, 0.06]$.  In this work, the Amati correlation \\citep{amati2008} has been used, given that it requires only two parameters whose determination can be inferred and increasing the number of useful GRB events available. Progress in the calibration and on the theoretical framework of this calibration would be invaluable, reducing the error margins in the GRB data and considerably improving the constraints on the parameters."
    },
    "1004/1004.1171_arXiv.txt": {
        "abstract": "{} {In this paper we show expectations on the radio--X-ray luminosity correlation of radio halos at 120 MHz. According to the ``turbulent re-acceleration scenario'', low frequency observations are expected to detect a new population of radio halos that, due to their ultra-steep spectra, are missed by present observations at $\\sim$ GHz frequencies. These radio halos should also be less luminous than presently observed halos hosted in clusters with the same X-ray luminosity.} {Making use of Monte Carlo procedures, we show that the presence of these ultra-steep spectrum halos at 120 MHz causes a steepening and a broadening of the correlation between the synchrotron power and the cluster X-ray luminosity with respect to that observed at 1.4 GHz.} {We investigate the role of future low frequency radio surveys, and find that the upcoming LOFAR surveys will be able to test these expectations.} {} ",
        "introduction": "Radio halos are diffuse synchrotron sources from the intra--cluster medium (ICM) extended on mega-parsec scale (\\eg Feretti 2005; Ferrari et al. 2008). They provide the most important evidence of non-thermal components (relativistic particles and magnetic fields) mixed with the hot ICM. Galaxy clusters hosting radio halos are always characterized by a non-relaxed dynamical status suggestive of recent or ongoing merger events (\\eg Buote 2001; Schuecker et al 2001; Govoni et al. 2004; Venturi et al. 2008; Giacintucci et al. 2009). Furhermore, the halo radio power at 1.4 GHz increases with the cluster X-ray luminosity, mass and temperature (\\eg Liang et al. 2000; En\\ss lin \\& R\\\"ottgering 2002; Bacchi et al. 2003; Clarke 2005; Dolag et al. 2005; Cassano et al. 2006, 2007; Brunetti et al. 2009; Rudnick \\& Lemmerman 2009; Giovannini et al. 2009). These correlations and the radio halo-merger connection suggest that gravity provides the reservoir of energy to generate the non-thermal components (\\eg Kempner \\& Sarazin 2001). Cluster mergers drive shocks and turbulence in the ICM that may amplify the magnetic fields (\\eg Carilli \\& Taylor 2002; Dolag et al. 2002; Br\\\"uggen et al. 2005; Subramanian et al. 2006; Ryu et al. 2008) and accelerate high energy particles (\\eg Fujita et al. 2003; Hoeft \\& Br\\\"uggen 2007; Brunetti \\& Lazarian 2007; Pfrommer et al. 2008; Vazza et al. 2009).  Two main scenarios have been proposed to explain the origin of relativistic particles in radio halos, namely {\\it i)} the {\\it turbulent re-acceleration} model, whereby relativistic electrons are re-energized {\\it in situ} due to the interaction with MHD turbulence generated in the ICM during cluster mergers (\\eg Brunetti et al. 2001; Petrosian et al. 2001), and {\\it ii)} the {\\it secondary electron} models, whereby the relativistic electrons are secondary products of the collisions between cosmic rays and thermal protons in the ICM (\\eg Dennison 1980; Blasi \\& Colafrancesco 1999, Pfrommer \\& En\\ss lin 2004).  Observations provide support to the idea that turbulence may play a role in the particle re-acceleration process (\\eg Brunetti et al. 2008; Ferrari et al. 2008; Cassano 2009; Giovannini et al. 2009), in which case, the population of radio halos is predicted to be a mixture of sources with different spectral properties, with halos having steeper spectra being more common (Cassano et al. 2006a, hereafter C06; Cassano et al. 2009, hereafter C09). In this respect, since very steep spectrum halos should glow up at low radio frequency, upcoming observations with the Low Frequency Array (LOFAR) and the Long Wavelength Array (LWA) will be crucial.  In this paper we discuss how such a predicted population is expected to affect the properties of the radio--X-ray luminosity correlation at low radio frequency, and investigate the potential of LOFAR surveys. A $\\Lambda$CDM cosmology ($H_{o}=70\\,\\rm km\\,\\rm s^{-1}\\,\\rm Mpc^{-1}$, $\\Omega_{m}=0.3$, $\\Omega_{\\Lambda}=0.7$) is adopted. ",
        "conclusions": "The observed correlations between the halo radio power (at 1.4 GHz) and the cluster X-ray luminosity, mass and temperature, and the observed connection between radio halos and cluster mergers, suggest a link between the gravitational process of cluster formation and the generation of radio halos. Radio halos are likely generated during cluster-cluster mergers where a fraction of the gravitational energy dissipated is channelled into the acceleration of relativistic particles. A crucial expectation of the {\\it turbulent re-acceleration} scenario, put forward to explain radio halos, is that the synchrotron spectrum of halos is characterized by a cut-off at frequency $\\nu>\\nu_s$ with $\\nu_s$ determined by the efficiency of the acceleration process. The presence of this cut-off causes a bias, so that present radio observations at $\\sim$ GHz frequencies are expected to detect only the most efficient radio phenomena in clusters, leaving unexplored a large population of radio halos characterized by spectral cut-off at lower frequencies (\\eg C06, Brunetti et al. 2008; C09). Future low frequency radiotelescopes as LOFAR and LWA are expected to unveil the populations of ultra steep spectrum radio halos, with $\\nu_s<1$ GHz, in clusters, providing to test the idea of {\\it turbulent re-acceleration}.  One may wonder whether halos with $\\nu_s<1$ GHz could be detected by present radiotelescopes at 1.4 GHz. To address this point in Fig.~\\ref{Fig.NH_z0p6_1p4GHz} we report the flux distribution at 1.4 GHz of halos with $600\\leq\\nu_s<1400$ MHz (that in our model have $\\alpha\\approx 1.7$ between 120 MHz and 1400 MHz) that are expected to be detected by LOFAR at 120 MHz assuming $\\xi\\,F\\simeq 0.25$ mJy/beam. Calculations are derived by assuming: {\\it i)} at $z<0.3$ the X-ray flux limit and sky coverage of the extended {\\it ROSAT} Brightest Cluster Sample (eBCS, Ebeling et al. 1998, 2000) and of the ROSAT-ESO Flux Limited X-ray Galaxy Cluster Survey (REFLEX, B\\\"oringher et al. 2004), and {\\it ii)} at $z=0.3-0.6$ the X-ray flux limit and sky coverage of the Massive Cluster Survey (MACS, Ebeling et al. 2001). Calculations show that potentially very deep pointed observations at 1.4 GHz of all these clusters may lead to the detection of a few of these ultra steep spectrum halos (those with $600\\leq\\nu_s<1400$ MHz); the ultra steep spectrum halo detected in the cluster Abell 521 (Brunetti et al. 2008; Dallacasa et al. 2009) belong to this class of halos, although it is among the flatter spectrum objects in this class, with $\\nu_s\\approx 1200$ MHz.    In this paper, we discuss the consequence of this new population of radio halos on the slope of the radio--X-ray luminosity correlation at low frequency. According to homogeneous models, ultra-steep spectrum halos are expected to be less luminous than halos with larger $\\nu_s$ associated with clusters of the same mass. Also, radio halos with smaller $\\nu_s$ should be statistically generated in clusters with smaller mass (and $L_X$). The combination of these two expectations implies that the radio-- X-ray luminosity correlation should be broader and steeper at lower frequencies.   \\begin{figure} \\centerline{ \\includegraphics[width=0.36\\textwidth]{NH_z0p6_0p678_1p4GHz_new.ps}} \\caption[]{Integrated number counts at 1.4 GHz of halos with $600\\leq\\nu_s<1400$ MHz that are expected to be detected by LOFAR at 120 MHz assuming $\\xi\\,F\\simeq0.25$ mJy/beam. Calculations are derived by combining {\\it i)} at $z<0.3$, the sky coverage and X-ray flux limit of eBCS and REFLEX clusters, and {\\it ii)} at $z=0.3-0.6$, the sky coverage and X-ray flux limit of MACS clusters. Dashed and solid lines account for the uncertainty due to the finite frequency range of $\\nu_s$ ($600-1400$ MHz) assumed in the calculations. The dashed blue region shows the typical minimum flux of halos detectable with present radio facilities (\\eg the VLA considering C+D configurations) at 1.4 GHz (approximatively $f(1.4)>3-4$ mJy).} \\label{Fig.NH_z0p6_1p4GHz} \\end{figure}  Based on this model, we perform Monte Carlo simulations of the distribution of radio halos in the $P(120)-L_X$ plane. We find that halos distribute in the $P(120)-L_X$ plane according to a correlation which is steeper ($\\Delta \\alpha\\approx0.4$) and broader than that observed at 1.4 GHz, with ultra-steep spectrum halos broadening the scatter in the region of low luminosity. We find that the number of ultra-steep spectrum halos increases with increasing the survey sensitivity and this further steepens the correlation. The forthcoming LOFAR surveys should constrain the expected steepening of the correlation and test our expectations.  Although a self-consistent treatment of turbulence acceleration and amplification of the magnetic field in clusters is mandatory and deserve future efforts, we show that the main ingredients in the adopted scenario are two ``fast'' processes: particle acceleration and particle cooling that follow the decay of turbulence. Being a ``slow'' process, we show that the possible decay of the field with turbulence is not expected to affect the modeling of halo statistics significantly."
    },
    "1004/1004.2458_arXiv.txt": {
        "abstract": "We present results of direct imaging observations for HAT-P-7 taken with the Subaru HiCIAO and the Calar Alto AstraLux. Since the close-in transiting planet HAT-P-7b was reported to have a highly tilted orbit, massive bodies such as giant planets, brown dwarfs, or a binary star are expected to exist in the outer region of this system. We show that there are indeed two candidates for distant faint stellar companions around HAT-P-7. We discuss possible roles played by such companions on the orbital evolution of HAT-P-7b. We conclude that as there is a third body in the system as reported by Winn et al. (2009, ApJL, 763, L99), the Kozai migration is less likely while planet-planet scattering is possible. ",
        "introduction": "The discovery of over 400 extrasolar planets and the diversity of their orbital distributions dramatically changed our perception of planetary systems in the last 15 years. Especially, the existence of exoplanets in very close-in or eccentric orbits stimulated theorists to develop various models for planetary migration during the epoch of planet formation. To explain the orbital distribution of known exoplanets, a number of planetary migration models have been proposed, including disk-planet interaction models (i.e., Type I and Type II migration models; e.g., \\cite{1985prpl.conf..981L, 1996Natur.380..606L, 2002A&A...385..647D, 2004ApJ...616..567I}), planet-planet scattering models considering gravitational interaction among multiple giant planets (i.e., jumping Jupiter models; e.g., \\cite{1996Sci...274..954R, 2002Icar..156..570M, 2008ApJ...678..498N, 2008ApJ...686..580C}),  or Kozai migration models considering perturbation by a distant massive companion and coinstantaneous tidal evolution (e.g., \\cite{2003ApJ...589..605W, 2005ApJ...627.1001T, 2007ApJ...669.1298F, 2007ApJ...670..820W}).   These planetary migration models can now be tested by observations of the Rossiter-McLaughlin effect (hereafter the RM effect: \\cite{1924ApJ....60...15R, 1924ApJ....60...22M}) for transiting planetary systems, which is an anomalous shift in observed radial velocities due to the occultation of a rotating star. Measurements of the RM effect enable us to estimate the sky projection angle of the planetary orbital axis relative to the stellar spin axis (i.e., the spin-orbit alignment angle; \\cite{2005ApJ...622.1118O, 2007ApJ...655..550G, 2010ApJ...709..458H}). The information of the spin-orbit alignment angle is very useful to differentiate the planet-planet scattering models and the Kozai migration models from the disk-planet interaction models, because the former models predict a wider range of spin-orbit alignment angles for migrated planets, while the latter models predict that migrated planets would have only small spin-orbit alignment angles. Thus an observation of a highly tilted orbit for a specific planet is strong evidence of a planet-planet scattering process or perturbation by an outer companion during its migration history.   Indeed, very recently several transiting exoplanets have been reported to show such highly tilted orbits via measurements of the RM effect. These observations lead to an interesting prediction: around these spin-orbit misaligned exoplanets, other massive bodies (e.g., massive planets, brown dwarfs, or a low-mass stellar companion) should be present. In addition, since migration mechanisms for such exoplanets cannot be distinguished by the spin-orbit alignment angles alone, direct imaging of such massive bodies gives us important additional information to distinguish between the two migration mechanisms of highly tilted orbit planets. Motivated by these facts, we have initiated to search for such outer massive bodies around known transiting planetary systems with the Subaru HiCIAO (High Contrast Instrument for the Subaru next generation Adaptive Optics; \\cite{2006SPIE.6269E..28T, 2008SPIE.7014E..42H}; Suzuki et al. 2010, in prep.), as part of the SEEDS project (Strategic Explorations of Exoplanets and Disks with Subaru, PI: Motohide Tamura). The Subaru HiCIAO is a powerful instrument to search for outer faint bodies around stars, proven by the detection of a massive planet or a brown dwarf around GJ~758 \\citep{2009ApJ...707L.123T}.   In this paper, we targeted the transiting planetary system HAT-P-7, which was reported to have a planet (HAT-P-7b) on an orbit highly inclined relative to the stellar equatorial plane (\\cite{2009PASJ...61L..35N}; hereafter NSH09, \\cite{2009ApJ...703L..99W}; hereafter WJA09). Consequently, we report two candidates of faint stellar companions to HAT-P-7 based on the Subaru HiCIAO and the Calar Alto AstraLux data. Although our data alone cannot distinguish whether or not the candidate companion stars are physically associated with HAT-P-7 at this point, the findings are useful to constrain the mechanism of planetary migration in this system. We summarize the properties of our target in section~2, and report our observations, analyses, and results in section~3. We present theoretical discussions of the migration mechanism of HAT-P-7b in section~4. Finally, section~5 summarizes the findings in this paper.   \\begin{table}[t] \\caption{Summary of stellar and planetary parameters.} \\begin{center} \\begin{tabular}{l|ccc} \\hline Parameter & Value & Error & Source \\\\ \\hline Star & & & \\\\ $M_s$ [$M_{\\odot}$] & $1.520$ & 0.036 & CKB10 \\\\ $R_s$ [$R_{\\odot}$] & $1.991$ & 0.018 & CKB10 \\\\ Age [Gyr] & $2.14$ & 0.26 & CKB10 \\\\ Distance [pc] & $320$ & $^{+50}_{-40}$ & PBT08 \\\\ app. $H$ mag. & 9.344 & 0.029 & \\citet{2003tmc..book.....C} \\\\ \\hline Planet & & & \\\\ $M_p$ [$M_{Jup}$] & $1.82$ & 0.03 & WOS10 \\\\ $R_p$ [$R_{Jup}$] & $1.50$ & 0.02 & WOS10 \\\\ $i$ [$^{\\circ}$] & $83.1$ & 0.5 & WOS10 \\\\ $a$ [AU] & 0.0386 & 0.0001 & WOS10 \\\\ $P$ [days] & $2.204733$ & 0.000010 & WOS10 \\\\ \\hline \\multicolumn{4}{l}{\\hbox to 0pt{\\parbox{80mm}{\\footnotesize }\\hss}} \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "In this section, we discuss realizable migration mechnisms of the highly tilted orbit planet HAT-P-7b. First of all, as a simple case, if neither of the candidate companions we discovered is physically associated with HAT-P-7, only planet-planet scattering can explain the tilted orbit of HAT-P-7b. We thus examine a Kozai migration scenario for HAT-P-7b first, based on an assumption that either of the candidate companion stars is a true binary of HAT-P-7. We present conditions required for the Kozai migration of HAT-P-7b with a binary companion in section 4.1. We show a restricted area of a third body for the Kozai migration in this system in section 4.2., and describe the impact of the possible third body (HAT-P-7c) reported by WJA09 in section 4.3.   \\subsection{Required Condition for the Kozai Migration of HAT-P-7b}   According to \\citet{1962AJ.....67..591K}, the angular momentum \\begin{equation} L_Z \\equiv \\sqrt{G (M_p + M_s) a (1 - e^2)} \\cos \\Psi_m \\end{equation} should be conserved in a planetary system with a binary star during the Kozai mechanism, where $G$ is the gravitational constant, $M_p$ and $M_s$ are the mass of the planet and its host star, $a$ and $e$ are the semi-major axis and the orbital eccentricity of the planet, and finally $\\Psi_m$ (domain: [0$^{\\circ}$, 180$^{\\circ}$]) is the mutual inclination between the orbital inclination of the planet and the binary star. In addition, given that the angular momentum is also conserved during tidal orbital evolution, $a (1-e^2) \\cos^2 \\Psi_m$ should be conserved through the Kozai migration.   Using the conservation relation above, we constrain the initial mutual inclination to initiate the Kozai migration in this system as follows. First we assume that HAT-P-7b was born in the snow line with the initial eccentricity $e_0$ and the initial mutual inclination $\\Psi_{m,0}$. The distance of the snow line from the host star is roughly estimated as $a_0 = 2.7 (M_s/M_\\odot)^2 = 6.24$ AU (i.e. 20 mas). We note that although the position of the snow line is somewhat uncertain, this makes little impact on the following discussions. Then using the conservation relation rewritten as \\begin{equation} a_0 (1-e_0^2) \\cos^2 \\Psi_{m,0} = a_n (1-e_n^2) \\cos^2 \\Psi_{m,n}, \\end{equation} where the indices $0$ and $n$ indicate values of the initial state and those as of now, we obtain \\begin{equation} |\\cos \\Psi_{m,0}| \\le \\sqrt{\\frac{0.0386}{6.24} \\frac{1}{1-e_0^2}} \\,\\, |\\cos \\Psi_{m,n}|, \\end{equation} as a necessary initial condition for the orbit of HAT-P-7b. If this condition is satisfied, the eccentricity would be excited over the critical value $\\sqrt{1-0.0386/6.24}=0.997$ and the planet would initiate tidal evolution.   We note that the timescale of the Kozai migration under consideration is approximated as \\citep{2007ApJ...670..820W}, \\begin{equation} P_{\\tiny{\\textrm{K}}} \\sim \\frac{M_s}{M_B} \\frac{P_B^2}{P_{0}} (1-e_B^2)^{3/2}, \\end{equation} where $M_B$, $P_B$, and $e_B$ are the mass, orbital period, and eccentricity of the binary star, and $P_{0}$ is the orbital period of HAT-P-7b at the initial position. Assuming that $M_B = 0.20 M_{\\odot}$ (as a typical mass of M star) and $e_B = 0$, we obtain $P_{\\tiny{\\textrm{Kozai}}} \\sim$ 300 Myr, which is sufficiently short relative to the age of this system ($\\sim2$ Gyr). In addition, the timescale of general relativity is estimated as \\citep{2003ApJ...589..605W}, \\begin{equation} P_{\\tiny{\\textrm{GR}}} \\sim \\frac{2 \\pi c^2 (1-e_0^2) a_0^{5/2}}{3 (GM_s)^{3/2}} \\sim 2 \\textrm{Gyr}, \\end{equation} where $c$ is the speed of light. Thus general relativity would not disturb the Kozai migration of HAT-P-7b at an early stage.   Based on equation (3), if HAT-P-7b was born with $e_0 = 0$ and if $\\Psi_{m,n}=0^{\\circ}$, namely if HAT-P-7b and the binary star are coplanar now, the initial mutual inclination $\\Psi_{m,0}$ needs to be within $85.5^{\\circ}-94.5^{\\circ}$ to initiate the Kozai migration. Even if we assume that the initial eccentricity is large (e.g., $e_0=0.8$) and the current mutual inclination is zero, $\\Psi_{m,0}$ needs to be within $82.5^{\\circ}-97.5^{\\circ}$. This is a very optimistic case; if the eccentricity is lower and $\\Psi_{m,n}$ is not zero, the required condition becomes more stringent. These required conditions are very tight, but still possible (a few suggestive circumstellar disk observations of nonzero $\\Psi_{m,0}$ for young binary stars were reported; e.g., \\cite{1998ApJ...505..358A,2005ApJ...628..832D,2009PASJ...61.1271H}).   \\subsection{Restricted Area of a Third Body for the Kozai Migration Scenario}   As discussed by \\citet{2003ApJ...589..605W} for HD~80606b, a hypothetical additional body HAT-P-7c in the HAT-P-7 system could destroy the Kozai migration process \\citep{1997AJ....113.1915I}, if the timescale of orbital precession of HAT-P-7b caused by the gravitational perturbation from HAT-P-7c ($P_{{\\rm{G}},c}$) is shorter than that caused by the Kozai mechanism due to the binary companion ($P_{{\\rm{K}},B}$). We calculated a conditional equation for a restricted area of an outer third body at the initial stage as, \\begin{equation} M_c > \\frac{3}{2} M_s \\frac{a_c^2 a}{a_B^3} \\frac{1}{b_{3/2}^{(1)}}, \\end{equation} where $a_c$ and $M_c$ are the semi-major axis and mass of the additional planet, $a_B$ is the semi-major axis of the binary star, and $b_{3/2}^{(1)}$ is the Laplace coefficient (see e.g., \\cite{2000ssd..book.....M, 2003ApJ...589..605W}).   The boundary of the restricted area is plotted in figure~3 by solid (for $a_B=1000$ AU) and dotted (for $a_B=2000$ AU) lines. The horizontal axis and the vertical axis represent $a_c$ and $M_c$, respectively. More specifically, the upper region of the solid (dotted) line is the restricted area where HAT-P-7c cannot exist initially for the Kozai migration caused by the binary companion. This constraint is very stringent, and even analogies of Saturn ($a_c = 9.6$ AU, $M_c = 0.3 M_{Jup}$) or Uranus ($a_c = 19.2$ AU, $M_c = 0.04 M_{Jup}$) cannot exist.   \\subsection{Impact of HAT-P-7c Reported by Winn et al. (2009)}   On the other hand, WJA09 reported that there is indeed a possible third body HAT-P-7c in the HAT-P-7 system (hereafter, just ``c''). As a constraint on the mass and semi-major axis of the additional body, WJA09 reported the following equation; \\begin{equation} \\frac{M_c \\sin i_c}{a_c}^2 \\sim (0.121\\pm0.014)\\,\\, M_{Jup}\\,\\, \\textrm{AU}^{-2}, \\end{equation} where $i_c$ is the orbital inclination of ``c'' relative to the line of sight. We plotted equation (7) (assuming $\\sin i_c = 0$ for simplicity) in figure~3 using a dashed-dotted line for reference. Obviously, the Kozai migration scenario is totally excluded if ``c'' existed in the outer region (beyond the snow line) at the initial stage. Thus in the presence of ``c'', it is impossible to explain the tilted orbit of HAT-P-7b by the Kozai migration caused by the distant binary companion only.   However, there is another chance of a ``sequential'' Kozai migration scenario for a 2-body system with a binary star, as introduced by \\citet{2008ApJ...683.1063T} and \\citet{2010arXiv1003.0633K}. Namely, an inclined binary companion induces the Kozai mechanism for an outer body first, and then the inclined outer body leads the Kozai mechanism for an inner body. In this case, ``c'' could play an important role for the migration of HAT-P-7b. If this is the case, ``c'' could have a tilted orbital axis relative to both the stellar spin axis and the orbital axis of HAT-P-7b. We cannot discuss this possibility in detail at this point, because the orbital parameters of ``c'' have not yet been determined. If orbital parameters of ``c'' are firmly determined, it would be interesting to discuss the possibility of a sequential Kozai migration scenario. We also note that if the semi-major axis of ``c'' turns out to be large, further direct imaging of this inner body would be also interesting in the future (e.g., with TMT or E-ELT).   From the above discussions, we found that the Kozai migration scenario caused by a distant binary star could be realized only in a very limited situation. In addition, if ``c'' existed, the Kozai migration of HAT-P-7b caused directly by the binary star could not have occured, although we could not refute the possibility of a sequential Kozai migration for HAT-P-7b at this point. In addition, if neither of the candidate stars is a physical companion of HAT-P-7, only planet-planet scattering can explain the tilted orbit of HAT-P-7b. Thus with a few exceptions above, we conclude that planet-planet scattering is a more plausible explanation for the migration mechanism of HAT-P-7b."
    },
    "1004/1004.2491_arXiv.txt": {
        "abstract": "We use a series of Monte Carlo simulations to investigate the theory of galaxy-galaxy lensing by non-spherical dark matter haloes.  The simulations include a careful accounting of the effects of multiple deflections on the galaxy-galaxy lensing signal. In a typical observational data set where the mean tangential shear of sources with redshifts $z_s \\simeq 0.6$ is measured with respect to the observed symmetry axes of foreground galaxies with redshifts $z_l \\simeq 0.3$, we find that the signature of anisotropic galaxy-galaxy lensing differs substantially from the simple expectation that one would have in the absence of multiple deflections. In general, the observed ratio of the mean tangential shears, $\\gamma^+ (\\theta) / \\gamma^- (\\theta)$, is strongly suppressed compared to the function that one would measure if the intrinsic symmetry axes of the foreground galaxies were known.  Depending upon the characteristic masses of the lenses, the observed ratio of the mean tangential shears may be consistent with an isotropic signal (despite the fact that the lenses are non-spherical), or it may even be reversed from the expected signal (i.e., the mean tangential shear for sources close to the observed minor axes of the lenses may exceed the mean tangential shear for sources close to the observed major axes of the lenses).  These effects are caused primarily by the fact that the images of the lens galaxies have, themselves, been lensed and therefore the observed symmetry axes of the lens galaxies differ from their intrinsic symmetry axes. We show that the effects of lensing of the foreground galaxies on the observed function $\\gamma^+ (\\theta) / \\gamma^- (\\theta)$ cannot be eliminated simply by the rejection of foreground galaxies with very small image ellipticities, nor by simply focusing the analysis on sources that are located very close to the observed symmetry axes of the foreground galaxies. We conclude that any attempt to use a measurement of $\\gamma^+ (\\theta) / \\gamma^- (\\theta)$ to constrain the shapes of dark matter galaxy haloes must include Monte Carlo simulations that take multiple deflections properly into account. ",
        "introduction": "Galaxy-galaxy lensing is a form of weak gravitational lensing in which background galaxies are systematically lensed by foreground galaxies. Brainerd, Blandford \\& Smail (1996; BBS) published the first statistically-significant ($4\\sigma$) detection of this effect using a small data set that consisted of 439 foreground galaxies, 506 background galaxies, and 3202 foreground-background galaxy pairs.  Since this early work, galaxy-galaxy lensing has been detected with high precision using various data sets, most of which contain millions of foreground-background galaxy pairs.  These high-precision detections have allowed direct constraints to be placed on the nature of the dark matter haloes that surround the lens galaxies, as well as on the bias between mass and light in the universe (see, e.g., Fischer et al.\\ 2000; Guzik \\& Seljak 2002; Hoekstra, Yee \\& Gladders 2004; Hoekstra et al.\\ 2005; Sheldon et al.\\ 2004; Heymans et al.\\ 2006; Kleinheinrich et al.\\ 2006; Mandelbaum et al.\\ 2006ab; Mandelbaum, Seljak \\& Hirata 2008; Limousin et al.\\ 2007; Parker et al.\\ 2007; Natarajan et al.\\ 2009; Tian et al.\\ 2009).  Observations of galaxy-galaxy lensing by field galaxies have shown: [1] at fixed luminosity, the haloes of red (early-type) galaxies are more massive by a factor of $\\sim 2$ than the haloes of blue (late-type) galaxies (e.g., Guzik \\& Seljak 2002; Kleinheinrich et al.\\ 2006; Sheldon et al.\\ 2004; Mandelbaum et al.\\ 2006a), [2] the haloes of high-luminosity galaxies are more massive than the haloes of low-luminosity galaxies (e.g., Sheldon et al.\\ 2004; Mandelbaum et al.\\ 2006a), and [3] the dark matter profiles of the haloes are consistent with the spherically-averaged Navarro, Frenk \\& White (NFW) profile (Navarro, Frenk \\& White 1995, 1996, 1997; e.g., Heymans et al.\\ 2006; Hoekstra et al.\\ 2004, 2005; Kleinheinrich et al.\\ 2006; Mandelbaum et al.\\ 2008). In other words, observations of galaxy-galaxy lensing by field galaxies have yielded a picture of luminous galaxies and their dark matter haloes that is broadly consistent with the expectations of galaxy formation in the context of the cold dark matter (CDM) model.  Despite the popularity of the spherically-averaged NFW density profile, CDM haloes are not spherical.  Rather, CDM haloes are triaxial and the degree of flattening increases with halo viral mass (e.g.,  Warren et al.\\ 1992; Jing \\& Suto 2002; Bailin \\& Steinmetz 2005; Kasun \\& Evrard 2005; Allgood et al.\\ 2006). In principle, galaxy-galaxy lensing should be able to provide constraints on the shapes of the dark matter haloes of field galaxies, since a non-spherical weak lens will produce an anisotropic shear pattern.  Consider an isolated weak galaxy lens with a non-spherical dark matter halo (i.e., a halo that, in projection on the sky, has an elliptical surface mass density). For fixed source redshift and fixed angular distance from the lens, sources that are located closer to the major axis of the lens will experience greater shear than sources that are located closer to the minor axis of the lens.  If the halo of the lens can be approximated as a singular isothermal ellipsoid with projected ellipticity $\\epsilon_{\\rm halo} = 0.3$, the shear experienced by sources nearest the minor axis of the lens will be $\\sim 80$\\% that of the shear experienced by sources nearest the major axis of the lens (see, e.g., Brainerd \\& Blandford 2002).  Although small, such an anisotropy in the galaxy-galaxy lensing signal should be observable provided that, in projection on the sky, mass and light are reasonably well aligned within the lens galaxies.  Weak lensing by galaxy clusters in the Sloan Digital Sky Survey (SDSS; e.g., Abazajian et al.\\ 2009 and references therein) has shown that the dark mass associated with galaxy clusters is non-spherical and has a projected axis ratio of $b/a = 0.48^{+0.14}_{-0.09}$ (Evans \\& Bridle 2009). The detection of non-spherical haloes by galaxy-galaxy lensing has, however, proven to be more problematical.  In a study of galaxy-galaxy lensing by galaxies in the Red-Sequence Cluster Survey, Hoekstra et al.\\ (2004) modeled the projected shapes of the haloes as $\\epsilon_{\\rm halo} = \\lambda~ \\epsilon_{\\rm light}$, where $\\epsilon_{\\rm light}$ is the ellipticity of the image of the luminous galaxy within the halo.  Here $\\lambda=1$ indicates that the projected shapes of the haloes are identical to the shapes of the galaxy images, and $\\lambda=0$ indicates that the haloes are perfectly circular in projection on the sky. From their analysis, Hoekstra et al.\\ (2004) concluded that the haloes of their galaxies were somewhat rounder than the images of the galaxies: $\\lambda = 0.77^{+0.18}_{-0.21}$. Using the same parametrization of the relationship between the ellipticities of the haloes and the images of the galaxies, Mandelbaum et al.\\ (2006b) found $\\lambda = 0.1 \\pm 0.06$ for red SDSS lens galaxies and $\\lambda = -0.8 \\pm 0.4$ for blue SDSS lens galaxies.  Here the negative sign indicates an apparent anti-alignment of mass and light for blue SDSS lens galaxies.  Finally, Parker et al.\\ (2007) computed the galaxy-galaxy lensing signal using data from the Canada-France-Hawaii Telescope Legacy Survey.  When Parker et al.\\ (2007) averaged the signal over all lens galaxies, they found a weak ($2\\sigma$) preference for the haloes of the lens galaxies to be non-spherical with a projected ellipticity of $\\sim 0.3$. When Parker et al.\\ (2007) restricted their analysis to elliptical galaxies, the mean halo ellipticity and the significance of the result was found to increase somewhat.  Here we construct a series of Monte Carlo simulations in order to explore the theory of weak galaxy-galaxy lensing by non-spherical dark matter haloes. Using these simulations we demonstrate that, in practice, it is challenging to {\\it interpret} the results of an observational effort to detect anisotropic galaxy-galaxy lensing.  This is because, in general, the observed signature of anisotropic galaxy-galaxy lensing is strongly affected by the fact that the central, ``lens'' galaxies have, themselves been weakly lensed.  As a result, the observed symmetry axes of the central, lens galaxies differ from their intrinsic symmetry axes. In our work below we pay particular attention to the effects of multiple weak deflections on the galaxy-galaxy lensing signal. As was first pointed out by BBS, galaxy-galaxy lensing is inherently a multiple deflection problem.  That is, it is common for a source galaxy located at redshift $z_s$ to be weakly lensed by a galaxy located at $z_{l1} < z_s$.  Oftentimes these two galaxies are then subsequently lensed by another galaxy at redshift $z_{l2} < z_{l1}$.  In other words, the galaxy at $z_{l1}$ serves simultaneously as a lens for the galaxy at $z_s$ and a source for the galaxy at $z_{l2}$.  In addition, the galaxy at $z_s$ is lensed by two different foreground galaxies.  Neglecting such multiple deflections when modeling an observed galaxy-galaxy lensing signal will give rise to incorrect conclusions about the underlying properties of the haloes of the lens galaxies.  For a detailed discussion of the frequency and relative strengths of multiple deflections in a deep galaxy-galaxy lensing data set, the reader is referred to Brainerd (2010).  Below, the haloes of the lens galaxies will be modeled as truncated singular isothermal ellipsoids.  This choice is motivated by two considerations.  Firstly, the singular isothermal ellipsoid gives rise to a gravitational lensing shear that can be computed analytically (e.g., Kormann, Schneider \\& Bartelmann 1994). Secondly, at the present time the observational galaxy-galaxy lensing data are not of sufficiently high quality to allow one to distinguish between singular isothermal ellipsoid haloes and those that are triaxial CDM haloes.  The outline of the paper is as follows.  In Section 2 we present the basic theory of gravitational lensing by singular isothermal ellipsoids and we introduce a shorthand notation that we will use throughout the paper. In Section 3 we outline the construction of Monte Carlo simulations of galaxy-galaxy lensing by non-spherical haloes, where the locations and apparent magnitudes of the Monte Carlo galaxies are taken from a large observational data set. In Section 4 we present the signature of galaxy-galaxy lensing by non-spherical haloes that one should expect to obtain from a realistic observational data set.   In Section 5 we explore the effects of galaxy-galaxy lensing on the images of relatively nearby galaxies (i.e., galaxies that are ordinarily be considered to be ``lenses'' but are not always considered to be ``sources''). In Section 6 we construct a second suite of Monte Carlo simulations in order to determine the effect of multiple weak deflections on observations of anisotropic galaxy-galaxy lensing.  In Section 7 we demonstrate that the effects of lensing of foreground galaxies on the observed signature of anisotropic galaxy-galaxy lensing cannot be eliminated by selective rejection of either lens or source galaxies.  We summarize our results and present our conclusions in Section 8. Throughout, we will refer to the weak lensing of a background galaxy by a single foreground galaxy as a ``deflection'', and we will adopt a flat $\\Lambda$-dominated cosmology with $H_0 = 70$~km~sec$^{-1}$~Mpc$^{-1}$, $\\Omega_{m0} = 0.3$ and $\\Omega_{\\Lambda 0} = 0.7$. ",
        "conclusions": "We have investigated the theory of galaxy-galaxy lensing by non-spherical dark matter haloes, which should give rise to an anisotropy in the tangential shear experienced by distant source galaxies. If each distant source is lensed by only one foreground elliptical lens, and if the observed symmetry axes of the elliptical lens correspond to the intrinsic symmetry axes of its projected dark matter halo, one would expect the signature of anisotropic galaxy-galaxy lensing to manifest as $\\gamma^+ (\\theta) > \\gamma^- (\\theta)$ over a wide range of angular scales. Here $\\gamma^+ (\\theta)$ is the angular dependence of the mean tangential shear experienced by sources whose azimuthal coordinates place them close to the major axis of the lens, and $\\gamma^- (\\theta)$ is the angular dependence of the mean tangential shear experienced by sources whose azimuthal coordinates place them close to the minor axis of the lens.  Using an observational data set (observed coordinates and $I$-band apparent magnitudes) as a framework for a set of Monte Carlo simulations, we have demonstrated that the actual signature that one should expect to observe for anisotropic galaxy-galaxy lensing is far from the above idealised case. Because galaxies are broadly distributed in redshift space, it is common for a distant source galaxy located at redshift $z_s$ to be lensed by another galaxy located at redshift $z_{l1} < z_{s}$.  In turn, this original lens-source pair may then be lensed by yet another galaxy (or galaxies) located at redshift $z_{l2} < z_{l1}$.  Such instances of ``multiple deflections'' cause the observed signature of anisotropic galaxy-galaxy lensing to deviate from the expected signature.  The degree to which the observed signature of galaxy-galaxy lensing deviates from the expected signature is a strong function of the characteristic velocity dispersion of the haloes of $L^\\ast$ galaxies.  In the case of low characteristic velocity dispersions, $\\sigma_v^\\ast = 100$~km~sec$^{-1}$, the observed ratio of mean tangential shears, $\\gamma^+ (\\theta) / \\gamma^- (\\theta)$, exceeds a value of unity on all scales $\\theta < 100''$ and is only slightly lower than the function one would obtain if the intrinsic symmetry axes of the foreground galaxies were used to perform the calculation.  In the case of moderate velocity dispersions, $\\sigma_v^\\ast = 150$~km~sec$^{-1}$, the observed ratio of mean tangential shears shows little to no anisotropy on scales $\\theta > 20''$.  In the case of high velocity dispersions, $\\sigma_v^\\ast = 200$~km~sec$^{-1}$, the observed function is actually reversed from the expected function (i.e., $\\gamma^+ (\\theta) < \\gamma^- (\\theta)$) on scales $20'' < \\theta < 70''$, and is consistent with no anisotropy on scales $70'' < \\theta < 120''$.  In summary, our simulations show that if one observes $\\gamma^+ (\\theta) = \\gamma^- (\\theta)$ in a large galaxy-galaxy lensing data set, the observation cannot be simply interpreted as proof that the haloes of the lens galaxies are spherically-symmetric.  That is, although the measured signal appears to be isotropic, it is entirely possible that anisotropic galaxy-galaxy lensing by non-spherical haloes may have taken place. Further, our simulations show that if one observes $\\gamma^+ (\\theta) < \\gamma^- (\\theta)$ in a large galaxy-galaxy lensing data set, the observation cannot be simply interpreted as proof that mass and light are ``anti-aligned'' in the lens galaxies.  That is, although the measured signal appears to be reversed from the expected signal, the reversal may occur when mass and light are, in fact, perfectly aligned within the lens galaxies.  The primary reason that the observed signature of anisotropic galaxy-galaxy lensing differs from the expected signature is that the foreground galaxies that are used as centres to compute the mean tangential shear have, themselves, been weakly lensed.  The expectation that $\\gamma^+ (\\theta)$ will exceed $\\gamma^- (\\theta)$ over a wide range of angular scales is based upon a picture in which the observed symmetry axes of the lenses are identical to the intrinsic symmetry axes of their projected dark matter haloes.  However, when one computes $\\gamma^+ (\\theta)$ and $\\gamma^- (\\theta)$ in an observational data set, one cannot directly view the intrinsic symmetry axes of the bright, central galaxies.  Instead, one is forced to use their observed symmetry axes and, in general, these will differ from the intrinsic symmetry axes.  Our simulations show that, even in the limit of multiple deflections being experienced by the distant source galaxies, if one could use the intrinsic symmetry axes of the lenses to define the geometry of the problem, one would expect to observe $\\gamma^+ (\\theta) > \\gamma^- (\\theta)$.   That is, multiple deflections experienced by the source galaxies have little effect on the intrinsic signature of anisotropic galaxy-galaxy lensing by non-spherical haloes.  However, weak lensing of the bright, central foreground galaxies causes their observed symmetry axes (which are used to define the geometry for the calculation of $\\gamma^+ (\\theta)$ and $\\gamma^- (\\theta)$) to differ from their intrinsic symmetry axes (i.e., the unlensed symmetry axes, which define the geometry for the actual lensing of the distant galaxies).  It is this change in the symmetry axes of the bright, foreground galaxies that gives rise to the suppression of the observed function, $\\gamma^+ (\\theta) / \\gamma^- (\\theta)$, compared to the function that would be obtained if the intrinsic symmetry axes were used for the calculation.  The effects of weak lensing of the bright, foreground galaxies on an observation of $\\gamma^+ (\\theta) / \\gamma^- (\\theta)$ cannot be eliminated simply by rejecting foreground galaxies with very small image ellipticities, or by using sources that are particularly close to the observed symmetry axes of the foreground galaxies.  We conclude, therefore, that in order to properly interpret any observed galaxy-galaxy lensing signal (be it isotropic or anisotropic), it is vital that full, multiple-deflection Monte Carlo simulations be used. Especially important is accounting for the fact that the images of the bright, foreground centres are likely to have been weakly lensed.  If the effects of multiple deflections are not taken into account when interpreting an observed galaxy-galaxy lensing signal, there is a high probability that incorrect conclusions will be drawn about the nature of the haloes surrounding the lens galaxies."
    },
    "1004/1004.0388_arXiv.txt": {
        "abstract": "We present  {\\it Hubble Space Telescope} ACS  and STIS FUV/NUV/optical imaging of the radio galaxy 3C~236, whose relic $\\sim 4$ Mpc radio jet lobes and inner 2 kpc CSS radio source are evidence of multiple epochs of AGN  activity. Consistent with  previous results, our  data confirm the presence of four bright knots  of FUV emission in an arc along the edge of the inner circumnuclear  dust disk in the galaxy's nucleus, as well as FUV emission cospatial  with the nucleus itself.  We interpret these to  be sites  of recent or  ongoing star formation.   We present photometry of  these knots,  as well as  an estimate for  the internal extinction in the source  using line ratios from archival ground-based spectroscopy.   We estimate  the ages  of the  knots by  comparing our extinction-corrected  photometry  with  stellar  population  synthesis models. We  find the four  knots cospatial with  the dusty disk  to be young, of order $\\sim10^7$ yr old. The FUV emission in the nucleus, to which  we do not  expect scattered  light from  the AGN  to contribute significantly, is likely due to an episode of star formation triggered $\\sim10^9$ yr  ago.  We argue that  the young $\\sim10^7$  yr old knots stem from  an episode of star  formation that was  roughly coeval with the event resulting in reignition  of radio activity, creating the CSS source.  The $\\sim10^9$ yr old  stars in the nucleus may be associated with the  previous epoch  of radio activity  that generated the  4 Mpc relic source, before being cut off by exhaustion or interruption.  The ages of  the knots,  considered in the  context of both  the disturbed morphology of the nuclear dust and the double-double morphology of the ``old''  and   ``young''  radio  sources,  present   evidence  for  an AGN/starburst  connection that  is  possibly episodic  in nature.   We suggest that  the AGN fuel supply  was interrupted for  $\\sim 10^7$ yr due to a  minor merger event and has now  been restored. The resultant non-steady flow of gas in the  disk is likely responsible for both the new  episode of infall-induced  star formation  and also  the multiple epochs of radio activity. ",
        "introduction": "\\begin{figure*} \\plottwo{overlay.eps}{csssource.ps} \\caption{The  two radio sources  associated with  3C~236. (a)  326 MHz WSRT  radio  contours  (in  green,  from  \\citealt{mack97})  of  the ``relic'' radio  emission associated with 3C~236,  overlayed on SDSS imaging of the same region of  sky.  The deprojected size of the jet is $\\sim 4.5$ Mpc, making it  the largest known radio galaxy and one of the  largest objects in the  universe.  (b) Global  VLBI 1.66 GHz radio contours  of the  central 2 kpc  CSS (``young'')  radio source from \\citet{schilizzi01}, overlayed on {\\it HST}/STIS NUV imaging of the star forming knots described by \\citet{odea01}.  The jet axes of the Mpc- and kpc- scale radio sources are aligned on nearly the same position angle.  } \\label{fig:doubledouble} \\end{figure*}   Galaxies  occupy  a heavily  bimodal  distribution in  color-magnitude space, wherein young,  predominantly disk-dominated galaxies reside in a  `blue  cloud'  and  evolve  onto  a  characteristically  quiescent, bulge-dominated `red  sequence' (e.g., \\citealt{bell04,faber07}).  The underdensity  of  galaxies  in  the `green  valley'  separating  these populations  implies   that  cloud-to-sequence  evolution   is  swift, requiring  a cessation  of star  formation  more rapid  than would  be expected  in  passively  evolving systems  (e.g.,  \\citealt{cowie96}). Quasar-  and   radio-mode  feedback  models  have   been  proposed  as mechanisms by which star formation may be truncated by the heating and expulsion of gas \\citep{silk98,hopkins05,croton06,schawinski06}, as it is now  known that  quasar activity was  two orders of  magnitude more common  at  redshifts  $z\\sim2$   than  at  the  present  time  (e.g., \\citealt{schmidt91}).   This, considered in  the context  of declining star  formation   rates  in   massive  galaxies  at   $z\\sim2$  (e.g., \\citealt{perezgonzalez07}),  along with  the  emerging consensus  that most  populations of galaxies  harbor quiescent  black holes  at their centers (hereafter BHs, e.g,  \\citealt{kormendy95}), has given rise to questions of whether all bright galaxies go through one or more active phases  (e.g., \\citealt{haehnelt93,cavaliere89}).   In  this scenario, the  quenching of  star formation  via feedback  from  active galactic nuclei (AGN)  may be one of  the primary drivers  of cosmic downsizing (e.g., \\citealt{cowie96,scannapieco05}, and references therein).    The relationship between the AGN duty cycle and the regulation of host galaxy stellar evolution  is far more complicated, however,  as it can play competing  roles at successive stages of  galactic evolution. AGN activity has been associated not only with quenching star formation on large scales,  but also triggering  it via ISM cloud  compression from the propagating relativistic jets associated with radio galaxies (e.g, the             so-called             ``alignment            effect'', \\citealt{rees89,baum89a,mccarthy93,best00,  privon08}).   Moreover, it is  natural  to  expect  a  correlated (but  not  necessarily  causal) relationship  between  AGN activity  and  star  formation.  The  tight relationship between BH mass and host galaxy bulge velocity dispersion \\citep{magorrian98,ferrarese00,gebhardt00} implies  that the growth of the   BH   and  the   galaxy   bulge   are   tightly  coupled   (e.g., \\citealt{kauffmann00,ciotti01}).   It   is  therefore  expected  that, throughout the  process of  hierarchical galaxy formation,  gas infall due to major mergers or  tidal stripping from a gas-rich companion can fuel not  only AGN,  but also  the growth of  the host  galaxy stellar component via  infall-induced starbursts (e.g., \\citealt{dimatteo05}). A significant fraction ($\\sim 30\\%$) of nearby powerful radio galaxies exhibit  evidence  of  infall-induced  starbursts  near  their  nuclei \\citep{smith89,allen02,baldi08,tremblay09},    suggesting   that   the phenomenon is both common and  comperable to the lifetime of the radio source    ($\\sim10^7-10^8$   yr,   e.g.,    \\citealt{parma99}).    The AGN/starburst connection  is therefore likely real  and fundamental to galaxy  evolution itself, and  its characterization  has been  a major pursuit      of       the      past      two       decades      (e.g., \\citealt{barnes91,silk98,fabian99,dimatteo05,hopkins05,springel05,silverman08,quillen08a}).       A key  discriminant in understanding  the nature and evolution  of the AGN/starburst  connection may be  found in  some radio  galaxies whose morphology  is clear  evidence for  multiple epochs  of  AGN activity. Several   such  examples   have   been  observed   (e.g.,  3C~219   -- \\citealt{bridle86, clarke92}, 0108+388  -- \\citealt{baum90}), and have come  to constitute a  new class  of ``double-double''  radio sources, representing  $\\sim 5-10$\\%  of predominantly  large ($>1$  Mpc) radio galaxies       (e.g.,      \\citealt{schoenmakers00a,schoenmakers00b}). Double-doubles  are   characterized  by  outer   (`older')  and  inner (`younger') radio sources propagating outwards amidst the relic of the previous epoch  of activity.   This apparently repetitive  activity is thought  to  be a  consequence  of the  AGN  fuel  supply having  been interrupted,  whether by  exhaustion, smothering,  or  disturbance, at some  time in the  past \\citep{baum90}.   This scenario  is consistent with     models     of     radio     galaxy     propagation     (e.g., \\citealt{kaiser00,brocksopp07}).   The  relative  ages  of  the  radio sources (and therefore the timescale over which the engine was cut off and re-ignited) can be estimated  using size estimates from radio maps coupled with a dynamical model  for the jets and radio spectral energy distributions      (SEDs)     of      the      radiating     electrons \\citep{schoenmakers00a,schoenmakers00b,odea01}.     \\subsection{An important test case: 3C~236}  The  nearby  ($z=0.1005$)  double-double  radio galaxy  3C~236  is  an important test case in studies of the AGN/starburst connection, and is the basis of  both this paper and a  previous study by \\citet{odea01}. 3C~236 is a powerful  double-double with a relic edge-brightened FR~II \\citep{fanaroff74}  radio  source   whose  deprojected  linear  extent exceeds 4 Mpc,  making it the second largest  known radio galaxy (only J1420-0545  is larger,  \\citealt{machalski07}),  and even  one of  the largest objects in the  universe \\citep{schilizzi01}.  Its inner young Compact  Steep  Spectrum  (CSS)   source,  whose  apparent  origin  is cospatial  with  the  nucleus,  is   only  2  kpc  in  extent  and  is morphologically reminiscent of a young classical double.  Anecdotally, 3C~236 was  initially classified  as a pure  CSS source before  it was associated years later with  the massive relic FR~II source (R.~Laing, private communication).  The jet propagation axes of both the Mpc- and kpc- scale sources  are aligned to within $\\sim10\\deg$  of one another (as projected on the  sky).  See Fig.~\\ref{fig:doubledouble} for radio contour overlays  of both sources, using 326  MHz Westerbork Synthesis Radio  Telescope  (WSRT)  and  1.66  GHz  Global  Very  Long  Baseline Interferometry  (VLBI) Network radio  mapping from  \\citet{mack97} and \\citet{schilizzi01} for the relic and CSS sources, respectively.   In addition  to its  rare radio morphology,  3C~236 is also  unique in that its  nuclear dust  complex is made  up of an  inner circumnuclear dusty  disk that is  somewhat misaligned  with an  apparently separate outer dust lane  \\citep{martel99,dekoff00,tremblay07}.  The total dust mass in the complex is estimated to be $\\sim 10^7$ M$_\\odot$, based on {\\it  Hubble Space  Telescope} ({\\it  HST}) absorption  maps  and IRAS luminosities \\citep{dekoff00}.   In Fig.~\\ref{fig:colormap} we present a  $1.6$ $\\mu$m  /  0.7 $\\mu$m  absorption  map of  the dust  complex, originally presented  in \\citet{tremblay07}  and made via  division of {\\it   HST}/NICMOS   and   WFPC2   data  from   \\citet{martel99}   and \\citet{madrid06}, respectively. \\begin{figure} \\plotone{3c236_colormap.eps} \\caption{$1.6$  $\\mu$m /  0.7 $\\mu$m  colormap of  the outer  lane and inner dusty disk in the nucleus of 3C~236, made via division of {\\it HST}/NICMOS   and    WFPC2   data   from    \\citet{martel99}   and \\citet{madrid06}, respectively.  This absorption map  was originally presented in \\citet{tremblay07}. } \\label{fig:colormap} \\end{figure}        \\begin{deluxetable*}{ccccccccc} \\tablecaption{Summary of Observations of 3C~236} \\tablewidth{0pc} \\tablehead{ \\colhead{Observatory} & \\colhead{Instrument} & \\colhead{Aperture} & \\colhead{Filter/Config.} & \\colhead{Waveband/Type} & \\colhead{Exp. Time [Orbits]} & \\colhead{Reference}  & \\colhead{Obs. Date} & \\colhead{Comment} \\\\ \\colhead{(1)} & \\colhead{(2)} & \\colhead{(3)} & \\colhead{(4)} & \\colhead{(5)} & \\colhead{(6)} & \\colhead{(7)} & \\colhead{(8)} & \\colhead{(9)} } \\startdata \\cutinhead{{\\sc New Observations}} {\\it HST}  & ACS & HRC   &  F330W  & $U$-band Imaging&  2516s [1] &  {\\it HST} 9897 & 21 Oct 2003 & SF Knots \\\\ {\\it HST}  & ACS  & HRC  &  F555W & $V$-band Imaging &  2612s [1] & {\\it HST} 9897   & 22 Oct 2003 & Dust Lanes   \\\\ {\\it HST}  & ACS  & SBC  &  F140LP &   FUV Imaging & 6900s [3]  & {\\it HST} 9897 & 21 Oct 2003 & SF Knots \\\\ {\\it HST}  & STIS & NUV-MAMA    &  F25SRF2  & NUV Imaging & 2520s[1]  &  {\\it HST} 9897 & 19 Oct 2003 & SF Knots \\\\ \\cutinhead{{\\sc Archival Observations}} {\\it HST}  & WFPC2 & PC1    &  F702W  & $R$-band Imaging&  4$\\times$140s &  {\\it HST} 5476  & 7 May 1995 & Galaxy \\\\ {\\it HST}  & WFPC2 & PC1    &  F555W  & $V$-band Imaging&  2$\\times$300s &  {\\it HST} 6384  & 12 Jun 1996 & Galaxy \\& Dust \\\\ {\\it HST}  & STIS &  NUV-MAMA   &  F25SRF2  & NUV Imaging &  1440s &  {\\it HST} 8275  & 03 Jan 1999 & SF Knots \\\\ {\\it HST}  & NICMOS &  NIC2-FIX   &  F160W  & NIR Imaging &  1152s &  {\\it HST} 10173  & 02 Nov 2004  & Host Isophotes\\cr \\enddata \\tablecomments{ A summary of the new and archival observations used in our analysis. (1) Facility name; (2) instrument used for observation; (3) configuration of instrument used; (4) filter used; (5) corresponding waveband and specification of whether the observation was imaging or spectroscopy; (6) exposure time (if the observatory is {\\it HST}, the corresponding number of orbits also appears in brackets); (7) corresponding reference for observation. If the observatory is {\\it HST}, the STScI-assigned program number is listed; (8) date of observation; (9) comment specific to observation. } \\label{tab:tab1} \\end{deluxetable*}     \\begin{figure*} \\plottwo{hrc_color.eps}{sbc_combined_color.eps} \\caption{({\\it  left}) 2500  s  $V$-band exposure  of  the nucleus  of 3C~236, using {\\it HST}/ACS HRC with the broadband F555W filter. The outer dust lane  is seen in white, while three of  the four knots of star formation are  seen to the south and west  of the nucleus along the inner dusty  disk, whose position angle is  slightly offset from that of  the outer.({\\it right})  combined 6900 s {\\it  HST}/ACS SBC FUV  (F140LP) image  of the  star forming  knots observed  along the inner dust structure of 3C~236.   The image has been smoothed with a two  pixel Gaussian  kernel. North  is up,  east is  left.   The two images  are  on the  same  scale.  At  a  redshift  of $z\\sim  0.1$, 1\\arcsec corresponds to $\\sim 1.8$ kpc.} \\label{fig:acs} \\end{figure*}        The work by  \\citet{odea01} studied {\\it HST} NUV  and optical imaging of  the central  few  arcsec of  3C~236,  finding four  knots of  blue emission  arranged in  an  arc along  the  dust lane  in the  galaxy's nucleus.  Their original NUV data is presented in greyscale with black contours  in  Fig.~\\ref{fig:doubledouble}({\\it  b}).  The lack  of  an obvious  spatial relationship  between the  knots and  the  CSS source suggests   that  the   starbursts  are   infall-induced   rather  than jet-induced. \\citet{martel99} had also  detected the knots of emission in their {\\it HST} $R$-band imaging,  albeit to a lesser degree as the knots  are  very  blue.    \\citet{odea01}  used  their  photometry  in comparison with stellar population  synthesis models to estimate upper limits to  the ages  of the individually  resolved star  forming knots (seen  in Fig.~\\ref{fig:doubledouble}{\\it  b}).  They  found disparate ages  between the  clumps of  emission, finding  two to  be relatively young with  ages of  order $\\sim  10^7$ yr, while  the other  two were estimated at  $\\sim 10^8 - 10^9$  yr old, comparable  to the estimated age  of the  giant  relic radio  source.   \\citet{odea01} argued  that 3C~236  is an  ``interrupted''  radio galaxy,  and  has undergone  two starburst episodes  approximately coeval with the two  epochs of radio activity  observed  on Mpc-  and  kpc-  scales.   That work  motivated follow-up  observations  with  {\\it  HST} at  higher  sensitivity  and spatial resolution, the results of which we present in this paper.   We organize  this work as follows.   In section 2 we  describe the new and archival data  presented in this paper, as  well as the associated data  reduction. In  section 3  we  present our  results, including  a comparison of our photometry with stellar population synthesis models, following (in the interests of consistency) the analysis strategy used in \\citet{odea01}.  We  discuss our results in section  4, focusing on the  role  played  by  the  AGN/starburst connection  in  the  special test-case environment  of 3C~236.  We summarize this  work and provide some concluding  remarks in section  5.  Throughout this paper  we use $H_0   =  71$  km   s$^{-1}$  Mpc$^{-1}$,   $\\Omega_M  =   0.27$,  and $\\Omega_{\\Lambda} = 0.73$. ",
        "conclusions": "We have presented follow-up {\\it HST} ACS and STIS observations of the radio galaxy 3C~236, described by \\citet{odea01} as an ``interrupted'' radio source.  The galaxy is  associated with a massive relic $\\sim 4$ Mpc FR~II  radio source (making it  one of the largest  objects in the universe), as well as an inner 2 kpc CSS ``young'' radio source.  This ``double-double'' radio morphology is  evidence for multiple epochs of AGN  activity, wherein  the BH  fuel supply  is thought  to  have been exhausted or cut  off at some time in the past,  and has only recently been reignited.   We present {\\it HST} FUV, NUV, $U$-, and $V$-band imaging of four star forming  knots,  previously  described  by  \\citet{odea01},  that  are arranged in an arc along  the outer edge of the galaxy's circumnuclear dust  disk   (which  itself  is  surrounded  by   a  misaligned  outer filamentary dust lane). We  have also detected blue emission cospatial with the nucleus itself. We  describe these observations in detail, as well as the steps taken to  reduce the data.  We present photometry of the  blue knots,  and  discuss our  efforts  to correct  the data  for internal extinction to the source using the Balmer decrement available from archival SDSS spectroscopy.    We compare the  measured four-color SEDs of the  star forming knots to synthetic  SEDs from  {\\sc starburst99}  stellar  population synthesis models, with the  ultimate goal of roughly estimating  the ages of the knots. We  find that the four  knots cospatial with the  outer edge of the dusty  disk are likely $\\sim10^7$  yr old, while  the FUV emission cospatial with the nucleus is likely older, at $\\sim10^9$ yr old (with the caveat that undercorrection for internal extinction in the nucleus would lower this limit). We argue that the ages of the young knots are suggestive of a causal connection with the young central radio source.   We  frame these  results in  the context  of 3C~236  as  an apparently ``interrupted''  radio galaxy.  Our  results are  generally consistent with those  of \\citet{odea01}, and  we argue along similar  lines that the transport  of gas in the  nucleus of 3C~236  is nonsteady, wherein the active phase giving rise to the  4 Mpc relic source was cut off by exhaustion or disturbance of the AGN fuel supply. We suggest that this lead to a  dormant period punctuated by a minor  merger event, and the subsequent  infalling gas  triggered not  only a  new episode  of star formation, but also ushered the  galaxy into a new active phase giving rise to the young CSS radio source.   Results such as these support  the natural argument that infalling gas largely fosters  the relationship  between the growth  of the  AGN and that of its host galaxy  stellar component. In time, that relationship may evolve from  mutual growth to the regulation of  the latter by the former through  quenching from  AGN feedback.  As  infalling gas  is a critical   component  in   merger-driven   hierarchical  models,   the triggering of AGN  may be fundamental to galaxy  evolution itself.  In this  way, the  relevance of  3C~236  might extend  beyond studies  of episodic activity in  radio galaxies to studies of  AGN in the context of galaxy evolution as a whole."
    },
    "1004/1004.0985.txt": {
        "abstract": "The Great Observatories All-sky LIRG Survey (GOALS) consists of a complete sample of 202 Luminous Infrared Galaxies (LIRGs) selected from the IRAS Revised Bright Galaxy Sample (RBGS).  The galaxies span the full range of interaction stages, from isolated galaxies to interacting pairs to late stage mergers.  We present a comparison of the UV and infrared properties of 135 galaxies in GOALS observed by GALEX and Spitzer.  For interacting galaxies with separations greater than the resolution of GALEX and Spitzer ($\\sim2-6^{\\prime\\prime}$), we assess the UV and IR properties of each galaxy individually. The contribution of the FUV to the measured SFR ranges from 0.2\\% to 17.9\\%, with a median of 2.8\\% and a mean of $4.0\\pm0.4$\\%. The specific star formation rate of the GOALS sample is extremely high, with a median value ($3.9\\times10^{-10}~{\\rm yr}^{-1}$) that is comparable to the highest specific star formation rates seen in the Spitzer Infrared Nearby Galaxies Survey sample.  We examine the position of each galaxy on the IR excess--UV slope (IRX-$\\beta$) diagram as a function of galaxy properties, including IR luminosity and interaction stage.  The LIRGs on average have greater IR excesses than would be expected based on their UV colors if they obeyed the same relations as starbursts with $L_{IR}<10^{11}L_{\\odot}$ or normal late-type galaxies. The ratio of $L_{IR}$ to the value one would estimate from the IRX-$\\beta$ relation published for lower luminosity starburst galaxies ranges from 0.2 to 68, with a median value of 2.7.  A minimum of $19\\%$ of the total IR luminosity in the RBGS is produced in LIRGs and ULIRGs with red UV colors ($\\beta>0$). Among resolved interacting systems, $32\\%$ contain one galaxy which dominates the IR emission while the companion dominates the UV emission. Only $21\\%$ of the resolved systems contain a single galaxy which dominates both wavelengths. ",
        "introduction": "The Infrared Astronomical Satellite (IRAS) provided the first unbiased survey of the sky at mid and far-infrared wavelengths, giving us a comprehensive census of the infrared emission properties of galaxies in the local Universe.  A major result of this survey was the discovery of a large population of luminous infrared galaxies (LIRGs) which emit a large majority of their bolometric luminosity in the far-infrared, and have $10^{11}\\leq{\\rm L_{IR}} [8-1000\\mu{\\rm m}] <10^{12} L_{\\odot}$.  LIRGs are a mixture of single galaxies, disk galaxy pairs, interacting systems, and advanced mergers.  They exhibit enhanced star-formation rates and a higher fraction of Active Galactic Nuclei (AGN) compared to less luminous and non-interacting galaxies (Sanders \\& Mirabel 1996 and references therein). At the highest luminosities, ultraluminous infrared galaxies (ULIRGs: ${\\rm L_{IR}} \\geq 10^{12} L_{\\odot}$) may represent an important evolutionary stage in the formation of QSOs \\citep{sanders88a,sanders88b} and massive ellipticals \\citep[e.g.,][]{genzel01,tacconi02}. Since LIRGs comprise the bulk of the cosmic infrared background and dominate the star-formation activity between $0.5 < z < 1$ \\citep{lefloch05,caputi06}, they may also play a key role in our understanding of the general evolution of galaxies and black holes \\citep[e.g.,][]{magorrian98}.  The Great Observatories All-sky LIRG Survey \\citep[GOALS;][]{armus09} contains a complete sample of low-redshift LIRGs and ULIRGs with observations across the electromagnetic spectrum.  The GOALS targets are drawn from the IRAS Revised Bright Galaxy Sample (RBGS; Sanders et al. 2003), a complete sample of 629 galaxies with IRAS $60~\\mu$m flux densities ${\\rm S}_{60} > 5.24$~Jy, covering the full sky above Galactic latitudes $|b| > 5$ degrees.  The 629 galaxies have a median redshift of $z = 0.008$ and a maximum redshift of $z=0.088$. There are 181 LIRGs and 21 ULIRGs in the RBGS, and these galaxies define the GOALS sample.  In LIRGs and ULIRGs, UV radiation is produced by young stars and AGN.   A fraction of the UV radiation is absorbed by dust and re-radiated in the far-infrared.  To understand the power sources in these galaxies, it is essential to fully characterize the energy budget by measuring both the emerging UV and the infrared flux.  The relationship between the UV continuum slope and the infrared excess (the IRX-$\\beta$ correlation) provides a useful parameterization of this energy budget.  \\citet{charlot00} showed that the IRX-$\\beta$ relation is a sequence in effective optical depth for star forming systems. However this relation does not hold in all systems.  While lower luminosity starbursts follow the correlation, ULIRGs do not \\citep*{meurer,goldader}. The GOALS sample allows us to explore the IRX-$\\beta$ correlation precisely over the luminosity range where it breaks down. A detailed study of LIRGs may indicate the luminosity threshold or the time during the merger when the UV slope becomes decoupled from the IR emission.  Being a flux limited sample of the nearest and most well studied LIRGs and ULIRGs, GOALS provides an important local benchmark against which to compare the observed visual properties of high redshift galaxies.  This paper looks at global UV and IR properties.  Future work will address nearby spatially resolved LIRGs.  This paper is divided into five sections.  The data are discussed in \\S~2.  Analysis of the sample is presented in \\S~3, results are discussed in \\S~4, and conclusions are given in \\S~5.  A cosmology of $\\Omega_{\\Lambda}=0.72$, $\\Omega_m = 0.28$, with $H_0 = 70~{\\rm km~s^{-1}~Mpc^{-1}}$ is adopted throughout. ",
        "conclusions": "We present a comparison of the UV and infrared properties of 135 LIRGs and ULIRGs in the GOALS sample observed by GALEX and Spitzer.  We find that: \\begin{itemize} \\item{LIRGs have larger IR excesses than lower luminosity galaxies of similar UV color.  On average, more luminous LIRGs and ULIRGs have larger IRX and redder colors.} \\item{The contribution of the FUV to the measured SFR is on average $4\\%$; UV emission alone is not a reliable indicator of the SFR for LIRGs.} \\item{The median SSFR of the GOALS sample ($3.9\\times10^{-10}~{\\rm yr}^{-1}$, corresponding to a mass doubling timescale of 2.6~Gyr) is approximately equal to the maximum SSFR seen in lower luminosity galaxies, however the median IR/UV ratio (39) for GOALS galaxies is more than an order of magnitude greater.} \\item{Deviations from the starburst IRX-$\\beta$(GALEX) relation $\\Delta$IRX increase with IR luminosity for $L_{IR}\\gtrsim10^{10}~L_{\\odot}$, with considerable scatter.  LIRG systems with IRAC colors that may indicate the presence of an AGN have average IRX ratios a factor of six larger than the rest of the sample.  $\\Delta$IRX is not strongly correlated with IRAS $25~\\mu$m/$60~\\mu$m color, IRAS $60~\\mu$m/$100~\\mu$m color, Spitzer $8~\\mu$m/$24~\\mu$m color, $L_{FUV}$, or $8~\\mu$m concentration (1~kpc/Total).} \\item{A minimum of 19\\% of the total $L_{IR}$ of the RBGS sample is produced in LIRGs and ULIRGs with $\\beta>0$, sources that are typically absent from UV-selected samples at high redshift.  A minimum of 11\\% of the total $L_{IR}$ of the RBGS sample is produced in LIRGs and ULIRGs with $\\Delta{\\rm IRX}>1$, an order of magnitude above the starburst relation.} \\item{Using the starburst IRS-$\\beta$ relation to estimate $L_{IR}$ from rest-frame UV observations of LIRGs and ULIRGs would underestimate $L_{IR}$ by a factor of three on average, with a wide range (factors of 0.2--68) of possible under- or over estimates, particularly for red UV colors (large values of $\\beta$) where $L_{IR}$ could be overestimated by as much as a factor of 2400 using a linear extrapolation of the starburst relation.} \\item{The UV and IR properties of GOALS systems are qualitatively consistent with an evolutionary picture in which some galaxies transition from LIRGs to ULIRGs over the course of a major merger event.  More luminous galaxies, mergers, and galaxies with high SSFR are more heavily obscured than less luminous galaxies, non-mergers, and galaxies with lower SSFR.} \\item{Among LIRG systems resolved into individual interacting galaxies, pairs in which one galaxy dominates the IR emission while the companion dominates UV emission (such as the well-studied LIRG system VV~114) are more common than pairs in which one galaxy dominates both wavelengths (32\\% and 21\\% of the sample, respectively).  On average, galaxies which dominate both wavelengths have $\\Delta$IRX values four times larger than an IR-dominant galaxy in a ``VV~114-like\" system.  The large fraction of ``VV~114-like\" systems has important implications for observations of interacting galaxies at high redshift in that the IR and UV properties of the unresolved systems can differ by over an order of magnitude from the properties of the component galaxies.} \\end{itemize}"
    },
    "1004/1004.0207_arXiv.txt": {
        "abstract": "The detection of non-Gaussianity in the CMB data would rule out a number of inflationary models. A null detection of non-Gaussianity, instead, would exclude alternative models for the early universe. Thus, a detection or non-detection of primordial non-Gaussianity in the CMB data is crucial to discriminate among inflationary models, and to test alternative scenarios. However, there are various non-cosmological sources of non-Gaussianity. This makes important to employ different indicators in order to detect distinct forms of non-Gaussianity in CMB data. Recently, we proposed two new indicators to measure deviation from Gaussianity on large angular scales, and used them to study the Gaussianity of the \\emph{raw} band WMAP maps with and without the \\emph{KQ75} mask. Here we extend this work by using these indicators to perform similar analyses of deviation from Gaussianity of the \\emph{foreground-reduced}  Q, V, and W band maps. We show that there is a significant deviation from Gaussianity in the considered full-sky maps, which is reduced to a level consistent with Gaussianity when the \\emph{KQ75} mask is employed. ",
        "introduction": "Cosmic Microwave Background (CMB) data from the Wilkinson Microwave Anisotropy Probe (WMAP)\\cite{WMAP1} is under rigorous scrutiny for the possible deviations from Gaussianity\\cite{non-Gauss}\\cdash\\cite{BRb} and for statistical isotropy\\cite{RELATED} in the CMB temperature field. A detection or non-detection of primordial non-Gaussianity in the CMB data is crucial to discriminate among inflationary models, and to test alternative scenarios.\\cite{Bartolo} However, since non-cosmological effects can produce non-Gaussianity in the CMB data, the extraction of a possible \\emph{primordial} non-Gaussianity is not a simple enterprise. Besides, one does not expect that any single statistical estimator can be sensitive to all possible forms of non-Gaussianity. It is therefore important to test CMB data for deviations from Gaussianity by using different statistical tools, since they can potentially provide information about multiple forms of non-Gaussianity that may be present in CMB data. In a previous work\\cite{BR} we introduced two indicators, based on skewness and kurtosis of CMB data, and used them to study the deviation from Gaussianity in the WMAP \\emph{raw} band (uncleaned) maps. In the present work, we complement this study of WMAP five-year data by performing a similar analysis but now using the \\emph{foreground-reduced} Q, V, and W band maps with and without the \\emph{KQ75} mask. We show that there is a significant deviation from Gaussianity in these \\emph{foreground-reduced} full-sky band maps, which is reduced to a level that is consistent with Gaussianity when the \\emph{KQ75} mask is used. ",
        "conclusions": "In this section we report the results of our Gaussianity analysis performed with the indicators discussed in the previous section.  To minimize the statistical noise, in the calculations of $S$ and $K$ maps from the foreground-reduced Q, V, W input maps, we have scanned the celestial sphere with spherical caps of aperture $\\gamma = 90^{\\circ}$, whose centers are located at $12\\,288$ points homogeneously distributed on the celestial sphere. The input maps we have used in our analyses have the resolution HEALPix parameter $N_{\\mbox{\\footnotesize side}}= 256$,\\cite{Gorski-et-al-2005} which corresponds to a partition of the celestial sphere into $786\\,432$ pixels. The calculations of the $S$ and $K$ maps by scanning the CMB masked maps sometimes include caps whose center is within or close to the \\emph{KQ75} masked region. In these cases, the calculations are made with a smaller number of pixels $N_{\\mbox{\\footnotesize p}}$. Finally, we note that examples of $S$ and $K$ maps can be found in Refs.~\\refcite{BR} and \\refcite{BRb}.  In Fig.~\\ref{fig1} we show the angular power spectra of the $S$ maps and $K$ maps calculated from the WMAP foreground-reduced Q, V, and W CMB band maps by considering the data in the full-sky and \\emph{KQ75} masked maps. These angular power spectra show that the \\emph{KQ75} mask lowers the non-Gaussianity to a level below  $95.4 \\,\\%$~CL.  An overall quantitative measure of deviation from Gaussianity for $\\ell=1-10$  is obtained through the $\\chi^2$ goodness-of-fit test by calculating the $\\chi^2/\\text{dof}$ (where dof stands for degrees of freedom) of the spectra values $S_{\\ell}$ and $K_{\\ell}$ calculated from the Q, V, and W band maps as compared to the mean power spectra $\\overline{S}_{\\ell}$ and $\\overline{K}_{\\ell}$ obtained from $S$ and $K$ maps calculated from MC CMB maps. The results regarding the masked and the full-sky maps are shown in Table~\\ref{table1}. The values obtained in the masked maps analysis ($\\chi^2 \\lesssim 10$) indicate that these may be seen as consistent with Gaussianity, while the values  for full-sky foreground-reduced maps ($\\chi^2 > 10^{10}$) shows a large deviation from Gaussianity.  \\begin{figure} \\mbox{\\hspace{-1.0cm} \\epsfig{file=fig1a.ps,height=7.5cm,width=7.cm}}  \\vspace{-7.55cm} \\mbox{\\hspace{6.0cm} \\epsfig{file=fig1b.ps,height=7.5cm,width=7.cm}} \\caption{Low $\\ell $ \\emph{differential} power spectra of skewness $|S_{\\ell} - \\overline{S}_{\\ell}|$ (left) and kurtosis (right) $|K_{\\ell} - \\overline{K}_{\\ell}|$ calculated from the WMAP \\emph{foreground-reduced} band maps. The values for the full-sky and \\emph{KQ75} masked maps are shown. The $95.4\\%$ confidence level (obtained from $S$ and $K$ maps calculated from MC Gaussian CMB maps) is indicated by the dashed line. More details in the text. \\label{fig1}} \\end{figure}   \\vspace{0.3cm} \\begin{table} \\tbl{Values for the ratio $\\chi^2/\\text{dof}$ that measure the goodness-of-fit of the $S_{\\ell}$ and $K_{\\ell}$ spectra from the WMAP \\emph{foreground-reduced} band maps as compared to the mean spectra $\\overline{S}_{\\ell}$ and $\\overline{K}_{\\ell}$, respectively. The values for the full-sky and \\emph{KQ75} masked maps are shown.} {\\begin{tabular}{@{}lcccccc@{}} \\toprule   % $\\chi^2$  $\\backslash$ CMB maps & Q [full-sky]  & V [full-sky] & W [full-sky] & Q [\\emph{KQ75}] & V [\\emph{KQ75}] & W [\\emph{KQ75}] \\\\ \\colrule $S_{\\ell}$  & $9.4 \\times 10^{11}$  & $1.2 \\times 10^{12}$ & $7.3 \\times 10^{10}$ & 3.1 & 4.4 & 3.4 \\\\ $K_{\\ell}$  & $1.2 \\times 10^{21}$  & $3.3 \\times 10^{21}$ & $1.0 \\times 10^{20}$ & 6.3 & 6.5 & 6.0  \\\\ \\botrule     % \\end{tabular} \\label{table1}} \\end{table}"
    },
    "1004/1004.2728_arXiv.txt": {
        "abstract": "We present multiwavelength broadband photometry and $V, I$ time resolved photometry for two variable bright stars in the SMC, OGLE004336.91-732637.7 (\\va) and \\vbb~ (\\vb). The light curves span 12 years and  show long-term periodicities (SMC-SC3) and modulated eclipses (SMC-SC4) that are discussed in terms of wide-orbit intermediate mass interacting binaries and associated  envelopes. SMC-SC3 shows a primary period of 238.1 days  along with a complicated waveform suggesting ellipsoidal variablity influenced by an eccentric orbit. This star also shows a secondary variability with an unstable periodicity that has a mean value of 15.3 days. We suggest this could be associated with nonradial pulsations. ",
        "introduction": "The evolution of massive stars is an important topic of stellar astrophysics because it provides clues to the mechanisms that feed the Galactic medium with the building blocks of future generations of stars. Many of the evolutionary stages of massive stars are short-lived and hence challenge our ability to find enough examples of a common group to characterize them. Detecting objects in these brief phases of evolution is of great use in testing current theories of massive star evolution, especially if they represent populations with low metal abundances, a property of stars in the Magellanic Clouds. A complete census of pre- and just post-main sequence population of massive stars in the Clouds has not yet been compiled, and thus attributes of these subpopulations are not well known, but they are known to be generally variable. Catalogs of B stars with unusual variable light curves in the Magellanic Clouds (e.g., Mennickent et al. 2002, hereafter M02 and Sabogal et al. 2005), based on OGLE photometry (Udalski et al. 1997, Szyma\\'nski 2005), provide excellent material for analysis of massive stars near the main sequence. In an effort to  characterize some of these stars, follow up spectroscopy was conducted for two stars in the Type\\,3 M02 sample\\footnote{Type-3 variables are SMC Be star candidates with $I$-band light curves varying periodically or quasi-periodically.}, that will be given in a future paper (Mennickent \\& Smith 2010, hereafter MS10).  Both these stars turn out to be probable binaries and to have highly peculiar characteristics. In this paper we present the photometric results for these stars.   The stars chosen were taken from an initial  sample of 8 Type-3 variables in the M02 catalog satisfying the arbitrary criterion of having visual magnitudes brighter than 14.2. No other criteria were imposed on their selection. The stars selected were \\vaa~  ($\\equiv$ SMC-SC3-63371, MACHO ID 213.15560; hereafter \\va) and \\vbb~ ( $\\equiv$ SMC-SC4-67145, MACHO ID 212.15735.6; hereafter \\vb).  Exploratory  optical spectra  exhibited substantial H$_{\\alpha}$ emissions in both objects, and in the case of SMC-SC3, multiperiodic variability (Mennickent et al. 2006, hereafter M06). \\va~ was recently included in the slitless survey of H$\\alpha$ emission line objects by Martayan et al. (2010, their star in cluster SMC-17).   The goals of this paper are to characterize the photometric properties of these two stars with the longer time baseline available and gain further insights on the formation and nature of these systems.   \\clearpage  \\begin{deluxetable}{cccccccccc} \\tabletypesize{\\scriptsize} \\tablewidth{0pt} \\tablecaption{$UBVR$ magnitudes from Massey (2002) and OGLE photometry are given. f Dereddened $B-V$ colors and derived spectral types are also included. } \\tablehead{ \\colhead{Object} & \\colhead{U} & \\colhead{B} & \\colhead{V} & \\colhead{R}&  \\colhead{V$_{OGLE}$} & \\colhead{(B-V)$_{OGLE}$} & \\colhead{(V-I)$_{OGLE}$}  & \\colhead{(B-V)$_{0}$} &Sp. type } \\startdata \\va~ & -&  13.63&  13.48&  13.30& 14.18&  0.181& 0.331& 0.08  &A4\\\\ \\vb~ & 14.34& 14.15&  13.94 & 13.70 & 14.06&  0.206& 0.385& 0.11&A5  \\\\ \\enddata \\end{deluxetable}  \\clearpage    \\clearpage  \\begin{deluxetable}{cccccccc} \\tabletypesize{\\scriptsize} \\tablewidth{0pt} \\tablecaption{Infrared magnitudes for program stars. Phases refers to ephemeris given in Equations 1 and 4.} \\tablehead{ \\colhead{Object} & \\colhead{I} & \\colhead{J} & \\colhead{H} & \\colhead{K}&  \\colhead{JD/Date} & \\colhead{Phase} & \\colhead{Source}  } \\startdata \\va~  &13.701(9) &13.346(22) &             -&12.954(116) &1998-08-12  &0.52   & cds.u-strasbg.fr/denis.html\\\\ \\va~  &13.847(30)&13.472(90) &             -&13.050(180) &2450414.6148&0.90 & cds.u-strasbg.fr/denis.html \\\\ \\va~  &13.781(40)&13.560(110)&             -&13.134(160) &2451048.7763&0.56   & cds.u-strasbg.fr/denis.html \\\\ \\va~  &-         &13.545(42) &13.341(50)    &13.275(40)  &2451034.7109&0.50   & www.ipac.caltech.edu/2mass/\\\\ \\va~  &-   &13.540(20)& 13.380(10)& 13.180(20)&-&-&pasj.asj.or.jp/v59/n3/590315/\\\\ \\vb~  &13.655(30)&13.388(90) &             -&13.316(210) &2450418.5524&0.42& cds.u-strasbg.fr/denis.html \\\\ \\vb~  &13.616(30)&13.383(130)&             -&13.079(150) &2451039.7991&0.79& cds.u-strasbg.fr/denis.html \\\\ \\vb~  &-         &13.403(29)&13.236(34)     &13.054(35)  &2451034.7134&0.76& www.ipac.caltech.edu/2mass/\\\\ \\vb~  &-   &13.380(10)& 13.260(10)& 13.170(20)&-&-&pasj.asj.or.jp/v59/n3/590315/\\\\ \\enddata \\end{deluxetable}  \\clearpage ",
        "conclusions": "We have presented the analysis of OGLE light curves of  two SMC bright systems showing novel photometric properties. These properties are consistent with the interpretation that these stars are long-period interacting binaries with an evolved most-luminous stellar component. We find eclipses in SMC-SC4 with P$_{o}$= 184 days modulated in depth and  perhaps  shape on time scales of hundreds of days, suggesting the presence of a variable and non-stellar eclipsing region. In addition, we discovered an unusually strong reflection effect  in the orbital light curve  of SMC-SC3 (P$_{o}$= 238 days) and a short variability with quasi-period 15.3 days changing on time scales of 3800 days. We note the possibility that the short-term fluctuations observed in both stars are  signatures of  nonradial pulsation.  This may explain the ``drifting\" of a single 15-day periodicity in the light curve of SMC-SC3."
    },
    "1004/1004.1617_arXiv.txt": {
        "abstract": "In this paper we discuss two approximate methods previously suggested for modeling hyperfine spectral line emission for molecules whose collisional transitions rates between hyperfine levels are unknown. Hyperfine structure is seen in the rotational spectra of many commonly observed molecules such as HCN, HNC, NH$_3$, \\diaz , and C$^{17}$O. The intensities of these spectral lines can be modeled by numerical techniques such as $\\Lambda-$iteration that alternately solve the equations of statistical equilibrium and the equation of radiative transfer. However, these calculations require knowledge of both the radiative and collisional rates for all transitions.  For most commonly observed radio frequency spectral lines, only the net collisional rates between rotational levels are known. For such cases, two approximate methods have been suggested. The first method, hyperfine statistical equilibrium ({\\it HSE}), distributes the hyperfine level populations according to their statistical weight, but allows the population of the rotational states to depart from local thermodynamic equilibrium (LTE). The second method, the {\\it proportional} method approximates the collision rates between the hyperfine levels as fractions of the net rotational rate apportioned according to the statistical degeneracy of the final hyperfine levels. The second method is able to model non-LTE hyperfine emission. We compare simulations of \\diaz\\ hyperfine lines made with approximate and more exact rates and find that satisfactory results are obtained. ",
        "introduction": "\\label{introduction}  The rotational spectra of many commonly observed molecules such as HCN, HNC, NH$_3$, \\diaz , and C$^{17}$O exhibit hyperfine structure from the splitting of the rotational energy levels by electric quadrupole and magnetic dipole interactions induced by the nuclear moments of atoms such as N or $^{17}$O with non-zero spin. Hyperfine lines reduce the effective optical depth of the rotational transition by spreading the emission out over a wider bandwidth. Because estimates of the density, temperature, and molecular abundance depend on the optical depth, the hyperfine structure should be taken into account in analyzing spectral line observations Properly treated, the hyperfine structure is quite useful. The observed relative intensities of pairs of hyperfine lines constrain the optical depth independently of the molecular abundance and independently of the spatial coupling of the telescope beam with the cloud structure (beam filling factor). In contrast, optical depth determination from the brightness ratios of spectral lines of isotopologues such as $^{12}$CO and $^{13}$CO requires knowledge of the isotopic abundance ratios, and furthermore the lines may be at sufficiently different frequencies that the observing beam may be differently coupled to the structure of the cloud.  Numerical techniques such as $\\Lambda-$iteration that alternately solve the equations of statistical equilibrium to determine the level populations and the equation of radiative transfer to determine the mean radiation field are able to predict line intensities over a broad range of conditions including varying temperature and density and non-LTE excitation. However, these calculations require knowledge of both the radiative and collisional rates for all transitions.  This presents a problem in the case of the hyperfine lines. For most molecules, the radiative rates, Einstein $A_{ij}$, are known for all the transitions including hyperfine transitions, but the collisional rates are known only as the net rates between rotational levels. These net rates represent the weighted sum of the rates of all the individual hyperfine transitions between the rotational levels. Collisional rates between the individual hyperfine levels themselves have been calculated for only three molecules: HCN (Monteiro \\& Stutzki 1986), NH$_3$ (Chen, Zhang \\& Zhou 1998), and \\diaz \\ (Daniel et al 2005), and even then for only a limited number of hyperfine levels.  Two approximations have been suggested for modeling the emission from molecules with unknown hyperfine collisional rates. The first approximation, \"hyperfine statistical equilibrium\" ({\\it HSE}), assumes that the the hyperfine levels within each rotational level are populated in proportion to their statistical weights \\citep{Keto1990, Keto2004}. The second approximation, the {\\it proportional} approximation, assumes that the collisional rates between the individual hyperfine levels are proportional to the total rate between their rotational levels and the statistical degeneracy of the final hyperfine level of the transition \\citep{GB1991,Daniel2006}.  In this paper, we discuss and evaluate these two approximations, and compute sample \\diaz\\ spectra from each method. Because the collisional rates for the hyperfine transitions of \\diaz\\ are known \\citep{Daniel2005} we can compare \\diaz\\ spectra produced using the approximate collisional rates of the {\\it proportional} method against spectra produced using the ``exact\" rates determined from the numerical quantum mechanical calculations. We also show how the collisional rates for the elastic ($\\Delta J = 0$) rotational transitions may be determined by extrapolation from the inelastic rates. These elastic rates are required in the {\\it proportional} approximation in order to determine the collisional rates for hyperfine transitions within the same rotational state. However, the elastic rates are generally not included in compilations of calculated rate coefficients. ",
        "conclusions": "The modeling of molecular spectra with hyperfine splitting by ALI or Monte Carlo methods has been hampered by the lack of collisional rate coefficients for the hyperfine transitions. Two approximations previously suggested, the approximation of hyperfine statistical equilibrium ({\\it HSE}) and the {\\it proportional} approximation, both provide satisfactory results in tests modeling \\diaz\\ spectra. The {\\it HSE} approximation, based on a modified line profile function, is simpler to implement, faster to compute, and models the non-LTE distribution in the rotational levels but cannot model non-LTE distributions of the hyperfine levels themselves. The {\\it proportional} approximation uses easily computed approximate hyperfine collision rates, and is able to model non-LTE hyperfine emission with an accuracy comparable to calculations using the exact hyperfine collision rates. These results suggest that these two methods could also be useful for other molecules with hyperfine splitting."
    },
    "1004/1004.4518_arXiv.txt": {
        "abstract": "In the early days of April 2010, the blazar 3C 454.3 ($z=0.859$) underwent a strong $\\gamma-$ray outburst, reaching fluxes ($E>100$~MeV) in excess of $10^{-5}$~ph~cm$^{-2}$~s$^{-1}$. The \\emph{Fermi Gamma ray Space Telescope} performed a $200$~ks long pointed observation starting from 5 April 2010 19:38 UTC. This allowed us to try probing the variability of the $\\gamma-$ray emission on time scales of hours or less. We found the variability on a few hours time scale. On sub-hour time scale we found no evidence of significant variability, although the present statistics is not yet conclusive and further observations are needed. ",
        "introduction": "Blazars are one type of active galactic nuclei (AGN) with relativistic jet. The small viewing angle of the jet makes it possible to observe strong effects of the special relativity, such as a boosting of the emitted power and a shortening of the characteristic time scales. Although the variability in blazars has been observed at all the frequencies (see, for example, the reviews by Ulrich et al. 1997 and Wagner 2008), the part at high energies deserves a specific interest, since most of the bolometric power of these sources is emitted at $\\gamma$ rays.  The shortest time scales measured by the EGRET experiment onboard the \\emph{Compton Gamma-Ray Observatory} are $\\sim 4$ hours for PKS~1622-29 (Mattox et al. 1997) and $8$ hours for 3C~279 (Wherle et al. 1998). Early results from the Large Area Telescope (LAT) onboard the \\emph{Fermi Gamma-ray Space Telescope} (hereafter \\emph{Fermi}) indicated similar results: from about half-day for PKS~1454-354 (Abdo et al. 2009a) and PKS~1502+106 (Abdo et al. 2010a) to $5-6$ hours in the cases of 3C~454.3 and PKS~B1510-089 (Tavecchio et al. 2010) and 3C~273 (Abdo et al. 2010b).  These time scales are in agreement with the common paradigm that the characteristic spatial scale of the emitting region is of the order of the gravitational radius $r_{\\rm g}=GM/c^2$ of the central spacetime singularity (Begelman et al. 2008). The zone where most of the dissipation occurs is located at distances greater than $\\Gamma^2r_{\\rm g}$, where $\\Gamma$ is the bulk Lorentz factor (for most blazars $\\Gamma \\sim 10$, see Ghisellini et al. 2010). The size of the emitting region is in turn linked to the variability time $t_{\\rm var}$ by the relationship $R < ct_{\\rm var}\\delta(1+z)$ (where $\\delta$ is the Doppler factor). Obviously, it is expected that $R > r_{\\rm g}$. Moreover, the emitting region must be optically thin and sufficiently far from the central black hole to allow the $\\gamma$ rays to escape, otherwise they convert into electron-positron pairs.  An order-of-magnitude estimation of the above parameters results in an overall agreement with the observed variability on scales of days-hours. However, the recent detection by ground-based Cerenkov telescopes (HESS and MAGIC) of fast variability (minutes time scale) at $\\gamma$ rays in some TeV BL Lac Objects severely threatens the above scenario. In 2006, during an exceptional outburst (average flux 7 Crab with peaks of more than 14 Crab), PKS 2155-304 displayed variability with doubling flux time scale of about $200$~s (Aharonian et al. 2007). In 2005, Mkn 501 changed its flux within a few minutes (Albert et al. 2007).  Given the masses of these two blazars of the order of $10^9M_{\\odot}$, the measured variability is more than one order of magnitude smaller than the minimum allowed. Several solutions have been proposed, from a ``simple'' increase of the Doppler factor ($\\delta > 50$) to invoking an internal structure of the jet (Begelman et al. 2008, Ghisellini \\& Tavecchio 2008, Giannios et al. 2009).  As noted by Begelman et al. (2008), while these explanations could fit reasonably with BL Lac subclass of blazars, the same is not appropriate for flat-spectrum radio quasars (FSRQ) subclass. Indeed, this type of blazars generate $\\gamma$ rays by inverse Compton on a population of seed photons external to the jet (external Compton, EC). Therefore, the pair photosphere can be at fairly large distances ($>10^4r_{\\rm g}$), which in turn means that the blob has a size so large to determine an insufficient energy density to develop strong and fast flares.  It is evident that the search for sub-hour time scales in FSRQ can have strong impact on the current knowledge. If measured, such very short variability would call for a strong revision of the present day models and understanding of relativistic jets. What we need is a strong outburst -- like those occurred in the BL Lac Objects PKS~2155-304 and Mkn~501 -- and a highly performing instrument.  To date, the shortest time scale observed at high energies in a FSRQ is $\\sim 2000$~s, but in the hard X rays. It is the case of NRAO~530 as observed by \\emph{INTEGRAL} ($20-40$~keV energy band) during an exceptional flare in February 2004 (Foschini et al. 2006). Despite of the good performance of \\emph{INTEGRAL}, that episode consists of a single $5\\sigma$ detection, while the source is below the detection limit of the instrument during the remaining of the time. No multiwavelength observations were possible at that time and, specifically, no $\\gamma-$ray satellites were operating. Therefore, this remains an isolated exceptional case and, as for every exceptional claim without a strong observational support, there is always an ``halo'' of doubts.  The launch of \\emph{Fermi}/LAT in June 2008 could give more opportunities to search for sub-hour variability. The Large Area Telescope (LAT, Atwood et al. 2009) could have the necessary sensitivity to probe the shortest time scales at $\\gamma$ rays in quasars. If we consider a flux of $10^{-5}$~ph~cm$^{-2}$~s$^{-1}$ ($E>100$~MeV) and an effective area of $\\sim 5000$~cm$^2$ (see Fig.~14 of Atwood et al. 2009, by taking into account the $0.1-100$~GeV energy range), it results that for normal incidence there are about $15$~photons per $5$~minutes (about a $4\\sigma$ detection, by assuming that on such short time scales the background is almost absent). However, LAT operates almost always in scanning mode. This, in turn, has the great advantage to offer an efficient monitoring of all the sky (a complete coverage every three hours or two orbits), but it hampers the possibility to probe time scales shorter than three hours (the source is not always at the boresight of the instrument and, therefore, the effective area is smaller).  A first tentative to bypass this issue has been recently made by performing a pointed observation during the giant outburst of 3C 454.3 ($z=0.859$) occurred in early April 2010 (Wallace et al. 2010), when the source reached fluxes in excess of $10^{-5}$~ph~cm$^{-2}$~s$^{-1}$ ($E>100$~MeV). Such flux levels were already reached by 3C 454.3 during the first ten days of December 2009 (see Bonnoli et al. 2010 and references therein), but it was not possible at that time to perform a pointed observation and hence to search for sub-hours time scales.  \\begin{figure*} \\centering \\includegraphics[angle=270,scale=0.35]{global3hrs.ps} \\includegraphics[angle=270,scale=0.35]{global1ora.ps} \\caption{Light curves of 3C 454.3 ($E>100$~MeV) with 3 hours \\emph{(left panel)} and 1 hour \\emph{(right panel)}, respectively. Time starts on 31 March 2010 (MJD 55286), so that the number of the days corresponds also to days of April. The horizontal dashed lines in each panel correspond to the weighted average plus/minus one standard deviation.} \\label{fig:curva3p1} \\end{figure*}  Here we report the study of the variability of the $\\gamma-$ray emission of 3C 454.3 as observed by \\emph{Fermi}/LAT during the pointed observation performed from 5 to 8 April 2010. ",
        "conclusions": "The whole data set of $200$~ks observation has been fitted with a broken power-law model in the form of $F(E)\\propto (E/E_{\\rm break})^{\\Gamma_1}$ for $E<E_{\\rm break}$ and $F(E)\\propto (E/E_{\\rm break})^{\\Gamma_2}$ otherwise. This results in these parameters: the photon indices below and above the break energy are $\\Gamma_1=2.22\\pm 0.04$ and $\\Gamma_2=2.80\\pm 0.08$, respectively. The break energy is $0.9\\pm 0.2$~GeV and the flux integrated over the $0.1-100$~GeV energy band is $(12.7\\pm 0.3)\\times 10^{-6}$~ph~cm$^{-2}$~s$^{-1}$. The recorded counts are $3044$ and $TS=10592$ ($>100\\sigma$). We compare these values with the early observation with \\emph{Fermi} referring to the first month of data (Abdo et al. 2009b). In August 2008, the flux was about a factor 4 smaller than in April 2010, with a value of $(3.0\\pm 0.1)\\times 10^{-6}$~ph~cm$^{-2}$~s$^{-1}$, the break energy was greater ($2.4\\pm 0.3$~GeV), and the photon index above the break quite softer, with a value of $\\Gamma_2=3.5\\pm 0.2$. $\\Gamma_1$ is consistent within the measurement errors. It seems that the higher flux needs of a harder spectrum.  The $\\gamma-$ray emission of 3C 454.3 in Fig.~\\ref{fig:curva3p1} (left panel) is variable on a few hours time scale. A fit with a constant-flux line gave $\\chi^2/dof=\\tilde{\\chi}^2\\sim 2.7$ (with 18 degrees of freedom, this means a chance probability of 0.99). However, we note that the last point, which is mainly driving the global variability of the light curve, is not a statistical fluctuation, but the beginning of a real flux increase. This is confirmed by the fact that on 8 April there was the peak of the $\\gamma-$ray emission of the April 2010 outburst, with a day average of $\\sim 16\\times 10^{-6}$~ph~cm$^{-2}$~s$^{-1}$. On the other hand, no strong variability of the photon index has been measured, although some trend is visible between $\\sim 6.5$ and $\\sim 7$ April: $\\Gamma$ changed from $2.58\\pm 0.13$ on $\\sim 6.63$~April to $2.21\\pm 0.08$ ($\\sim 3\\sigma$ significance) on $\\sim 6.88$~April. Similar spectral changes have been observed in $\\gamma-$ray blazars by EGRET (Nandikotkur et al. 2007) and LAT in the case of 3C 273 (Abdo et al. 2010b).  \\begin{figure} \\centering \\includegraphics[angle=270,scale=0.35]{global30min.ps} \\caption{\\emph{(top panel)} Flux light curve of 3C 454.3 ($E>100$~MeV) with 30 minutes time bin; \\emph{(bottom panel)} photon index. Time starts on 31 March 2010 (MJD 55286), so that the number of the days corresponds also to days of April. The horizontal dashed lines in each panel correspond to the weighted average plus/minus one standard deviation.} \\label{fig:curva30min} \\end{figure}  The light curve with 1 hour time bin (Fig.~\\ref{fig:curva3p1}, right panel) has larger error bars and hence the slight variability seen in the 3-hours-bin curve is now smoothed. The fit with a constant flux line gives $\\tilde{\\chi}^2=1.6$, but the visible trends are consistent with those observed in Fig.~\\ref{fig:curva3p1} (left panel). The most significant trend is the drop in flux observed between $\\sim 6.7-6.8$~April.  To evaluate the time scales of these trends, we calculated the time of exponential rise or decay as defined by:  \\begin{equation} F(t) = F(t_0)\\exp[-(t-t_0)/\\tau] \\end{equation}  where $F(t)$ and $F(t_0)$ are the fluxes at the time $t$ and $t_0$, respectively, and $\\tau$ is the characteristic time scale.  In this case, the drop on $\\sim 6.7-6.8$~April visible in Fig.~\\ref{fig:curva3p1} (right panel), has $\\tau = 4.8\\pm 4.6$ hours (significance of the flux variation $3.1\\sigma$). This value is confirmed in the 30-minutes time bin light curve shown in Fig.~\\ref{fig:curva30min} ($\\tau = 4.7\\pm 4.4$~hrs, $3.1\\sigma$). Moreover, the rise at the end of the curve between $\\sim 8.0-8.1$ has $\\tau = 2.7\\pm 1.0$~hrs ($3.3\\sigma$).  It is possible to estimate the minimum Doppler factor $\\delta$ (e.g. Dondi \\& Ghisellini 1995, Mattox et al. 1997)\\footnote{In the present work we used the most recent value for the Hubble-Lema\\^{i}tre constant $H_0=73$~km~s$^{-1}$~Mpc$^{-1}$ (Freedman \\& Madore 2010).}, although during this pointed observation it was not possible to perform multiwavelength observations, because the blazar was too close to the Sun. Therefore, we adopted for the X-ray flux and spectrum, the values measured during the December 2009 outburst, when 3C~454.3 reached similar $\\gamma-$ray fluxes: $\\alpha=0.4$ and $F_{\\rm 1~keV}\\sim 30$~$\\mu$Jy (Bonnoli et al. 2010). We obtain $\\delta \\geq 14$, for 1~GeV $\\gamma$ rays and $\\tau = 2.7$~hours, a high value for a lower limit, but not unlikely (cf Ghisellini et al. 2010).  We noted a few cases of sudden drop in the flux with $\\sim 30$~min time scale on $\\sim 6.1-6.2$ and $\\sim 7.3$~April. A closer inspection of the data revealed that these bins had small source on time ($\\sim 25$\\% of the whole bin). Therefore, it is likely that it is a fake drop, caused by the not good reconstruction of the flux in presence of too few events, although sufficient for a high $TS$."
    },
    "1004/1004.1421_arXiv.txt": {
        "abstract": "The discovery of multi-planet extrasolar systems has kindled interest in using their orbital evolution as a probe of planet formation.  Accurate descriptions of planetary orbits identify systems which could hide additional planets or be in a special dynamical state, and inform targeted follow-up observations. We combine published radial velocity data with Markov Chain Monte Carlo analyses in order to obtain an {\\it ensemble} of masses, semimajor axes, eccentricities and orbital angles for each of 5 dynamically active multi-planet systems: HD 11964, HD 38529, HD 108874, HD 168443, and HD 190360. We dynamically evolve these systems using 52,000 long-term N-body integrations that sample the full range of possible line-of-sight and relative inclinations, and we report on the system stability, secular evolution and the extent of the resonant interactions. We find that planetary orbits in hierarchical systems exhibit complex dynamics and can become highly eccentric and maybe significantly inclined.  Additionally we incorporate the effects of general relativity in the long-term simulations and demonstrate that can qualitatively affect the dynamics of some systems with high relative inclinations.  The simulations quantify the likelihood of different dynamical regimes for each system and highlight the dangers of restricting simulation phase space to a single set of initial conditions or coplanar orbits. ",
        "introduction": "Currently, over 30 multi-planet exosystems are each known to include 2-5 known planets. The orbital architectures and formation scenarios for these systems have become subjects of numerous investigations. The recent discovery of a fifth planet around 55 Cnc \\citep{fisetal2008} has prompted a flurry of follow-up studies \\citep{gayetal2008,rayetal2008a,jietal2009} that aim to better describe the dynamics and evolution of that system. Despite sparse data, the directly-imaged triple system HR 8799 \\citep{maretal2008} has been subject to intense scrutiny \\citep{fabmur2008,clomal2009,fuketal2009,gozmig2009,lafetal2009,reietal2009,sudetal2009}. These explorations demonstrate that major open questions regarding multi-planet systems remain, including: What is the timescale for instability in these systems?  Could they admit additional, currently undetectable planets?  How tightly ``packed'' are such systems (e.g. \\citealt*{rayetal2009})?  These questions can be addressed by investigating the orbital evolution of individual planetary systems.  How do the orbital eccentricities and semimajor axes of planets change with time?  How does additional physics (e.g. tidal forces and general relativity, henceforth referred to as GR) affect the orbits of short-period planets and indirectly other planets in the system?  Theoretical investigations addressing these questions should account for the uncertainties in orbital parameters obtained from observations.  In order to study dynamics, we need to establish initial conditions.  Unfortunately, the accuracy of such initial conditions are limited due to measurement uncertainties and degeneracies inherent to the radial velocity (RV) exoplanet discovery technique. Occasionally, the best-fit RV data yields parameters which indicate that the timescale for such a system to undergo instability is much less than the age of the system (e.g., in less than $10^5$ yr, for HD 82943; \\citealt*{mayetal2004,fabmur2008}).  These short timescales (relative to system lifetimes) suggest that the best-fit orbital model is unlikely and motivate investigations to find the plausible and stable solutions. Further, both $i_{rel}$ and $i_{LOS}$  for planets in most multi-planet systems have not yet been measured. Noteworthy exceptions include a recent estimate of the relative inclination ($i_{rel}$) of GJ 876 planets \\citep{beasei2009} and recent estimates for the line-of-sight inclinations ($i_{LOS}$) of the planets in HR 8799 \\citep{lafetal2009}.  In an effort to remedy some of these issues and better describe system properties, investigators have been developing techniques to model RV data and to describe their orbital evolution.  \\cite{gozetal2005} and \\cite{gozetal2008} utilize stability constraints and optimization techniques partly derived from the MEGNO (Mean Exponential Growth Factor for Nearby Orbits) method \\citep{cinsim2000}.  The advantages to MEGNO include simultaneous multi-planet fitting and efficient identification of quasi-periodic or irregular (chaotic) motion.  A disadvantage is that the arbitrary choice of the ``penalty'' parameter determining this timescale prevents a rigorous interpretation of the results. Some authors \\citep{ford2005,ford2006,gregory2007a,gregory2007b} use Markov Chain Monte Carlo (MCMC) techniques to generate a sample of initial conditions from the posterior probability distribution. The strength of these techniques relies on rigorous Bayesian calculation and interpretations, and a weakness is that stability must be individually tested in each of a large number of models.  For datasets that result in many unstable models, this procedure may result in a time-consuming and perhaps inefficient computational effort.  Here, we follow the methodology of \\cite{verfor2009}, which relies on MCMC-based simulations to derive ensembles of initial conditions.  We assume that the motion is described by a sum of Keplerian (i.e. non-interacting) orbits.  Each ensemble of initial conditions consists of a list of masses, semimajor axes, eccentricities, arguments of pericenter, inclinations and nodes, assuming an initial mean anomaly for each planet.  We then assign line-of-sight and relative inclinations to planets with these ensembles of orbital elements in order to sample from the entire phase space of possible initial conditions.  When planetary orbits are integrated forward in time, the system may exhibit a variety of evolutionary paths. Unrestricted planetary inclinations admit a wide variety of phase space regimes, often including those where apsidal and resonant angles circulate, librate and exhibit chaotic motions \\citep{micetal2006b}. In this work, because we consider systems containing two massive exoplanets with a wide range of eccentricities and inclinations, we must appeal to numerical integrations.  Here, we investigate the dynamical evolution of five two-planet systems which contain sufficiently numerous and accurate radial velocity orbital data suitable for a wide-ranging study. In \\S \\ref{results}, we show that the stability and dynamical properties of these systems can be significantly influenced by the initial values of $i_{rel}$ and $i_{LOS}$. Four of these systems are ``hierarchical'', meaning they contain large ratios of orbital distances, which often include a close-in planet whose evolution could be affected by the general relativistic precession of the pericenter. We aim to characterize the eccentricity variation, stability, secular effects and resonant signatures of planets in these systems as a function of $i_{rel}$ and $i_{LOS}$, while taking into account the uncertainties of the measured orbital parameters for each planet. We combine investigations of the hierarchical HD 12661 with the system studied here to draw inferences for effective methods of future dynamical investigations of hierarchical systems, and to determine their general evolutionary trends (such as the propensity, or lack thereof, for planets to periodically attain circular orbits; \\citealt*{bargre2006a,bargre2006b}).  We describe our methodology in \\S \\ref{methods}, and present the results of the N-body simulations in \\S \\ref{results}. Tables \\ref{Tabe11964}-\\ref{Tablib190360} present summary statistics, two or three tables per system, arranged by row according to the binned relative inclination between both planets; the accompanying figures help the reader visualize the data. We discuss the results in \\S \\ref{discussion} and conclude in \\S \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}   Due to limitations of real astronomical observations, often a significant range of planetary masses is consistent with observations. Therefore, it is necessary to investigate the dynamics of ensembles of planetary orbits and masses to accurately model the orbital evolution of exoplanets. Hierarchical multi-planet systems demonstrate a wide variety of dynamical behaviors depending on the $i_{LOS}$ and $i_{{\\rm rel}}$ values, which are only weakly constrained by observations.  Inclusion of GR in simulations of multi-planet systems with a Hot Jupiter may crucially affect the long-term stability, extent of eccentricity variation, and apsidal configuration directly. The eccentricity and inclination evolution of stable highly inclined systems are often dominated by Kozai-like oscillations, but can be limited by precession due to other planets or GR."
    },
    "1004/1004.3879_arXiv.txt": {
        "abstract": "\\noindent Recent observations of sunspot light-bridges have shed new light on the fact that they are often associated with significant chromospheric activity. In particular chromospheric jets \\citep{shimizu2009} persisting over a period of days have been identifies, sometimes associated with large downflows at the photospheric level \\citep{louis2009}. One possible explanation for this activity is reconnection low in the atmosphere. Light-bridges have also been associated with a constant brightness enhancement in the 1600 \\AA\\ passband of TRACE, and the heating of 1 MK loops. Using data from EIS, SOT and STEREO EUVI we investigate the response of the transition region and lower corona to the presence of a light-bridge and specific periods of chromospheric activity. ",
        "introduction": "Sunspot light-bridges are bright lanes of material that divide the umbra. Their appearance signifies the reestablishment of the granulation within the spot, and often indicates the beginning of fragmentation of the spot itself \\citep{vazquez1973}. Observations from the Swedish Solar Telescope \\citep{scharmer2003} by \\cite{berger2003} showed for the first time that dark central lanes are common features of strong light-bridges. The magnetic field within light-bridges has been observed to be both weaker and more inclined than in the surrounding umbra \\citep[e.g.][]{leka1997} and their increased brightness relative to the surrounding umbra is a clear indication that the plasma temperature in this region is higher. It has also been noted that light-bridges often show enhanced chromospheric activity, with H$\\alpha$ surges and chomospheric jets reported in a number of cases \\citep{roy1973,asai2001,bharti2007,shimizu2009}. \\cite{berger2003} also found a constant brightness enhancement above a light-bridge in TRACE 1600 \\AA\\ observations, while \\cite{katsukawa2007} found that light-bridge  formation was spatially and temporally coincident with the heating of $\\approx$ 1 MK loops as observed by TRACE.  Light-bridges thus seem important for releasing magnetic energy stored in the spot as well as in its decay. In this work we investigate the extent to which the presence of a sunspot light-bridge affects the overlying transition region and corona. ",
        "conclusions": "We find evidence for increased intensity above the light-bridge in STEREO EUVI 171\\AA\\ data, in EIS He II ( 256 \\AA\\ ) and Fe XII (195 \\AA\\ ) rasters, and also in the EIS 266$^{\\prime\\prime}$ Fe XV 284 \\AA\\ slot images, but not in Si VII 275 \\AA\\ raster images.  There are some puzzling inconsistencies, but nevertheless indications that enhanced heating above light-bridges extends higher than previously thought into the corona. Enhanced upflows are seen with EIS in the regions above the light-bridge.These upflows are strongest in the He II 256 \\AA\\ line, and are also located in the vicinity of the outflows that have been identified as potential contributors to the slow solar wind \\citep[e.g.][and references therein]{baker2009}. A more detailed presentation of this work will appear in a forthcoming paper."
    },
    "1004/1004.3562_arXiv.txt": {
        "abstract": "The gravitational waves (GWs) emitted by inspiraling binary black holes, expected to be detected by the {\\it Laser Interferometer Space Antenna (LISA)}, could be used to determine the luminosity distance to these sources with the unprecedented precision of $\\lsim 1$\\%.  We study cosmological parameter constraints from such standard sirens, in the presence of gravitational lensing by large--scale structure. Lensing introduces magnification with a probability distribution function (PDF) whose shape is highly skewed and depends on cosmological parameters.  We use Monte-Carlo simulations to generate mock samples of standard sirens, including a small intrinsic scatter, as well as the additional, larger scatter from lensing, in their inferred distances. We derive constraints on cosmological parameters, by simultaneously fitting the mean and the distribution of the residuals on the distance {\\it vs} redshift ($d_L-z$) Hubble diagram. We find that for standard sirens at redshift $z\\approx1$, the sensitivity to a single cosmological parameter, such as the matter density $\\Omega_{m}$, or the dark energy equation of state $w$, is $\\sim 50\\%-80\\%$ tighter when the skewed lensing PDF is used, compared to the sensitivity derived from a Gaussian PDF with the same variance. When these two parameters are constrained simultaneously, the skewness yields a further enhanced improvement (by $\\sim 120\\%$), owing to the correlation between the parameters.  The sensitivity to the amplitude of the matter power spectrum, $\\sigma_8$ from the cosmological dependence of the PDF alone, however, is $\\sim 20\\%$ worse than that from the Gaussian PDF.  The improvements for $\\Omega_m$ and $w$ arise purely from the non-Gaussian shape of the lensing PDF; the dependence of the PDF on these parameters does not improve constraints relative to those available from the mean $d_L-z$ relation.  At higher redshifts, the PDF resembles a Gaussian more closely, and the effects of the skewness become less prominent. These results highlight the importance of obtaining an accurate and reliable PDF of the lensing convergence, in order to realize the full potential of standard sirens as cosmological probes. ",
        "introduction": "\\label{sec:introduction}  The proposed space-based GW detector {\\it LISA}, sensitive to frequencies between $\\sim 10^{-5}$ to $\\sim 0.1$ Hz, will be able to detect massive black hole binary (MBHB) mergers out to redshifts $z\\gsim 5$.  In addition to providing information on black hole physics and general relativity, observations of GWs by {\\it LISA} could be used as a probe of cosmology.  As pointed out in a pioneering paper by Schutz (1986), GW observations of a binary system could yield an accurate estimate of the luminosity distance to the source, independent of assumptions about the masses and orbital parameters of the binary members.  Recent analyses in the context of {\\it LISA} show that for many binaries, the luminosity distance could be measured to percent--level precision (see Arun et al. 2009b for a review and references).  Although a typical {\\it LISA} source will be relatively poorly localized on the sky (to $\\sim 0.1$deg$^2$), when the spatial information along the line of sight is taken into account, the number of galaxies in the three-dimensional error volume will be reduced significantly (Holz \\& Hughes 2005; Kocsis et al. 2006).  Combining this with possible tell-tale time-variable signatures (see Haiman et al. 2009 for a review of several possibilities), it may become feasible to identify an electromagnetic (EM) counterpart.  By measuring the redshift of the counterpart, a gravitational version of the Hubble diagram could be constructed.  This Hubble diagram, with a relatively small intrinsic scatter, and spanning a large range in redshift, can possibly impose tight constraints on cosmological parameters (see, e.g., Arun et al. 2009a for a recent review and references).  Unfortunately, however, gravitational lensing significantly changes this picture. Standard sirens, just as type Ia supernovae (SNe Ia), are (de)magnified by inhomogeneities in the matter distribution in the foreground, which introduces an uncertainty in the measured distance-redshift relation by up to $\\sim 10$\\% for high--redshift ($z\\gsim 2$) events (e.g. Holz \\& Hughes 2005; Kocsis et al. 2006). In the case of SNe Ia, proposed missions, such as the Supernova/Acceleration Probe (SNAP), are expected to find a few thousand useful sources. The random lensing magnification errors then average out, and even if they are unaccounted for, they have a relatively modest ($\\lsim1\\%$) impact on cosmological parameter-estimation, which can be ignored (Holz \\& Linder 2005; Sarkar et al. 2008).  The same strategy is unlikely for standard sirens: although the expected {\\it LISA} MBHB event rate is highly uncertain, most models predict that it is significantly below the SNe Ia rate, with perhaps tens of detections per year (e.g. Menou et al. 2001; Sesana et al. 2007; Lippai et al. 2009; Arun et al. 2009b). Furthermore, the EM counterpart may be identifiable for only a fraction of these events.\\footnote{The {\\it Big Bang Observer (BBO)}, a concept for a space mission to succeed {\\it LISA}, could detect a more than sufficient number of compact stellar binaries; this would even allow a useful measurement of the spatial power spectrum of the lensing convergence, which can provide additional cosmological constraints (Cutler \\& Holz 2009; see also Cooray et al. 2006 for the same idea with Type Ia SNe).  In this paper, we will not consider the possibility of such a large number ($\\gsim 10^3$) of detectable events.}  Motivated by the tremendous potential, in the absence of lensing, of standard sirens for cosmology, there have been proposals to correct for the effects of lensing of individual sources on a case-by-case basis.  These proposals include measuring or constraining the magnification using either photometric and spectroscopic properties of foreground galaxies (e.g. J\\\"onsson et al. 2007), or the combination of arcminute--scale shear and flexion maps (Shapiro et al. 2009; the earlier work of Dalal et al. 2003, which only considered the shear, concluded that only modest, $\\sim 20\\%$, corrections were feasible). Both of these methods could reduce the lensing--induced distance errors, in idealized cases, by a factor of up to $\\sim$two.  Alternatively, several authors have investigated the possibility of using lensing as a signal, rather than as noise, in probing cosmology (Dodelson \\& Vallinotto 2006, hereafter DV06; Linder 2008; see also Wang et al. 2009). In particular, the variance of the lensing magnification probability distribution function (PDF) depends on the amplitude of density fluctuations and on the growth function, and therefore could be used to constrain cosmological parameters such as $\\sigma_8$ and $\\Omega_m$.  More specifically, DV06 showed that $\\sigma_8$ could be constrained to an accuracy of $\\approx 5\\%$ by observations of 2000 SNe Ia.  The overall shape of the lensing convergence distribution contains additional cosmological information, beyond the variance (Wang et al. 2009). For example, Linder (2008) emphasized that the theoretically possible minimum (de)magnification, which occurs along an empty beam, depends on cosmology. Several authors have indeed proposed fitting formulae for the lensing PDF, derived from numerical simulations, which are self-similar, and depend on cosmology only through the variance and the minimum magnification (see below).  In this paper, we study the cosmological parameter constraints from standard sirens, in the presence of lensing magnification.  In general, lensing causes significant degradation of the constraints, which includes increased uncertainties, as well as a possible bias, in the parameters inferred from a given finite source sample.  If the lensing PDF was Gaussian, and did not depend on the cosmological parameters, this degradation could be estimated simply by the increase in the variance of the inferred distance error to individual sources. However, the lensing PDF is highly skewed, and it does depend on the cosmological parameters.  Our focus here is to quantify the extent to which these two features affect (and hopefully, mitigate) the degradation of the constraints expected from {\\it LISA}.  In the context of SNe, these questions have already been addressed in detail by several authors, including the impact of lensing on inferred dark energy parameters (e.g. Holz \\& Linder 2005; Sarkar et al. 2008) and on the normalization of the matter power--spectrum, $\\sigma_8$ (e.g. DV06).  Here we consider, instead, a relatively small sample ($\\sim$tens) of sources, with otherwise very small ($\\lsim 1\\%$) distance errors, as expected from {\\it LISA} standard sirens.  The conclusions are not necessarily the same for such a standard siren sample as for the SNe Ia.  This is because the probability distribution of the inferred distances, which is the convolution of intrinsic and lensing probability distributions, is much more skewed in the case of standard sirens, due to the dominance of the highly skewed lensing distribution. Additionally, depending on the statistic being used, having fewer events can increase the impact of non-Gaussianity on the inferred parameters.  A few works have also studied the cosmological utility of standard sirens, and have included the effect of lensing, noting the significant degradation they cause in the constraints (e.g. Holz \\& Hughes 2005; Dalal et al. 2006; Linder 2008).  Our present study adds to these existing papers in the following ways: (i) in our analysis, we include the dependence of the convergence PDF on cosmological parameters, thus treating lensing as a potential source of signal, rather than as pure noise; (ii) we include $\\sigma_8$ among our parameters, since the lensing PDF is particularly sensitive to this parameter; (iii) we perform multiple Monte Carlo realizations of mock standard-siren samples, incorporating the non-Gaussian shape of the PDF, to obtain accurate estimates of the confidence intervals on the inferred parameters; and (iv) we use more recent fitting formulae for the lensing PDF, which improve the fit to numerical simulations.  The remainder of this paper is organized as follows. In \\S~\\ref{sec:physics}, we briefly summarize the basic background material required for our study, including information on both standard siren distance measurements and on gravitational lensing. \\S~\\ref{sec:mc} outlines the details of our Monte-Carlo simulation procedure. In \\S~\\ref{sec:results}, we present and discuss our main results, and contrast these with the analogous results in the case of SNe Ia. Finally, we summarize our main conclusions in \\S~\\ref{sec:conclusion}. Throughout this paper, we adopt standard ${\\rm \\Lambda CDM}$ as our fiducial cosmological model, with parameter values consistent with the {\\it WMAP} fifth year results (Komatsu et al. 2009), i.e., $\\{\\Omega_m, \\Omega_{\\Lambda},\\Omega_b, h, \\sigma_8, n_s\\}=\\{0.279, 0.721, 0.0462, 0.701, 0.817, 0.96\\}$.  These values are in general agreement with most recent, seventh year results, as well (Komatsu et al. 2010). ",
        "conclusions": "\\label{sec:conclusion}   In this paper, we have shown that the sensitivity of low--redshift ($z\\sim 1$) standard sirens to the parameters $\\Omega_m$ and $w$ are tightened, by a factor of $1.5-1.8$, by the non--Gaussian shape of the lensing magnification PDF, relative to a Gaussian PDF with the same variance.  When these two parameters are constrained simultaneously, the improvement of the constraints, attributable to the skewness alone, is further enhanced, owing to the correlation between the parameters.  We expect similar conclusions to hold for other parameters whose constraints are driven primarily by the changes they induce in the luminosity distance $d_L(z)$.  Interestingly, the constraint on $\\sigma_8$, which comes from the shape of the PDF itself, is, however, degraded by a factor of $\\approx 0.8$.  In our study, we relied on the shape of the lensing PDF described by the fitting formulae in Das \\& Ostriker (2006). This shape is somewhat less skewed than the simulation results, suggesting that the improvement / degradation of the constraints has likely been underestimated.  The improvements for $\\Omega_m$ and $w$ are comparable to those that may be available from correlations with lensing measurements on larger scales, or from directly subtracting the lensing contribution of foreground galaxies. However, these corrections from non-Gaussianities are ``automatically'' available, as long as the lensing PDF can be modeled ab--initio.  We also found that at higher redshift, the effects from non--Gaussianity are less pronounced due to the reduced skewness of the lensing PDF.  Overall, our results highlight the importance of obtaining an accurate and reliable PDF of the point--source lensing magnification, in order to realize the full potential of standard sirens as cosmological probes.  \\vspace{-0.5\\baselineskip}"
    },
    "1004/1004.3048_arXiv.txt": {
        "abstract": "{The solar irradiance is known to change on time scales of minutes to decades, and it is suspected that its substantial fluctuations are partially responsible for climate variations. } {We are developing a solar atmosphere code that allows the  physical modeling of the entire solar spectrum composed of quiet Sun and active regions. This code is a tool for modeling the variability of the solar irradiance and understanding its influence on Earth. } {We exploit further development of the radiative transfer code COSI that now incorporates the calculation of  molecular lines. We validated COSI under the conditions of  local thermodynamic equilibrium (LTE) against the synthetic spectra calculated with the ATLAS code. The synthetic solar spectra were also calculated in non-local thermodynamic equilibrium (NLTE) and compared to the available measured spectra. In doing so we have defined the main problems of the modeling, e.g., the lack of opacity in the UV part of the spectrum and the inconsistency in the calculations of the visible continuum level, and we describe a solution to these problems. } {The improved version of COSI allows us to reach good agreement between the calculated and observed solar spectra as measured by SOLSTICE and SIM onboard the SORCE satellite  and ATLAS 3 mission operated from the Space Shuttle.  We find that NLTE effects are very important for the modeling of the solar spectrum even in the visual part of the spectrum and for its variability over the entire solar spectrum. In addition to the strong effect on the UV part of the spectrum, NLTE effects influence the concentration of the negative ion of hydrogen, which results in a significant change of the visible continuum level and the irradiance variability.} {} ",
        "introduction": "\\label{sec:intro} The solar radiation is the main source of the input of energy to the terrestrial atmosphere, so that it determines Earth's thermal balance and climate.  Although it has been known  since 1978 that the solar irradiance is not constant but instead varies on  scales from several minutes to decades \\citep[cf.][]{froehlich2005,krivovasolanki2008}, the influence of this variability on the climate is not yet fully understood. Nowadays several datasets for the past spectral solar irradiance (SSI) based on the different reconstruction approaches \\citep[e.g.][]{lean2005, krivovaetal2009} and satellite measurements are available.  However, the remaining disagreements between these data  lead to different atmospheric responses when they are used in the climate models  \\citep{shapiroetal2009}. The task of constructing a self-consistent physical model in order to reconstruct the past solar spectral irradiance (SSI) remains of high importance.  Modern reconstructions of the SSI are based on the assumption that the irradiance changes are determined by the evolution of the solar surface magnetic field \\citep[][]{foukallean1988, krivovaetal2003, domingoetal2009}. Areas of the solar disk are associated to several components (e.g. quiet Sun, bright network, plage, and sunspot) according to the measured surface magnetic field and the contrast  of the features. These components are represented by corresponding atmosphere structures \\citep[cf.][]{kurucz1991,fontenlaetal1999}. The SSI is calculated by weighting the irradiance from each model with the corresponding filling factor.    The calculation of the emergent solar radiation, even from the atmosphere with known thermal structure, is a very sophisticated problem because consistent models have to account for the NLTE effects in the solar atmosphere. As the importance of these effects has become clear over the past several decades, several numerical codes have been developed. One of the first NLTE-codes, LINEAR, was published by \\citet{aueretal1972}, who used a complete linearization method developed by \\citet{auermihalas1969,auermihalas1970}.  Later, the MULTI code was published by \\citet{carlsson1986}. This code is based on the linearization technique developed by \\citet{scharmer1981} and  \\citet{scharmercarlsson1985a, scharmercarlsson1985b}. More recently, the RH code,  which is based on the MALI (Multi-level Approximate Lambda Iteration) formalism of \\citet{rybickihummer1991, rybickihummer1992} has been developed by \\citet{uitenbroek2001}.   Combining the radiative transfer code  by  \\citet{hamannschmutz1987,  schmutzetal1989} and  spectral synthesis code SYNSPEC by \\citet{hubeny1981}, \\citet{haberreiter2008} has developed the 1D spherical symmetrical COde for Solar Irradiance (COSI). The NLTE opacity distribution function (ODF) concept, which was implemented in this code, allows an indirect account for the NLTE effects in several million lines. This makes COSI especially suitable for calculating the overall energy distribution in the solar spectrum. In this paper we introduce a new version of COSI (version 2), describe the main modifications, and present first results.   In Sect.~\\ref{sec:obs} we describe the datasets of the measured spectral irradiance, which were used for comparison with our calculated spectrum. In Sect.~\\ref{sec:molecules} we introduce the integration of molecular lines into COSI and show that it can solve the discrepancies to observations and to other codes. In Sect.~\\ref{sec:NLTEproblems} we present the resulting solar spectrum from NLTE calculations. Due to missing opacity in the UV, the flux appears to be significantly higher than measured \\citep{busaetal2001, shorthauschildt2009}. We solve this problem by introducing  additional opacity to the ODF for selected spectral ranges (Sect.~\\ref{subsec:NLTEUV}),  while the problem of the NLTE visible continuum due to the deviations in the concentration of the hydrogen negative ion is addressed in Sect.~\\ref{subsec:NLTEcont}. In Sect.~\\ref{sec:active} we present the synthetic spectra of the active regions and its implications for the solar variability study. Finally, we summarize the main results in Sect.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} We have presented a further development of the radiative transfer code COSI.  The code accounts for the NLTE effects in several hundred lines, while the NLTE effects in the several million other lines are indirectly included via iterated opacity distribution function. The radiative transfer is solved in spherical symmetry. The main conclusions can be summarized as follows. \\begin{itemize} \\item The inclusion of the molecular lines in COSI allowed us to reach good agreement with the SORCE observations in the main molecular bands (especially in the CN violet system and CH G band). It has also solved the previous  discrepancies  between the LTE calculations with the  COSI code and ATLAS 12  calculations. We showed that their strong temperature sensitivity allows molecular lines to significantly contribute to the solar irradiance variability. \\item We introduced additional  opacities into the opacity distribution function. It allowed us to solve the well-known problem of overestimating the UV flux in the synthetic spectrum.  We should emphasize, however, that the magnitude and behavior of the additional opacity strongly depend on the applied model. \\item We have shown that the concentration of negative hydrogen is strongly affected by NLTE effects as explained in Sect.~\\ref{subsubsec:Hminus}. It decreases the level of the visible and infrared continuum and leads to a discrepancy with the measured level if the calculations are done with the current models of the quiet Sun atmosphere. \\item We presented  calculations of the total and spectral solar irradiance changes due to the presence of the active regions and showed that NLTE effects can strongly affect both of these values.   \\end{itemize}"
    },
    "1004/1004.5491_arXiv.txt": {
        "abstract": "We investigate baryogenesis in the $\\nu$MSM, which is the Minimal Standard Model (MSM) extended by three right-handed neutrinos with Majorana masses smaller than the weak scale.  In this model the baryon asymmetry of the universe (BAU) is generated via flavour oscillation between right-handed neutrinos.  We consider the case when BAU is solely originated from the CP violation in the mixing matrix of active neutrinos.  We perform analytical and numerical estimations of the yield of BAU, and show how BAU depends on mixing angles and CP violating phases.  It is found that the asymmetry in the inverted hierarchy for neutrino masses receives a suppression factor of about 4\\% comparing with the normal hierarchy case.  It is, however, pointed out that, when $\\theta_{13}=0$ and $\\theta_{23} = \\pi/4$, baryogenesis in the normal hierarchy becomes ineffective, and hence the inverted hierarchy case becomes significant to account for the present BAU. ",
        "introduction": "The origin of the baryon asymmetry of the universe (BAU) is one of the most mysterious problems in particle physics and cosmology, since the Minimal Standard Model (MSM) and the Big Bang cosmology cannot answer it.  So far various mechanisms for generating BAU have been proposed~\\cite{Riotto:1999yt}.  One promising possibility is the so-called leptogenesis scenario~\\cite{Fukugita:1986hr} (see also Ref.~\\cite{Buchmuller:2005eh}), where the origins of neutrino masses as well as BAU can be addressed at the same time by introducing right-handed neutrinos with superheavy Majorana masses.  The observed tiny masses of neutrinos can be naturally understood by the seesaw mechanism~\\cite{Seesaw}.  Further, the lepton asymmetry generated by decays of right-handed neutrinos can be a source of BAU.  In the simplest thermal leptogenesis, the required Majorana masses is heavier than about $10^9$ GeV~\\cite{Giudice:2003jh}.  It should be, however, noted that the connection between the origins of neutrino masses and BAU can be obtained even when Majorana masses are below the weak scale.  One interesting possibility is the so-called $\\nu$MSM~\\cite{Asaka:2005pn,Asaka:2005an}, which is the MSM extended by three right-handed neutrinos with masses smaller than the weak scale.  In this model the problems of neutrino masses, BAU and also dark matter can be solved simultaneously.  One attractive advantage of the $\\nu$MSM lies in the fact that the direct detection of right-handed neutrinos becomes possible~\\cite{Gorbunov:2007ak}.  In the $\\nu$MSM BAU can be generated by invoking the mechanism via flavour oscillation of right-handed neutrinos~\\cite{Akhmedov:1998qx}. (See also Ref.~\\cite{Asaka:2005pn,Shaposhnikov:2008pf}.)  In this mechanism the lepton asymmetry is not generated for temperatures of interest because of the smallness of Majorana masses, which is very different from the leptogenesis scenario.  The lepton asymmetry is separated into left-handed and right-handed leptonic sectors due to the CP violations in the production and oscillation of right-handed neutrinos.  Then, the asymmetry stored in the left-handed sector is partially transferred into the baryon asymmetry through the rapid sphaleron transitions~\\cite{Kuzmin:1985mm}.  One of right-handed neutrinos in the $\\nu$MSM, which is a candidate of dark matter, plays no essential role to generate BAU as well as neutrino masses observed in the oscillation experiments, since its Yukawa coupling constants should be very suppressed.  The rest two are responsible to BAU via their flavour oscillation, but also induce the masses of active neutrinos through the seesaw mechanism. Therefore, physics of these two right-handed neutrinos connects BAU with the neutrino parameters of active neutrinos, \\ie, mass hierarchy, mixing angles, and CP violating phases.  In this letter we would like to extend the analysis in Ref.~\\cite{Asaka:2005pn}.  Under the considering situation there are three CP violating phases in the leptonic sector which can be a source of the asymmetry.  Especially, we concentrate here on the case when BAU is originated only from the CP violation in the mixing matrix $U$ of active neutrinos, namely, the Dirac phase $\\delta$ and Majorana phase $\\eta$ in $U$.  We then present the analytical expression of BAU shows explicitly how BAU depends on these CP phases and the mixing angles of active neutrinos. Moreover, we also perform the numerical estimation of BAU and justify the validity of the analytical expression. ",
        "conclusions": "\\label{sec:Conc} We have studied baryogenesis in the $\\nu$MSM via flavour oscillation between right-handed neutrinos $N_2$ and $N_3$.  In particular, the case when BAU comes solely from the CP violating phases in the mixing matrix of active neutrinos has been investigated.  We have presented the analytical expressions of BAU for both the NH and IH cases of active neutrino masses, and have demonstrated how the present value of BAU depends on the Dirac and Majorana phases as well as the neutrino mixing angles.  We have shown that BAU depends on the neutrino mass hierarchy and $Y_B$ in the IH case receives the suppression factor of $S_{m_\\nu} \\simeq 4\\%$. It has been found that $Y_B$ is very sensitive to the mixing angles $\\theta_{23}$ and $\\theta_{13}$. When $\\theta_{23} = \\pi/4$, the leading contribution to $Y_B$ is proportional to $\\sin (\\delta + \\eta)$ for the NH case while to $\\sin \\eta$ for the IH case. Moreover, when $\\theta_{23}$ is almost maximal and $\\theta_{13}$ is very small, the CP asymmetry parameter in $Y_B$ vanishes and no baryon asymmetry is generated (at least the leading ${\\cal O}(F^6)$ contribution) in the NH case. In this case, the IH case is required to explain the observed BAU."
    },
    "1004/1004.2600_arXiv.txt": {
        "abstract": "We use recent observations of the He~I~$\\lambda10830 \\rm{\\AA}$ absorption line and 3D hydrodynamical numerical simulations of the winds collision, to strengthen the case for an orientation of the semimajor axis of the massive binary system $\\eta$ Carinae where the secondary star is toward us at periastron passage. Those observations show that the fast blue absorption component exists for only several weeks prior to the periastron passage. We show that the transient nature of the fast blue absorption component supports a geometry where the fast secondary wind, both pre and post-shock material, passes in front of the primary star near periastron passage. ",
        "introduction": "\\label{sec:intro}  $\\eta$ Car is a very massive stellar binary system, with an orbital period of $5.54 \\yr$ (Damineli 1996), as observed in all wavelengths (e.g., radio, Duncan \\& White 2003; IR, Whitelock et al. 2004; visible, van Genderen et al. 2006, Fernandez-Lajus et al. 2009; UV, Smith et al. 2004; emission and absorption lines, Nielsen et al. 2009, Damineli et al. 2008a, b; X-ray, Corcoran 2005, 2010, Hamaguchi et al. 2007). The high eccentricity of $e \\simeq 0.9$ results in rapid changes in emission and absorption lines, as well as in the continuum, near each periastron passage. The several weeks of the rapid changes occurring every orbital period is termed the spectroscopic event. These lines might originate in different places in the binary system: the primary star, the secondary star, their respective winds, and the colliding winds structure which is termed the conical shell.  As $\\eta$ Car is the best studied binary luminous blue variable (LBV), it holds the key to our understanding of other LBVs. It is particulary important to understand the behavior near periastron passage, where the strongest binary interaction takes place, and for that is crucial to know the orientation of the binary system. Namely, the direction of the primary more massive LBV star relative to its less massive but hotter companion at periastron passage. While it is agreed that the inclination of the binary system is $i\\simeq 41^\\circ$ (Davidson et al. 2001; Smith 2006), there is no agreement on the direction of the semimajor axis, termed periastron longitude. The orientation is measured by the angle $\\omega$: $\\omega=0^\\circ$ for the case when the secondary moves toward us before periastron passage and the semimajor axis is perpendicular to the line of sight, $\\omega=90^\\circ$ when the secondary is toward us at periastron passage, and $\\omega=270^\\circ$ when the primary is toward us at periastron passage. Several properties of the binary system have been used to deduce the orbital orientation, with contradicting results (for details see Kashi \\& Soker, 2008b, 2009c).  One of the properties that led to a contradicting conclusion on the orientation is the behavior of the blue absorption wing of the He~I~$\\lambda10830 \\rm{\\AA}$ line. In Kashi \\& Soker (2009b) we constructed a toy model where the material responsible for the blue absorption wing was assumed to reside in the colliding winds region -- the conical shell -- close to the binary system. This model is able to account for the transient appearance of the blue absorbing wing and to the finding that the maximum absorbing velocity is reached several days before periastron passage, only if the secondary is toward us near periastron passage, i.e., $\\omega \\simeq 90^\\circ$. In a recent paper Groh et al. (2010; hereafter G2010) reached an opposite conclusion. Comparing their observations with a model based on 3D numerical simulations of the colliding winds structure, G2010 suggested that the orientation is that of $\\omega=243^\\circ$. Namely, the primary is closer to us just before periastron passage. In this paper we critically reexamine both models. ",
        "conclusions": "\\label{sec:summary}  The He~I~$\\lambda10830 \\rm{\\AA}$ line of $\\eta$ Car has been observed across the 2009 periastron passage by G2010. We identify three main absorbing components in the He~I~$\\lambda 10830 \\rm{\\AA}$ line profile (Fig. \\ref{fig:profile}). Two components can be identified with the dense primary stellar wind; the P-component ($500$--$650 \\km \\s^{-1}$) and the M-component ($-v = 620$--$720 \\km \\s^{-1}$). The third, more interesting F-component shows a flat part in the velocity range $-v \\simeq 1150$--$1400 \\km \\s^{-1}$, with a tail to $-v = 1900 \\km \\s^{-1}$. The tail results from faster moving gas, and not from line broadening. There is no absorption in this line at $-v > 1900 \\km \\s^{-1}$.  G2010 present their interpretations of their observations, concluding that the orientation of the binary system is such that $\\omega=243^\\circ$ (the primary closer to the observer just before periastron). In section \\ref{sec:problem} we point out some problems in the model of G2010, in particular that they cannot account for the transient nature of the line. The model of G2010 fails to reproduce, even qualitatively, the observed absorption profiles of the He~I~$\\lambda10830 \\rm{\\AA}$ line, and cannot account for the absence of the F-component after periastron passage. In analyzing our (Kashi \\& Soker 2008b, 2009b) preferred orientation of $\\omega=90^\\circ$, G2010 use a point-like central continuum source for calculating the column density of the absorbing material. Based on that, they claim that the $\\omega=90^\\circ$ orientation is not compatible with their observations. In section \\ref{sec:omega90} we find that the usage of a central point source does not do justice to the $\\omega=90^\\circ$ case. This is the main reason why G2010 fail to reproduce results of a model where the absorption occurs a few $\\AU$ from the stars, like in our model.  We improve our pervious model (Kashi \\& Soker 2009b) which assumed that the conical shell is the main absorber of the high velocity component of the He~I~$\\lambda10830 \\rm{\\AA}$ line (section \\ref{sec:omega90}). Our model and results can be summarized as follows: \\begin{enumerate} \\item We assume that the central star is a Gaussian weighted extended continuum source as given by equation (\\ref{eq:Gaussian}). This assumption stands on a solid ground (G2010; Weigelt et al. 2007).  \\item We assume that the increase in the $1.083 \\mum$ continuum near periastron passage comes from the conical shell (the collision region of the two winds). The main source is the shocked primary wind. This is based on analysis we performed in previous papers.  \\item We assume that the main absorbing gas of the He~I~$\\lambda10830 \\rm{\\AA}$ line is both pre and post-shock material in the fast secondary wind. This assumption stands on a solid ground. Based on recent numerical simulations (Akashi \\& Soker 2010), we showed in section \\ref{sec:omega90} that the column density of the fast moving gas is as high as found by G2010. Some segments of the wind are at a temperature of $\\sim 10^4 \\K$, and there are enough recombining He atoms. This will be calculated in a future paper based on high resolution numerical simulations.  \\item We concentrate on the absorption of the F-component. We assume that in the flat velocity range of the F-component (see Fig. \\ref{fig:profile}), the absorbing gas in the conical shell absorbs most of the conical shell emission. We therefore consider two limiting values for this parameter $f_s = 0.5$ and $f_s = 1$.  \\item We assume that the optical depth of the conical shell in the flat velocity range of the F-component is very high, such that it absorbs all the radiation of the central source it hides from our line of sight (in the flat part).  \\item We assume that the colliding winds shell has an hyperbolic shape. We consider its tilt due to the orbital motion. In calculating the shape and tilt the primary wind acceleration zone is considered.  \\item We assume, based on our previous papers (Kashi \\& Soker 2008b,2009c), that the secondary is toward us near periastron ($\\omega =90^\\circ$). We take an inclination angle of $i=41^\\circ$, and the other commonly used binary parameters (semimajor axis, eccentricity, etc.).  \\item We note that after periastron passage some of the assumptions break down, because accretion is likely to occur (Kashi \\& Soker 2009a; Akashi \\& Soker 2010).  \\item We calculate the part of the central source covered by the conical shell, for $\\omega=90^\\circ$.  \\item We calculate the covered part of the $1.083 \\mum$ central continuum emission source (equation \\ref{eq:Fobs}). By that we show that for $\\omega=90^\\circ$ our model explains the observed increase in absorption strength of the F-component close to periastron passage (Fig. \\ref{fig:Fobs}), and explains its transient nature.  \\item We therefore conclude that the orientation is indeed that the secondary is toward us near periastron ($\\omega =90^\\circ$).  \\item In addition, using hydrodynamic simulations (Fig. \\ref{fig:simulation} and \\ref{fig:N_H_vs_v}) we find that the He~I in the conical shell and in the pre-shocked secondary wind has a substantial column density in the velocity range $-v = 1150$--$1400 \\km \\s^{-1}$. We assume that the fraction (out of total He atoms and ions) of the He~I in the $2~^3S$ level is high enough to absorb most radiation in the $-v < 1400 \\km \\s^{-1}$ (with a tail up to $-v \\simeq 1900 \\km \\s^{-1}$). \\end{enumerate}  The finding that the secondary is closer to us near periastron requires that some properties near periastron passage be explained by the accretion of the primary wind onto the secondary star (Kashi \\& Soker 2008a, 2009d). For accretion to occur onto the secondary star, that has a strong wind of his own, the binary must be close (Akashi \\& Soker 2010) and to interact strongly with the primary. The accretion that occurs near periastron passage is crucial to the understanding not only the present behavior of $\\eta$ Car, but also its behavior during the major eruptions it undergone in the 19th century, the Great Eruption (GE) and Lesser Eruption.  The debate on the absorbing source of the blue wing has implications far beyond the specific question on the orientation of the major axis of the binary system. The essence of the debate is the nature of the binary interaction process, which has implications on the nature of LBV major eruptions, and mass loss by very massive stars (e.g., Soker 2001, 2005, 2007; Kashi \\& Soker 2010; Smith et al. 2010; Smith 2010a,b).  The GE is the best studied example for a major LBV eruption, and serves as a test case for theories and models. A high rate accretion during the GE could have supplied the extra luminosity for 20 years, and the accreting secondary star could have launched two jets that shaped the bipolar structure -- the Homunculus (Soker 2001). Most likely other LBV major eruptions are also related to binary interaction (Kashi \\& Soker 2010), as was argued for the 17th century eruption of P~Cygni (Kashi 2010), rather than a single star phenomena as suggested by, e.g., Smith (2007). Smith \\& Owocki (2006) suggested that LBVs lose most of their envelope mass during major eruptions. Therefore, the presence of a strongly interacting companion can play a major role in the evolution of very massive stars. As suggested by Kashi et al. (2010), major mass transfer events in LBVs are related to optical transient objects and have a common powering mechanism -- accretion onto a companion star. Understanding the binary interaction in $\\eta$ Car will shed light on other objects where binary interaction is thought to shape circumstellar nebulae, like planetary nebulae, symbiotic systems, and related objects such as the Red Rectangle."
    },
    "1004/1004.0119_arXiv.txt": {
        "abstract": "{The Galactic microquasar SS\\,433 possesses a circumbinary disk most clearly seen in the brilliant Balmer H$\\alpha$ emission line. The orbital speed of the glowing material is an important determinant of the mass of the binary system. The circumbinary disk may be fed through the L2 point and in turn may feed a very extended radio feature known as the ruff.} {We present (i) an analysis of spectroscopic optical data from H$\\alpha$ and He I spectral lines which reveal the circumbinary disk (ii) comparisons of the rather different signals, to better understand the disk and improve estimates of the rotational speed of the inner rim (iii) a simple model that naturally explains some apparently bizarre spectral variations with orbital phase.} {Published spectra, taken almost nightly over two orbital periods of the binary system, show H$\\alpha$  and He I lines. These were analysed as superpositions of Gaussian components and a simple model in terms of a circumbinary disk was constructed. The possible contributions to the signal of an outflow through the L2 point were considered. } { The data can be understood in terms of a hot spot, generated in proximity to the compact object and rotating round the inner circumbinary disk with a period of 13 days. The glowing material fades with time, quite slowly for the H$\\alpha$ source but more rapidly for the He I spectral lines. The orbital speed of the inner rim is approximately 250 km s$^{-1}$. It may be that absorption lines attributed to the atmosphere of the companion are in fact formed in this circumbinary material.} { The mass of the binary system must exceed 40 $M_\\odot$ and the compact object must be a rather massive stellar black hole. The corollary is that the orbital speed of the companion must exceed 130 km s$^{-1}$.} ",
        "introduction": "The Galactic microquasar  SS 433 is famous for its continual ejection of plasma in two opposite jets at approximately one quarter the speed of light. Precession of the jet axis gives rise to the famous moving spectral lines but the so-called stationary lines are more intense. The system is a binary with a period of 13.08 days (Crampton, Cowley and Hutchings 1980) and eclipses at both oppositions (e.g. Goranskii et al 1998).  He II 4686 \\AA\\  emission has been observed (Crampton and Hutchings 1981, Fabrika and Bychkova 1990) , attributed to the base of the jets (Fabrika 1997). C II lines orbiting with the compact object have been detected (Gies et al 2002, K. M. Blundell private communication). The orbital speed of the compact object about the binary centre of mass is now well established as 176 km s$^{-1}$ and the mass function as 7.7 $M_\\odot$. There is no consistency among the many reports of Doppler speeds for the companion. Relatively recent observations of absorption lines, attributed to the atmosphere of the companion, have yielded an orbital velocity for the companion about the binary centre of mass of 132 km s$^{-1}$ (Cherepashchuk et al 2005), which implies a system mass of 42 $M_\\odot$, and 58 km s$^{-1}$ (Hillwig \\& Gies 2008, Kubota et al 2010), the latter value implying a system mass of 17 $M_\\odot$. Observations  in H$\\alpha$ interpreted in terms of a circumbinary disk imply a system mass of approximately 40 $M_\\odot$ or greater (Blundell, Bowler \\& Schmidtobreick 2008).  A rotating ring of glowing gas, viewed almost edge on, will produce a spectrum dominated by radiation from regions to which the line of sight is tangential. A suitable spectral line thus appears split; two horns Doppler shifted by the rotational speed of the ring. Just such a split appears in various stationary features of the spectrum of the microquasar SS 433; Filippenko et al (1988) reported the H Paschen series to be split by approximately 290 km s$^{-1}$ and in the blue various unblended Fe II lines split by 250-300 km s$^{-1}$. Balmer H$\\beta$ was split by approximately 480 km s$^{-1}$. These observations covered only a few consecutive days but were interpreted as evidence for a disk structure. The authors evidently thought radiation from the accretion disk most likely, but also considered the possibility of radiation from a circumbinary disk, fed from the L2 point in the SS 433 system. This was taken up by Fabrika (1993) in a prescient paper.  The same two horned structure was observed in H$\\alpha$, He I, O I 8446 \\AA\\ and the Paschen sequence through a campaign of nightly observations of SS 433 with the 3.6-m telescope on La Silla, Chile (Schmidtobreick \\& Blundell 2006a,b). The relevant observations covered more than two orbits during a period when SS 433 was quiescent. The two horned signature can be followed night by night in the spectra displayed in Fig.2 of Schmidtobreick \\& Blundell (2006b), in both H$\\alpha$ and He I. The least noisy signal is to be found in the brilliant H$\\alpha$ line and in Blundell, Bowler \\& Schmidtobreick (2008) this stationary H$\\alpha$ line was fitted as a superposition of Gaussian profiles. It consists of a broad component shown to be associated with the wind from the accretion disk and two narrow components separated by about 400 km s$^{-1}$. These narrow lines run almost railroad straight and do not shift much in position or separation over 30 days (see Fig.1 of Blundell, Bowler \\& Schmidtobreick 2008 and Fig.1 of the present paper). The orbital plane of the binary system is almost edge on and this sequence of pairs of narrow lines is the classic signature of the inner rim of a circumbinary disk, radiating strongly all the time from the regions to which the line of sight is tangent. Nonetheless, the intensities vary in antiphase with a period of 13 days, the orbital period of the binary. The interpretation was that the glowing inner rim fades and is refreshed as the binary rotates, possibly through ejection of material through the L2 point or perhaps by ultra violet and X rays from the accretion disk. The observations of SS 433 over an extended period (Blundell, Bowler $\\&$ Schmidtobreick 2007,2008) did not reach into the blue, but published data contain in addition to the brilliant Balmer H$\\alpha$ line the stronger He I lines at 6678 and 7065 \\AA\\ , which are also split, have narrow components which fluctuate with opposite phase but fade much faster than H$\\alpha$.  Both position and separation contain marked 13 day periodicities, so that it is not immediately obvious that they share the same origin as the simple H$\\alpha$ structure. In this paper I present those relevant data and discuss their interpretation in terms of a simple model for the stimulated inner circumbinary disk.  In this model the data are explained by the ring of fire orbiting with a speed of approximately 250 km s$^{-1}$, greater than half the separation of the centres of the fitted Gaussians, the brightest spot rotating close to the compact object and its accretion disk.  There are more recent observations concerning the putative circumbinary disk. First, it has been observed in Brackett $\\gamma$ (Perez \\& Blundell 2009) over about one orbital period. The extracted rotational velocity is again $\\sim$ 200 km s$^{-1}$ but the signal is squeezed between probable accretion disk lines, which complicates its extraction. Secondly, observations in both H$\\alpha$ and H$\\beta$ suggest that the apparent circumbinary disk lines are not attenuated by the wind from the accretion disk and hence their source is indeed {\\it{circumbinary}} (Perez \\& Blundell 2010).  I consider such other possible models for the origin of these split lines as have occurred to me. They are not plausible because of the marked degree to which the red and blue narrow components of H$\\alpha$ in Blundell, Bowler \\& Schmidtobreick (2008) are unmoving over more than two orbits, which is naturally explained by the disk model. ",
        "conclusions": "The great stability of the narrow red and blue components of the stationary H$\\alpha$ line in the spectrum of SS 433 is easily understood in terms of an orbiting circumbinary ring, presumed to be the inner rim of a larger disk. This stability is very much at odds with a source in an outward flow through the L2 point. The more complicated behaviour of the He I lines is consistent with an origin in the circumbinary disk, provided only that the more rapid fading can be accomodated. Thus the very simple model in which radiation from the circumbinary disk decays exponentially behind a leading edge, convoluted with a Gaussian function, accounts astonishingly well for the narrow components found within the stationary H$\\alpha$ and He I lines. H$\\alpha$ at least is contributed by radiation from the inner circumbinary disk, orbiting the binary at very approximately 250 km s$^{-1}$. The apparent systemic velocity of the ring is approximately 70 km s$^{-1}$.  H$\\alpha$ emission fades on a timescale of 14 days whereas He I has a fading time of about 4 days. In this simple model the leading edge of the H$\\alpha$ emission is found close to the passage of the compact object and its disk but the leading edge for He I is perhaps a day or so earlier. The irradiation of a given point on the circumbinary disk by the source of intense radiation in the vicinity of the compact object varies by a factor of almost three over the orbit, from geometric effects alone. When the compact object is furthest from a point on the disk, additionally the companion eclipses that portion. This suggests that intense radiation from the vicinity of the compact object periodically refreshes emission from the circumbinary disk - the effect might be augmented by arrival of material from L2 - and a decay scale of about one period is not unreasonable. It is perhaps curious that the He I signal decays much faster than H$\\alpha$ and the hot spot is ahead of the compact object. It seems entirely possible that some phases at least of the He I lines are dominated by radiation from the stream leaving the L2 point, which feeds the circumbinary disk and eventually the wider environment. For the remainder of this paper I shall suppose that the circumbinary disk is real and that the inner edge has at least approximately the orbital speed extracted from the model. The remaining uncertainty is the radius at which the ring of fire orbits.  The rotational speed of the inner circumbinary disk provides an important constraint on the mass of the system and hence on the mass of the compact object. If the radius at which the material orbits with speed $v$ is $fA$, $A$ being the separation of the two members of the binary, then the mass of the system {\\it$M_{\\rm S}$} is given by  \\begin{equation} M_S = 1.35 f^{3/2}(v/100)^3 \\end{equation}  in units of {\\it$M_{\\odot}$}, {\\it$v$} being specified in km s$^{-1}$ [Eq. 3 of Blundell, Bowler $\\&$ Schmidtobreick (2008)]. The innermost stable orbit about the binary system corresponds to {\\it$f$} approximately 2, estimates varying between 1.8 and 2.3. However, in a system such as SS 433 where material is almost certainly being added to the circumbinary disk via the L2 point, the glowing inner rim may be within the innermost stable orbit, depending on the residence time and the rate at which matter is added. I show in Table 1 the mass of the system {\\it$M_S$}, the companion {\\it$M_C$} and the compact object (including the mass of the accretion disk) {\\it$m_X$} as a function of the value of {\\it$f$}.  \\begin{table} \\centering \\vspace{1cm} \\begin{tabular}{llrr} \\hline $f$        &$M_{\\rm S}$       &$M_{\\rm C}$        &$m_{\\rm X}$  \\\\ 1.5       &38.8                     &22.6                       &16.1                \\\\ 1.6       &42.7                     &24.1                       &18.6                \\\\ 1.7       &46.8                     &25.6                       &21.0                \\\\ 1.8       &51.0                     &27.2                       &23.9                \\\\ 1.9       &55.3                     &28.7                       &26.7                \\\\ 2.0       &59.7                     &30.2                       &29.6                \\\\ 2.1       &64.2                     &31.7                       &32.6               \\\\ 2.2       &68.8                     &33.2                       &35.8               \\\\ 2.3       &73.6                     &34.7                       &38.9                \\\\  \\end{tabular} \\caption{\\label{tab:two} Masses in the binary system SS 433 as a function of the radius parameter $f$. $M_{\\rm S}$ is the total mass, $M_{\\rm C}$ is the mass of the companion and $m_{\\rm X}$ is the mass of the compact object, in units of $M_\\odot$ . ( From Eq.3, assuming $\\it v $ equal to 250 km s$^{-1}$. The masses $M_{\\rm S}$ scale with the cube of the orbital speed.)} \\end{table}  Table 1 has been composed assuming that the material in the ring of fire is orbiting the centre of mass of the binary at 250 km s$^{-1}$. This value is model dependent; to the extent that the model is accurate the radiating He is orbiting slower than H. It is most unlikely that bulk material radiating in the inner rim of the circumbinary disk is orbiting slower than 200 km s$^{-1}$. For this speed, the system mass is 40 $M_\\odot$ for a value of $f$ of 2.3, as assumed in Blundell, Bowler \\& Schmidtobreick (2008); material orbiting in the ring of fire at 250 km s$^{-1}$ would lie further inwards, $f$ of 1.5. Perhaps the H$\\alpha$ radiation is coming from increasing amounts of material close to joining the circumbinary disk rather than in stable orbits. In any event, the evidence from the circumbinary disk is that the system is massive, almost certainly exceeding 40 $M_{\\odot}$, and the compact object is a rather massive stellar black hole.    The Doppler shifts of a number of lines reported by Cherepashchuk et al. (2005) and attributed to the companion yield an orbital velocity of the companion about the binary centre of mass of 132 km s$^{-1}$ and hence a system mass of 42 $M_{\\odot}$. This is not in disagreement with data on the circumbinary disk, although $f$ would have to be a little less than 1.7 for disk material moving as fast as 250 km s$^{-1}$. The observations of Hillwig $\\&$ Gies (2008), Kubota et al (2010), interpreted as absorption in the atmosphere of the companion, are not consistent with the circumbinary disk modelled in this paper. They infer an orbital velocity for the companion of 58 km s$^{-1}$ and hence a system mass of about 17 M$_{\\odot}$. For disk material orbiting at 250 km s$^{-1}$ this yields a value of $f$ of 0.85 - unbelievably close in because the L2 radius corresponds to $f=1.26$. (For material at 220 km s$^{-1}$ $f$ would be 1.1 and even for 200 km s$^{-1}$ an $f$ of 1.33.) However, it seems possible that the absorption lines of Hillwig \\& Gies (2008) and Kubota et al (2010) are in fact produced in circumbinary material. The arguments against a circumbinary origin for such absorption lines set out in Hillwig et al (2004) do not obviously apply to material in an orbiting disk. The observations of a sinusoidal oscillation of amplitude $\\sim$ 60 km s$^{-1}$, sharing the orbital phase of the companion, are explained quantitatively if the origin of these lines is absorption of continuum light from the companion in circumbinary material orbiting at $\\sim$ 240 km s$^{-1}$. The reason is that the companion presents an orbital radius projected on the sky of $R_{\\rm C}{\\sin\\phi}$, where here $\\phi$ is the orbital phase, and at this elongation is viewed through disk material moving with a radial component of velocity $V_{\\rm r}$ given by  \\begin{equation} V_ r = \\frac{R_C}{R_D}V_D{\\sin\\phi} \\end{equation}  where the radius at which the circumbinary material orbits at speed $V_D$ is $R_D$. The radius of the companion orbit is $\\sim A/2$ and the disk material orbits at $\\sim 2A$. This is most easily illustrated by considering the configuration at an orbital phase of 0.25, when the companion is receding. At extreme elongation any circumbinary material through which it  is seen has, in the centre of mass system of the binary, a recessional velocity of $\\sim$ 60 km s$^{-1}$. Perhaps it is just a coincidence, but the numerical agreement is as good as it could be. The absorption line data of Hillwig \\& Gies (2008) and Kubota et al (2010) might thus be reconciled with the data attributed to the circumbinary disk which, as analysed in this paper, imply an orbital velocity for the companion exceeding 130 km s$^{-1}$. If those absorption lines do in fact originate in the circumbinary disk then identification of the spectral type of the companion as mid A is premature."
    },
    "1004/1004.0928_arXiv.txt": {
        "abstract": "We have used the {\\it Spitzer} satellite to monitor the mid-IR evolution of SN 1987A over a 5 year period spanning the epochs between days $\\sim$~6000 and 8000 since the explosion. The supernova (SN) has evolved into a supernova remnant (SNR) and its radiative output is dominated by the interaction of the SN blast wave with the pre-existing equatorial ring (ER). The mid-IR spectrum is dominated by emission from $\\sim 180$~K silicate dust, collisionally-heated by the hot X-ray emitting gas with a temperature and density of $\\sim 5\\times10^6$~K and $\\sim 3\\times 10^4$~\\cc, respectively. The mass of the radiating dust is $\\sim 1.2\\times 10^{-6}$~\\msun\\ on day 7554, and scales linearly with IR flux. Comparison of the IR data with the soft X-ray flux derived from \\chan\\ observations shows that the IR-to-Xray flux ratio, \\irx, is roughly constant with a value of 2.5. Gas-grain collisions therefore dominate the cooling of the shocked gas. The constancy of \\irx\\ is most consistent with the scenario that very little grain processing or gas cooling have occurred throughout this epoch. The shape of the dust spectrum remained unchanged during the observations while the total flux increased by a factor of $\\sim 5$ with a time dependence of $t'^{0.87\\pm0.20}$, $t'$ being the time since the first encounter between the blast wave and the ER. These observations are consistent with the transitioning of the blast wave from free expansion to a Sedov phase as it propagates into the main body of the ER, as also suggested by X-ray observations. The constant spectral shape of the IR emission provides strong constraints on the density and temperature of the shocked gas in which the interaction takes place. Silicate grains, with radii of $\\sim 0.2$~\\mic\\ and temperature of $T\\sim 180$~K, best fit the spectral and temporal evolution of the $\\sim 8 - 30$~\\mic\\ data. The IR spectra also shows the presence of a secondary population of very small, hot ($T \\gtrsim 350$~K), featureless dust. If these grains spatially coexist with the silicates, then  they must have shorter lifetimes. The data show slightly different rates of increase of their respective fluxes, lending some support to this hypothesis. However, the origin of this emission component and the exact nature of its relation to the silicate emission is still a major unsolved puzzle. ",
        "introduction": "About 10 years after its explosion on February 23, 1987, supernova (SN) 1987A has evolved from a supernova, when its radiative output was dominated by the release of radioactive decay energy in the ejecta, into a supernova remnant (SNR), when its radiative output became dominated by the interaction of its blast wave with the inner equatorial ring (ER). The ER is located at a distance of about 0.7 lyr from the center of the explosion, and could have been produced by mass loss from a single rotating supergiant  \\citep{heger98} or by a merger event in a binary system that also formed the two outer rings \\citep{morris09}. The transition from SN to SNR was observed at wavelengths ranging from radio to X-rays, and is summarized in Figure~18 in \\cite{bouchet06}.  In this paper we report on the continuing evolution of the $\\sim 5-30$~\\mic\\ spectrum and the 3.6, 4.5, 5.8, 8.0, and 24~\\mic\\ photometric fluxes from SN 1987A, spanning the $\\sim 5$ year period from day $\\sim$~6000 until day $\\sim$~8000 after the explosion. Initial reports and analysis of the IR evolution were presented by \\cite{bouchet04, bouchet06} and \\cite{dwek08a}. In \\S2 we present the IR data obtained by the \\spitz\\ satellite. The evolution of the IR emission and the dust composition are presented in \\S3. In \\S4 we derive the plasma conditions from the IR observations. The comparison of the IR emission with the X-ray emission, and the evolution of their flux ratio are discussed in \\S5. A brief summary of the paper is presented in \\S6.  \\begin{figure*}[ht!] % \\plotone{fig1.eps} \\caption{\\footnotesize{{\\it Spitzer} images of SN 1987A at Day 7975 (3.6 - 8~\\mic) and Day 7983 (24 - 70~\\mic). These reverse grayscale images use logarithmic scaling.}} \\label{grey} \\end{figure*}  \\begin{figure*}[ht!] % \\plotone{fig2.eps} \\caption{\\footnotesize{{\\it Spitzer} images of SN 1987A at Day 7975 (3.6 - 8~\\mic) and Day 7983 (24 - 70~\\mic).}} \\label{color} \\end{figure*} ",
        "conclusions": "In this paper we presented the mid-IR evolution of SN 1987A over a 5 year period spanning the epochs between days $\\sim$ 6000 and 8000 since the explosion. Its radiative output during this epoch is dominated by the interaction of the SN blast wave with the pre-existing equatorial ring (ER). The main results of this paper can be briefly summarized as follows: \\begin{enumerate} \\item The $\\sim 8-30$~\\mic\\ mid-IR spectrum is dominated by emission from $\\sim 180$~K silicate dust, collisionally-heated by the hot X-ray emitting gas with a temperature and density of $\\sim 5\\times10^6$~K and $\\sim (2-4)\\times 10^4$~\\cc, respectively. The mass of the radiating dust is $\\sim 1.2\\times 10^{-6}$~\\msun\\ on day 7554, and scales linearly with IR flux. \\item A secondary emission component dominates the spectrum in the $\\sim 5-8$~\\mic\\ region. Its intensity and spectral shape rule out any possible gas or synchrotron emission mechanism as the source of this emission. It must therefore attributed to a secondary dust component radiating at temperatures above $\\sim 350$~K. \\item The overall shape of the $\\sim 5-40$~\\mic\\ dust spectrum has not changed during the observations, suggesting that the density and temperature of the soft X-ray emitting gas have not significantly changed during the more than 5 years of IR observations. The constancy in the spectral shape of the IR emission also suggests that the mass ratio of the silicate to the secondary dust component remained roughly constant during this period. \\item The evolution of the IRAC and MIPS fluxes can be described by a power law in time since the first shock-ER encounter. The silicate emission increases as $t^{0.87}$, consistent with X-ray observations, suggesting that the blast wave has transitioned from a free expansion to the Sedov phase, and is now expanding into the main body of the ER. \\item The infrared-to-X-ray flux ratio, \\irx, is constant with a value of $\\sim 2.5$ throughout this epoch. The magnitude of \\irx\\ shows that the cooling of the shocked gas is dominated by IR emission from the collisionally-heated dust with radii $\\gtrsim 0.2$~\\mic, and that a significant fraction of the refractory elements in the ER should be depleted onto dust. \\item The constancy of \\irx\\ is consistent with the premise that neither grain destruction by sputtering nor cooling of the shocked gas have played a significant role in the evolution of the IR and X-ray emission. This scenario is consistent with the sputtering rates currently used in the literature as expressed in eq. (1). \\item The presence of a secondary dust component, radiating at significantly higher temperatures than the silicate dust, suggests that the grain radii or IR emissivities of this component must be significantly smaller than those of the silicates. Their sputtering lifetime could therefore be significantly shorter than that of the silicate grains and their evolution distinctly different from that of the silicates, especially around the $\\sim 6100$~d time interval. However, a significantly shorter grain lifetime contradicts the nearly constant or slightly decreasing flux ratio between the hotter secondary component and the silicate grain component. Hence, the nature of the secondary dust component remains a mystery. \\end{enumerate}  Five years of continued monitoring of SN1987A have lead us to conclude there is no current evidence for grain destruction. This conclusion stands in contrast to the results of our earlier analysis \\cite{dwek08a}, which were based only on two epochs of data at early times, and which we now see to be somewhat anomalous. Continuing observations will reveal evidence for grain destruction in the ER, and may elucidate the nature of the mysterious secondary dust component.   {"
    },
    "1004/1004.0575_arXiv.txt": {
        "abstract": "We investigate the implications of energy-dependence of the speed of photons, one of the candidate effects of quantum-gravity theories that has been most studied recently, from the perspective of observations in different reference frames. We examine how a simultaneous burst of photons would be measured by two observers with a relative velocity, establishing some associated conditions for the consistency of theories. For scenarios where the Lorentz transformations remain valid these consistency conditions allow us to characterize the violations of Lorentz symmetry through an explicit description of the modification of the quantum-gravity scale in boosted frames with respect to its definition in a preferred frame. When applied to relativistic scenarios with a deformation of Lorentz invariance that preserves the equivalence of inertial observers, we find an insightful characterization of the necessity to adopt in such frameworks non-classical features of spacetime geometry, e.g. events that are at the same spacetime point for one observer cannot be considered at the same spacetime point for other observers. Our findings also suggest that, at least in principle (and perhaps one day even in practice), measurements of the dispersion of photons in relatively boosted frames can be particularly valuable for the purpose of testing these scenarios. ",
        "introduction": "Over the past few years there has been a growing interest in the investigation of possible high-energy quantum-gravity-induced deviations from Lorentz invariance, that would induce modifications of the energy-momentum dispersion (on-shell) relation \\cite{grbgac,gampul,billetal,schaefer,kifune,mexweave,ita,aus,gactp}. This hypothesis finds support in preliminary results obtained in some popular approaches to the study of the quantum-gravity problem, most notably approaches based on ``spacetime noncommutativity\" \\cite{gacmajid,kowaNCSTdisp} or inspired by ``loop quantum gravity\"~\\cite{gampul,urrutiaPRD}. It is expected that the scale governing such deviations from Lorentz symmetry is the scale where the same frameworks predict a breakdown of the familiar description of spacetime geometry, the ``quantum-gravity energy scale\", here denoted by $E_{QG}$, which should be roughly of the order of the Planck scale ($E_{pl}=\\sqrt{\\hbar c^5/G}\\sim1.2\\times10^{28}$ eV). \\par  While the study of particles with energies close to the Planck scale is far beyond our reach in particle-physics laboratories and even in astrophysical observatories, it is possible to look for the minute effects that Planck-scale deviations from Lorentz symmetry produce for particles with energies much lower than the Planck scale. From this perspective of the search of small leading-order corrections, some astrophysical phenomena could provide meaningful insight, by either producing signals of the modified dispersion, or alternatively, providing constraints on the relevant models. Indeed, astronomical observations have already set valuable limits on these leading-order corrections and as such on some of the alternative scenarios for ``quantum gravity\". Most insightful are the studies based on the high-energy modifications of particle speeds \\cite{grbgac,schaefer,neutNature,magicPLB,hessPRL,emnPLB2009,gacSMOLINprd,fermiNATURE} and on the modifications to high-energy reaction thresholds (as relevant for the ultra-high-energy cosmic-rays or very-high-energy $\\gamma$-rays) \\cite{kifune,ita,aus,gactp,liberatiUHECR2009,steckeNEW,absorptionPRD}. The rapid development of astronomical instrumentation over the last years has improved significantly our ability to test these scenarios. In particular the recent results of the Fermi telescope collaboration \\cite{fermiNATURE}, using a time-of-flight analysis of a short gamma-ray burst GRB090510, achieved a sensitivity to Planck-scale effects and set for the first time a limit of $E_{QG}\\gtrsim E_{pl}$ (for the case of effects introduced linearly in the quantum-gravity scale, defined below as $n=1$). \\par  The conceptual perspective that guides these studies is consistent with the history of other symmetries in physics, which were once thought to be fundamental but eventually turned out to be violated. There is no essential reason to believe that Lorentz symmetry will be spared from this fate. If the symmetry is broken at a scale $E_{QG}$, we should expect leading-order corrections to arise even at much lower energies, and hence it is just natural to explore their possible implications. However, in view of the importance of Lorentz symmetry in the logical consistency of our present formulation of the laws of physics, one should ask which modified dispersion relations can be placed into a viable and consistent theory? The analysis presented here intends to contribute in this direction. We investigate this question from a rather general perspective, aware of the fact that it is quantum-gravity research that provides the key motivation for these studies, but also in principle open to the possibility that the conjectured modifications of the dispersion relation might have different origin. \\par  We consider a generic ``low-energy\", $E\\ll E_{QG}$, leading-order modification of the dispersion relation of the form:\\footnote {The factor $\\frac{2}{n+1}$ is here introduced only for the consistency of $E_{QG}$ in the description of speeds with the analogous parameter most commonly used in the relevant literature.} \\begin{equation} \\label{LIDform} E^2-p^2c^2-m^2c^4 \\simeq \\pm \\frac{2}{n+1}E^2\\left(\\frac{E}{E_{QG}}\\right)^n . \\end{equation} The power of this leading correction, $n$, and the parameter $E_{QG}$ are model-dependent and should be determined experimentally. The conventional speed of light constant, $c$, remains here the low-energy limit of massless particles' speed, and we put hereafter $c=1$ (with the exception of only a few formulas where we reinstate it for clarity). Notice that the fact that we work in leading order in the Planck-scale corrections allows us to exchange the (modulus of) momentum of a photon with its energy in all Planck-scale suppressed terms. The modification of the dispersion relation produces a difference between energy and momentum of a photon, but this is itself a first-order correction and taking it into account in terms that are already suppressed by the smallness of the Planck scale would amount to including subleading terms. \\par  The approach based on (\\ref{LIDform}), which is adopted here, has been considered by many authors as the natural entry point to the phenomenology of Lorentz invariance violation \\cite{grbgac,kifune,ita,aus,gactp,neutNature,magicPLB,hessPRL,liberatiUHECR2009}. We are mainly concerned with the conceptual implications of such modifications of the dispersion relation for the way in which the same phenomenon is observed in different reference frames, and for our exploratory purposes it is sufficient to focus on (\\ref{LIDform}). Our results apply to all cases in which (\\ref{LIDform}) is satisfied to leading order. The findings should apply also to scenarios with birefringence, which could be induced by Planck-scale effects (see, {\\it e.g.}, Ref.~\\cite{gampul}). Our results may provide a first level of intuition even for the possibility of Planck-scale induced ``fuzziness\", which is the case of scenarios in which there is no systematic modification of the dispersion relation but modifications roughly of the form (\\ref{LIDform}) occur randomly, affecting different particles in different ways depending on the quantum fluctuations of spacetime that they experience (for details see, {\\it e.g.}, Refs.~\\cite{ngfuzzy,gacSMOLINprd}). \\par  Assuming, as commonly done in the related literature~\\cite{grbgac,gampul,billetal,schaefer,urrutiaPRD}, that the standard relation $v= \\partial E /\\partial p$ holds, the dispersion relation (\\ref{LIDform}) leads to the following energy dependence of the speed of photons: \\begin{equation} \\label{LawForSpeed} v(E) = 1 \\pm \\left( \\frac {E}{E_{QG}} \\right)^n . \\end{equation} This is an important phenomenological prediction of the modified-dispersion scenario. This effect is the basis for the most generic tests of the Lorentz-violation theories. We will consider the case of subluminal motion, corresponding to the `-' sign in Eq. (\\ref{LawForSpeed}). A similar analysis can be performed for the case with superluminal speeds. We here contemplate an experiment where two relatively boosted observers detect two photons, which are emitted simultaneously at the source but arrive separated by a delay due to their different energies according to (\\ref{LawForSpeed}). The basic idea of our study is to compare, under different theoretical scenarios, the time delay measured in the different reference frames. This primarily serves as a gedanken experiment that establishes features which modified-dispersion models must include in order to have a consistent scenario. \\par  We show that if in all reference frames photons have a dispersion relation as in (\\ref{LIDform}) and a speed law as in (\\ref{LawForSpeed}) (with fixed frame-independent parameters), then assuming the transformations between frames are governed by the standard Lorentz laws would lead to a contradiction between the measurements of the different observers. This is expected, but we use our explicit derivation to deduce other requirements the models must fulfill for a solution. We consider two classes of modified-dispersion models. The first class consists of theories where the Lorentz Symmetry is Broken (LSB) with the existence of a preferred inertial frame (sometimes attributed to the Cosmic Microwave Background frame). The small-scale structure and hence high-energy behavior is defined in this reference frame. Therefore, the dispersion relation is allowed to take different forms in other frames. In the following we will consider only the most common LSB theories in which the Lorentz transformations are still applicable when going from one frame to another. The second class of models describes a scenario where the Lorentz symmetry is merely deformed, such that the equivalence of all inertial observers is maintained. This means that the Lorentz symmetry makes way for a more complex symmetry, but the theory is still relativistic (there is no preferred reference frame). These are called ``Doubly Special Relativity\" (DSR) models. Within this scenario the laws of transformation are necessarily modified from the standard Lorentz transformations. This is done in such a way that all observers agree on the physical laws, including the energy-momentum dispersion relation. Not all models of DSR predict energy-dependent velocities of photons, but we deal with the analysis of the common DSR framework where the velocities behave as in (\\ref{LawForSpeed}). This is to say that we will have two complementary scenarios - LSB with standard Lorentz transformations but allowing for departures between the dispersion relations in different frames, and DSR where the theory is frame-independent but the transformations between reference frames differ from the Lorentz ones. \\par  We observe that the large effort recently devoted to the phenomenology of modifications to Lorentz invariance has focused on analyses performed in a single frame. We believe that deeper insight on the fate of the Lorentz symmetry can be gained by also investigating how the same phenomenon is viewed by two different observers. We argue that our study sheds light on several conceptual issues, which in turn may well provide guidance even for the ongoing effort utilizing the standard ``laboratory frame\" tools. We obtain definitive results for our LSB scenario, and for the DSR case we find that some non-classical features of spacetime are required (while the phenomenological DSR results are severely conditioned by our restrictive assumption of a classical-spacetime implementation). In principle (and perhaps one day in practice, after a Lorentz-violating effect is observed in one frame) our description of the effects of the LSB scenario for observations of the same phenomenon in two reference frames with a relative boost could be even exploited experimentally. In spite of the limitations associated with the assumption of a classical-spacetime implementation for the DSR case, our findings for the comparison between the LSB scenario and the DSR scenario provide some encouragement for the possibility that our scheme could also be exploited to discriminate between these alternative spacetime-propagation models. ",
        "conclusions": ""
    },
    "1004/1004.4951_arXiv.txt": {
        "abstract": "\\vspace{1cm} \\centerline{\\bf ABSTRACT}\\vspace{2mm} In the present work, by the help of the newly released Union2 compilation which consists of 557 Type Ia supernovae (SNIa), we calibrate 109 long Gamma-Ray Bursts (GRBs) with the well-known Amati relation, using the cosmology-independent calibration method proposed by Liang {\\it et al.}. We have obtained 59 calibrated high-redshift GRBs which can be used to constrain cosmological models without the circularity problem (we call them ``Hymnium'' GRBs sample for convenience). Then, we consider the joint constraints on 7 cosmological models from the latest observational data, namely, the combination of 557 Union2 SNIa dataset, 59 calibrated Hymnium GRBs dataset (obtained in this work), the shift parameter $R$ from the WMAP 7-year data, and the distance parameter $A$ of the measurement of the baryon acoustic oscillation (BAO) peak in the distribution of SDSS luminous red galaxies. We also briefly consider the comparison of these 7 cosmological models. ",
        "introduction": "\\label{sec1}  Since the discovery of current accelerated expansion of our universe~\\cite{r1}, Type Ia supernovae (SNIa) have been considered to be a powerful probe to study this mysterious phenomenon. However, SNIa are plagued with extinction from the interstellar medium, and hence the current maximum redshift of SNIa is only about $z\\simeq 1.755$. On the other hand, the redshift of the last scattering surface of cosmic microwave background (CMB) is about $z\\simeq 1090$. There is a wide ``desert'' between the redshifts of SNIa and CMB. So, the observations at intermediate redshift are important to distinguish cosmological models.  Recently, Gamma-Ray Bursts (GRBs) have been proposed to be a complementary probe to SNIa (see e.g.~\\cite{r2} and references therein). So far, GRBs are the most intense explosions observed in our universe. Their high energy photons in the gamma-ray band are almost immune to dust extinction, in contrast to supernovae. Up to now, there are many GRBs observed at $0.1<z\\leq 8.1$, whereas the maximum redshift of GRBs is expected to be $10$ or even larger~\\cite{r3}. Therefore, GRBs are considered to be a hopeful probe to fill the ``desert'' between the redshifts of SNIa and CMB. We refer to e.g.~\\cite{r2,r4,r5,r6} for some comprehensive reviews on the so-called GRB cosmology. As is well known,  there is a circularity problem in the direct use of GRBs~\\cite{r7}, mainly due to the lack of a set of low-redshift GRBs at $z<0.1$ which are cosmology-independent. To alleviate the circularity problem, some statistical methods have been proposed, such as the scatter method~\\cite{r8}, the luminosity distance method~\\cite{r8}, and the Bayesian method~\\cite{r9}. Other methods trying to avoid the circularity problem have been proposed in e.g.~\\cite{r10,r11}. Recently, a new idea of the distance ladder to calibrate GRBs in a completely cosmology-independent manner has been proposed in~\\cite{r12,r13} independently. Similar to the case of calibrating SNIa as secondary standard candles by using Cepheid variables which are primary standard candles, we can also calibrate GRBs as standard candles with a large amount of SNIa. And then, the calibrated GRBs can be used to constrain cosmological models without the circularity problem. We refer to e.g.~\\cite{r12,r14,r15} for some relevant works.  It is worth noting that in the literature most relevant works mainly used the 69 GRBs compiled in~\\cite{r2} or the 70 GRBs compiled in~\\cite{r16}. As of the end of 2009, there are 109 long GRBs with measured redshift and spectral peak energy compiled in~\\cite{r17}. The number of usable GRBs is significantly increased~\\cite{r60}. Among these 109 GRBs, the data of 70 GRBs are taken from~\\cite{r16}; the data of 25 GRBs are taken from~\\cite{r18}; the data of remaining 14 GRBs (090516, 090618, 090715B, 090812, 090902B, 090926, 090926B, 091003, 091018, 091020, 091024, 091029, 091127, 091208B) are provided by L.~Amati in private communication~\\cite{r19}. On the other hand, the cosmology-independent method to calibrate GRBs proposed in~\\cite{r12} needs a large amount of SNIa at $z<1.4$. Thus, the SNIa dataset plays an important role in the calibration of GRBs. Previously, the relevant works (e.g.~\\cite{r12,r14,r15}) used the 192 Davis SNIa dataset or the 307 Union SNIa dataset to calibrate GRBs. Very recently, the Supernova Cosmology Project (SCP) collaboration released their Union2 compilation which consists of 557 SNIa~\\cite{r20}. The Union2 compilation is the largest published and spectroscopically confirmed SNIa sample to date. The number of usable SNIa is also significantly increased. Obviously, a larger SNIa dataset could bring a better interpolation for the cosmology-independent method to calibrate GRBs proposed in~\\cite{r12}. Motivated by both the significant improvements in GRBs and SNIa, in the present work, we update the calibration of GRBs with the well-known Amati relation, using the cosmology-independent method proposed in~\\cite{r12}, and obtain 59 calibrated high-redshift GRBs which can be used to constrain cosmological models without the circularity problem (we call them ``Hymnium'' GRBs sample for convenience).  On the other hand, very recently the WMAP Collaboration also released their 7-year CMB data (WMAP7) in~\\cite{r21}. Therefore, it is natural to consider the joint constraints on cosmological models from these latest observational data, namely, the combination of 557 Union2 SNIa dataset~\\cite{r20}, 59 calibrated Hymnium GRBs dataset (obtained in this work), the shift parameter $R$ from the WMAP7 data~\\cite{r21}, and the distance parameter $A$ of the measurement of the BAO peak in the distribution of SDSS luminous red galaxies~\\cite{r22,r23}. In this work, we consider 7 cosmological models and obtain the observational constraints on them. Notice that we assume the universe to be spatially flat throughout this work.  This paper is organized as followings. In Sec.~\\ref{sec2}, by the help of the newly released Union2 compilation which consists of 557 SNIa, we calibrate the 109 GRBs with Amati relation, using the cosmology-independent method proposed in~\\cite{r12}. In Sec.~\\ref{sec3}, we briefly introduce the methodology to constrain cosmological models with the combined latest observational data. Then, we obtain the corresponding constraints on 7 cosmological models in Sec.~\\ref{sec4}. Further, we consider the comparison of these 7 models in Sec.~\\ref{sec5}. Finally, we give the brief conclusions in Sec.~\\ref{sec6}. ",
        "conclusions": "\\label{sec6}  In the present work, by the help of the newly released Union2 compilation which consists of 557 SNIa, we calibrated 109 long GRBs with the well-known Amati relation, using the cosmology-independent calibration method proposed in~\\cite{r12}. We have obtained 59 calibrated high-redshift GRBs which can be used to constrain cosmological models without the circularity problem (we call them Hymnium GRBs sample for convenience). One can directly read off the numerical data of these 59 Hymnium GRBs from Table~\\ref{tab2} of this paper. They are also available upon request to our email address. Considering their unique ability to extend our vision up to redshift $z=8.1$, we recommend to use these 59 Hymnium GRBs in the relevant works.  We considered the joint constraints on 7 cosmological models from the latest observational data, namely, the combination of 557 Union2 SNIa dataset~\\cite{r20}, 59 calibrated Hymnium GRBs dataset (obtained in this work), the shift parameter $R$ from the WMAP7 data~\\cite{r21}, and the distance parameter $A$ of the measurement of the BAO peak in the distribution of SDSS luminous red galaxies~\\cite{r22,r23}. Note that in general it is better to see how the parameters are constrained both with and without GRBs. However, in fact the cosmological constraints on the same models without GRBs have been considered in e.g.~\\cite{r52}, although they used slightly earlier SNIa and CMB data. So, we do not consider the constraints without GRBs in the present work. Comparing with the previous results in the literature (e.g.~\\cite{r52}), one can find that the constraints obtained in this work are tighter, thanks to these newly improved observational data. We also briefly considered the comparison of these 7 cosmological models. The $\\Lambda$CDM model is still the best one from the perspective of the latest cosmological observations."
    },
    "1004/1004.3789_arXiv.txt": {
        "abstract": "Fully cosmological, high resolution N-Body + SPH simulations are used to investigate the chemical abundance trends of stars in simulated stellar halos as a function of their origin. These simulations employ a physically motivated supernova feedback recipe, as well as metal enrichment, metal cooling and metal diffusion. As presented in an earlier paper, the simulated galaxies in this study are surrounded by stellar halos whose inner regions contain both stars accreted from satellite galaxies and stars formed in situ in the central regions of the main galaxies and later displaced by mergers into their inner halos.  The abundance patterns ([Fe/H] and [O/Fe]) of halo stars located within 10 kpc of a solar-like observer are analyzed. We find that for galaxies which have not experienced a recent major merger, in situ stars at the high [Fe/H] end of the metallicity distribution function are more [$\\alpha$/Fe]-rich than accreted stars at similar [Fe/H]. This dichotomy in the [O/Fe] of halo stars at a given [Fe/H] results from the different potential wells within which in situ and accreted halo stars form. These results qualitatively match recent observations of local Milky Way halo stars. It may thus be possible for observers to uncover the relative contribution of different physical processes to the formation of stellar halos by observing such trends in the halo populations of the Milky Way, and other local $L^{\\star}$ galaxies. ",
        "introduction": "Observational evidence has been mounting that the stellar halo of the Milky Way (MW) was likely assembled through a combination of both hierarchical accretions and in situ star formation. The properties of local halo stars suggest that the Galaxy's halo is composed of at least two distinct but overlapping stellar populations \\citep{Carollo2010}. These two halo components exhibit different spatial distributions, orbits and metallicities, most likely because each formed through a different mechanism, and on a different timescale. While the outer halo component, which dominates at distances greater than $\\sim 20$ kpc, with [Fe/H] $\\sim -2.5$, is consistent with having been accreted from merged dwarf galaxies, the stars in the more metal rich ([Fe/H] $\\sim -1.5$) inner halo likely formed predominantly in situ \\citep{Norris1994,Chiba2000,Carollo2007, Carollo2010}.  Studies of the detailed chemical abundances of MW halo stars have sought to place further constraints on the formation of the Galactic halo \\citep{Nissen1997,Stephens2002,Venn2004,Pritzl2005,Ishigaki2009}. These works rely on the behavior of the [$\\alpha$/Fe] ratio to indicate the timescale of the star formation history of the halo. Alpha elements are primarily produced in core-collapse supernovae, which occur after $\\sim 10^7$ years, while the majority of iron is produced in supernovae Type Ia, whose timescale of $\\leq 10^9$ years is significantly longer \\citep{Maoz2010}. In the metallicity range of $-2<[Fe/H]<-0.5$, inner halo stars are $\\sim 0.1$ dex more alpha rich than outer halo stars, implying that the Galaxy's halo was not uniformly enriched by core-collapse supernovae \\citep{Ishigaki2009}. In a local sample of stars, those with [$\\alpha$/Fe] $\\sim 0.3$ show a constant [$\\alpha$/Fe] over the above metallicity range, while stars with lower $\\alpha$-ratio, of about 0.15, exhibit a decrease in [$\\alpha$/Fe] toward higher metallicity \\citep{Nissen2010}. Similar declines to lower [$\\alpha$/Fe] have been observed in dwarf galaxies, though at lower [Fe/H] values \\citep{Tolstoy2009}. These trends strengthen the argument that the Milky Way's halo likely hosts two populations of stars, with distinct formation origins.  In the past year, theoretical results from galaxy simulations have been used to investigate the contribution of in situ stars to the formation history of stellar halos. Using fully cosmological N-Body + Smooth Particle Hydrodynamic (SPH) simulations of four stellar halos surrounding disk galaxies \\citet[][hereafter Paper I]{Zolotov2009} showed that while the formation of most stellar halos is dominated by accretion events, a population of in situ stars resides in the inner regions of halos as well.  These in situ halo stars form in the inner most regions of the simulated galaxy and are displaced into the stellar halo through major mergers. In situ stars may also end up in stellar halos, along side accreted stars, through minor accretion events, which heat up disks and eject a small fraction of their stars into kinematic halo orbits \\citep{Purcell2009}.  Previous work on the chemical abundance trends of stellar halos has relied on simulations that only allowed for the formation of stellar halos via hierarchical buildup, using dark-matter only simulations, combined with semi-analytic models \\citep{Bullock2005,Robertson2005, Font2006a,Font2006c,Johnston2008, DeLucia2008}. These simulations, however, do not have the capability to examine possible chemical abundance trends for in situ halo stars.   In this paper, we use four high resolution cosmological SPH+N-Body simulations of disk galaxies to investigate the chemical abundance trends of both accreted and in situ halo stars. While past numerical work on chemical abundance trends in halo stars have focused solely on accreted stars, the goal of this work is to test whether in situ and accreted stars are distinguishable using chemical abundance patterns. ",
        "conclusions": "In this paper, we have studied the chemical abundance trends of inner halo stars in four SPH+N-Body simulations of approximately L* disk galaxies. We investigate the possibility of using chemical abundance trends to discern the relative importance of in situ and accretion processes in building up the inner stellar halos of Milky Way-mass galaxies.  For each simulated galaxy we concentrated on a sample of halo stars located within 10 kpc of a simulated sun-like observer. This sample included stars from two populations, one formed in situ, and one accreted. We find that in the inner halos of galaxies where a recent binary merger did not occur, high [Fe/H] in situ halo stars are more [$\\alpha$/Fe] rich than accreted stars at similar [Fe/H].  The bimodal distribution of [$\\alpha$/Fe] at high [Fe/H] for in situ and accreted stars is primarily due to the different mass galaxies in which they formed. The in situ halo populations (in all but h258) were formed early on deep in the potential wells of the primary dark matter halos, where a high star formation rate ensured that only core-collapse SNe dominated the chemical enrichment up to high [Fe/H]. The accreted stars, which formed later on average than the in situ stars, and in shallower potential wells, underwent a slower chemical evolution, where SNe Type Ia contributed iron at lower [Fe/H], resulting in a lower [$\\alpha$/Fe] at a given [Fe/H]. Such trends are expected given a mass-metallicity relation \\citep{Brooks2007}. Though the satellite galaxies in our simulations are too bright, and hence too metal rich, we have found that the relative chemical abundance patterns between in situ and accreted stars is a qualitative trend that is robust to the details of the simulated halos' MDFs.   We predict that a large systematic survey of the detailed chemical abundances of inner halo stars in the Milky Way will exhibit such dual trends in [Fe/H] vs [$\\alpha$/Fe]. At higher metallicities, [Fe/H]$\\geq -1.0$, Milky Way stars in the inner several kpc of the halo should exhibit a bimodal distribution in [$\\alpha$/Fe]. Unless the Galaxy has experienced a recent major major, high [$\\alpha$/Fe] stars in this metallicity range will have formed in the Milky Way's disk or bulge, and then been displaced into halo orbits. Low [$\\alpha$/Fe] stars in this regime will likely have formed in smaller potential wells than the Milky Way, and were accreted from satellite companions at high redshift. In fact, such trends in $\\alpha$ and Fe have already been observed in a small sample of 94 local Milky Way halo dwarfs \\citep{Nissen2010}. Larger planned surveys, like HERMES, which will obtain high resolution spectra of MW halo stars will be able to observe such patterns, and quantify the importance of different physical processes to the MW's halo formation."
    },
    "1004/1004.1329_arXiv.txt": {
        "abstract": "TeV photons from blazars at relatively large distances, interacting with the optical--IR cosmic background, are efficiently converted into electron--positron pairs. The produced pairs are extremely relativistic (Lorentz factors of the order of $10^6$--$10^7$) and promptly loose their energy through inverse Compton scatterings with the photons of the microwave cosmic background, producing emission in the GeV band. The spectrum and the flux level of this reprocessed emission is critically dependent on the intensity of the intergalactic magnetic field, $B$, that can deflect the pairs diluting the intrinsic emission over a large solid angle. We derive a simple relation for the reprocessed spectrum expected from a steady source. We apply this treatment to the blazar 1ES 0229+200, whose intrinsic very hard TeV spectrum is expected to be approximately steady. Comparing the predicted reprocessed emission with the upper limits measured by the {\\it Fermi}/Large Area Telescope, we constrain the value of the intergalactic magnetic field to be larger than $B \\simeq 5\\times 10^{-15}$ Gauss, depending on the model of extragalactic background light. ",
        "introduction": "Although cosmic magnetic fields are present at all scales and likely play a non negligible role in shaping the Universe, our knowledge of them is still rather poor. In particular, magnetic fields at cosmological scales ($>1$ Mpc) remain elusive (e.g. Kronberg 2001, Neronov \\& Semikoz 2009). From the theoretical point of view, intergalactic space can contain the traces of magnetic fields produced in the initial phases of the Universe (e.g., Grasso \\& Rubinstein 2001) and can also be polluted by magnetic field ejected by galaxies and quasars (e.g. Furlanetto \\& Loeb 2001). Methods based on the rotation measure in the radio band constrain the intensity of the intergalactic magnetic field (IGMF) below $B<10^{-11}$--$10^{-9}$ G (e.g. Kronberg 2001 and references therein) but these values depends on the highly uncertain correlation length of the field.  As early proposed by Plaga (1995), IGMF could be probed by using indirect methods based on the effects of the IGMF on the electron--positron pairs produced through the interaction of a TeV photon from a cosmological source with a low energy photon of the optical--IR cosmic background. The produced pairs promptly loose their energy through inverse Compton (IC) scattering, in which a photon of the cosmic microwave background (CMB) is boosted at $\\gamma$--ray (GeV) energies. The existence of a magnetic field, even with the tiny intensities suggested by the measures mentioned above, modifies the properties of the $\\gamma$--ray emission, causing observable effects (time delays, e.g. Dai \\& Lu 2002; formation of extended $\\gamma$--ray halos, Aharonian, Coppi \\& Voelk 1994, Neronov et al. 2010) that can be exploited to infer the value of $B$ (e.g. Neronov \\& Semikoz 2009 and references therein). The reprocessing of multi--TeV emission from blazars into GeV--MeV emission has been also proposed has a possible contribution to the observed $\\gamma$--ray background (Coppi \\& Aharonian 1997, Venters 2010).  In this paper we use a method based on the observed level of the GeV emission of an highly absorbed BL Lac, 1ES 0229+200, whose intrinsic TeV spectrum is expected to be almost steady and very hard (Aharonian et al. 2007, Tavecchio et al. 2009). Several papers discussed a similar method applied to the emission from gamma--ray bursts (e.g. Dai \\& Lu 2002, Razzaque et al. 2004, Takahashi et al. 2008). A similar method for blazars have been already outlined in Dai et al. (2002) and Murase et al. (2008), but these works were focused on the discussion of the {\\it variable} GeV emission, by--product of the reprocessing of the primary TeV emission of a rapidly flaring blazar. Here, instead, we present a simplified treatment of the case of stationary emission of the primary TeV source, suitable to describe the case of 1ES 0229+200. Basically, our approach is based on the fact that for increasing values of the IGMF the level of expected reprocessed emission in the MeV-GeV band  decreases, since the total amount of the primary flux is spread over larger solid angles. The comparison between the expected level of the reprocessed flux and the flux at GeV energies observed by the {\\it Fermi}/Large Area Telescope thus provide a value (or a lower limit) on the intensity of the IGMF.  We use the notation $Q=Q_X 10^X$ in cgs units.  After the completion of this work a paper on same argument have been published (Neronov \\& Vovk 2010). The authors derive a value of the magnetic field very close to the value derived here, using a full treatment considering also the development of a cascade. However, they implicitly assume that the source is emitting TeV radiation {\\it isotropically}. In this case the total flux of the reprocessed emission does not depends on the IGMF. The IGMF has the only effect to reduce the expected surface brightness of the reprocessed emission, lower for larger values of the magnetic field. Instead, in our work we assume that the primary high-energy photons from the blazar, as commonly assumed, are strongly beamed since they are produced in a narrow relativistic jet. In this case the effect of the IGMF is to spread the reprocessed emission in solid angles larger than the original beaming angle and thus reduce the total observed flux, even if the source remains unresolved.   \\begin{figure} \\hskip -1.7truecm \\psfig{file=cartoon.eps,height=8.9cm,width=10.9cm}  % \\vskip 0.5 true cm \\caption{ Cartoon of the reprocessing of the absorbed TeV radiation. The source illuminates the regions inside a cone with semi--aperture $\\theta _{\\rm c}$. TeV photons are converted into e$^{+/-}$ pairs at a typical distance of hundreds of Mpc. Pairs cool rapidly through inverse Compton scattering on the CMB. For a given intensity of the (perpendicular component of the) extragalactic magnetic field, pairs with very high energy (upper panel) cool without changing their direction (arrows) and thus the reprocessed GeV emission from the IC is beamed within the same angle $\\theta _{\\rm \\,\\, c}$. The emission of progressively lower energy of the pairs is instead spread over larger angles, due to the curved trajectories of the pairs, in turn due to their longer lifetimes. As detailed in the text the resulting overall spectrum (right panels) can be approximates by three power laws, each branch due to electrons emitting within the original cone (upper), within a cone larger than the original (middle) and almost isotropically (lower). } \\label{cartoon} \\end{figure} ",
        "conclusions": "The lower limit on the value of the magnetic field derived here, $B>5\\times 10^{-15}$ G can be considered one of the most stringent value ever derived for the IGMF. The value is mainly constrained by the LAT upper limit above 10 GeV, in which the source is tentatively detected at the 4$\\sigma$ level. We remark that, since we derive a lower limit on $B$, even if this {\\it Fermi}/LAT measure is considered as an upper limit the conclusion does not change.  The main assumptions adopted in this paper are: (i): the amount of reprocessed energy is derived from the level of the observed spectrum measured by H.E.S.S. (Aharonian et al. 2007); (ii) we use the {\\it Low SFR} Kneiske et al. (2004) model to calculate the optical depth; (iii) we assume that the maximum energy of the intrinsic absorption is 11 TeV.  Assumption (i) is supported by the extremely small variability observed by H.E.S.S. during the observations in 2005--2006. Of course, any variations of the intrinsic emission of the blazars, will be reflected into the reprocessed emission. Ideally, simultaneous TeV and GeV observations would be required to take into account variability and exploit the information carried by the delay between the intrinsic TeV and the reprocessed GeV emission (see Dai et al. 2002, Murase et al. 2008 for discussions). However, this approach is prevented in the specific case of 1ES 0229+200 by the very small GeV flux, still at the detection limit of LAT after 18 months of observations. Also in other TeV BL Lacs more bright in the TeV band, integration over few days are necessary to detect the GeV emission (e.g. Abdo et al. 2010b). Moreover, in these objects, the intrinsic GeV emission is much more luminous than the expected reprocessed radiation. These features make difficult to reveal the reprocessed radiation even during bright flares.  The use of a specific model for the absorption is unavoidable. However, several of recent EBL models converge at the wavelengths above 10 $\\mu$m, the ones determining the opacity for $\\gamma$--rays of energy below 10 TeV. In particular, the recent models of Franceschini et al. (2008), Gilmore et al. (2009), Finke et al. (2010) agree well (see e.g. the discussion of Finke et al. 2010) and are consistent with the low level of the EBL suggested by recent observations of Cherenkov telescopes (e.g. Aharonian et al. 2006, Mazin \\& Raue 2007; see also Kneiske \\& Dole 2010). The model adopted here ({\\it LowSFR} of Kneiske et al. 2004) provides an optical depth similar to that of all the other updated models up to energies of 4--5 TeV. For larger energies the predicted optical depth is {\\it lower} than that of the other models (see also Tavecchio et al. 2009). We again stress that this implies to minimize the luminosity of the reprocessed radiation, and hence the derived lower limit of $B$ is a conservative value.   The assumed maximum energy of the intrinsic emission, point (iii), is somewhat critical, since it determines the total amount of energy reprocessed and then re-emitted in the GeV band, and the possible development of cascades. The first point is clear: since the intrinsic spectrum is rather hard, the total luminosity of reprocessed radiation (and its maximum energy) depends on $E_{\\rm max}$. Under the assumption that the spectrum is proportional to $E^{1/3}$, the normalization of the reprocessed emission is (see Eq. \\ref{finale}) $k \\propto E_{\\rm max}^{4/3}/\\epsilon_{\\rm max}\\propto E_{\\rm max}^{2/3}$. Therefore all the curves reported in Fig.\\ref{sed} will shift upward by this factor. We consider unlikely that the spectrum extends unbroken at energies well above few tens of TeV and therefore the correction factor cannot be much larger than a factor of a few. In any case we remark again that, since  a larger $E_{\\rm max}$ would have the effect to {\\it increase} the value of the derived lower limit (a larger $B$ would be required to downshift the curves), our limit can be considered also in this respect, again, as a conservative value. The same argument applies when considering the possible role of electromagnetic cascades. For values of $E_{\\rm max}$ larger than $\\sim 30$ TeV a sizeable amount of the reprocessed flux would be emitted above $\\sim 500$ GeV (the energy above which the optical depth for pair conversion exceeds 1) and then reabsorbed, in turn producing new pairs and thus initiating a cascade. In this case, apart the different resulting spectrum, the total flux of the reprocessed component would be larger than what simply estimated here: therefore, also in this case, our value of $B$ can considered a conservative value.  Finally, we note that in our derivation we are implicitly assuming that the magnetic field is oriented perpendicularly to the direction of the relativistic pairs. In reality the IGMF is probably randomly oriented, maintaining its coherence in domains with a size $\\sim$1 Mpc (e.g. Neronov \\& Semikoz 2009). Including the geometry of the field in our calculations would result in a limit to the IGMF slightly larger than that derived here. Again, our inferred value of $B$ is then a conservative one."
    },
    "1004/1004.3606_arXiv.txt": {
        "abstract": "We report on a confirmed galaxy cluster at $z=1.62$. We discovered two concentrations of galaxies at $z\\sim1.6$ in the Subaru/XMM-Newton deep field based on deep multi-band photometric data. We made a near-IR spectroscopic follow-up observation of them and confirmed several massive galaxies at $z=1.62$. One of the two is associated with an extended X-ray emission at $4.5\\sigma$ on a scale of $0'.5$, which is typical of high-$z$ clusters. The X-ray detection suggests that it is a gravitationally bound system. The other one shows a hint of an X-ray signal, but only at $1.5\\sigma$, and we obtained only one secure redshift at $z=1.62$. We are not yet sure if this is a collapsed system. The possible twins exhibit a clear red sequence at $K<22$ and seem to host relatively few number of faint red galaxies. Massive red galaxies are likely old galaxies -- they have colors consistent with the formation redshift of $z_f=3$ and a spectral fit of the brightest confirmed member yields an age of $1.8_{-0.2}^{+0.1}$ Gyr with a mass of $2.5_{-0.1}^{+0.2}\\times10^{11}\\rm M_\\odot$. Our results show that it is feasible to detect clusters at $z>1.5$ in X-rays and also to perform detailed analysis of galaxies in them with the existing near-IR facilities on large telescopes. ",
        "introduction": "The galaxy evolution is closely linked to the structure formation of the Universe. Isolated galaxies and those in clusters evolve in different ways and the differential evolution results in the strong environmental dependence of galaxy properties observed in the local Universe. The environmental dependence of galaxy properties is already strong at $z=1$.  Intensive studies of $z\\sim1$ clusters have shown that red galaxies have already become the dominant population in clusters at $z=1$ (e.g., \\citealt{blakeslee03,nakata05,lidman08,mei09}). The spectroscopically confirmed highest redshift X-ray cluster known until now is located at $z=1.45$ \\citep{stanford06}, but red galaxies are abundant even in this highest-$z$ cluster \\citep{hilton09}. A few authors reported a possible clusters at $z>1.5$. For example, \\citet{kurk09} presented an over-density of galaxies at $z=1.6$, which may later collapse to a cluster. It is not detected in X-rays down to a limit of $<1\\times10^{-16}\\rm ergs\\ s^{-1}\\ cm^{-2}$. \\citet{andreon09} presented a cluster at $z_{phot}\\sim1.9$, but its redshift has been questioned by Bielby et al. (in prep) --- it is likely a complex system with one confirmed at $z=1.1$ and another one possibly at $z\\sim1.5$ (see Bielby et al. for details). It is challenging to identify clusters at $z>1.5$, but one has to explore this redshift regime to identify the epoch when red galaxies become the dominant population in clusters and how the environmental dependence of galaxy properties is established.   We are conducting a distant X-ray cluster survey in the Subaru/XMM-Newton Deep Field (SXDF). We refer the reader to \\citet{finoguenov09b} for details of our survey and the construction of the deep multi-band photometric catalog\\footnote{ We note that we have revised the X-ray catalog presented in \\citet{finoguenov09b} by adding the 0.3-0.5 keV band data from XMM. We extract the images separately using the single events, produce a corresponding background estimation, and add them to the mosaic only at the very last stage. This increases signals from distant low-mass clusters. }. As part of the survey, we identified two concentrations of red galaxies first by their red sequence and then we found one of the sources is also securely detected in X-rays. We estimated their redshifts to be $z\\sim1.6$ both from the location of the cluster red sequence and from photometric redshifts. Not only they have very similar redshifts, but they are also close to each other on the sky; they are separated only by $\\sim2.5$ arcmin ($\\sim1.3$ Mpc at $z=1.6$). Followed by the initial photometric identification, we carried out a near-IR spectroscopic follow-up observation of the possible twin clusters and we report on the results in the paper. Recently, \\citet{papovich10} presented a completely independent study of one of the clusters. Readers are referred to their paper for a similar, but independent analysis.   The layout of the paper is as follows.  We first describe the near-infrared spectroscopic follow-up observation and present a discovery of the most distant X-ray cluster from the observation in section 2. Section 3 discusses our results and summarizes the paper. Unless otherwise stated, we adopt H$_0=72\\rm km\\ s^{-1}\\ Mpc^{-1}$, $\\Omega_{\\rm M}=0.25$, and $\\Omega_\\Lambda =0.75$. Magnitudes are on the AB system. ",
        "conclusions": "The cluster red sequence is a ubiquitous feature of galaxy clusters. We draw color-magnitude diagrams in Fig. \\ref{fig:cmd} using galaxies at $<1$ Mpc from the twins. We refer the readers to \\citet{papovich10} for a similar analysis. The photo-$z$ selected galaxies form a clear red sequence at $K<22$, while the sequence is not very clear at fainter magnitudes. We model the location of the red sequence with the updated version of the \\citet{bruzual03} code, which takes into account effects of thermally pulsating AGB stars, using the procedure described in \\citet{lidman08}. Here we assume the Chabrier IMF \\citep{chabrier03}. A biweight fit to the red galaxies gives a sequence very close to the $z_f=3$ model sequence, suggesting that they are relatively old galaxies. The cosmic time between $z=3$ and 1.6 is 2.0 Gyr. We note that we obtain $z_f=3$ or higher if we use the $z-J$ color as used in \\citet{papovich10} who obtained $z_f=2.4$. We do not know the cause of the difference in $z_f$ at this point, but it could be due to the rather old UKIDSS catalog they used. Their paper is still under review at the time of this writing and we do not pursue this point further. The $z-K$ color is more sensitive to $z_f$ at this redshift\\footnote{ The $z-K$ color of the model red sequence formed at $z_f=5$ is redder than the $z_f=2$ sequence by 0.74, while the difference is only 0.26 for the $z-J$ color. Thus, the $z-K$ color is more sensitive to $z_f$ than $z-J$. } and we prefer to use it in this paper. As a sanity check, we perform a statistical subtraction of fore-/background galaxies without using photo-$z$ following the recipe by \\citet{tanaka05} in the right panel. A clump of red galaxies is clearly visible in the right panel as well. Again, the sequence vanishes at $K>22$. There are a number of proto-clusters at even higher redshift that have a confirmed over-density of low-mass star forming galaxies \\citep{miley08}, but Fig \\ref{fig:cmd} shows that a near-IR spectroscopic observation is essential to confirm the dominant population of massive galaxies in $z>1.5$ clusters.   To further quantify the red sequence, we plot luminosity functions (LFs) of red galaxies in Fig. \\ref{fig:lf}. Here we define red galaxies as those having $\\Delta |z-K|<0.2$ from the best fitting red sequence. We confirmed that a small change in the definition does not alter the result below. But, we should note that the current statistics is very poor. We use galaxies at $1.4<z_{phot}<1.8$ and statistically subtract the remaining contamination using the entire SXDF field as a control field sample. We compare the $z=1.6$ LF with those at lower redshifts taken from \\citet{tanaka08}. We plot LFs relative to $m_K^*$ at each redshift to correct for the passive evolution. The LF of the red galaxies in the twins is similar to the group LF at $z=1.1$, which suggests a deficit of faint red galaxies.  On the other hand, it seems that clusters at $z=1.6$ already host abundant massive galaxies. As seen in Fig. \\ref{fig:cmd}, these massive galaxies are mostly red galaxies and the clusters host few bright blue galaxies. It is hard to generalize the trend we see here with only a few clusters, but the implication  would be that massive red galaxies might have become the dominant population in clusters already at $z=1.6$.   Finally, we present a spectral analysis of the brightest red galaxy with secure redshift in Fig. \\ref{fig:specfit}. We fit the observed spectra with model spectra generated by the updated version of the \\citet{bruzual03} code.  We assume the exponentially decaying star formation histories and use age, dust extinction, star formation time scale, and metallicity as free parameters. We adopt the Chabrier IMF \\citep{chabrier03}\\footnote{ If we use the Salpeter IMF, we obtain 1.8 times higher stellar mass. The other properties are nearly the same. } and fix the redshift at $z_{spec}$. The fit gives an age of $1.8_{-0.2}^{+0.1}$ Gyr and a small extinction of $A_V=0.2_{-0.1}^{+0.2}$ mag. Its star formation time scale is rather short, $0.1_{-0.1}^{+0.2}$ Gyr. Assuming the exponential decay, it assembled 80\\% of its stars at $z_f\\sim2.5$ that it would have at $z=0$, which is in agreement with that inferred from the $z-K$ color. Note that $z_f=3$ obtained from the red sequence gives an average formation redshift of the red galaxies, while we are specifically discussing the brightest red galaxy with secure redshift here. The current star formation rate is $<0.1\\rm M_\\odot\\ yr^{-1}$, which is consistent with the absence of any emission lines in the observed spectrum. Its mass is estimated to be $2.5_{-0.1}^{+0.2}\\times10^{11}\\rm  M_\\odot$, revealing the presence of such a massive galaxy in a $z=1.6$ cluster. As seen in lower redshift clusters, massive cluster galaxies form early and evolve passively  (e.g., \\citealt{thomas05}).   This study is the first result on cluster galaxies at $z>1.5$ using near-IR spectrograph. Our observation suggests that $z=1.6$ clusters exhibit a red sequence at bright magnitudes and are likely luminous in X-rays, which may have an implication for future high-$z$ cluster surveys. The cluster we discovered will evolve to a 4--5 keV cluster at $z=0$ \\citep{vandenbosch02}. Such clusters are a sensitive probe of cosmology at high redshifts and yet its estimated temperature preclude a detection in a Sunyaev-Zeldovich observation. This demonstrates the power of X-ray observations in finding distant clusters. The redshift of $z>1.5$ is the epoch when early dark energy is important and large-scale structures at $z>1.5$ are a robust tracer of the primordial local non-Gaussianity (e.g., \\citealt{bartelmann06,sadeh07,grossi09}). More massive clusters than the one reported in this paper may be found in shallower X-ray observations down to $10^{-14}\\rm ergs\\ s^{-1}\\ cm^{-2}$, but one needs to survey an order of $5,000$ square degrees \\citep{finoguenov09b}, which may be difficult to follow-up photometrically. However, a similar yield in the number of clusters can be achieved by reaching $10^{-15}\\rm ergs\\ s^{-1}\\ cm^{-2}$ over 100 square degrees. Therefore, deep X-ray surveys at high spatial resolution have a unique window for cosmological studies at the time when Universe was only a third of its present age. Our observation also shows that detailed analysis of galaxy populations at this redshift regime is feasible with the current near-IR facilities. This is an encouraging result and motivates us to push our X-ray cluster survey forward."
    },
    "1004/1004.0952_arXiv.txt": {
        "abstract": "We present low- and high-resolution mid-infrared spectra and photometry for eight compact symmetric objects (CSOs) taken with the Infrared Spectrograph on the \\textit{Spitzer Space Telescope}. The hosts of these young, powerful radio galaxies show significant diversity in their mid-IR spectra. This includes multiple atomic fine-structure lines, \\htwo~gas, PAH emission, warm dust from T~=~50 to 150~K, and silicate features in both emission and absorption. There is no evidence in the mid-IR of a single template for CSO hosts, but 5/8 galaxies show similar moderate levels of star formation ($<10M_\\sun$~yr$^{-1}$ from PAH emission) and silicate dust in a clumpy torus. The total amount of extinction ranges from $A_V\\sim$~10~to~30, and the high-ionization \\neV~14.3 and 24.3~\\um~transitions are not detected for any galaxy in the sample. Almost all CSOs show contributions both from star formation and active galactic nuclei (AGNs), suggesting that they occupy a continuum between pure starbursts and AGNs. This is consistent with the hypothesis that radio galaxies are created following a galactic merger; the timing of the radio activity onset means that contributions to the infrared luminosity from both merger-induced star formation and the central AGN are likely. Bondi accretion is capable of powering the radio jets for almost all CSOs in the sample; the lack of \\neV~emission suggests an advection-dominated accretion flow (ADAF) mode as a possible candidate. Merging black holes (BHs) with $M_{BH}>10^8M_\\sun$ likely exist in all of the CSOs in the sample; however, there is no direct evidence from this data that BH spin energy is being tapped as an alternative mode for powering the radio jets. ",
        "introduction": "An enduring mystery in the study of radio galaxies is the nature of their power source. Specifically, most galaxies which possess a radio-luminous active galactic nucleus (AGN) and rapidly advancing jets show little evidence for a central accreting disk. For example, in Centaurus A and Virgo A (M87) there is no blue continuum or broad emission lines, but also no IR-emitting screen large enough to obscure a direct view of a typical broad emission line region \\citep{per01a,why04,rad08}. While this modern paradigm unifies all AGN as supermassive black holes (BH) fed by an accretion disk orbiting the BH, many AGN show little evidence for the large amounts of circumnuclear gas needed to power them. These and other considerations led to the proposal that some AGN could be dominated by ``advection-dominated accretion flows'' \\citep[ADAFs; e.g.][]{nar95}.  An alternative is that the active phase may be triggered by a major merger of disk galaxies and, concurrently, a merger of their supermassive BHs. It is now understood on the basis of detailed numerical simulations that the merger of disk galaxies can create a relatively normal elliptical galaxy \\citep{bar92}. Since most (if not all) bright disk galaxies contain supermassive BHs, \\citet{wil95} proposed a scenario for the creation of a radio-loud AGN in which the merger of two disk galaxies and their BHs can spin up the resulting BH, thus powering the AGN and its jets. The detection of X-ray point sources in radio-loud ellipticals provides observational evidence for the merging BH scenario \\citep{kom03,bia08}. BH spin accounts naturally for the diminished output of successive radio outbursts in the brightest cluster galaxy (BCG) because the spin energy is diminished by each outburst and the accretion rate is too low for replenishment. At high redshift a merger of disk galaxies (a so-called ``wet merger'') and their individual supermassive BHs creates a rapidly-spinning supermassive BH and a luminous quasar.  As the cluster or group of galaxies in which the BCG is imbedded develops a hot, dense ICM, a large fraction of cluster galaxies are stripped of their cold and warm gas. As a result, all such mergers are ``dry'' and provide no additional accretion power. The BH spin hypothesis thus predicts the correct sign of radio-loud AGN evolution, which is not explained by the simplest accretion scenarios.  An ideal class of galaxies in which to test these competing hypotheses are compact symmetric objects (CSOs). CSOs are radio-loud galaxies characterized by jet and/or hot spot activity on both sides of a central engine. They are morphologically similar to classical double radio sources, but with size scales of $<1$~kpc \\citep{phi82,wil94}. The radio emission typically peaks near frequencies of a few GHz, which can be explained in most CSOs by synchrotron self-absorption in the galaxies \\citep{dev09a,dev09}. Kinematic ages of the radio jets show they are extremely young, with ages of $\\leq4000$~yr \\citep{ows98,gug05}. As a result, it has been suggested that CSOs represent the early stages of an AGN evolving into an FR~II radio galaxy \\citep{rea96a}, akin to a ``mini-Cygnus A'' \\citep{beg96}. Given the young ages of these galaxies, they offer a much better chance for detecting signs of a recent merger that may be powering the radio activity.  HST imaging of three nearby CSO host galaxies \\citep{per01} show them to be nearly normal ellipticals; however, disturbed isophotes at the outer edges suggest that a past merger has taken place in all objects \\citep{per01}. Simulations suggest that the radius of the photometric irregularities determines the timescale since the merger due to virialization of the stellar distribution \\citep{mih95,mih96}. This is evidence for a recent merger in all three CSOs which commenced $\\sim100$~million years ago, the time delay required for supermassive BHs to merge \\citep{beg84}. This is also comparable to the timescale for driving gas to the center of the merger suggested by other simulations \\citep{dim05,dim08,hop05,spr05,spr05a}. While this timescale supports the Wilson-Colbert hypothesis, it offers no concrete proof that radio-loud AGN are powered by BH spin. The accretion model requires large amounts of cold or warm gas close to the CSO, however, which is not seen - therefore, the power source for the jets is still an open question.  Since distinguishing between a model for radio-loud AGN fueled by BH spin or one fueled by accretion will significantly affect our understanding of jet production and the differences between radio-loud and radio-quiet AGN, CSOs provide a unique opportunity to investigate the nuclear conditions in radio galaxies recently triggered by some mechanism into activity. One key question is whether large amounts of nuclear obscuration unresolved in the HST optical images could hide the accreting gas. Mid-IR observations should detect this gas if it exists, which may provide strong supporting evidence for accretion flows leading to AGN outbursts in CSOs. Observations like the HST imaging and the mid-IR results presented in this paper are the first that address such a hypothesis.  To this end we observed a sample of nearby CSOs with the {\\it Spitzer Space Telescope} with the goal of distinguishing between these two hypotheses and characterizing the properties of CSO host galaxies, including gas and dust content and their star formation rates. If accretion is the power source for these galaxies, then mid-IR photons should penetrate any extinction screen that might obscure the AGN at shorter wavelengths; accreting gas should also emit via high-ionization atomic transitions due to the proximity of the AGN.  Absence of this would suggest a ``naked'' AGN with no luminous accreting gas, supporting a model in which radio-loud AGN are powered by BH spin due to either rapidly spinning BHs or binary BHs created in a merger.  \\begin{deluxetable*}{lcclcrclccrc} \\tabletypesize{\\scriptsize} \\tablecaption{Properties of CSOs observed with Spitzer IRS \\label{tbl-csosum}} \\tablewidth{0pt} \\tablehead{ \\colhead{Object} & \\colhead{RA} & \\colhead{Dec} & \\colhead{z$_\\sun$} & \\colhead{D$_L$} & \\colhead{H{\\scriptsize{I}} column} & \\colhead{H{\\scriptsize{I}}} & \\colhead{Optical} & \\colhead{Optical} & \\colhead{log~P$_{radio}$} & \\colhead{log~L$_{IR}$} & \\colhead{log~L$_{X}$} \\\\ \\colhead{} & \\colhead{J2000.0} & \\colhead{J2000.0} & \\colhead{} & \\colhead{[Mpc]} & \\colhead{[cm$^{-2}$]} & \\colhead{ref.} & \\colhead{spec.} & \\colhead{ref.} & \\colhead{[W/Hz]} & \\colhead{[L$_\\odot$]} & \\colhead{[L$_\\odot$]} } \\startdata 4C~+31.04                   &  01 19 35.0 & +32 10 50   & 0.0602 & 264  & $1.08\\times10^{21}$    & (1)    & WLRG   & (8)    & 25.34 & 10.60  & \\ldots \\\\ 4C~+37.11                   &  04 05 49.2 & +38 03 32   & 0.055  & 242  & $1.8\\times10^{20}$     & (2)    & NLRG   & (9)    & 25.05 & \\ldots & \\ldots \\\\ 1146+59                     &  11 48 50.3 & +59 24 56   & 0.0108 & 48.2 & $1.82\\times10^{21}$    & (1)    & LINER  & (10)   & 23.10 & 9.09   & \\ldots \\\\ 1245+676\\tablenotemark{a}   &  12 47 33.3 & +67 23 16   & 0.1073 & 495  & $6.73\\times10^{20}$    & (3)    & WLRG   & (8)    & 25.02 & \\ldots & \\ldots \\\\ 4C~+12.50                   &  13 47 33.3 & +12 17 24   & 0.1217 & 571  & $T_s 6.2\\times10^{18}$ & (4)    & NLRG   & (11)   & 26.30 & 12.15  & 43.3   \\\\ OQ~208                      &  14 07 00.4 & +28 27 15   & 0.0766 & 349  & $1.83\\times10^{20}$    & (5)    & BLRG   & (12)   & 25.08 & 11.33  & 42.7   \\\\ PKS~1413+135                &  14 15 58.8 & +13 20 24   & 0.2467 & 1244 & $4.6\\times10^{22}$     & (6)    & BL Lac & (13)   & 26.31 & 11.97  & 44.4   \\\\ PKS~1718-649                &  17 23 41.0 & $-$65 00 37 & 0.0142 & 62.6 & \\ldots                 & \\ldots & LINER  & (14)   & 24.25 & \\ldots & \\ldots \\\\ 1946+70                     &  19 45 53.5 & +70 55 49   & 0.1008 & 460  & $2.2\\times10^{23}$     & (7)    & \\ldots & (15)   & 25.39 & \\ldots & \\ldots \\\\ \\enddata \\tablenotetext{a}{Undetected with the spectral modules on the IRS; see \\S\\ref{ssec-no1245}.} \\tablerefs{ (1) - \\citet{van89}; (2) - \\citet{man04}; (3) - \\citet{sai07}; (4) - \\citet{mir89}; (5) - \\citet{ver03a}; (6) - \\citet{per02}; (7) - \\citet{pec99}; (8) - \\citet{mar96}; (9) - \\citet{rod06}; (10) - \\citet{kim89}; (11) - \\citet{baa98}; (12) - \\citet{pet07}; (13) - \\citet{sto92}; (14) - \\citet{fil85}; (15) - \\citet{sne99}} \\end{deluxetable*}  In \\S\\ref{sec-obs}, we describe the CSO sample and the observations taken with {\\it Spitzer}; we describe the data reduction process in \\S\\ref{sec-reduction}. \\S\\ref{sec-results} presents results of the observations, including full spectra and data tables for all measured features. We derive physical properties for the CSO hosts in \\S\\ref{sec-diagnostics} and classify them in the mid-IR. In \\S\\ref{sec-discussion} we discuss the evidence for the power source of CSOs and their role in the evolution of radio galaxies. An Appendix summarizes the jet properties and multi-wavelength observations of the host galaxies from the literature. We assume the WMAP5 cosmology with $H_0=70.5$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_M=0.27$, and $\\Omega_\\Lambda=0.73$. ",
        "conclusions": "We observed a sample of eight low-$z$ CSOs with the IRS, representing very young radio galaxies and possible progenitors to FR~IIs. {\\it Spitzer} investigations have revealed:  \\begin{enumerate} \\item The mid-IR SEDs for CSOs show significant variety. This includes silicate features seen in both emission and absorption, spectral indices indicating dust at a wide range of temperatures, and fine-structure and PAH emission. We sort the galaxies into four categories based on their IRS spectra, with the most common features including weak 9.7~\\um~dust absorption, moderate PAH and atomic emission, and \\htwo~temperatures of 200--400~K. There is no evidence in the mid-IR, however, of a single ``family'' for CSOs. \\item PKS~1413+135 is a unique object among CSOs, showing no emission or absorption features associated with gas in the ISM. The only identifiable features in the mid-IR are the silicate dust troughs at 9.7 and 18~\\um. While the dust redshifts are roughly consistent with the near-IR obscuration, it does not address the possibility that the AGN lies along the line of sight and may not be physically associated with the optical counterpart \\citep{per02}. \\item Numerical models indicate that a clumpy dust torus can fit the silicate features in most CSO spectra. Features common to most models have 1--10 dust clouds along the line of sight, with optical depths reaching as high as a few hundred in $\\tau_V$. \\item Multi-wavelength evidence (including X-ray fluxes, silicate/PAH ratios, and radio-imaging of SMBHs) indicate that significant fractions of the sample have AGN at their centers. Evidence for continuing star formation, however, suggests that CSOs occupy a continuum between starburst galaxies and AGN. This supports the scenario that the elliptical hosts were formed by a merger of disk galaxies with a delay before the activation of an AGN. \\item Based on the work necessary to expand the radio jet cavities, accretion is a viable power source for 6/8 CSOs in the sample. The non-detection of high-ionization lines and absence of large extinction support an ADAF model, if this is the case. Although BHs with $M>10^8M_\\sun$ likely exist in all CSOs, there is no direct evidence that BH spin energy provides the power source for the radio jets. \\end{enumerate}"
    },
    "1004/1004.2166_arXiv.txt": {
        "abstract": "Over the last two decades the Andromeda Galaxy (M31) has been something of a test-bed for methods aimed at obtaining accurate time-domain relative photometry within highly crowded fields. Difference imaging methods, originally pioneered towards M31, have evolved into sophisticated methods, such as the Optimal Image Subtraction (OIS) method of \\cite{ala98}, that today are most widely used to survey variable stars, transients and microlensing events in our own Galaxy. We show that modern difference image (DIA) algorithms such as OIS, whilst spectacularly successful towards the Milky Way bulge, may perform badly towards high surface brightness targets such as the M31 bulge. Poor results can occur in the presence of common systematics which add spurious flux contributions to images, such as internal reflections, scattered light or fringing. Using data from the Angstrom Project microlensing survey of the M31 bulge, we show that very good results are usually obtainable by first performing careful photometric alignment prior to using OIS to perform point-spread function (PSF) matching. This separation of background matching and PSF matching, a common feature of earlier M31 photometry techniques, allows us to take full advantage of the powerful PSF matching flexibility offered by OIS towards high surface brightness targets. We find that difference images produced this way have noise distributions close to Gaussian, showing significant improvement upon results achieved using OIS alone. We show that with this correction light-curves of variable stars and transients can be recovered to within $\\sim 10$ arcseconds of the M31 nucleus. Our method is simple to implement and is quick enough to be incorporated within real-time DIA pipelines. We also demonstrate that OIS is remarkably robust even when, as in the case of the central regions of the M31 bulge, the sky density of variable sources approaches the confusion limit. ",
        "introduction": "Difference Image Analysis (DIA) is now used routinely to provide very accurate relative photometry of variable stars, transient objects and microlensing events in the Milky Way and other nearby galaxies \\citep{woz08}. DIA permits very accurate relative photometry even within extremely dense stellar fields where conventional photometric methods can fail or suffer from serious bias.  Most DIA pipelines currently in use derive from the Optimal Image Subtraction (OIS) algorithm of \\cite{ala98} and \\cite{ala00}. The OIS algorithm has found widest use in providing accurate relative photometry of stars in the Milky Way and the Magellanic Clouds. Similar schemes that pre-date OIS were originally employed to look for microlensing towards the Andromeda Galaxy \\citep{tom96,ans97}. Since the stars in the Andromeda Galaxy (M31) are two orders of magnitude further away than those typically monitored in our Galaxy, difference imaging towards M31 throws up several additional challenges to the standard technique.  In this paper we show that towards the bulge of M31, and similarly towards other targets where diffuse background surface brightness dominates the total flux, DIA pipelines based on the OIS algorithm can often yield poor results due to common image systematics such as internal reflections, scattered light or fringe effects. Systematics, which may appear at a low level ($\\sim 1\\%$) on the original exposures, can give rise to large-amplitude differential background residuals on difference image frames. Since OIS minimizes mismatches in both the point spread function (PSF) and the differential background simultaneously, poor background matching often results in poor PSF matching and therefore substantial systematic errors in differential photometry. We propose a straightforward remedy to allow the effects of such systematics to be minimized.  The structure of the paper is as follows. In Section~\\ref{dialimits} we briefly describe how M31 has been used as a test-bed in developing time-domain photometry towards crowded stellar fields. The evolution of these methods culminated in the Optimal Image Subtraction (OIS) algorithm of \\cite{ala98} and \\cite{ala00}, which forms the basis for most DIA pipelines currently in use.  We use images obtained by the Angstrom Project \\citep{ker06} of the bulge region of M31 to show how the OIS algorithm may not perform optimally in the presence of bright backgrounds. In Section~\\ref{moddia} we show how difference images with noise levels close to the photon noise limit can be recovered by separating the photometric alignment and PSF matching stages. In Section~\\ref{examplelc} we show some example periodic variable light-curves from the Angstrom Project dataset, illustrating the impact of the correction on their photometric quality and on the ability to characterise variable stars at a range of distances from the M31 core. We discuss the findings in Section~\\ref{discuss}. ",
        "conclusions": "\\label{discuss}  In this paper we have shown that optimal difference imaging in regions of very high background levels, such as the bulge regions of galaxies, can be severely compromised due to the presence of systematics such as internal reflections, scattered light,  or fringe effects. In these cases difference images created using the Optimal Image Subtraction (OIS) algorithm \\citep{ala98,ala00}, which is the most widely used difference image algorithm, may exhibit large amplitude background residuals that make reliable relative photometry difficult if not impossible. Fortunately, we have shown that OIS is able to give very good results provided the images to be differenced are first photometrically aligned prior to difference imaging.  Using the photometric alignment procedure described in this paper we find we can produce difference images of the M31 bulge that are close to photon noise limited. Not only does it minimize or eliminate large amplitude background residuals but it also noticeably improves the quality of the PSF kernel transformation. We have tested the modified pipeline by producing several examples of periodic variable light-curves from the Angstrom Project survey of the M31 bulge. The results allow characterisation of periodic signals even to within $\\sim 10$ arcseconds of the M31 nucleus.  The problem we highlight is specific to targets in which the background brightness is high and subject to systematic variations which have complex spatial signatures. The OIS method is well known to cope admirably for stellar fields within our Galaxy, as imaged by current optical surveys, where the image flux is dominated by resolved or semi-resolved stars rather than by the diffuse background light. However, future near-infrared time-domain surveys of the Galactic Centre could also conceivably benefit from a separation of the photometric and kernel matching stages."
    },
    "1004/1004.3998_arXiv.txt": {
        "abstract": "\\noindent A generic prediction in the paradigm of weakly interacting dark matter is the production of relativistic particles from dark matter pair-annihilation in regions of high dark matter density. Ultra-relativistic electrons and positrons produced in the center of the Galaxy by dark matter annihilation should produce a diffuse synchrotron emission. While the spectral shape of the synchrotron dark matter haze depends on the particle model (and secondarily on the galactic magnetic fields), the morphology of the haze depends primarily on (1) the dark matter density distribution, (2) the galactic magnetic field morphology, and (3) the diffusion model for high-energy cosmic-ray leptons. Interestingly, an unidentified excess of microwave radiation with characteristics similar to those predicted by dark matter models has been claimed to exist near the galactic center region in the data reported by the WMAP satellite, and dubbed the ``WMAP haze''. In this study, we carry out a self-consistent treatment of the variables enumerated above, enforcing constraints from the available data on cosmic rays, radio surveys and diffuse gamma rays. We outline and make predictions for the general morphology and spectral features of a ``dark matter haze'' and we compare them to the WMAP haze data. We also characterize and study the spectrum and spatial distribution of the inverse Compton emission resulting from the same population of energetic electrons and positrons. We point out that the spectrum and morphology of the radio emission at different frequencies is a powerful diagnostics to test whether a galactic synchrotron haze indeed originates from dark matter annihilation. ",
        "introduction": "A compelling paradigm for the particle nature of the dark matter is that of Weakly Interacting Massive Particles, or WIMPs \\cite{Bertone:2004pz,Bergstrom:2009ib}. Although the Standard Model of particle physics does not encompass a viable particle dark matter candidate, WIMPs are predicted to exist in several well motivated extensions. These include weak-scale supersymmetry \\cite{Jungman:1995df}, models with universal extra dimensions \\cite{Hooper:2007qk}, and many others (for reviews see \\cite{Bertone:2004pz,Bergstrom:2009ib}). WIMPs are massive particles with masses near the electro-weak scale, and are typically charged under weak interactions. General arguments indicate that WIMPs in thermal equilibrium in the very early universe would {\\em freeze-out}, decoupling from the thermal bath when the temperature dropped below a fraction (typically $\\sim$1/20) of their mass \\cite{Bertone:2004pz}. The remaining dark matter particles would then populate the universe with a relic density which is of the same order as dark matter on cosmological scales \\cite{Lee:1977ua,Krauss:1983ik,Kolb:1985nn,Scherrer:1985zt}. This ``WIMP miracle'' \\cite{Steigman:1979kw} warrants extensive investigation, due to the possibility of observing the particle debris stemming from the occasional dark-matter pair-annihilation event in today's cold universe.  The probability of two dark matter particles ($\\chi$) pair-annihilating into observable standard model particles is proportional to the thermally averaged pair-annihilation cross section times the relative particle velocity, $\\langle\\sigma v\\rangle$, multiplied by the local particle dark matter number density squared, $n_\\chi^2$. The first quantity, $\\langle\\sigma v\\rangle$,  can be inferred from the requirement of having a relic abundance $\\Omega_\\chi\\sim1/\\langle\\sigma v\\rangle$ on the same order as the universal dark matter density, i.e. $\\Omega_{\\rm DM}\\sim0.24$ \\cite{Komatsu:2008hk}. The latter quantity is given by $n^2_\\chi=(\\rho_{\\rm DM}/m_\\chi)^2$, where $m_\\chi$ indicates the mass of the dark matter particle $\\chi$. Most dark matter annihilation processes are therefore predicted to occur in regions with a large dark matter density. Both intuition and detailed results from N-body simulations (see e.g. \\cite{Diemand:2006ik, Diemand:2008in,Springel:2008by}) indicate that the center of the Milky Way galaxy is likely the brightest local dark matter annihilation site (barring the possibility of highly concentrated local dark matter clumps \\cite{Brun:2009aj}). Among the possible signatures of dark matter annihilation, extensive studies have focused on gamma rays (see e.g. \\cite{Baltz:2008wd,Jeltema:2008hf} and references therein) and neutrinos (for a recent study see \\cite{Spolyar:2009kx}), particle species which have the benefit of carrying directional and spectral information.  Other stable standard model particles produced in the pair annihilation of dark matter include charged cosmic rays such as electrons and positrons ($e^\\pm$) as well as (anti-)protons and (anti-)deuterons. These charged species scatter off of magnetic field irregularities in the Galaxy, losing energy and diffusing before reaching the Earth. The random-walk propagation of charged cosmic rays in the Galaxy is usually described with a diffusion-loss equation and solved numerically \\cite{Strong:1998pw} or semi-analytically \\cite{Baltz:1998xv,Regis:2009qt}. A pedagogical review of cosmic-ray propagation and interactions in the Galaxy is given in Ref.~\\cite{Strong:2007nh}.  The possibility of detecting an anomalous spectral feature in the flux of leptons, which could be traced back to the pair-annihilation of particle dark matter, has been long discussed (for early studies see e.g. Ref.~\\cite{Rudaz:1987ry,Ellis:1988qp,Kamionkowski:1990ty,Profumo:2004ty}). This scenario has recently gained great momentum after results from the Pamela space-based antimatter detector reported an excess of high-energy (10-100 GeV) positrons over the assumed background of secondary positron production by inelastic cosmic-ray interactions \\cite{Adriani:2008zr}. Tantalizingly, an excess (namely an anomalous ``bump'') in the flux of electrons plus positrons was also recently reported by the balloon-borne experiments ATIC \\cite{2008Natur.456..362C} and PPB-BETS \\cite{Torii:2008xu} at energies in the range of several hundred GeV. This excess was subsequently not confirmed by data from the Fermi Large Area Telescope (LAT), which, however, did indicate a much harder spectrum for the $e^\\pm$ flux at high energy than previously assumed \\cite{Abdo:2009zk,Grasso:2009ma}. This implies that the positron deficit reported by Pamela is at an even greater contrast with the standard expectation from cosmic ray models. The Fermi results agree with the low-energy range of other determinations of the $e^\\pm$ flux \\cite{Collaboration:2008aa}.  Although astrophysical sources such as pulsars \\cite{Hooper:2008kg,Yuksel:2008rf,Profumo:2008ms,Malyshev:2009tw,Grasso:2009ma,Gendelev:2010fd} and supernova remnants \\cite{Shaviv:2009bu,Blasi:2009hv} have been shown to provide a possible explanation for the positron fraction anomaly, dark matter models have been formulated which can fit both cosmic ray data and account for the positron excess (for a list of references see e.g. \\cite{Profumo:2008ms}). In general, dark matter models that account for the Pamela anomaly need to have a large pair-annihilation cross section compared to the estimates from the standard thermal-relic calculation.  Furthermore, a mechanism or conservation law must be invoked to ensure that excess antiprotons are not produced at a detectable level, as the Pamela data place stringent constraints on any antiproton excess \\cite{Adriani:2008zq}. Dark matter models which match both the Pamela and Fermi-LAT data commonly feature pair-annihilation modes dominated by positron production, and have rather large masses ($\\sim$~1~TeV). Recently, Ref.~\\cite{Kane:2009if} also pointed out that a wino-like candidate with a somewhat lighter mass can also account for the noted cosmic ray anomalies. We will entertain this possibility for one of the particle dark matter models considered below (although recent data on gamma-ray observations from local dwarf galaxies put stringent constraints on this category of models \\cite{Collaboration:2010ex}). Winos dominantly pair-annihilate into $W^+W^-$ pairs, with a significant subsequent production of energetic $e^\\pm$.  While the source term for the $e^\\pm$ injected by dark matter pair annihilation traces the dark matter density profile squared, the resulting non-thermal population of electrons and positrons settles to an equilibrium configuration only after losing energy and spatially diffusing through scattering off of both the regular and turbulent components of the galactic magnetic fields. At high energy, the dominant energy loss processes are synchrotron and inverse Compton scattering, which create secondary radiation at radio frequencies as well as X-ray to gamma-ray energies, respectively (for a recent review of multi-wavelength emissions from dark matter annihilation see Ref.~\\cite{Profumo:2010ya}). This secondary emission is a common denominator to any WIMP model. In particular, dark matter models that also explain the Pamela positron excess generically give very large signals, which are typically barely compatible with the extragalactic gamma-ray background, as shown e.g. in \\cite{Profumo:2009uf,Huetsi:2009ex,Belikov:2009cx}. Constraints on dark matter models from the secondary radio and X-ray emission from $e^\\pm$ produced in dark matter annihilation in extragalactic objects such as dwarfs and galaxy clusters has been studied in a number of recent analyses, see e.g. \\cite{Totani:2004gy,Colafrancesco:2005ji,Baltz:2006sv,Colafrancesco:2006he,Profumo:2008fy,PerezTorres:2008ug,Jeltema:2008ax}.  \\subsection{The role of synchrotron emission}  WIMP annihilation into charged leptons should produce synchrotron emission throughout a large spherical halo around the Milky Way. Placing the injection energy on the weak scale should create synchrotron emission at radio frequencies, since the galactic magnetic fields are of $\\mathcal{O}\\sim1-10\\mu$G. This fact was envisioned long ago, see e.g \\cite{Gondolo:2000pn,Bertone:2001jv,Aloisio:2004hy,Bergstrom:2006ny}, and, in conjunction with available radio data, was recently used to put severe constraints on the detectability of a gamma-ray signal from the galactic center \\cite{Regis:2008ij} as well as on the viability of dark matter models that could explain the Pamela excess \\cite{Bergstrom:2008ag}.  Interestingly, Finkbeiner pointed out in 2004 that a residual, unaccounted-for radio ``haze'' could actually be present in the data reported by the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{Finkbeiner:2004us}, and that this radio emission could in principle be explained precisely with the radio haze predicted in typical WIMP models.\\footnote{Throughout this paper we refer to the observed WMAP signal as the ``WMAP Haze\" and our simulated results as the ``dark matter haze\" or ``synchrotron haze\"} This point was further elaborated upon in Ref.~\\cite{Hooper:2007gi,Hooper:2007kb,Cholis:2008vb,Bottino:2008sv,Caceres:2008dr,Cumberbatch:2009ji, McQuinn:2010ju}, which essentially confirmed that the morphology and intensity of the haze from WIMP annihilation might match the observed radio excess in the WMAP data. Other studies have shown that pulsars can create an additional lepton component which adequately traces the morphology of the WMAP haze \\cite{Kaplinghat:2009ix}. We note, however, that recent analyses by the WMAP team finds no evidence for a haze emission from the galactic center region in the polarization data \\cite{Gold:2010fm}. Furthermore, the extraction of a haze may depend sensitively on the assumption that the 408 Mhz synchrotron map is morphologically similar to skymaps in the WMAP band \\citep{Mertsch:2010ga}. Due to both the anomalous positron results, and in view of the recent successful launch of the Planck satellite \\cite{Bouchet:2009tr}, there is considerable interest in models with a large annihilation rate into leptons. Thus, we consider this a timely moment to critically re-assess the diffuse radio emission from WIMP annihilation in the central regions of the Galaxy and to outline possible diagnostics that will allow to disentangle it from other possible astrophysical sources.  The secondary radio emission from dark matter annihilation depends on both the injection spectrum and the pair-annihilation rate of dark matter - quantities which are set by specifying the particle dark matter model at hand. In addition, the intensity and spatial distribution of the WIMP synchrotron halo depend on three astrophysical quantities: (1) the dark matter density distribution, (2) the diffusion setup, and (3) the magnetic field setup. While information on (1) can be gained from N-body simulations, with educated estimates for the role of baryons and stars in shaping the dark matter density distribution, a self-consistent guess for (2) and (3) is a difficult task. Thus, the production of accurate predictions involves not only factoring in as many astronomical and cosmic-ray data as possible, but also treating with accuracy the distribution of astrophysical cosmic-ray sources and the diffusion and energy losses of galactic cosmic rays.  The importance of the diffusion setup, including in particular the geometry of the diffusion zone, the magnitude of the diffusion coefficient and its rigidity dependence, as well as the spatial distribution of Galactic magnetic fields, cannot be understated in the context of calculating the radio emission from dark matter annihilation. For instance, the height of the diffusion zone determines at which galactic latitude the $e^\\pm$ population responsible for the radio emission leaks out and is artificially truncated by the boundary conditions imposed to the diffusion-loss equation; similarly, the spatial distribution of magnetic fields determines the intensity of the synchrotron emission, which scales quadratically with the magnitude of the magnetic fields.  In the present work we study the morphology of the synchrotron halo produced by electrons and positrons created via dark matter annihilation in the context of {\\em self-consistent diffusion models}. In particular, we interface the most complete publicly available cosmic-ray propagation code, {\\tt Galprop} \\cite{Strong:1998pw}, with the spectra $e^\\pm$ from dark matter annihilation calculated by the {\\tt DarkSUSY} package \\cite{Gondolo:2004sc}. Testing various WIMP annihilation pathways, we compare our observations to constraints from cosmic-ray observations (e.g. Boron to Carbon and $^{10}$Be to $^9$Be). Imposing this set of observational constraints on the diffusion setup forces the morphology of the dark matter ``haze'' to specific configurations, which we outline in detail and compare to the WMAP data residual as analyzed in Ref.~\\cite{Dobler:2007wv}.  Finally, we explore the diffuse inverse Compton emission from the up-scattering of the inter-stellar radiation field, which stems from the same population of dark-matter-induced $e^\\pm$ that would produce the radio haze. We then compare the predicted emission with the Fermi diffuse gamma-ray data, and note the ability for observations of $\\gamma$-ray emission to put strong constraints on models of the WMAP haze.   The outline of this study is as follows: in the next section we describe the ``baseline'' setup adopted for the benchmark particle dark matter models  (Sec.~\\ref{sec:dmmodels}) and for the set of galactic dark matter density profiles (Sec.~\\ref{sec:dmdensity}). In addition, we provide details on the diffusion setup for cosmic ray propagation we employ in this analysis (Sec.~\\ref{sec:diffusion}). We study how the morphology of the synchrotron emission from dark matter annihilation depends on the diffusion setup in Sec.~\\ref{sec:morpho}, on the dark matter density profile in Sec.~\\ref{sec:morphodm}, and on the spatial distribution of magnetic fields in Sec.~\\ref{sec:morphob}. We then study the impact of matching the adopted diffusion scenarios to cosmic ray data in Sec.~\\ref{sec:cosmicrays}, and outline the expected spectrum of radio emission and inverse Compton emission in the context of the self-consistent scenarios in Sec.~\\ref{sec:radio} and~\\ref{sec:invco}, respectively. Finally, we give an overview and a discussion of our results and predictions in Sec.~\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  WIMP annihilation in the Galaxy unavoidably yields high-energy electrons and positrons (barring the possibility of pair annihilations in neutrino pairs or particles outside the Standard Model). In turn, this implies a diffuse galactic emission, whose morphology is affected by both the distribution of dark matter and charged cosmic-ray transport in the Galaxy. In this study, we outlined in detail how the form and nature of the WIMP radio halo depends upon the particle model for dark matter as well upon astrophysical parameters such as the galactic magnetic fields and the parameters entering the propagation of cosmic rays. We pointed out how cosmic-ray data very significantly constrain the morphology of the WIMP radio haze, and pointed out the importance of other secondary emission channels, in particular Inverse Compton.  Given the interest in a possible excess emission at radio frequencies -- the WMAP haze -- we have shown that standard dark matter annihilation pathways and cosmic ray diffusion models do not provide a reasonable match for the morphology of the excess. Dark matter density profiles with high densities near the galactic center produce synchrotron morphologies which are dominated by the diffusion of charged leptons away from the galactic center. However, changes in the diffusion setup do not allow one to reproduce the morphology of the WMAP haze while remaining consistent with local observations of cosmic ray primary to secondary ratios. Models which lack a dense dark matter distribution in the galactic center (e.g. the Burkert profile) require large cross-section enhancements which are ruled out by Fermi observations.  We also found that the particle physics properties of dark matter models which seek to explain both the WMAP haze as well as higher energy observations such as those from the Fermi-LAT are beginning to be constrained. Low energy dark matter decay spectra such as our {\\bf Soft} model are unable to reproduce the intensity of the WMAP haze without overproducing the Fermi signal with $\\gamma$-rays produced via direct annihilation.  However many uncertainties still exist in the magnetic field morphology to allow viable matches for the observed excess in synchrotron emission. The magnetic field morphology primarily influences the synchrotron intensity and morphology, with very little effect on $\\gamma$-ray or cosmic ray constraints. The best constraints on galactic magnetic fields would come from careful examination of the WMAP dataset which would determine the morphology of any WMAP haze resulting from various magnetic field models. Such an analysis is beyond the scope of this study, but has been examined in both \\citet{Gold:2010fm} and \\citet{Dobler:2007wv}, with contrasting results. Furthermore, many parameters not mentioned in this work may be fine tuned to produce a reasonable match to all known constraints. However, the formidable size of this parameter space makes the existence of such matches unconvincing evidence that the WMAP haze arises from a dark matter source.  Further inquiry in the nature of an excess diffuse radio emission is certainly warranted, especially since e.g. any additional IC signal  in the high energy sky \\citep{Dobler:2009xz,Linden:2010ea} implies the production of some microwave residual detectable in the microwave sky. Diffuse radio emission, and constraints on it, can thus be regarded as a powerful multi-wavelength diagnostic of any claims of new anomalous signals. Lastly, future data collected by the Planck telescope \\cite{Bouchet:2009tr} will be imperative in determining the intensity of the WMAP haze, and in extending its spectrum to the higher energy regime."
    },
    "1004/1004.4430_arXiv.txt": {
        "abstract": "{Access to astronomical data through archives and VO is essential but does not solve all problems.  Availability of appropriate software for analyzing the data is often equally important for the efficiency with which a researcher can publish results.  A number of legacy systems (e.g. IRAF, MIDAS, Starlink, AIPS, Gipsy), as well as others now coming online are available but have very different user interfaces and may no longer be fully supported.  Users may need multiple systems or stand-alone packages to complete the full analysis which introduces significant overhead.  The OPTICON Network on `Future Astronomical Software Environments' and the USVAO have discussed these issues and have outlined a general architectural concept that solves many of the current problems in accessing software packages.  It foresees a layered structure with clear separation of astronomical code and IT infrastructure. By relying on modern IT concepts for messaging and distributed execution, it provides full scalability from desktops to clusters of computers.  A generic parameter passing mechanism and common interfaces will offer easy access to a wide range of astronomical software, including legacy packages, through a single scripting language such as Python.  A prototype based upon a proposed standard architecture is being developed as a proof-of-concept.  It will be followed by definition of standard interfaces as well as a reference implementation which can be evaluated by the user community.  For the long-term success of such an environment, stable interface specifications and adoption by major astronomical institutions as well as a reasonable level of support for the infrastructure are mandatory.  Development and maintenance of astronomical packages would follow an open-source, Internet concept.  }  \\FullConference{Accelerating the Rate of Astronomical Discovery, SpS5\\\\ August 11-14, 2009\\\\ Rio de Janeiro, Brazil}  \\begin{document} ",
        "introduction": "The efficiency and speed with which one can obtain new astronomical results depend upon a long chain of facilities and tools such as telescopes, instruments, data archives, and software packages.  Although many discoveries are made using new, more capable telescopes and instruments, one should not underestimate the importance of easy access to both archival data and state-of-the-art software.  The combination of multi-wavelength data (e.g. obtained through data archives) and application of new algorithms for analysis of data can lead to new understanding of astronomical phenomena.  In this paper, we focus on how access to and sharing of astronomical software can be made easier and thereby increase the speed with which new results are obtained. The prime concern is the efficiency of turning raw data into astronomically relevant information. First an overview of the current situation is given, followed by considerations on possible ways to improve it.  An architectural concept for a future astronomical software environment is then presented as discussed by OPTICON and the USVAO (formerly NVO).  Finally, steps needed to provide the astronomical community with an efficient software environment for future data challenges are reviewed. ",
        "conclusions": "The current situation with a multitude of largely incompatible legacy or more modern but narrowly focused systems for data processing and analysis is not satisfactory since it makes access to software packages more difficult and requires significant resources for maintenance. To remedy these problems, a new approach is required based upon the concept of a common astronomical software framework which allows flexible integration of both new and legacy code, which is scalable to leverage modern computational hardware architectures while meeting the challenge of exponentially increasing data volumes, and which separates astronomical code from the IT infrastructure to provide stability for critical  astronomical code while allowing rapid uptake of new IT technologies.  The requirements and architectural concept for such an environment have been discussed within an OPTICON Network with US and EU participation. It is proposed to adopt a modular, layered model based upon a framework supporting messaging and execution in a distributed system. Major parts of the framework can be taken directly from general open-source IT projects while other areas, like parameter passing, need to be customized for astronomical applications.  The isolation of IT infrastructure will safeguard the astronomical code against the rapid changes in the IT world.  The generic, modular structure would allow easy integration of both new and legacy packages.  This would significantly increase the efficiency with which astronomers can access a wide variety of available astronomical software.  An efficient access to software tools would directly lead to accelerating the rate of astronomical discovery."
    },
    "1004/1004.0276_arXiv.txt": {
        "abstract": "We model the time variability of $\\sim$9,000 spectroscopically confirmed quasars in SDSS Stripe 82 as a damped random walk. Using 2.7 million photometric measurements collected over 10 years, we confirm the results of Kelly et al.~(2009) and Koz{\\l}owski et al.~(2010) that this model can explain quasar light curves at an impressive fidelity level (0.01-0.02 mag). The damped random walk model provides a simple, fast [O($N$) for $N$ data points], and powerful statistical description of quasar light curves by a characteristic time scale ($\\tau$) and an asymptotic rms variability on long time scales ($SF_{\\infty}$).  We searched for correlations between these two variability parameters and physical parameters such as luminosity and black hole mass, and rest-frame wavelength. Our analysis shows $SF_{\\infty}$ to increase with decreasing luminosity and rest-frame wavelength as observed previously, and without a correlation with redshift. We find a correlation between $SF_{\\infty}$ and black hole mass with a power law index of 0.18$\\pm$0.03, independent of the anti-correlation with luminosity. We find that $\\tau$ increases with increasing wavelength with a power law index of 0.17, remains nearly constant with redshift and luminosity, and increases with increasing black hole mass with power law index of 0.21$\\pm$0.07.  The amplitude of variability is anti-correlated with the Eddington ratio, which suggests a scenario where optical fluctuations are tied to variations in the accretion rate. However, we find an additional dependence on luminosity and/or black hole mass that cannot be explained by the trend with Eddington ratio. The radio-loudest quasars have systematically larger variability amplitudes by about 30\\%, when corrected for the other observed trends, while the distribution of their characteristic time scale is indistinguishable from that of the full sample. We do not detect any statistically robust differences in the characteristic time scale and variability amplitude between the full sample and the small subsample of quasars detected by ROSAT.  Our results provide a simple quantitative framework for generating mock quasar light curves, such as currently used in LSST image simulations. ",
        "introduction": "The optical variability of quasars has been recognized since they were first identified (Matthews \\& Sandage 1963). Indeed, most quasars are variable ($\\sim$90\\% at the 0.03 mag rms level; Sesar et al.~2007), and the variations in brightness are aperiodic and on the order of 20\\% on time scales of months to years (e.g., Hook et al.~1994; Vanden Berk et al.~2004). Furthermore, the smooth power spectra suggest a chaotic, or stochastic, origin for the variability. A range of models have been advanced to describe quasar variability, including supernova bursts, microlensing, and accretion disk instabilities (Aretxaga et al.~1997; Hawkins~1993; Kawaguchi et al.~1998; Tr{\\`e}vese \\& Vagnetti~2002). These models are discussed and compared in Hawkins~(2007). Reverberation mapping studies (e.g., Peterson et al.~2005) show that the broad emission lines respond to continuum fluctuations, therefore providing strong evidence that the variability is intrinsic to the quasars. A number of studies have utilized standard accretion disk models to demonstrate that the optical-UV variability of quasars could be driven by a variable accretion rate (e.g., Pereyra et al.~2006; Li \\& Cao 2008; Liu et al.~2008). Blackburne \\& Kochanek (2010) find evidence in the light curves of microlensed quasars that the optical variability is caused by a change in the effective area of the accretion disk.  Recently, Kelly et al.~(2009, hereafter KBS09) proposed a model where the optical variability is described by a damped random walk (a self-correcting term added to a random walk model that acts to push any deviations back towards the mean value). They proposed that the variability time scale might be identified with the thermal time scale of accretion disks, as also proposed by Collier \\& Peterson~(2001). A thermal origin of the variability would explain why quasars become bluer as they brighten (e.g., Giveon et al.~1999; Tr{\\`e}vese et al.~2001; Geha et al.~2003).  Although the physical causes have yet to be proven, it has been established by KBS09 and Koz{\\l}owski et al.~(2010a, hereafter Koz10) that a damped random walk can statistically explain the observed light curves of quasars. Using 100 well-sampled single-band light curves compiled from the literature, KBS09 show that this stochastic process is capable of modeling complex quasar light curves at an impressive fidelity level (0.01-0.02 mag). Koz10 applied the model to the OGLE light curves (Udalski et al.~1997; Udalski et al.~2008) of mid-infrared-selected quasars behind the Magellanic Clouds from Koz{\\l}owski \\& Kochanek (2009). Their analysis shows that this stochastic model is robust enough to efficiently select quasars from other variable sources (see Schmidt et al.~2010 for a different method of selecting quasars based on variability). The model has only three free parameters: the mean value of the light curve, the driving amplitude of the stochastic process, and the damping time scale. The predictions are only statistical, and the random nature reflects our uncertainty about the details of the physical processes.  Instead of applying a model to observed light curves for individual quasars, numerous studies have looked at the ensemble variability of quasars, particularly in samples where individual light curves are not available. Significant progress in the description of quasar variability has been made by employing the Sloan Digital Sky Survey (SDSS) data (Vanden Berk et al.~2004, hereafter VB04; Ivezi\\'{c} et al.~2004, hereafter I04; de Vries et al.~2005; Wilhite et al.~2005, 2006, 2008; Sesar et al.~2006). For example, the size and quality of the sample analyzed by VB04 (two-epoch photometry for 25,000 spectroscopically confirmed quasars) allowed them to constrain how quasar variability in the rest frame optical/UV regime depends upon rest-frame time lag (up to $\\sim$2 years), luminosity, rest wavelength, redshift, the presence of radio and X-ray emission, and the presence of broad absorption line systems. Using repeated SDSS photometric observations, Wilhite et al.~(2008) confirmed the result of Wold et al.~(2007) that variability is correlated with black hole mass, and show that this is independent of the anti-correlation between variability and luminosity established by many studies. This led them to suggest that the amplitude of variability may be driven by the quasar's Eddington ratio, implying differences in accretion rate.  These studies typically quantify the observed optical variability of quasars using a structure function (SF) analysis (see also Hughes et al.~1992; Collier \\& Peterson 2001; Bauer et al.~2009; Koz{\\l}owski et al.~2010b), where the SF is the root-mean-square (rms) magnitude difference as a function of the time lag ($\\Delta t$) between measurements. This autocorrelation-like function is less sensitive to aliasing and other time-sampling problems than a power spectral distribution. By studying the magnitude difference distribution for appropriately chosen subsamples with fixed values of absolute i-band magnitude ($M_i$), rest-frame time lag ($\\Delta t_{RF}$, in days), and wavelength ($\\lambda_{RF}$, in \\AA), the mean dependence of the SF on these quantities was inferred by I04 to be \\begin{equation} SF_{model}= A[1 + B\\,M_i] \\left(\\frac{\\Delta t_{RF}}{\\lambda_{RF}}\\right)^C \\rm mag, \\label{eq:I04} \\end{equation} with $A=1.00 \\pm 0.03$, $B=0.024 \\pm 0.04$, and $C=0.30 \\pm 0.05$. A qualitatively similar result was obtained by VB04. Koz{\\l}owski et al.~(2010b), in the first large study of the mid-IR structure functions of quasars, also found lower variability for higher luminosities and longer wavelengths, but the temporal slope of the ensemble structure functions were significantly steeper than in the optical. In addition, there is evidence for a turnover in the SF on long time lags (I04; Rengstorf et al.~2006; Wold et al.~2007). Studies by de Vries et al.\\ (2005) and Sesar et al.\\ (2006) using SDSS combined with earlier Palomar Observatory Sky Survey measurements for 40,000 SDSS quasars constrained quasar continuum variability on time scales of 10 to 50 years in the observer's frame. They report that the characteristic time scale, which in this context is the time lag above which the SF flattens to a constant value, is of order 1 year in the quasar rest frame. Using a shot-noise light curve model, de Vries et al.~(2005) found evidence for multiple variability time scales in long-term ensemble variability measurements, while Collier \\& Peterson (2001) found a wide range of different time scales in their analysis of individual light curves, and even evidence for multiple time scales in a single active galactic nucleus (AGN).  These analyses of ensemble variability are based on a fundamental assumption that photometric observations at two epochs for a large number of quasars will reveal the same statistical properties as well-sampled light curves for individual objects.  This assumption has been tested by MacLeod et al.~(2008) using light curves for spectroscopically confirmed quasars observed roughly 50 times over 8 years in SDSS Stripe 82 (S82). They found that while the mean SF for individual sources is consistent with Eq.~\\ref{eq:I04}, the contribution of the mean trends to the observed dispersion in variability properties is minor compared to an intrinsic stochasticity of unknown origin. Further investigation of this stochastic behavior is one of the main goals of this study.  In order to better understand the relationship between the two types of data analyses (individual versus ensemble quasar variability), and to begin linking to physical models, we apply the damped random walk model to the $ugriz$ light curves of $\\sim$9,000 spectroscopically confirmed SDSS S82 quasars.  This large sample greatly benefits from the robust, accurate, five-band SDSS photometry. We estimate the variability parameters following Koz10, who demonstrated that their approach is more statistically powerful than the forecasting methods used by KBS09. We also note that both the Koz10 and KBS09 approaches are much faster than that used by Schmidt et al.~(2010), requiring only O($N$) rather than O($N^2$) operations to determine the model parameters for a light curve with $N$ data points.  In Section~\\ref{sec:methodology}, we describe the model, define our variability parameters, and demonstrate their relationship to those utilized in previous studies. In Section~\\ref{sec:data}, we introduce the S82 data set and outline our initial light curve selection. In Section~\\ref{sec:results}, we present the best-fit variability parameters for our final sample of light curves and estimate their scatter due to the limited time sampling of SDSS. We also estimate the sensitivity of our results to variations in the slope of the model power spectrum on long time scales. In Section~\\ref{sec:trends}, we describe the relationship between the long-term variability parameters and physical parameters such as wavelength, absolute magnitude, black hole mass, and Eddington ratio. Using these results, we also provide a prescription for simulating mock quasar light curves. In Section~\\ref{sec:radio}, we explore the variability properties of subsamples detected at radio and X-ray wavelengths. Finally, we summarize our results in Section~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion} We have used the damped random walk model of KBS09 and Koz10 to model the optical/UV variability of $\\sim$9000 SDSS Stripe 82 quasars with the $ugriz$ light curves. The dataset includes 2.7 million photometric measurements collected over 10 years.  We confirm that this is a good model of quasar variability, and quantify the dependence of two variability parameters, the long-term rms variability $SF_{\\infty}$ and the damping time scale $\\tau$, on physical parameters such as wavelength, luminosity, black hole mass and Eddington ratio. Our main results are the following: \\begin{enumerate} \\item A stochastic process with an exponential covariance function characterized by an amplitude and time scale provides a good fit to observed quasar light curves, as shown by KBS09 and Koz10, using smaller samples with less wavelength coverage, but better time sampling. \\item The long-term rms variability, $SF_{\\infty}$, has a mode at $\\sim$0.2 mag and characteristic time scales, $\\tau$, are roughly 200 days in the rest frame, as found previously by KBS09 and Koz10. These time scales are consistent with thermal time scales, but simple accretion disk models fail to reproduce the observed scaling of $\\tau$ with physical parameters. \\item Quasars with similar physical parameters can have different characteristic time scales for variability. It is now clear that the distribution of $\\tau$ accounts for most of the scatter in the structure function on short time scales for quasars with similar luminosity, rest wavelength, and time lag, which explains the puzzling results from MacLeod et al.~(2008). Results from fitting a power law to observed ensemble structure functions should be interpreted with caution. \\item The variability time scale is correlated with the long-term rms variability with a slope of $1.30\\pm0.01$ dex/dex. Quasars that have large long-term variability amplitudes generally vary on longer characteristic time scales. The amplitude of short term variations is also correlated with $\\tau$. This conclusion is unaffected by any time sampling issues in the S82 dataset. \\item The damped random walk corresponds to a PSD proportional to $1/f^2$ at frequencies $f>(2\\pi\\tau)^{-1}$, flattening to a constant at lower frequencies. At large $f$, the data are in great agreement with $PSD \\propto 1/f^2$. In terms of the structure function, this means that $SF \\propto (\\Delta t)^{1/2}$. Whereas previous analyses of the SF obtained a power law slope of $\\beta = 0.3$, here we demonstrated that this would be a consequence of fitting the data around the ``knee'' (turn-over) of the SF. Our constraints for small $f$ are much weaker. As discussed in Section~\\ref{sec:PSD}, due to a lack of sufficient long-time scale information, we are unable to distinguish between a $1/f^{0}$ or a $1/f$ PSD at frequencies $f<(2\\pi\\tau)^{-1}$ using the data and computational technique described here. \\item The rest-frame variability parameters show a negligible trend with redshift, suggesting that they are intrinsic to the quasars, and these properties do not evolve over cosmic time for fixed physical parameters of the quasar ($M_{BH}$, $M_i$, and $\\lambda_{RF}$). \\item For fixed luminosity and black hole mass, $\\tau$ increases with increasing rest-frame wavelength with a power law index of $0.17$, and $SF_{\\infty}$ decreases with a power law index of $-0.48$. The latter result is similar to previous findings (e.g., MacLeod et al.~2008, VB04). Koz10 also observed that the variability increases to shorter $\\lambda$, but they kept $\\tau$ fixed in their fits. If wavelength is a proxy for radius in the accretion disk, this implies that the characteristic time scales are longer and the variability amplitudes are smaller in the outer regions than in the inner regions. \\item The long-term variability $SF_{\\infty}$ is strongly anti-correlated with luminosity as found in previous studies such as VB04, Wilhite et al.~(2008), and references therein. By studying the median $SF_{\\infty}$ in the plane of absolute magnitude and black hole mass, we can separate the anti-correlation of amplitude with luminosity from the positive correlation with black hole mass. As suggested in Wilhite et al.~(2008), these trends may be largely explained if the amplitude of variability is tied to changes in the accretion rate in the disk, and is simply related to the Eddington ratio. However, despite the strong anti-correlation between $SF_{\\infty}$ and $L/L_{Edd}$ (which accounts for most of the dependence on $M_i$ and $M_{BH}$), the exact dependence with $M_i$ and $M_{BH}$ is not consistent with $L/L_{Edd}$ as the sole driver of quasar variability. \\item The damping time scale $\\tau$ appears to be nearly independent of luminosity and correlated with $M_{BH}$ with a power law index of $0.21\\pm 0.07$.  The mild discrepancy with the KBS09 result (1.0$\\pm$0.4) may be due to the different range of sampled luminosity and black hole mass, as well as contamination by host galaxy emission in many of the very low luminosity systems they consider (see Koz10). \\item While the mean variability parameters can be related to physical parameters, for fixed values of $M_i$, $\\lambda_{RF}$, and $M_{BH}$,  there is still a large scatter around the mean values, similar to the variance of the observed distributions. Some of that scatter can be attributed to measurement and fitting errors ($\\sim$60\\% for $\\tau$  and $\\sim$70\\% for $SF_{\\infty}$), but there is enough evidence for residual stochastic nature of quasar variability. Therefore, it cannot be assumed that quasars with similar $M_i$, $\\lambda_{RF}$, and $M_{BH}$ will necessarily have similar variability properties. \\item The radio-loudest quasars have systematically larger variability amplitude by about 30\\%, while the distribution of their characteristic time scale is indistinguishable from that of the full sample. There are no statistically robust differences in the characteristic time scale and variability amplitude between the full sample and a small subsample of quasars detected by ROSAT. \\end{enumerate}  With this paper, and results from KBS09 and Koz10, the ability of the damped random walk model to quantitatively describe quasar variability is well-established.  As emphasized by Koz10, this means that variability studies can become fully quantitative because the entire process of identifying and assigning parameters to quasars can be simulated to allow estimates of completeness and parameter biases.  In particular, an important next step is to determine the variability equivalents of luminosity functions, i.e., the intrinsic distributions of the variability parameters.  While our results represent a good first step in this direction, we caution that a non-negligible fraction of the S82 quasars have indeterminately long time scales.  If, following KBS09, we identify the characteristic time scale with the thermal time scale, then the next question is whether there are additional time scales (such as the dynamical or viscous time scale), or sources of variability. For example, Blackburne \\& Kochanek~(2010) recently found evidence for changes in disk size with changes in luminosity using gravitational microlensing. Such searches for additional sources of variability will likely require better sampled light curves and over a longer time scale.  The prospect of advancing these studies of quasar variability now faces a bottleneck. The S82 quasars have the advantage of sample size, wavelength coverage, and spectroscopy, but the light curves have poor sampling and modest overall lengths, leading to significant problems for accurately estimating $\\tau$ when it is long.  The quasars behind the Magellanic Clouds (Koz{\\l}owski \\& Kochanek 2009) are a smaller sample without spectroscopic confirmation, but have superb, long term light curves that continue to be extended because of the continuing microlensing projects.  Improving on our present results in the short term depends on either reviving the monitoring of S82 or spectroscopically confirming the Magellanic Cloud quasars.  Resuming the monitoring of Stripe 82 is challenging with the decommissioning of the SDSS imaging system.  The best short-term prospects are the Pan-STARRS project (Kaiser et al.~2002), the Palomar-QUEST project  (Schweitzer et al.~2006), or using the DECam being built for the Dark Energy Survey (Honscheid et al.~2008).  Since the challenge is to constrain long time scales, the presence of a multiple-year gap is mainly a complication for ensuring that problems in matching photometric bands are not interpreted as a form of long-term variability. Obtaining spectra of the Magellanic Cloud quasars is in some ways easier, because the quasar magnitudes and densities are well-suited to the AAOmega fiber spectrograph on the AAT and, to a lesser degree, the IMACS spectrograph on Magellan.  The more general, Northern monitoring projects such as Pan-STARRS and Palomar-QUEST will slowly build light curves for essentially all the SDSS quasars, but at present their cadences are not ideal (see Schmidt et al.~2010) and it will take nearly a decade to build the long duration light curves needed for the analysis. In the long term, observations will be significantly improved with the advent of next-generation sky surveys. Most notably, the Large Synoptic Survey Telescope  (LSST, Ivezi\\'{c} et al. 2008) will obtain accurate, well-sampled light curves for millions of AGN. The observed distribution of rest-frame characteristic time scales for S82 quasars spans the range from about 10 days to 1000 days (c.f.\\ Figure~\\ref{fig:pardist}).  To probe the time scales as short as $0.1\\tau$, and assuming a characteristic redshift of 2, the light curves should be sampled every 3 days in the observer's frame, which is in good agreement with the baseline cadence of LSST. With a 10 year-long survey, the length of the light curves will be in the range (1-200)$\\tau$. A combination of the SDSS, Pan-STARRS, DES, and LSST data for $\\sim$10,000 Stripe 82 quasars would span well over two decades, with multi-band photometry obtained for hundreds of epochs, and would represent the best sample to date for studying the optical continuum variability of quasars. In particular, such a dataset would enable a robust measurement of the low-frequency behavior of their PSD (c.f.\\ point 5 above). For illustration, the LSST photometric errors in the $r$ band will be $<0.02$ mag for $r<22$, and there are roughly 2-3 million AGN with $r<22$ in the 20,000 sq.\\ deg.\\ covered by the main LSST survey (see Table 10.2 in the LSST Science Book; Abell et al.~2009). Each of these objects will be observed about 1000 times, yielding a database of over 2 billion photometric measurements. This data set, roughly a thousand times larger than that analyzed here, will enable a significant improvement in our understanding of quasar variability."
    },
    "1004/1004.5386_arXiv.txt": {
        "abstract": "We study the properties of simulated high-redshift galaxies using cosmological N-body/gasdynamical runs from the OverWhelmingly Large Simulations ({\\small OWLS}) project. The runs contrast several feedback implementations of varying effectiveness: from no-feedback, to supernova-driven winds to powerful AGN-driven outflows.  These different feedback models result in large variations in the abundance and structural properties of bright galaxies at $z=2$. In agreement with earlier work, models with inefficient or no feedback lead to the formation of massive compact galaxies collecting a large fraction (upwards of $50\\%$) of all available baryons in each halo. Increasing the efficiency of feedback reduces the baryonic mass and increases the size of simulated galaxies. A model that includes supernova-driven gas outflows aided by the energetic output of AGNs reduces galaxy masses by roughly a factor of $\\sim 10$ compared with the no-feedback case. Other models give results that straddle these two extremes. Despite the large differences in galaxy formation efficiency, the net specific angular momentum of a galaxy is, on average, roughly half that of its surrounding halo, independent of halo mass (in the range probed) and of the feedback scheme. Feedback thus affects the baryonic mass of a galaxy much more severely than its spin. Feedback induces strong correlations between angular momentum content and galaxy mass that leave their imprint on galaxy scaling relations and morphologies.  Encouragingly, we find that galaxy disks are common in moderate-feedback runs, making up typically $\\sim 50\\%$ of all galaxies at the centers of haloes with virial mass exceeding $\\sim 10^{11} \\, M_{\\odot}$. The size, stellar masses, and circular speeds of simulated galaxies formed in such runs have properties in between those of large star-forming disks and of compact early-type galaxies at $z=2$. Once the detailed abundance and structural properties of these rare objects are well established it may be possible to use them to gauge the overall efficacy of feedback in the formation of high redshift galaxies. ",
        "introduction": "\\label{sec:intro}  The established paradigm for structure formation offers a clear road map for galaxy formation. Primordial fluctuations in the dominant cold dark matter (CDM) component of the Universe grow via gravitational instability, sweeping baryons into an evolving hierarchy of dark matter haloes that grow through mergers of preexisting units as well as through the accretion of material from the intergalactic medium \\citep{White1978}. On galaxy mass scales, baryons caught in a halo are able to radiate away the gravitational energy gained through the collapse, sink to the center of the halo, and assemble into the dense aggregations of gas and stars that we call galaxies \\citep{Blumenthal1985}.  The structure and morphology of a galaxy results from the complex interplay between the time of collapse, the mode of assembly, the efficiency of cooling, and the rate of transformation of gas into stars \\citep[see, e.g.,][]{SteinmetzNavarro2002}. Where cooling dominates and outpaces star formation, baryons collect into thin, rotationally-supported disks \\citep{Fall1980}. Stars formed in these disks inherit these morphological features, but are vulnerable to swift transformation into dispersion-supported spheroids during subsequent merger events \\citep{Toomre1977}. Disks may re-form if mergers or accretion bring fresh supplies of cooled gas, making morphology a constantly evolving rather than an abiding feature of a galaxy \\citep{Cole2000,Robertson2006}.  The galaxy formation scenario driven by gravitational collapse and radiative losses outlined above is compelling, but incomplete. Indeed, cooling is so effective at early times that, unless impeded somehow, most baryons would be turned into stars in early-collapsing protogalaxies, which would then merge away to form by the present time a majority of spheroid-dominated remnants, in vehement disagreement with observations \\citep{White1978,Cole1991,White1991}. The problem is compounded by the fact that, during mergers, cooled gas tends to transfer its angular momentum to the surrounding dark matter halo. As a result, even in cases where disks could form, their structural properties would be at odds with those of spiral galaxies \\citep{Navarro1991,Navarro1995,NavarroSteinmetz1997}.  A gas heating mechanism that prevents runaway cooling and that regulates the formation of stars in step with mergers and accretion is widely believed to be the most likely solution to these problems. The energetic output from evolving stars and supernovae is a natural candidate. It scales directly with star formation and, in a typical galaxy, the total energy released by supernovae can be comparable to the binding energy of the baryons. Thus, if channeled properly, feedback energy from supernovae may temper the gravitational deposition of cooled gas into a galaxy and effectively self-regulate its star formation history \\citep{White1991}.  The even standing of gravity, feedback and cooling may thus help reconcile the observed galaxy population with hierarchical clustering models, but it comes at the price of complexity: the main structural properties of a galaxy, such as stellar mass, rotation speed, and morphology, are then expected to depend on details of its assembly history, such as the exact timing, geometry and mass spectrum of accretion events \\citep[see, e.g.,][]{Abadi2003a,Abadi2003b,Meza2003,Governato2007,Zavala2008, Scannapieco2009,Governato2010}.  Such sensitivity to feedback has held back progress in direct simulation of the process of galaxy formation. As recent work demonstrates, different but plausible implementations of feedback within the {\\it same} dark halo lead to galaxies of very different mass, morphology, dynamics, and star formation history \\citep[see, e.g.,][]{Okamoto2005}. Numerical parameters may thus be tuned to reproduce some properties of individual galaxies, but at the expense of wider predictability in the modeling.  These results suggest that further progress in the subject requires the testing of different feedback schemes on a statistically significant sample of dark haloes formed with representative assembly histories. The viability of each feedback implementation may then be assessed by contrasting the statistics of such samples with observational constraints such as the stellar mass function, clustering, color distribution, and scaling laws.  We take a step in this direction here by analyzing a subset of cosmological N-body/gasdynamical simulations from the OverWhelmingly Large Simulations ({\\small OWLS}) project \\citep{Schaye2010}. We present results regarding the morphology, stellar mass, and angular momentum content of galaxies assembled at $z=2$, and compare them with the few observational constraints available at that epoch. We limit our analysis to the $z=2$ galaxy population because most high-resolution {\\small OWLS} runs follow volumes too small to be evolved until $z=0$.  In future papers, we plan to extend this analysis to the present-day galaxy population using samples drawn from the closely-related {\\small GIMIC} project, designed to follow a few representative volumes selected from the Millennium Simulation \\citep{Crain2009}.  The paper is organized as follows. In Sec.~\\ref{sec:numsim} we present a short overview of the simulations and feedback models. We then present our main numerical results in Sec.~\\ref{sec:numres} and analyze them in the context of available observational constraints in Sec.~\\ref{sec:obsdiag}. We end with a brief summary in Sec.~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We study the effects of various feedback implementations on the structure and morphology of simulated galaxies at $z=2$.  Our analysis uses nine runs from the OverWhelmingly Large Simulations ({\\small OWLS}) project, and probe a variety of possible feedback implementations, from ``no feedback'' to supernova-driven wind feedback to strong outflows aided by the contribution from AGNs. Except for the no-feedback and AGN-feedback cases, all other runs assume that the {\\it same amount} of feedback energy (per mass of stars formed) is devolved by supernovae to the interstellar medium: the main difference is {\\it how} this energy is coupled to the medium, which in turn determines the overall effectiveness of the feedback.  Each run follows the evolution of the {\\it same} $25 \\, h^{-1}$ Mpc box up to $z=2$, with $512^3$ dark matter particles and $512^3$ particles for the baryonic component. All other simulation parameters (star formation algorithm, stellar initial mass function, etc) are kept constant, so any differences between runs may be traced solely to feedback. In total, we analyze for each run $\\sim 150$ galaxies formed at the centers of haloes with virial mass in the range $10^{11} \\, h^{-1} \\, M_\\odot< M_{\\rm vir} < 3\\times 10^{12} h^{-1} \\, M_\\odot$. Our main results may be summarized as follows.  \\begin{itemize}  \\item Varying the feedback implementation can lead to dramatic differences in the mass of galaxies formed in a given dark matter halo. The galaxy formation efficiency, $\\eta_{\\rm gal}=M_{\\rm gal}/(f_{\\rm bar} M_{\\rm vir})$, varies by roughly an order of magnitude when comparing the no-feedback run (NoF, where $\\eta_{\\rm gal}\\sim 0.5$) to the AGN+supernova feedback run (AGN, where $\\eta_{\\rm gal} \\sim 0.05$), the two extremes probed by our simulations.  \\item The ability of feedback to regulate the efficiency of galaxy formation in haloes of different mass varies according to the details of the adopted numerical implementation of the feedback. Weak or ineffective feedback leads to a decrease in galaxy formation efficiency with mass, whereas strong feedback curtails preferentially the formation of galaxies in low-mass haloes. The mass dependence is, however, modest, with variations in $\\eta_{\\rm gal}$ of less than a factor of $\\sim 2$ over the (factor of $\\sim 30$) mass range spanned by haloes in our sample.  \\item Feedback results in strong correlations between galaxy mass and angular momentum.  This leaves an imprint on galaxy morphologies and on the scaling laws relating mass, size, and circular velocity.  \\item Weak feedback minimizes disturbances to the settling of gas in rotationally-supported structures, and favors the formation and survival of quiescent {\\it gaseous} disks. However, weak feedback also allows much of the gas to form stars early in dense protogalactic clumps that are later disrupted in mergers as the final galaxy assembles. Such mergers also transfer angular momentum from the baryons to the halo. The net result is a predominance of dense, spheroid-dominated stellar components and a scarcity of spatially-extended star-forming disks.  \\item Strong feedback, on the other hand, promotes the formation of large, extended galaxies. Indeed, the more efficient the feedback the more massive (and therefore, larger) the halo inhabited by a galaxy of given stellar mass. It is thus possible to have fairly large galaxies of modest stellar mass because, when feedback is strong, they inhabit large, massive haloes. The size, mass, and rotation speeds of these extended galaxies compare favorably with those reported by the SINS survey. This, however, comes at the expense of inhibiting the survival of rotationally-supported disks of quiescent kinematics and of preventing the formation of compact stellar spheroids.  \\item Moderate-feedback runs result in galaxies that follow scaling laws that are intermediate between large star-forming disks, such as those studied by the SINS collaboration \\citep{Forster2009}, and the compact, quiescent early-type systems analyzed by \\citet{vanDokkum2008}. Disk-like morphologies in both gas and stars are common in these runs, in numbers that appear commensurate with current constraints.  \\end{itemize}  Although far from definitive, the results outlined above are encouraging. Properly calibrated, simple feedback recipes such as the ones we explore here seem able to produce galaxies with properties in broad agreement with observation. One should be aware, however, of the numerical sensitivity of the results to details of feedback implementation. Nevertheless, if developed in step with observational progress in the characterization of the high-redshift galaxy population, simulations are likely to become more and more reliable tools, useful when trying to make sense of the striking diversity of high-z galaxies in terms of the current paradigm of structure formation.   \\begin{center} \\begin{figure} \\includegraphics[width=84mm]{figs/fig_angmom_ekinet.ps} \\caption{ Upper panel shows the histogram of $\\kappa_{\\rm rot}=K_{\\rm rot}/K$, the fraction of kinetic energy of star-forming gas particles in ordered rotation for all galaxies in the WF2 run. The shaded red histogram is the same, but only for the half most massive, and therefore best numerically resolved, systems. The similarity between the two suggests that numerical resolution does not play a significant role in the statistics. $\\kappa_{\\rm rot}$ should be approximately unity for a disk where all particles are in circular orbits and much smaller for systems where ordered rotation plays a less important role. The large number of systems around $\\kappa_{\\rm rot} \\sim 0.8$ indicates that systems where star formation occurs in well-defined disks are quite common in this run (see Fig.~\\ref{fig:RotMorph} for examples). The cumulative fraction of systems as a fraction of $\\kappa_{\\rm rot}$ for the four different feedback implementations are shown in the bottom panel. A trend for gaseous disks becoming more prevalent as feedback efficiency decreases is clearly seen.} \\label{fig:HistErot} \\end{figure} \\end{center}"
    },
    "1004/1004.2335_arXiv.txt": {
        "abstract": "We present results of an imaging observation of the central region of a giant radio galaxy B1358+305. The classical, standard scenario of Fanaroff-Riley II radio galaxies suggests that shock produced hot electrons contained in a radio galaxy are a good reservoir of the jet-supplied energy from active nuclei. The aim of our observation is to search for the Sunyaev-Zel'dovich effect induced by these hot electrons. The observation was performed at 21 GHz with the Nobeyama 45-m telescope. Deep imaging observation of a wide region of size $6.7^{\\prime}\\times6.7^{\\prime}$  with the beam size $\\theta_{\\rm HPBW}=81.2^{\\prime\\prime}$ enables the most detailed examination of the possible thermal energy of electrons contained in a radio galaxy. The resultant intensity fluctuation is 0.56 mJy beam$^{-1}$ (in terms of the  Compton $y$-parameter, $y=1.04\\times 10^{-4}$) at a 95 \\% confidence level. The intensity fluctuation obtained with imaging analysis sets the most stringent upper limit on the fluctuations in the central region of a giant radio galaxy obtained so far, and our results will be a toehold for future plans of SZE observation in a radio galaxy. ",
        "introduction": "Extended lobes of radio galaxies are interesting as a probe of the energetics of active galaxies. In the standard model of radio galaxies of Fanaroff-Riley type II \\citep[FR-II;][]{fanaroff1974}, radio lobes are formed by shock interactions of the jets with the surrounding intergalactic medium at the jet extremities \\citep{scheuer1974,blandford1974}. In this scenario, radio lobes consist mainly of the jet-supplied matter that passed the termination shock. The gas density in the radio lobes is sufficiently low that radiative cooling is ineffective, resulting in the energy supplied to the radio lobes to be well conserved for the lifetime of the radio galaxy. The lobe has a higher pressure than the ambient intergalactic medium (IGM) due to shock-compression, and it can expand supersonically into the ambient IGM (``overpressured cocoon\": \\citealp{blandford1974,begelman1989, nath1995, heinz1998, yamada1999}). It is expected to form a shell of IGM matter compressed by the external shock exterior to the radio lobes.   Detections of inverse Compton scattered photons in the X-ray energy bands of radio galaxies can drastically improve our understanding of the properties of hotspots and radio lobes. The distribution of energy among the magnetic field component and in the non-thermal electrons has been investigated by examining the spectral energy distribution from the radio to X-ray bands (e.g., for Cygnus A,  \\citealt{harris1994,wilson2000}; for 3C123, \\citealt{hardcastle2001}; for 3C295, \\citealt{harris2000}; for Pictor A, \\citealt{wilson2001}; for 3C120,  \\citealt{harris1999}; for 3C390.3, \\citealt{harris1998}).  The results of these studies revealed that high energy electrons of $\\gamma\\approx 10^3 - 10^5$ have an energy comparable to or greater than that in the magnetic field \\citep[e.g.,][]{tashiro1998, grandi2003, croston2005, isobe2006,isobe2009}. If the energy spectrum of non-thermal electrons is given by a power-law ($dN(E)/dE \\propto E^{-\\eta}$), most of the non-thermal electron energy is carried by those electrons characterized by energies near the lower cut-off energy of $N(E)$ for $\\eta>1$. However, direct measurements of the thermal and/or low energy ($\\gamma\\approx 10^3$ or less) electrons are quite difficult, since the plasma density in a radio galaxy is too low to emit detectable radiation. The standard overpressured cocoon models of FR-II radio galaxies are characterized by two types of shocks, i.e., the external and the internal (jet-terminal) shock, where the latter is believed to be observed as hotspots. In this model, radio galaxies emit synchrotron radiation by shock accelerated electrons with an acceleration efficiency $\\xi_e\\equiv U_\\mathrm{syn}/U_\\mathrm{tot}$, which is inferred to be low ($U_\\mathrm{syn}$ and $U_\\mathrm{tot}$ are energy of synchrotron-emitting electrons and total internal energy, respectively). If we adopt the standard model and assume the low acceleration efficiency in both shocks (low $\\xi_e$), we can expect that a large amount of thermal electron energy, as well as non-thermal, synchrotron-emitting electron energy. % In order to investigate the energetics of active galaxies, it is quite important to measure all of the energy in the electron component, including those that emitting strongly.  In this paper we employ the Sunyaev-Zel'dovich effect (SZE) as a tool to measure the energy of the electrons in a radio galaxy.  The SZE represents the spectral deformation of the Cosmic Microwave Background (CMB) radiation due to the inverse Compton scattering of these photons by the energetic electrons \\citep{zeldovich1969} in the galaxy.  The intensity change of the thermal SZE is classically described as follows \\citep{zeldovich1969}; \\begin{eqnarray} \\frac{\\Delta I_{x}}{I_{x}} &= &\\frac{xe^x}{(e^x-1)} \\left[ x\\left(\\frac{e^x+1}{e^x-1}\\right)-4\\right]  y,  \\label{eq:SZ} \\\\ y &=& \\int \\frac{k_BT_e}{m_ec^2}\\sigma_Tn_e dl \\propto\\int p_edl,  \\label{eq:press} \\end{eqnarray} where $x\\equiv h\\nu/k_BT_r$ is the non-dimensional frequency, $T_r$ is the temperature of the CMB, $k_B$ is the Boltzmann constant, $T_e$ is the electron kinetic temperature, $\\sigma_T$ is the cross section of Thomson scattering, $n_e$ is the electron number density, $m_e$ is the electron mass, and $p_e$ is the electron thermal pressure, respectively, integrated along the line of sight. Equation (\\ref{eq:press}) shows that the Compton parameter $y$ is proportional to the sum of the thermal pressure of electrons along the line of sight. Therefore we can estimate the thermal energy of the electrons contained in a radio galaxy with the thermal SZE. Similarly, the decrease in CMB intensity by SZE in the Rayleigh-Jeans regime reflects all of the electron energy, not only the high energy electrons which generate X-ray photons, but also the lower energy electrons and the thermal electrons. Although the energy distribution of the electrons is lost in the SZE (Eq.[\\ref{eq:press}]),  it measures the total energy deposited in the radio galaxy. In this paper, we ignore the non-thermal and kinetic SZE as well as relativistic corrections and focus mainly on thermal SZE as the first trial observation \\citep{birkinshaw1999,ensslin2000,yamada2001}.  The study of the SZE has been directed toward understanding the thermal properties of the intra cluster medium (ICM), and it has been detected in dozens of clusters of galaxies (see for  recent reviews, \\citealt{birkinshaw1999}; \\citealt{carlstrom2001}; \\citealt{rephaeli2002}). Among the numerous SZE detections in clusters of galaxies, \\citet{mckinnon91} first tried to detect the SZE in the radio lobes with the NRAO 12-m telescope at 90 GHz using the double-subtraction method. They observed 4 FR-II radio galaxies whose lobe sizes were smaller than the beam size,  and obtained upper limits on the antenna temperature fluctuation ($\\approx 0.1$ mK, or $y\\approx 10^{-4}$). In this paper, we report on the results of the imaging observation of a giant radio galaxy B1358+305, extending much further in angular size than the beam size ($\\sim 80^{\\prime\\prime}$ in our observation) using the Nobeyama 45-m telescope at 21 GHz. A two dimensional imaging study of a region in a radio galaxy is expected to provide the most reliable limit on the SZE in a radio galaxy by resolving the substructure within and the emission sources in the field of view.  The organization of the paper is as follows. In Section 2 we briefly review the overpressured cocoon model \\citep[see for detail,][]{yamada1999}. In Section 3 we describe the features of our target B1358+305, and the observation and analysis procedures. In Section 4 the results are presented.  We discuss the possible uncertainties in the estimation of the SZE amplitude in B1358+305 and their implications for radio galaxy models in section 5. Finally, we summarize our conclusions and future prospects in the last section. ",
        "conclusions": "We have reported the results of a trial observation of SZE in a giant radio galaxy B1358+305 with the Nobeyama 45-m telescope at 21 GHz.  By performing the imaging observation, we have obtained the most stringent upper limit achieved for the Compton $y$ parameter in a radio galaxy.  The obtained upper limit is close to the expected value derived for a low acceleration efficiency of synchrotron electrons at the shock (the low value $\\langle \\xi_e\\rangle$, see Eq.\\,[\\ref{eq:expect}]), but detailed analysis shows that the obtained intensity fluctuation is likely to be caused by the excess atmospheric noise. If we assume that the obtained intensity fluctuations in the observed region were due to the thermal electrons in the cocoon of B1358+305, our results are qualitatively consistent with the supersonic expansion of B1358+305 but for the original assumption $\\langle\\xi_e\\rangle$ being too small. Alternatively, if the true pressure of thermal electrons is much lower than that derived from the obtained upper limit, our observation is % not sensitive enough to derive any definitive conclusions on the pressure balance of B1358+305. Since the SZE is sensitive only to electron energy, it would serve as a probe of the pressure components in radio galaxies as well as galaxy clusters.  Future high sensitive multi-frequency SZE observations of multi-regions will provide an important information about the contributions of thermal or low-energy non-thermal electrons to the total pressure of a radio galaxy.  This may provide a clue for disentangling the pressure discrepancy between the surrounding ICM derived by X-ray observations and the minimum energy of radio galaxies at the centers of clusters of galaxies \\citep{hardcastle2000,leahy2002}. Though we have failed to obtain a definitive SZE signal with the Nobeyama 45-m telescope, higher frequency observations using either a large field-of-view (e.g., $\\gtrsim 5^{\\prime}$), multi-beam receivers on large (e.g., $\\gtrsim$ 30 m) single dish telescopes, or interferometers with a large number of small dishes (e.g., $\\lesssim 50$ cm) similar to the AMiBA project\\footnote{See {\\tt http://amiba.asiaa.sinica.edu.tw/} for the AMiBA project.} might lead to the detection of SZE of a cocoon. The SZE would provide a new observational tool to probe the energetics of radio galaxies, along with the projects to detect low frequency radio emission from non-thermal low energy electrons like the Long Wavelength Array (LWA: \\citealp{harris2005})."
    },
    "1004/1004.2429_arXiv.txt": {
        "abstract": "We consider dynamical scales in magnetized GRB outflows, using the solutions to the Riemann problem of expanding arbitrarily magnetized outflows (Lyutikov 2010). For high ejecta magnetization, the   behavior of the forward shock closely resembles the so-called thick shell regime of  the hydrodynamical expansion.  The exception is at small radii, where the motion of the  forward shock is determined by the dynamics of {\\it subsonic} relativistic outflows.  The  behaviors of the reverse shock is different in fluid and magnetized cases:  in the latter case, even for  medium magnetization,  $\\sigma \\sim 1$,  the reverse shock  forms at fairly large distances,  and may never form in a wind-type external density profile. ",
        "introduction": "Magnetic fields  may play an important dynamical role in the GRB  outflows \\cite[\\eg][]{LyutikovJPh,Lyutikov:2009}. They may power the relativistic outflow through \\cite{BlandfordZnajek} process \\citep[\\eg][]{Komissarov05}, and contribute to particle acceleration in the emission regions. In this paper we discuss the dynamics of the relativistic, strongly magnetized   ejecta. The results  are based on an exact solution of  a one-dimensional Riemann problem of expansion of a cold,  strongly magnetized into vacuum and into external medium of density $\\rho_{\\rm ex}$ (Lyutikov, submitted); they are   reviewed in \\S \\ref{Riemann}.  In application to GRBs, we assumes that the central engine produces   jet  with density  $\\rho_{0}$ and  magnetization $\\sigma$ ($\\sigma=B_0^2/\\rho_{0}$; \\Bf\\ is normalized by $\\sqrt{4 \\pi} $), moving with Lorentz factor  $\\gamma_w\\gg 1$. In fact, parameters $\\gamma_w$ and $\\sigma$ are not always independent quantities: at small radii, when the motion of the ejecta is subsonic, they should be determined together with the motion of the boundary, see \\S  \\ref{subsonic}. In a supersonic regime, relation between $\\gamma_w$ and $\\sigma$ depends on the details of the flow acceleration (\\eg\\ in conical flows we expect $\\gamma_w \\sim \\sqrt{\\sigma}$). For generality, we do not assume any relationship between $\\sigma$ and $\\gamma_w$. The ejecta is moving into external density $\\rho_{\\rm ex}$. ",
        "conclusions": "In this paper we discuss the dynamics of strongly magnetized   outflows in GRBs.  We find that the evolution of the forward shock driven by  strongly magnetized outflows are qualitatively the same as in the case of fluid shocks. The   definitions of radii  $r_N,\\,r_E$ and $r_\\Delta$   involve only the total energy of the ejecta, it's thickness and initial Lorentz factor, and {\\it not} the information about it's content, \\eg, parameter $\\sigma$. The  typical radii (\\ref{xi}) are the same for two flows \\citep[\\cf\\ Eq.  (\\ref{xi}) of the present paper and Eq. (9-10) of ][]{Sari95}. These similarities may be understood, first, by noting that jump conditions in perpendicular magnetized shocks may be reduced to fluid shock jump conditions, with an appropriate choice of the equation of state, and, second, by the fact that the thin shell approximation is applicable in our case (so that the global  conservation of the toroidal magnetic flux, which modifies the global flow dynamics \\citep{kc84},  is not important). Another   reason for this similarity is that \\Bf\\ behaves in many respects as a fluid with internal pressure. The only difference in the dynamics of the forward shocks driven by magnetized and fluid flows occurs for supersonic flows, $\\gamma_w > \\sqrt{\\sigma} \\gamma_{CD}$,  at very early stages  $r\\leq r_N$ or $r\\leq r_E$, see Fig. (\\ref{xiless1}). Qualitatively, magnetized outflows are similar to thick shell hydrodynamic outflow, $\\xi< 1$  at $r> r_N$.   Only at very early times, at $r< r_N$,  the forward shock bears information about anergy content: forward shock is coasting with $\\gamma_w =$const in the fluid case and decelerating $\\gamma \\propto r^{-1/2}$ in the magnetized case. Dynamics of the reverse shock is quite different in case of high magnetization. First, the reverse shock forms at a finite distance from the source (Eq. \\ref{rN}), and may not form at all in a wind environment,  (Eq. \\ref{rss1}).  This fact may be related to observed paucity  of optical flashes in the  {\\it Swift} era \\citep{Gomboc}. (The standard model had a clear prediction, of a bright optical flare with a definite decay properties \\citep{1996ApJ...473..204S,1997ApJ...476..232M}. Though a flare closely resembling the predictions was indeed observed \\citep[GRB990123,][]{1999MNRAS.306L..39M}, this was an exception.)  In addition, at distances  $r_N< r< \\sigma r_N$, where the RS is weak, the formation of  the RS shock depends on the details of the flow: RS forms if the flow is strictly radial, but need not to form if the the flow pattern is more complicated. We suggest that optical variability often seen in GRBs (\\eg\\ GRB021004 and most notoriously GRB080916C) is a reflection of the non-trivial flow patterns and the corresponding non-steady RS formation. Also, a recent detection of high polarization in optical \\citep{2009Natur.462..767S} indicates a presence of an ordered \\Bf\\ in the ejecta.   I am greatly  thankful to  Dimitros Gianios, Sergey Komisarov and Alexandre Tchekhovskoy."
    },
    "1004/1004.0937_arXiv.txt": {
        "abstract": "We propose a novel mechanism for dark matter to explain the observed annual modulation signal at DAMA/LIBRA which avoids existing constraints from every other dark matter direct detection experiment including CRESST, CDMS, and XENON10.    The dark matter consists of at least two light states with mass $\\sim$ few GeV and splittings $\\sim 5$ keV.  It is natural for the heavier states to be cosmologically long-lived and to make up an $\\OrderOne$ fraction of the dark matter.  Direct detection rates are dominated by the exothermic reactions in which an excited dark matter state down-scatters off of a nucleus, becoming a lower energy state.  In contrast to (endothermic) inelastic dark matter, the most sensitive experiments for exothermic dark matter are those with light nuclei and low threshold energies.  Interestingly, this model can also naturally account for the observed low-energy events at CoGeNT.  The only significant constraint on the model arises from the DAMA/LIBRA unmodulated spectrum but it can be tested in the near future by a low-threshold analysis of CDMS-Si and possibly other experiments including CRESST, COUPP, and XENON100. ",
        "introduction": "\\label{Sec:Intro}  A variety of experiments are currently probing the interaction between dark matter and the standard model. These dark matter direct detection experiments are at sensitivities where they can observe weak scale interactions between nuclei and dark matter. The DAMA collaboration has observed an annual modulation of the event rate in their NaI-target detectors at the 8.9$\\sigma$ level \\cite{DAMAResults}. This can be interpreted as modulation of the rate of dark matter (DM) scatters in the detector due to the changing velocity of the Earth through the DM halo. However, if one assumes an elastically scattering WIMP and standard quenching of nuclear recoils then this signal is in conflict with null results from other experiments, particularly CDMS and XENON10 \\cite{Kopp:2009qt}.  The choice of target nuclei and the technique employed to detect  dark matter is unique to each experiment.  Consequently, the comparison of results from different experiments requires the specification of an underlying model that determines the interaction between the dark matter and the nucleus. A change in this underlying model could alter the expected event rate and the nuclear energy recoil spectrum in different experiments. In particular, many  experiments are optimized to look for the elastic scattering of dark matter with nucleons. This  leads to the familiar exponentially falling nuclear recoil spectrum. However, by modifying the nature of the dark matter - nucleus interaction, this spectrum could be modified, thus changing the sensitivity of different experiments  \\cite{IDMOriginal,IDMinLightofDAMA, Khlopov:2008ty}.  This strategy has been used to resolve the conflict between the observations of DAMA and the null results of other experiments. One possibility is upscattering inelastic dark matter (iDM) \\cite{IDMOriginal,IDMinLightofDAMA}: dark matter of $O(100 \\GeV)$ mass that scatters inelastically off nuclei to a higher mass state. For a mass splitting of $\\delta \\sim 100 \\keV$, comparable to typical WIMP kinetic energies, this produces a nuclear recoil energy spectrum that is peaked at high energies, in contrast to the falling exponential expected from elastic scattering. Such spectra can avoid constraints from experiments that focus on low-energy recoils.  iDM is most constrained by experiments that observe spectra at high recoil energies. The XENON10 collaboration has analyzed their data up to recoil energies of 75 keV and found no dark matter signal \\cite{Xe10Results}. Additionally, iDM preferentially scatters off heavy nuclei. The null results from  CRESST-II \\cite{CRESSTIICommissioning,CRESSTTalk} and ZEPLIN-III \\cite{Akimov:2010vk} rule out most of the parameter space under the assumption of a standard Maxwell-Boltzmann distribution of DM halo velocities. Other halo models may relax some of these constraints and open parts of parameter space \\cite{MarchRussell:2008dy,Non-Maxwellian}.  Another avenue that has been explored to explain the DAMA signal is light dark matter (LDM) \\cite{Petriello:2008jj,Fairbairn:2008gz,Chang:2008xa,ModerateChanneling}. This relies on the possibility of a channeling effect in DAMA. DAMA can only measure the electronic energy (denoted by units of ``keVee\") deposited in the detector by a recoiling nucleus, which is typically a small fraction of the total recoil energy (in units of ``keVr\"). Some recoils however may be ``channeled\" through the NaI crystal and deposit almost 100\\% of their energy electronically. While the DAMA collaboration has not reported experimental measurements of the channeling effect in their detector, they have theoretically estimated it to be $\\sim 30\\%$ for recoil energies of $\\sim 3 \\keV$ \\cite{DAMAChanneling}. If channeling occurs then DAMA can observe nuclear recoils at lower energies than previously anticipated. This would give DAMA a low energy threshold compared to most other experiments, allowing it to probe light ($\\sim 10 \\GeV$) dark matter that other experiments may be blind to. Recently it has been noted that light dark matter can also explain the excess of low-energy events observed by the CoGeNT experiment \\cite{CoGeNT}, for a slightly different parameter space than the fit to DAMA \\cite{Fitzpatrick:2010em,Interpretations}.   However, this scenario is severely constrained by the null results from XENON10 and CDMS Silicon, which also have low thresholds. In particular, we find that incorporation of all the CDMS Silicon datasets strongly disfavors this explanation of DAMA (see Figure \\ref{Fig:ElasticMSigma}).   In this paper we explore the possibility of explaining the DAMA signal through exothermic dark matter (exoDM); i.e. dark matter that can exist in two states with a small mass splitting, just as in conventional iDM, but which scatters from an excited state to a lower state to produce the signal observed by DAMA. (Although this model may also be described as ``inelastic dark matter\", we will use that term to refer to the upscattering scenario). In this exothermic process, the energy of the recoiling nucleus is peaked around a scale that is proportional to the splitting between the dark matter states and is inversely proportional to the nuclear mass. Consequently, the nuclear recoils caused by exoDM are more visible in experiments with light nuclei and low thresholds.   The approach of this model to reconciling the DAMA results with other experiments is similar to earlier proposals of light dark matter: we consider a parameter space which produces a modulation signal at DAMA while scattering below the recoil energy detection threshold in other experiments. The assumption of downscattering allows us to fit DAMA with lower mass WIMPs ($2 - 5 \\GeV$) which are less constrained by other existing experiments. There is some tension in this model between fitting the DAMA modulation and not exceeding the unmodulated rates observed in the same experiment; however as we will discuss these constraints rely on assumptions about the response of the DAMA detectors at very low energies which has not been well-measured. Similarly, XENON10 could also constrain this parameter space, but these constraints also depend upon uncertainties in the very-low-energy response of XENON10. The CDMS experiment does not constrain this parameter space.  The CoGENT excess may be accounted for with the same parameters that fit DAMA if some fraction of events are channeled in germanium as well.  The model-building aspects of the exoDM scenario are similar to conventional iDM. Generically,  most iDM scenarios lead to a cosmologically long lived relic population of excited states \\cite{MetastableWIMPs,pospelov}. As noted previously, downscattering of such excited states with $O(100 \\keV)$  splittings can produce dramatic signals at high energies in direct detection experiments \\cite{MetastableWIMPs,pospelov,Lang:2010cd}, though these can strongly constrain inelastic upscattering explanations of DAMA. Our parameter space in contrast has splittings of a few keV and produces peaked signals at very low energies. Because this signal is below threshold for most experiments, we evade current bounds; however because of the very high cross sections and rates we predict striking signals for experiments with sufficiently low thresholds. ",
        "conclusions": "The nature of dark matter is one of the great cosmological mysteries of our time. The observed complexity in the low energy dynamics of the standard model, with its multitude of cosmologically stable states, makes it plausible that the dark sector could also exhibit similar complexity. Interactions between such a dark sector and the standard model will  generically include exothermic interactions wherein a metastable dark matter particle dumps energy into a nucleus.  The phenomenology of such interactions in a direct detection experiment is different from that of the conventional elastic scattering between dark matter and a nucleus. The energy deposited in the recoiling nucleus peaks around $\\sim \\delta \\frac{m_\\chi}{m_N}$, with a spread around the peak determined by the kinetic energy of the dark matter. The modulation of this kinetic energy spread with the dark matter velocity gives rise to a modulation in the event rate at any given energy, even though the total rate is constant.  Such a modulation could explain the observations of DAMA while remaining consistent with null observations at a variety of other experiments. Exothermic dark matter preferentially deposits more energy into lighter nuclei than heavier nuclei. This is in stark contrast to endothermic ({\\it i.e.} upscattering) dark matter interactions that preferentially scatter off heavier nuclei. The absence of such events in  CRESST-II \\cite{CRESSTIICommissioning,CRESSTTalk}, which uses  tungsten (a heavy nucleus), severely constrains these explanations. Other explanations such as form factor dark matter \\cite{FormFactor} and momentum dependent dark matter \\cite{Chang:2009yt} that also preferentially scatter off of heavy nuclei are also severely constrained by CRESST-II.  The overall energy scale of nuclear recoils caused by exothermic dark matter is determined by the mass splittings in the dark sector and not by the dark matter kinetic energy. This enables exoDM to fit the DAMA signal with very light WIMPs (2 - 5 keV), while ensuring that the tail of the recoil spectrum remains below threshold in other experiments. This is in contrast to light elastic dark matter  \\cite{Fitzpatrick:2010em} explanations of DAMA where the somewhat larger masses needed to accommodate the DAMA signal causes events in low threshold experiments like CDMS Silicon. The absence of such events severely constrains the light elastic dark matter explanations of DAMA.  Currently, the most stringent constraints on exothermic dark matter arise from the DAMA event rate at $\\sim $ keV. Uncertainties in the detector efficiency and energy resolution at these energies makes it difficult to impose stringent limits on this scenario. A better understanding of the response of dark matter detectors at around $\\sim$ keV could significantly constrain exothermic dark matter. In particular the installation of new photomultiplier tubes in DAMA may allow an even lower energy threshold \\cite{DAMAResults}, leading to additional constraints from the low energy modulated and unmodulated rates. The constraints (or lack thereof) from xenon experiments could be clarified by future measurements of $\\Leff$ for liquid xenon, such as those planned by the XENON100 collaboration \\cite{Collaboration:2010er}.  Exothermic dark matter can explain the DAMA modulation using a large cross section without running afoul of other experimental bounds because the nuclear recoils typically occur below the threshold energies of these experiments. Low-threshold experiments and analyses are required to probe this scenario. Studies of low-energy channeling are necessary to determine the true sensitivities of crystal detectors such as CoGeNT to exoDM, and in particular to determine whether the model may simultaneously fit the DAMA and CoGeNT signals. In experiments with heavy nuclei such as XENON10 and CRESST-II, the energy spectrum of exoDM peaks at very low energies, but the event rate is enhanced by the $\\sim A^2$ scaling in the cross-section for spin-independent interactions. Because of the very high rate the tail of the spectrum may be visible for sufficiently low thresholds. Experiments with light nuclei such as CDMS Silicon, COUPP and CRESST can be directly sensitive to the peak region for thresholds of $\\sim 1 - 2$ keV. Low threshold analyses of these experiments could confirm or rule out this explanation of the DAMA signal."
    },
    "1004/1004.2932_arXiv.txt": {
        "abstract": "We have designed and built the first band-limited coronagraphic mask used for ground-based high-contrast imaging observations. The mask resides in the focal plane of the near-infrared camera PHARO at the Palomar Hale telescope and receives a well-corrected beam from an extreme adaptive optics system. Its performance on-sky with single stars is comparable to current state-of-the-art instruments: contrast levels of $\\sim10^{-5}$ or better at 0.8\" in $K_s$ after post-processing, depending on how well non-common-path errors are calibrated. However, given the mask's linear geometry, we are able to conduct additional unique science observations. Since the mask does not suffer from pointing errors down its long axis, it can suppress the light from two different stars simultaneously, such as the individual components of a spatially resolved binary star system, and search for faint tertiary companions. In this paper, we present the design of the mask, the science motivation for targeting binary stars, and our preliminary results, including the detection of a candidate M-dwarf tertiary companion orbiting the visual binary star HIP 48337, which we are continuing to monitor with astrometry to determine its association. ",
        "introduction": "The band-limited mask (BLM) is an occulter located in the image plane of the Lyot coronagraph. It suppresses telescope diffraction by manipulating the electric-field amplitude of incident starlight \\citep{kuchner_traub_02}. Compared to other coronagraphic designs, the BLM makes efficient use of photons \\citep{guyon_06}, is robust to low-order optical aberrations \\citep{kcg_05,sg_05,crepp_06} and finite stellar size (Crepp et al. 2009), and has generated the deepest contrast in the lab to date \\citep{trauger_traub_07}. It is a viable option for directly detecting terrestrial exoplanets from space \\citep{levine_09} and is currently slated for use with the James Webb Space Telescope NIRCAM \\citep{krist_09}.  Although the BLM shows great promise for conducting large dynamic-range observations, one has never been tested on an astrophysical source. Moreover, the most common type of BLM used for numerical simulations and lab experiments has a linear structure -- as opposed to an azimuthally symmetric (radial) profile -- and can be used to search for tertiary companions by blocking the light from two stars at the same time. This simple geometric feature provides access to an important and as yet unexplored observational parameter space: visual binary stars as high-contrast imaging targets.  To justify use of a BLM (of any shape) from the ground, strehl ratios in excess of $\\approx0.88\\:S_{qs}$ are required, where `qs' stands for quasi-static and $S_{qs}<1$ is the Strehl ratio provided by the optical system in the absence of atmospheric turbulence \\citep{crepp_07}. This level of correction corresponds to the `extreme' adaptive optics (AO) regime at near-infrared wavelengths. We have access to an AO system and a 1.6m diameter\\footnote{Telescope diameter as projected onto the primary mirror.} unobscured and well-corrected off-axis subaperture at Palomar that generates Strehl ratios as high as 96\\% in the $K_s$-band under good seeing conditions \\citep{serabyn_07}. Upon satisfying the aforementioned criterion, we have built a BLM for on-sky tests and installed it in the Palomar High Angular Resolution Observer (PHARO, \\cite{hayward_01}) camera on the Hale telescope. The aim of this paper is to demonstrate the general utility of the BLM and to highlight the importance of studying the immediate vicinity of binaries. ",
        "conclusions": "We have designed and built the first band-limited mask (BLM) used for ground-based high-contrast imaging observations. The BLM is installed in the PHARO camera at Palomar and is capable of generating $5\\sigma$ contrast levels of order $10^{-5}$ at subarcsecond separations in the $K_s$ band when operating in tandem with an extreme AO system and performing PSF reference subtraction. We have demonstrated this technology on the star $\\epsilon$ Eridani, achieving a mass sensitivity of $\\approx24.0M_J$ at 0.8\" and $\\approx14.7M_J$ at 4.0\". These observations are the most sensitive in $K_s$ yet reported for the star inside of 5\" to date, even though the aperture used is much smaller than in previous studies. Further improvements to calibration of non-common-path errors between the AO system and PHARO will provide even deeper contrast.  We have outlined two science cases that motivate binary star high-contrast imaging. The search for tertiary companions in P-type orbits can significantly improve our understanding of: (i) star formation, by improving the statistics of the Multiple-Star Catalog \\citep{tokovinin_97,tokovinin_04}, which is currently incomplete beyond $\\sim10$ pcs, and (ii) planet formation, by complementing the efforts of radial velocity teams searching for planets in S-type orbits. We have already detected a candidate M-dwarf tertiary orbiting the visual binary HIP 48337, showing that a linear BLM is a viable tool for exploring this observational parameter space. It is possible to speculate that the companion, if gravitationally bound, may be exciting the eccentricity of HIP 48337 A,B via the Kozai mechanism: their eccentricity is $\\sim 0.93$ according to the USNO WDS catalog.  We thank Karl Stapelfeldt for support at the JPL microdevices lab to cut and clean the BLM prior to installation and Dimitri Mawet for helpful discussions at the observatory. J. Carson acknowledges support from the NASA postdoctoral program. This work was funded in part by the UCF-UF-SRI program and NASA grant NNG06GC49G.  \\begin{footnotesize}"
    },
    "1004/1004.1336_arXiv.txt": {
        "abstract": "Acoustic waves and pulses propagating from the solar photosphere upwards may quickly develop into shocks due to the rapid decrease of atmospheric density. However, if they propagate along a magnetic flux tube, then the nonlinear steepening may be balanced by tube dispersion effects. This may result in the formation of sausage soliton. The aim of this letter is to report an observational evidence of sausage soliton in the solar chromosphere. Time series of Ca II H line obtained at the solar limb with the Solar Optical Telescope (SOT) on the board of Hinode is analysed. Observations show an intensity blob, which propagates from 500 km to 1700 km above the solar surface with the mean apparent speed of 35 km s$^{-1}$. The speed is much higher than expected local sound speed, therefore the blob can not be a simple pressure pulse. The blob speed, length to width ratio and relative intensity correspond to slow sausage soliton propagating along a magnetic tube. The blob width is increased with height corresponding to the magnetic tube expansion in the stratified atmosphere. Propagation of the intensity blob can be the first observational evidence of slow sausage soliton in the solar atmosphere. ",
        "introduction": "Energy transport from the solar photosphere towards the corona, which eventually may lead to coronal heating, is still an open problem. There are several possible ways of the energy transport: waves, pulses or electric currents. The energy transport by the waves has been recently observed through the oscillatory motions of plasma in the chromosphere  \\citep{kukhianidze,zaqarashvili,De Pontieu,Jess,zaqarashvili1}. On the other hand, the dynamic photosphere may excite pulses due to convective shootings and/or magnetic reconnections, which then may propagate upwards. Several kinds of impulsive events are frequently observed on the solar disc: chromospheric bright grains  \\citep{lites},  blinkers \\citep{Harrison} and explosive events \\citep{Porter}. Recent observations by Hinode spacecraft revealed various types of energetic events such as chromospheric jet-like structures \\citep{Katsukawa,Shibata,Nishizuka} and type II spicules \\citep{De Pontieu1}. However, direct observational evidence of pulse propagation at the solar limb from the photosphere upwards, to our knowledge, was not reported yet.  Upward propagating pressure pulses may quickly steepen into shocks due to the rapid decrease of density. However, if the pulses propagate along magnetic flux tubes, then tube dispersive effects may prevent the nonlinear steepening. This may lead to the formation of a soliton, which is a stable structure propagating without significant change of shape. The formation of sausage solitons in magnetic tubes first has been suggested by \\citet{Roberts1}. Since that, numerous papers addressed the  soliton formation problem \\citep{Roberts,Roberts2,Merzljakov,Merzljakov2,Sahyouni,Ofman,zhugzhda,Nakariakov,ruderman,ballai,erdelyi,ryutova}. Most of the studies consider a sausage soliton ($m=0$ mode in magnetic tubes), but no observational support to the theory was reported yet. On the other hand, some observations suggest the propagation of nonlinear soliton-like kink waves ($m=1$ mode in tubes) identified with moving magnetic features around sunspots \\citep{ryutova}.  Here we report the upward propagation of a pressure blob in time series of Ca II H line obtained by Hinode/SOT \\citep{Tsuneta}. Estimated parameters of the blob fit with a solution of slow sausage soliton propagating along a magnetic tube. Therefore, we suggest that this is the first observational evidence of sausage soliton propagation in the lower solar atmosphere.   \\begin{figure} \\begin{center} \\includegraphics[width=9.3cm]{fig1.eps} \\end{center} \\caption{Corrected Ca II H image of quiet Sun obtained by Hinode/SOT. The image was rotated by 90$^0$, therefore the $x$-axis corresponds to the Solar-$Y$ and the $y$-axis corresponds to the Solar-$X$. The white arrow shows the place of intensity blob propagation. } \\end{figure} ",
        "conclusions": "\\begin{enumerate} \\item Time series of Ca II H line obtained by Hinode/SOT at the solar limb shows upward propagating intensity blob. The blob appears at 500-600 km height above the surface and reaches to the height of $\\sim$ 1700 km after 35 s. Therefore, the mean apparent propagation speed is 35 km s$^{-1}$. The blob has elongated form and the length to width ratio is $\\sim$ 3 in average. The length to width ratio, the relative intensity and the propagation speed change slightly during the propagation. \\item The observed parameters fit with theoretically expected properties of slow sausage soliton propagating along a magnetic flux tube, which has the Alfv\\'en speed of $\\sim$ 70 km s$^{-1}$ and is in temperature balance with surroundings (note, that an inclination of the tube along the line of sight may slightly increase the value of Alfv\\'en speed). Therefore, we suggest that this is the first observational evidence of slow sausage soliton in the solar atmosphere. \\item The width of the blob increases with height, which coincides with the expected expansion of magnetic tubes in the stratified atmosphere. \\end{enumerate}"
    },
    "1004/1004.1046_arXiv.txt": {
        "abstract": "{We study the problem of searching for cosmic string signal patterns in the present high resolution and high sensitivity observations of the Cosmic Microwave Background (CMB). This article discusses a technique capable of recognizing Kaiser-Stebbins effect signatures in total intensity anisotropy maps from isolated strings. We derive the statistical distributions of null detections from purely Gaussian fluctuations and instrumental performances of the operating satellites, and show that the biggest factor that produces confusion is represented by the acoustic oscillation features of the scale comparable to the size of horizon at recombination. Simulations show that the distribution of null detections converges to a $\\chi^2$ distribution, with detectability threshold at $99\\%$ confidence level corresponding to a string induced step signal with an amplitude of about 100 $\\muK$ which corresponds to a limit of roughly $G\\mu \\sim 1.5\\times 10^{-6}$. We implement simulations for deriving the statistics of spurious detections caused by extra-Galactic and Galactic foregrounds. For diffuse Galactic foregrounds, which represents the dominant source of contamination, we construct sky masks outlining the available region of the sky where the Galactic confusion is sub-dominant, specializing our analysis to the case represented by the frequency coverage and nominal sensitivity and resolution of the Planck experiment. As for other CMB measurements, the maximum available area, corresponding to 7\\%, is reached where the foreground emission is expected to be minimum, in the 70-100 GHz interval.} ",
        "introduction": "Current theories of particle physics predict spontaneous symmetry breaking in the early Universe, with the consequent creation of coherent structures in the spatial distribution of the fields involved in the process; these structures, known as topological defects, represent genuine tracers and unique remnants of those processes, occurring at energies which are not achievable with the current terrestrial particle colliders, see \\cite{copelandkibble} and references therein. Among the most known and studied relics of this kind, cosmic strings correspond to an uni-dimensional region in which a (scalar) field is trapped by true minima of its potential, storing energy in that region. The energy density is parametrized by the dimensionless quantity $G\\mu$ where $G$ is the Newton's constant, and $\\mu$ the string tension, see \\cite{copelandkibble} for review.\\\\ Years ago, while it was still unclear if strings could dominate the structure formation process, their signal was searched in the main statistical indicators of structure formation itself, like the power spectra of anisotropies in the Cosmic Microwave Background (CMB) or Large Scale Structure (LSS). The conclusion of these studies was that strings should contribute to cosmic structure formation by no more than a few percent \\cite{pogosian.et.al,wu,fraisse2,daviskibble,wyman.et.al,seljakslosar}. Indeed, the evidence for coherent acoustic oscillations in the CMB and LSS power spectra, dominated by density fluctuations, favored the scenario of isentropic or adiabatic Gaussian fluctuations like predicted in the simplest inflationary scenarios, and indicated that cosmic strings, if any, played a minor role in providing the initial conditions, in a statistical sense, to the Universe we see today. Early in their study, it was evident that strings produced perturbations that were not coherent as the observed acoustic oscillations and that they tend to produce equal power in scalar (density), vector (vorticity) and tensor (gravitational waves) perturbation modes. Recently, \\cite{seljakslosar} proposed to look at sub-dominant, but most interesting components of the power spectrum of CMB anisotropies, namely the B modes activated by cosmological gravitational waves, to constrain cosmic strings.\\\\ The hypothesis is then that these objects are washed away by inflation itself motivating the direct search of the signal from isolated and rare cosmic strings. Among these studies, those on CMB are based on the well known Kaiser-Stebbins effect \\cite{KaiserStebbins} resulting in a step like feature in the CMB temperature caused by a string which is orthogonal and moving relativistically with respect to the line of sight. These studies produce upper limits on the abundance of string in our own observable Universe of course, but more quantitatively on the string density parameter $G\\mu$ \\cite{hindmarsh1,perivolaropoulossimatos,jeongsmoot1,lowright,pogosian.et.al2,kuijken.et.al,gasparini.et.al,christiansen.et.al,morganson.et.al}. On the other hand, studies in the optical band look for strong lensing events, in which the strings bend the light from a distant Galaxy, causing a mirror image of that\\cite{vilenkinshellard}.\\\\ In this paper, we specifically search for signals imprinted on CMB temperature observations by long cosmic strings of cosmological scale, which was predicted by theory and also supported by simulations\\footnote{http://www.damtp.cam.ac.uk/research/gr/public/cs\\_evol.html}. The long strings between the last scattering surface and us would be back-lighted by CMB photons and leave signatures in CMB temperature map via the Kaiser-Stebbins effect. While the works mentioned above are dealing with the effects of cosmic string network on statistical observables, this work aims at the direct search of signals produced by isolated cosmic strings left on CMB temperature anisotropy maps, having a coherence length corresponding to at least the angle subtended by an Hubble horizon at decoupling. The Planck survey\\footnote{www.rssd.esa.int/planck} of CMB temperature and polarization anisotropies is ongoing \\cite{mandolesi.et.al}, promising a leap forward in many research fields, including constraints on cosmic strings abundance and tensions. The unprecedented frequency coverage, sensitivity and angular resolution requires a dedicated discussion targeted to the applicability and expectations from Planck direct searches of cosmic strings. In this paper we work out two missing issues concerning the direct search of KS effects of cosmic strings. The first one is represented by the statistics of null detections from pure Gaussian CMB anisotropies, with nominal noise and angular resolution of the operating satellites, which, needs to be quantified prior to data analysis. Consequently, as the second step, one can use the available models and data of Galactic and extra-Galactic foreground signal simulations at Planck frequencies, to select the sky area where false detections provide a negligible effect with respect to the main contaminant represented by the background, Gaussian and uniformly distributed, CMB anisotropies. In this paper we discuss and specialize the technique discussed in \\cite{jeongsmoot1} to deal with these issues. In Section 2 we review in detail the effects of cosmic strings on the CMB temperature anisotropies and outline the string search algorithm. In Section 3, we apply the technique to simulated Planck data sets, quantifying the statistics of null detection based on the nominal performances of the satellite, as well as selecting the sky region which is available for searches based on the current knowledge of Galactic and extra-Galactic foreground emission. Concluding remarks are in Section 4. ",
        "conclusions": "We studied a direct cosmic string search based on the Cosmic Microwave Background (CMB) anisotropies in total intensity, in the context of the existing satellite observations, focusing on the case of the ongoing \\PLANCK survey. The search algorithm uses an algebraic method to detect step like structures in CMB maps, caused by the Kaiser-Stebbins effect induced by strings moving on the orthogonal plane with respect to the line of sight. We quantified two important aspects towards the application of this algorithm to the actual data.\\\\ First, we derived the statistics of null detections due to instrumental noise and purely Gaussian CMB anisotropies as predicted within the $\\Lambda$ Cold Dark Matter ($\\Lambda$CDM) cosmology. We have shown that the biggest cause of confusion for the detection algorithm is caused by the CMB acoustic oscillations at wavelengths comparable to the size of the horizon at decoupling, while noise and angular resolution which are typical of the operating satellites, are less important. The distribution of null detection converges to a $\\chi^{2}$-distribution. We derived the confidence levels of detection of cosmic strings in \\PLANCK maps by estimating the 68\\%, 95\\% and 99\\% of cumulative occurrence of null detection, thresholds give in Table \\ref{table2}, which corresponds to $G\\mu\\sim 1.5\\times 10^{-6}$ for 95\\% of occurrences in terms of cosmic string tension, according to the relation between the string tension and temperature shift of CMB photons in \\ref{eq1}.\\\\ On the basis of these results, we were able to evaluate the effect of extra-Galactic and Galactic foreground emissions on cosmic string searches, by comparison the foreground induced string signals with the detectability threshold found in the first part of this work. As expected, we find that after removal of the brightest sources exceeding $5\\sigma$ times the CMB rms, residual confusion due to unresolved point sources does not have a significant impact on the string detection algorithm. On the other hand, the diffuse Galactic emission from our own Galaxy do cause false detection in a significant fraction of the sky, not because of their overall intensity, but rather due to their gradients, reaching intermediate Galactic latitudes. By using the current data and models of the diffuse Galactic emissions based on off microwave band data as well as the measurements of the \\WMAP satellite, we determine the area of the sky where false detection are expected due to the Galactic emission, at the level of 68\\%, 95\\% and 99\\% of cumulative occurrences for each of seven \\PLANCK frequencies, where the thresholds are given in Table \\ref{table2}. We specialize the discussion to seven frequency channels on-board \\PLANCK, constructing for each one the sky mask to be used for string searches. The maximum available area in which the Galactic contamination is sub-dominant with respect to the contribution from Gaussian CMB fluctuations is found in the frequency interval 70-100 GHz, and represents about 7\\% of the entire sky. These criteria and findings can be used in the actual application of string searches in CMB data.\\\\ For the future work on the pattern search of cosmic strings in CMB temperature data, it is desirable to exploit the correlated signals of discrete steps lined up close to each other, since the pattern search is mainly pointed at the long strings of cosmological scale rather than isolated segments or loops of cosmic strings. The methology provided in this paper on finding step signal by a cosmic string segment will be the basis for the long string search.  \\subsection"
    },
    "1004/1004.0836.txt": {
        "abstract": "We present a comparative study of the thermal emission of the transiting exoplanets WASP-1b and WASP-2b using the {\\it Spitzer Space Telescope}. The two planets have very similar masses but suffer different levels of irradiation and are predicted to fall either side of a sharp transition between planets with and without hot stratospheres. WASP-1b is one of the most highly irradiated planets studied to date. We measure planet/star contrast ratios in all four of the IRAC bands for both planets (3.6--8.0\\,\\um), and our results %We find planet/star contrast ratios that indicate the presence of a strong temperature inversion in the atmosphere of WASP-1b, particularly apparent at 8\\um, and no inversion in WASP-2b. In both cases the measured eclipse depths favor models in which incident energy is not redistributed efficiently from the day side to the night side of the planet. We fit the \\spitzer\\  light curves simultaneously with the best available radial velocity curves and transit photometry in order to provide updated measurements of system parameters. We do not find significant eccentricity in the orbit of either planet, suggesting that the inflated radius of WASP-1b is unlikely to be the result of tidal heating. Finally, by plotting ratios of secondary eclipse depths at 8\\um\\ and 4.5\\um\\ against irradiation for all available planets, we find evidence for a sharp transition in the emission spectra of hot Jupiters at an irradiation level of $2\\times 10^9\\,\\rm erg\\,s^{-1}\\,cm^{-2}$. We suggest this transition may be due to the presence of TiO in the upper atmospheres of the most strongly irradiated hot Jupiters. ",
        "introduction": "%One of the great achievements of the {\\it Spitzer Space Telescope} has been %the first measurements of the thermal emission of exoplanets orbiting %main sequence stars \\citep{Deming05,Charbonneau05}. Such measurements %exploit the secondary eclipses of transiting exoplanets, where the %planet is eclipsed by the star. The depth of the secondary eclipse provides %a direct measurement of the brightness of the planet.  %One of the great achievements of the {\\it Spitzer Space Telescope} %\\citep{Werner04} has been the The {\\it Spitzer Space Telescope} \\citep{Werner04} has been used to carry out the first photometry and emission spectroscopy of exoplanets that orbit main sequence stars \\citep{Deming05,Charbonneau05,Richardson07,Grillmair07}. This %has been was achieved by observing transiting planets at secondary eclipse (when the planet is eclipsed by the star) which allows the %thermal emission of the planet to be separated from that of the star. Photometry and spectroscopy are the key measurements needed to determine the physical properties of any astronomical object, and secondary eclipse observations allow us to consider the temperature structure and chemical composition of exoplanet atmospheres and the redistribution of energy from the day side to the night side of the planet \\citep[e.g.][]{Burrows05,Fortney05,Seager05,Barman05}.   \\begin{table*} \\begin{center} \\caption{Log of \\spitzer\\/  observations of WASP-1b and WASP-2b.} \\label{tab-log} \\begin{tabular}{crcccrrrrrc} %\\hline Target & Prog. & Date & Start time & Duration & Bary.\\ corr. &\\multicolumn{4}{c}{No.\\ of frames $\\times$ effective exposure per frame}&Pipeline\\\\ &          &      & UTC        & s  & s      & 3.6\\,$\\rm \\mu m$ & 4.5\\,$\\rm \\mu m$ & 5.8\\,$\\rm \\mu m$ & 8.0\\,$\\rm \\mu m$ & version\\\\ & & & & & \\\\ WASP-1b & 30129 & 2007-09-08 & 14:02:56 & 27\\,586 & $+238.2$ & $2072 \\times 10.4$\\,s & & $2072 \\times 10.4$\\,s & & S16.1.0\\\\ &   282 & 2006-12-30 & 14:15:27 & 27\\,146 & $+252.5$ & & $2055 \\times 10.4$\\,s & &$2055 \\times 10.4$\\,s & S15.0.5\\\\   WASP-2b & 30129 & 2007-07-01 & 13:44:18 & 12\\,427 & $+240.3$ & $1818 \\times \\phantom{0}1.2$\\,s & & $909 \\times 10.4$\\,s & & S16.1.0 \\\\ &   282 & 2006-11-28 & 08:23:11 & 12\\,032 & $-11.8$ & & $910 \\times 10.4$\\,s & & $910 \\times 10.4$\\,s & S15.0.5 \\\\ %\\hline \\end{tabular} \\end{center} \\end{table*}     %The secondary eclipse signal is weak even in the best cases %(of order 0.1 per cent in hot Jupiters) and secondary eclipse detections %from the ground, while possible \\citep{}, are extremely challenging \\citep{}.  Secondary eclipse detections have been made from the ground \\citep[e.g.][]{Mooij09,Sing09} but since %The secondary eclipse the signal is weak even in the best cases (of order 0.1 per cent in hot Jupiters) the bulk of measurements to date are from space. Together with detections in the optical with \\corot\\ and \\kepler\\ \\citep{Snellen09,Alonso09,Borucki09} and the near infrared with \\hst\\ \\citep{Swain09a,Swain09b}, \\spitzer\\ is providing an increasingly clear %a tantalizing picture of the thermal emission of exoplanets.  %and near infra-red detections from the %ground have remained elusive. %% (e.g. refs ??). %A marginal detection has been made of OGLE-TR-113b in the K band %\\citep{Snellen07}, but \\spitzer\\ remains the only instrument proven to be %capable of secondary eclipse detections of large numbers of exoplanets.  \\spitzer\\ detections of thermal emission have been reported in various combinations of wavebands between 3.6 and 24\\,\\um. The bulk of the newly-discovered hot Jupiters are detectable only in the shorter wavelength \\spitzer\\ bands of the IRAC instrument \\citep[3.6, 4.5, 5.8 \\& 8.0\\,$\\rm \\mu m$;][]{Fazio04} but fortunately this is a range in which strong molecular bands are expected, providing good constraints on atmospheric conditions. The IRAC observing modes also allow more efficient observations of these fainter systems, to some extent compensating for their lower brightness. The rapidly growing number of moderately-bright transiting hot Jupiters therefore provide an excellent opportunity to improve our understanding of exoplanet atmospheres. % during the remaining cryogenic and warm phases %of the \\spitzer\\ mission.  %To date, \\spitzer\\ detections of thermal emission have been reported for six %exoplanets in various combinations of \\spitzer\\ bands between %3.6 and 24\\,$\\rm \\mu m$. These consist of: the three brightest transiting hot %Jupiters, %HD209458b \\citep{Deming05,Knutson08}, %HD189733b \\citep{Deming06,Knutson07,Charbonneau08,Knutson09a} %and HD149026 \\citep{Harrington07}; %two fainter hot Jupiters, %TrES-1 \\citep{Charbonneau05} %and XO-1b \\citep{Machalek08}; %and the Neptune-sized transiting planet, %GJ\\,436 \\citep{Deming07,Demory07}. The fainter hot Jupiters %are only detectable in the shorter wavelength \\spitzer\\ bands of the IRAC %instrument \\citep[3.6, 4.5, 5.8 \\& 8.0\\,$\\rm \\mu m$;][]{Fazio04} %but fortunately this is the range in which the strongest molecular bands are %expected, %providing the best constraints on atmospheric conditions. The %rapidly growing number of moderately bright transiting hot Jupiters %therefore provide an excellent opportunity to improve our understanding of %exoplanet atmospheres during the remaining cryogenic and warm phases %of the \\spitzer\\ mission.  The existing \\spitzer\\ observations show that hot Jupiters are strongly heated by their parent stars, with typical brightness temperatures in the range 1000--2000\\,K, and that their spectra deviate strongly from black bodies. Preliminary theoretical calculations predicted strong molecular absorption in the IRAC bands \\citep{Burrows05,Fortney05,Seager05,Barman05}, for which there is evidence in some systems %Some systems have infra-red brightnesses that appear to be consistent with %expected molecular absorption bands %\\citep[e.g. HD189733b and TrES-1;][]{Charbonneau08,Charbonneau05}, \\citep[e.g. HD189733b;][]{Charbonneau08}, however other systems have measured brightness temperatures in excess of expectations, indicating emission features in the IRAC bands %\\citep[e.g. HD209458b and HD149026b;][]{Knutson08,Harrington07}. \\citep[e.g. HD209458b;][]{Knutson08}. In these cases it is thought that an opacity source high in the atmosphere results in a hot stratosphere and that the molecular bands are driven into emission by the temperature inversion \\citep{Hubeny03, Fortney06, Harrington07, Burrows07, Sing08}. \\citet{Fortney08} and \\citet{Burrows08} suggest that %It has been suggested that irradiated exoplanets may fall into two distinct classes, those with and those without hot stratospheres, depending on the level of incident stellar flux \\citep[dubbed pM and pL class planets respectively by][]{Fortney08}. The brightest and best studied systems, HD209458b and HD189733b, fall either side of the predicted transition between these classes \\citep{Fortney08}, and their IRAC fluxes support the presence of a hot stratosphere in HD209458b \\citep{Burrows07,Knutson08} and its absence in HD189733b \\citep{Charbonneau08}. Most of the other systems that have been observed in all four IRAC bands also seem to support this overall picture \\citep{Knutson09b,Machalek09,Todorov10,O'Donovan10,Machalek10,Campo10}, %HD149026b and TrES-1 also seem consistent with this picture %(class pM and pL respectively), but XO-1b and TrES-3 do not, with XO-1b presenting evidence for a temperature inversion despite low irradiation \\citep{Machalek08,Machalek09}, and TrES-3 not exhibiting evidence for a temperature inversion despite high irradiation \\citep{Fressin10}. It may be that additional parameters dictate the presence of a temperature inversion, although to some extent the picture is confused by the different models and criteria for inversion detection applied by different authors. \\citet{Gillon10} show that the planets studied to date cover a wide region in color-color space, and do not fall clearly into two groups.  %Systems with hot stratospheres, dubbed pM class planets by \\citet{Fortney08}, %should have relatively high brightness temperatures and thus deep eclipses %in the IRAC bands, due to molecular bands in emission, %whereas systems without hot stratospheres, pL class, should have weaker %eclipses due to molecular bands in absorption. pM systems should also have %higher day/night contrast than pL planets since heat should be transported less %efficiently around the planet. This is seen in HD189733b \\citep{Knutson...?}.  In this paper we present \\spitzer\\ IRAC secondary eclipse detections of the transiting planets WASP-1b and WASP-2b, which were discovered by the Wide Angle Search for Planets (WASP) project \\citep{Cameron07a,Pollacco06}. These planets make an interesting pair for comparative study since they have near identical mass %($\\sim0.85\\,M_{\\rm J}$) ($0.9\\,M_{\\rm J}$) and yet WASP-1b is highly irradiated and expected to have a hot stratosphere, while WASP-2b is not \\citep{Fortney08}. Indeed, WASP-1b is one of the most highly irradiated planets studied with \\spitzer\\  to date (incident flux of $2.5\\times10^9\\,\\rm erg\\,s^{-1}\\,cm^{-2}$). It is also one of the group of oversized hot Jupiters that have radii larger than can be explained with canonical models \\citep{Charbonneau07}. %as is the confirmed pM class exoplanet HD209458b.  In addition to presenting and discussing \\spitzer\\ detections of WASP-1b and WASP-2b, we present revised parameters for both planets based on simultaneous %MCMC fits of all available photometry and spectroscopy. %(Sect.\\,\\ref{sec-mcmc}).  %WASP-1b and WASP-2b are transiting exoplanets discovered by the %Wide Angle Search for Planets (WASP) project \\citep{Cameron07a,Pollacco06}. %They make an interesting pair for comparative studies, %since they have near-identical masses %($\\sim0.85\\,M_{\\rm J}$), %%$0.79\\pm^{0.10}_{0.08}\\,M_{\\rm J}$ and %%$0.88\\pm^{0.06}_{0.09}\\,M_{\\rm J}$ respectively (** this paper **), %and yet WASP-1b has a much larger radius than WASP-2b: %$1.4\\,R_{\\rm J}$ compared with $1.0\\,R_{\\rm J}$ %%$1.44\\pm0.08\\,R_{\\rm J}$ compared with $1.04\\pm0.06\\,R_{\\rm J}$ %\\citep{Cameron07a,Charbonneau07}. %WASP-1b is also more heavily irradiated than WASP-2b, %since it orbits an F7V star at a separation of 0.04\\,AU, %while WASP-2b orbits a K1V star at a separation of 0.03\\,AU %\\citep{Cameron07a}. %Note that revised parameters for both planets are given in %Sect.\\,\\ref{sec-mcmc}. %%, which orbits an ** star at a separation of **, %%also suffers much stronger insolation than WASP-2b, which orbits an ** star %%as a separation of ** \\citep{}.  %%Due to this difference in incident flux, %Due to the difference in irradiation, %\\citet{Fortney08} predicted that WASP-1b should be in the %temperature inversion regime, %with emission features in the \\spitzer\\  IRAC bands (pM class), while %WASP-2b should lie just below the transition into the pL class, %expected to have absorption features in the \\spitzer\\  IRAC bands.  %In this paper we present \\spitzer\\  secondary eclipse observations of both %WASP-1b and WASP-2b, resulting in direct detections of the planetary %thermal emission in both cases. We discuss whether these detections are c %consistent the distinct classes of irradiated planetary %atmosphere predicted by \\citet{Fortney08}.   \\begin{figure*} \\begin{center} \\plotone{wasp1_all_raw_v4_lin.ps} %\\plottwo{wasp1_model.ps}{wasp2_model.ps} %\\includegraphics[width=8.5cm]{wasp1_model.ps} %\\includegraphics[width=8.5cm]{wasp1_model.ps}\\\\ \\caption[]{ Upper panels show the \\spitzer\\ IRAC light curves of the star WASP-1 during the expected times of secondary eclipse of the planet WASP-1b. Points show the measurements from individual images, crosses show the data binned into five hundred bins per orbital period. Solid lines show our best fits to the eclipse light curves, including linear decorrelation and other instrumental effects described in Sects.~\\ref{sec-insb} \\& \\ref{sec-sias}. The lower panels show the radial velocity curve of WASP-1 (measured with SOPHIE at OHP) and the z-band optical light curve during primary transit (from Keplercam), each plotted on different phase ranges. These data were fitted simultaneously with the secondary eclipse observations. } \\label{fig-w1-raw} \\end{center}  \\end{figure*} \\begin{figure*} \\begin{center} \\plotone{wasp2_all_raw_v4_lin.ps} %\\plottwo{wasp1_model.ps}{wasp2_model.ps} %\\includegraphics[width=8.5cm]{wasp1_model.ps} %\\includegraphics[width=8.5cm]{wasp1_model.ps}\\\\ \\caption[]{ Upper panels show the \\spitzer\\ IRAC light curves of the star WASP-2 during the expected times of secondary eclipse of the planet WASP-2b. Points show the measurements from individual images, crosses show the data binned into five hundred bins per orbital period. Solid lines show our best fits to the eclipse light curves, including linear decorrelation and other instrumental effects described in Sects.~\\ref{sec-insb} \\& \\ref{sec-sias}. The lower panels show the radial velocity curve of WASP-2 (measured with SOPHIE, Coralie and HARPS) and the z-band optical light curve during primary transit (from Keplercam), each plotted on different phase ranges. These data were fitted simultaneously with the secondary eclipse observations. } \\label{fig-w2-raw} \\end{center} \\end{figure*}    % ",
        "conclusions": "The measured secondary-eclipse depths of WASP-1b and WASP-2b presented here indicate a strong temperature inversion in the atmosphere of WASP-1b, but no temperature inversion in WASP-2b. This difference is likely to be related to the much higher level of irradiation of the atmosphere of WASP-1b. The eclipse depths of both planets also favor models in which incident energy is not redistributed efficiently from the day side to the night side of the planet. We do not find significant eccentricity in the orbits of either planet, suggesting that the inflated radius of WASP-1b is unlikely to have arisen through tidal heating. Finally, we find evidence for a sharp transition in the properties of planetary atmospheres with irradiation levels above $2\\times10^9\\,\\rm erg\\,s^{-1}\\,cm^{-2}$, and we suggest that this transition might be due to the presence of TiO in the upper atmospheres of the most strongly irradiated hot Jupiters.   %% If you wish to include an acknowledgments section in your paper, %% separate it off from the body of the text using the"
    },
    "1004/1004.5045_arXiv.txt": {
        "abstract": "{Massive stars have high-multiplicity fractions, and many of them have still undetected components, thus hampering the study of their properties.} {I study a sample of massive stars with high angular resolution to better characterize their multiplicity.} {I observed 138 fields that include at least one massive star with AstraLux, a lucky imaging camera at the 2.2 m Calar Alto telescope. I also used observations of 3 of those fields with ACS/HRC on HST to obtain complementary information and to calibrate the AstraLux data. The results were compared with existing information from the Washington Double Star Catalog, Tycho-2, 2MASS, and other literature results.} {I discover 16 new optical companions of massive stars, the majority of which are likely to be physically bound to their primaries. I also improve the accuracy for the separation and magnitude difference of many previously known systems. In a few cases the orbital motion is detected when comparing the new data with existing ones and constraints on the orbits are provided.} {The analysis indicate that the majority of the AstraLux detections are bound pairs. For a range of separations of 0\\farcs1-14\\arcsec\\ and magnitude differences lower than 8, I find that the multiplicity fraction for massive stars is close to 50\\%. When objects outside those ranges are included, the multiplicity fraction should be considerably higher.} ",
        "introduction": "Massive stars play a crucial role in the dynamical and chemical evolution of galaxies. They are the major source of ionizing and UV radiation and, through their huge mass-loss rates, they have a strong mechanical impact on their surroundings. Massive stars are also important because they are critical contributors to stellar and explosive nucleosynthesis. The nuclear products are ejected into space in the stellar winds and in the final supernova explosions that put an end to the massive stars' lives. Despite their importance, our understanding of these objects and of their evolution is still fragmentary due to their relatively small numbers, the existence of unresolved multiple systems, and to their concentration along the Galactic plane, where extinction affects their detection and the measurement of their distances and other properties.  A major ongoing project, the Galactic O-Star Spectral Survey, GOSSS (PI: J. Ma\\'{\\i}z Apell\\'aniz, see \\citealt{Walbetal10a,Gameetal08a}, Sota et al. in prep.), is currently obtaining ground-based spectroscopy of all known Galactic O stars with $B< 13$ with the purpose of remedying some of the gaps in our knowledge of Galactic massive stars. All stars are being observed at $R\\sim 3000$ to obtain accurate spectral classifications, and subsamples are being observed at $R\\sim 40\\,000$ at multiple epochs to detect spectroscopic binaries and at $R\\sim 1500$ to accurately measure their spectral energy distributions. The detection of spectroscopic binaries is especially important for massive stars because the multiplicity fraction among them is especially high \\citep{Masoetal98}, a fact that the preliminary GOSSS results are confirming and even increasing in value. Undetected binaries complicate the study of samples of stars because they introduce biases in the results and lead to incorrect conclusions. The presence of spatially unresolved components can alter the observed spectral type and, in some circumstances, yield line-ratio combinations that cannot be present in single O stars (e.g. different spectral types may be deduced from the He and the N lines, see e.g. \\citealt{Walbetal02b}). Therefore, correctly characterizing the multiplicity can shed light on objects with such composite spectra.  Time-resolved spectroscopy is crucial to detecting short-period multiple systems. However, many massive binaries are known to have long periods \\citep{Masoetal98} and it has been proposed that their period distribution may follow \\\"Opik's law \\citep{Opik24}, which would lead to essentially all massive stars being born in multiple systems. To verify that assertion for large-separation systems, spectroscopy is of little use because the periods involved are thousands of years or more and the velocity changes quite small. Therefore, one needs to use high angular-resolution methods in order to visually detect companions. Even so, a gap still exists between the two ranges easily detected by spectroscopic and high angular-resolution methods. For example, a 30 \\Ms/ + 20 \\Ms/ in a circular edge-on orbit with $a = 50$ AU at a distance of 1 kpc would have a maximum separation of 50 mas and a period of 50 years, both near or beyond the limit of current capabilities. Making the orbit characteristics, mass ratio, or orientation more unfavorable or placing the system at longer distances will make it even harder to detect, to the point that most LMC massive binaries are likely to be currently unknown.  The likely distribution of binary periods (or separations) among massive stars led me to start a project to observe as many massive stars as possible using high-resolution imaging in order to complement the multiple-epoch spectroscopy obtained with GOSSS. Such a systematic project is needed to eliminate biases in our knowledge of massive stars. Different imaging techniques have been attempted to observe binary systems: speckle interferometry, imaging from space, and adaptive optics, among them. Some of the recent attempts are finding previously undetected pairs with large magnitude differences, \\deltam\\ \\citep{Turnetal08,Maizetal10a}, thus opening a new part of the parameter space. In this paper I explore the use of lucky imaging, a technique that, to my knowledge, has not been systematically applied to massive stars. In future papers I plan to extend the sample to include several hundred more stars. ",
        "conclusions": ""
    },
    "1004/1004.0726_arXiv.txt": {
        "abstract": "We study the stability regions and families of periodic orbits of two planets locked in a co-orbital configuration. We consider different ratios of planetary masses and orbital eccentricities, also we assume that both planets share the same orbital plane. Initially we perform numerical simulations over a grid of osculating initial conditions to map the regions of stable/chaotic motion and identify equilibrium solutions. These results are later analyzed in more detail using a semi-analytical model.  Apart from the well known quasi-satellite (QS) orbits and the classical equilibrium Lagrangian points $L_4$ and $L_5$, we also find a new regime of asymmetric periodic solutions. For low eccentricities these are located at $(\\Delta \\lambda,\\Delta \\varpi) = (\\pm 60^\\circ,\\mp 120^\\circ)$, where $\\Delta \\lambda$ is the difference in mean longitudes and $\\Delta \\varpi$ is the difference in longitudes of pericenter. The position of these {\\it Anti-Lagrangian} solutions changes with the mass ratio and the orbital eccentricities, and are found for eccentricities as high as $\\sim 0.7$.  Finally, we also applied a slow mass variation to one of the planets, and analyzed its effect on an initially asymmetric periodic orbit. We found that the resonant solution is preserved as long as the mass variation is adiabatic, with practically no change in the equilibrium values of the angles. ",
        "introduction": "In the restricted three-body problem there are different domains of stable motion associated to co-orbital motion. Each can be classified according to the center of libration of the critical argument, $\\sigma = \\lambda -\\lambda$', where $\\lambda$ denotes the mean longitude of the minor body and $\\lambda$' the same variable for the disturbing planet. These types of motion are known as: {\\it (i)} Tadpole orbits, corresponding to a libration of $\\sigma$ around $L_4$ or $L_5$; {\\it (ii)} Horseshoe orbits, where motion occurs around $\\sigma = 180^{\\circ}$ and encompasses both equilateral Lagrangian Points, and {\\it (iii)} quasi-satellite (QS) orbits, where $\\sigma$ oscillates around zero.  The term ``quasi-satellites'' was originally introduced by Mikkola and Innanen (1997) and can be viewed as an extension of retrograde periodic orbits in the circular restricted three-body problem (e.g. Jackson 1913, H\\'enon 1969). Although not present for circular orbits, they exist for moderate to high eccentricities of the particle. In a reference frame rotating with the planet, QS orbits circle the planet like a retrograde satellite, although at distances so large that the particle is not gravitationally bounded to the planetary mass (Mikkola et al 2006).  The first object confirmed in a QS configuration was the asteroid $2002 VE68$ (Mikkola et al., 2004) with Venus as the host planet. The Earth has one temporary co-orbital object, ($3753 \\, Cruithne$, Namouni et al. 1999), and one alternating horseshoe-QS object ($2002 AA29$, Connors et al. 2002). The co-orbital asteroidal population in the inner Solar System was studied in Brasser et al. (2004) by numerical integrations. All QS orbits appear to be temporary, escaping in timescales of the order of $10^2-10^4$ years  Wiegert et al. (2000) numerically investigated the stability of QS orbits around the giant planets of the Solar System. Although no stable solutions were found for Jupiter and Saturn, some initial conditions around Uranus and Neptune lead to QS orbits that survive for timescales of the order of $10^9$ yr. It thus appears that a primordial population of such objects may still exist in the Solar System. Kortenkamp (2005) used N-body simulations to model the combined effects of solar nebula gas drag and gravitational scattering of planetesimals by a protoplanet. He showed that a significant fraction of scattered planetesimals could become trapped into QS trajectories. It then seems plausible that this trapped-to-captured transition may be important not only for the origin of captured satellites but also for continued growth of protoplanets.  At variance with these results, in the case of the general (non-restricted) three-body problem, although equilateral solutions and horseshoe orbits are well known, quasi-satellite configurations have only been studied very recently. Hadjidemetriou et al. (2009) performed a detailed study of periodic orbits in the 1/1 MMR for fictitious planetary systems with different mass ratios. They found that stable QS solutions occur for $\\sigma = \\Delta \\lambda = \\lambda_2 - \\lambda_1 = 0$ and $\\Delta\\varpi = \\varpi_2 - \\varpi_1 = 180^{\\circ}$, where the subscripts identify each planet. Unstable trajectories were found at $\\sigma = 180^{\\circ}, \\Delta\\varpi = 0$. Although at present there are no confirmed cases of exoplanets in quasi-satellite configurations, Go\\'zdziewski \\& Konacki (2006) found that the radial velocity curves of the HD82943 and HD128311 planets could correspond to co-orbital motion in highly inclined orbits. Numerical simulations of both systems show QS trajectories, instead of Trojan orbits as initially believed.  In the present work we aim to revisit the 1/1 mean-motion resonance (MMR) in the planar planetary three-body problem, trying to identify possible domains of stable solutions and their location in the phase space. Section \\ref{regular} presents several dynamical maps constructed from numerical simulations for different initial conditions. These maps allow us to identify stable fixed points and periodic orbits, as well as the domains of regular motions. In Section \\ref{model} we develop a semi-analytical model for co-orbital planets, which is then applied in Section \\ref{op} to calculate the families of stable periodic orbits. In the same section we also present a brief study of the effects of an adiabatically slow mass variation in one of the planetary bodies. Finally, conclusions close the paper in Section \\ref{conclusions}.  \\begin{figure} \\centerline{\\includegraphics*[width=18pc]{fig1.eps}} \\caption{Results of numerical integrations of initial conditions in a grid in the $(\\sigma, \\Delta \\varpi)$ plane. Planetary masses were taken equal to $m_1=m_2= m_{\\rm Jup}$, and initial semimajor axes equal to $a_1 = a_2 = 1$ AU. Regions of regular motion are shown in white, while the dashed regions correspond to chaotic and unstable trajectories.} \\label{fig1} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics*[width=18pc]{fig2.eps}} \\caption{Semi-amplitude maps. The Left (Right) column shows the amplitude variation for $\\sigma$ ($\\Delta\\varpi$) in gray scale. Light domains correspond to near zero amplitude families, darker regions indicate oscillation amplitudes up to $90^{\\circ}$, and dashed regions correspond to unstable orbits. Initial values of eccentricities are indicated in each panel. Color scale is indicated at bottom and ACR solutions are marked on the right panels.} \\label{fig2} \\end{figure} ",
        "conclusions": "We studied the stability regions and families of periodic orbits of two-planet systems in the vicinity of a 1/1 mean-motion resonance (i.e. co-orbital configuration). We considered different ratios of planetary masses and orbital eccentricities, also we assumed that both planets share the same orbital plane (coplanar motion).  As result we identified two separate regions of stability, each with two distinct modes of motion:  \\begin{itemize} \\item \\textbf{Quasi-Satellite region:} Originally identified by Hadjidemetriou et al. (2009) for the planetary problem, QS orbits correspond to oscillations around an ACR located at $(\\sigma,\\Delta\\varpi) = (0,180^{\\circ})$. Although not present for quasi-circular trajectories, they fill a considerable portion of the phase space in the case of moderate to high eccentricities.  We also found a new regime, associated to stable orbits displaying oscillations around $(\\sigma,\\Delta\\varpi) = (0,0)$, even though this point is unstable and corresponds to a collision between the two planets.  \\item \\textbf{Lagrangian region:} Apart from the previous symmetric solutions, we also found two distinct types of asymmetric ACR orbits in which both $\\sigma$ and $\\Delta\\varpi$ oscillate around values different from $0$ or $180^\\circ$. The first is the classical equilateral Lagrangian solution associated to local maxima of the averaged Hamiltonian function. Independently of the mass ratio $m_2/m_1$ and their eccentricities, these solutions are always located at $(\\sigma,\\Delta\\varpi) = (\\pm 60^\\circ,\\pm 60^\\circ)$. However, the size of the stable domain decreases rapidly for increasing eccentricities, being practically undetectable for $e_i > 0.7$.  The second type of asymmetric ACR correspond to local minima of the averaged Hamiltonian function. We have dubbed them Anti-Lagrangian solutions ($AL_4$ and $AL_5$). For low eccentricities, they are located at $(\\sigma,\\Delta\\varpi) = (\\pm 60^\\circ,\\mp 120^\\circ)$. Each is connected to the classical $L_4$ and $L_5$ solution through the $\\sigma$-family of periodic orbits in the averaged system. Contrary to the classical equilateral Lagrangian solution, their location in the plane $(\\sigma,\\Delta\\varpi)$ varies with the planetary mass ratio and eccentricities. Although their stability domain also shrinks for increasing values of $e_i$ they do so at a slower rate than the classical Lagrangian solutions, and are still appreciable for eccentricities as high as $\\sim 0.7$. \\end{itemize}  Finally, we also applied an ad-hoc adiabatically slow mass variation to one of the planetary bodies, and analyzed its effect on the $AL_4$ configuration. We found that the resonant co-orbital solution was preserved, with practically no change in the equilibrium values of the angles. The eccentricities, however, varied with the larger planet approaching a quasi-circular orbit as the smaller planet had its eccentricity increased. These solution still exist in the limit of the restricted three-body problem (i.e. $m_2 \\rightarrow 0$), although both types of asymmetric solutions ($L_4$ and $AL_4$) have different geometries. While the first are true stationary solutions in the unaveraged system, the latter are periodic orbits around the classical equilateral Lagrangian points."
    },
    "1004/1004.2209_arXiv.txt": {
        "abstract": "{Current generation millimeter wavelength detectors suffer from scaling limits imposed by complex cryogenic readout electronics.  These instruments typically employ multiplexing ratios well below a hundred. To achieve multiplexing ratios greater than a thousand, it is imperative to investigate technologies that intrinsically incorporate strong multiplexing.  One possible solution is the kinetic inductance detector (KID).  To assess the potential of this nascent technology, a prototype instrument optimized for the 2 mm atmospheric window was constructed.  Known as the N\\'{e}el IRAM KID Array (NIKA), it has recently been tested at the Institute for Millimetric Radio Astronomy (IRAM) 30-meter telescope at Pico Veleta, Spain.} {There were four principle research objectives:  to determine the practicality of developing a giant array instrument based on KIDs, to measure current in-situ pixel sensitivities, to identify limiting noise sources, and to image both calibration and scientifically-relevant astronomical sources.  } {The detectors consisted of arrays of high-quality superconducting resonators electromagnetically coupled to a transmission line and operated at $\\sim$100 mK.  The impedance of the resonators was modulated by incident radiation; two separate arrays were tested to evaluate the efficiency of two unique optical-coupling strategies.  The first array consisted of lumped element kinetic inductance detectors (LEKIDs), which have a fully planar design properly shaped to enable direct absorbtion.  The second array consisted of antenna-coupled KIDs with individual sapphire microlenses aligned with planar slot antennas.  Both detectors utilized a single transmission line along with suitable room-temperature digital electronics for continuous readout.} {NIKA was successfully tested in October 2009, performing in line with expectations.    The measurement resulted in the imaging of a number of sources, including planets, quasars, and galaxies.  The images for Mars, radio star MWC349, quasar 3C345, and galaxy M87 are presented.  From these results, the optical NEP was calculated to be around $1 \\times 10^{-15}$ W$/$Hz$^{1/2}$.  A factor of 10 improvement is expected to be readily feasible by improvements in the detector materials and reduction of performance-degrading spurious radiation. } {} ",
        "introduction": "Millimeter and sub-millimeter wavelength observations are at the forefront of cosmology and astrophysics.  They have proven to be an indispensable tool for understanding the early stages of star and galaxy formation.  Furthermore, continuum measurements provide a direct probe of the anisotropy of the cosmic microwave background, the Sunyaev-Zel'dovich effect in clusters of galaxies and the dust emission in the Galaxy and external galaxies.  Low-temperature bolometers have historically been utilized as detectors for these measurements. For this reason, bolometers have enjoyed decades of technical improvements and single pixels are now reaching fundamental limits.  Once intrinsic device limits have been achieved, greater sensitivity can only be realized by increasing the focal plane area and pixel count.  Current bolometer arrays such as MAMBO, LABOCA, SABOCA, APEX-SZ, SPT, SCUBA-2, Herschel PACS, and SPIRE now typically employ up to a few thousand pixels.  This scaling has resulted in increasingly complex readout electronics that must implement some form of multiplexing to reduce cryogenic wire-counts.  For example, time-domain multiplexing, or rapid switching, can be achieved using quantum point contact high-mobility transistors (\\cite{Benoit:940}) or MOSFET (\\cite{Reveret:32}).  In frequency-domain multiplexing, instead of rapid switching between separate wires, the bandwidth of a single wire is subdivided between multiple pixels.  This technique has been previously demonstrated by integrating superconducing quantum interference devices (SQUIDs) with a type of bolometer known as the transition edge sensor (TES) (\\cite{yoon:371}, \\cite{irwin:2107}, \\cite{irwin:63}).  However, to realize the very large multiplexing ratios required for many-kilopixel arrays, it is necessary to develop detectors intrinsically adapted to strong frequency-domain multiplexing.  One potential device that achieves high sensitivity and is fundamentally compatible with frequency-domain multiplexing is the kinetic inductance detector (KID).  Based on high-quality, low-volume superconducting microwave resonators, KIDs were first proposed less than ten years ago for millimeter radiation detection (\\cite{Day2003}).  Benefiting from intense research in quantum computing and fundamental physics (\\cite{Hofheinz2009}; \\cite{Grolosvsky6462}), development proceeded quickly and resulted in the first telescope measurement employing KIDs at the Caltech Submillimeter Observatory (\\cite{Schlaerth2008}).  In view of a future large instrument and in order to test the feasibility of a receiver array based on KIDs, we have recently constructed a prototype instrument, known as the N\\'{e}el IRAM KID Array (NIKA).  NIKA was recently tested at the 30-meter Institute for Millimeteric Radio Astronomy (IRAM) telescope at Pico Veleta, Spain.  Two complementary technologies were tested during this run: a 30-pixel lumped element kinetic inductance detector (LEKID) array (\\cite{doyle:156}) and a 42 pixel antenna-coupled KID array.  Here we present an overview of the operating principles of the detector arrays, a summary of the principal instrument design elements, and the results from the first measurement run. ",
        "conclusions": ""
    },
    "1004/1004.4891.txt": {
        "abstract": "{The CoRoT mission is in its third year of observation and the data from the second long run in the galactic centre direction are being analysed. The solar-like oscillating stars that have been observed up to now have given some interesting results, specially concerning the amplitudes that are lower than predicted. We present here the results from the analysis of the star HD~170987.} {The goal of this research work is to characterise the global parameters of HD~170987. We look for global seismic parameters such as the mean large separation, maximum amplitude of the modes, and surface rotation because the signal-to-noise ratio in the observations do not allow us to measure individual modes. We also want to retrieve the stellar parameters of the star and its chemical composition.} {We have studied the chemical composition of the star using ground-based observations performed with the NARVAL spectrograph. We have used several methods to calculate the global parameters from the acoustic oscillations based on CoRoT data. The light curve of the star has been interpolated using inpainting algorithms to reduce the effect of data gaps.} {We find power excess related to p modes in the range [400 - 1200]\\,$\\mu$Hz with a mean large separation of 55.2\\,$\\pm$\\,0.8\\,$\\mu$Hz with a probability above 95\\,\\% that increases to 55.9 $\\pm$\\,0.2\\,$\\mu$Hz in a higher frequency range [500 - 1250] \\,$\\mu$Hz and a rejection level of 1\\,$\\%$. A hint of the variation of this quantity with frequency is also found. The rotation period of the star is estimated to be around 4.3 days with an inclination axis of $i$\\,=\\,$50\\degr \\;^{+20}_{-13}$. We measure a bolometric amplitude per radial mode in a range [2.4 - 2.9]~ppm around 1000~$\\mu$Hz. Finally, using a grid of models, we estimate the stellar mass, M\\,=\\,1.43~$\\pm$\\,0.05~$M_\\odot$, the radius, R\\,=\\,1.96~$\\pm$\\,0.046~$R_\\odot$, and the age $\\sim$2.4~Gyr. } {} ",
        "introduction": "During the present decade the number of confirmed solar-like pulsators -- those with acoustic modes excited by turbulent motions in the near-surface convection \\citep[e.g.][and reference therein]{2004SoPh..220..137C} -- has increased enormously thanks, first, to the growing number of ground-based observing campaigns \\citep[e.g.][]{2007CoAst.150..106B,2008ApJ...687.1180A}, and second, to the high-precision photometry measurements provided by space instrumentation such as WIRE \\citep[Wide-Field Infrared Explorer, e.g.][]{2007A&A...461..619B}, MOST \\citep[Microvariability and Oscillations of STars,][]{2003PASP..115.1023W} and CoRoT \\citep{2006ESASP.624E..34B}. The latter has been providing data with an unprecedented quality both in terms of photometric precision and in terms of uninterrupted observation lengths. CoRoT has already observed several main-sequence solar-like pulsators \\citep{2008Sci...322..558M} while it has enabled to resolve the individual modes of the oscillations spectra of several F stars  \\citep{2008A&A...488..705A,2009A&A...506...51B,2009A&A...506...41G,2009A&A...506...33M} and G stars \\citep[][Ballot et al. in prep.]{2010arXiv1003.4368D}. At least, it allowed to derive a large spacing for the faintest targets \\citep{2009A&A...506...41G,2009A&A...506...33M}. The measurement of these seismic parameters already offer a valuable tool for accurate determinations of radii \\citep[e.g.][]{stello09} and ages \\citep[e.g.][]{2008arXiv0810.2440C} of stars, which are specially interesting for better understand stellar evolution as well as to characterise stars hosting planets. All of these CoRoT observations are the starting point for a better understanding of the structure \\citep{2009A&A...506..175P,2009Ap&SS.tmp..241D} and the surface dynamics \\citep{2009A&A...506..167L,2009arXiv0910.4027S,2009arXiv0910.4037S} of this class of stars.  %%% ICI, ou dans la conclusion ?? The Kepler mission, successfully launched  in March 7, 2009 \\citep{2009IAUS..253..289B}, will also contribute to this field by observing stars on very long runs (4 years). The quality of the first data  on stars showing solar-like oscillations \\citep{2010ApJ...713L.176B, 2010ApJ...713L.169C, 2010ApJ...713L.187H, 2010ApJ...713L.182S} promises the asteroseismology a bright future on the study of stellar interiors and dynamical processes \\citep[][]{2009arXiv0911.4629C,2009arXiv0912.0817S,2009A&A...506..811M}.  In this paper we present results about a star recently observed by CoRoT, HD~170987 (or  HIP 90851). This target is a well-known double star where components are separated by $0.7\\arcsec$ \\citep[e.g.][]{2002yCat.1274....0D}. The main star is a F5 dwarf star with a magnitude $m_V$ ranging from 7.4 to 7.7 in the literature, while the second component has a magnitude around 8.5.   \\begin{figure*}[!ht] \\begin{center} \\includegraphics[width=15cm, height=6cm, trim=0cm 1.5cm 0cm 9.5cm]{Figures/HDNAR170987a_vwaview.pdf}   \\caption{Comparison of the observed (black) and computed (blue) spectrum of HD~170987. The four solid vertical lines mark the lines used in the spectral analysis while the vertical dotted lines show other well-known lines.} \\label{fig:spec} \\end{center} \\end{figure*}    We start by reporting the latest spectroscopic results observed by the NARVAL spectrograph (in Sect.~2), which shows that this star is very similar to Procyon \\citep[e.g.][]{allende02}. Then we describe the observations done with CoRoT during 149 days and the interpolation done in the data gaps of the light curve (Sect.~3). In Sect.~4 we infer the surface-rotation period of the star from the detailed analysis of the low-frequency region of the power spectrum and then, we obtain the global properties of the acoustic modes and of the star, respectively in Sect.~5 and Sect.~6. We finish in Sect.~7 with a discussion of the results and in Sect.~8 with the conclusions of the paper. ",
        "conclusions": "We have analysed the 149~days of the light curve of the CoRoT target HD~170987, allowing us to determine its global parameters. We find that the rotation period of the star is around 4.3 days and the v $\\sin i$ of 19\\,km/s, leading to an inclination angle $i$ of $\\sim$\\,50$\\degr \\;^{+20}_{-13}$. By fitting the background, we obtain that the time-scale for the granulation is around 383~$\\pm$\\,28~s either assuming that we have some contribution of faculae or of modes.  We confirm that power excess between 400 and 1200~$\\mu$Hz is due to acoustic modes, which are characterised by a mean large separation $\\langle \\Delta\\nu \\rangle$~=~55.2~$\\pm$\\,0.8~$\\mu$Hz with more than 95\\,\\% of probability and a posteriori probability of 68\\,\\%. Using the EACF method, we find  $\\langle \\Delta\\nu \\rangle$~=~55.9~$\\pm$\\,0.2\\,$\\mu$Hz in a higher frequency range [500 - 1250] $\\mu$Hz with a 1\\,\\% rejection. However, we recognize that all these confidence levels are dependent on our assumption that our spectrum is dominated by white noise statistics. It may therefore be that we should be less confident as there could be other sources of noise embedded in the signal. However, at such a level, the detection seems to be unambiguous. Moreover, this value is in agreement with the expected one derived from the stellar fundamental parameters, which gives 52.7\\,$\\pm$\\,4\\,$\\mu$Hz, the large error bar being due to the large uncertainty on the parallax. Now, thanks to the detection of this pattern, we can be confident that the bump observed around 1000~$\\mu$Hz is very likely to be related to acoustic modes.  We characterise the power excess with a maximum amplitude of 2.7~$\\pm$\\,0.6~ppm at $\\sim$1070~$\\mu$Hz, where the uncertainty on the amplitude has been obtained by taking into account the scatter of the smoothed power spectrum about the background fit  outside the oscillation range.%{\\bf TO BE VERIFIED/CHANGED DEPENDING ON DANIEL\"S VALUES}  Because of the low signal-to-noise ratio, we cannot fit the acoustic modes individually. However, from the global seismic parameters found and the results form the NARVAL spectrohraph, we have estimated that HD~170987 has a mass, M~=~1.43~$\\pm$\\,0.05~$M_\\odot$, a radius, R~=~1.96~$\\pm$\\,0.046~$R_\\odot$, and an age $\\sim$2.4~Gyr.  Further studies are needed to have a better comprehension of the propagation of waves in stellar atmosphere of stars located at this position in the HR diagram. It would also enable us to understand why the oscillation-amplitudes measured in intensity fluctuations are so small.  %TBD"
    },
    "1004/1004.1516_arXiv.txt": {
        "abstract": "\\noindent  It is known that most of the craters on the surface of the Moon were created by the collision of minor bodies of the Solar System. Main Belt Asteroids, which can approach the terrestrial planets as a consequence of different types of resonance, are actually the main responsible for this phenomenon.  Our aim is to investigate the impact distributions on the lunar surface that low-energy dynamics can provide. As a first approximation, we exploit the hyberbolic invariant manifolds associated with the central invariant manifold around the equilibrium point $L_2$ of the Earth -- Moon system within the framework of the Circular Restricted Three -- Body Problem. Taking transit trajectories at several energy levels, we look for orbits intersecting the surface of the Moon and we attempt to define a relationship between longitude and latitude of arrival and lunar craters density. Then, we add the gravitational effect of the Sun by considering the  Bicircular Restricted Four -- Body Problem.  In the former case, as main outcome we observe a more relevant bombardment at the apex of the lunar surface, and a percentage of impact which is almost constant and whose value depends on the Earth -- Moon distance $d_{EM}$ assumed. In the latter, it seems that the Earth -- Moon and Earth -- Moon -- Sun relative distances and the initial phase of the Sun $\\theta_0$ play a crucial role on the impact distribution. The leading side focusing becomes more and more evident as $d_{EM}$ decreases and there seem to exist values of $\\theta_0$ more favorable to produce impacts with the Moon. Moreover, the presence of the Sun make some trajectories to collide with the Earth. The corresponding percentage floats between 1 and 5 $\\%$.  As further exploration, we assume an uniform density of impact on the lunar surface, looking for the regions in the Earth -- Moon neighbourhood these colliding trajectories have to come from. It turns out that low-energy ejecta originated from high-energy impacts are also responsible of the phenomenon we are considering. ",
        "introduction": "\\label{sec:1}  The surface of the Moon is constellated by impact craters of various sizes, mainly generated from the collision of objects coming from the Main Asteroid Belt. Indeed, the Inner Solar System can be reached by such minor bodies as a consequence of different types of resonance \\cite{BMJPLMM}. The intense lunar bombardment took place between 3.8 and 4 Gy ago, being at the present day the meteroidal flux about $10^3$ lower \\cite{H}.  The cratering process is interesting for several branches of science. First of all, by comparing densities of craters on different surfaces it is possible to derive the relative age of the corresponding terrains (see, for instance, \\cite{NIH,SR,MMCMM}). Roughly speaking, the greater the density the older the surface. Also, the geological chronology of the terrestrial planets is now becoming more and more accurate thanks to the space missions that provide radiometric age estimates for different regions. This is especially true if we take as reference case the Moon, for which a great amount of data is now available. From this kind of analysis, new insight on the Solar System evolution can be obtained. As further aspect, the flux of impacts offers information on the Solar System minor bodies population.  The main problem in all these studies resides on the fact that the crater formation is a phenomenon not fully understood yet. There does not exist a predictive, quantitative model of crater formation, that is, a reliable methodology that can be applied to all situations. The size of the crater that forms at the end of the excavation stage depends on the asteroid's size, speed and composition, on the collision angle, on the material and structure of the surface in which the crater forms and on the surface gravity of the target \\cite{M}. The problem in the determination of the crater's dimension concerns with the poorness of the experimental or observational data. This difficulty is usually overcome by extrapolating beyond experimental knowledge through scaling laws.  This work regards the paths that impacting asteroids might have followed. In particular, we will deal with low-energy trajectories, first derived in the Circular Restricted Three -- Body Problem (CR3BP) framework applied to the Earth -- Moon system and then analysed also accounting for the Sun gravitational attraction by means of the Bicircular Restricted Four -- Body Problem (BR4BP). We assume the minor bodies to have already left the Main Asteroid Belt and we consider as main entrance to the Earth -- Moon neighbourhood the stable invariant manifold associated with the central invariant manifold corresponding to the $L_2$ equilibrium point.  We will look for the distribution of impacts that such orbits can create, paying attention to the fact that the Moon is locked in a $1:1$ spin--orbit resonance. In particular, we wonder if, for the range of energy under consideration, the Moon acts as a shield for the Earth or if the greatest concentration of collisions still takes place on the leading side of the surface, as other authors have pointed out with different approaches. See, for example, \\cite{HN,MF,LW}.  In a second step, from a backward integration, we attempt at discovering any other gate that can lead to a lunar impact within low-energy regimes.  We recall that due to the small values of energy we consider, the impacts obtained can yield to at most $40$ km in diameter craters. This value has been computed by applying the scaling laws of Melosh \\cite{M} to the Moon's surface with an impact velocity corresponding to the escape lunar velocity (about $2.4$ km$/$s). ",
        "conclusions": "\\label{sec:7}  The main purpose of this paper is to establish a relationship between low-energy trajectories in the Earth -- Moon system and lunar impact craters. This is actually a quite wide and challenging topic, which involves knowledge related to mathematics, astronomy and geology.  As primary goal, we define some tools which are effective in the determination of the dynamics pushing a massless particle under low-energy regimes. We exploit invariant objects within the Circular Restricted Three -- Body Problem approximation, in particular transit trajectories lying inside the stable invariant manifold associated with the central invariant manifold of the $L_2$ equilibrium point. We implement a method that allows to reproduce the behaviour associated with the unstable component of any central orbit and does not need to distinguish between them. This fits with our investigation, because we are interested in minor bodies collisions that take advantage of the channels represented by the whole hyperbolic manifolds. To this purpose, we adopt well-known procedures to compute periodic Lyapunov orbits together with their hyperbolic invariant manifolds.  With this approach, we perform extensive numerical simulations to determine both the lunar region of heavier impact and the sources of a potential uniform craters distribution. We also look for the influence of the Sun on these paths, by means of the Bicircular Restricted Four -- Body Problem. Several outcomes can be highlighted, even if they have to be seen as patterns that require a more robust proof: further calculations with different dynamical and astronomical models are in progress. The investigation carried out is promising from many points of view, as it indicates future developments that are worth to be considered.  Without the effect of the Sun, we get a confirmation that the neighbourhood of the apex of the surface of the Moon is the region where most collisions take place. We remark that the impact trajectories simulated reach the surface of the Moon with the lowest possible velocity: this point does not corrupt the apex concentration that other authors discovered without this restriction. The total time span considered (60 years) is sufficient to describe the general behaviour of the massless particles and no distinctions among different values for the Earth -- Moon distance were observed.  However, the gravitational force exerted by the Sun seems to blur the above phenomenon. Changing the ratio between the Earth -- Moon -- Sun distance and the Earth -- Moon one, we notice different patterns. From our computations, it turns out that in more ancient epochs the low-energy lunar impacts were focused on the Moon leading side, but this is not true going further in time. Moreover, we get evidence that the position of the Sun with respect to the Earth -- Moon barycenter affects the distribution of lunar impacts. These are the first aspects we plan to study with more detail. For instance, it would be deserving to integrate the BR4BP equations of motion for a longer interval of time and for a greater number of values of the initial phase of the Sun to understand the real nature of these numerical observations.  On the other hand, we realize that small craters can also be generated by the impact of dust arising from more energetic collisions than the ones investigated here. Such phenomenon comes from the existence of periodic orbits that cross the surface of the Moon, that is, double collision orbits. In the future, we would like to see how they are transformed by the perturbation of the Sun.  A natural step would be to add the gravitational attractions of other planets to see their consequences on the orbits simulated. This will be done by means of a Restricted n -- Body Problem, using position and velocity of the primaries given by the JPL ephemerides (for instance the DE405 ones) and taking several initial epochs to compare the whole outcome.  Moreover, we would like to link our methodology with real observational data, concerning either the existing lunar craters and the orbital parameters at a certain epoch of a given set of Near Earth Objects. This information would affect especially the way we generate the initial conditions corresponding to transit orbits.  Finally, to apply the same kind of analysis to the terrestrial planets would be of large interest. Starting from the CR3BP approximation, we mean to study the density of impact provided by Sun -- planet low-enery orbits and then to add further gravitational effects, trying to figure out the orbital elements and also the regions in the phase space which more likely lead to collision."
    },
    "1004/1004.4105_arXiv.txt": {
        "abstract": "We use $N$-body simulations to find the effect of neutrino masses on halo properties, and investigate how the density profiles of both the neutrino and the dark matter components change as a function of the neutrino mass. We compare our neutrino density profiles with results from the $N$-one-body method and find good agreement. We also show and explain why the Tremaine-Gunn bound for the neutrinos is not saturated. Finally we study how the halo mass function changes as a function of the neutrino mass and compare our results with the Sheth-Tormen semi-analytic formulae. Our results are important for surveys which aim at probing cosmological parameters using clusters, as well as future experiments aiming at measuring the cosmic neutrino background directly. ",
        "introduction": "Massive neutrinos are known to have a significant effect on cosmic structure formation~\\cite{Bond:1980ha,Doroshkevich:1980zs}. In the early universe they contribute to the relativistic energy density and influence the transition from radiation to matter domination. At late times they contribute to the dark matter density, and therefore also to cosmic structure formation. However, as opposed to Cold Dark Matter (CDM), they do not contribute to structure formation on physical scales smaller than the free-streaming scale, roughly equal to the distance traversed before the neutrinos become non-relativistic. This suppression of small-scale structure leaves a very distinct imprint on large-scale structure observables such as the matter power spectrum, which can in turn be used to probe neutrino physics.  Many studies have by now been devoted to this topic, most of which focussing on constraining the neutrino mass, $m_\\nu$. At present an upper limit on the neutrino mass can be derived from observations of the Cosmic Microwave Background (CMB) anisotropies alone or in conjunction with various large-scale structure data sets, such as the SDSS-DR7 LRG catalog, and falls in the range $\\sum m_\\nu \\lwig 0.4-0.7$ eV, depending both on the complexity of the model space and the combination of data sets used (e.g., \\cite{Komatsu:2010fb,Reid:2009nq,Hamann:2010pw}). In the future the sensitivity of large-scale structure observations to the neutrino mass will increase significantly. For example it has been estimated that the combination of CMB data from Planck and a weak lensing survey from the LSST will push the 1$\\sigma$ sensitivity to better than 0.05~eV, close to the minimum $\\sum m_\\nu$ allowed by oscillation data~\\cite{Hannestad:2006as}.  While most neutrino mass constraints at present have been derived using large-scale structure correlation functions (or power spectra), there are other observables that are potentially just as interesting. One prime example is cluster number counts which are in principle very sensitive to the neutrino mass \\cite{Wang:2005vr}. However, in order to fully utilise such data it is necessary to have accurate theoretical predictions, which so far do not exist for $\\Lambda$CDM models extended with massive neutrinos (see, however, \\cite{Kofman:1995ds} for an early calculation based on the old mixed dark matter scenario). In the present paper we calculate the halo mass function in $\\Lambda$CDM cosmologies with massive neutrinos included for a variety of neutrino masses. However, before proceeding to this and a discussion of other observables related to halo properties, let us briefly review how neutrinos affect structure formation in the linear regime.   \\subsection{The effect of neutrinos} The effect of neutrinos on structure formation in linear theory has been studied numerous times in the literature (see, e.g., \\cite{Lesgourgues:2006nd}). In general the amount of fluctuations at a given wavenumber $k$ is represented by the power spectrum, $P(k) = |\\delta_k|^2$, which can be split in the following form \\begin{equation} P(k,z)=D(z) T^2(k,z) P_0(k), \\end{equation} where $D(z)$ is a scale independent growth factor and $P_0$ is the initial power spectrum. $T(k)$ is the transfer function (TF) which is both time and scale dependent in general. The effect of massive neutrinos is embedded entirely in the TF and is separated into two regimes. On scales much larger than the free-streaming scale, \\begin{equation} k_{\\rm FS} \\sim 0.8 \\frac{m_\\nu}{\\rm eV} \\, h \\, {\\rm Mpc}^{-1}, \\end{equation} where $m_\\nu$ is the one-particle neutrino mass, neutrinos behave essentially like CDM, while on smaller scales they suppress structure formation. Very na\\\"{\\i}vely one might expect the suppression arising from replacing a fraction of the CDM component with neutrinos to be of order $\\Delta P/P \\sim -\\Omega_\\nu/\\Omega_m$, where $\\Omega_{m} = \\Omega_c+\\Omega_b+\\Omega_\\nu$ is the total matter density, because neutrinos do not cluster. However, this grossly underestimates the true effect because massive neutrinos also influence the background expansion around the time of matter-radiation equality. The final result in linear theory is that the suppression is approximately given by \\begin{equation} \\frac{\\Delta P}{P} \\sim - 8 \\frac{\\Omega_\\nu}{\\Omega_m}. \\end{equation} This shows that most of the effect actually comes from the modification to the background, i.e., sub-eV to eV scale neutrinos lead to a longer radiation era. This effect is also much larger than the effect of  replacing a fraction $\\Delta \\Omega_m$ of the CDM energy density with $\\Lambda$. In this case the change in the matter power spectrum on small scales (with the large-scale normalisation held constant, i.e., ignoring the effects of $\\Delta \\Omega_m$ on the growth factor) is very approximately given by \\begin{equation} \\frac{\\Delta P}{P} \\sim \\left(\\frac{\\Omega_{m}^{'}}{\\Omega_{m}}\\right)^{7/2} \\sim - 3.5 (1-\\Omega_{m}^{'}/\\Omega_{m}), \\label{eq:cdm} \\end{equation} where $\\Omega_{m}^{'} = \\Omega_{m}+\\Delta \\Omega_m$, for small changes in $\\Delta \\Omega_m \\ll \\Omega_m$. The effect here is approximately two times smaller than that due to assigning $\\Delta \\Omega_m$ to massive neutrinos.  Since neutrinos have such a strong effect on the power spectrum even in linear theory it is natural to expect a similarly strong effect in the non-linear regime. This was tested in detail for the power spectrum in a number of papers \\cite{Brandbyge1,Brandbyge2,Brandbyge3}, and a significant enhancement in the power spectrum suppression was indeed found: The maximum suppression is increased from $-8 \\Omega_\\nu/\\Omega_{m}$ to approximately $-9.8 \\Omega_\\nu/\\Omega_{m}$,\\footnote{This finding was confirmed in the very recent paper \\cite{Viel:2010bn}.} with a pronounced feature at $k \\sim 0.7 \\, h \\, {\\rm Mpc}^{-1}$.  Another issue which has so far not been addressed with precision $N$-body simulations is how the presence of massive neutrinos affect halo formation. Here we study how CDM halo properties are altered by the presence of massive neutrinos, and we also present detailed results for the corresponding neutrino halos. This last point is important for example for understanding the prospects for a direct experimental detection of the cosmic relic neutrino background.  The paper is organised as follows: In Section \\ref{sec:numericalsetup} we present the numerical setup required for the analysis. In Section \\ref{sec:halostructure} we present results on halo profiles for both the neutrino and matter components. In Section \\ref{sec:massfunction} we discuss how the halo mass function is altered in models with massive neutrinos, and finally Section \\ref{sec:conclusion} contains our conclusions. ",
        "conclusions": "\\label{sec:conclusion}  We have performed a detailed study of halo properties in $\\Lambda$CDM cosmologies with massive neutrinos included. An important goal was to study the neutrino density profiles in dark matter halos. To this end we employed detailed $N$-body simulations across a wide range of scales to test halo masses from Milky Way size ($10^{12} \\, {\\rm M}_\\odot$) to large clusters ($10^{15} \\, {\\rm M}_\\odot$), as well as the $N$-one-body method developed to solve the neutrino Boltzmann equation approximately around existing CDM halos. In general we found good agreement between the full $N$-body and the $N$-one-body results. The difference between the $N$-body and $N$-one-body methods arise from the fact that the latter assumes the CDM halo to be monolithic and at all times describable in terms of a NFW profile, i.e., it does not take into account halo substructure and larger merger events. It also assumes an analytic evolution of the concentration parameter. We also discussed in some detail how the density profiles of neutrino halos can be understood in terms of the Tremaine-Gunn bound, i.e., the bound coming from the fact that a coarse-grained distribution can never attain values exceeding the maximum of the original fine-grained distribution.  For smaller halo masses, the neutrino profiles in isolated halos are in excellent agreement with the prediction from the $N$-one-body method. This result is not too surprising since this is exactly the case where the infall on an existing spherical NFW halo is most realistic. However, many smaller mass halos are embedded in larger cluster halos and for the smaller neutrino masses the local neutrino profile in such a halo is dominated by the background of neutrinos bound in the much larger cluster halos.  In terms of the local neutrino density enhancement, which is relevant for possible future attempts at direct C$\\nu$B detection, a Milky Way-size galaxy halo is too small to have a significant overdensity, even when taking a possible cluster background into account.  We also briefly studied how neutrinos impact on the density profiles of the CDM halos. While neutrinos contribute very little to the total density in the halo, the presence of massive neutrinos in the model leads to slightly later formation of halos with a given mass and consequently to generally lower concentration parameters, $c$.  Finally, we calculated halo mass functions for $\\Lambda$CDM models with massive neutrinos. Since large cluster surveys will become available in the coming years, the halo mass function is an important cosmological observable. As expected, we find a very strong suppression of halo formation with increasing neutrino mass. As noted in previous analytic or semi-analytic studies the suppression is particularly marked for massive halos because the suppression in linear theory power from massive neutrinos shifts the maximum cluster mass down, i.e.\\ the scale beyond which the halo mass function is exponentially suppressed.   We then compared the halo mass functions from simulations with halo mass functions calculated using the semi-analytic method developed by Sheth and Tormen. If used na\\\"{\\i}vely, i.e.\\ just processing the linear theory power spectrum without any adjustment to the method, the agreement is poor. However, it is easy to see that the disagreement arises because the ST method implicitly assumes that all matter clusters in the same way (the value of $\\Omega_m$ used is $\\Omega_c+\\Omega_b+\\Omega_\\nu$). However, even large clusters bind relatively few neutrinos and for all halos it is true that neutrinos make a negligible contribution to the halo mass. If the ST formalism is corrected for this by using $\\Omega_m=\\Omega_c+\\Omega_b$, i.e.\\ taking into account only the clustering species (but of course using the correct initial power spectrum and the correct background evolution), the agreement between the modified ST and the $N$-body results is remarkable. On all measurable scales it is better than 2-3\\%. This is important for analysing future cluster surveys because it means that existing semi-analytic methods can be used instead of having to perform time consuming simulations for all neutrino masses.  Alternatively, neglecting the neutrino perturbations in the $N$-body simulation will also be a very accurate approximation for $\\sum m_\\nu \\lesssim 0.5 \\, {\\rm eV}$ as long as only matter halo properties are considered. This approximation is not valid for a precise calculation of the matter power spectrum \\cite{Brandbyge1}. In general the accuracy of the approximation is determined by contrasting the neutrino free-streaming length with the physical extent of the scales simulated: Considering halo properties and realistic neutrino masses, this approximation is very good."
    },
    "1004/1004.5580.txt": {
        "abstract": "Cosmological Gravitational Waves (GWs) are usually associated with the transverse-traceless part of the metric perturbations in the context of the theory of cosmological perturbations. These modes are just the usual polarizations `+' and `$\\times$' which appear in the general relativity theory. However, in the majority of the alternative theories of gravity, GWs can present more than these two polarization states. In this context, the Newman-Penrose formalism is particularly suitable for evaluating the number of non-null GW modes. In the present work we intend to take into account these extra polarization states for cosmological GWs in alternative theories of gravity. As an application, we derive the dynamical equations for cosmological GWs for two specific theories, namely, a general scalar-tensor theory which presents four polarization states and a massive bimetric theory which is in the most general case with six polarization states for GWs. However, the mathematical tool presented here is quite general, so it can be used to study cosmological perturbations in all metric theories of gravity. ",
        "introduction": "The future detection of gravitational waves (GWs) of cosmological origin will strongly constrain the possible inflationary scenarios which have been proposed in the last decades. Also, GWs will be useful to distinguish between the standard inflationary model and the alternative early Universe cosmologies, like the Pre-Big-Bang scenario, since the predicted power spectrum of each model can present very different features.  In the general relativity theory (GRT) the usual procedure in order to evaluate the power spectrum of cosmological GWs is to expand small metric perturbations around the spatially homogeneous and isotropic Friedmann-Robsertson-Walker (FRW) metric. The next step is to identify GWs with the transverse-traceless (TT) part of the metric perturbations which do not couple with the perturbations of the perfect-fluid. Thus, once it was generated, this radiative gravitational field can freely travel through the space and reach an observer today.  The observational effect of GWs is to generate relative tidal accelerations between test particles. The Riemann tensor determines these relative accelerations and it is the only locally observable imprint of gravity. To see that, consider a freely falling observer at any fiducial point $P$ in the region. Let the observer set up an approximately Lorentz, normal coordinate system $\\{x^\\mu\\} = \\{t,x^i\\}$, with $P$ as origin. For a particle with spatial coordinates $x^i$ at rest, the relative acceleration with respect to $P$ is: \\begin{equation}\\label{relative accel} a_i = -R_{i0j0}x^j, \\end{equation} where $R_{i0j0}$ are the so-called `electric' components of the Riemann tensor due to waves or other external gravitational influences. When the linearized theory is considered, the Riemann tensor can be split in six algebraically independent components, but for the vacuum field equations of GRT they reduce to two, which represent the two polarization states of free GWs. These are the `TT-modes', also called the $+$ and $\\times$ polarizations.  Although these are the most studied modes in the GWs physics, when the framework of an alternative theory of gravity is considered, the number of non-null components of the Riemann tensor can be greater than two and the theory presents not only the TT-modes, but also other polarization states can appear. This is a direct consequence of the new field equations which can generate other radiative modes. Thus, for a generic theory of gravity, GWs can present up to six polarization states corresponding to the six independent components of $R_{i0j0}$.  Therefore, we should state that in order to work out the cosmological metric perturbations in an alternative theory of gravity, the first step is to find the number of independent polarizations of GWs in such a theory, i.e., the number of non-null components of the Riemann tensor. This can be done by a very elegant method which consists in the evaluation of the non-null Newman-Penrose (NP) quantities \\cite{Eardley1973,Newman} of a given theory. These quantities are the irreducible parts of the Riemann tensor written in a complex tetrad basis which make them very useful to evaluate the polarizations of GWs in an unambiguous way.  Thus, the present paper intends to consider the general formalism of cosmological perturbations in the context of alternative theories of gravity, focusing on the dynamical equations of the GW modes. The theory of cosmological perturbations in GRT has been largely studied in the literature. Some classical examples are the works by Lifshiftz \\cite{Lifshiftz1946}, Bardeen \\cite{Bardeen1980}, Peebles \\cite{Peebles1993}, and Mukhanov, Feldman and Brandenberger \\cite{Mukhanov1992}. For a recent review on cosmological dynamics see, e.g.,\\cite{Durrer2004,Malik2009}. In the usual approach, the TT-modes of GWs are consistently described by the superadiabatic, or parametric, amplification mechanism in the GRT \\cite{Grishchuk1974}. Furthermore, it was shown by Barrow and de Garcia Maia \\cite{Barrow1993_2,Maia1994} that the same mechanism applies for modified theories of gravity as scalar-tensor theories and the so-called $f(R)$ theories. Although they have analyzed only the evolution of the two TT-modes, it is known since the work by Eardley et al. \\cite{Eardley1973} that scalar-tensor theories present in addition at least one scalar GW polarization (more general scalar-tensor theories present two scalar GW polarizations) as a consequence of the additional degree of freedom included by the scalar field. The relic scalar GW production which arises from this kind of theories was recently discussed by Capozziello et al. \\cite{Capozziello2007}, and an upper limit was obtained from the amplitude of scalar perturbations in the Wilkinson Microwave Anisotropy Probe (WMAP) data. But in their analysis, Capozziello et al. have considered only the vacuum field equations, this is a limitation of their results since a coupling between scalar GWs and the scalar perturbations of the cosmological perfect fluid are expected as will be clear in our derivations.  In the case of $f(R)$ theories, using the NP formalism, it was shown that a particular class of functions $f(R)$ presents two scalar GW modes in addition to the $+$ and $\\times$ modes, thus totalizing four independent polarizations of GWs \\cite{Alves2009B}. However, when the Palatini approach is used in the derivation of the field equations, the theory reveals only the usual two TT polarizations. It was also found that the scalar longitudinal mode which appears in $f(R)$ theories is a massive mode which is potentially  detectable by the future space GW interferometer LISA \\cite{Capozziello2008}. Again by considering only the vacuum equations, the production of the relic GWs of this particular mode was also considered and constraints using the WMAP data was established \\cite{Corda2008}.  Moreover, the study of extra polarization states of cosmological GWs in the context of alternative theories of gravity can reveal new interesting features of these theories which do not appear in GRT. A remarkable example is the presence of vector longitudinal polarization modes of GWs in some theories. These modes give rise to a non usual Sachs-Wolf effect which leaves a vector signature on the CMB polarization \\cite{Bessada2008}. Otherwise, vector perturbations in GRT decay too fast and it would not leave any signature on CMB polarization.  Therefore, it is clear that the future detection of GWs, and the corresponding determination of the number of polarization modes, are powerful tools to test the underlying gravity theory. Thus, the goal of the present paper is to furnish a general formalism to find the evolution equations of all the possible polarization modes which could appear in a generic theory of gravity. Once the number of independent polarization modes are found and the corresponding evolutionary equations could be obtained, one is in a position to obtain the power spectrum of each mode, finding the CMB signatures and constraining the additional modes. In order to show the application of the formalism, and in order to find some new features of the current studied theories, we have chosen to obtain the dynamical equations for GWs in the context of two particular theories, namely, a general scalar-tensor theory and a bimetric massive theory of gravity.  First proposed by Brans and Dicke \\cite{Brans1961} in the aim of making the theory of gravity compatible with the Mach's principle, the scalar-tensor theories are of a great interest since, as pointed out by several authors, a coupling between a scalar field and gravity seems to be a generic outcome of the low-energy limit of string theories (see, e.g., \\cite{Casas1991}). Another interest in the scalar-tensor models is that the $f(R)$ theories can be written as the Einstein equations plus a scalar field, and thus we could in principle extend the same formalism applied for the scalar-tensor theories to the $f(R)$ field equations. The bimetric massive theory we consider was proposed by Visser \\cite{vis1998} in the aim to obtain general covariant field equations with massive gravitons. His method was based on the introduction of a non-dynamical metric $(g_0)_{\\mu\\nu}$ besides the physical metric $g_{\\mu\\nu}$. The resulting equations appear as a small modification of the Einstein field equations for which the massive gravitons and the metric $(g_0)_{\\mu\\nu}$ are present only in an additional energy-momentum tensor. Furthermore, our past studies have shown that the Visser's theory is a potential explanation for the current acceleration of the expansion of the Universe \\cite{Alves2006,Alves2009}.  In deriving the equations for GWs in the two theories we will first review how to obtain the number of independent polarization modes for any theory following the Eardley et al. approach \\cite{Eardley1973}. In the case of the scalar-tensor models the theory present four polarization states in the more general case. Otherwise, the Visser's theory is a simple example of how a weak modification of gravity can produce six polarization modes. The subsequent analysis show that all the polarization modes, apart from the usual $+$ and $\\times$ polarizations, are dynamically ``coupled'' to the perturbations of the cosmological perfect fluid. We argue that this kind of coupling and the existence of additional polarization states could furnish distinguishable signatures of alternative theories in the power spectrum of the relic GWs.  The paper is organized as follows: in the section \\ref{sec 0} we present an overview of the NP formalism starting from the definition of the NP quantities which define the six possible polarization states for GWs. Then we find the non-vanishing parameters for the GRT, scalar-tensor theories and for the Visser's model. In the section \\ref{sec 2}, considering a generic theory, we find general expressions for the perturbed Einstein tensor and for the energy-momentum tensor in the generalized harmonic coordinates. In the section \\ref{sec 3} we introduce a decomposition scheme which depends on the number of non-vanishing polarization modes of GWs which could appear in the various alternative theories. In the sections \\ref{sec 4} and \\ref{sec 5} we apply the formalism of the preceding sections for two particular theories, the scalar-tensor theory and the Visse's bimetric model. Finally, we present our conclusions and discussions in the section \\ref{sec 6}.  Throughout the paper we use units such that $c=1$ unless otherwise mentioned. ",
        "conclusions": ""
    },
    "1004/1004.4920_arXiv.txt": {
        "abstract": "We investigate the recently quantified misalignment of $\\alpha_{mis} \\approx 20^\\circ-40^\\circ$ between the 3-D geometry of stereoscopically triangulated coronal loops observed with STEREO/EUVI (in four active regions) and theoretical (potential or nonlinear force-free) magnetic field models extrapolated from photospheric magnetograms. We develop an efficient method of bootstrapping the coronal magnetic field by forward-fitting a parameterized potential field model to the STEREO-observed loops. The potential field model consists of a number of unipolar magnetic charges that are parameterized by decomposing a photospheric magnetogram from MDI. The forward-fitting method yields a best-fit magnetic field model with a reduced misalignment of $\\alpha_{PF} \\approx 13^\\circ-20^\\circ$. We evaluate also stereoscopic measurement errors and find a contribution of $\\alpha_{SE}\\approx 7^\\circ-12^\\circ$, which constrains the residual misalignment to $\\alpha_{NP}=\\alpha_{PF}-\\alpha_{SE}\\approx 5^\\circ -9^\\circ$, which is likely due to the nonpotentiality of the active regions. The residual misalignment angle $\\alpha_{NP}$ of the potential field due to nonpotentiality is found to correlate with the soft X-ray flux of the active region, which implies a relationship between electric currents and plasma heating. ",
        "introduction": "The STEREO mission provides us an unprecedented view of the solar corona, enabling us for the first time to fully constrain the three-dimensional (3-D) geometry of the coronal magnetic field. Stereoscopic triangulation of coronal loops has been conducted at small STEREO spacecraft separation angles ($\\alpha_{sep} \\lapprox 10^\\circ$), for several active regions observed with STEREO A(head) and B(ehind) in April and May 2007 (Aschwanden et al.~2008a,b; 2009). The reconstructed 3-D geometry of STEREO-observed coronal loops has been compared with theoretical magnetic field models based on extrapolations from photospheric magnetograms, using nonlinear force-free field (NLFFF) models (DeRosa et al.~2009), as well as potential and stretched potential field models (Sandman et al. 2009), but surprisingly it turned out that the two types of magnetic field lines exhibited an average misalignment angle of $\\alpha_{mis} \\approx 20^{\\circ}-40^{\\circ}$, regardless of what type of theoretical magnetic field model was used. From this dilemma it was concluded that a more realistic physical model is needed to quantify the transition from the non-force-free photospheric boundary condition to the nearly force-free field at the base of the solar corona (DeRosa et al.~2009).  At this junction, it is not clear what a viable method is to obtain a force-free boundary of the magnetic field at the coronal base, or how to correct the non-force-free magnetograms. However, the stereoscopic triangulation supposedly provides the correct 3-D directions of the magnetic field ${\\bf B}({\\bf x})$, which together with Maxwell's equation of divergence-freeness ($\\nabla {\\bf B}=0)$, constrain also the absolute values of the field strengths. In this Paper I we choose a magnetic field model that is defined in terms of multiple unipolar charges. An approach in terms of multiple dipoles is employed in Paper II (Sandman and Aschwanden 2010). Since both unipolar or dipolar magnetic fields represent potential magnetic fields that fulfill the divergence-free condition, the superposition of multiple unipolar and dipolar magnetic field components fulfill the same condition. We develop a numerical code of such a parameterized divergence-free magnetic field that can be forward-fitted to the 3-D geometry of stereoscopically triangulated coronal loops. So, the simple goal of this study is to evaluate how closely the stereoscopically observed loops can be modeled in terms of potential fields, a goal that was already attempted with Skylab observations (Sakurai and Uchida 1977). Modeling with non-potential fields, such as nonlinear force-free field (NLFFF) models, will be considered in future studies.  This paper is organized as follows: The definition of a parameterized potential field is described in Section 2, the development and tests of a numeric magnetic field code and the results of forward-fitting to stereoscopically triangulated loops is presented in Section 3, and conclusions follow in Section 4. An alternative approach with dipolar magnetic fields is the subject of Paper II (Sandman and Aschwanden 2010). ",
        "conclusions": "  \\begin{enumerate} \\item{The amount of misalignment can be reduced to about half of the values for potential-field models optimized by a bootstrapping method that minimizes the field directions with the stereoscopically triangulated loops. Our potential-field model is parameterized with $\\approx 200$ unipolar charges per active region, whose positions and field strengths are approximately derived from a gaussian decomposition of a photospheric magnetogram, and then varied until a best fit is obtained. The best-fit potential field model has an improved misalignment of $\\alpha_{PFU} \\approx 13^\\circ-20^\\circ$. Because the best-fit potential field model defines an improved magnetic field boundary condition at the bottom of the corona, the difference to the observed photospheric magnetogram contains information on the currents between the photosphere and the base of the force-free corona.} \\item{We estimate the misalignment contriubtion caused by stereoscopic correlation errors from self-consistency measurements between the magnetic field misalignments of adjacent loops. We find contributions in the order of $\\Delta \\alpha_{SE} \\approx 7^\\circ-12^\\circ$.} \\item{We estimate the contributions to the field misalignment due to non-potentiality caused by electric currents from the residuals between the best-fit potential field and the stereoscopic triangulation errors and find misalignment contributions in the order of $\\alpha_{NP}\\approx 5^\\circ-9^\\circ$.} \\item{The overall average misalignemnt angle between potential field models and stereoscopic loop directions, as well as the contribution to the misalignment due to non-potentiality, are found to correlate with the soft X-ray flux of the active region, which suggests a correlation between the amount of electric currents and the amount of energy dissipation in form of plasma heating in an active region.} \\end{enumerate}  In this study we identify for the first time the contributions to the misalignment of the magnetic field, in terms of optimized potential field models, non-potentiality due to electric currents, and stereoscopic triangulation errors. These results open up a number of new avenues to improve theoretical modeling of the coronal magnetic field. First of all, optimized potential field models can be found that represent a suitable lower boundary condition at the base of the force-free corona, which provides a less computing-expensive method than nonlinear force-free codes. Second, methods can be developed that allow us to localize electric currents in the non-force-free photophere and chromosphere. Third, the misalignment angle can be used as a sensitive parameter to probe the evolution of current dissipation, energy build-up in form of non-potential magnetic energy in different quiescenct and flaring zones of active regions. The high-resolution magnetic field data from Hinode and {\\sl Solar Dynamics Observatory} provide excellent opportunities to obtain better theoretical models of the coronal magnetic field using our bootstrapping method, which is not restricted to stereoscopic data only, but can also be applied to single-spacecraft observations.  \\medskip Acknowledgements: We are grateful to helpful discussions with Marc DeRosa and Allen Gary. This work was partially supported by the NASA contract NAS5-38099 of the TRACE mission and by NASA STEREO under NRL contract N00173-02-C-2035. The STEREO/SECCHI data used here are produced by an international consortium of NRL, LMSAL, RAL, MPI, ISAS, and NASA. The MDI/SoHO data were produced by the MDI Team at Stanford University and NASA."
    },
    "1004/1004.5321_arXiv.txt": {
        "abstract": "We compare mid-infrared emission-line properties, from high-resolution {\\it Spitzer} spectra of a  hard X-ray (14 -- 195 keV) selected sample of nearby (z $< 0.05$) AGN detected by the Burst Alert Telescope (BAT) aboard {\\it Swift}. The luminosity distribution for the mid-infrared emission-lines, [O~IV] 25.89 $\\mu$m, [Ne~II] 12.81 $\\mu$m, [Ne~III] 15.56 $\\mu$m and  [Ne~V] 14.32/24.32 $\\mu$m, and hard X-ray continuum show no differences between Seyfert 1 and Seyfert 2 populations, however six newly discovered BAT AGNs are  under-luminous in [O~IV], most likely the result of dust extinction in the  host galaxy. The overall tightness of the  mid-infrared correlations and BAT fluxes and luminosities suggests that the emission lines primarily arise in gas ionized by the AGN. We also compare  the mid-infrared emission-lines in the BAT AGNs with those from published studies of ULIRGs, PG~QSOs, star-forming galaxies and LINERs. We find that the  BAT AGN sample fall into a distinctive region when comparing  the [Ne~III]/[Ne~II] and the [O~IV]/[Ne~III] ratios. These line ratios are lower in sources that have been previously classified in the mid-infrared/optical as AGN  than those found for the BAT AGN,  suggesting that, in our X-ray selected sample, the AGN represents  the main contribution to the  observed line  emission. These  ratios  represent a new  emission line diagnostic  for distinguishing  between AGN and star forming galaxies. ",
        "introduction": "Active galactic nuclei (AGN) span over seven orders of magnitude in bolometric luminosity ($L_{bol}$) \\citep{1999PASP..111....1K} and yet are all believed to be powered by the same physical mechanism: accretion of matter onto supermassive black holes \\citep[e.g.,][]{1984ARA&A..22..471R,2004ApJ...613..682P}. One way to approach the study of AGN is to concentrate on those in the local Universe (e.g. $z < 0.05$), which permits us, among other things, to determine the properties of the host galaxy. Such studies tend to focus on Seyfert galaxies, which are modest luminosity AGN ($L_{bol} \\lesssim 10^{45}$ erg s$^{-1}$), but bright enough, due to their proximity, to be studied across the full electromagnetic spectrum.    Although Seyfert galaxies and other AGN have been traditionally defined in terms of their optical properties \\citep[e.g., classification into  Type I and Type II,][]{1974ApJ...192..581K}, sample selection of AGN via a single waveband can lead to observational bias \\citep[e.g.,][]{1994ApJ...436..586M}. For example, most AGN are obscured from our line of sight by dust and gas \\citep{2000A&A...355L..31M} and any selection based on optical (or UV) properties would miss many objects or could highly skew a sample towards unobscured objects \\citep[e.g.,][]{2005AJ....129..578B}. The soft X-ray properties of Seyfert galaxies generally follow the same dichotomy as their optical properties. The X-ray continuum source in Seyfert~1s can be observed directly  \\citep[e.g.,][]{1998ApJS..114...73G}, while the central X-ray source is sometimes undetectable in Seyfert 2s, due to material with ${\\rm N_H} >$ 10$^{22.5}$~cm$^{-2}$ along our line of sight. There is additionally a wide range of effective IR colors \\citep{2003ApJ...590..128K,2005ApJ...632L..13L} which can introduce selection bias. A comparison of infrared and X-ray data \\citep{2005AJ....129.2074F} shows a factor of 30 range in the IR 24~$\\micron$ to $\\sim$4 keV X-ray flux ratio for X-ray selected AGN, suggesting a range of geometries and optical depths for dust reprocessing, and probably variance in the intrinsic power law AGN continuum. Obscuration and star formation in the host galaxy can  also  dominate and introduce confusion   in the IR  \\citep[e.g.,][]{2004A&A...418..465L,2006ApJ...642..126B}. In fact, virtually all surveys for AGN based purely on IR, optical, UV or soft X-ray data have been biased \\citep{2004ASSL..308...53M}.  Even Sloan surveys \\citep{2004ApJ...613..109H} or IR surveys \\citep{2005AJ....129.2074F} have required indirect AGN indicators which are known to be not necessarily robust \\citep{2008ApJ...682...94M}.    To understand the intrinsic properties of AGN as a class, it is critical to start with a survey where we can be as certain as possible that we are viewing the AGN-only parts of these galaxies. At X-ray energies of E $> 10-20$ keV, the obscuring material is relatively optically thin for column densities less than $\\sim 3 \\times 10^{24}$ cm$^{-2}$ (Compton-thin objects). Even if an AGN is well buried within its host galaxy there is an unaffected view of the central power source. A hard X-ray survey should thus find all Compton thin AGN in a uniform fashion and is the most representative, since at present, there are very few, if any, known X-ray ``quiet\" AGN. Such a hard X-ray survey is now available from the {\\it Swift} Burst Alert Telescope (BAT). The Swift BAT is sensitive over $\\sim$85\\% of the sky to a flux threshold of $2 \\times 10^{-11}$ ${\\rm ergs~cm^{-2} s^{-1}}$ in the 14$-$195~keV band \\citep{2005ApJ...633L..77M}. The BAT data are about 10 times more sensitive than the previous hard X-ray all sky survey \\citep{1984ApJS...54..581L}. The BAT detects all bright AGN, whether they are obscured or not. Moreover,  several of the BAT sources are newly discovered AGN or have been poorly studied, if at all, at other wavelengths \\citep{2008ApJ...674..686W,2008ApJ...681..113T,2010ApJS..186..378T}.    Nevertheless, although all of the BAT-detected objects are true AGN, in order to fully explore the properties of these AGN one needs to take a multi-wavelength approach. For example, studying the IR properties of the BAT AGN will provide insight into the IR/X-ray scatter and thus determine the true distribution of IR properties. There have been a large number of studies of the mid-infrared emission line properties of active galaxies using both {\\it Infrared Space Observatory} \\citep{1996A&A...315L..27K} and {\\it Spitzer Space Telescope} \\citep{2004ApJS..154....1W}. The ratios of high- and low-ionization mid-infrared emission lines have been widely used to separate the relative contribution of the  AGN and star formation \\citep[e.g.,][]{1998ApJ...498..579G,2002A&A...393..821S,2006ApJ...646..161D,2007ApJ...656..148A,2007ApJ...667..149F,2008ApJ...689...95M,2010ApJ...710..289B}. More recently, \\cite{2009ApJ...704.1159H} (H09) used new high-resolution {\\it Spitzer} spectroscopy to probe the utility of mid-infrared emission line diagnostics as a way to separate active galaxies from star forming galaxies. In our first  study  of mid-infrared properties of the BAT AGNs \\citep{2008ApJ...682...94M}, we found the [O~IV]25.89$\\mu$m to be an accurate indicator of the AGN luminosity, with an uncertainty of $\\sim$0.3~dex; this result has been   confirmed  using   larger samples  \\citep{2009ApJ...700.1878R,2009ApJ...698..623D}. Using a complete, volume-limited,   sample of galaxies \\cite{2009MNRAS.398.1165G}  (GA09) demonstrated the utility of high-ionization mid-infrared emission lines, such as [Ne~V]14.32$\\mu$m,  to identify AGN including those that were not identified as AGN in optical studies \\citep[see also,][]{2007ApJ...669..109L,2008ApJ...678..686A,2009ApJ...691.1501D,2009ApJ...704..439S}. Similar results have been found by \\cite{2009ApJS..184..230B} (B09) in their study of starburst galaxies.    This paper is the first in a series seeking to understand the nature of the observed mid-infrared luminosities in AGN and their wide variety of spectral forms. Here we report results from the portion of our sample that have high-resolution {\\it Spitzer} spectra. This work complements the extensive optical imaging and spectroscopy of the AGN population \\citep[Koss, accepted in ApJL;][]{2010ApJ...710..503W} and the detailed analysis of the X-ray properties of the BAT AGN sample \\citep[e.g.,][]{2008ApJ...674..686W,2009ApJ...690.1322W,2009ApJ...701.1644W}. In following papers we will report on the results from our analysis of the low-resolution {\\it Spitzer} spectra which will include the study of polycyclic aromatic hydrocarbon (PAH) features, silicate absorption and mid-infrared continuum properties of the BAT AGN  sample. In order to calculate the luminosities presented in this work  we assumed a flat universe with a Hubble constant $H_o=71{\\rm kms^{-1}Mpc^{-1}}$, $\\Omega_\\Lambda=0.73$ and ${\\rm \\Omega_M=0.27}$, with redshift values taken from NASA's ExtraGalactic Database (NED), except for sources with redshift values of $z<0.01$ were distances are take from The Extragalactic Distance Database (EDD) \\citep{1988ngc..book.....T,2009AJ....138..323T}. ",
        "conclusions": "Using high-resolution {\\it Spitzer}  IRS spectra, we have examined the mid-infrared emission-line properties of a sample of hard X-ray selected AGN, detected by {\\it Swift}/BAT. Our principle conclusions are as follows.  1. The luminosity distribution for the mid-infrared emission lines and BAT continuum luminosities show no differences between Seyfert~1 and Seyfert~2 populations for the BAT sample.  The correlations between all the mid-infrared emission lines and BAT in both flux--flux and luminosity--luminosity are statistically significant, even when factoring the distance effect in luminosity-luminosity correlations.  The dispersion/tightness in these correlations is due to differences in the X-ray absorbing column densities, dust extinction and/or nuclear star formation activity. Moreover, the tight correlation found in the [Ne~III]-BAT relationship suggests that, on average,  there is no strong enhancement due to  star formation in the [Ne~III] emission in the BAT sample. Also, the slopes for the [Ne~III],[Ne~V] and [O~IV] versus BAT luminosities relationships are smaller in Seyfert~1 galaxies than in Seyfert~2s (which are around unity), which suggests that, while the amount of extinction towards the NLR is similar in both types, the X-ray absorbing columns are large enough in Seyfert~2s to affect the hard X-ray band, confirming the results of \\cite{2008ApJ...682...94M} and \\cite{2009ApJ...700.1878R}. This result  is in agreement with the  fact that the BAT/[O~IV] ratio statistically separates Seyfert~1s and Seyfert~2s.   2.  Although all of the correlations among the mid-infrared emission lines are strong, the worst correlations are for [Ne~V]-[Ne~II] and [O~IV]-[Ne~II], because of enhancement of the [Ne~II] from nuclear stellar activity \\citep[see also][]{2008ApJ...689...95M}. While the tightness of these mid-infrared correlations suggests that dust extinction is not the driving physical process behind the mid-infrared relationships, approximately $\\sim$40$\\%$ (including upper limits) of the sample have values for the ratio of the [Ne~V] emission lines below the low-density theoretical limit, suggesting dust extinction as the physical process responsible. Exploring this, we found that all of the newly discovered BAT AGNs in our sample, which are under-luminous in [O~IV] and [Ne~V]14/24$\\mu$m, are found on  inclined host galaxies, and all but one have [Ne~V] ratios below the critical density limit. Hence, it is likely that the newly found BAT  AGN in our sample lack optical AGN signatures because of host galaxy extinction towards their NLRs. However the lack of correlation between host galaxy inclination and the neon ratios suggest that extinction along the plane of the host galaxy cannot be responsible for  the observed extinction in all the BAT AGN sample.   3. We compared the BAT AGNs with different starburst and H~II galaxies, so-called [Ne~V] active galaxies, and LINERs \\citep{2009MNRAS.398.1165G,2009ApJS..184..230B,2009ApJ...704.1159H}. We found that the BAT AGN   fall into a distinctive region based on the [Ne~III]/[Ne~II] and [O~IV]/[Ne~III] ratios. Using [Ne~III] and [O~IV] emission, previously connected with  AGN power \\citep[e.g.,][]{2007ApJ...655L..73G,2008ApJ...682...94M}, does not unambiguously identify AGNs as an stand alone diagnostic because Wolf-Rayet stars or another energetic phenomena (perhaps ULXs) could enhance the observed emission. While it is likely that detection of [Ne~V] indicates the presence of an AGN, the strongest of the [Ne~V] lines have $\\sim$1/3  less flux than [O~IV] an thus will be more difficult to detect in weak or faint AGN. Therefore, our composite method  using the  [Ne~II], [Ne~III] and [O~IV], represents a strong  and simple diagnostic by using only three emission lines to identify an AGN. Based on this, we found that   NGC~520, NGC~1614, NGC~4536, NGC~4676 could harbor an AGN, although these starburst galaxies don't show optical evidence of such (B09;GA09). Of these, there is X-ray evidence for a Compton thick AGN in NGC~1614 \\citep{2000A&A...357...13R}, [Ne~V]~14.31\\micron~emission detected in NGC~4536 \\citep{2008ApJ...677..926S} and NGC~4676 appears in the multi-wavelength LINER catalogue compiled by \\cite{1999RMxAA..35..187C}.  Such line  diagnostic will be particularly useful to analyze spectra from new IR missions, such as the {\\it James Webb Space Telescope} \\citep{2006SSRv..123..485G}.    We also found that ULIRGs and PG QSOs occupy two distinctive regions in our emission line diagnostic.  Most ULIRGs fall into  the gap between the  BAT AGN and the SB/HII branch, in agreement with the idea that ULIRGs are composite systems mainly powered by stellar activity \\citep{2007ApJ...656..148A,2009ApJS..182..628V}. On the other hand,  PG QSOs overlap with the BAT AGN branch, in agreement with the high, typically larger than $\\sim 80\\%$, AGN contribution to their  bolometric luminosities \\citep{2009ApJS..182..628V}.    Finally, most of the non-BAT AGNs presented in our study, AGN that have been selected because their optical and mid-infrared emission line properties, have smaller mid-infrared ratios than that found for the BAT AGN. In this regard, half of the BAT sample can be uniquely distinguished  from SB/HII/BCD galaxies by having  both [Ne~III]/[Ne~II] and [O~IV]/[Ne~III] ratios greater than unity. Moreover, when comparing  the  12\\micron~and our BAT selected AGN we found that the  [Ne~III]/[Ne~II] ratio distribution between the samples  is statistically different with   sources in the 12\\micron~sample having on average   lower ratios than that found in the BAT AGN, or alternatively higher recent stellar activity. This mild  contamination due to star formation becomes noticeable when comparing the 12\\micron~sample with the 14--195~keV sample, the latter of which is less  biased towards star-forming systems.    Despite the fact that  both samples have a strong AGN contribution to their observed narrow line emission, this result suggests that the BAT sample represents a unique opportunity to study high ionization  AGN, sources in which their optical/mid-infrared  emission signatures are dominated by the AGN, thus, providing  the most representative  sample in terms of galaxy population  and stellar content."
    },
    "1004/1004.5098_arXiv.txt": {
        "abstract": "%  Gravitational waves from the final stages of inspiralling binary neutron stars are expected to be one of the most important sources for ground-based gravitational wave detectors.  The masses of the components are determinable from the orbital and chirp frequencies during the early part of the evolution, and large finite-size (tidal) effects are measurable toward the end of inspiral, but the gravitational wave signal is expected to be very complex at this time.  Tidal effects during the early part of the evolution will form a very small correction, but during this phase the signal is relatively clean.  The accumulated phase shift due to tidal corrections is characterized by a single quantity related to a star's tidal Love number.  The Love number is sensitive, in particular, to the compactness parameter $M/R$ and the star's internal structure, and its determination could provide an important constraint to the neutron star radius.  We show that the Love number of normal neutron stars are much different from those of self-bound strange quark matter stars.  Observations of the tidal signature from coalescing compact binaries could therefore provide an important, and possibly unique, way to distinguish self-bound strange quark stars from normal neutron stars. ",
        "introduction": "Gravitational waves from the final stages of inspiralling binary neutron stars are expected to be one of the most important sources for ground-based gravitational wave detectors \\citep{C1993}.  To date, LIGO observations have only been able to set an upper limit to the neutron star-neutron star coalescence rate of 0.039 yr$^{-1} L_{10}^{-1}$ \\citep{A2009}, where $L_{10}$ is the blue luminosity in units of $10^{10} {\\rm~L}_\\odot$, which translates to about 0.075 events per year in the Milky Way.  This is a thousand times larger than the predicted rates \\citep{K2004}.  Nevertheless, the observed neutron star-neutron star inspiral rate from the universe is expected to be about 2 per day in LIGO II \\citep{K2004}.  The masses of the components will be determined to moderate accuracy, especially if the neutron stars are slowly spinning, during the early part of the evolution \\citep{CF1994,LIGO}.  Mass measurements from inspiralling binaries will be useful, especially in constraining the equation of state through limits to the neutron star maximum and minimum masses, but constraints to the radius would be much more effective in constraining the nuclear equation of state \\citep{LP2001}.  Large finite-size effects, such as mass exchange and tidal disruption, are measurable toward the end of inspiral \\citep{BC1992}, but the gravitational wave signal is expected to be very complex during this period. \\citet{FH2008} have recently pointed out that tidal effects are also potentially measurable during the early part of the evolution when the waveform is relatively clean. The tidal fields induce quadrupole moments on the neutron stars. This response of each star to external disturbance is described by the Love number $k_2$ \\citep{L1909}, which is a dimensionless coefficient given by the ratio of the induced quadrupole moment $Q_{ij}$ and the applied tidal field $E_{ij}$ \\begin{equation} \\label{k2} Q_{ij}=-k_2\\frac{2 R^5}{3 G}E_{ij}\\equiv-\\lambda E_{ij} \\, , \\end{equation} where $R$ is the radius of the star and $G$ is the gravitational constant. The tidal Love number $k_2$, which is dimensionless, depends on the structure of the star and therefore on the mass and the equation of state (EOS) of dense matter.  The quantity $\\lambda$ is the induced quadrupole polarizability.  Tidal effects will form a very small correction in which the accumulated phase shift can be characterized by a single quantity $\\bar\\lambda$ which is a weighted average of the induced quadrupole polarizabilities for the individual stars, $\\lambda_1$ and $\\lambda_2$.  Since both neutron stars have the same equation of state, the weighted average $\\bar\\lambda({\\cal M})$, as a function of chirp mass ${\\cal M}=m_1^{3/5}m_2^{3/5}/(m_1+m_2)^{1/5}$, is relatively insensitive to the mass ratio $m_1/m_2$, as is shown by \\citet{Hind2009}.  We therefore focus on the behavior of the quadrupole polarizability $\\lambda$ of individual stars.  These are related to the dimensionless tidal Love number $k_2$ for each star by $k_2=(3/2)G\\lambda R^{-5}$.  The Love number $k_2$ is sensitive to the neutron star equation of state, in particular to the compactness parameter $M/R$ as shown by \\citet{DN2009} and the overall compressibility of the equation of state.  In particular, the tidal Love numbers of strange quark matter stars are qualitatively different from those of normal matter stars.  In a fashion similar to moment of inertia measurements from relativistic binary pulsars \\citep{LS2005}, an important constraint to the neutron star radius might become possible from gravitational wave observations. Detection of the tidal signature from coalescing compact binaries might provide an important, and possibly unique, way to distinguish self-bound strange quark matter stars from normal neutron stars.  Our paper is organized as follows.  In Sec. I, a new technique for the computation of tidal Love numbers is described. The influence of density discontinuities and phase transitions on Love numbers is discussed in Sec. II. Results of Love numbers for polytropic equations of state are presented in Sec. IV. Sec. V contains results for select analytic solutions of Einstein's equations  in spherical symmetry.  Love numbers for proposed model equations of state for normal stars with hadronic matter and self-bound stars with strange quark matter with and without crusts are given in Sec. VI,  wherein a comparison of results between these two distinct classes of stars are also made. In Sec VII, we discuss the role of a solid crust on Love numbers. Our results and conclusions are summarized in Sec. VII. Relevant parameters required for the computation of Love numbers for analytic solutions of Einstein's equations (discussed in Sec. V) are  to be found in Appendix A. ",
        "conclusions": "The quadrupole polarizabilities of normal neutron stars and self-bound quark matter stars have been calculated for a wide class of proposed equations of state of dense matter for both normal and strange quark matter stars. The quadrupole polarizabilities $\\lambda=2 R^5 k_2/(3 \\, G)$ are characterized by the dimensionless Love number $k_2$ and both are sensitive to the equation of state, in particular to the compactness parameter $M/R$ and the overall compressibility of the equation of state. For normal neutron stars, $k_2$ and $\\lambda$ exhibit pronounced maxima for configurations with masses close to a solar mass for most equations of state. The maximum value of $k_2$ is not very sensitive to the EOS, lying in the range 0.1--0.14.  In each case, maximum mass configurations have significantly lower values of $k_2$ and $\\lambda$ than their solar mass counterparts.  Love numbers for self-bound strange quark matter stars with or without crusts are qualitatively different than those of normal neutron stars. The maxima in the value of $k_2$ for strange quark matter stars without crusts occurs for masses less than 0.1 M$_\\odot$, and maximum values of order 0.8 are achieved.  As in the normal matter case, the maxima in quadrupole polarizabilities occurs for configurations near 1 M$_\\odot$.  In contrast, the magnitudes of quadrupole polarizabilities of strange quark matter stars are usually much less than those of normal stars, owing to the larger radii of the latter.  Our investigations also point the need to examine the core-crust interface region of both normal and self-bound quark matter stars more closely.  The important issue that bears close scrutiny is the precise nature (first or second order) of possible phase transitions.  In the case that strong discontinuities exist near the core-crust interface of strange quark matter stars, dimensionless Love numbers are suppressed for low mass stars relative to the cases for which there is no crust.  However, for stars of order 1 M$_\\odot$ or larger, the presence or absence of a crust has little influence on Love numbers.  The strength of the tidal signatures from coalescing compact binaries is proportional to $\\lambda$, and is therefore quite sensitive to the radii of the stars.  For stellar configuratons with radii of order 11 km or less, the tidal response might be too small to observe, implying that a positive detection might be sufficient to rule out the presence of a self-bound star, such as a strange quark matter star, in the observed system."
    },
    "1004/1004.0959_arXiv.txt": {
        "abstract": "The strong variability of magnetic central engines of AGN and GRBs may result in highly intermittent strongly magnetized relativistic outflows. We find a new magnetic acceleration mechanism for such impulsive flows that can be much more effective than the acceleration of steady-state flows. This impulsive acceleration results in kinetic-energy-dominated flows that are conducive to efficient dissipation at internal MHD shocks on astrophysically relevant distances from the central source.  For a spherical flow, a discrete shell ejected from the source over a time $t_0$ with Lorentz factor $\\Gamma\\sim 1$ and initial magnetization $\\sigma_0 = B_0^2/4\\pi\\rho_0c^2\\gg 1$ quickly reaches a typical Lorentz factor $\\Gamma \\sim \\sigma_0^{1/3}$ and magnetization $\\sigma\\sim\\sigma_0^{2/3}$ at the distance $R_0\\approx ct_0$.  At this point the magnetized shell of width $\\Delta \\sim R_0$ in the lab frame loses causal contact with the source and continues to accelerate by spreading significantly in its own rest frame. The expansion is driven by the magnetic pressure gradient and leads to relativistic relative velocities between the front and back of the shell. While the expansion is roughly symmetric in the center of momentum frame, in the lab frame most of the energy and momentum remain in a region (or shell) of width $\\Delta\\sim R_0$ at the head of the flow. This acceleration proceeds as $\\Gamma \\sim (\\sigma_0R/R_0)^{1/3}$ and $\\sigma \\sim \\sigma_0^{2/3} (R/R_0)^{-1/3}$ until reaching a coasting radius $R_c \\sim R_0\\sigma_0^2$ where the kinetic energy becomes dominant: $\\Gamma \\sim \\sigma_0$ and $\\sigma \\sim 1$ at $R_c$. Then the shell starts coasting and spreading (radially), its width growing as $\\Delta\\sim R_0(R/R_c)$, causing its magnetization to drop as $\\sigma\\sim R_c/R$ at $R>R_c$.  Given the typical variability time-scales of AGN and GRBs, the magnetic acceleration in these sources is a combination of the quasi-steady-state collimation acceleration close to the source and the impulsive (conical or locally quasi-spherical) acceleration further out.  The interaction with the external medium, which can significantly affect the dynamics, is briefly addressed in the discussion. ",
        "introduction": "\\label{sec:introduction}  The first questions raised by the discovery of astrophysical jets are how they are powered, collimated, and accelerated.  Most of them -- jets from young stars, Active Galactic Nuclei (AGN), Galactic X-ray Binaries, and Gamma Ray Bursts (GRBs), are associated with disk accretion\\footnote{The only exceptions are the jets of Pulsar Wind Nebulae as there are no indications of accretion disks around their pulsars.  These jets are most likely not produced directly by the pulsars but instead form downstream of the termination shock of pulsar winds~\\citep{L02,KL04}.}, and this suggests that accretion disks are essential for jet production. The astrophysical jets seem to be highly supersonic as many of their features are nicely explained by internal shocks.  In the laboratory, highly collimated supersonic jets are normally produced when a high pressure (and temperature) gas escapes from a chamber via a finely designed nozzle.  However, it seems highly unlikely that such refined ``devices'' are formed naturally in astrophysical systems. They would require cold and dense gas to form the walls of the chamber with a massive compact object in the center \\citep{BR74}, but such configurations are highly unstable \\citep{NSSW81,SSNW83}. This has lead to the idea that the collimation of astrophysical jets may have a completely different mechanism involving a strong magnetic field.  Although this magnetic field still needs to be confined within a channel, the conditions on its geometry are less restrictive. If this field is anchored to a rotating object, such as an accretion disk, then it naturally develops an azimuthal component. The hoop stress associated with this magnetic field component creates additional collimation of the flow within the channel. Moreover, this leads to a magnetic torque being applied to the rotating object and thus a natural way of powering outflows by tapping the rotational energy of the central object.  In order to produce a relativistic flow this way, the magnetic energy per particle must exceed its rest energy. Thus, the jet plasma must be highly rarefied. Such rarefied plasma is naturally produced only in the magnetospheres of black holes and neutron stars.  Moreover, the strong magnetic field shields these magnetospheres and prevents them from being contaminated by the much denser surrounding plasma.  In contrast, young stars can eject a lot of mass from their surface and this seems to explain why their jets are not relativistic.  Magnetospheres of accretion disks are likely to be heavily mass-loaded and are not able to produce relativistic jets for the same reason.  It has to be stressed that magnetic flows must still be collimated externally until they become super-fast-magnetosonic. The magnetic hoop stress can result in self-collimation of the inner core but cannot prevent sideways expansion of the outer sheath. However, when the flow becomes super-fast-magnetosonic, the speed of this lateral expansion becomes smaller than the flow speed along the jet direction, and the jet remains collimated.  For non-relativistic jets the condition of passing through the fast-magnetosonic surface also implies almost completed acceleration of the flow (50\\% conversion of magnetic energy into kinetic energy). In contrast, the relativistic jets still remain Poynting-flux dominated at this point and the acceleration process may continue well into the super-fast-magnetosonic regime.  The issue of the efficiency of energy conversion (from magnetic to kinetic form) is related to the issue of subsequent energy dissipation, which is required in order to explain the observed electromagnetic emission from both the jets and the structures they create when they collide with the external medium. Traditionally, one of the most favorite channels of dissipating the energy of supersonic flows has been the formation of shock waves. However, in the case of relativistic flows this mechanism can be much less efficient if the flow is Poynting-flux dominated. First of all, it is the kinetic energy of the flow that is dissipated\\footnote{This applies to fast magnetosonic shocks. At a slow magnetosonic shock, the magnetic energy dissipates as well and the kinetic energy can actually increase. However, slow shocks are much less robust and harder to generate compared to the fast ones.}, and if only a small fraction of the total energy is in the kinetic form then this already severely limits the efficiency of dissipation. Secondly, the compression ratio and hence the fraction of kinetic energy that dissipates also decrease with increasing magnetization. Thus, in order to dissipate a significant fraction of the available energy the flow should not only become super-fast-magnetosonic, but it should also become dominated by kinetic energy before it is shocked \\citep{Leis05,MGA09,MA10}.  The magnetic acceleration of relativistic flows has been the subject of theoretical research for decades. The main focus of this research has been on the models of steady-state axisymmetric dissipation-free flows (the ``standard model''). The main reason behind this is simplicity. Only in this case was there a hope of building a rigorous theory. Yet, even this idealized model is rather complex, and solutions could be found only if an additional symmetry, e.g. self-similarity, or other simplifying condition was introduced \\citep[e.g.][]{BL92,VK03,BN06}.  More recently the problem was approached using numerical methods \\citep{KBVK07,KVKB09}.  There are a number of problems with the standard model, which are most severe in the case of a spherical wind. In this case the theory predicts an asymptotic Lorentz factor of $\\Gamma \\sim \\sigma_0^{1/3}$, where $\\sigma_0 = B_0^2/4\\pi\\rho_0c^2 \\gg 1$ is the initial magnetization parameter, which determines the maximum possible Lorentz factor corresponding to a total conversion of the Poynting flux into the bulk motion kinetic energy in a steady-state flow \\citep[e.g.,][]{GJ70}. This is in conflict with the observations of many astrophysical sources. In particular, the high observed values of $\\Gamma$ in many sources would require an extremely large initial magnetization $\\sigma_0$ that would in turn imply a very high asymptotic magnetization, $\\sigma \\sim \\sigma_0^{2/3} \\gg 1$, making it impossible to achieve efficient shock dissipation within the outflow.  A potential way to overcome this problem is by resorting to collimated outflows. This can increase the asymptotic value of $\\Gamma$ and reduce that of $\\sigma$ by up to a factor of $\\sim\\theta_{\\rm jet}^{-2/3}$, where $\\theta_{\\rm jet}$ is the asymptotic half-opening angle of the jet.  The collimation has to be strong enough to preserve causal connectivity across the flow (in the lateral direction). The faster the flow and the higher its fast-magnetosonic Mach number becomes, the smaller its opening angle should be. By the time one half of the Poynting flux is converted into kinetic energy ($\\sigma\\sim 1$), the jet half-opening angle $\\theta_{\\rm jet}$ should not exceed $\\theta_{\\rm max} = 1/\\Gamma$, where $\\Gamma \\sim \\sigma_0$ is the jet Lorentz factor at that time.  Observations of AGN jets do indeed show that $\\theta_{\\rm jet}< 1/\\Gamma$ ~\\citep{P09}.  However, for GRB jets with $\\Gamma\\simeq 400$ (or $10^2\\la\\Gamma \\la 10^{3.5}$) this constraint gives $\\theta_{\\rm max}\\simeq 0.14^\\circ$ (or $0.018^\\circ\\la\\theta_{\\rm max} \\la 0.57^\\circ$), which is much smaller compared to generally accepted values of the half-opening angle, $2^\\circ \\la\\theta_{\\rm jet} \\la 30^\\circ$ \\citep{FWK00,PK01}.  In addition, the standard theory of GRB afterglow emission can explain the jet-break in their light curves only if $\\theta_{\\rm jet}\\Gamma \\gg 1$ \\citep{Rhoads99,SPH99}. Although the Swift observations show that clear jet breaks are not as common as we used to think \\citep[e.g.,][]{Liang08}, this might be at least partly due to observational selection effects (Swift GRBs are dimmer on average as Swift is more sensitive than previous missions), and there are still some clear cases for jet breaks in the Swift era. Finally, late time radio afterglow observations, when the flow becomes sub-relativistic, provide fairly robust (no longer susceptible to strong relativistic beaming) lower limits \\citep[e.g.,][]{EW05} on the true energy that remains in the afterglow blast wave at that time, of a few to several times $10^{51}\\;$ergs \\citep{BKF04,Frail05}. Such a large true energy, together with the inferred energy per solid angle in the prompt gamma-ray emission and in the afterglow shock at early times imply that the initial jet half-opening angle cannot be too small (typically not much less than a few degrees).  It turns out that a transition from laterally confined to ballistic flow is accompanied by a relatively short phase of acceleration of a different kind \\citep{KVK09,TNM09}.  Such a transition may occur in the collapsar model at the stellar surface.  A sudden loss of lateral pressure support causes a sideways expansion of the jet.  If the jet is highly relativistic at the stellar surface the corresponding increase in the jet opening angle is negligible.  However, a rarefaction wave propagates into the jet and brings it out of lateral balance. The magnetic pressure force accelerates the flow in the lateral direction, which results in a significant increase of the jet Lorentz factor, particularly in the outer layers of the jet. This may alleviate the $\\theta_{\\rm jet}\\Gamma\\simeq 1$ problem of the magnetic model. However, as soon as the rarefaction crosses the jet it is well in the ballistic regime and the acceleration is over.\\footnote{This is in contrast with the highly robust mechanism of thermal acceleration, where for an adiabatic index of $\\gamma=4/3$ the jet Lorentz factor grows linearly with the jet radius, $\\Gamma\\propto R$, even in the ballistic regime.} Moreover, it does not ensure full conversion of electromagnetic to kinetic energy. Should, it happen a bit too soon and the jet remains Poynting-dominated.  Even under the best of circumstances the resultant jet magnetization is still close to $\\sigma\\simeq 1$, which is too high for effective shock dissipation \\citep{Leis05,MGA09,MA10}.  Given the problems with this basic case, other ideas have been put forward. The most radical idea is to assume that relativistic astrophysical jets do not become kinetic energy dominated but remain Poynting dominated on all scales and that the observed emission comes not from shocks but from magnetic dissipation cites \\citep{B02,Lt06}. In the context of the present work this may potentially serve as an alternative to internal shocks in cases where for some reason the magnetization remains high at large distances from the source.  Others propose various ways of increasing the efficiency of magnetic acceleration compared to the basic model, e.g., via allowing non-axisymmetric instabilities and randomization of magnetic field \\citep{HB00}. In fact, the magnetic dissipation may also help the transition from Poynting dominated to kinetic energy dominated states \\citep{Dre02,DS02}.  In this work we focus on the acceleration of an impulsive (strongly time-dependent) highly magnetized relativistic outflow, which has received relatively little attention so far. \\citet{C95} was first to consider the non-relativistic case of impulsive magnetic acceleration and dubbed it an ``astrophysical plasma gun''. The relativistic version presents a number of qualitatively different properties. In \\S~\\ref{sec:test_case} we present a detailed study of a simplified test case featuring a cold and initially highly magnetized ($\\sigma_0\\gg 1$) one dimensional finite shell (of initial width $l_0$) initially at rest (at $t = 0$), whose back end leans against a ``wall'' and with vacuum in front of it. The initial evolution (\\S~\\ref{sec:ss-phase} and Appendix~\\ref{app:self-sim}) is described by a self-similar rarefaction wave traveling toward the wall and accelerating the Poynting-dominated flow away from the wall. At the end of this phase, at time $t_0\\approx l_0/c$ when the rarefaction wave reaches the wall, the mean Lorentz factor of the flow is $\\mean{\\Gamma} \\sim \\sigma_0^{1/3}$.  Soon after $t_0$ the shell separates from the wall and moves away from it (\\S~\\ref{sec:ref-phase}). The shell continues to accelerate and keeps an almost constant width of $\\sim 2l_0$. Using both numerical (\\S~\\ref{sec:num}) and analytical (\\S~\\ref{sec:ref-phase},~\\S~\\ref{sec:after_separation} and Appendixes~\\ref{app:an-int}, \\ref{sec:acc2}) methods, we find that during the second phase the mean Lorentz factor grows as $\\mean{\\Gamma} \\sim (\\sigma_0 t/t_0)^{1/3}\\propto t^{1/3}$.  This phase ends at time $t_c=t_0\\sigma_0^2$, when the acceleration slows down and the shell starts coasting. At this point $\\mean{\\Gamma} \\sim \\sigma_0$ and $\\sigma \\sim 1$.  In \\S~\\ref{sec:back_envelope} we present crude but simple derivations of the main results of \\S~\\ref{sec:test_case} that allow us to understand the underlying physics and show that the results are robust -- not very sensitive to the exact initial configuration. The analysis of the coasting phase (\\S~\\ref{sec:coast+summary}) shows that at $t>t_c$ the shell width increases as $\\Delta \\sim 2l_0t/t_c \\propto t$ while its magnetization decreases as $\\sigma \\sim t_c/t \\propto t^{-1}$, resulting in a kinetic energy-dominated flow.  In \\S~\\ref{sec:BM-effect} we address the apparent paradox of self-acceleration -- how can the shell keep accelerating after it separates from the wall? We analyze a variation of our simple test case in which the wall is removed when the rarefaction wave reaches it (at $t_0$). At subsequent times there are no external forces on the system, implying that the center of momentum (CM) velocity or Lorentz factor ($\\Gamma_{\\rm CM}$) remain constant and there is no global acceleration at $t>t_0$ in this strict sense. Nevertheless, even though we find that $\\Gamma_{\\rm CM}\\sim\\sigma_0^{1/2}$ remains constant, the more relevant astrophysical quantity is the mean value of $\\Gamma$ weighted over the energy in the lab frame, $\\langle\\Gamma\\rangle_E$, and it indeed increases as $\\langle\\Gamma\\rangle_E \\sim (\\sigma_0t/t_0)^{1/3}$ at $t_0<t<t_c$. In \\S~\\ref{sec:dis} we discuss the connection between our test case and relativistic astrophysical flows and study the possible implications of our impulsive acceleration mechanism for the dynamics of GRB and AGN jets. We also briefly address the interaction of the magnetized flow with the external medium for GRBs. Our main results and conclusions are presented in \\S~\\ref{sec:conclusions}.  Soon after the first version of our paper had appeared on the electronic archive (http://arxiv.org/archive/astro-ph/), an independent study of impulsive magnetic acceleration was published there as well \\citep{Lt10a,Lt10b,LL10}, indicating growing interest in this mechanism.  Where the covered topics overlap, the results of both studies agree very well. As to the differences, their study focuses on the initial phase of fast acceleration (at $t<t_0$) and shock formation (when instead of pure vacuum the shell expands into a rarefied plasma), whereas the main subject of our paper is the operation of the impulsive acceleration mechanism after the shell separates from the ``wall'' (at $t>t_0$). ",
        "conclusions": "\\label{sec:conclusions}  In this paper we investigated the properties of magnetic acceleration of relativistic impulsive flows. As a first step, we focused on a relatively simple test case where a uniform cold and highly magnetized ($\\sigma_0\\gg 1$) shell of initial width $l_0$, whose back end leans against a conducting ``wall'' and whose head faces vacuum.  The evolution of the flow that develops in this test case splits into three distinct phases.  The first phase can be described as a formation of a plasma pulse (or a moving shell). During this phase, which lasts for the time $\\sim t_0\\equiv l_0/c_{\\rm ms,0} \\approx l_0/c$, a self-similar rarefaction wave develops at the interface with vacuum and travels towards the wall.  At the end of this phase, the mean Lorentz factor of the outflow is only $\\mean{\\Gamma}\\sim\\sigma_0^{1/3}$ and, apart from the very thin layer at the vacuum interface, the shell of plasma is still highly magnetized, with a mean magnetization parameter of $\\mean{\\sigma} \\sim \\sigma_0^{2/3}$.  The first phase ends when the rarefaction wave reaches the wall. At this point a secondary rarefaction wave forms that propagates from the wall into the back of the shell and decelerates the material that passes through it so that the shell quickly separates from the wall and moves away from it.  During this second phase, the center of momentum Lorentz factor of the shell remains fairly constant ($\\Gamma_{\\rm CM} \\sim \\sigma_0^{1/2}$). However, the leading part of the plasma shell, ahead of the secondary rarefaction, continues to accelerate at the same rate as in the self-similar solution.  It contains most of the shell energy and its mean Lorentz factor grows as $\\mean{\\Gamma}\\propto t^{1/3}$.  At the end of the second phase, which lasts up to $\\sim t_c \\equiv \\sigma_0^2 t_0$, the magnetization of the shell drops down to $\\sigma \\sim 1$, one half of the electromagnetic energy is converted into the bulk motion kinetic energy of the plasma, and the growth of the mean Lorentz factor begins to saturate at $\\mean{\\Gamma} \\sim \\sigma_0$. Thus, the flow enters a phase of coasting.  During the coasting phase the pulse width grows faster, approaching $l \\propto t$. The decrease of the magnetization parameter also accelerates, approaching $\\sigma\\propto t^{-1}$, and the pulse soon becomes kinetic-energy dominated. This property of impulsive magnetic acceleration is most valuable in astrophysical context as the efficiency of relativistic MHD shock dissipation decreases dramatically with magnetization. In contrast to an impulsive flow, a steady-state magnetized jet either remains highly magnetized ($\\sigma \\gg 1$) all the way, or approaches $\\sigma \\approx 1$, depending on the efficiency of external collimation.  This implies at best only modest shock dissipation efficiency.  In some cases of truly explosive phenomena, such as magnetar flares, our impulsive magnetic acceleration mechanism can be solely responsible for the flow acceleration. In most other cases, such as GRB and AGN jets, strong variability of their central engines is not expected on time scales below the viscous time-scale of the inner accretion disc around a black hole, which powers relativistic outflow. This gives plenty of time to establish a quasi-steady super-fast-magnetosonic flow near the source where it is accelerated via the collimation mechanism. The observed strong collimation of these jets supports our conclusion that the collimation mechanism plays a part in their acceleration.  The impulsive acceleration mechanism comes in force further out, where an individual ejecta element starts being accelerated after the head rarefaction crosses it and creates conditions similar to those of our test case flow in phases two and three. The mean Lorentz factor of the shell, however, starts increasing significantly above the value achieved by the quasi-steady collimation acceleration only when the tail rarefaction wave crosses about half of the shell. Provided the central engine variability is sufficiently strong, so that the flow can be described as individual ejecta shells separated by long gaps, the impulsive acceleration mechanism can complete the acceleration process and produce kinetic energy dominated relativistic flows on astrophysically relevant distances from the central engine.  For short GRBs this may still work well even if the ejecta effectively form a single uniform shell.  Our analysis of GRBs show that a combination of the collimation and impulsive mechanisms can accelerate GRB jets up to $\\Gamma \\gtrsim 10^3$, as has been inferred recently for several bright GRBs detected by the Fermi Large Area Telescope, for both long~\\citep{080916C,090902B} and short~\\citep{090510-phys} duration GRBs.\\footnote{We do note, however, that these lower limits on $\\Gamma$ from pair opacity are somewhat model dependent and a fully self consistent calculation appropriate for an internal shock origin of the gamma-ray emission gives limits that are a factor of $\\sim 3$ lower \\citep{Granot08,GRB090926A}, $\\Gamma \\gtrsim 10^{2.5}$, which are significantly easier to satisfy.}  Moreover, their jets can become kinetic energy dominated before the interaction with the interstellar or stellar wind gas begin to decelerate the ejecta at $R_{\\rm dec}\\sim 10^{16}-10^{17}\\;$cm.  The dissipation at internal shocks can become efficient on scales $R\\gtrsim R_c \\approx 10^{13}(\\sigma_0/300)^2(t_v/4\\,{\\rm ms})\\;$cm.  The large variation of Lorentz factor at the coasting phase, $\\Delta\\Gamma\\sim\\Gamma$, insures that the internal shock will be strong and can dissipate and radiate of the order of $\\sim 10\\%$ or so of the flow kinetic energy, leading to a possibility of strong prompt emission.  The AGN jets are likely to be accelerated up to their observed Lorentz factors already during the collimation acceleration phase. However, the impulsive acceleration phase remains important, providing effective conversion of remaining electromagnetic energy and producing kinetic energy dominated flows. Our estimates show that efficient shock dissipation region, analogous to the prompt emission region of GRBs, is located around $\\sim 1-10\\;$pc, where VLBI observations reveal the presence of super-luminal ``blobs''."
    },
    "1004/1004.5267_arXiv.txt": {
        "abstract": "X-ray Flash (XRF) 100316D, a nearby super-long under-luminous burst with a peak energy $E_{\\rm p}\\sim 20$ keV, was detected by {\\it Swift} and was found to be associated with an energetic supernova SN 2010bh. Both the spectral and the temporal behavior are rather similar to XRF 060218, except that the latter was associated with a ``less energetic\" SN 2006aj, and had a prominent soft thermal emission component in the spectrum. We analyze the spectral and temporal properties of this burst, and interpret the prompt gamma-ray emission and the early X-ray plateau emission as synchrotron emission from a dissipating Poynting flux dominated outflow, probably powered by a magnetar with a spin period of $P \\sim 10$ ms and the polar cap magnetic field $B_{\\rm p} \\sim 3\\times 10^{15}$ G. The energetic supernova SN 2010bh associated with this burst is however difficult to interpret within the slow magnetar model, and we suspect that the nascent magnetar may spin much faster with an initial rotation period $\\sim 1$ ms. It suggests a delay between the core collapse and the emergence of the relativistic magnetar wind from the star. The diverse behaviors of low-luminosity GRBs and their associated SNe may be understood within a unified picture that invokes different initial powers of the central engine and different delay times between the core collapse and the emergence of the relativistic jet from the star. ",
        "introduction": "After four years of waiting since the detection of X-Ray Flash (XRF) 060218/SN 2006aj, another pair of low-luminosity (LL) GRB - supernova (SN) association, XRF 100316D/SN 2010bh at redshift $z=0.059$ \\citep{Vergani10}, was captured by {\\em Swift} \\citep{Gehrels04} on March 16, 2010 \\citep{Stamatikos10,Wiersema10,Chornock10,Rau10}, with a detection rate fully consistent with the population studies of these nearby LL-GRB events \\citep{Coward05,Soderberg06,Liang07,Guetta07}. Before this event, four pairs of nearby ($z<0.2$) secure GRB(XRF)-SN associations have been identified. These are GRB 980425/SN 1998bw at $z=0.0085$ \\citep[e.g.,][]{Galama98}, GRB 030329/SN 2003dh at $z=0.168$ \\citep[e.g.,][]{Hjorth03}, GRB 031203/SN 2003lw at $z=0.105$ \\citep[e.g.,][]{Malesani04}, and XRF 060218/SN 2006aj at $z=0.0331$ \\citep[e.g.,][]{Campana06}. The nature of the GRB/SN connection and the interplay between the GRB and the SN components are still poorly understood. In this paper, we analyze and interpret the {\\em Swift} BAT and XRT data of XRF 100316D, paying special attention to the similarities and differences between the XRF 100316D/SN 2010bh and XRF 060218/SN 2006aj. ",
        "conclusions": "So far the nearby supernova-associated GRBs, except GRB 030329, are found to be intrinsically under-luminous. They share the similarities such as low isotropic energies and smooth light curves, but differ in some aspects. For example, GRB 980425 and GRB 031203 have shorter durations and higher $E_{\\rm p}$'s than XRF 060218 and XRF 100316D. The underlying physical processes that result in these differences are not well understood. Through supernova modeling, it is found that GRB 980425 and GRB 031203 have a progenitor star massive enough to form a black hole \\citep{Deng05,Mazzali06a}, while XRF 060218 has a less massive progenitor that most plausibly produces a neutron star \\citep{Mazzali06b}. The luminosity and the duration of XRF 100316D are consistent with the radiation from a neutron star with a dipole magnetic field $B_{\\rm p} \\sim 3\\times 10^{15}$ G and a rotation period $P \\sim 10$ ms. This seems to point towards a hypothesis that two types of central engines define the apparent dichotomy of the SN-associated LL-GRBs, i.e. black hole engines give rise to ``shorter\" and ``harder\" GRBs such as GRB 980425 and GRB 031203, while magnetar engines give rise to very long and soft XRFs such as XRF 060218 and XRF 100316D \\footnote{We caution that such a scenario is not robust, since a magnetar engine may also drive SN 1998bw \\citep{Woosley09}, and a black hole engine may be also able to reproduce the XRF 100316D-like light curves through tuning the parameters of fall-back materials and arguing for a Poynting-flux-dominated outflow from a highly magnetized black hole engine \\citep[e.g.][]{MacFadyen01,ZhangW08}.}.  Adding in SN data makes the scenario more complicated. Although SN 2006aj associated with XRF 060218 does not conflict with a slow magnetar central engine, SN 2010bh associated with XRF 100316D may be too energetic to be interpreted with a slow magnetar central engine. If it is confirmed that the kinetic energy of SN 2010bh is in excess of $10^{52}$ erg, neither the neutrino energy nor the magnetar spin energy ($\\sim 10^{50}$ erg) are adequate to power the SN. A salient feature of the dipole spindown formula (Eq.[\\ref{eq:E_inj}]) is that if one shifts the time zero point (e.g. $t'=t - t_0$), the spindown law still applies, with the initial angular frequency re-defined as $\\Omega'=\\Omega(t'=0)$, and the characteristic spindown time scale re-defined as $\\tau'_0=1.6\\times 10^4 B_{\\rm p,14}^{-2} {\\Omega'_4}^{-2} I_{45} R_{s,6}^{-6}$. This suggests that the observed plateau feature can be still interpreted if the initial period is much shorter than 10 ms, say, $P_0 \\sim 1$ ms, if the time zero point is much earlier than $T_{\\rm trig}-500$ s (i.e. $t=0$). This is because a power-law decay light curve may show an artificial plateau in the log-log space, if the zero time is mis-placed to a later epoch \\citep{Yamazaki09,Liang09}. Within such a scenario, a nascent magnetar was born with an initial period $P_0 \\sim 1$ ms at $t \\sim -5\\times 10^3$ s. Its intial dipole radiation was trapped by the envelope of the progenitor and could not escape. This spindown energy gives enough impetus to explode the star and power the energetic SN 2010bh. After a significant delay ($\\sim 5\\times 10^3$ s to spin down from 1 ms to 10 ms for $B_{\\rm p} \\sim 3\\times 10^{15}$ G), the magnetar wind finally managed to escape as a relativistic Poynting-flux-dominated outflow. An observer noticed the jet emission only around $t=0$. The above argument also applies to the model of fallback accretion onto a nascent black hole.  With such a hypothesis, one may envision a unified picture to understand the diversity of GRB/SN associations, by invoking a variety of initial powers and the delay times between the core collapse and the emergence of the relativistic jet from the star. The speculation is the following: \\begin{itemize} \\item To produce an energetic SN/luminous GRB (e.g. GRB 030329/SN 2003dh), the central engine is powerful (a black hole with an accretion disk or a rapidly spinning magnetar) and the relativistic outflow can break out the progenitor soon enough when the engine is still working effectively. \\item To produce an energetic SN / underluminous GRB (e.g. GRB 980425/SN 1998bw, GRB 031203/SN 2003lw, and XRF 100316D/SN 2010bh), the central engine is initially powerful, but it takes time for the relativistic wind to emerge from the star. As it breaks out the star, the central engine already fades down with a decreased power. The longer, softer XRFs are probably powered by a magnetar, while the shorter, harder GRBs are probably powered by a black hole. \\item To produce a less-energetic SN / underluminous GRB (e.g. XRF 060218/SN 2006aj), the central engine is a slow magnetar with an initial rotation energy less than $10^{51}$ ergs. The emergence of the relativistic outflow can be prompt or somewhat (but not significantly) delayed. \\end{itemize}  Finally, a straightforward expectation from the speculation that XRF 100316D outflow is Poynting-flux-dominated is that the prompt emission should be linearly polarized \\citep[e.g.,][]{Fan05}. The polarimetry measurements of events such as XRF 100316D and XRF 060218 would provide a criterion to differentiate this model from the shock breakout model, which does not predict a strong polarization signal."
    },
    "1004/1004.2671_arXiv.txt": {
        "abstract": "{ We estimated black hole masses and Eddington ratios ($L/L_\\mathrm{Edd}$) for a well defined sample of local ($z < 0.3$) broad line AGN from the Hamburg/ESO Survey (HES), based on the H$\\beta$ line and standard recipes assuming virial equilibrium for the broad line region. The sample represents the low-redshift AGN population over a wide range of luminosities, from Seyfert 1 galaxies to luminous quasars.  From the distribution of black hole masses we derived the active black hole mass function (BHMF) and the Eddington ratio distribution function (ERDF) in the local universe, exploiting the fact that the HES has a well-defined selection function. While the directly determined ERDF turns over around $L/L_\\mathrm{Edd} \\sim 0.1$, similar to what has been seen in previous analyses, we argue that this is an artefact of the sample selection. We employed a maximum likelihood approach to estimate the \\emph{intrinsic} distribution functions of black hole masses and Eddington ratios simultaneously in an unbiased way, taking the sample selection function fully into account. The resulting ERDF is well described by a Schechter function, with evidence for a steady increase towards lower Eddington ratios, qualitatively similar to what has been found for type~2 AGN from the SDSS.  Comparing our best-fit active BHMF with the mass function of inactive black holes we obtained an estimate of the fraction of active black holes, i.e.\\ an estimate of the AGN duty cycle. The active fraction decreases strongly with increasing black hole mass. A comparison with the BHMF at higher redshifts also indicates that, at the high mass end, black holes are now in a less active stage than at earlier cosmic epochs. Our results support the notion of anti-hierarchical growth of black holes, and are consistent with a picture where the most massive black holes grew at early cosmic times, whereas at present mainly smaller mass black holes accrete at a significant rate. } ",
        "introduction": "The observed relations between the black hole mass and the properties of the spheroidal galaxy component imply a close connection between the growth of supermassive black holes (SMBH) and the evolution of their host galaxies. For local galaxies a strong correlation between the mass of the SMBH and the luminosity or mass of the bulge component \\citep{Magorrian:1998,Marconi:2003,Haering:2004}, as well as with the stellar velocity dispersion \\citep[e.g.][]{Ferrarese:2000,Gebhardt:2000, Tremaine:2002,Gultekin:2009} have been established. Semi-analytical and numerical simulations also show the importance of black hole activity and their corresponding SMBH feedback for galaxy evolution \\citep[e.g.][]{DiMatteo:2005,Springel:2005,Croton:2006,Cattaneo:2006,Khalatyan:2007,Booth:2009}. It became clear that the central SMBH of a galaxy and especially its growth is an important ingredient for our understanding of galaxy formation and evolution.  Therefore a complete census of the black hole population and its properties is required. Active black holes that will be observable as AGN are particularly important to study black hole growth. A useful tool to study the AGN population is the luminosity function (AGNLF). The observed evolution of the AGNLF has been used to gain insight into the growth history of black holes  \\citep[e.g.][]{Soltan:1982,Yu:2002,Marconi:2004,Merloni:2004,Shankar:2007}, and it became clear that most of the accretion occurs during bright QSO phases. But, using the AGNLF alone usually requires some additional assumptions, e.g. for the mean accretion rate, and thus is affected by uncertainties and degeneracies. Disentangling the AGNLF into the underlying distribution functions, namely the active black hole mass function (BHMF) and the distribution function of Eddington ratios (ERDF), is able to provide additional essential constraints on the growth of SMBHs.  To understand the influence of black hole growth on galaxy evolution over cosmic time, first the properties of growing black holes in the local universe have to be well understood. Thus, it is important to derive black hole masses and accretion rates for large, well defined samples of AGN. However, measuring black hole masses is much more  difficult than measuring luminosities. Black hole masses for large samples of AGN can not be measured directly, but only estimated, using locally established scaling relations.  The best method to estimate \\mbh for type~1 AGN is reverberation mapping of the broad line region \\citep{Blandford:1982,Peterson:1993}. Assuming virial equilibrium black hole masses can be estimated by $M_{\\mathrm{BH}} = f R_{\\mathrm{BLR}} \\Delta V^2 / G$, where $R_{\\mathrm{BLR}}$ is the size of the broad line region (BLR), $\\Delta V$ is the broad line width in km/s and $f$ is a scaling factor of order unity, which depends on the structure, kinematics and orientation of the BLR. Although the physics of the BLR is still not well understood and thus a source of uncertainty \\citep[e.g.][]{Krolik:2001}, the validity of the virial assumption has been shown by the measurement of time lags and line widths for different broad lines in the same spectrum \\citep{Peterson:2000,Onken:2002,Kollatschny:2003}.  However, reverberation mapping requires extensive and meticulous observations and thus is not appropriate for large samples. Fortunately, a scaling relationship has been established between $R_{\\mathrm{BLR}}$ and continuum luminosity of the AGN, $R_{\\mathrm{BLR}} \\varpropto L^\\gamma$  \\citep{Kaspi:2000,Kaspi:2005,Bentz:2006}. Thus it became possible to estimate \\mbh from  single-epoch spectra for large samples, and has been used extensively in the previous years for large AGN samples \\citep[e.g.][]{McLure:2004,Vestergaard:2004,Kollmeier:2006,Netzer:2007a,Shen:2007,Fine:2008,Gavignaud:2008, Trump:2009}.  For the measurement of the line width, different measures are commonly used, and it is unclear if one is superior to the others for estimating black hole masses. Most commonly used is the FWHM, but it has been suggested that the line dispersion $\\sigma_{\\mathrm{line}}$, i.e. the second central moment of the line profile, is a better measure of the line width \\citep{Peterson:2004,Collin:2006}. The line dispersion is more sensitive to the wings of a line and less to the core, whereas for the FWHM the opposite is the case. An additional measure of line width used is the inter-percentile value \\citep[IPV,][]{Fine:2008}.  The application of the virial method to large AGN samples allowed the estimation of the active BHMF \\citep{McLure:2004,Shen:2007,Greene:2007a,Vestergaard:2008,Kelly:2008b,Vestergaard:2009}. A dataset that is perfectly suited to study especially low redshift AGN is provided by the Hamburg/ESO Survey (HES). In this paper we use a local AGN sample, drawn from the HES, to estimate their black hole masses and Eddington ratios, and construct the active black hole mass function as well as the distribution function of Eddington ratios.  We first present our data and our treatment of the spectra. We estimate black hole masses and Eddington ratios from the spectra using the virial method. Next, we determine the active BHMF, taking care to account for sample selection effects, inducing a bias on the BHMF. Thereby, we not only constrain the local active BHMF but also put constraints on the intrinsic underlying distribution function of Eddington ratios. Finally, we discuss our results in the context of the local quiescent BHMF as well as that of other surveys.  Thoughout this paper we assume a Hubble constant of $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$ and cosmological density parameters $\\Omega_\\mathrm{m} = 0.3$ and $\\Omega_\\Lambda = 0.7$. ",
        "conclusions": "We have presented a study of the low-redshift active black hole population, residing in broad-line active galactic nuclei. We estimated black hole masses and Eddington ratios, and from it estimated the local active black hole mass function and the Eddington ratio distribution function. Our sample was drawn from the Hamburg/ESO Survey and contains 329 quasars and Seyfert~1 galaxies with $z<0.3$, selected from surveying almost 7000~deg$^2$ in the southern sky.  We estimated black hole masses from single-epoch spectra, measuring the line dispersion of the broad H$\\beta$ line and the continuum luminosity at 5100\\,\\AA{}\\, $L_{5100}$, using the common virial method. We took care to avoid contamination of the line measurement by neighbouring emission lines and roughly estimated the degree of host galaxy contribution to $L_{5100}$. This has been found to be negligible for the most luminous AGN and not dominant even at the low luminosity end of our sample. We applied a rough statistical correction to the continuum luminosities to take into account the host contribution. The bolometric luminosity and thus the Eddington ratio $\\er$, has been estimated from $L_{5100}$.  The observed black hole masses cover a range $10^6 - 2\\cdot10^9 \\,M_\\odot$ and the Eddington ratio is roughly confined between $0.01 -1$. The observed distributions of these quantities are understood by the underlying distribution functions of black hole mass and Eddington ratio, in combination with the survey selection function, as we explicitly demonstrated by Monte Carlo simulations.  We made an attempt to determine these two distribution functions in an unbiased way. First of all, when we want to determine the \\textit{active} BHMF, we have to make clear what we mean by an \\textit{active} black hole, due to the wide distribution of accretion rates. We used a lower Eddington ratio cut of $\\log \\er = -2$, in agreement with the observed range of Eddington ratios. Using a different cut for $\\er$ will preserve the shape of the BHMF, but change their normalisation, due to our assumption of an uncorrelated BHMF and ERDF. This is already shown in the left panel of Fig.\\ref{fig:acf}. The normalisation and therefore the space density clearly depend of the chosen definition of an active black hole.  Next we have to be aware that our sample is selected on AGN luminosity, not on black hole mass. Therefore, we have to make sure that we properly account for active black holes below the flux limit of the survey. We presented a method that determines the active BHMF as well as the ERDF at the same time, by a maximum likelihood fit. Here, the bivariate probability distribution of black hole mass and Eddington ratio is fitted to the observations. This probability distribution is given by an assumed BHMF, ERDF and the selection function of the survey. We also corrected for evolution within our redshift range, transforming the distribution functions to $z=0$. This maximum likelihood method also ensures the consistency of the derived BHMF and ERDF with the AGN luminosity function. We were able to put tight constraints on both the active black hole mass function and the Eddington ratio distribution function.  The Eddington ratio distribution function is well described by a Schechter function with low $\\er$ slope $\\alpha_{\\er} \\approx -1.9$. The data are consistent with no decrease of the ERDF at low $\\er$, within the constrained range. Using a log-normal distribution, we found a maximum at $\\log \\er = -1.8$, what can be taken as an upper limit for a potential turnover in the ERDF. Our results clearly show a wide distribution of Eddington ratios, in contrast to a single value or to a narrow log-normal distribution, which is based on the observed distribution, without accounting for the underlying selection effects. While we also observe a narrow log-normal distribution of Eddington ratios, this is in agreement with the constrained Schechter function or wide log-normal distribution for the Eddington ratio distribution function, when survey selection effects are properly accounted for, because low-$\\er$ objects will be systematically missed in flux limited samples.  The active BHMF is well described by different analytic models. In general, it strongly decreases at the high mass end and follows a power law at the low mass end with slope of $\\alpha \\approx -2$. A good fit to the data is achieved by a function similar to a Schechter function, but modified by an extra parameter that determines the steepness of the high mass decrease. We found no evidence for a decrease of the BHMF toward low mass, as indicated by \\citet{Greene:2007a} for $\\mbh \\lesssim 10^{6.5} M_\\odot$. However, our sample is not very sensitive in this low mass range.  We compared our local active BHMF with the local quiescent BHMF from \\citet{Marconi:2004}, determining the active fraction, or duty cycle, of local black holes. This active fraction is decreasing with increasing black hole mass, consistent with a power law with slope $\\sim-0.86$. Thus, the most massive black holes in the present universe are less active than their lower mass companions. At the highest \\mbh only $10^{-4}$ of all black holes are currently in an active stage, i.e. accreting above 0.01 of the Eddington rate. This supports the general picture of anti-hierarchical growth of black holes. This mass dependence of the active fraction indicates that our assumption of an uncorrelated BHMF and ERDF cannot be sustained up to low values of $\\er$ and thus we caution to extrapolate the distribution functions into the low $\\er$ regime. Investigating a mass dependence of the ERDF would especially require a wider luminosity coverage of the sample.  By comparing our low~$z$ BHMF with the BHMF of a higher $z$-bin, presented by \\citet{Vestergaard:2008} and \\citet{Vestergaard:2009}, we found an indication that the most massive black holes are currently in a less active stage than at earlier cosmic times, also in general agreement with anti-hierarchical  black hole growth.  Recently, \\citet{Marconi:2008} proposed a modified method to estimate \\mbh, taking into account the effect of radiation pressure. So far, it is still unknown if radiation pressure has an important effect on the BLR or not \\citep[see e.g.][]{Netzer:2008}. If we take into account radiation pressure and apply their \\mbh estimation formula to our sample, the major effect is an increase of \\mbh especially for the low~\\mbh objects. In total, the dispersion of the \\mbh distribution decreases from 0.65~dex to 0.63~dex. In the BHMF the space density at median \\mbh increases, while at high \\mbh the space density slightly decreases. This would strengthen even further the evidence for anti-hierarchical black hole growth. On the other hand it would change our observed \\mbh, and especially our $\\er$, distributions, and thereby our constrained BHMF and the Eddington ratio distribution function.  Our work strengthens the scenario of anti-hierarchical  growth of black holes, also seen in other studies \\citep{Merloni:2004,Heckman:2004,Greene:2007a,Shankar:2007,Vestergaard:2009}, at least at low redshift. The observation of 'cosmic downsizing' in the X-ray luminosity function \\citep[e.g.][]{Ueda:2003,Hasinger:2005}, as well as in the optical, radio and IR luminosity function \\citep[e.g.][]{Hunt:2004,Cirasuolo:2005,Matute:2006,Croom:2009}, i.e. the flattening of the faint end slope of the luminosity function towards higher redshift, is explained by the shift of the typical black hole mass of an active accreting black hole toward lower mass.  The presented local active black hole mass function and Eddington ratio distribution function serve as a local anchor for future studies of both distribution functions. These will provide further information on the cosmic history of growth and activity of supermassive black holes."
    },
    "1004/1004.3261_arXiv.txt": {
        "abstract": "We   present  VLT/FORS2   spectroscopy  and   GROND  optical/near-IR photometry  of   the  afterglow   of  the  bright   {\\it  Fermi}/LAT GRB~090926A. The spectrum  shows prominent Lyman-$\\alpha$ absorption with $N_{\\rm HI} =  10^{21.73\\pm0.07}$\\,cm$^{-2}$ and a multitude of metal lines  at a common  redshift of $z=2.1062\\pm0.0004$,  which we associate with the redshift of the GRB. The metallicity derived from SII is $\\log (Z/Z_\\odot)\\approx -1.9$, one of the lowest values ever found in  a GRB Damped  Lyman-$\\alpha$ (DLA) system.   This confirms that  the spread  of metallicity  in GRB-DLAs  at $z\\approx2$  is at least  two  orders of  magnitude.   We  argue  that this  spread  in metallicity does  not require a  similar range in abundances  of the GRB progenitors, since the neutral interstellar medium probed by the DLA is expected  to be at a significant  distance from the explosion site.   The hydrogen  column  density derived  from {\\it  Swift/XRT} afterglow  spectrum  (assuming  $\\log (Z/Z_\\odot)\\approx  -1.9$)  is approx.  $\\approx100$  times higher  than the $N_{\\rm  HI}$ obtained from  the Lyman-alpha  absorptions.   This suggests  either a  large column  density  of ionized  gas  or  a  higher metallicity  of  the circum-burst medium compared to the  gas traced by the DLA.  We also discuss  the afterglow  light curve  evolution and  energetics.  The absence of a  clear jet-break like steeping until  at least 21\\,days post-burst   suggests  a   beaming  corrected   energy   release  of $E_{\\gamma}>3.5\\times10^{52}$\\,erg, indicating that GRB~090926A  may have been one of the most energetic bursts ever detected. \\\\ ",
        "introduction": "\\label{sec:intro}  For the brief moments of  their existence, Gamma-ray Bursts (GRBs) and their  X-ray/optical counterparts  are  the brightest  beacons in  the Universe.   The afterglow  luminosities  and simple  power law  shaped spectra make  them ideal background lights for  probing the conditions in their  host galaxies and in intervening  systems through absorption line spectroscopy.  Although the  afterglows fade away within hours to days, rapid follow-up observations have been performed for a number of GRBs  and have  revealed tell tale  features of  the circum-stellar medium  around  the  progenitor  and  the  interstellar  medium  (ISM) \\citep[e.g.,][]{Castro:2003aa}.   These   studies  provided  otherwise hidden           details          of           the          kinematics \\citep[e.g.,][]{Klose:2004aa,Fox:2008aa},                     excitation \\citep[e.g.,][]{Vreeswijk:2007aa},     dust     and     gas     content \\citep[e.g.,][]{Savaglio:2004aa,Prochaska:2006aa},   and  the  chemical abundances    \\citep[e.g.,][]{Savaglio:2003aa,Fynbo:2006aa}    in   the star-forming,      low-mass      galaxies      that     host      GRBs \\citep[e.g.,][]{Le-Floch:2003aa,Christensen:2004aa,Rau:2005ab,Savaglio:2009aa}.   In  this  paper we  present  photometric  and spectroscopic  follow-up observations of  the bright \\GRB,  concentrating on the  energetics of the burst and the chemical enrichment traced by the optical afterglow. \\GRB\\    was    discovered   by    the    Gamma-ray   Burst    Monitor \\citep[GBM;][]{Meegan:2009aa} onboard the  {\\it Fermi} Gamma-ray Space Telescope on  2009 September 26  at $T_0$=04:20:27\\,UT and  belongs to the                         long-duration                        class \\citep[T$_{90}=20\\pm2$\\,s,][]{Bissaldi:2009aa}.   Further detetections of   the  prompt   emission   were  reported   from  {\\it   Fermi}/LAT \\citep{Uehara:2009aa,Bissaldi:2009ab},        {\\it        Suzaku}-WAM, \\citep{Noda:2009aa}  and {\\it  Konus}-Wind, \\citep{Golenetskii:2009aa} and  the X-ray  \\citep[{\\it Swift}/XRT,][]{Vetere:2009aa}  and optical ({\\it Swift}/UVOT, Gronwall et al. 2009; Skynet, Haislip et al. 2009a) afterglows were  also detected. A redshift of  $z=2.1062$ was measured with VLT/X-Shooter observations \\citep{Malesani:2009aa}  Throughout  the  paper, we  adopt  concordance $\\Lambda$CDM  cosmology ($\\Omega_{\\rm  M}=0.27$, $\\Omega_{\\rm \\Lambda}=0.73$,  H$_{0}=71$\\ (km s$^{-1}$) Mpc$^{-1}$), and the convention that the flux density of the GRB   afterglow    can   be    described   as   $F_\\nu\u03bd    \u221d   \\propto \\nu^{-\\beta}t^{-\\alpha}$. ",
        "conclusions": "\\GRB\\ was  an event of  two extremes.  It  was likely one of  the most energetic explosions detected so  far and simultaneously showed one of the most  metal poor  GRB-DLA found until  now.  The  large luminosity $E_{\\gamma}>3.5\\times10^{52}$\\,erg,  coupled   with  a  bright  slowly decaying afterglow, allowed a detailed study of the temporal evolution of the optical transient.  It furthermore enabled high signal-to-noise spectroscopy  of the  afterglow as  late  as one  day post-burst  with VLT/FORS2.  With  the spectrum we  confirmed the  redshift of  the host  galaxy of $z=2.1062\\pm0.0004$ and discovered  two intervening absorption systems at $z=1.946\\pm0.001$ and  $z=1.748\\pm0.001$. Furthermore, we derived a neutral    hydrogen    column     density    of    $N_{\\rm    HI}    = 10^{21.73\\pm0.07}$\\,cm$^{-2}$  and a  metallicity of  the  neutral ISM along the line of sight in the host galaxy of $\\log (Z/Z_\\odot)\\approx -1.9$.   We close this paper with one reminder. The DLAs found in GRB afterglow spectra   are  predominantly  probing   the  diverse   conditions  and sub-structures  within  the   host  galaxies.   They  therefore  allow important  insight  into  the  gas  and  metallicity  distribution  in galaxies at  high redshift.  The  evolution of metallicity  and column densities with  redshift is, however,  challenging to access  with the current number of  GRB- DLAs.  A much larger sample  of sight lines is required   to  first   characterize   the  intrinsic   gas-metallicity dispersions,  and before  GRB-DLAs  can be  used  reliably for  cosmic chemical evolution studies."
    },
    "1004/1004.1925_arXiv.txt": {
        "abstract": "We consider the issue of characterizing the coherent large-scale patterns from CMB temperature maps in globally anisotropic cosmologies. The methods we investigate are reasonably general; the particular models we test them on are the homogeneous but anisotropic relativistic cosmologies described by the Bianchi classification. Although the temperature variations produced in these models are not stochastic, they give rise to a ``non--Gaussian'' distribution of temperature fluctuations over the sky that is a partial diagnostic of the model. We explore two methods for quantifying non--Gaussian and/or non-stationary fluctuation fields in order to see how they respond to the Bianchi models. We first investigate the behavior of phase correlations between the spherical harmonic modes of the maps. Then we examine the behavior of the multipole vectors of the temperature distribution which, though defined in harmonic space, can indicate the presence of a preferred direction in real space, i.e. on the 2-sphere. These methods give extremely clear signals of the presence of anisotropy when applied to the models we discuss, suggesting that they have some promise as diagnostics of the presence of global asymmetry in the Universe. ",
        "introduction": "\\label{secIntro} Observations of the Cosmic Microwave Background (CMB) provide some of the most compelling support for the currently favored $\\Lambda$CDM, or \\emph{concordance}, cosmological model. The concordance framework predicts that the CMB should posses temperature fluctuations which are both statistically isotropic (i.e. stationary over the celestial sphere) and Gaussian \\citep{Guth1982,Starobinskij1982,Bardeen1983}. Measurements by the Wilkinson Microwave Anisotropy Probe (WMAP) \\citep{Bennett2003,Hinshaw2009} have undergone extensive statistical analysis, much of which has confirmed the concordance model but with some indications of departures that may be significant; see for example \\cite{Yadav2008}. More specifically, there is some evidence for hemispherical power asymmetry \\citep{Eriksen2004a,Park2004,Eriksen2007,Hoftuft2009,Hansen2009} and also a Cold Spot has been identified \\citep{Vielva2004,Cruz2005}. In other words there is some evidence of an anisotropic universe, i.e. one in which the background cosmology may not be described by the standard Friedman-Robertson-Walker (FRW) metric. Of course the background cosmology for a non-isotropic universe may still be described by the FRW metric, but this would require a non-standard topology which we do not consider in this analysis.  The Bianchi classification provides a complete characterization of all the known homogeneous but anisotropic exact solutions to General Relativity. The classification was first proposed by Bianchi and later applied to General Relativity \\citep{Ellis1969}. Initial studies used the lack of large-scale asymmetry in the CMB temperature to put strong constraints on the possible Bianchi models \\citep{Barrow1985, Bunn1996, Kogut1997}. However, simulations of the CMB from Bianchi universes not only show a preferred direction, but models with negative spatial curvature (such as the types V and VII$_h$) can produce localized features \\citep{Barrow1985}. So more recently attention has shifted to reproducing a Cold Spot such as that claimed to exist in the WMAP data. Initially, Type VII$_h$ was the favored model to best reproduce the anomaly \\citep{Jaffe2005, Jaffe2006a,Jaffe2006b}, and this has subsequently been investigated quite thoroughly \\citep{McEwen1,McEwen2,Pontz1,Pontz2,Sung2010}, although more recent work has also looked at the Bianchi Type V which also produces localized features \\citep{Sung2009}.  The most interesting range of anisotropic structures is produced in Bianchi Types VII$_h$, VII$_0$ and V. These different Bianchi types have the effect of focusing and/or twisting the initial quadrupole over time (see Figure \\ref{figBT}). In this paper we study the behavior of these Bianchi models so as to identify characteristics of the radiation fields they produce and develop methods that can be used to identify more general forms of anisotropy. Understanding the characteristics identified in these particular cases will hopefully help us find better and more systematic ways of constraining the level of global symmetry present in the real Universe. Note we consider just characteristics observable in the CMB temperature; we shall return to a study of the polarization radiation component in later work.  We consider two statistical measures of anisotropy in some detail in this paper. Neither of these is entirely new and both have previously been applied to observed CMB maps. However, the general philosophy behind previous applications of these methods has been simply to look for departures from the (composite) null hypothesis of statistical isotropy and Gaussianity (or more recently they have been developed to look at universes with multiply-connected topologies \\citep{Bielewicz2009}). In other words, they have been used to construct hypothesis tests with the concordance cosmology but their performance has not hitherto been evaluated on models with built-in anisotropy.  For example, if the concordance model is correct, the {\\em phases} of the spherical harmonic coefficients of the CMB should be independently random and uniformly distributed. Recent studies have suggested some deviation from this \\citep{Coles2004,sc2005,dc2005,Chiang2007,ccno7} but it is not clear whether they indicate global anisotropy or departures from Gaussianity, let alone whether these are of cosmic or instrumental origin. Here we examine the use of phase correlations in quantifying the temperature patterns generated in models with known levels of global inhomogeneity.  Multipole vectors were first introduced over a  century ago \\citep{Maxwell1891}. There have since been attempts to understand the multipole vectors in order to explain the CMB anomalies reported at large angular scales \\citep{Katz2004, Schwarz2004, Copi2004, Land2005a, Land2005b, Land2005c, Land2005d, Land2005e, Copi2006, Copi2007} since is not clear how to quantify and verify such properties from the CMB anomalies in spherical harmonics. They have been used in a number of studies to show anomalies, such as alignments of multiples \\citep{Abramo2006} in a similar plane to the axis of evil \\citep{Land2005d,Land2007}. Our aim here is to examine the behavior of the multipole vectors in cases where the form of anisotropy is known {\\em priori} in order to assess their potential to act as more general descriptors.  Two points are worth making before we continue. First,  any realistic cosmology (whether of FRW or Bianchi type) will possess random fluctuations on top of a smooth background. If these fluctuations are stationary Gaussian then they will add correlated ``noise'' to any signal arising from the background model and will thus hamper the performance of any statistical analysis method, especially at smaller angular scales. This Gaussian ``noise'' (which is equivalent to stationary Gaussian fluctuations, and not to be confused with instrumental noise) is completely characterized by second-order statistical quantities (i.e. the power spectrum in harmonic space or the autocorrelation function in pixel space). The statistical descriptors we explore are {\\em independent} of the power-spectrum, so adding Gaussian noise will not produce any systematic response in them.  We also restrict ourselves to looking at just the large-scale features because the patterns in the temperature maps resulting from the Bianchi models is over large scales. Therefore, by looking at large scales only, there is more chance of detecting the anisotropy. However, it goes without saying we are not claiming that these Bianchi models are in themselves complete alternatives to the concordance cosmology. Rather we think of them as representing possible perturbation modes of the FRW background.  The layout of this article is as follows. In Section \\ref{secPixel} we look at pixel  distributions of the CMB maps to show how the statistical anisotropy present in these models produces a form of non-Gaussianity in the pixel distribution over the celestial sphere. We then introduce phase correlations in Section \\ref{secPhase} to provide characterization of the anisotropy displayed by the models. In Section \\ref{secMultipole} we look at the behavior of the multipole vectors as characteristics of the anisotropy of the same maps. Finally, Section \\ref{secConclu} summarizes the conclusions.  \\begin{figure} \\begin{centering} \\includegraphics[width=58mm]{MapVz000.ps} \\includegraphics[width=58mm]{MapVII0z000.ps} \\includegraphics[width=58mm]{MapVIIhz000.ps}. \\caption{Simulated maps of the the CMB temperature, at redshift z = 0, using Bianchi type cosmologies. From left to right the Bianchi types are: V, VII$_0$ and VII$_h$. The colour scale is marked in milliKelvin. All the maps started as a quadruple at z = 500. The Bianchi V map shows a focused feature, the Bianchi VII$_0$ map has a twisted feature and the Bianchi VII$_h$ map has both focusing and twisting in the resulting temperature pattern.} \\label{figBT} \\end{centering} \\end{figure} ",
        "conclusions": "\\label{secConclu} The aim of this article was to explore some simple ways of characterizing the large-scale temperature patterns in CMB maps generated in anisotropic Bianchi type V, VII$_h$ and VII$_0$ universes. The ultimate purpose of investigating this behavior is to find ways of quantifying the global properties of the pattern produced in order to isolate the effect of anisotropy from that of non-Gaussianity. We repeat that when we talk about non-Gaussianity here is not related to a stochastic field; there are no fluctuations in the Bianchi maps.  We first discussed perhaps the simplest and perhaps the most obvious possible  descriptive statistics, the histogram of the pixel values, primarily with the aim of demonstrating how non-Gaussianity of a sort can arise from asymmetry. We evaluated the pixel distribution functions for each of the maps and compared them to results expected in a universe consistent with the concordance model. The type VII$_0$ maps show the strongest deviation from the null hypothesis; but types V and VII$_h$ behaved in a similar fashion to each other, and closer to that of the null hypothesis. The reason these two gave lesser indications of the presence of anomalies was because the focussing effect produces a pattern that covers only a smaller part of the celestial sphere, which tends to get lost when averaged over the whole sky. This method is therefore useful to characterize coherent signals extended over a large region, such as a spiral pattern, but not if they are concentrated.  Phase analysis is a relatively new technique, and has consequently not been used to quantify many alternative situations to the concordance model. The phases of the spherical harmonic coefficients provide a generic way of looking at correlations in harmonic space that could arise from non-stationarity or non-Gaussianity. While this is a potential strength of the approach - while phase correlations will not just be useful for identifying anisotropies specific to the Bianchi models, but in theory any isotropy introduced to the CMB - it could also prove a weakness, in that more general methods may lack the power to discriminate very specific models.  The phase correlations identified in our Bianchi maps using this technique were much stronger than we at first expected; given the generic nature of the metric it was not expected to yield good results. In addition to this, the strong correlations were found to be robust to both rotation and moderate noise. Significant correlations in both twisted and focusing features were also identified. However using the same methods on the WMAP 5 year data shows little evidence of non-Gaussianity. Given that the diagnostics are identified in harmonic space, it is difficult to say whether any of the anomalies identified this way are down to isotropy or homogeneity.  The analysis of multipole vectors is also a relatively new technique.  It has been used to identify non-Gaussianities in the WMAP data, and has been particularly successful in identifying anisotropies (i.e. asymmetries and/or preferred directions). The multipole vectors are calculated from spherical harmonic coefficients which, as we have already shown, themselves provide a very effective way of identifying correlations in Bianchi (and presumably other anisotropic) patterns. The multipole vectors must include at least some of the information needed to describe these mode correlations. The advantage of multipole vectors over the spherical harmonics themselves, however, is that they give results in real (i.e. pixel) space which is much more informative to the user. The results when applied to the Bianchi maps show very strong correlations between the directions of the multipole vectors for low $\\l$, often with them entirely overlapping, and hence showing preferred directions. Since these vectors would not be aligned in the case of a stationary stochastic field over the sky, these results demonstrate that they are sensitive to departures from the standard cosmological model.  It remains the case that the standard cosmological model is a  good fit to a huge range of observational data. Nevertheless, it is important that tools are developed that are sufficiently sensitive to hunt efficiently for possible anomalies in the next generation of observations. There are many ways that the CMB temperature pattern could be anomalous other than through the presence of Bianchi perturbation modes such as those we have studied here. Just as there are many ways a distribution can be non-Gaussian, so are there also many ways a fluctuation field can be non-stationary. Testing for departures from the standard model will require not one but a battery of statistical techniques each sensitive to particular aspects of the distribution.  This has been a very preliminary analysis, aimed at  establishing whether the diagnostics described in this paper are {\\em in principle} capable of uncovering evidence of underlying anomalies in CMB data. Of course these patterns represent somewhat extreme departures from the standard framework so it is no real surprise that they register strongly in the descriptors used. However, in all cases our analysis has involved only a relatively small number of quantities, so the fact that we see quantifiable effects emerging is very encouraging."
    },
    "1004/1004.5501_arXiv.txt": {
        "abstract": "Using observations from the \\textit{Chandra} X--ray Observatory and Giant Metrewave Radio Telescope, we examine the interaction between the intracluster medium and central radio source in the poor cluster AWM~4. In the \\textit{Chandra} observation a small cool core or galactic corona is resolved coincident with the radio core. This corona is capable of fuelling the active nucleus, but must be inefficiently heated by jet interactions or conduction, possibly precluding a feedback relationship between the radio source and cluster.  A lack of clearly detected X--ray cavities suggests that the radio lobes are only partially filled by relativistic plasma. We estimate a filling factor of $\\phi$=0.21 (3$\\sigma$ upper limit $\\phi<0.42$) for the better constrained east lobe. We consider the particle population in the jets and lobes, and find that the standard equipartition assumptions predict pressures and ages which agree poorly with X--ray estimates. Including an electron population extending to low Lorentz factors either reduces ($\\gamma_{min}=100$) or removes ($\\gamma_{min}=10$) the pressure imbalance between the lobes and their environment.  Pressure balance can also be achieved by entrainment of thermal gas, probably in the first few kiloparsecs of the radio jets. We estimate the mechanical power output of the radio galaxy, and find it to be marginally capable of balancing radiative cooling. ",
        "introduction": "\\label{sec:intro}  X--ray observations of clusters and groups of galaxies over the last decade have led to a significant revision of our models of the intergalactic medium in these systems. The \\chandra\\ and \\xmm\\ observatories have provided strong evidence that despite cooling times being significantly shorter than the Hubble time \\citep[e.g.,][]{Sandersonetal06}, relatively little gas actually cools below $\\sim$1~keV \\citep{Petersonetal03,Kaastraetal04}.  It is now widely accepted that excessive cooling is in many systems prevented by a feedback mechanism in which the AGN of the central dominant galaxy, fuelled by cooling intra--cluster gas, can reheat the gas through a variety of mechanisms \\citep[e.g.,][and references therein]{PetersonFabian06,McNamaraNulsen07}.  X--ray and radio images provide numerous examples of interactions between AGN and the surrounding intra--cluster medium (ICM). Deep multiwavelength observations of the brightest nearby clusters have revealed complex structures associated with the radio jets and lobes, including shocks, sound waves, individual and linked chains of cavities, uplifted material and cooling filaments \\citep[e.g.,][]{Fabianetal05,Fabianetal06,Formanetal07,Wiseetal07,Blantonetal09}. Much of this work has concentrated on the cavities in the ICM which radio lobes produce as they inflate. The enthalpy of the cavities can be used as a measure of the mechanical power output of the radio jets, and has been shown to be sufficient to prevent or greatly reduce cooling in many systems \\citep{Birzanetal04}, provided the energy can be transfered into the intracluster medium and distributed relatively isotropically.  The disturbed structures produced by AGN jet/ICM interactions are relatively short--lived, and increasingly difficult to detect as they age. The radio lobes which inflate cavities fade rapidly once the AGN outburst ceases, and the X--ray cavities, which are detected by contrast with their surroundings, become less visible once they move beyond the dense group or cluster core. It is therefore considerably more difficult to study older AGN outbursts. However, since the mechanism by which cavities heat their surroundings is still a matter of debate, it is desirable to observe systems with as wide a range of ages as possible, so as to understand clearly the interaction between the radio lobes and their environment.  Most observations of radio galaxies to date have been based on observations at frequencies $>$1~GHz. Radio lobes may be studied over a wider range of timescales by observations at lower radio frequencies, which probe lower energy electrons less affected by spectral aging. Deep X--ray imaging is needed to complement such observations, and in this paper we discuss one example where this combination is available, the poor cluster AWM~4.  A previous \\xmm\\ observation found AWM~4 to be approximately isothermal to a radius of $\\sim$150~kpc, with no evidence of a central cool core \\citep[hereafter referred to as OS05]{OSullivanetal05_special}. Comparison with MKW~4, a cluster of similar temperature and galaxy population, but which hosts a large cool core and lacks a central radio source \\citep{OSullivanetal03}, leads to the suggestion that AWM~4 had been strongly heated by its central radio galaxy, 4C+24.36. However, examination of the \\xmm\\ data showed no evidence of cavities or shocks, and no high resolution images showing the lobes were available in the literature or radio archives. The existing VLA 1.4~GHz data were interpreted as evidence against the presence of lobes of sufficient volume to be responsible for reheating a large cool core \\citep{Gastaldelloetal08}. The ICM and galaxy distribution both appear relaxed with no significant substructure, with a strong concentration of early--type galaxies toward the core \\citep{KoranyiGeller02}. A cluster merger, which could also have heated the ICM, therefore appears unlikely. The central elliptical, NGC~6051, shows no signs of recent interactions \\citep{Schombert87} and is considerably more luminous than its neighbours, with a difference in magnitude above the second--ranked galaxy of $M_{12}$=1.6 (SDSS $g$-band).  An analysis of deep GMRT radio observations at 235, 327 and 610~MHz was presented in \\citep[hereafter referred to as GVM08]{Giacintuccietal08_special}. These data provided much new information about 4C+24.36, revealing radio emission from the jets and lobes extending $\\sim$75~kpc from the AGN. The source was shown to be a wide--angle--tail radio galaxy with inner jets oriented close to the plane of the sky, probably moving southward with a velocity of $\\la$120\\kmps.  From modelling of the progressive steepening of the spectral index $\\alpha$ (defined as $S\\propto\\nu^{-\\alpha}$ where $S$ is flux and $\\nu$ frequency) along the jets the radiative age of the electron population was estimated as 160-170~Myr.  In this paper, we use a new \\chandra\\ ACIS-S observation of AWM~4, in combination with the GMRT and archival VLA observations, to study the structure of the ICM and the interaction of the AGN, radio jets and lobes with the surrounding hot gas. The general properties of the cluster are summarised in Table~\\ref{tab:intro}, along with the position, distances and angular scale of the system. Throughout the paper we assume \\Ho=70, $\\Omega_M=0.3$, and $\\Omega_{\\Lambda}=0.7$. Uncertainties are generally quoted at the 1$\\sigma$ level, except in the case of X--ray spectral fitting, where 90 percent uncertainties were estimated. Section~\\ref{sec:obs} describes the observation and data reduction, and Sections~\\ref{sec:img} and \\ref{sec:spec} our imaging and spectral analysis. In Section~\\ref{sec:coregas} we examine the properties of the gas in the core of NGC~6051, immediately surrounding the AGN, and in Section~\\ref{sec:PB} we discuss the interaction between the radio jets and ICM, and place limits on the timescale of the outburst and the particle content of the radio lobes. We discuss our results in Section~\\ref{sec:discuss} and list our conclusions in Section~\\ref{sec:con}.    \\begin{table} \\caption{\\label{tab:intro} General properties of the AWM~4 system} \\begin{tabular}{llc} \\hline AWM~4    & z & 0.0318 \\\\ & D$_A$ (Mpc) & 130.9 \\\\ & D$_L$ (Mpc) & 139.3 \\\\ & angular scale (kpc/\\arcs) & 0.63 \\\\%0.634 & $\\sigma_v$ (\\kmps) & 400 \\\\ NGC~6051 & RA$_{\\rm J2000}$ (h m s) & 16 04 56.8 \\\\ & DEC$_{\\rm J2000}$ (\\degree\\ \\arcm\\ \\arcs) & +23 55 56 \\\\ & \\LB\\ (\\LBsol) & 6.94$\\times10^{10}$ \\\\ 4C+24.36 & S$_{1.4 GHz, NVSS}$ (mJy) & 608 \\\\ & S$_{235 MHz}$ (mJy) & 2750 \\\\ & log P$_{1.4 GHz}$ (W Hz$^{-1}$) & 24.14 \\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "\\label{sec:con}  We have used a deep, $\\sim$75~ks \\chandra\\ observation of the poor cluster AWM~4 to examine its structure and properties, and the relationship between the central radio galaxy and the ICM. Previous studies of AWM~4 found the cluster to have a number of unusual and conflicting features. GMRT observations showed that its dominant galaxy hosts an old, active FR-I radio galaxy, but \\xmm\\ found no evidence of cooling in the cluster core to fuel this AGN.  Conversely, heating a cool core to produce the approximately isothermal $\\sim$2.6~keV ICM observed required more energy than was estimated to be available from the radio source. Our analysis provides solutions to some of these problems, as well as insights which may be applicable to other clusters and cluster central radio sources. Our results can be summarised as follows:  \\begin{enumerate}  \\item The \\chandra\\ observation reveals a small cool core located at the centre of NGC~6051 and coincident with the core radio source. This meets the criteria for a galactic corona \\citep{Sunetal07,Sun09}. It is compact (radius $\\sim$1-2~kpc), significantly cooler than the surrounding cluster halo ($k_BT$=1.0$^{+0.35}_{-0.19}$~keV compared to $\\sim$2.6 keV for the ICM), and has a short cooling time (181$^{+108}_{-57}$~Myr) and moderate luminosity (L$_\\mathrm{X,0.5-2}=1.76\\times10^{40}$\\ergps). Heat conduction at the Spitzer rate would be sufficient to heat the core to the temperature of the surrounding ICM in 10-20~Myr.  This suggests that conduction is strongly suppressed. Similarly, a few percent of the mechanical energy of the radio jets would be sufficient to have heated the corona over the lifetime of the AGN outburst, and we conclude that any interaction between the jets and corona must be extremely inefficient. VLA 4.9~GHz radio maps do not resolve the jets inside the corona, suggesting that they are collimated and narrow, only broadening at its outer edge.  \\item We estimate the mass deposition rate through radiative cooling of the corona gas to be $\\dot{M}_{cool}$=0.067\\Msol yr$^{-1}$. This would be sufficient to power the AGN, requiring an efficiency in converting the cool gas to energy of only 0.1 per cent. Mass loss from stars within the corona appears sufficient to approximately balance cooling losses from the corona.  Direct accretion from the 1~keV gas at the Bondi rate could fuel the AGN, though the accretion rate is rather uncertain owing to the large extrapolation in radius required.  Magnetic separation of the corona from the ICM would prevent gas cooling from the ICM reaching the central engine, and the AGN jets do not significantly heat the corona. These factors appear to preclude a feedback relationship between AGN and ICM. However, the corona itself is capable of fuelling the AGN for long periods, and this may explain the unusually long outburst timescale estimated for the radio source. It may also explain the lack of a large cool core in the system, since without a feedback relationship, AGN heating seems likely to have exceeded cooling over the recent history of the cluster.   \\item Imaging shows the gaseous halo of AWM~4 to be generally relaxed, in agreement with previous observations. There are weak indications of structures associated with the radio source, including a broad bay-like structure around the west jet and lobe. However, the only statistically significant surface brightness feature is a decrement near the centre of the east lobe. There is no evidence of spectrally hard emission associated with the lobes, and the expected level of inverse Compton emission is below our detection threshold. If the lobes contained only relativistic plasma, we would expect to detect the cavities with high statistical significance. We interpret these results as indicating that the lobes are only partially filled. This is supported by the clumpy, filamentary appearance of the lobes in radio images. Based on the surface brightness, we can place limits on the fraction of ICM plasma in the lobes. Assuming the remaining volume is occupied by radio--emitting relativistic plasma, we find filling factors for this component of $\\phi=0.24$ and $\\phi=0.21$ for the east and west lobes respectively, with 3$\\sigma$ upper limits of $\\phi<0.43$ and $\\phi<0.76$.   \\item We measure the pressure profile of the ICM, and compare this with minimum energy pressure estimates for the jets and lobes of the radio source. Under the most conservative conditions, assuming contributions from particles emitting between 10~MHz and 100~GHz, we find a strong pressure imbalance between the lobes and their environment, with the lobes apparently underpressured by a factor $\\sim$160. However, these conditions imply an age for the source which is considerably longer than the timescale for the lobes to buoyantly rise to their current position. Estimates which include less energetic particles greatly reduce the pressure imbalance, to a factor $\\sim$15 for $\\gamma_{min}$=100, or to a factor $\\sim$4 for $\\gamma_{min}$=10, with the latter case consistent with pressure equilibrium within uncertainties. The radiative age estimated for $\\gamma_{min}$=100 is roughly consistent with the buoyant timescale of the lobes, while the age estimated for $\\gamma_{min}$=10 is significantly shorter and could imply a supersonic expansion phase for the jets. We consider the effects of bending in the jets on our pressure estimates, but find that they are unlikely to have a significant impact on our results.  \\item From these measurements and the filling factor estimates described above, we estimate the required ratio of energy in non--radiating particles to the observed energy in electrons, $k$, for each lobe. Assuming $\\gamma_{min}=100$, we estimate $k$=37.5 and 24.8 for the east and west lobes, with 3$\\sigma$ upper limits of $k<741.6$ and $k<517.0$ respectively. For $\\gamma_{min}=10$, we estimate $k$=10.4 and $k$=6.8 for the east and west lobes, with large uncertainties consistent with $k$=1 (an electron--proton plasma) or $k$=0 (an electron--positron plasma), or with the $\\gamma_{min}=100$ values. This indicates that in principle the apparent pressure imbalance in the lobes can be resolved by the inclusion of these lower energy relativistic particles. Alternatively, entrainment and heating of thermal plasma (either from stars or the ICM) could provide the necessary additional pressure. However, such entrained material must have a low density and very high temperature, rendering it undetectable with the current data.   \\item We estimate the enthalpy of the radio lobes and find that for the filling factors estimated above it is $\\sim0.3-7.0\\times10^{58}$~erg. This is lower than the estimated energy required to reheat a large cool core such as that seen MKW~4, a cluster of similar temperature and galaxy population. The mechanical power output of the jet depends on the timescale of the outburst; for our best estimate of $\\sim$170~Myr the jet power is $\\sim0.6-1.3\\times10^{43}$\\ergps. This is lower than or comparable to the bolometric X--ray luminosity of the ICM within the radius of the lobes, $\\sim1.3\\times10^{43}$\\ergps, suggesting that in the absence of other forms of heating, the energy available from the radio lobes is at best just sufficient to balance cooling in the ICM.  However, the relative isothermality of the ICM and lack of any significant cooling region outside the corona strongly suggests that the AGN has provided enough energy to balance or exceed cooling losses in the past. This could be achieved through additional heating mechanisms (e.g., weak shocks, sound waves, cosmic rays), larger filling factors, or a shorter outburst timescale. If we instead assume the shorter radiative timescale ($\\sim$66~Myr) estimated from the spectra of the lobes and assuming $\\gamma_{min}$=10, the jet mechanical power is 2.6$\\times10^{43}$\\ergps, in excess of the cooling rate. This timescale would also require supersonic expansion of the jets, providing additional heating through weak shocks.   \\end{enumerate}   \\noindent{\\textbf{ACKNOWLEDGEMENTS}}\\\\ The authors thank M. Sun and P. Mazzotta for informative discussions, and the anonymous referee for a number of useful suggestions. Support for this work was provided by the National Aeronautics and Space Administration through Chandra Award Number GO8-9127X-R 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. E.  O'Sullivan acknowledges the support of the European Community under the Marie Curie Research Training Network. We thank the staff of the GMRT for their help during the observations. GMRT is run by the National Centre for Radio Astrophysics of the Tata Institute of fundamental Research. We acknowledge the usage of the HyperLeda database (http://leda.univ-lyon1.fr)."
    },
    "1004/1004.0324_arXiv.txt": {
        "abstract": "In searches for gravitational waves emitted by known isolated pulsars in data collected by a detector one can assume that the frequency of the wave, its spindown parameters, and the position of the source in the sky are known, so the almost monochromatic gravitational-wave signal we are looking for depends on at most four parameters: overall amplitude, initial phase, polarization angle, and inclination angle of the pulsar's rotation axis with respect to the line of sight. We derive two statistics by means of which one can test whether data contains such gravitational-wave signal: the $\\G$-statistic for signals which depend on only two unknown parameters (overall amplitude and initial phase), and the $\\F$-statistic for signals depending on all four parameters. We study, by means of the Fisher matrix, the theoretical accuracy of the maximum-likelihood estimators of the signal's parameters and we present the results of the Monte Carlo simulations we performed to test the accuracy of these estimators. ",
        "introduction": "We study the detection of almost monochromatic gravitational waves emitted by known single pulsars in data collected by a detector. Several such searches were already performed with data collected by the LIGO and GEO600 detectors \\cite{LSC04,LSC05,LSC07,LSC08,LSC10}. We thus assume that the frequency of the wave (together with its time derivatives, i.e.\\ the spindown parameters) and the position of the source in the sky are known. The gravitational-wave signal we are looking for depends on at most four (often called amplitude) parameters: overall amplitude, initial phase, polarization angle, and inclination angle (of the pulsar's rotation axis with respect to the line of sight).  In Sec.\\ 2 we introduce three statistics by means of which one can test whether data contains a gravitational-wave signal: the $\\mathcal{H}$-statistic for completely known signals, the $\\G$-statistic for signals which depend on only two unknown parameters (overall amplitude and initial phase), and the $\\F$-statistic suitable for signals depending on all four amplitude parameters. Both statistics $\\G$ and $\\F$ are derived from the maximum likelihood (ML) principle, and the statistic $\\G$ is independently obtained using Bayesian approach and the composite hypothesis testing. In Sec.\\ 3 we study, by means of the Fisher matrix, the theoretical accuracy of the ML estimators of the signal's parameters and in Sec.\\ 4 we present the results of the Monte Carlo simulations we performed to test the accuracy of the ML estimators. ",
        "conclusions": ""
    },
    "1004/1004.0781_arXiv.txt": {
        "abstract": "In the last decade, numerous Lyman-alpha emitters (LAEs) have been discovered with narrow-band filters at various redshifts. Recently, multi-wavelength observations of LAEs have been performed and revealed that while many LAEs appear to be young and less massive, a noticeable fraction of LAEs possess much older populations of stars and larger stellar mass. How these two classes of LAEs are concordant with the hierarchical galaxy formation scenario has not been understood clearly so far. In this paper, we model LAEs by three-dimensional cosmological simulations of dark halo merger in a $\\Lambda$CDM universe. As a result, it is shown that the age of simulated LAEs can spread over a wide range from $2\\times 10^6$yr to $9\\times 10^8$yr. Also, we find that there are two types of LAEs, in one of which the young half-mass age is comparable to the mean age of stellar component, and in the other of which the young half-mass age is appreciably shorter than the mean age. We define the former as Type 1 LAEs and the latter as Type 2 LAEs. A Type 1 LAE corresponds to early starburst in a young galaxy, whereas a Type 2 LAE does to delayed starburst in an evolved galaxy, as a consequence of delayed accretion of a subhalo onto a larger parent halo. Thus, the same halo can experience a Type 2 LAE-phase as well as a Type 1 LAE-phase in the merger history. Type 1 LAEs are expected to be younger than $1.5 \\times 10^8$yr, less dusty, and less massive with stellar mass $M_{\\rm star} \\la 5 \\times 10^8 \\rm ~M_{\\odot}$, while Type 2 LAEs are older than $1.5 \\times 10^8$yr, even dustier, and as massive as $M_{\\rm star} \\sim 5 \\times 10^8 - 3\\times 10^{10} \\rm ~M_{\\odot}$. The fraction of Type 2s in all LAEs is a function of redshift, which is less than 2 percent at $z \\ga 4.5$, $\\sim$30 percent at redshift $z=3.1$, and $\\sim$70 percent at $z=2$. Type 2 LAEs can be discriminated clearly from Type 1s in two color diagram of z'-H vs J-K. We find that the brightness distribution of Ly$\\alpha$ in Type 2 LAEs is more extended than the main stellar component, in contrast to Type 1 LAEs. This is not only because delayed starbursts tend to occur in the outskirts of a parent galaxy, but also because Ly$\\alpha$ photons are effectively absorbed by dust in an evolved galaxy. Hence, the extent of Ly$\\alpha$ emission may be an additional measure to distinguish Type 2 LAEs from Type 1 LAEs. The sizes of Type 2 LAEs range from a few tens to a few hundreds kpc. At lower redshifts, the number of more extended, older Type 2 LAEs increases. Furthermore, it is anticipated that the amplitude of angular correlation function for Type 2 LAEs is significantly higher than that for Type 1 LAEs, but comparable to that for Lyman break galaxies (LBGs). This implies that LBGs with strong Ly$\\alpha$ line may include Type 2 LAEs. ",
        "introduction": "To explore the early evolutionary phases of galaxies is important to understand galaxy formation. \\citet{PP1967} predicted that the starbursts in primeval galaxies emit significant Ly$\\alpha$ emission through the recombination of ionized hydrogen in interstellar matter. Although many surveys attempted to discover such Ly$\\alpha$ emitting galaxies (Ly$\\alpha$ emitters: hereafter LAEs), but did not succeed to find them for a long time. In late 1990's, \\citet{CH1998} discovered LAEs with narrow-band filters for the first time. Currently, numerous LAEs have been discovered at high redshifts $3 < z < 7$ by $8 \\sim 10 {\\rm ~m}$ class telescopes with narrow-band filters \\citep{Hu98, Hu99, Hu02,Kodaira03,Shimasaku03,Shimasaku06,Ha2004,Ou04,Ou05,Taniguchi05,Matsuda04,Matsuda05}. Although the number of observed LAEs increases constantly, the nature of LAEs is still veiled. Recently, surveys of LAEs in the various wavelength bands (optical, infrared, sub-mm, etc) have been performed actively \\citep{Fin2007, Lai2008, Matsuda2007, Uchimoto2008, Fin2009, SMGLAE}, and have revealed that while many LAEs appear to be young and less massive, a noticeable fraction of LAEs possess much older stellar populations and larger stellar mass. We have not well understood how such two classes of LAEs are concordant with the hierarchical galaxy formation scenario. As for the physical origin of Ly$\\alpha$ emission, the cooling radiation from a primordial collapsing cloud \\citep{Haiman00,Fardal01}, from a galactic wind-driven shell \\citep{TS00}, or from star-forming clouds in a young starburst galaxy \\citep{MUF04} has been considered.  Recently, \\citet{MU2006} proposed a galaxy evolution scenario from LAEs to LBGs, based on a supernova-dominated starburst galaxy model. In this scenario, LAEs correspond to an early evolutionary phase of $<3 \\times 10^8$yr. Also, \\citet{S2007} have constructed an analytic model of LAEs in a $\\Lambda$CDM universe, and found that if LAEs form in relatively low density regions of the universe and the duration of starburst is as short as $0.7 \\times 10^8$~yr, the spatial distributions match the weak angular correlation function of LAEs observed at $z=3.1$. The spectral energy distribution (SED) fitting for observed LAEs has shown that LAEs mostly have young average age ($\\sim 10^8$yr) and low stellar mass ($10^8 \\sim 10^9 \\rm ~M_{\\odot}$), and are less dusty or dust free \\citep{Ga2006, Fin2007, Lai2008}. These young LAEs are consistent with the picture by \\citet{MU2006} and \\citet{S2007}. Very recently, deep surveys of LAEs allow us to study detailed properties of individual LAEs. As a result, it has been revealed that LAEs have a wide range of age ($10^6 \\sim 10^9$yr), stellar mass ($10^6 \\sim 10^{10} {\\rm ~M_{\\odot}}$), and dust extinction with ${\\rm A_{\\rm V}}$ up to $1.3 \\rm ~mag$ \\citep{Fin2007, Lai2008, Fin2009}. LAEs detected in rest-frame optical/near infrared (NIR) bands tend to have older age, larger stellar mass, and stronger dust extinction than LAEs undetected in those bands. Thus, the picture of purely young starburst galaxies are not always reconciled with observed LAEs. So far, the physical reason has not been clarified for the existence of an old, massive, and dusty population of LAEs. The previous study has shown that a starburst-dominated galaxy can emit strong Ly$\\alpha$ radiation in dust-free or less dusty environments \\citep{MU2006}. However, starburst galaxies cannot be always LAEs in dusty environments (e.g. ultra-luminous infrared galaxies). Hence, what physical state corresponds to an old population of LAEs is an issue of great significance. Some authors argue that the clumpy distributions of dusty gas is important for the transfer of Ly$\\alpha$ photons \\citep{Neufeld1991, HansenOh2006, Fin2009c}. Since Ly$\\alpha$ photons undergo resonant scatterings on the surface of gas clumps, photons can easily escape from the clumpy media. Such an effect provides the possibility of old, massive and dusty LAEs.  In this paper, we explore how a young and old population of LAEs are concordant with a hierarchical galaxy formation paradigm. For the purpose, we perform tree-dimensional cosmological simulations of dark halo merger in a $\\Lambda$CDM universe, incorporating the prescriptions of star formation, spectral evolution, and dust extinction. Throughout this paper, we adopt $\\Lambda$CDM cosmology with the matter density $\\Omega_{\\rm{M}} = 0.3$, the cosmological constant $\\Omega_{\\Lambda} = 0.7$, the Hubble constant $h = 0.7$ in units of $H_0 = 100 \\rm{~km ~s^{-1} ~Mpc^{-1}}$, the baryon density $\\Omega_{\\rm B}h^2 = 0.02$, and $\\sigma_8 = 0.92$ \\citep{WMAP}. ",
        "conclusions": "To explore the origin of two populations of LAEs recently found, we have performed three-dimensional cosmological $N$-body simulations of subhalo merging history in a $\\Lambda$CDM universe. We have incorporated star formation history, SED evolution, and dust extinction. As a result, we have found that the age of simulated LAEs can spread over a wide range from $2\\times 10^6$yr to $9\\times 10^8$yr. Also, we have revealed that there are two types of LAEs. We have defined LAEs younger than $1.5 \\times 10^8$yr as Type 1s, and older ones as Type 2s. In Type 1 LAEs early coeval starbursts occur in small parent halos, while in Type 2 LAEs delayed starbursts take place in evolved massive haloes as a consequence of delayed accretion of subhalos. A parent halo can experience repeatedly a Type 2 LAE-phase after a Type 1 LAE-phase.  The stellar mass of Type 1 LAEs is $M_{\\rm star} \\la 5 \\times 10^8 \\rm ~M_{\\odot}$, while Type 2 LAEs are as massive as $M_{\\rm star} \\sim 5 \\times 10^8 - 3\\times 10^{10} \\rm ~M_{\\odot}$. The physical properties of Type 1 and Type 2 LAEs are concordant with those of two populations of LAEs observed with multi-wavelengths \\citep{Fin2007, Lai2008, Matsuda2007, Uchimoto2008, Fin2009}. The fraction of Type 2s in all LAEs is a function of redshift, which is less than 2 percent at $z \\ga 4.5$, $\\sim$30 percent at redshift $z=3.1$, and $\\sim$70 percent at $z=2$. This trend is consistent with two populations of LAEs found by \\citet{Nilsson2009}. Type 2 LAEs are expected to be discriminated clearly from Type 1 LAEs in two color diagram of z'-H vs J-K. We find that the brightness distribution of Ly$\\alpha$ in Type 2 LAEs is more extended than the main stellar component, in contrast to Type 1 LAEs. This is not only because delayed starbursts tend to occur in the outskirts of a parent galaxy, but also because Ly$\\alpha$ photons are effectively absorbed by dust in an evolved galaxy. The sizes of Type 2 LAEs range from a few tens to a few hundreds kpc. At lower redshifts, the number of more extended, older Type 2 LAEs increases. Small Type 2 LAEs are as compact as Type 1 LAEs, while large Type 2 LAEs exceeding $100~ \\rm kpc$ are comparable to Ly$\\alpha$ blobs (LABs) \\citep{Matsuda04}.  Moreover, we have found that the clustering of Type 2 LAEs are even stronger than Type 1 LAEs. The amplitude of angular correlation function of Type 2 LAEs is comparable to that of Lyman break galaxies (LBGs) \\citep{Gia1998}. This suggests that LBGs with strong Ly$\\alpha$ line can be Type 2 LAEs. The two-point angular cross-correlation function is still weaker than that of all LAEs. If many Type 2 LAEs can be detected as SMGs, this result is consistent with recent observation by \\citet{SMGLAE}.  Interestingly, in a low redshift universe at $0.2<z<0.35$, the {\\it Galaxy Evolution Explorer (GALEX)} have found many LAEs older than $2 \\times 10^8$yr, more massive than $10^9 \\rm ~M_{\\odot}$ in stellar component, and having small escape fractions of Ly$\\alpha$ photons \\citep{Deh08, Fin2009c}. These properties are quite similar to those of Type 2 LAEs we defined in this paper. Hence, local LAEs could be a good sample to study the detailed physical states of Type 2 LAEs.  Furthermore, in an evolved population of LAEs, active galactic nucleus (AGN) events may be anticipated, since supermassive black holes reside evolved galaxies in a local universe \\citep[e.g.,][and references therein]{MH03}. Recently, \\citet{Yamada2009} have found that AGNs are exclusively $(90\\%)$ associated with the massive objects with stellar mass larger than $10^{10.5}$ $\\rm M_{\\odot}$. Also, {\\it GALEX} found that in low-redshift LAEs an AGN fraction is as high as forty percent \\citep{Fin2009b}. These findings may imply that Type 2 LAEs could have a higher fraction of AGNs. Hence, it seems worth investigating the contribution of AGNs to Ly$\\alpha$ emission especially for Type 2 LAEs (Shimizu \\& Umemura, in preparation)."
    },
    "1004/1004.1186_arXiv.txt": {
        "abstract": "We present new far-ultraviolet observations of the young M8 brown dwarf 2MASS J12073346-3932539, which is surrounded by an accretion disk.  The data were obtained using the $Hubble$~$Space$~$Telescope$-Cosmic Origins Spectrograph.  Moderate resolution spectra ($R$~$\\approx$~17,000~--~18,000) obtained in the 1150~--~1750~\\AA\\ and 2770~--~2830~\\AA\\ bandpasses reveal H$_{2}$ emission excited by \\ion{H}{1} Ly$\\alpha$ photons, several ionization states of carbon (\\ion{C}{1}~--~\\ion{C}{4}), and hot gas emission lines of \\ion{He}{2} and \\ion{N}{5} ($T$~$\\approx$~10$^{4}$~--~10$^{5}$ K). Emission from some species that would be found in a typical thermal plasma at this temperature (\\ion{Si}{2}, \\ion{Si}{3}, \\ion{Si}{4}, and \\ion{Mg}{2}) are not detected. The non-detections indicate that these refractory elements are depleted into grains, and that accretion shocks dominate the production of the hot gas observed on 2MASS J12073346-3932539. We use the observed \\ion{C}{4} luminosity to constrain the mass accretion rate in this system. We use the kinematically broadened H$_{2}$ profile to confirm that the majority of the molecular emission arises in the disk, measure the radius of the inner hole of the disk ($R_{hole}$~$\\approx$~3$R_{*}$), and constrain the physical conditions of the warm molecular phase of the disk ($T$(H$_{2}$)~$\\approx$ 2500~--~4000~K).  A second, most likely unresolved H$_{2}$ component is identified.  This feature is either near the stellar surface in the region of the accretion shock or in a molecular outflow, although the possibility that this Jovian-like emission arises on the day-side disk of a 6~$M_{J}$ companion (2M1207b) cannot be conclusively ruled out.  In general, we find that this young brown dwarf disk system is a low-mass analog to classical T Tauri stars that are observed to produce H$_{2}$ emission from a warm layer in their disks, such as the well studied TW Hya and DF Tau systems. ",
        "introduction": "Brown dwarfs form the low-mass, low-temperature end of the galactic stellar mass distribution.  Optical and near-IR surveys carried out over the past two decades have detected hundreds of objects with masses below the sustained hydrogen burning threshold (about 0.07~M$_{\\odot}$; Riaz \\& Gizis 2007).~\\nocite{riaz07} These brown dwarfs have effective temperatures $T_{eff}$~$\\lesssim$~2500 K in the spectral sequence from late M-stars through the L and T dwarfs. Although models of such low $T_{eff}$ brown dwarfs predict essentially no photospheric emission in the far-ultraviolet (far-UV; 1100~$\\leq$~$\\lambda$~$\\leq$~1700~\\AA), the far-UV spectra of of these objects are mostly unexplored (see Hawley \\& Johns-Krull 2003 and Gizis et al. 2005 for UV spectra of 4 late M-stars).  \\nocite{hawley03,gizis05}  More massive M dwarfs (spectral types earlier than $\\sim$~M5) have strong magnetic fields that, through processes still not fully understood, produce hot plasma and nonthermal particles. The rich phenomenology includes chromospheres ($T$~$\\sim$~10$^{4}$ K), transition regions ($T$~$\\sim$~10$^{5}$ K), coronae ($T$~$\\sim$~10$^{6}$ K), flares, and other transient activity. We observe this phenomenology through X-ray, far-UV, optical, and nonthermal radio emission lines and continua. It is unclear if far-UV emission is common in brown dwarfs.  H$\\alpha$ compared to L$_{bol}$ is observed to decrease in late M dwarfs (e.g., Berger et al. 2010), which would suggest less hot gas capable of producing far-UV emission.\\nocite{berger10} Alternatively, magnetically powered activity may extend through the late M dwarfs (and into the L and T classes), possibly with enhanced flaring relative to their more massive counterparts~\\citep{fleming00,berger08}.  Using $HST$-STIS, Hawley \\& Johns-Krull (2003) found that M7-M9 stars show bright emission in C~IV $\\lambda$1550 and other lines formed in high temperature plasmas. Also, Welsh et al. (2006) found both near-UV and far-UV emission and flares on several late-M dwarfs observed by $GALEX$.~\\nocite{hawley03,welsh06}  Far-UV emission from accretion shocks is another mechanism that could be important for young brown dwarfs where gas-rich accretion disks exist.  2MASS J12073346-3932539 (2M1207) is an M8 brown dwarf ($M$~$\\approx$~0.024~$M_{\\odot}$; Riaz \\& Gizis 2007) located in the $\\sim$ 10 Myr TW Hya association~\\citep{kastner97,webb99}, at a distance of 52.4 pc (V = 20.15; Ducourant et al. 2008).\\nocite{riaz07,ducourant08} Initial evidence for a circumstellar disk in this system was inferred from the association with the TW Hya group and the detection of strong H$\\alpha$ emission, indicative of active accretion~\\citep{gizis02}.  More recently, photometric and spectroscopic observations of 2M1207 from the far-red to the mid-infrared (IR) have allowed various groups to confirm the existence of a circumstellar dust disk~\\citep{riaz07,morrow08}.  The accretion from this system has been shown to vary on timescales ranging from hours to weeks~\\citep{scholz05,scholz06}, although the absolute level of the mass accretion rate ($\\dot{M}$$_{acc}$) is somewhat unclear~\\citep{herczeg09}.  While 2M1207 is an active accretor, searches for magnetic fields have returned only a surprisingly low 3-$\\sigma$ upper limit.  Reiners et al. (2009) constrain the total magnetic flux, $Bf$~$<$~1 kG, where $B$ is the unsigned photospheric magnetic field, and $f$ is the magnetic field filling factor.~\\nocite{reiners09} This object has the highest quality existing far-UV dataset of any brown dwarf, obtained with $HST$-STIS~\\citep{gizis05}, where several lines produced in hot gas (most notably \\ion{He}{2} and \\ion{C}{4}) were observed.  The observation of these lines, combined with a non-detection of \\ion{Si}{4} emission, led these authors to conclude that the UV emission is produced by accretion, and that the silicon in the 2M1207 disk has depleted into dust grains.  If the hot gas lines were produced in a thermal plasma due to stellar activity, all astrophysically abundant species with similar emissivities should be present (for reasonable metallicities), the lack of \\ion{Si}{4} implies that the hot gas is not primarily created in the stellar atmosphere, but at the shock interface between the accretion disk and the stellar surface.   \\begin{figure*} \\begin{center} \\hspace{+0.0in} \\epsfig{figure=f1.eps,width=4.5in,angle=90} \\caption{\\label{cosovly} The far-UV spectrum of 2M1207.  The top plot is a weighted coaddition of the $HST$-COS G130M and G160M observations. The bottom plot highlights the 1480~--~1580~\\AA\\ region with relatively strong lines of warm (H$_{2}$; $T$~$\\gtrsim$~2500~K.) and hot (\\ion{C}{4}; $T$~$\\approx$~ 1~$\\times$~10$^{5}$ K) gas. } \\end{center} \\end{figure*}  \\begin{deluxetable}{cccc} \\tabletypesize{\\footnotesize} \\tablecaption{2M1207 COS observing log. \\label{cos_obs}} \\tablewidth{0pt} \\tablehead{ \\colhead{Dataset} & \\colhead{COS Mode} & \\colhead{Central Wavelength} & \\colhead{T$_{exp}$ (s)} } \\startdata lb4p02fm\t& \tG160M \t& \t1600\t\t&\t 1410\t \\\\ lb4p02ft\t& \tG160M \t& \t1611 \t&\t 2975\t \\\\ lb4p02gn\t& \tG160M \t& \t1623 \t&\t 2975\t \\\\ lb4p02gr\t& \tG130M \t& \t1291 \t&\t 2967\t \\\\ lb4p02gx\t& \tG130M \t& \t1300 \t&\t 3000\t \\\\ lb4p02h6\t& \tG130M \t& \t1309 \t&\t 3005\t \\\\ lb4p01s1\t& \tG285M \t& \t2676 \t&\t 3050  \t \\\\ lb4p01rx\t& \tG140L \t& \t1230 \t&\t 3015  \t \\\\ lb4p01rv\t& \tG140L \t& \t1230 \t&\t 3015  \t \\\\ lb4p01rt\t& \tG140L \t& \t1230 \t&\t 200  \t \\\\ \\enddata  \\end{deluxetable}   Molecular hydrogen (H$_{2}$) emission was also detected in the STIS observations, although the low spectral resolution of the G140L mode ($\\Delta$$v$~$\\sim$~240 km s$^{-1}$) prevented a conclusive determination of the location of the molecular gas in the system.  The most likely origin of the H$_{2}$ emission is in a warm molecular layer of the circumstellar disk, in analogy to the warm disks seen around more massive classical T Tauri stars (CTTSs; Herczeg et al. 2006).  It is worth noting in introducing the 2M1207 system that another reservoir of H$_{2}$ resides at the extrasolar giant planet companion, 2M1207b ($M_{b}$~$\\sim$~6 $M_{J}$, secondary/primary mass ratio $\\sim$~1:4; Chauvin et al. 2004; Song et al. 2006).~\\nocite{chauvin04,song06}  This object is at a radial distance of $\\sim$~40 AU, although the orbit is poorly constrained due to the long temporal baseline necessary for orbital monitoring.  The nature of 2M1207b is the subject of some debate~\\citep{mohanty07,mamajek07,ducourant08}.      In this paper, we present new far-UV observations of this interesting low-mass system. These data cover a similar spectral bandpass as the $HST$-STIS observations presented by~\\citet{gizis05}, but with order of magnitude increases in sensitivity and resolving power. We describe the COS observations and data reduction in \\S2.  A quantitative analysis of the far-UV spectrum is presented in \\S3, with a focus on the properties of the H$_{2}$ in the circumstellar disk and the hot gas produced in the accretion shock.  In \\S4, we put these results in the context of other young objects in a stage of active disk accretion and argue that the similarities are evidence that accreting brown dwarfs are low-mass analogs to CTTSs.  We conclude with a brief summary in \\S5. ",
        "conclusions": "\\begin{figure*} \\begin{center} \\hspace{+0.0in} \\epsfig{figure=f8.eps,width=6.5in,angle=0} \\caption{\\label{cosovly} A cartoon representation of the inner disk region of 2M1207, based on interpretations in the literature and the new $HST$-COS observations presented here. The distances are not to scale.  In this cartoon, we see hot gas emission (Ly$\\alpha$ and \\ion{C}{4}, $T_{s}$~$\\sim$~10$^{5}$K) produced where the accretion stream makes contact with the brown dwarf surface, dust and depleted species (Si and Mg) distributed near the mid-plane, while the Ly$\\alpha$-pumped H$_{2}$ traces a warmer ($T(H_{2})$~$\\approx$~2500~--~4000 K), extended surface layer of the disk.  The dust sublimation point is assumed to be $T_{sub}$~=~1600 K~\\citep{riaz07}. } \\end{center} \\end{figure*}   \\subsection{Mass Accretion Rate: \\ion{C}{4} Luminosity }  We do not have contemporaneous observations of 2M1207 in the optical or NUV, where traditional accretion diagnostics are located (Herczeg et al. 2009 and references therein).~\\nocite{herczeg09} \\citet{krull00} present empirical relations for determining the $\\dot{M}$$_{acc}$ of CTTSs, and while these relations were created for higher mass objects with larger mass accretion rates, it is interesting to compare an extrapolation of the CTTS relation to a more standard technique.  For this purpose, we use Equation 2 of~\\citet{krull00}. In addition to 2M1207 having over an order of magnitude smaller mass than any stars considered by~\\citet{krull00}, we note they assume that the accretion emission is in excess of a saturated magnetic component that produces a surface \\ion{C}{4} flux level of $F_{CIV}$ $>$ 10$^{6}$ ergs cm$^{-2}$ s$^{-1}$.  Taking a distance of 52.4 pc~\\citep{ducourant08} and a stellar radius of 0.24~$R_{\\odot}$, we find a total \\ion{C}{4} surface flux of $F_{CIV}$ = 1.28 $\\times$ 10$^{5}$ ergs cm$^{-2}$ s$^{-1}$ (where no saturated magnetic component has been subtracted). While this is lower than the saturation threshold suggested by~\\citet{krull00}, the combination of the low magnetic field at 2M1207 ($<$ 1 kG; Reiners et al. 2009) and the non-detection of Si and Mg species (Section 3.2.1) lead us to assert that essentially all of the \\ion{C}{4} emission from 2M1207 is produced by accretion.   The empirical relation between the \\ion{C}{4} luminosity ($L_{CIV}$, in units of ergs s$^{-1}$) and $\\dot{M}$$_{acc}$ depends strongly on the method and values used for dereddening the observed fluxes, particularly at the wavelength of \\ion{C}{4} (1550 \\AA), where the effects of interstellar extinction are large~\\citep{ccm}. The mass accretion rate is then \\begin{equation} log_{10}(\\dot{M}_{acc}) = 0.753 \\ log_{10}(L_{CIV}) - 29.89 \\end{equation} The $L_{CIV}$ is calculated to be 4.49 $\\times$ 10$^{26}$ ergs s$^{-1}$ We note that the 2M1207 sightline is generally assumed to suffer no interstellar extinction ($A_{V}$~=~0.0; Herczeg et al 2009), and no correction was applied to the \\ion{C}{4} line fluxes presented in Table 3.  We find log$_{10}$ $\\dot{M}$$_{acc}$~$\\approx$~-9.8 [$\\dot{M}$$_{acc}$ = 1.6~$\\times$~10$^{-10}$ $M_{\\odot}$ yr$^{-1}$] from the \\ion{C}{4} observations of the 2M1207 system.  This value is is consistent with the accretion level of 2M1207 derived from H$\\alpha$ observations (log$_{10}$ $\\dot{M}$$_{acc}$~=~-10.1~$\\pm$~0.7) obtained in the ``high'' state~\\citep{scholz05}.  \\citet{krull00} note that alternative calibrations produce accretion rates that about 10 times lower than those given in Equation 1 above.  If that scaling is applied, we find that the $\\dot{M}$$_{acc}$ derived from the \\ion{C}{4} line strengths is consistent with the lower values observed by Scholz et al. (2005; log$_{10}$ $\\dot{M}$$_{acc}$ = -10.8~$\\pm$~0.5).   Interestingly, while we find the \\ion{C}{4}-based accretion rate to be in excellent agreement with that derived from H$\\alpha$ observations, our low value are approximately an order of magnitude greater than those measured using deep, low-resolution observations of the Balmer continuum (log$_{10}$ $\\dot{M}$$_{acc}$~=~-11.9; Herczeg et al. 2009).  The accretion rate in the 2M1207 system is known to vary by at least an order of magnitude, and since none of observations were acquired simultaneously, it is plausible that variability causes the discrepancy between the mass accretion rates measure by H$\\alpha$, \\ion{C}{4}, and Balmer continuum observations.  Alternatively, absorption of Balmer continuum emission by the edge-on disk may lead to a lower estimation of the accretion rate by this method.        \\subsection{Physical Origin of the Narrow H$_{2}$ Component}  In Figure 2, we displayed the coadded H$_{2}$ emission line profile of the six lines with the highest S/N in the COS M-grating data.  In Section 3.1, we discussed the dominant broad component and identify it as emission from a pile-up of material at the inner wall of the circumstellar disk, approximately at the disk sublimation radius (Figure 8).  The velocity width of the second component is poorly constrained as the fit is dominated by the broader, stronger component.  The velocity width is consistent with being an unresolved feature.  There are several possible physical origins for an unresolved H$_{2}$ population.  The most likely scenario seems to be that this additional emission arises at the stellar surface.  The photospheric temperature ($T_{*}$) is 2550 K~\\citep{riaz07}, ideal for maintaining an H$_{2}$ population that is capable of being pumped by Ly$\\alpha$ photons in thermal equilibrium.  If a photospheric origin is the correct interpretation, this would argue that the emitting molecules are near the accretion hotspot created at the interface of the infalling material from the disk, seen in our COS observations through several ionization states of He, C, and N.  The relaxation time for the electronic transitions of H$_{2}$ is very short ($A_{TOT}$ for the (1~--~2) R(6) transition coincident with Ly$\\alpha$ is 1.68~$\\times$~10$^{9}$ s$^{-1}$; Abgrall et al. 1993), and the UV transitions of the H$_{2}$ molecules would not be visible if they were not being actively excited.~\\nocite{abgrall93a}  The $\\approx$ 15 km s$^{-1}$ blueshift of this component relative to the bulk of the H$_{2}$ emission from the disk suggests a possible outflow origin.  The CTTSs T Tau and RU Lupi show narrow, blueshifted H$_{2}$ emission that is thought to be indicative of a bipolar outflow~\\citep{herczeg06}.  The blueshift of the outflow emission in these objects is roughly the same ($v$~=~-12 km s$^{-1}$) as that found for 2M1207. If an outflow is the correct interpretation, the 15 km s$^{-1}$ relative velocity of the narrow H$_{2}$ component in 2M1207 is surprising because T Tau and RU Lupi host nearly face-on disks, where the outflow jet is pointed more directly at the observer.  It seems unlikely that the edge-on orientation of the 2M1207 disk would permit the same magnitude of blueshift produced in more massive, face-on disks, however, 2M1207 is observed to have [\\ion{O}{1}] emission that is consistent with an outflow~\\citep{whelan07}.  One final possibility is that the weak H$_{2}$ emission originates in the dayglow or aurorae of the 6~$M_{J}$ companion, 2M1207b (Chauvin et al. 2004; and see France et al. 2010 for a detailed discussion of the predicted UV emission properties of extrasolar giant planets).~\\nocite{chauvin04,france10a}  The far-UV spectrum of Jupiter is dominated by H$_{2}$ emission, where the excitation is caused by electron-impact where the magnetic field lines connect to the planetary surface near the poles and solar-induced Ly$\\beta$ fluorescence in the equatorial regions~\\citep{feldman93,wolven98}.  In the instance of an additional energy source (in this case the Shoemaker Levy 9 impact), the Jovian atmosphere supports Ly$\\alpha$ pumped H$_{2}$ emission~\\citep{wolven97} similar to that observed in 2M1207.  The velocity shift due to the orbital motion of the planet would be undetectable at the COS resolution ($v_{orb}$~$\\sim$~0.7 km s$^{-1}$ at 40 AU, assuming a circular orbit), and this scenario would require both a mechanism to heat the 2M1207b atmosphere to $T$~$\\gtrsim$~2500 K, and produce a 15 km s$^{-1}$ outflow.  While we favor a photospheric origin for the narrow H$_{2}$ component in 2M1207, we cannot conclusively rule out an outflow or the extrasolar giant planet companion as possible sources of the observed H$_{2}$.  \\begin{figure} \\begin{center} \\hspace{+0.0in} \\epsfig{figure=f9.eps,width=2.5in,angle=90} \\caption{\\label{cosovly} Time-tagged fluxes from emission lines tracing the warm (H$_{2}$) and hot (\\ion{C}{4} and \\ion{N}{5}) components of the 2M1207 system, in 200 second time intervals. The line fluxes are essentially constant, with most of the variability in H$_{2}$ and \\ion{N}{5} caused by a time variable background level. } \\end{center} \\end{figure}   \\begin{figure} \\begin{center} \\hspace{+0.0in} \\epsfig{figure=f10.eps,width=2.5in,angle=90} \\caption{\\label{cosovly} A direct comparison of the background subtracted line fluxes ($\\Delta$$t$~=~200 s) of \\ion{C}{4} and H$_{2}$.  \\ion{C}{4} emission is representative of the strength of the \\ion{H}{1} Ly$\\alpha$ line, which drives the observed flux level of H$_{2}$.  The non-variable nature of the lines leads to a Pearson correlation coefficient of 0.15, essentially uncorrelated. } \\end{center} \\end{figure}   \\subsection{Young Brown Dwarfs: Low-Mass Classical T-Tauri Analogs}  As mentioned in the previous subsection, 2M1207 displays metal depletions consistent with those seen in some CTTSs.  The H$_{2}$ disk emission is also reminiscent of that observed around more massive young stars.   We therefore argue that 2M1207 is a low-mass analog to these systems. While TW Hya is a somewhat atypical pre-main sequence object (with respect to the ages, accretion rates, and abundances of other CTTSs; we refer the reader to Section 1 of Herczeg et al. (2002) for a concise review), we use it for comparison with 2M1207 based on its well-studied far-UV spectrum~\\citep{herczeg02}.  The H$_{2}$ emission seen in our COS observations is qualitatively similar to that of TW Hya, however there are quantitative differences in the far-UV spectra of these objects.  The first is the wealth of lines observed in the spectrum of TW Hya compared to 2M1207.  While the 2M1207 observations are at a lower S/N than the STIS observations of TW Hya,  there are numerous emission lines that would have been detected if they were present with the relative strengths seen in TW Hya (in particular, emission lines pumped by (0~--~2) R(0) 1217.21~\\AA\\ and  (0~--~2) R(1) 1217.64~\\AA).  This implies that the Ly$\\alpha$ emission profile in 2M1207 is considerably narrower than that observed in higher-mass CTTSs. These ``missing'' fluorescent progressions are pumped by the wings of a broad stellar/shock Ly$\\alpha$ emission profile, which are mostly inaccessible to COS due to contamination by geocoronal Ly$\\alpha$.  The lack of a broad Ly$\\alpha$ component in 2M1207 may be further evidence that Ly$\\alpha$ is created in the accretion shock in this object~\\citep{herczeg06}.  We can make a quantitative comparison of the H$_{2}$ flux from TW Hya and 2M1207.  The total flux ratio between the two ($R^{TW}_{2M}$($TOT$)~$\\equiv$~$I_{H2}$(TW Hya)/$I_{H2}$(2M1207)) is not the appropriate measure as TW Hya produces many more emission lines based on the broad stellar Ly$\\alpha$ profile.  We compare the total H$_{2}$ emission from specific states observed in 2M1207, namely, those pumped by (1~--~2) R(6) 1215.73~\\AA\\ and (1~--~2) P(5) 1216.07~\\AA.  The distance corrected flux ratios for the emission produced by pumping in those two lines are $R^{TW}_{2M}$(1~--~2 R(6))~=~391 and $R^{TW}_{2M}$(1~--~2 P(5))\\footnote{Adding up the flux from the individual lines in TW Hya (Table 2 of Herczeg et al. 2002), we found a total flux of 369.6 $\\times$~10$^{-15}$ ergs cm$^{-2}$ s$^{-1}$, in slight disagreement with the value of 350 quoted in their Table 6} =~350, respectively.  The $R^{TW}_{2M}$(1~--~2 R(6)) ratio is more susceptible to the effects of self-absorption by \\ion{H}{1} in the circumstellar environment, though the ratios for both lines are similar. This implies that there is more Ly$\\alpha$ flux per H$_{2}$ in the disk of TW Hya compared to the disk of 2M1207, assuming that the disk masses are proportional to the mass of the primary ($M_{TW}$/$M_{2M}$~=~0.7 $M_{\\odot}$/0.024 $M_{\\odot}$~$\\approx$~30).   The excess disk H$_{2}$ emission in TW Hya can be interpreted as a stronger local Ly$\\alpha$ radiation field, which we propose is due to the higher mass accretion rate in TW Hya ($\\sim$~2~$\\times$ 10$^{-9}$ $M_{\\odot}$ yr$^{-1}$ as compared to $\\sim$~1~--~150~$\\times$ 10$^{-12}$ $M_{\\odot}$ yr$^{-1}$ for 2M1207; Herczeg et al. 2006, Scholz et al. 2005; Herczeg et al. 2009; this work) as well as a larger surface flux contribution from the magnetic, nonaccreting component on TW Hya.\\nocite{herczeg06,scholz05,herczeg09}  While these differences may reflect lower mass accretion rates in lower mass objects, the general trends connecting CTTSs and 2M1207 seem clear.  2M1207 is actively accreting from its disk, retains a warm (2500~--~4000 K) layer of H$_{2}$ in the inner disk, and shows evidence for depletion of Si and Mg into grains.  Given the edge-on geometry of the 2M1207 system, a more direct comparison would be to the edge-on CTTS DF Tau.  DF Tau was observed by COS as part of the $HST$ Cycle 17 Guaranteed Time program, and a comparison with 2M1207 will be presented in a future work.  If the additional H$_{2}$ component described in \\S4.2 is attributable to an outflow, a better comparison might be made with the edge-on CTTS system DG Tau. \\\\ \\\\ \\\\ \\\\"
    },
    "1004/1004.4306_arXiv.txt": {
        "abstract": "{} {Investigation of relationships between dust and gas, and study of the star formation law in M~31. } {We derive distributions of dust temperature and dust opacity across M\\,31 at 45$\\arcsec$ resolution using the Spitzer data. With the opacity map and a standard dust model we de-redden the H$\\alpha$ emission yielding the first de-reddened H$\\alpha$ map of M~31. We compare the emissions from dust, H$\\alpha$, HI and H$_2$ by means of radial distributions, pixel-to-pixel correlations and wavelet cross-correlations. We calculate the star formation rate and star formation efficiency from the de-reddened H$\\alpha$ emission. } {The dust temperature steeply decreases from 30\\,K near the center to 15\\,K at large radii. The mean dust optical depth at the H$\\alpha$ wavelength along the line of sight is about 0.7.  The radial decrease of the dust-to-gas ratio is similar to that of the oxygen abundance. Extinction is about linearly correlated with the total gas surface density within limited radial intervals. On scales $<$\\,2\\,kpc, cold dust emission is best correlated with that of neutral gas and warm dust emission  with that of ionized gas. H$\\alpha$ emission is slightly better correlated with  emission at 70\\,$\\mu$m than at 24\\,$\\mu$m. The star formation rate in M\\,31 is low. In the area 6\\,kpc\\,$<$\\,$R$\\,$<$\\,17\\,kpc, the total SFR is $\\simeq$\\,0.3\\,${\\rm M}_{\\odot} {\\rm yr}^{-1}$. A linear relationship exists between surface densities of SFR and H$_2$. The Kennicutt-Schmidt law between SFR and total gas has a power-law index of 1.30$\\pm$0.05  in the radial range of $R$\\,=\\,7-11\\,kpc increasing by about 0.3 for $R$\\,=\\,11-13\\,kpc. } { The better 70\\,$\\mu$m--H$\\alpha$ than 24\\,$\\mu$m--H$\\alpha$ correlation plus an excess in the 24\\,$\\mu$m/70\\,$\\mu$m intensity ratio indicates that other sources than dust grains, e.g. of stellar origin, contribute to the 24\\,$\\mu$m emission.  The lack of H$_2$ in the central region could be related to the lack of HI and the low opacity/high temperature of the dust. Since neither SFR nor SFE is well correlated with the surface density of H$_2$ or total gas, other factors than gas density must play an important role in  the formation of massive stars in M~31.  The molecular depletion time scale of 1.1\\,Gyr indicates that M~31 is about three times less efficient in forming young massive stars than M~33. ",
        "introduction": "Dust, neutral gas and ionized gas are the major components of the interstellar medium (ISM) in galaxies. Observations of their properties and inter-relationships can give important clues to the physics governing star formation. Relationships between components in the ISM are to be expected. Observations have shown that in the Galaxy dust and neutral gas are well mixed. In dense clouds of molecular gas mixed with cold dust most of the stars are formed. They subsequently heat the dust and gas in their surroundings and ionize the atomic gas. As the major coolants of the ISM are continuum emission and line emission at various frequencies, a close comparison of these emissions could shed light on spatial and physical connections between the emitting components. Present-day IR and radio telescopes have produced sensitive high resolution maps of several nearby galaxies, which are ideal laboratories to study the interplay between the ISM and star formation  \\citep[e.g. ][]{Kennicutt_07,Bigiel_08,Verley_09}.  The spiral galaxy nearest to us, the Andromeda Nebula (NGC224), is a highly inclined Sb galaxy of low surface brightness. Table 1 lists the positional data on M~31. Its proximity and large extent on the sky ($> 5^{\\circ} \\times 1^{\\circ}$) enable detailed studies of the ISM over a large radial range.  Surveys of M~31 at high angular resolution ($< 1\\arcmin$) are available at many wavelengths. In the HI line the galaxy was mapped by \\cite{Brinks} at $24\\arcsec \\times 36\\arcsec$ resolution, the northeastern half by \\cite{Braun_90} at 10$\\arcsec$ resolution and, most recently, the entire galaxy with high sensitivity by \\cite{Braun_09} at a resolution of 15$\\arcsec$. \\cite{Nieten} made a survey in the $^{12}$CO(1-0) line at a resolution of 23$\\arcsec$. \\cite{Devereux_etal_94b} observed M~31 in the H$\\alpha$ line to obtain the distribution of the ionized gas. The dust emission from M~31 was recently observed by the multiband imaging photometer Spitzer \\citep[MIPS, ][]{Rieke} with high sensitivity at 24\\,$\\mu$m, 70\\,$\\mu$m, and 160\\,$\\mu$m at resolutions~$\\leq 40\\arcsec$.  The relationships between gas and dust as well as between gas and star formation in M~31 have been studied in the past at resolutions of several arcminutes. \\cite{Walterbos_87} derived a nearly constant dust temperature across M~31 using the IRAS 100\\,$\\mu$m and 60\\,$\\mu$m maps. They also found a strong increase in the atomic gas- to-dust surface density ratio with increasing radius. This increase was confirmed by \\cite{Walterbos_88} who used optical extinction as dust tracer, and by \\cite{Nieten} using the ISO map at 175\\,$\\mu$m \\citep{Haas}. Interestingly, the latter authors did not find a radial increase in the molecular gas-to-dust ratio.  The dependence of star formation on HI surface density in M~31 has been studied by a number of authors \\citep{Emerson_74,Berkhuijsen_77,Tenjes, Unwin,Nakai_82,Nakai_84} using the number density of HII regions or of OB stars as star formation tracers. They obtained power-law indices between 0.5 and 2, possibly depending on the region in M~31, the star formation tracer and the angular resolution. \\cite{Braun_09} plotted the star formation density derived from the brightnesses at 8$\\mu$m, 24$\\mu$m and UV against the surface densities of molecular gas, HI and total gas, but did not fit power laws to their data.  The high-resolution data available for M~31 show the morphologies of the emission from dust and gas components in detail. We apply a 2-D wavelet analysis technique \\citep{Frick_etal_01} to the MIPS IR data \\citep{Gordon_06} and the gas (HI, H$_2$, and H$\\alpha$) maps to study the scale distribution of emission power and to separate the diffuse emission components from compact sources. We then compare the wavelet-decomposed maps at various spatial scales. We also use pixel-to-pixel (Pearson) correlations to derive quantitative relations not only between different ISM components but also between them and the present-day star formation rate.    Following \\cite{Walterbos_87} and \\cite{Haas}, we derive the dust temperature assuming a $\\lambda^{-2}$ emissivity law for the MIPS bands at which the emission from the big grains and hence the LTE condition is relevant, and present a map of the dust color temperature. We also obtain  the distribution of the optical depth and analyze the gas-to-dust surface-density ratio at a resolution of 45$\\arcsec$ (170~pc\\,$\\times$\\,660~pc  along the major and minor axis, respectively, in the galaxy plane), 9 times higher than before \\citep{Walterbos_87}. We use the optical depth map to de-redden the H$\\alpha$  emission observed by \\cite{Devereux_etal_94b} yielding the distribution of the absorption-free emission from the ionized gas, and use this as an indicator of massive star formation. We compare it with the distributions of neutral gas to obtain the dependence of the star formation rate on gas surface density.  The paper is organized as follows: The relevant data sets are described in Sect. 2. In Sect. 3 we derive maps of the dust color temperature and optical depth, and correct the H$\\alpha$ emission for absorption by dust. Radial profiles of the dust and gas emission and of the various gas-to-dust ratios are obtained in Sect. 4. Sect. 5 is devoted to wavelet decompositions and wavelet spectra of the dust and gas distributions, and their cross correlations. Complementary, we discuss in Sect. 6 classical correlations between gas and dust. In Sect. 7 the dependence of the star formation rate on the gas surface density is presented. Finally, in Sect. 8 we summarize our results. \\begin{table} \\begin{center} \\caption{Positional data adopted for M~31.} \\begin{tabular}{ l l } \\hline \\hline Position of nucleus    & RA\\,=\\,$00^{h}42^{m}45.97^{s}$      \\\\ \\,\\,\\,(J2000)  &  DEC\\,=\\,$41^{\\circ}16\\arcmin11.64\\arcsec$\\\\ Position angle of major axis   & 37$^{\\circ}$ \\\\ Inclination$^{1}$   & 75$^{\\circ}$  (0$^{\\circ}$=face on)\\\\ Distance$^{2}$   & 780$\\pm$40\\,kpc$^3$\\\\ \\hline \\noalign {\\medskip} \\multicolumn{2}{l}{$^{1}$ \\cite{Berkhuijsen_77} and \\cite{Braun_91}}\\\\ \\multicolumn{2}{l}{$^{2}$ \\cite{Stanek} }\\\\ \\multicolumn{2}{l}{$^{3}$ 1$\\arcmin$=\\,227$\\pm$12\\,pc along major axis}\\\\ \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "In this paper, we studied the emission from dust, neutral gas and ionized gas in the disk of M~31, and the relationships between these components on various linear scales. We compared the Spitzer MIPS maps at 24\\,$\\mu$m, 70\\,$\\mu$m and 160\\,$\\mu$m \\citep{Gordon_06} to the distributions of atomic gas seen in the HI line \\citep{Brinks}, molecular gas as traced by the $^{12}$CO(1-0) line \\citep{Nieten} and ionized gas observed in H$\\alpha$ \\citep{Devereux_etal_94b}. All data were smoothed to an angular resolution of 45$\\arcsec$ corresponding to 170\\,pc\\,$\\times$\\,660\\,pc in the plane of the galaxy.  For each of the dust and gas maps, we calculated the mean intensity distribution as a function of radius (Fig.~\\ref{fig:surfir}), separately for the northern and the southern half of M~31. Using wavelet analysis, we decomposed the dust and gas distributions in spatial scales and calculated cross-correlations as a function of scale. We also used classical correlations to derive quantitative relations between the various dust and gas components.  Using the MIPS 70\\,$\\mu$m and 160\\,$\\mu$m maps, we derived the distributions of the  dust temperature and optical depth. The dust optical depth at the  H$\\alpha$ wavelength was used to  a) investigate the dust-to-gas ratio, b) derive scaling relations between extinction and neutral gas emission, and c) de-redden the H$\\alpha$ emission in order to estimate the recent star formation rate. We also presented the Kennicutt-Schmidt law indices obtained  for the bright emission ring {\\bf near $R$=\\,10\\,kpc} in M~31. We summarize the main results and conclusions as follows.\\\\ \\\\ 1. Dust temperature and opacity:\\\\ $\\bullet$ The dust temperature steeply drops from about 30\\,K in the center to about 19\\,K near $R=$\\,4.5\\,kpc, and stays between about 17~K and 20~K beyond this radius (Fig.~2). The mean dust temperature in the area studied is about 18.5~K. This is 3~K less than the temperature obtained by \\cite{Walterbos_87} between the IRAS maps at 60\\,$\\mu$m and 100\\,$\\mu$m that both trace warmer dust than the MIPS maps at 70\\,$\\mu$m and 160\\,$\\mu$m used here.\\\\ $\\bullet$ The dust optical depth at H$\\alpha$ along the line of sight varies in a range between about 0.2 near the center and about 1 in the `10~kpc ring' (Fig.~4) with a mean value of 0.7$\\pm$0.4 (the error is standard deviation) and a most probable value of $\\simeq$\\,0.5, indicating that M~31 is mostly optically thin to the H$\\alpha$ emission. The total flux density of the H$\\alpha$ emission increases by 30\\% after correction for extinction.\\\\ \\\\ 2. Radial distributions:\\\\ $\\bullet$ The radial scale lengths between the maximum in the `10~kpc ring' and $R=$~15~kpc of the warm dust are smaller than that of the cold dust, as is expected if the warm dust is mainly heated by UV photons from star forming regions and cold dust by the ISRF. With the largest scale length, atomic gas has the largest radial extent of the dust and gas components considered here.\\\\ $\\bullet$ The radial gradient of the total gas-to-dust ratio is consistent with that of the oxygen abundance in M~31. The gas-to-dust ratios observed in the solar neighborhood \\citep{Bohlin} occur near $R=$~8.5~kpc in the disk of M~31 where N(gas)/$\\tau_{{\\rm H}\\alpha}$= $2.6 \\times 10^{21}$ at\\,cm$^{-2}$.\\\\ \\\\ 3. Properties as a function of scale:\\\\ $\\bullet$ Spatial scales larger than about 8~kpc contain most of the emitted power from the cold dust and the atomic gas, whereas the emissions from warm dust, molecular gas and ionized gas are dominated by scales near 1~kpc, typical for complexes of star forming regions and molecular clouds in spiral arms (Fig.~11). \\\\ $\\bullet$ Dust emission is correlated ($r_w \\ge 0.6$) with both neutral and ionized gas on scales~$>1$\\,kpc. \\\\ $\\bullet$ On scales\\,$<1$\\,kpc, ionized gas is best correlated with warm dust and neutral gas (both HI and H$_2$) with cold dust. On the smallest scale of 0.4\\,kpc, an HI--warm dust correlation hardly exists ($r_w \\simeq 0.4$) because not much HI occurs on the scale of star forming regions (see Fig.~11).\\\\ \\\\ 4. Relationships between gas and dust:\\\\ $\\bullet$ H$\\alpha$ emission is slightly better correlated with the emission at 70\\,$\\mu$m than at 24\\,$\\mu$m (Fig.~13, Table 6), especially on scales~$<$\\,2 kpc (Fig.~12). As in M~33 the 24\\,$\\mu$m--H$\\alpha$ correlation is best, this suggests that in early-type galaxies like M~31 the contribution from evolved AGB stars to the 24\\,$\\mu$m emission is larger than in late-type galaxies like M~33.\\\\ $\\bullet$ Dust extinction A$_{H\\alpha}$ is not well correlated with N(2H$_2$) indicating that dust mixed with molecular clouds does not contribute much to the total extinction. Although the correlation with N(HI) is better, A$_{H\\alpha}$ is best correlated with N(HI+2H$_2$).\\\\ $\\bullet$ Dust opacity is proportional to the square root of N(2H$_2$) but about linearly related to N(HI), as was also found by \\cite{Nieten} at 90$\\arcsec$ resolution. This is an indirect indication of a balance between the formation and destruction rates of H$_2$ in cool, dusty HI clouds.\\\\ $\\bullet$ In the central 2 kpc both the dust opacity and the HI column density are very low and the dust temperature is high. This combination may explain the lack of H$_2$ in this region.\\\\ \\\\ 5. SFR and SFE:\\\\ $\\bullet$ The SFR in M~31 is low. The total SFR in the observed field between $R=$~6~kpc and $R=$~17~kpc is $0.27\\,{\\rm M}_{\\odot} {\\rm yr}^{-1}$ and the star formation efficiency is 0.9\\,Gyr$^{-1}$, yielding a molecular depletion time scale of 1.1~Gyr. This is about three times longer than observed in the northern part of M~33 (Gardan et al. 2007). The radial distribution of $\\Sigma_{\\rm SFR}$ in 0.5\\,kpc-wide rings in the plane of the galaxy (Fig.~17) varies between about 0.1 and 1\\,M$_{\\odot}$\\, Gyr$^{-1}$\\, pc$^{-2}$, values that are about 10 times smaller than in  the northern part of M~33 \\citep{Gardan}.  Between $R=$~6~kpc and $R=$~15~kpc, SFE varies between about 0.5\\,Gyr$^{-1}$ and 1.5\\,Gyr$^{-1}$, whereas the efficiency with respect to the total gas surface density slowly decreases from about 0.18\\,Gyr$^{-1}$ to about 0.03\\,Gyr$^{-1}$. \\\\ $\\bullet$ SFR is not well correlated with  neutral gas and worst of all with molecular gas in the radial range $30\\arcmin-50\\arcmin$ containing the `10\\,kpc ring' (Fig.~15, Table 7). In spite of this, the power-law fits are statistically significant. We find a linear relationship between the surface densities of SFR and molecular gas (power-law exponent 0.96~$\\pm$~0.03),  and a power law with index 1.30~$\\pm$~0.05 between the surface densities of SFR and total gas. These results agree with the average relationship for 7 nearby galaxies much brighter than M~31 \\citep{Bigiel_08}. While in these galaxies molecular hydrogen is the dominant gas phase, most of the neutral gas in M~31 is atomic. Thus, the surface density of SFR depends linearly on that of molecular gas irrespective of the fraction of molecular gas or the absolute value of the total gas surface density in a galaxy.   Some important implications of this study are: \\begin{itemize} \\item[-] Precaution is required in using the total IR luminosity (TIR) as an indicator of recent SFR or to derive dust opacity for an early-type galaxy like M~31, because the cold dust is mainly heated by the ISRF and the warm dust emission at 24\\,$\\mu$m is partly due to evolved stars (especially in the bulge of the galaxy).  \\item[-] Neither the present-day SFR nor SFE is well correlated with the surface density of molecular gas or total gas. Therefore, other factors than gas density must play an important role in the process of star formation in M~31. \\end{itemize}"
    },
    "1004/1004.0239_arXiv.txt": {
        "abstract": "We present the historic photographic light curves of three little known Blazars (two BL Lacs and one FSRQ), BZB J1058+5628, BZQ J1148+5254 and BZB J1209+4119 spanning a time interval of about 50 years, mostly built using the Asiago plate archive.  All objects show evident long-term variability, over which short-term variations are superposed. One source, BZB J1058+5628, showed a marked quasi-periodic variability of 1 mag on time scale of about 6.3 years, making it one of the few BL Lac objects with a quasi- periodic behavior. ",
        "introduction": "Long-term ($\\geq$ 20 years) optical light curves are still available only for a relatively small number of Active Galactic Nuclei (AGN) and Blazars. Pioneering work in this field was made by several authors (Pica et al. 1988; Webb et al. 1988) on relatively large samples of sources over a time window of about 20 years. For some sources these curves show long-term trends occurring over time scales of decades, while others show only short-term variability.  A few sources had dedicated papers to their long-term optical variability. Among the best studied sources we recall: OQ~530, whose optical brightness faded at 0.035 mag/yr for about one century (Massaro et al. 2004); S5~0716+71, showing a monotonic brightening trend of 0.11 mag/yr of its mean luminosity over the last 40 years (Nesci et al. 2005); ON~231, with a long term decreasing trend of 0.023 mag/yr followed by an increasing one of 0.07 mag/yr (Massaro et al. 2001); WGA 0447.9-0322, with a long monotonic trend of 0.11 mag/yr, similar to S5~0716+71 (Nesci et al. 2007);  5C 3.178 declining at 0.03 mag/yr over 30 years (Sharov 1995).  In some cases a periodicity of the outbursts has also been found, the best case by far being OJ~287 (Sillampaa et al 1988, Valtonen et al. 2009). Other such sources are AO~0235+164, with a possible 5 years period of strong outbursts (Raiteri et al. 2008 and references therein) and S5~0716+71 (Raiteri et al. 2003) with recurrence of about 3 years. Also in the radio band a few sources have shown periodic outbursts (e.g. 3C 454.3, Ciaramella et al. 2004, Qian et al. 2009).  The origin of the fast variations in Blazars is generally explained by the relativistic boosting of perturbations moving down a jet pointing close to the line of sight. The relevant quantity for the boosting is the beaming factor $\\delta = 1 \\times (\\Gamma \\times (1-\\beta~cos\\theta)^{-1}$, where $\\Gamma$ and $\\beta$ are the Lorentz factor and the velocity (in units of the speed of light) of the perturbations' bulk motion and $\\theta$ is the angle between the jet and the line of sight.  On the other hand the nature of secular variations is unclear. A suggestive possibility is that they can be associated with changes in the structure and/or direction of the inner jet (see e.g. Kadler et al. 2006). It is difficult, however, to obtain a clear evidence of such changes because it requires long and accurate multifrequency campaigns  with VLBI angular resolution on a sample of several sources: such a study has been done e.g. by Nesci et al. (2005) in the case of S5~0716+71, suggesting that precession of the jet may indeed explain its observed light curve behavior, or Massaro et al. (2004) in the case of OQ~530.  With the successful launch of the Fermi satellite for Gamma ray astronomy, an all-sky monitoring of the Blazars emission has started, which will probably bring to the attention of astronomers a number of poorly known sources. Lists of potential Gamma ray Blazars have been prepared in the framework of the GLAST project by Massaro et al. (2009), containing about 2800 objects, and Sowards-Emmerd et al. (2005), listing about 770 Northern sky sources. For most of these sources very few data exist, basically those allowing their detection and the classification as an AGN. A better knowledge of the properties of these sources will be useful for the interpretation of the Gamma-Ray data now available from the Fermi mission.  To determine the historic light curve of AGNs, an effective way is to use survey plates taken with wide angle instruments, like the Schmidt telescopes, in fields covered over a large time span for patrol of other targets, like Supernovae or variable stars. A good mine of such material is the archive of the Asiago Observatory (http://dipastro.pd.astro.it/asiago/), with its two Schmidt instruments, the 67/92cm, operative between 1965 and 1998, and the 40/50 cm, operative between 1958 and 1992.  We selected therefore, from the Massaro et al. (2009) catalogue, those sources for which no historic optical light curve is still published, nominally bright enough to be well measurable on the plates of the 67/92 cm telescope (B$\\le$17.5), without an obvious strong host galaxy around, and for which a large ($\\ge$50) number of plates is available spread over a long ($\\ge$10 years) time interval, so that a meaningful historic light curve can be derived. Unfortunately, a very small number of sources matched these conditions, mainly due to their optical faintness.  In this paper we report the results of our photometric measurements of these plates and the first optical historic light curve for three sources, BZBJ1058+5611, BZQJ1148+5254 and BZBJ1209+4119. ",
        "conclusions": "Pica et al. (1988) and Webb et al. (1988) studied the long-term behavior of several tens of AGNs with photographic plates for 15 to 20 years. For 61 sources they had enough data to morphologically classify their light curve in four types: Class I flickering without long term trend; Class II long term trend with small flickering; Class III long term and short term variability of comparable amplitude; Class IV rare outburst with stable flux level. In the Blazar sample of 22 sources of Webb et al. (1988) there is no significant difference in the frequency of the light curve classes between FSRQ and BL Lac objects, while a marked difference between Blazars and QSO exists in the Pica et al. (1988) sample of 39 sources, which includes a good number of radio steep spectrum Quasars and radio quiet QSO.  To have some physical insight of our sources we report in Table 7 some basic data: column 1 is the name, column 2 the Log(power) at 1.4 GHz, column 3 the absolute R magnitude computed from the USNO B1 catalogue magnitude and literature redshift, column 4 the Radio/Optical flux ratio, column 5 the NIR spectral slope from the JHK magnitudes in the 2MASS catalogue, column 6 the optical spectral slope from the SDSS data. We remark that, at variance with the other two sources, the spectral slopes computed for BZQJ1148+52 have a very poor \\chr ~and are therefore marked with \":\". Actually the spectrum of this source cannot be well fitted with a power law in neither of the two explored ranges, probably due to the strong emission lines (e.g. C IV equivalent width is 77 \\AA) and of the UV bump which falls in the optical due to the source redshift.  It is apparent from this Table that our three sources are flat-spectrum radio-loud objects (Radio/Optical flux ratio $\\ge$10) but have different absolute luminosities and show substantially different behaviors in their optical light curve. The strong-lined object is the brightest both in the radio and optical bands. Only BZBJ1058+56 was detected in Gamma-rays by Fermi-LAT (Abdo et al. 2009) and possibly also by EGRET (Bloom et al. 2000).  From the point of view of the overall Spectral Energy Distribution, a much used tool for the classification of Blazars is the $\\alpha_{ro}-\\alpha_{ox}$ diagram (Padovani and Giommi 1995). On this diagram the Blazars mainly occupy two areas: a horizontal branch and a diagonal branch; a diagonal line of negative slope -1 is a line of constant Radio/X-ray flux ratio. The line at $\\alpha_{rx}$=0.75 is the formal border between HBL and LBL sources  (Padovani and Giommi 1995). Extreme HBL sources are located at the left side of the horizontal branch, extreme LBL at the upper side of the diagonal branch.  For a Synchrotron Self Compton emission model, as the peak of the synchrotron emission of a Blazar moves from lower ($10^{13}$Hz) to higher ($10^{17}$Hz) frequencies the location of the source on this diagram moves from the upper left corner to the lower right one along the diagonal line and then back to the left along the horizontal branch.  We report in Fig. 7 this diagram for the sources in the Roma BZCat, with the positions of our three objects marked. BZBJ1209+41 and BZQJ1148+52 have $\\alpha_{rx}$ larger than 0.75, and are located in the diagonal branch, with BZBJ1209 being the most radio loud. BZBJ1058+56 is already on the horizontal branch and is an HBL, as discussed by Donato et al. (2005) also on the basis of BeppoSAX X-ray spectra. None of them is however an extreme case.  The long-term optical light curves of our three sources are rather different. As a general remark, the overall amplitude variability is anticorrelated with the intrinsic power.  BZBJ1058+56 showed regular oscillations of about 1 mag amplitude, with timescale of $\\sim$2300 days, over a monotonic decreasing trend of 0.07 mag/year; it can be classified as Class III and is the source with the larger variability.  Its historic light curve contains 5 outbursts sampled by the Asiago plate archive; this source seems therefore an interesting case of quasi-periodic BL Lac. Further multiwavelength monitoring of this source, should be performed.  BZQJ1148+52 showed a monotonic decreasing trend with a slope similar to 1058+56 but without the oscillating behavior: the detected short-term variability is comparable to our photometric uncertainty, so no firm conclusions can be derived on their time scale. It can be put into Class II. It showed a substantially smaller variability than the other two BL Lacs, both on short and long time scales.  BZBJ1209+41 showed a slight increasing trend (0.04 mag/year) with large dips in its light curve: it is therefore quite unusual (the opposite of Class I) and deserves further monitoring. Unfortunately it is not bright enough to be easily followed with small telescopes.  When a time interval of about 50 years in considered, all the long term trends detected in the time window of 27 years sampled by the Asiago plates do not seem to hold, so that they might be considered just as part of longer variability trends.  Which processes can be behind these secular trends? Both physical processes and geometrical effects can be at work: in the first case one can imagine a monotonic variation of the number of radiating electrons, or of the average ambient magnetic field; in the second case, a change in the Doppler boosting factor along our line of sight due to the jet precession. The latter possibility can be considered as an indicator of a massive black hole binary system in the nuclear region (see e.g. Romero et al. 2003).  If we interpret the long term trends of our three sources as due to a slow precession of the jet, as in was supposed to be in the case of S5 0716+71 (Nesci et al. 2005), then the periodicity should be of several 10$^4$ days and therefore comparable to (or even larger than) the human lifetime. This poses a strong challenge because observations must be accumulated for several tens of years before any firm conclusion can be reached. A further difficulty for the data interpretations, if the monitoring is not dense enough, could be the occurrence of fast and/or large occasional outbursts/dips, which can mask the long-term trends.  A strong support to confirm the precession model could come from imaging at high radio frequencies with VLBI techniques, which could detect monotonic variations in the jet direction and/or at lower frequencies showing residuals of radio emission in regions involved by the crossing of the jet in the past (see e.g. Massaro et al. 2004).  Finally, we remark that the detection of the quasi-periodicity of BZB J1058+56 from the Asiago plates suggests that further discoveries could be made using other, still unexplored, photographic plate archives."
    },
    "1004/1004.5152.txt": {
        "abstract": "Scattered lights from terrestrial exoplanets provide valuable information about the planetary surface. Applying the surface reconstruction method proposed by \\cite{Fujii2010} to both diurnal and annual variations of the scattered light, we develop a reconstruction method of land distribution with both longitudinal and latitudinal resolutions. We find that one can recover a global map of an idealized Earth-like planet on the following assumptions: 1) cloudless, 2) a face-on circular orbit, 3) known surface types and their reflectance spectra 4) no atmospheric absorption, 5) known rotation rate 6) static map, and 7) no moon. Using the dependence of light curves on the planetary obliquity, we also show that the obliquity can be measured by adopting the $\\chi^2$ minimization or the extended information criterion. We demonstrate a feasibility of our methodology by applying it to a multi-band photometry of a cloudless model Earth with future space missions such as the occulting ozone observatory (O3). We conclude that future space missions can estimate both the surface distribution and the obliquity at least for cloudless Earth-like planets within 5 pc. ",
        "introduction": "Recent progress in observational techniques has revealed various physical properties of exoplanets beyond orbital parameters and planetary mass. Detections of the atmospheric components have been reported for several systems using spectroscopy at the planetary transit and secondary eclipse \\citep[e.g.][]{charbonneau2002, vidal2003, vidal2004, tinetti2007, swain2008, swain2009}. Interior compositions can be inferred from planetary mass and radius \\citep[e.g.][]{2009A&A...506..287L, 2009Natur.462..891C}. Constructions of thermal maps of the planetary atmosphere have been proposed by \\citep{2006ApJ...649.1020W, 2008ApJ...678L.129C}. A longitudinal thermal map of HD 189733b has been constructed by \\cite{2007Natur.447..183K} based on the method proposed by \\cite{2008ApJ...678L.129C}.  Nevertheless, an identification of planetary surface components still remains an ambitious challenge. One of the promising approaches is to use the scattered light of exoplanets through the direct imaging observation \\citep[e.g.][]{2000ApJ...540..504S, 2001Natur.412..885F, 2005ApJ...627..520S}. \\cite{2001Natur.412..885F} focused on the inhomogeneity of the Earth surface which causes diurnal variation of the scattered light. They computed the scattered light from a model Earth observed at a distance of 10 pc and showed that time variations of the scattered light in different photometric bands highly depend on the geological and biological features on the planetary surface such as ocean, land, and even vegetation. More detailed characterizations (including spectroscopy) of the scattered light of the Earth and its time variations are discussed both via Earth-shine observations \\citep[e.g.][]{woolf2002, arnold2002,2006ApJ...651..544M} and simulations \\citep{2006AsBio...6...34T, 2006AsBio...6..881T, 2006ApJ...651..544M}. These studies have suggested a future possibility to investigate the surface of Earth-like exoplanets by the scattered light curves. We note that such time variations of the scattered light are also applicable to determine the rotation period from as shown by \\cite{2008ApJ...676.1319P}.  A variety of inversion techniques of the planetary surface from the scattered light curves have been proposed. Surprisingly, the first theoretical study of scattered light curves to make albedo maps has been performed at the beginning of the twentieth century although the author assume asteroid and satellites for the target \\citep{1906ApJ....24....1R}. \\citet{2009ApJ...700..915C} performed principal component analysis (PCA) on multi-band photometric data of the Earth observed by EPOXI (Extrasolar Planet Observation and Deep Impact Extended Investigation) mission, and extracted spectral features which roughly correspond to land and ocean. They also checked the time variation of these components and translated it to the longitudinal distribution of these components based on  the formulation by \\citet{2008ApJ...678L.129C}. \\cite{2009ApJ...700.1428O} paid attention to the gap of reflectivity between ocean and land, and reproduced a longitudinal map of land. \\cite{Fujii2010} (hereafter F10) have developed a methodology to estimate the areas of ocean, soil, vegetation, and snow from multi-band photometry, and showed that the area of these components can be recovered from mock observations of a cloudless Earth.  Since these authors focused on diurnal variations in mapping the surface, the resultant maps have only longitudinal resolution with little information of the latitudinal distribution.  One of the goals of the present paper is to develop a method to map inhomogeneous surfaces of exoplanets with both longitudinal and latitudinal resolutions using both diurnal and annual variations of the scattered lights.  We also consider the determination of the planetary obliquity from time variation of planetary light. The obliquity is an important property of Earth-like planets with its strong implications for climate, habitability \\citep[{\\it e.g.},][]{1997Icar..129..254W, 2003IJAsB...2....1W} and planetary formation. N-body simulations of the final stage of terrestrial planet formation indicated that the distribution of the obliquity $\\zeta$ is isotropic \\citep{1999Icar..142..219A,2001Icar..152..205C,2007ApJ...671.2082K}. \\cite{2004NewA...10...67G} modeled the infrared light curves of exoplanets and showed how the obliquity affects an annual variation. \\citet{2009ApJ...700.1428O} also pointed out a possibility to determine the planet's obliquity. In this paper, we demonstrate that the obliquity is estimated simultaneously  with the global map of the planet by analyzing the scattered light curves over its orbital period.  The rest of the paper is organized as follows. We first review the estimation method of the weighted area from multi-band photometry proposed by F10, and describe our methodology to reconstruct the planetary surface and the obliquity measurement in \\S 2. Assuming a future satellite mission for the direct imaging of Earth-like planets, we apply our methodology to mock observations based on real data of the scattering properties of the Earth in \\S 3. Finally we summarize our results in \\S 4. ",
        "conclusions": "We have developed the reconstruction method of the two-dimensional planetary surface via diurnal and annual variation of the scattered light. Applying the method to the mock photometric data, we have demonstrated that our method works for the mock Earth model, while this model has a lot of simplifying assumptions as follows: 1) cloudless, 2) a face-on circular orbit, 3) known reflectance spectra 4) no atmospheric absorption, 5) known rotation rate 6) static map, and 7) no moon. We also found that the planetary obliquity can be estimated by this method. With our method, future satellite missions such as the occulting ozone observatory \\citep{kasdin2010} might provide ``a global map'' of Earth-like exoplanets. While only the terrestrial planets have been considered in this paper, our method might be applicable to any planets with an inhomogeneous surface, including Jupiter-like exoplanets.  In this paper, we ignored the effect of clouds and expected that clouds affect the estimation as like a statistical noise because of relatively short time variation of clouds. The PCA performed by \\cite{2009ApJ...700..915C} is one of promising approaches because they could separate the land and ocean compositions even though they used the EPOXI data that contains the cloud effect. The effect of clouds is discussed elsewhere (Fujii et al. in preparation).\u00a1\u00a1We also assumed a face-on circular orbit in this paper.  This assumption might be too severe for practical applications. We will generalize our method in the next paper."
    },
    "1004/1004.2520_arXiv.txt": {
        "abstract": "In this work, we find exact gravastar solutions in the context of noncommutative geometry, and explore their physical properties and characteristics. The energy density of these geometries is a smeared and particle-like gravitational source, where the mass is diffused throughout a region of linear dimension $\\sqrt{\\alpha}$ due to the intrinsic uncertainty encoded in the coordinate commutator. These solutions are then matched to an exterior Schwarzschild spacetime. We further explore the dynamical stability of the transition layer of these gravastars, for the specific case of $\\beta=M^2/\\alpha<1.9$, where $M$ is the black hole mass, to linearized spherically symmetric radial perturbations about static equilibrium solutions. It is found that large stability regions exist and, in particular, located sufficiently close to where the event horizon is expected to form. ",
        "introduction": "Recently, an alternative picture for the final state of gravitational collapse has emerged \\cite{Mazur}. The latter, denoted as a gravastar ({\\it grav}itational {\\it va}cuum {\\it star}), consists of an interior compact object matched to an exterior Schwarzschild vacuum spacetime, at or near where the event horizon is expected to form. Therefore, these alternative models do not possess a singularity at the origin and have no event horizon, as its rigid surface is located at a radius slightly greater than the Schwarzschild radius. More specifically, the gravastar picture, proposed by Mazur and Mottola \\cite{Mazur}, has an effective phase transition at/near where the event horizon is expected to form, and the interior is replaced by a de Sitter condensate. This new emerging picture consisting of a compact object resembling ordinary spacetime, in which the vacuum energy is much larger than the cosmological vacuum energy, is also denoted as a ``dark energy star'' \\cite{Chapline}. In fact, a wide variety of gravastar models have been considered in the literature \\cite{gravastar1,gravastar2} and their observational signatures have also been explored \\cite{gravastar3}. In this work, we consider a further extension of the gravastar picture in the context of noncommutative geometry. The dynamical stability of the transition layer of these gravastars to linearized spherically symmetric radial perturbations about static equilibrium solutions is also explored. The analysis of thin shells \\cite{linear-thinshell} and the respective linearized stability analysis of thin shells has been recently extensively considered in the literature, and we refer the reader to Refs. \\cite{linearstability,linear-WH} for details.  In the context of noncommutative geometry, an interesting development of string/M-theory has been the necessity for spacetime quantization, where the spacetime coordinates become noncommuting operators on a $D$-brane \\cite{Witten}. The noncommutativity of spacetime is encoded in the commutator $\\left[ \\mathbf{x}^{\\mu},\\mathbf{x}^{\\nu}\\right] =i\\,\\theta^{\\mu\\nu}$, where $\\theta^{\\mu\\nu}$ is an antisymmetric matrix which determines the fundamental discretization of spacetime. It has also been shown that noncommutativity eliminates point-like structures in favor of smeared objects in flat spacetime \\cite{Smailagic:2003yb}. Thus, one may consider the possibility that noncommutativity could cure the divergences that appear in general relativity. The effect of the smearing is mathematically implemented with a substitution of the Dirac-delta function by a Gaussian distribution of minimal length $\\sqrt{\\alpha}$. In particular, the energy density of a static and spherically symmetric, smeared and particle-like gravitational source has been considered in the following form \\cite{Nicolini:2005vd} \\begin{equation} \\rho_{\\alpha}(r)=\\frac{M}{(4\\pi\\alpha)^{3/2}}\\;\\mathrm{exp}\\left( -\\frac{r^{2}}{4\\alpha}\\right)  \\,, \\label{NCGenergy} \\end{equation} where the mass $M$ is diffused throughout a region of linear dimension $\\sqrt{\\alpha}$ due to the intrinsic uncertainty encoded in the coordinate commutator.  The Schwarzschild metric is modified when a non-commutative spacetime is taken into account \\cite{Nicolini:2005vd, Esposito}. The solution obtained is described by the following spacetime metric \\begin{equation} ds^{2}=-f(r)\\,dt^{2}+\\frac{dr^{2}}{f(r)}+r^{2} \\,(d\\theta^{2}+\\sin^{2}{\\theta }\\,d\\phi^{2})\\,,\\label{NCW} \\end{equation} with $f(r)=1-2m(r)/r$, where the mass function is defined as \\begin{equation} m(r)=\\frac{2M}{\\sqrt{\\pi}}\\gamma\\left( \\frac{3}{2},\\frac {r^{2}}{4\\alpha}\\right)  \\,, \\label{massfunction} \\end{equation} and \\begin{equation} \\gamma\\left(  \\frac{3}{2},\\frac{r^{2}}{4\\alpha}\\right) =\\int\\limits_{0} ^{r^{2}/4\\alpha}dt\\sqrt{t}\\exp\\left(-t\\right) \\end{equation} is the lower incomplete gamma function \\cite{Nicolini:2005vd}. The classical Schwarzschild mass is recovered in the limit $r/\\sqrt{\\alpha}\\rightarrow\\infty$. It was shown that the coordinate noncommutativity cures the usual problems encountered in the description of the terminal phase of black hole evaporation. More specifically, it was found that the evaporation end-point is a zero temperature extremal black hole and there exist a finite maximum temperature that a black hole can reach before cooling down to absolute zero. The existence of a regular de Sitter at the origin's neighborhood was also shown, implying the absence of a curvature singularity at the origin. Recently, further research on noncommutative black holes has been undertaken, with new solutions found providing smeared source terms for charged and higher dimensional cases \\cite{newNCGbh}. Furthermore, exact solutions of semi-classical wormholes \\cite{Garattini2} in the context of noncommutative geometry were found \\cite{Garattini1}, and their physical properties and characteristics were analyzed.   This paper is outlined in the following manner. In Section \\ref{sec:II}, we present the generic structure equations of gravastars, and specify the mass function in the context of noncommutative geometry. In Section \\ref{sec:III}, the linearized stability analysis procedure is briefly outlined, and the stability regions of the transition layer of gravastars are determined. Finally in Section \\ref{sec:IV}, we conclude. We adopt the convention $G=c=1$ throughout this work. ",
        "conclusions": "\\label{sec:IV}  In this work, we have found exact gravastar solutions in the context of noncommutative geometry, and briefly explored their physical properties and characteristics. The energy density of these geometries is a smeared and particle-like gravitational source, where the mass is diffused throughout a region of linear dimension $\\sqrt{\\alpha}$ due to the intrinsic uncertainty encoded in the coordinate commutator.  We further explored the dynamical stability of the transition layer of these dark energy stars to linearized spherically symmetric radial perturbations about static equilibrium solutions. It was found that large stability regions do exist, which are located sufficiently close to where the event horizon is expected to form, so that it would be difficult to distinguish the exterior geometry of the gravastars, analyzed in this work, from a black hole."
    },
    "1004/1004.2699_arXiv.txt": {
        "abstract": "{ Observations of $^{12}$CO at high redshift indicate rapid metal enrichment in the nuclear regions of at least some galaxies in the early universe. However, the enrichment may be limited to nuclei that are synthesized by short-lived massive stars, excluding classical ``secondary'' nuclei like $^{13}$C. Testing this idea, we used the IRAM Interferometer to tentatively detect the $^{13}$CO $J$=3$\\rightarrow$2 line at a level of 0.3\\,Jy\\,km\\,s$^{-1}$ toward the Cloverleaf Quasar at $z$ = 2.5. This is the first observational evidence for $^{13}$C at high redshift. The $^{12}$CO/$^{13}$CO $J$=3$\\rightarrow$2 luminosity ratio is with 40$^{+25}_{-8}$ much higher than ratios observed in molecular clouds of the Milky Way and in the ultraluminous galaxy Arp220, but may be similar to that observed toward NGC~6240. Large Velocity Gradient models simulating seven $^{12}$CO transitions and the $^{13}$CO line yield $^{12}$CO/$^{13}$CO abundance ratios in excess of 100. It is possible that the measured ratio is affected by a strong submillimeter radiation field, which reduces the contrast between the $^{13}$CO line and the background. It is more likely, however, that the ratio is caused by a real deficiency of $^{13}$CO. This is already apparent in local ultraluminous galaxies and may be even more severe in the Cloverleaf because of its young age ($\\la$2.5\\,Gyr). A potential conflict with optical data, indicating high abundances also for secondary nuclei in quasars of high redshift, may be settled if the bulk of the CO emission is originating sufficiently far from the active galactic nucleus of the Cloverleaf. ",
        "introduction": "There is evidence for solar or super-solar metallicities in the circumnuclear environments of quasars out to redshifts $z$$>$4 (e.g., Hamann \\& Ferland 1999; Kurk et al. 2007; Jiang et al. 2007; Juarez et al. 2009; Matsuoka et al. 2009). This evidence, mainly from optical lines, is supported by millimeter detections of CO and dust in high-redshift sources, indicating rapid metal enrichment due to starbursts in the circumnuclear regions of at least some galaxies in the early universe (e.g., Solomon \\& Vanden Bout 2005). This enrichment, however, might apply mainly to atomic nuclei that are synthesized in short-lived massive stars, and not so much to ``secondary'' nuclei like $^{13}$C that are thought to be mainly synthesized in longer-lived, less-massive stars (but see, e.g., Hamann et al. 2002 for the mainly secondary element nitrogen).  In the local universe, $^{12}$C/$^{13}$C abundance ratios are sometimes considered to be a diagnostic of deep stellar mixing and a measure of ``primary'' vs.\\ ``secondary'' nuclear processing (e.g., Wilson \\& Rood 1994). While $^{12}$C is produced by He burning on rapid time scales in massive stars, $^{13}$C is mainly synthesized by CNO processing of $^{12}$C seed nuclei from earlier stellar generations. This processing occurs more slowly, during the red giant phase in low- and intermediate-mass stars or novae. The $^{12}$C/$^{13}$C ratio may therefore depend on the nucleosynthesis history. It could be much higher in high-$z$ galaxies that are too young to have synthesized large amounts of secondary nuclei like $^{13}$C.  At optical, near-IR, and UV wavelengths it is difficult to discriminate between an element's isotopes because their atomic lines are blended (e.g., Levshakov et al. 2006). The prospects are better with radio lines from isotopic substitutions in molecules, which are well separated by a few percent of their rest frequency from the main species. This separation allows both the main and rare species to be easily identified, and to be observable with the same radio receivers and spectrometers.  The Cloverleaf Quasar (H1413+117), partly because of amplification by gravitational lensing, is a high-$z$ source with exceptional peak flux densities in $^{12}$C$^{16}$O (hereafter $^{12}$CO; see Appendix~2 of Solomon \\& Vanden Bout 2005). This source is therefore one of the best candidates to search for $^{13}$C$^{16}$O (hereafter $^{13}$CO) to try to test models of ``chemical'' evolution over a Hubble time. In this paper we report on a search for $^{13}$CO(3--2) emission in the Cloverleaf at $z$=2.5579, when the universe was 2.5\\,Gyr old. ",
        "conclusions": "In order to further evaluate our observational result, we have to discuss the correlation between molecular $^{12}$CO/$^{13}$CO and atomic $^{12}$C/$^{13}$C abundance ratios and to summarize relevant observational data from low-redshift galaxies which are, like the Cloverleaf, ultraluminous in the infrared. Finally, we will address some fundamental problems, which are related to the still poorly known morphology of the gas surrounding the Cloverleaf QSO.   \\subsection{Chemical fractionation and isotope selective photodissociation}  Observed isotope ratios may be affected by fractionation. The $^{12}$CO/$^{13}$CO abundance ratio is likely influenced by the reaction $$ ^{13}{\\rm C}^+ + ^{12}{\\rm CO} \\rightarrow\\ ^{12}{\\rm C}^+ + ^{13}{\\rm CO} + \\Delta E_{\\rm 35K} $$ (Watson et al. 1976). The process enhances $^{13}$CO relative to $^{12}$CO in the more diffuse C$^+$ rich parts of molecular clouds. This may be compensated by isotope selective photodissociation. $^{12}$CO and $^{13}$CO need similar amounts of self-shielding to survive in a hostile interstellar environment. This favors the more abundant isotopologue (e.g., Sheffer et al. 2007).  For the Galaxy, such effects can be quantified. Milam et al. (2005) summarized $^{12}$C/$^{13}$C ratios from the galactic disk, obtained with the three molecules CO, CN, and H$_2$CO. These molecular species are synthesized by quite different chemical reactions. The good agreement between their $^{12}$C/$^{13}$C ratios and a lack of correlation with kinetic temperature suggests that chemical fractionation as well as isotope selective photodissociation do not greatly affect the determined isotope ratios.  Whether this result is also valid in the case of the Cloverleaf QSO may not be obvious at first sight. The ultraviolet radiation field in the vicinity of the quasar might be exceptionally strong, favoring $^{12}$CO over $^{13}$CO and thus leading to an enhanced molecular abundance ratio with respect to $^{12}$C/$^{13}$C. However, such a scenario is not likely.  Firstly, most of the galactic data were obtained toward prominent sites of massive star formation, where the UV radiation field is also exceptionally intense. Secondly, judging from C{\\sc i}, in the Cloverleaf the excitation of the molecular gas is intermediate between conditions found for the starburst galaxy M\\,82 ($T_{\\rm ex,CI}$ $\\sim$50\\,K) and the central region of the Milky Way ($T_{\\rm ex,CI}$ $\\sim$ 22\\,K) (Stutzki et al. 1997; Wei{\\ss} et al. 2003). Thirdly, polycyclic aromatic hydrocarbon (PAH) features are as strong as expected with respect to the far infrared luminosity when compared with more nearby ultraluminous star-forming galaxies, favoring ``normal'' conditions and a predominantly starburst nature of the Cloverleaf's huge FIR emission (Lutz et al.  2007). Finally, the CO emission from the Cloverleaf appears to be more extended than the effective radius out to which the quasar could dominate the UV field.  Modeling both the source and the lens of the Cloverleaf QSO, Venturini \\& Solomon (2003) find a characteristic radius of $r$ $\\sim$ 800\\,pc for the CO $J$=7--6 line, which is higher excited and thus possibly less widespread than the $J$=3$\\rightarrow$2 transition considered here. If the Cloverleaf's intrinsic far infrared luminosity ($L_{\\rm FIR}$ $\\sim$ 5$\\times$10$^{12}$\\,L$_{\\odot}$, Lutz et al. 2007) would entirely originate from 6.2--13.6\\,eV photons emitted by the active nucleus, we would obtain, at a radius of 800\\,pc, a UV photon illumination of $\\chi$ $\\sim$ 10$^5$\\,$\\chi_0$ with respect to the local galactic radiation field, $\\chi_0$ = 2$\\times$10$^{-4}$\\,\\,erg\\,cm$^{-2}$\\,s$^{-1}$\\,sr$^{-1}$ (see Draine 1978). The Cloverleaf QSO is a Broad Absorption Line (BAL) quasar which permits at least a partial view onto its nuclear engine. Therefore, taking the Cloverleaf's UV luminosity from Fig.\\,1 of Barvainis et al. (1995) and accounting for a gravitational amplification by a factor of 11 (Solomon \\& Vanden Bout 2005), we obtain accordingly $\\chi$ $\\sim$ 2.5$\\times$10$^4$\\,$\\chi_0$. Both $\\chi$ values are consistent with those encountered in prominent galactic sites of massive star formation and may be upper limits if the Cloverleaf posseses a self-shielding rotating disk. To summarize, physical conditions in the Cloverleaf host galaxy appear to be sufficiently normal so that the $^{12}$C/$^{13}$C isotope ratio should not strongly deviate from the $^{12}$CO/$^{13}$CO molecular abundance ratio.   \\subsection{$^{12}$CO/$^{13}$CO ratios in $z$$<$1 galaxies}  {\\it In our Galaxy}, the $^{12}$CO/$^{13}$CO line intensity ratios from molecular clouds are typically about 5, probably corresponding to true $^{12}$C/$^{13}$C abundance ratios of $\\sim$25 in the galactic Center, $\\sim$50 in the inner galactic disk and the LMC, $\\sim$70 at the Sun's galactocentric radius, and  $\\ga$100 in the outer Galaxy. The solar system ratio of 89 may have been typical of the galactic disk at the Sun's galactocentric radius 4.6\\,Gyr ago (e.g., Wilson \\& Rood 1994; Wouterloot \\& Brand 1996; Wang et al. 2009). Within the framework of ``biased infall'', where the galactic disk developed from inside out (Chiappini et al. 2001), there {\\it might} be a future chance to use $^{12}$C/$^{13}$C ratios as a chronometer for nucleosynthesis.  {\\it In nearby galaxies}, the $^{12}$CO/$^{13}$CO line intensity ratios are usually measured in the $J$=1--0 line and have typical values of $\\sim$10. They are higher than the values for individual molecular clouds in the Galaxy because they are mostly observed with larger beams. These include not only the dense clouds, where both species are (almost) optically thick, but also the molecular intercloud medium, where $^{13}$CO is optically thin. Like the better-resolved CO line ratios in our Galaxy, the ratios in nearby galaxies probably correspond to true $^{12}$C/$^{13}$C abundance ratios between 40 and 90 (e.g., Henkel et al. 1993).  In a presumably ``normal'' spiral {\\it galaxy at redshift 0.89}, in the lens of the background source PKS\\,1830-211, Wiklind \\& Combes (1998),  Menten et al. (1999), and  Muller et al. (2006) derive, from the optically thin wings of the absorption lines of HCO$^+$, HCN, and HNC, a $^{12}$C/$^{13}$C abundance ratio of 27$\\pm$2. Apparently, even at an age of the universe of $\\sim$6.5\\,Gyr, it appears that $^{13}$C is as abundant with respect to $^{12}$C as in the center of our Galaxy at the present epoch.  {\\it Some low-redshift (ultra)luminous infrared galaxies} ((U)LIRGs), however, show peculiarities, which may be relevant to the Cloverleaf. Local (U)LIRGs are known to reveal $^{12}$CO/$^{13}$CO $J$= 1$\\rightarrow$0 line intensity ratios which tend to be higher than the canonical value of 10 for ``normal'' galaxies (see, e.g., Aalto et al. 1991; Casoli et al. 1992; Henkel \\& Mauersberger 1993). According to Taniguchi \\& Ohyama (1998), there is a tight correlation between $L$($^{12}$CO $J$=1$\\rightarrow$0) and $L_{\\rm FIR}$. However, when comparing ``normal'' galaxies with those with a high $^{12}$CO/$^{13}$CO $J$=1$\\rightarrow$0 ratio, the $^{13}$CO luminosities show a deficiency by an average factor of $\\sim$3, This $^{13}$CO deficiency is readily explained by metallicity gradients in the progenitor galaxies and strong interaction- or merger-induced inflow of gas into the luminous cores (e.g., Rupke et al. 2008). Apparently, for ultraluminous galaxies the common luminosity - metallicity correlation is not valid. Ultraluminous galaxies are characterized by a lower metallicity, likely yielding higher $^{12}$C/$^{13}$C isotope ratios. In the early universe, gas from outside the cores of the merging progenitors may have been particularly metal poor, leading to extreme carbon isotope ratios.  For $T_{\\rm kin}$ $\\ga$ 20\\,K, the $^{12}$CO $J$=3$\\rightarrow$2 line is more opaque, typically by a factor of 3, than the corresponding 1$\\rightarrow$0 line. Thus our conservatively estimated $J$=3--2 $^{12}$CO/$^{13}$CO line intensity ratio of $\\ga$40$^{+25}_{-8}$ corresponds to a 1$\\rightarrow$0 ratio well in excess of 40. So far, only few $^{12}$CO/$^{13}$CO $J$=3$\\rightarrow$2 line ratios have been measured in luminous mergers of low redshift. Greve et al. (2009) find 8$\\pm$2 for the ULIRG Arp~220 and $\\ga$30 for the LIRG NGC~6240. The latter value {\\it might} be consistent with that of the Cloverleaf.   \\subsection{Are there alternatives to a $^{13}$C deficiency in the Cloverleaf?}  Sects.\\,5.1 and 5.2 suggest, that our measured $^{12}$CO/$^{13}$CO line intensity ratio (or its lower limit) require a significant $^{13}$C deficiency in the Cloverleaf. Are there caveats we may have overlooked when reaching this conclusion?  If the bulk of the CO emission would not arise, as suggested by Venturini \\& Solomon (2003), from a molecular disk but from a large scale outflow, such gas would not be in virial equilibrium and could arise predominantly from a diffuse gas phase. While this would yield (within the LVG approach) a higher velocity gradient and a lower [$^{12}$CO]/([H$_2$](d$v$/d$r$)) value than what is needed for virialized clouds, required densities would then be well in excess of 10$^4$\\,cm$^{-3}$, in contradiction with our assumption of predominantly diffuse gas. Furthermore, as long as $T_{\\rm kin}$ remains moderate ($\\la$50\\,K; see Figs.\\,\\ref{fig4} and \\ref{fig5}), $^{12}$C/$^{13}$C ratios remain larger than those encountered in the galactic disk (Sect.\\,5.1).  Following White (1977), radiative transfer models with simple geometry, either based on microturbulence or on systematic motions, lead to peak and integrated intensities which agree within the differences (up to a factor of three) caused by an uncertain cloud geometry. A full 3-D model of a rotating circumnuclear disk, computing the radiative transfer through many lines of sight, calculating the LVG level populations within each pixel of the simulated source, and also including continuum radiation from dust (e.g., Downes \\& Solomon 1998) may be worth doing. In the Cloverleaf, however, the distribution of the molecular gas is still poorly known.  A large $^{12}$C/$^{13}$C ratio, implying an underabundance of $^{13}$C, appears to be in direct conflict with optical data. As already mentioned in Sect.\\,1, solar or super-solar metallicities are common in quasars up to high redshifts. This does not only refer to so-called ``$\\alpha$-elements'' being rapidly synthesized in short-lived massive stars but also to iron (e.g., Iwamuro et al. 2004; Kurk et al. 2007; Sameshima et al. 2009), carbon (e.g., Jiang et al. 2007; Juarez et al. 2009), and, even more importantly, nitrogen (Hamann \\& Ferland 1999; De Breuck et al. 2000; Vernet et al. 2001; Hamann et al. 2002; Nagao et al. 2006; Matsuoka et al. 2009), with $^{14}$N being mainly a secondary nucleus produced by CNO burning just like $^{13}$C. A possible explanation for the contradictory results obtained at or near optical wavelengths and the microwave data presented here may be different locations. It is well possible that mainly secondary nuclei like $^{13}$C and $^{14}$N are enriched close to the quasar, in the Broad and Narrow Line Regions and in outflows originating from the active galactic nucleus (AGN). However, CO $J$=7$\\rightarrow$6 may arise hundreds of pc away from the AGN (Venturini \\& Solomon 2003) and some of the $J$=3$\\rightarrow$2 photons may be emitted from locations even farther away.  There exists, however, also the possibility that our measured high $^{12}$CO/$^{13}$CO luminosity ratio is misleading and does {\\it not} imply a large $^{12}$C/$^{13}$C ratio. As a consequence of different optical depths, $^{12}$CO lines are almost thermalized and are characterized by excitation temperatures well above the level of the cosmic microwave background even at $z$=2.5. $^{13}$CO is less thermalized. In our best fitting models, its $J$=3$\\rightarrow$2 excitation temperature lies in the range 20--30\\,K. This is significantly above the 9.7\\,K of the CMB. However, an extreme (and therefore unlikely) enhancement of the background level by dust radiation could reduce the contrast between line and background for $^{13}$CO far more efficiently than for $^{12}$CO (see Papadopoulos et al. 2010 for the case of Arp\\,220), thus establishing an apparent $^{13}$CO deficiency."
    },
    "1004/1004.0463_arXiv.txt": {
        "abstract": "{In this paper, we present the consistent evolution of short-period exoplanets coupling the tidal and gravothermal evolution of the planet. Contrarily to previous similar studies, our calculations are based on the {\\it complete} tidal evolution equations of the Hut (1981) model, valid at any order in eccentricity, obliquity and spin. We demonstrate, both analytically and numerically, that, except if the system was {\\it formed} with a nearly circular orbit ($e\\simle 0.2$), solving consistently the complete tidal equations is mandatory to derive correct tidal evolution histories. We show that calculations based on tidal models truncated at second order in eccentricity, as done in all previous studies, lead to quantitatively, and sometimes even qualitatively, erroneous tidal evolutions. As a consequence, tidal energy dissipation rates are severely underestimated in all these calculations and the characteristic timescales for the various orbital parameters evolutions can be wrong by up to three orders in magnitude. Such discrepancies can by no means be justified by invoking the uncertainty in the tidal quality factors.  Based on these complete, consistent calculations, we revisit the viability of the tidal heating hypothesis to explain the anomalously large radius of transiting giant planets. We show that, even though tidal dissipation does provide a substantial contribution to the planet's heat budget and can explain some of the moderately bloated hot-Jupiters, this mechanism can not explain alone the properties of the most inflated objects, including HD 209\\,458\\,b. Indeed, solving the complete tidal equations shows that enhanced tidal dissipation and thus orbit circularization occur too early during the planet's evolution to provide enough extra energy at the present epoch. In that case, either a third, so far undetected, low-mass companion must be present to keep exciting the eccentricity of the giant planet, or other mechanisms, such as stellar irradiation induced surface winds dissipating in the planet's tidal bulges and thus reaching the convective layers, or inefficient flux transport by convection in the planet's interior must be invoked, together with tidal dissipation, to provide all the pieces of the abnormally large exoplanet puzzle. } ",
        "introduction": "\\label{sec:intro}  Gravitational tides have marked out the history of science and astrophysics since the first assessment by Seleucus of Seleucia of the relation between the height of the tides and the position of the moon and the Sun in the second century BC. Modern astrophysics extended the study of gravitational tides in an impressive variety of contexts from the synchronization of the moon and other satellites to the evolution of close binary stars and even the disruption of galaxies.  The recent discoveries of short period extrasolar planetary systems and the determination of the anomalously large radius of some giant close-in exoplanets revived the need for a theory of planetary tides covering a wider variety of orbital configurations than previously encountered for the case of our own solar system planets. In particular, the orbital evolution of planetary systems such as HD 80\\,606, with an orbital eccentricity of 0.9337 \\citep{NLM01}, and XO-3, with a stellar obliquity $\\es\\gtrsim37.3\\pm3.7$ deg \\citep{WJF09}, cannot be properly treated with tidal models limited to the case of zero or vanishing eccentricity and obliquity such as in the models of e.g. \\citet{GS66}, \\citet{JGB08} and \\citet{FRH08}.  Following \\citet{BLM01} and \\citet{GLB03}, attempts have been made to explain the observed large radius of some transiting close-in gas giant exoplanets - the so-called \"Hot Jupiters\" - by means of tidal heating \\citep{JGB08,MFJ09,ISB09}. All these models, however, use tidal models truncated to low (second) order in eccentricity, in spite of initial eccentricities, as determined from the tidal evolution calculations, which can be as large as $e=0.8$! According to these calculations, a large eccentricity can remain long enough to lead to tidal energy dissipation in the planet's gaseous envelop (assuming a proper dissipation mechanism is at play in the deep convective layers) at a late epoch and then can explain the actual bloated radius of some observed planets.  In the present paper, we revisit the viability of this tidal heating hypothesis, using an extended version of the \\citet{Hut81} tidal evolution model, solving consistently the {\\it complete} tidal equations, to any order in eccentricity and obliquity, and coupling these latter with the gravothermal evolution of the irradiated planet. As will be shown in the paper, properly taking into account the full nature of the tidal equations severely modifies the planet's tidal and thermal evolution, compared with the aforementioned truncated calculations, leading to significantly different tidal heat rates and thus planet contraction rates.  After introducing our model in \\S\\ref{sec:hyp}, we examine in detail in \\S\\ref{sec:q} the relation between the \\textit{constant time lag} ($\\Delta t$) in Hut's (and thus our) model and the usual tidal quality factor ($Q$) widely used in the literature. Constraints on $\\Delta t$ from the study of the galilean satellites are also derived. In \\S\\ref{sec:2ndOrder}, we demonstrate, with {\\it analytical arguments}, that truncating the tidal equations at $2^{\\mathrm{nd}}$ order in eccentricity leads to wrong tidal evolution histories, with sequences drastically differing from the ones obtained when solving the complete equations. In \\S\\ref{sec:comp}, we compare our full thermal/orbital evolution calculations with similar studies based on a truncated and constant $Q$ tidal model. These numerical comparisons confirm and quantify the conclusions reached in \\S\\ref{sec:2ndOrder}, namely that low order eccentricity models substantially underestimate the tidal evolution timescales for initially eccentric systems and thus lead to incorrect tidal energy contributions to the planet's energy balance. For instance, we show that tidal heating can not explain the radius of HD~209\\,458\\,b, for the present values of their orbital parameters, contrarily to what has been claimed in previous calculations based on truncated eccentricity models \\citep{ISB09}. Finally, in \\S\\ref{sec:global}, we apply our model to the case of some of the discovered bloated planets. We show that, although tidal heating can explain the presently observed radius of some {\\it moderately bloated} hot Jupiters, as indeed suggested in some previous studies, tidal heating alone cannot explain {\\it all} the anomalously large radii. Indeed, in these cases, eccentricity damping occurs too early in the system's tidal evolution (assuming a genuine two-body planetary system) to lead to the present state of the planet's contraction. ",
        "conclusions": "\\label{sec:disc}   In this paper, we have demonstrated that the quasi-circular approximation ($e\\ll 1$, i.e. tidal equations truncated at the order $e^2$) usually made in tidal calculations of transiting planet systems and valid for our Solar system planets, is not valid for the exoplanetary systems that have - or were born with - an even modestly large ($e\\simgr 0.2$) eccentricity. As shown in \\S\\ref{sec:2ndOrder}, although the real frequency dependence of the tidal effect remains uncertain, there are dimensional evidences that for eccentric orbits, most of the tidal effect is contained in the high order terms and that truncating the tidal equations at $2^\\mathrm{nd}$ order in eccentricity can overestimate the characteristic timescales of the various orbital parameters by up to three orders of magnitude. Therefore, truncating the tidal equations at the second order can by no means be justified by invoking the large uncertainty in the dissipative processes and their frequency dependence. Therefore, high order tidal equations should be solved to derive reliable results for most of the existing exoplanet transiting systems. This need to solve the complete equations is met by any tidal model. In this context, even though no tidal model can claim describing perfectly a two body evolution, we recall that the Hut model is at least exact in the weak friction viscous approximation (see \\S\\ref{sec:q}).  We have tested our complete tidal model on several inflated planets to find out whether or not tidal heating can explain the large radius of most of the observed transiting systems. Although this mechanism is indeed found to be sufficient to explain moderately bloated planets such as OGLE-TR-211\\,b (see Fig.\\,\\ref{fig:tr211}), we have been \\textit{unable} to find evolutionary paths that reproduce both the measured radius and the orbital parameters of HD 209\\,458\\,b, WASP-12\\,b, TrES-4\\,b, and WASP-4\\,b (see Fig.\\,\\ref{fig:hd} and Fig.\\,\\ref{fig:w12_t4}) for their inferred age range. The main reason is the early circularization of the orbit of these systems. As demonstrated in the paper, this stems from the non-polynomial terms in eccentricity in the complete tidal equations, which are missing when truncating the equations at small $e$-order. The present results, based on complete tidal equations, show that tidal heating, although providing an important contribution to the planet's internal heat budget during the evolution, cannot explain {\\it alone} the observed properties of all exoplanets.  This is in contrast with some of the conclusions reached in previous studies. Based on truncated tidal models, \\citet{IB09} and \\citet{ISB09} find evolutionary tracks that match observed parameters for HD 209\\,458\\,b, WASP-12\\,b, and WASP-4\\,b and thus suggest that the tidal heating is the principal cause of the large radii of Hot Jupiters.  These particular properties of Hot Jupiters, including the extreme cases of the most severely bloated planets, can only be explained if the following explanations/mechanisms occur during the system lifetimes: \\begin{itemize} \\item Early spin up of the star: simulations of the rotational evolution of solar-like stars \\citep{BFA97} show that after the dispersion of the accretion disk, the rotation rate of the contracting star increases due to angular momentum conservation, until magnetic braking takes over. Considering Eq.\\,(\\ref{evol_e}), we see that stellar tides act as an eccentricity source if $\\frac{\\os}{n}\\geqslant\\frac{18}{11}\\frac{N_e(e)}{\\Omega_e(e)}$. Investigating whether the duration of this phase lasts long enough and whether the magnitude of this effect is large enough to drive enough eccentricity requires performing consistent star/planet thermal/tidal calculations and will be investigated in a forthcoming paper. \\item Presence of a third body: as proposed by \\citet{Mar07}, a low mass terrestrial planet can drive the eccentricity of a massive giant planet during up to Gyr timescales. Accurate enough observations are necessary to support or exclude the presence of such low-mass companions. \\item As mentioned earlier, combining tidal heat dissipation with other mechanisms such as surface winds, due to the stellar insolation, dissipating deep enough in the tidal bulges, or layered convection within the planet's interior may provide the various pieces necessary to completely solve the puzzle. \\end{itemize}  In conclusion, the suggestion that tidal heating is the main mechanism responsible to solve the problem of anomalously large short-period planets, as sometimes claimed in the literature, must be reformulated more rigorously: although providing a non-negligible contribution to hot-Jupiter heat content, tidal dissipation does not appear to provide the whole explanation. Further studies are thus necessary to eventually nail down this puzzling issue."
    },
    "1004/1004.3289_arXiv.txt": {
        "abstract": "{} { We present a new semi-analytical model of galaxy formation, GECO (Galaxy Evolution COde), aimed at a better understanding of when and how the two processes of star formation and galaxy assembly have taken place, by comparison with a wide variety of recent data on the evolutionary galaxy mass functions and star-formation histories.} {Our model is structured into a Monte Carlo algorithm based on the Extended Press-Schechter theory, for the representation of the merging hierarchy of dark matter halos, and a set of analytic algorithms for the treatment of the baryonic physics, including classical recipes for the gas cooling, the star formation time-scales, galaxy mergers and SN feedback. Together with the galaxies, the parallel growth of BHs is followed in time and their feedback on the hosting galaxies is modelled. We set the model free parameters by matching with data on local stellar mass functions and the BH-bulge relation at $z=0$.} {Based on such local boundary conditions, we investigate how data on the high-redshift universe constrain our understanding of the physical processes driving the evolution, focusing in particular on the assembly of stellar mass and on the star formation history. Since both processes are currently strongly constrained by cosmological near- and far-IR surveys with the Spitzer Space Telescope, the basic physics of the $\\Lambda CDM$ hierarchical clustering concept of galaxy formation can be effectively tested by us by comparison with the most reliable set of observables using a minimal number of free parameters.} {Our investigation shows that when the time-scales of the stellar formation and mass assembly are studied as a function of dark matter halo mass and the single galaxy stellar mass, the 'downsizing' fashion of star formation appears to be a natural outcome of the model, reproduced even in the absence of the AGN feedback. On the contrary, the stellar mass assembly history turns out to follow a more standard hierarchical pattern progressive in cosmic time, with the more massive systems assembled at late times mainly through dissipationless mergers. } ",
        "introduction": "In the past decade, several observational evidences were accumulating in favour of the $\\Lambda CDM$ paradigm for structure formation, now quite a successful rendition of the hierarchical clustering scenario for cosmic structure formation. In its standard form \\citep{Blum:84}, it predicts that structures formed from primordial fluctuations of the density field amplified during inflation and then collapsed to form the virialized structures that we see nowadays. The most compelling support in favour of this paradigm comes from the measurements of the temperature anisotropies of the cosmic microwave background \\citep{Spergel:03, Spergel:07}. Further evidences are due to the measurements of the power spectrum of galaxy clustering from large surveys of the local universe \\citep{Perc:02, Tegm:04}, the evidence for an accelerated expansion of the universe as inferred from high-redshift type Ia supernovae observations \\citep{Riess:98,Perl:99} and the baryon fraction observed in rich clusters \\citep{White:93}.  Devised to study galaxy evolution within this cosmological framework, the semi-analytical approach favours a relatively simple handling of the main physical parameters and understanding of their possible role in driving the evolution. This modelling has its root in the work of \\citet{WR:78}, where it was proposed that galaxy formation is a two-stage process, with dark matter halos forming in a dissipationless gravitational collapse and galaxies forming inside them following the radiative cooling of baryons. Although \\citet{WR:78} and,  after,  \\citet{WF:91}, based their work only on the analytic Press-Schechter  theory \\citep{PS:74}, predicting only average  quantities, subsequently  a number  of works followed their prescriptions using Monte Carlo (MC) merger trees. The MC approach allows to obtain several realizations of the merging history of individual dark matter halos. This approach was pioneered by \\citet{LC:93, kauff:93a, Cole:94}, and then followed by a number of authors \\citep{SK:99,  SL:99, Zen:07}. The great advantage of the semi-analytical method (SAM), apart from being computationally very fast and flexible, is the fact that it is possible to compute merging histories with arbitrary mass resolution. The alternative approach that can be followed in order to track the evolution of the dark matter halos, is through the use of large cosmological N-body simulations. Their large computational requirements are compensated by the amount of information that can be achieved. For instance, in the Millennium Simulation \\citep{Spring:01}, the evolution of substructures in massive halos can be followed in time, with the results of a more detailed information about the galaxy dynamics and the influence  of the cosmic environment on the process \\citep{DeL:04}. In the literature, various examples of this ``hybrid'' approach, which make use of N-body simulation for the dark matter evolution and the SAM technique for the baryonic physics, have been published \\citep{Hatton:03, DeL:06, Cr:06}.  In the present work we employ a MC merger tree, mainly because this allowed us to test the parameter space of the semi-analytical model with much more flexibility than using N-body simulations, and allows to compute merging histories down to arbitrary low mass resolution.  While the treatment of the evolution of dark matter structures is relatively simple, as being determined only by gravity, the physics of the baryons inside halos is much more complex to describe. In the most classical models \\citep{kauff:93b, Baugh:96, kauff:99, SP:99, Cole:00, Menci:02}, gas cooling, star formation, SN feedback and galaxy mergers are included. In recent years, it has become clear that some other form of highly energetic feedback is needed to prevent  star formation in massive galaxies at recent epochs, where the SN feedback is ineffective. The need of such form of feedback is required in order to avoid the overcooling in massive halos, and hence the overabundance of galaxies at the bright-end of the luminosity function. This source of feedback is commonly found in the AGN energy production. This effect, supported by the observational findings of a tight correlation between the BH and the bulge size \\citep{Fer:00, HR:04}, was implemented in different ways by \\citet{KH:00, Bower:06, Cr:06, Menci:06, Som:08}.  Alternative to the AGN feedback, the shutdown of star formation above a critical halo mass has been implemented as a quenching mechanism in massive galaxies \\citep{Cattaneo:06, Cattaneo:08}, motivated by the prediction of stable shock heating for halos more massive than this threshold \\citep{Dekel:06}.  Following the prescriptions of these models, we have built a new semi-analytical models, the Galaxy Evolution COde (GECO), whose aim is to identify a few key physical parameters and modify them by comparing with several basic properties of the galaxy population at $z=0$, as well as at high redshifts. Our main observational reference in this paper is the redshift-dependent stellar mass function of galaxies, which, from a suitable choice for the stellar Initial Mass Function (IMF), is a robust descriptor of the star formation history and the mass assembly history of galaxies.  The structure of the paper is as follows. In \\S\\ref{mtree} we describe the Monte Carlo merger tree used in the model. In \\S\\ref{model} we introduce the ingredients of the baryonic model. In \\S\\ref{param} we explain how the free parameters are set and provide a table for them. In  \\S\\ref{localUn} the basic results for the local universe are presented, while in  \\S\\ref{Highz} we focus on the high redshift predictions. We conclude in \\S\\ref{concl}. Throughout the rest of the paper we assume a ``concordance'' cosmological model, with $\\Omega_m=0.3$, $\\Omega_{\\Lambda}=0.7$, $h=0.7$, $\\sigma_8=0.9$ and $n=1$ (power index of the primordial power spectrum). However, when needed, we also show the dependence on the cosmological parameters showing the results for the new WMAP5 dataset, which are: $\\Omega_m=0.258$, $\\Omega_{\\Lambda}=0.742$, $h=0.719$, $\\sigma_8=0.796$ and $n=0.963$ \\citep{Dunkley:09}. ",
        "conclusions": "We have presented a new semi-analytical model of galaxy formation, the Galaxy Evolution Code, GECO which appears to reproduce several key statistical properties of local and high-redshift galaxies.  GECO  is based on a state-of-the-art Monte Carlo algorithm for the representation of the dark matter halo merging history, based on the Extended Press-Schechter formalism. GECO includes detailed implementations for gas cooling, star formation, feedback from SN and galaxy mergers, due to both dynamical friction and random collisions. Moreover, the parallel growth of BHs is followed in time and the subsequent AGN feedback is modelled.  We specifically tested our results directly on the observables involving the stellar mass and star-formation rate more then luminosities, as usually done by other published models. This is motivated on one hand by the fact that the stellar mass functions are the most direct outcome of the model. On the other hand, stellar masses in galaxies have recently become a rather straightforward observable thanks to rest-frame near-infrared data by the Spitzer Space Telescope, directly probing the stellar mass content in high-redshift galaxies. At the same time, Spitzer is also probing with deep far-infrared photometric imaging the rate of stellar formation in distant objects. Therefore, we can compare the outcomes of our model with the most reliable set of observables and a minimal number of free parameters. We thus believe that the basic physics of the $\\Lambda CDM$ hierarchical clustering concept of galaxy formation can be tested by us in a very effective way.   The main results obtained in the present work are summarised in the following.   \\begin{enumerate}  \\item The local stellar mass function results in a remarkably good agreement with the determination of \\citet{Cole:01} and  \\citet{Bell:03} (Fig. \\ref{geco_smfz0}). At the high-mass end the total mass function is dominated by the contribution of bulges, while discs dominate at the faint-end. When the total stellar mass is splitted into the contributions of early-type and late-type galaxies, the former populate the bright-end side, while the latter mainly contribute at low masses. The number densities of the two morphological types cross each other at $M_{star}\\sim 3 \\times 10^{10} \\Msunh$, as observed. Although we reproduce the general trend of the morphological mass functions, our model fails in matching the low-mass end of spheroids, showing an excess of low-mass systems. Likely, this is due to an oversimplification of the satellite population, since satellite galaxies loose  their hot gas reservoir as soon as they are incorporated in a more massive halo and the star formation is  quenched soon after.  \\item The co-evolution of galaxies and BHs is modelled following the prescriptions of \\citet{Cr:06}. A first mode of accretion onto BHs considered is the so called 'radio-mode', that inhibits the quiescent star formation, while the second one is the `QSO-mode', that is triggered only during galaxy mergers and constitutes a major channel of BH accretion. As a consequence of mergers, a starburst is induced as well, feeding the galactic bulge component (and destroying the disc in the case of a major merger). This leads to a parallel growth of BH and bulges with the two masses very well correlated, in agreement with observational data \\citep{HR:04} and to the local black-hole mass function in a remarkable agreement with the observations of \\citep{Shankar:04}.    \\item We compare the stellar mass functions resulting from the model with various observational determinations up to $z\\sim 3.5$ (Fig. \\ref{geco_ev_all}) and found a reasonably fair agreement over the whole redshift range considered. Nevertheless, the observed ratio between the evolution of the faint- and the bright-end of the stellar mass function is not very accurately reproduced: there is too much evolution in the model at the bright-end and too little at the faint-end compared to observations. Various sets of observables indicate a large increase in the number density of low-mass objects between $z\\sim 2$ and the present-day and a lower rate of evolution for massive objects. However, we mentioned that the completeness and robustness of the observational mass function are to be proven there, before claiming more definite conclusions. In the case of  WMAP5 cosmology we observe a delay in the formation of cosmic structure, that lead to a further reduction of high-mass systems.  \\item The bolometric quasar luminosity function is compared with \\citet{Hopk:07a}, showing a good level of agreement at low and intermediate redshift, but a tendency to underpredict the number of bright quasars at high redshift. A mechanism for enhancing the cooling rate at high redshift might simultaneously increase the fuel for BHs and enhance the SF at high redshift, as seems to be required to improve the match with the stellar mass functions.  \\item The integrated star formation rate density (Fig. \\ref{geco_madau}) shows an high level of star formation at high redshifts, a peak at $z\\sim 1.5-3$ and then a sharp decline below $z\\sim 1$. When compared with the determination of the SFH derived using various tracers (UV, optical, radio, IR), our predictions are in very close agreement with these observations. At very high redshift ($z \\simeq 7$) our model is able to correctly reproduce the recent determination of the SFR density  by HST-WFC3 data.   \\item We analysed in detail the SFH in simulated sets of galaxies, to gain insights into how the model treats star formation and how it depends on the galaxy or halo mass. We computed the averaged SFH for the central galaxies living in halos of different sizes (Fig. \\ref{geco_sfhav}). We identified two main trends with halo mass. First, going to high-mass systems, the contribution of the starburst mode to the total SF becomes increasingly important, and indeed predominant in very high mass objects. Second, the formation redshift, defined as that when half of the present-day stellar mass is formed, increases, leading to older stellar populations in massive systems. Hence galaxies in our model form their stars following a \\emph{downsizing} pattern, consistent, for instance, with the dating of stellar populations in local galaxies \\citep{Thomas:05}. The naive expectation of early versions of hierarchical galaxy formation models was that, since massive halos are assembled later than their lower-mass counterparts, the most massive galaxies, hosted in the largest halos, should form their stellar content at the same late cosmic time. Actually, as shown in Fig. \\ref{zform}, this is clearly not the case in our refined model. According to it, \\emph{downsizing} in star-formation is an intrinsic feature of semi-analytical models (see also \\citealt{Nei:06}). The present-day massive galaxies were formed through the assembly of a number of smaller progenitors, that collapsed at high redshifts from the highest density peaks of the primordial density field. According to this scenario, also named biased galaxy formation \\citep{Dekel:86}, bright and massive systems started to form stars early on. This feature is a natural outcome of the merger tree formalism: progenitors of high-mass halos fall below the resolution mass imposed to the merger tree after several time-steps back in cosmic time, so the leaves of the tree are found at high redshift. On the contrary, smaller systems, closer to the resolution mass, take only a few time-steps back to reach this minimum mass. Note that this not merely a computational artifact, but indeed corresponds to the fact that we expect quite negligible star formation to have occurred below such threshold mass, on consideration of the SF quenching by the UV photoionizing background (Sect. \\ref{UV}). Since baryons are put into halos starting from the leaves, in high mass halos star formation took place at early times. Moreover, at high redshift both mechanisms of star formation were more effective. Thanks to the efficient cooling of the gas, the quiescent mode of star formation occurs at enhanced rate. Also, the frequency of mergers at these times is high, allowing an efficient conversion of gas in stars via a starburst. The exhaustion of cold gas is, therefore, very rapid, leaving the galaxy devoid of fuel, and preventing further star formation to occur.  A further reason for the star formation quenching in massive systems may be ascribed to AGN feedback. In order to check the importance of such mechanism, we run a test simulation with the AGN emission switched off. The comparison between the two versions of the model are presented in Appendix A. We found that AGN feedback has some effect in increasing the average age of stellar populations in the most massive galaxies, as it was point out by previous works \\citep{Cr:06,DeL:06,Cattaneo:08}. However, in our model the \\emph{downsizing} trend is still obtained, although slightly weaker, even in the absence of an AGN.  Therefore, in GECO, the AGN feedback has not a dramatic effect in producing the local galaxy properties (see also \\citealt{Menci:06} and  \\citealt{Monaco:07}) and it is not the main reason for the \\emph{downsizing} pattern of galaxy evolution. In our modelling, the latter is rather due to a higher efficiency of SF back in time for the most massive galaxies in rich environments, as explained above. Although we have shown that the \\emph{downsizing} is quite a natural feature of  our model, it is not the case for all the published models and some of them fail in reproducing such trend with mass, as it is shown in \\citet{Fontanot:09}.    \\item Finally, we compared the star formation history of simulated galaxies with the mass assembly history, that is the history of the mass assembled into the main progenitor at each time-steps (Fig. \\ref{sfh_ass1}). As expected, the two processes may occur on very different time-scales, especially in high-mass systems, where the star formation took place at high redshift in many distinct progenitors, which assemble at low redshift. The late assembly of these systems occurs in the majority of the cases through \\emph{dry mergers}, i.e. mergers between spheroidal systems with little or no gas. In these cases, the merger is not accompanied by any event of star-formation, neither quiescent nor bursting. An intriguing question is then whether galaxies in the model assemble their stellar mass in a \\emph{downsizing} way, as it occurs for the star formation. Possible downsizing effects in the mass assembly were indicated by several authors \\citep{Bundy:06, Cimatti:06, Hopk:07} as inferred from the lack of evolution in the high-mass end of the mass function and by the evolution of the blue-to-red galaxy crossover mass (the mass for which the early-type mass function intersects that of late-types). In our model, we found that although the assembly time shows a shallow dependence on the host halo mass, on average galaxies living in massive halos assemble their stars before galaxies in less massive hosts (Fig. \\ref{zform}). Instead, the assembly times are almost constant with stellar mass, and decrease for very-high mass systems ($M_{star} > 10^{11} \\Msun$), hence leading to an \\emph{upsizing} trend with time in the high-mass end. In our modelling, dry mergers are the main reason for this late assembly of massive galaxies. These findings agree with other previously published semi-analytical models \\citep{DeL:06, Cattaneo:08, Cattaneo:10}. The importance of dry mergers in the formation of the most massive galaxies that we observe in the nearby universe, is implied also by the recent observations of the size evolution of massive spheroids \\citep{Trujillo:07, Cimatti:08}. Indeed, if such compact galaxies are absent in the local universe as suggested by \\citet{Trujillo:09}, the expected mechanisms that move the high-redshift compact galaxies to the local relation are dissipationless mergers. Anyway, other works \\citep{Valentinuzzi:09} suggest that such superdense galaxies in the local universe are not as rare as previously claimed.  How far this might be incongruous with the observational indication of downsizing in mass assembly mentioned above will become clear with further observational confrontation. For example, \\citet{Cattaneo:08} argue that the upsizing in mass assembly can coexist with a downwards trend in the transition mass, which then turns out to be a poor indicator of downsizing.   \\end{enumerate}   In conclusion, GECO presents an encouraging level of agreement with a wide range of observational data, at low and high-redshifts. We focus on comparing GECO with data on the two main phases of the galaxy formation process, that are the star formation and the mass assembly. On one hand, we confirmed that the observed \\emph{downsizing} in star formation is natural part of our scheme of  hierarchical growth of structures. On the other hand, the stellar mass assembly process remains more difficult to understand from both a theoretical and an observational point of view. The times of galaxy assembly in our model, related with both the galaxy merger time-scales and the star formation efficiency, strongly depend on the details of the implementations of galaxy dynamics (dynamical friction, satellite collisions and tidal stripping): further work in this sense remains to be done in order to have a deeper insight into the galaxy assembly process, and certainly a comparison with N-body simulations will be helpful.  The most striking conclusion is that, despite the simplicity of the prescriptions adopted, and the small number of free parameters used, the main features of the evolving galaxy population are reproduced. In particular, the AGN feedback is needed only to improve the match of the local stellar mass function, but its effect on stellar ages is not determinant.   We have to keep in mind in any case that this paper includes just a preliminary and partial confrontation with the data, showing at least no obvious clash. Much more extensive analyses and tighter constraints will be obtained as soon as refined data on the evolutionary mass functions and stellar birthrates will become available."
    },
    "1004/1004.1877.txt": {
        "abstract": "We study possible correlations between ultrahigh energy cosmic rays (UHECRs), observed by Auger, AGASA, and Yakutsk, and nearby active galactic nuclei (AGNs) and $Fermi$ sources. We consider the deflection effects by a Galactic magnetic field (GMF) model {constrained by} the most updated measurements. We found that the average deflection angles of UHECRs by the Galactic magnetic fields are less than $4^\\circ$. A correlation between the Auger cosmic-ray events and nearby AGNs with a significance level of $\\sim 4\\sigma$ was found for the Auger UHECR data sets with or without deflection correction. No correlation was found between the AGASA/Yakutsk events with nearby AGNs. Marginal correlations between the Auger events and the $Fermi$ sources, and between AGASA events and $Fermi$ AGNs were found when the deflections calculated by the GMF model were considered. However, no correlation was found between the Yakutsk data and $Fermi$ sources. Some $Fermi$ sources are close to the arrival directions of UHECR events detected by Auger, AGASA, and Yakutsk, most of which are probably chance coincidence rather than objects producing UHECRs in the nearby Universe. Four $Fermi$ sources, NGC 4945, ESO 323-G77, NGC 6951, and Cen A, within 100~Mpc have UHECR events within $3_.^\\circ1$ from their positions, which could potentially be cosmic-ray accelerators. However, the association can only be confirmed if more UHECRs are preferably detected in these directions ",
        "introduction": "The spectrum, origin, and composition of ultrahigh energy cosmic rays (UHECRs) with energies $\\geqslant$ 10$^{19}$~eV {(=10 EeV)} are a long standing mystery in high-energy astrophysics \\citep{ham84}. \\citet{gk66} and \\citet{zak66} showed a theoretical distant limit for the cosmic rays {with} energies of order 10$^{20}$~eV traveling through the microwave background radiation field, which is called the GZK effect. Because of the GZK effect, particles with energies above 10 EeV are able to reach our Earth only from nearby sources {within about} 100~Mpc. Another barrier in the investigation of the UHECR origin is the deflections of UHECRs by the magnetic fields. Due to the poor knowledge of the extragalactic and intergalactic magnetic fields, the deflections of UHECRs have not yet understood. \\citet{dgst04,dgst05} suggested that the deflections by extragalactic magnetic fields are generally less than 1$^{\\circ}$, while \\citet{rdk10} and \\citet{sme03} claimed that could be larger than 10$^{\\circ}$. The Galactic magnetic fields (GMFs) are relatively better known \\citep[e.g.][]{hml06,srwe08} and are widely discussed in the studies of UHECR origin \\citep[e.g.][]{st97,tt02,ps03,nam09}. \\citet{kst07} concluded that the deflections of UHECRs by the GMFs cannot be neglected even for the protons of $E$ = 10$^{20}$eV, since the deflection angles are comparable with the angular resolution of current experiments. \\citet{nam09} tried {seven} GMF models to study the correlations between UHECRs and source population(s). However, no halo component was included in the four GMF models they used and another three GMF models adopted from \\citet{srwe08} have a strong halo component about 7 $\\mu$G. Observational constraints on the Galactic magnetic field strength \\citep{hq94,hmq99,mos96} and the configuration of disk magnetic fields \\citep{hml06,han09} should be carefully considered in the GMF model.   Since the discovery of UHECRs \\citep{lin63}, many equipments have been used to search for these events, including Fly's Eye \\citep{fly94}, Yakutsk Extensive Air Showers Array \\citep{ikp03a,pgi05}, Akeno Giant Air Shower Array \\citep[AGASA;][]{hhik00,tsh03}, High Resolution Fly's Eye cosmic-ray detector \\citep[HiRes;][]{agasa04,hir08b} and Pierre Auger Observatory \\citep[PAO;][]{pao04,pao07,pao08}. The existence of the GZK cutoff has been observed by the HiRes and Auger \\citep{hir08b,pao08}. Some objects have been suggested to be possible sources of UHECRs, e.g., pulsars \\citep{beo00}, active galactic nuclei (AGNs) and subclasses of AGNs \\citep{prs92,fab98,tt01,tt01b,vbj02,gtt02,gtt04,abb06,fzb09}, radio lobes of FR II galaxies \\citep{rab93,hcf09}, and $\\gamma$-ray bursts \\citep{wax95,miu05}. {However, the real sources of UHECRs are not known yet. AGNs are favored as the most probable sources for accelerating particles to the extreme energies \\citep{ham84} for a long time.  Recently, \\citet{pao07,pao08} studied the correlation between the arrival directions of UHECRs and the positions of nearby AGNs in the \\citet{vcv06} AGN catalog (hereafter VCV catalog). They concluded that the arrival directions of cosmic rays with energies above $\\sim$ 60 EeV are anisotropic and UHECRs have a good correlation with the positions of nearby AGNs ($z$~$<$~0.018). The intriguing result attracted much attention. \\citet{iaa08} found the correlation between Yakutsk UHECRs and the nearby VCV AGNs ($\\lesssim$~100~Mpc). \\citet{gfb08} investigated the correlation between the {\\it Swift} Burst Alert Telescope AGN catalog with the Auger UHECR events, and found a correlation at a significance level of 98$\\%$ when the AGNs were weighted by their hard X-ray flux and the Auger experiment exposure. However, some associated AGNs {of Auger events} may not have enough energy to accelerate particles to ultrahigh energies \\citep{msp09}. The High Resolution Fly\u00a1\u00afs Eye Collaboration searched for {possible} correlation between the HiRes UHECRs and AGNs located in the northern hemisphere; however, no significant correlation was found.  The $Fermi$ high energy $\\gamma$-ray sources are also possible UHECR sources. The recently released $Fermi$ Large Area Telescope First Source Catalog (1FGL) contains 1451 $\\gamma$-ray point sources \\citep{fermi10} with nearly uniform sky coverage \\citep{fermi09b}. \\citet{mio10} first investigated the correlation {between} Auger UHECRs and 1FGL sources without considering the deflection by the GMFs and redshifts of $Fermi$ AGNs, and concluded that the UHECRs are not associated with $Fermi$ sources. The possible correlation of UHECRs and $Fermi$ sources should be re-examined after the GZK cutoff and the UHECR deflection by the GMFs are considered.   In this work, we construct a new GMF model based on the updated measurements of the Galactic magnetic fields and investigate the deflections of UHECRs by the GMFs. Considering the GZK cutoff and the deflection correction through our GMF model, we re-examine the possible correlation between UHECRs and nearby AGNs and $Fermi$ sources. In Section 2, we discuss available data of UHECRs detected by Auger, AGASA, and Yakutsk and possible astrophysical objects. The deflections of UHECRs by the GMFs are discussed in Section 3. The correlation studies are given in Section 4. Discussions and conclusions are presented in Section 5.    %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%2 %************************************************************************ ",
        "conclusions": "}  We collected 135 published UHECR events including 57 UHECRs recorded by AGASA with energy $E>40$ EeV, {51 events observed by Yakutsk, both} located in the northern hemisphere, and 27 events with energy $E\\geqslant 57$ EeV detected by Auger located in the southern hemisphere. We use a new GMF toy model constrained by updated measurements to evaluate the deflection effects on the arrival directions of UHECRs. Considering the possible deflection correction by our toy model and the PS model, as well as the different magnetic field components in our model, we search for the possible correlations of UHECRs with nearby AGNs extracted from the new 13{th} VCV AGN catalog of \\citet{vcv10} and the $Fermi$/LAT First Source Catalog of $\\gamma$-ray sources. We found a correlation between the Auger UHECR events and nearby VCV AGNs with a chance probability of $2\\times10^{-5}$, and a significance level of $\\sim 4\\sigma$. Using the same data as \\citet{pao08}, we found fewer UHECR-AGN pairs when deflection is considered, which implies the weakened correlation. A marginal correlation was found between the Auger events and the first year $Fermi$ $\\gamma$-ray sources with a significance level of $\\sim 4\\sigma$ if the deflection by the GMF model is considered. Some $Fermi$ sources of nearby AGNs, NGC 4945, ESO 323-G77, NGC 6951, and Cen A, may be related to UHECRs within $3_.^{\\circ}1$. For AGASA and Yakutsk UHECRs, no evidence of significant correlation is found for the nearby AGNs or the $Fermi$ sources because the matched pairs can be reproduced by the simulated random isotropic UHECR samples, though some $\\gamma$-ray point sources are coincident with the UHECR events within 2$^\\circ$.   The correlations of UHECRs with some astrophysical objects suggest that at least some of the UHECRs are protons \\citep{pao07,pao08}. However, most UHECRs seem to come from various directions and do not associate with known astrophysical objects, which indicates that the majority of UHECRs might suffer larger deflections in the trajectory, due to either the unknown extragalactic magnetic fields or the heavy nuclei component of UHECRs \\citep{piran10}. The deflection of heavy UHECRs by the GMF models is proportional to the charge of nuclei, which leads to a very large deflection angle, for example tens of degrees for iron, and then any correlation discussed in this work can be diminished (Gureev \\& Troisky 2010). If the primaries of the UHECRs are heavy nuclei, instead of proton, the identification of UHECR sources would be very difficult. Obviously, the understanding of UHECR origin will strongly depend on our knowledge about the strength and configuration of the Galactic and extragalactic magnetic fields, which definitely needs more measurements \\citep{ha08}."
    },
    "1004/1004.1629_arXiv.txt": {
        "abstract": "{The energy balance of cataclysmic variables with strong magnetic fields is a central subject in understanding accretion processes on magnetic white dwarfs. With XMM-Newton, we perform a spectroscopic and photometric study of soft X-ray selected polars during their high states of accretion.} {On the basis of X-ray and optical observations of the magnetic cataclysmic variable AI~Tri, we derive the properties of the spectral components, their flux contributions, and the physical structure of the accretion region in soft polars.} {We use multi-temperature approaches in our \\textsc{xspec} modeling of the X-ray spectra to describe the physical conditions and the structures of the post-shock accretion flow and the accretion spot on the white-dwarf surface. In addition, we investigate the accretion geometry of the system by a timing analysis of the photometric data.} {Flaring soft X-ray emission from the heated surface of the white dwarf dominates the X-ray flux during roughly 70\\% of the binary cycle.  This component deviates from a single black body and can be described by a superimposition of mildly absorbed black bodies with a Gaussian temperature distribution between $kT_\\mathrm{bb,low} := 2\\,\\mathrm{eV}$ and $kT_\\mathrm{bb,high} = 43.9^{+3.3}_{-3.2}\\,\\mathrm{eV}$, and $N_\\mathrm{H,ISM} = 1.5^{+0.8}_{-0.7}\\times10^{20}\\,\\mathrm{cm}^{-2}$. In addition, weaker hard X-ray emission is visible nearly all the time. The spectrum from the cooling post-shock accretion flow is most closely fitted by a combination of thermal plasma \\textsc{mekal} models with temperature profiles adapted from prior stationary two-fluid hydrodynamic calculations. The resulting plasma temperatures lie between $kT_{\\mathsc{mekal},\\mathrm{low}} = 0.8^{+0.4}_{-0.2}\\,\\mathrm{keV}$ and $kT_{\\mathsc{mekal},\\mathrm{high}} = 20.0^{+9.9}_{-6.1}\\,\\mathrm{keV}$; additional intrinsic, partial-covering absorption is on the order of $N_\\mathrm{H,int} = 3.3^{+2.5}_{-1.2}\\times 10^{23}\\,\\mathrm{cm}^{-2}$. The soft X-ray light curves show a dip during the bright phase, which can be interpreted as self-absorption in the accretion stream. Phase-resolved spectral modeling supports the picture of one-pole accretion and self-eclipse. One of the optical light curves corresponds to an irregular mode of accretion. During a short XMM-Newton observation at the same epoch, the X-ray emission of the system is clearly dominated by the soft component.} {} ",
        "introduction": "\\object{AI~Tri} (RX\\,J0203.8+2959) was first described within a sample of ROSAT-discovered bright soft X-ray sources by \\citet{beuermann:93}. Their classification of AI~Tri as an AM~Her type binary (also called a polar) was later confirmed in a multiwavelength study by \\citet{schwarz:98}, who identified cyclotron humps in optical spectra obtained during a high state of accretion. The orbital period of $P_\\mathrm{orb} = 4.6\\,\\mathrm{hrs}$ is one of the longest known among polars, whereas the magnetic field strength of $B = 38\\pm2\\,\\mathrm{MG}$ and the amplitude of the long-term brightness variations between $V = 18\\fm0 - 15\\fm5$ \\citep{schwarz:98} lie in the typical parameter range of this class.  Based on the wavelength dependence of the $UBVRI$ light curve minima and the variations in both the linear and the circular polarization, \\citet{katajainen:01} suggest that the system has a high inclination of $i\\approx 70\\degr\\pm 20\\degr$ and accretes onto two almost equally fed magnetic poles.  On the other hand, \\citet{schwarz:98} propose that a single dominating accretion region is active at the epoch of their observations.  AI~Tri belongs to a significantly large group of AM~Her systems that were found to emit almost entirely at X-ray energies below 0.5\\,keV during the ROSAT All-Sky Survey \\citep{beuermann:93,thomas:98,beuermann:99}. Although these systems could play an important role in investigating the energy balance of polars, only a few of them have been studied using high-resolution X-ray spectroscopy \\citep{ramsay:03,ramsay:04shortper}. We, therefore, initiated dedicated observations with XMM-Newton to perform a detailed study of the spectral components, their flux contributions, and the physical structure of the accretion region of polars selected by their distinct soft X-ray fluxes during high-states of accretion.  In the following, we present an analysis of the magnetic cataclysmic variable AI~Tri based on new XMM-Newton and optical data, and archival ROSAT data.    \\begin{table*} \\caption{\\footnotesize{Log of the ROSAT and XMM-Newton observations and of the optical photometry of AI~Tri.}} \\label{tab:obslog} {\\centering \\begin{tabular}{c@{\\qquad}l@{\\qquad\\quad}c@{\\qquad\\quad}c@{\\qquad}r @{\\qquad}r@{\\qquad}c@{\\qquad}l} \\hline\\hline Date & Telescope\\tablefootmark{a} & Instrument & Filter & $t_\\mathrm{exp} [\\mathrm{s}]$ & $t_\\mathrm{cycle} [\\mathrm{s}]$ & Duration [h] & Observer\\\\ \\hline 1998 Jan 15$-$Feb 05 & ROSAT & HRI & 0.1$-$2.4\\,keV & & & 9.6 & PI Schwarz\\\\ 2005 Aug 15 & XMM-Newton & EPIC/pn & 0.1$-$10\\,keV & & & 0.3 & PI Reinsch\\\\ 2005 Aug 15 & XMM-Newton & EPIC/MOS & 0.1$-$10\\,keV & & & 1.4 & PI Reinsch\\\\ 2005 Aug 15 & XMM-Newton & OM & UVM2 & & & 1.1 & PI Reinsch\\\\ 2005 Aug 15 & XMM-Newton & RGS & 0.3$-$2.5\\,keV  & & & 1.3 & PI Reinsch\\\\ 2005 Aug 22 & XMM-Newton & EPIC & 0.1$-$10\\,keV & & & 5.6 & PI Reinsch\\\\ 2005 Aug 22 & XMM-Newton & OM & UVM2 & & & 5.6 & PI Reinsch\\\\ 2005 Aug 17 & AIP 70\\,cm & TK1024-01 & V & 60 & 66 & 6.5 & Schwarz\\\\ 2005 Aug 29 & AIP 70\\,cm & TK1024-01 & V & 60 & 66 & 7.6 & Schwarz\\\\ 2006 Nov 09 & IAG 50\\,cm & STL-6303E & WL & 180 & 191 & 6.6 & Traulsen\\\\ 2006 Nov 15 & IAG 50\\,cm & STL-6303E & V & 240 & 251 & 4.7 & Traulsen\\\\ 2006 Nov 16 & IAG 50\\,cm & STL-6303E & V & 240 & 247 & 1.8 & Traulsen\\\\ 2007 Jan 14 & IAG 50\\,cm & STL-6303E & V &  90$-$180 &  97$-$187 & 2.2 & Traulsen\\\\ 2007 Jan 25 & AIP 70\\,cm & TK1024-01 & V &  90 & 95 & 4.6 & Schwarz\\\\ 2007 Jan 29 & \\textsc{Monet}/North & Alta E47+ & V &  15 &  18 & 1.8 & Hessman\\\\ 2007 Jan 30 & \\textsc{Monet}/North & Alta E47+ & V &  15 &  18 & 2.0 & Hessman\\\\ 2007 Mar 14 & IAG 50\\,cm & STL-6303E & V & 150 & 157 & 1.8 & Traulsen\\\\ 2007 Mar 14 & AIT 80\\,cm & STL-1001E & V &  20$-$30 &  23$-$33 & 1.5 & Nagel\\\\ \\hline \\end{tabular} } \\tablefoottext{a}{AIP: Astrophysikalisches Institut Potsdam. IAG: Institut f\\\"ur Astrophysik G\\\"ottingen. AIT: Institut f\\\"ur Astronomie und Astrophysik T\\\"ubingen, Abt.~Astronomie (Kepler Center for Astro and Particle Physics).} \\end{table*} ",
        "conclusions": "By analyzing our phase-resolved X-ray spectroscopy with XMM-Newton, we have confirmed a strong soft X-ray excess of AI~Tri. Both the emission from the heated accretion spot and the emission from the post-shock accretion flow depart from single temperature models. From our multi-temperature approaches, we have found a bolometric flux ratio of ${F_\\mathrm{bb}}/{F_\\textsc{mekal}} \\ga 5.7^{+5.6}_{-2.5}$ between the two components during the bright phase of AI~Tri on August 22, 2005. An even higher flux ratio of $F_\\mathrm{soft}/F_\\mathrm{hard} \\ga 130^{+95}_{-60}$ in the XMM-Newton energy bands was observed on August 15, 2005, when the total source flux was by a factor of at least ten higher. The distinct soft X-ray excess is probably related to inhomogeneous, blobby accretion as indicated by the high variability of the optical and X-ray flux on short timescales.  From the light-curve characteristics and the spectral variation in the intrinsic absorption with phase, we have inferred the most likely geometry of AI~Tri. One main soft X-ray emitting accretion region undergoes a self-eclipse and a short period of accretion-stream absorption during the orbital revolution of the system. In this interpretation, the maxima of the UV and optical light curves occur at the moment that the accretion region passes the line of sight. The finding that hard X-ray emission is present at an almost constant, very low level over the whole orbital cycle, including the faint phase, may imply that a second accretion region with weak hard X-ray emission exists.  Changes in the accretion geometry may occasionally occur, possibly in response to significant variations in the mass accretion rate. Signatures of these changes are the irregular light curves on October/ November 1992 and on August 17, 2005, and the new minimum dip seen in the $V\\!$-band light curves from November 2006 onward.  Assuming synchronous rotation of the white dwarf with the binary orbit, the main accretion region appears to be located around an unusual longitude of $\\psi=275\\degr$. Since we extrapolated the spectroscopic ephemeris from data of \\citet{schwarz:98}, obtained more than ten years before our multiwavelength observations, we cannot exclude there being a slightly asynchronous rotation of AI~Tri within the accuracy reached. Future optical spectroscopy is necessary to clarify the ephemeris and to finally decide on the location of the accretion pole(s)."
    },
    "1004/1004.5228_arXiv.txt": {
        "abstract": "We consider dark masses measured from kinematic tracers at discrete radii in galaxies for which baryonic contributions to overall potentials are either subtracted or negligible.  Recent work indicates that rotation curves due to dark matter (DM) halos at intermediate radii in spiral galaxies are remarkably similar, with a mean rotation curve given by $\\log_{10}[V_{c,\\mathrm{DM}}/(\\mathrm{km s^{-1}})]=1.47_{-0.19}^{+0.15}+0.5\\log_{10}[r/\\mathrm{kpc}]$.  Independent studies show that while estimates of the dark mass of a given dwarf spheroidal (dSph) galaxy are robust only near the half-light radius, data from the Milky Way's (MW's) dSph satellites are consistent with a narrow range of mass profiles.  Here we combine published constraints on the dark halo masses of spirals and dSphs and include available measurements of low surface brightness galaxies for additional comparison.  We find that most measured MW dSphs lie on the extrapolation of the mean rotation curve due to DM in spirals.  The union of MW-dSph and spiral data appears to follow a mass-radius relation of the form $M_{\\mathrm{DM}}(r)/M_{\\odot}=200_{-120}^{+200}(r/\\mathrm{pc})^2$, or equivalently a constant acceleration $g_{\\mathrm{DM}}=3_{-2}^{+3}\\times 10^{-9}\\mathrm{cm s^{-2}}$, spanning $0.02\\la r \\la 75$ kpc.  Evaluation at specific radii immediately generates two results from the recent literature: a common mass for MW dSphs at fixed radius and a constant DM central surface density for galaxies ranging from MW dSphs to spirals.  However, recent kinematic measurements indicate that M31's dSph satellites are systematically less massive than MW dSphs of similar size.  Such deviations from what is otherwise a surprisingly uniform halo relation presumably hold clues to individual formation and evolutionary histories. ",
        "introduction": "Considered within the framework of Newtonian mechanics, the observed motions of stars and gas imply that the luminous components of galaxies are embedded within extended halos of unseen material.  Absent the ability to measure invisible masses directly, for a given galaxy one measures dynamical and luminous masses independently and then identifies dark matter simply as the difference between the two: $M_{\\mathrm{DM}}=M_{\\mathrm{dyn}}-M_{\\mathrm{lum}}$.  Here we shall investigate the properties of dark matter halos only insofar as they are accessible via the direct application of this equation, without adopting any particular halo model.  This task is made difficult by the fact that there are no galaxies for which both $M_{\\mathrm{dyn}}$ and $M_{\\mathrm{lum}}$ are easily measured.  While the circular motions of stars and gas in spiral galaxies provide a relatively clean measure of the dynamical mass profile, subtraction of the luminous contribution suffers from systematic uncertainties, primarily those related to stellar mass-to-light ratios.  One sidesteps this problem, or at least exchanges it for others, by considering dwarf spheroidal (dSph) galaxies, for which baryonic components contribute negligibly to internal gravitational potentials.  However, because dSphs are supported by stellar velocity dispersions instead of rotation, estimation of their individual mass profiles requires large kinematic data sets and typically employs strong modeling assumptions (see discussion by, e.g., \\citealt{pryor90,wilkinson02,gilmore07}).  Larger elliptical galaxies (see \\citealt{napolitano10}) combine the most challenging aspects of both regimes, as their significant stellar masses are supported in large part by random motions.  Nevertheless, recent work indicates that dark masses are constrained reasonably well at intermediate radii of both spiral and dSph galaxies.  \\citet[``M07'']{mcgaugh07} use rotation curve data for a sample of 60 spirals and discard data points interior to $r<1$ kpc, where dynamical masses can be affected by noncircular motions.  After subtracting baryonic masses, M07 find that the rotation curves due exclusively to dark matter halos lie nearly on top of each other, with mean rotation curve \\begin{equation} \\log_{10} [V_{c,\\mathrm{DM}}/(\\mathrm{km s^{-1}})]=1.47_{-0.19}^{+0.15}+0.5\\log_{10} [r/\\mathrm{kpc}]. \\label{eq:rotation} \\end{equation}  Working at smaller scales, \\citet[``W09'']{walker09d} use kinematic data for eight Milky Way (MW) dSphs to show that for a wide range of plausible halo shapes and velocity anisotropies, the estimated mass within the projected half-light radius is approximately (subject to validity of the assumptions of spherical symmetry, dynamical equilibrium, flat velocity dispersion profiles $\\langle v^2\\rangle^{1/2}(R)=\\sigma$, and negligible contributions to measured velocity dispersions from unresolved binary motions) \\begin{equation} M(r_{\\mathrm{half}})\\approx \\frac{5r_{\\mathrm{half}}\\sigma^2}{2G}. \\label{eq:w09} \\end{equation} Subsequently \\citet{wolf09} have provided an analytic argument for why such an estimate is insensitive to anisotropy.  Since dSph kinematics tend to be dominated by dark matter at all radii (Section \\ref{subsec:dsphs}; see also \\citealt{mateo98} and references therein; \\citealt{walker07b}), Equation \\ref{eq:w09} provides a crude estimate of the dark mass within $r_{\\mathrm{half}}$ for any dSph for which measurements of $r_{\\mathrm{half}}$ and $\\sigma$ are available.  W09 apply Equation \\ref{eq:w09} to data for 28 classical dSphs and ultrafaint satellites in the Local Group and find that, allowing for scatter by a factor of less than two in normalization, the ensemble of dSph data is consistent with a mass profile of the form $M\\propto r^{1.4\\pm 0.4}$.  Thus M07 and W09 independently note the apparent self-similarity of spiral and dSph dark matter halo profiles, respectively.  Here we combine data from those studies in order to examine whether the mean rotation curve attributable to dark matter in spiral galaxies extends to the small radii characteristic of dSphs.  For further comparison we also include recent measurements by \\citet{kuzio08} of rotation curves in low surface brightness (LSB) galaxies.  The combined data set samples dark matter halos over the range $0.2 \\la r \\la 75$ kpc. ",
        "conclusions": "\\label{sec:discussion}  M07 and W09 independently identify similarities among spiral and dSph dark mass profiles, respectively.  M07 find that the rotation curves due to dark matter in spirals lie approximately on top of each other, and W09 show that the hypothesis of a ``universal'' dSph mass profile is as well-supported by kinematic data as independent claims of a common dSph mass scale (\\citealt{mateo93,strigari08}, see below).  Upon joining data sets, we find that spirals and MW dSphs---excluding outliers Her and Sgr and taking published velocity dispersions of the faintest satellites at face value---trace what appears to be a single relation under which dark mass scales with radius.  This relation can be stated in terms of a rotation curve as in M07 (Equation \\ref{eq:rotation}), mass profile (Equation \\ref{eq:mass}), or constant acceleration (Equation \\ref{eq:acceleration}).  Before discussing implications of the new data available for M31 dSphs, we first consider the apparent scaling of MW dSph and spiral dark matter halos in the context of other reported scaling relationships involving these objects.  If the M07 relation encodes a fundamental property of dark matter halos, then along with the appearance of a universal dSph mass profile (W09), it contains and generalizes two additional results reported in the recent literature.  First, \\citet{strigari08} model dSph dark matter halos and find that if they evaluate mass profiles at a fixed radius of 300 pc (requiring extrapolation of the mass profiles of the smallest dSphs to radii of $\\sim 10r_{\\mathrm{half}}$), MW dSphs all have $M_{300}\\sim 10^7 M_{\\odot}$.  If we simply evaluate the M07 mass profile at 300 pc, we obtain $M_{300}=1.8_{-1.1}^{+1.8}\\times 10^7 M_{\\odot}$, instantly reproducing the \\citet{strigari08} result.  Here it is important to note that the M07 relation does \\textit{not} imply that all dSph dark matter halos actually extend to 300 pc, or to any other radius that may be larger than that of the observed tracer population.  Second, \\citet{kormendy04}, \\citet{spano08} and \\citet{donato09} find that if they adopt cored dark matter halos with core radius $r_0$ and central density $\\rho_0$, spiral rotation curves imply a small range in central halo surface density.   \\citet{donato09} extend this analysis to MW dSphs and show that over 14 magnitudes in luminosity, cored dark matter halo models imply central dark matter surface densities of $r_0\\rho_0=140_{-30}^{+80} M_{\\odot}$ pc$^{-2}$.  The apparent universality of $r_0\\rho_0$ would then imply universality of the acceleration generated by the dark matter at $r_0$, with $g_{\\mathrm{DM}}(r_0)\\approx 0.5 G\\pi r_0\\rho_0=3.2_{-1.2}^{+1.8}\\times 10^{-9}$ cm s$^{-2}$ \\citep{gentile09}.  This is the same acceleration implied by the M07 relation for dark matter halos at all radii (Equation \\ref{eq:acceleration}).  Thus the M07 relation unites under a single scaling relation the appearances of a ``universal mass profile'' (W09) and ``common mass scale'' for MW dSphs \\citep{strigari08}, and a ``universal'' central surface density of dark matter halos \\citep{donato09,gentile09}.  The new kinematic data for M31 dSphs indicate that these objects have masses systematically smaller than their MW counterparts, placing them below the extrapolation of the M07 relation.  In order to bring the M31 dSphs onto the M07 relation, either their velocity dispersions would need to be revised upward by a factor of $\\sim 2.5$ on average or their half-light radii (perhaps involving revision of distances) would need to be revised downward to $\\sim 0.3$ times the current estimates.  Absent a compelling reason to suppose that the present M31 data are grossly affected by systematic errors, it seems that the formation and/or evolution of the MW and M31 systems have differed sufficiently to produce measurable differences in the inferred properties of the dark matter halos of their dSph satellites.  For example, \\citet{penarrubia10} use simulations to demonstrate that tidal interactions with the baryonic disk of a host galaxy reduce the masses of dSph satellites at all radii.  These simulations reproduce the observed \\textit{offset} in mass between MW and M31 dSphs by invoking a more-massive disk for M31.  Note that since the degree of mass loss depends on the parent's disk mass, this scenario requires no fine tuning of the orbital distributions of MW and M31 satellites.  Left unexplained is the puzzling circumstance that the dSph satellites of one but not both the MW and M31 appear to follow the same dark matter halo scaling relation as spiral galaxies.  Here we have emphasized the consistency of MW dSph data with extrapolation of the M07 relation in order to generalize recent and independent claims of uniformity among these galaxies (\\citealt{strigari08,donato09,gentile09}, W09).  To the extent that the new data for M31 dSphs undermine the extrapolation of the M07 relation to dSphs, they also undermine each of these previous claims of uniformity.  On the other hand, given the susceptibility of the smallest dark matter halos to evolution driven by environment, it is feasible that one of the Local Group satellite populations has evolved significantly more than the other, altering what may have been more similar conditions initially.  In any case, it is likely that the emerging contrast between MW and M31 dSphs relates important details of the processes that shaped these two populations.  This work follows directly from conversations that took place at the workshop \\textit{Extreme Star Formation in Dwarf Galaxies} (Ann Arbor, MI, July 2009).  We are grateful to O.\\ Gnedin for organizing the workshop, and to J. Pe\\~narrubia, S. Koposov and J.\\ Wolf for helpful discussions.  MGW acknowledges support from the STFC-funded Galaxy Formation and Evolution program at the Institute of Astronomy, Cambridge.  SSM acknowledges support from NSF grant AST-0908370.  MM and EWO acknowledge support from NSF grants AST-0808043 and AST-0807498, respectively."
    },
    "1004/1004.4642_arXiv.txt": {
        "abstract": "In this work, we present results for the photometric and clustering properties of galaxies that arise in a $\\Lambda$ cold dark matter hydrodynamical simulation of the local universe. The present-day distribution of matter was constructed to match the observed large scale pattern of the {\\it IRAS} 1.2-Jy galaxy survey. Our simulation follows the formation and evolution of galaxies in a cosmological sphere with a volume of $\\sim130^3$ $h^{-3}$ Mpc$^3$ including supernova feedback, galactic winds, photoheating due to an uniform meta-galactic background and chemical enrichment of the gas and stellar populations. However, we do not consider AGNs. In the simulation, a total of $\\sim 20000$ galaxies are formed above the resolution limit, and around $60$ haloes are more massive than $\\sim10^{14}$ M$_{\\odot}$. Luminosities of the galaxies are calculated based on a stellar population synthesis model including the attenuation by dust, which is calculated from the cold gas left within the simulated galaxies. Environmental effects like colour bi-modality and differential clustering power of the hydrodynamical galaxies are qualitatively similar to observed trends. Nevertheless, the overcooling present in the simulations lead to too blue and overluminous brightest cluster galaxies (BCGs). To overcome this, we mimic the late-time suppression of star formation in massive halos by ignoring recently formed stars with the aid of a simple post-processing recipe. In this way we find luminosity functions, both for field and group/cluster galaxies, in better agreement with observations. Specifically, the BCGs then follow the observed luminosity-halo mass relation. However, in such a case, the colour bi-modality is basically lost, pointing towards a more complex interplay of late suppression of star formation than what is given by the simple scheme adopted. ",
        "introduction": "\\label{intro}  The observational study of galaxy populations has seen an outstanding progress in recent years. With the advent of large galaxy redshift surveys, such as the Sloan Digital Sky Survey (SDSS; York et al. 2000) and the Two-degree Field Galaxy Redshift Survey (2dFGRS; Colless et al. 2001) it was possible to extend our knowledge of the local Universe to a new level of accuracy.  \\begin{figure*} \\begin{center} {\\includegraphics[width=0.8\\textwidth]{figures/FullSky_Ne_4seb.eps} \\caption{Full-sky map of the simulated local volume in supergalactic coordinates. The map displays the projected gas density distribution up to a distance of $\\sim 80$ $h^{-1}$ Mpc from the observer where main galaxy clusters are shown. Note that Virgo cluster, being the closest one to the observer, is particularly prominent.} } \\end{center} \\label{full_sky} \\end{figure*}   In particular, it has been possible to carry out a robust determination of the luminosity function (LF) for galaxies in the field (Norberg et al. 2002; Blanton et al. 2003) in different spectral bands and to better establish it for those galaxies populating denser environments, such as groups and clusters (e.g. Popesso et al. 2004, 2006). Consistently with previous work (e.g. Lin et al. 1996; Colless et al. 1999) it has been found that the field LF is well described by a single Schechter function requiring, however, a higher value for the luminosity density of the Universe indicating that previous surveys suffered from selection effects and systematics due to photometry (Blanton et al. 2001). For higher density environments, and by means of the RASS-SDSS galaxy cluster survey, Popesso et al. (2004, 2006) found that a double Schechter function better fits the composite cluster LF accounting for a possible upturn in the number of faint objects as the luminosity of member galaxies decrease.  Moreover, the large number of galaxies with measured spectroscopic redshifts, compared to that obtained in the past, has made possible the determination of their two-point clustering properties in a reliable way out to scales of order of tens of Mpc. These studies have shown that the real, projected and redshift-space correlations can be well described by a decreasing power-law function that depends on the sample colour, having a correlation length that increases with absolute magnitude (e.g. Norberg et al. 2001, 2002; Zehavi et al. 2002; Madgwick et al. 2003; Hawkins et al. 2003).   Both the LFs and the clustering properties of the galaxies are tools of fundamental importance in the study of galaxy formation since they provide a way to describe the most basic galaxy statistics. A successful model for structure formation must account for these observations.   From the theoretical point of view, the widely spread {\\it semi-analytic} models (SAMs) of galaxy formation (e.g. Kauffmann et al. 1999; Springel et al. 2001; Mathis et al. 2002; De Lucia, Kauffmann \\& White 2004; Springel et al. 2005; Bower et al. 2006; Cattaneo et al. 2006; Croton et al. 2006; Lagos, Cora \\& Padilla 2008; Fontanot et al. 2009; Guo \\& White 2009) provide a way to study the properties of galaxy populations with the advantage of a rapid exploration of the parameter space. In this approach, the galaxy population is followed within the skeleton provided by a parent dark matter simulation with the aid of physically motivated recipes to describe the different baryonic processes involved. These studies have pointed out the need of limiting excessive gas condensation in massive haloes to avoid the formation of very bright central galaxies (also known as brightest cluster galaxies; BCGs) that are inconsistent with observation. The usually invoked channel responsible for the star formation (SF) quenching in massive haloes is the active galactic nucleus (AGN) phenomenon. Similarly, it has been noted that in order to prevent an excess luminosity in the faint end of the galaxy field LF, a relatively strong SN feedback would be needed in smaller systems.   On the other hand, cosmological hydrodynamical simulations of galaxy samples, without including the effects of AGN feedback, have also been used to study the building up of the structure and their resulting properties in periodic boxes (e.g. Pearce et al. 2001; White et al. 2001; Yoshikawa et al. 2001; Oppenheimer \\& Dav\\'e 2006, 2008; Dav\\'e \\& Oppenheimer 2007; Ocvirk et al. 2008). In a similar way, using the dubbed {\\it zooming} technique, several authors simulated hydrodynamical galaxies within high-resolution regions in a cosmological framework (e.g. Dolag et al. 2005; Saro et al. 2006, 2008, 2009; Crain et al. 2009). In particular, Saro et al. (2006) resimulated a set of galaxy clusters studying the composite cluster LF and the environmental induced galaxy properties in these high density regions. These authors show that it is possible to reproduce the general observed trends of the cluster galaxy population, including, e.g. the colour-magnitude relation, the age and colour cluster-centric distance dependence and the cluster LF (although the bright-end is affected by the presence of very bright BCGs).  \\begin{figure} {\\includegraphics[width=85mm]{figures/att.ps}} \\caption{Estimated optical depth profile in the $V$-band for one of the most massive BCGs in the simulation including extended intra-cluster light ($M_*\\sim10^{13}$ M$_{\\sun}$) averaged over radial projected bins (top panel) and its correspondingly smoothed attenuation fraction due to dust (bottom panel) for the $b_{\\rm j}$- and $K$-bands (open and filled circles respectively).} \\label{tauv} \\end{figure}    The present work makes use of the same {\\it zooming} technique to simulate the observed local volume following the chemical enrichment of gas and stars until the present epoch. This enables us to estimate in a consistent way the luminosities and colours of galaxies formed in hydrodynamical simulations. The simulated volume is big enough to compute LFs, both for field and group/cluster galaxies, as well as galaxy correlation functions. In order to infer the way feedback acts in massive haloes (e.g. AGN feedback), we use a similar approach as SAMs, applying a simple recipe to the post-processed data of the simulation at $z=0$. In this way, we ignore a fraction of the late-formed stars which would not appear in such a scenario and study the resulting effects on the luminosity-dependent properties of the simulated galaxies. It is important to note that, within the simple post-processing scheme adopted here, the simulations are no longer fully self-consistent, as also stars which are quenched by our post-processing procedure still interact with the surrounding medium. However, this effect is not expected to change our results significantly and can only be accessed with the next generation of hydrodynamical simulations including sub-scale models for AGN feedback more directly.  The paper is organized as follows. In Section~\\ref{simul} we describe the hydrodynamical cosmological simulation, together with the method used to compute galaxy luminosities and the associated dust-obscuration. In Section~\\ref{results} we show the main results, presenting luminosities (for field and group/cluster galaxies), galaxy colours and correlation functions. We also discuss the implemented recipe to {\\it suppress} SF in massive haloes. Finally, we close the paper with a summary and our conclusions in Section~\\ref{concl}.  \\begin{table} \\begin{center} \\caption{Simple SF suppression recipe: $z_{\\rm cut}=(V_{\\rm max}-V_0)/V_1$, where $z_{\\rm cut}$ represents the redshift since which we do not consider star formation in a given galaxy. Systems having $V_{\\rm max}<V_0$ are not affected by the suppression.} \\vspace{0.1cm} \\label{table1} \\begin{tabular}{cccc} \\hline \\hline Model & $V_0$ & $V_1$\\\\ & [km s$^{-1}$] &   [km s$^{-1}$]\\\\ \\hline 1   & 100 & 150 \\\\ 2   & 100 & 200 \\\\ 3   & 100 & 250 \\\\ 4   & 150 & 150 \\\\ 5   & 150 & 200 \\\\ 6   & 150 & 250 \\\\ 7   & 200 & 150 \\\\ 8   & 200 & 200 \\\\ 9   & 200 & 250 \\\\ \\hline \\hline \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": "\\label{concl}  In this paper, we have focused on the luminosity, colour and clustering properties of galaxies that arise naturally in a hydrodynamical cosmological simulation of structure formation. The spherical local volume of $\\sim130^3$ $h^{-3}$ Mpc$^3$ was simulated using the {\\it zooming} technique with more than $\\sim 50$ million gas and dark matter particles. In particular, our simulation code include different physical mechanisms such as supernova feedback, galactic winds, an UV pervading background since $z\\sim6$ and chemical enrichment of the gas and stellar populations as time elapses. However, we have not explicitly included a physical mechanism capable of quenching SF in massive haloes (like AGNs). In this way, by $z=0$ a total of $\\sim20000$ objects having more than 32 star particles are formed in the volume. The luminosities of the hydrodynamical galaxies were computed using the differential contribution provided by every star particle present inside the optical radius of the systems. As a way of estimating the dust-attenuated luminosity in the bands of interest, we implemented a simple model for dust extinction that makes use of the cold gas distribution in each galaxy as a tracer of its dust content. In order to mimic the required SF suppression in massive (i.e. $V_{\\rm vir} \\gtrsim 150$ km s$^{-1}$) haloes we also applied a simple recipe to the present day output of the simulation to reconcile the bright-end of the simulated LFs with observations and to further investigate its inclusion on other properties of the simulated galaxies. In the following we summarize the main conclusions of the present work:  \\begin{itemize} \\item As shown earlier by semi-analytic work, it is essential to suppress SF in massive haloes to describe the bright-end of the field LF at $z\\sim0$. For instance, haloes with masses corresponding to $V_{\\rm vir}\\sim300$ km s$^{-1}$ are required to quench their SF since $z\\sim1$ in the {\\it best-fitting} SF suppression case. In this way, it is possible to avoid the formation of very bright blue BCGs in the simulation. However, after a small renormalization to $L_*$, the faint-end of the simulated field LFs tends to overpredict the number density of galaxies per magnitude bin in the $b_{\\rm j}$-band for galaxies with luminosities lower than $\\sim L_*$. To better study this effect higher resolution simulations are needed.\\\\  \\item The group/cluster bright-end luminosity distributions are better described when SF suppression is considered. This is due to the fact that massive cluster satellites used to be BCGs in the past, and also require some level of SF suppression. However, as a consequence of our relatively high-mass galaxy resolution limit we cannot confirm the upturn of faint cluster satellites in the cluster LF function as previously obtained by Saro et al. (2006, 2009) using high-resolution resimulations of galaxy clusters. On the other hand, the $K$-band luminosity-halo mass relation that naturally arises in the simulation also needs further SF quenching in massive haloes to better describe observations given by Brough et al. (2008).\\\\  \\item The stellar masses of the hydrodynamical galaxies in the plain simulation display a clear $B-V$ colour bi-modal behavior as suggested by observations (e.g. Baldry et al. 2006), although the {\\it active} blue galaxies present very massive systems, due to overcooling, contrary to what data shows. When taking into account the SF suppression scheme, the $B-V$ galaxy colours turn redder, displaying a colour-stellar mass distribution that is not able to fully account for observations. In this case, a bi-modal distribution is observed for stellar masses lower than $\\sim2\\times10^{10}$ $h^{-1}$ M$_{\\odot}$, while for higher masses most of the galaxies are red. This is a manifestation of the rough nature of the simple parametrization adopted to quench SF in massive haloes.\\\\  \\item The clustering properties of different hydrodynamical galaxy populations that naturally arise in the simulation are in qualitative agreement with observational results. The scale-dependent correlation function of our $\\sim L_*$ objects in the $b_{\\rm j}$-band displays a power-law behavior up to $\\sim10~h^{-1}$ Mpc that is in good agreement with data given by Norberg et al. (2001). As expected, higher luminosity systems show a stronger clustering pattern, where galaxies having luminosities around $\\sim 4L_*$ display an excellent agreement with observations. On the other hand, it is also possible to disentangle the galaxy correlation properties when the sample is splitted by galaxy colours. The correlation function of {\\it active} blue ($B-V<0.8$) systems is in rough agreement with the results given by Madgwick et al. (2003). However, {\\it passive} red systems ($B-V>0.8$) present an excess of power in comparison with observations, showing a clear departure from the power-law behavior. As shown when using the SF suppression scheme, this discrepancy is alleviated when more galaxies are able to populate the {\\it red sequence} of the colour-magnitude diagram. This is due to the fact that, typically, the added systems tend to reside in less clustered haloes, thus lowering the correlation length value.\\\\   In summary, we conclude that analyzing the global properties of the galaxy population within hydrodynamical, cosmological simulations, start to be a promising tool to study galaxy evolution. Current simulations already fairly represent the underlying hydrodynamical effects (at least in a global sense) and, in general, describe the star formation process well enough to qualitatively reproduce observed environmental trends. However, similar than for the widely used SAMs, overcooling in  massive halos has to be quenched by additional feedback effects. It has to be seen in future simulations, which directly include such additional processes, if such a quenching happens mildly enough not to destroy some of general trends already captured in the current generation of hydrodynamical simulations.   \\end{itemize}"
    },
    "1004/1004.3548_arXiv.txt": {
        "abstract": "The linear growth factor of density perturbations is generally believed to be a powerful observable quantity of future large redshift surveys to probe physical properties of dark energy and to distinguish among various gravity theories. We investigate systematic effects on determination of the linear growth factor $f$ from a measurement of redshift-space distortions. Using a large set of high-resolution $N$-body simulations, we identify dark matter halos over a broad mass range. We compute the power spectra and correlation functions for the halos and then investigate how well the redshift distortion parameter $\\beta\\equiv f/b$ can be reconstructed as a function of halo mass both in Fourier and in configuration space, where $b$ is the bias parameter.  We find that the $\\beta$ value thus measured for a fixed halo mass generally is a function of scale for $k>0.02\\hmpci$ in Fourier space or $r<80\\himpc$ in configuration space, in contrast with the common expectation that $\\beta$ approaches a constant described by Kaiser's formula on the large scales. The scale dependence depends on the halo mass, being stronger for smaller halos. It is complex and cannot be easily explained with the exponential distribution function in configuration space or with the Lorentz function in Fourier space of the halo peculiar velocities. We demonstrate that the biasing for smaller halos has larger nonlinearity and stochasticity, thus the linear bias assumption adopted in Kaiser's derivation become worse for smaller halos. Only for massive halos with the bias parameter $b\\geq 1.5$, the $\\beta$ value approaches the constant predicted by the linear theory on scales of $k<0.08\\hmpci$ or $r>30\\himpc$. Luminous red galaxies (LRGs), targeted by the Sloan Digital Sky Survey (SDSS) and the SDSS-III's Baryon Oscillation Spectroscopic Survey (BOSS), tend to reside in very massive halos. Our results indicate that if the central LRG sample is used for the measurement of redshift-space distortions, fortunately the linear growth factor can be measured unbiasedly.  On the other hand, emission line galaxies, targeted by some future redshift surveys such as the BigBOSS survey, are inhabited in halos of a broad mass range.  If one considers to use such galaxies, the scale dependence of $\\beta$ must be taken into account carefully; otherwise one might give incorrect constraints on dark energy or modified gravity theories. We also find that the $\\beta$ reconstructed in Fourier space behaves fairly better than that in configuration space when compared with the linear theory prediction. ",
        "introduction": "The presence of dark energy, which changes the gravitational assembly history of matter in the universe, explains observed acceleration of the cosmic expansion well within the framework of general relativity \\citep[see Komatsu et al. 2010 for the latest constraints]{Riess1998, Perlmutter1999, Spergel2003}.  There are also many attempts to explain the acceleration without dark energy by modifying general relativity on cosmological scales \\citep[see, e.g.,][]{Dvali2000, Carroll2004}.  Cosmological models in different gravity theories that predicts a similar expansion rate $H(z)$, can have the different cosmic growth rate $f(z)$. The $f(z)$ is often parameterized as $f(z)=\\Omega_m^\\gamma(z)$ where $\\Omega_m(z)$ is the mass density parameter at a given redshift $z$ and the growth index $\\gamma\\simeq 0.55 $ in the $\\Lambda$CDM model \\citep{Linder2005}.  Thus the precise measurement of the growth rate enables us to investigate the deviation of gravity from the general relativity.  Recent analysis which focused on such deviations using weak gravitational lensing data, cosmic microwave background data, and type Ia supernova data, showed a good agreement with the pure $\\Lambda$CDM model \\citep[e.g., see][for the latest work]{Daniel2010}.  One of the most promising tools to investigate modified gravity is redshift-space distortion effects caused by peculiar velocities in galaxy redshift surveys.  In linear theory and under the plane-parallel approximation, \\citet{Kaiser1987} derived a formula to relate the observed power spectrum of galaxies $P^{(s)}(k,\\mu_{\\bf k})$ and the true power spectrum of dark matter $P^{(r)}_m(k)$ through \\begin{equation} P^{(s)}(k,\\mu_{\\bf k})=b^2(1+\\beta\\mu_{\\bf k}^2)^2P^{(r)}_m(k), \\label{eq:kaiser} \\end{equation} where $(r)$ and $(s)$ respectively denote quantities in real and redshift space, $\\mu_{\\bf k}$ is the cosine of the angle between the line of sight and the wavevector ${\\bf k}$, $\\beta$ is the linear redshift distortion parameter related to the growth rate as $\\beta=f/b$, and $b$ is the bias parameter \\citep{Kaiser1984}.  Thus the measurement of the redshift-space distortions allows one to directly probe deviations from general relativity, although the determination of the biasing is another important issue.  The Kaiser's formula (equation (\\ref{eq:kaiser})) is modified on small scales because the nonlinear random velocities smear the clustering along the line of sight known as the `finger-of-god' effect \\citep{Jackson1972}. However such a nonlinear model still relies on Kaiser's formula on large scales \\citep{Peacock1994}.  For the importance of nonlinearity on such scales, see \\citet{Scoccimarro2004} \\citet{Taruya2009}, \\citet{Desjacques2010}, and \\citet{Jennings2010}  Constraints on $\\beta$ have been reported in various surveys \\citep[e.g.,][]{Peacock2001, Zehavi2002, Hawkins2003, Tegmark2004, Tegmark2006, Ross2007, Guzzo2008, da Angela2008, Cabre2009}. \\citet{Okumura2008} also showed using a luminous red galaxy sample from the Sloan Digital Sky Survey (SDSS) that simultaneously analyzing redshift-space distortions and anisotropy of the baryon acoustic scales allows one to give a strong constraint on dark energy equation-of-state, as was theoretically predicted \\citep{Hu2003, Seo2003, Matsubara2004} and this fact was explicitly emphasized by \\citet{Amendola2005}. \\citet{Guzzo2008} considered constraints on $f$ to test the deviation from general relativity using the observations at different redshifts \\citep[see also,][] {Di Porto2008, Nesseris2008, Yamamoto2008, Cabre2009}. We note that all these previous studies have used linear theory prediction of redshift-space distortions to compare with their measurements on scales presumably large enough for the linear theory to be valid.  \\citet{Nakamura2009} adopted an alternative approach and constrained the growth factor by measuring the damping of the baryon acoustic oscillations. \\citet{Reyes2010} gave a strong constraint on modified gravity theory and confirmed general relativity using the method proposed by \\citet{Zhang2007} which can eliminate the uncertainty of the galaxy biasing by combining weak gravitational lensing, galaxy clustering, and redshift-space distortions \\citep[for similar theoretical attempts, see e.g.,][]{Percival2009, McDonald2009, Song2010}.  There are many ongoing and upcoming large galaxy surveys, such as the SDSS-III's Baryon Oscillation Spectroscopic Survey \\citep[BOSS;][]{Schlegel2009a}, the Fiber Multiobject Spectrograph \\citep[FMOS;][]{Sumiyoshi2009}, the Hobby-Eberly Dark Energy Experiment \\citep[HETDEX;][]{Hill2004}, the WiggleZ \\citep{Glazebrook2007}, the BigBOSS \\citep{Schlegel2009b}, and so on. These observations are expected to enable us to distinguish among gravity theories with high precision through measurement of redshift-space distortions as well as that of baryon acoustic oscillations. However, it is not clear if the accuracy of predicting redshift-space distortions is better than or comparable to the precision required from future surveys.  In addition, we do not know how large the deviation from true cosmology is if any. Precision of the constraint may depend on galaxy types, such as luminosity and host halo mass.  There were many attempts to investigate the validity to use the redshift-space distortions to extract the cosmological information \\citep[e.g.,][]{Hatton1998, Hatton1999, Berlind2001, Tinker2006}.  \\citet{Tinker2006} found that $\\beta$ can be estimated accurately using linear theory if the finger-of-god effect is removed perfectly.  In this paper, we present a detailed study on this aspect using a large set of $N$-body simulations. We measure the power spectra and correlation functions of dark matter halos. Using them, we estimate the redshift distortion parameter $\\beta$ from the monopole-to-real-space ratio and the quadrupole-to-monopole ratio, both of which are related to $\\beta$ in linear theory.  Then we examine whether $\\beta$ measured in these ways can give true cosmological information.  We also investigate the dependence of the precision of reconstructed $\\beta$ on halo mass. Particularly we will clearly show that the $\\beta$ value obtained from the small-halo clustering does not approach a constant even on large scales as linear theory predicts.  In addition such small halos are shown to be more stochastic tracers of the underlying density field than massive halos. We also discuss in detail on which scale and with which method one can get the correct $\\beta$ or $f$ from the redshift-space distortions.  The paper is organized as follows. In Section \\ref{sec:nbody}, we describe the $N$-body simulations and the halo occupation distribution model used to populate them with mock galaxies.  The basic two-point statistics used in our analysis are also presented.  In Section \\ref{sec:theory} we briefly review linear theory of redshift distortions and how to determine the redshift distortion parameter $\\beta$ from the power spectrum and the correlation function.  Section \\ref{sec:result} is devoted to the analysis of redshift distortion effects in simulations to determine $\\beta$ and the growth rate $f$. Our conclusions are given in Section \\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion} We have investigated how accurately the redshift-space distortions can be used to measure the linear growth factor $f$. The growth factor is a powerful observable targeted by future large redshift surveys to probe dark energy and to distinguish among different gravity theories. For this purpose, we constructed a large set of $N$-body simulations, dividing each dark matter halo catalog into the subsamples with narrow mass ranges.  As an example of a galaxy sample, mock SDSS LRG catalogs were constructed by applying the HOD modeling to the simulated halos. Then we have measured the two-point statistics, power spectra and correlation functions, for dark matter halos and LRGs. The dark matter halos were analyzed as a function of halo mass in order to see dependence of the $\\beta$ measurement on the halo mass.  We have determined the $\\beta$ values as a function of halo mass and scale using four methods. First, we found that $\\beta$ reconstructed from the ratio of the monopole to the real-space power spectra $P^{(0/r)}=P^{(s)}/P^{(r)}$ (equation (\\ref{eq:pk_mono})) asymptotically approaches the true value.  In particular for the massive halos and LRGs, the prediction from linear theory known as Kaiser's formula is applicable to give a correct constraint on the growth rate. However, for less massive halos, the ratio approaches the true value only at a very large scale $k<0.02\\hmpci$.  Second, $\\beta$ reconstructed from the ratio of the monopole to the real-space correlation function $\\xi^{(0/r)}=\\xi^{(s)}/\\xi^{(r)}$ (equation (\\ref{eq:xi_mono})) approaches neither the true value nor a constant even on large scales. This statement is valid especially for small halos with the bias parameter $b\\leq 1$. Because the growth rate is assumed to be a constant when modified gravity theories are tested, the ratio $\\xi^{(s)}/\\xi^{(r)}$ cannot be used in a simple way for this purpose.  Third, the quadrupole-to-monopole ratio in Fourier space $P^{(2/0)}=P_2/P_0$ (equation (\\ref{eq:pk_quad})) gives almost the same value of $\\beta$ as $P^{(0/r)}$ but with larger error bars as expected.  Finally, we found that when the quadrupole-to-monopole ratio in configuration space $\\xi^{(2/0)}=\\xi_2/\\xi_0$ (equation (\\ref{eq:xi_quad})) is used, a similar conclusion is reached to that of $P^{(2/0)}$ when $r=\\lambda=\\pi/k$ is adopted.  For small halos with $b\\leq 1.3$, the reconstructed $\\beta$ values do not approach a constant in most of measurable regions, particularly those from $\\xi^{(0/r)}$ in the configuration space.  No method can provide a reliable estimator for the determination of the growth factor from the clustering of such small halos on the large range of scales probed.  Using the halo catalogs with different box sizes, we confirmed that such a behavior is not caused by the resolution effect of small dark matter halos.  While the scale dependence changes with the halo mass, the peculiar velocity of halos does not change much with the mass \\citep{Hamana2003}.  Using the simple dispersion model, we demonstrated that the different scale dependence of $\\beta$ among small and large halos cannot be simultaneously explained. Also there are two types of velocity biases which affect the redshift distortion; the dynamical bias caused by dynamical friction and the spatial bias caused by the difference between the distribution of halos and that of dark matter. There is no dynamical velocity bias because the halo velocities are determined from the mean velocities of dark matter within the halos in our analysis. The spatial velocity bias should exist, which is coupled with the nonlinear stochastic bias discussed in the text.  On the other hand, it is known that the clustering of small dark matter halos depends not only on their mass but also on their assembly history, so called the assembly bias \\citep[e.g.,][]{Gao2005}. \\citet{Wang2007} showed that old and low-mass halos that are preferentially associated with a high density field are more strongly clustered than young halos with the same mass, and consequently have higher velocities.  Besides, the stochasticity between halos and dark matter is \\citep{Dekel1999} can be a source of the systematic errors in the $\\beta$ reconstruction.  Using a method introduced by \\citet{Taruya2000} and applied to simulation data by \\citet{Yoshikawa2001}, we have found that both the nonlinearity and the stochasticity of small halos become larger than massive halos. Particularly the stochastic bias monotonically increases as the mass of halos decreases, as was found in real space by \\citet{Hamaus2010} using the two-point statistics.  Thus the strong scale dependence of $\\beta$ for low mass halos could be caused by the assembly and/or nonlinear stochastic bias. However fortunately, the scale dependence of the measured $\\beta$ weakens with the increase of halo mass. For massive halos with $b>1.5$, the measured $\\beta$ approaches the constant predicted by Kaiser's formula on scales $k<0.08 \\hmpci$ or $r>30 \\himpc$.  Because the analysis of redshift-space distortions is powerful to investigate not only properties of dark energy but also modified gravity theories, it will keep on playing a key role in ongoing and upcoming large redshift surveys, such as BOSS, FMOS, HETDEX, WiggleZ, and BigBOSS. Galaxies targeted by the BOSS survey are luminous red galaxies, which reside in massive halos. In this work we demonstrate the $\\beta$ value can be well reconstructed with a redshift distortion analysis of LRGs. On the other hand, one of the samples targeted by the BigBOSS, for example, is that of emission-line galaxies, which reside in halos with a broad range of halo mass. One needs to be careful in using such a sample to constrain the growth rate from the redshift distortion, because it can be a scale-dependent function. While one might be able to obtain a result consistent (or inconsistent) with the prediction from general relativity, it could be just a coincidence after the scale-dependent growth rate is averaged over some separation or wavenumber ranges. We will use semi-analytical modeling or a halo occupation model to investigate this issue in future work.  Recently, an interesting method was proposed by \\citet{Seljak2009} to suppress the shot noise in power spectrum measurement.  They considered an optimal weighting function $f(M)$ in measuring the galaxy overdensity, where they give higher weights on higher mass halos.  Compared with our results presented here, such a mass weighting scheme is useful not only for suppressing the shot noise but also obtaining the true value of the growth rate. This scheme can be naturally incorporated into our method and such a study will be presented in our future paper.  Another theoretical improvement to be applied for observation is evading the cosmic variance limit, which is one of the most important tasks for precise measurement of the redshift-space distortions, as we have already seen above. \\citet{McDonald2009} showed using multiple tracers of density with different biases suppresses the noise for measurement of $\\beta$ on large scales dramatically compared to the traditional single tracer method \\citep[see also][]{White2009, Gil-Marin2010}. But the different scale-dependent properties of $\\beta$ for different halo masses found in Figure 4 imply that the real situation might be more complex, and realistic models of galaxies must be adopted to investigate if the method of multiple tracers works."
    },
    "1004/1004.1856_arXiv.txt": {
        "abstract": "Considerable progress has been made in determining the Hubble constant over the past two decades.  We discuss the cosmological context and importance of an accurate measurement of the Hubble constant, and focus on six high-precision distance-determination methods: Cepheids, tip of the red giant branch, maser galaxies, surface-brightnes fluctuations, the Tully-Fisher relation and Type Ia supernovae. We discuss in detail known systematic errors in the measurement of galaxy distances and how to minimize them. Our best current estimate of the Hubble constant is 73 $\\pm$2 (random) $\\pm$4 (systematic)~km~s$^{-1}$~Mpc$^{-1}$. The importance of improved accuracy in the Hubble constant will increase over the next decade with new missions and experiments designed to increase the precision in other cosmological parameters. We outline the steps that will be required to deliver a value of the Hubble constant to 2\\% systematic uncertainty and discuss the constraints on other cosmological parameters that will then be possible with such accuracy. ",
        "introduction": "In 1929 Carnegie astronomer, Edwin Hubble, published a linear correlation between the apparent distances to galaxies and their recessional velocities. This simple plot provided evidence that our Universe is in a state of expansion, a discovery that still stands as one the most profound of the twentieth century (Hubble 1929a).  This result had been anticipated earlier by Lema{\\^i}tre (1927), who first provided a mathematical solution for an expanding universe, and noted that it provided a natural explanation for the observed receding velocities of galaxies.  These results were published in the Annals of the Scientific Society of Brussels (in French), and were not widely known.  Using photographic data obtained at the 100-inch Hooker telescope situated at Mount Wilson CA, Hubble measured the distances to six galaxies in the Local Group using the Period-Luminosity relation (hereafter, the Leavitt Law) for Cepheid variables. He then extended the sample to an additional 18 galaxies reaching as far as the Virgo cluster, assuming a constant upper limit to the brightest blue stars (HII regions) in these galaxies. Combining these distances with published radial velocity measurements (corrected for solar motion) Hubble constructed Figure 1. The slope of the velocity versus distance relation yields the Hubble constant, which parameterizes the current expansion rate of the Universe.  The Hubble constant is usually expressed in units of kilometers per second per megaparsec, and sets the cosmic distance scale for the present Universe.  The inverse of the Hubble constant has dimensions of time.  Locally, the Hubble law relates the distance to an object and its redshift: cz = H$_0$d, where d is the distance to the object and z is its redshift.  The Hubble law relating the distance and the redshift holds in any Friedman-Lemaitre-Robertson-Walker cosmology (see \\S \\ref{sec:cosmology}) for redshifts less than unity.  At greater redshifts, the distance-redshift relationship for such a cosmology also depends on the energy densities of matter and dark energy.  The exact relation between the expansion age and the Hubble constant depends on the nature of the mass-energy content of the Universe, as discussed further in \\S \\ref{sec:cosmology} and \\S \\ref{sec:age}. In a uniformly expanding universe, the Hubble parameter, H(t), changes as a function of time; H$_\\circ$, referred to as the Hubble constant, is the value at the current time, $t_\\circ$.  Measurement of the Hubble constant has been an active subject since Hubble's original measurements of the distances to galaxies: the deceptively simple correlation between galaxy distance and recession velocity discovered eighty years ago did not foreshadow how much of a challenge large systematic uncertainties would pose in obtaining an accurate value for the Hubble constant.  Only recently have improvements in linear, solid-state detectors, the launch of the Hubble Space Telescope (HST), and the development of several different methods for measuring distances led to a convergence on its current value.  Determining an accurate value for H$_o$ was one of the primary motivations for building HST. In the early 1980's, the first director of the Space Telescope Science Institute, Riccardo Giacconi, convened a series of panels to propose observational programs of significant impact requiring large amounts of Hubble observations. He was concerned that in the course of any regular time allocation process there would be reluctance to set aside sufficient time to complete such large projects in a timely manner.  For decades a `factor-of-two' controversy persisted, with values of the Hubble constant falling between 50 and 100 km s$^{-1}$ Mpc$^{-1}$.  A goal of 10\\% accuracy for H$_o$ was designated as one of HST's three ``Key Projects''. (The other two were a study of the intergalactic medium using quasar absorption lines, and a ``medium-deep'' survey of galaxies.)  This review is organized as follows: We first give a brief overview of the cosmological context for measurements of the Hubble constant.  We discuss in some detail methods for measuring distances to galaxies, specifically Cepheids, the tip of the red giant branch (TRGB), masers, the Tully-Fisher relation and Type Ia supernovae (SNe~Ia).  We then turn to a discussion of H$_o$, its systematic uncertainties, other methods for measuring H$_o$, and future measurements of the Hubble constant.  Our goal is to describe the recent developments that have resulted in a convergence to better than 10\\% accuracy in measurements of the Hubble constant, and to outline how future data can improve this accuracy. For wide-ranging previous reviews of this subject, readers are referred to those of Hodge (1982), Huchra (1992), Jacoby et al. (1992), van den Bergh (1992), Jackson (2007), and Tammann, Sandage \\& Reindl (2008).  An extensive monograph by Rowan-Robinson (1985) details the history of the subject as it stood twenty-five years ago. ",
        "conclusions": "(1) Several nearby distance determination methods are now available that are of high precision, having independent systematics. These include Cepheid variables, the tip of the red giant branch (TRGB) stars, and the geometrically determined distances to maser galaxies.  (2) The Cepheid Period-Luminosity relation (Leavitt Law) now has an absolute calibration based on HST trigonometric parallaxes for Galactic Cepheids. This calibration and its application at near-infrared wavelengths significantly reduces two of the four leading systematic errors previously limiting the accuracy of the Cepheid-based distance scale: zero-point calibration and metallicity effects.  (3) The maser galaxy distances, TRGB distances and Cepheid distances agree to high precision at the one common point of contact where they can each be simultaneously intercompared, the maser galaxy NGC 4258, at a distance of 7.2  Mpc.  (4) Galactic Cepheid parallax and NGC 4258 maser calibrations of the distance to the LMC agree very well.  Based on these measurements and other independent measurements, we adopt a true, metallicity-corrected distance modulus to the LMC of 18.39 $\\pm$ 0.06 mag.  (5) HST optical and near-infrared observations of Cepheids in SNe Ia galaxies calibrated by the maser galaxy, NGC 4258, have decreased systematics due to calibration, metallicity and reddening in  the SNe Ia distance scale, and increased the number of well-observed SN calibrators to six.  (6) The current calibration of the Cepheid and maser extragalactic distance scales agree to within their quoted errors, yielding a value of $H_\\circ$ = 73 $\\pm$ 2 (random) $\\pm$ 4 (systematic) km~s$^{-1}$~Mpc$^{-1}$.  (7) Within a concordance cosmology (that is, $\\Omega_{matter}$ = 0.27 and $\\Omega_{vacuum}$ = 0.73) the current value of the Hubble constant gives an age for the Universe of 13.3 $\\pm$ 0.8~Gyr. Several independent methods (globular cluster ages, white dwarf cooling ages, CMB anisotropies within a concordance model) all yield values in good agreement with the expansion age.  (8) Further reductions of the known systematics in the extragalactic distance scale are anticipated using HST, {\\it Spitzer}, GAIA and JWST. A factor of two decrease in the currently identified systematic errors is within reach, and an uncertainty of 2\\% in the Hubble constant is a realistic goal for the next decade.  (9) A Hubble constant measurement to a few percent accuracy, in combination with measurements of anisotropies in the cosmic microwave background from Planck, will yield valuable constraints on many other cosmological parameters, including the equation of state for dark energy, the mass of neutrinos, and the curvature of the universe."
    },
    "1004/1004.3853_arXiv.txt": {
        "abstract": "{During the evolution of rotating first stars, which initially consisted of only hydrogen and helium, CNO elements may emerge to their surface. These stars may therefore have winds that are driven only by CNO elements.} {We study weak wind effects (Gayley-Owocki heating and multicomponent effects) in stellar winds of first generation stars driven purely by CNO elements.} {We apply our NLTE multicomponent models and hydrodynamical simulations.} {The multicomponent effects (frictional heating and decoupling) are important particularly for low metallicity winds, but they influence mass loss rate only if they cause decoupling for velocities lower than the escape velocity. The multicomponent effects also modify the feedback from first stars. As a result of the decoupling of radiatively accelerated metals from hydrogen and helium, the first low-energy cosmic ray particles are generated. We study the interaction of these particles with the interstellar medium concluding that these particles easily penetrate the interstellar medium of a given minihalo. We discuss the charging of the first stars by means of their winds.} {Gayley-Owocki heating, frictional heating, and the decoupling of wind components occur in the winds of evolved low-metallicity stars and the solar metallicity main-sequence stars.} ",
        "introduction": "Many aspects of chemical evolution in our Universe remain unclear. The first elements, helium and a very small amount of, e.g.,~lithium, were almost certainly synthesised during the era of primordial nucleosynthesis \\citep[e.g.,][]{coc}. The first stars in the Universe can therefore be considered to be purely hydrogen-helium stars \\citep[see][for a review]{lesdiablerets}.  Subsequent chemical evolution is less clear, partly because it is difficult to test the theoretical predictions observationally. For example, stars with a very low abundance of iron are observed, which are expected to be relics from ancient times \\citep[e.g.,][]{starka}. However, their relevance as observational testbeds to the theory of evolution of chemical composition may be hampered by several secondary effects \\citep{torna}. Several hypotheses have been developed to explain the chemical composition of these stellar relics \\citep[cf.,][]{sty,medaci,mee}.  Hot star winds are supposed to play an important role in the chemical evolution of our Universe. Since they remove material from the outer stellar envelopes, they only affect the stellar mass during the early phases of stellar evolution and do not contribute to the change in the chemical composition of the interstellar medium. On the other hand, as soon as freshly synthesised elements emerge at the stellar surface during later phases of stellar evolution \\citep{mee,samhir}, the hot star winds may contribute to the chemical evolution of the interstellar medium even before the star possibly explodes as a supernova.  Hot star winds are studied mainly by assuming a solar mixture of elements and information about winds of more exotic composition is scarce \\citep{vikowr,grahamz,uncno}. An interesting mixture of heavier elements, which is uncommon in contemporary Universe, is represented by a pure CNO composition. This composition may be typical of later phases in the evolution of the first stars. This is connected to the possibility that the envelopes of the first stars in later evolutionary phases are enriched by the products of helium burning \\citep{mee,samhir}. A chemical mixture rich in CNO elements and underabundant in iron is typical for one group of low-metallicity stars \\citep[e.g.,][]{pannorris,starka}.  The study of CNO driven winds is important not only for early stellar generations \\citep[cf.,][]{uncno}. The low density winds of present stars are also accelerated mostly by CNO lines because the contribution of other heavier elements is relatively small \\citep[e.g.,][]{vikolamet}.%  To understand the role of CNO driven winds in hot evolved first stars, \\citet[hereafter \\citetalias{cnovit}]{cnovit} calculated wind models of these stars. They concluded that CNO elements do not drive winds as efficiently as iron peak elements because of the lower number of their strong lines. Therefore, the total amount of mass lost by these winds does not significantly affect stellar evolution.  For subsequent stellar generations, the wind enrichment of primordial halos by heavier elements does not overcome the metallicity threshold for the formation of very massive stars. On the other hand, the enrichment could be large enough to change the behavior of primordial stars during their formation -- a mass fraction of CNO higher than about $10^{-10}$ is sufficient to enable hydrogen burning via CNO cycle and preclude initial helium burning by means of the 3$\\alpha$ reaction \\citep{vittorio}.  Some  CNO driven winds (especially the low-metallicity ones) may be subject to weak wind effects. For weak winds, two effects that are negligible for high-density winds may become important, namely the Gayley-Owocki (Doppler) heating/cooling \\citep[hereafter \\citetalias{go}]{go} and multicomponent effects \\citep[hereafter \\citetalias{kkii}]{kkii}.  The GO heating/cooling is caused by a frequency difference between photons entering and escaping the Sobolev resonance zone. Multicomponent effects are connected to the momentum transfer between heavier elements (accelerated by line absorption) and bulk wind material, i.e., hydrogen and helium. In low-density winds, momentum transfer may become inefficient, causing frictional heating or even decoupling of wind components \\citep{cak76,treni,kkii,op,ufo,uncno}.  To understand the role of weak wind effects in the CNO winds of massive first stars, we calculate models of the multicomponent winds of these stars for which GO heating is taken into account. ",
        "conclusions": "We have studied the effect of the Gayley-Owocki (Doppler) heating and multicomponent flow structure in CNO driven winds of hot stars. The parameters of these stars were selected to represent massive initially pure hydrogen-helium (Pop~III) stars.  For the first time, we have included the GO heating term directly using atomic linelist and NLTE calculations. We have shown that GO heating is important especially for winds of CNO enriched first stars with high metallicities ($Z\\approx0.01$). In these winds, the GO heating can compete with radiative cooling because both the number of strong lines and the wind terminal velocity are large. On the other hand, for stars with low metallicities ($Z\\lesssim0.001$) there are an insufficient number of strong lines and the adiabatic cooling dominates.  The effects of multicomponent flows are important especially at low metallicities ($Z\\lesssim0.001$) in the case of evolved stars, and for relatively high metallicities ($Z\\approx0.01$) for main-sequence stars. The frictional heating itself does not influence the wind mass-loss rate. On the other hand, decoupling probably leads to a zero mass-loss rate of hydrogen and helium if it occurs at velocities lower than the escape one. We have developed an approximate formula that estimates the minimum metallicity above which hydrogen and helium leave the star.  The decoupling of radiatively accelerated metals from hydrogen and helium leads to generation of particles with typical energies of the order of 1\\,MeV, i.e., the first stars may be the first sources of low-energy cosmic rays. We have shown that these particles easily penetrate the interstellar medium of a given minihalo. We have discussed the possibility of charging of first stars via their multicomponent winds.  Wind models presented here can also be used to describe the winds of possible subsequent generations of CNO rich stars and low-luminosity stars of solar chemical composition."
    },
    "1004/1004.2158_arXiv.txt": {
        "abstract": "{} We calculate and analyze the longevity of magnetohydrodynamic (MHD) wave modes that occur in the plane of a magnetic thin sheet. Initial turbulent conditions applied to a magnetically subcritical cloud are shown to lead to relatively rapid energy decay if ambipolar diffusion is introduced at a level corresponding to partial ionization primarily by cosmic rays. However, in the flux-freezing limit, as may be applicable to photoionized molecular cloud envelopes, the turbulence persists at ``nonlinear'' levels in comparison with the isothermal sound speed $\\cs$, with one-dimensional rms material motions in the range of $\\approx 2\\,\\cs - 5\\,\\cs$ for cloud sizes in the range of $\\approx 2\\,\\pc - 16\\,\\pc$. These fluctuations persist indefinitely, maintaining a significant portion of the initial turbulent kinetic energy. We find the analytic explanation for these persistent fluctuations. They are magnetic-tension-driven modes associated with the interaction of the sheet with the external magnetic field. The phase speed of such modes is quite large, allowing residual motions to persist without dissipation in the flux-freezing limit, even as they are nonlinear with respect to the sound speed. We speculate that long-lived large-scale MHD modes such as these may provide the key to understanding observed supersonic motions in molecular clouds. ",
        "introduction": "Nonthermal linewidths are ubiquitous in molecular clouds \\citep{sol87} and are interpreted to represent highly supersonic random internal motions \\citep[see][for a recent review]{mck07}. Principal component analysis \\citep{bru09} reveals that most of the energy is contained in modes that span the largest scale of the cloud.  Since molecular clouds are threaded by large-scale magnetic fields, an attractive suggestion has been that the turbulence represents supersonic but sub-\\Alfvc magnetohydrodynamic (MHD) waves, with the noncompressive shear \\Alf mode identified as a possible long-lived component \\citep{aro75}. This was intended to bypass the usual problem of rapid dissipation of supersonic hydrodynamic turbulence through shocks. A compilation of available Zeeman measurements of the (line-of-sight) magnetic field strength, gas density $\\rho$, and one-dimensional velocity dispersion $\\sigma_v$ shows that cloud fragments obey a direct linear correlation between $\\sigma_v$ and the mean \\Alf speed $V_{\\rm A}$ \\citep{mye88,bas00}, with a proportionality coefficient $\\approx 0.5$. This correlation may be attributed to a rough equality of the magnitudes of gravitational, magnetic, and turbulent energies, and was interpreted by \\citet{mou95} and \\citet{mou06} to mean that the motions are \\Alfvc disturbances in which the perturbed magnetic field is comparable in strength to the background magnetic field. Furthermore, the MHD motions in a molecular cloud may represent long-wavelength standing waves, as argued by \\citet{mou75,mou87}. This brings out the possibility of ``global'' effects (e.g. due to cloud boundaries and external interaction) being important in understanding cloud turbulence, and the need to go beyond comparing observations with models of wave propagation in an infinite medium.  In direct contrast to the scenario of long-lived motions, numerical simulations of molecular cloud turbulence using a three-dimensional simulation cube with periodic boundary conditions have revealed that supersonic MHD turbulence decays away rapidly, like its hydrodynamic counterpart, on about a sound crossing time of the driving scale \\citep{sto98,mac98,mac99,ost01}. This happens in either the case of sub-\\Alfvc or super-\\Alfvc turbulence, and in both cases, turbulence is maintained for long periods only by constant driving of velocity perturbations in Fourier space. When interpreting the above results, we should keep in mind that periodic box simulations represent a ``local'' patch of uniform background density that is embedded within a larger cloud, and are equivalent to studying an infinite uniform medium. By comparison, a 1.5-dimensional global model including vertical stratification \\citep{kud03,kud06} found that the decay of turbulence could be delayed, but only mildly, by some transfer of internal kinetic energy from small to large scale modes along the magnetic field direction. The rapid turbulence dissipation in all of these models is due to the presence of shocks and takes place under the assumption of magnetic flux freezing, without any contribution from magnetic field dissipation, e.g., by ambipolar diffusion.  The bottom line from the above studies is that all previous numerical modeling of MHD turbulence leads to rapid dissipation, in about a crossing time, with a logical conclusion that matching observations requires vigorous driving of turbulence from unspecified sources. The alternate possibility of maintaining some long-lived global modes is appealing but remained largely unexplored quantitatively.  In a recent paper, \\citet{bas09b} carried out an extensive parameter survey of fragmentation initiated by nonlinear turbulent flows, employing the magnetic thin-sheet approximation and also including the effect of ambipolar diffusion \\citep[see also][]{bas04,li04,nak05,cio06,bas09a}. In this approximation, the sheet interacts at its upper and lower surfaces with an external magnetic field, and can be considered a global model in the $z$-direction (parallel to the mean background magnetic field), although it is a local (periodic) model in the $x$- and $y$-directions. \\citet{bas09b} found that initial turbulent fluctuations decayed away quite rapidly in all models with supercritical mass-to-flux ratio, as well as for subcritical models that included the effect of ambipolar diffusion. However, a surprising result was that subcritical clouds evolving under flux-freezing were able to maintain a substantial portion of their initial kinetic energy to indefinitely large times.  In this paper, we analyze this unique instance of a turbulent MHD simulation that yields long-lived nonlinear motions. We perform a suite of numerical simulations to test its generality, and also establish an analytic explanation for this very interesting result. ",
        "conclusions": "\\label{disc}  The study of the decay of MHD turbulence has generally been based on the modeling of \\alf, slow MHD, and fast MHD modes in media that have a uniform background. The study of more complex global MHD (including magnetogravitational) modes for molecular clouds remains to be explored. In this paper, we have analyzed turbulent decay in a magnetically subcritical sheet-like cloud. It is an idealized model of a molecular cloud that is tied to a magnetic field anchored in the interstellar medium. A large fraction of the initial input kinetic energy is retained by the deformed magnetic field, and then persists in the cloud as large-scale oscillations. These represent linear waves of large extent which are nevertheless supersonic since the phase speed of the magnetic-tension-driven modes is up to $\\approx 10$ times the sound speed for typical cloud sizes.  Our model may approximate the situation of molecular clouds that are embedded in a low-density warm H {\\sc I} halo, or even the case of molecular cloud clumps that may be embedded in a matrix of H {\\sc I} gas \\citep[see][]{hen06}. Flux freezing is a good approximation for molecular cloud envelopes (as opposed to dense cores), due to significant photoionization by background starlight \\citep{mck89,cio95}. Observations also reveal that the low-column-density molecular cloud envelopes actually contain most of the cloud mass \\citep{kir06,gol08}. These envelopes may have a subcritical mass-to-flux ratio, as implied by their lack of star formation \\citep{kir06}, velocity data \\citep[e.g., in Taurus,][]{hey08}, and the subcritical state of the H {\\sc I} clouds \\citep{hei05} from which molecular clouds are presumably assembled.  Continuous driving of turbulent motions in molecular clouds is often invoked because the canonical numerical result of decay in a crossing time \\citep[e.g.,][]{sto98,mac99} is inconsistent with estimated cloud lifetimes that are at least a few crossing times \\citep{wil97}. \\citet{bas01} have argued that continuous driving of turbulence is consistent with observational constraints only if the driving occurs on the largest scale in the cloud, i.e., most of the energy is contained on that scale. Furthermore, continuous driving may not even be required if the decay time of the large-scale modes is greater than or equal to the estimated cloud lifetimes. Our models suggest that large-scale modes that are coupled to the external magnetic field can persist for very long times, thus reducing the need for continuous driving in order to explain observations. These modes preferentially span the largest scales in the model cloud, in agreement with analysis of observed cloud turbulence \\citep{bru09}. Future spectral line modeling of the large-scale cloud oscillations in our model cloud may make for useful comparison with observations, as has been done in a previous study of motions in the vicinity of dense cores \\citep{kir09}.  A counter-effect to the maintenance of large-scale modes is the loss of energy to the external medium. This can be accomplished by a coupling of the magnetic-tension-driven modes to MHD modes in the external medium. In a related example, \\citet{eng02} found significant energy loss to an external medium during core contraction using an approximate treatment of the transmission of transverse \\Alf waves through the bounding surfaces of a thin sheet. Some form of MHD wave coupling is certainly at work between any molecular cloud and its environment, although one may also argue that a clump embedded in a larger complex may reach a steady state in which it gains as much energy from its exterior as it loses. In any case, the study of the propagation of waves outside the cloud is outside the scope of our model. Future three-dimensional global MHD models of molecular clouds, which include the effect of an external medium, can address this point."
    },
    "1004/1004.4312_arXiv.txt": {
        "abstract": "{ A data and computation center for helioseismology has been set up at the Max Planck Institute for Solar System Research in Germany to prepare for the SDO mission. Here we present the system infrastructure and the scientific aims of this project, which is funded through grants from the German Aerospace Center and the European Research Council. } ",
        "introduction": "The Solar Dynamics Observatory (SDO), launched in \\linebreak February of this year, is the most important helioseismology mission of the coming decade, forming part of NASA's \\linebreak Living With a Star program. The main objective of SDO is to better understand solar variability and its inevitable \\linebreak impacts on the Earth and near-Earth environment. To \\linebreak achieve this goal, SDO carries a payload consisting \\linebreak of 3 scientific instruments: the Helioseismic and Magnetic Imager (HMI), the Atmospheric Imaging Assembly (AIA), and the Extreme ultraviolet Variability Experiment (EVE). SDO has a continuous downlink data rate of 130~Mbits per second.  Headquartered at Stanford University, the Joint Science Operations Center (JSOC) is in charge of collecting, processing, and archiving the HMI and AIA data. JSOC will make the data available to the scientific community (as well as basic science data products). The expected data volume of  $1.5$~Tbytes per day will make the analysis and processing of SDO data extremely challenging.  The German Science Center for SDO, hosted by the Max Planck Institute for Solar System Research (MPS), is a scientific IT infrastructure built around two projects: \\begin{itemize} \\item[--] The German Data Center for SDO (GDC-SDO), funded by the German Aerospace Center (DLR), will collect, manage, and store all the calibrated HMI data as well as selected AIA data sets. It will be a master European distribution center for HMI data. \\item[--] The Seismic Imaging of the Solar Interior (SISI) project, supported by a European Research Council (ERC) Starting Grant, will deliver specific helioseismic analyses of HMI data. \\end{itemize}  Here we present the scientific aims and the system infrastructure of the German Science Center for SDO. ",
        "conclusions": "SDO has launched and is currently undergoing a commissioning phase where, thus far, all mission operations have been completed successfully. The GDC-SDO is ready to receive HMI and AIA data in the coming months as they become available. An artificial HMI data series, manufactured and distributed by JSOC, was used to validate and verify the automatic transfer software (netDRMS), as well as to ensure that the GDC-SDO has sufficient network capabilities. For the short-term full cadence AIA data, LTO-4 tapes may be posted to Stanford and will remain on standby, ready for when the data comes online. The SDO workflow, and pipeline modules, are under development and are being tested with MDI data. Improvements in the methods of data analysis and their implementation will be required to extract the most scientific insight from the observations."
    },
    "1004/1004.3278_arXiv.txt": {
        "abstract": "{The image degradation produced by atmospheric turbulence and optical aberrations is usually alleviated using post-facto image reconstruction techniques, even when observing with adaptive optics systems.} {These techniques rely on the development of the wavefront using Zernike functions and the non-linear optimization of a certain metric. The resulting optimization procedure is computationally heavy. Our aim is to alleviate this computationally burden.} {To this aim, we generalize the recently developed extended Zernike-Nijboer theory to carry out the analytical integration of the Fresnel integral and present a natural basis set for the development of the point spread function in case the wavefront is described using Zernike functions.} {We present a linear expansion of the point spread function in terms of analytic functions which, additionally, takes defocusing into account in a natural way. This expansion is used to develop a very fast phase-diversity reconstruction technique which is demonstrated through some applications.} {This suggest that the linear expansion of the point spread function can be applied to accelerate other reconstruction techniques in use presently and based on blind deconvolution.} ",
        "introduction": "Atmospheric turbulence degrades astronomical images by introducing aberrations in the wavefront. During the last years, functional adaptive optics systems have been developed with the aim of partially cancelling these aberrations and generating a diffraction limited image of any astronomical object \\citep[e.g.,][]{beckers_ao93}. The development of such systems has been particularly challenging for solar telescopes because the reference object (the solar surface) has spatial structure. In spite of the great success of adaptive optics systems \\citep{rimmele_ao00,scharmer_ao00}, the limited number of degrees of freedom of the deformable mirror and the speed at which it has to work, especially in solar observations, impedes a full correction of the wavefront. Therefore, it has become customary in solar imaging to apply post-facto reconstruction techniques to the observed images to improve the wavefront correction. The advantage of these techniques is that the time limitation is much less restrictive since they do not need to work in real-time. Hence, one can use lots of computation power to reconstruct the images. One of the methods considered first for the reconstruction of solar images was based on speckle techniques \\citep{vonderluhe93,vonderluhe94} which produced images of extraordinary quality. However, it suffers from two fundamental problems. On the one hand, the number of images that one needs to acquire is large (although comparable with the number of images needed in complex blind deconvolution algorithms) because the missing information on the wave phases is statistically estimated from a succession of many short exposures. On the other hand, it relies on a theoretical description of the effect of the atmospheric turbulence and adaptive optics on the image \\citep{woger07}, which might be inaccurate. However, recent efforts have shown that almost real-time processing \\citep{denker01} and photometric precision \\citep{woger08} can be achieved routinely.  A complementary approach to speckle reconstruction appeared with the systematic application of techniques based on phase diversity \\citep[e.g.,][]{lofdahl94,lofdahl98}, developed some years before by \\cite{gonsalves79} and later extended by \\cite{paxman92}. These techniques rely on a maximum likelihood simultaneous estimation of the wavefront and the original image (before being affected by  atmospheric+telescope aberrations) using a reference image and a second one that contains exactly the same aberrations as the reference one plus a known artificially-induced aberration. Because of the mechanical simplicity, a defocus typically constitutes the added aberration. The main drawback of the phase diversity technique is the heavy computational effort needed to compute the maximum likelihood solution. Many Fourier transforms have to be used to estimate the point spread function (PSF) from the wavefront at the pupil plane.  Even more computationally expensive is the blind deconvolution technique developed by \\cite{vannoort05}. This technique, termed multi-object multi-frame blind deconvolution (MOMFBD), is able to reconstruct estimates of the object and the wavefront from the observation of different objects (same location in the Sun observed at different wavelengths) with some constraints, which are essentially used to augment the amount of information introduced in the likelihood function. As a consequence, the estimation of the wavefront and the original image is less dominated by noise and hence more stable. In close analogy with phase diversity, this enhancement in stability comes at the price of a higher computational burden since many Fourier transforms have to be computed for the maximization of the likelihood. In conclusion, all these techniques suffer from time-consuming computations. As modern solar telescopes increase their flux of high quality data, observers are condemned to struggle with the bottleneck of reconstruction algorithms.  During the 1940's, \\cite{nijboer42} made it possible to estimate the PSF of an aberrated optical device in terms of the generalized complex pupil function. The resulting integral expressions were very involved and remained impossible to evaluate until recent years. What is presently known as the extended Nijboer-Zernike theory \\citep[ENZ;][]{janssen_enz02,braat_enz02} presents analytical expressions for the complex field at the focal volume that depend on the aberration through the coefficients of the Zernike expansion of the wavefront. The expressions are closed in the sense that they only depend on computable functions and the coefficients of the Zernike expansion.  Our purpose in this paper is to extend the ENZ theory to arbitrary aberrations exploiting mathematical properties of the Zernike functions recently found. As a consequence, we find a general expression for the PSF in the focal volume as a linear combination of given basis functions. Starting from this expression, we propose to use it in a phase diversity reconstruction algorithm. Due to the analytical character of the PSF, we are able to avoid the calculation of Fourier transforms during the iterative process, thus inducing an important gain in computational time. This shall be our fundamental result: that through this formalism we can drastically diminish the computational burden of reconstruction algorithms and therefore the time needed to perform the reconstruction. We illustrate it using phase diversity. But in principle any reconstruction algorithm can benefit from our formalism, since we only change the way in which the PSF is described.  The outline of the paper is the following. Section \\ref{sec:wavefront_expansion} gives a description of the wavefront expansion in Zernike functions and how the field distribution can be analytically obtained for general aberrations. The expansion of the point spread function is presented in \\S\\ref{sec:psf} with an statistical analysis of that expansion for atmospheric Kolmogorov turbulence in \\S\\ref{sec:statistical}. Section \\ref{sec:phase_diversity} presents the application of the formalism to the fast restoration of images using phase-diversity. Finally, the conclusions are presented in \\S\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  The use of the recently developed Extended Nijboer-Zernike theory together with some mathematical results on the multiplication of Zernike polynomials has allowed us to rewrite the image formation integrals with an aberrated wavefront in terms of linear combinations of analytic functions. Generalizing the usual ENZ theory, we are able to do so   \\textit{a priori} independently of the amplitude of the aberrations. With such mathematical tool we are able to rewrite the techniques of post-facto image reconstruction taking advantage primarily of the fact that the Fourier transforms of those analytic functions can be precomputed once and for all. Building different PSFs in the minimization process at the core of these reconstruction algorithms does not require to recompute any other Fourier transform but just modify the scalar, although complex, coefficients multiplying the functions.  \\begin{figure}[!t] \\centering \\includegraphics[width=1.4\\columnwidth]{themis_reconstructed.eps} \\caption{Example of the phase-diversity reconstruction at 850 nm. We represent the focused image (left) and the reconstructed sub-image indicated with a box (right). Images have been acquired with THEMIS.} \\label{fig:pd_example_themis} \\end{figure}  The gain in speed is enormous. We illustrate it through two observations readied for phase diversity. Image reconstruction with quality identical to the standard algorithms is performed in question of seconds per patch, and one to five minutes for wide-field images, in standard desktop computers. If image reconstruction is a must in present and future solar observations, the bottleneck of computationally expensive and time consuming algorithms was almost a showstopper for future instruments. The present improvement through the use of analytical PSFs solves the core of that problem and suggests that image reconstruction can even be implemented as a routine on-line procedure on the data pipelines of the telescope instruments themselves.  Since we have only modified the description of the PSF with respect to previous approaches, all common reconstruction schemes (with any desired complexity), and not just phase diversity, can be extended to use our formalism. Among them, we find methods like multi-image phase diversity \\citep[e.g.,][]{paxman92} or multi-object multi-frame blind deconvolution \\citep{vannoort05} that use many images to estimate information about frequencies that has been destroyed by the presence of noise. In essence, all these schemes require writing an error metric (equivalently, a likelihood function) like Eq. (\\ref{eq:metric}) that takes into account the presence of the additional information. The larger amount of information helps regularize the problem and reduces the influence of noise. All of them are penalized by the computing time involved and all of them can potentially take profit of the analytical approach we present to reduce the computational burden. The final goal is not to slow down the flows of images from instruments because of the lack of capability to handle the image reconstruction problem.  Beyond those known techniques, we also want to point out that it is possible to introduce regularization by using a fully Bayesian approach in which prior information is introduced in the problem, eliminating a priori the noise nuisance in the minimization of the metric."
    },
    "1004/1004.4344_arXiv.txt": {
        "abstract": "We consider large-scale collective motion of flat edge-on spiral galaxies from the Revised Flat Galaxy Catalogue (RFGC) taking into account the curvature of space-time in the Local Universe at the scale $100h^{-1}\\,\\rmn{Mpc}$. We analyse how the relativistic model of collective motion should be modified to provide the best possible values of parameters, the effects that impact these parameters and ways to mitigate them. Evolution of galactic diameters, selection effects, and difference between isophotal and angular diameter distances are inadequate to explain this impact. At the same time, measurement error in \\mbox{H\\,{\\sc i}} line widths and angular diameters can easily provide such an impact. This is illustrated in a toy model, which allows analytical consideration, and then in the full model using Monte Carlo simulations. The resulting velocity field is very close to that provided by the non-relativistic model of motion. The obtained bulk flow velocity is consistent with $\\Lambda$CDM cosmology. ",
        "introduction": "At present time the Universe is essentially inhomogeneous on the scales of about 10--100 Mpc. The development of initial fluctuations led to an observable large-scale structure. The regions with increased matter density provide an additional attraction of surrounding galaxies. The regions with decreased density, e.g.  voids, also make an input to the collective large-scale motion of galaxies on the background of Hubble expansion. Investigation of such motion on one side allows to map the matter density, including dark matter, in the Local Universe, and on the other side its parameters are linked with cosmological parameters.  All of this makes the study of collective galaxy motions important.  In recent years a number of articles was published claiming that typical velocities of large-scale collective motions are inconsistent with the standard $\\Lambda$CDM model. For example, \\citet{ref:WFH09} obtained the value $407 \\pm 81\\,\\rmn{km\\,s}^{-1}$ at the scale $100h^{-1}\\,\\rmn{Mpc}$, whereas the $\\Lambda$CDM model gives about $250\\,\\rmn{km\\,s}^{-1}$. However, our estimation of $210 \\pm 86\\,\\rmn{km\\,s}^{-1}$ at the same scale, obtained in the article \\citep{ref:APSS09}, is consistent with the $\\Lambda$CDM predictions. Additionally, in the same article we obtained from the peculiar velocities the constraints on the cosmological parameters $\\Omega_m$ and $\\sigma_8$ and their combinations, which match the other more precise constraints like baryonic acoustic oscillations or WMAP observations.  In the article \\citep{ref:APSS09} we used a sample of RFGC galaxies with measured redshifts and \\mbox{H\\,{\\sc i}} line widths. The Revised Flat Galaxy Catalogue (RFGC) \\citep{ref:RFGC} and its previous version Flat Galaxy Catalogue (FGC) \\citep{ref:FGC} contain the information about Right Ascension and Declination for the epochs J2000.0 and B1950.0, galactic longitude and latitude, major and minor blue and red diameters in arcminutes in the POSS-I diameter system, morphological type of the spiral galaxies according to the Hubble classification, index of the mean surface brightness and some other parameters, which are not used in this article. The RFGC contains data about 4236 flat edge-on spiral galaxies, almost uniformly covering the celestial sphere and satisfying the conditions $a_b/b_b\\ge7$ and $a_b>0\\farcm 6$. Here $a_b$ and $b_b$ are the major and minor axial diameters in the $a_{25}$ system. The original goal of this catalogue was to estimate the distance to galaxies according to the Tully-Fisher relation in the ``\\mbox{H\\,{\\sc i}} line width -- linear diameter'' version without using their redshifts. The data about the redshifts and \\mbox{H\\,{\\sc i}} line widths or gas rotation velocities $V_{rot}$ were taken from different sources. There were a number of gradually increasing samples of galaxies with such data \\citep{ref:K00,ref:Par01,ref:ParTug04}. The latest version of this sample including 1623 galaxies was compiled and described by \\citet{ref:APSS09}. A list of peculiar velocities based upon this list in the non-relativistic model of motion was assembled by \\citet{ref:arxiv09}.  In this article we use the same sample, but with different model of collective motion of galaxies \\citep{ref:KudAlex02,ref:KudAlex04}, based upon the general theory of relativity (GTR). This model was applied earlier to the previous version of the sample by \\citet{ref:ParGayd05}; however, the present article offers a much more in-depth analysis. ",
        "conclusions": "We applied the relativistic model of motion supplied with the generalised Tully-Fisher relation (\\ref{eqn:TFR}) to the sample of 1623 flat edge-on spiral galaxies from the RFGC catalogue. The analysis of results prompted us to switch first to the semirelativistic model, and then to the semirelativistic model with fixed $\\gamma$. The parameters of the collective motion obtained in the framework of this model appeared to be close to that obtained in the non-relativistic case. We analysed certain reasons behind the decrease of $\\gamma$ in the semirelativistic model. Evolution of galactic diameters, selection effects, and difference between isophotal and angular diameter distances appeared to be inadequate to explain this effect. At the same time, measurement error in \\mbox{H\\,{\\sc i}} line widths and angular diameters can easily provide such a decrease. This was illustrated in a toy model, which allows analytical consideration, and then in the full model using Monte Carlo simulations. The obtained bulk flow velocity is consistent with $\\Lambda$CDM cosmology."
    },
    "1004/1004.4805_arXiv.txt": {
        "abstract": "We study evolution of isolated neutron stars on long time scale and calculate distribution of these sources in the main evolutionary stages: Ejector, Propeller, Accretor, and Georotator. We compare different initial magnetic field distributions taking into account a possibility of magnetic field decay, and include in our calculations the stage of subsonic Propeller.  It is shown that though the subsonic propeller stage can be relatively long, initially highly magnetized neutron stars ($B_0\\ga 10^{13}$~G) reach the accretion regime within the Galactic lifetime if their kick velocities are not too large. The fact that in previous studies made $>$10 years ago, such objects were not considered results in a slight increase of the Accretor fraction in comparison with earlier conclusions. Most of the neutron stars similar to the Magnificent seven are expected to become accreting from the interstellar medium after few billion years of their evolution. They are the main predecestors of accreting isolated neutron stars. ",
        "introduction": "Accreting isolated neutron stars (AINS) were predicted  40 years ago by \\cite{s71} % and independently by \\cite{ors70}. % In early 90s there was some enthusiasm due to the launch of the ROSAT satellite, which was expected to find many sources of this kind \\citep{tc91}. Several populational studies have been made \\citep{blaes90,br91,bm93,mb94,blaes95,manning96}. However, it came out that AINS, if they exist, are very elusive (Colpi et al. 1998). The main reason is that initial (kick) velocities of NSs appeared to be significantly larger, than it have been thought before \\citep{ll94}. % Initial guess that the number of Accretors is small due to low luminosity of high-velocity NSs was shown to be wrong. In a detailed study by \\cite{census2000} % (hereafter Paper I) it was shown that INS with Crab-like initial parameters and constant magnetic fields spend all their lives as Ejectors (we follow the classification summarized in \\citealt{l92}% ) if their initial velocities are $\\ga 100$~km~s$^{-1}$.  Then, the fraction of Accretors was mainly determined by the fraction of low-velocity NSs.  Up to the very end of 90s, it was believed that the wast majority of NSs are born similar to the Crab pulsar. I.e., that they have short initial spin periods (from milliseconds to few tens of millisecond) and magnetic fields $B\\sim 10^{12}$~G. Now it is believed, that about one half of NSs have different initial properties \\citep{ptp2006,keane2008}. There are at least three groups of sources with distinct parameters: compact central objects (CCOs) in supernova remnants (SNR), magnetars (anomalous X-ray pulsars - AXPs, and soft gamma-ray repeaters - SGRs), and cooling radioquiet NSs dubbed the Magnificent seven (M7) \\citep{zoo08}. CCOs have low initial fields $\\sim 10^{11}$~G \\citep{h07,gh2009} and relatively long spin periods (hundreds of millisecond). AXPs and SGRs have large fields $\\sim 10^{14}$~G (see a review in \\citealt{Mereghetti2008}). The Magnificent seven-like NSs have fields slightly above $10^{13}$~G \\citep{haberl2007, kaplan2008}. Probably, some of rotating radio transients (RRATs, \\citealt{rrats2006}) are similar to the M7. This variety in initial properties deserves new studies of evolving NSs using the population synthesis technique (see a review in \\citealt{pp2007}). In this paper we present the first step.  We describe two models. At first, we discuss a simple semianalytical approach, which is used to illustrate the main features of the scenario.  In this model velocities and ambient densities are not changing. Then we present a detailed numerical model, which takes into account spatial movements of NSs in the Galactic potential and realistic 3D distribution of the interstellar medium (ISM). Our main results are based on this model.  In the next section we present basic concepts used in both models, and describe each of them. Then, in Sec. 3, we present results. Discussion is given in Sec.4. In the last section we present our conclusions. ",
        "conclusions": "After the first of the M7  have been discovered \\citep{walter96}, several authors proposed and discussed that they can be AINSs \\citep{walter96, kp97, nt99}. Though, it appeared that it is not so. The M7 are young NSs with relatively large fields. Probably, they are related to evolved magnetars \\citep{popov10}. % Here we demonstrate that in future the M7 and similar sources are expected to become AINS if their magnetic fields do not decay significantly. Even a relatively long stage of subsonic Propeller \\citep{ikhsanov2001} cannot prevent accretion. This is a good news for observers. Probably, telescopes like eROSITA aboard Spektr-RG will be able to detect AINS, soon. However, the question of the accretion efficiency is still on the list \\citep{toropina03}.  The distribution over evolutionary stages strongly depends on kick velocity distribution, initial magnetic field distribution and field evolution. Because of that precise predictions are not possible now. This shows how important is to detect old isolated NSs as Accretors (or, less probable, other stages) to learn more about initial properties and evotuion of INSs."
    },
    "1004/1004.2081_arXiv.txt": {
        "abstract": "GRB afterglows offer a probe of the intergalactic medium out to high redshift which complements observations along more abundant quasar lines-of-sight. Although both quasars and GRB afterglows should provide a-priori random sight-lines through the intervening IGM, it has been observed that strong Mg-II absorbers are twice as likely to be found along sight-lines toward GRBs. Several proposals to reconcile this discrepancy have been put forward, but none has been found sufficient to explain the magnitude of the effect. In this paper we estimate the effect of gravitational lensing by galaxies and their surrounding mass distributions on the statistics of Mg-II absorption. We find that the multi-band magnification bias could be very strong in the spectroscopic GRB afterglow population and that gravitational lensing can explain the discrepancy in density of absorbers, for plausibly steep luminosity functions. The model makes the prediction that approximately 20\\%-60\\% of the spectroscopic afterglow sample  (i.e. $\\sim5-15$ of 26 sources) would have been multiply imaged, and hence result in repeating bursts. We show that despite this large lensing fraction it is likely that none would yet have been identified by chance owing to the finite sky coverage of GRB searches. We predict that continued optical monitoring of the bright GRB afterglow locations in the months and years following the initial decay would lead to identification of lensed GRB afterglows. A confirmation of the lensing hypothesis would allow us to constrain the GRB luminosity function down to otherwise inaccessibly faint levels, with potential consequences for GRB models. ",
        "introduction": "At moderate redshifts ($0.3\\la z\\la2.2$) the dispersal of metal enriched gas through the IGM via galactic winds is most readily traced via Mg-II absorption systems in optical spectra of bright  background sources. Although the detailed origin of Mg-II absorbers remains uncertain it is thought that they are associated with galaxies and galactic outflows. For example, Mg-II absorbers have been shown to be associated with neutral hydrogen absorbers over a range of column densities, including damped Lyman-$\\alpha$ absorbers \\citep[e.g.][]{rao2006}.  Moreover the host halo masses associated with MgII absorbers have been estimated at $z\\sim0.5$ via cross-correlation with luminous red galaxies in the Sloan Digital Sky Survey Data Release 3 \\citep[][]{bouche2006}, yielding a host mass of $M\\sim10^{12}M_\\odot$ (i.e. massive galaxies). Indeed the large number of MgII absorbers and galaxies available for cross-correlation yields a statistical accuracy of a factor of two in host halo mass.  Obtaining an unbiased census of the density and distribution of Mg-II in the IGM requires that the back-ground sources be uncorrelated with the foreground absorbers under study. Most current knowledge regarding the census of Mg-II absorption systems comes from spectra of quasars \\citep[e.g.][]{prochter2006b}. More recently, observations of GRB afterglows have begun to offer a new probe of the intergalactic medium out to high redshift which complements the more abundant quasar lines-of-sight. Indeed, since they are associated with star formation rather than supermassive black holes (which, given their long assembly times, become extremely rare at high redshift), GRBs could potentially be seen out to much higher redshift than quasars; the current record holder is at $z\\sim8.1$ \\citep{salvaterra09}. Interestingly, although both quasars and GRB afterglows should provide a-priori random sight-lines through the intervening IGM it has been observed that strong Mg-II absorbers are several times as likely to be found along sight-lines to GRBs as along quasar sightlines \\citep[][]{prochter2006,tejos2009,vergani2009}, indicating that one or both of these samples are biased relative to the underlying population of absorbers. Several proposals to reconcile this discrepancy have been put forward, with a detailed discussion of possible biases presented by \\citet[][]{porciani2007}. For example, the incidence of  Mg-II systems in quasars could be lowered because of dust obscuration associated with the absorbing systems; this turns out to be too small to explain the discrepancy. Alternatively, it has been argued that different sizes of the source could lead to different absorber incidence between GRB afterglows and quasars \\citep[][]{frank2007}. However, from the similarity of the equivalent width distributions in GRBs and quasars, \\citet[][]{porciani2007} show that the absorbers must be larger than the beamsize, and hence it is not possible to explain the difference in this way.  A third potential bias is provided through gravitational lensing of GRBs by foreground galaxies associated with the Mg-II absorber. There are several lines of circumstantial evidence for this. Firstly, imaging studies of the host galaxies of GRBs with early time afterglow spectra \\citep[][]{chen2009} show that additional galaxies are found within $2''$ of all GRB host fields where the line-of-sight contains a Mg-II absorber, while no additional galaxies are found within $2''$ of GRB lines-of-sight that do not have any Mg-II absorbers. Moreover, GRB afterglows that have more than one absorber are found to be a factor of 1.7 brighter than the others \\citep[][]{porciani2007}. Under the assumption that gravitational lensing magnification of GRBs is responsible for the excess number of absorbers in GRB spectra relative to quasars, \\citet[][]{tejos2009} estimated the fraction of GRB afterglows that  need to have been included in the sample owing to magnification. They find a value of 60\\% which indicates a large gravitational lensing induced magnification bias. On the other hand, no strongly lensed GRBs have been discovered.\\footnote{The expected rate of strongly lensed GRBs was discussed by \\citet[][]{porciani2001}, who found that although a few strongly lensed GRBs should be present in a 3 year survey with $Swift$, the partial sky coverage means that it is unlikely that lensed pairs would have been identified as recurrent bursts.} At first sight this suggests that gravitational lensing could not provide the required explanation for the excess absorbers in GRB afterglow spectra. However, as noted by \\citet[][]{porciani2007} the selection of GRB afterglows in two (apparently) independent bands (namely gamma-rays and the optical) leads to the possibility of an increased multi-band magnification bias \\citep[][]{wyithe2003}. One goal of this paper is to determine under which circumstances a sufficiently large magnification bias can be obtained.   The most recent collection of GRB afterglow spectra from which strong ($>1$\\AA) Mg-II absorption can be studied is summarised in the work of \\citet[][]{vergani2009} who compiled a list of 26 spectra (with a combined redshift path-length of $\\Delta z=31.55$) containing 22 strong Mg-II absorbers among 15 of the 26 lines-of-sight.    From that paper we take the following values for observables describing the absorber population. The fraction of lines-of-sight that contain one or more absorbers is $F^{\\rm obs}_{\\rm Mg}\\sim0.6\\pm0.15$. By contrast, the results of high resolution spectroscopy from \\citet[]{prochter2006b} yielded 22 absorbers along 91 quasar lines-of-sight, implying that $F_{\\rm Mg,q}=0.25$. Assuming the same redshift pathlength distribution as the GRB sample, these values imply that the relative incidence of the number of strong systems in GRB afterglows relative to quasars is $R^{\\rm obs}_{\\rm GRB,q}\\sim2.1\\pm0.6$. Where required we take the redshifts of GRBs and associated absorbers, as well as the probed redshift path-lengths from this paper. Our goal is to make a quantitative comparison of a simple parameterised model for the expected number of absorbers (which includes the bias introduced by gravitational lensing), with these observations.  Our paper is presented as follows. Section~\\ref{model} describes the basics of the lensing model with which we interpret the absorber statistics. In \\S~\\ref{statistics} and \\S~\\ref{comparison} we discuss the statistics of Mg-II absorbers in GRB afterglow spectra, and their comparison with statistics from quasar lines-of-sight within the context of our model parameters. We then discuss the predicted distribution of absorber redshifts (\\S~\\ref{absorber})  and observed source magnifications (\\S~\\ref{magnification}). In \\S~\\ref{strong} we discuss the rate of multiple imaging among the spectroscopic GRB afterglow sample before presenting our summary in \\S~\\ref{summary}. In our numerical examples, we adopt the standard set of cosmological parameters \\citep[][]{komatsu2009}, with values of $\\Omega_{\\rm b}=0.04$, $\\Omega_{\\rm m}=0.24$ and $\\Omega_Q=0.76$ for the matter, baryon, and dark energy fractional density respectively, and $h=0.73$, for the dimensionless Hubble constant. ",
        "conclusions": "\\label{summary}  In recent years the rapid optical imaging and spectroscopic followup of GRBs and their afterglows has began to offer a new probe of the intergalactic medium out to high redshift which complements the more abundant quasar lines-of-sight. Studies of the incidence of Mg-II absorption systems along these lines-of-sight have revealed that although both quasars and GRB afterglows should provide a-priori random sight-lines through the intervening IGM, strong Mg-II absorbers are several times as likely to be found along sight-lines to GRBs \\citep[][]{prochter2006,tejos2009,vergani2009}. Several proposals to reconcile this discrepancy have been put forward \\citep[see e.g.][]{porciani2007}, but none have been quantitatively successful. In this paper we have described a simple model that associates Mg-II absorbers with the foreground galaxy population, and includes a 2-band luminosity function to describe the number densities of GRBs and their afterglows. We have used this model to estimate the effect of gravitational lensing by galaxies and their surrounding mass distributions on the statistics of Mg-II absorption.  Our model leads to two main findings. Firstly, we show that the multi-band magnification bias could be very strong in the GRB afterglow population. Gravitational lensing can explain the discrepancy in the incidence of absorbers between quasar and GRB lines-of-sight for GRB afterglow luminosity functions with cumulative slopes $\\alpha_{\\rm A}\\ga3.5$. Secondly our model makes the prediction that approximately 20\\%-60\\% (i.e. between $\\sim5$ and 15) of the existing afterglow sample would have been multiply imaged, and hence repeating sources. We show that despite this large lensing fraction it is likely that none would yet have been identified by chance owing to the finite sky coverage of GRB searches. As a result we predict that continued optical monitoring of the bright GRB afterglow locations in the months and years following the initial decay would lead to identification of lensed GRB afterglows. A confirmation of the lensing hypothesis would allow us to constrain the GRB luminosity function down to otherwise inaccessibly faint levels, with potential consequences for GRB models.   {\\bf Acknowledgments} The research was supported by the Australian Research Council (JSBW). SPO was supported by NSF grant AST 0908480. JSBW and BP thank the Physics department at UCSB for hospitality during this work. SPO thanks the UCSB Wednesday 'Gastrophysics' group for lively discussions.  \\newcommand{\\noopsort}[1]{}"
    },
    "1004/1004.2048_arXiv.txt": {
        "abstract": "The statistics of strongly lensed arcs in samples of galaxy clusters provide information on cluster structure that is complementary to that from individual clusters. However, samples of clusters that have been analyzed to date have been either small, heterogeneous, or observed with limited angular resolution. We measure the lensed-arc statistics of 97 clusters imaged at high angular resolution with the Hubble Space Telescope, identifying lensed arcs using two automated arc detection algorithms. The sample includes similar numbers of X-ray selected (MACS) and optically selected (RCS) clusters, and spans cluster redshifts in the range $0.2 < z < 1$. We compile a catalogue of $42$ arcs in the X-ray selected subsample and $7$ arcs in the optical subsample. All but five of these arcs are reported here for the first time. At $0.3 \\leq z \\leq 0.7$, the X-ray selected clusters have a significantly higher mean frequency of arcs, $1.2\\pm 0.2$ per cluster, versus $0.2\\pm 0.1$ in the optical sample. The strikingly different lensing efficiencies indicate that X-ray clusters trace much larger mass concentrations, despite the similar optical luminosities of the X-ray and optical clusters. The mass difference is supported  also by the lower space density of the X-ray clusters, and by the small Einstein radii of the few arcs in the optical sample. Higher-order effects, such as differences in concentration or substructure, may also contribute. ",
        "introduction": "Galaxy clusters are natural laboratories for studying a variety of astrophysical processes and for testing cosmological models. In particular, the masses and mass profiles of clusters have proved to be useful for constraining cosmological parameters (e.g. Bridle et al. 1999; Reiprich \\& B{\\\" o}hringer 2002; Voit 2005; Allen et al. 2008; Vikhlinin et al. 2009). Gravitational lensing is frequently used to map the evolution of cluster mass profiles, ellipticities, and substructure. One approach is to perform detailed modeling of individual clusters using strong and weak lensing (e.g., Abdelsalam et al. 1998; Broadhurst et al.  2005; Leonard et al.  2007; Limousin et al. 2007; Richard et al. 2007). However, since this kind of approach requires deep data for individual clusters that exhibit numerous lensed images, the results may not be representative of the vast majority of clusters. A complementary approach is to measure the statistics of lensed arcs in large samples of clusters.  Lensing statistics thus provide another means to study clusters as a population.  For the past decade there has been debate concerning theoretical lensing statistics predictions and their confrontation with observations. Bartelmann et al. (1998; B98) performed lensing simulations using artificial sources at redshift $z=1$ by ray tracing through the five most massive clusters formed in a cosmological N-body dark matter simulation (Kauffmann et al. 1999). The observed number of giant arcs, with length-to-width ratio $l/w \\geq 10$ and $R<21.5$ mag, present over the whole sky was estimated by extrapolating from observations of a subsample of X-ray selected clusters from the Einstein Extended Medium Sensitivity Survey (EMSS), and compared to the theoretical calculation. B98 found that the estimated number of observed arcs is larger by almost an order of magnitude than the number predicted by the now-standard $\\Lambda {\\rm CDM}$ model. Later estimates of lensed arcs statistics in clusters from both the Las Campanas Distant Cluster Survey (Zaritsky \\& Gonzales 2003; arcs with $l/w \\geq 10$ and $R<21.5$ mag) and the Red-Sequence Cluster Survey (RCS; Gladders et al. 2003) confirmed the estimates of the observed number of arcs derived by B98. Most recently, Hennawi et al. (2008) analyzed a sample of 240 clusters, optically selected from the Sloan Digital Sky Survey (SDSS), and found that $10\\% - 20\\%$ of them are strong lenses, similar to the findings of Gladders et al. (2003). The largest catalogue of arcs to date was compiled by Sand et al. (2005) who found $104$ arcs in $128$ clusters.  However, their systematic search for arcs was performed on a largely heterogeneous cluster sample.  The apparent overproduction of arcs by real clusters has stimulated further theoretical studies of arc statistics. Meneghetti et al. (2000) studied numerically the effect of the masses of the individual cluster galaxies on a cluster's lensing cross section, and found it to be negligible, as also found in a study by Flores, Maller, \\& Primack (2000). However, the increase in lensing cross section due to the central cluster cD galaxy may be as high as $\\sim 50\\%$ (Meneghetti, Bartelmann, \\& Moscardini 2003) and the increase in cross section due to the intra-cluster gas could perhaps be by a factor of a few (Puchwein et al.  2005, Rozo et al.  2008). Oguri, Lee, \\& Suto (2003) argued that halo triaxiality could also play an important role in increasing cluster lensing cross sections.  Torri et al. (2004) raised the possibility that X-ray selection of clusters may favor merging systems, which may be more efficient lenses. Wambsganss, Bode, \\& Ostriker (2004) pointed out that since lensing cross section is a steep function of source redshift, the conflict between theory and observations could be the result of the assumed source redshifts in the simulations. Similarly, Dalal, Holder, \\& Hennawi (2004) performed a lensing simulation using artificial background sources at different redshifts and a large sample of simulated clusters.  They found that their prediction for the number of lensed arcs was consistent with an observed number that they derived from a sample of X-ray selected EMSS clusters. The difference between this result and that of B98 was explained by the combination of three main effects: the inclusion of sources at different redshifts; the use of a higher source density in the Dalal et al.  simulation; and an observed cluster number density lower than the one used by B98 for estimating the all-sky number of arcs.  A more observationally oriented approach to lensing statistics simulations was introduced by Horesh et al. (2005; H05) in order to test specifically the lensing efficiency of individual clusters, independent of the separate question of the number density of clusters. H05 repeated the B98 simulations using the same simulated clusters, but using background sources from the Hubble Deep Field (HDF), each at a redshift based on its actual photometric redshift. Observational effects including background, photon noise, and the light of cluster galaxies were added to the simulated lensed images. A mass-matched sample of $10$ X-ray-selected clusters (Smith et al. 2005) observed at high angular resolution with the Hubble Space Telescope (HST) was used for comparison with the simulated sample. Finally, an automated objective arc-detection algorithm was applied to both the observed and the simulated samples. This procedure permitted measuring and comparing the frequency of arcs over a larger range in magnitudes (down to $R \\leq 24$ mag). H05 found that the lensing efficiency of their simulated clusters at $z\\approx 0.2$ was consistent, to within Poisson errors, with that of their observed sample. While the analysis suggested that the observed clusters could be somewhat more efficient lenses by up to a factor of two, this conclusion was limited by the small size of both the observed and the simulated samples, as well as the parameters assumed in the simulations.  Indeed, an important parameter that affects all theoretical studies of arc statistics is $\\sigma_{8}$, the overdensity within an $8~{\\rm Mpc}$ radius comoving sphere. Past simulations have used diverse values: 0.9 (B98; Dalal et al. 2004; H05) or 0.95 (Wambsganss et al. 2004; Hennawi et al.  2007). Fedeli et al. (2008) have recently analyzed the effect of $\\sigma_{8}$ on the arc statistics question, and pointed out that the most recent values of $\\sigma_{8}$ from WMAP5 ($0.796\\pm 0.036$; Dunkley et al. 2009) revive and reinforce the discrepancy between theory and observations of arc statistics.  A possibly related debate has emerged recently on the subject of the size of the Einstein radius in clusters. Broadhurst and Barkana (2008) calculated the distribution of Einstein radii in clusters with a spherical Navarro, Frenk, \\& White (NFW; 1996) profile, and with a concentration distribution according to Neto et al. (2007). They compared their prediction with the observed Einstein radii of three clusters, among them Abell 1689, and found that the observed radii are significantly larger than the theoretical expectation. Yet another cluster with a large Einstein radius was recently reported by Zitrin et al. (2009). Sadeh \\& Rephaeli (2008) have calculated the concentration distribution of clusters based on the distribution of cluster formation times. They too find a discrepancy, albeit weak, between the observed Einstein radius of Abell 1689 and its expected value. Oguri \\& Blandford (2009), however, argue that the Einstein radius they obtain using a generalized triaxial form of the NFW profile (Jing \\& Suto 2002) is consistent with that observed in Abell 1689. In addition, they provide a prediction for the distribution of Einstein radii, which can be tested with a large statistical cluster sample.  Clearly, resolution of these problems requires, on the theoretical side , improved simulations, incorporating the most realistic cosmological parameters, source parameters, and observational effects; and from the observational perspective, large, well-understood samples of clusters at various redshifts, selected by diverse methods and uniformly observed at the high depth and resolution needed for the clear detection of large arcs.  In this paper, we address this observational perspective. We explore the observed statistical properties of 97 galaxy clusters imaged with HST. This cluster sample is large enough to be separated into several subsamples based on redshift and selection type. We apply two different arc detection algorithms to the clusters, and compile a high-resolution arc catalogue. We then study the arc statistics in the various subsamples. In a forthcoming publication, we will compare the observed statistics of this sample to new, improved, calculations of matched simulated samples. Throughout this paper we adopt a $\\Lambda {\\rm CDM}$ cosmology with parameters $\\Omega_{\\rm m}=0.3$, $\\Omega_{\\rm \\Lambda}=0.7$, and $H_{0}=70~{\\rm km~s}^{-1}~{\\rm Mpc}^{-1}$. Magnitudes are in the Vega system. ",
        "conclusions": "The results of our arc survey, presented above, can serve as a new and improved observational basis for future arc statistic studies. However, our survey also shows clearly that the arc-production efficiency of X-ray-selected clusters such as MACS and XBACS is higher by a factor of $5-10$ than that of RCS clusters. In this section, we carry out additional analysis and discussion of the meaning of this result.  At a given redshift, the cross section for lensed arc formation depends primarily on mass, although mass profile, ellipticity and substructure are also important. The mass dependence weakens towards the high mass end at ${\\rm M}_{\\rm 200}\\sim 10^{15} {\\rm M}_\\odot$ (Dalal et al. 2004; Hennawi et al. 2007). The stark difference in the arc frequency between the X-ray selected and optically selected clusters immediately raises the possibility that they probe different mass ranges. Based on their X-ray luminosities, the X-ray selected clusters have masses of ${\\rm M}_{200} > 10^{15} {\\rm M}_{\\odot}$. Unfortunately, there is scant information of the X-ray properties of the RCS clusters, and hence on their masses. For example, Hicks et al. (2008) recently observed with {\\it Chandra} a sample of 13 RCS clusters, of which detailed analysis was possible for nine. They found significant differences in the mass-temperature-luminosity relations of X-ray selected and RCS clusters, X-ray underluminosity in some RCS clusters, and evidence that RCS clusters have a larger fraction of their baryons in stars. Nevertheless, since optical flux is one of the few observables we do have available for the RCS clusters, we begin by comparing the optical luminosities of the MACS and RCS subsamples. The HST/ACS field of view covers only the central core regions of the clusters, and therefore we examine several proxies for the optical luminosity.  As a first proxy for optical luminosity, we examine the luminosities of the brightest cluster galaxies (BCGs) of the RCS and MACS samples. In SDSS clusters, Hansen et al. (2005) have found a correlation between cluster mass and BCG luminosity. The BCG magnitudes were measured using SExtractor by including the light from pixels which belong to the BCG and are above the detection threshold. Since the low-redshift subsamples are observed through different filters, we first calculate the F606W$-$F814W color for each RCS cluster redshift using an elliptical galaxy spectral template from Kinney et al.  (1996), and convert the RCS cluster BCG F814W magnitudes to F606W. In this comparison we exclude the following clusters (four MACS and five RCS) due to the uncertainty in determining their centres and in identifying the dominant BCGs: MACSJ0916.1$-$0023, MACSJ1354.6$+$7715, MACSJ2243.3$-$0935, MACSJ0257.1$-$2325, RCS131912$-$0206.9, RCS022403$-$0227.7, RCS110104$-$0351.3, RCS033414$-$2824.6, and RCS110814$-$0430.8.  We find that the BCG magnitudes are more uniformly distributed in the RCS subsample than in the MACS subsample, and the BCGs span a wider magnitude range. Nevertheless, in the low-redshift MACS and RCS subsamples, the median BCG absolute magnitudes are, ${\\rm M}_{606}=-21.9$ and ${\\rm M}_{606}=-22$, respectively. Likewise, in the the medium-redshift MACS and RCS subsamples the median BCG magnitudes are ${\\rm M}_{814}=-23.8$, and ${\\rm M}_{814}=-23.6$, respectively. A Kolmogorov-Smirnov (KS) test indicates that for both the low- and the medium-redshift subsamples, the null hypothesis that both the RCS and MACS BCG magnitudes are derived from the same parent distribution cannot be confidently rejected (probabilities of 0.12 and 0.35, respectively, for the null hypothesis). These numbers are summarized in Table 8.  For a second comparison of optical luminosities, we measure integrated optical luminosity of the brightest galaxies within the cluster cores. We measure the total light of galaxies inside a physical aperture of radius 270 kpc (at low $z$) and 370 kpc (at medium $z$). The contribution of foreground and background galaxies to the light is determined statistically in annuli of $400-530$ kpc and $550-730$ kpc, for the low- and medium-redshift subsamples, respectively, and subtracted from the core light. The area in which we measure the ``background'' is still well within the cluster, and hence, our cluster core luminosities are underestimated due to background over-subtraction. Nevertheless, barring large profile differences (see below), these biased estimates of cluster luminosity can still be compared meaningfully between the X-ray and optical samples. We include only the light from objects with magnitudes fainter than the cluster BCG magnitude, but brighter than 24 mag. We convert the MACS low-redshift subsample's F606W luminosities to F814W luminosities assuming, again, the Kinney et al. (1996) elliptical galaxy template and the filter transmission curves for the two bands. The resultant optical luminosity distributions (Fig. 9) of both the RCS and MACS cluster are consistent with being drawn from the same parent distribution ($0.12$ and $0.37$ probabilities for the null hypothesis, see Table 8). \\begin{figure} \\center \\includegraphics[width=8cm, angle=0]{radial_arcs2.eps}  \\caption{$33'' \\times 28''$ image sections of two radially distorted objects (marked as $1$ and $2$) in the cluster MACSJ2129.4$-$0741.}  \\end{figure} \\begin{figure} \\center \\includegraphics[width=8cm, angle=0]{high_z_arc2.eps}  \\caption{$33'' \\times 28''$ image section of a large arc found in the field of view of the high-redshift cluster RCS025242.5$-$150024, around a foreground galaxy.}  \\end{figure}  Finally, as a third method of comparing optical luminosities, we simply count the light from {\\it {\\bf all}} the pixels inside the above apertures and annuli (but still leaving out the light from objects brighter than the BCG). This method takes into account the light from all the stars in the cluster cores, including stars in galaxies below the detection limit and diffuse intracluster light. As in the previous method, the core luminosity is underestimated due to background over-subtraction, but in a consistent way for the X-ray and optical clusters. In contrast to the two previous methods, where the dominant galaxies in the cluster core are early types, in this case the color correction, applied in order to convert F606W fluxes to F814W, is less clear-cut, since fainter and undetectable dwarf galaxies may well be blue. For the range in color terms from early-type to late-type galaxies, the luminosity distributions are either consistent with being drawn from the same parent distribution (0.05 probability for the null hypothesis, ``blue'' color correction), to marginally consistent (0.01 probability for the null hypothesis, ``red'' color correction). At most (in the case of the low-redshift subsamples, and assuming the reddest color correction) the medians of these two distributions differ only by a factor of $1.6$.  Overall, as shown in Table 8, the medians of the RCS and MACS cluster luminosity distributions and the results of the KS tests we applied to these distributions suggest that the RCS and MACS cluster samples have similar optical luminosities. We note also that the MACS clusters that actually display arcs (shaded histograms in Fig. 9) are not necessarily the most luminous ones, and that there is a large overlap of their luminosities with those of RCS clusters that are much less efficient arc producers. We have also measured and compared the optical light profiles of the two samples and, within the limited range of the cluster cores covered by the ACS data, we find no significant differences.  \\begin{table*} \\caption{Comparison of MACS and RCS cluster luminosities} \\smallskip \\begin{tabular}{lcc} \\hline \\noalign{\\smallskip} Subsample & Optical luminosity measure & KS Probability  \\\\ \\noalign{\\smallskip} \\hline \\multicolumn{3}{c}{Cluster BCG absolute magnitudes}\\\\ \\hline MACS ($0.3\\leq z < 0.5$) & $-21.9$ & \\multirow{2}{*}{$0.12$} \\\\ RCS  ($0.3\\leq z < 0.5$) & $-22$ & \\\\ MACS ($0.5\\leq z < 0.7$) & $-23.8$ & \\multirow{2}{*}{$0.35$} \\\\ RCS  ($0.5\\leq z < 0.7$) & $-23.6$ &\\\\ \\hline \\multicolumn{3}{c}{Cluster core luminosities (light of bright galaxies only)}\\\\ \\hline  MACS ($0.3\\leq z < 0.5$) & $1.5$ & \\multirow{2}{*}{$0.12$} \\\\ RCS  ($0.3\\leq z < 0.5$) & $1.3$ & \\\\ MACS ($0.5\\leq z < 0.7$) & $2.4$ & \\multirow{2}{*}{$0.37$} \\\\ RCS  ($0.5\\leq z < 0.7$) & $2.2$ & \\\\ \\noalign{\\smallskip} \\hline \\multicolumn{3}{c}{Cluster core luminosities (total light)}\\\\ \\hline MACS ($0.3\\leq z < 0.5$) & $3.8$ & \\multirow{2}{*}{$0.01$} \\\\ RCS  ($0.3\\leq z < 0.5$) & $2.3$ & \\\\ MACS ($0.5\\leq z < 0.7$) & $4.5$ & \\multirow{2}{*}{$0.07$} \\\\ RCS  ($0.5\\leq z < 0.7$) & $3.1$ & \\\\  \\noalign{\\smallskip} \\hline \\smallskip \\end{tabular}  Note- Optical luminosity measure indicates F606W or F814W absolute magnitudes for the BCGs, and luminosities in units of $10^{45}~{\\rm erg}~{\\rm s}^{-1}$ for the two measures of core optical luminosity. Probability is for the null hypothesis that a pair of distributions are not different.  \\end{table*}     However, stars, let alone the small fraction of the stars that dominate the optical luminosity, are a tiny component of the total cluster mass, and it is therefore plausible that the masses of the two samples are very different, despite the similar optical luminosities, with masses significantly below $10^{15} {\\rm M}_\\odot$ for the RCS clusters. A strong argument for such a mass difference is the difference in the space densities of the two samples. From the numbers of clusters and the area surveyed (see \\S~2), the projected density of MACS clusters is $\\sim 0.01 {\\rm ~deg}^{2}$. Assuming the cluster mass function is probed correctly by X-ray surveys, only about one MACS-like massive cluster  is expected in the $\\sim 100 {\\rm ~deg}^{2}$ search area of the RCS survey. Based on the cluster mass function (e.g. Reiprich \\& B{\\\" o}hringer 2002), the $\\sim 1000$ clusters found in the RCS search area imply that the vast majority of these clusters have masses   of ${\\rm M}_{\\rm 200} = 10^{14} {\\rm M}_\\odot$, an order of magnitude lower than MACS clusters.  This picture is further supported by the distributions of arc separations from the cluster centres. Since arcs occur near critical curves, the separations can roughly represent the Einstein radii of the clusters. As noted in \\S3 and seen in Fig.~4, the MACS clusters have arcs at $10''-50''$, with a median at $24''$, while the RCS arcs are generally much closer in, with a median of $10''$. The small Einstein radii of most of the RCS clusters with arcs are similar to those of rich groups.  A puzzling corollary of the above arguments, however, is the fact that in the small subsample of 52 RCS clusters imaged with ACS, which constitute  just 5 per cent of the full RCS sample, there are as many as two clusters (RCS 141910$+$5326.1 and RCS 215609.1$+$012319) with arcs at separations implying $20''$ Einstein radii, and hence MACS-like masses, in contrast to the expectation that of order just one such cluster exists in the {\\it entire} RCS survey. Furthermore, despite their mass, only about half of the MACS clusters display arcs in our survey, because a galaxy in a suitable position in the source plane is required in order to produce an arc. The two large-separation RCS clusters in the HST sample would thus imply about 4 massive RCS clusters in the HST sample, and $\\sim 100$ in the full RCS sample, as opposed to the $\\sim 1$ expected from the X-ray-derived mass function. A possible explanation is that the HST RCS sample is not a fully representative selection of the RCS. Indeed, the wide arcs of the cluster RCS 141910$+$5326.1 above were already noted in ground-based images by Gladders et al. (2003), and it may have been included in the HST sample for this reason. Thus, the HST RCS sample could be a representative subsample of the RCS, {\\it plus} a few of the most massive RCS clusters, and would thus be pre-biased in favor of lensing. Since, despite this bias, the X-ray-selected clusters are still much more efficient lenses, the observed arc occurrence frequency in the RCS clusters imaged with HST provides an upper limit on the RCS arc frequency as a whole.  On the other hand, Gladders et al. (2003) discussed eight potential lenses out of the full RCS sample of about 1000 clusters. With random selection, one would expect 1.1 of these eight lenses to be included among the 150 RCS clusters in the HST Snapshot sample, out of which the actually observed targets were chosen by HST schedulers. In fact, there are two of the eight Gladders et al. (2003) potential lenses among the 150 HST targets. This could be the result of some slight bias in favor of known lenses, as described above, but could of course be due just to chance. We reiterate that, if the RCS sample is unbiased, our conclusion about the relatively low lensing inefficiency of RCS clusters hold. If the RCS sample was pre-biased, this conclusion is only strengthened.  The simplest explanation for our measurement of a low lensing efficiency among RCS clusters, compared to X-ray-selected clusters, is a typical RCS cluster mass that is lower by an order of magnitude. This leaves open the question of what stands behind the similarity of X-ray and RCS clusters, in terms of stellar luminosity, optical profiles, numbers of galaxies, and general optical appearance. These similarities cannot be due to chance line-of-sight projections (RCS clusters are chosen based on redshifts of the early-type galaxies that characterise dense environments, so they are, in fact, real associations), nor due to projection effects along large-scale structure ``filaments'' -- simulations have shown that large scale structure may contribute only about 10 per cent to the cluster surface mass density (Wambsganss, Bode, \\& Ostriker 2005; Hilbert et al. 2007). Instead, there is a real and large variation in the total-mass-to-optical-light ratio among clusters. The low mass-to-light ratio of RCS cluster cores may be caused by a bias in favour of line-of-sight mergers in the optical selection process, a prominent and spectacular example of which is Cl0024+24 (Czoske et al. 2002). Indeed, extensive spectroscopic follow-up of RCS clusters has uncovered several cases of close projection effects of possibly physically associated systems as well as line-of-sight substructure (Gilbank et al. 2007; Cain et al. 2008).  A further effect to consider is the question of whether X-ray selection may favour the inclusion of clusters that are in the process of merging. Torri et al. (2004) have found that, during a merger, the lensing cross section is increased by a factor of $5-10$ for a duration of a couple of hundred million years, while the X-ray luminosities of merging clusters are increased by a factors of $\\sim 5$. A similar conclusion regarding the X-ray luminosity of clusters during mergers was reached by Randall, Sarazin, \\& Ricker (2002). If X-ray-selected cluster samples indeed have a larger fraction of merging clusters, one could thus expect a larger fraction of highly efficient lenses in those samples. Thus, the masses of the X-ray-selected clusters may be systematically overestimated as well.  An interesting question is whether comparable optically and X-ray-selected cluster samples at $z > 0.7$ also differ in their arc production efficiencies. We did not find giant arcs in any of the high-redshift RCS clusters we analyzed, even though their optical luminosities are comparable to those of the RCS clusters at low and medium redshifts. This contrasts with the results of Gladders et al. (2003) who found RCS clusters to be more efficient lenses at high redshift.  Finally, the many arcs found in the MACS low- and medium-redshift subsamples provide a statistically improved handle on the angular distribution of arcs in clusters. Our results show that arcs do form at large angular separations from cluster centres, at up to $60''$, in some cases. Thus, the large Einstein radius of Abell 1689 is probably not unique."
    },
    "1004/1004.3344_arXiv.txt": {
        "abstract": "We use the Markov Chain Monte Carlo method to  investigate a global constraints on  the generalized Chaplygin gas (GCG) model as the unification of dark matter and dark energy from  the latest observational data: the Constitution dataset of type supernovae Ia (SNIa), the observational Hubble data (OHD),  the cluster X-ray gas mass fraction, the baryon acoustic oscillation (BAO), and the cosmic microwave background (CMB) data. In a non-flat universe,  the constraint results for GCG model are, $\\Omega_{b}h^{2}=0.0235^{+0.0021}_{-0.0018}$ ($1\\sigma$) $^{+0.0028}_{-0.0022}$ $(2\\sigma)$, $\\Omega_{k}=0.0035^{+0.0172}_{-0.0182}$ ($1\\sigma$) $^{+0.0226}_{-0.0204}$ $(2\\sigma)$, $A_{s}=0.753^{+0.037}_{-0.035}$ ($1\\sigma$) $^{+0.045}_{-0.044}$ $(2\\sigma)$, $\\alpha=0.043^{+0.102}_{-0.106}$ ($1\\sigma$) $^{+0.134}_{-0.117}$ $(2\\sigma)$, and $H_{0}=70.00^{+3.25}_{-2.92}$ ($1\\sigma$) $^{+3.77}_{-3.67}$ $(2\\sigma)$, which is more stringent than the previous results for constraint on GCG model parameters. Furthermore, according to the information criterion, it seems that the current observations much support $\\Lambda$CDM model relative to the GCG model. ",
        "introduction": "$}   { Recently, mounting cosmic observations suggest that the expansion of present universe is speeding up rather than slowing down \\cite{SNeCMBLSS}.  And they indicates that baryon matter component is about 5\\% for total energy density, and about 95\\% energy density in universe is invisible. Considering the four-dimensional  standard cosmology, this accelerated expansion for universe predict that dark energy (DE) as an exotic component with negative pressure is filled in universe. And it is shown that DE takes up about two-thirds of the total energy density from cosmic observations. On the other hand, in theory many kinds of DE models have already been constructed in order to explore the DE properties. For a review on DE models, please see Refs. \\cite{DEmodels}.  It is well known that the generalized Chaplygin gas (GCG) model have been widely studied for interpreting the accelerating universe \\cite{GCG}\\cite{GCGpapers}. The most interesting property for this scenario is that, two unknown dark sections in universe--dark energy and dark matter  can be unified by using an exotic equation of state. In this paper, we use the Markov Chain Monte Carlo (MCMC) technique to constrain the GCG model from  the latest observational data: the Constitution dataset \\cite{397Constitution} including 397 type Ia supernovae (SNIa), the observational Hubble data (OHD) \\cite{OHD},  the cluster X-ray gas mass fraction \\cite{ref:07060033},  the measurement results of baryon acoustic oscillation (BAO) from Sloan Digital Sky Survey (SDSS) and Two Degree Field Galaxy Redshift Survey (2dFGRS) \\cite{SDSS}\\cite{ref:Percival2}, and the current cosmic microwave background (CMB) data from five-year WMAP \\cite{5yWMAP}. ",
        "conclusions": "$}  The constraints on the flat and non-flat GCG model as the unification of dark matter and dark energy are studied in this paper by using the latest observational data: the  Constitution dataset including 397 SNIa, the Hubble parameter data, the cluster X-ray gas mass fraction, the baryon acoustic oscillation and  the five-year WMAP data. The constraint  on GCG model parameters are more stringent than the previous papers \\cite{GCG-As-alpha}\\cite{constraintGCG}. According to the constraint results, since the best fit values of parameters $\\alpha$ and $\\Omega_{k}$ are near to zero, it seems that the current observations tends to make the GCG model reduce to  the flat $\\Lambda$CDM model. Furthermore, according to the IC, we can get the same result. In addition, we also make a stringent constraint on $\\Lambda$CDM model, and in a flat  $\\Lambda$CDM model  it is shown that the cosmic age is about, $t_{age}(Gyr)=13.725^{+0.099}_{-0.141}$ ($1\\sigma$) $^{+0.134}_{-0.165}$ $(2\\sigma)$, for using a tophat prior as 10 Gyr $< t_{age} <$ 20 Gyr in our calculation. At last,  for the  CBF constraint method, as a reference we list the best-fit values of the parameters in $f_{gas}$: $K=0.9919, \\eta=0.2089, \\gamma=1.0299, b_0=0.7728, \\alpha_b=-0.0582, s_0=0.1656, \\alpha_s=0.1128$ for the flat universe and $K=0.9871, \\eta=0.2114, \\gamma=1.0507, b_0=0.7749, \\alpha_b=-0.0950, s_0=0.1741, \\alpha_s=0.0194$ in the non-flat case."
    },
    "1004/1004.1177_arXiv.txt": {
        "abstract": "\\medskip \\noindent We propose a new mechanism for leptogenesis, which is naturally realized in models with a flavor symmetry based on the discrete group $A_4$, where the symmetry breaking parameter also controls the Majorana masses for the heavy right handed (RH) neutrinos. During the early universe, for $T\\gtrsim \\TeV$, part of the symmetry is restored, due to finite temperature contributions, and the RH neutrinos remain massless and can be produced in thermal equilibrium even at temperatures well below the most conservative gravitino bounds. Below this temperature the phase transition occurs and they become massive, decaying out of equilibrium and producing the necessary lepton asymmetry. Unless the symmetry is broken explicitly by Planck-suppressed terms, the domain walls generated by the symmetry breaking survive till the quark-hadron phase transition, where they disappear due to a small energy splitting between the $A_4$ vacua caused by the QCD anomaly. ",
        "introduction": "Leptogenesis provides a natural scenario to explain the observed baryon abundance, where an asymmetry in the leptonic sector is initially generated by the out-of-equilibrium decay of heavy right-handed (RH) neutrinos and then distributed into baryons by so-called sphaleron processes \\cite{LG}. The most attractive feature of leptogenesis is that it relies uniquely on the addition to the Standard Model of three RH neutrinos, addition which is already very welcome to explain the smallness of light neutrino masses via the see-saw mechanism \\cite{seesaw}. This synergy between neutrino physics and the origin of matter underlines the need for studying leptogenesis in models with realistic neutrino mixing.  While our understanding of the origin of fermion masses and mixing angles remains at a primitive level, flavor symmetries provide a practical mean to replicate, within experimental uncertainties, the observed data. As originally proposed by Froggat and Nielsen \\cite{Froggatt:1978nt}, a broken $U(1)$ flavor symmetry can be responsible for the small ratios between masses in the quark sector. Similarly, tri-bimaximal neutrino mixing, in excellent agreement with observations from neutrino oscillation experiments, can be reproduced within models with broken discrete non-abelian symmetries in the leptonic sector where sub-leading corrections in the symmetry breaking parameter account for small deviations from exact tri-bimaximal mixing \\cite{Altarelli:2005yp}.   In this article, % we study leptogenesis in models with natural tri-bimaximal mixing from a flavor symmetry based on the discrete group $A_4$ \\cite{Altarelli:2005yx}. The arguments we use, however, can be easily extended to a more general class of models where the scale of lepton number violation is introduced by a spontaneous symmetry breaking mechanism. Symmetries that are spontaneously broken today might be restored during the early universe, due to finite temperature effects. We find that this symmetry restoration provides a natural scenario for thermal leptogenesis to work efficiently even at low reheat temperatures $T_{RH}\\lesssim 10^{4\\div5}\\GeV$, well below the most stringent bounds from gravitino overproduction, with no need for the addition of  structure or fine-tuning. Indeed, a crucial aspect of the $A_4$-based models  is the structure of the breaking sector and the particular vacuum alignment that reproduces tri-bimaximal mixing in the neutrino sector. The symmetry is partially broken (down to a subgroup isomorphic to $Z_2$) by the vacuum expectation value (VEV) of a flat direction that obtains a potential  only in the presence of soft supersymmetry (SUSY) breaking terms. During the early universe, the finite energy density of the inflationary vacuum or the different occupation numbers of fermions and bosons in the thermal bath after reheating, break SUSY and contribute a (positive\\footnote{For particular forms of the K\\\"ahler potential is it also possible to arrange for a negative mass of the flat-direction during inflation. In this case, the peculiar dynamics of flat directions in the early universe \\cite{Basboll:2007vt,Giudice:2008gu} can reproduce the observed baryon asymmetry in a non-thermal way, as in \\cite{Giudice:2008gu} or via Affleck-Dine baryogenesis \\cite{Affleck:1984fy}.}) mass-squared term for flat directions \\cite{Dine:1995uk}, restoring part of the flavor symmetry. The relevant aspect for leptogenesis is that, in these models, the lepton number violating Majorana masses of RH neutrinos are also proportional to the flavor symmetry breaking parameter and therefore vanish in the early universe. Thus, at early times, the would-be heavy RH neutrinos are effectively massless states and remain in equilibrium with the thermal bath down to temperatures of order the SUSY breaking soft masses, $\\tilde{m}\\sim {\\cal O}(10^2-10^3)~\\GeV$. Below this temperature, the thermal corrections to the flat direction potential become too small and its origin unstable: a smooth phase transition takes place in which the flat direction VEV rolls to large values, breaking the flavor group and lepton number and giving large Majorana masses $M_i$ to the RH neutrinos, which suddenly find themselves out of equilibrium at $T\\ll M_i$. This is the perfect starting point for standard thermal Leptogenesis: an equilibrium abundance of RH neutrinos, ready to decay via CP-violating interactions producing a large lepton asymmetry. At temperatures above $\\sim 100~\\GeV$, sphaleron processes convert this asymmetry into the observed baryon asymmetry \\cite{kuzmin}.  A potential problem of discrete symmetries broken spontaneously at low temperatures is the creation of domain walls. Their energy density red-shifts slower than that of matter or radiation and eventually, independently of their initial abundance, they come to dominate the energy density of the whole universe, leading to cosmological consequences incompatible with observations, such as imprinting large signatures in the cosmic microwave background \\cite{DW}. We will show that, in the $A_4$-based models, the putative discrete symmetry is not one in fact: it is broken explicitly by the QCD anomaly that lifts the degeneracy between vacua separated by domain wall and drives them to collapse. Thus, unless the $A_4$ symmetry is also broken explicitly by the gravitational interactions (in which case the domain walls do not form at all), the domain wall network disappears at the QCD scale and standard cosmology is recovered.  In the next section we review models based on the $A_4$ flavor symmetry to reproduce tri-bimaximal mixing,  paying particular attention to the structure of the symmetry breaking sector. In section \\ref{SecThermal} we discuss the dynamics of the flavor symmetry breaking during the early universe and the thermal production of RH neutrinos. In section \\ref{LG} and \\ref{secDW} we conclude with a study of leptogenesis and a note on the domain wall problem. ",
        "conclusions": "We have analyzed thermal leptogenesis in the framework of models with an $A_4$ flavor symmetry that naturally reproduce tri-bimaximal neutrino mixing, although our findings apply to the more general case where the RH neutrino masses arise from a mechanism of spontaneous symmetry breaking. We found that leptogenesis can be successful, without any fine-tuning, even at low energies well below the gravitino bound, independently of the masses of RH neutrinos. This is possible thank to a mechanism of symmetry restoration, already present in the $A_4$-based models, that ensures that the RH neutrino masses vanish at temperatures $T\\gtrsim 1 \\TeV$. Hence, during the early universe, the RH neutrinos can be produced copiously by thermal scattering, without requiring very high temperatures that would lead to the overproduction of gravitinos, in contrast with the standard cosmological scenario. The symmetry breaking discussed here takes place along a supersymmetric flat direction, lifted only by SUSY breaking effects. Its potential is therefore naturally almost flat and induces very large field VEVs, despite its scale being of order the soft SUSY breaking masses, $\\tilde{m}\\sim {\\cal O}(10^2 - 10^3) \\GeV$. Hence the phase transition that gives a mass to the RH neutrinos takes place only at temperatures much smaller than the would-be RH neutrino masses. Eventually the heavy neutrinos decay out of equilibrium via their CP-violating interactions producing an abundance of leptons over anti-leptons which is then transformed into the observed baryon asymmetry by sphaleron interactions.  The $A_4$-based models are very predictive in terms of the neutrino mixing angles. Indeed all 3 mixing angles are predicted to depart from their pure tri-bimaximal values, eq. (\\ref{TB}) by the same amount. Since measured deviation of $\\theta_{12}$ are very small, see eq. (\\ref{data}), the same must be true for $\\theta_{13}$. So, observation of $\\theta_{13}\\neq 0$ in future experiments will be able to exclude models with tri-bimaximal mixing in the neutrino sector.  Models with discrete symmetry spontaneously broken at energies below the inflationary scale, generally suffer from a domain wall problem. We showed that the $A_4$-based models, when extended to account for quark masses, posses a QCD anomaly, in the sense that the $A_4$ acts non-trivially on the quark mass matrix. In this case, the same mechanism that gives a mass to the axion, here lifts the degeneracy between vacua on opposite sides of the domain walls. We found that for neutrinos with inverted hierarchy this solution of the domain wall problem is compatible with our scenario of leptogenesis. For neutrinos with a normal hierarchy, on the contrary, this effect is not enough to avoid a domain wall dominated universe if we insist on the mechanism of leptogenesis as described above. In both cases, if the discrete symmetries are broken explicitly by the gravitational interactions, there remains no barrier separating different vacua and domain walls never form.     \\vspace{2cm} \\centerline{\\bf Ackowledgements} \\noindent I'm particularly indebted with Ferruccio Feruglio for several useful discussions and with Antonio Riotto for his comments. I also thank Ben Gripaios and Luca Merlo for interesting conversations. This work was supported by the Fondazione Cariparo Excellence Grant \\emph{LHCosmo} and in part by the European Programme \\emph{Unification in the LHC Era}, contract PITN-GA-2009-237920 (UNILHC), and the Research and Training Network UniverseNet."
    },
    "1004/1004.2032_arXiv.txt": {
        "abstract": "The Extragalactic Background Light (EBL) is the integrated light from all the stars that have ever formed, and spans the IR-UV range. The interaction of very-high-energy (VHE: $E>100\\,$GeV) $\\gamma$-rays, emitted by sources located at cosmological distances, with the intervening EBL results in $e^-e^+$ pair production that leads to energy-dependent attenuation of the observed VHE flux. This introduces a fundamental ambiguity in the interpretation of measured VHE\\,$\\gamma$-ray spectra: neither the intrinsic spectrum, nor the EBL, are separately known -- only their combination is. In this paper we propose a method to measure the EBL photon number density. It relies on using simultaneous observations of BL\\,Lac objects in the optical, X-ray, high-energy (HE: $E>100\\,$MeV) $\\gamma$-ray (from the Fermi telescope), and VHE\\,$\\gamma$-ray (from Cherenkov telescopes) bands. For each source, the method involves best-fitting the spectral energy distribution (SED) from optical through HE\\,$\\gamma$-rays (the latter being largely unaffected by EBL attenuation as long as $z \\mincir 1$) with a Synchrotron Self-Compton (SSC) model. We extrapolate such best-fitting models into the VHE regime, and assume they represent the BL\\,Lacs' intrinsic emission. Contrasting measured versus intrinsic emission leads to a determination of the $\\gamma\\gamma$ opacity to VHE photons. Using, for each given source, different states of emission will only improve the accuracy of the proposed method. We demonstrate this method using recent simultaneous multi-frequency observations of the high-frequency-peaked BL\\,Lac object PKS\\,2155-304 and discuss how similar observations can more accurately probe the EBL. ",
        "introduction": "The Extragalactic Background Light (EBL), in both its level and degree of cosmic evolution, reflects the time integrated history of light production and re-processing in the Universe, hence the history of cosmological star-formation. Roughly speaking, its shape must reflect the two humps that characterize the spectral energy distributions (SEDs) of galaxies: one arising from starlight and peaking at $\\lambda \\sim 1\\,\\mu$m (optical background), and one arising from warm dust emission and peaking at $\\lambda \\sim 100\\,\\mu$m (infrared background). However, direct measurements of the EBL are hampered by the dominance of foreground emission (interplanetary dust and Galactic emission), hence the level of EBL emission is uncertain by a factor of several.  One approach to evaluate the EBL emission level has been modeling the integrated light that arises from an evolving population of galactic stellar populations. However, uncertainties in the assumed galaxy formation and evolution scenarios, stellar initial mass function, and star formation rate have led to significant discrepancy among models (e.g., Salamon \\& Stecker 1998; Malkan \\& Stecker 1998 and Stecker \\& de\\,Jager 1998; Kneiske et al. 2002 and 2004; Stecker, Malkan \\& Scully 2006; Razzaque, Dermer \\& Finke 2009 and Finke, Razzaque \\& Dermer 2010). These models have been used to correct observed VHE spectra and deduce (EBL model dependent) 'intrinsic' VHE\\,$\\gamma$-ray emissions.  The opposite approach, of a more phenomenological kind, deduces upper limits on the level of EBL attenuation making basic assumptions on the intrinsic VHE\\,$\\gamma$-ray shape of AGN spectra. Specifically, it was assumed that the latter are described by a power-law photon index $\\Gamma \\geq 1.8$ (Schroedter 2005), $\\Gamma \\geq 1.5$ (e.g., Aharonian et al. 2006; Mazin \\& Goebel 2007; Mazin \\& Raue 2007), and $\\Gamma \\geq 1$ (Finke \\& Razzaque 2009). These assumptions correspond to various possibilities of producing TeV spectra. Shock-accelerated electrons are unikely to produce VHE\\,$\\gamma$-rays with $\\Gamma < 1.5$ from Compton scattering (e.g., Blandford \\& Eichler 1987). However, either internal $\\gamma \\gamma$ absorption (Aharonian et al. 2008), or harder electron spectra at the highest energies in relativistic shocks (Stecker, Baring \\& Summerlin 2007), or Compton scattering of CMB photons (B\\\"ottcher et al. 2008), or top-heavy power-law energy distributions of the emitting electrons (Katarzynski et al. 2006) could lead to harder intrinsic TeV spectra -- not to mention that pion decay from a hadronic source would produce a very hard TeV component, irrespective of the lower-energy electron synchrotron spectrum (M\\\"ucke et al. 2003). Another proposed approach to derive EBL upper limits involves assuming that a same-slope extrapolation of the observed {\\it Fermi}/LAT HE spectrum into the VHE domain exceeds the intrinsic VHE spectrum there (Georganopulos, Finke \\& Reyes 2010). An approach to exploring the redshift evolution of the EBL exploits the GeV-TeV connection for blazar spectra (Stecker \\& Scully 2010).  A different, but related, approach to constraining the EBL rests on evaluating the collective blazar contribution to the extragalactic $\\gamma$-ray background when the inescapable electromagnetic cascades of lower-energy photons and electrons, initiated by the interaction of VHE photons with the EBL, are accounted for. The collective intensity of a cosmological population of VHE\\,$\\gamma$-ray sources will be attenuated at the highest energies through interaction with the EBL and enhanced at lower energies by the resulting cascade: the strength of the effect depends on the source $\\gamma$-ray luminosity function and spectral index distribution, and on the EBL model (Venters 2010). The extragalactic $\\gamma$-ray background, resulting from the contributions of different classes of blazars, can then be used to constrain the EBL (Kneiske \\& Mannheim 2008; Venters, Pavlidou \\& Reyes 2009) -- even though the amount of energy flux absorbed and reprocessed is probably only a small fraction of the total extragalactic $\\gamma$-ray background energy flux (Inoue \\& Totani 2009).  The only unquestionable constraints on the EBL are model-independent lower limits based on galaxy counts (Dole et al. 2006; Franceschini, Rodighiero \\& Vaccari 2008). It should be noted, however, that the EBL upper limits in the 2--80$\\mu$m obtained by Mazin \\& Raue (2007) combining results from all known TeV blazar spectra (based on the assumption that the intrinsic $\\Gamma \\geq 1.5$) are only a factor $\\approx$2--2.5 above the absolute lower limits from source counts. So it would appear that there is little room for additional components like Pop\\,III stars (Raue, Kneiske \\& Mazin 2009; Aharonian et al. 2006), unless we miss some fundamental aspects of blazar emission theory (which have not been observed in local sources, however).  An attempt to measure the EBL used the relatively faraway blazar 3C\\,279 as a background light source (Stecker, de\\,Jager \\& Salamon 1992), assuming that the intrinsic VHE spectrum was known from extrapolating an apparently perfect $E^{-2}$ power-law differential energy spectrum, known in the interval from 70\\,MeV to $>$5\\,GeV from EGRET data, by a couple of further decades in energy into the VHE regime. However, blazars are highly variable sources, so it's almost impossible to determine with confidence the intrinsic TeV spectrum -- which itself can be variable.  In this paper we propose a method to measure the EBL that improves on Stecker et al. (1992) by making a more realistic assumption on the intrinsic TeV spectrum. Simultaneous optical/X-ray/HE/VHE (i.e., eV/keV/GeV/TeV) data are crucial to this method, considering the strong and rapid variability displayed by most blazars. After reviewing features of EBL absorption (sect.\\,2) and of the adopted BL\\,Lac emission model (sect.\\,3), in sect.\\,4 we describe our technique, in sect.\\,5 we apply it to recent multifrequency observations of PKS\\,2155-304 and determine the $\\gamma \\gamma$ optical depth out to that source's redshift. In sect.\\,6 we discuss our results. ",
        "conclusions": "The method for measuring the EBL we have proposed in this paper is admittedly model-dependent. However, its only requirement is that all the sources used as background beamlights should have one same emission model. In the application proposed here, we have used a one-zone SSC model where the electron spectrum was a (smoothed) double power law applied to the SED of the HBL object PKS\\,2155-304. While this choice was encouraged by the current observational evidence that HBLs seem to have, with no exception, single-slope {\\it Fermi}/LAT spectra, we could have as well adopted the choice (Aharonian et al. 2009) of a triple power law electron spectrum in our search for the best-fit SSC model of PKS\\,2155-304's SED. Should the latter electron distribution, or any other (e.g., curved) distribution, be shown to generally provide a better fit to high-quality {\\it Fermi}/LAT spectra of HBLs, then that would become our choice. In general, what matters to the application of this method, is that {\\it all} source SEDs be fit with one same SSC model.  Another assumption implicit in our method is that there is an absolute minimum in the $\\chi^2$ manifold of BL\\,Lac emission modeling, and that our $\\chi^2$-minimization procedure is actually able to find it. Had that not been the case for PKS\\,2155-304, we would have checked whether the derived $\\tau_{\\gamma\\gamma}$s are appreciably different for different model fits."
    },
    "1004/1004.2803_arXiv.txt": {
        "abstract": "{ The Crab pulsar is well-known for its anomalous giant radio pulse emission. Past studies have concentrated only on the very bright pulses or were insensitive to the faint end of the giant pulse luminosity distribution. With our new instrumentation offering a large bandwidth and high time resolution combined with the narrow radio beam of the Westerbork Synthesis Radio Telescope (WSRT), we seek to probe the weak giant pulse emission regime.  The WSRT was used in a phased array mode, resolving a large fraction of the Crab nebula. The resulting pulsar signal was recorded using the PuMa II pulsar backend and then coherently dedispersed and searched for giant pulse emission. After careful flux calibration, the data were analysed to study the giant pulse properties. The analysis includes the distributions of the measured pulse widths, intensities, energies, and scattering times. The weak giant pulses are shown to form a separate part of the intensity distribution. The large number of giant pulses detected were used to analyse scattering and scintillation in giant pulses. We report for the first time the detection of giant pulse emission at both the main- and interpulse phases within a single rotation period. The rate of detection is consistent with the appearance of pulses at either pulse phase as being independent.  These pulse pairs were used to examine the scintillation timescales within a single pulse period.  } ",
        "introduction": "\\label{intro}  Identified as the supernova remnant that resulted from SN 1054, the Crab nebula is one of the strongest radio sources in the sky, and it harbours the young neutron star PSR B0531+21. The pulsar is visible across the entire observable electromagnetic spectrum, and at radio wavelengths it is the second brightest pulsar in the northern sky. PSR B0531+21 was discovered by \\citet{sr68}, soon after the discovery of pulsars. This pulsar is noted for several features including the near orthogonal alignment of the magnetic and rotational axis that gives rise to the observed interpulse emission. The average emission profile of the pulsar, obtained by averaging the radio emission from many rotations of the star, exhibits a number of features that change quite remarkably with radio frequency \\citep{mh94}. The single pulses show a large variation in amplitude and duration as a function of time. The most enigmatic of these are its occassional intense bursts known as giant pulses \\citep{hcr70,ss70}. The giant pulses can be extremely narrow, of the order of 0.4 $n$s \\citep{he07} and the pulse flux can be several 1000 times the average pulse flux. The ultrashort durations of the giant pulses imply very high equivalent brightness temperatures \\citep{hkwe03} indicating that they originate from nonthermal, coherent emission processes. In this work, we define giant pulses as the pulses with a significantly narrower width than the average emission and contain a flux of at least 10 times the mean flux density of the pulsar.  The Crab pulsar is one of just a handful of pulsars that have been shown to have giant pulse emission. Some other pulsars, like the young Vela pulsar, also show narrow, bursty emission called giant micropulses \\citep{jvkb01}. The fluxes of these micropulses are within a factor of 3 times the average pulse flux. In the pulsars that show giant pulse emission, the pulse intensity and energy distributions exhibit power-law statistics \\citep{ag72}, while the giant micropulses give rise to log-normal distributions \\citep{cjd01}. In contrast, the bulk of the pulsar population have pulse intensities and energies that follow either a normal or an exponential distribution \\citep{hw74,rit76}. This indicates that the giant pulses and micropulses may form a different emission population.  The Crab giant pulses have been studied by different groups, yet the nature of the emission process remains elusive. In the very early studies at low sky frequencies, the data were afflicted by dispersion smearing and scattering \\citep{hcr70,ga72}, but the power-law nature of the intensity distribution of giant pulses was identified. In the next major study, \\citet{lcu+95} discuss a multi-wavelength observation of giant pulse emission, and note the possibility of a weak giant pulse emission population at radio wavelengths, which they are unable to resolve owing to insufficient sensitivity. \\citet{sbh+99} found that the Crab giant pulses are broad band at radio wavelengths. They also determine giant pulse spectral indices in the range of -2.2 to -4.9 using their widely spaced observation bands and 29 simultaneously detected giant pulses. Observations by \\citet{hkwe03} revealed that giant pulses at 5.5 GHz contain nanosecond wide subpulses and the presence of such narrow features has been predicted in numerical modelling by \\citet{wea98}. At these frequencies the radio emission character of the Crab pulsar changes, with the interpulse emission becoming dominant. A multi-wavelength radio observation of Crab giant pulses with widely spaced frequency bands (0.43 GHz and 8.8 GHz) is presented by \\citet{cbh+04}, who discuss the effects of scintillation over a wide range of frequencies. \\citet{ps07} and \\citet{eah+02} investigated pulse width distributions and find that narrow pulses tend to be brighter. \\citet{btk08} carried out a similar analysis in addition to scattering and dispersion variations in the nebula. All of these studies point to the peculiarity of the Crab pulsar and its puzzling emission process, and motivates further study in finer detail using a large number of pulses. For the work discussed in this paper, we utilised the wide band capabilities of the new pulsar machine, PuMa--II \\citep{kss08} and the Westerbork Synthesis Radio Telescope (WSRT) in the coherent tied-array mode.  At small hour angles, the synthesised beam of the WSRT effectively resolves out the Crab nebula, reducing the nebular contribution to the system temperature. Thus the WSRT and PuMa--II combination makes this study much more sensitive in terms of signal-to-noise ratio achieved, and in number of pulses than was possible in the past. The rest of the paper is organised as follows: in \\S{2} we describe the observational set up and data reduction, flux calibration is discussed in \\S{3}, the giant pulse characteristics are discussed in \\S{4}.  We report detections of double giant pulses in \\S{5}, and the scattering analysis is presented in \\S\\ref{scatter}.  \\begin{figure}[htbp] \\includegraphics[scale=0.95]{13279fg1.ps} \\caption{Total intensity of a coherently dedispersed giant pulse at the main pulse phase detected in all recorded bands at 4.1 $\\mu$s resolution. The total dispersion delay of 24.9 $m$s across the seven bands was removed for this plot. The lower most panel shows the pulse after combining the signal in all seven bands. The pulses displayed here are scaled relative to the pulse at 1330 MHz.} \\label{GPs} \\end{figure} ",
        "conclusions": "The large collection of single pulses we gathered has allowed us to perform a range of statistics with the data.  After careful flux calibration, a detailed analysis of the pulse intensities, energies, widths, and separation times was done by computing distributions of these quantities. In the single-pulse intensity distributions, we find a clear evidence of two distinct populations in the giant pulses. The giant pulse separation times show a Poission distribution, and the rate of occurrence of giant pulses was determined. Spectral indices for a large number of giant pulses were computed with the narrowly spaced multi band data. Significant dispersion in the spectral indices was found and a small negative average spectral index was found for the main- and interpulse giants, and they are flatter than the average pulse emission. We also note that in some cases there is evidence for intensity modulation with bandwidths that are smaller than the full band but not consistent with scintillation effects. The previously undetected double giant pulses were presented and we find that they are not more frequent than would be expected by chance. The scatter time for a large number of giant pulses was determined by modelling the scatter broadening as an exponenial function and the distribution of scatter times was computed. The double giant pulses were reported for the first time and it is found that they are not very different from the normal giant pulses. Using multiple emission components either at the main- or interpulse phase and the double giant pulses, we find no evidence of variation of the scattering material on timescales shorter than 14 $m$s based on the correlation coefficient computed for emission within a single-pulse period."
    },
    "1004/1004.0492_arXiv.txt": {
        "abstract": "\\vspace{1cm} \\centerline{\\bf ABSTRACT}\\vspace{2mm} In this work, we consider the cosmological constraints on the interacting dark energy models. We generalize the models considered previously by Guo {\\it et al.}~\\cite{r15}, Costa and Alcaniz~\\cite{r16}, and try to discuss two general types of models: type~I models are characterized by $\\rho_{_X}/\\rho_m=f(a)$ and $f(a)$ can be any function of scale factor $a$, whereas type~II models are characterized by $\\rho_m=\\rho_{m0}\\,a^{-3+\\epsilon(a)}$ and $\\epsilon(a)$ can be any function of $a$. We obtain the cosmological constraints on the type~I and~II models with power-law, CPL-like, logarithmic $f(a)$ and $\\epsilon(a)$ by using the latest observational data. ",
        "introduction": "\\label{sec1}  The dark energy has been one of the most active fields in modern cosmology since the discovery of the accelerated expansion of our universe (see e.g.~\\cite{r1} for reviews). Among the conundrums in the dark energy cosmology, the so-called cosmological coincidence problem is the most familiar one. This problem is asking why are we living in an epoch in which the densities of dark energy and matter are comparable? Since their densities scale differently with the expansion of our universe, there should be some fine-tunings. To alleviate the cosmological coincidence problem, it is natural to consider the possible interaction between dark energy and dark matter in the literature (see e.g.~\\cite{r2,r3,r4,r5,r6,r7,r8,r9,r31,r32}). In fact, since the nature of both dark energy and dark matter are still unknown, there is no physical argument to exclude the possible interaction between them. On the contrary, some observational evidences of this interaction have been found recently. For example, in a series of papers by Bertolami {\\it et al.}~\\cite{r10}, they shown that the Abell Cluster A586 exhibits evidence of the interaction between dark energy and dark matter, and they argued that this interaction might imply a violation of the equivalence principle. On the other hand, in~\\cite{r11}, Abdalla {\\it et al.} found the signature of interaction between dark energy and dark matter by using optical, X-ray and weak lensing data from 33 relaxed galaxy clusters. Therefore, it is reasonable to consider the interaction between dark energy and dark matter in cosmology.  We consider a flat Friedmann-Robertson-Walker (FRW) universe. In the literature, it is usual to assume that dark energy and dark matter interact through a coupling term $Q$, according to \\bea &&\\dot{\\rho}_m+3H\\rho_m=Q,\\label{eq1}\\\\ &&\\dot{\\rho}_{_X}+3H\\rho_{_X}(1+w_{_X})=-Q,\\label{eq2} \\eea where $\\rho_m$ and $\\rho_{_X}$ are densities of dark matter and dark energy (we assume that the baryon component can be ignored); $w_{_X}$ is the equation-of-state parameter (EoS) of dark energy and it is assumed to be a constant; a dot denotes the derivative with respect to cosmic time $t$; $H\\equiv\\dot{a}/a$ is the Hubble parameter; $a=(1+z)^{-1}$ is the scale factor (we have set $a_0=1$; the subscript ``0'' indicates the present value of corresponding quantity; $z$ is the redshift). Notice that Eqs.~(\\ref{eq1}) and (\\ref{eq2}) preserve the total energy conservation equation \\be{eq3} \\dot{\\rho}_{tot}+3H\\rho_{tot}(1+w_{\\rm eff})=0, \\ee where $\\rho_{tot}=\\rho_{_X}+\\rho_m$ is the total energy; $w_{\\rm eff}$ is the total (effective) EoS. Since there is no natural guidance from fundamental physics on the coupling term $Q$, one can only discuss it to a phenomenal level. The most familiar coupling terms extensively considered in the literature are $Q=\\alpha\\kappa\\rho_m\\dot{\\phi}$, $Q=3\\beta H\\rho_{tot}$, and $Q=3\\eta H\\rho_m$. The first one arises from, for instance, string theory or scalar-tensor theory (including Brans-Dicke theory)~\\cite{r4,r5,r6}. The other two are phenomenally proposed to alleviate the coincidence problem in the other dark energy models~\\cite{r7,r8,r9}. In the usual approach, one should priorly write down the coupling term $Q$, and then obtain the evolutions of $\\rho_m$ and $\\rho_{_X}$ from Eqs.~(\\ref{eq1}) and (\\ref{eq2}), respectively. In fact, this is the common way to study the interacting dark energy models in the literature.  However, there is also an alternative way in the literature~\\cite{r12,r13,r14,r15,r16}. One can reverse the logic mentioned above. Due to the interaction $Q$, the evolutions of $\\rho_m$ and $\\rho_{_X}$ should deviate from the ones without interaction, i.e., $\\rho_m\\propto a^{-3}$ and $\\rho_{_X}\\propto a^{-3(1+w_{_X})}$, respectively. If the deviated evolutions of $\\rho_m$ and/or $\\rho_{_X}$ are given, one can find the corresponding interaction $Q$ from Eqs.~(\\ref{eq1}) and (\\ref{eq2}). Naively, the simplest example has been considered by Wang and Meng~\\cite{r12}, namely \\be{eq4} \\rho_m=\\rho_{m0}\\,a^{-3+\\epsilon}, \\ee where $\\epsilon$ is a constant which measures the deviation from the normal $\\rho_m\\propto a^{-3}$. Substituting into Eq.~(\\ref{eq1}), it is easy to find the corresponding interaction $Q=\\epsilon H\\rho_m$~\\cite{r9,r12,r13,r17}. Alternatively, one can consider another type of interacting dark energy model which is characterized by~\\cite{r14,r15} \\be{eq5} \\frac{\\rho_{_X}}{\\rho_m}=\\frac{\\rho_{_{X}0}}{\\rho_{m0}}\\,a^\\xi, \\ee where $\\xi$ is a constant which measures the severity of the coincidence problem. From Eqs.~~(\\ref{eq1}), (\\ref{eq2}) and~(\\ref{eq5}), one can find that the corresponding interaction is given by~\\cite{r15} \\be{eq6} Q=-H\\rho_m\\Omega_X\\left(\\xi+3w_{_X}\\right) =-H\\rho_{_X}\\Omega_m\\left(\\xi+3w_{_X}\\right), \\ee where $\\Omega_i\\equiv 8\\pi G\\rho_i/(3H^2)$ for $i=m$ and $X$, which are the fractional energy densities of dark matter and dark energy, respectively. In fact, Guo {\\it et al.}~\\cite{r15} considered the cosmological constraints on the interacting dark energy model characterized by Eq.~(\\ref{eq5}) with the 71 SNLS Type~Ia supernovae (SNIa) dataset, the shift parameter $R$ from the Wilkinson Microwave Anisotropy Probe 3-year (WMAP3) data, and the distance parameter $A$ of the measurement of the BAO peak in the distribution of SDSS luminous red galaxies. On the other hand, the interacting dark energy model characterized by Eq.~(\\ref{eq4}) has been extended in~\\cite{r16}. It is more realistic that $\\epsilon$ is a function of time. Costa and Alcaniz~\\cite{r16} considered the interacting dark energy model characterized by \\be{eq7} \\rho_m=\\rho_{m0}\\,a^{-3+\\epsilon(a)}, \\ee in which $\\epsilon(a)$ was chosen to be \\be{eq8} \\epsilon(a)=\\epsilon_0\\,a^{\\epsilon_1}, \\ee where $\\epsilon_0$ and $\\epsilon_1$ are constants. They obtained the constraints on this model by using the 307 Union SNIa dataset, the CMB constraint $\\Omega_{m0}h^2=0.109\\pm 0.006$ from the Wilkinson Microwave Anisotropy Probe 5-year (WMAP5) data, and the distance ratio from $z_{BAO}=0.35$ to $z_{LS}=1089$ measured by SDSS, namely $R_{BAO/LS}=0.0979\\pm 0.0036$.  In the present work, we generalize the interacting dark energy models considered in~\\cite{r15,r16}, and we call them type~I and II models, respectively. The type~I models are characterized by \\be{eq9} \\frac{\\rho_{_X}}{\\rho_m}=f(a), \\ee where $f(a)$ can be any function of $a$, beyond the special case in Eq.~(\\ref{eq5}). The type~II models are characterized by Eq.~(\\ref{eq7}) but $\\epsilon(a)$ can be any function of $a$, beyond the special case in Eq.~(\\ref{eq8}). In the present work, we consider the constraints on these models by using the latest cosmological observations, namely, the 397 Constitution SNIa dataset~\\cite{r18}, the shift parameter $R$ from the newly released Wilkinson Microwave Anisotropy Probe 7-year (WMAP7) data~\\cite{r19}, and the distance parameter $A$ of the measurement of the BAO peak in the distribution of SDSS luminous red galaxies~\\cite{r20,r21}. In the next section, we briefly introduce these observational data. In Sec.~\\ref{sec3} and Sec.~\\ref{sec4}, we discuss the type~I and~II models, and consider their cosmological constraints, respectively. A brief summary is given in Sec.~\\ref{sec5}. ",
        "conclusions": "\\label{sec5}  In this work, we considered the cosmological constraints on the interacting dark energy models. We generalized the models considered previously by Guo {\\it et al.}~\\cite{r15}, Costa and Alcaniz~\\cite{r16}, and we have discussed two general types of models: type~I models are characterized by $\\rho_{_X}/\\rho_m=f(a)$ and $f(a)$ can be any function of scale factor $a$, whereas type~II models are characterized by $\\rho_m=\\rho_{m0}\\,a^{-3+\\epsilon(a)}$ and $\\epsilon(a)$ can be any function of $a$. We obtained the cosmological constraints on the type~I and~II models with power-law, CPL-like, logarithmic $f(a)$ and $\\epsilon(a)$ by using the latest observational data.  Some remarks are in order. Firstly, here we briefly justify the interaction forms considered in the present work. We take type~I models as examples. For the power-law case with $f(a)=f_0\\,a^\\xi$ in Eq.~(\\ref{eq23}), noting that in the case without interaction $\\rho_{_X}\\propto a^{-3(1+w_{_X})}$ and $\\rho_m\\propto a^{-3}$, from definition Eq.~(\\ref{eq9}), it is reasonable to parameterize $f(a)=\\rho_{_X}/\\rho_m\\propto a^\\xi$, where $\\xi$ measures the severity of the coincidence problem~\\cite{r14,r15}. For the CPL case with $f(a)=f_0+\\xi(1-a)$ in Eq.~(\\ref{eq26}) and the logarithmic case with $f(a)=f_0+\\xi\\ln a$ in Eq.~(\\ref{eq29}), noting that the Taylor series expansion of any function $F(x)$ is given by $F(x)=F(x_0)+F_1\\,(x-x_0)+(F_2/\\,2!)\\,(x-x_0)^2+ (F_3/\\,3!)\\,(x-x_0)^3+\\dots$, the CPL and logarithmic cases can be regarded as the Taylor series expansion of $f$ with respect to the scale factor $a$ and the $e$-folding time ${\\cal N}=\\ln a$ up to first order (linear expansion), similar to the well-known EoS parameterizations $w(a)=w_0+w_a(1-a)$ and $w(z)=w_0+w_1\\,z$.  Secondly, we would like to briefly consider the comparison of these models. For convenience, we also consider the well-known $\\Lambda$CDM model in addition. Fitting $\\Lambda$CDM model to the observational data considered in the present work, it is easy to find the corresponding best-fit parameter $\\Omega_{m0}=0.278$, whereas $\\chi^2_{min}=466.317$. A conventional criterion for model comparison in the literature is $\\chi^2_{min}/dof$, in which the degree of freedom $dof=N-k$, whereas $N$ and $k$ are the number of data points and the number of free model parameters, respectively. We present the $\\chi^2_{min}/dof$ for all the 7 models in Table~\\ref{tab1}. On the other hand, there are other criterions for model comparison in the literature, such as Bayesian Information Criterion (BIC) and Akaike Information Criterion (AIC). The BIC is defined by~\\cite{r33,r35} \\be{eq39} {\\rm BIC}=-2\\ln{\\cal L}_{max}+k\\ln N\\,, \\ee where ${\\cal L}_{max}$ is the maximum likelihood. In the Gaussian cases, $\\chi^2_{min}=-2\\ln{\\cal L}_{max}$. So, the difference in BIC between two models is given by $\\Delta{\\rm BIC}=\\Delta\\chi^2_{min}+\\Delta k \\ln N$. The AIC is defined by~\\cite{r34,r35} \\be{eq40} {\\rm AIC}=-2\\ln{\\cal L}_{max}+2k\\,. \\ee The difference in AIC between two models is given by $\\Delta{\\rm AIC}=\\Delta\\chi^2_{min}+2\\Delta k$. In Table~\\ref{tab1}, we also present the $\\Delta$BIC and $\\Delta$AIC of all the 7 models considered in this work. Notice that $\\Lambda$CDM has been chosen to be the fiducial model when we calculate $\\Delta$BIC and $\\Delta$AIC. From Table~\\ref{tab1}, it is easy to see that the rank of models is coincident in all the 3 criterions ($\\chi^2_{min}/dof$, BIC and AIC). The $\\Lambda$CDM model is the best one, whereas ICPLw model is the worst one. This result is consistent with the one obtained in e.g.~\\cite{r35}. However, it is well known that $\\Lambda$CDM model is plagued with the cosmological constant problem and the coincidence problem (see e.g.~\\cite{r1}). On the other hand, as mentioned in the beginning of Sec.~\\ref{sec1}, there are some observational evidences for the interaction between dark energy and dark matter, and the coincidence problem can be alleviated in the interacting dark energy models. Therefore, it is still worthwhile to study the interacting dark energy models."
    },
    "1004/1004.0989_arXiv.txt": {
        "abstract": "Recent cosmic ray, gamma ray, and microwave signals observed by Fermi, PAMELA, and WMAP indicate an unexpected primary source of $e^+e^-$ at 10-1000 GeV.  We fit these data to ``standard backgrounds'' plus a new source, assumed to be a separable function of position and energy. For the spatial part, we consider three cases: annihilating dark matter, decaying dark matter, and pulsars.  In each case, we use GALPROP to inject energy in log-spaced energy bins and compute the expected cosmic-ray and photon signals for each bin. We then fit a linear combination of energy bins, plus backgrounds, to the data. We use a non-parametric fit, with no prior constraints on the spectrum except smoothness and non-negativity.  In addition, we consider arbitrary modifications to the energy spectrum of the ``ordinary'' primary source function, fixing its spatial part, finding this alone to be inadequate to explain the PAMELA or WMAP signals.  We explore variations in the fits due to choice of magnetic field, primary electron injection index, spatial profiles, propagation parameters, and fit regularization method. Dark matter annihilation fits well, where our fit finds a mass of $\\sim$1 TeV and a boost factor times energy fraction of $\\sim$70. While it is possible for dark matter decay and pulsars to fit the data, unconventionally high magnetic fields and radiation densities are required near the Galactic Center to counter the relative shallowness of the assumed spatial profiles. We also fit to linear combinations of these three scenarios, though the fit is much less constrained. ",
        "introduction": "Several apparent anomalies in recent astrophysical data hint at a new source of high energy electrons, positrons, and possibly gamma rays, at the 10 GeV to 1 TeV scale.  The cosmic ray signals observed by Fermi \\cite{Latronico:2009uw,Latronico:FermiSymposium,PesceRollins:2009af,Fermi1999} and PAMELA \\cite{Adriani:2008zr} are direct evidence for these energetic electrons and positrons ($\\epp$), which would lose their energy primarily through synchrotron radiation and inverse Compton scattering (IC).  If the number density of these $\\epp$ rises towards the Galactic Center (GC), then this synchrotron and IC could explain the WMAP microwave ``haze\" \\cite{Dobler:2007wv} and the Fermi diffuse gamma ray ``haze\" \\cite{Dobler:2009xz}, respectively.  It is difficult to explain these signals within the conventional diffusive propagation model and with standard assumptions about the interstellar medium (ISM). In this framework, the positron signal arises from secondary production from spallation of proton cosmic rays on the ISM. Assuming that 1. positrons and electrons have the same energy losses, 2. primary electrons and protons have the same production spectrum, and 3. the proton escape time decreases with energy, then the predicted positron fraction generically falls with energy, in contrast to the rising fraction observed by PAMELA.  Katz et al. \\cite{Katz:2009yd} point out these assumptions can be wrong, and explore alternative scenarios. Indeed, secondary production at shock fronts could explain the $e^+$ excess \\cite{Blasi:2009hv, Blasi:2009bd}, but this would also imply an excess of anti-protons, which is not observed. We will not consider these alternatives further.  We examine here whether a new primary source of $\\epp$ is a viable explanation of the signals.  First, the rise in the positron fraction measured by PAMELA suggests the presence of a new hard source of positrons \\cite{Serpico:2008te}. Second, the WMAP ``haze'' is consistent with a hard synchrotron signal in the inner galaxy, in addition to a soft-spectrum synchrotron component traced by Haslam. Though this decomposition is not unique, it is a good fit to the WMAP data. Third, the Fermi gamma ray ``haze'' similarly extends to $|b| > 30^\\circ$ above and below the plane in the inner galaxy. Neither haze correlates with the morphology of any known astrophysical objects or the ISM. (See Fig.~\\ref{fig:haze}.)  Many attempts to explain the data operate by including a new component of high energy particles and gamma rays originating from one of the following sources: \\begin{enumerate} \\item Annihilation of TeV-scale dark matter, \\item Decay of TeV-scale dark matter, or \\item An astrophysical source such as pulsars. \\end{enumerate} These sources can produce energetic electrons, positrons, and gamma rays.  In addition, the dark matter distribution in the Galaxy is expected to be roughly spherical, providing at least qualitative agreement with the morphology of the gamma-ray and microwave haze. Nevertheless, each explanation above has drawbacks.  While annihilating dark matter may seem natural given a weak-scale WIMP which has a thermal freeze-out annihilation cross section, this vanilla scenario cannot explain the observed signals.  Boost factors in the annihilation rate, arising from substructure or particle physics enhancement, of order 100-1000 are typically needed, depending on the annihilation channels and dark matter mass.  Significant model-building effort is also required to explain the lack of excess in the observed $\\bar p/p$ flux \\cite{Adriani:2008zq}.  For examples, see \\cite{Cirelli:2008id,Cirelli:2008pk,ArkaniHamed:2008qn,Cholis:2008wq, Mardon:2009rc}.  In the decaying dark matter scenario, dark matter has the freeze-out annihilation cross section but also decays with lifetime $\\tau_\\chi \\sim 10^{26}$ s. These models also must explain why there is no excess in $\\bar p/p$, though no boost factors are required. Examples include \\cite{Nardi:2008ix, Mardon:2009gw, Arvanitaki:2008hq, Ibarra:2009dr, Ruderman:2009ta, Chen:2008yi, Yin:2008bs}.  The pulsar explanation is the least exotic, but there are significant astrophysical uncertainties in pulsar distributions and $e^+e^-$ emission spectra.  The Fermi cosmic ray signals can be explained by the presence of one or more nearby pulsars with hard $\\epp$ emission spectra \\cite{Hooper:2008kg, Profumo:2008ms, Yuksel:2008rf, Malyshev:2009tw, Gendelev:2010fd}. However, pulsars are generally expected to be concentrated in the disk and it can be difficult to explain the shape of the WMAP and Fermi ``haze\" signals, which are much more spherical.  See also \\cite{Kaplinghat:2009ix, Harding:2009ye}.  In this paper we quantify how well each of these three scenarios can explain the data described above without resorting to model-dependent details of the particle physics or pulsars.  Rather we use the data to determine the best-fit injection spectrum of electrons and positrons produced by each new source. We also show that a simple modification to the background electron injection can explain the Fermi $e^+ + e^-$ spectrum and the Fermi gamma ray spectrum but not the rest of the data.    \\begin{table}[tb] \\begin{center} \\begin{tabular}{|c|c|c|c|} \\hline & $K_0 \\ [\\kpc^2/\\text{Myr}]$  & $\\delta$ & L [kpc] \\\\ \\hline Default\\ & 0.097 & \\ 0.43  \\ & 4 \\\\ M1\t&  0.0765 & \\  0.46 \\  & 15 \\\\ MED & 0.0112 & \\ 0.70  \\ & 4 \\\\ M2 & 0.00595 &  \\ 0.55  \\ & 1  \\\\ \\hline \\end{tabular} \\end{center} \\caption{Typical propagation parameters consistent with low-energy cosmic ray data \\cite{Delahaye:2007fr}.  We use the ``Default\" parameters and show the effect of using M1 and MED in Fig~\\ref{fig:errors}.} \\label{tab:prop} \\end{table}    The standard procedure to analyse whether a model can explain the astrophysical signals is: \\begin{center} pulsar or particle physics model \\\\ $\\Downarrow\t$ \\\\ spectrum of particles produced by the source\\\\ $\\Downarrow\t$ \\\\ propagation (e.g., GALPROP) \\\\ $\\Downarrow\t$ \\\\ comparison with data \\end{center} Often, one fits a specific dark matter or pulsar model to a subset of the ``anomalous'' signals described above. For dark matter, the particle physics model is usually processed through Pythia \\cite{Sjostrand:2006za} to generate a spectrum of $\\epp$.  The injection spectrum is the spectrum of $\\epp$ produced per unit source times the rate of production of $\\epp$ per source and the spatial distribution of the source. These injected $\\epp$ are propagated through the Galaxy to obtain a steady-state solution.  The signals are then compared with data.  While some analyses have studied the cases above in a less model-dependent way, the injection spectrum is assumed to have one of a few common forms \\cite{Cirelli:2008pk,Barger:2009yt,Zhang:2009ut}.    \\begin{figure*}[t] \\begin{center} \\begin{align} \\text{(a)} & \\includegraphics[width=.43\\textwidth]{wmap_haze_23GHz.ps} \\text{(b)} & \\includegraphics[width=.43\\textwidth]{fermi_haze_5_10GeV.ps} \\nn \\\\ \\text{(c)} & \\includegraphics[width=.4\\textwidth]{synch-pulsar-65-stripe.ps} \\hspace{.6cm} \\text{(d)} & \\includegraphics[width=.4\\textwidth]{synch-ann-6522-stripe.ps} \\hspace{.6cm} \\nn \\\\ \\text{(e)} & \\includegraphics[width=.4\\textwidth]{synch-pulsar-45-stripe.ps} \\hspace{.6cm} \\text{(f)} & \\includegraphics[width=.4\\textwidth]{synch-ann-4522-stripe.ps} \\hspace{.6cm}  \\nn \\end{align} \\end{center} \\caption{Maps of the (a) WMAP haze at 23 GHz and (b) Fermi gamma-ray haze at 5-10 GeV for the region $-90<\\ell<90$ and $-45<b<45$. The gamma-ray haze is obtained by subtracting the Fermi diffuse model from the data. All maps are centered on the GC. The data are compared to normalized maps of (c) pulsar synchrotron at 23 GHz and (d) synchrotron at 23 GHz from dark matter annihilation with an Einasto profile. The magnetic field has the form of Eq.~\\ref{eq:Bfield} with $r_B = $ 6.5 kpc. The morphology of the haze signals more closely resembles the signals from dark matter than from pulsars.  We show the corresponding results for $r_B = $ 4.5 kpc in (e) and (f). Choosing a steeper magnetic field can change the morphology, but this is not preferred by the Haslam data (see Fig.~\\ref{fig:haslam}). The region $-5 \\le b \\le 5$ is masked out because we only fit the region $b \\le -5$.} \\label{fig:haze} \\end{figure*}     In this paper we effectively reverse the arrows in the procedure above. We fit the data from Fermi, PAMELA, and WMAP to expected backgrounds plus a new source which produces positrons and electrons. We assume the injection of the new source is separable in position and space. Rather than specifying the spectrum of $e^+e^-$ injected by the new source, we fit for this injection spectrum as a function of energy. The shape of the spectrum is the same everywhere, and the spatial distribution is varied over several conventional models. Therefore, for our purposes, the three scenarios listed above differ only in their spatial distributions.  In the pulsar case, the assumption that the pulsar injection is a separable function of position and energy is a crude approximation that allows us to fit the data without specifying the details of pulsar physics, since the position dependence of pulsar populations and their $\\epp$ injection spectra is very uncertain.  In our fits of the injection spectra, we simultaneously account for possible variations in the conventionally assumed spectrum and spatial distribution of the background injected electrons, as well as propagation parameters, magnetic field, and starlight densities. This takes into account the uncertainties in current models of the Galactic backgrounds.  We describe the signals and their expected backgrounds in more detail in the next section.  We then present the overall framework of the analysis. Predictions are computed using GALPROP, and we allow for variations in the background model.  We then present the best-fit injection spectrum for each of the three scenarios above, as well as the best-fit of the data to an arbitrary modification of the background electron injection spectrum. Finally we present injection spectra for linear combinations of these scenarios. ",
        "conclusions": "We have thoroughly examined the annihilating dark matter, decaying dark matter, and pulsar explanations of the recent anomalous cosmic ray, gamma ray, and synchrotron signals. We investigated whether each scenario can fit all of these data simultaneously. Our analysis is independent of the particle physics or pulsar model details of each scenario and only depends on the spatial profiles and background models. We determined the necessary injection spectrum of electrons and positrons in each case in order to reproduce the data, including the effects of different background models, propagation models, and solar modulation.  Though decaying dark matter is the best fit, the large normalization factors suggest that it will be difficult to find a fully self-consistent model with decaying dark matter that can explain the data, without changing some aspect of our model by a large amount. In particular, it may be necessary to find either a radiation field model with roughly twice as much starlight to produce enough IC, a rather large magnetic field of $33 \\mu G$ in the Galactic center, enormous amounts of low energy electrons or gamma rays injected, a much steeper dark matter profile, or a combination of these.  Pulsars give the worst fit; the disk-like profile makes it nearly impossible to produce both the gamma ray and synchrotron signals. Much like the decaying dark matter case, this suggests that dramatic re-assessments of backgrounds and models are necessary to find a self-consistent interpretation of the data.  Annihilating dark matter, however, has self-consistent fits with conventional astrophysical background models. Though we had to choose a somewhat shallower dark matter halo profile with Einasto $\\alpha = 0.22$, it is still within the current range of profiles found by simulations. Furthermore, we can satisfy the gamma ray constraints from the GC. The boost factors are $\\sim 70/f_E$, which at first seems much lower than the boost factors of $\\sim 1000$ often used in the literature. Several factors enter in this difference: our use of the updated $\\rho_0 = .4$ GeV/cm$^3$ rather than $\\rho_0 = .3$ GeV/cm$^3$ \\cite{Catena:2009mf}, the relatively hard spectrum allowed by the fit, and our assumption that the cutoff of the spectrum is $m_\\chi$.  Given these factors, our result of $\\sim 70/f_E$ is typical of the models discussed in the introduction of this paper.  However, the shape of the spectrum, combined with a lack of $\\pi^0$ or $\\bar p$ production, may still be difficult to achieve in current particle physics models of dark matter.  Our results should be regarded as qualitative guidelines to injection spectra. The specifics will necessarily change as both Galactic models and data are refined. The WMAP ``haze'' data will be superseded by data from Planck \\cite{Planck:2006uk}, while data from Fermi and PAMELA will improve. In addition, cosmic ray data from AMS-02 \\cite{AMS02} may also soon be available.  If the data does not change substantially, and if current models indeed describe Galactic propagation and interactions, then the qualitative results of this paper will remain valid."
    },
    "1004/1004.5576_arXiv.txt": {
        "abstract": "Over the last few years, needlets have emerged as a useful tool for the analysis of Cosmic Microwave Background (CMB) data. Our aim in this paper is first to introduce in the CMB literature a different form of needlets, known as Mexican needlets, first discussed in the mathematical literature by Geller and Mayeli (2009a,b). We then proceed with an extensive study of the properties of both standard and  Mexican needlets; these properties depend on some parameters which can be tuned in order to optimize the performance for a given application. Our second aim in this paper is then to give practical advice on how to adjust these parameters for WMAP and Planck data in order to achieve the best properties for a given problem in CMB data analysis. In particular we investigate localization properties in real and harmonic space and propose a recipe on how to quantify the influence of galactic and point source masks on the needlet coefficients. We also show that for certain parameter values, the Mexican needlets provide a close approximation to the Spherical Mexican Hat Wavelets (whence their name), with some advantages concerning their numerical implementation and the derivation of their statistical properties. ",
        "introduction": "Over the last decade, wavelet systems have grown as one of the most important tools in the analysis of Cosmological and Astrophysical data. A lot of proposals for wavelet systems on the sphere have been advanced in the mathematical literature, see for instance \\cite{freeden}, \\cite{jfaa1}, \\cite% {dahlke}, \\cite{poisson}, \\cite{rosca}, \\cite{jfaa3}, \\cite{starck} and the references therein. Some of these attempts have been explicitly motivated by Astronomy and/or Cosmology (see for instance \\cite{jfaa2} for a review). In particular in the area of Cosmic Microwave Background (CMB) data analysis, wavelets have been used for a large number of applications (see references in the next paragraph). The interest for wavelets in this area is very easily understood; predictions from CMB theory are typically cast in the Fourier domain, however exact Fourier analysis cannot be entertained because of the presence of foreground and masked regions. The double-localization properties of wavelet systems (in real and harmonic domain) hence turn out to be most valuable. Moreover, addressing important issues such as the possible existence of features and asymmetries in CMB maps is nearly unfeasible without ideas which are broadly related to the wavelet literature.  Among spherical wavelets, particular attention has been recently devoted to so-called needlets, which were introduced into the Functional Analysis literature by \\cite{npw1,npw2}; their statistical properties were first considered by \\cite{bkmpb,bkmp}. Needlets enjoy several properties that make them worth of attention for Cosmological data analysis. In particular, they are computationally very simple, and naturally adapted to standard packages such as HealPix \\citep{healpix}; they do not require any form of tangent plane approximation, but they are naturally embedded into the manifold structure of the sphere; they are compactly supported in the harmonic domain, i.e. they depend only on a finite number of multipoles which are explicitly known and can be controlled by the data analysts; they are quasi-exponentially localized in real space, i.e. their tails decay faster than any polynomial; and finally, it has been shown in \\cite{bkmpb} that random needlet coefficients enjoy a very useful uncorrelation property: namely, for any fixed angular distance, random needlets coefficients are asymptotically uncorrelated as the frequency parameter grows larger and larger. As well-known, uncorrelation entails independence in the Gaussian case: as a consequence, from the above-mentioned property it follows that needlet coefficients from a CMB map can be seen as nearly independent at high frequencies, making thus possible the introduction of a variety of statistical procedures for testing non-Gaussianity, estimating the angular power spectrum, testing for asymmetries, implementing bootstrap techniques, testing for cross-correlation among CMB\\ and Large Scale Structure data, and many others, see for instance \\cite{bkmpb,bkmp}, \\cite{guilloux}, \\cite{ejslan}, % \\cite{pbm}, \\cite{mpbb08}, \\cite{dela08}, \\cite{rudjord1,rudjord2}, \\cite{cap}. More recently, the needlet construction has also been extended to the case of spin/polarization data, see for instance \\cite{ghmkp,glm}, \\cite{gm1} and \\cite{gelmar}.  The first purpose of this paper is to introduce a new kind of needlets to the field of CMB analysis following an approach which has been very recently advocated in mathematics by \\cite{gm4,gm2,gm3}. This approach (which we shall discuss in Section 2) can be labeled Mexican needlets. As we shall discuss below, a special case of the Mexican needlets provides at high frequencies a close approximation to the widely used Spherical Mexican Hat Wavelet (SMHW, see for instance \\cite{cayon,vielva}), with some advantages in terms of their numerical implementation and the investigation of their localization and statistical properties. As such, the investigation of their properties in this case will allow to understand the stochastic properties of SMHW and compare them with standard needlets. Mexican needlets depend on a parameter $p$, and we shall show how this parameter can be tuned to improve the localization properties in the real or in the harmonic domain.  The second purpose of this paper is to provide a practical description of properties of different needlet types important for CMB analysis. Proper knowledge of the localization properties on the pixelized sphere as well as in multipole space is crucial for selecting and applying the proper type of needlet to a specific problem. Although the exact mathematical properties of the needlets are well known, we have studied the properties which are of high importance for CMB analysis and which are too complicated to be easily deduced from the mathematical results. In particular, in the presence of foregrounds and masks, it is important to know their influence on the needlet coefficients. For the SMHW it has been shown for several applications \\citep{vielva,masks3,masks4,masks2} that an extended scale dependent mask must be used when analyzing masked data with wavelets. Here we will study this in detail for the different needlet types with different parameter values.  We then provide a very thorough comparison between different needlets. The previous discussion leads naturally to the issue about their optimal construction, i.e. how to devise numerical recipes which will enhance their localization properties. Here, we shall compare the numerical recipe implemented by \\cite{bkmpb,pbm} with an alternative proposal based on Bernstein polynomials (see also \\cite{guilloux} for a related numerical investigation). The latter entail weight functions with just a finite number of bounded derivatives; to distinguish it from the previous construction we will label this procedure Bernstein needlets. We stress, however, that the underlying mathematical theory presents no real novelty as compared to the results by \\cite{npw1,npw2}. We then go on to provided vast numerical evidence on the various forms of localization, by means of a number of different indicators. In particular, the role of the different parameters in the determination of the various properties is fully exploited.  As well-known, there is usually a trade-off between localization properties in the frequency and real domains, as a consequence of the Uncertainty Principle (\\textquotedblleft It is impossible for a non-zero function and its Fourier transform to be simultaneously very small\\textquotedblright , see for instance \\cite{narco,havin}); the main purpose of this paper is to show how Mexican and standard needlets jointly provide a flexible set of tools where each user can optimize this trade-off according to the needs of a specific data analysis problem.  The plan of the paper is as follows: in Section \\ref{sec:intro_on_need}, we review the standard needlet construction, introduce Mexican needlets and compare their respective properties from a mathematical point of view. In the following sections, we provide numerical evidence on the localization properties of these procedures by means of several different figures of merit; we discuss at length the interplay among the different properties and the trade-off to face when choosing which procedure to adopt for a given Astrophysical problem. In section \\ref{sec:summary}, we summarize the main indicators used and the properties of the different needlets measured in terms of these indicators. In the Appendix, we provide some details on the numerical recipes we adopted, some short discussion on mathematical properties and analytic fits. ",
        "conclusions": "We have performed the following set of tests of localization properties of needlets when used with CMB data: \\begin{itemize} \\item The distance of influence from a galactic cut on needlet coefficients was found from simulations and a range from conservative (allowing no influence) to  less conservative 'distances of safety' (expressed by the critical angle $\\theta_{crit}$) from the mask. We express the resulting distances (critical angles $\\theta_{crit}$) as functional fits of the form $\\beta/\\ell^*$ and have obtained analytical fits for $\\beta$ for standard and Mexican needlets (see expressions \\ref{eq:betaB} and \\ref{eq:beta_mex}). Results for some values are shown in figure \\ref{fig:galfits} (standard and Mexican) and listed in tables \\ref{tab:beta_mex_gal} (Mexican) and \\ref{tab:beta_bern} (Bernstein).  \\item In a similar manner, the 'distance of safety' from a point source hole (one WMAP-like, one Planck-like) was found; analytical fits for the resulting critical angles  (see again expressions \\ref{eq:beta_a0_fit}, \\ref{eq:betaB} and \\ref{eq:beta_mex}) were obtained. Some values are shown in figures \\ref{fig:wmapfits} (WMAP-hole) and \\ref{fig:planckfits} (Planck-hole) and listed in table \\ref{tab:beta_mex_hol} (Mexican) and \\ref{tab:beta_bern_hol} (Bernstein).  \\item We have then also shown that our idealized masks and the results obtained from there can be used and are valid for more general and realistic masks. In particular, the results are found valid both for the WMAP KQ85 mask and for a smaller Planck-like galactic mask. The same holds for point source holes: The results obtained for one hole is shown to be correct also for point source masks with many holes.  \\item We have also studied the correlation distance of needlet coefficients on the sphere. We use an analytical formula which we have found to agree well with simulations (equation \\ref{eq:cor}) to find the degree of correlation as a function of distance between coefficients. Results for some distances are shown in figure \\ref{fig:corr_4th}. \\item We have defined a measure $\\sigma_h$ (see equation \\ref{eq:sigharm2}) of localization in harmonic space. This characterizes the number of multipoles contained in one single needlet scale $j$ and is found to be proportional to $\\ell^*$. As a result (remembering that $\\theta_{crit}$ is proportional to $1/\\ell^*$ and that $\\theta_{crit}$ depends heavily on real space localization), we find that, as expected, $\\theta_{crit}\\times\\sigma_\\mathrm{pix}$ is roughly a constant such that when the localization in real space is improved, the localization in harmonic space worsens and vice versa. \\end{itemize}  For a certain type of needlet, one can thus obtain the critical angle by (1) calculating $\\ell^*$ for the given scale $j$, using equation \\ref{eq:lstar}, (2a) reading the critical angle off figures \\ref{fig:galfits}, \\ref{fig:wmapfits} and \\ref{fig:planckfits} OR (2b) use the analytic fits presented in appendix \\ref{app:galacticbeta} and \\ref{app:holebeta}.  Our results for the different types of needlets can be summarized as follows: \\begin{itemize} \\item {\\bf Standard needlets:} The standard needlets are much better localized in harmonic space than the Mexican needlets (but similar to Bernstein needlets). In fact, the contribution to a certain scale $j$ is coming from a limited number of multipoles with no influence from multipoles outside this range. However, the penalty for high localization in harmonic space is that the real space localization is lower than for other needlets. The parameter $B$ controls the localization properties: The higher the $B$, the higher is the real space localization and the larger are the multipole ranges included in each scale $j$. \\item {\\bf Mexican needlets:} The Mexican needlets are much better localized in real space than any of the other types. In harmonic space however, a large (in principle infinite) multipole range contributes to each scale $j$. The Mexican needlets depend on the parameter $p$, larger $p$ have worse real space localization properties but better harmonic space localization. For $p=3-4$, the real space localization properties of the Mexican needlets approach those of standard needlets for high values of $B$. For $p=1$ and high $j$, the Mexican needlets are almost identical to the Spherical Mexican Hat Wavelets. \\item {\\bf Bernstein needlets:} Bernstein needlets are in between the Mexican and standard needlets. Their real and harmonic space localization is similar to the standard needlets, sometimes it is slightly better, sometimes it is slightly worse. The Bernstein needlets depend both on the parameter $B$ and (weakly) on a parameter $k$. As for the standard needlets, a higher value of $B$ increases the number of multipoles included in each scale $j$ and improves the real space localization. The parameter $k$ may slightly improve or worsen the localization properties depending on the exact measure used. As an example we find that for the critical angle for a galactic cut, for $k\\ge2$ the angle for the less stringent thresholds increases for increasing $k$, whereas the angle for the more stringent thresholds decreases for increasing $k$. \\end{itemize}"
    },
    "1004/1004.2998_arXiv.txt": {
        "abstract": "In the \\swift era, two-component jet models were introduced to explain the complex temporal profiles and the diversity of early afterglows. In this paper, we concentrate on the two-component jet model; first component is the conventional afterglow and second is the emission due to the late internal dissipation such as the late-prompt emission. We suggest herein that the two-component jet model can be probed by the existence of two optical peaks for orphan GRB afterglows. Each peak is caused by its respective jet as its relativistic beaming cone widens to encompass the off-axis line of sight. Typically, the first peak appears at $10^4-10^5$ s and the second at $10^5-10^6$ s. Furthermore, we expect to observe a single, bright X-ray peak at the same time as the first optical peak. Because orphan afterglows do not have prompt emission, it is necessary to monitor the entire sky every $10^4$ s in the X-ray regime. We can test the model with orphan afterglows through the X-ray all-sky survey collaboration and by using ground-based optical telescopes. ",
        "introduction": "The complex temporal profiles of early X-ray afterglows discovered by the \\swift on-board X-ray telescope (XRT) were unexpected in the pre-\\swift era. Explaining these results is one of the most difficult problems in the field of gamma-ray bursts (GRBs) (Evans et al. 2009). Most X-ray afterglows have the following three phases: steep decay, plateau, and normal decay. During the steep-decay phase, the X-ray afterglow decays with a slope of approximately $\\sim -3$ or steeper: this phase extends up to $10^2-10^3$ s. The most popular interpretation of this phase attributes it to the tail of the prompt emission (Kumar \\& Panaitescu 2000; Zhang et al. 2006; Yamazaki et al. 2006). The plateau phase follows the steep-decay phase and begins with a  slope of approximately $\\sim -0.5$ or flatter and covers the time range $10^3-10^4$ s, during which a temporal break occurs. The normal decay phase then begins; the afterglow breaks and decays with a slope of approximately $\\sim -1.2$, which is in agreement with the predictions of the afterglow model (M{\\'e}sz{\\'a}ros \\& Rees 1997; Sari et al. 1998).  However, two major problems exist with this interpretation of X-ray afterglows. The first is that the plateau phase is not completely explained by the external-shock models. To address this discrepancies, various models have been proposed, such as the energy-injection model (Nousek et al. 2006; Zhang et al. 2006; Granot \\& Kumar 2006), the inhomogeneous or two-component jet model (Toma et al. 2006; Eichler \\& Granot 2006; Granot et al. 2006), the time-dependent microphysics model (Ioka et al. 2006; Granot et al. 2006; Fan \\& Piran 2006), and the prior activity model (Ioka et al. 2006; Yamazaki 2009). The second problem is that the conventional afterglow models (M{\\'e}sz{\\'a}ros \\& Rees 1997; Sari et al. 1998) cannot explain the difference between the epoch of the breaks in light curves in the X-ray and optical observations.  Ghisellini et al. (2009) reported that the observed light curves are well explained by a two-component (early and late) jet model. Early component is responsible for the conventional afterglow emission and late component for the plateau phase. When the latter dominates the former, the X-ray light curve exhibits a plateau phase. However, there is little evidence that the early and late jet models are correct and no way to distinguish these models from other models such as the energy-injection model or the time-dependent micro-physics model. We show that observations of orphan afterglows are crucial for testing a part of the early and late jet models such as the late-prompt emission model (Ghisellini et al. 2007).  In this study, we calculate the light curves of orphan afterglows for the late-prompt emission model. We find that the two components can be distinguished in the optical regime by observing from off-axis. Therefore, if the orphan afterglows present two optical peaks, this fact becomes a strong evidence that GRB afterglows consist of two components. Moreover, we estimate the detection rate of the orphan afterglows on the basis of the late-prompt emission model and find that it would be possible to detect them and thereby test this model.  This article is organized as follows: In Section 2, we explain the motivation for considering the two-component (early and late) jet models by reviewing the results of Ghisellini et al. (2009) and introduce the late-prompt emission model. In Section 3, we describe the numerical method to calculate the off-axis light curves. The results and their qualitative interpretations are presented in Section 4. Finally, in Section 5, we summarize and discuss the possibility of observing orphan afterglows. ",
        "conclusions": "In this study, we consider the late-prompt emission model and calculate the off-axis observed light curves. Then, we find that the off-axis observer is able to see the emission with two peaks in the optical (Figure 4) and with a single peak in the X-ray (Figure 3). We here suggest a method to test whether two components actually exist or not by applying the characteristics of off-axis light curves.  In this section, we discuss the possibility of observing the emission from off-axis. As seen in Section 4, in the late-prompt model, the single peak in the X-ray regime is roughly as bright as the flux amplitude at the transition time $T_a$. Then, we expect that the emission with a single peak in the X-ray is observable. We can easily estimate the detection rate of the X-ray flux viewed off-axis for the late-prompt model by using the same samples as in Ghisellini et al. (2009). Assuming that the properties of GRBs are not affected by the accuracy with which the redshift is determined, we can extrapolate the flux distribution of all \\swift bursts from the above samples. Here, we suppose that the average jet angle of the GRBs is 5$^\\circ $. Next, the real number $N_{\\rm{real}}(F)$ of GRBs per year for flux $F$ is not the number $N_{\\rm{obs}}(F)$ of GRBs inferred from observations, but $2N_{\\rm{obs}}(F)/(1-\\cos \\theta _j)$. From our calculation, we find the relationship between the viewing angle $\\theta _v$ and the peak flux for the late-prompt emission to be $F_{\\rm{peak}}(\\theta _v)\\propto \\theta _v ^{-2.09}$ ($\\theta _v >\\theta _j$). The detection rate $N_{\\rm{exp}}$ is then estimated as \\begin{equation} N_{\\rm{exp}}=\\int \\frac{d\\phi d\\cos \\theta _v}{4\\pi } N_{\\rm{real}}(F_{\\rm{peak}}(\\theta _v)>F_{\\rm{lim}}), \\end{equation} where $F_{\\rm{lim}}$ is the instrument sensitivity.  Because the GRB observed off-axis do not have the prompt emission, the all-sky survey that monitors the whole sky is necessary to detect the single peak in the X-ray. For example, such peak with duration time of $10^4$ s can be detected by the current X-ray survey missions with Monitor of All-sky X-ray Image (MAXI) (Matsuoka et al. 2009). The sensitivity of MAXI is 20 mCrab ($7\\times 10^{-10}$ erg~cm$^{-2}$~s$^{-1}$ over the energy band  $2-30$ keV) for observations during a single orbit. Because MAXI observes a point of sky every $90$ min, it is suitable for detecting the peak. In this case, the detection rate by MAXI is estimated to be $\\sim 0.41$ events per year. As seen Figure 3, the peak amplitudes of the orphan late-prompt emission are a little lower than the MAXI sensitivity. In case of  the orphan late-prompt emission, however, the number of GRBs whose relativistic beaming cones encompass the observer increases by a factor of $(1-\\cos \\theta _v)/(1-\\cos \\theta _j)$. The event rate we estimate is as much as the event rate that MAXI detects on-axis afterglows without prompt emissions. This forecast is an important outcome of the late-prompt model. Because the angular resolution of MAXI is about 0.1 arcminute, we can follow up the afterglows with ground-based optical telescopes (e.g., Gamma-Ray Burst Optical/Near-Infrared Detector; GROND) and observe emission with two peaks. Typical amplitudes of the first peak are about $24$ to $25$ in R-band magnitude. Because this sensitivity is of the same order as that of GROND (Greiner et al. 2008), it should be possible to observe the first optical peak. Although the second peak is slightly dimmer than the first, we consider that after a few days, the second peak can be observed by GROND or the larger optical telescopes. Seven bands of GROND observation help to distinguish the afterglows from other variables. The second peak, which has a synchrotron spectrum because of the external shock, can be distinguished from supernovae.  In this study, we calculate the off-axis light curves in the late-prompt model by integrating the analytical expressions. However, in recent study, van Eerten, Zhang \\& MacFadyen (2010) show that the off-axis light curves calculated analytically are different from the results with a two-dimensional axisymmetric hydrodynamics simulation. Then, it is significant to calculate the off-axis late-prompt emission in order to discuss whether two peaks are distinguishable. Furthermore, the late-prompt emission model assumes the radius $R$ to be constant. If the radius expands of the emission area, the peak time $T_{pL}$ for the late-prompt emission more delays in the same way as the peak time $T_p$ for the early emission as the viewing angle increases. In this case, it may be, therefore, difficult to distinguish the optical two peaks in the optical with the observation of orphan afterglows.  Finally, we mention other late-internal dissipation models. In these models, the transition time $T_a$ does not necessarily correspond to the jet break but to the time scale for the matter accretion (Kumar et al. 2008) or the spin-down time scale for the magnetar scenarios (Zhang \\& M{\\'e}sz{\\'a}ros 2001; Thompson et al. 2004). If this is the case, then the Lorentz factor evolution history of the late outflow is unpredictable, and the predicted orphan afterglow light curves would be subject to even larger uncertainties. In future works, we should discuss other possibilities for testing two-component jet models. Nevertheless, if we observe a characteristic light curve, which has a bright single peak in the X-ray and two peaks in the optical, we obtain an evidence that GRBs light curves are composed of the early and late jet, and may restrict other models than the early and late jet model.    \\bigskip We would like to thank T. Nakamura for his continuous encouragement and K. Murase, K. Yagi, M. Hayashida, S. Inoue, Y. Inoue, and Y. Doleman for useful discussions. This work is supported in part by the Grant-in-Aid for the global COE program {\\it The Next Generation of Physics, Spun from Universality and Emergence} at Kyoto University. RT is supported by a Grant-in-Aid for the Japan Society for the Promotion of Science (JSPS) Fellows and is a research fellow of JSPS."
    },
    "1004/1004.1739_arXiv.txt": {
        "abstract": "We study the magnetosphere of a slowly rotating magnetized neutron star subject to toroidal oscillations in the relativistic regime. Under the assumption of a zero inclination angle between the magnetic moment and the angular momentum of the star, we analyze the Goldreich-Julian charge density and derive a second-order differential equation for the electrostatic potential. The analytical solution of this equation in the polar cap region of the magnetosphere shows the modification induced by stellar toroidal oscillations on the accelerating electric field and on the charge density. We also find that, after decomposing the oscillation velocity in terms of spherical harmonics, the first few modes with  $m=0,1$ are responsible for energy losses that are almost linearly dependent on the amplitude of the oscillation and that, for the mode $(l,m)=(2,1)$, can be a factor $\\sim8$ larger than the rotational energy losses, even for a velocity oscillation amplitude at the star surface as small as $\\eta=0.05 \\ \\Omega \\ R$. The results obtained in this paper clarify the extent to which stellar oscillations are reflected in the time variation of the physical properties at the surface of the rotating neutron star, mainly by showing the existence of a relation between $P\\dot{P}$ and the oscillation amplitude. Finally, we propose a qualitative model for the explanation of the phenomenology of intermittent pulsars in terms of stellar oscillations that are periodically excited by star glitches. ",
        "introduction": "The theoretical study of radio pulsars dates back to the work of \\cite{Goldreich1969} who first suggested the existence of a magnetosphere with a charge-separated plasma around rotating magnetized neutron stars. A spinning magnetized neutron star generates huge potential differences between different parts of its surface. The cascade generation of electron-positron plasma in the polar cap region already proposed by \\citet{Sturrock1971} and \\citet{Ruderman1975} requires that the magnetosphere of a neutron star is filled with plasma, thus screening the longitudinal electric field and bringing the plasma into co-rotation with the neutron star. Because co-rotation is not possible outside the light cylinder (the radius $R_{LC}=c/\\Omega$ at which the co-rotation speed equals the speed of light), essentially two different topologies of the magnetic field lines are naturally produced: closed lines, namely those returning to the stellar surface, and open lines, i.e. those crossing the light cylinder and going to infinity. As a result, plasma may leave the neutron star along the open field lines and it is generally thought that pulsar radio emission is produced in the region of open field lines well inside the light cylinder and within a given angle from the polar axis.  Beside the seminal papers by   \\cite{Goldreich1969}, \\cite{Sturrock1971},  \\cite{Ruderman1975}, \\cite{Mestel1971} and \\cite{Arons1979}, pulsar magnetospheres have been investigated by a large number of authors over the years. We only mention here the reviews by \\cite{Arons1991}, \\cite{Mestel1992} and \\cite{Muslimov1997}, where subsequent achievements and some new ideas have been presented. Thorough description of known magnetosphere properties may be found, for example, in the book of \\cite{Beskin2009}. It should also be mentioned that in the last few years time dependent numerical simulations of neutron star magnetospheres have been proposed as a new promising tool for investigating the complex physics of these systems. At least qualitatively, the numerical approach has confirmed the most fundamental features of what was expected from the stationary solution of the Grad-Shavranov equation~\\citep{Contopoulos1999, Gruzinov2005}, such as the existence of closed magnetic field lines up to the light cylinder~\\citep{Komissarov2006, McKinney2006a}, or the scaling of the spin down luminosity on the angular velocity and on the inclination angle of the neutron star angular momentum with respect to its magnetic moment~\\citep{Spitkovsky2006}. In spite of this spectacular progress, however, the numerical approach still suffers from some serious limitations, such as the lack of a unified scheme in which both the force free regime and the plasma regime of magnetohydrodynamics are simultaneously taken into account, or the lack of a consistent treatment of resistive effects in the current sheet.  The analytic approach, on the other hand, can still provide a deep understanding of pulsar physics. In particular, a lot of attention has  been paid to the existence of a strong electric field induced by the rotation of the star, as already noticed by \\cite{Deutsch1955}. More recently,~\\cite{Beskin1990} and, independently, \\cite{Muslimov1990} were the first to find that the frame dragging induced by general relativistic effects provides a source of additional electric field contributing to particle acceleration in the polar cap region. The accelerating component (parallel to the magnetic field) of the  electric field is driven by deviations of the space density charge from the Goldreich-Julian (GJ) charge density, which is determined by the magnetic field geometry. As noted by several authors ~\\citep{Beskin1990, Muslimov1990, Muslimov1997, Dyks2001, Mofiz2000, Morozova2008}, the corrections of general relativity in the plasma magnetosphere of rotating neutron stars are first-order in the angular velocity of the dragging of inertial frames and have to be carefully included in any self-consistent model of pulsar magnetosphere,  especially when computing the resulting electromagnetic radiation.  Tightly related to this aspect is the possibility that neutron star oscillations, most likely excited during a glitch phenomenon (sudden change of the rotational period), propagate into the magnetosphere, thus affecting the acceleration properties in the polar cap region. The first attempt to generalize the Goldreich-Julian formalism to the case of an oscillating neutron star was made by~\\cite{Timokhin2000}, who developed a general procedure for calculating the GJ charge density in the near zone of an oscillating neutron star. Using this procedure, the GJ charge density and the electromagnetic energy losses were computed for the case of toroidal oscillations at the neutron star surface. A similar approach has been recently extended to the general relativistic context by~\\cite{Abdikamalov2009} who, just like ~\\cite{Timokhin2000}, based their results on the so called low current density approximation, i.e. on the assumption that the magnetic field is mainly produced by volume currents inside the neutron star and by surface currents on its surface, while the magnetic field due to magnetospheric currents can be neglected. In the paper of~\\cite{Abdikamalov2009} the influence of oscillations to the magnetosphere electrodynamics was considered for the case of a non-rotating Schwarzschild star. In the present paper we apply some of the results of~\\cite{Abdikamalov2009} to investigate how oscillations, produced at the star surface, reflect in the energy losses from the polar cap region of the magnetosphere of slowly rotating neutron star. In this respect we extend the work of \\cite{Muslimov1997} by performing a local analysis in the domain of open magnetic field lines in the inner magnetosphere and taking into account the effects of toroidal oscillations excited at the star surface.  The plan of the paper is as follows. In Sec.~\\ref{Goldreich-Julian_relativistic_charge_density} we provide the minimum general relativistic formalism for understanding neutron star electrodynamics and we perform a detailed analysis of the GJ charge density of slowly rotating and oscillating neutron star. In Sec.~\\ref{Poisson_equation} we derive a version of the Poisson equation that takes into account both general relativistic effects and the oscillating behavior of the magnetosphere of the rotating star. Sec.~\\ref{Energy_losses}, on the other hand, is devoted to the computation of the energy losses induced by oscillations together with rotation. In Sec.~\\ref{Connection_to_the_phenomenology_of_part-time_pulsars} we propose and motivate a suggestive idea to explain the phenomenology of intermittent pulsars in terms of the excitation of stellar oscillations. Finally, Sec.~\\ref{concl} contains the conclusions of our work.  Throughout, we assume a signature $\\{-,+,+,+\\}$ for the space-time metric and we use Greek letters (running from $0$ to $3$) for four-dimensional space-time tensor components, while Latin letters (running from $1$ to $3$) will be employed for three-dimensional spatial tensor components. Moreover, we set $c = G = 1$ (however, for those expressions with an astrophysical application we have written the speed of light explicitly). ",
        "conclusions": "\\label{concl}  In this paper we have studied the astrophysical processes in the polar cap region of the magnetosphere of an oscillating neutron star. The background spacetime is given by the metric of ~\\citet{Hartle1968} within the slow rotation approximation. The novelties of our analysis consists in quantifying the contributions of stellar oscillations in a general relativistic framework. In particular, we have computed the general-relativistic corrections to the Goldreich-Julian charge density, to the electrostatic scalar potential and to the component of the electric field parallel to the magnetic field lines in the polar cap region when toroidal stellar oscillations are present. As already remarked by~\\citet{Timokhin2007}, the effective electric charge density \\ie the difference between the Goldreich-Julian charge density $\\rho_{\\rm{GJ}}$ (proportional to $\\vec{\\Omega}\\cdot \\vec{B}$ for rotating stars and to $\\vec{\\omega}\\cdot \\vec{B}$ for oscillating stars in the flat space-time case) and the electric charge density (proportional to $\\vec{B}$) in the oscillating star magnetosphere is responsible for the generation of an electric field parallel to the magnetic field lines. Such difference vanishes only at the surface of the star while in general it becomes significantly large at some distance $r$ from the surface, due to the fact that $\\rho$ can not compensate $\\rho_{\\rm{GJ}}$. As already pointed out by~\\citet{Muslimov1992}, general relativistic terms arising from the dragging of inertial frames give very important additional contribution to this difference. These terms depend on the radial distance from the star as $1/r^3$ and have important influence on the value of accelerating electric field generated in the magnetosphere near the surface of the neutron star.  Our solutions for the accelerating electric field for the oscillating and rotating magnetized neutron stars may have some significant implications for pulsar polar cap models. These models assume that charged particles are accelerated above the polar caps, initiating pair cascades through one-photon pair creation of photons.  The electric field induced by the stellar oscillations becomes therefore very important. Thus, the potential drop at the pair formation front, and the total energy gained by particles in the open field region is larger for the oscillating star. Since the contribution from the stellar oscillations to the electric field depends on the amplitude of stellar oscillations, pulsars having larger $K=\\tilde{\\eta}(1)/\\Omega R$ will have larger accelerating potential drops. Our main conclusions about oscillating and rotating neutron stars can be summarized as follows:  \\begin{enumerate}  \\item The oscillation regime of particle ejection from the stellar surface increases the total power carried away by relativistic primary particles relative to the purely rotating regime. Moreover, the fluctuation of the charge density of particles ejected from the stellar surface modulates the particle energy along a field line.  \\item The energy losses along the open magnetic field lines in the polar cap region and due to toroidal oscillations are significantly larger than the rotational energy losses for the $m=1$ modes of oscillation. In particular, the energy losses of the mode $(l,m)=(2,1)$ can be a factor $8$ larger than the rotational energy losses, even for an oscillation amplitude at the star surface as small as $\\eta=0.05 \\ \\Omega \\ R$.   \\item The oscillation-induced inhomogeneity of the physical conditions at the stellar surface may substantially affect the global electrodynamics within the inner magnetosphere of a neutron star.  \\item The new dependence obtained for the energy losses on the oscillating behavior reflects in a new relation, namely Eq.~\\eqref{PP}, between the product $P\\dot{P}$ and the amplitude of the oscillation at the star surface. In cases when the moment of inertia of the star is known with good accuracy, such a relation will allow to fully appreciate the effects of oscillations on pulsar magnetospheres.    \\end{enumerate}  Finally, we have proposed a connection between the phenomenology of intermittent pulsars, characterized by the periodic transition from active to dead periods of radio emission in few observed sources, with the presence of an oscillating magnetosphere. In particular, we propose that, during the active state, star oscillations induced by periodic glitches of the neutron star create relativistic wind of charged particles by virtue of the additional accelerating electric field. After a timescale of the order of tens of days stellar oscillations are damped, and the pulsar shifts below the death line in the  $P - B$ diagram, thus entering the OFF invisible state of intermittent pulsars. This seminal idea, proposed here on a qualitative level, will be further explored in a future work."
    },
    "1004/1004.2789_arXiv.txt": {
        "abstract": "High-resolution spectroscopy in the near-infrared could become the leading method for discovering extra-solar planets around very low-mass stars and brown dwarfs. To help to achieve an accuracy of $\\sim$\\,m/s, we are developing a gas cell which consists of a mixture of gases whose absorption spectral lines span all over the near-infrared region. We present the most promising mixture, made of acetylene, nitrous oxide, ammonia, chloromethans and hydrocarbons. The mixture is contained in a small size 13\\,cm long gas cell and covers most of the H and K-bands. It also shows small absorptions in the J-band but they are few and not sharp enough for near infrared wavelength calibration. We describe the working method and experiments and compare our results with the state of the art for near infrared gas cells. ",
        "introduction": "In the last few years, a new generation of near infrared (NIR) spectrographs with high spectral resolution and radial velocity accuracy ($\\sim$ m/s) for exoplanet detection is under development \\citep{2004SPIE.5492.1274O,2005AN....326.1015M,2008PASP..120..887R}. For this purpose, a high precision calibration system with a stable and very high number of lines spanning all over the instrumental wavelength range is mandatory.  Available Thorium-Argon (Th-Ar) emission lamps \\citep{2007A&A...468.1115L} provide a good coverage in the optical. However, there is a lack of such lines in the near infrared, and those available are relatively faint and unstable (primary Ar, as they are the dominant brightness lines in the $1-1.8 \\mu$m region; Wahling et al. 2002). Laser frequency combs have the advantage of generating series of equally spaced very narrow lines \\citep{2007MNRAS.380..839M,2008Sci...321.1335S} but stability and repeatability in long timescales has not been achieved yet. \\\\ In the optical regime, iodine gas cell has proved to be a very good method for simultaneous calibration of echelle spectra \\citep{1996PASP..108..500B}. This method has several advantages compared to others. The cell is located along the stellar beam in front of the spectrograph. Thus, the stellar spectrum is superimposed to the absorption spectrum for simultaneous calibration so that unstabilities and differences in illumination of the spectrograph can be measured and modelled to remove these effects. It is also very cheap (specially compared to laser combs) and of easy implementation and maintenance but since there are no pure gases showing a wide spectral domain in the near infrared, as iodine in the optical, the main problem is to get a suitable gas cell with strong and well distributed absorption lines, stable with time and temperature.\\\\ \\cite{2009ApJ...692.1590M} have considered series of commercially available absorption cells of H$^{13}$C$^{14}$N, $^{12}$C${_2}$H${_2}$, $^{12}$CO and $^{13}$CO for the H-band and \\cite{2008SPIE.7014E.126D} have simulated a mixture of HCl, HBr and HI gas cell for the GIANO spectrograph in the 0.9--2.5\\,$\\micron$ range using the HITRAN database \\citep{2005JQSRT..96..139R}.  Also a couple of instruments are already using absorption gas cells that cover small parts of the near infrared spectrum. A N$_2$O gas cell \\citep{2007ASPC..364..461K} has been developed for the CRIRES spectrograph.  Also, recently, an ammonia gas cell has been used by \\cite{2009arXiv0911.3148B} in the K-band, and they have proven that such a cell can achieve precisions of $\\sim$ 5 m/s over long timescales with CRIRES, which proves the feasibility of this technique also in the near infrared. \\cite{2003A&A...403.1077K} reported a radial velocity precision of 2.65\\,m/s for Barnard's star (spectral type M4) and they found that the radial velocity is limited by stellar noise (activity and convection). \\\\   Over the last years we have been working on a gas cell for simultaneous calibration of spectra from the Y, J, H and K-bands. In this paper we present the experimental aproach and the promising results obtained so far. The paper is organized as follows. In section~\\ref{2}, we describe the gas cell properties and laboratory experiments to construct the gas mixtures. In section~\\ref{3} we describe the laboratory measurements and present the results obtained. We conclude in section~\\ref{4} with a discussion. ",
        "conclusions": "\\label{4} We have presented new results of gas mixtures for wavelength calibration echelle spectrographs in the near infrared. We have worked on different gas cells including several new gases. The working method and the properties of the gas cells have been described. We have obtained several mixtures and we have presented a compact and manageable gas cell which covers the widest wavelength range to date in the H and K-bands, with a potentially high number of lines than for currently available gas cells, stable in time scales of months under atmospheric temperature conditions of our laboratory, and which can be useful for high precision radial velocity measurements. We work on the improvement of the gas cell using different partial pressures of the individual gases in order to solve the pressure broadening of some absorption bands. Some of these gases have been recently tested with real observations and obtained promising results with few m/s accuracy. Such gas cells can be of interest for several new generation high resolution near infrared spectrographs under development (NAHUAL, GIANO, PRVS, SPIROU, CARMENES).\\\\"
    },
    "1004/1004.5160_arXiv.txt": {
        "abstract": "Scalar fields with a ``chameleon\" property, in which the effective particle mass is a function of its local environment, are common to many theories beyond the standard model and could be responsible for dark energy. If these fields couple weakly to the photon, they could be detectable through the ``afterglow\" effect of photon-chameleon-photon transitions. The ADMX experiment was used in the first chameleon search with a microwave cavity to set a new limit on scalar chameleon-photon coupling $\\beta_\\gamma$ excluding values between $2\\times10^{9}$ and $5\\times10^{14}$ for effective chameleon masses between 1.9510 and 1.9525 $\\mu$eV. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.5144.txt": {
        "abstract": "We propose a nonlocal scalar-tensor model of gravity with pseudodifferential operators inspired by the effective action of $p$-adic string and string field theory on flat spacetime. An infinite number of derivatives act both on the metric and scalar field sector. The system is localized via the diffusion equation approach and its cosmology is studied. We find several exact dynamical solutions, also in the presence of a barotropic fluid, which are stationary in the diffusion flow. In particular, and contrary to standard general relativity, there exist solutions with exponential and power-law scale factor also in an open universe, as well as solutions with sudden future singularities or a bounce. Also, from the point of view of quantum field theory, spontaneous symmetry breaking can be naturally realized in the class of actions we consider. ",
        "introduction": "Many proposals for modified gravity have been invoked in the hope of finding new insights into the open issues of the standard cosmological model. Among them, theories with pseudodifferential operators have been favored with particular attention. A reason is that nonlocal theories can have very different ultraviolet properties with respect to ordinary second- or higher-order actions (including popular Gauss--Bonnet extensions) and, hence, could play a role near the big bang and as spacetime effective formulations of nonperturbative quantum gravity. String field theory (SFT) is a concrete realization of this notion where pseudodifferential operators of the form \\be\\label{eb} \\rme^{r_*\\B} \\ee decorate the effective target action of the fields, where $r_*$ is a constant and $\\B$ is the spacetime d'Alembertian. The imprint of nonlocal dynamics in the history of the early universe, or even as dark energy models, has motivated the study of cosmological models inspired by open SFT \\cite{are04,AJ,AKV1,cutac,AK,AV,kos07,AJV,AV2,cuta2,Jou07,Jo081,ArK,Jo082,NuM,BMNR,KV,cuta6,Ver09}, the $p$-adic string \\cite{Jo082,NuM,BBC,lid07,BC,cuta4,BK2}, or other nonlocal effective actions featuring the operators \\Eq{eb} \\cite{BMS,kho06} or inverse powers of the d'Alembertian \\cite{SW,DW,NO,Jhi08,Koia,Koib,CENO,DeW,CENOZ}. When nonlocality is of the type \\Eq{eb}, it can be conveniently manipulated with the diffusion equation approach, which has been developed and employed, in analytic and numerical fashion, under different formulations \\cite{cuta2,Jou07,Jo081,Jo082,NuM,cuta6,cuta4,vol03,FGN,vla05,roll,cuta3,MuN,cuta5,cuta7}.  A rather common assumption in the literature of nonlocal fields in cosmology is that nonlocality is confined only to one sector of the model, while the others are \\emph{local}. In the case of SFT-motivated actions, the nonlocal sector is matter (a scalar field) and gravity is local and with Einstein--Hilbert action: \\be\\label{lnl} S=S_{\\rm nonloc}(\\phi)+\\frac{1}{2\\k^2}\\int \\rmd^Dx\\sqrt{-g}\\,R\\,, \\ee where $D$ is the topological dimension of spacetime, $g$ is the determinant of the metric $g_{\\mu\\nu}$, $\\mu=0,\\dots,D-1$, $\\k^2=8\\pi G$ is Newton's constant, and $R$ is the Ricci curvature scalar. This ansatz has been dictated mainly by the urgency of understanding, in the broadest sense, (i) the dynamics of the yet-unclear nonlocal scalar field theories, (ii) the combined effect of curvature and nonlocality, and (iii) its possible consequences for phenomenology, in particular, in relation to cosmology (inflation, dark energy) and the modification of flat open SFT solutions (can cosmological friction damp the wild oscillations of the OSFT solution with marginal deformations? \\cite{cutac,Jo082,BMNR,cuta6,FGN}).  Now that robust analytical and numerical methods have been established to solve nonlocal equations of motion, it would be highly desirable to address the conceptual inconsistency subjacent to Eq.~\\Eq{lnl}. Not only would we like to define a model with nonlocality implemented in all sectors (and reproducing standard general relativity in the limit of weak nonlocality), but we want also to find nontrivial cosmological solutions. Such is the twofold objective of this paper.  The problem of nonlocal gravity can be faced under three independent perspectives, one motivated by string field theory, one purely phenomenological, and another a hybrid approach. In the first case, the Einstein--Hilbert action in Eq.~\\Eq{lnl} is introduced by hand as an educated guess on ``how the effective SFT action of tachyon might look like in the presence of gravity.'' Gravity is minimally coupled with a tachyon-type or $p$-adic scalar field whose action is dictated or inspired by concrete Minkowski calculations. Obviously, a fully consistent effective tachyonic action should be derived from first principles in all its sectors. As far as gravity is concerned, the natural framework is closed SFT \\cite{SZ,KKS,KuS,KS3,Zwi92,SZ1,SZ2,SZ3,OZ,YZ1,YZ2,Mic06,Moe1,Moe2,Moe3}, which features the same nonlocal operator \\Eq{eb} of open SFT. The subject is rather intricate and, unfortunately, effective gravitational nonlocal actions are known only at the linear level \\cite{KS3,Moe3}.\\footnote{On the other hand, the \\emph{local} low-energy effective field theory of the closed string tachyon-dilaton-graviton system is well understood also in its cosmological properties \\cite{YZ3,Swa08}.}  Instead of facing the rigors of closed SFT some toy models have been considered, in particular the open-closed $p$-adic tachyonic action \\cite{BF2,MoS,Vla06} and a closed SFT-inspired tachyon-tachyon model \\cite{Ohm03}. However, the graviton is not included, thus leaving the (third) possibility to consider phenomenological actions where the matter sector is as close as possible to SFT or the $p$-adic string \\cite{BCK}, while a nonlocal gravitational sector is built from reasonable requirements (mainly, that it contains the same type of pseudodifferential operators as the matter sector). Here we shall follow the third path.  Nonlocal gravity sectors have been constructed with inverse powers of the $\\B$ operator \\cite{SW,DW,NO,Jhi08,Koia,Koib,CENO,DeW,CENOZ} or more general kinetic functions \\cite{kho06}, while keeping matter local. In \\cite{BMS} a nonlocal total action has been proposed with nonminimal coupling between gravity and a scalar field, but the dynamical analysis therein does not go beyond cosmological solutions when the matter sector is switched off.  In Sec.~\\ref{act} we adopt the diffusion equation method to infer the form of a solvable scalar-tensor nonlocal action with pseudodifferential operators of exponential type, Eq.~\\Eq{eb}. This approach is chosen by virtue of its nonperturbative character, which does not require truncating the theory in order to find solutions, exact or asymptotic. Exact nonvacuum solutions of the equations of motion of a $p$-adic-like system will be found in Sec.~\\ref{sols} for cosmological backgrounds; their classical stability is checked. The exact solutions are stationary along the diffusion flow (i.e., the diffusion equation is trivially satisfied) but the scalar and Hubble profiles as well as their dynamics are nontrivial. Notably, there exist solutions with exponential and power-law scale factor when intrinsic curvature is negative (Sec.~\\ref{cos2}), as well as a most general class of explicit solutions for actions with conformal operators (Sec.~\\ref{cos3}). We make some general remarks on nonstationary asymptotic solutions in Sec.~\\ref{cos4}, showing that, for natural choices of the potential, the system realizes spontaneous symmetry breaking.  We are interested in a scalar-tensor theory where nonlocal effects are dominant, and such a system finds a natural application only in the early, inflationary, or bouncing universe (here we ignore applications of nonlocal models to dark energy). For this reason, in most of the paper we do not assume the presence of any cosmological fluid such as dust or radiation, but nevertheless inclusion of extra matter components is briefly discussed in Sec.~\\ref{matte}. In Sec.~\\ref{slac} the analysis is extended to another action with a kinetic operator similar to the one of the SFT tachyon; de Sitter and power-law exact solutions are found. Section \\ref{disc} is devoted to discussion. In the Appendix we recall flat and curved exponential and power-law solutions in standard general relativity.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{disc}  In this paper we have constructed and solved, on cosmological backgrounds, an effective nonlocal model of gravity nonminimally coupled with a scalar field. The actions \\Eq{act1} and \\Eq{actt} are nonperturbative both in the order of curvature invariants,\\footnote{Actions of the form \\Eq{loca} are sometimes dubbed ``exponential gravity'' \\cite{CENOS,Lin09}.} \\be\\label{loca} \\rme^{R}\\sim 1+R+\\frac12R^2+\\dots\\,, \\ee and in the number of derivatives acting on the metric, \\be\\label{trunc} \\rme^\\B R\\sim R+\\B R+\\frac12\\B^2R+\\dots\\,. \\ee Nonlocal cosmology is radically different from higher-order cosmological models of $f(R)$, Gauss--Bonnet, or $f(\\textrm{Gauss--Bonnet})$ gravity. This is because truncation of a nonlocal model in spacetime derivatives produces an order $n$ Ostrogradski-like problem with altogether different physical properties. The well-known fact that theories with an infinite number of derivatives are not the large $n$ limit of finite-order actions has been also invoked to question the relevance of higher-order cosmologies in the early universe \\cite{CDD}. Rather than a finite-order truncation of the gravitational action, near the big bang curvature effects should be consistently taken into account only within a fully nonperturbative framework in the sense of the left-hand side of Eqs.~\\Eq{loca} and \\Eq{trunc}. The right-hand side of Eq.~\\Eq{trunc} might not be even well defined on a general nonlocal solution of the system \\cite{cuta2}. The diffusion equation method allows one to deal with the full nonlocal operators and bypass the problems of a series expansion.  The diffusion structure we have explored is asymmetric in the gravity and scalar sector; in fact the former does not diffuse at all. The only solutions we have been able to find do not diffuse even in the matter sector (more precisely, they are stationary along the diffusion flow), but in general a nonstationary diffusion structure is necessary to solve the system with a self-interacting (higher-order potential) scalar field. For the purpose of finding analytic solutions, this should exclude the \\emph{a priori} assumption that, preferring a ``symmetric'' formulation of the model, also the scalar sector does not diffuse. In this case, geometry through the curvature term $f(R)$ would replace diffusion along $r$. Therefore, the theory of diffusion associated with nonlocal actions would be simply defined differently: Diffusion always takes place through geometry, but in the case of trivial geometry (Minkowski background), this is realized by an auxiliary higher-dimensional structure. If this was really the case, however, it would be probably difficult to find analytic or semianalytic solutions with nonlinear self-interaction (nonquadratic $V$).  The actions we have studied are structurally similar to the one advanced in \\cite{BMS} for the following reason. On one hand, the proposal of \\cite{BMS} aimed at an ultraviolet-finite action for quantum gravity which would address the big bang singularity problem. On the other hand, we wanted an action which would be nonlocal in both matter and gravity sectors and be endowed with a diffusion structure allowing one to reduce the dynamics to a set of local equations with both a second-order differential structure (in spacetime) and an algebraic structure (in the diffusion direction \\cite{cuta2,cuta3,cuta5}). These questions, however, are implicitly related: ghost and asymptotic freedom are determined by the specific choice of pseudodifferential operators, in this case one with a natural diffusion structure. So diffusion and good ultraviolet properties are tied together, as expected in string field theory \\cite{cuta7}.  There is, anyway, a caveat in this comparison. We not only stressed the importance of solving a fully nonlocal action with both gravity and matter cosmological nontrivial profiles, but in doing so it was also shown how these profiles can differ, even considerably, with respect to local scenarios.\\footnote{Deviations from local cosmology is not limited to background solutions. It would be interesting to study the inflationary spectra stemming from the inhomogeneous perturbation of the Einstein equations.} The simplest cosmological profiles (de Sitter and power law) are exact solutions of the nonlocal dynamics. There are a couple of remarkable facts associated with that. First, we needed only to look at stationary solutions along the diffusion flow. Second, contrary to standard general relativity these profiles correspond to exact dynamics even when the intrinsic curvature $\\textsc{k}$ is negative definite. In particular, de Sitter is an exact solution for a nonconstant scalar field profile, also in an open universe. When the nonlocal operators are chosen to be conformal, for models with $f(R)=-\\a_* R$ we have found the general solution for \\emph{any} flat or open FRW background, embodied by Eq.~\\Eq{geso}; these results have been extended to the inclusion of a barotropic fluid, Eq.~\\Eq{vptau12}, also for stringlike actions. Within this class there are solutions without big bang singularity, but there also exist an infinite number of solutions with big bang. Therefore we incline not to link nonsingular solutions with the ultraviolet structure of the nonlocal action.  At any rate, the space of solutions is likely to be much larger than the portion we have explored here. All our exact solutions have a quadratic potential. Highly nonlinear equations of motion are of great interest, especially in string theory, but the exact solutions can give some indication of the behaviour for general potentials near a local minimum. Cosmological friction modifies the dynamics of nonlocal scalars with respect to Minkowski and, in particular, should drastically change the rolling of the tachyon in string field theory.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%  %"
    },
    "1004/1004.2430_arXiv.txt": {
        "abstract": "Magnetically-driven non-stationary acceleration of  jets in AGNs  results in the leading parts of the flow been accelerated to much higher Lorentz factors than in the case of steady state acceleration with  the same parameters. The higher Doppler-boosted parts of the flow may dominate the high energy emission of blazar jets. We suggest  that  highly variable  GeV and  TeV  emission in blazars is produced by the faster moving  leading edges of  highly magnetized non-stationary ejection blobs, while the radio data trace the slower-moving bulk flow.  Model predictions compare favorably with the latest   Fermi $\\gamma$-ray and MOJAVE radio VLBI results. ",
        "introduction": "\\subsection{ Bulk Lorentz Factor Crisis in AGNs}  One of the defining characteristics of many AGNs is their flux variability in all spectral bands \\citep{Krolik:1999}. In particular, extremely short time scales of TeV variability are challenging to the models.  The rapid  flares reported for Mrk 501 and PKS 2155$-$304, on timescales of 3-5 minutes \\citep{2007ApJ...669..862A, 2007ApJ...664L..71A} imply  an emitting size smaller than the gravitational radius $t_{lc}\\sim$hours of the supermassive black holes of these blazars. There are two contradictory issues related to short time scale variability. First, it implies  a small emission size, which poses a problem for efficiency of energy conversion into radiation. Secondly, there is the compactness problem \\citep{Guilbert:1983}.  If variability is detected in $\\gamma$-ray photons of energies exceeding the electron rest mass energy, then the emission region contains photons which can pair produce.  If the number density of these photons is too high, then none of the photons will escape the region. The solution to both  problems is bulk relativistic motion towards the observer, which  reduces the intrinsic luminosity, decreases the implied energy of the photons, and  increases the internal  time scales.   The required  Doppler factor  then exceeds $\\delta \\geq 100$.    While highly relativistic motion may appear to be a cure-all, in AGNs  the  bulk Lorentz factor $\\gamma$ can be directly constrained by VLBI observations of bright blobs moving with apparent speeds on the sky, $\\beta_{app}$, that appear to be superluminal.  This type of motion occurs when the emitting region is moving relativistically and close to the line of sight \\citep{Rees:1966}.  The apparent transverse motion can exceed $c$ due to propagation effects.  If a blob is moving along with the bulk flow of a jet and its velocity vector makes an angle, $\\theta_{ob}$, with the line of sight, then its apparent motion transverse to the line of sight will be: $ \\beta_{app}={\\beta \\sin{\\theta_{ob}}}/({1-\\beta \\cos{\\theta_{ob}}}) $. The maximum $\\beta_{app}$ can reach is $\\beta \\gamma$ when $\\theta_{ob}\\cong1/\\gamma$.  Thus, if the blob motion corresponds to the underlying bulk motion of the jet, measuring $\\beta_{app}$ can constrain the possible bulk Lorentz factor, $\\gamma$.  \\subsection{MOJAVE results}  The latest MOJAVE VLBI results do support  the interpretation of moving jet features as physical entities, as opposed to patterns  \\citep{2009AJ....138.1874L}. Observations of  bidirectional motions, the near-absence of inward moving features, ejections of  multiple blobs in the same jet  with the same speed, and tight  correlations of jet speeds with other properties, such as  $\\gamma$ ray emission, apparent $\\gamma$-ray luminosity, brightness temperature,  and even optical classification, all support the notion that the blob motion reflects the underlying flow.  The MOJAVE survey of compact, highly beamed radio-loud AGN has analyzed the motion of emitting blobs in 127 jets and found that the observed superluminal speed distribution peaks at $\\beta_{app}\\sim10$ and tapers off at $\\beta_{app}\\sim 50$ \\citep{Lister:2009}.  This suggests that the bulk Lorentz factors of such objects are typically around $\\sim10$, and extend up $\\sim50$, making the estimated values of $\\gamma \\geq 50$ for PKS 2155$-$304 and Mrk 501  rather difficult to reconcile with the radio data.  Furthermore, direct VLBI observations of these sources on parsec scales have not even detected superluminal motion \\citep{Piner:2004,Giroletti:2004}.  VLBI observations of blazars such as PKS 2155$-$304 and Mrk 501 are not the only data which imply a low $\\gamma$.  Another way of investigating blazars is to search for their AGN counterparts whose jets are not directed along the line of sight, which are presumed to be radio galaxies (there are actually two distinct types of radio galaxies, FRI and FRII, that are thought to correspond to the two categories of blazars, BL Lacs and optically violent variable quasars \\citealt*{Urry:1995}.)  However, studies comparing the relative fluxes and numbers of radio galaxies and blazars point towards Lorentz factors of $\\gamma <10$ \\citep{Henri:2006}. Indeed, preliminary results of MOJAVE observations of low luminosity jets from radio galaxies imply speeds of $c$ or less. Thus, there is apparent contradiction between measured superluminal velocities and bulk Lorentz factors required by radiation modeling. This is known as  the  {\\it \"Blazars'  Bulk Lorentz Factor Crisis\"} \\citep{Henri:2006}. ",
        "conclusions": "In this paper we discuss  the effects associated with non-stationarity of the jet ejection. In particular, we argue that  the leading edge of a non-stationary {\\it magnetized} outflow  can achieve a bulk Lorentz factor much larger than would be inferred for a steady state outflow given similar conditions. In the case of the expansion of a highly magnetized plasma, the ratio  of Lorentz factors of  the bulk flow and that of the leading edge can be as high as $2 \\sigma^{2/3}$ ($\\sigma$ is plasma magnetization). This ratio can reach tens for highly magnetized flows with $\\sigma \\sim 10$. We suggest that the Doppler factor crisis in AGNs (a difference of the Doppler factors  inferred from radiation modeling, especially of short time scale TeV flares and from the observations of radio blobs) may be resolved by non-stationary outflows: highly variable emission is produced by the fast-moving leading edge of an expansion, while the radio data trace the slower-moving bulk flow.  The suggested model is qualitatively different from the internal shock models, where non-stationarity is invoked to produce shocks and dissipate the energy of the relative bulk motion of the colliding media. Collision of  strongly magnetized plasma blobs results in  only weakly dissipative  internal shocks. The fast leading expansion edges will  generate powerful  shocks in the surrounding medium that may produce the high energy emission. We do not address the question of how magnetic and bulk energy is converted into radiation.  A somewhat similar {\\it continuous} acceleration mechanism was proposed by \\cite{Tchekhovskoy} \\citep[see also][]{2009arXiv0912.0845K}.  It relies on {\\it sideways} expansion of the jet after the  break out. Sideways expansion of unconfined magnetically dominated plasma proceeds with Lorentz factor $1+2\\sigma$, so that the total Lorentz factor (of a plasma near the edge, affected by the rarefaction wave) is $(1+2\\sigma) \\gamma_w$, two times smaller than for the case of expansion wave propagating along the direction of motion.  This factor of 2 may have an important effect on escape of high energy radiation, since the optical depth to pair production scales approximately as $\\Gamma^{-6 }$, a difference in a factor of 2 in  $\\Gamma $ will result in a difference of $64$ in the optical depth.   (\\cite{Tchekhovskoy} cannot treat  parallel acceleration akin to breaking into vacuum since the rotation of the central engine is turned on gradually for numerical reasons.)   In summary, we believe that magnetically-driven non-stationary jet acceleration can provide a potential resolution of the longstanding bulk Lorentz factor crisis in blazars. The rich Fermi-VLBA dataset that is currently being gathered on a broad set of AGN should provide an excellent means of testing our proposed scenario.  Acknowledgments:  This research has made use of data from the MOJAVE database that is maintained by the MOJAVE team. The MOJAVE program is supported under National Science Foundation grant 0807860-AST and NASA-Fermi grant NNX08AV67G. The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. We would like to thank Roger Romani for discussions."
    },
    "1004/1004.0435_arXiv.txt": {
        "abstract": "{Doppler tomography of emission line profiles in low mass X-ray binaries allows us to disentangle the different emission sites and study the structure and variability of accretion disks.} {We present UVES high-resolution spectroscopic observations of the black hole binary \\axb\\ at quiescence.} {These spectroscopic data constrain the orbital parameters $P_{\\rm orb}=0.32301405(1)$ d and $K_2=437.1\\pm2.0$ \\kmso. These values, together with the mass ratio $q=M_2/M_1=0.062\\pm0.010$, imply a minimum mass for the compact object of $M_1 \\sin^3 i = 3.15 \\pm 0.10$~\\Msuno, consistent with previous works.} {The H$\\alpha$ emission from the accretion disk is much weaker than in previous studies, possibly due to a decrease in disk activity. Doppler imaging of the H$\\alpha$ line shows for the first time a narrow component coming from the secondary star, with an observed equivalent width of $1.4\\pm0.3$~{\\AA}, perhaps associated to chromospheric activity. Subtracting a K-type template star and correcting for the veiling of the accretion disk yields to an equivalent width of $2.8\\pm0.3$~{\\AA}. A bright hot-spot is also detected at the position where the gas stream trajectory intercepts with the accretion disk.} {The H$\\alpha$ flux associated to the secondary star is too large to be powered by X-ray irradiation. It is comparable to those observed in RS CVn binaries with similar orbital periods and, therefore, is probably triggered by the rapid stellar rotation. } ",
        "introduction": "\\label{introduction}  The black hole binary \\axb\\ (V616 Mon) is one of the most studied low mass X-ray binaries (LMXBs), considered as the prototype soft X-ray transient (SXT). It was discovered in 1975 by the \\emph{Ariel 5} \\citep{elv75} satellite during a X-ray outburst where the optical brightness of the system increased by roughly 6~mag in few days. One year and several months after the system returned to its quiescent state at $m_V\\sim18.35$~mag. The spectrum of a stellar counterpart was then identified and classified as a K5V-K7V star \\citep{oke77,mur80}. Further spectroscopic observations allowed the determination of the orbital period of the secondary star at $\\sim7.75$~hr \\citep{mar86} which implied the presence of a black hole of  minimun mass $\\sim 3.1$~\\Msun \\citep{mrw94}.  The orbital inclination of the system was later estimated from IR light curves at $\\sim41^{\\circ}$ implying a primary black hole mass of $11.0\\pm1.9$~\\Msun (Gelino et al. 2001, see also Shahbaz, Naylor \\& Charles 1994). However, this study adopts a K4V stellar component with \\teff$\\sim4600$~K, 300~K cooler than the effective temperature \\teff$=4900\\pm100$~K derived by \\citet{gon04} from high-resolution optical spectroscopic observations. This may affect the required veiling in the IR, and consequently, the derived inclination and black hole mass \\citep{hyn05}. Furthermore, \\citet{can10} present evidence for substantial disk contamination in their IR light curves and find $i=51\\pm1^{\\circ}$ which translates into a lower black hole mass of $6.6\\pm0.3$~\\Msun. On the other hand, contradictory results have been reported using low-resolution IR spectra. \\citet{har07} find a very small or negligible disk veiling at IR wavelengths whereas \\citet{fro07} conclude that it can be 18\\% in the H band. The latter also argue that \\teff$\\lesssim 4600$~K is needed to fit their observations. However, we must note here that most studies just adopt a \\teff based on spectral classification, derived through comparison with low resolution template spectra, without determining the \\emph{true} \\teff of the template \\citep{hyn05}. A similar inconsistency has been found for the black hole binary \\mbox{XTE J1118+480} for which \\citet{gel06} adopted a \\teff$=4250$~K where \\citet{gon06,gon08b} derive a spectroscopic \\teff of~$4700\\pm100$~K.  The accretion disk of \\axb\\ has been studied in the UV \\citep{mhr95} and the optical \\citep{mrw94,oro94}, allowing the investigation of the inner and outer disk. Both works in the UV and optical seem to agree that the accretion disk is in a \\emph{true} quiescent state. However, this does not mean that the disk is inactive and probably its variability appears to be relevant \\citep{sha04,can08,can10}. \\citet{sha04} also suggested that the accretion disk in \\axb\\ could be eccentric which may have been confirmed by more recent observations reported by \\citet{nsv08}. As pointed out by these authors, to determine the definite mass of the compact object, it is very important to understand the structure and variability of the accretion disk as exemplified by \\cite{can10}.  Here we present high-resolution spectroscopy of \\axb\\ where we detect clear emission arising from the secondary star in the H$\\alpha$ Doppler map. This feature has only been observed before in the systems GU Mus \\citep{cas97}, Nova Oph 77 \\citep{har97}, Cen X-4 \\citep{tor02,dav05} and Nova Scorpii 1994 \\citep{sha99}. These data also allows us to revisit the orbital parameters of the system, which we find to be consistent with previous studies. ",
        "conclusions": "\\citet{str90} measured the H$\\alpha$ EWs of a sample of F6-M2 single and binary stars, showing that the H$\\alpha$ EW increases towards shorter rotation periods, $P$. Extrapolation of this trend EW vs. $P$ at the orbital period of \\axb\\ provides an expected EW of $\\sim 2.3$~{\\AA}, assuming that the secondary star is a dwarf main-sequence star and its rotation is sychronized with the orbital motion. This EW is marginally consistent (at 1.5$\\sigma$) with the observed EW of the secondary star. However, the extrapolation might not be adequate since chromospheric features powered by rotation typically saturates for period shorter than 1--3 days \\citep[e.g.][for chromospheric \\ion{Mg}{ii} lines]{car07}.  \\begin{figure}[ht!] \\centering \\includegraphics[width=8.5cm,angle=0]{14088f5.ps} \\caption{\\footnotesize{Chromospheric flux ratios vs Rossby numbers of chromospherically active single stars from \\citet[][in press, filled circles]{lop10} and chromospherically active binary systems (RS CVn and BY Dra classes) from \\citet[][open circles]{mon95}. Low mass X-ray binaries are also depicted: A0620--00 (diamond), Nova Muscae 1991 (square), Cen X-4 (filled triangle, from Torres et al. 2002; open triangle, from D'avanzo et al. 2005), and Nova Scorpii 1994 (inverted triangle).}} \\label{figrharo} \\end{figure}  We can also compare the H$\\alpha$ EW, converted into flux following the approach of \\citet{sod93}, with the observations of field stars and binaries. Thus, the H$\\alpha$ flux at the stellar surface, is $F_{\\rm H\\alpha}=8.8\\pm1.1\\times10^6$ erg~cm$^{-2}$~s$^{-1}$ (see Table~\\ref{tabfluxha}). We use a similar prescription than that of Equation~2 in \\citet{sod93}, i.e. $F_{\\rm H\\alpha} = {\\rm EW (H\\alpha)} F_c$, where is $F_c$ is the continuum flux at H$\\alpha$ and is derived using the flux calibration of \\citet{hal96}, $\\log F_c = 7.538 - 1.081 (B-V)_0$. Here we have used $(B-V)_0=0.965$ estimated from theoretical colours \\citep{bes98} according to the stellar parameters of the secondary star \\citep{gon04}. The surface flux is usually normalized to the bolometric flux, i.e. $R_{\\rm H\\alpha}=F_{\\rm H\\alpha}/\\sigma T_{\\rm eff}^4$. The connection between chromospheric activity and rotation is obtained by studying the correlation $R_{\\rm H\\alpha}$ with the {\\em Rossby} number $R_0=P/\\tau_c\\lesssim 2\\pi R_\\star/\\tau_c v \\sin i$, where $\\tau_c$ is the convective turnover time, $P$, the rotation period of the secondary star and $R_\\star$, its radius. Using the {\\em Rossby} number is usually preferred over $P$ and $v \\sin i$ since it does not depend on the mass of the star.  For the case of \\axb, we adopt $T_{\\rm eff}=4900$~K \\citep{gon04}, $P=P_{\\rm orb}=0.32$~d and $\\tau_c\\sim23$~d, computed from Eq.~(4) of \\citet{noy84}. We derive $\\log~R_{\\rm H\\alpha}=-3.57\\pm0.06$ and $\\log~R_0=-1.85$, and this is listed in Table~\\ref{tabfluxha}, together with values derived for other quiescent X-ray transients. In Fig.~\\ref{figrharo} we compare these values with those of chromospherically active single stars \\citep[][in press]{lop10} and binary systems (RS CVn and BY Dra classes) from \\citet[][]{mon95}. Our value of the {\\em Rossby} number places the secondary star of \\axb\\ in the region of activity saturation, where all measurements tend to the same average value of $\\log~R_{\\rm H\\alpha}\\sim-3.7$, and is consistent with the general trend. The other X-ray binaries show similar results, except for the black hole X-ray binary Nova Scorpii 1994 which displays a too large $R_{\\rm H\\alpha}$ value for its relatively low Rossby number. This is even more evident when comparing the H$\\alpha$ fluxes, with $\\log F_{\\rm H\\alpha}\\sim7.9$ for Nova Scorpii 1994, significantly higher, by almost one order of magnitude, than the saturation level at $\\log F_{\\rm H\\alpha}\\sim6.9$. In addition, \\citet{dav05} also suggested that the H$\\alpha$ EW of the secondary star in Cen X-4 is correlated with the veiling of the accretion disk (see Table~\\ref{tabfluxha}), by comparing their values with those given by \\citet{tor02}. Although this is not expected in a chromospheric activiy scenario, the two Cen X-4 points in Fig.~\\ref{figrharo} fall in the saturation region, together with other X-ray and chromospherically active binaries. The behaviour seen in Nova Scorpii 1994 and Cen X-4 suggests that rapid rotation might not be the only explanation for the narrow H$\\alpha$ feature, at least in these LMXBs, but perhaps a combination of rotation and reprocessing of X-ray flux from the accretion disk into H$\\alpha$ photons in the secondary star. Hence, it is worth investigating in \\axb\\ if X-ray heating could be an alternative explanation for this feature.  The system \\axb\\ has been observed in quiescence with the {\\em Chandra} X-ray satellite, providing a 0.5-10~KeV unabsorbed flux $F_{X,0}=6.7^{+0.8}_{-2.3}\\times10^{-14}$~erg~cm$^{-2}$~s$^{-1}$ \\citep{gal06}.  The X-ray flux at the stellar surface can be computed as $F_{X,\\star}=F_{X,0}(d/a)^2=9.5\\times10^6$~erg~cm$^{-2}$~s$^{-1}$, where we have adopted an orbital separation $a=4.47$~\\Rsun and a distance $d=1.2$~kpc. This means that almost 92\\% of the incident X-ray radiation would have to be reprocessed to H$\\alpha$ photons in order to power the observed H$\\alpha$ emission. Following \\citet{hyn02}, $F_{\\rm H\\alpha,\\star}=f_1f_2F_{X,\\star}$, where $f_1$ is the fraction of X-ray emission intercepted by the companion, i.e. the solid angle subtended by the companion from the compact object ($f_1=[R_\\star/(2a)]^2$), and $f_2\\lesssim 0.3$ is the fraction of input energy emitted in H$\\alpha$ \\citep[][and references therein]{hyn02}. Adopting $R_\\star=1.1$~\\Rsun from \\citet{gon04}, we obtain $F_{\\rm H\\alpha,\\star} \\lesssim 4.5 \\times 10^{-3} F_{X,\\star}$. This number is significantly lower than the observed value, what indicates that the incident X-ray irradiation is not enough to produce the narrow H$\\alpha$ line in the secondary star. \\citet{dav05} also derived these quantities for the case of the neutron star binary Cen X-4 and found both estimates to be consistent, due to the fact that the X-ray flux in Cen X-4 is $F_{X,\\rm Cen X-4}=5\\times 10^8$~erg~cm$^{-2}$~s$^{-1}$, i.e. almost two order of magnitude higher than in \\axb.  There is still a remote posibility that the source of irradiating photons is hidden away in the EUV (Extreme UV) energy range, between 100--1200~{\\AA}. Although this energy range is not directly observed, we can roughly guess how much flux is involved through interpolating the nearby soft X-ray and Far-UV (FUV) emission. The Far-UV flux (in the range 1350--2200~{\\AA}) has been determined at $F_{{\\rm FUV},0}=0.2-1.4\\times10^{-13}$~erg~cm$^{-2}$~s$^{-1}$ \\citep{mhr95}, i.e. similar to the X-ray flux. This together with the absence of \\ion{He}{ii}~$\\lambda$4686~{\\AA} line emission in the optical spectrum \\citep{mrw94} suggests that flux in the EUV should be of the order of $10^{-13}$~erg~cm$^{-2}$~s$^{-1}$. Even if we consider all the ionizing photons (X-ray$+$EUV$+$FUV), the total flux would be roughly three times the X-ray flux, i.e. $\\sim 2\\times10^{-13}$~erg~cm$^{-2}$~s$^{-1}$. Therefore, if we assume that the incident irradiation is $F_{X-UV,\\star}= 3\\,F_{X,\\star}$, then the 31\\% of the incident radiation would have to be reprocessed to H$\\alpha$ photons in order to power the observed H$\\alpha$ emission, which still is a too large fraction compared to the fraction previously estimated to be bellow 1\\%."
    },
    "1004/1004.3600_arXiv.txt": {
        "abstract": "{}  % {Nova Scorpii 2008 was the target of our Directory Discretionary Time proposal at VLT+UVES in order to study the evolution, origin and abundances of the heavy-element absorption system recently discovered in 80\\% of classical novae in outburst.} % {The early decline of Nova Scorpii 2008 was monitored with high resolution echelle spectroscopy at 5 different epochs. The analysis of the absorption and the emission lines show many unusual characteristics.}  % {Nova Scorpii 2008 is confirmed to differ from a common Classical Nova as well as a Symbiotic Recurrent Nova, and it shows characteristics which are common to the so called, yet debated, red-novae. The origin of this new nova remains uncertain.}  % {} ",
        "introduction": "Nova Scorpii 2008 was independently discovered at mag $\\sim$9.5-10 (unfiltered CCD) on Sep 2.5 UT by K. Nishiyama, F. Kabashima and Sakurai in  Japan and by Guoyou Sun and Xing Gao in China (CBET 1496). Nakano (2008) reports precise coordinates of the new object, noting that it was invisible on earlier observations taken by the same authors on Aug 20.5 and 21.5 UT (limit magnitude 12.8 and 12, respectively) and that is only 1.14 arcsec away from the USNO-B1.0 star 0592-0608962 of magnitude B=16.9 and R=14.8. The first low resolution (R$\\sim$500 and 1200) optical spectra were secured by Naito and Fujii (2008) who observed V1309 Sco on 3 consecutive days (Sep 3 to 5) reporting a smooth continuum, some absorption lines, and strong Balmer lines. Rudy et al. (2008a) observed V1309 Sco in NIR spectroscopy on Sep 5 and confirmed it to be a slow nova at very early stages with some emission lines (early H Paschen series and FeII), but mostly absorption lines (H Paschen and Brackett as well as CaII, NI and CI). The lines were narrow with FWHM not exceeding 300 km/s. One month later the same group obtained additional NIR spectroscopy between 1-5 microns that showed a strong continuum resembling that of a late M giant star, upon which were superposed strong molecular absorption from CO, H\\_2O, and weaker features of TiO and VO (see Rudy et al. 2008b).  Recently, Williams et al. (2008) have shown that the majority of classical novae (CNe) in outburst show a short lived absorption system from heavy element (Transient Heavy Element Absorption -THEA- system) which is external to the primary expanding ejecta. The numerous early reports of absorption lines in the spectra of V1309 Sco motivated us to apply for Director\u2019s Discretionary Time (DDT) on the VLT+UVES to better characterize the THEA system in classical novae.  In this paper we present the spectra we obtained and their analysis. ",
        "conclusions": "Inspection of the AAVSO V-band data shows that V1309 Sco reached maximum light on Sep 6 2008 and declined with a relatively smooth light curve in the following 1.5 months. The same light curve shows that the nova $t_2$ time is $\\sim$20 days, while its t$_3$ is about 25 days. However, the maximum of Sep 6, could have been a secondary one, the first having occurred on Sep 1 at visual magnitude $\\sim$7.5 (see Fig.~\\ref{fig0}) The scarce monitoring of the following days shows a sudden drop with a possible t$_2$ of only $\\leq$3 days.  The postion of V1309~Sco is nearly coincident with that of a red USNO-B1 star ($\\sim$1$\\sigma$) of magnitudes B=16.88 and R=14.80 (year of observation: 1966). Within $\\leq$2$\\sigma$ from the nova position there is a 2MASS object of J, H and K magnitudes equal to 13.282, 12.373 and 11.099, respectively. On a 1958 POSS-I~E/red plate,  the V1309~Sco progenitor has been identified with an object fainter than 19 mag (Jaques and Pimentel 2008, IAUC 8972; the mag limit of the POSS-I~E survey is 20). Depending on the correct progenitor, V1309~Sco outburst amplitude could either have been $\\sim$7 mag or as large as 12 mags.  The V1309~Sco light curve and spectral evolution described in the previous section are peculiar in the sense that the nova does not follow the prototype of any single class of mid-large outburst amplitude objects. Its  rapid spectral evolution toward a red continuum  with possible development of red-giant signatures after only a few weeks from the outburst (Rudy et al. 2008b, IAUC 8997) make it similar to the symbiotic recurrent novae  V745 Sco/89 and V3890 Sgr/90 (Williams et al. 1991). The postoutburst luminosities of the red continua for all three novae are many magnitudes greater than the preoutburst brightnesses, and require a photospheric radius that is orders of magnitude larger than that of a normal late-type giant, and substantially larger than the size of the Roche lobe of a CV with a period of order one day.  However, both V745~Sco and V3890~Sgr showed a faster decline and a somewhat different spectral evolution. Sekiguchi et al. (1990) report t$_2$=5 and t$_3$=9 days for V745~Sco; while Anupama and Sethi (1994) measured t$_2$=12 and t$_3$=17 days, in the case of V3890~Sgr. Their outburst amplitude (A$\\sim$7.3 mag for V745~Sco, Sekiguchi et al. 1990; 7$\\leq$A$\\leq$9 mag for V3890~Sgr, Wenzel 1990) are consistent with those typically observed in symbiotic recurrent novae such as RS-Oph and T~CrB.  V745~Sco spectra were characterized by large velocities (FWHM$\\sim$1000 km/s and extended wings up to 4000 km/s, Sekiguchi et al. 1990), and early development of high ionization energy emission lines (e.g. the 4640\\AA \\ blend and the HeII$\\lambda$4686) and forbidden transitions ([OII](1), Sekiguchi et al. 1990, as well as [FeX], [FeVII] and [FeXI], Williams et al. 1991). Similarly Anupama and Sethi (1994) observations showed that in less than one month V3890~Sgr has developed forbidden coronal lines from [FeX], [FeXIV], [AX] and [AXI]. Rapid evolution and early development of high ionization potential emission lines and forbidden lines was also observed by Williams et al. (1991) within their survey and monitoring program for CN in outburst at CTIO.  V1309 Sco spectra has not developed any high ionization coronal forbidden transition, yet, though the presence of the [CaII] doublet tends to be observed in those variables that display a late-type stellar continuum in the red, and therefore it might be the signature of emission from an extended chromosphere of the secondary star rather than ejecta from the surface of the white dwarf. The fact that the [CaII] $\\lambda\\lambda$7292,7324 lines are typically not observed in classical novae or nova-like variables is indicative of a density regime different from that of classical novae.  In addition, V1309 Sco velocities as measured from the emission lines FWHM  and their extended wings never exceeded 150 km/s and 1000 km/s, respectively\\footnote{Higher velocities reported in the IAUCs reflect the lower spectral resolution of the instruments.}.  We should further note that it is difficult to fit a symbiotic binary in the progenitor of V1309 Sco because of the missing giant companion. The NIR 2MASS colors measured for the nearby star mentioned above, dereddened using the measured E(B-V) and Cardelli et al. (1989) reddening law, provide J-H and K-H colors which are consistent with a M1 type giant (Frogel and Whitford 1987) at a distance of $\\sim$11kpc. V1309~Sco is in the direction of the galactic center and hence, at a distance $<$8kpc. Though the two distance do not appear significantly different, we believe that V1309~Sco is much closer than 8~kpc because of the relatively small E(B-V) we have estimated from the maximum spectra. Hence, a cool giant companion in V1309~Sco progenitor should have brighter NIR magnitudes than those reported by 2MASS.  It should be added that the presence of heavy element Fe-peak narrow absorption line systems in the early post outburst spectra is not completely unrelated to classical novae. These transient heavy element absorption systems have recently been observed in almost all novae studied at high spectral resolution (Williams et al. 2008), and their prominence in V1309~Sco is remarkably by far the most extensive such system observed so far.  At the same time the red continuum, the narrow Balmer emission lines and the heavy-element absorptions are characteristic of the so called 'red-novae' such as V838~Mon/02 and the less well observed V4332 Sgr/94 and M31-RV/88. These objects all evolved in relatively short time toward M and K giant spectra, with V838~Mon developing the first L-giant spectrum ever claimed  (e.g. Munari et al. 2007). They have never shown evidence either of high ionization potential element emission lines nor they entered the nebular or coronal phases typically observed in CNe (e.g. Munari et al. 2007, Barsukova et al. 2007, Rushton et al. 2005, Banerjee  and Ashok 2002, Rich et al. 1989, Mould et al. 1990, Martini et al. 1999). Forbidden transitions from [OI] and [FeII] have been reported in the late spectra of V838~Mon ($>+$7 months since outburst, Wagner and Starrfield 2002, Munari et al. 2007, Kaminski et al. 2009) and V4332~Sgr ($\\geq$+5 months since outburst, Martini et al. 1999). Large luminosities have been derived for M31-RV (M$_{bol}$=-10 mag, Rich et al. 1989, see also Mould et al. 1990) and V838~Mon (M$_V$=-9.8 mag, Sparks et al. 2008, see also Bond et al. 2003, Tylenda 2004 and Tylenda 2005 and reference therein, Soker and Tylenda 2003 and Tylenda et al. 2005); while the distance of V4332~Sgr is uncertain. In addition, the outburst light curve of the 3 objects differ in their time scale and (within the number of data points available for each of them) morphology, though Munari et al. (2007) have noted that they are ''remarkably similar'' once scaled by the time of their optical brightness free-fall. Munari et al. (2007) also noticed that all the three objects displayed the whole range of M type giants during such 4 mags free fall. We cannot do exactly the same comparison for V1309~Sco, but note that the presence of multiple/secondary maxima in the object light curve, makes it similar to V838~Mon (e.g. Goranskij et al. 2007).  V1309~Sco seems to share few peculiarities with symbiotic recurrent novae and more characteristics with the yet un-understood class of red-novae. Yet it cannot firmly be classified as belonging to this latter type of objects due to the lack of later epochs spectra and the possibly low luminosity. By placing V1309~Sco at the distance of 8 kpc and assuming m$_V$(max)=7.9 mag (see Fig~\\ref{fig0}), we derive the upper limit of M$_V$=-8.3 mag, which is almost 2 mag fainter than the absolute magnitude derived for V838~Mon (Sparks et al. 2008) and M31-RV (Rich et al. 1989). However, should the red nova be explained by stellar merging phenomena, the maximum luminosity is not necessarily a stringent constraint. At the same time, V1309~Sco might represent a link (an intermediate case) between the two classes of symbiotic and red-novae, should the red-nova class be caused by a thermo-nuclear reaction on a small mass accreting white dwarf (Shara et al.2009, see also Iben and Tutukov 1992 and Idan et al. 2009 for a discussion about CNe outburst on small mass accreting white dwarfs) rather than to stellar merging (Tylenda and Soker 2006 and reference therein). The most recent models (Shara et al. 2009) for classical nova outburst on low mass white dwarf ($\\sim$0.5M$_\\odot$) accreting at a low rate (a few 10$^{-11}$\\.M/yr) predict that these type of novae will accumulate large amount of mass on the white dwarf surface, before the TNR ignition. The outburst will result in  massive ($\\sim$10$^{-3}$M$_\\odot$) cool, red ejected envelopes. In addition multiple maxima, tremendous absolute magnitudes (up to M$_V\\sim$-9$\\div$-13 mag or luminosity L$\\sim$10$^6$L$_\\odot$), low expansion velocities and oxygen rich and shocked spectra are predicted too, thus fitting the main characteristics  observed in the red-novae.  However, recent observation in high resolution spectroscopy of V838 Mon (Kaminski et al. 2009) favor the young planet merging in a triple/multiple system within a open cluster. Hence, whether V838 Mon could be explained by the CN outburst on a small mass WD remains uncertain, doubtful and highly debated. In addition, whether it is the prototype of the red novae variables or just a peculiar object has to be established as well.  A class of novae  hosting a small mass white dwarf should exist and may already have been observed. Whether V1309~Sco belong to such a subclass or instead is a red-nova fitting the star-merging scenario can only be established through further spectroscopic observations. Optical and NIR spectroscopy should enable the identification of 1) the late epoch evolution of the object, possibly toward later giant spectral types, 2) the development of high excitation coronal lines and forbidden transitions, 3) the possible presence of a blue companion 4) radial velocity shifts which could be ascribed to orbital motion. In the case of a binary symbiotic-like system, the orbital periods are expected to be of the order of a few hundreds of days. Significantly shorter orbital periods (a few hours to day) would imply a dwarf donor companion."
    },
    "1004/1004.4957_arXiv.txt": {
        "abstract": "The binding energy parameter $\\lambda$ plays an important role in common envelope (CE) evolution. Previous works have already pointed out that $\\lambda$ varies throughout the stellar evolution, though it has been adopted as a constant in most of the population synthesis calculations. We have systematically calculated the binding energy parameter $\\lambda$ for both Population I and Population II stars of masses $1-20 M_{\\sun}$, taking into account the contribution from the internal energy of stellar matter. We present fitting formulae for $\\lambda$ that can be incorporated into future population synthesis investigations. We also briefly discuss the possible applications of the results in binary evolutions. ",
        "introduction": "One of the key stages in the evolution of close binary stars is the common  envelope (CE) evolution. When the primary star fills its Roche-lobe (RL) and the following mass transfer is dynamically unstable so that the secondary cannot accrete all the transferred material, the transferred matter will form an envelope embedding both stars and the binary enters the CE evolution. The secondary star orbits inside the CE, and experiences  a drag force by the envelope. This causes the orbital decay and spiral-in of the star with possible ejection of the envelope \\citep[for reviews see][]{iben93,taa00}. Two main kinds of theories have been developed to describe the envelope ejection process. One is the $\\alpha$ formalism, which considers the conversion of orbital energy to overcome the envelope binding energy \\citep{web84,dek90,dew00}. If the orbital energy of the secondary is large enough to eject the envelope, the system will survive and become a compact binary containing the core of the lobe-filling star and the secondary in a much smaller orbit; if the orbital energy is not enough, the two stars will coalesce. The other is the $\\gamma$ formalism, which considers angular momentum transformation during the spiral-in \\citep[see][and reference therein]{nel05}. Despite extensive three dimensional hydrodynamical simulations of CE evolution \\citep[for recent reviews, see][]{taa00,taa06}, the physics of the CE phase remains poorly understood, including the efficiency with which the CE is ejected from the system. Alternative evolutionary paths have been proposed to avoid the CE evolution, such as those by \\citet{kin99} and \\citet{bee07}, who suggested that super-Eddington accretion may lead to mass ejection from the system instead of entering the CE stage.  According to \\citet{han94, han95} and \\citet{dew00}, the total binding energy of the envelope is described by: \\begin{equation} \\label{equ1} E_{\\rm bind} = \\int^{M_{\\rm donor}}_{M_{\\rm core}} -\\frac{GM(r)}{r} {\\rm d}m+\\alpha_{\\rm th}  \\int^{M_{\\rm donor}}_{M_{\\rm core}} U {\\rm d}m, \\end{equation} where $M_{\\rm donor}$ is the mass of the donor star and $M_{\\rm core}$ is its core mass, $-GM(r)/r$ and $U$ are the gravitational and internal energy of the stellar matter respectively, $G$ is the gravitational constant, and $\\alpha_{\\rm th}$ the percentage of the internal energy contributing to ejection of the envelope, normally taking the value between 0 and 1.  For convenience the total binding energy of the envelope is usually expressed as \\citep{web84, dek90, dew00}: \\begin{equation} \\label{equ2} E_{\\rm bind} = -\\frac{G M_{\\rm donor} M_{\\rm env}}{\\lambda a_{\\rm i} r_{\\rm L}}, \\end{equation} where $M_{\\rm env}$ is the envelope mass, $\\lambda$ is the binding energy  parameter, $r_{\\rm L} = R_{\\rm L}/a_{\\rm i}$ is the ratio of the RL radius and the orbital separation at the onset of CE, and $a_{\\rm i} r_{\\rm L}$ is normally taken as the stellar radius once a star fills its RL and starts to transfer matter. If we assume that part of the change of the orbital energy is used to eject the envelope, as described by \\citet{web84}, then \\begin{equation} E_{\\rm bind} = \\alpha_{\\rm CE} (\\frac{GM_{\\rm core}M_2}{2a_{\\rm f}} - \\frac{GM_{\\rm donor}M_2}{2a_{\\rm i}}) \\end{equation} where  $\\alpha_{\\rm CE}$ is the efficiency of converting the orbital energy to kinetic energy to eject the CE, $M_2$ the mass of the secondary, and $a_{\\rm i}$ and $a_{\\rm f}$ refer to the initial and final orbital separation of the CE phase, respectively.  Combine Eqs.~(2) and (3), the final orbital separation is related to the pre-CE separation by the following formula, \\begin{equation} \\frac{a_{\\rm f}}{a_{\\rm i}} = \\frac{M_{\\rm core} M_2}{M_{\\rm donor}} \\frac{1}{M_2+2M_{\\rm env}/{\\alpha_{\\rm CE}\\lambda r_{\\rm L}}}. \\end{equation} Equation (4) indicates that the orbital separation after CE is sensitively dependent on the product of $\\alpha_{\\rm CE}$ and $\\lambda$, and the convenient way is to treat both $\\alpha_{\\rm CE}$ and $\\lambda$ as constants of the order of unity. In most of the population synthesis investigations concerning the evolution of close binaries in the literature, $\\lambda$ was usually taken to be a constant, normally around $0.5$ throughout the evolution of the system. However, there have lots of works \\citep[e.g.][]{han94,dew00,dew01, pod03,web07} suggesting that $\\lambda$ is a variable. \\citet{dew00} and \\citet{dew01} found that the value of $\\lambda$ changes as the star evolves, and reaches far more than $0.5$ in late evolutionary stages for stars with mass between $3$ and $6 M_{\\odot}$. In recent population synthesis on post-CE binaries, \\citet{dav10} used linear interpolations of the tabular data in \\citet{dew00} to calculate $\\lambda$. We notice that \\citet{dew00} and \\citet{dew01}'s calculations did not cover Pop. I stars less massive than $3 M_{\\odot}$ and Pop. II stars, and their results in  tabular form may not be easily used. Obviously more advanced population synthesis requires a simple description of the parameter $\\lambda$ for better constraints on the initial stellar population. In this work we are attempting to investigate the binding energy parameter $\\lambda$ by considering the contribution of  internal energy of stellar matter for both Pop. I and Pop. II stars, and to find easy-to-use fitting formulae of $\\lambda$ for further research.  This paper is organized as follows. We describe the stellar evolution code and assumptions adopted in Sect.~\\ref{sec2}. In Sect.~\\ref{sec3} we present the calculated results. We summarize our results and briefly discuss their possible implications in Sect.~\\ref{sec4}. ",
        "conclusions": "\\label{sec4}  Our calculations show that stars with different masses have different values of $\\lambda$, and $\\lambda$ is not constant for the same star in different evolutionary stages. From Figs.~\\ref{fig1} and \\ref{fig2}, $\\lambda_{\\rm b}$ diverges from $0.5$ for most $M$ and $R$, and it is obvious that assuming a constant $\\lambda = 0.5$ in population synthesis calculations is far from fact and therefore lack of reliability. It is also interesting to note that the range of $\\lambda_{\\rm g}$ is much narrower than that of $\\lambda_{\\rm b}$. Of course the actual value of $\\lambda$ should lie between $\\lambda_{\\rm g}$ and $\\lambda_{\\rm b}$ since not all of the internal energy contributes to the ejection of the envelope \\citep{dew00}. As we have given the fitting formulae for $\\lambda$, the results can be useful for future population synthesis works. However, this approach should be adopted carefully, especially in the cases when the binding energies are positive\\footnote{Although the gravitational binding energy remains negative for all the stars throughout their evolution, the total binding energy, on the other hand, can turn to be positive for stars with mass ranging from $\\sim 3 - 5 M_ {\\odot}$ (specifically, $3$ to $5.6 M_{\\odot}$ for Pop. I stars and $2.8$ to $4.5 M_{\\odot}$ for Pop. II stars at the end of their AGB phase when their radii expand to be more than hundreds of solar radius, and ionizaiton energy dominates the total binding energy).}, and where the efficiency of conversion of ionization energy to kinetic energy may differ from the efficiency of conversion of other forms of internal energy. As we can see from Eqs. (\\ref{equ1}) and (\\ref{equ2}), the calculated values of $\\lambda$ depend on the stellar mass, the core mass, the envelope mass distribution and the percentage of internal energy that contributes to envelope ejection. Factors like metallicity and stellar wind mass loss also affect the stellar evolution and the resulting values of $\\lambda$. Additionally, the different definitions of core-envelope boundary lead to different stellar core mass and therefore affect the calculated results.  Our results agree reasonably with that of \\citet{han94}, \\citet{dew00}  and \\citet{dew01}, although there exist some differences. \\citet{dew00} proposed that $\\lambda$ is almost independent of the stellar chemical composition. From our calculations we find that given the same mass, Pop. I and Pop. II stars have different $\\lambda$-values (see Figs.~\\ref{fig1} and \\ref{fig2}). For example, the mass range in which the values of $\\lambda_{\\rm b}$  can become negative is narrower for Pop. II stars than for Pop. I ones. In Fig.~\\ref{fig3} we compare the evolution of the two $\\lambda$s for a $5 M_\\odot$ star with different metallicity. It is seen that the $\\lambda$ evolutions show little difference at the beginning, but later they diverge from each other, and the values of $\\lambda$ for the Pop. I star always lie above those of the Pop. II star in late evolutionary stages. The reason is as follows. The binding energy (absolute value) of the Pop. I star is either comparable with (in early evolutionary stage) or lower than (due to the contribution of ionization energy) that of the Pop. II star. Since the Pop. I star evolves slower than the Pop. II one, as a result the Pop. I star has a smaller H-exhausted core and higher-mass envelope than the Pop. II star when they have the same radii, thus higher $\\lambda$ values.  We used the 15\\% H abundance layer as the stellar core-envelope boundary to calculate the binding energy and $\\lambda$, since the core mass from such definition increases in most of the time and is suitable for calculating the binding energy (Eggleton 2009). Note that \\citet{dew00} used the 10\\% H abundance layer as the core-envelope boundary, and in our calculation we found that the two definitions give similar results. There are other criteria to define this boundary, such as the maximum energy generation layer or the density gradient criteria \\citep{han94, dew00}. According to \\citet{tau01} and \\citet{dew01}, most of these criteria give outer core boundary and higher core mass. In general, if the total stellar mass is a constant, a more massive core will decrease the total binding energy of the envelope and result in a larger $\\lambda$-value \\citep{dew00, tau01,dew01}. In other words, the $\\lambda$-values we calculated may be regarded as the lower limit of the actual ones. Stars more massive than $10 M_\\odot$ lose $10\\%$ to $30\\%$ of its total mass during its evolution, which can affect the stellar structure and the values of $\\lambda$ \\citep{dew00, dew01, pod03}. \\citet{dew01} and \\citet{pod03} have also found that $\\lambda$ is sensitive to the wind mass loss for massive stars, and including stellar wind leads to higher binding energy (absolute value) and lower value of $\\lambda$.  Our results suggest that $\\lambda_{\\rm b}$ can take values significantly larger than $0.5$ during the late-stage evolution of stars with mass between $\\sim 2$ and $10 M_\\odot$, which will significantly increase the efficiency of envelope ejection. If we consider an evolved giant or AGB star in this mass range with $\\lambda>>0.5$, the envelope binding energy may be much smaller compared to that with $\\lambda = 0.5$. It is then much easier for its companion to eject the CE, and to produce wider binaries after the CE phase. It is also worth noticing that both $\\lambda_{\\rm g}$ and $\\lambda_{\\rm b}$ are less than $1$ in the late evolutionary stages of massive stars. For stars massive than $\\sim 8 M_{\\odot}$, $\\lambda$ even lies below $0.2$. As a result, the actual envelope ejection efficiency for these stars decreases considerably. Taking a $20M_\\odot$ Pop. I giant star as an example. As seen from Fig.~\\ref{fig1}, both $\\lambda_{\\rm b}$ and $\\lambda_{\\rm g}$ are $\\ll 1$, which means that the envelope binding energy become much higher. The secondary has to convert more orbital energy to eject the envelope, which probably leads to a much tighter final orbit, or even a merger. These may have important implications for the formation and evolution of related binaries that produce compact stars, which will be discussed in future works in more detail."
    },
    "1004/1004.4166_arXiv.txt": {
        "abstract": "Globular clusters were thought to be simple stellar populations, but recent photometric and spectroscopic evidence suggests that the clusters' early formation history was more complicated. In particular, clusters show star-to-star abundance variations, and multiple sequences in their colour-magnitude diagrams. These effects seem to be restricted to globular clusters, and are not found in open clusters or the field.  In this paper, we combine the two competing models for these multiple populations and include a consideration of the effects of stellar collisions. Collisions are one of the few phenomena which occur solely in dense stellar environments like (proto-)globular clusters. We find that runaway collisions between massive stars can produce material which has abundances comparable to the observed second generations, but that very little total mass is produced by this channel. We then add the contributions of rapidly-rotating massive stars (under the assumption that massive stars are spun up by collisions and interactions), and the contribution of asymptotic giant branch stars. We find that collisions can help produce the extreme abundances which are seen in some clusters. However, the total amount of material produced in these generations is still too small (by at least a factor of 10) to match the observations. We conclude with a discussion of the additional effects which probably need to be considered to solve this particular problem. ",
        "introduction": "Globular clusters have long been viewed as the epitome of simple stellar populations. Their stars have a common age, a common distance, and a common metallicity; there is no interstellar gas and little else to get in the way of studying the stars directly. These systems are the closest we can come to a ``controlled experiment'' in stellar astrophysics, and as such they have proved incredibly valuable for studies of both stellar evolution and stellar dynamics in the past.  However, in recent years, cracks have been appearing in this simple picture. Both photometric and spectroscopic studies of clusters have started to unearth puzzles and inconsistencies. At first, these problems were thought to be an oddity in one particular cluster, or evidence of some details of stellar physics that we didn't quite understand. Over the last five years or so, however, it is becoming clear that our picture of a globular cluster needs to change. They cannot have formed instantly out of a single molecular cloud, removing all their leftover gas immediately, and then evolved passively for the next 10 billion years. Their history is more complicated.  \\subsection {Observational Background}  Hints that something strange was going on in globular clusters came first from spectroscopic studies of their red giants. For an excellent review, see \\citet{2004ARA&A..42..385G}. The general results have not changed since that review was written, although we now have observations of more stars per cluster, and more stars observed with high-resolution spectra. Most globular clusters have a constant iron and iron-peak element abundances, with the notable exception of $\\omega$ Centauri \\citep{1975ApJ...201L..71F} and hints of a very small spread in M22 and M92 \\citep{2009A&A...505.1099M, 1998AJ....115..685L}. However, it has been known since the late 1970s that lighter metals, particularly carbon, nitrogen and oxygen, do vary from star to star in many clusters \\citep{1978ApJ...223..487C}.  Other light elements also show star-to-star variations in clusters, including Na, Al and Mg.  The most striking piece of evidence to date that this abundance variation phenomenon occurs in all clusters is the spectroscopic study of sodium and oxygen by \\citet{2009A&A...505..117C,2009A&A...505..139C}. A general anti-correlation is seen, with oxygen-depleted stars having higher sodium abundances.  Aluminum and magnesium have been studied in fewer clusters, but they also show similar trends \\citep{1996AJ....112.1517S,2009A&A...505..139C}, with a large range in aluminum, a smaller spread in magnesium, and some evidence for high-aluminum stars having lower magnesium abundances.  \\citet{2008arXiv0811.3591C} also looked for correlations between the extent of the Na-O anti-correlation and cluster properties. The strongest correlation was between the extent of the correlation and the maximum temperature of stars on the horizontal branch, in the sense that clusters with very hot (or blue) horizontal branches had a large spread in Na and O. They also found a weaker trend with the total mass (or magnitude) of the cluster, and a trend with galactic orbit, in the sense that clusters which have spent more of their lifetime in the outskirts of the halo and not interacting with the galaxy, have a larger spread in abundances.  One of the main reasons that globular clusters have been so useful to stellar astrophysics is that their colour-magnitude diagrams are very clean. Other than the blue straggler stars and a few other unusual objects, the stars fall onto a single isochrone \\citep[see e.g.][]{2007AJ....133.1658S}. Again, the exception has been $\\omega$ Centauri. Its colour-magnitude diagram shows a large amount of structure beyond a single age/composition isochrone. The spread of the giant branch of this cluster was understood in the context of the measured differences in iron abundance of the stars, and so it did not come as much of a surprise to the community when the main sequence of this cluster was found to contain more than one sequence as well \\citep{2004ApJ...605L.125B}. However, the most intriguing result from this detailed study of $\\omega$ Centauri was the determination that the iron abundance of the bluest main sequence was not, as one would expect, the lowest in the cluster, but was in fact the highest \\citep{2005ApJ...621..777P}. The only possible way to reconcile the spectroscopic abundances with the photometric information was to infer a high helium content for these stars, perhaps even as high as Y=0.4. Because the abundance anomalies are found in light elements only but not iron, researchers have been casting this as a problem of 'pollution' from an early generation of stars in the cluster.  The globular cluster community finally came to realize that multiple populations were ubiquitous when \\citet{2005ApJ...631..868D} discovered an intrinsic spread in the main sequence NGC 2808, which was later confirmed to be {\\em three} separate main sequences \\citep{2007ApJ...661L..53P}. The turnoff region in this cluster is quite tight without much evidence for a spread, but starting about one magnitude below the turnoff, the sequences become obviously separated.  At about the same time, we saw the first CMDs showing multiple subgiant branches (e.g.the ACS observation of NGC 1851 \\citep{2008ApJ...673..241M}). Similar observations have since been seen in M22, NGC 6388, and M54 \\citep{2009arXiv0902.1422P} and many intermediate age clusters in the LMC and SMC \\citep{2008ApJ...681L..17M,2009A&A...497..755M}. NGC 1851 is an interesting case, is that it does {\\em not} show any evidence for a splitting of the main sequence, even in very careful observations including proper-motion cleaning of the CMD. It does, however, show a split in its giant branch when observed in $U-I$ \\citep{2009ApJ...707L.190H} which is not seen in $V-I$, highlighting the importance of observing in multiple bands.  Understanding the horizontal branch (HB) morphology in globular clusters has been a problem for decades, commonly referred to as the `second parameter problem.' Because the position of HB stars in the CMD is sensitive to both helium and metal abundances, they are a very useful population with which to discuss both pollution and multiple populations \\citep[e.g.][]{2008MNRAS.390..693D}.   \\subsection{Possible Explanations}   Many of the models to date have attempted primarily to understand the source of the pollution. The work of \\citet{2007A&A...470..179P} showed that the general abundance patterns could be explained by invoking the products of hot hydrogen burning, at a temperature of 70-80 $\\times 10^6$ K. At these temperatures, the hydrogen is burned in a series of cycles -- the standard CNO cycle, as well as the neon-sodium cycle and the aluminum-magnesium cycle. Conveniently, the second two cycles work to give the abundance patterns that are needed: higher sodium, lower Mg and higher Al. This processed material must be removed from the star before helium burning can proceed. Otherwise, the helium will be converted to carbon and oxygen. We will not have the large amount of helium that is required at the surface, and the constraint of (almost) constant C+N+O abundance will be violated.  A number of polluters have been proposed. The leading contender is a population of intermediate-mass (3-10 \\msun) asymptotic giant branch (AGB) stars \\citep[for a review, see][]{2008MNRAS.391..354R}. These stars can reach the appropriate high temperatures in their hydrogen burning shells, and the processed material is brought to the surface by the outer convection zone as it reaches into the burning shell during the hydrogen burning portion of the thermal pulse phase. AGB stars have strong but low-velocity winds, which means that any mass that is removed from the star has a good chance of remaining in the cluster. Indeed, there is recent evidence that giant stars in this mass range do show surface abundances which are consistent with the necessary pollution \\citep{2009A&A...504..845V}.  An alternative source of pollution is a population of fast-rotating massive stars \\citep{2007A&A...464.1029D}. Massive stars reach sufficiently high temperatures in their hydrogen-burning cores. These stars also have substantial winds, but under normal circumstances, the core regions are not exposed until very late in the stars' lives, well past the helium-burning phase. If a star is rotating rapidly, however, then meridional circulation and other rotational instabilities will mix material from the core to the surface, bringing these hydrogen-burning products up to a region where the wind can take them away from the star. Also, if the star is rotating rapidly enough that it is close to its break-up velocity, material can escape from the equator of the star in a slow outflowing disc or slow wind. This material in particular has a slow enough velocity that it will stay in the cluster, and there are even indications that low mass stars could form in situ in the disc.  Both AGB and fast-rotating massive star models have some difficulties in explaining the pollution of globular cluster gas and the formation of the second generation of stars. In both cases, the amount of material that is required to form the second generation is quite large. Observations of clusters such as NGC 2808 suggest that the mass of the second generation is approximately equal to the first (within a factor of 2-3 or so), which puts a limit on the amount of ejecta that the first generation must produce. Early versions of both models \\citep{2004ApJ...611..871D,2007A&A...464.1029D} suggested non-standard IMFs, heavily weighted towards the polluter in question, but these IMFs are difficult to justify. Current versions suggest that in fact the first generation needed to be significantly larger than what is observed today, and that the cluster needs to have lost approximately 90\\% of its first generation, while retaining all of the gas and the second generation of stars \\citep{2009A&A...499..835V,2008A&A...492..101D}. Both models also require a certain amount of dilution of the polluted ejecta with primordial gas, in order to match the observed abundances of light elements, including lithium.   The work of \\citet{2009ApJ...697..275P} looked at the effect of both enhanced helium abundance, and enhanced C+N+O abundance, on isochrones appropriate for globular clusters. They confirm that multiple main sequences are caused by a change in helium abundance at constant metal abundance (both iron and light elements). Multiple subgiant branches, on the other hand, are caused by a difference in total C+N+O abundance at constant helium, constant iron, and constant age. \\citet{2009ApJ...707L.190H} show that the split subgiant and giant branches of NGC 1851 are best explained in both $U-I$ and $V-I$ by a small change in both helium and metal abundance. The alternative explanation is a difference in age of the two populations of $\\sim$1 Gyr and constant abundances. It is, however, difficult to explain a delay in star formation of such a long time, and difficult to understand where the gas for the second generation came from.  Only a few groups have tried to put together the entire scenario. The most successful is the work of \\citet{2008MNRAS.391..825D}, which incorporated hydrodynamic simulations to study the flow of the AGB and supernovae ejecta in the vicinity of the proto-cluster. They find that the ejecta collects in a cooling flow and returns to the core of the cluster. They also look at a scenario in which pristine gas, which was pushed out of the cluster vicinity by the supernovae and massive stars, returns to the cluster after a few million years and mixes with the ejecta. Similar to \\citet{2007A&A...470..179P, 2007MNRAS.379.1431D, 2009A&A...499..835V}, they find that the combination of pristine gas and ejecta is necessary. They also find that much of the first generation in the cluster may be lost if the cluster is tidally limited.  Finally, both \\citet{2008MNRAS.391..825D} and \\citet{2008A&A...492..101D} use N-body dynamical models to study the subsequent evolution of this two-generation cluster and the mixing of the two populations. They concur with the general results of \\citet{2007ApJ...662..341D} that, at the current time, a two-generation cluster will not be dramatically different from a single-generation cluster in terms of the dynamics and distribution of stars. ",
        "conclusions": "In this paper, we take another theoretical look at the problem of multiple populations in globular clusters. In particular, we tried to address the observational constraint that this is specifically a {\\em dense old cluster} phenomenon. We concentrated on the effects of stellar collisions, which is one of the only physical mechanisms which affects clusters much more strongly than anywhere else in the universe. We also combined two previous models for multiple populations -- the fast-rotating massive stars and the asymptotic giant branch stars -- to determine if having more polluters would help with either abundances or with the total mass in the subsequent generations.  We included the effects of collisions in three ways: the impact of a runaway collision, an increase in the number of fast-rotating massive stars, and an increase in the number of intermediate mass AGB stars. We also allowed for two subsequent generations. One is formed within a few million years of the first cluster formation, from the ejecta of runaway collisions and fast-rotating massive stars. The other is formed approximately 100 million years later, from the ejecta of the AGB stars.   Our results are not substantially different from those of other groups who are modeling the early evolution of multiple population clusters. We still find that the mass produced by the polluters alone is insufficient. Loss of a large fraction of the initial cluster, dilution by primordial gas, or both, is required. However, stellar collisions do produce material with abundances which are interestingly extreme in all elements, and may well be a substantial contribution to the extreme populations in the most massive clusters.  It is clear, from this work and others, that we do not yet have a complete understanding of the multiple populations issue. Here, we try to highlight a number of the largest remaining problems, starting with those specific to our model and then discussing the puzzle more generally.     So far, there have been no simulations of star formation in the environment(s) envisioned here. Star formation typically occurs in cores in molecular clouds, an environment which is cold, dark, dense, and shielded from the outside universe. In the centre of a dense cluster, a cloud of ejecta will be subject to the radiation field of the first generation stars. No simulations of star formation in a region with this kind of external radiation bath have been done, and so we do not really understand the properties of the second generation. We have assumed a Salpeter mass function for our subsequent generations, for example, with a maximum mass of 120 \\msun. It may be reasonable to postulate that high mass stars could never form in the second generation (as was suggested by \\citet{2008MNRAS.391..825D}), which would mean that the second generations have more low-mass stars then in our toy model. For a Salpeter mass function between 0.1 and 120 \\msun, 44\\% of the mass is found in stars more massive than 0.8 \\msun. Under the assumption that only stars of mass 0.8 \\msun~ or less form in the subsequent generations, we can increase the predicted mass of these generations by approximately a factor of two.  If the cluster needs to be substantially more massive initially than assumed in this paper, then we should consider how the different polluter channels scale with cluster mass. For fast-rotating stars and AGB stars it is reasonable to assume that the yields scale with the total cluster mass. For the collision runaway this is not so clear. The mass of the collisional runaway as a function of the cluster mass has not been considered in detail in the literature. If the mass of the collision runaway scales with cluster mass, then it behaves as the other two polluter models. However, if the mass of the collision runaway depends more steeply on the mass of the cluster, then it will become relatively more important with increasing initial cluster mass. In that case the number of stars that need to be lost from the cluster can be smaller than when either of the other two scenarios alone are considered. If the total mass of the runaway is only a weak function of cluster mass the other polluter scenarios will become relatively more important for increasing cluster mass.     In these models, we have neglected the effects of binary stars. We know that globular clusters do have binaries, although work suggests that the fraction may be lower than in the field and open clusters \\citep{2008AJ....135.2155D, 1997ApJ...474..701R}. Binary stars can modify our toy model in a number of ways. Collisions between binary stars are more likely than between single stars \\citep{1989AJ.....98..217L} because of the larger cross section of the binary, and many of those interactions can result in more than two stars merging. Binaries can also increase the likelihood of having more than one runaway collision \\citep{2006ApJ...640L..39G}.  There are indications that massive stars have a higher binary fraction than low mass stars in the field \\citep{2006ApJ...640L..63L} and so it may not be unreasonable to expect that interacting binary stars may have a significant impact on the mass lost from massive stars. Following suggestions by \\citet{2009A&A...507L...1D}, \\citet{2009arXiv0909.3431V} goes as far to suggest that interacting massive binaries are responsible for all the pollution in globular clusters, not fast-rotating massive stars or AGB stars.  All models of multiple populations need to be refined substantially in order to explain the cluster-to-cluster variations. While it is clear that something is going on in almost every well-studied cluster, it is not clear that we understand how that effect depends on the cluster properties. In our scenario, we would argue that the cluster-to-cluster variations were caused by different collision rates early in the clusters' lives, perhaps driven by slightly different formation conditions and initial densities. In fact, it may be that we can use the extreme abundances produced in runaway collisions to determine which clusters hosted a runaway all those years ago. Only some clusters show evidence for extremely high Y values (Y $\\approx$ 0.4) and very extreme sodium-oxygen anti-correlations, such as $\\omega$ Centauri and NGC 2808. Perhaps only those clusters were formed with sufficiently high initial density to induce a runaway. A more detailed study of the abundance patterns in the individual pollution mechanisms may help disentangle these processes.  It has become clear that the epochs of globular cluster formation and very early evolution are crucial pieces of the puzzle when disentangling the present-day properties of these ancient objects. It is also clear that our understanding of the dominant processes and effects during these epochs is not as strong as we would like. Learning about the early lifetimes of dense clusters is not easy -- the stellar archaeology required is quite intricate. We are looking at small changes in current surface abundances and brightnesses of very old, low mass stars, and inferring a significant amount of action involving more massive stars over 10 billion years ago. However, we have learned a great deal about this phenomenon since it first was identified only a few years ago, and progress will continue to be made."
    },
    "1004/1004.4350_arXiv.txt": {
        "abstract": "The cosmological model of dark energy interacting with cold dark matter without coupling to the baryonic matter, is studied in the background of both classical Einstein and loop quantum cosmology. We consider two types of interacting models. In the former model, the interaction is a linear combination of the densities of two dark sectors, while in the latter model, the interaction with a constant transfer rate depends only on the density of cold dark matter. It is shown that the dynamical results in loop quantum cosmology are different from those in classical Einstein cosmology for both two kinds of interacting models. Moreover, the form of the interaction affects significantly the dynamical results in both kinds of cosmology. ",
        "introduction": "Recently, the discovery of the acceleration of cosmological expansion at present epoch has been the most principal achievement of observational cosmology. Numerous cosmological observations, such as Type Ia Supernovae (SNIa)~\\cite{price1}, Cosmic Microwave Background Radiation (CMBR)~\\cite{price2} and Large Scale Structure~\\cite{price3}, strongly suggest that the universe is spatially flat with about $4\\%$ ordinary baryonic matter, $20\\%$ dark matter and $76\\%$ dark energy. The accelerated expansion of the present universe is attributed to the dominant component of the universe, dark energy, which has a large negative pressure but not cluster. In fact, it has not been detected directly and there is no justification for assuming that dark energy resembles known forms of matter or energy. A large body of recent work has focussed on understanding the nature of dark energy. However, the physical origin of dark energy as well as its nature remain enigmatic at present.  The simplest model of dark energy is the cosmological constant $\\Lambda$~\\cite{price4}, whose energy density remains constant with time $\\rho_{\\Lambda}=\\Lambda/8\\pi G$ (natural units $c=\\hbar=1$ is used throughout the paper) and whose equation of state (defined as the ratio of pressure to energy density) remains $w=-1$ as the universe evolves. Unfortunately, the model is burdened with the well-known cosmological constant problems, namely the fine-tuning problem: why is the energy of the vacuum so much smaller than its estimation? and the cosmic coincidence problem: why is the dark energy density approximately equal to the matter density today? These problems have led many researchers to try different approaches to the dark energy issue. A possible method is to assume the equation of state (EoS) $w$ is a dynamical variable, and thus the dynamical scenario of dark energy is investigated. The most popular model among them is dubbed quintessence~\\cite{price5}. Besides, other scalar-field dark energy models have been studied, including phantom~\\cite{price6}, tachyon~\\cite{price7}, quintom~\\cite{price8}, ghost condensates~\\cite{price9}, etc. Also, there are other candidates, for example, Chaplygin gas which attempt to unify dark energy and dark matter~\\cite{price10}, braneworld model~\\cite{price11} and 5-dimensional gravity model~\\cite{phys1} which explain the acceleration through the assumption that spacetime has five dimensions instead of the usual four. In addition, since the cosmological scaling solution (i.e., the energy densities of dark energy and cold dark matter remain proportional) could probably alleviate the coincidence problem, interacting dark energy models are also proposed~\\cite{price12}.  As we all know, observations at the level of the solar system severely constrain non-gravitational interactions of baryons, namely, non-minimal coupling between dark energy and ordinary matter fluids is strongly restricted by the experimental tests in the solar systems~\\cite{phys23}, we therefore neglect this possibility. However, since the nature of dark sectors remains unknown, it is possible to have non-gravitational interactions between dark energy and dark matter. So we focus on dark energy interacting with dark matter alone.  Actually, many dark energy models are considered in the framework of classical Einstein cosmology. However, an outstanding problem in classical Einstein cosmology is the big bang singularity which is expected to be solved by quantum gravity. As a background independent quantization of general relativity, loop quantum gravity (LQG) is one of the best candidate theories of quantum gravity~\\cite{phys71}. It has been applied in cosmology to analyze our universe, known as Loop Quantum Cosmology (LQC)~\\cite{phys72}. In LQC, non-perturbative effects lead to $-\\rho^{2}/\\rho_{c}$ corrections to the standard Friedmann equation and thus allow us the possibility of resolving any past and future singularities~\\cite{phys72,phys73}. The modification becomes important when energy density of the universe becomes to be the same order of a critical density $\\rho_{c}$. When the correction term $-\\rho^{2}/\\rho_{c}$ dominates during the evolution of our universe, it will cause the quantum bounce and hence avoid the singularity. Recently, more and more researchers have taken their attention to LQC for the appealing features: avoidance of various singularities~\\cite{phys74}, inflation in LQC~\\cite{phys75}, large scale effect~\\cite{phys2} and so on. Concretely, some dark energy models are investigated in the background of LQC, such as phantom~\\cite{phys76}, coupling phantom~\\cite{phys77}, quintom and hessence~\\cite{phys78}, interacting dark energy model~\\cite{phys79}, etc.  In this paper, we study the dynamical evolution of two classes of interacting dark energy models in classical Einstein and Loop Quantum Cosmology. Here some questions naturally arise as follows. Can these models alleviate the coincidence problem in classical Einstein cosmology? Are there scaling solutions arising from the effect of loop quantum cosmology? Can the future singularities be resolved in LQC? By our analysis, it turns out that in the former model, there are two attractors in classical Einstein cosmology and LQC. One is an accelerated scaling solution and the other is a baryon dominant solution. However, in the latter model, there exists one attractor in classical Einstein cosmology, which is a dark energy dominated solution rather than a scaling solution, whereas in LQC all fixed points are unstable. Thus, there exists no scaling solution in the latter case, namely, this kind of interacting dark energy model can not be regarded as a candidate to alleviate the coincidence problem. Also, we find that dynamical results in LQC are different from those in classical Einstein cosmology for both two kinds of interacting models. Our universe finally enters an oscillating phase in LQC. Moreover, the oscillating frequencies are significantly different for varied parameters of models. These results are different from the those obtained in classical Einstein cosmology. Thus, LQC allow us the possibility of resolving future singularities. Hence, the quantum gravity effect may be manifested in large scale in the interacting dark energy models.  In Sec. II, we study dynamical properties for the general case in classical Einstein cosmology and LQC. Then the dynamical results of two types of interacting models are respectively studied in Secs. III and IV. In Sec. V, the numerical results are presented. Finally, the conclusions are summarized in Sec. VI. ",
        "conclusions": "In previous sections, we have studied the cosmological evolution of two interacting dark energy models in classical Einstein and Loop Quantum Cosmology. We consider two kinds of interaction term between dark energy and cold dark matter. Note that observations at the level of the solar system severely constrain non-gravitational interactions of baryons. So the baryonic matter solely satisfies the energy conservation equation. By our analysis, we find that dynamical results in LQC are different from those in classical Einstein cosmology for both two kinds of interacting models.  In the interacting model I, namely, $Q=3H(c_{1}\\rho_{m}+c_{2}\\rho_{d})$, there are two attractors in both classical Einstein cosmology and LQC. One is a baryon dominated solution and the other is an accelerated scaling solution. Since the same results are obtained in both classical Einstein cosmology and LQC for the baryon dominated attractor, we only focus on the accelerated scaling solution. Interestingly, we find that if $w_{d}>-1$, the stable region in LQC is the same as that in classical Einstein cosmology, while when $w_{d}<-1$, the stable region in LQC is smaller than that in classical Einstein cosmology. The total EoS $w$ approaches finally to a constant, which depends on EoS $w_{d}$ and the coupling constants $c_{1}$ and $c_{2}$, but is independent of the theory describing our universe. When we select the parameters in the unstable region in LQC, the universe experience bouncing, which can resolve the singularity problem. The bounce in scale factor occurs later for greater value of $w_{d}$, $c_{1}$ or $c_{2}$. Furthermore, the oscillating frequencies are distinct for different parameters.  However, in the interacting model II with $Q=\\Gamma\\rho_{m}$, there exists one attractor in classical Einstein cosmology for $w_{d}<-1$, which is a dark energy dominated solution rather than a scaling solution, whereas in LQC, all fixed points is not stable. Thus, there exists no scaling solutions in the interacting model II. So this kind of interacting dark energy model can not be regarded as a candidate to alleviate the coincidence problem in both classical Einstein and loop quantum cosmology. In classical Einstein cosmology, the final state $w_{*}$ is a constant, which equals to $w_{d}$ and is independent of the coupling constant $\\beta$. The bounce in scale factor occurs later for greater value of $w_{d}$ or $\\beta$. Our universe finally enters an oscillating phase in LQC. Moreover, the oscillating frequencies are significantly different for varied parameters.  In summary, the interacting model I may alleviate the coincidence problem in both classical Einstein and loop quantum cosmology, depending on the values of the parameters selected in the model. However, the interacting model II can not be regarded as a candidate to alleviate the coincidence problem in both kinds of cosmology. Thus, dynamical results are different not only in different theories describing the universe but also in different interacting models. In addition, the results that our universe finally enters an oscillating phase in LQC, which are different from the those obtained in classical Einstein cosmology, show that LQC allow us the possibility of resolving future singularities. Therefore, the quantum gravity effect may be manifested in large scale in the interacting dark energy models."
    },
    "1004/1004.0691_arXiv.txt": {
        "abstract": "TeV-mass dark matter charged under a new GeV-scale gauge force can explain electronic cosmic-ray anomalies.  We propose that the CoGeNT and DAMA direct detection experiments are observing scattering of light stable states --- ``GeV-Matter'' --- that are charged under this force and constitute a small fraction of the dark matter halo.  Dark higgsinos in a supersymmetric dark sector are natural candidates for GeV-Matter that scatter off protons with a universal cross-section of $5\\times 10^{-38} \\cm^2$ and can naturally be split by 10--30 keV so that their dominant interaction with protons is down-scattering.  As an example, down-scattering of an $\\OO(5)$ GeV dark higgsino can simultaneously explain the spectra observed by both CoGeNT and DAMA.  The event rates in these experiments correspond to a GeV-Matter abundance of $0.2\\!-\\!1\\%$ of the halo mass density. This abundance can arise directly from thermal freeze-out at weak coupling, or from the late decay of an unstable TeV-scale WIMP. Our proposal can be tested by searches for exotics in the BaBar and Belle datasets. ",
        "introduction": "For decades the WIMP --- a particle with TeV-scale mass that freezes out through weak-strength couplings --- has been the canonical dark matter candidate. The electronic cosmic ray excesses observed by PAMELA~\\cite{Adriani:2008zr}, Fermi~\\cite{Abdo:2009zk}, and HESS~\\cite{Collaboration:2008aaa,Aharonian:2009ah}, along with the discovery of the WMAP and Fermi ``haze''\\cite{Finkbeiner:2003im,Finkbeiner:2004us,Hooper:2007kb,Dobler:2007wv,Dobler:2009xz}, have both bolstered and complicated the WIMP picture.  Together these experiments suggest that WIMP dark matter either annihilates or decays into lepton-rich final states (see e.g.~\\cite{Cirelli:2008pk}), with a rate 10-1000 times larger than expected from $s$-wave annihilation.  Dark matter charged under a new GeV-scale $U(1)_D$ `dark force' that kinetically mixes with hypercharge can explain these results through either Sommerfeld-enhanced annihilation \\cite{Finkbeiner:2007kk, ArkaniHamed:2008qn, Pospelov:2008jd, Cholis:2008qq, Cholis:2008wq} or decay \\cite{Ruderman:2009tj} into light dark-sector states, which are in turn forced by kinematics to decay into leptons.  Such a GeV-scale sector is naturally realized in models with supersymmetry-breaking near the weak scale, where $D$-term mixing dynamically generates a GeV-scale mass for the $U(1)_D$ gauge boson $A'$  \\cite{Dienes:1996zr,Cheung:2009qd,Katz:2009qq,Morrissey:2009ur,Cui:2009xq} (summarized in Section \\ref{sec:model}).  The spontaneous breaking of this $U(1)_D$ gauge symmetry requires a higgs sector, and this sector can naturally contain stable GeV-scale states such as dark higgsinos. These and other stable ``GeV-Matter'' particles charged under $U(1)_D$ interact with protons with a cross-section \\be \\label{DDCrossSection} \\sigma_{\\chi, p} \\approx 5\\times 10^{-38} {\\rm~cm}^2 \\ee that is {\\bf completely fixed by Standard Model couplings}~\\cite{Cheung:2009qd}.   Because this cross section is so large, even a small abundance of GeV-Matter can be observed at direct detection experiments.  In fact, we should expect GeV-Matter states to comprise only a small fraction of the dark matter halo because their annihilation cross-section is large ($\\sim \\alpha^2/m_{\\mathrm{GeV}}^2$).  Two experiments have reported anomalous events that may arise from scattering of a light dark matter species. DAMA \\cite{Bernabei:2000qi,Bernabei:2005hj,Bernabei:2008yi,Bernabei:2010mq} has reported a $9 \\sigma$ annual modulation signal at 2--5 keVee in their NaI crystal detector. Early this year, the CoGeNT collaboration reported about 100 events from an unknown source in a low-threshold, high-resolution Ge detector \\cite{Aalseth:2010vx}; these events are consistent with light dark matter scattering with a mass and cross-section similar to those needed to explain the DAMA modulation.  As shown in Figure \\ref{fig:DD1} (see Section~\\ref{sec:DirectDetection} for details), both signals can be  explained by 3--12 GeV dark matter scattering with a cross-section $\\sigma_{\\rm eff} \\sim (1-5)\\times 10^{-40}$ cm$^2$, or by a fraction of the halo scattering with a larger cross-section (see also \\cite{Aalseth:2010vx,Fitzpatrick:2010em,Andreas:2010dz}).  If the events reported by these experiments arise from light dark matter scattering, they provide further evidence for a new sector at the GeV scale.  In light of the cross section (\\ref{DDCrossSection}), dark sector higgsinos at 3-12 GeV naturally explain the CoGeNT and DAMA excesses if they comprise a fraction \\be f_{\\rm light} \\equiv \\frac{\\rho_{\\rm light}}{\\rho_{\\rm DM}} \\approx (2 {\\rm - } 10)\\times 10^{-3} \\label{flight} \\ee of the halo mass density.   We propose two natural origins for an abundance of this size: \\begin{description} \\item[Thermal freeze-out] generates an abundance (assuming $s$-wave annihilation) \\be \\label{eq:annscaling} f_{\\rm light} \\sim 3 \\times 10^{-3} \\parf{m_{\\rm light}}{6 \\GeV}^2 \\parf{1/150}{\\alpha}^2, \\ee which yields the desired GeV-Matter relic density for small couplings, as we discuss in Section \\ref{sec:annihilation1}. \\item[Late Decays] of TeV-mass WIMP-sector states generate a density \\be \\frac{n_{\\rm light}}{n_{\\rm heavy}} \\sim 1, \\qquad f_{\\rm light} \\sim \\frac{m_{\\rm light}}{m_{\\rm heavy}} \\frac{\\rho_{\\rm heavy}}{\\rho_{\\rm DM}} \\sim \\rm{few} \\times 10^{-3} \\ee of GeV-Matter that can easily dominate over its thermal abundance. This sharp prediction is insensitive to the details of the decay process, and agrees with the halo fraction required to explain the CoGeNT and DAMA signals!  The decaying TeV-mass particle can be the scalar superpartner of a stable fermionic WIMP, as discussed in Section \\ref{sec:decay}. \\end{description} We therefore propose a new unified framework for interpreting both the direct detection and cosmic ray anomalies from the same dynamics.  \\begin{figure*}[!] \\begin{center} \\includegraphics[width=.485\\textwidth]{cogentAndDAMAV2.pdf} \\; \\includegraphics[width=.485\\textwidth]{cogentAndDAMAbmQQ25V2Paper.pdf} \\caption{ {\\bf Left:} Regions of scattering cross-section per nucleon times GeV-Matter mass fraction in the halo versus mass favored by CoGeNT (filled regions), DAMA with channeling (solid-line contours), and DAMA without channeling (dashed contours) for light dark matter scattering coherently off protons.  We show favored regions (from right to left) for elastic scattering (green) and inelastic splittings $|\\delta|=15$ keV (orange), $25$ keV (red), and $35$ keV (blue), with equal fractions of up- and down-scattering.  The contours correspond to ``$1.6\\sigma$,'' neglecting systematic uncertainties. The black dot and upside-down triangle correspond to benchmark points we consider in Figure \\ref{fig:DAMACogentfit}. {\\bf Right:} A close-up of the $|\\delta|=25$ keV allowed region with ``$3\\sigma$'' contours added and constraints from XENON10 (black lines) and CDMS-Si (gray lines).  Thick (thin) lines denote 5.0 (2.3) events expected, again neglecting systematics.  More details are given in \\S\\ref{sec:DirectDetection}.} \\label{fig:DD1} \\end{center} \\end{figure*}  $U(1)_D$-breaking mass splittings of ${\\cal O}(\\mbox{few keV}-\\mbox{few GeV})$ for both WIMPs and GeV-Matter arise from generic interactions with TeV-scale states.  In the presence of such mass splittings, dark matter scatters through inelastic up- and down-scattering, with elastic processes highly suppressed.  These splittings have two important effects: when the GeV-Matter states are split by ${\\cal O}(10\\mbox{ keV})$, direct detection is dominated by down-scattering, which improves the agreement between DAMA and CoGeNT and permits lower masses of the light dark matter than in the elastic case, as discussed in Section \\ref{sec:DirectDetection} and more generally in \\cite{Graham:2010ca}.  A larger $\\gtrsim$ MeV splitting of the TeV-scale WIMP allows it to explain the INTEGRAL 511 keV excess and renders it practically invisible to direct detection experiments \\cite{Strong:2005zx, Finkbeiner:2007kk}.  The success of these models is particularly striking in light of the difficulty of generating such high direct detection rates through Standard Model processes, even when a large abundance of light dark matter is generated through an asymmetry \\cite{Kaplan:2009ag}.  Dark matter coupled to the $Z$ boson can explain the reported rate, but is quite constrained by LEP measurements of the $Z$ invisible width~\\cite{Kuflik:2010ah}.  Standard Model Higgs exchange, in contrast, leads to typical cross-sections $\\lesssim 10^{-43} \\cm^2$ (see e.g.~\\cite{Giedt:2009mr}) because the Higgs couples only weakly to nuclei.  Explanations of the CoGeNT signal through scalar interactions posit a new light scalar that mixes significantly with the Higgs, but remains relatively light through a percent-level fine-tuning \\cite{Fitzpatrick:2010em,Andreas:2010dz}. In contrast, the low mass of our dark matter and mediator arise from simple dynamics.  \\subsection*{Implications for $B$-factories}  Analysis of existing low energy collider data is crucially important for testing the GeV-Matter scenarios in this paper (see e.g.~\\cite{Essig:2009nc,Batell:2009yf,Bossi:2009uw,AmelinoCamelia:2010me,Freytsis:2009ct,McDonald:2010iq}). The strongest constraint on $U(1)_D$ gauge boson decays to Standard Model matter is from a BaBar search for a resonance in $\\gamma \\mu^+\\mu^-$ in $\\Upsilon(3s)$ data \\cite{Aubert:2009cp}, which is sensitive to the ``radiative return'' process $e^+e^- \\rightarrow \\gamma A'$. The parameter space consistent with the GeV-Matter explanation of the CoGeNT and DAMA excesses is just below the limit from this search, but {\\bf can be tested by a search in the full Belle and BaBar datasets}, which together comprise about a factor of 30 higher statistics than the $\\Upsilon(3s)$ data. The ``higgs$'$-strahlung'' process \\cite{Batell:2009yf} may also be accessible at $B$-factories, leading to a striking six-lepton final state. Searches at $J/\\psi$- and $\\phi$-factories can also explore the low-mass parameter space \\cite{Yin:2009mc,Li:2009wz,Bossi:2009uw,AmelinoCamelia:2010me}.  \\subsection*{Down-Scattering and Lighter Dark Matter} Three flaws have been noted with elastically scattering light dark matter as an explanation of the DAMA annual modulation signal: the predicted modulation spectrum grows exponentially at low energies while the DAMA spectrum appears to fall near threshold \\cite{Chang:2008xa}.  Moreover, the mass range favored by both the DAMA and CoGeNT spectra, $\\sim \\! 10$ GeV, is disfavored by  XENON10 \\cite{Angle:2009xb, Manzur:2009hp, Angle:2007uj} and especially CDMS Silicon \\cite{CDMS:talk} data (see \\cite{Fitzpatrick:2010em}).  Finally, the scattering rates required to explain the CoGeNT and DAMA signals differ by about a factor of 5 (though this final discrepancy depends on the precise channeling properties of Iodine \\cite{Fitzpatrick:2010em,Drobyshevski:2007zj, Bernabei:2007hw, MarchRussell:2008dy}).  However, there is no reason to expect the light dark matter scattering to be purely elastic.  Mass splittings of order 10-30 keV for GeV-Matter are readily generated by TeV-suppressed higher-dimension operators.  Because the dark sector kinetically decouples from the Standard Model at temperatures high compared to this splitting, both states will be equally populated and the heavy state will be cosmologically long-lived \\cite{Finkbeiner:2009mi, Alves:2009nf, Alves:2010dd}.  The excited states interact with nuclei through down-scattering, which pushes the nuclear recoil spectrum to higher energies for a given dark matter mass.  As we discuss in \\S\\ref{sec:DirectDetection}, this effect improves the agreement of light dark matter with DAMA modulation data (relative to the elastic case), brings the CoGeNT and DAMA expected cross-sections together, and is consistent with the XENON10 null search.  The residual tension of this scenario with the 5-tower CDMS Silicon analysis~\\cite{CDMS:talk} is extremely sensitive to the experimental selection efficiency near threshold and small energy-scale systematics.  There is also tension with the DAMA unmodulated spectrum in \\cite{Bernabei:2008yi}.  A general analysis of down-scattering light dark matter \\cite{Graham:2010ca} also finds a lower-mass region consistent with DAMA and CDMS Silicon results. ",
        "conclusions": "If supersymmetry stabilizes the weak scale, then kinetic mixing with a new gauge force can naturally generate a GeV-scale sector. This remarkable possibility could open the door to exploring supersymmetry at low energies. In simple supersymmetric dark sectors, stable GeV-Matter can arise in the dark higgs sector, and these may comprise a portion of the dark matter with new and interesting phenomenology. This possibility is well-motivated by cosmic ray data suggestive of GeV-scale forces.  In this paper, we have proposed simple models of GeV-Matter within a supersymmetric dark sector that couples to the Standard Model through kinetic mixing. The low-energy effective Lagrangian we have studied has already appeared in the literature to realize inelastic up-scattering of $\\OO(100 \\GeV)$ mass dark matter as an explanation of the DAMA signal. Our interpretation of the data is different: \\begin{itemize} \\item The CoGeNT and DAMA direct detection anomalies are simultaneously explained by inelastic {\\it down-scattering} of $\\OO(\\GeV)$ mass dark-sector states (GeV-Matter), for which the higgsinos of a supersymmetric dark sector are an excellent candidate. These states comprise a very sub-dominant fraction of the dark matter. Tension with existing experiments is reduced, and the CoGeNT and DAMA rates readily agree --- in contrast to elastic scattering explanations. \\item Electroweak symmetry breaking triggers breaking of the dark sector $U(1)_D$, giving the vector multiplet a mass $m_{A'}\\sim \\sqrt{\\epsilon g_D}m_W\\sim \\GeV$. The  direct detection cross-section of GeV-Matter is predicted (equation~\\eqref{eq:DtermScatter}) in terms of Standard Model parameters with no dependence on $\\epsilon$ or  $\\alpha_D$.  \\item The correct CoGeNT and DAMA event rates are obtained for GeV-Matter comprising 0.2--1\\% of the halo mass density.  This abundance can be explained by thermal freeze-out, and the couplings required for this scenario are readily tested with existing $B$-factory data. \\item Another possibility is that a metastable WIMP decays to GeV-Matter after thermal GeV-Matter annihilation has frozen out. In this case, $\\frac{\\rho_{\\rm light}}{\\rho_{\\rm WIMP}}\\approx \\frac{m_{\\rm light}}{m_{\\rm WIMP}}$ robustly produces the desired relic abundance for metastable WIMP masses of 2--3 TeV.  We have exhibited a model where the metastable WIMP is the bosonic superpartner of the TeV-mass dark matter whose decays can explain the cosmic ray data. \\end{itemize}  Thus, in our unified and predictive framework, the GeV-scale dark sector plays a crucial role in explaining cosmic ray excesses {\\it and} contains stable GeV-Matter responsible for direct detection signals. The most important test of our interpretation of DAMA and CoGeNT will come from further searches with direct detection experiments. In particular, further studies using CDMS Silicon and Germanium data, especially with lower threshold energy, will be vital. The power of these experiments depends sensitively on their threshold energies and signal efficiencies.  Future studies should include these uncertainties when quoting constraints.  Low-energy flavor factories play an equally important role in testing our proposal.  A resonance search in $e^+e^-\\rightarrow \\gamma \\mu^+\\mu^-$ in the full BaBar or Belle datasets and higher lepton-multiplicity searches should be powerful enough to discover or exclude the light vector $A'$ over most of the parameter space that can explain CoGeNT and DAMA.  Additional searches in $J/\\Psi$-factories \\cite{Yin:2009mc}, $\\phi$-factories \\cite{Bossi:2009uw, AmelinoCamelia:2010me}, and fixed-target experiments \\cite{Bjorken:2009mm, Reece:2009un, Batell:2009di, Essig:2010xa, Freytsis:2009bh, Heinemeyer:2007sq, Schuster:2009au} can provide evidence for (or constrain) a GeV-scale gauge sector coupled through kinetic mixing.  Finally, if Standard Model super-partners are light enough to be directly produced at colliders, then supersymmetric dark sectors can give rise to spectacular lepton jet signals at the Tevatron or LHC \\cite{ArkaniHamed:2008qp,Baumgart:2009tn, Cheung:2009su, Abazov:2009hn}.  \\vspace{5mm} \\noindent {\\bf Acknowledgements} \\vspace{3mm}  We thank Clifford Cheung, Liam Fitzpatrick, Michael Peskin, and Jay Wacker for many useful discussions.  NT thanks Peter Graham, Roni Harnik, Surjeet Rajendran, and Prashant Saraswat for early discussions of inelastic down-scattering at DAMA. We thank Spencer Chang, Jia Liu, Aaron Pierce, Neal Weiner, and Itay Yavin for bringing to our attention their related work \\cite{Chang:2010yk}.  We also thank N. Weiner for alerting us to a quenching factor error in the first version of this paper, and for providing us with \\cite{cogentNote}.  RE, JK, and PS are supported by the US DOE under contract number DE-AC02-76SF00515.   \\appendix"
    },
    "1004/1004.5411_arXiv.txt": {
        "abstract": "The coalescence of a supermassive black hole binary (SMBHB) is thought to be accompanied by an electromagnetic (EM) afterglow, produced by the viscous infall of the surrounding circumbinary gas disk after the merger.  It has been proposed that once the merger has been detected in gravitational waves (GWs) by the {\\it Laser Interferometer Space Antennae} ({\\it LISA}), follow-up EM observations can search for this afterglow and thus help identify the EM counterpart of the {\\it LISA} source.  Here we study whether the afterglows may be sufficiently bright and numerous to be detectable in EM surveys alone.  The viscous afterglow is characterized by an initially rapid increase in both the bolometric luminosity and in the spectral hardness of the source.  For binaries with a total mass of $10^{5}-10^{8}\\Msol$, this phase can last for years to decades, and if quasar activity is triggered by the same major galaxy mergers that produce SMBHBs, then it could be interpreted as the birth of a quasar.  Using an idealized model for the post-merger viscous spreading of the circumbinary disk and the resulting light curve, and using the observed luminosity function of quasars as a proxy for the SMBHB merger rate, we delineate the survey requirements for identifying such birthing quasars.  If circumbinary disks have a high disk surface density and viscosity, an all-sky soft X-ray survey with a sensitivity of $F_{\\rm X}\\ltsim 3\\times 10^{-14}~{\\rm erg~s^{-1}~cm^{-2}}$ which maps the full sky at least once per several months, could identify a few dozen birthing quasars with a brightening rate $d\\ln F_{\\rm X}/dt > 10\\% \\yr^{-1}$ maintained for at least several years.  If $>1\\%$ of the X-ray emission is reprocessed into optical frequencies, several dozen birthing quasars could also be identified in optical transient surveys, such as the {\\it Large Synoptic Survey Telescope}. Distinguishing a birthing quasar from other variable sources may be facilitated by the monotonic hardening of its spectrum, but will likely remain challenging.  This reinforces the notion that observational strategies based on joint EM-plus-GW measurements offer the best prospects for the successful identification of the EM signatures of SMBHB mergers. ",
        "introduction": "Observational evidence robustly indicates that all or nearly all galaxies harbor a supermassive black hole in their nucleus \\citep[SMBH; e.g.,][]{Maggor+98}.  Since cosmological structure formation models predict a hierarchy of galaxy mergers, if nuclear SMBHs were indeed common at earlier times, then these mergers should result in the formation of SMBH binaries \\citep[SMBHBs;][]{Begelman+80}, and these binaries should then be common throughout cosmic time \\citep{Haehnelt94, Menou+01, VHM03, WL03, Sesana+07, Lippai+09, TH09}.  It has also long been known, both observationally (e.g., \\citealt{Sanders+88}) and theoretically (e.g., \\citealt{BarnesHernquist+91}) that galaxy mergers can drive gas to the nucleus of the merger remnant, which could facilitate the merger of the nuclear SMBHs on one hand, while also providing fuel for quasar activity on the other.  Mergers are therefore generically also believed to trigger quasar activity; the rate of major galaxy mergers can indeed provide an explanation for the observed evolution of the quasar population as a whole (\\citealt{Carlberg+90}; for more recent work, see, e.g., \\citealt{Hopkins+07a} and \\citealt{WyitheLoeb+09} and references therein).  Despite their expected ubiquity, observational evidence for SMBHBs is scarce, and the precise timing of any quasar activity, and when it occurs relative to the merger of the nuclear SMBH binary, remains unclear \\citep{Kocsis+06}. A handful of pairs of active SMBHs in the same galaxy have been resolved directly, at $\\sim$kpc separation in X--ray \\citep{Komossa+03} and optical \\citep{Comerford+09} images, and at $\\sim$10pc separation in the radio \\citep{Rodriguez+06}, confirming that gas is present around the SMBH binary, and that quasar activity can, at least in some systems, commence prior to their coalescence. However, there has been at least one suggestion that luminous activity can be occurring later, at the time of the merger, as well -- momentarily interrupted by the coalescence of the SMBHs and reactivated after-wards \\citep{Liu+03}.  While there are many more observed SMBHB candidates with small separations \\citep[e.g.,][]{Roos+93, Schoen+00, Merritt+02, Sudou+03, Liu04, Boroson+09}, the evidence for these tighter binaries is indirect, and each candidate system has alternative explanations.  The expectation is that at large separations, the binaries rapidly lose orbital angular momentum through dynamical friction with background stars and through tidal--viscous interaction with the surrounding gas \\citep{IPP99, AN02, Escala+05b, MM05, Dotti+07, SHM07, Cuadra+09, Callegari+09, Colpi+09, HKM09, Chang+09}.  Once sufficiently compact, gravitational wave (GW) emission rapidly shrinks the orbit, culminating in a merger.  How long this process lasts, and at what stage(s) the SMBHs light up as luminous quasars, is, however, also poorly understood theoretically.  Apart from the cosmological context, interest in EM signatures of SMBH mergers surged recently \\citep[e.g.,][]{MP05, BP07, Lippai+08, SK08, SB08, O'Neill+09, Chang+09, Megevand+09, Corrales+10, Rossi+10, Anderson+10, Krolik10, TM10, Shapiro10}, driven by (i) the prospect that the {\\it Laser Interferometer Space Antennae}\\footnote{http://lisa.nasa.gov/} ({\\it LISA}) will detect the mergers in GWs and provide a tractable list of (perhaps as low as a few hundred; e.g., \\citealt{Kocsis+08}) EM candidates for SMBHBs and (ii) the breakthrough in numerical general relativistic calculations of BH mergers (e.g., \\citealt{Pretorius05, Campan+06, Baker+06}), which led to robust predictions of significant mass--loss and recoil that can significantly perturb the ambient gas.  A simultaneous observation of the merger in gravitational and EM waves would enable new scientific investigations in cosmology and BH accretion physics \\citep{Cutler98, HH05, Kocsis+06, Kocsis+07, LH08, Phinney09, Bloom+09}.  In this paper, we focus on one particular signature of SMBHB coalescence, which we will hereafter refer to as the ``viscous afterglow''.  The physics of this model was discussed by \\citet{Liu+03} in the context of the interruption of jets in double-double radio galaxies, and later by \\citeauthor{MP05} (2005; hereafter MP05) in the context of EM counterparts of {\\it LISA} sources.  Prior to merger, the SMBHB torques open and maintain a cavity in the center of a thin circumbinary gas disk \\citep{ArtymowiczLubow94}.  When the binary becomes sufficiently compact, GW emission causes the binary orbit to shrink faster than the gas just outside the cavity can viscously respond.  The merger takes place inside the cavity, which is subsequently filled as the disk viscously spreads inward.  Because the refilling inner disk produces higher-energy photons than the outer regions, the disk is predicted to transition from an X-ray-dim state to an X-ray-bright one, with its bolometric luminosity increasing by a factor of $\\sim 10$ during this time.  This transition is expected to take place on humanly tractable timescales, with the cavity filling in $\\sim 10 (1+z) (M/10^{6}\\Msol)^{1.3} \\yr$, where $M$ is the total mass of the binary. A study of an optically selected sample by \\cite{Gibson+08} found X-ray-dim AGN to be rare ($\\ltsim 2\\%$ at $z\\sim2$), suggesting that it would be tractable to catalog and monitor such systems for possible observational signatures of a merger afterglow.  In the observational scenario originally proposed by MP05, {\\it LISA} would detect the GWs from the merger and determine its approximate location in the sky to within $\\sim 0.1$ deg, triggering a follow--up search to identify the EM counterpart and host galaxy.  A natural question to ask, however -- and the subject of the present paper -- is {\\em whether the viscous afterglows may be sufficiently bright and numerous to be detectable in EM surveys alone, even before {\\it LISA} is launched}.  The identification of mergers by their EM signatures alone could, in fact, be valuable for several reasons.  First, {\\it LISA} will be sensitive to GWs from relatively low--mass SMBHBs, with total masses of $\\sim (10^{4}-10^{7})/(1+z)\\Msol$.  EM studies could, in principle, detect coalescing SMBHBs outside this mass range, and therefore complement the {\\it LISA} binary population.  Second, while many models for the cosmological evolution of SMBHs predict that {\\it LISA} will detect dozens or hundreds of mergers (if ``seed'' black holes are abundant and merge often; e.g., \\citealt{Sesana+07}), there are some SMBH assembly scenarios that may result in far fewer {\\it LISA} events (i.e. if seeds are rare and grow primarily through rapid accretion or are very massive already at formation; \\citealt{TH09,Lippai+09}).  It is therefore plausible that EM surveys could deliver a larger SMBH binary sample than available from GWs. Third, several transient EM surveys are already under way, or are planned to be completed before the expected launch date of the {\\it LISA} mission around 2020.  If luminous quasar activity is triggered by major mergers of galaxies, as argued above, then the viscous afterglow could plausibly be interpreted as the signature of the birth of a quasar.  In this paper, we estimate the number of identifiable afterglow sources, i.e. birthing quasars, in the sky, by (i) adopting an idealized time-dependent model \\citep[][hereafter \\citetalias{TM10}]{TM10} of the evolution of the disk structure, to calculate photometric light curve and variability of the afterglow, and (ii) by using the observed luminosity function of quasars as a proxy for the SMBHB merger rate. Our two main goals are:  \\begin{enumerate}  \\item To assess whether there is any hope of detecting and identifying the viscous afterglows with conventional EM telescopes alone.  \\item To see how the identifiability of the afterglows depends on theoretical parameters and to delineate the ideal survey attributes (wavelength, angular coverage and depth).  We compare the derived attributes to those similar to planned large surveys of the transient sky: a soft-X-ray survey with specs similar to those that were proposed recently, unsuccessfully, for the the {\\it Lobster-Eye Wide-Field X-ray Telescope} \\footnote{http://www.star.le.ac.uk/lwft/} ({\\it LWFT}) mission; and the {\\it Large Synoptic Survey Telescope}\\footnote{http://www.lsst.org/lsst} ({\\it LSST}) in the optical. \\end{enumerate}  We find that the detectability of the afterglow is sensitive to the properties of the circumbinary disk, in particular to the ratio of the viscous stress to the gas pressure, and to the surface density of the disk.  We conclude that purely EM identification of the afterglows by the planned surveys are unlikely, unless the surface density and the viscosity in the circumbinary disk are at the high end of the expected range.  In this latter, optimistic scenario, several dozen birthing quasars could be identified in a soft X-ray transient survey.  We also find that if $\\gsim 1\\%$ of the X-ray radiation emitted in the central regions is reprocessed into the optical frequencies by dust surrounding the source, or by warps or geometric splaying in the disk itself \\citepalias{TM10}, several dozen afterglows could be detected in an optical transient survey, such as {\\it LSST}.  This paper is organized as follows.  In \\S~2, we summarize the viscous afterglow model, and describe our methods for estimating the identifiable population of AGN harboring a recently merged SMBHB.  In \\S~3, we present estimates for the number of identifiable afterglow sources in the sky.  We summarize our results and offer our conclusions in \\S~4. ",
        "conclusions": "Using an idealized model for the population of coalescing SMBHBs, and for the light curve of the afterglow produced by the viscously spreading post-merger circumbinary disk, we have shown that ongoing afterglows of SMBHB mergers may be present in the data sets of wide X-ray and optical surveys.  In soft X-ray bands, this requires that the surface density and the viscosity in the circumbinary disk be at the high end of the expected range, while afterglows could only be found in optical surveys if the X-ray emission is promptly and significantly reprocessed into optical frequencies.  Despite the highly approximate nature of our analysis and other model uncertainties, our calculations provide a proof-of-concept for a very general hypothesis: {\\it SMBHB mergers may exhibit identifiable, steady brightening rate for a period of the order of decades, and such afterglows could be detected serendipitously in a large survey that revisits the sky at least every few months for several years.}  Our more specific findings can be summarized as follows:  \\begin{itemize}  \\item For optimistic parameter values, several birthing quasars, brightening by at least $d\\ln L_{\\rm X}/dt_{\\rm obs}>30\\% \\yr^{-1}$ for several years, could be identified in the 0.1 - 3.5 keV soft X-ray band by an all-sky survey with specifications comparable to those proposed for the {\\it LWFT} mission.  \\item At any given time, there could be up to $N_{\\rm ag}\\sim 100$ sources in the sky that exhibit a brightening at or above $d\\ln L_{\\rm X}/dt_{\\rm obs}>10\\% \\yr^{-1}$, with soft X-ray luminosities $L_{\\rm X}\\gtrsim 10^{42}\\erg\\,\\s^{-1}$.  The most luminous sources typically spend $t_{\\rm ag, obs}\\gtrsim 10\\yr$ in this state, and thus can be monitored on humanly tractable timescales.  These numbers depend weakly on most system parameters.  \\item To have any hope of detecting birthing quasars, a survey has to reach a depth of at least a few $\\times 10^{-13}\\erg\\,\\s^{-1}\\cm^{-2}$.  However, the slopes of our calculated $\\log N - \\log S$ distributions at fluxes just below this threshold are relatively shallow (Figures \\ref{fig:fluxX}, \\ref{fig:fluxu} and \\ref{fig:fluxrep}), implying that surveys should favor large angular sky coverage over depth, once they reach this flux threshold.  \\item If identified, candidate sources can be followed up by pointed observations at higher frequencies, where they are expected to continue both their monotonic brightening and their spectral hardening.  \\item Most birthing quasars that are identifiable have, coincidentally, SMBH masses lying in the middle of {\\it LISA}'s sensitivity window ($M\\sim 10^{6}\\Msol$), and are thus members of the same population that would be probed with GW detections. However, a minority ($\\gtrsim {\\rm few}~\\%$ for $S\\gtrsim 3$) of the detectable X-ray variables have masses of $\\gsim 10^7~{\\rm M_\\odot}$, probing a population above {\\it LISA}'s range.  \\item These sources may be identifiable by {\\it LSST} if a fraction as low as $\\sim1\\%$ of the X-ray flux is promptly reprocessed into the optical frequencies.  \\end{itemize}  Our calculations are contingent on theoretical caveats of the afterglow scenario we have considered.  The two primary uncertainties regarding the post-merger evolution of the circumbinary cavity are related to the viscous and advective properties of the disk.  As stated in \\S~\\ref{subsec:curve}, the viscosity of accretion flows, including the possibility of viscous instability, are not well understood when radiation pressure dominates gas pressure, which is the relevant regime for the gas refilling the circumbinary cavity. Additionally, the disk may be geometrically thick \\citepalias{MP05, TM10}, either right at decoupling or later during the afterglow phase, suggesting that horizontal advection may play a significant role in determining the surface density evolution and the disk net emission properties.  The importance of viscous instabilities in radiation-dominated accretion flows remains a general open question, and the role of advection in a viscously spreading accretion flow remains a largely unexplored regime.  More detailed studies of the circumbinary cavity will be needed to address how these effects may affect the emission predicted by simple analyses based on a thin disk formalism such as ours.  Another major uncertainty is the validity of our assumption that quasar activity can be associated with SMBH coalescence.  In reality, there may not be a one-to-one relation: it is possible that for at least some AGN, gas accretion or changes in radiative efficiency are triggered by mechanisms other than SMBH mergers; conversely, some SMBH mergers may not trigger prolonged quasar activity.  If the former is true, our analysis overestimates the number of identifiable afterglow sources; if the latter is true, then our results could in principle be an underestimate.  For completeness, we note that while we focused here on the viscous afterglows, other SMBHB merger--related signatures could also be looked for in EM surveys.  For example, the GW--emission--induced mass--loss and recoil can cause strong disturbances in the circumbinary disk, which can produce a detectable afterglow \\citep{Lippai+08, SK08, SB08, O'Neill+09, Megevand+09, Corrales+10, Rossi+10}.  For the low SMBH masses of $\\sim10^6\\Msol$ relevant for {\\it LISA}, these signatures are expected to have a short duration $\\sim$ few years (e.g., \\citealt{Corrales+10}) and would be too rare to be found serendipitously, without a trigger from {\\it LISA}. However \\citet{SK08} and \\citet{SB08} focused on these signatures in disks around more massive SMBHs, which occur on longer ($\\sim 10^4$yr) time--scales, and proposed detecting a flare by monitoring a population of AGN in the infrared or X-rays bands.  Another possibility is that the binary is activated, and produces periodic emission, tracking the orbital frequency, prior to the merger. \\citet{HKM09} argued that as long as this emission is at a few percent of the Eddington luminosity, a population of these variable sources, with periods of tens of weeks, may be identifiable in optical or X-ray surveys.  To conclude, the concomitant observation of a SMBHB merger based on GW and EM signals remains by far the most promising scenario for the unambiguous detection of such systems. The precision with which {\\it LISA} would determine the masses, spins, and luminosity distances of coalescing binaries can not be replicated by current or planned EM telescopes.  However, detections based on EM signatures alone could still help identify SMBHB mergers before {\\it LISA} is launched, and perhaps more importantly, possibly outside {\\it LISA}'s mass sensitivity window.  Detecting the EM signatures from the mergers of the most massive SMBHs would complement the synergistic EM-plus-GW observations of lower-mass systems, and help provide a more complete picture of the accretion physics and cosmological evolution history of SMBHBs.  $~$  TT thanks Joshua Peek and Jennifer Sokoloski, and ZH thanks Stefanie Komossa, Jules Halpern, and Richard Mushotzky, for useful conversations on AGN surveys.  This work was supported by the Pol\\'anyi Program of the Hungarian National Office for Research and Technology (NKTH) and by NASA ATFP grant NNXO8AH35G."
    },
    "1004/1004.0773.txt": {
        "abstract": "We consider spherically symmetric inhomogeneous dust models with a positive cosmological constant, $\\Lambda$, given by the Lema\\^\\i tre--Tolman--Bondi metric. These configurations provide a simple but useful generalization of the $\\Lambda$--CDM model describing cold dark matter (CDM) and a $\\Lambda$ term, which seems to fit current cosmological observations. The dynamics of these models can be fully described by scalar evolution equations that can be given in the form of a proper dynamical system associated with a 4--dimensional phase space whose critical points and invariant subspaces are examined and classified. The phase space evolution of various configurations is studied in detail by means of two 2--dimensional subspaces: a projection into the invariant homogeneous subspace associated with $\\Lambda$--CDM solutions with FLRW metric, and a projection into a subspace generated by suitably defined fluctuations that convey the effects of inhomogeneity. We look at cases with perpetual expansion, bouncing and loitering behavior, as well as configurations with ``mixed'' kinematic patters, such as a collapsing region in an expanding background. In all cases, phase space trajectories emerge from and converge to stable past and future attractors  in a qualitatively analogous way as in the case of the FLRW limit. However, we can identify in both projections of the phase space various qualitative features absent in the FLRW limit that can be useful in the construction of toy models of astrophysical and cosmological inhomogeneities. ",
        "introduction": "From a phenomenological and empiric point of view, cosmological observations are usually tested as a first attempt by fitting them to the so--called $\\Lambda$--CDM model, which is a FLRW spacetime with flat spacelike sections, whose source is dust (cold dark matter CDM) and a $\\Lambda$ field (dark energy) (see \\cite{reviewDE1,reviewDE2} for comprehensive reviews). Although the universe may be nearly homogeneous at scales larger than the so--called homogeneity scale (over 150--300 Mpc), thus justifying the use of linear perturbations in its dynamical study, it is clearly inhomogeneous at smaller scales in which structure formation mostly involving CDM has taken  place. Hence, the study of inhomogeneous sources made of dust and a nonzero ``$\\Lambda$ field'' is a very relevant topic. In particular, spacetimes of this type of source with spherical symmetry provide simple but non--trivial inhomogeneous generalizations of the $\\Lambda$--CDM model.  Inhomogeneous spherically symmetric dust solutions of Einstein's equations with $\\Lambda=0$,  described by the Lema\\^\\i tre--Tolman--Bondi (LTB) metric have been widely used to construct simple models of cosmological inhomogeneities (see \\cite{kras1,kras2} for a comprehensive discussion and review). However, the literature on LTB models with $\\Lambda> 0$ (which we will denote by $\\Lambda$--LTB models) is much less abundant \\cite{LTBL1,LTBL2}, dealing mostly on singularities and censorship \\cite{LTBLS}. See \\cite{WA} for a recent study of the dynamics of LTB models (considering also $\\Lambda\\ne 0$), while Szekeres models with $\\Lambda\\ne 0$ have been examined in \\cite{barrow}.  In the present paper we attempt to study the $\\Lambda$--LTB models under the qualitative and numerical phase space techniques known as ``dynamical systems''. These techniques have been applied successfully to a wide variety of spacetimes (FLRW, Kantowski--Sachs, Bianchi models, self--similar and static spacetimes, etc), for which evolution equations for the ``expansion normalized'' phase variables can be reduced to a system of autonomous ODE's that define a self--consistent phase space (see \\cite{DS1,DS2,DS3,DS4} for a comprehensive elaboration and review). A dynamical system analysis has been been successfully applied to LTB dust solutions with $\\Lambda=0$ in reference \\cite{suss08}. We generalize here this study to the case $\\Lambda> 0$. The plan of this article is summarized below.  The class of $\\Lambda$--LTB spacetimes and their corresponding covariant objects are introduced in section 2, while in section 3 we show that their dynamics can be completely determined by ``quasi--local'' (QL) variables \\cite{sussPRD,sussLTB1,sussLTB2} that lead to a set of ``fluid flow'' evolution equations \\cite{1plus3} and an initial value parametrization that is useful for carrying on numeric work on the models. These evolution equations  are transformed in section 4 into a 4--dimensional dynamical system constructed by ``expansion normalized'' variables  \\cite{DS1,DS2,DS3,DS4}. In section 5 we list the critical points (past and future attractors and saddle points) of this dynamical system and classify all the  invariant subspaces of the phase space: the homogeneous FLRW subset, the spatially flat subspace, as well as the Kottler vacuum solution \\cite{kottler} (Schwarzschild--de Sitter) and ``pure'' dust with $\\Lambda=0$ subspaces. We examine in section 6 the different patterns of kinematic evolution of $\\Lambda$--LTB models. The admissible topological classes of the space slices orthogonal to the 4--velocity are summarized in section 7. Since the phase space is 4--dimensional, we decompose it in section 8 in two 2--dimensional projections: the ``homogeneous'' and ``inhomogeneous'' subsystems. These projections convey the full dynamical information, and thus are used in section 9 to study in detail the phase space trajectories of representative configurations associated with each one of the kinematic patterns and topologies of the space slices listed in sections 6 and 7. For this purpose, we solve numerically the dynamical system obtained in section 4 for each configuration, together with the QL evolution equations derived in section 3 (which is necessary when the phase space variables diverge in collapsing or bouncing configurations). Section 10 summarizes the results and provides a final discussion. ",
        "conclusions": "We have conducted a comprehensive and detailed dynamical systems study of $\\Lambda$--LTB spacetimes by means of covariant quasi--local (QL) variables, generalizing a previous dynamical systems study \\cite{suss08} on LTB dust models with $\\Lambda=0$. Since the phase space for the dynamical system (\\ref{DSa})--(\\ref{DSd}) is 4--dimensional, we have examined and plotted phase space trajectories in two 2--dimensional subsets: the ``homogeneous'' and ``inhomogeneous'' subspaces, $\\textrm{H}_0$ and $\\textrm{I}$, defined by the projections (\\ref{Hproj}) and (\\ref{Iproj}), and respectively displayed in panels (a) and (b) of figures \\ref{fig4}--\\ref{fig11}. The decomposition of the phase space into $\\textrm{H}_0$ and $\\textrm{I}$ provides an important theoretical connection between the dynamical system study of $\\Lambda$--LTB spacetimes and the perturbation formalism presented in \\cite{sussPRD,sussLTB1}. In this formalism, the fluctuations $\\{\\Dm,\\,\\Dz\\}$ (which parametrize $\\textrm{I}$) are rigorously characterized as covariant and gauge invariant non--linear perturbations on a FLRW background defined by quasi--local scalars, which satisfy FLRW dynamics, and hence corresponding to the projected homogeneous subspace $\\textrm{H}_0$ parametrized by $\\{\\Omm,\\,\\Oml\\}$.  In LTB dust models with $\\Lambda=0$ all configurations in which the space slices $\\T[t]$ have closed ($\\mathbb{S}^3$) or wormhole ($\\mathbb{S}^3\\times \\mathbb{R}$ or $\\mathbb{S}^3\\times \\mathbb{S}^1$) topologies, the regularity conditions (\\ref{grr1})--(\\ref{grr2}) imply that spatial curvature must be positive and thus, perpetual expansion is not possible \\cite{sussLTB1,ltbstuff,suss02}. However, this restriction is no longer true when $\\Lambda>0$. While (\\ref{grr1})--(\\ref{grr2}) still requires positive spatial curvature for configurations with these topologies, positive spatial curvature no longer forbids perpetual expansion (see section 7). Hence, perpetually expanding and fully regular $\\Lambda$--LTB models with closed or wormhole topologies are possible.  The homogeneous projections in panels (a) of figures \\ref{fig4}--\\ref{fig11} show that phase space trajectories are qualitative analogous to those of dust--$\\Lambda$ FLRW models. However, each FLRW spacetime configuration would correspond to one (and only one) phase space trajectory in a $\\{\\Omm,\\,\\Oml\\}$ diagram (because each spacetime configuration is uniquely determined by initial conditions given by constant $\\Ommi,\\Omli$). As a contrast, since initial conditions for any single $\\Lambda$--LTB configuration depend on $r$, all displayed curves in the diagrams in the panels (a) of each one of these figures  corresponds to a single spacetime configuration. In other words, in terms of the projection (\\ref{Hproj}) each $\\Lambda$--LTB configuration is formally equivalent to a superposed one--parameter family of dust--$\\Lambda$ FLRW models.  In all $\\Lambda$--LTB configurations (except those like the one in figure \\ref{fig8}) the phase space trajectories initiate in the same source or past attractor {\\bf{PA}}, defined in (\\ref{PA}) and associated with an initial singularity. This critical point can be projected into the past attractor of the homogeneous subspace, and can be identified with a spatially flat FLRW dust cosmology ($\\Omm=1,\\,\\Oml=0$). Since $\\Dm=\\Dz=-1$, this critical point exactly coincides with the past attractors of LTB dust solutions with $\\Lambda=0$ and spatially flat models with $\\Lambda>0$, thus indicating that the effects of $\\Lambda$ and of spatial curvature are negligible in the early stages of the evolution near the big bang. The past attractor {\\bf{PA}} can also be identified with self--similar conditions \\cite{suss08}, which is consistent with the fact that both spatially flat FLRW and LTB dust spacetimes with $\\Lambda=0$ are compatible with a homothetic Killing vector. The past attractor (\\ref{PA}) becomes the future attractor or sink for phase space trajectories of layers that collapse to a curvature singularity (big crunch). This is so, either for a collapse that is a trivial time reversal of a perpetual expansion (figures \\ref{fig4}--\\ref{fig6}), or for configurations where all layers expand and collapse (figure \\ref{fig7}), or with mixed dynamics in which some layers expand perpetually and some collapse (figure \\ref{fig10}).  Phase space trajectories of perpetually expanding layers ($z_q>0$) terminate in the same future attractor (sink) {\\bf{FA}}, defined in (\\ref{FA}): figures \\ref{fig4}--\\ref{fig6} and in the perpetually expanding trajectories of figures \\ref{fig8}, \\ref{fig10} and \\ref{fig11}. This critical point can be projected into the future attractor of the homogeneous subspace ($\\Omm=0,\\,\\Oml=1$), and is also the future attractor for spatially flat configurations. Hence, it can be identified with de Sitter spacetime and spatially flat conditions, indicating that the effects of $\\Lambda$ are dominant in the future late stages of the evolution of $\\Lambda$--LTB configurations in which the worldlines of expanding dust layers are infinitely inextensible. As a consequence, with the exception of the unstable loitering models, all $\\Lambda$--LTB models endowed with a late time evolution are compatible with the ``{\\it cosmic no hair}'' conjecture~\\cite{barrow,cnh}, with the asymptotic de Sitter state identified with the critical point {\\bf{FA}}. For dust layers collapsing from infinity ($L\\to\\infty$), the phase space trajectories emerge from a past attractor which exactly coincides with the point {\\bf{FA}}.  Seen from the inhomogeneous projection (\\ref{Iproj}) (panels (b)), the evolution towards {\\bf{FA}} of phase space trajectories of perpetually expanding layers contains extra information not available in the projection (\\ref{Hproj}). Trajectories of layers with negative spatial curvature (figures \\ref{fig4}--\\ref{fig5}) approach the Minkowskian saddle {\\bf{SP2}}, whereas layers with positive curvature approach the saddle {\\bf{SP4}} (figures \\ref{fig6}, \\ref{fig7}, \\ref{fig10} and \\ref{fig11}) associated with spatially flat conditions. The approach to {\\bf{SP2}} clearly indicates an intermediate low density state, and also occurs in trajectories of the FLRW phase space (notice that {\\bf{SP2}} is projected by (\\ref{Hproj}) into the single saddle of this phase space). However, the approach to {\\bf{SP4}} has no equivalent in trajectories of the FLRW phase space of layers evolving towards the future attractor.  In configurations containing dust layers that bounce and collapse ($z_q$ changes sign: figures \\ref{fig7}--\\ref{fig11}) we used the system (\\ref{sys1a})--(\\ref{sys1d}) to obtain the full phase space trajectories. In all cases the phase space trajectories in the homogeneous projection  start in the past attractor {\\bf{FA}}, reach infinite values $\\Omm,\\Oml\\to\\infty$ as $z_q=0$ and return along the same trajectories to the same attractor (which is now a future attractor). However, in the inhomogeneous projection (panels (b) of figures \\ref{fig7} and \\ref{fig8}) the curves do not return to the attractor along the same trajectories: if the pattern is expansion/collapse (figure \\ref{fig7}) they reach $\\Dz\\to\\infty$ as $z_q=0$ and return to the attractor from $\\Dz\\to-\\infty$. If the pattern is collapse/bounce (figure \\ref{fig8}), the trajectories reach $\\Dz\\to-\\infty$ and return to the attractor from $\\Dz\\to\\infty$. The same effect would happen in the mixed pattern configurations displayed in figures \\ref{fig10} and \\ref{fig11}, but was not included in the plots as the full curves would make a very messy pattern. This difference in the behavior of curves emerging/returning from/to an attractor is an effect inherent in the inhomogeneity of LTB solutions.  The bouncing and loitering models deserve a separate mention, as they have no equivalence in the case $\\Lambda=0$. In the bouncing models (figure \\ref{fig8}), layers initially collapse from infinity, hence their past attractor coincides with the future attractor of perpetually expanding configurations (figures  \\ref{fig4}--\\ref{fig6}).  In the loitering configuration that we examined (figure \\ref{fig9}), the phase space trajectories emerge from the past attractor {\\bf{PA}}, but as the layers become asymptotically static $z_q\\to 0$ the trajectories evolve towards $\\Omm,\\Oml,\\Dz\\to\\infty$. In the homogeneous projection (panel (a) of figure \\ref{fig9}) the trajectories are indistinguishable from those of figure \\ref{fig7}, but in the inhomogeneous projection (panel (b) of figure \\ref{fig9}) the curves reach $\\Dz\\to\\infty$ along different paths and converge towards $\\Dm\\to 0$. This effect is an inherent feature of inhomogeneity and does not occur in FLRW phase space diagrams for loitering models. It is worthwhile remarking that loitering models do not evolve towards an asymptotic de Sitter state associated with the future attractor {\\bf{FA}}.  The dynamics of $\\Lambda$--LTB models is not much different (qualitatively speaking) from that of the homogeneous FLRW models that they generalize. Essentially, the past attractor is a state close to spatially flat FLRW and self--similarity, while the future attractor is a de Sitter state dominated by $\\Lambda$ (in agreement with the ``cosmic no hair'' conjecture). The two most significant effects that arise from the inhomogeneity of $\\Lambda$--LTB models are: % \\begin{enumerate} % \\item  the possibility of accommodating wholly different kinematic patterns in the same spacetime configuration (see figures \\ref{fig10}--\\ref{fig11}). However, even when the kinematic patterns of all layers are qualitatively similar, each unique $\\Lambda$--LTB spacetime can be understood as an inhomogeneous model made up of some sort of superposed FLRW models (as the phase space trajectory of each comoving worldline can be projected by (\\ref{Hproj}) into the phase space trajectory of a unique FLRW model). % \\item the asymmetry in the evolutions from a past attractor and a bouncing point and from that point to the same attractor, which becomes a future attractor (see panels (b) of figures \\ref{fig7} and \\ref{fig8}) % \\end{enumerate} % However, in spite of the simplification involved in assuming spherical symmetry and geodesic motion, the present dynamical system study clearly shows that even this idealized level of inhomogeneity does provide extra degrees of freedom that can be be very handy for constructing useful toy models of astrophysical and cosmological inhomogeneities."
    },
    "1004/1004.4284_arXiv.txt": {
        "abstract": "We describe the results of a search for the remnants of the Sun's birth cluster among stars in the Hipparcos Catalogue. This search is based on the predicted phase space distribution of the Sun's siblings from simple simulations of the orbits of the cluster stars in a smooth Galactic potential. For stars within $100$ pc the simulations show that it is interesting to examine those that have small space motions relative to the Sun. From amongst the candidate siblings thus selected there are six stars with ages consistent with that of the Sun. Considering their radial velocities and abundances only one potential candidate, HIP 21158, remains but essentially the result of the search is negative. This is consistent with predictions by \\cite{SPZ2009} on the number of siblings near the Sun. We discuss the steps that should be taken in anticipation of the data from the Gaia mission in order to conduct fruitful searches for the Sun's siblings in the future. ",
        "introduction": "The Sun's life history has long been a subject of interest not just in astrophysics but also in fields such as solar system studies, the history of the earth's climate, and understanding the causes of mass extinctions. The possible birth environment of the Sun was discussed extensively by \\cite{Adams2010} who shows how inferences about this environment can be made by considering its impact on the formation and morphology of our planetary system, the removal of the solar nebula, and the presence of short-lived radioactive nuclei in meteorites. The subsequent life and times of the Sun as it travels through our Galaxy have attracted attention in the context of trying to understand climate change and mass extinctions as the consequences of astronomical impacts. The evidence for and against this idea was reviewed by \\cite{BailerJones2009}, who points out problems in the methodology of the various studies into climate change or mass extinctions and also the uncertainties in the details of the Sun's path through our Galaxy even over the past 545 Myr.  As discussed by \\cite{SPZ2009} the Sun is likely to have been born in a bound open cluster consisting of a few thousand stars. This cluster probably had a radius of a few pc and as pointed out by \\cite{Adams2010} the Sun was located not too far from the cluster centre ($\\sim0.2$ pc) as inferred from the necessity of a nearby supernova explosion. The fact that the Sun thus has a large `family' prompted \\cite{SPZ2009} to ask the question: can we find the Sun's siblings? The answer to this question is important as the inferences about the Sun's birth environment all come from considering the Sun and its planets, there is as yet no direct observational constraint on the birth cluster itself. Identifying even a small number of the Sun's siblings would put constraints on the number of stars in the cluster, by extrapolation for a plausible IMF, and possibly even on the IMF itself if siblings were found over a range of stellar masses. Reconstructing the orbits of the siblings in the Galaxy would lead to a more accurate determination of the Sun's birth location as well as the subsequent path to its present day position. This information could be used, for example, to investigate whether the Sun's relatively high metallicity \\citep[cf.][]{Adams2010} can be explained by its birth at a different radius in the Galaxy. In addition we would obtain a determination of the Sun's motion through the Galaxy independent from the geological record, which was listed by \\cite{BailerJones2009} as an important goal for the study of the history of the earth's climate and mass extinctions. \\cite{SPZ2009} proceeded by considering the constraints on the Sun's birth cluster and performing simple simulations of the evolution of a cluster of stars initially confined to a $1$ pc virial radius and orbiting our Galaxy along the presumable path the Sun followed in the past. Depending on how quickly the cluster became unbound \\cite{SPZ2009} concluded that $\\sim10$--$40$\\% of the Sun's siblings should still be located with $1$ kpc of the present day location of the Sun.  Thus we can expect to find about $\\sim 100$--$1000$ of the Sun's siblings with $1$ kpc from the present day position of the Sun. This will make a search for the siblings extremely challenging as they will have to be weeded out from among the $\\sim 10^8$ stars within $1$ kpc. Nevertheless we set out in this paper to make a first attempt at identifying candidate siblings of the Sun by searching in the Hipparcos Catalogue and adding complementary data from the Geneva-Copenhagen survey of the Solar neighbourhood \\citep{GCS2009}. Our motivation is to carry out a first exploration of kinematic searches for the Sun's siblings. We describe our search methodology in section \\ref{sec:method} and present our results in section \\ref{sec:results}. We discuss the results in section \\ref{sec:discussion} and outline the steps needed to carry out a thorough future search for the Sun's siblings in section \\ref{sec:future}. ",
        "conclusions": "\\label{sec:future}  Motivated by the desire to find the remnants of the Sun's birth cluster we have conducted a preliminary search in the Hipparcos Catalogue for stars that could have been born in the same cluster. This search was based on the predicted phase space distribution of the Sun's siblings from simple simulations of the orbits of the cluster stars in a smooth Galactic potential. For nearby stars the simulations show that it is interesting to examine those that have small space motions relative to the Sun. From amongst the candidate siblings thus selected there are six stars with ages consistent with that of the Sun. Of these six candidate siblings 5 can be excluded on the basis of their radial velocity or metallicity, leaving only one plausible candidate sibling, HIP 21158.  However, the latter still has a radial velocity somewhat higher than predicted from our simulations. This means we have not found a single convincing solar sibling within $100$ pc from the Sun which is consistent with the predictions by \\cite{SPZ2009} and the fact that only a small fraction of the stars near the Sun was examined.  Now, even if a stronger case could have been made for the candidate siblings in table \\ref{tab:siblings} based on their age and value of \\feh, this would not have proven that these stars are truly siblings of the Sun. We discuss below what steps need to be taken in future searches for the Sun's siblings.  The process of cluster disruption in the Galactic potential was simulated in a simplified manner in both this work and in \\cite{SPZ2009}. A better understanding of the expected distribution of the siblings in phase space is essential for an efficient search for them in future large surveys. Hence it is important to do simulations of cluster disruption that are as realistic as possible. Effects to be included are the self-gravity of the cluster stars, non-axisymmetric structures in the Galactic potential, such as the bar and spiral arms, and the collisions of the cluster with molecular clouds in the Galaxy. The resulting phase space distribution is expected to be less orderly than depicted in figures \\ref{fig:xyprojections} and \\ref{fig:pmparallax} but to what extent is unknown at the moment. In these simulations it will be important to ensure that the resolution is comparable to realistic cluster and molecular cloud mass scales. To properly understand selection effects in surveys it is also necessary to include a realistic initial mass function and stellar evolution in the cluster simulations.  The observational challenge is equally daunting. On the one hand a large scale survey of phase space is needed, covering a large volume of the Galactic disk. Only the Gaia mission \\citep{Gaia2008} will provide this data at the precision needed to probe for siblings far away from the Sun. The above simulations will have to be exploited to develop efficient search methods that can weed out the candidate siblings from among the billion stars in the Gaia catalogue. However a search in phase space only is not sufficient as the stars from different clusters on similar orbits as the Sun's birth cluster could be confused with the genuine siblings. In addition it is known that clustering of stars in phase space can also be caused by dynamical effects \\citep[see for example,][]{Antoja2009}. The phase space search will have to be complemented by a very detailed astrophysical characterization of candidate siblings of the Sun. The age and overall abundance of the stars will not be enough in this respect as many clusters with abundances similar to the Sun's birth cluster may have formed around the same time. In addition the errors on individual stellar ages ($\\sim0.5$--$1$ Gyr in table \\ref{tab:siblings}) are likely to remain larger than any plausible age spread within the birth cluster or the lifetime of the molecular cloud from which the cluster formed. However, the Sun's siblings are expected to have the same detailed chemical composition as the Sun and true siblings can thus be identified through the analysis of high resolution spectra. The latter will have to be collected in a dedicated follow-up programme.  This `chemical tagging' of stars has been proposed as a powerful method for associating them with their formation sites \\citep{Freeman2002}. So far it has been demonstrated for the Hyades, the HR1614 moving group, and Cr 261 that these groups of stars indeed have unique chemical signatures and promising elements have been identified that can be used to chemically identify groups of stars \\citep{DeSilva2007}. However, a number of important studies regarding this technique remain to be done: \\begin{itemize} \\item No attempt has been made so far to {\\em identify} (new) moving groups or clusters on the basis of abundance patterns, so the feasibility of this important aspect of the chemical tagging method remains to be demonstrated. \\item It has not yet been definitively established to what accuracy the abundance patterns of stars have to be measured in order to identify them with their birth sites. Although \\cite{DeSilva2007} conclude that $\\sim0.05$ dex accuracy on individual abundance measurements may be enough to do so, it is not clear what accuracy is needed to distinguish formation sites at the same Galactic radius. \\item If higher accuracy is needed differential abundance analyses offer the possibility of reaching $\\sim0.01$--$0.02$ dex accuracies as demonstrated by \\cite{Melendez2009} and \\cite{Ramirez2009} for solar analogs. Can these accuracies also be reached over wider ranges in the effective temperatures of stars? As these two papers suggest, at this level of accuracy the abundance patterns in stars may be affected by the presence or absence of a planetary system. This would then have to be accounted for in the search for the Sun's siblings. \\end{itemize}  With the Gaia survey starting in a few years from now, the questions above will be actively pursued in order to ensure that precision Galactic archaeology can be done by combining the accurate distances and kinematics from Gaia with accurate abundances for large samples of stars throughout the Galaxy. The results will open up the exciting prospect of further unravelling the birth environment and life and times of the solar system through the identification of the Sun's lost siblings."
    },
    "1004/1004.1142_arXiv.txt": {
        "abstract": "The \\emph{SAGE-Spec} \\emph{Spitzer} Legacy program is a spectroscopic follow-up to the \\emph{SAGE-LMC} photometric survey of the Large Magellanic Cloud carried out with the \\emph{Spitzer Space Telescope}. We present an overview of SAGE-Spec and some of its first results. The \\emph{SAGE-Spec} program aims to study the life cycle of gas and dust in the Large Magellanic Cloud, and to provide information essential to the classification of the point sources observed in the earlier \\emph{SAGE-LMC} photometric survey. We acquired 224.6 hours of observations using the \\emph{InfraRed Spectrograph} and the SED mode of the \\emph{Multiband Imaging Photometer for Spitzer}. The \\emph{SAGE-Spec} data, along with archival \\emph{Spitzer} spectroscopy of objects in the Large Magellanic Cloud, are reduced and delivered to the community. We discuss the observing strategy, the specific data reduction pipelines applied and the dissemination of data products to the scientific community. Initial science results include the first detection of an extragalactic ``21 $\\mu$m'' feature towards an evolved star and elucidation of the nature of disks around RV\\,Tauri stars in the Large Magellanic Cloud.  Towards some young stars, ice features are observed in absorption.  We also serendipitously observed a background quasar, at a redshift of $z\\approx 0.14$, which appears to be host-less. ",
        "introduction": "A photometric survey in the infrared of the Large Magellanic Cloud (LMC) was performed by the \\emph{Spitzer} Legacy Program \\emph{Surveying the Agents of Galaxy Evolution} \\citep[SAGE-LMC;][]{2006AJ....132.2268M}, which charts the budget of gas and dust contributing to the cycle of star formation and stellar death in the Magellanic Clouds. Here we discuss \\emph{SAGE-Spec}, a spectroscopic follow-up to \\emph{SAGE-LMC}, which is also a \\emph{Spitzer} Legacy Program. For \\emph{SAGE-Spec} we observed a variety of circumstellar and interstellar environments with the \\emph{Infrared Spectrograph} \\citep[IRS;][]{2004ApJS..154...18H} aboard \\emph{Spitzer} \\citep{2004ApJS..154....1W}, as well as the SED mode available on the \\emph{Multiband Imaging Photometer for Spitzer} \\citep[MIPS;][]{2004ApJS..154...25R}.  The \\emph{SAGE-Spec} dataset is exceptionally suited to address the following issues: First, it allows us to trace the lifecycle of dust and molecular gas on its journey through the galaxy, from dust production sites (AGB stars, red supergiants, post-AGB-objects, planetary nebulae), to the ISM (atomic and molecular clouds) to star forming regions (\\ion{H}{2} regions, young stellar objects); and, second, it allows us to develop a photometric color-color and color-magnitude classification scheme to increase the legacy of the larger \\emph{SAGE-LMC} database.  In addition, a large number of smaller astrophysical questions can be addressed using the data set provided here, and a rich harvest in scientific results is expected.  This paper gives an overview of the \\emph{SAGE-Spec} Legacy program. We outline the observing strategy, describe the data reduction process and discuss the data products that are currently publicly available to the astronomical community \\citep{SAGESpec-delivery2}, or will become publicly available in the near future.  Existing surveys targeting gas and dust in the LMC are discussed in Sect.~\\ref{sec:surveys}. This section also includes a description of existing publications of infrared spectroscopy on LMC targets.  The observing strategy of the \\emph{SAGE-Spec} project is discussed in Sect.~\\ref{sec:observations}, along with a description of the data reduction.  This paper finishes with some first scientific results of the \\emph{SAGE-Spec} project (Sect.~\\ref{sec:results}), and conclusions and an outlook to the future (Sect.~\\ref{sec:conclusions}). ",
        "conclusions": "\\label{sec:conclusions}  The \\emph{SAGE-Spec} program provides useful data for understanding the life cycle of gas and dust in galaxies. The extensive dataset of \\emph{IRS} and \\emph{MIPS SED} spectroscopy, obtained within the \\emph{SAGE-Spec} program, complemented with archival data, sample the relevant environments and ultimately provides insights on dust mineralogy and gas properties in these environments. Feeding the results back to the original \\emph{SAGE-LMC} data leads to conclusions on stellar populations, and allow us to study the mineralogical dust cycle in the Large Magellanic Cloud, in combination with the global star formation rate \\citep[][]{2008AJ....136...18W,2009ApJS..184..172G} and injection rate of stellar mass loss into the ISM \\citep[e.g.~][]{2009MNRAS.396..918M,2009AJ....137.4810S}.  An important outcome of the \\emph{SAGE-Spec} program is contributing distinguishing diagnostics to classify sources in the \\emph{SAGE-LMC} point source catalog.  The initial results discussed in the paper include the first extragalactic detection of the 21 $\\mu$m feature; the study of crystalline silicates in the disks around RV Tauri stars; the possible detection of a host-less quasar; the analysis of ices towards massive YSOs; and investigations into feature and line ratios in atomic and \\ion{H}{2} regions to probe physical conditions, such as radiation field and ionization fraction.  True to its legacy status, the \\emph{SAGE-Spec} team has delivered a significant fraction of its reduced data to the scientific community already, with further data deliveries planned in the near future. The unprocessed data have been available to the community in the \\emph{Spitzer} archive from the date of observing. We will also deliver enhanced data products, particularly spectral feature maps and source and spectral classifications in those future deliveries."
    },
    "1004/1004.1973_arXiv.txt": {
        "abstract": "A general analytic procedure is developed for the post-Newtonian limit of $f(R)$-gravity with metric approach in the Jordan frame by using the harmonic gauge condition. In a pure perturbative framework and by using the Green function method a general scheme of solutions up to $(v/c)^4$ order is shown. Considering the Taylor expansion of a generic function $f$ it is possible to parameterize the solutions by derivatives of $f$. At Newtonian order, $(v/c)^2$, all more important topics about the Gauss and Birkhoff theorem are discussed. The corrections to \"standard\" gravitational potential ($tt$-component of metric tensor) generated by an extended uniform mass ball-like source are calculated up to $(v/c)^4$ order. The corrections, Yukawa and oscillating-like, are found inside and outside the mass distribution. At last when the limit $f\\rightarrow R$ is considered the $f(R)$-gravity converges in General Relativity at level of Lagrangian, field equations and their solutions. ",
        "introduction": "The study of possible modifications of Einstein's theory of gravity has a long history which reaches back to the early 1920s \\cite{Weyl:1918,Pauli:1919,Bach:1921}. Corrections to the gravitational Lagrangian, leading to higher-order field equations, were already studied by several authors \\cite{Weyl:1921,Eddington:1924,Lanczos:1931} shortly after General Relativity (GR) was proposed. Developments in the 1960s and 1970s \\cite{Buchdahl:1962,DeWitt:1965,Bicknell:1974,Havas:1977,Stelle:1978}, partly motivated by the quantization schemes proposed at that time, made clear that theories containing {\\it only} a $R^2$ term in the Lagrangian were not viable with respect to their weak field behavior. Buchdahl, in 1962 \\cite{Buchdahl:1962}, rejected pure $R^2$ theories because of the non-existence of asymptotically flat solutions.  In  recent years, the effort to give a  physical explanation to the today observed cosmic acceleration \\cite{sneIa,lss,cmbr} has attracted a good amount of interest in $f(R)$-gravity, considered as a viable mechanism to explain the cosmic acceleration by extending the geometric sector of field equations without the introduction of dark matter and dark energy. There are several physical and mathematical motivations to enlarge GR by these theories. For comprehensive review, see \\cite{GRGreview,OdintsovLadek,farhoudi}.  Other issues as, for example, the observed Pioneer anomaly problem \\cite{anderson} can be framed into the same approach \\cite{bertolami} and then, apart the cosmological dynamics, a systematic analysis of such theories urges at short scale and in the low energy limit.  While it is very natural to extend Einstein's gravity to theories with additional geometric degrees of freedom, recent attempts focused on the old idea of modifying the gravitational Lagrangian in a purely metric framework, leading to higher-order field equations. Due to the increased complexity of the field equations in this framework, the main body of works dealt with some formally equivalent theories, in which a reduction of the order of the field equations was achieved by considering the metric and the connection as independent objects \\cite{Francaviglia}.  In addition, many authors exploited the formal relationship to scalar-tensor theories to make some statements about the weak field regime \\cite{olmo}, which was already worked out for scalar-tensor theories \\cite{Damour:Esposito-Farese:1992}. Also a Post-Newtonian parameterization with metric approach in the Jordan Frame has been considered \\cite{clifton}.  In this paper, we study the Post Newtonian limit of $f(R)$ in the harmonic gauge. We are going to focus on the small velocity and weak field limit within the metric approach. In principle, any alternative or extended theory of gravity should allow to recover positive results of General Relativity. It will be very important to check at any level of our modified theory we can cover the outcomes of GR.  The plan of the paper is the following. In the Sec.\\ref{fieldequations}, we report the complete scheme of Newtonian and post-Newtonian limit of field equations for $f(R)$-gravity and their formal solutions in the harmonic gauge condition. General comments about the mathematical properties of equations, their relative solutions (Gauss and Birkhoff theorem) and Minkowskian behavior of metric tensor are reported. In the Sec. \\ref{ball_like} we show the complete solutions (Newtonian and post-Newtonian level) when an uniform mass ball-like source is considered. The point-like source limit of the newtonian solution is considered and the compatibility of $f(R)$-gravity with respect to GR is shown. Concluding remarks are drawn in Sec. \\ref{conclusions}. ",
        "conclusions": "In this paper, we have generally reformulated the Newtonian limit of $f(R)$-gravity by applying the Green function method. Moreover the post-Newtonian limit has been studied in the harmonic gauge condition and the relative spatial behaviors (Yukawa-like and oscillating) of gravitational potential in the matter and in the vacuum have been shown. The Taylor expansion of a generic $f$ has been considered obtaining general solutions in term of the derivatives up to third degree when an uniform mass ball-like source is considered. All metric potentials, however, depend strictly on the coupling parameters appearing indirectly in the Lagrangian of the theory (the derivatives of $f$ in $R\\,=\\,0$ are arbitrary constants).  A detailed discussion has been developed for systems presenting spherical symmetry. In this case, the role of corrections to the Newtonian potential is clearly evident. This means that one of the effects introducing a generic function of scalar curvature is to select a characteristic scale length which could have physical interests.  Furthermore, it has been shown that the Birkhoff theorem is not a general result for $f(R)$-gravity. This is a fundamental difference between GR and fourth order gravity. While in GR a spherically symmetric solution is, in any case, stationary and static, here time-dependent evolution can be achieved depending on the order of perturbations.  Hypothesizing a nonlinear Lagrangian we obtained a gravitational attraction stronger than in GR. The hypothesis of dark matter is needed in GR to have more gravitational attraction. Here without hypothesis of alternative matter and without modifying the gravitational constant we have qualitatively the same outcome. This occurrence could be particularly useful to solve the problem of missing matter in large astrophysical systems like galaxies and clusters of galaxies. In fact dark matter could be nothing else but the effects that GR, experimentally tested only up to Solar System scales, does not work at extragalactic scales and then it has to be corrected.  Then it is worth pointing out that $f(R)$-gravity seem good candidate to solve several shortcomings of Modern Astrophysics and Cosmology. Taking into account the results presented, it is clear that only GR presents directly the Newtonian potential in the weak field limit while corrections appear as soon as the theory is non-linear in the Ricci scalar.  In forthcoming researches, there is the intention to confront such solutions with experimental data in order to see if large self-gravitating systems could be modelled by them and if the experimental test of GR in the Solar System are compatible with $f(R)$-gravity."
    },
    "1004/1004.5141_arXiv.txt": {
        "abstract": "We combine new high sensitivity ultraviolet (UV) imaging from the Wide Field Camera 3 (WFC3) on the Hubble Space Telescope (HST) with existing deep HST/Advanced Camera for Surveys (ACS) optical images from the Great Observatories Origins Deep Survey (GOODS) program to identify UV-dropouts, which are Lyman break galaxy (LBG) candidates at \\uvdrops. These new HST/WFC3 observations were taken over 50~\\sqmin\\ in the GOODS-South field as a part of the Early Release Science program. The uniqueness of these new UV data is that they are observed in 3 UV/optical (WFC3 UVIS) channel filters (F225W, F275W and F336W), which allows us to identify three different sets of UV-dropout samples.  We apply Lyman break dropout selection criteria to identify F225W-, F275W- and F336W-dropouts, which are $z\\!\\simeq\\!1.7$, $2.1$ and $2.7$ LBG candidates, respectively.  We use multi-wavelength imaging combined with available spectroscopic and photometric redshifts to carefully access the validity of our UV-dropout candidates. Our results are as follows: (1) these WFC3 UVIS filters are very reliable in selecting LBGs with $z\\!\\simeq\\!2.0$, which helps to reduce the gap between the well studied $z\\!\\gtrsim\\!3$ and $z\\!\\sim\\!0$ regimes; (2) the combined number counts with average redshift $z\\!\\simeq\\!2.2$ agrees very well with the observed change in the surface densities as a function of redshift when compared with the higher redshift LBG samples; and (3) the best-fit Schechter function parameters from the rest-frame UV luminosity functions at three different redshifts fit very well with the evolutionary trend of the characteristic absolute magnitude, $M^*$, and the faint-end slope, $\\alpha$, as a function of redshift.  This is the first study to illustrate the usefulness of the WFC3 UVIS channel observations to select $z\\!\\lesssim\\!3$ LBGs.  The addition of the new WFC3 on the HST has made it possible to uniformly select LBGs from $z\\!\\simeq\\!1$ to $z\\!\\simeq\\!9$, and significantly enhance our understanding of these galaxies using HST sensitivity and resolution. ",
        "introduction": "The Lyman break `dropout' technique was first applied to select Lyman break galaxies (LBGs) at $z\\!\\simeq\\!3$ \\citep{guha90,stei96,stei99}, and since then it has been extensively used to select LBG candidates at $z\\!\\simeq\\!3\\!-\\!9$ \\citep[e.g.,][]{sawi06,bouw07,redd08,rafe09,oesc10,bunk10,yan10}. This dropout technique has generated large samples of star-bursting galaxy candidates at $z\\!\\simeq\\!3\\!-\\!9$, but there is only one major study \\citep{ly09} that investigates LBGs at \\uvdrops\\ based on dropout selection criteria.  The primary reason for this is that we need highly sensitive space-based cameras to observe the mid- to near-ultraviolet (UV) wavelengths required to select LBGs at \\uvdrops. The new Wide Field Camera 3 (WFC3) on the refurbished Hubble Space Telescope (HST) with its superior sensitivity --- compared to the Wide-Field Planetary Camera 2 (WFPC2) or the Galaxy Evolution Explorer (GALEX) --- and filters below the atmospheric cut-off wavelength (e.g., F225W and F275W), allows us to photometrically identify and study lower redshift (\\uvdrops) LBGs. The improved sensitivity/depth allows us to probe the lower luminosity systems at these redshifts.  There are two important reasons to understand these LBGs. First, to study the star formation properties of these LBGs, because they are at redshifts corresponding to the peak epoch of the global star formation rate \\citep[e.g.,][]{ly09,bouw10a,bouw10b,yan10}, and, secondly, they are likely lower redshift analogs of the high redshift LBGs --- because of the similar dropout selection at all redshifts --- whose understanding will help shed light on the process of reionization in the early Universe \\citep[e.g.,][]{labb10,star10,yan10}. The major advantage of identifying and studying various properties --- including star formation properties --- of lower redshift LBGs is that these LBGs can be investigated in rest-frame UV \\emph{as well as} rest-frame optical filters.  The high redshift LBGs have very little information on their rest-frame \\emph{optical} properties, so a detailed understanding of lower redshift LBGs is very important to get insight into the physical and morphological nature of high redshift LBGs.  The new UV observations of the WFC3 Science Oversight Committee (SOC) Early Release Science extragalactic program (PID: 11359, PI: O'Connell; hereafter ``ERS2''), covers approximately 50~\\sqmin\\ in the north-western part of the Great Observatories Origins Deep Survey \\citep[GOODS;][]{giav04a} South field. Here we use the high sensitivity of the new WFC3 UVIS channel data, along with existing deep optical data obtained with the Advanced Camera for Surveys (ACS) as part of the GOODS program, to search for LBG candidates at \\uvdrops. We use dropout color selection criteria based on color-color plots, obtained with the WFC3 UVIS and ACS filters to find three unique sets of UV dropouts --- F225W-dropouts, F275W-dropouts and F336W-dropouts --- which are LBG candidates at $z\\!\\simeq\\!1.7$, $2.1$ and $2.7$, respectively (as shown in \\figref{fig:lybreak}).  This paper is organized as follows: In \\secref{data} we summarize the WFC3 ERS2 observations, and in \\secref{sample} we discuss the selection, and in \\secref{reliable} the reliability of our color selected \\uvdrops\\ LBG sample.  In \\secref{results} we discuss the data analysis, which includes measuring their number counts and surface density (\\secref{ncounts}), and compare these with other surveys at higher redshifts, and estimate rest-frame UV luminosity functions (\\secref{lfs}) for these samples.  In \\secref{conclusion} we conclude with a summary of our results.  In the remaining sections of this paper we refer to the HST/WFC3 F225W, F275W, F336W, filters as \\wfcfuv, \\wfcnuv, \\wfcuv, and HST/ACS F435W, F606W, F775W, F850LP filters as \\acsb, \\acsv, \\acsi, \\acsz, respectively, for convenience.  We assume a \\emph{Wilkinson Microwave Anisotropy Probe} (WMAP) cosmology with $\\Omega_m$=0.274, $\\Omega_{\\Lambda}$=0.726 and \\Ho=70.5~km~s$^{-1}$~Mpc$^{-1}$, in accord with the 5 year WMAP estimates of \\citet{koma09}.  This corresponds to a look-back time of 10.37~Gyr at $z\\!\\simeq\\!2$.  Magnitudes are given in the AB$_{\\nu}$ system \\citep{oke83}. ",
        "conclusions": "\\label{results} \\subsection{Number Counts}\\label{ncounts}  The observed raw number counts of LBG candidates at \\uvdrops\\ at a rest-frame wavelength of 1700~\\AA\\ are shown in \\figref{fig:ncounts}. When we combine all three dropout samples, the average photometric redshift is $z\\!\\simeq\\!2.2$. For proper comparison, these number counts are \\emph{not} corrected for incompleteness or cosmic variance, and therefore, the counts start to drop at fainter magnitudes ($\\gtrsim26.0$~mag).  \\figref{fig:ncounts} (top panel) shows number counts (in number per arcmin$^2$ per 0.5 mag bin) of \\emph{all} dropouts (\\uvdrops) in our sample compared with other ground-based and space-based LBG surveys \\citep{stei99,noni09,ly09} at $z\\!\\simeq\\!2\\!-\\!3$. We have also plotted $z\\!\\simeq\\!4\\!-\\!6$ number counts from \\citet{bouw07} to show the change in the surface densities (number per arcmin$^2$) as a function of redshift.  \\citet[$z\\!\\simeq\\!3$]{stei99} used ground-based imaging in $\\sim$14 fields, with each field observed for many kilo-seconds (ks), followed by ground-based spectroscopy to confirm many of their color selected candidates. The \\citet{stei99} selection was based on LBG color criteria down to AB$\\sim$25~mag. \\citet{noni09} observed the GOODS-South field using VLT/VIMOS to get deep $U$-band imaging (AB$\\sim$27~mag). Their number counts for LBG candidates at $z\\!\\simeq\\!3$ shown in \\figref{fig:ncounts} come from the deepest part of the VIMOS field, which covers $\\sim$88~\\sqmin\\ with exposure time of $\\sim$20~hours (72~ks). On the other hand, \\citet{ly09} observed the Subaru Deep Field \\citep{kash04} using deep ($>$100~ks) near-UV imaging from the space-based GALEX observations (with $\\sim$5\\arcsec\\ FWHM resolution) to select LBG candidates at $z\\!\\simeq\\!2.2$ down to AB$\\sim$25~mag, and used ground-based spectroscopy to confirm many LBGs at $z\\!\\simeq\\!2.2$.  From \\figref{fig:ncounts} (top panel) we note three major points. First, there is only one space-based --- GALEX --- LBG survey at $z\\!\\simeq\\!2.2$ \\citep{ly09}, which clearly shows that the WFC3 UV observations --- with better sensitivity and resolution --- can play a vital role in identifying LBGs at $z\\!\\lesssim\\!2.5$.  Secondly, all surveys mentioned above use deep UV imaging with $\\gtrsim70$~ks exposures, while our WFC3 UV observations are only $\\lesssim5$~ks (1 to 2 orbits), and still we find that our observations are $\\sim$0.5--1.0~mag deeper compared to some of these surveys.  Finally, our numbers agree very well with the decreasing trend of LBG surface densities as a function of redshift from $z\\!\\simeq\\!2.0$ to $6.0$, which we will address quantitatively in the next section.  The bottom panel of \\figref{fig:ncounts} shows number counts for each dropout sample. The \\wfcnuv- and \\wfcuv-dropout samples show comparable number counts and agree generally with surveys at higher redshifts, but the \\wfcfuv-dropouts show lower number counts.  Given the trend with redshift shown by other samples in the upper panel of \\figref{fig:ncounts}, we would expect more \\wfcfuv-dropouts than other dropouts at higher redshifts. The numbers are smaller than expected because of the conservative selection criteria we applied owing to the absence of a second filter below the Lyman break (see \\secref{sample}) to confirm our dropout selection.  This approach led us to small numbers of \\wfcfuv-dropouts in a relatively narrow redshift range around $z\\!\\simeq\\!1.7$. Hence, we don't have a fully representative \\wfcfuv-dropout sample, but with the future deeper observations we can use a somewhat more liberal selection criteria to get better statistics for this sample.  \\subsection{Determination of the UV Luminosity Function}\\label{lfs}  We calculated the rest-frame UV luminosity functions (LF) using the $V_{\\rm eff}$ method \\citep[e.g.,][]{stei99,sawi06,ly09} in 0.5~mag wide bins.  The absolute magnitudes of LBG candidates were measured in the observed bands that are equivalent to rest-frame 1500~\\AA\\ to minimizes $k$-corrections, and using the average redshift for each object in each sample ($z\\!\\simeq\\!1.7$, $2.1$, $2.7$, respectively).  These absolute magnitudes are uncorrected for internal dust absorption.  We compute LFs for the three dropout samples: \\wfcfuv-dropouts ($z\\!\\simeq\\!1.7$ LBG candidates), \\wfcnuv-dropouts ($z\\!\\simeq\\!2.1$ LBG candidates), and \\wfcuv-dropouts ($z\\!\\simeq\\!2.7$ LBG candidates).  \\figref{fig:lf} shows the LFs for these three dropout samples. We model these LFs with a standard Schechter function \\citep{sche76}, which is parametrized by the characteristic absolute magnitude ($M^{*}$), the normalization ($\\phi^{*}$), and the faint-end slope ($\\alpha$).  The shaded gray regions in \\figref{fig:lf} show the uncertainty in the LF based on 1-$\\sigma$ uncertainty in $M^*$ and $\\alpha$. \\tabref{tab:lfs} lists the best-fit Schechter function parameters $M^*$, $\\alpha$ and $\\phi^*$ for these three dropout samples.  To investigate incompleteness in each redshift bin, we ran simulations to calculate $P(m,z)$, which is the probability that a galaxy of apparent magnitude $m$ and at redshift $z$ will be detected in the image \\emph{and} will meet our color selection criteria. In these simulations, large numbers of artificial objects with a range of redshifts and magnitudes were added to the real ERS2 images, and then recovered using exactly the same method and selection criteria that were employed for the real observations. For these simulations, we used BC03 models assuming Salpeter Initial Mass Function (IMF), constant SFR, solar metallicity, $E(B-V)=0.0-0.3$, an age of 1~Gyr with different redshift range for each sample and varying magnitudes. These models were used to generate color and extinction properties of our artificial objects. We chose artificial objects to be point-like sources. The selection function obtained from this exercise (adding and recovering artificial objects) is similar in shape as the distributions in \\figref{fig:redshifts}, and the mean redshift value obtained from these simulations (for each sample) is within 1-sigma of the mean value obtained in \\figref{fig:redshifts}.  These $P(m,z)$ estimates were used to determine $V_{\\rm eff}$ for the LF.  We did not make the corrections for interlopers in our LF estimates. There are five main reasons for this.  First, we have checked our LFs by boosting the errors by 10\\%, and we find that the best fit values remain the same, while the uncertainties on these values increases slightly.  Second, the limited number of spectroscopic redshifts ($\\lesssim30$\\% candidates have spectroscopic redshifts) does not give us a correct estimate of interlopers in our color selected sample. Third, the total fraction of spectroscopic interlopers is very small ($\\lesssim9$\\%), and when we subdivide them as a function of magnitude it is even smaller. Fourth, the estimate of interlopers based on photometric redshifts is not very accurate because of uncertainty in the photometric redshifts. Though the ERS2 photometric redshifts (Cohen~et~al.~2010, in~prep) are better than some of the publicly available photometric redshifts, they are still uncertain by a few percent. Finally, for the \\wfcnuv- and \\wfcuv-dropouts the faintest bin is most affected by the interlopers, but that data point is already uncertain because of very few objects in that bin.  \\subsubsection{Luminosity Functions}  The leftmost panel of \\figref{fig:lf} shows the resulting LF for \\wfcfuv-dropouts. The three brightest bins contain on average 3 objects per bin, and hence they are more uncertain. That leaves us with only three data points with a statistically significant number of objects. It is not possible to fit a Schechter function to three data points by keeping all three parameters free. In the absence of deeper data, we fix the faint-end slope, $\\alpha$, based on the best-fit observed trend between redshifts and $\\alpha$ for LBGs at $z\\!\\simeq\\!1.5\\!-\\!8$ (\\figref{fig:mstr}). The best-fit parameters for this dropout sample are meant to be mostly illustrative due to low number statistics.  The middle panel of \\figref{fig:lf} shows the LF for the \\wfcnuv-dropouts and the rightmost panel shows the LF for the \\wfcuv-dropouts.  It is difficult to estimate the faint-end slope from these observations, as we can see from the uncertain faintest point in the LFs of \\wfcnuv- and the \\wfcuv-dropouts. Though our best-fit estimates are very close to what we expect at these redshifts (\\uvdrops) from other studies \\citep[e.g.,][]{stei99,ly09} at nearby redshifts, we will need deeper ($\\sim$1--2 mag) UV observations to properly constrain the faint-end slope for these three dropout samples.  \\subsubsection{Redshift Evolution of $M^*$ and $\\alpha$}  In general, it is not straightforward to directly compare our LFs with those from previous studies. First, our redshift range is different, and this is the first time that this camera and filter set have been used to select LBGs.  Secondly, in some cases the adopted cosmologies are slightly different.  It is well known \\citep[e.g.,][]{sawi06} that the derived LFs strongly depends on the assumed cosmological models, but the evolutionary trends seen in the LFs in our three redshift bins are virtually independent of the assumed cosmology. \\figref{fig:mstr} shows the evolutionary trends in our three redshift bins, as well as comparisons to other studies on LBGs at different redshifts.  The top panel of \\figref{fig:mstr} shows the faint-end LF slope, $\\alpha$, as function of redshift.  The \\citet{arno05} $z\\!\\lesssim\\!1.5$ sample is based on the spectroscopically confirmed galaxies with the GALEX near-UV detection ($\\lesssim24.5$~mag), and the $z\\!>\\!1.5$ sample is based on the photometric redshifts.  The \\citet{arno05} samples are not selected based on Lyman break color criteria but because of the lack of LBG candidates at $z\\!\\lesssim\\!2.0$, we have used this star-forming galaxies sample for comparison.  The \\citet{redd09} and the \\citet{ly09} samples are dropout selected LBGs at $z\\!\\simeq\\!3$ and $z\\!\\simeq\\!2.2$, respectively.  The black line is the best-fit observed trend between $\\alpha$ and $z$ for LBGs at $z\\!\\simeq\\!1.5\\!-\\!8$, which is very similar to that of \\citet{ryan07}. The observed trend is that as the redshift increases, the faint-end slope, $\\alpha$, becomes steeper (more negative), illustrating that lower luminosity dwarf galaxies dominate the galaxy population at higher redshifts. Excluding the fixed $\\alpha$ data point at $z\\!\\simeq\\!1.7$, our data points at $z\\!\\simeq\\!2.1$ and $z\\!\\simeq\\!2.7$ agree very well --- within the current uncertainties --- with the black line, as well as with other data points in close redshift proximity.  The bottom panel of \\figref{fig:mstr} shows the characteristic absolute magnitude, $M^*$, as a function of redshift. Again, the general observed trend is that as redshift increases, the characteristic absolute magnitude, $M^*$, becomes brighter (more negative) until $z\\!\\simeq\\!3.5$.  This trend is considered as an evidence of `downsizing' galaxy formation scenario \\citep[e.g.,][]{cowi96}, where luminous massive galaxies form at higher redshifts. Our first data point at $z\\!\\simeq\\!1.7$ follows this general trend, but it is more uncertain due to the limited statistics in this dropout sample. The other two data points at $z\\!\\simeq\\!2.1$ and $2.7$ fit very well within the evolutionary trend seen at these and higher redshifts.  \\figref{fig:mstr} shows rapid decline of $M^*$ between $z\\!\\simeq\\!3$ and extending to $z\\!\\simeq\\!1.5$. This turnover is well defined in our and \\citet{ly09} samples. It is important for future surveys to exploit the special capabilities of the WFC3 in the near-UV to obtain larger samples to understand the relation between this critical transition in $M^*$ and physical processes in LBGs at $z\\!<\\!3$.  \\citet{redd09} have used deep ground-based imaging data to constrain the UV LF of the `BX' \\citep[e.g.,][]{adel04} population at $1.9\\!<\\!z\\!<\\!2.7$, which selects star-forming galaxies based on $U_{n}GR$ colors. When we compare our LFs with that of the `BX' population, we find some differences in $M^*$ and $\\alpha$ values. First, our \\wfcfuv-dropout sample has lower redshift ($z\\!\\sim\\!1.7$) compared to the `BX' population ($z\\!\\sim\\!2.3$), so our $M^*$ value (--19.43~mag) is fainter than their value of --20.70~mag, and agrees with the general trend discussed above. Second, $M^*$ values of our \\wfcnuv-dropout sample and the `BX' sample agree within our 1-$\\sigma$ uncertainty, while the $\\alpha$ is little steeper for the `BX' sample.  We believe that complete agreement between our LBGs sample and the `BX' sample is not possible, because although the `BX' selection selects star-forming galaxies, it is very likely that the dropout selected sample at a similar redshift might not be same as the `BX' selected sample.  Some galaxies which are selected through the `BX' color selection criteria might not be in the dropout selected sample, and vice versa. Therefore, it is difficult to directly compare the `BX' and the LBG samples, and the differences in these samples could cause the LF parameters at similar redshift to differ \\citep[see also][]{ly09}.  Therefore, for both $M^*$ and $\\alpha$ our results agree very well with the expected observed trends (\\figref{fig:mstr}) as a function of redshift. At lower redshifts ($z\\!<\\!3$), our data points will help to reduce the gap between the well studied $z\\!\\gtrsim\\!3$ and $z\\!\\sim\\!0$ regimes.  The agreement with the observed evolutionary trend of M$^*$ and $\\alpha$ also show the reliability of our LFs, which can be improved with the future deeper and wider WFC3 UV observations (e.g., CANDELS Multi-Cycle Treasury program \\# 12060-12064, PI: S.~Faber)."
    },
    "1004/1004.0849_arXiv.txt": {
        "abstract": "{Cygnus X-1 (Cyg X-1) is a high mass X-ray binary system, known to be a black hole candidate and one of the brightest sources in the X-ray sky, which shows both variability on all timescales and frequent flares. The source spends most of the time in a hard spectral state, dominated by a power-law emission, with occasional transitions to the soft and intermediate states, where a strong blackbody component emerges.} {We present the observation of Cyg X-1 in a hard spectral state performed during the AGILE science verification phase and observing cycle 1 in hard X-rays (with SuperAGILE) and gamma rays (with the gamma ray imaging detector) and lasting for about 160 days with a live time of $\\sim 6$ Ms.} {We investigated the variability of Cyg X-1 in hard X-rays on different timescales, from $\\sim 300$ s up to one day, and we applied different tools of timing analysis, such as the autocorrelation function, the first-order structure function, and the Lomb-Scargle periodogram, to our data (from SuperAGILE) and to the simultaneous data in soft X-rays (from RXTE/ASM). We concluded our investigation with a search for emission in the energy range above 100 MeV with the maximum likelihood technique.} {In the hard X-ray band, the flux of Cyg X-1 shows its typical erratic fluctuations on all timescales with variations of about a factor of two that do not significantly affect the shape of the energy spectrum. {From the first-order structure function, we find that the X-ray emission of Cyg X-1 is characterized by \\textit{antipersistence} (anticorrelation in the time series, with an increase in the emission likely followed by a decrease), indicative of a negative feedback mechanism at work.} In the gamma ray data a statistically significant point-like source at the position of Cyg X-1 is not found, and the upper limit on the flux is $\\mathrm{5 \\times 10^{-8} \\; ph \\; cm^{-2} \\; s^{-1}}$ over the whole observation (160 days). Finally we compared our upper limit in gamma rays with the expectation of various models of the Cyg X-1 emission, both of hadronic and leptonic origin, in the GeV -- TeV band.} {The time history of Cyg X-1 in the hard X-ray band over 13 months (not continuous) is shown. Different analysis tools do not provide fully converging results of the characteristic timescales in the system, suggesting that the timescales found in the structure function are not intrinsic to the physics of the source. While Cyg X-1 is not detected in gamma rays, our upper limit is a factor of two lower than the EGRET one and is compatible with the extrapolation of the flux measured by COMPTEL in the same spectral state.} ",
        "introduction": "Cygnus X-1 (Cyg X-1) is one of the brightest X-ray sources in the sky. It is a binary system composed of a compact object and the O9.7 Iab supergiant star HDE 226868, filling $97 \\; \\%$ of its Roche Lobe, with a mass ranging between $\\sim 15$ and $\\sim 30 \\;\\MSun$ \\citep[see for example][]{Gierlinski_et_al_1999,Caballero_et_al_2009}. The measurement of the mass of the compact object, with a range between 4.8 $\\MSun$ and 14.7 $\\MSun$ by \\citet{Herrero_et_al_1995} and $8.7 \\pm 0.8 \\; \\MSun$ from \\citet{Shaposhnikov_Titarchuk_2007}, suggests identification with a black hole. The distance to the binary system is measured as $2.1 \\pm 0.1$ kpc by \\citet{Ziolkowski_2005}.  A characteristic feature of the black hole binaries (as Cyg X-1) in the X-ray band, discussed for example by \\citet{Frontera_et_al_2001} between 0.5 and 200 keV, is the existence of two well distinct emission states: ``low/hard'' and ``high/soft''. The typical energy spectrum in the low/hard state, in which the source spends most of its time, is described well by a power-law ($E^{-\\Gamma}$) with photon index $\\Gamma \\sim$1.7 and a high-energy cutoff at $\\sim 150$ keV. Instead, in the high/soft state the source is characterized by a strong blackbody component with $kT \\sim 0.5$ keV and a soft power-law tail with $\\Gamma$ usually ranging between 2 and 3. An ``intermediate'' spectral state also exists, discovered by \\citet{Belloni_et_al_1996} in observations with the Proportional Counter Array (PCA) aboard the \\textit{Rossi X-Ray Timing Explorer} (RXTE), in which the flux is higher of about a factor of two with respect to the low/hard state and the spectrum is softer (with a photon index of $\\sim 2.1$ and a blackbody component of $\\sim 0.36$ keV temperature).  It is useful to note that the definition of low/hard and high/soft states derives from the observations at soft X-rays. When observing in hard X-rays, the condition reverses: the source is a factor of $\\sim$2 brighter in the low/hard state than in the high/soft one. The typical flux in the low/hard state is $\\sim 8 \\times 10^{-9} \\; \\CGS$ in 20 -- 40 keV and $\\sim 3 \\times 10^{-8} \\; \\CGS$ in 1 -- 10 keV. Although these definitions may be misleading when applied to the observations in hard X-rays reported in this paper, throughout the text we use the classification low/hard versus high/soft to comply with the classical literature on this source.  Cyg X-1 is a highly variable source in X-rays. Its variability is observed on any timescale, from months to seconds \\citep[see e. g.][]{Brocksopp_et_al_1999,Pottschmidt_et_al_2003,Ling_et_al_1997}. In particular, in the hard X-ray range the experiments of the \\textit{Interplanetary Network} detected seven episodes of giant flaring in the 15 -- 300 keV energy range, with a duration of 0.9 to 28 ks, peak flux of order of $10^{-7} \\; \\mathrm{erg \\; cm^{-2} \\; s^{-1}}$ and fluence ranging from $5 \\times 10^{-5} \\; \\mathrm{erg \\; cm^{-2}}$ to $8 \\times 10^{-4} \\; \\mathrm{erg \\; cm^{-2}}$ \\citep[see][]{Golenetskii_et_al_2003}. These outbursts were detected during both low/hard and high/soft spectral states, and, in general, the giant bursting events seem to maintain the spectral parameters (and likely the emission mechanisms) of the underlying state.  Recently, Cyg X-1 has been observed by \\citet{Albert_et_al_2007} above 150 GeV energy with the \\textit{Major Atmospheric Gamma Imaging Cherenkov} (MAGIC) telescope. A TeV excess of 4.1 $\\sigma$ compatible with a point-like source and spatially consistent with Cyg X-1 was observed simultaneously with a hard X-ray flare taking place in the low/hard state ($\\sim$1.5 Crab - $1.2 \\times 10^{-8} \\; \\mathrm{erg \\; cm^{-2} \\; s^{-1}}$ - in 20 -- 40 keV with INTEGRAL \\citep{Malzac_et_al_2008} and $\\sim$1.8 Crab in 15 -- 50 keV of Swift/BAT and $\\sim$0.6 Crab in 2 -- 12 keV of RXTE/ASM). The TeV excess was detected at the rising edge of the hard X-ray peak, one day before its maximum, while no variation is found in the soft X-ray emission. The source does not show any steady emission in the TeV band and the upper limits above $\\sim 150$ GeV energy at the 95 \\% confidence level reach 2 \\% of the Crab Nebula flux.  Although Cyg X-1 is a well-known source in X-rays, very little information is known about its emission in gamma rays. During the first three cycles of observation (1991 -- 1994) of the \\textit{Compton Gamma Ray Observatory} (CGRO), the source, in hard state, was detected by COMPTEL only between 2 and 5 MeV \\citep[see][for details]{McConnell_et_al_2000}. EGRET did not detect the source during that observation and the upper limit to the flux is of the order of $10 \\times 10^{-8} \\; \\mathrm{ph \\; cm^{-2} \\; s^{-1}}$, posing no need for a high-energy cut-off. In the soft state, the spectrum of Cyg X-1 can be modelled with a power-law that extends with the same photon index ($\\sim 2.5$ -- 3) beyond 1 MeV and up to about 10 MeV as detected by COMPTEL \\citep{McConnell_et_al_2002}. Unfortunately in this case, no EGRET measurement is available.  In steady conditions, the radio emission of Cyg X-1 is stable during low/hard states \\citep[see][]{Gleissner_et_al_2004}, except for rarely observed flares \\citep{Fender_et_al_2006}, and appears to be quenched below a detectable level during the high/soft state \\citep{Brocksopp_et_al_1999}. A non-thermal radio jet, extending up to $\\sim15 \\times 10 ^{-3}$ arcsec with an opening angle less than 2$\\degree$, was detected in VLBA observations by \\citet{Stirling_et_al_2001}.  The X-ray flares and a radio-emitting jet allow the classification of Cyg X-1 as a microquasar, as described by \\citet{Mirabel_Rodriguez_1999}. The high-energy particles in the radio-emitting jet are likely to produce gamma rays as well \\citep[see][]{Dubus_2007}. Following the model by \\citet{Zdziarski_Gierlinski_2004}, another possible source of gamma rays, especially in the low/hard spectral state, is the high-energy tail of the electron distribution in the corona. Instead in the high/soft state, the energy spectrum does not show an energy cutoff, as confirmed by the COMPTEL detection up to $\\sim 10$ MeV \\citep{McConnell_et_al_2002} and a gamma ray emission above this energy is expected.  AGILE \\citep{Tavani_et_al_2008} is the first satellite mission sensitive to gamma rays (in 30 MeV -- 30 GeV) flown after EGRET. It observed Cyg X-1 for $\\sim 160$ days in four different pointings during its first observing cycle. Thanks to the X-ray monitor SuperAGILE \\citep{Feroci_et_al_2007}, the source is observed at the same time in the hard X-ray band (between 20 and 50 keV). AGILE can continuously monitor the source, with a duty cycle of about 50 \\%. Other scanning instruments, such as RXTE/ASM or Swift/BAT, observe each source many times a day but for a shorter duration, and the duty cycle is usually shorter than 10 \\%.  In this paper we report the probably longest uninterrupted observation of Cyg X-1 in the hard X-ray and gamma ray bands from the AGILE data, complemented and extended at lower energy (2 -- 12 keV) with the information from the public web archive\\footnote[1]{\\texttt{http://heasarc.nasa.gov/xte\\_weather/}} \\footnote[2]{\\texttt{http://xte.mit.edu/ASM\\_lc.html}} of the All Sky Monitor (ASM) aboard RXTE. After the summary of the observations (given in Sect. 2) and the description of the methods of analysis (Sect. 3), we report the results of the data analysis in X-rays (Sect. 4) and gamma rays (Sect. 5). Finally, in Sect. 6 we discuss our findings and draw our conclusions in Sect. 7. ",
        "conclusions": "We reported the first campaign of observation of the hard spectral state of Cyg X-1 in gamma rays ($>100$ MeV) on various timescales, with exposures ranging from one day up to $\\sim 160$ days. We monitored the source simultaneously in hard (20 -- 50 keV) and soft (2 -- 12) X-rays using SuperAGILE and RXTE/ASM. The observation in hard X-rays by SuperAGILE shows the well-known erratic variability of the flux of Cyg X-1. The analysis of both the hardness ratio, estimated from the SuperAGILE flux in the two energy bands 20 -- 25 keV and 25 -- 50 keV, and of the colour-colour diagram, obtained from the public data of RXTE/ASM following the method reported by \\citet{Reig_et_al_2002}, does not show any transition of spectral state. We adopted exposure times ranging from $\\sim 300$ s up to one day and found the typical short time flickering of Cyg X-1 with intensity variations by about a factor of two.  From the study of the first-order structure function, we did not find any short lag plateau (``horizontal branch''). A power-law behaviour is found, with the Hurst exponent smaller than 0.5, denoting that the emission mechanism of Cyg X-1 is \\textit{antipersistent}, i. e. dominated by negative feedback. We also found timescales from the minima in the structure function but, from the analysis of the autocorrelation function and the Lomb-Scargle periodogram, we derive that these timescales are more probably related to particular patterns in the specific lightcurve, such as the distance between the repeated minima or the peaks of the emission, rather than to intrinsic properties of the source.  Cyg X-1 is not detected as a point like source above 100 MeV, and we find values of the upper limit ranging from $\\sim 100 \\times 10^{-8} \\; \\mathrm{ph \\; cm^{-2} \\; s^{-1}}$ in one day down to $\\mathrm{\\sim 5 \\times 10^{-8} \\; ph \\; cm^{-2} \\; s^{-1}}$ for the whole observation, about a factor of two lower than from the EGRET data. We compared the luminosity derived from the AGILE upper limit above 100 MeV (assuming a distance of 2 kpc to the source) with various models of GeV -- TeV emission. The predictions of the pairs' cascade model \\citep[proposed by][to explain the TeV flaring emission detected by MAGIC]{Zdziarski_et_al_2009} are compatible with our upper limit only for flaring emission (e. g. one-day timescale), while they are about one order of magnitude higher when compared to the year-long upper limit. We compared our results with the predictions of two alternative models, involving hadronic interactions between the cold matter of the stellar wind and the relativistic jet, proposed by \\citet{Araudo_et_al_2009} and \\citet{Orellana_et_al_2007}. The first model is more suitable for flaring emission, while the second one can explain the persistent emission better. We find that the luminosity in the $\\sim 100$ MeV range is compatible with the results of the AGILE monitoring for either short transient or for selected model parameters."
    },
    "1004/1004.2885_arXiv.txt": {
        "abstract": "We compute maps of CMB temperature fluctuations seeded by cosmic strings using high resolution simulations of cosmic strings in a Friedmann-Robertson-Walker universe.   We create full-sky, 18$^\\circ$ and $3^\\circ$ CMB maps, including the relevant string contribution at each resolution from before recombination to today.  We extract the angular power spectrum from these maps, demonstrating the importance of recombination effects.  We briefly discuss the probability density function of the pixel temperatures, their skewness,  and kurtosis. ",
        "introduction": "Despite improving observational limits, interest in cosmic strings has remained durable (for a review, see~\\cite{vilenkin94}).  Strings are a generic phenomena in fundamental theories and they can emerge in macroscopic form in braneworld cosmologies, for example, at the end of inflation~\\cite{sarangi2002,copeland2004}.   They are also common to cosmologically viable supersymmetry grand unified theory models~\\cite{jeannerot2003}. Stringent constraints on strings are important, therefore, in restricting the latitude available for cosmological model building.  The detection of cosmic strings would be a watershed for high energy theory.  Despite the potential significance, the investigation of cosmic strings and their observational consequences faces many numerical and analytic challenges, not least in creating accurate realizations of string imprints in the cosmic microwave sky.  In this paper, we take this study a step further forward by presenting full-sky and small-angle CMB maps of temperature fluctuations seeded by cosmic string networks using high resolution simulations in an Friedmann-Robertson-Walker expanding universe (with the longest dynamic range to date).  This work includes all the relevant recombination physics and can be used not only to determine the angular power spectrum of string CMB anisotropies but also the higher order correlators such as the bispectrum, trispectrum, and beyond.  Current constraints on cosmic strings result from line-of-sight CMB power spectrum calculations sourced either by unequal-time correlators obtained from field theory string simulations~\\cite{pen97,durrer99,bevis2004} or semianalytic models of Nambu strings~\\cite{battye98,pogosian99}.  Qualitatively these two approaches produce consistent spectra, that is, without the strong coherent acoustic peaks associated with inflation.  However, quantitatively there is a mismatch between the two approaches in both the shape of the primary peak and its amplitude, which differs by a factor of 2--3.  This disparity arises primarily from a difference in string network densities, which has been discussed at some length elsewhere~\\cite{martins2004} (see also \\cite{battye2008}). Nevertheless, there is general agreement that the relative amplitude of string induced CMB fluctuations cannot exceed more than 10\\% of those arising from adiabatic inflationary perturbations \\cite{battye98,bevis2004}.  There have also been a number of studies going beyond the power spectrum through map making with cosmic strings~\\cite{bennett88, allen96, lacmb, fraisse2008} in order to study the degree of Gaussianity of the resulting CMB signatures. However, this work has generally only included late-time gravitational effects, ignoring the recombination physics which makes an important contribution to the signal over a wide range of multipoles $l\\approx 200$-$2000$.  The motivation for the present work, then, is twofold: first, to include all recombination effects in the string CMB maps, so that we can ultimately characterize their primary statistical properties, and secondly, to match the accuracy of future experiments such as Planck~\\cite{bluebook}, AMI~\\cite{AMI2008} and QUIET~\\cite{QUIET2007} which will impose considerably more stringent constraints on cosmic strings through improving precision, resolution, and added polarization information. ",
        "conclusions": "We have computed maps of cosmic string induced CMB fluctuations at various resolutions and we have extracted their angular power spectra.  We have demonstrated the importance of recombination effects for the power spectrum over a broad range of multipoles $200< l<2000$.    We have also shown that the resulting maps are remarkably Gaussian, though with potential deviations which are worthy of closer investigation as testable string signatures in the CMB."
    },
    "1004/1004.5377_arXiv.txt": {
        "abstract": "In recent work (Seljak, Hamaus and Desjacques 2009) it was found that weighting central halo galaxies by halo mass can significantly suppress their stochasticity relative to the dark matter, well below the Poisson model expectation. This is useful for constraining relations between galaxies and the dark matter, such as the galaxy bias, especially in situations where sampling variance errors can be eliminated. In this paper we extend this study with the goal of finding the optimal mass-dependent halo weighting. We use $N$-body simulations to perform a general analysis of halo stochasticity and its dependence on halo mass. We investigate the stochasticity matrix, defined as $C_{ij}\\equiv\\langle(\\delta_i -b_i\\delta_m)(\\delta_j-b_j\\delta_m)\\rangle$, where $\\delta_m$ is the dark matter overdensity in Fourier space, $\\delta_i$ the halo overdensity of the $i$-th halo mass bin, and $b_i$ the corresponding halo bias. In contrast to the Poisson model predictions we detect nonvanishing correlations between different mass bins. We also find the diagonal terms to be sub-Poissonian for the highest-mass halos. The diagonalization of this matrix results in one large and one low eigenvalue, with the remaining eigenvalues close to the Poisson prediction $1/\\bar{n}$, where $\\bar{n}$ is the mean halo number density. The eigenmode with the lowest eigenvalue contains most of the information and the corresponding eigenvector provides an optimal weighting function to minimize the stochasticity between halos and dark matter. We find this optimal weighting function to match linear mass weighting at high masses, while at the low-mass end the weights approach a constant whose value depends on the low-mass cut in the halo mass function. This weighting further suppresses the stochasticity as compared to the previously explored mass weighting. Finally, we employ the halo model to derive the stochasticity matrix and the scale-dependent bias from an analytical perspective. It is remarkably successful in reproducing our numerical results and predicts that the stochasticity between halos and the dark matter can be reduced further when going to halo masses lower than we can resolve in current simulations. ",
        "introduction": "The large-scale structure (LSS) of the Universe carries a wealth of information about the physics that governs cosmological evolution. By measuring LSS we can attempt to answer such fundamental questions as what the Universe is made of, what the initial conditions for the structure in the Universe were, and what its future will be. Traditionally, the easiest way to observe it is by measuring galaxy positions and redshifts, which provides the 3D spatial distribution of LSS via so-called redshift surveys (e.g., \\cite{SDSS4,SDSS7}).  However, dark matter dominates the evolution and relation to fundamental cosmological parameters, while galaxies are only biased, stochastic tracers of this underlying density field. On large scales, this bias is expected to be a constant offset in clustering amplitude relative to the dark matter, which can be removed to reconstruct the dark matter power spectrum~\\cite{Bias}. Nevertheless, this reconstruction is hampered due to a certain degree of randomness in the distribution of galaxies, which is based on the nonlinear and stochastic relation between galaxies and the dark matter. In the simplest model one describes this stochasticity with the Poisson model of \\emph{shot noise}. Shot noise constitutes a source of error in the power spectrum \\cite{Stochasticity} and therefore limits the accuracy of cosmological constraints. The Poisson model predicts it to be determined by the inverse of the galaxy number density, assuming galaxies to be random and pointlike tracers. However, galaxies are born inside dark matter halos and for these extended, gravitationally interacting objects, the shot noise model is harder to describe. It is thus desirable to develop estimators that are least affected by this source of stochasticity.  In Fourier space the stochasticity of galaxies is usually described by the \\emph{cross-correlation coefficient} \\begin{equation} r_{gm}\\equiv \\frac{P_{gm}}{\\sqrt{\\hat{P}_{gg}P_{mm}}}\\;, \\end{equation} where $\\hat{P}_{gg}$ is the measured galaxy autopower spectrum, $P_{mm}$ the dark matter autopower spectrum, and $P_{gm}$ the cross-power spectrum of the two components. The cross-correlation coefficient $r_{gm}$ can be related to the shot noise power $\\sigma^2$, which is commonly defined via the decomposition $\\hat{P}_{gg}=P_{gg}+\\sigma^2$, with $P_{gg}=b^2P_{mm}$ and the bias defined as $b=P_{gm}/P_{mm}$. This yields \\begin{equation} \\frac{\\sigma^2}{P_{gg}}=\\frac{1-r_{gm}^2}{r_{gm}^2} \\;. \\label{sigma-r} \\end{equation} Thus, the lower the shot noise, the smaller the stochasticity, i.e., the deviation of the cross-correlation coefficient from unity. Minimizing this stochasticity is important if one attempts to determine the relation between galaxies and the underlying dark matter. One example of such an application is correlating the weak lensing signal, which traces dark matter, to properly radially weighted galaxies \\cite{Pen}: an accurate determination of the galaxy bias can be combined with a 3-dimensional galaxy redshift survey to greatly reduce the statistical errors relative to the corresponding 2-dimensional weak lensing survey.  The ultimate precision on how accurate the galaxy bias can be estimated from such methods is determined by the cross-correlation coefficient and previous work has shown that it can deviate significantly from unity for uniformly weighted galaxies or halos \\cite{Stochasticity}. However, it was demonstrated recently that weighting halos by mass considerably reduces the stochasticity between halos and the dark matter \\cite{sn_letter}. The purpose of this paper is to explore this more systematically and to develop an optimal weighting method that achieves the smallest possible stochasticity.  Our definition of the shot noise above is relevant for the methods that attempt to cancel sampling variance (or \\emph{cosmic variance}) \\cite{fnl_cv,beta_cv} and this will be our primary motivation in this paper. Alternatively, the shot noise is often associated with its contribution to the error in the power spectrum determination, this error usually being decomposed into sampling variance and shot noise. Sampling variance refers to the fact that in a given volume $V$ the number $N_k$ of observable Fourier modes of a given wave vector amplitude is finite. In the case of a Gaussian random field the relative error in the measured galaxy power spectrum $\\hat{P}_{gg}$ due to the sum of the two errors is given by $\\sigma_{\\hat{P}_{gg}}/\\hat{P}_{gg}=1/\\sqrt{N_k}$ (each complex Fourier mode has two independent realizations and we only count modes with positive wave vector components). Using the above decomposition of the measured power $\\hat{P}_{gg}$ into intrinsic power $P_{gg}$ and shot noise $\\sigma^2$, one finds \\begin{equation} \\frac{\\sigma_{P_{gg}}}{P_{gg}}=\\frac{1}{\\sqrt{N_k}}\\left(1+\\frac{\\sigma^2}{P_{gg}}\\right) \\;. \\label{sigma_P} \\end{equation} This definition is ambiguous, since it leaves the decomposition of the measured power into shot noise and shot noise subtracted power unspecified. Most of the analyses so far have simply assumed the Poisson model, where the shot noise is given by the inverse of the number density of galaxies, $\\sigma^2=1/\\bar{n}$. A second possibility is to define the shot noise such that the galaxy bias estimator~$\\sqrt{P_{gg}/P_{mm}}$ becomes as scale independent as possible. The third way is to define it via the stochasticity between halos and the dark matter, i.e., the cross-correlation coefficient $r_{gm}$ as in Eq.~(\\ref{sigma-r}). We choose the third definition, but will comment on the relations to the other two methods as well.  It is important to emphasize here that the first two definitions are not directly related to the applications where the sampling variance error can be eliminated, since they do not include correlations between tracers (where the dark matter itself can also be seen as a tracer). While in this paper we focus on minimizing the error on the bias estimation using the sampling variance canceling method, there are other applications where correlating dark matter and galaxies, or two differently biased galaxy samples, allows us to reduce the sampling variance error \\cite{fnl_cv,beta_cv,Gil-Marin}. In such cases the stochasticity, or the shot noise to power ratio as defined in Eq.~(\\ref{sigma-r}), is the dominant source of error and methods capable of reducing it offer the potential to further advance the precision of cosmological tests. Indeed, since the error on the power spectrum as in Eq.~(\\ref{sigma_P}) contains two contributions, in the past there was not much interest in investigating the situation where the shot noise is much smaller than sampling variance. It is the situations where the sampling variance error vanishes that are most relevant for our study.  In this paper we will focus on the relation between halos and the underlying dark matter, using two-point correlations in Fourier space (i.e. the power spectrum) as a statistical estimator. A further step to connect halos to observations of galaxies can be accomplished by specification of a halo occupation distribution for galaxies \\cite{HOD}, but we do not investigate this in any detail. Alternatively, one can think of the halos as a sample of central halo galaxies from which satellites have been removed. ",
        "conclusions": "In a previous paper \\cite{sn_letter} it was shown that weighting dark matter halos by their mass can lead to a suppression of stochasticity between halos and the dark matter relative to naive expectations. In this work we investigated the shot noise matrix, defined as the two-point correlator $C_{ij}\\equiv\\langle(\\delta_i -b_i\\delta_m)(\\delta_j-b_j\\delta_m)\\rangle$ in Fourier space, split into equal number density mass bins. The eigensystem of this matrix reveals two nontrivial eigenvalues, one of them being enhanced, the other suppressed compared to the Poisson model expectation. It is the latter that leads to a reduced stochasticity. The optimal estimator of the dark matter and the eigenvector corresponding to the lowest eigenvalue are very similar and the latter dominates the signal-to-noise ratio of the halo density field. We fit both vectors by a smooth function of mass which we denote \\emph{modified mass weighting}. It is proportional to halo mass at the high-mass end and approaches a constant towards lower masses which is determined by the minimum halo mass resolved in the simulations. This constant is roughly 3 times the minimum halo mass over the range of masses we explored.  Applying this function to weight the halo density field results in a field that is more correlated with the dark matter with a suppressed shot noise component, improving upon previous results \\cite{sn_letter} by a factor of 2 in signal-to-noise. We investigate the effect of uncertainty in halo mass, finding that it does not change our fundamental conclusions, even if it weakens the strength of the method: a realistic amount of log-normal scatter in mass at the level of 0.5 increases the shot noise by a about a factor of 2. Our results can directly be applied to methods that attempt to eliminate sampling variance by investigating the relation between galaxies and the dark matter both tracing the same LSS. In this case the error is determined by the stochasticity between the two and reducing it can improve the ultimate reach of these methods \\cite{Pen,Stochasticity}.  Considering the halo model as a theoretical approach to describe the shot noise matrix, we find analytical expressions for its eigenvalues and eigenvectors. In particular, the two nontrivial eigenvectors can be written as a linear combination of halo bias and halo mass, which yields a considerable agreement with our simulation results. Furthermore, the two estimators of the scale-dependent bias, $\\langle\\delta_h\\delta_m\\rangle/\\langle\\delta_m^2\\rangle$ and $\\sqrt{\\langle\\delta_h^2\\rangle/\\langle\\delta_m^2\\rangle}$, are well reproduced. However, our model suffers from the lack of mass and momentum conservation: its implementation, together with higher-order perturbation theory and halo exclusion, further improves the agreement and will be presented elsewhere.  The halo model suggests the stochasticity between modified mass-weighted halos and the dark matter to decrease linearly with mass resolution below $M\\simeq10^{12}h^{-1}\\mathrm{M}_{\\odot}$, yielding a suppression by almost 2 orders of magnitude at $M_{\\mathrm{min}}\\simeq10^{10}h^{-1}\\mathrm{M}_{\\odot}$ as compared to uniformly weighted halos. While we focused on the question of how well halos can reconstruct the dark matter, our analysis is also applicable to the study of stochasticity between halos themselves. Indeed, reducing the stochasticity between different halo tracers by optimal weighting techniques, while at the same time canceling sampling variance, should be possible even if the dark matter field is not measured. This will be addressed in more detail in a future work.  Specific applications are the best way to test the efficiency of our method. There is probably not much advantage in applying it to the standard power spectrum determination, where the sampling variance error dominates the error budget in the limit of small stochasticity, while in the opposite limit of rare halos, when the shot noise power is comparable to the intrinsic halo power, we do not see much gain (as demonstrated by the fact that all the points in Fig.~\\ref{sn_hm} overlap for the highest $M_{\\mathrm{min}}$, which corresponds to halos with the lowest number density). More promising applications are those where the sampling variance error is eliminated and the error budget is dominated by stochasticity, or the ratio of shot noise power to the halo power.  In this paper we have focused on the bias determination from galaxy and dark matter correlations \\cite{Pen} as a specific application, but other applications are possible, such as constraining $f_{\\mathrm{NL}}$ from non-Gaussianity \\cite{fnl_constraint,Slosar} and the redshift-space parameter $\\beta$ from redshift-space distortions \\cite{beta,beta_cv}, to name a few. Upcoming surveys like SDSS-III \\cite{SDSS-III}, JDEM/EUCLID \\cite{JDEM,EUCLID} or BigBOSS \\cite{BigBoss} and LSST \\cite{LSST} will increase the available number of galaxies significantly, providing both 3D galaxy maps and 2D to 3D dark matter maps (via weak lensing techniques, enhanced by lensing tomography).  Our results suggest that correlating modified halo mass-weighted galaxies against the dark matter has the potential to lead to dramatic improvements in the precision of cosmological parameter estimation. We will explore more explicit demonstrations of the above mentioned applications in future work."
    },
    "1004/1004.2627_arXiv.txt": {
        "abstract": "{ The Atacama Pathfinder Experiment (APEX) 12 m telescope was used to observe the $N=1-0, J=0-1$ ground-state transitions of \\OHP\\ at 909.1588~GHz with the CHAMP+ heterodyne array receiver.  Two blended hyperfine structure transitions were detected in absorption against the strong continuum source Sagittarius B2(M) and in several pixels offset by 18\\arcsec.  Both absorption from Galactic center gas and absorption from diffuse clouds in intervening spiral arms in a velocity range from --116 to 38.5~\\kms\\ is observed.  The total \\OHP\\ column density of absorbing gas is $2.4\\times 10^{15}$~\\cmsq. A column density local to Sgr B2(M) of $2.6 \\times 10^{14}$~\\cmsq\\ is found. On the intervening line-of-sight, the column density per unit velocity interval is in the range of 1 to $40 \\times 10^{12}$~\\cmsq/(\\kms). \\OHP\\ is found to be on average more abundant than other hydrides, such as \\SHP\\ and \\CHP. Abundance ratios of OH and atomic oxygen to \\OHP\\ are found in the range of $10^{1-2}$ and $10^{3-4}$, respectively. The detected absorption of a continuous velocity range on the line-of-sight shows \\OHP\\ to be an abundant component of diffuse clouds. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.1624_arXiv.txt": {
        "abstract": "{ Phase referencing is a standard calibration procedure in radio interferometry. It allows to detect weak sources by using quasi-simultaneous observations of closeby sources acting as calibrators. Therefore, it is assumed that, for each antenna, the optical paths of the signals from both sources are similar. However, atmospheric turbulence may introduce strong differences in the optical paths of the signals and affect, or even waste, phase referencing for cases of relatively large calibrator-to-target separations and/or bad weather. The situation is similar in wide-field observations, since the random deformations of the images, mostly caused by atmospheric turbulence, have essentially the same origin as the random astrometric variations of phase-referenced sources with respect to the phase center of their calibrators. In this paper, we present the results of a Monte Carlo study of the astrometric precision and sensitivity of an interferometric array (a realization of the Square Kilometre Array, SKA) in phase-referenced and wide-field observations. These simulations can be extrapolated to other arrays by applying the corresponding corrections. We consider several effects from the turbulent atmosphere (i.e., ionosphere and wet component of the troposphere) and also from the antenna receivers. We study the changes in dynamic range and astrometric precision as a function of observing frequency, source separation, and strength of the turbulence. We find that, for frequencies between 1 and 10\\,GHz, it is possible to obtain images with high fidelity, although the atmosphere strongly limits the sensitivity of the instrument compared to the case with no atmosphere. Outside this frequency window, the dynamic range of the images and the accuracy of the source positions decrease. We also find that, even if a good model of the atmospheric turbulence (with an accuracy of 99\\%) is used in the imaging, residual effects from the turbulence can still limit the dynamic ranges of deep, high-contrast ($10^5 - 10^6$), images. } ",
        "introduction": "It is well-known that ground-based astronomical observations are affected by the atmosphere. Changes in the atmospheric opacity produce a bias in the source flux density, while changes in the refraction index distort the shape of the electromagnetic frontwave of the source. Such a distortion translates into a deformation of the observed source structure and/or a variation of the relative positions of all sources observed in a given field. In the case of astronomical devices based on interferometry, atmospheric effects can be well modelled if the atmosphere above each element of the interferometer (hereafter, {\\em station}) remains unchanged over the whole portion of the sky being observed. In such cases, the observed visibilities can be calibrated using {\\em station-based} algorithms, which are relatively simple and computationally inexpensive (e.g. Readhead \\& Wilkinson \\cite{Readhead1978}).   However, when the spatial variations of the atmosphere are significant within the observed portion of the sky, as it happens if there is atmospheric turbulence, the opacity and dispersive effects cannot be modeled as a single time-dependent station-based complex gain over the field of view.  Unless more complicated calibration algorithms are used (e.g., van der Tol, Jeffs, \\& van der Veen \\cite{vdT07}), the effect of these errors on the image are difficult to correct. In this paper, we report on a study of the effects that a turbulent atmosphere may introduce in interferometric observations. We focus our study on the effects produced by turbulence in the dynamic range and astrometric accuracy after a phase-referenced calibration between a strong (calibrator) source and a weak source, located a few degrees away. This study is numerically equivalent to the study of the deformation of a wide-field interferometric image at any point located at a given distance from the center of the field (i.e., the {\\em phase center} of the image). In both cases, the phases introduced by the atmosphere in the signal of each antenna for the different pointing directions are the same, so the effects of the atmosphere in Fourier space (and therefore on the sky plane) will also be the same.  The results here reported are an extended version of those previously reported in the SKA memo by Mart\\'i-Vidal et al. (\\cite{MartiVidal2009}). In the next section, we describe the details of the array distribution used, as well as the characteristics of the simulated observations. In Sect. \\ref{NoiseModel}, we describe how the noise from the atmosphere and the receivers was added to the visibilities and in Sect. \\ref{algorithm} describe the procedures followed in our Monte Carlo analysis. In Sect. \\ref{Results}, we present the main results obtained; in Sect. \\ref{Summary}, we summarize our conclusions. ",
        "conclusions": "\\label{Summary}  We report on Monte Carlo estimates of the sensitivity and astrometric precision of an interferometric array, with a station distribution similar to that of the planned SKA, as a function of observing frequency, flux density, and source separation. These results can also be applied to other array distributions by taking into account the corresponding correction factors. Our estimates are based on simulations of snapshot phase-referenced observations, in which we take into account several effects from the turbulent atmosphere and the finite temperature of the receivers. We find that the astrometric uncertainty strongly depends on the observing frequency and smoothly increases as the source separation increases. For frequencies below $\\sim$1\\,GHz, ionospheric effects dominate and the astrometry uncertainties (when the source is detectable) can be as large as $\\sim$1\\,as. For frequencies between 1 and 10\\,GHz (these values slightly depend on the source flux density) atmospheric effects are minimum and we roughly reach the theoretical precision of the interferometer. Above these frequencies, the wet troposphere begins to dominate and the astrometric uncertainty increases to $\\sim$10\\,mas for the highest simulated frequency (25\\,GHz). The dynamic range of the images is strongly limited by atmospheric turbulence at all frequencies and for all flux densities (it can decrease, in the worse cases, several orders of magnitude).  We propose analytical models for the loose of dynamic range and astrometric accuracy as a function of distance between calibrator and target source. These expressions could also be used to estimate the deformations and local dynamic ranges of wide-field images as a function of distance to the image phase center (i.e., the point in the sky where the data correlation is centered)."
    },
    "1004/1004.1647.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary %  {} % aims heading (mandatory) { {\\sffamily\\itshape Aims.} A significant fraction of early-type galaxies (ETGs) show emission lines in their optical spectra. We aim at understanding the powering mechanism and the origin of the ionized gas in ETGs, and its connection with the host galaxy evolution. %{  {\\sffamily\\itshape Methods.} We analyzed intermediate-resolution optical spectra of a sample of 65 ETGs, mostly located in low density environments and biased toward the presence of ISM traces, for which we already derived in the previous papers of the series the stellar population properties. To extract the emission lines from the galaxy spectra, we developed a new fitting procedure that properly subtracts the underlying stellar continuum, and that accounts for the uncertainties due to the age-metallicity degeneracy. %The procedure makes use of new SSP models based on the MILES %spectral library, which represents a substantial improvement over previous libraries used in population synthesis models. The emission line luminosities derived in annuli of increasing galacto-centric distance were used to constrain the excitation mechanism and the metallicity of the ionized gas. % results heading (mandatory) %{  {\\sffamily\\itshape Results.} Optical emission lines are detected in 89\\% of the sample. The detection fraction drops to 57\\% if only the galaxies with EW(H$\\alpha$ $+$ [NII])$>$3 \\AA \\ are considered. The incidence and strength of emission do not correlate either with the E/S0 classification, or with the fast/slow rotator classification. Comparing the nuclear r$<$r$_e$/16 emission with the classical  [OIII]/H$\\beta$ vs [NII]/H$\\alpha$ diagnostic   diagram, the galaxy activity is so classified: 72\\% of the galaxies with emission are LINERs, 9\\% are Seyferts, 12\\% are Composite/Transition objects, and 7\\% are non-classified. Seyferts have young luminosity-weighted ages ($\\lesssim$ 5 Gyr), and appear, on average, significantly younger than LINERs and Composites. Excluding the Seyfert galaxies from our sample, we find that the spread in the ([OIII], H$\\alpha$ or [NII]) emission strength increases with the galaxy central velocity dispersion $\\sigma_c$, low-$\\sigma_c$ galaxies having all weak emission lines, and high-$\\sigma_c$ galaxies displaying both weak and strong emission lines. Furthermore, the [NII]/H$\\alpha$ ratio tends to increase with  $\\sigma_c$. A spatial analysis of the emission line properties within the individual galaxies reveals that the [NII]/H$\\alpha$ ratio decreases with increasing galacto-centric distance, indicating either a decrease of the nebular metallicity, or a progressive ``softening'' of the ionizing spectrum. The average nebular oxygen abundance is slightly less than solar.  A comparison with the stellar metallicities derived in Paper~III shows that the gas oxygen abundance is $\\approx$ 0.2 dex lower than that of stars. % conclusions heading (optional), leave it empty if necessary % {  {\\sffamily\\itshape Conclusions.} The stronger nuclear (r$<r_e/16$) emission can be explained with photoionization by PAGB stars alone only for $\\approx$ 22\\% of the LINERs/Composite sample. On the other hand, we can not exclude an important role of PAGB star photoionization at larger radii. For the major fraction of the sample, the nuclear emission is consistent with excitation from a low-accretion rate AGN, fast shocks (200 -500 km/s) in a relatively gas poor environment (n $\\lesssim100$ cm$^{-3}$) , or coexistence of the two. The derived [SII]6717/6731 ratios are consistent with the low gas densities required by the shock models. % The natural explanation for  the observed trends of the [NII] intensity and of the [NII]/H$\\alpha$ ratio with $\\sigma_c$ is a scenario in which the shock mechanical energy comes from turbulent motions of clouds within the potential well of the galaxy. %A scenario in which the shock mechanical energy comes from turbulent motions of clouds within the potential well of the galaxy is very appealing, since it naturally explains the observed trends of the [NII] intensity and of the [NII]/H$\\alpha$ ratio with $\\sigma_c$. The derived nebular metallicities suggest either an external origin of the gas, or an overestimate of the oxygen yields by SN models. % } ",
        "introduction": "Early-type galaxies (ETGs) have long been considered to be inert stellar systems, essentially devoid of gas and dust. However, this view has radically changed since a number of imaging and spectroscopy studies from both the ground and space have revealed the presence of a multiphase interstellar medium (ISM): a hot ($T \\sim 10^6 - 10^7$ K), X-ray emitting halo (Forman \\& Jones~\\cite{fj85}; Fabbiano et al.~\\cite{fab92}; O'Sullivan et al.~\\cite{sul01}); a warm ($T \\sim 10^4$ K) gas (often referred to as ''ionized gas``) (e.g. Phillips et al.~\\cite{ph86}; Sadler \\& Gerhard~\\cite{sg85}; %van Dokkum \\& Franx \\cite{vdf95}; see also Goudfrooij~\\cite{gou99} for a review); and colder components detected in the mid and far infrared, CO and HI (see review of Sadler et al.~\\cite{sad02}; Temi et al.~\\cite{temi04}; Bressan et al.~\\cite{bress06}; Sage \\& Welch~\\cite{sw06}; Morganti et al.~\\cite{mo06}; Helmboldt~\\cite{hel07}; Oosterloo et al.~\\cite{oos07}; Panuzzo et al.~\\cite{panu07}). The hot gas component dominates the ISM in E/S0 galaxies, and only comparatively small quantities are detected in the warm and cold phases (Bregman et al.~\\cite{bre92}). The derived masses of the warm ISM range from $\\sim 10^3$ to $\\sim 10^6 M_{\\odot}$ (Phillips et al.~\\cite{ph86}; Trinchieri \\& di Serego Alighieri~\\cite{tds91}; Macchetto et al.~\\cite{mac96}; Goudfrooij et al.~\\cite{gou94a,gou94b,gou94c}; Goudfrooij \\& de Jong~\\cite{gouj95}), orders of magnitude below the X-ray emitting hot ISM and the stellar mass. Ionized gas is detected in 40-80\\% of early-type galaxies via its optical emission lines (Caldwell~\\cite{cal84}; Phillips et al.~\\cite{ph86}; Macchetto et al.~\\cite{mac96}; Sarzi et al.~\\cite{sar06,sar09}; Yan et al.~\\cite{yan06}; Serra et al.~\\cite{serra08}). The ionized gas is also usually morphologically associated with dust (Kim~\\cite{kim89}; Goudfrooij et al.~\\cite{gou94c}).  Despite the number of studies, several issues remain still open. The first question is the origin of the ISM in ETGs. Narrow band imaging centered around H$\\alpha + $[NII]$\\lambda$6584 shows that the ionized gas presents a variety of morphologies, from regular, disk-like structures, to filamentary structures (Demoulin-Ulrich et al.~\\cite{du84}; Buson et al.~\\cite{bus93}; Macchetto et al.~\\cite{mac96}; Zeilinger et al.~\\cite{zei96}; Martel et al.~\\cite{mar04}; Sarzi et al.~\\cite{sar06}). Evidence for the acquisition of external gas comes from kinematical studies showing that the angular momentum of the gas is often decoupled from that of the stars (Bertola et al.~\\cite{ber92}; van Dokkum \\& Franx~\\cite{vdf95}; Caon et al.~\\cite{caon00}), even if according to Sarzi et al.~(\\cite{sar06}) the angular momenta are inconsistent with a purely external origin for the gas. The fact that the ionized gas emission is always associated with dust (e.g. Tran et al.~\\cite{tr01}) tends to exclude ``cooling flows'' as the origin for the warm gas in E/S0 galaxies (Goudfrooij~\\cite{gou99}). In the scenario proposed by Sparks et al.~(\\cite{spa89}) and de Jong et al.~(\\cite{dej90}), the dusty filaments can arise from the interaction of a small gas rich galaxy with the giant elliptical, or from a tidal accretion event in which the gas and dust are stripped from a passing spiral (see e.g. Domingue et al.~\\cite{Domingue03}; Tal et al.~\\cite{Tal09}).  The history of star formation and evolution leaves its chemical signature in the ISM of ETGs. Historically, X-ray abundance measurements of the hot ISM have been problematic, as typified by the so-called Fe discrepancy (Arimoto et al.~\\cite{ari97}). More recently, {\\it Chandra} and {\\it XMM} measurements of X-ray bright galaxies at the center of groups have supported roughly solar or even slightly supersolar abundances for the  ISM (e.g. Gastaldello \\& Molendi~\\cite{gm02}; Buote et al.~\\cite{buo03}; Tamura et al.~\\cite{tamu03}), and have claimed no evidence for very subsolar Fe abundances (Humphrey \\& Buote~\\cite{hb06}; Ji et al.~\\cite{ji09}). These studies also report anomalously low oxygen abundances compared to iron and magnesium. This is in conflict with results from SN stellar yields if one assumes that the source of ISM is the material injected into interstellar space by evolved stars. Up to now, the only study of metallicity in the warm ISM is that of Athey \\& Bregman~(\\cite{ab09}), who derived oxygen abundances for seven ETGs from optical emission lines. The authors find an average solar metallicity, favoring an internal origin of the warm ISM. However, this study is limited by the small sample. Determining the metallicity of the warm ISM is fundamental to create a link between the hot gas phase and the galaxy stellar population, and to discriminate between an internal and an external origin of the interstellar matter in E/S0 galaxies.  The second still open issue concerns the ionizing source of the warm gas. Optical spectroscopic studies show that ETGs are typically classified as LINERs (Low-Ionization Nuclear Emission-line Regions, Heckman~\\cite{heck80}) according to their emission line ratios (e.g. Phillips et al.~\\cite{ph86}; Yan et al.~\\cite{yan06}). However, there is still strong debate about the ionization mechanism in LINERs. At present, the most viable excitation mechanisms are: a low accretion-rate AGN (e.g., Ho~\\cite{ho99b}; Kewley et al.~\\cite{kew06}, Ho~\\cite{ho09a}), photoionization by old post-asymptotic giant branch (PAGB) stars (e.g.  Trinchieri \\& di Serego Alighieri~\\cite{tds91}; Binette et al.~\\cite{bin94}; Stasi{\\'n}ska et al.~\\cite{sta08}), fast shocks (Koski \\& Osterbrock~\\cite{ko76}; Heckman~\\cite{heck80}; Dopita \\& Sutherland~\\cite{dosu95}; Allen et al.~\\cite{allen08}). The AGN mechanism is strongly supported by the detection of broad emission lines in optical spectra (e.g. Ho et al.~\\cite{ho97a}; Wang \\& Wei~\\cite{ww08}), by the properties of the X-ray emission (Terashima et al.~\\cite{tera02}; Flohic et al.~\\cite{flo06}; Gonzalez-Martin et al.~\\cite{go09}), by the presence of UV and X-ray variability (e.g Maoz et al.~\\cite{maoz05}, Pian et al.~\\cite{pian10}), and by the evidence that massive black holes (BHs) appear to be a generic component of galaxies with a bulge (Magorrian et al.~\\cite{mago98}; Ho~\\cite{ho99a}; Kormendy~\\cite{kor04}). PAGB stars are good candidates because their radiation, much harder than that from young stars, is able to reproduce the observed emission-line ratios in LINERs (Trinchieri \\& di Serego Alighieri~\\cite{tds91}; Binette et al.~\\cite{bin94}; Stasi{\\'n}ska et al.~\\cite{sta08}). The most compelling evidence in support of the PAGB mechanism is the finding that the emission-line flux correlates very well with the host galaxy stellar luminosity within the emission-line region (Macchetto et al.~\\cite{mac96}), and that the emission-line flux distribution follows that of the stellar continuum (Sarzi et al.~\\cite{sar06,sar09}). These correlations suggest that the source of ionizing photons is distributed in the same way as the stellar population. LINERs emission has also been observed in extranuclear regions associated with large-scale outflows and related shocks (L{\\'{\\i}}pari et al.~\\cite{li04}), or regions shocked by radio jets (Cecil et al.~\\cite{cec00}).  Evidence is growing that more than one mechanism may be at work in LINERs. In a recent study, Sarzi et al.~(\\cite{sar09}) showed that the role of AGN activity is confined to the central regions, whereas  at larger radii the stellar and nebular fluxes follow each other, thus suggesting a stellar ionizing source. Eracleous et al.~(\\cite{era10}) showed that low accretion-rate AGNs do not produce enough ionizing photons to explain the observed H$\\alpha$ luminosities within a 2'' $\\times$ 4'' aperture, so that other mechanisms are needed.    With the aim of understanding the formation history of ETGs, and the connection between the stellar populations and the ionized gas, we have analyzed intermediate-resolution optical spectra for a sample of 65 field E/S0 galaxies %biased toward the presence of ISM emission (Rampazzo et al.~\\cite{ramp05}, hereafter Paper~I; Annibali et al.~\\cite{anni06}, hereafter Paper~II). The stellar population properties (i.e. ages, metallicities, and elements abundance ratios) were derived at different radii in Annibali et al.~(\\cite{anni07}) (hereafter Paper~III) through the analysis of the Lick indices. In this paper we present the study of the optical emission lines in annuli of increasing galacto-centric distance. The purpose is to characterize the properties of the ionized gas, its origin, and the possible excitation mechanisms at different radii.  The paper is organized as it follows. In section~2 we provide a brief overview of the sample. In section~3 we describe the method used to extract the emission-line fluxes from the optical spectra. In section~4 we provide a spectral classification through standard diagnostic diagrams, and investigate possible correlations with the host galaxy properties. In section~5 we determine the oxygen metallicity. In section~6 we compare our data with the models. A summary and the conclusions are given in section~7.  %----------------------------------------------Table 1------------------------------------------------------ \\begin{table*} \\caption{Sample overview} \\label{sample} \\centering {\\tiny{ \\begin{tabular}{lccccccccc} \\hline \\hline Galaxy Id  & RSA & RC3 & $\\sigma_c$  & $r_e$  & $V_{hel}$    &   $\\rho_{xyz}$ & Age & Z &  [$\\alpha$/Fe] \\\\ &          &         &  km s$^{-1}$ & arcsec & km s$^{-1}$ & gal Mpc$^{-3}$  & Gyr &  &  \\\\ \\hline NGC~128  & S02(8) pec      & S0 pec sp     &  183    & 17.3  &  4227 &        &    9.7   $\\pm$    1.7   &    0.024   $\\pm$   0.004   &   0.16    $\\pm$   0.03    \\\\ NGC~777  & E1                   & E1                 &  317    & 34.4  & 5040  &        &    5.4   $\\pm$    2.1   &    0.045   $\\pm$   0.020   &    0.28    $\\pm$   0.10    \\\\ NGC~1052 & E3/S0             & E4                 &  215   & 33.7   &1475  & 0.49 &   14.5  $\\pm$    4.2   &    0.032   $\\pm$   0.007   &    0.34    $\\pm$     0.05    \\\\ NGC~1209 & E6                  & E6:                &  240    &18.5  & 2619 & 0.13  &    4.8    $\\pm$    0.9  &     0.051   $\\pm$   0.012   &    0.14    $\\pm$     0.02      \\\\ NGC~1297 & S02/3(0)         & SAB0 pec:    &  115    &  28.4 & 1550  & 0.71 &   15.5   $\\pm$   1.2   &    0.012   $\\pm$    0.001   &    0.29    $\\pm$     0.04    \\\\ NGC~1366 & E7/S01(7)      & S0 sp             & 120    &  10.6 & 1310 & 0.16  &    5.9    $\\pm$     1.   &     0.024   $\\pm$    0.004   &    0.08    $\\pm$     0.03     \\\\ NGC~1380 & S03(7)/Sa      & SA0               &  240   &  20.3 & 1844 & 1.54   &    4.4    $\\pm$   0.7   &     0.038   $\\pm$    0.006   &    0.24    $\\pm$     0.02   \\\\ NGC~1389 & S01(5)/SB01  & SAB(s)0-:      & 139    & 15.0 & 986 & 1.50      &    4.5     $\\pm$   0.6   &     0.032   $\\pm$    0.005   &    0.08    $\\pm$     0.02    \\\\ NGC~1407 & E0/S01(0)      & E0                  &  286   & 70.3 & 1766 & 0.42   &    8.8     $\\pm$   1.5   &      0.033   $\\pm$    0.005   &    0.32    $\\pm$     0.03    \\\\ NGC~1426 & E4                  & E4                  &  162   & 25.0 & 1443 & 0.66   &    9.0     $\\pm$    2.5   &     0.024   $\\pm$    0.005    &   0.07    $\\pm$     0.05     \\\\ & & & & \\\\ NGC~1453 & E0                  &    E2         & 289 &  25.0 & 3906 &  &    9.4         $\\pm$       2.1   &      0.033   $\\pm$      0.007     &    0.22    $\\pm$     0.03      \\\\ NGC~1521 & E3                  &    E3         & 235 & 25.5 & 4165 &   &    3.2         $\\pm$       0.4   &      0.037   $\\pm$      0.006     &    0.09    $\\pm$     0.02       \\\\ NGC~1533 & SB02(2)/SBa & SB0-          & 174  & 30.0 & 773 & 0.89  &   11.9   $\\pm$       6.9   &      0.023   $\\pm$      0.020     &    0.21    $\\pm$     0.10        \\\\ NGC~1553 & S01/2(5)pec  & SA(r)0        & 180 & 65.6 & 1280 & 0.97 &    4.8     $\\pm$       0.7   &      0.031   $\\pm$      0.004     &    0.10    $\\pm$     0.02       \\\\ NGC~1947 & S03(0) pec    & S0- pec     &  142 & 32.1 & 1100  & 0.24   &    5.9     $\\pm$       0.8   &      0.023   $\\pm$      0.003     &    0.05    $\\pm$     0.02   \\\\ NGC~2749 & E3                  & E3           &  248 & 33.7 & 4180 &  &   10.8     $\\pm$       2.3   &      0.027   $\\pm$      0.006     &    0.25    $\\pm$     0.04    \\\\ NGC~2911 & S0p or S03(2)& SA(s)0: pec  & 235  & 50.9 &  3131  & &    5.7     $\\pm$       2.0   &      0.034   $\\pm$      0.019     &    0.25    $\\pm$     0.10      \\\\ NGC~2962 & RSB02/Sa     & RSAB(rs)0+ &   & 23.3 & 2117 & 0.15  & & &  \\\\ NGC~2974 & E4                  & E4           & 220 & 24.4 & 1890  & 0.26   &   13.9     $\\pm$       3.6   &      0.021   $\\pm$      0.005     &    0.23    $\\pm$     0.06      \\\\ NGC~3136 & E4                  & E:           & 230 & 36.9 & 1731  & 0.11 &    1.5     $\\pm$       0.1   &      0.089   $\\pm$      0.004     &    0.36    $\\pm$     0.02     \\\\ & & & & \\\\ NGC~3258 & E1                  & E1          & 271 & 30.0 & 2778 & 0.72  &    4.5     $\\pm$       0.8   &      0.047   $\\pm$      0.013     &    0.21    $\\pm$     0.03      \\\\ NGC~3268 & E2                  & E2          & 227 & 36.1 & 2818 & 0.69  &    9.8     $\\pm$       1.7   &      0.023   $\\pm$      0.004     &    0.34    $\\pm$     0.04    \\\\ NGC~3489 & S03/Sa           & SAB(rs)+   & 129   & 20.3 & 693  & 0.39 &    1.7     $\\pm$       0.1   &      0.034   $\\pm$      0.004     &    0.05    $\\pm$     0.02         \\\\ NGC~3557 & E3                  & E3          & 265 & 30.0 & 3038 & 0.28 &    5.8     $\\pm$       0.8   &      0.034   $\\pm$      0.004     &    0.17    $\\pm$     0.02     \\\\ NGC~3607 & S03(3)            & SA(s)0:     & 220 & 43.4 &  934 & 0.34 &    3.1     $\\pm$       0.5   &      0.047   $\\pm$      0.012     &    0.24    $\\pm$     0.03       \\\\ NGC~3818 & E5                  & E5            & 191  & 22.2  & 1701 & 0.20   &    8.8     $\\pm$       1.2   &      0.024   $\\pm$      0.003     &    0.25    $\\pm$     0.03    \\\\ NGC~3962 & E1                  & E1          & 225 & 35.2 & 1822 & 0.32  &   10.0     $\\pm$       1.2   &      0.024   $\\pm$      0.003     &    0.22    $\\pm$     0.03      \\\\ NGC~4374 & E1                  & E1      & 282 & 50.9 & 1060 & 3.99  &    9.8     $\\pm$       3.4   &      0.025   $\\pm$      0.010     &    0.24    $\\pm$     0.08     \\\\ NGC~4552 & S01(0)            & E           &  264 & 29.3 &  322 & 2.97  &    6.0     $\\pm$       1.4   &      0.043   $\\pm$      0.012     &    0.21    $\\pm$     0.03      \\\\ NGC~4636 & E0/S01(6)      & E0-1        &  209 & 88.5 & 937 & 1.33  &   13.5     $\\pm$       3.6   &      0.023   $\\pm$      0.006     &    0.29    $\\pm$     0.06   \\\\ & & & & \\\\ NGC~4696 &(E3)            & E+1 pec  & 254  & 85.0 & 2958 & 0.00   &   16.0     $\\pm$       4.5   &      0.014   $\\pm$      0.004     &    0.30    $\\pm$     0.10   \\\\ NGC~4697 & E6              & E6      & 174  & 72.0 & 1241 & 0.60  &   10.0     $\\pm$       1.4   &      0.016   $\\pm$      0.002     &    0.14    $\\pm$     0.04       \\\\ NGC~5011 & E2              & E1-2    & 249 & 23.8 & 3104 & 0.27  &    7.2     $\\pm$       1.9   &      0.025   $\\pm$      0.008     &    0.25    $\\pm$     0.06    \\\\ NGC~5044 & E0              & E0      & 239 & 82.3 & 2704 & 0.38  &   14.2     $\\pm$       10.   &      0.015   $\\pm$      0.022     &    0.34    $\\pm$     0.17     \\\\ NGC~5077 & S01/2(4)     & E3+        & 260 & 22.8  &  2764 & 0.23  &   15.0     $\\pm$       4.6   &      0.024   $\\pm$      0.007     &    0.18    $\\pm$     0.06       \\\\ NGC~5090 & E2              & E2      & 269 & 62.4 &  3421 & &   10.0     $\\pm$       1.7   &      0.028   $\\pm$      0.005     &    0.26    $\\pm$     0.04       \\\\ NGC~5193 & S01(0)        & E pec  & 209 & 26.7 & 3711 &   &    6.8     $\\pm$       1.1   &      0.018   $\\pm$      0.002     &    0.26    $\\pm$     0.04     \\\\ NGC~5266 & S03(5) pec & SA0-: & 199 & 76.7 & 3074 & 0.35 &    7.4     $\\pm$       1.4   &      0.019   $\\pm$      0.003     &    0.15    $\\pm$     0.05    \\\\ NGC~5328 & E4              & E1:          & 303 & 22.2 & 4671 &  &   12.4     $\\pm$       3.7   &      0.027   $\\pm$      0.006     &    0.15    $\\pm$     0.05     \\\\ NGC~5363 & [S03(5)]      & I0:          &  199 & 36.1 & 1138 & 0.28 &   12.1     $\\pm$       2.3   &      0.020   $\\pm$      0.004     &    0.16    $\\pm$     0.05    \\\\ & & & & \\\\ NGC~5638 & E1                  & E1   &  168 & 28.0 & 1676 & 0.79   &    9.1     $\\pm$       2.3   &      0.024   $\\pm$      0.008     &    0.24    $\\pm$     0.05       \\\\ NGC~5812 & E0                  & E0   & 200 & 25.5 & 1930 & 0.19  &    8.5     $\\pm$       2.1   &      0.027   $\\pm$      0.008     &    0.22    $\\pm$     0.05    \\\\ NGC~5813 & E1                  & E1-2 & 239 & 57.2 & 1972 & 0.88   &   11.7     $\\pm$       1.6   &      0.018   $\\pm$      0.002     &    0.26    $\\pm$     0.04      \\\\ NGC~5831 & E4                  & E3   & 164 & 25.5 & 1656 & 0.83  &    8.8     $\\pm$       3.5   &      0.016   $\\pm$      0.011     &    0.21    $\\pm$     0.09      \\\\ NGC~5846 & S01(0)            & E0+         &250 & 62.7  & 1709 & 0.84  &    8.4     $\\pm$       1.3   &      0.033   $\\pm$      0.005     &    0.25    $\\pm$     0.03     \\\\ NGC~5898 & S02/3(0)         & E0           & 220  & 22.2 & 2267& 0.23  &    7.7     $\\pm$       1.3   &      0.030   $\\pm$      0.004     &    0.10    $\\pm$     0.03      \\\\ NGC~6721 & E1                  & E+:          & 262 & 21.7 &  4416 & &    5.0     $\\pm$       0.8   &      0.040   $\\pm$      0.007     &    0.24    $\\pm$     0.02       \\\\ NGC~6758 & E2 (merger)   & E+:  & 242 & 20.3 & 3404 &  &   16.0     $\\pm$       2.5   &      0.016   $\\pm$      0.002     &    0.32    $\\pm$     0.05      \\\\ NGC~6776 & E1 pec           & E+pec & 242 & 17.7 & 5480 & &  2.7     $\\pm$       0.5   &      0.033   $\\pm$      0.010     &    0.21    $\\pm$     0.05   \\\\ NGC~6868 & E3/S02/3(3)   & E2           & 277 & 33.7 & 2854 & 0.47  &    9.2     $\\pm$       1.8   &      0.033   $\\pm$      0.006     &    0.19    $\\pm$     0.03      \\\\ & & & & \\\\ NGC~6875 & S0/a(merger)    & SAB(s)0- pec: & & 11.7  &  3121 &  & & &    \\\\ NGC~6876 & E3                     & E3          & 230 & 43.0 & 3836 &  &    9.8     $\\pm$       1.6   &      0.023   $\\pm$      0.003     &    0.26    $\\pm$     0.03        \\\\ NGC~6958 & R?S01(3)          & E+         & 223 & 19.8 & 2652 & 0.12  &    3.0     $\\pm$       0.3   &      0.038   $\\pm$      0.006     &    0.20    $\\pm$     0.03      \\\\ NGC~7007 & S02/3/a             & SA0-:      & 125 & 15.4 & 2954 & 0.14   &    3.4     $\\pm$       0.6   &      0.031   $\\pm$      0.010     &    0.15    $\\pm$     0.05       \\\\ NGC~7079  & SBa                  & SB(s)0     & 155 & 19.8 & 2670 & 0.19  &    6.7     $\\pm$       1.1   &      0.016   $\\pm$      0.003     &    0.21    $\\pm$     0.05       \\\\ NGC~7097 & E4                     & E5          & 224 & 18.1 & 2404 & 0.26   &   10.5     $\\pm$       2.4   &      0.024   $\\pm$      0.005     &    0.30    $\\pm$     0.05     \\\\ NGC~7135 & S01 pec            & SA0- pec   & 231 & 31.4 & 2718 &  0.32  &    2.2     $\\pm$       0.4   &      0.047   $\\pm$      0.010     &    0.46    $\\pm$     0.04    \\\\ NGC~7192 & S02(0)               & E+:         & 257 & 28.6 & 2904 &  0.28   &    5.7     $\\pm$       2.0   &      0.039   $\\pm$      0.015     &    0.09    $\\pm$     0.05    \\\\ NGC~7332 & S02/3(8)            & S0 pec sp & 136 & 14.7 &  1207 &  0.12 &    3.7     $\\pm$       0.4   &      0.019   $\\pm$      0.002     &    0.10    $\\pm$     0.03       \\\\ NGC~7377 & S02/3/Sa pec    & SA(s)0+  & 145 & 36.9 & 3291 & &    4.8     $\\pm$       0.6   &      0.020   $\\pm$      0.002     &    0.10    $\\pm$     0.03       \\\\ & & & & \\\\ IC~1459    & E4                      & E          & 311 & 34.4 & 1659 & 0.28  &    8.0     $\\pm$       2.2   &      0.042   $\\pm$      0.009     &    0.25    $\\pm$     0.04     \\\\ IC~2006    & E1                      & E          & 122 & 28.6 & 1350 & 0.12   &    8.1     $\\pm$       0.9   &      0.026    $\\pm$     0.003     &    0.12     $\\pm$    0.02    \\\\ IC~3370    & E2 pec               & E2+       & 202 & 38.6 & 2934 & 0.20   &    5.6     $\\pm$       0.9   &      0.022   $\\pm$      0.004     &    0.17    $\\pm$     0.04    \\\\ IC~4296    & E0                      & E          & 340 & 41.4 & 3762  & &    5.2     $\\pm$       1.0   &      0.044   $\\pm$      0.008     &    0.25    $\\pm$     0.02      \\\\ IC~5063   & S03(3)pec/Sa    & SA(s)0+: & 160 & 26.7 & 3402 & & & &  \\\\ \\hline \\end{tabular} }} \\end{table*}  %-------------------------end Table 1-----------------------------------------------------  %macro spettroxpaper/xpaper.sm , command fits2 %----------------------    Figure 1-------------------------------------------------- \\begin{figure*} \\centering %  \\includegraphics[width=17cm]{spettropaper.eps} \\includegraphics[width=17cm]{annibali.fig1.eps} \\caption{Central (r $\\leq$r$_e$/16) spectra for a high emission galaxy (IC~1459) and for a weak emission galaxy (NGC 5813) in the wavelength ranges where the most important emission lines are measured. For each galaxy, we show in the top panels the observed spectrum (black line), the underlying stellar spectrum model (red), and the fit to the observed spectrum obtained by adding emission lines (Gaussian curves) to the underlying stellar continuum (blue). In the bottom panels we show the subtracted spectrum, where the emission lines emerge. Notice that in this figure the subtracted stellar continuum is an SSP of age, Z, and extinction equal to the average of all the solution with ${\\rm \\chi^2 < 2}$. However, this is only for illustration purposes, and it is slightly different from the procedure described in Section~3.1. } \\label{spec} \\end{figure*} %----------------------  end  Figure 1-------------------------------------------------- ",
        "conclusions": "Optical emission lines are detected in 89\\% of our sample. This large fraction is somehow expected, since our sample was selected among a compilation of galaxies showing traces of ISM, and it is thus biased toward the presence of emission lines. This fraction is higher than that derived by Phillips et al.~(\\cite{ph86}) (55\\%-60\\%), Macchetto et al.~(\\cite{mac96}); (72\\% - 85\\%),  Sarzi et al.~(\\cite{sar06}) (75\\%), Yan et al.~(\\cite{yan06}) ( 52\\%), and Serra et al.~(\\cite{serra08}) (60\\%), for other samples of early-type galaxies. If we consider only the galaxies in our sample with EW(H$\\alpha$ $+$ [NII]6584)$>$3 \\AA, the detection fraction drops to  57\\%.   The incidence and the strength of emission in our sample is independent on the elliptical or lenticular morphological classes. These results are in agreement with Phillips et al.~(\\cite{ph86}), while on the other hand Macchetto et al.~(\\cite{mac96}) and Sarzi et al.~(\\cite{sar06})  found a higher incidence of emission in lenticular than in elliptical galaxies. Following the studies of Cappellari et al.~(\\cite{capp07}) and Emsellem et al.~(\\cite{Emsellem07}), we attempted a classification between fast and slow rotators in our sample. Fast and slow rotators account for $\\sim$ 70\\%  and  $\\sim$ 30\\% of the sample, respectively. This is interesting given the fact that the sample is composed of $\\sim$ 70\\% ellipticals and  $\\sim$ 30\\% lenticulars. Emission does not depend on the fast and slow rotator classification in our sample.  The EW of the lines tends to decrease from the center outwards. The [NII] line presents a steeper decrease than the H$\\alpha$. However, given the uncertainties, we consider significant only the decrease from the center to  $\\sim$ 0.1 r$_e$. Deeper observations are needed to trace with higher confidence the EW out to larger radii. Some galaxies deviate from the general trend, either because they have a flatter emission profile (NGC~3489, NGC~7007, NGC~7377, NGC~5898), or because their emission is nuclear concentrated (IC~4296, NGC~4374, NGC~5090). Previous narrowband imaging and integral field spectroscopy studies have revealed extended distribution for the ionized gas (Caldwell~\\cite{cal84}; Goudfrooij et al.~\\cite{gou94b}; Macchetto et al.~\\cite{mac96}; Sarzi et al.~\\cite{sar06,sar09}).   Extinction was derived from the observed H$\\alpha$/H$\\beta$ flux ratio, assuming an intrinsic value of $\\approx$ 3.1 for AGN-like objects (Osterbrock~\\cite{Osterbrock89}). This choice was driven by the fact that our galaxies show AGN-like [NII]/H$\\alpha$ ratios. For the majority of the galaxies, the derived ${\\rm E(B-V)}$ values are lower than 0.3. Some galaxies display however very large reddenings, up to ${\\rm E(B-V)\\sim 1.5}$ or even more. Since the observed continuum is incompatible with such large values, our result suggests that in these galaxies the dust has a patchy distribution. This is consistent with narrowband imaging studies revealing the presence of dust lanes and patches in early-type galaxies (e.g. Goudfrooij et al.~\\cite{gou94b}).  We performed a spectral classification for our sample through the classical [OIII]/H$\\beta$ versus [NII]/H$\\alpha$ diagnostic diagram (BPT diagram). The classification  based on the central r$<$r$_{e}$/16 emission lines is provided in Table ~4. About 72\\% of the galaxies present a  central LINERs emission activity. Seyferts account for $\\sim$ 9\\% of the emission sample. However, only IC~5063 can be classified as a {\\it bona fide} Seyfert  galaxy. The other 4 Seyferts (NGC~3489, NGC~777, NGC~7007, and NGC~6958), have very large errors in [OIII]/H$\\beta$ (because of the faint emission in H$\\beta$), and the classification is less robust. We notice however that NGC~3489 and NGC~777 were classified as Seyferts also by other authors (Ho et al.~\\cite{ho97b}, Sarzi et al.~\\cite{sar06}). 7 galaxies ($\\sim$ 12\\% of the emission line sample) are  ``Composites'' (or Transition objects): NGC~3258, NGC~4552, NGC~5193, NGC~5328, NGC~6721, NGC~6876, and IC~2006. They have [NII]/H$\\alpha$ ratios intermediate between HII regions and LINERs/Seyferts. In our sample, these galaxies display the weakest emission lines. For the remaining 7\\% of the emission line galaxies we can not provide a classification, either because the lines are too weak, or because the spectra are too noisy.    We derive spatial trends in the emission line ratios: moving from the center toward annuli of increasing galacto-centric distance, we detect a clear decrease in the [NII]/H$\\alpha$ ratio. The decrease in [OIII]/H$\\beta$ is less clear. It results that galaxies classified as LINERs from to their nuclear line ratios, shift toward the region of ``Composites'' if one consider the emission from more external regions. This could be interpreted either with a decrease in the nebular metallicity, or with a progressive ``softening'' of the ionizing spectrum.    We have investigated possible relations between the central line emission properties and the host galaxy properties. The main results are:  \\begin{enumerate}  \\item Seyferts have young luminosity weighted ages ($\\lesssim$ 5 Gyr), and are on average younger than LINERs.  LINERs span a wide age range, from a few Gyrs to a Hubble time. ``Composite'' have ages older than 5 Gyr, at variance with the idea that they may originate from a combined contribution of AGN and star formation.   \\item Excluding the Seyferts, the spread in the emission equivalent width ([NII], H$\\alpha$ or [OIII]) increases with the galaxy central velocity dispersion. Low-$\\sigma$ galaxies have weak emission lines, while high-$\\sigma$ galaxies show both weak and strong emission lines. Equivalent widths in [NII]$\\lambda$6584 larger than $\\sim$ 5 \\AA \\ are found only in galaxies with $\\sigma_c >$ 200 km s$^{-1}$.  \\item If we consider only the high emission subsample (H) (see Table~4), where the emission line ratios are more reliable, a positive correlation exists between the [NII]/H$\\alpha$ ratio and  $\\sigma_c$.  \\item Excluding the Seyferts, which have larger [OIII]/H$\\beta$ and younger ages, no relation between the emission line ratios and the galaxy age, metallicity, or $\\alpha$/Fe enhancement is found.   \\end{enumerate}  \\subsection{The AGN-starburst connection in Seyferts}  The first result is in agreement with previous findings that LINERs are older than Seyferts (e.g., Kewley et al.~\\cite{kew06}), and supports the idea that the star formation and the AGN phenomena co-exist (Terlevich et al.~\\cite{terl90}; Cid Fernandes \\& Terlevich~\\cite{cf95}; Heckman et al.~\\cite{heck97}; Gonz{\\'a}lez Delgado et al.~\\cite{gd01}; Kauffmann et al.~\\cite{Kauff03}; Cid Fernandes at al.~\\cite{cf05}; Davies et al.~\\cite{davies07}; Riffel et al.~\\cite{riff09}). The young luminosity weighted ages in the Seyferts of our sample are likely the result of recent star formation episodes superimposed to a several Gyrs old stellar population (see Paper~III). Davies et al.~(\\cite{davies07}) provide evidence for a delay of 50-100 Myr between the onset of star formation and the subsequent fueling of the black hole. Given the short duration of the AGN phenomenon ($\\lesssim$ 100 Myr), this suggests that star formation has occurred a few hundreds Myrs ago in the Seyfert galaxies of our sample. Because of the short duration of the AGN phase, we also expect to observe galaxies with relatively young luminosity-weighted ages but no signatures of current Seyfert activity. This is in agreement with the fact that LINERs span a wide luminosity-weighted age range, from a few Gyrs up to a Hubble time.  \\subsection{The ionization mechanism in LINERs}   Several studies have demonstrated that a minimum level of ionizing photons are produced in any evolved stellar systems. Thus, if cold gas is present, some level of nebular emission is  always expected in early-type galaxies. However, the critical question is: can PAGB stars provide {\\it all} the ionization observed in early-type galaxies? In answering this question, we should first of all distinguish between the central LINERs activity, and the more extended LINERs-like emission. Indeed, through the BPT classification carried out in annuli of increasing galacto-centric distance, we have demonstrated that ETGs have both a central and an extended LINERs-like emission. When comparing the stronger  {\\underline {nuclear}} r$<$r$_e$/16 emission with the results of the photoionization models from Binette et al.~(\\cite{bin94}), we obtain that only 11 out of the 49 galaxies classified as LINERs/Composites in our sample (i.e. in $\\approx$ 22\\% of the LINERs/Comp sample) can be explained with photoionization by old stars alone. Of these 11 galaxies, 6 are classified as LINERs and 5 as Composites.  For the other 78\\% of LINERs/Composites, some mechanism different than photoionization by PAGB stars must be at work. We warn however that this fraction is very uncertain, since there is no precise description yet for the evolution and number of PAGB stars (e.g., Brown et al.~\\cite{br08}).  On the other hand, we can not rule out the importance of hot old stars as a photoionization source at larger radii. Indeed, the emission strength progressively decreases toward larger galacto-centric distances, so that and an increasing number of LINERs become compatible with pure PAGB stars photoionization.  {\\it Summarizing, we can not exclude a scenario in which more than a source of ionizing photons are present, with different roles at different radii. We can think of a transition region from the center, where PAGB stars fail in accounting alone for the ionizing photon budget, to more extended regions, where they may do well the job, in agreement with recent findings (Sarzi et al.~\\cite{sar09}, Eracleous et al.~\\cite{era10}). }   Studies of LINERs based on large sample, such as the SDSS, have been recently presented (Stasi{\\'n}ska et al.~(\\cite{sta08}, Cid Fernandes et al.~(\\cite{cf09}). However, we warn against direct comparisons of our findings with the results from SDSS for at least two reasons: the 3''-wide aperture of the SDSS usually encompasses large Kpc-scale regions where diffuse emission contribute to the nebular fluxes; and our sample is biased toward the presence of emission lines. %Here we just recall that Stasi{\\'n}ska et al.~(\\cite{sta08}) showed that only $\\sim$ 1/4 of the LINERs in the SDSS can be explained with photoionization by old stars. Cid Fernandes et al.~(\\cite{cf09}) showed that this fraction is higher if very weak emission galaxies (S/N in H$\\beta$ $<$ 3) are also included.}   Once ascertained the limited role of PAGB stars in the nuclear LINERs region, we investigated other possible scenarios, more specifically photoionization by sub-Eddington accretion rate AGN and fast shocks. Direct evidence for the presence of AGNs comes from the detection of compact nuclear X-ray/ radio sources. High resolution X-ray and/or radio data are available for 18 galaxies in our sample, 16 of which are classified as LINERs. Among the 16 LINERs, evidence of an AGN is found in 10. Even if the statistics are small, this result suggests the presence of low accretion rate AGNs in $\\sim$ 62\\% of LINERs. However, the presence of an AGN does not exclude the coexistence of shock heating: indeed, jet driven flows and dissipative accretion disks in active galaxies with strongly sub-Eddington accretion have ben proposed as possible sources of shocks (Dopita \\& Sutherland~(\\cite{dosu95}), Dopita et al.~\\cite{dop97}), and bubbles likely produced by an AGN outburst have been observed in LINERs (e.g. Baldi et al.~\\cite{bald09}).   Both AGN and shock models reproduce the central emission line ratios in our sample. In the dusty AGN models of Groves et al.~(\\cite{gro04}), LINERs require a ionization parameter ${\\rm \\log U < -3}$, significantly lower than for Seyferts. This is consistent with a scenario in which LINERs are AGNs accreting at a much lower rate than Seyferts. LINERs can also be modeled in terms of fast shocks (200-500 km s$^{-1}$)  in a relatively gas poor environment (pre-shock densities $n \\approx 0.01 -100$ cm$^{-3}$, depending on the magnetic field), with gas metallicity $\\sim$ solar. These low pre-shock densities are consistent with the observed [SII]6717/[SII]6731 ratios. A critical point concerns the high shock velocities required to explain the emission line spectrum in LINERs. Jet driven flows or accretion into  a massive BH can well do the job (Nagar et al.~\\cite{nag05}, Dopita et al.~\\cite{dop97}). We suppose that an alternative source of mechanical energy comes from turbulent motion of gas clouds within the potential well of the galaxy. This scenario can not be excluded since several studies revealed a gaseous component of high velocity dispersion, in the range 150- 250 Km/s, in the center of early-type galaxies (e.g. Bertola et al.~\\cite{ber95}; Zeilinger et al.~\\cite{zei96}, Sarzi et al.~\\cite{sar06}), which translates into higher intrinsic 3D velocity dispersions. This scenario is very appealing since it would naturally explain the observed trends between the emission strength, and the [NII]/H$\\alpha$ ratio, with the galaxy stellar central velocity dispersion.   A solid result from our study is that the [NII]/H$\\alpha$ ratio decreases with increasing galacto-centric distance. This trend may reflect metallicity gradients in the gas, similarly to what observed for the stars. However, this is difficult to fit in a scenario in which the gas has an external origin (see Section~7.3). If not due to metallicity variations, the [NII]/H$\\alpha$ gradient should reflect variations in the properties of the ionization mechanism. In the shock scenario, this requires decreasing shock velocities  with increasing galacto-centric distances. Unfortunately, we lack the spectral resolution necessary to determine the gas velocity dispersion and, ultimately, to ascertain the presence of a radial gradient in the shock velocity.  A large fraction (40/65) of our sample was observed with {\\it Spitzer}-IRS. In a forthcoming paper, we will use MIR ionic and molecular emission lines, and Polycyclic Aromatic Hydrocarbon features, to further investigate the powering mechanism in LINERs, and better disentangle the contributions from AGN photoionization and shock heating.    \\subsection{Origin of the gas}  We derived the oxygen abundance using two calibrations present in the literature: the Kobulnicky et al.~(\\cite{Kobulnicky99}) calibration, valid for HII regions, and the Storchi-Bergmann et al.~(\\cite{Storchi98}) calibration, derived in the assumption of photoionization by a typical AGN continuum (Mathews \\& Ferland~\\cite{mat87}). The first approach was also adopted by  Athey \\& Bregman~(\\cite{ab09}) to derive nebular metallicities in  7 early-type galaxies, and is probably reasonable in the case of photoionization by PAGB stars. The second calibration is valid in the case of AGN excitation, and may be reasonable if LINERs are low-accretion rate AGNs. Our result is the following: irrespective of the adopted calibration, the nebular oxygen metallicities are significantly lower than the stellar ones. The main caveat is that power-law ionizing spectra tend to provide abundances up to $\\sim$0.5 dex larger compared to a typical AGN continuum. A calibration in the case of shock heating is not available in the literature. However, we can see that shock models indicate nebular metallicities not higher than solar in our sample, in agreement with the results obtained with the Kobulnicky et al.~(\\cite{Kobulnicky99}) and  Storchi-Bergmann et al.~(\\cite{Storchi98}) calibrations.  Our result has an important implication for the origin of the gas, suggesting that it may have at least in part an external origin (e.g., from a cooling flow or from accretion from a companion galaxy). This is in agreement with photometric and kinematical studies  in ETGs showing in many cases gas/star misalignment and gas/star angular momentum decoupling (Bertola et al.~\\cite{ber92}; van Dokkum \\& Franx~\\cite{vdf95}; Caon et al.~\\cite{caon00}; Sarzi et al.~\\cite{sar06}). Indeed, kinematical and morphological peculiarities  are present in 50\\% of our sample, suggesting that accretion is frequent. Accretion of fresh gas can feed the central supermassive BH, and activate the AGN emission. This is consistent with the fact that a major fraction of our galaxies studied in the X-ray and/or radio show nuclear compact sources.   Finally, we can not exclude that the low oxygen abundance is unrelated to the gas origin, and reflects the fact that oxygen does not vary on lockstep with the other $\\alpha$-elements. In a survey of 27 Milky Way bulge giants, Fulbright, McWilliam, \\& Rich~(\\cite{fulbr06}) found that, although Mg appears to be enhanced at all [Fe/H], [O/Fe] declines with increasing [Fe/H] and is solar or mildly sub-solar at [Fe/H] $\\ge 0$. Also, X -ray studies have derived subsolar [O/Fe] values, but supersolar [Mg/Fe] values, in the hot gas of elliptical galaxies (e.g., Humphrey \\& Buote~\\cite{hb06}; Ji et al.~\\cite{ji09}). This may be the evidence of an overestimate of O yield by SNII models, which do not consider  significant mass loss at the late stage of massive  progenitor stars (Ji et al.~\\cite{ji09})."
    },
    "1004/1004.1412_arXiv.txt": {
        "abstract": "We present the result of a study of the X-ray emission from the Galactic Centre (GC) Molecular Clouds (MC), within 15 arcmin from Sgr A*. We use \\xmm\\ data spanning about 8 years. We observe an apparent super-luminal motion of a light front illuminating a MC. This might be due to a source outside the MC (such as Sgr A* or a bright and long outburst of a X-ray binary), while it can not be due to low energy cosmic rays or a source located inside the cloud. We also observe a decrease of the X-ray emission from G0.11-0.11, behaviour similar to the one of Sgr B2. The line intensities, clouds dimensions, columns densities and positions with respect to Sgr A*, are consistent with being produced by the same Sgr A* flare. The required high luminosity (about 1.5$\\times10^{39}$ erg s$^{-1}$) can hardly be produced by a binary system, while it is in agreement with a flare of Sgr A* fading about 100 years ago. ",
        "introduction": "In the central few hundreds parsecs of the Milky Way a high concentration of MC is present (Morris \\& Serabyn 1996; Bally et al. 1987; Tsuboi et al. 1999). Sunyaev et al. (1993; 1998) first realised that these clumps of material can behave like mirrors of past bright X-ray events occurring in the GC. Thus, several authors have studied the X-ray bright MC (Koyama et al. 1996; Murakami et al. 2001) to constrain the past activity of the GC and, in particular, of Sgr A*, the counterpart of the supermassive black hole at the GC (Sch\\\"odel et al. 2002; Gillessen et al. 2009). On the other hand, cosmic ray irradiation can explain the high X-ray emission from these MC equally well (Valinia et al. 2000; Yusef-Zadeh et al. 2002; 2007; Dogiel et al. 2009; Bykov 2003). Here we study the MC emission during the 8 years \\xmm\\ monitoring of the 15 arcmin around Sgr A*. ",
        "conclusions": ""
    },
    "1004/1004.4692_arXiv.txt": {
        "abstract": "\\noindent If dark matter in the galactic halo is composed of bosons that form a Bose-Einstein condensate then it is likely that the rotation of the halo will lead to the nucleation of vortices.  After a review of the Gross-Pitaevskii equation, the Thomas-Fermi approximation and vortices in general, we consider vortices in detail.  We find strong bounds for the boson mass, interaction strength, the shape and quantity of vortices in the halo, the critical rotational velocity for the nucleation of vortices and, in the Thomas-Fermi regime, an exact solution for the mass density of a single, axisymmetric vortex. ",
        "introduction": "The observational evidence for the existence of non-baryonic dark matter is impressive.  Precision measurements of anisotropies in the cosmic microwave background reveal a universe composed of 4\\% baryons, 22\\% non-baryonic dark matter and 74\\% dark energy \\cite{wmap}.  Observation of galaxies and galaxy clusters from rotation curves \\cite{rubin}, gravitational lensing \\cite{refregier} and X-ray spectra \\cite{lewis} has lead to a picture in which galaxies are composed of a luminious galactic disk surrounded by a spherical galactic halo of dark matter which comprises roughly 95\\% of the total mass of the galaxy \\cite{freese}.  Various explanations for dark matter have been proposed.  Since we are interested in a particle explanation we will mention the two most prominent proposals from particle physics: WIMPs (weakly interacting massive particles) \\cite{feng} and axions \\cite{pqww}.  WIMPs have a mass around a TeV and interact via the weak force.  Weak interactions automatically lead to the correct relic density \\cite{feng}.  A well known WIMP candidate comes from supersymmetry and is known as the lightest supersymmetric partner (LSP).  In many scenarios this is the lightest neutralino, which is a linear combination of the fermionic superpartners of the Higgs scalars and neutral electroweak gauge bosons.  Axions are scalar bosons first proposed to solve the strong CP problem in QCD and may be produced, for example, by an initial misalignment of its vacuum angle in the early universe \\cite{pww}.  An intriguing possibility is that if dark matter is composed of bosons (which includes, but is not limited to, the axion \\cite{sikyan}) then at sufficiently low temperatures the bosons will form a Bose-Einstein condensate (BEC).  Galactic halos could then be gigantic BECs.  Such an idea, that dark matter may be in the form of a BEC, has been proposed by a number of authors \\cite{sin,leekoh,hubak,silmal,wang,mfs} and has been shown to alleviate known problems with the standard cold dark matter picture, such as cuspy central densities and an abundance of small scale structure in halos  \\cite{hubak}.  The standard approach to treating a non-relativistic BEC is to use the Gross-Pitaevskii equation, also known as a nonlinear Schr\\\"ondinger equation, which includes a self-interaction term and a confining, trap potential.  For gravitating BECs the trap potential is simply the Newtonian gravitational potential, so that the Gross-Pitaevskii equation is coupled with the Poisson equation \\cite{wang, BH, brook}.  A rotating BEC is known to give rise to a lattice of quantized vortices \\cite{expBEC, Fetter}, without which the condensate would be irrotational.  The formation of vortices leads to lower energy states and thus is energetically preferred.  If we are imagining the galactic halo to be made up of bosonic dark matter in the form of a BEC the rotation of the halo may give rise to the formation of vortices \\cite{silmal, yumor, brook, shapiro}.  If vortices form in the halo, they may affect the mass distribution and other properties.  In this paper we consider vortices in BEC dark matter in detail. In the next section we review the Gross-Pitaevskii equation, the Thomas-Fermi approximation and vortices.  These equations will be the framework in which all of our work is done.  In section \\ref{massbound} we consider the condensate as a whole, populated by a lattice of vortices, and make a model independent argument that for a stable condensate with vortices the boson should have a mass centered around $m\\sim10^{-60}-10^{-58} \\text{ kg} \\sim 10^{-24}-10^{-22} \\text{ eV}$.  In section \\ref{1vortex} we analyze a single, axisymmetric vortex and, in the Thomas-Fermi regime, find an exact solution for the mass density.  Our results lead to a comprehensive picture for both the condensate and vortices, which we present in section \\ref{vorstruc}.  Using our vortex solution, in section \\ref{criticalvelocity} we compute the critical rotational velocity of the halo such that vortices would form and find that it is likely to have happened.  We conclude in section \\ref{conclusion}.  Although much of this paper is theoretical, it is important that we connect our results to data.  When doing so we will use the following numbers from M31 (the Andromeda galaxy) \\cite{M31}, which we take to be not unreasonably characteristic of dark matter halos: A total mass of $M\\sim 10^{12} M_\\odot \\sim 10^{42}$ kg, an average radius of roughly $R \\sim 100 \\text{ kpc} \\sim 10^{21}$ m and an approximate rotational velocity of 125 km/s at 30 kpc.  From this we estimate an average mass density of $\\rho_\\text{avg}\\sim 10^{-23}$ kg/m$^3$ and a rotational angular velocity for the halo of $\\Omega \\sim 10^{-16}$ rad/s. ",
        "conclusions": "\\label{conclusion} We considered the possibility that dark matter in a galactic halo is composed of bosons in the form of a Bose-Einstein condensate.  Rotation of the halo is likely to nucleate vortices and we studied such vortices in detail.  We found, in the Thomas-Fermi regime, an exact solution for the mass density of a single, axisymmetric vortex and, by fitting our equations to astrophysical data from the Andromeda galaxy, developed a comprehensive picture for the condensate and vortices and strong bounds for the parameters.  If dark matter in the galactic halo is composed of bosons in the form of a BEC with vortices then model independent arguments require a boson mass centered around $m \\sim 10^{-60} - 10^{-58} \\text{ kg} \\sim 10^{-24} - 10^{-22} \\text{ eV}$ and a radius around $10^{20} - 10^{21} \\text{ m}$, which corresponds to roughly $1-100$ vortices in the halo.  Combining this with our solution, the scatting length is centered around $a \\sim 10^{-80} - 10^{-76} \\text{ m}$.  Using our solution we also determined the size of the vortex core and calculated the critical rotational velocity for the nucleation of vortices, showing that there exists large portions of parameter space where vortices would form.  While these results are specific to the Andromeda galaxy, since the Andromeda galaxy is not atypical we expect them to hold more generally."
    },
    "1004/1004.3767_arXiv.txt": {
        "abstract": "{} {We present a catalog of sources of very high energy ($E>100$~GeV) \\gr s detected by \\textit{ Fermi} telescope at Galactic latitudes $|b|\\ge 10^\\circ$.} {We cross correlate the directions of individual photons with energies above 100~GeV detected by \\textit{ Fermi} with the catalog of sources detected at lower energies. We find significant correlation between the arrival directions of the highest energy photons and positions of \\textit{ Fermi} sources, with the possibility of chance coincidences at the level of  $10^{-38}$. We present a list of \\textit{ Fermi} sources contributing to the correlation signal.  A similar analysis is done for cross-correlation of the catalog of BL Lac objects with the highest energy photons detected by {\\it Fermi}. } {We produce a catalog of high Galactic latitude \\textit{ Fermi} sources visible above 100~GeV. The catalog is split onto two parts. First part contains a list of 46 higher significance sources among which there can be 2 or 3 possible false detections.  Second part of the catalog contains a list of 21 lower significance sources, among which 5 or 6 are possibly false detections. Finally we identify 7 additional sources from the cross-correlation analysis with the BL Lac catalog. The reported sources of $E>100$~GeV \\gr s span a broad range of redshifts, up to $z\\sim 1$. Most of the sources are BL Lac type objects. Only 16 out of 74 objects in our list were previously reported as VHE \\gr\\ sources.} {} ",
        "introduction": "Ground-based Cherenkov \\gr\\ telescopes HESS, MAGIC and VERITAS have discovered a population of sources of Very High Energy (VHE) ($E\\ge 100$~GeV) \\gr s. Except for the sources discovered in the Galactic Plane survey by HESS \\citep{HESS_survey_science,HESS_survey}, most of the sources were discovered in dedicated pointed observations.  Surveys of large regions of the VHE \\gr\\ sky with the existing  Cherenkov telescopes are difficult because of the too narrow size of the field of view. Wide field of view ground-based \\gr\\ detectors MILAGRO \\citep{milagro} and Tibet \\citep{tibet} arrays have produced a systematic survey of the VHE \\gr\\ sky. The energy threshold of the air shower arrays like MILAGRO and TIBET is  rather high (in the multi-TeV band) so that only sources with spectra extending well above 1 TeV could be detected.  All-sky monitoring of \\gr\\ sources at the energies $E\\sim 100$~GeV became possible with the start of operation of {\\it Fermi} telescope. Space-based  {\\it Fermi} telescope has a much smaller effective collection area, compared to the ground based \\gr\\ telescopes ($\\sim 1$~m$^2$, compared to $\\sim 10^5$~m$^2$ for the ground-based \\gr\\ telescopes). At the same time, at the energies above 100~GeV {\\it Fermi} signal is almost background-free (contrary to the ground based telescopes in which the signal has to be identified on top of strong background created by cosmic rays and optical/UV night sky background).  \\citet{paper1}  have searched for  the point sources  of $E\\ge 100$~GeV \\gr s  using   {\\it Fermi} data and found 8 significant excesses with at least 3 photons within $0.1^\\circ$ circle corresponding to the 68\\% containment radius of point spread function (PSF) of the Large Area Telescope (LAT) on board of {\\it Fermi}. Seven excesses were associated with known VHE \\gr\\ sources, while the remaining one was identified with head-tail radio galaxy IC 310. Detection of IC 310 in the VHE band was later confirmed by MAGIC telescope \\citep{IC310_MAGIC}. Due to the moderate collection area of {\\it Fermi}, only the brightest VHE \\gr\\ sources were detected individually in the {\\it Fermi} VHE \\gr\\ sky survey. All other known VHE \\gr\\ sources at Galactic latitudes $|b|\\ge 10^\\circ$ gave $\\le 2$ photons within the LAT PSF  circle and could not be found  from the analysis of the data above $100$~GeV alone.  A complementary method to identify the sources of $E\\ge 100$~GeV photons detected by {\\it Fermi} is to use a prior knowledge of source positions on the sky and to verify which of the already known sources could have produced the highest energy  \\gr s detected by the LAT telescope. In other words, sources of $E\\ge 100$~GeV \\gr s could be identified also via cross-correlation of arrival directions of the $E\\ge 100$~GeV \\gr s with the source positions on the sky. Such an approach to the identification of the sources was previously applied to the analysis of EGRET data above $10$~GeV by \\citet{dingus01} and \\citet{10GeV_EGRET}.  In what follows we perform the cross-correlation analysis of the arrival directions of \\gr s with energies above 100~GeV detected by {\\it Fermi} at Galactic latitudes $|b|\\ge 10^\\circ$ with  the 1-st year  {\\it Fermi}  source catalog \\citep{fermi_catalog}. Our analysis results in a catalog of 46 high Galactic latitude sources of $E\\ge 100$~GeV \\gr s. Although each individual source in the catalog is detected with significance around $3\\sigma$ in this energy band, overall significance of detection of the entire source set is very high. Simple analysis shows that most of the sources contributing to the correlation signal are real VHE \\gr\\ sources, only 2 or 3 of them are expected to be false detections. Taking into account the fact that most of the sources in the list are BL Lac type objects, we extend our cross-correlation analysis to the catalog of BL Lacs \\citep{veron13} and find 7 more sources which correlate with the arrival directions of $E\\ge 100$~GeV \\gr s and are not listed in the 1-st year {\\it Fermi} catalog. Finally, we list for completeness sources from the 1-st year {\\it Fermi} catalog for which the $E\\ge 100$~GeV \\gr s are found within the 95\\%, rather than 68\\% containment circle of LAT PSF. There are 21 such sources, 5 or 6 of them are expected to be false detections due to the chance coincidence of the arrival direction of the VHE \\gr\\ with the source position.  The plan of the paper is like follows. In Section \\ref{sec:data} we discuss data selection and data analysis methods. In Section \\ref{sec:correlation} we present the results of the correlation analysis of the arrival directions of $E\\ge 100$~GeV \\gr s with the sources of the 1-st year {\\it Fermi} catalog. In Section \\ref{sec:catalog} we give the list of sources contributing to the correlation signal. In Section \\ref{sec:BL} we perform the correlation analysis with the BL Lac catalog of \\citet{veron13} and give the list of additional BL Lacs correlating with the highest energy {\\it Fermi} photons, but not present in the 1-st year {\\it Fermi} catalog. In Section \\ref{sec:individual} we comment upon individual sources. Finally, in Section \\ref{sec:discussion} we discuss the results. ",
        "conclusions": ""
    },
    "1004/1004.3551_arXiv.txt": {
        "abstract": "We investigate particle production near extra species loci (ESL) in a higher dimensional field space and derive a  speed limit in moduli space at weak coupling.  This terminal velocity is set by the characteristic ESL-separation and the coupling of the extra degrees of freedom to the moduli, but it is independent of the moduli's potential if the dimensionality of the field space is considerably larger than the dimensionality of the loci, $D\\gg d$. Once the terminal velocity is approached, particles are produced at a plethora of nearby ESLs, preventing a further increase in speed via their backreaction. It is possible to drive inflation at the terminal velocity, providing a generalization of trapped inflation with attractive features: we find that more than sixty e-folds of inflation for sub-Planckian excursions in field space are possible if ESLs are ubiquitous, without fine tuning of initial conditions and less tuned potentials. We construct a simple, observationally viable model with a slightly red scalar power-spectrum and suppressed gravitational waves; we comment on the presence of additional observational signatures originating from IR-cascading and individual massive particles. We also show that moduli-trapping at an ESL is suppressed for $D\\gg d$, hindering dynamical selection of high-symmetry vacua on the landscape based on this mechanism. ",
        "introduction": "The presence of many light fields, or moduli, is a common feature of string theory. At late times, the expectation values of these fields determine low energy observables; hence, their evolution is heavily constraint from the time of nucleosynthesis, although they are expected to be dynamical in the very early universe. For these reasons, moduli trapping is an important aspect of inflationary model building in string theory, see \\cite{HenryTye:2006uv,Cline:2006hu,Burgess:2007pz,McAllister:2007bg,Baumann:2009ni} for reviews; often all but one degree of freedom are stabilized by construction during inflation, but how they stabilize is seldom addressed.  A possible dynamical stabilization process is the String Higgs effect \\cite{Bagger:1997dv,Watson:2004aq}, which singles out certain locations in moduli space. At these loci additional degrees of freedom become light, often due to the presence of an enhanced symmetry, and are produced \\cite{Kofman:2004yc,Watson:2004aq} if the moduli approach the location. The process of particle production is identical to the one examined in preheating \\cite{Traschen:1990sw,Kofman:1997yn}, see \\cite{Bassett:2005xm,Kofman:2008zz} for reviews. Backreaction of these new states on the moduli can be dramatic: moduli can be trapped \\cite{Kofman:2004yc,Watson:2004aq,Patil:2004zp,Patil:2005fi,Cremonini:2006sx}, or, if moduli drive inflation, the inflaton velocity can decrease temporarily \\cite{Chung:1999ve}, as in trapped inflation \\cite{Kofman:2004yc,Green:2009ds,Silverstein:2008sg}. A concrete realization of the string Higgs effect can be found in a system of parallel $D$-branes \\cite{Witten:1995im}, whose separation is a modulus field from the four-dimensional point of view. If they come close to each other, strings stretching between the branes become massless, gauge symmetries are enhanced and the branes stick to each other due to backreaction -- the modulus is trapped. For a small selection of other examples, see \\cite{Seiberg:1994rs,Seiberg:1994aj,Intriligator:1995au,Strominger:1995cz,Witten:1995ex,Katz:1996ht,Bershadsky:1996nh,Witten:1995gx}.  In this paper, we examine the consequences of particle production near such extra species loci (ESL) if the dimensionality of moduli space $D$ is large compared to the dimensionality of the locus, $d$. $D\\gg 1$ is natural in string theory, leading to the notion of a landscape \\cite{Susskind:2003kw}, yet the effect of ESLs has primarily been investigated for $D-d=1,2$ \\cite{Kofman:2004yc,Watson:2004aq,Greene:2007sa}. Based on geometric arguments, we show in Sec.~\\ref{sec:particleproduction} that trapping is suppressed if the dimensionality of moduli space is larger than the dimensionality of a given locus ($D>d+1$). This result is expected, since there is no classical attraction towards ESLs and it is improbable to run head on into an ESL if $D\\gg d$. This means  that the mere presence of ESLs does not guarantee a dynamical preference of high symmetry states for moduli \\cite{Kofman:2004yc,Dine:2000ds,Dine:1998qr}. However, in Sec.~\\ref{sec:particpleproductionwithimpactparameter} we show that the presence of many loci with a characteristic inter-ESL distance $x$ leads to a general speed limit, or terminal velocity, on moduli-space. At strong coupling, speed limits are known \\cite{Silverstein:2003hf}, leading to DBI-inflation \\cite{Alishahiha:2004eh}. Here, we derive a speed limit at weak coupling caused by the combined backreaction of particles produced at many ESLs in the vicinity of the trajectory.  We take a bottom up approach, assuming the viability of low-energy effective field theory and treating the characteristic distance between ESLs, as well as their dimensionality, as free parameters. Furthermore, we model the additional light degrees of freedom by a massless scalar field that couples to the moduli via interactions of the type $g^2\\chi^2\\varphi^2$, as in \\cite{Kofman:2004yc}. In this notation, the speed limit takes the simple form $|\\dot{\\vec{\\varphi}}|<gx^2\\equiv v_t$, given that the classical trajectory is reasonably straight (we allow for a classical potential for the moduli)and $D$ is large. The presence of a terminal velocity offers the opportunity to drive inflation by moduli with hitherto unsuitable potentials.  This type of inflation is a generalization of trapped inflation \\cite{Kofman:2004yc,Green:2009ds}, which recently re-surfaced as monodromy inflation \\cite{Silverstein:2008sg}, but with several crucial differences: firstly, particle production at any given ESL is minor because ESLs are not approached too closely, in contrast to the one-dimensional case in \\cite{Green:2009ds}; there, extra species points (ESPs) are encountered head on and all particle production occurs at a single ESP at any given time. Secondly, in the large $D$ limit the velocity at which the trajectory is traversed is independent of the slope of the potential as long as the slope is steep enough. Again, this differs from \\cite{Green:2009ds}, where the slope still determines the speed and observational parameters, such as the scalar spectral index.  We investigate this type of trapped inflation in higher dimensions  in Sec.~\\ref{sec:trappedinfllargeD}, where we encounter several attractive features: focusing on a model with quadratic potentials for the moduli, more than sixty e-folds of inflation can be achieved with sub-Planckian field excursions if $v_t$ is small enough. The inflationary scale is set by the COBE normalization, but the dynamics is not sensitive to the actual shape of the potential as long as a single, reasonably straight classical trajectory with large slope is present. This is, in a sense, opposite to setups needed for slow-roll inflation, where the potential needs to be shallow over long ranges in field space. As a consequence, the $\\eta$-problem is alleviated. Furthermore, no fine-tuning of the initial speed is needed: if the initial speed is large, particle production is strong, causing strong backreaction and temporary trapping; after the produced particles are diluted away due to the expansion of the universe, the speed picks up again and the terminal velocity is approached from below. In this sense, trapped inflation is an attractor solution.  The ongoing particle production during inflation has several additional observational consequences, Sec.~\\ref{sec:IR-cascading} and \\ref{sec:massive}. Since $\\chi$-particles scatter off the inflaton condensate, they cause IR-cascading and additional contributions to the power-spectrum \\cite{Barnaby:2009mc,Barnaby:2009dd}; furthermore, $\\chi$-particles dilute and become massive once the trajectory moves away from their ESL of production. Thus, one needs to consider the effect of individual massive particles that are coupled to the moduli, which could lead to additional circular cold-spots in the CMBR \\cite{Itzhaki:2008ih,Fialkov:2009xm,Kovetz:2010kv}. We plan to investigate both effects, as well as the generation of non-Gaussianities, in a future publication, accompanied by a more rigorous string theoretical implementation of trapped inflation in higher dimensions.  Readers familiar with the notion of particle production, as discussed in \\cite{Kofman:2004yc}, and a primary interest in the terminal velocity and consequences for inflation may skip Sec.~\\ref{sec:particleproduction} and go directly to the main part of this work, Sec.~\\ref{sec:moduliareinflatons}. On the other hand, readers with a primary interest on a dynamical selection principle on the landscape will find Sec.~\\ref{sec:particleproduction} and Appendix \\ref{A:trappingrw} useful (these sections review \\cite{Kofman:2004yc} and contain minor extensions and consequence which are implied to in \\cite{Kofman:2004yc} but not fully spelled out). ",
        "conclusions": "Whether or not moduli get trapped at extra species loci (ESL), which are often associated with additional symmetries, has important consequences for selecting vacua on the landscape. If the trajectory in field space comes close to such a locus, particles are produced and backreaction can trap the trajectory. In this paper we gave a simple geometric argument to highlight that this type of moduli trapping is unlikely if the dimensionality of moduli space is considerably larger than the dimensionality of the loci; for $D>d+1$ (or $D>d+2$ if the trajectory is intertwined), the probability of getting trapped is suppressed by $(\\mu/x)^{D-d-1}$ (or $(\\mu/x)^{D-d-2}$ if the trajectory is intertwined), where $x$ the characteristic inter-ESL distance, $\\mu=\\sqrt{v/g}<x$ is an effective trapping radius, $v$ the velocity in field space, and $g$ a coupling constant. We concluded that the sole presence of ESLs does not guarantee a dynamical preference of high symmetry vacua on the landscape, but since the effective trapping radius increases with the velocity, ESLs are always crucial at large speeds, for instance if moduli are responsible for inflation.  For $D\\gg 1$ and $d=0$ (extra species points or ESPs, a generalization to ESLs is straightforward), we derived a terminal velocity $v_t=g^2x$ on the landscape: once the speed approaches $v_t$, the trajectory comes within reach of a plethora of ESPs at which $\\chi$-particles are produced. Their combined backreaction acts opposite to the velocity if the classical trajectory is straight over distances covered in a few Hubble times (we allow for a classical potential of the moduli). Since backreaction increases sharply for $v>v_t$ if the dimensionality of moduli space is large, $v_t$ acts as a speed limit. The classical trajectory is followed with this speed as long as the slope along the classical potential is large and the path is reasonably straight.  The presence of a terminal velocity offers the intriguing possibility to drive inflation with potentials that do not otherwise support slow-roll inflation. This type of inflation is a generalization of trapped inflation. The main difference is the independence of the speed from the slope of the potential, as long as $D$ is large enough.  We then provided a phenomenological model of trapped inflation in higher dimensions, assuming a simple quadratic potential and ubiquitous ESPs ($x\\lesssim 10^{-3.5}$, $g\\sim 0.1$). For these parameters, more than sixty e-folds for inflation result for sub-Planckian excursions in field space and observations are within observational bounds (the power-spectrum of curvature perturbations is slightly red and gravitational waves are suppressed). We would like to highlight that no fine tuning of initial conditions is needed and constraints on the potential are lessened (the COBE normalization still sets the inflationary scale, but the $\\eta$-problem is alleviated). Additional observational signatures are expected, such as an extra contribution to the power-spectrum from IR-cascading (less than a few percent if $g<0.2$), potentially large non-Gaussianities and additional circular cold-spots in the CMBR caused by individual $\\chi$-particles that became heavy before being diluted away. We plan to come back to these interesting signatures in a future study.  In this paper, we took a bottom-up approach and parametrized the undoubtedly rich structure of the landscape and the distribution of ESLs by a few characteristic parameters, such as $x$, that we left undetermined. Concrete implementations within string theory are needed to see whether or not the distribution of ESLs and the potential on the landscape supports trapped inflation in higher dimensions. It appears that ubiquitous ESLs and steep potentials are common, but we leave an in depth analysis to future studies."
    },
    "1004/1004.3545_arXiv.txt": {
        "abstract": "We make a detailed investigation of the properties of Lyman-break galaxies (LBGs) in the $\\Lambda$CDM model. We present predictions for two published variants of the \\GALFORM\\ semi-analytical model: the \\citet{Baugh05} model, which has star formation at high redshifts dominated by merger-driven starbursts with a top-heavy IMF, and the \\citet{Bower06} model, which has AGN feedback and a standard Solar neighbourhood IMF throughout. We show predictions for the evolution of the rest-frame far-UV luminosity function in the redshift range $z=3-20$, and compare with the observed luminosity functions of LBGs at $z=3-10$. We find that the \\citeauthor{Baugh05} model is in excellent agreement with these observations, while the \\citeauthor{Bower06} model predicts too many high-luminosity LBGs. Dust extinction, which is predicted self-consistently based on galaxy gas contents, metallicities and sizes, is found to have a large effect on LBG luminosities. We compare predictions for the size evolution of LBGs at different luminosities with observational data for $2\\lsim z \\lsim 7$, and find the \\citeauthor{Baugh05} model to be in good agreement. We present predictions for stellar, halo and gas masses, star formation rates, circular velocities, bulge-to-disk ratios, gas and stellar metallicities and clustering bias, as functions of far-UV luminosity and redshift. We find broad consistency with current observational constraints. We then present predictions for the abundance and angular sizes of LBGs out to very high redshift ($z \\leq 20$), finding that planned deep surveys with \\JWST\\ should detect objects out to $z \\lsim 15$. We predict that the effects of dust extinction on the far-UV luminosity density should be large ($\\sim 2$~mag), even out to high redshifts. The typical UV luminosities of galaxies are predicted to be very low at high redshifts, which has implications for detecting the galaxies responsible for reionizing the IGM; for example, at $z=10$, 50\\% of the ionizing photons are expected to be produced by galaxies fainter than $\\MUV \\sim -15$. ",
        "introduction": "The discovery of Lyman-break galaxies at $z \\sim 3$ by \\citet{Steidel96} was a breakthrough in observational studies of galaxy formation.  It provided, for the first time, a significant sample of normal galaxies at high redshift whose properties and population statistics could then be investigated observationally and compared to the predictions of theoretical models \\citep{Baugh98}.  Lyman-break galaxies (LBGs) are star-forming galaxies which are identified through the Lyman-break feature in their spectra. This feature is produced by absorption by neutral hydrogen in the atmospheres of massive stars, in the interstellar medium (ISM) of the galaxy and in the intergalactic medium (IGM) \\citep{Steidel92}. For ground-based observations, detection of the Lyman break is restricted to redshifts $z \\gsim 3$.  Since the first successful demonstration by \\citet{Steidel96} at $z \\sim 3$, the technique has been extended to identify galaxies at both higher redshifts and lower luminosities, using ground-based telescopes and the Hubble Space Telescope (HST) \\citep[e.g.][]{Madau96,Steidel99,Bouwens03,Shimasaku05,Yoshida06}. Finding LBGs at $z \\gsim 6$ requires observing in the near-IR, which was first done using NICMOS on HST, and led to the identification of a small number of $z \\sim 7-8$ objects \\citep{Bouwens04b}. More recently, the WFC3/IR camera on HST has allowed the discovery of much larger samples of LBGs at $z\\sim 7-8$ \\citep{Bouwens10a,McLure10,Oesch10a,Bunker10,Yan09} and even a few candidates at $z\\sim 10$ \\citep{Bouwens10b}. By observing in the UV from space, the \\GALEX\\ satellite has also been used to find LBGs at $z \\sim 1$ \\citep{Burgarella06}.  Follow-up observational investigations on LBGs have included estimates of their star formation rates (SFRs), sizes, morphologies, stellar and dynamical masses, galactic outflows, metallicities, dust extinctions, gas masses, IR/sub-mm emission from dust and clustering \\citep{Steidel96,Giavalisco96,Adelberger98,Chapman00,Pettini01,Shapley01,Ferguson04}. Since LBGs are selected on the basis of their rest-frame far-UV emission, which is dominated by massive young stars, LBG samples at different redshifts also provide a means to trace the cosmic star formation history, a key component in our picture of galaxy formation \\citep{Madau96}, although important uncertainties remain due to the effects of dust extinction. Observations of LBGs at $z=3-10$ probe galaxy evolution over the first 3--15\\% of the age of the universe.  Other observational techniques have also been used to find normal galaxies at high redshift. Searches for star-forming galaxies via their Ly$\\alpha$ emission line \\citep[e.g.][]{Hu98} cover a similar redshift range to LBGs. The main drawback with this technique is that some star forming-galaxies show Ly$\\alpha$ absorption rather than emission \\citep{Shapley03}, and consequently are missed in narrowband surveys. A further complication is that inferring the SFR from the Ly$\\alpha$ emission line is much more uncertain than inferring it from the far-UV continuum, since the effects of dust extinction are amplified by resonant scattering of Ly$\\alpha$ photons by hydrogen atoms. Another technique is to search for sub-mm or IR emission from dust in high-$z$ star-forming galaxies \\citep{Smail97,Hughes98}. This method is currently limited by source confusion at faint fluxes due to the relatively poor angular resolution of current IR/sub-mm telescopes, which restricts searches to redshifts $z \\lsim 3$ and mostly to the galaxies with the highest SFRs. Other techniques for selecting high-$z$ galaxies, which are sensitive also to non-star-forming galaxies (such as ERO and DRG colour selection) are limited to even lower redshifts $z \\lsim 2$. Overall, the Lyman-break technique still seems the most effective  for finding large samples of star-forming galaxies at  $z \\gsim 3$ that cover a wide range of luminosities and SFRs.    The theoretical significance of the discovery of LBGs was highlighted early on using semi-analytical models of galaxy formation. \\citet{Baugh98} showed that the abundance and observed properties of the $z \\sim 3-4$ LBGs found by \\citet{Steidel96} and \\citet{Madau96} could be explained in the framework of CDM, and that they fitted into a picture in which the cosmic SFR density peaked around a redshift $z \\sim 2$. This evolution of the cosmic SFR density was driven by the combined effects of the build-up of dark matter halos, gas cooling and supernova feedback. Their model had star formation occuring mostly in quiescent disks, and neglected the effects of dust extinction. Subsequent observational studies found evidence from UV continuum slopes for significant dust extinction in LBGs \\citep{Meurer99,Steidel99}. \\citet{Somerville01} proposed a different semi-analytical model in which star formation bursts triggered by galaxy mergers played an important role, and combined this with an empirical prescription for dust extinction, tuned to match observational estimates of the extinction for $z\\sim 3$ LBGs.  More recent studies of LBGs in semi-analytical models include \\citet{Guo09} and \\citet{LoFaro09}, {which investigated the effects on inferred luminosity functions and other properties of applying observational LBG colour selections to model galaxies. The former used a phenomonelogical model for dust extinction, while the latter used a physical model similar to that in the present paper.} LBGs were also studied in gas-dynamical simulations of galaxy formation \\citep{Nagamine02,Weinberg02}, but these simulations had the drawback that they did not predict properties for the present-day galaxy population consistent with observations, unlike the semi-analytical models. Furthermore, none of these models were able to explain the number counts and redshifts of faint sub-mm galaxies discovered in surveys at 850$\\mum$ \\citep{Smail97}, which were subsequently shown to be dusty starbursts at $z\\sim 2-3$ \\citep{Chapman03}.  In order to explain within a single framework the sub-mm and Lyman-break galaxies at high redshift, together with a wide range of galaxy properties at $z=0$ (including optical and near- and far-IR luminosity functions, gas fractions, metallicities and galaxy sizes), \\citet{Baugh05} introduced a new semi-analytical model in which the gas consumption timescale in disks at high redshifts was increased, with the result that starbursts triggered by galaxy mergers played a more significant role at high redshift. They assumed, further, that stars formed in these bursts with a top-heavy initial mass function (IMF). Unlike previous models of high-redshift galaxies, this model included a fully self-consistent treatment of both absorption and emission of radiation by dust, with dust extinction calculated from radiative transfer based on the predicted gas masses, metallicities and sizes of galaxies, and the spectrum of the dust emission calculated by solving for the temperature distribution of the dust grains in each galaxy. In subsequent papers, we have explored other predictions from the same model, including stellar and gas metallicities \\citep{Nagashima05a,Nagashima05b}, galaxy colours, sizes and morphologies in the local universe \\citep{Almeida07,Gonzalez09}, the evolution of Ly$\\alpha$-emitters \\citep{LeDelliou05,LeDelliou06,Orsi08}, and the evolution of galaxies at mid- and far-IR wavelengths \\citep{Lacey08,Lacey10}, finding generally good agreement with observational data.  In \\citet{Baugh05}, we made only a limited comparison with observational data on LBGs, focussing on their rest-frame far-UV luminosity function at $z=3$. Since then, there has been a huge increase in the amount and quality of observational data on LBGs, in particular enabling measurements of the luminosity function of LBGs out to $z\\sim 10$. Therefore in this paper we return to studying LBGs, making detailed predictions for the evolution of their luminosity functions over a wide redshift range ($z = 3-20$) and for many other properties. We consider two variants of the \\GALFORM\\ semi-analytical model \\citep{Cole00}, those of \\citet{Baugh05} and \\citet{Bower06}. The two models differ in a number of significant ways, the most important being that the \\citeauthor{Bower06} model includes AGN feedback, while the \\citeauthor{Baugh05} model has a variable IMF, as already mentioned. We focus here on far-UV-selected galaxies in the redshift range $z\\gsim 3$, where they are observationally detected using their Lyman-break features. We investigate the present-day descendants of LBGs in a companion paper \\citep{Gonzalez10b}, and make predictions for the reionization of the IGM from the same models in \\citet{Raicevic10a,Raicevic10b}. The properties and evolution of far-UV-selected galaxies at lower redshifts will be the topic of a separate paper.  The plan of this paper is as follows: In \\S\\ref{sec:GALFORM}, we briefly review the main features of the two models. In \\S\\ref{sec:lf-evoln}, we compare predictions for the evolution of the far-UV luminosity function with observational data from LBGs, and investigate the sensitivity of these predictions to various model parameters. In \\S\\ref{sec:props}, we investigate the sizes and other physical properties of UV-selected galaxies, and carry out a detailed comparison with the observed sizes of LBGs. In \\S\\ref{sec:high-z}, we present predictions for LBGs at very high redshifts, which may be accessible with future telescopes such as \\JWST\\ and ELTs. In \\S\\ref{sec:UVdens}, we show how our predictions for LBGs fit into the wider picture of the evolution of the cosmic star formation and UV luminosity densities. We briefly consider the contribution of LBGs to the reionization of the IGM. Finally, we present our conclusions in \\S\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have made a detailed investigation of the properties and evolution of Lyman-break galaxies (LBGs) predicted by hierarchical models of galaxy formation. We followed the galaxy formation process in the framework of the $\\Lambda$CDM cosmology using the \\GALFORM\\ semi-analytical model, which includes physical treatments of the hierarchical assembly of dark matter halos, shock-heating and cooling of gas, star formation, feedback from supernova explosions, AGN and photoionization of the IGM, galaxy mergers and chemical enrichment. The luminosities of galaxies are calculated from a stellar population synthesis model, and dust extinction is then included using a self-consistent theoretical model based on the results of radiative transfer calculations. The dust mass is calculated from the predicted mass and metallicity of the cold gas component, and this is combined with the predicted galaxy radius to calculate the dust extinction optical depth. The far-UV dust extinction is a critical component in any model for LBGs.  We have presented predictions for two variants of the \\GALFORM\\ model. In the \\citeauthor{Baugh05} (2005, \\Bau) model, the formation of very massive galaxies is inhibited by supernova-driven superwinds which eject gas from halos, star formation at high redshifts is dominated by starbursts triggered by galaxy mergers, and stars form in these bursts with a top-heavy IMF. The top-heavy IMF was motivated by the need to explain the number counts and redshift distributions of the faint sub-mm galaxies. This model also matches a wide range of other data on local galaxies (such as gas masses and disk sizes). In the \\citeauthor{Bower06} (2006, \\Bow) model, the formation of very massive galaxies is instead inhibited by AGN feedback which heats the gas in halos, starbursts play a much smaller role in star formation, and all stars form with a Solar neighbourhood IMF. The \\Bow\\ model underpredicts the sub-mm number counts by more than an order of magnitude, due to having too few very luminous and dusty star-forming galaxies at high redshifts.  This shortcoming might be remedied by introducing a top-heavy IMF in bursts, as we will explore in a future paper. Both models match the present-day optical and near-IR luminosity functions.  We first considered the evolution of the rest-frame far-UV luminosity function (\\S\\ref{sec:lf-evoln}). Both models predict modest evolution over the redshift range $3 \\lsim z \\lsim 8$, but a more rapid evolution at higher redshifts, $z \\gsim 8$, driven by the build-up of dark matter halos. However, the models differ in that the \\Bow\\ model predicts a more extended high-luminosity tail than the \\Bau\\ model, once dust extinction is included. The effects of dust extinction on the far-UV luminosity function are much larger in the \\Bau\\ model ($\\sim 2$~mag) because the bright end of the luminosity function is dominated by starbursts in which the dust content is enhanced by metal production with the top-heavy IMF. We made a detailed comparison of the predictions of both models with observed far-UV luminosity functions of LBGs over the redshift range $z=3-10$. We found that the \\Bau\\ model, without any modification of its parameters, predicts a far-UV luminosity function in excellent agreement with current observational data over this whole range. On the other hand, the \\Bow\\ model conflicts with the LBG observations at $z=3-7$ because it predicts too many high-luminosity galaxies. We then investigated the effect on the predicted luminosity functions of varying some of the model parameters from their default values. Assuming a Solar neighbourhood, rather than top-heavy, IMF in bursts has only a modest impact on the far-UV luminosity function, because the effects of lower intrinsic stellar luminosities are partly compensated by lower dust extinctions. However, such a model predicts far too few sub-mm galaxies. We find that the luminosity function over the range covered by observational data is fairly sensitive to the assumed star formation timescale in bursts, especially at higher redshifts, but is less sensitive to the strength of supernova feedback, when these parameters are varied over physically reasonable ranges. The inability of the \\Bow\\ model to match the observed LBG luminosity function data appears to be caused mainly by the short assumed star formation timescale in bursts, rather than by the AGN feedback model or the assumed IMF.  We next investigated a wide range of other physical properties of LBGs predicted by the models (\\S\\ref{sec:props}). We first considered the sizes of galaxies in the rest-frame far-UV, and compared to recent observational measurements at $z \\sim 2-7$. We found that both models predicted sizes in reasonable agreement with observations at higher UV luminosities, but only the \\Bau\\ model is consistent with observed sizes at lower luminosities. We then presented predictions of the \\Bau\\ model for stellar, halo and gas masses, clustering bias, circular velocity, burst fractions, bulge-to-disk ratios, star formation rates and gas and stellar metallicities, and made brief comparisons with relevant observational data. The model predictions appear to be broadly compatible with current observational constraints (many of which are rather uncertain) in most cases. A particularly interesting issue is the stellar masses - our predicted values are well below observational estimates based on fitting stellar population models to broad-band photometry. However, the observational estimates all assume a Solar neighbourhood IMF, while in the \\Bau\\ model the LBG population is dominated by starbursts forming stars with a top-heavy IMF. The observationally inferred stellar masses therefore cannot be directly compared with the values from the model. When instead we compare the IR fluxes (which drive the observational stellar mass estimates) directly, the model is much closer to the observations.  We will investigate this important issue in more detail in a future paper. We will also make a more detailed study of LBG clustering in future work, since this provides constraints on the masses of the dark matter halos hosting LBGs.  In \\S\\ref{sec:high-z}, we showed predictions for LBGs at very high redshifts ($z=7-20$) {in the \\Bau\\ model,} including surface densities of objects down to very faint apparent magnitudes ($m_{AB}=32$), relevant for observations with future telescopes such as \\JWST\\ and ELTs. We find that deep surveys planned with \\JWST\\ should be able to detect a few LBGs at $z\\sim 15$ and $m_{AB}\\sim 30-31$, and many more at lower redshifts. LBGs detected at $m_{AB}\\sim 31$ are predicted to have angular radii $\\sim 0.02-0.05$~arcsec, depending only weakly on redshift over this range, and to have circular velocities $\\sim 30-100\\kms$, again only weakly dependent on redshift.  In \\S\\ref{sec:UVdens}, we showed the predicted evolution of the far-UV luminosity density, and its relation to the cosmic SFR history, {again in the \\Bau\\ model}. The unextincted 1500\\AA\\ luminosity density tracks the SFR density in high-mass stars ($m\\gsim 5\\Msol$) more closely than the total SFR density, since the relative contributions of quiescent and burst star formation (with Solar neighbourhood and top-heavy IMFs respectively) change with redshift. The effect of dust extinction on the far-UV luminosity density is predicted to be large, $\\approx 2$~mag at 1500\\AA\\ in the range $3 \\lsim z \\lsim 15$, dropping to $\\approx 1$~mag at $z=0$. Finally, we considered the predicted contribution of galaxies to the emissivity of ionizing photons which can reionize the IGM. For a constant escape fraction of ionizing photons from galaxies, this emissivity falls by a factor $\\sim 100$ from its peak at $z\\sim 5$ to $z=15$. At high redshift, most of the ionizing photons are predicted to come from very low luminosity galaxies, so that, for example, to detect the galaxies responsible for $>50\\%$ of the ionizing emissivity at $z=10$ would require an LBG survey probing fainter than $\\MUV \\sim -15$, {corresponding to an apparent magnitude $m_{AB}\\sim 32$ at the same rest-frame wavelength.} The predictions of our model for reionization of the IGM are discussed in much greater detail in \\citet{Raicevic10a,Raicevic10b}.  In conclusion, we find that the \\citet{Baugh05} model, which was originally constructed to match the far-UV luminosity function of LBGs only at $z=3$, predicts results in remarkably good agreement with subsequent observations of LBGs out to $z=10$. Further exploration of whether this model provides a physically accurate description of LBGs and other high-redshift galaxy populations will require more detailed comparisons between the model predictions and observational data (e.g. for stellar masses, clustering, and colours), but also new and more sensitive observations."
    },
    "1004/1004.1776_arXiv.txt": {
        "abstract": "Ghost inflation predicts almost scale-invariant primordial cosmological perturbations with relatively large non-Gaussianity. The bispectrum is known to have a large contribution at the wavenumbers forming an equilateral triangle and the corresponding nonlinear parameter $f_{NL}^{equil}$ is typically of order $O(10^2)$. In this paper we calculate trispectrum from ghost inflation and show that the corresponding nonlinear parameter $\\tau_{NL}$ is typically of order $O(10^4)$. We investigate the shape dependence of the trispectrum and see that it has some features different from DBI inflation. Therefore, our result may be useful as a template to distinguish ghost inflation from other models of inflation by future experiments. ",
        "introduction": "\\label{sec:intro}  Almost scale-invariant primordial cosmological perturbations predicted by inflation fits observational data very well~\\cite{Komatsu:2010fb}. While this is certainly a great success of the general idea of inflation, there still remain many unanswered important questions about inflation.   One of those important questions is how to distinguish different models of inflation observationally. There are many models of inflation which are consistent with observational data. It is often thought that tensor mode fluctuations~\\cite{GW} and non-Gaussianity~\\cite{Maldacena:2002vr,squeeze,equilateral, ArkaniHamed:2003uz,Holman:2007na,trispectrum,Chen:2009bc,Huang:2006eha,Arroja:2009pd, Mizuno:2009mv} will be useful to distinguish some of them. Non-Gaussianity is the main subject of the present paper.   While single-field, simple slow-roll inflation predicts negligibly small non-Gaussianity~\\cite{Maldacena:2002vr}, it is well known that there are many ways to generate non-Gaussianities large enough to be detected by near future experiments. They can be categorized into three by epochs in which non-Gaussianities are generated: (i) super-horizon, (ii) horizon-crossing and (iii) sub-horizon epochs. Each of these three types has a bispectrum with characteristic dependence on shapes of triangle formed by wave-vectors. The bispectrum for each type has a large contribution at (i) squeezed-~\\cite{squeeze}, (ii) equilateral-~\\cite{Chen:2006nt,equilateral, ArkaniHamed:2003uz} and (iii) folded-triangle~\\cite{Holman:2007na}, respectively. Among them, our interest in the present paper is on the type (ii), in which large non-Gaussianity is typically due to higher derivative terms whose importance is enhanced by smallness of the sound speed. Concrete examples of this type includes k-inflation~\\cite{Chen:2006nt}, DBI inflation~\\cite{Chen:2006nt,equilateral} and ghost inflation~\\cite{ArkaniHamed:2003uz,ArkaniHamed:2003uy,ArkaniHamed:2005gu,Senatore:2004rj}.   While bispectrum is the leading deviation from Gaussian statistics and thus a useful tool to distinguish some models of inflation from others, trispectrum can also provide additional information about inflation~\\cite{trispectrum,Chen:2009bc,Huang:2006eha,Arroja:2009pd,Mizuno:2009mv,RenauxPetel:2009sj}. In this paper we shall calculate trispectrum from ghost inflation, hoping to find ways to distinguish ghost inflation from other inflationary models which predict similar bispectra. We shall find that the shape-dependence of trispectrum in ghost inflation shows some difference from that in DBI inflation~\\cite{Chen:2009bc,Huang:2006eha,Arroja:2009pd,Mizuno:2009mv,RenauxPetel:2009sj}. Therefore, the result of the present paper may be useful as a template to distinguish different models of inflation by future experiments.   The rest of this paper is organized as follows. In Sec.~\\ref{review} we  review the ghost inflation and show the powerspectrum and the bispectrum in this model. In Sec.~\\ref{trispectrum} we calculate the trispectrum in the ghost inflation. In Sec.~\\ref{Sum} is devoted to a summary of this paper and discussion. In Appendix~\\ref{in-in} we give a brief review of the in-in formalism. In Appendix~\\ref{interactionhamiltonian} we present the method to construct the interaction Hamiltonian. In Appendix~\\ref{calculations} we show the details of the calculations. ",
        "conclusions": "\\label{Sum}  In this work we have calculated and investigated the trispectrum of curvature perturbation generated during ghost inflation.  The analysis of scaling dimensions of operators makes it possible to identify the leading diagrams contributing to the trispectrum in ghost inflation. Actually, there are two leading-order contributions. One is represented by a diagram with one four-point vertex. This contribution is called contact term contribution. The other is represented by a diagram with two three-point vertices and called scalar exchange contribution. We have analyzed these two contributions separately.   We have obtained general expressions for the two contributions as functions of six independent parameters. The six parameters are amplitudes of four $3$-momenta and two angles between momenta. (Note that the sum of four $3$-momenta must vanish because of the momentum conservation.)   In order to calculate the concrete values, we have focused on the equilateral case where all momenta has the same amplitude and where there remain two independent angular parameters as well as an overall amplitude of $3$-momenta. Then we have calculated the non-linear parameters $\\tau_{NL}^{cc}$ and $\\tau_{NL}^{se}$ for the contact term contribution and the scalar exchange contribution, respectively.   In the case of local-type non-Gaussianity, it was forecasted that PLANCK will give the constraint $|\\tau_{NL}|\\sim 560$~\\cite{Kogo:2006kh}. In the present paper we have shown that $\\tau_{NL}^{cc}$ and $\\tau_{NL}^{se}$ are typically of order $O(10^4)$. Therefore, the trispectrum from ghost inflation is probably detectable by PLANCK. Note, however, that the meaning of the non-linear parameter $\\tau_{NL}$ is different for different types of non-Gaussianities. (The same is true for the non-linear parameter $f_{NL}$ of bispectra.) Therefore, as a future work, it is important to investigate detectability of the trispectrum predicted by ghost inflation in more detail.   Now let us compare our results with the trispectrum from DBI inflation calculated in Refs.~\\cite{Arroja:2009pd,Mizuno:2009mv}. The overall behaviors of trispectra are indeed similar. The trispectrum from ghost inflation has a peak at the equilateral configurations, i.e. when all four $3$-momenta have the same amplitude, because non-Gaussianity is mainly generated in the horizon-crossing epoch. This feature of trispectrum is shared with DBI inflation. Moreover, the dependence of the equilateral trispectrum on the angular variables $C_2$ and $C_3$ also has similarities. In both DBI inflation and ghost inflation, the scalar exchange contribution has the maximum value at the most symmetric point $C_2=C_3=C_4=-1/3$, and the absolute value of the contact term contribution becomes minimum at that point. The contact term contribution to the equilateral trispectrum has similar dependence on $C_2$ and $C_3$ in the two models of inflation. This can be easily seen by comparing Fig.~\\ref{fig:Tcc.eps} in the present paper and the right figure of Fig.1 in Ref.~\\cite{Mizuno:2009mv}.   There are also some differences between the trispectrum from DBI inflation and that from ghost inflation. In DBI inflation, the value of the equilateral trispectrum is almost constant except for the edge region near the boundaries defined by $C_i=-1$ ($i=2,3,4$). This feature can be seen in, e.g., the left figure of Fig.1 of Ref.~\\cite{Mizuno:2009mv}. The trispectrum rapidly decreases near the boundaries $C_i=-1$ and the plateau looks like a triangle. On the other hand, as we can see in Fig.\\ref{fig:seT.eps} of the present paper, in ghost inflation the value of the equilateral trispectrum smoothly decreases towards the boundaries $C_i=-1$ ($i=2,3,4$). As a result the shape of the plateau looks different from that in DBI inflation.   If we look into actual values, we find an important difference. (Note that those figures mentioned above are normalized by the values at $C_2=C_3=C_4=-1/3$ and thus do not tell the actual values of trispectra.) The sign and magnitude of the contact term contribution depend on the sign and magnitude of the quartic coupling constant. (This is in contrast to the scalar exchange contribution, whose sign does not depend on the sign of the cubic coupling constant.) In DBI inflation, the quartic coupling constant is determined by the sound speed. As a result, the equilateral trispectrum from DBI inflation has a positive value at the most symmetric configuration $C_2=C_3=C_4=-1/3$. On the other hand, in ghost inflation the dimensionless quartic coupling constant is an arbitrary parameter of order unity and thus the equilateral trispectrum at $C_2=C_3=C_4=-1/3$ can be either positive or negative.  Therefore, if equilateral-type non-Gaussianity is detected either by bispectrum or by trispectrum and if the negative equilateral trispectrum at the most symmetric point is observed, then it will support the ghost inflation scenario.    \\vspace{1cm} \\textit{Note added}: While we were preparing the present paper, ref.~\\cite{Huang:2010ab} appeared on the arXiv. Before that, we had finished all calculations and presented our results at the workshop  \"The non-Gaussian universe\" at Yukawa Institute for Theoretical Physics on March 25, 2010. The presentation file has been available from the workshop website at \\verb+http://www2.yukawa.kyoto-u.ac.jp/~nlg/2010_3/program.htm+ since March 27, 2010."
    },
    "1004/1004.4365_arXiv.txt": {
        "abstract": "We present the results of a baryonic Tully-Fisher relation (BTFR) study for a local sample of relatively isolated disk galaxies. We derive a BTFR with a slope near $3$ measured over about $4$ dex in baryon mass for our combined \\textrm{H\\,\\scriptsize{I}} and bright spiral disk samples. This BTFR is significantly flatter and has less scatter than the TFR (stellar mass only) with its slope near $4$ reported for other samples and studies.  A BTFR slope near 3 is in better agreement with the expected slope from simple $\\Lambda$CDM cosmological simulations that include both stellar and gas baryons. The scatter in the TFR/BTFR appears to depend on $W_{20}$: galaxies that rotate slower have more scatter. The atomic gas--to--stars ratio shows a break near $W_{20} = 250$ \\kms\\, probably associated with a change in star formation efficiency. In contrast the absence of such a break in the BTFR suggests that this relation was probably set at the main epoch of baryon dissipation rather than as a product of later galactic evolution. ",
        "introduction": "\\label{sec:intro} The baryonic Tully-Fisher relation (BTFR) for disk galaxies relates the total baryon disk mass to the disk rotational velocity \\cite[$\\eg$][]{kcf99, mcg2000,bel01}.  It has long been recognized that the (luminous) TFR implies a coupling between the luminous and dark components of disk galaxies \\citep[$\\eg$][]{pie92}. Simple cosmological arguments \\citep[$\\eg$][]{whi97} predict that the slope of the BTFR should be close to 3.  In this approach, the galaxy mass is calculated within its virial radius, taken to be the radius $r_{200}$ within which the mean baryon mass surface density is $200$ times the critical density of the universe. In its simplest form, the dark halo is modeled as a singular isothermal sphere with a density distribution $\\rho(r) = V^{2}/(4\\pi G r^{2})$. The only dimensional parameter is the rotational velocity $V$.  It follows that $r_{200} = V/(10$\\hub) where \\hub is the Hubble constant, and the halo mass within $r_{200}$ is $\\mathcal{M}_{r_{200,\\rm{halo}}} = V^{3}/(10G$\\hub).  If some fraction $f_{\\rm d}$ of the halo mass is in the form of gas which becomes the exponential disk of the galaxy, then ${M}_{\\rm disk} = f_{\\rm d} V^{3}/(10GH_{\\mathrm{0}}\\,)$. We would then expect a BTFR with a slope of 3.  In this argument, the virial radius within which the mass was estimated is not a structural scalelength of the system in the sense of the scalelength of an exponential disk: it depends on the rotational velocity.  This predicted BTFR slope near $3$ is also seen in semi-analytic and numerical simulations of galaxy formation within the $\\Lambda$CDM framework: see \\citet{mo00, vand00, nav00, krav04}.  In reality, it appears that the rotational velocity $V$ of disk galaxies depends on the gravitational fields of both the baryons and the dark matter. $V$ is affected by the structure of the dark matter halo, the initial angular momenta of the baryons and dark matter, the structural evolution of the baryons and the adiabatic compression of the halo by the disk. The stellar and gas baryon masses are affected by baryon loss via winds and other feedback processes, the star formation efficiency and history, all of which vary, possibly in a systematic way, with galaxy mass and environment. Therefore, the slope, zero-point and possible departures from linearity of the BTFR should be sensitive to the many evolutionary processes that go on during galaxy formation from the main epoch of hierarchical assembly until the present time.  We should stress that all empirical TF/BTFR studies suffer from the caveat that there is no way yet to measure the rotational velocity at the virial radius, which may be larger or smaller than the velocity inferred from $W_{20}$ measurements \\citep[e.g.][]{bat05}. Also, the baryonic mass may or may not be proportional to the virial mass.  In this paper we derive the TFR and BTFR for a sample of relatively isolated disk galaxies covering a large range in mass and rotational velocity. Our rotational velocity measure is the width $W$ at twenty percent of the peak of the integrated $\\textrm{H\\,\\scriptsize{I}}$ profile. In \\S 2 we describe the two samples of galaxies used in this study.  In \\S 3 we present the observations and data reduction, and in \\S 4 outline the method of our analysis. \\S 5 contains the results of the observed and derived quantities. In \\S 6 our empirical disk scaling relations are presented and some astrophysical implications are discussed. In \\S 7, we conclude with a summary of the main results of this study. The Appendix includes an overview of the results of similar studies by other authors along with some discussion. ",
        "conclusions": "\\label{sec:conc} We choose a sample of isolated disk galaxies ranging from faint dwarfs to bright spirals. We construct TF and BTF relations and explore the difference between the theoretically predicted BTFR slope of 3 and the TFR slope of 4 obtained by many observers.  Regarding this difference, \\cite{vand00} argued that ``the physics regulating star formation and feedback, coupled with the mass dependence of halo densities and stellar populations has to tilt the TF relation to its observed slope. The introduction of a stability-related star formation threshold density increases the slope of the TF relation $\\ldots$''. Our results are entirely consistent with this argument. As \\velw decreases, the increasing gas--to--stars ratio and decreasing mean baryon mass surface density, possibly associated with a decreasing trend in star formation efficiency, generate the ``tilt\" between the TFR and the BTFR.  We show the gas--to--stars ratio for our combined sample, and the break at baryon masses near $1$ x $10^{10} \\mathcal{M}_{\\odot}$. Because the BTFR shows no such break, and its slope is close to that expected from cosmological arguments, one could argue that the total baryon content of isolated disk galaxies (as measured by stellar + \\textrm{H\\,\\scriptsize{I}} mass) has not been much affected by galaxy evolution, including star formation history. In this sense, the BTFR would be a fundamental relation relating back to the main epoch of galaxy assembly.  There are some systematic uncertainties which affect any discussion of the BTFR. (1) A problem inherent to any TF study is the change in \\textrm{H\\,\\scriptsize{I}} profile shape with $W_{20}$ \\citep{nor07} and therefore an uncertainty in how to relate $W_{20}$ to the rotational velocity $V$ across the whole range of $W_{20}$ values. (2) We have not included ionized or molecular gas in the total baryonic masses. The idea of large amounts of molecular gas in the dwarfs seems unlikely \\cite[$\\eg$][]{pil04,rea2005}. However, a larger fraction of ionized undetected baryons in the more massive galaxies would steepen the slope of the true BTFR. This ionized (warm) gas in the more massive galaxies \\cite[$\\eg$][]{mall04, fuku06} may turn out to be more significant in this respect."
    },
    "1004/1004.3290_arXiv.txt": {
        "abstract": "We examine the thirteen most luminous sources in the WMAP free-free map using the Spitzer GLIMPSE and MSX surveys to identify massive star formation regions, emitting one-third of the Galactic free-free luminosity. We identify star forming regions by a combination of bubble morphology in 8 $\\micronm$ (PAH) emission and radio recombination line radial velocities. We find 40 star forming regions associated with our WMAP sources, and determine unique distances to 31. We interpret the bubbles as evidence for radial expansion. The radial velocity distribution for each source allows us to measure the intrinsic speed of a region's expansion. This speed is consistent with the size and age of the bubbles. The high free-free luminosities, combined with negligible synchrotron emission, demonstrate that the bubbles are not driven by supernovae. The kinetic energy of the largest bubbles is a substantial fraction of that measured in the older superbubbles found by Heiles. We find that the energy injected into the ISM by our bubbles is similar to that required to maintain the turbulent motion in the gas disk inside 8 kpc. We report a number of new star forming regions powered by massive ($\\textrm{M}_{*} > 10^4 \\textrm{M}_\\sun$) star clusters. We measure the scale height of the Galactic O stars to be $h_{\\textrm{*}} = 35 \\pm 5 \\pc$. We determine an empirical relationship between the PAH and free-free emission of the form $F_{\\textrm{PAH}} \\propto F^2_{\\textrm{ff}}$. Finally, we find that the bubble geometry is more consistent with a spherical shell rather than a flattened disk. ",
        "introduction": "Massive star forming regions provide a unique laboratory to study both the process of massive star formation and feedback within the Galaxy. These clusters are home to the most massive stars in the Galaxy, including the vast majority of O and B stars. The stars in turn produce most of the ionizing luminosity and stellar winds that inject energy and momentum into the interstellar medium (ISM), blowing bubbles and producing shell structures. Finally, the most massive stars explode as supernovae, which also inject energy and momentum into the ISM.  The most massive star clusters (M$_{*}> 10^4 \\textrm{ M}_{\\sun}$), referred to as {\\it super star clusters}, have been regularly observed in extragalactic star forming regions \\citep{ho97}, but until recently such massive clusters have evaded detection in our own Galaxy. This has largely been due to the heavy dust obscuration within our own Galactic disk. Over the last decade a number of young, massive clusters have been found in the Galaxy, including the Arches and Quintuplet clusters near the Galactic centre region \\citep{figer99}, Westerlund 1 \\citep{clark05}, and RSGC 1, 2, \\& 3 \\citep{figer06, davies07, clark09}, either by directly imaging the stars, or by the stellar radio emission.  Another possible method of locating these young, massive star forming regions is to look for the environmental effects caused by such clusters, such as \\hii regions, or shell and bubble structures. These effects can be observed in wavebands where extinction through the Galactic plane becomes less of an issue, such as the radio and infrared. Extensive surveys of \\hii regions both in the northern and southern sky have been conducted with limited success in finding these massive star forming regions \\citep[for a recent census, refer to][]{conti04}.  In \\citet[][hereafter Paper I]{paperI}, we used the Wilkinson Microwave Anisotropy Probe (WMAP) maximum entropy method free-free foreground emission map \\citep{bennett03, gold09} to determine the star formation rate in the Galaxy. We found that the 18 most luminous regions (located within 13 WMAP sources) produce over half the total ionizing luminosity of the Galaxy, implying that half the total number of O stars reside in these regions. These sources contain bubble structures having radii ranging from 5 to 100 pc. Most of the ionizing photons from the embedded stellar population escape the bubble and ionize the surrounding material, creating the Extended Low-Density regions (ELD) \\citep{lockman76, anan85, anan85a}. Examination of the Spitzer Galactic Legacy Infrared Midplane Survey Extraordinaire (GLIMPSE) 8 $\\mu$m images shows that many known \\hii regions appear on shells or walls of bubbles. These known \\hii regions are generated either by the illumination of swept up material by the central cluster, or by star formation triggered in the swept up material. In either case, the luminosity of the central cluster is larger than that given by summing the total emission from all of the \\hii regions in the area.  Further, a large number of the known \\hii regions in a WMAP source had not previously been associated with one another, a result of their disparet radial velocities. The range of radial velocities had been translated into different distances along the line of sight. We, on the other hand, interpret the differences in radial velocity as the result of bubble expansion; the \\hii regions lie on shells, with velocity differences of order $15\\kms$, consistent with the expected expansion speed of a bubble in the ISM, e.g., \\citet{harperclark09}.  In this paper we analyze the most luminous WMAP sources using the GLIMPSE and MSX surveys, and previously known \\hii region velocities. We describe our procedure for associating the \\hii and PAH emission using an intrinsic expansion velocity criteria in \\S \\ref{dataanalysis}. We discuss the general properties of the star forming regions (which we define in \\S \\ref{sec:SFR}) in \\S \\ref{sfr}. We discuss each of the star forming regions in depth in Section \\ref{individual}. In Section \\ref{discussion}, we determine the scale height of O stars in the Galactic disk, quantify the relationship between the PAH and free-free emission, discuss the expansion of the star forming regions as a turbulent driving mechanism of the Galaxy's molecular gas, and comment on the three-dimensional geometry of observed bubble structures. Finally, we summarize our results in Section \\ref{summary}. ",
        "conclusions": "\\label{discussion} It has been noted before that regions with elevated free-free emission are associated with elevated PAH emission \\citep{cohen01, paperI}. In Paper I, we used this correlation to argue that, based on the higher resolution 8 $\\micronm$ images, both the free-free and PAH emission are powered by a central source. In this paper, we have shown that bubbles are associated with all of the sources we have examined. The bubbles are identified by their bright rims as seen in PAH emission.  In addition, the bubbles are surrounded by an elevated background of PAH emission. The total luminosity in the $8\\micronm$ band of our WMAP sources is dominated by the elevated background as opposed to the high-surface brightness emission from the bubble walls. We conclude that the bulk of the ionizing photons from the central source escape through the bubble walls to reach distances of 100-200 $\\pc$ or more, confirming the results of Paper I.  Previous estimates of stellar masses associated with giant \\hii regions have been underestimates of the total stellar mass at those locations. This follows from the fact that the \\hii regions reprocess only a small fraction of the ionizing photons in a given WMAP source.  We suggest that massive clusters inhabit the central regions of the bubbles. These clusters have evaded detection due to the lack of high surface brightness \\hii or radio continuum emission. We suggest that this results from a lack of gas and PAH particles in the immediate vicinity of the young clusters. The clusters have evacuated their surroundings, leaving little material to reprocess the ionizing photons emitted by the cluster.  \\subsection{Scale Height of O Stars } \\label{scaleheight} As discussed above, the WMAP sources studied in this paper are expected to contain approximately one-third of the O stars in the Galaxy. Using those objects in our catalogue with unique distances, we estimate the scale height of the O stars. As a first approximation, we assume that each of the star forming regions contain a similar number of O stars.  Using this assumption, we construct the cumulative distribution function of star forming regions as a function of their height above the galactic plane. The scale height of the distribution is  \\begin{equation} h_{\\textrm{*}} = 35\\pm 5 \\textrm{pc.} \\end{equation}  This value is consistent with that obtained by \\citet{elias06} for the local galactic disk (within 1 kpc) of $h_{\\textrm{LGD}}= 34 \\pm 3$ pc for their O-B2 subsample. Both are smaller than the value obtained by \\citet{reed00} of $h_{\\textrm{LGD}}= 45 \\pm 20$ pc for all OB stars within a distance of 4 kpc, but easily fall within their error bars. We note that the scale height of the molecular gas is $h \\sim 40$ pc \\citep{malhotra94}.  \\subsection{Free-Free to PAH Emission Relationship } \\label{pah} A morphological correlation between the MSX 8 micron emission and the radio continuum \\citep{cohen01}. We investigate this correlation using the WMAP free-free emission maps and the PAH emission from the GLIMPSE 8 micron mosaics.  For each WMAP source, we summed the free-free emission from the inside the Source Extractor ellipse from Paper I. We summed the $8\\micronm$ from the same ellipse in the Spitzer GLIMPSE mosaic. The result is shown in Figure \\ref{fluxfigure}. From a least-square fit, we find the following relationship:  \\begin{equation} F_{\\textrm{PAH}}\\propto F_{\\textrm{ff}}^{2.0\\pm0.34} \\label{obsrel} \\end{equation}  where $F_{\\textrm{PAH}}$ is the integrated 8 micron GLIMPSE flux and $F_{\\textrm{ff}}$ is the integrated 90 GHz free-free emission from the WMAP free-free foreground emission map. We find a similar relationship using the MSX Band A integrated flux in place of the GLIMPSE measurements.  \\subsection{Velocity Dispersion of the Molecular Gas}  Both molecular and atomic gas in the disk of the Milky Way are seen to have supersonic velocity dispersions.  These dispersions are normally interpreted as being due to turbulence, although their origin is uncertain. If the motions are due to turbulence, they must be driven, since undriven turbulence decays on roughly a dynamical time \\citep{maclow98}. Furthermore, the turbulence must be driven on $100 \\pc$ scales, since three dimensional turbulence cascades from large scales to small scales and not the other way around. A number of driving mechanisms have been proposed \\citep{maclow04, miesch94}, including supernovae, stellar winds, and gravitational instabilities, with no conclusive evidence for any particular mechanism. We investigate the kinetic energy that is injected into the ISM by the expansion of the massive bubbles that we identify. We note that all of the star forming regions we have identified are likely to be in giant molecular clouds---a preliminary search shows that more than thirty are in fact in molecular clouds.  We calculate the mechanical luminosity in the expansion of each of the star forming regions using \\begin{equation} L_{mech} = \\frac{\\pi}{2}\\Sigma_{0} \\Delta v_{c}^{3} r, \\end{equation} where $\\Sigma_{0} = 170 M_{\\sun}$ pc$^{-2}$ \\citep{solomon87} is the surface density of a GMC, and $r$ and $\\Delta v_{c}$ are taken from Table \\ref{sfrlist}. For SFRs where the half-spread velocity was not measured, we used the mean half-spread velocity, $\\Delta v$ = 12 km s$^{-1}$, of the known sources. For cases where the kinematic distance remains ambiguous, we take the mean radius of the region determined from both the near and far distances. Since the ratio of the two distances is less than two, this introduces an error of less than a factor of two for 9 out of our 40 SFRs. For each of the SFRs, we present the calculated dynamical properties in Table \\ref{dynprop} with the columns as follows: column (1) the catalogue number, column (2) the mass swept up in the shell in solar masses, column (3) the kinetic energy of the swept up shell in erg, and column (4) the mechanical luminosity produced by the motion of the shell in $\\textrm{erg s}^{-1}$.  The total mechanical luminosity being injected into the ISM inside the solar radius due to the expansion of these SFRs is the sum of the individual luminosities.  We find \\begin{equation} L_{mech} \\approx 6.7 \\times 10^{38} \\left(\\frac{\\Sigma_{0}}{170 M_{\\sun}\\pc^{-2}}\\right) \\textrm{ erg s}^{-1}. \\end{equation} This sum encompases only one-third of the star formation in the Galaxy. The other two-thirds should supply a proportionate amount of mechanical luminosity, for a total kinetic luminosity of $\\sim2\\times10^{39}\\erg\\s^{-1}$.  We compare this luminosity to that required to maintain the velocity dispersion in the molecular gas within the the solar circle:  \\begin{equation} L_{turb} \\equiv \\frac{1}{2}\\frac{ M v^3}{2h}, \\end{equation} where $h = 40$ pc is the scale height of the molecular disk, the molecular gas mass inside the solar circle $M = 1.0 \\times 10^{9} \\textrm{M}_{\\sun}$, and a molecular gas velocity $v = \\sqrt{2 \\ln{2}} \\sigma_{mol}$ with $\\sigma_{mol} = 7\\kms$ \\citep{malhotra94}.  The turbulent luminosity is \\begin{equation} L_{turb} \\approx 2.4 \\times 10^{39} \\left(\\frac{v}{8 \\kms} \\right)^3 \\textrm{ erg s}^{-1}. \\end{equation}  We conclude that the mechanical luminosity we see in the bubbles, multiplied by a factor of three to account for the other two thirds of Galactic star formation, is sufficient to power the turbulent luminosity seen in the molecular gas.  We note that the kinetic energies calculated for most of the bubbles (refer to Table \\ref{dynprop}) range from $10^{48}$ to just below $10^{52}\\textrm{ erg s}^{-1}$. The kinetic energies of our most energetic bubbles are similar to those measured in the superbubbles identified by \\citet{heiles79}, which range from $4\\times10^{51}\\erg\\s^{-1}$ to $2\\times10^{53}\\erg\\,\\s^{-1}$.  We note that our bubbles are selected by their ionizing luminosity, i.e., we require a very young stellar population. In addition, we do not see significant synchrotron radiation from these areas, indicating the absence of supernova remnants. This implies that the clusters are too young to have had more than a few supernova explode, and many of our sources likely have not had any supernovae go off. We infer that the superbubbles in the Galaxy are not initially driven by the energy from supernovae, but rather by the energy injected into the medium by the massive stars during their lifetime. It may well be, however, that supernovae contribute significantly to the kinetic energy later in the evolution of a superbubble.  \\subsection{Bubbles: Spherical Shells or Flattened Rings?}  The three-dimensional geometry of the bubbles identified by \\citet{churchwell06} has been contested recently. \\citet{beaumont10} propose that Churchwell et al.'s bubbles are flattened rings. \\citet{beaumont10} suggest that the aspect ratios of these rings may be anywhere from a few to as much as 10. Nineteen of our forty eight bubbles are in the GLIMPSE bubble catalogs. We test the assertion of \\citet{beaumont10} statistically using the bubbles in Table \\ref{bubblelist}.  We model our bubbles as ellipsoids, following the procedure of \\citet{noumeir99} with semi-axes $a, b, c$ aligned along the x, y and z axes such that $a \\ge b \\ge c$. The matrix equation for the ellipsoid is \\begin{equation} u^{\\intercal} X u = 1 \\end{equation} where $u$ is an arbitrary position vector. The matrix $X$ is a diagonal matrix with entries $a^2, b^2, c^2$, corresponding to the semi-axes of the ellipsoid. We rotate the ellipsoid along the z-axis by an angle $\\theta$, and along the y-axis by an angle $\\phi$, producing the matrix  \\begin{equation} X' = \\left[ \\begin{array}{ccc} {\\scriptstyle \\alpha \\cos^2 \\phi + c^2 \\sin^2 \\phi } & {\\scriptstyle \\beta \\cos \\phi \\cos \\theta \\sin \\theta }& {\\scriptstyle (c^2 - \\alpha)\\sin \\phi \\cos \\phi} \\\\ {\\scriptstyle \\beta \\cos \\phi \\cos \\theta \\sin \\theta }& {\\scriptstyle a^2 \\sin^2 \\theta + b^2 \\cos^2 \\theta }& {\\scriptstyle -\\beta \\sin \\phi \\cos \\theta \\sin \\theta }\\\\ {\\scriptstyle -(\\alpha + c^2) \\sin \\phi \\cos \\phi }& {\\scriptstyle -\\beta \\sin \\phi \\cos \\theta \\sin \\theta }& {\\scriptstyle \\alpha \\sin^2 \\phi + c^2 \\cos^2 \\phi} \\end{array} \\right] \\end{equation} where we define $\\alpha = a^2 \\cos^2 \\theta + b^2 \\sin^2 \\theta $, and $\\beta = a^2 - b^2 $.  To determine the observed ellipse on the sky resulting from the rotation of the ellipsoid, we produce an orthographic projection of the ellipsoid onto the y-z plane. The resulting projection is \\begin{equation} X_P = \\left[ \\begin{array}{cc} a^2 \\sin^2 \\theta + b^2 \\cos^2 \\theta & -\\beta \\sin \\phi \\cos \\theta \\sin \\theta \\\\ -\\beta \\sin \\phi \\cos \\theta \\sin \\theta & \\alpha \\sin^2 \\phi + c^2 \\cos^2 \\phi\\\\ \\end{array} \\right] \\end{equation}  To minimize the observed axis ratio, we assume the bubbles are circular rings with $a = b$, implying the semi-major axis of the projected bubble is $a$, and the projected semi-minor axis is $\\sqrt{a^2 \\sin^2 \\phi + c^2 \\cos^2 \\phi}$. The resulting axis ratio, a function of only the rotation along the y-axis, is given as \\begin{equation} \\label{eqn: aspect} R_{ax}(\\phi) = \\frac{a}{\\sqrt{a^2 \\sin^2 \\phi + c^2 \\cos^2 \\phi}} \\end{equation}  We compute the expectation value of the axis ratio $\\langle R_{ax}(a/c)\\rangle$, for $a/c=10$ and $a/c=4$. Using eqn. (\\ref{eqn: aspect}) we find $\\langle R_{ax}(10)\\rangle = 2.4$ and $\\langle R_{ax}(4)\\rangle = 1.8$.  We compare this to our sample of $48$ bubbles, for which we find $\\langle R_{ax}\\rangle = 1.3\\pm0.3$, suggesting that the intrinsic aspect ratio $a/c\\lesssim4$ at the two sigma level. In fact $\\langle R_{ax}\\rangle = 1.3$ corresponds to a mean intrinsic aspect ratio $a/c=1.7$.  A more sensitive test is that of the maximum aspect ratio; the maximum observed aspect ratio is 2.2, corresponding to a maximum inclination angle of $31^{\\circ}$ from the face-on position of a ring with $a/c=4$. For a simple random distribution of inclination angles, we would expect $48\\%$, or $23\\pm7$ of the bubbles to have an aspect ratio greater than $2.2$.  We note that if the observed maximum aspect ratio ($2.2$) is in fact the true mean aspect ratio, then $\\langle R_{ax}(2.2)\\rangle=1.4$, consistent with the observed mean aspect ratio.  We find an insufficient number of observed bubbles with large aspect ratios to support the geometry of a flattened ring and the original picture of nearly-spherical shells is more likely to be physically correct."
    },
    "1004/1004.1406_arXiv.txt": {
        "abstract": "We develop a simple model of planetary formation, focusing our attention on those planets with masses less than $10 M_{\\oplus}$ and studying particularly the primordial spin parameters of planets resulting from the accretion of planetesimals and produced by the collisions between the embryos. As initial conditions, we adopt the oligarchic growth regime of protoplanets in a disc where several embryos are allowed to form. We take different initial planetary system parameters and for each initial condition, we consider an evolution of $2x10^7$ $years$ of the system. We perform simulations for $1000$ different discs, and from their results we derive the statistical properties of the assembled planets. We have taken special attention to the planetary obliquities and rotation periods, such as the information obtained from the mass and semi major axis diagram, which reflects the process of planetary formation. The distribution of obliquities was found to be isotropic, which means that planets can rotate in direct or indirect sense, regardless of their mass. Our results regarding the primordial rotation periods show that they are dependent on the region where the embryo was formed and evolved. According to our results, most of the planets have rotation periods between $10$ and $10000 \\; hours$ and there are also a large population of planets similar to terrestrial planets in the Solar System. ",
        "introduction": "Following the first discovery of an extrasolar planet around 51 Peg \\citep{b20}, the number of exoplanets known has risen to 429. Although most of them are giant planets, the improvements in observational techniques have ensured that planets with masses less than $15 M_{\\oplus}$ have started being detected with radial velocity survey (e.g., \\citet{b22,b23,b24,b25,b21}) and gravitational microlensing survey \\citep{b26}.  Although most of extrasolar planets so far discovered are giant planets, several statistical models for planetary growth presented in the last years suggest that a large number of small planets who fail to have enough mass to start the gas accretion onto the core exists \\citep{b9,b2,b28}, and has still not been able to be discovered \\citep{b27}. At the time, several projects are in progress to detect terrestrial planets, we expect that they may find more Earth-size planets in  a close-future, but today, the sample is not enough and we also have to rely on what we know from our own Solar System, and through computational models of planetary formation.  This evidence supports the standard scenario, where terrestrial planets are formed through the next different stages: 1) agglomeration of dust particles through physical collisions and setting in the protoplanetary disc, 2) planetesimal formation from grains in a thin midplane \\citep{b31,b32}, 3) runaway (e.g., \\citet{b33}) and oligarchic \\citep{b4,b15} accumulation of planetesimals to form protoplanets and 4) giant impact stage, where the embryos formed by oligarchic growth collide with one another to form planets \\citep{b30}.  The final stage of terrestrial planetary formation is the particular importance as it has a deep effect on the final characteristics of the planets: mass, orbital and spin parameters. After this stage of planetary formation, the spin parameters of the planets change and evolve due mainly to tidal interactions with their satellite and host star. All of the terrestrial planets in our Solar System do not maintain their primordial spin state and this is the reason why we unknown what primordial planetary spin would be expected to find in a protoplanet. So questions as, what are the typical obliquity and rotation period that characterise the primordial planets? and how many collisions suffers a planetary embryo along its firsts years of formation? remain uncertain.  A few works dealing with the study of planetary spins have been presented. \\citet{b34} have examined the accretion rate of spin angular momentum by a planet immersed in a differentially rotating disc of planetesimals. They determined the mass and spin accreted by the embryos as a function of the velocity dispersion of the disc particles and the ratio of the planetary radius to the Hill radius. They found that if a protoplanet grows by accreting a large number of small planetesimals the spin angular momentum of the planet will be determined by the called ``ordered component'', but if a few giant impacts occur, most of the spin will be contributed by the ``stochastic component''. \\citet{b35} have investigated the spin of a planet which accreted in a disc of planetesimals with non uniform spatial distribution. They results show that the ordered component can dominate the final spin of the planet only if half of the size of the planet was acquire by the accretion of small planetesimals and the size of the impactors is not too large.  On the other hand \\citet{b13} and \\citet{b19} have studied through N-body simulations the last stages of the terrestrial planet formation. They analysed the planetary obliquities as those found only considering the impacts between large embryos and have shown that this obliquities are expected to be represented by an isotropic distribution, result that was confirmed and generalized by \\citet{b14}, who also considered an N-body code, but analysed a larger sample of embryos considering the standard disc model.  Our principal aim is to make a statistical study of the primordial spin parameters of planets (obliquity and rotation period), resulting from the accretion of planetesimals and also due to the collisions between the emerging embryos. To this end we take different initial conditions, meaning different discs, stars, initial number of embryos, and study the primordial planetary spins in different systems with the intention of obtain a better understanding of what we should expect to find in the Universe. We also analyse what are the consequences of planetary impacts in the mass and semi major axis diagram, considering embryos with masses less than $10 M_{\\oplus}$. Our semi-analytical model takes as initial condition the oligarchic growth regime of protoplanets and allows them to migrate, fact that has a huge influence on the number of collisions suffered by an embryo. We adopt a perfect accretion in collisions, supposition that was also considered by other authors \\citep{b13,b19,b14}, but which says that the results should be interpreted cautiously.  Each one of the 1000 systems considered, evolves for $2 x 10^7$ $years$ and we analyse the results statistically, finding an isotropic distribution of obliquities and where most of the planets rotate with a period between $10$ and $10000 \\; hours$. We also found a large population of planets with the characteristics of terrestrial planets in the Solar System. ",
        "conclusions": "In the process of planetary formation protoplanets collide with one another to form planets. We have investigated the final assemblage of terrestrial planets from protoplanets using a simple model which consider the oligarchic growth regime of protoplanets as initial condition in a disc where several embryos are allowed to form. As explained in our previous work the formation of several cores simultaneously in the disc has a strong influence on the dynamic of the planetesimal disc, which influences directly the growth of the embryos' cores and the final assemblage of planets found.  In our model we also have included the interaction between the protoplanets and the disc, which leads to a planetary migration. When a embryo is migrating towards the central star it could perturbate the cores placed in its path, causing the accretion of the core in most cases, this collisions affect the spin state of the embryos.  As collision among giant planets are poorly understood we have focused our attention on planets with masses less than $10 M_{\\oplus}$ where a very simple model for planetary impacts has been considered. We suppose that when two embryos are a distance less than $3.5$ $R_{H}$ the merger between both protoplanets occurs, which leads to the union of two embryos to form a single body. This perfect accretion model produces spin rates that are too high and when the acceleration produced by the rotation is greater than those of gravity the body overcome the critical spin angular velocity for rotational instability and is fragmented. This simple model allows us to obtain some interesting results regarding the final properties of terrestrial planets.  We also have considered the acquisition of angular momentum due to accretion of planetesimals. The accretion of a large amount of planetsimals produces an ordered spin that adds angular momentum to that acquired during collisions, so the final spin of the planets is a result of this two effects.  In order to analyse the statistical properties of the assembled planets we take different initial planetary system parameters, considering $1000$ different discs, where each planetary system evolves $2x10^7$ $years$.  As in our previous works we have analysed the information provided by the mass and semi major axis diagram, which reflects the process of planetary formation. We observe fewer planets with masses less than $1 M_{\\oplus}$ considering the fragmentation by collisions that those found in the without this effect. This means that the effect of fragmentation by collision has a strong influence on the final population of terrestrial planets formed and should be considered when these planets are involved.  We also have studied the effects produced by the collisions between the embryos, where we find that most of the planets suffer less than 5 impacts during its formation, which means that in most of the cases primordial spins of planets are randomly determined by a very few impacts suffered during accretion.  We also take special attention to final spin state, which means planetary obliquities and rotation periods, where we found that the distribution of obliquities of final planets is well expressed by an isotropic distribution, result that confirms those obtained previously by other authors \\citet{b13,b14} and is independent on the planetary mass. This fact is in marked contrast to the terrestrial planets in our own Solar System, whose current spin axes are more or less perpendicular to their orbital planes (except for Venus). However, the spin axis of the terrestrial planets strongly depends on the gravitational perturbations from the other planets of the Solar System that create a large chaotic zone for their obliquities. So all of the terrestrial planets could have experienced large, chaotic variations in obliquity in their history, and this is why their obliquities can not be considered as primordial \\citep{b36}. So the fact that the terrestrial planets in our Solar System present obliquities $\\sim 0^{\\circ}$ does not necessarily indicate a problem with the model considered here. Other studies such as body and atmospheric tides and core-mantle friction among others, must be taken into account for explaining the present obliquities of the terrestrial planets.  Regarding the findings on the rotation period, we found that the primordial rotation periods of terrestrial planets are dependent on the semi major axis, which means on the region where the embryos were formed and evolved.  On the one hand we note a very small population of planets with small rotation periods (less than $\\sim 0.5 \\; hours$), which are very rare planets, because at that rotation periods the spin angular velocities are high enough to overcome the critical rotation angular velocity for rotation instability.  On the other hand there are a large population of embryos with rotation periods until $\\simeq 10000 \\; hours$. These planets with large rotation periods probably acquired them mainly by the accretion of planetesimals, while those with shorter periods need one or more impacts for acquire that spin.  Another important result is that we have found a large population of planets with the characteristics of the Terrestrial Planets, and our results suggest that they did not acquire their rotation period only by the accretion of planetesimals, but during one or more impacts during their formation."
    },
    "1004/1004.3918_arXiv.txt": {
        "abstract": "{ Helicity is a fundamental property of magnetic fields, conserved in ideal MHD. In flux rope topology, it consists of twist and writhe helicity. Despite the common occurrence of helical structures in the solar atmosphere, little is known about how their shape relates to the writhe, which fraction of helicity is contained in writhe, and how much helicity is exchanged between twist and writhe when they erupt. } { Here we perform a quantitative investigation of these questions relevant for coronal flux ropes. } { The decomposition of the writhe of a curve into local and nonlocal components greatly facilitates its computation. We use it to study the relation between writhe and projected S shape of helical curves and to measure writhe and twist in numerical simulations of flux rope instabilities. The results are discussed with regard to filament eruptions and coronal mass ejections (CMEs). } { (1) We demonstrate that the relation between writhe and projected S shape is \\emph{not} unique in principle, but that the ambiguity does not affect low-lying structures, thus supporting the established empirical rule which associates stable forward (reverse) S shaped structures low in the corona with positive (negative) helicity. (2) Kink-unstable erupting flux ropes are found to transform a far smaller fraction of their twist helicity into writhe helicity than often assumed. (3) Confined flux rope eruptions tend to show stronger writhe at low heights than ejective eruptions (CMEs). This argues against suggestions that the writhing facilitates the rise of the rope through the overlying field. (4) Erupting filaments which are S shaped already before the eruption and keep the sign of their axis writhe (which is expected if field of one chirality dominates the source volume of the eruption), must reverse their S shape in the course of the rise. Implications for the occurrence of the helical kink instability in such events are discussed. (5) The writhe of rising loops can easily be estimated from the angle of rotation about the direction of ascent, once the apex height exceeds the footpoint separation significantly. } { Writhe can straightforwardly be computed for numerical data and can often be estimated from observations. It is useful in interpreting S shaped coronal structures and in constraining models of eruptions. }  \\keywords { Magnetic fields -- Magnetohydrodynamics (MHD) -- Sun: corona -- Sun: filaments, prominences -- Sun: coronal mass ejections (CMEs) } ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.0363.txt": {
        "abstract": "We present a detailed investigation of the Cepheid distance scale by using both theory and observations. Through the use of pulsation models for fundamental mode Cepheids, we found that the slope of the Period-Luminosity ($P$-$L$) relation covering the entire period range (0.40$\\le$log$P\\le$2.0) becomes steeper when moving from optical to near-infrared (NIR) bands, and that the metallicity dependence of the slope decreases from the $B$- to the $K$-band. The sign of the metallicity dependence for the slopes of the $P$-$L_V$ and $P$-$L_I$ relation is at odds with some recent empirical estimates.  We determined new homogeneous estimates of $V$- and $I$-band slopes for 87 independent Cepheid data sets belonging to 48 external galaxies with nebular oxygen abundance 7.5$\\le$12+$\\log$(O/H)$\\le$8.9. By using Cepheid samples including more than 20 Cepheids, the $\\chi^2$ test indicates that the hypothesis of a steepening of the $P$-$L_{V,I}$ relations with increased metal content can be discarded at the 99\\% level. On the contrary, the observed slopes agree with the metallicity trend predicted by pulsation models, i.e. the slope is roughly constant for galaxies with 12+$\\log$(O/H)$<$8.17 and becomes shallower in the metal-rich regime, with a confidence level of 62\\% and 92\\%, respectively. The $\\chi^2$ test concerning the hypothesis that the slope does not depend on metallicity gives confidence levels either similar ($PL_V$, 62\\%) or smaller ($PL_I$, 67\\%).  We investigated the dependence of the Period-Wesenheit ($P$-$W$) relations on the metal content and we found that the slopes of optical and NIR $P$-$W$ relations in external galaxies are similar to the slopes of Large Magellanic Cloud (LMC) Cepheids. They also agree with the theoretical predictions suggesting that the slopes of the $P$-$W$ relations are independent of the metal content. On this ground, the $P$-$W$ relations provide a robust method to determine distance moduli relative to the LMC, but theory and observations indicate that the metallicity dependence of the zero-point in the different passbands has to be taken into account.  To constrain this effect, we compared the independent set of galaxy distances provided by Rizzi et al.\\ (2007) using the Tip of the Red Giant Branch (TRGB) with our homogeneous set of extragalactic Cepheid distances based on the $P$-$W$ relations. We found that the metallicity correction on distances based on the $P$-$WBV$ relation is $\\gamma_{B,V}$=$-0.52$ mag dex$^{-1}$, whereas it is vanishing for the distances based on the $P$-${WVI}$ and on the $P$-${WJK}$ relations. %Point C These findings fully support Cepheid theoretical predictions. ",
        "introduction": "The cosmic distance scale and the estimate of the Hubble constant, $H_0$, are tightly connected with the Period-Luminosity ($P$-$L$) relation of Classical Cepheids and the distances to external galaxies are traditionally determined by using a universal $P$-$L$ linear relation based on the Large Magellanic Cloud (LMC) variables.  However, theoretical predictions based on nonlinear, convective Cepheid models computed by our group (see Caputo 2008 for a comprehensive list of references) indicate that the instability strip boundaries in the log$L$-log$T_e$ plane are almost linear, but when transformed into the different Period-Magnitude planes they are better described by quadratic $P$-$L$ relations. In particular, we found that at fixed metal content: (1) the predicted optical $P$-$L$ relations can be properly fit with quadratic relations, (2) a discontinuity around log$P\\sim$ 1.2 should be adopted to constrain the theoretical results into linear approximations, and (3) the predicted $P$-$L$ relations become more and more linear and tight when moving from optical to near-infrared bands. Moreover, we suggested that the metal-poor Cepheids follow $P$-$L$ relations which are steeper and brighter than the metal-rich ones, with the amount of this metallicity effect again decreasing from the $B$ to the $K$ band. Furthermore, we drew attention to the evidence that the metallicity effect on the predicted Period-Wesenheit ($P$-$W$) relations, which present several advantages when compared with the $P$-$L$ relations, significantly depends on the adopted Wesenheit function.  With a few exceptions, these theoretical results have been considered with a certain skepticism. Only during the last few years, several observational investigations disclosed the nonlinearity of the $P$-$L$ relation (Tammann, Sandage, \\& Reindl 2003; Ngeow et al.\\ 2005; Ngeow \\& Kanbur 2006), as well as the evidence that the Cepheid $P$-$L$ relation cannot be universal and that both the slope and the zero-point might change from galaxy to galaxy. Quoting Sandage, Tammann \\& Reindl (2009), \"{\\it the existence of a universal $P$-$L$ relation is an only historically justified illusion\".}  According to these new empirical evidence, which might imply severe limits in the precision of Cepheid distances, we examine the available observations by using the theoretical framework provided by the pulsation models and we address three main issues concerning the Cepheid distance scale: the intrinsic features of the $P$-$L$ and the $P$-$W$ relations, the dependence of the $P$-$L$ and of the $P$-$W$ slope on the Cepheid metal content and the impact of the metallicity effect on the Cepheid distances.   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "We performed a comprehensive investigation of the Cepheid distance scale by taking into account both theory and observations. In particular, we addressed the intrinsic features of both optical and NIR $P$-$L$ relation. Here are results:  {\\em i)} {\\em Filter wavelength} Theory and observations indicate that the slopes of the $P$-$L$ relation become steeper when moving from optical to NIR bands.  {\\em iii)} {\\em Nonlinearity} The slopes of the observed optical $P$-$L$ relations of Magellanic Cepheids are nonlinear. No firm conclusion was reached concerning the nonlinearity of the $P$-$L$ relations based on Galactic and M31 Cepheids.  {\\em ii)} {\\em Period range} The slopes of NIR $P$-$L$ relations are less sensitive to the period range covered by Cepheids than the slopes of optical $P$-$L$ relation.  {\\em iv)} {\\em Metal content} The derivative $\\partial b_{all}/\\partial$log$(Z/X)$ of the predicted slopes covering the entire period range decreases by more than a factor of two when moving from the $V$ to the $J$-band and by almost one order of magnitude when moving from the $V$ to the $K$-band. Moreover, the observed slopes of Magellanic, Galactic and M31 Cepheids agree quite well with predicted ones. In particular, they suggest a flattening of the slope when moving from metal-poor to metal-rich Cepheids. This finding is at odds with the steepening recently suggested by TSR08, by STR08 and by ST08.  In order to provide an empirical estimate of the dependence of the  $P$-$L$ relation on metal content, we also adopted Cepheids in external Galaxies. To avoid possible deceptive uncertainties in the adopted metallicity scale, we derived a new relation to transform the old nebular oxygen abundances given by  Zaritsky et al.\\ (1994) into the new metallicity scale provided by Sakai et al.\\ (2004).  Moreover, we provided new homogeneous estimates of $V$- and $I$-band slope for 87 independent Cepheid data sets available in the literature and 57 of them include more than 20 Cepheids. They are hosted in 48 external galaxies and for 27 of them two or more independent data sets are available. Four galaxies with multiple data sets (NGC~598, NGC~3031, NGC~4258, NGC~5457) have Cepheids located in an inner and in an outer galactic field. The galaxies with more than 20 Cepheids cover a wide metallicity range (12+log(O/H)$\\sim$7.7 [WLM], 12+log(O/H)$\\sim$8.9 dex [NGC~3351, NGC~4548]) Note that the quoted range is approximately a factor of five larger than the metallicity range covered by SMC (12+log(O/H)$\\sim$8) and Galactic (12+log(O/H)$\\sim$8.6) Cepheids. By using Cepheid data sets larger than 20, we tested three hypotheses concerning the dependence of the $P$-$L$ relation on metal content:  {\\em a)} {\\em Correlation between the slope of the $P$-$L_{V,I}$ relations and the metallicity.}  The $\\chi^2$ test on $V$- and $I$-band slopes indicates that this hypothesis can be discarded at the 99\\% confidence level.  {\\em b)} {\\em No dependence of the $P$-$L_{V,I}$ relations on the metallicity.} The $\\chi^2$ test on $V$- and $I$-band slopes indicates that this hypothesis can be accepted, but only at the 62\\% and 67\\% confidence level.  {\\em c)} {\\em Pulsation models predict that the slope of the $P$-$L_{V,I}$ relations becomes shallower in the metal-rich regime and constant in the metal-poor regime.} The outcome of the $\\chi^2$ test on observed $V$- and $I$-band slopes is that the predicted trend can be accepted at the 62\\% and 92\\% confidence level.  The main result of the above analysis based on external galaxies with sizable Cepheid samples is that the observed slopes of the  $P$-$L_I$ relation show the same metallicity trend predicted by pulsation models, while the slopes of the $PL_V$ relation either follow theory or do not depend on metallicity.  Together with the $P$-$L_{V,I}$ relations we also investigated the reddening independent $P$-$W$ relations and the results are the following:  {\\em i)} {\\em Dependence of the slope of the $P$-$W$ relations on metal content} Empirical estimates indicate that the slopes of optical ($P$-$WBV$,$P$-$WVI$) and NIR ($P$-$WJK_s$) relations in metal-poor and in metal-rich galaxies agree quite well with the slope of LMC Cepheids. This finding supports previous results by % Point B Benedict et al. (2007), Pietrzynski et al.\\ (2007), van Leeuwen et al. (2007), and by Riess et al.\\ (2009). % Moreover, it brings forward the evidence that the $P$-$W$ relations provide accurate estimates of LMC-relative true distance moduli. However, the metallicity dependence of the zero-point of the $P$-$W$ relations, if present, has to be taken into account.  {\\em ii)} {\\em Use of the $P$-$W$ relations as a metallicity diagnostic} Current predictions indicate that LMC-relative true distance moduli based on the $P$-$WBV$ relation strongly depend on the metal content, whereas those based on the $P$-$WVI$ and on the $P$-$WJK_s$ relation minimally depend on metallicity. The difference between the quoted distances can provide estimates of individual Cepheid metallicities.  The above findings further support the evidence that distances based on different $P$-$W$ relations should not be averaged, since the metallicity effect strongly depends on the adopted bands.  Furthermore, we adopted the true distance moduli based on the TRGB method (Ri07) to validate the predicted metallicity corrections of the Cepheid distance scale. We found that the metallicity correction --$\\gamma$-- obtained using the TRGB distances agree quite well with pulsation predictions, namely $\\gamma(WBV)$=$-0.52\\pm$0.09 mag dex$^{-1}$, $\\gamma(WVI)$=$-0.03\\pm$0.07 mag dex$^{-1}$ and $\\gamma(WJK)$=$-0.05\\pm$0.06 mag dex$^{-1}$.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "1004/1004.5363_arXiv.txt": {
        "abstract": " ",
        "introduction": "%  The corner-stone of modern cosmology is that, at least on large scales, the visible universe seems to be the same in all directions around us and around all points, i.e. the Universe is almost homogeneous and isotropic. This is borne out by a variety of observations, particulary observations of cosmic microwave background (CMB); this radiation has been traveling to us for about 14000 million years (see \\hbox{Fig. \\ref{cmbe}}), supporting the conclusion that the Universe at sufficiently large distances is nearly the same. On the other hand, it is apparent that nearby regions of the observable Universe are at present highly inhomogeneous, with material clumped into stars, galaxies and galaxy clusters. It is believed that these structures have formed over the time via gravitational attraction, from a distribution that was more homogeneous in the past.  The large-scale behavior of the Universe can be described by assuming a homogeneous background. On this background, we can superimpose the short scale irregularities. For much of the evolution of the observable Universe, these irregularities can be considered to be small perturbations on the evolution of the background (unperturbed) Universe. The metric of unperturbed Universe is called the Friedman-Leimatre-Roberson-Walker metric, and its line element can be to written as: \\be ds^2=-dt^2+a^2(t)\\(dr^2+r^2(d\\theta^2+\\sin\\phi^2d\\phi^2\\), \\ee where $a(t)$ is the scale factor and $r,\\theta,\\phi$ are the spherical comoving coordinates\\footnote{A particle in this metric have fixed-coordiantes.}  The model described by the above metric is known as the standard cosmological model (known also as Big-Bang cosmological model) \\cite{friedman1,robertson1,robertson2,robertson3,walker} and is the successful framework that describes the observed properties of the Universe: homogeneity and isotropy at large scales, Hubble expansion, almost 14 billion years of evolution in agreement with globular clusters and radioactive isotopes dating, cosmic microwave background radiation (CMB) confirmed by Penzias and Wilson's discovery in 1965 \\cite{dicke,penziaswilson}, and the relative abundances of light elements \\cite{alpher1,alpher2,gamow,hoyle,olive,wagoner,walkeretal} in full agreement with observation.  \\begin{figure}[!h] \\begin{center} \\includegraphics[scale=0.55]{CMBTimeline.eps} \\end{center} \\caption[A representation of the evolution of the universe over 13.7 billion years.]{A representation of the evolution of the universe over 13.7 billion years. The far left depicts the earliest moment we can now probe, when a period of `` inflation \" produced a burst of exponential growth in the universe. (Size is depicted by the vertical extent of the grid in this graphic.) For the next several billion years, the expansion of the universe gradually slowed down as the matter in the universe pulled on itself via gravity. More recently, the expansion has begun to speed up again as the repulsive effects of dark energy have come to dominate the expansion of the universe. The afterglow light seen by WMAP was emitted about 380,000 years after inflation and has traversed the universe largely unimpeded since then. The conditions of earlier times are imprinted on this light; it also forms a backlight for later developments of the universe (Courtesy of the NASA/WMAP Science Team \\cite{wmap}).} \\label{cmbe} \\end{figure}  The introduction of a period of exponential expansion (called inflationary) \\cite{albste,guth,linde82a}, prior to the Big-Bang, brought an elegant solution to the horizon, flatness, and unwanted relics problems that were present in the original standard cosmological model \\cite{albste,guth,kolb,linde82a,riotto1}. In spite of its success at solving the above mentioned problems, the inflationary period became perhaps more important because of its ability to stretch the quantum fluctuations of the fields living in the FRW spacetime \\cite{bst,guthpi82,hawking,linde82a,mukhanov1,mukhanovrep,riotto1,starobinsky1}, making them classical \\cite{albrecht1,Burgess:2006jn,grishchuk,guthpi,Kiefer:2008ku,lombardo,lyth84,lythbook,Lyth:2006qz,Martineau:2006ki,Nambu:2008my} and almost constant soon after horizon exit. They correspond to small inhomogeneities in the energy density and are responsible, via gravitational attraction, of the large-scale structure seen today in the Universe. If this scenario turned to be correct, the energy density inhomogeneities should have left their trace in the CMB released at the time of recombination. Indeed, the Cosmic Background Explorer (COBE) in 1992 \\cite{cobe} found and mapped small anisotropies in the CMB temperature of the order of 1 part in $10^5$ (with average temperature \\mbox{$T_0 = 2.725 \\pm 0.002$ K} \\cite{bennett}), on scales of order thousands of Megaparsecs. With 30 times better angular resolution and sensitivity than COBE, the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{wmap} confirmed this picture (see \\hbox{Fig. \\ref{wmapsky}}), measuring in turn the cosmological parameters with a $1\\%$ order precision \\cite{wmap5} on scales of order tens of Megaparsecs. The PLANCK satellite \\cite{planck,planck1}, launched in may 2009, will be able to refine these observations (see Fig. \\ref{plancksky} and \\ref{plancksky2}). With 10 times better angular resolution and sensitivity than WMAP, PLANCK promises to determine the temperature anisotropies with a resolution of the order of 1 part in $10^6$, and the cosmological parameters with a $0.1\\%$ order precision.  \\begin{figure*}[!h] \\begin{center} \\includegraphics[scale=1.2]{5yrFullSkyWMAPW.eps} \\end{center} \\caption[CMB temperature anisotropies as seen by the WMAP satellite.]{CMB temperature anisotropies as seen by the WMAP satellite (five years resuls) \\cite{wmap5}. The oval shape is a projection to display the whole sky. The temperature anisotropies are found to be of the order of 1 part in $10^5$. The background temperature is \\hbox{$T_0 = 2.725 \\pm 0.002$ K}; regions at that temperature are in very light blue. The hottest regions (in red) correspond to $\\Delta T \\simeq 200 \\mu{\\rm K}$. The coldest regions (in very dark blue) correspond to $\\Delta T \\simeq -200 \\mu{\\rm K}$ (Courtesy of the NASA/WMAP Science Team \\cite{wmap}).} \\label{wmapsky} \\end{figure*}  \\begin{figure*}[!h] \\begin{center} \\includegraphics[scale=0.80]{SkyPlanck_UIS.eps} \\end{center} \\caption[Simulation of the CMB temperature anisotropies as seen by the PLANCK satellite.]{Simulation of the CMB temperature anisotropies as seen by the PLANCK satellite. PLANCK will provide a map of the CMB field at all angular resolutions greater than 10 arcminutes and with a temperature resolution of the order of 1 part in $10^6$ (ten times better than WMAP) (Courtesy of ESA's PLANCK mission \\cite{planck}).} \\label{plancksky} \\end{figure*}  \\begin{figure*}[!h] \\begin{center} \\includegraphics[scale=0.5]{Planck.EPS} \\end{center} \\caption[The CMB temperature anisotropies as seen by the PLANCK satellite.]{A map of the area of the sky mapped by PLANCK during the first light survey. The colours indicate the magnitude of the deviations of the temperature of the Cosmic Microwave Background from its average value (red is hotter and blue is colder). (Courtesy of ESA's PLANCK mission \\cite{planck}).} \\label{plancksky2} \\end{figure*}  The anisotropies in the CMB temperature\\footnote{From now on, and unless otherwise stated, the perturbation $\\delta y$ in any quantity $y$ will be regarded as first-order in cosmological perturbation theory. Unperturbed quantities will be denoted by a subscript 0 unless otherwise stated.} $\\delta T/T_0$ are directly related to the perturbation in the spatial curvature $\\zeta$ (Sachs-Wolfe effect), whose primarily origin is the stretched quantum fluctuations of one or several scalar fields $\\phi_i$ that fill the Universe during inflation \\cite{lythbook,sachs}\\footnote{In this and the following expressions the subscripts $k$ stand for the Fourier modes with comoving wavenumber $k$.}: \\be \\left(\\frac{\\delta T}{T_0}\\right)_k = -\\frac{1}{5} \\, \\zeta_k \\,. \\label{connections1} \\ee The quantity $\\zeta$ is related to the perturbation in the intrinsic curvature of space-time slices with uniform energy density \\cite{malikwands}: \\be ^{(3)}R=\\frac{4}{a^2}\\;\\nabla^2\\psi\\;, \\ee where $\\psi$ is the first order scalar perturbation in the spacial metric.  Astronomers work with the observable quantity $\\delta T/T$ and theoretical cosmologists work with $\\zeta$. Therefore, we may study the statistical porperties of the observed $\\delta T/T$ through the spectral functions associated with the primordial curvature perturbation $\\zeta$, whose properties are in general model dependent. Knowing the statistical descriptors of $\\zeta$ for some particular and well motivated cosmological model proposed for the origin of large scale strucutre, we can reject the model or keep it, because some of the statistical descriptors for $\\delta T/T$ are known with good acuracy or at least have an upper bound \\cite{wmap5}.  The statistical properties of the CMB temperature anisotropies can be then described in terms of the spectral functions, like the spectrum, bispectrum, trispectrum, etc., of the primordial curvature perturbation $\\zeta$. This spectral functions are given in terms of other quantities, which have an observational value or an uppper bound. For example, the spectrum $\\pz$ is parametrized in terms of an amplitude $\\calpz\\half$, a spectral index $\\nz$ and the level of statistical anisotropy $\\gz$; the bispectrum $\\bz$ and trispectrum $\\tz$ are parametrized in terms of products of the spectrum $\\pz$ and the quantities $\\fnl$, and $\\tnl$ and $\\gnl$, respectively. As we will see in the next chapter, the statistical descriptors $\\fnl$, $\\tnl$ and $\\gnl$ are usually called levels of non-gaussianity, because non zero values for these quantities imply non-gaussianity in the primordial curvature perturbation $\\zeta$ as well in the constrast in the temperature of the CMB radiation $\\delta T/T$. The non-gaussian characteristics in the CMB are actually present in the observation \\cite{wmap5} as we will see in more detail in Section \\ref{observational}. The status of observation can be summarized as follows\\footnote{We are using values according of the five year of data from NASA's WMAP satellite \\cite{wmap5}}: the spectral amplitude $\\calpz\\half=(4.957\\pm 0.094)\\times 10^{-5}$ \\cite{bunn}, the spectral index $\\nz=0.960\\pm 0.014$ at $2 \\sigma$ \\cite{wmap5}, the level of non-gaussianity $\\fnl$ in the bispectrum is in the range $-9 < f_{NL} < 111$ at $2\\sigma$ \\cite{wmap5}; and there is no observational bound on the levels of non-gaussinity $\\tnl$ and $\\gnl$ in the trispectrum $\\tz$. The amount of statistical anisotropy $\\gz$ in the spectrum $\\pz$ is in the range $\\gz \\simeq 0.290 \\pm 0.093$ \\cite{gawe}.  Regarding the statistical descriptors, non-gaussianity in the primordial curvature perturbation $\\zeta$ is one of the subjects of more interest in modern cosmology, because the non-gaussianity parameters $\\fnl$ and $\\tnl$ together with the spectrum amplitude $A_\\zeta$ and spectral index $n_\\zeta$ allow us to discriminate between the different models proposed for the origin of the large-scale structure (see for example Refs. \\cite{alabidi3,alabidi1,alabidi2}). The most studied and popular models are those called the slow-roll models with canonical kinetic terms, because of their simplicity and because they easily satisfy the spectral index $\\nz$ requierements from observation. However, the usual predictions of these models is that the levels of non-gaussianity in the primordial curvature perturbation are expected to be unobservable  \\cite{battefeld,maldacena,seery7,vernizzi,yokoyama1}.  However, as we will show in chapters \\ref{chaptsca} and \\ref{chaptsca2}, there are some aditional issues that have not been taken into account in the current literature. We study these issues to show that it is possible to generate sizeable and observable levels of non-gaussianity in a subclass of small- field {\\it slow-roll} inflationary models with canonical kinetic terms; our main conclussion is that if non-gaussianity is detected, the aforementioned models could have strong possibilities to be the ones responsibles for the formation of the large-scale structure.  According to the usual assumption, one or more of these scalar field perturbations are responsible for the curvature perturbation. In that case, the $n$-point correlators of $\\zeta$ are translationally and rotationally invariants. However, violations of such invariances entail modifications of the usual definitions for the spectral functions in terms of the statistical descriptors \\cite{acw,armendariz,carroll}. These violations may be consequences either of the presence of vector field perturbations \\cite{armendariz,bdmr,vc,vc2,RA2,dklr,dkw,dkw2,go,gmv2,gmv,gvnm,himmetoglu3,himmetoglu,himmetoglu2,himmetoglu4,kksy,dkl,koh,ys}, spinor field perturbations \\cite{bohmer,shan}, or p-form perturbations \\cite{germani,germani2,kobayashi,koivisto,koivisto2}, contributing significantly to $\\zeta$, of anisotropic expansion \\cite{bamba,bohmer,dechant,gcp,himmetoglu4,kksy,koivisto,ppu1,ppu2,watanabe} or of an inhomogeneous background \\cite{armendariz,carroll,dklr}. Violation of the statistical isotropy (i.e. violation of the rotational invariance in the $n$-point correlators of $\\zeta$) seems to be present in the data \\cite{app,ge,hl,samal} and, although its statistical significance is still low, the continuous presence of anomalies in every CMB data analysis (see for instance Refs. \\cite{bunn1,dvorkin,dipole2,dipole1,hansen,dipole3,hoftuft,hou,land1,land2,oliveira,schwarz,tegmark}) suggests the evidence might be decisive in the forthcoming years. The presence of vector fields in the inflationary dynamics is not only important to be responsible of violations of the statistical isotropy, they also may generate sizeable levels of non-gaussianity described by $\\fnl$ and $\\tnl$; particularly we will show in Chapter \\ref{chaptvec}, that including vector fields allows us to get consistency relations between the statistical descriptors, more precisely between the non-gaussianity levels $\\fnl$ and $\\tnl$ and the amount of statistical anisotropy $\\gz$.  Because of the progressive improvement in the accuracy of the satellite measurements described above, it is pertinent to study the statistical descriptors of the primordial curvature perturbation $\\zeta$, for cosmological models of the origin of the large-scale structure in the Universe. It is very important because they could be a crucial tool to discriminate between some of most usual cosmological models \\cite{alabidi1,alabidi2}.  The layout of the thesis is the following: the Chapter \\ref{chaptgen} is devoted to study the statistical descriptors for a probability distribution function and its relation with the observational parameters, i.e., the spectrum amplitude, the spectral index the levels of non-gaussianity $\\fnl$ and $\\tnl$ in the bispectrum $\\bz$ and trispectrum $\\tz$ respectively and the level of statistical anisotropy in the power spectrum, $\\gz$. In this chapter, we also review some generalities of the $\\dn$ formalism, it has become the standard technique to calculate $\\zeta$ and its statistical descriptors. In Chapters \\ref{chaptsca} and \\ref{chaptsca2} we show that it is possible to attain very high, {\\it including observable}, values for the levels of non-gaussianity $f_{NL}$ and $\\tnl$, in a subclass of small-field {\\it slow-roll} models of inflation with canonical kinetic terms. Comparison with current observationally bounds is made. Chapter \\ref{chaptvec} is devoted to study the statistical descriptors of the primordial curvature perturbation $\\zeta$  when scalar and vector fields perturbations are present in the inflationary dynamicc. The levels of non-gaussianity $\\fnl$ and $\\tnl$ are calculated and related to the level of statisitcal anisotropy in the power spectrum, $\\gz$. We show that the levels of non-gaussianity may be very high, in some cases exceeding the current observationally limit. Finally we conclude in Chapter \\ref{chaptconclu}.   \\chapter{$\\dn$ FORMALISM AND STATISTICAL DESCRIPTORS FOR $\\zeta$}\\label{chaptgen}       %    Since COBE \\cite{cobe} discovered and mapped the anisotropies in the temperature of the cosmic microwave background radiation \\cite{smooth}, many balloon and satellite experiments have refined the measurements of such anisotropies, reaching up to now an amazing combined precision. The COBE sequel has continued with the WMAP satellite \\cite{wmap} which has been able to measure the temperature angular power spectrum up to the third peak with unprecedent precision \\cite{hinshaw}, and increase the level of sensitivity to primordial non-gaussianity in the bispectrum by two orders of magnitude compared to COBE \\cite{wmap1,wmap5}.  The next-to-WMAP satellite, PLANCK \\cite{planck}, which was launched in may of 2009, is expected to precisely measure the temperature angular power spectrum up to the eighth peak \\cite{planck1}, and improve the level of sensitivity to primordial non-gaussianity in the bispectrum by one order of magnitude compared to WMAP \\cite{komatsu}.  Because of the progressive improvement in the accuracy of the satellite measurements % described above, it is pertinent to study cosmological inflationary models that generate significant (and observable) levels of non-gaussianity. An interesting way to address the problem involves the $\\delta\\textit{N}$ formalism \\cite{dklr,lms,lr,ss,st,starobinsky}, which can be employed to give the levels of non-gaussianity $f_{NL}$ \\cite{lr} and $\\tau_{NL}$ \\cite{alabidi2,bl} in the bispectrum $B_\\zeta$ and trispectrum $T_\\zeta$ of the primordial curvature perturbation $\\zeta$ respectively. Such non-gaussianity levels are given, for slow-roll inflationary models, in terms of the local evolution of the universe under consideration, as well as of the $n$-point correlators, evaluated a few Hubble times after horizon exit, of the perturbations $\\delta\\phi_{i}$ in the scalar fields that determine the dynamics of such a universe during inflation.  In the $\\delta\\textit{N}$ formalism for slow-roll inflationary models, the primordial curvature perturbation $\\zeta(\\textbf{x},t)$ is written as a Taylor series in the scalar field perturbations $\\delta\\phi_{i}(\\textbf{x},t_\\star)$ evaluated a few Hublle times after horizon exit\\footnote{This equation is similar to one in \\eq{dNsc}, the difference is that here we are redefined it so that $\\langle \\zeta(t,\\bfx)=0\\rangle$.}, \\begin{eqnarray} \\zeta(t,\\textbf{x})&=&\\sum_{I}N_{I}(t)\\delta\\phi_{I}(\\textbf{x},t_\\star) - \\sum_{I}N_{I}(t) \\langle\\delta\\phi_{I} (\\textbf{x},t_\\star)\\rangle + \\nonumber \\\\ &&+\\frac{1}{2}\\sum_{IJ}N_{IJ}(t)\\delta\\phi_{I}(\\textbf{x},t_\\star)\\delta\\phi_{J}(\\textbf{x},t_\\star)-\\frac{1}{2}\\sum_{IJ}N_{IJ}(t) \\langle\\delta\\phi_{I}(\\textbf{x},t_\\star)\\delta\\phi_{J}(\\textbf{x},t_\\star)\\rangle + \\nonumber \\\\ &&+\\frac{1}{3!}\\sum_{IJK}N_{IJK}(t)\\delta\\phi_{I}(\\textbf{x},t_\\star)\\delta\\phi_{J}(\\textbf{x},t_\\star)\\delta\\phi_{k}(\\textbf{x},t_\\star)\\no\\\\& -& \\frac{1}{3!}\\sum_{IJK}N_{IJK}(t)\\langle\\delta\\phi_{I}(\\textbf{x},t_\\star)\\delta\\phi_{J}(\\textbf{x},t_\\star)\\delta\\phi_{K}(\\textbf{x},t_\\star)\\rangle + \\nonumber \\\\ &&+...\\;, \\end{eqnarray} It is in this way that the correlation functions of $\\zeta$ (for instance, $\\langle\\zeta_{\\bf k_{1}}\\zeta_{\\bf k_{2}}\\zeta_{\\bf k_{3}}\\rangle$) can be obtained in terms of series, as often happens in Quantum Field Theory where the probability amplitude is a series whose possible truncation at any desired order is determined by the coupling constants of the theory. A highly relevant question is that of whether the series for $\\delta N$ converges in cosmological perturbation theory and whether it is possible in addition to find some quantities that determine the possible truncation of the series, which in this sense would be analogous to the coupling constants in Quantum Field Theory.  In general such  quantities will depend on the specific inflationary model; the series then cannot be simply truncated at some order until one is sure that it does indeed converge, and besides, one has to be careful not to forget any term that may be leading in the series even if it is of higher order in the coupling constant. This issue has not been investigated in the present literature, and generally the series has been truncated to second- or third-order neglecting in addition terms that could be the leading ones \\cite{alabidi1,alabidi2,battefeld,bl,bsw1,lr,ss,seery3,vernizzi,yokoyama2,yokoyama1,zaballa}.  The most studied and popular inflationary models nowadays are those of the slow-roll variety with canonical kinetic terms \\cite{lyth6,lythbook,lyth5}, because of their simplicity and because they easily satisfy the spectral index requirements for the generation of large-scale structures. One of the usual predictions from inflation and the theory of cosmological perturbations is that the levels of non-gaussianity in the primordial perturbations are expected to be unobservably small when considering this class of models \\cite{battefeld,li,maldacena,seery3,seery5,seery4,seery7,vernizzi,yokoyama1,zaballa}\\footnote{One possible exception is the two-field slow-roll model analyzed in Ref. \\cite{alabidi1} (see also Refs. \\cite{bernardeu2,bernardeu1}) where {\\it observable, of order one, values for} $f_{NL}$ are generated for a reduced window parameter associated with the initial field values when taking into account only the tree-level terms in both $P_\\zeta$ and $B_\\zeta$. However, such a result seems to be incompatible with the general expectation, proved in Ref. \\cite{vernizzi}, of $f_{NL}$ being of order the slow-roll parameters, and {\\it in consequence unobservable}, for two-field slow-roll models with separable potential when considering only the tree-level terms both in $P_\\zeta$ and $B_\\zeta$.  The origin of the discrepancy could be understood by conjecturing that the trajectory in field space, for the models in Refs. \\cite{alabidi1,bernardeu2,bernardeu1}, seems to be sharply curved, being quite near a saddle point; such a condition is required, according to Ref. \\cite{vernizzi}, to generate $f_{NL} \\sim \\mathcal{O}(1)$. \\label{laila}}. This fact leads us to analyze the cosmological perturbations in the framework of first-order cosmological perturbation theory. Non-gaussian characteristics are then suppressed since the non-linearities in the inflaton potential and in the metric perturbations are not taken into account.  The non-gaussian characteristics are actually present and they are made explicit if second-order \\cite{lr1} or higher-order corrections are considered.  The whole literature that encompasses the slow-roll inflationary models with canonical kinetic terms reports that the non- gaussianity level $f_{NL}$ is expected to be very small, being of the order of the slow-roll parameters $\\epsilon_i$ and $\\eta_i$, ($\\epsilon_i, |\\eta_i| \\ll 1$) \\cite{battefeld,maldacena,seery7,vernizzi,yokoyama1}. These works have not taken into account either the convergence of the series for $\\zeta$ nor the possibility that loop corrections dominate over the tree level ones in the $n$-point correlators. Our main result in this chapter is the recognition of the possible convergence of the $\\zeta$ series, and the existence of some ``coupling constants'' that determine the possible truncation of the $\\zeta$ series at any desired order. When this situation is encountered in a subclass of small-field {\\it slow-roll} inflationary models with canonical kinetic terms, the one-loop corrections may dominate the series when calculating either the spectrum $P_\\zeta$, or the bispectrum $B_\\zeta$. This in turn {\\it may generate sizeable and observable levels of non-gaussianity} in total contrast with the general claims found in the present literature.  The layout of the chapter is the following:  Section \\ref{conver} is devoted to the issue of the $\\zeta$ series convergence and loop corrections in the framework of the $\\delta N$ formalism. The presentation of the current knowledge about primordial non-gaussianity in slow-roll inflationary models is given in Section \\ref{ngsr}.  A particular subclass of small-field slow- roll inflationary models is the subject of Section \\ref{model} as it is this subclass of models that generate significant levels of non- gaussianity. The available parameter space for this subclass of models is constrained in Section \\ref{rest} by taking into account some  observational requirements such as the COBE normalisation, the scalar spectral tilt, and the minimal amount of inflation. Another requirement of methodological nature,  the possible tree-level or one-loop dominance in $P_\\zeta$ and/or $B_\\zeta$, is considered in this section. The level of non-gaussianity $f_{NL}$ in the bispectrum $B_\\zeta$ is calculated in Section \\ref{fnl} for models where $\\zeta$ is generated during inflation; a comparison with the current literature is made. Section \\ref{seccou} is devoted to central issues in the consistency of the approach followed such as satisfying necessary conditions for the convergence of the $\\zeta$ series and working in a perturbative regime. Finally in Section \\ref{conclusca} we conclude. As regards the level of non-gaussianity $\\tau_{NL}$ in the trispectrum $T_\\zeta$, it will be studied in the following Chapter.  % The primordial curvature perturbation $\\zeta$ \\cite{dodelson,lythbook,mukhanov,weinberg3}, and its $\\delta N$ expansion\\footnote{By ``$\\delta N$ expansion'' we mean approximating $\\delta N$ by a power series expansion in the initial conditions. By ``$\\delta N$ formula'' we mean the statement that to lowest order in spatial gradients $\\zeta \\equiv \\delta N$. These conventions have been and will be used throughout the text.} \\cite{dklr,lms,lr,ss,st,starobinsky}, was the subject of study in a previous chapter (see also \\cite{cogollo}). We were interested in how well the convergence of the $\\zeta$ series was understood, and if the traditional arguments to cut out the $\\zeta$ series at second order \\cite{lr,zaballa}, keeping only the tree-level terms to study the statistical descriptors of $\\zeta$ \\cite{alabidi1,battefeld,byrnes3,bsw1,seery3,vernizzi,yokoyama1,yokoyama2,yokoyama3}, were reliable\\footnote{We follow the terminology of Ref. \\cite{byrnes1} to identify the tree-level terms and the loop contributions in a diagrammatic approach. The associated diagrams are called {\\it Feynman-like diagrams}.}. We argued that a previous study of the $\\zeta$ series convergence, the viability of a perturbative regime, and the relative weight of the loop contributions against the tree- level terms, were completely necessary and in some cases surprising. For instance, the levels of non-gaussianity $f_{NL}$ and $\\tau_{NL}$ in the bispectrum $B_\\zeta$ and trispectrum $T_\\zeta$ of $\\zeta$ respectively, for slow-roll inflationary models with canonical kinetic terms \\cite{lyth6,lythbook,lyth5}, are usually thought to be of order $\\mathcal{O} (\\epsilon_i,\\eta_i)$ \\cite{battefeld,vernizzi,yokoyama1}\\footnote{See however Refs. \\cite{alabidi1,byrnes3}.} and $\\mathcal{O} (r)$ \\cite{seery3,ssv}\\footnote{See however Refs. \\cite{slri,bch2}.} respectively, were $\\epsilon_i$ and $\\eta_i$ are the slow-roll parameters with $\\epsilon_i,|\\eta_i| \\ll 1$ \\cite{lyth5} and $r$ is the tensor to scalar ratio \\cite{lyth6} with $r < 0.22$ at the $95 \\%$ confidence level \\cite{wmap5}. However, in order to reach such a conclusion, generic models were used where the loop contributions are comparatively suppressed and, therefore, the truncated $\\delta N$ expansion may be used. Of course exceptions may occur, and in those cases it is crucial to check up to what order the truncated $\\delta N$ expansion may be used, and which loop contributions are larger than the tree-level terms.  In any of these cases, general models or exceptions, the question regarding the representation of $\\zeta$ by the $\\delta N$ expansion is a matter to discuss.  Refs. \\cite{alabidi1,byrnes3} show that large, {\\it and observable}, non-gaussianity in $B_\\zeta$ is indeed possible for certain classes of {\\it slow-roll} models with {\\it canonical} kinetic terms and special trajectories in field space, relying only on the tree-level terms.  Ref. \\cite{bch2} does the same for $B_\\zeta$ and $T_\\zeta$ but this time arguing that the loop corrections are always suppressed against the tree-level terms if the quantum fluctuations of the scalar fields do not overwhelm the classical evolution. Nonetheless, although the resultant phenomenology from papers in Refs. \\cite{alabidi1,byrnes3,bch2} is very interesting, the classicality argument used in Ref. \\cite{bch2} is very conservatively stated leading to too strong conclusions as we will argue later in this chapter. More research remains to be done to understand the role of the quantum diffusion and, being this beyond the scope of the present chapter, we will leave the discussion for a future research project. We addressed the $\\zeta$ series convergence and the existence of a perturbative regime in the previous chapter, showing how important the requirements to guarantee those conditions are. Moreover, we showed that for a subclass of small-field {\\it slow-roll} inflationary models with {\\it canonical} kinetic terms, the one-loop correction to $B_\\zeta$ might be much larger than the tree-level terms, giving as a result large, {\\it and observable}, non-gaussianity parameterised by $f_{NL}$.  The present chapter extends the analysis presented in the previous one to $T_\\zeta$ showing, for the first time, that {\\it large and observable} non-gaussianity parameterised by $\\tau_{NL}$ is possible in {\\it slow-roll} inflationary models with {\\it canonical} kinetic terms due to loop effects, in total contrast with the usual belief based on the results of Refs. \\cite{seery3,ssv}. In order to properly identify the non-gaussianity levels found in previous chapter and in the present one with those that are constrained by observation, we comment on the probability that an observer in an ensemble of realizations of the density field in our scenario sees a non-gaussian distribution. As we will show such a probability is non-negligible for the concave downward potential, making indeed the observation of the non-gaussianity studied in this chapter quite possible.  The layout of the chapter is the following: in Section \\ref{model2} we make some aditional comments about ot the slow-roll inflationary model that exhibits large levels of non-gaussianity when loop corrections are considered. This model was described in more detail in Section \\ref{model}. In Section \\ref{class} we study the impact of the quantum fluctions of the scalar fields on their classical evolution.  As a result we argue how the loop suppression proof given in Ref. \\cite{bch2} does not apply to our model. Section \\ref{prob} studies the probability of realizing the scenario proposed in this thesis for a typical observer.  Section \\ref{constraints} is devoted to the reduction of the available parameter window by taking into account some restrictions of methodological and physical nature. The level of non-gaussianity $\\tau_{NL}$ in the trispectrum $T_\\zeta$ is calculated in Section \\ref{taonl} for models where $\\zeta$ is, or is not, generated during inflation; a comparison with the current literature and the results found in the previous chapter for $f_{NL}$ is done.  In Section \\ref{fnlafter} the level of non-gaussianity $f_{NL}$ in the bispectrum $B_\\zeta$ is calculated for models where $\\zeta$ is not generated during inflation.  Finally, Section \\ref{conclusca2} presents the conclusions. %  The anisotropies in the temperature of the cosmic microwave background (CMB) radiation, which have strong connections with the origin of the large-scale structure in the observable Universe, is one of hottest topics in modern cosmology. The properties of the CMB temperature anisotropies are described in terms of the spectral functions, like the spectrum, bispectrum, trispectrum, etc., of the primordial curvature perturbation $\\zeta$ \\cite{cogollo}. In most of the cosmological models the $n$-point correlators of $\\zeta$ are supposed to be translationally and rotationally invariants. However, violations of such invariances entail modifications of the usual definitions for the spectral functions in terms of the statistical descriptors \\cite{acw,armendariz,carroll}. These violations may be consequences either of the presence of vector field perturbations \\cite{armendariz,bdmr,vc,vc2,RA2,dklr,dkw,dkw2,go,gmv2,gmv,gvnm,himmetoglu3,himmetoglu,himmetoglu2,himmetoglu4,kksy,dkl,koh,ys}, spinor field perturbations \\cite{bohmer,shan}, or p-form perturbations \\cite{germani,germani2,kobayashi,koivisto,koivisto2}, contributing significantly to $\\zeta$, of anisotropic expansion \\cite{bamba,bohmer,dechant,gcp,himmetoglu4,kksy,koivisto,ppu1,ppu2,watanabe} or of an inhomogeneous background \\cite{armendariz,carroll,dklr}. Violation of the statistical isotropy (i.e. violation of the rotational invariance in the $n$-point correlators of $\\zeta$) seems to be present in the data \\cite{app,ge,hl,samal} and, although its statistical significance is still low, the continuous presence of anomalies in every CMB data analysis (see for instance Refs. \\cite{bunn1,dvorkin,dipole2,dipole1,hansen,dipole3,hoftuft,hou,land1,land2,oliveira,schwarz,tegmark}) suggests the evidence might be decisive in the forthcoming years. Since the statistical anisotropy is observationally low, it entails a big problem when vector fields are present during inflation, because they generically lead to a high amount of statistical anisotropy, higher than that coming from observations \\cite{dklr,gmv,kksy}. To solve this problem, people use different mechanisms in order to make those models consistent with observation, for example using a triad of orthogonal vectors \\cite{armendariz,bento}, a large number of identical randomly oriented vectors fields \\cite{gmv}, or assuming that the contribution of vector fields to the total energy density is negligible \\cite{dklr,kksy}.  The amount of statistical anisotropy is quantified throught the parameter $g_\\zeta$, usually called the level of statistical anisotropy in the spectrum. \\eq{astadef} gives us the primordial power spectrum that takes into account the leading effects of violations of statistical isotropy by the presence of some vector field in the inflationary era. As we could see in Section \\ref{observational} the $\\gz$ parameter has observational bounds and works, together with the non-gaussianity parameters $\\fnl$, $\\tnl$, $\\gnl$, etc., as statistical descriptors for $\\zeta$. Therefore, it could be a crucial tool to discriminate between some of the more usual cosmological models.  Recent works point out the possibility that a vector field causes part of the primordial curvature perturbation and show that the particular presence of vector fields in the inflationary dynamics may generate sizeable levels of non-gaussianity described by $\\fnl$ \\cite{bdmr1,dkl,vrl} and $\\tnl$ \\cite{bdmr2,vr}. In such works the authors included both vector and scalar field perturbations, and asummed that the contributions to the spectrum from vector field perturbations were smaller than those coming from scalar fields and in an opposite way for bispectrum and trispectrum.  In this chapter we use the $\\dn$ formalism to calculate the tree-level and one-loop contributions to the bispectrum $\\bz$ and trispectrum $\\tz$ of $\\zeta$, including vector and scalar field perturbations. We then calculate the order of magnitude of the levels of non-gaussianity in $\\bz$ and $\\tz$ including the one-loop contributions and write down formulas that relate the order of magnitude of the levels of non-gaussianity $\\fnl$ and $\\tnl$ with the amount of statistical anisotropy in the spectrum $\\gz$. Finally, comparison with the expected observational bound from WMAP is done. ",
        "conclusions": "% Observational cosmology is in its golden age: current satellite and balloon experiments are working extremely well \\cite{hinshaw,wmap}, dramatically improving the quality of data \\cite{wmap5}. Moreover, foreseen experiments \\cite{planck,planck1} will take the field to a state of unprecedent precission where theoretical models will be subjected to the most demanding tests. Given such a state of affairs, it is essential to study the higher order statistical descriptors for cosmological quantities such as the primordial curvature perturbation $\\zeta$, which give us information about the non-gaussianity and about the possible violations of statistical isotropy in their corresponding probability distribution functions.  The slow-roll class of inflationary models with canonical kinetic terms are the most popular and studied to date. Inflationary models of the slow-roll variety predict very well the spectral index in the spectrum $P_\\zeta$ of $\\zeta$ but, if the kinetic terms are canonical, they seem to generate unobservable levels of non-gaussianity in the bispectrum $B_\\zeta$ and the trispectrum $T_\\zeta$ of $\\zeta$ making them impossible to test against the astonishing forthcoming data. Where does this conclusion come from? The answer relies on careful calculations of the levels of non-gaussianity $f_{NL}$ and $\\tau_{NL}$ by making use of the $\\delta N$ formalism \\cite{battefeld,seery3,vernizzi,yokoyama1}. In this framework, $\\zeta$ is given in terms of the perturbation $\\delta N$ in the amount of expansion from the time the cosmologically relevant scales exit the horizon until the time at which one wishes to calculate $\\zeta$.  Due to the functional dependence of the amount of expansion, $\\zeta$ is usually Taylor-expanded (see Eq. (\\ref{Nexp})) and truncated up to some desired order so that $f_{NL}$ and $\\tau_{NL}$ are easily calculated (see for instance Eq. (\\ref{fdnf})). Two key questions arise when noting that it is impossible to extract general and useful information from the $\\zeta$ series expansion in Eq. (\\ref{Nexp}) until one chooses a definite inflationary model and calculates explicitly the $N$ derivatives. First of all, when writing a general expression for $f_{NL}$ or $\\tau_{NL}$ in terms of the $N$ derivatives, how do we know that such an expression is correct if the series convergence has not been examined? Moreover, if the convergence radius of the $\\zeta$ series is already known, why is each term is the $\\zeta$ series supposed to be smaller than the previous one so that cutting the series at any desired order is thought to be enough to keep the leading terms? Nobody seems to have formulated these questions before and, by following a naive line of thinking, $f_{NL}$ and $\\tau_{NL}$ were calculated for slow-roll inflationary models with canonical kinetic terms without checking the $\\zeta$ series convergence and keeping only the presumably leading tree-level terms \\cite{battefeld,maldacena,seery3,seery7,vernizzi,yokoyama1}.  These two questions have been addressed in this thesis (see Chapt. 3 and 4) by paying attention to a particular quadratic small-field slow-roll model of inflation with two components and canonical kinetic terms (see Eq. (\\ref{pot})). Although the non-diagrammatic approach followed in Section \\ref{seccou} to find the necessary condition for the convergence of the $\\zeta$ series in our model might not be applicable to all the cases, we have been able to show that not being careful enough when choosing the right available parameter space could make the $\\zeta$ series, and therefore the calculation of $f_{NL}$ and $\\tau_{NL}$ from the truncated series (e.g. Eq. (\\ref{fdnf})), meaningless. We also have been able to show in our model that the one-loop terms in the spectrum $P_\\zeta$, the bispectrum $B_\\zeta$ and trispectrum $\\tz$ of $\\zeta$ could be bigger or lower than the corresponding tree-level terms, but are always much bigger than the corresponding terms whose order is higher than the one-loop order. If $B_\\zeta$ is dominated by the one-loop correction but $P_\\zeta$ is dominated by the tree-level term, {\\it sizeable and observable values for} $f_{NL}$ {\\it are generated}, so they can be tested against present and forthcoming observational data, a similar conclusion was reached when the trispectrum is dominated by one-loop corrections and the $P_\\zeta$ is dominated by the tree- level term. Finally, if both $P_\\zeta$ and $B_\\zeta$ or $\\tz$ are dominated by the tree-level terms, $f_{NL}$ or $\\tnl$ {\\it are slow-roll suppressed} (see Eqs. \\ref{fnlslowroll} and \\ref{taonlslowroll}) as was originally predicted in Refs. \\cite{battefeld,vernizzi,yokoyama1}. What these results teach us is that the issue of the $\\zeta$ series convergence and loop corrections is essential for making correct predictions about the statistical descriptors of $\\zeta$ in the framework of the $\\delta N$ formalism, and promising for finding high levels of non-gaussianity that can be compared with observations.  The above disccusion about $\\zeta$ was made assuming that the $n$-point correlators of $\\zeta$ are tranlationally and rotationally invariant. However as we could see in the section \\ref{obsvec}, violations of the translational (rotational) invariance (i.e. violations of the statistical homogeneity (isotropy)) seem to be present in the data \\cite{dipole2,dipole1,hansen,dipole3,hoftuft,hou} (\\cite{app,gawe,ge,hl,samal}); therefore it is pertinent to study theoretical models that include those violations. This is the reason why in the chapter \\ref{chaptvec} we studied the statistical descriptors for $\\zeta$ for models with vector field perturbations, which are responsible of violations of statistcal isotropy. We studied in that chapter the order of magnitude of the levels of non-gaussianity $\\fnl$ and $\\tnl$ in the bispectrum $\\bz$ and in the trispectrum $\\tz$, when statistical anisotropy is generated by the presence of one massive vector field. We have shown that it is possible to get an upper bound on the order of magnitude of $\\fnl$ (see \\eq{fnlm5}) and $\\tnl$ (see \\eq{tnl15}) if we assume that the tree-level contributions dominate over all higher order terms in both the vector field spectrum ($\\calp_{\\zeta_A}$), the bispectrum ($\\calbz_A$) and trispectrum ($\\calt_{\\zeta_A}$). We also show that it is possible to get high levels of non-gaussianity $\\fnl$ and $\\tnl$, easily exceeding the expected observational bounds from WMAP, if we assume that the one-loop contributions dominate over the tree-level terms in both the vector field spectrum ($\\calp_{\\zeta_A}$) and the bispectrum ($\\calb_{\\zeta_A}$) or in both the vector field spectrum ($\\calp_{\\zeta_A}$) and the trispectrum ($\\calt_{\\zeta_A}$). We could see that there are a consistency relations between the order of magnitude of $\\fnl$ and the amount of statistical anisotropy in the spectrum $\\gz$ [\\eq{fnl13}] and between the order of magnitude $\\tnl$ and $\\gz$ [\\eq{tnl23}]. Two other consistency relations are given by Eqs. (\\ref{fnlv}) and (\\ref{tnlfnl}), this time relating the order of magnitude of the non-gaussianity parameter $\\fnl$ in the bispectrum $\\bz$ with the amount of statistical anisotropy $g_\\zeta$ and the order of magnitude of the level of non-gaussianity $\\tnl$ in the trispectrum $\\tz$.  Such consistency relations let us fix two of the three parameters by knowing about the other one, i.e. if the non-gaussianity in the bispectrum (or trispectrum) is detected and our scenario is appropriate, the amount of statistical anisotropy in the power spectrum and the order of magnitude of the non-gaussianity parameter $\\tnl$ (or $\\fnl$) must have specific values, which are given by Eqs. (\\ref{fnlv}) (or (\\ref{tnl23})) and (\\ref{tnlfnl}). A similar conclusion is reached if the statistical anisotropy in the power spectrum is detected before the non-gaussianity in the bispectrum or the trispectrum is.  \\appendix \\chapter{TREE-LEVEL AND ONE-LOOP DIAGRAMS FOR $P_\\zeta$, $B_\\zeta$ AND $\\tz$ : SCALAR FIELDS} \\label{app}           % We show in this appendix the mathematical expressions for the tree-level and one-loop Feynman-like diagrams associated with the spectrum $P_\\zeta$, the bispectrum $B_\\zeta$ an trispectrum of $\\zeta$, following the set of rules presented in Ref. \\cite{byrnes1}. To this end we have taken into account the $N$ derivatives for our small-field slow-roll model given in Eqs. (\\ref{1d}), (\\ref{2d}), and (\\ref{3d}).  After presenting the mathematical expressions, we will estimate the order of magnitude of each diagram in order to determine the respective leading terms at tree-level and one-loop for both $P_\\zeta$ and $B_\\zeta$."
    },
    "1004/1004.2858_arXiv.txt": {
        "abstract": "{Blue compact dwarf (BCD) galaxies are low-luminosity, low-metal content dwarf systems undergoing violent bursts of star formation. They present a unique opportunity to probe galaxy formation and evolution and to investigate the process of star formation in a relatively simple scenario. Spectrophotometric studies of BCDs are essential to disentangle and characterize their stellar populations.} {We perform integral field spectroscopy of a sample of BCDs with the aim of analyzing their morphology, the spatial distribution of some of their physical properties (excitation, extinction, and electron density) and their relationship with the distribution and evolutionary state of the stellar populations.} {Integral field spectroscopy observations of the sample galaxies were carried out with the Potsdam  Multi-Aperture Spectrophotometer (PMAS) at the 3.5 m telescope at Calar Alto Observatory.  An area $16\\arcsec\\times 16\\arcsec$ in size was mapped with a spatial sampling of  $1\\arcsec\\times 1\\arcsec$.  We obtained data  in the 3590--6996 \\AA\\ spectral range, with a linear dispersion of 3.2 \\AA\\  per pixel. \\textnormal{From these data we built two-dimensional maps of the flux of the most prominent emission lines, of two continuum bands, of the most relevant line ratios, and of the gas velocity field. Integrated spectra of the most prominent star-forming regions and of whole objects within the FOV were used to derive their physical parameters and the gas metal abundances.} } {Six galaxies display the same morphology both in emission line and in continuum maps; only in two objects, Mrk~32 and Tololo~1434+032, the distributions of the ionized gas and of the stars differ considerably. In general the different excitation maps for a same object display the same pattern and trace the star-forming regions, as expected for objects ionized by hot stars; only the outer regions of Mrk~32, I~Zw~123 and I~Zw~159 display higher [\\ion{S}{ii}]/\\Ha\\ values, suggestive of shocks. Six galaxies display an inhomogeneous dust distribution. Regarding the kinematics, Mrk~750, Mrk~206 and I~Zw~159 display a clear rotation pattern, while in Mrk~32, Mrk~475 and I~Zw~123 the velocity fields are flat.} {} ",
        "introduction": "\\label{Section:Introduction}   Blue compact dwarf (BCD) galaxies are narrow emission-line objects, which undergo at the present time violent bursts of star formation \\citep{Sargent1970}. They are compact and low-luminosity objects (starburst diameter $\\leq 1$ kpc; $M_{B} \\geq -18$ mag), with a low-metal content ($Z_{\\sun}/50 \\leq Z \\leq Z_{\\sun}/2$) and high star-forming (SF) rates, able to exhaust their gas content on a time scale much shorter than the age of the Universe. Initially it was hypothesized that BCDs were truly young galaxies, forming their first generation of stars \\citep{Sargent1970, Lequeux1980,Kunth1988}, but the subsequent detection of an extended redder stellar host galaxy in the vast majority of them has shown that most BCDs are actually old systems \\citep{Loose1987,Telles1995,Papaderos1996a,Cairos2001II, Cairos2001I,Cairos2002,Cairos2003} undergoing recurrent star-formation episodes \\citep{Thuan1991,MasHesse1999}.   These galaxies present a unique opportunity to gain insights on central issues in contemporary galaxy research. Chemically unevolved nearby SF dwarfs like BCDs are an important link to the early Universe and the epoch of galaxy formation, as they have been regarded as the local counterparts of the distant subgalactic units (building blocks) from which larger systems are created at high {\\em redshifts} \\citep{Kauffmann1993,Lowenthal1997}; the study of these systems hence provides important insights into the star-formation process of distant galaxies. Moreover, even though most BCDs are not genuinely young galaxies, their metal deficiency makes them useful objects to constrain the primordial $^{4}$He abundance and to monitor the synthesis and dispersal of heavy elements in a nearly pristine environment \\citep{Pagel1992,Masegosa1994,Izotov1997,Kunth2000}. Blue compact dwarfs are also ideal laboratories for the study of the starburst phenomenon: as they are smaller and less massive than normal galaxies, they cannot sustain  a spiral density wave and do not suffer from disk instabilities, which considerably simplifies the study of the star formation process. Besides, the radiation emitted by their SF regions is less diluted by the stellar  continuum than in giant spiral galaxies, allowing for more precise  studies of element abundance ratios.  However, and in spite of the great effort done during the last two decades on the field of BCDs, fundamental questions like the mechanisms responsible for the ignition of their starburst, their evolutionary status or their SF histories are still far from well understood.  To answer these questions it is of paramount importance to first disentangle and characterize the different components that make up a BCD galaxy. This is a demanding and difficult task. At any location in the galaxy, the emitted flux is the sum of the emission from the local starburst, the flux produced by the nebula surrounding the young stars, and the emission from the underlying, old stellar population, all possibly modulated by dust \\citep{Cairos2002,Cairos2003,Cairos2007}. Substantial work in the field has shown that photometry alone does not allow us to distinguish the different components in BCDs (see \\citealp{Kunth2000,Cairos2002}). The properties of the SF knots in the same galaxy may vary widely: accounting for the flux in emission lines through broad-band filters and for the contribution of the stellar host is fundamental to derive the actual broad-band colors of the knots \\citep{Cairos2002,Cairos2007}. On the other hand, the dust content (usually assumed to be negligible in BCDs) turned out to be quite significant in several objects \\citep{Hunt2001,Cairos2003,VanziSauvage2004}. \\commentout{We think the correction is wrong, as \"accounting\" here   is the subject of \"is fundamental\". We rewrote as: \"... may vary widely: accounting for the flux ...\"}  The few spectrophotometric studies performed so far have shown indeed that they are the right way to tackle the problem: combining high resolution broad- and narrow-band images with high-quality spatially resolved spectra does allow us to distinguish the young stars from the older stars, derive the history of the SF knots and constrain the evolutionary status of BCDs \\citep{Cairos2002,Guseva2003SBS1129,Guseva2003HS1442,Guseva2003SBS1415,Cairos2007}. That very few spectrophotometric analyses can be found in the literature, and virtually all of them focused on one single object, is essentially due to the large amount of observing time that conventional observational techniques require. Acquiring images in several broad-band and narrow-band filters, plus a sequence of long-slit spectra sweeping the region of interest translates into observing times of two or more nights per galaxy.  Thus comprehensive analysis of a statistically meaningful sample of BCDs based on traditional imaging and spectroscopic techniques are in terms of observing time just not feasible. Moreover, these observations usually suffer from varying instrumental and atmospheric conditions, which makes combining all these data complicated. Long-slit spectroscopy has also the additional problem of the uncertainty on the exact location of the slit.   On the other hand, it has been recently shown \\citep{Izotov2006,GarciaLorenzo2008,Kehrig2008,Vanzi2008,Lagos2009,James2009} that the state-of-the-art observational technique of integral field spectroscopy (IFS) offers an alternative way to approach BCDs studies in a highly effective manner. IFS provides simultaneous spectra of each spatial resolution element under identical instrumental and atmospheric conditions. This is not only a more efficient way of observing,  but it also guarantees the homogeneity of the dataset. In terms of observing time, IFS observations of BCDs are one order of magnitude more efficient than traditional observing techniques. This implies that now, for the first time, spectrophotometric studies of substantial samples of BCD galaxies have become feasible.  Consequently, we have undertaken a long-term project, which aims to map an extensive and representative sample of BCDs by means of IFS. This galaxy sample, composed of about 40 objects, has been chosen so as to span the large range in luminosities and morphologies found among the galaxies classified as BCDs. The analysis of such a dataset will allow us to get insights into basic questions of BCDs research, i.e. how to effectively disentangle the old and young stellar populations, set constraints of the age and SF history of the galaxies, study the triggering and propagation mechanisms of the star formation and investigate the metal abundance patterns.   In the first two papers of this series \\citep{Cairos2009Mrk409, Cairos2009Mrk1418}, we illustrated the full potential of this study by showing results on two representatives BCDs, Mrk~1418 and Mrk~409, both observed with the Potsdam multi-aperture spectrophotometer (PMAS), attached at the 3.5m telescope at Calar Alto Observatory. In this paper, the remaining objects observed with PMAS are studied. The whole sample will be analyzed in a series of future publications.   This paper is structured as follows: In Sect.~\\ref{Section:Observations} we describe the observations, the data reduction process and the method employed to build the maps. In Sect.~\\ref{Section:Results} we present the main results of the work, that is, the flux, emission line and velocity maps, as well as the results derived from the analysis of the integrated spectra of the selected galaxy regions. These results are discussed in Sect.~\\ref{Section:Discussion} and summarized in Sect.~\\ref{Section:Conclusions}. ",
        "conclusions": "\\label{Section:Conclusions}  We present here what is to our knowledge the most extensive IFS analysis of a sample of BCDs. This study is based on PMAS data, which cover a wavelength range of 3590-6996 \\AA, with a linear dispersion of 3.2 \\AA per pixel, and map an area $16\\arcsec\\times 16\\arcsec$ with a spatial sampling of $1\\arcsec\\times 1\\arcsec$.  For all the sample galaxies we produced an atlas of two-dimensional maps: two continuum bands, the brightest emission lines (i.e. [\\ion{O}{ii}]~$\\lambda3727$,  \\Hb\\ , [\\ion{O}{iii}]~$\\lambda5007$, [\\ion{O}{i}]~$\\lambda6300$, \\Ha, [\\ion{N}{ii}]~$\\lambda6584$ and [\\ion{S}{ii}]~$\\lambda\\lambda6717,\\;6731$) and the most relevant line ratios (i.e. [\\ion{O}{iii}]/\\Hb, [\\ion{O}{i}]/\\Ha, [\\ion{N}{ii}]/\\Ha, [\\ion{S}{ii}]/\\Ha\\ and \\Ha/\\Hb) as well as the velocity field of the ionized gas. Integrated spectroscopic properties of the most prominent SF regions and of the whole galaxy have been also derived.   From this work we highlight the following results:  \\begin{enumerate}  \\item All the objects except Mrk~750 and Tololo~1434+032 exhibit a mostly regular morphology in the continuum, with one (or several for Mrk~32 and Tololo 1434+032) central SF regions placed atop a more extended host galaxy. The galaxy Mrk~750 reveals elongated outer isophotes, and Tololo~1434+032 displays a clumpy continuum morphology.  All the galaxies show a similar morphology in the different mapped emission lines, as expected for objects ionized by hot stars, and for most of the galaxies the emission line morphology traces also the stellar component. \\textnormal{Only for Mrk~32 and Tololo~1434+032 we found that the distribution of the gaseous emission differs considerably from that of the stellar component. Spatial discrepancies in the distribution of emission lines and continuum are interpreted as signs of a spatial migration of the SF over the history of the galaxies \\citep{Petrosian2002}. However, small spatial offsets between continuum and emission line peaks, as those seen in Tololo~1434+032, and which are indeed a common feature in compact starburst galaxies \\citep{ HunterThronson1995,MaizApellaniz1998,Lagos2007}, are likely related to the release of kinetic energy by massive stars and supernova explosions.}  \\item   The different excitation maps produced for the same galaxies display a similar pattern and trace the regions of star formation as expected in objects ionized by hot stars. In three out of the eight sample galaxies, namely Mrk~32, I~Zw~123 and I~Zw~159, higher values of [\\ion{S}{ii}]/\\Ha\\  in the outer galaxy regions suggest shocks.   \\item  \\textnormal{Six out of the eight objects display inhomogeneous extinction maps, with interstellar reddening values \\EBV\\ varying across the galaxy from $\\leq0.1$ up to 0.7. This result stresses the importance of performing a bidimensional study of the interstellar extinction even when dealing with the less luminous and more compact BCDs as those studied here. Assuming a single, spatially constant value for the extinction, as is usually done in long-slit or single-aperture spectroscopic studies, can lead to large errors in the derivation of fluxes and magnitudes in the different regions of the galaxy.}     \\item  All SF regions in the sample galaxies have low electron densities, ranging from $\\leq100$ to 320 cm$^{-3}$, typical of classical \\ion{H}{ii} regions.  \\item The oxygen abundances in the present objects range from $12+\\log(\\mathrm{O/H})=7.56$ to 8.44 ($Z=1/13Z_{\\sun}$ to $Z=0.6Z_{\\sun}$). We measured for the first time the oxygen abundances of Mrk~407 and Mrk~32. The galaxy Mrk~407 is found to be a relatively high metallicity BCD, while the oxygen abundance found for Mrk~32 \\textnormal{from the [\\ion{O}{iii}]~$\\lambda4363$ line flux would place} it in the list of extremely metal-poor galaxies. These systems, with $12+\\log(\\mathrm{O/H})\\leq7.6$, are  excellent laboratories for galaxy formation and evolution studies, as they allow us to study chemical compositions and stellar populations in conditions approaching those of distant protogalactic systems. However, they are also very difficult to find, and at the present time only about ~30 extremely metal-deficient BCDs are known \\citep{Kunth2000,Kniazev2004,Papaderos2008}.     \\item Wolf-Rayet features were measured in three out of the eight galaxies; a marginal detection was reported for Mrk~32.  \\item Three galaxies display a clear rotation pattern (Mrk~750, Mrk~206, I~Zw~159); for Mrk~407 and Tololo~1434+032, although the maps are noisier, both seem to indicate a low amplitude rotation around a preferred axis. For Mrk~32, Mrk~475 and I~Zw~123 the velocity fields are nearly flat.  \\end{enumerate}   This paper is part of a larger project that aims to map of the properties of an externsive and representative sample of BCDs by means of IFS. Results for five luminous BCDs were published in \\cite{GarciaLorenzo2008}, and results for the galaxies Mrk~409 and Mrk~1418, also observed with PMAS, have been shown in \\cite{Cairos2009Mrk409} and \\cite{Cairos2009Mrk1418} respectively. The global properties of the whole sample will be discussed in a forthcoming publication."
    },
    "1004/1004.5155_arXiv.txt": {
        "abstract": "We present a vector formulation of an interferometric observation of a star, including the effects of the barycentric motion of the observatory, the proper motions of the star, and the reflex motions of the star due to orbiting planets.  We use this model to empirically determine the magnitude and form of the signal due to a single Earth-mass planet orbiting about a sun-mass star.  Using bounding values for the known components of the model, we perform a series of expansions, comparing the residuals to this signal.  We demonstrate why commonly used first order linearizations of similar measurement models are insufficient for signals of the magnitude of the one due to an Earth-mass planet, and present a consistent expansion which is linear in the unknown quantities, with residuals multiple orders of magnitude below the Earth-mass planet signal.  We also discuss numerical issues that can arise when simulating or analyzing these measurements. ",
        "introduction": "Much study has been dedicated in recent years to the possibility of using an ultra-precise, space-based interferometer for the purpose of discovering extra-solar planets by their impact on the astrometric positions of their parent stars.  This has become one of the major science areas of the proposed Space Interferometry Mission (SIM) \\citep{sozzetti2002,sozzetti2003,catanzarite2006}, and has also been considered as an application for the European Space Agency's Gaia mission \\citep{casertano1996astrometry}.  Of particular interest to the exoplanet community is the possibility that interferometers capable of sub-$\\mu$as precision can be used to detect the presence of Earth-sized planets in Earth-like orbits---a goal which is many years away from being realized by any of the other currently employed or studied planet-finding methods.  A number of studies have been completed in order to assess the exact planet-finding capabilities of astrometric instruments \\citep{traub2009,casertano2008double,brown2009}.  One byproduct of these studies has been the realization that the classical description of astrometric observations (as described, for instance, in \\citet{green1985spherical}) makes approximations that are suitable only when dealing with levels of precision of 1 mas or higher.  Several more precise descriptions have been published, including a very thorough one in \\citet{konacki2002frequency}, but most of these take the classical approach of separately treating the effects of proper motion, parallax, and the stellar reflex due to companions, with separate expansions of each effect.  Furthermore, when demonstrating analysis techniques, these studies often still only use a first order expansion to simplify the required computations.  While it is possible to achieve the required numerical precision with these approaches, there is an added burden from having to separately consider the expansion of the direction vector and other effects.  We believe that a simpler approach is to linearize a single measurement equation to produce one unified expression.  Here, we derive the exact\\footnote{By exact we do not mean that all possible contributions to the measurement are included; for instance, we have not yet considered relativistic effects.  Rather, we mean that we are formulating the exact form of the nonlinear measurement for the given set of effects included: parallax, proper motion, and stellar reflex.} expression for an astrometric measurement, and then present several expansions to multiple levels of precision.  This exercise is important for two reasons.  First, if one wishes to evaluate an algorithm, it is crucial to ensure that any simulated test data does not contain biases or components not present in the true data stream.  Even if such structures are below the level of other simulated noise sources, they may have an effect on any processing algorithm which makes the assumption of white, gaussian (or pseudo-gaussian) noise.  The added signals will not be random, and, as shown below, may closely resemble the signal sought in planet-finding applications.  For these reasons, we believe that the correct way to simulate astrometric data is to use an exact representation of the physical system being modeled.  This removes the possibility of inadvertently introducing non-random noise sources, or otherwise creating an unfair test for the analysis algorithm.  Second, when analyzing astrometric data, while linearization of the signal is a very useful tool, we must always ensure that such manipulation does not produce a template that is measurably different from the data.  If data is generated using the same linearization as is assumed by the analysis method, and the linearization introduces measurable structure not present in the true signal being simulated, then use of the same linearization in both simulation and analysis does not constitute a fair test of the algorithm.  Therefore, the main focus of this paper will not be a specific analysis technique.   Rather, we seek to  develop an exact formulation of the astrometric measurement so that completely unbiased data can be produced, on which various analysis techniques can be tested, and to examine the simplifications that can safely be made to this exact form. ",
        "conclusions": "In this treatment, we have presented an exact vector formulation of a narrow-angle astrometric observation incorporating the effects of parallax, stellar reflex, and barycenter motion.  Because the derivation leading to equations (\\ref{eq:rhat_ssc}) and (\\ref{eq:rhat_norm}) does not represent a significant computational burden for modern computers, we believe that they (or something fairly similar) should be used for the simulation of data for the purpose of testing analysis techniques, rather than any linearization of any order of precision.  While it may be argued that using a simplified data set (i.e., one derived from approximate or linearized descriptions of the true observation) allows for initial testing of an analysis method, which can then be followed up with more refined tests, we do not believe that such experiments are scaleable.  The most important point to keep in mind when using astrometry for planet finding is that the expected signals will be of very low order, and will interact non-linearly with several unknown (or partially known) parameters.  Therefore, testing data produced by a linearization which may itself introduce residual signals close to the order of the planet signal, and which adds the planet signal without modeling its interaction with other terms, does not actually tell us anything about our analysis technique's ability to isolate the planet signal in a real data set.  Another argument which is often put forward is that such considerations are unimportant because the measurement noise will generally be higher than the errors discussed here.  However, most analysis techniques (be they Bayesian inference or nonlinear programming minimization), make some assumptions as to the structure of the noise.  Even if the noise is not described as additive white Gaussian, any autocorellation and non-zero mean components are carefully modeled based on our knowledge of the physics of the observed systems and instrument.  Linearization residuals are not random and often introduce patterns that can be quite similar to planet signatures, especially due to the parallax effect of any sun-orbiting observatory. Use of the exact observation description obviates the need for all such considerations.  Furthermore, the linearization presented in \\S\\ref{sec:expand_apriori} allows us to use a sufficiently precise linear expression for analysis, with only the assumption that we have measurements of parallax and barycenter motion of certain fidelity.  This should be very helpful for any methods reliant on fitting techniques, as the nonlinearities in the second order expansion from \\S\\ref{sec:expand_noapriori} are quite difficult for most least-squares algorithms.  Again, both the functional representation and computational requirement of this linearization are only slightly more demanding than those of the classical first order astrometric equations, and so there appears to be no reason not to use representations of the astrometric observation derived in a way similar to the one shown here."
    },
    "1004/1004.1967_arXiv.txt": {
        "abstract": "{Polycyclic aromatic hydrocarbons (PAHs) produce characteristic infrared emission bands that have been observed in a wide range of astrophysical environments, where carbonaceous material is subjected to ultraviolet (UV) radiation. Although PAHs are expected to form in carbon-rich AGB stars, they have up to now only been observed in binary systems where a hot companion provides a hard radiation field. In this letter, we present low-resolution infrared spectra of four S-type AGB stars, selected from a sample of 90 S-type AGB stars observed with the infrared spectrograph aboard the Spitzer satellite. The spectra of these four stars show the typical infrared features of PAH molecules. We confirm the correlation between the temperature of the central star and the centroid wavelength of the 7.9~\\um feature, present in a wide variety of stars spanning a temperature range from 3\\,000 to 12\\,000~K. Three of four sources presented in this paper extend this relation towards lower temperatures. We argue that the mixture of hydrocarbons we see in these S-stars has a rich aliphatic component. The fourth star, BZ~CMa, deviates from this correlation. Based on the similarity with the evolved binary TU~Tau, we predict that BZ~CMa has a hot companion as well.} ",
        "introduction": "Polycyclic aromatic hydrocarbon (PAH) molecules are large molecules containing carbon and hydrogen atoms. They carry typical infrared (IR) emission features that have been observed in many astrophysical environments \\citep{Tielens2008}. Since these features are generally attributed to the IR fluorescence of ultraviolet-pumped molecules, we expect PAH features to arise from regions where carbon-rich material is exposed to ultraviolet (UV) radiation \\citep{Leger1984, Cohen1985, Puget1989, Allamandola1989}. PAHs play a major role in photoelectric heating processes and the ionization balance of the interstellar material \\citep{Lepp1988}.  Although interstellar PAH molecules are thought to originate in the winds of carbon-rich asymptotic giant branch (AGB) stars, there is little observational support for this idea. In a complete ISO/SWS survey of 50 carbon-rich AGB stars, \\cite{Boersma2006} detected PAH emission in only one source, TU~Tau, a carbon-rich AGB star with a hot A2 companion providing UV radiation. Recently, \\cite{Sloan2007} has detected PAH emission in an R-type carbon-rich giant with a circumstellar disk, which is the likely region where the PAH emission originates. This lack of PAH detections in carbon-rich AGB stars is most likely the result of the high optical opacities of carbon-rich dust: when formed, the PAH molecules are shielded from optical and ultraviolet photons by the opaque carbonaceous circumstellar material \\citep{Jaeger1998}.  S-type stars are objects on the ascent of the AGB. They pass through the phase where the photosphere turns from oxygen-rich to carbon-rich ($\\mathrm{C}/\\mathrm{O} \\approx 1$). Because S-type stars are not yet forming carbon-rich dust, they have less opaque circumstellar shells than their carbon-rich successors \\citep{Jaeger2003}. Nevertheless, due to shock-induced non-equilibrium chemistry effects, S-type stars could display carbon-rich molecules \\citep{Cherchneff2006}. Recent studies of ISO/SWS spectra of S-type AGB stars indeed show the presence of molecules like HCN and C$_2$H$_2$ \\citep{Yang2007, Hony2009}.  We have studied the data obtained with the Spitzer Infrared Spectrograph \\citep[IRS,][]{Houck2004} of a sample of 90 intrinsic\\footnote{Intrinsic S-type stars are not enriched in C or s-process elements by accretion from an evolved companion, but through internal nucleosynthesis} S-type stars (Program ID 30737, P.I. S.\\,Hony). A description of the sample can be found in \\citet{Cami2009}. They present the detection of SiS absorption bands, in a large subset of the sample. In this paper, we present the detection of PAH emission around four cool S-type AGB stars. Three of the four objects extend the known correlation between the centroid wavelength of the 7.9~\\um PAH feature and the stellar effective temperature to lower temperatures and more redshifted features. We argue that the hydrocarbon molecules have a high aliphatic/aromatic ratio. Since AGB stars are important producers of the dust in the interstellar medium, we might be looking at the hydrocarbon mixture, as formed in the current day wind. ",
        "conclusions": "\\label{sec:conclusions} In this paper we present a positive detection and identification of PAH emission in 4 S-type AGB stars. In this small sample we see a clear difference between the strong PAH emission in BZ~CMa, which can be classified as Peeters Class B, and the much weaker PAH features seen in the other sources of Class C. We predict that BZ~CMa is a binary system with a hot, late-A-type companion.  Our data are consistent with the strong correlation found between the centroid wavelength of the 7.9~\\um feature and the temperature of the central star. They extend this correlation towards lower temperatures and more redshifted features. This is consistent with the hypothesis that Class C PAHs are hydrocarbon molecules with a high aliphatic/aromatic content ratio found around stars with weak UV radiation fields. The hydrocarbons around CSS~757, KR~Cam, and CSS~173 thus represent the composition as condensed in the AGB wind, before entering the interstellar medium where harsh UV radiation alters their chemical structure."
    },
    "1004/1004.4545_arXiv.txt": {
        "abstract": "Some established views of the solar magnetic cycle are discussed critically, with focus on two aspects at the core of most models:  the role of convective turbulence, and the role of the `tachocline' at the base of the convection zone. The standard view which treats the solar cycle as a manifestation of the interaction between convection and magnetic fields is shown to be misplaced. The main ingredient of the solar cycle, apart from differential rotation, is instead buoyant instability of the magnetic field itself. This view of the physics of the solar cycle was already established in the 1950s, but has been eclipsed mathematically by mean field turbulence formalisms which make poor contact with observations and have serious theoretical problems. The history of this development in the literature is discussed critically.  The source of the magnetic field of the solar cycle is currently assumed to be located in the `tachocline': the shear zone at the base of the convection zone. While the azimuthal field of the cycle is indeed most likely located at the base of the convection zone, it cannot be powered by the radial shear of the tachocline as assumed in these models, since the radiative interior does not support significant shear stresses. Instead, it must be the powered by the latitudinal gradient in rotation rate in the convection zone, as in early models of the solar cycle. Possible future directions for research are briefly discussed. ",
        "introduction": "For a star to generate a self-sustained magnetic field, it is sufficient that it rotate differentially. This differs from the traditional view of dynamos in stars, which holds that in addition to the shear flow due to differential rotation, a small scale velocity field has to be imposed in order to `close the dynamo cycle', thus enabling a selfsustained field independent of initial conditions. Convection can provide such a velocity field, and in fact convection has become such an integral part of thinking about dynamos in stars that the subject of `stellar magnetic fields' has been almost synonymous with `convective dynamos' for decades (for reviews see e.g. Weiss 1981 - 1997, R\\\"udiger and Hollerbach 2004, Tobias 2005, for recent texts Brandenburg 2009,  Jones et. al. 2009, \\bhl Charbonneau 2005\\ehl). Whether or not such a dynamo process can take place in principle is a separate matter. From the observations it is evident, however,  that it is not the way the solar cycle works. Instead, as I will argue below, the cycle operates on dynamic instability of the magnetic field itself. Convection plays only an indirect role, namely by maintaining the differential rotation of the envelope from which the cycle derives its energy.  \\subsection{Mechanism of the solar cycle as inferred from observations} \\label{obs}  The common ingredient in all dynamo models such as those for the Earth's magnetic field or the solar cycle is the generation of a toroidal (azimuthally directed) field by stretching (`winding-up') of the lines of a poloidal field (e.g. Elsasser 1956). This is `the easy part'. It produces a field that increases in strength linearly with time and is proportional to the imposed initial field. To produce a cyclic, self-sustained field as observed there must be a second step that turns some of the toroidal field into a new poloidal component, which is again wound up, completing a field-amplification cycle that becomes independent of initial conditions. The particular process by which the new poloidal field is generated distinguishes the models from each other. In early models of the solar cycle that were popular in their time (Babcock 1961, 1963, Leighton 1969) observations of the emergence of active regions were used to infer the nature of the process responsible for this key step in the dynamo cycle.  \\begin{figure} \\hfil \\includegraphics[width=0.6\\linewidth, clip]{dbdt.pdf}\\hfil \\caption{\\small Rate of increase of the azimuthal field strength as  a function of heliographic latitude, due to the observed differential rotation acting on an assumed uniform poloidal field.} \\label{dbdt} \\end{figure}  These models proposed that the increasing toroidal field eventually becomes unstable, erupting to the surface to form the observed active regions (Cowling 1953, Elsasser 1956, Babcock 1961, see sketches in Figs.\\ \\ref{sketchemerg}, \\ref{spotsketch}). The equatorward drift of the main zone of activity reflects the latitude dependence of the time taken for the toroidal field to reach the point of buoyant instability (Babcock 1961, 1963). This is illustrated with a simple model in Fig.\\ \\ref{dbdt}. In this sketch, a uniform poloidal field is assumed to be stretched  passively by the latitudinal differential rotation as observed on the surface of the Sun. Helioseismic observations (see review by Howe, 2009) show that this pattern of rotation also holds to a fair approximation inside the convection zone.  The azimuthal field becomes unstable to buoyant rise when a critical strength of $\\sim 10^5$ G is reached (Sch\\\"ussler et al. 1994). This happens first at the latitude where the rate of increase of the field is largest, around a latitude of $60^\\circ$ in the simple model of Fig.\\ \\ref{dbdt}. This agrees with observations (Altrock 2010), though initially only small-scale magnetic activity without sunspots is produced. As time progresses, the field also becomes unstable at lower latitudes, producing an equatorward drift of the zone of activity. For reasons unknown, sunspots form only below a latitude of around 40$^\\circ$. As Fig.\\ \\ref{dbdt} implies, Babcock's model also predicts a poleward propagating branch. Such a branch (but without sunspots)  is actually present on the Sun (the `poleward rush', Leroy \\& Trellis, 1974, Altrock 2010). Its observational status and interpretation are not entirely clear, however.  \\begin{figure} \\hfil\\includegraphics[width=0.9\\linewidth, clip]{sketchemerg.jpg}\\hfil \\caption{\\small Closing of the dynamo cycle by active region emergence. Left: sub-surface field produced by stretching of a poloidal field $ {\\bf B}_{\\rm p}$ by differential rotation (equator rotates faster). Coriolis forces during emergence of a stretch of the field (broken) to the surface causes displacements of the footpoints, observed at the surface as `tilt' of the active regions (circles). At depth, this produces a new poloidal field component of opposite sign.} \\label{sketchemerg} \\end{figure}  The process of emergence of an active region has been studied in great detail for more than a century. A small patch of fragmented magnetic fields with mixed polarities appears and expands as more flux emerges (Fig.\\ \\ref{trilobite}). The surroundings of this patch remain unaffected by this process. The mix of polarities then separates into two clumps, the polarities traveling in opposite directions to their destination, ignoring the convective flows in the region.  \\begin{figure} \\includegraphics[width=0.24\\linewidth, clip]{hinode0.jpg}\\hfil\\includegraphics[width=0.24\\linewidth, clip]{hinode1.jpg}\\hfil\\includegraphics[width=0.24\\linewidth, clip]{hinode2.jpg}\\hfil\\includegraphics[width=0.24\\linewidth, clip]{hinode3.jpg} \\caption{\\small Sequence (time from left to right) showing the emergence of an active region at the solar surface observed with the Hinode satellite. The opposite magnetic polarities (vertical component of the field) are shown in black and white. For a movie of this sequence see {http://science.nasa.gov/headlines/y2007/images/trilobite/Hinode\\_lower.mov} } \\label{trilobite} \\end{figure}  \\begin{figure}[h] \\hfil\\includegraphics[width=0.5\\linewidth, clip]{risingtree.jpg}\\hfil \\caption{\\small `Rising tree' sketch to explain the phenomenology of a active region emergence. (From Zwaan, 1978)} \\label{zwaan} \\end{figure}  This striking behavior is the opposite of diffusion. To force it into a diffusion picture, one would have to reverse the arrow of time. Instead of opposite polarities decaying by diffusing into each other, they segregate out from a mix. The MHD equations are completely symmetric with respect to the sign of the magnetic field, however. There are no flows (no matter how complex) that can separate fields of different signs out of a mixture. \\bhl This rules out a priori all models attempting to explain the formation of sunspots and active regions by turbulent diffusion. For recent such attempts, which actually ignore the observations they are trying to explain, see Kitiashvili  et al. (2010), Brandenburg et al. (2010).  The observations, instead, demonstrate that the orientation and location of the polarities seen in an active region must already be have been present in the initial conditions: in the layers \\ehl below the surface from which the magnetic field traveled to the surface.  The fragmented state near the surface in the early stages of the eruption process is only temporary. The intuitive `rising tree' picture (Zwaan, 1978) illustrates this (Fig.\\ \\ref{zwaan}). The observed fragmentation and subsequent formation of spots from a horizontal strand of magnetic field below the surface has recently been reproduced in striking realism in full 3-D radiative MHD simulations (Cheung et al. 2008, Rempel, this volume).  The axes of active regions are observed to be tilted with respect to the east-west direction. This was  attributed to the action of Coriolis forces during the emergnece process, and identified with the generation of the new poloidal field component that closes the dynamo cycle by Leighton (1969, see sketch in Fig.\\ \\ref{sketchemerg}).  \\begin{figure}[h] \\hfil\\includegraphics[width=0.8\\linewidth, clip]{spotsketch.jpg}\\hfil \\caption{\\small Vertical cut through an active region illustrating the connection between a sunspot at the surface and its origins in the toroidal field layer at the base of the convection zone. (From Spruit and Roberts 1983).} \\label{spotsketch} \\end{figure}  \\subsection{Later developments}  Models like Leighton's thus made a direct connection between observations of the active regions that make up the solar cycle and the functioning of the cycle as a whole. One might have expected that this natural state of affairs would have led to a further development of the theoretical ideas in continued contact with the observations. But this has not been the case.  Instead, the development of these ideas has been eclipsed for several decades by the parallel development of turbulent mean field formalisms for the solar cycle. These ideas postulated mathematically tractable equations which were claimed to represent the physics of the interaction between magnetic fields and convection in some statistical sense. They relied on theoretical assumptions like cascades in wavenumber space, correlation functions to represent the interaction between magnetic fields and flows, and an assumed separation of length scales between mean fields and fluctuations. Just looking at the data as described above, it is difficult see how a separation would be accomplished. What is more, the data themselves already contain more detailed and more critical information on the functioning of the cycle than is present in mean field models.  The dominance of these formalisms in the astrophysical literature (thousands of papers) has led to a particularly sterile theoretical view of the solar cycle, supported neither by a sound theoretical foundation of the equations used nor making much contact with the observations.  In addition, it has had the effect of obscuring an important fact, namely that no turbulence needs to be imposed at all for dynamo action to take place.  A system that is completely laminar in the absence of magnetic fields can produce dynamo action from shear and magnetic instability alone (cf. Spruit 2002). A well studied and very successful example of such a dynamo process is the MRI turbulence observed in numerical simulations of accretion disks (e.g. Hawley et al. 1996). The models by Babcock and Leighton are just  another example where magnetic instability is the key element in closing the field amplification cycle. These kinds of magnetic cycle are intrinsically non-linear (i.e. not `kinematic' in dynamo parlance): their functioning depends on the finite amplitude of the field generated. This is because the time scale of the magnetic instabilities that close the dynamo cycle depends on field strength.  The conditions for self-sustained field generation to occur by differential rotation and instability alone, the properties of the magnetic field produced in this way, and its observable consequences all reflect the nature of the magnetic instability involved. \\bhl In the case of the solar cycle: the properties of magnetic buoyancy. It is sometimes argued that such a process just  brings about an `alpha effect', so that one just has to  use \\ehl a set of equations that incorporate such an effect. Neither the fact that a poloidal field component can appear by a process changing the direction of an initially toroidal field, however, nor the fact that turbulent mean field equations contain a term describing such an effect, are justifications for using these equations. An understanding of the solar cycle, or any other dynamo process, requires physics to be found out first, rather than assumed in some parametrized form. The idea that insight about the solar cycle can be obtained from the solutions of such models has been an impediment to real progress, however tempting the equations may have looked.  Justification for this critical view is found in the history of ideas about the solar cycle; this is done in the following section. I briefly discuss there how mean field thinking has led to a systematic disconnect between theory and observations. In all likelihood this would not have been necessary if the observations and their interpretation in models such as Leighton's (1969), had been taken more serious. ",
        "conclusions": "Observations of active region phenomenology, most of them already old and well-established, show that the solar cycle operates on buoyant instability of the magnetic field itself rather than the conventional view based on interaction with convection. This puts us back to ideas developed half a century ago. Significant steps forward, however are the direct 3-D, radiative numerical MHD simulations which are now beginning to make contact with some of the classical observations. Though these simulations cannot deal with the cycle as a whole, their success in reproducing limited aspects such as the emergence of magnetic flux discussed above, or the observed structure of sunspots (\\bhl Heinemann et al. 2007, Scharmer et al. 2008, Rempel et al. 2009\\ehl) give confidence for the future. At the same time they clean the table by eliminating a number of dead-end views on the solar cycle, some of which considered well-established thus far.  At the same time, a number of unsolved questions appear that are specific for the picture of a magnetic cycle operating on buoyant instability. Some of these questions are unlikely to be answered from first principles or numerical simulations. Clues taken from observations may well play an important role in making progress in figuring out the physics relevant for these questions. As the history of the subject shows, however, taking observational clues serious will require one to jettison the turbulent mean field baggage that has impeded the development of a sensible theory of the solar cycle for so long. This process would be assisted by healthy skepticism on the part of the observational community. In fact, it is rather surprising how easily observers have acquiesced in the past to the treatment of their data by mean field theories (`sorry but your observations are just turbulence, they have to be averaged out').  \\medskip"
    },
    "1004/1004.4897_arXiv.txt": {
        "abstract": "We report new cm-wave measurements at five frequencies between 15 and 18\\,GHz of the continuum emission from the reportedly anomalous ``region 4'' of the nearby galaxy NGC~6946. We find that the emission in this frequency range is significantly in excess of that measured at 8.5\\,GHz, but has a spectrum from 15--18\\,GHz consistent with optically thin free--free emission from a compact {\\sc Hii} region. In combination with previously published data we fit four emission models containing different continuum components using the Bayesian spectrum analysis package {\\tt radiospec}. These fits show that, in combination with data at other frequencies, a model with a spinning dust component is slightly preferred to those that possess better-established emission mechanisms. ",
        "introduction": "The complete characterization of microwave emission from spinning dust grains is a key question in both astrophysics and cosmology as it probes a region of the electromagnetic spectrum where a number of different astrophysical disciplines overlap: it is important for CMB observations in order to correctly characterise the contaminating foreground emission (Gold et~al. 2010); for star and planetary formation it is important because it potentially probes a regime of grain sizes that is not otherwise easily observable (Rafikov 2000).  Although a number of objects have now been found to exhibit anomalous microwave emission, attributed to spinning dust, it is still unclear what differentiates those objects from the many other seemingly similar targets that do not show the excess. In order to investigate this question a number of Galactic observations have been made towards known star formation regions (see e.g. AMI Consortium: Scaife et~al. 2010 and references therein; Casassus et~al. 2008; Tibbs et~al. 2009). In addition Murphy et~al. (2010; hereinafter M10) made the first extra-galactic search for anomalous microwave emission within the star formation regions of the nearby galaxy NGC~6946 using the Caltech Continuum Back-end on the Green Bank Telescope. M10 found a significantly anomalous spectrum in only one of 10 star-forming regions: extra-nuclear region 4 (hereinafter NGC~6946-E4). The excess of emission was seen between 27--38\\,GHz relative to the continuum emission at 8.5\\,GHz measured using combined Effelsberg 100\\,m Telescope and VLA observations (Beck 2007).  In this letter we present follow-up observations of NGC~6946-E4 at frequencies from 15-18\\,GHz using the Arcminute Microkelvin Imager (AMI) Large Array (LA). In Section~\\ref{sec:obs} we present the details of these observations, in Section~\\ref{sec:results} we present the results of the AMI-LA observations and a comparison with other radio data, and in Section~\\ref{sec:conc} we discuss the implications of these results and form our conclusions. In what follows we use the convention: $S\\propto \\nu^{\\alpha}$, where $S$ is flux density, $\\nu$ is frequency and $\\alpha$ is the spectral index. All errors are quoted to 1\\,$\\sigma$. ",
        "conclusions": "\\label{sec:conc} Considered on their own, the AMI-LA data (after correction for flux loss) have a spectral index of $\\alpha_{\\rm{AMI}}=-0.11\\pm0.77$. Although this value is consistent with optically thin free--free emission, the error on this estimate is large and we cannot rule out other mechanisms. Since the spectral index between the Effelsberg-VLA measurement at 8.5\\,GHz and the AMI band is rising ($\\alpha_{8.5}^{16}=0.67\\pm0.08$, see Fig.~\\ref{fig:E4spec}), we need to consider the possibility that region E4 contains one or more compact {\\sc Hii} regions with their opacity reaching unity at approximately 12\\,GHz. Such an opacity would require an emission measure of $\\simeq 5\\times10^8$\\,pc\\,cm$^{-6}$, assuming $T_{\\rm{e}} = 10^4$\\,K, and would be appropriate for a compact {\\sc Hii} region.  We therefore examine two alternative hypotheses for the emission from region E4. The first is that the emission is due to the usual diffuse synchrotron and free--free mechanisms associated with star-formation, with an additional high-opacity free--free component (Hypothesis 1; H1). The second hypothesis is that there is a spinning dust component rather than high-opacity free--free (H2). A summary of how well these two hypotheses fit the observed data is shown in the form of fan-diagrams in Fig.~\\ref{fig:E4spec}. As can be seen from this figure, neither of the hypotheses can be ruled out, although the spinning-dust appears to somewhat better match the data. This is also confirmed by a simple comparison of the models: assuming flat priors and no a priori difference between the models, the logarithmic Bayes factor is $3.7\\pm0.3$ in favour of the spinning dust model. From the Jeffreys' scale of evidence (Jeffreys 1961; Kass \\& Raftery 1995; see e.g. Efron \\& Gous 2001 for further discussion of this scale) this would indicate a weak positive preference for the spinning dust model above the free--free model. The maximum likelihood parameters for these models are listed in Table~\\ref{tab:parms}.  For comparison we have also carried out a similar analysis on the region E8, shown in Fig~\\ref{fig:e8specfan}. In this case the two hypotheses are a simple diffuse synchrotron plus free--free model (H1) versus the same model with an additional spinning dust component (H2). In this case the logarithm of the Bayesian evidence ratio is $0.5\\pm0.3$ in favour of the simpler model without the spinning dust. A ratio of this size indicates no perceptible difference between the two models.  In region E4 where a spinning dust model is the preferred hypothesis, we can marginalise the posterior distribution of the model parameters to obtain an estimate of the gas mass containing the spinning dust, shown as a histogram in the left panel of Fig.~\\ref{fig:sdmass-margin} and in numerical form in Table~\\ref{tab:parms}, i.e., $10^{8.1\\pm0.1}\\,{\\rm M}_\\odot$. In region E8 spinning dust is not the preferred hypothesis but proceeding with this hypothesis anyway, we find an upper limit for the mass of the gas bearing the spinning dust, which is around $10^{7.5}\\,{\\rm M}_\\odot$. We draw a conclusion that if the conditions in E4 and E8 are similar, then the mass of any gas bearing spinning dust in E8 must be at least a factor of five smaller than in E4.  Region E4 is located on the dense rim of a ``remarkable'' {\\sc Hi} hole (Boomsma et~al. 2008) within NGC~6946. Such an association may be relevent to the differentiation of this star formation region from the eight others found to exhibit no anomalous emission by M10. The hole is remarkable for a number of reasons, notably the almost unbroken symmetry of its dense {\\sc Hi} rim, unusual in so large a hole, and the small scale high velocity gas complexes observed in connection with it.  As mentioned above, the spinning dust model is preferred by the evidence calculation for NGC~6946-E4, but not at a very high level. Definitive confirmation of the nature of the emission requires measurements at frequencies above 50\\,GHz where the spinning dust and compact H\\,{\\sc ii} region models have significantly different behaviour. For example, Fig.~\\ref{fig:E4spec} shows that at 100\\,GHz the difference between these two models should be at least a factor of two in brightness.  In the sub-mm there are data available from SCUBA at 850\\,$\\mu$m for this region by \\citealt{2008ApJS..175..277D}, which might be used to constrain the mass of NGC~6946-E4 and place constraints on the frequency at which the optical depth reaches unity. If a compact {\\sc Hii} region is present, with $\\tau=1$ at $\\nu>8.5$\\,GHz it should have a correspondingly high dust mass. From analysis of the SCUBA data we obtained a flux estimate for region E4 of $S_{850} = 11\\pm14$\\,mJy\\,beam$^{-1}$. However, the errors on this estimate are too high to allow any useful constraint on the properties of the thermal dust emission or to calculate a reliable dust mass estimate.  From the existing data the possibility of a spinning dust component in this region cannot be ruled out, but the evidence is not yet definitive. Further observations of this object at frequencies covering the higher frequency minimum between spinning dust emission and thermal dust emission ($\\approx 90$\\,GHz) would be most useful."
    },
    "1004/1004.3962_arXiv.txt": {
        "abstract": "The observable universe could be a 1+3-surface (the ``brane'') embedded in a 1+3+\\textit{d}-dimensional spacetime (the ``bulk''), with Standard Model particles and fields trapped on the brane while gravity is free to access the bulk. At least one of the \\textit{d} extra spatial dimensions could be very large relative to the Planck scale, which lowers the fundamental gravity scale, possibly even down to the electroweak ($\\sim$~TeV) level. This revolutionary picture arises in the framework of recent developments in M~theory. The 1+10-dimensional M~theory encompasses the known 1+9-dimensional superstring theories, and is widely considered to be a promising potential route to quantum gravity. At low energies, gravity is localized at the brane and general relativity is recovered, but at high energies gravity ``leaks'' into the bulk, behaving in a truly higher-dimensional way. This introduces significant changes to gravitational dynamics and perturbations, with interesting and potentially testable implications for high-energy astrophysics, black holes, and cosmology. Brane-world models offer a phenomenological way to test some of the novel predictions and corrections to general relativity that are implied by M~theory. This review analyzes the geometry, dynamics and perturbations of simple brane-world models for cosmology and astrophysics, mainly focusing on warped 5-dimensional brane-worlds based on the Randall--Sundrum models. We also cover the simplest brane-world models in which 4-dimensional gravity on the brane is modified at \\emph{low} energies -- the 5-dimensional Dvali--Gabadadze--Porrati models. Then we discuss co-dimension two branes in 6-dimensional models. ",
        "introduction": "\\label{section_1}  At high enough energies, Einstein's theory of general relativity breaks down, and will be superceded by a quantum gravity theory. The classical singularities predicted by general relativity in gravitational collapse and in the hot big bang will be removed by quantum gravity. But even below the fundamental energy scale that marks the transition to quantum gravity, significant corrections to general relativity will arise. These corrections could have a major impact on the behaviour of gravitational collapse, black holes, and the early universe, and they could leave a trace -- a ``smoking gun'' -- in various observations and experiments. Thus it is important to estimate these corrections and develop tests for detecting them or ruling them out. In this way, quantum gravity can begin to be subject to testing by astrophysical and cosmological observations.  Developing a quantum theory of gravity and a unified theory of all the forces and particles of nature are the two main goals of current work in fundamental physics. There is as yet no generally accepted (pre-)quantum gravity theory. Two of the main contenders are M~theory (for reviews see, e.g., \\cite{mtheory_1, mtheory_2, mtheory_3}) and quantum geometry (loop quantum gravity; for reviews see, e.g., \\cite{loop_1, loop_2}). It is important to explore the astrophysical and cosmological predictions of both these approaches. This review considers only models that arise within the framework of M~theory.  In this review, we focus on RS brane-worlds (mainly the RS 1-brane model) and their generalizations, with the emphasis on geometry and gravitational dynamics (see~\\cite{m2, rev_1, rev_2, rev_3, rev_4, rev_5, rev_6, rev_7, rev_8, rev_9, lan} for previous reviews with a broadly similar approach). Other reviews focus on string-theory aspects, e.g., \\cite{que_1, que_2, que_3, que_4}, or on particle physics aspects, e.g., \\cite{r_1, r_2, r_3, r_4, cav}. We also discuss the 5D DGP models, which modify general relativity at low energies, unlike the RS models; these models have become important examples in cosmology for achieving late-time acceleration of the universe without dark energy. Finally, we give brief overviews of 6D models, in which the brane has co-dimension two, introducing very different features to the 5D case with co-dimension one branes.  \\epubtkUpdateA{Extended a paragraph moved here from the end of Section~\\ref{section_1.2}.}    \\subsection{Heuristics of higher-dimensional gravity}  One of the fundamental aspects of string theory is the need for extra spatial dimensions\\epubtkFootnote{We do not consider timelike extra dimensions: see~\\cite{varun1} for an interesting example.}. This revives the original higher-dimensional ideas of Kaluza and Klein in the 1920s, but in a new context of quantum gravity. An important consequence of extra dimensions is that the  4-dimensional Planck scale $M_\\mathrm{p}\\equiv M_{4}$ is no longer the fundamental scale, which is $M_{4+d}$, where $d$ is the number of extra dimensions. This can be seen from the modification of the gravitational potential. For an Einstein--Hilbert gravitational action we have \\begin{eqnarray} S_\\mathrm{gravity} &=& {1\\over 2\\kappa_{4+d}^2}\\int d^4x\\, d^dy\\,\\sqrt{-^{(4+d)\\!}g} \\left[ {}^{(4+d)\\!}R- 2\\Lambda_{4+d} \\right], \\\\ {}^{(4+d)\\!}G_{AB} & \\equiv & \\;{}^{(4+d)\\!}R_{AB}-{1\\over2} \\;{}^{(4+d)\\!} R \\;{}^{(4+d)\\!}g_{AB} = -\\Lambda_{4+d} \\;{}^{(4+d)\\!}g_{AB}+ \\kappa_{4+d}^2 \\;{}^{(4+d)\\!}T_{AB}, \\label{defe} \\end{eqnarray}% where $X^A=(x^\\mu,y^1, \\dots, y^d)$, and $\\kappa_{4+d}^2$ is the gravitational coupling constant, \\begin{equation} \\kappa_{4+d}^2=8\\pi G_{4+d}={8\\pi\\over M_{4+d}^{2+d}}. \\end{equation} The static weak field limit of the field equations leads to the $4+d$-dimensional Poisson equation, whose solution is the gravitational potential, \\begin{equation} V(r) \\propto {\\kappa_{4+d}^2\\over r^{1+d}}. \\label{v} \\end{equation} If the length scale of the extra dimensions is $L$, then on scales $r\\lesssim L$, the potential is $4+d$-dimensional, $V\\sim r^{-(1+d)}$. By contrast, on scales large relative to $L$, where the extra dimensions do not contribute to variations in the potential, $V$ behaves like a 4-dimensional potential, i.e., $r\\sim L$ in the $d$ extra dimensions, and $V \\sim L^{-d}r^{-1}$. This means that the usual Planck scale becomes an effective coupling constant, describing gravity on scales much larger than the extra dimensions, and related to the fundamental scale via the volume of the extra dimensions: \\begin{equation} M_\\mathrm{p}^2 \\sim M_{4+d}^{2+d}\\,L^d. \\end{equation} If the extra-dimensional volume is Planck scale, i.e., $L\\sim M_\\mathrm{p}^{-1}$, then $M_{4+d}\\sim M_\\mathrm{p}$. But if the extra-dimensional volume is significantly above Planck scale, then the true fundamental scale $M_{4+d}$ can be much less than the effective scale $M_\\mathrm{p} \\sim 10^{19} \\mathrm{\\ GeV}$. In this case, we understand the weakness of gravity as due to the fact that it ``spreads'' into extra dimensions and only a part of it is felt in 4 dimensions.  A lower limit on $M_{4+d}$ is given by null results in table-top experiments to test for deviations from Newton's law in 4 dimensions, $V\\propto r^{-1}$. These experiments currently~\\cite{exp} probe sub-millimetre scales, so that \\begin{equation} L \\lesssim 10^{-1} \\mathrm{\\ mm} \\sim  (10^{-15} \\mathrm{\\ TeV})^{-1} \\quad \\Rightarrow \\quad M_{4+d}\\gtrsim 10^{(32-15d)/(d+2)} \\mathrm{\\ TeV}. \\label{tt} \\end{equation} Stronger bounds for brane-worlds with compact flat extra dimensions can be derived from null results in particle accelerators and in high-energy astrophysics~\\cite{cav, cheung, hanraf_1, hanraf_2}.     \\subsection{Brane-worlds and M~theory} \\label{section_1.2}  String theory thus incorporates the possibility that the fundamental scale is much less than the Planck scale felt in 4 dimensions. There are five distinct 1+9-dimensional superstring theories, all giving quantum theories of gravity. Discoveries in the mid-1990s of duality transformations that relate these superstring theories and the 1+10-dimensional supergravity theory, led to the conjecture that all of these theories arise as different limits of a single theory, which has come to be known as M~theory. The 11th dimension in M~theory is related to the string coupling strength; the size of this dimension grows as the coupling becomes strong. At low energies, M~theory can be approximated by 1+10-dimensional supergravity.  It was also discovered that p-branes, which are extended objects of higher dimension than strings (1-branes), play a fundamental role in the theory. In the weak coupling limit, p-branes ($p>1$) become infinitely heavy, so that they do not appear in the perturbative theory. Of particular importance among p-branes are the D-branes, on which open strings can end. Roughly speaking, open strings, which describe the non-gravitational sector, are attached at their endpoints to branes, while the closed strings of the gravitational sector can move freely in the bulk. Classically, this is realised via the localization of matter and radiation fields on the brane, with gravity propagating in the bulk (see Figure~\\ref{figure_01}).  \\epubtkImage{figure01.png}{% \\begin{figure}[htbp] \\def\\epsfsize#1#2{0.5#1} \\centerline{\\epsfbox{figure01.eps}} \\caption{Schematic of confinement of matter to the brane, while gravity propagates in the bulk (from~\\cite{cav}).} \\label{figure_01} \\end{figure}}  In the Horava--Witten solution~\\cite{hv}, gauge fields of the standard model are confined on two 1+9-branes located at the end points of an $S^1/Z_2$ orbifold, i.e., a circle folded on itself across a diameter. The 6 extra dimensions on the branes are compactified on a very small scale close to the fundamental scale, and their effect on the dynamics is felt through ``moduli'' fields, i.e., 5D scalar fields. A 5D realization of the Horava--Witten theory and the corresponding brane-world cosmology is given in~\\cite{low_1, low_2, low_3}.  \\epubtkImage{figure02.png}{% \\begin{figure}[htbp] \\def\\epsfsize#1#2{0.5#1} \\centerline{\\epsfbox{figure02.eps}} \\caption{The RS 2-brane model. (Figure taken from~\\cite{cheung}.)} \\label{figure_02} \\end{figure}}  These solutions can be thought of as effectively 5-dimensional, with an extra dimension that can be large relative to the fundamental scale. They provide the basis for the Randall--Sundrum (RS) 2-brane models of 5-dimensional gravity~\\cite{rs1} (see Figure~\\ref{figure_02}). The single-brane Randall--Sundrum models~\\cite{rs2} with infinite extra dimension arise when the orbifold radius tends to infinity. The RS models are not the only phenomenological realizations of M~theory ideas. They were preceded by the Arkani--Hamed--Dimopoulos--Dvali (ADD) brane-world models~\\cite{add_1, add_2, add_3, add_4, add_5, add_6, add_7, add_8}, which put forward the idea that a large volume for the compact extra dimensions would lower the fundamental Planck scale, \\begin{equation} M_\\mathrm{ew}\\sim 1 \\mathrm{\\ TeV} \\lesssim M_{4+d} \\leq M_\\mathrm{p} \\sim 10^{16} \\mathrm{\\ TeV}, \\label{scales} \\end{equation} where $M_\\mathrm{ew}$ is the electroweak scale. If $M_{4+d}$ is close to the lower limit in Equation~(\\ref{scales}), then this would address the long-standing ``hierarchy'' problem, i.e., why there is such a large gap between $M_\\mathrm{ew}$ and $M_\\mathrm{p}$.  In the ADD models, more than one extra dimension is required for agreement with experiments, and there is ``democracy'' amongst the equivalent extra dimensions, which are typically flat. By contrast, the RS models have a ``preferred'' extra dimension, with other extra dimensions treated as ignorable (i.e., stabilized except at energies near the fundamental scale). Furthermore, this extra dimension is curved or ``warped'' rather than flat: The bulk is a portion of anti-de Sitter ($\\mathrm{AdS}_5$) spacetime. As in the Horava--Witten solutions, the RS branes are $Z_2$-symmetric (mirror symmetry), and have a tension, which serves to counter the influence of the negative bulk cosmological constant on the brane. This also means that the self-gravity of the branes is incorporated in  the RS models. The novel feature of the RS models compared to previous higher-dimensional models is that the observable 3 dimensions are protected from the large extra dimension (at low energies) by curvature rather than straightforward compactification.  The RS brane-worlds and their generalizations (to include matter on the brane, scalar fields in the bulk, etc.) provide phenomenological models that reflect at least some of the features of M~theory, and that bring exciting new geometric and particle physics ideas into play. The RS models also provide a framework for exploring holographic ideas that have emerged in M~theory. Roughly speaking, holography suggests that higher-dimensional gravitational dynamics may be determined from knowledge of the fields on a lower-dimensional boundary. The AdS/CFT correspondence is an example, in which the classical dynamics of the higher-dimensional gravitational field are equivalent to the quantum dynamics of a conformal field theory (CFT) on the boundary. The RS model with its $\\mathrm{AdS}_5$ metric satisfies this correspondence to lowest perturbative order~\\cite{acft} (see also~\\cite{acftcosmo_1, acftcosmo_2, acftcosmo_8, acftcosmo_3, acftcosmo_4, acftcosmo_5, acftcosmo_6, acftcosmo_7} for the AdS/CFT correspondence in a cosmological context).   Before turning to a more detailed analysis of RS brane-worlds, We discuss the notion of Kaluza--Klein (KK) modes of the graviton.     \\subsection{Heuristics of KK modes}  The dilution of gravity via extra dimensions not only weakens gravity on the brane, it also extends the range of graviton modes felt on the brane beyond the massless mode of 4-dimensional gravity. For simplicity, consider a flat brane with one flat extra dimension, compactified through the identification $y\\leftrightarrow y+2\\pi n L$, where $n=0,1,2,\\dots$. The perturbative 5D graviton amplitude can be Fourier expanded as \\begin{equation} f(x^a,y)=\\sum_n e^{iny/L}\\,f_n(x^a), \\end{equation} where $f_n$ are the amplitudes of the KK modes, i.e., the effective 4D modes of the 5D graviton. To see that these KK modes are massive from the brane viewpoint, we start from the 5D wave equation that the massless 5D field $f$ satisfies (in a suitable gauge): \\begin{equation} {}^{(5)\\!}\\Box f=0 \\quad \\Rightarrow \\quad \\Box f+\\partial_y^2 f=0. \\end{equation} It follows that the KK modes satisfy a 4D Klein--Gordon equation with an effective 4D mass $m_n$, \\begin{equation} \\Box f_n=m_n^2\\,f_n, \\qquad m_n={n\\over L}. \\end{equation} The massless mode $f_0$ is the usual 4D graviton mode. But there is a tower of massive modes, $L^{-1},2L^{-1},\\dots$, which imprint the effect of the 5D gravitational field on the 4D brane. Compactness of the extra dimension leads to discreteness of the spectrum. For an infinite extra dimension, $L\\to\\infty$, the separation between the modes disappears and the tower forms a continuous spectrum. In this case, the coupling of the KK modes to matter must be very weak in order to avoid exciting the lightest massive modes with $m\\gtrsim 0$.  From a geometric viewpoint, the KK modes can also be understood via the fact that the projection of the null graviton 5-momentum ${}^{(5)\\!}p_A$ onto the brane is timelike. If the unit normal to the brane is $n_A$, then the induced metric on the brane is \\begin{equation} g_{AB}= {}^{(5)\\!}g_{AB}-n_An_B, \\qquad {}^{(5)\\!}g_{AB}n^An^B=1, \\qquad g_{AB}n^B=0, \\end{equation} and the 5-momentum may be decomposed as \\begin{equation} {}^{(5)\\!}p_A=mn_A+ p_A, \\qquad p_An^A=0, \\qquad m= {}^{(5)\\!}p_A\\, n^A, \\end{equation} where $p_A=g_{AB} \\;{}^{(5)\\!}p^B$ is the projection along the brane, depending on the orientation of the 5-momentum relative to the brane. The effective 4-momentum of the 5D graviton is thus $p_A$. Expanding ${}^{(5)\\!}g_{AB}{}{} \\;{}^{(5)\\!}p^A \\;{}^{(5)\\!}p^B=0$, we find that \\begin{equation} g_{AB}p^Ap^B=-m^2. \\end{equation} It follows that the 5D graviton has an effective mass $m$ on the brane. The usual 4D graviton corresponds to the zero mode, $m=0$, when ${}^{(5)\\!}p_A$ is tangent to the brane.  The extra dimensions lead to new scalar and vector degrees of freedom on the brane. In 5D, the spin-2 graviton is represented by a metric perturbation ${}^{(5)\\!}h_{AB}$ that is transverse traceless: \\begin{equation} {}^{(5)\\!}h^A{}_{A}=0=\\partial_B \\;{}^{(5)\\!}h_{A}{}^{B}. \\end{equation} In a suitable gauge, ${}^{(5)\\!}h_{AB}$ contains a 3D transverse traceless perturbation $h_{ij}$, a 3D transverse vector perturbation $\\Sigma_i$, and a scalar perturbation $\\beta$, which each satisfy the 5D wave equation~\\cite{durkoc}: \\begin{eqnarray} & {}^{(5)\\!}h_{AB} \\quad \\longrightarrow \\quad h_{ij}, \\Sigma_i, \\beta, & \\label{5dg} \\\\ & h^i{}_i = 0 = \\partial_j h^{ij}, & \\\\ & \\partial_i \\Sigma^i = 0, & \\\\ & (\\Box + \\partial_y^2) \\left( \\begin{array}{c} \\beta \\\\ \\Sigma_i \\\\ h_{ij} \\end{array} \\right) = 0. & \\end{eqnarray}% The other components of ${}^{(5)\\!}h_{AB}$ are determined via constraints once these wave equations are solved. The 5 degrees of freedom (polarizations) in the 5D graviton are thus split into 2 ($h_{ij}$) + 2 ($\\Sigma_i$) +1 ($\\beta$) degrees of freedom in 4D. On the brane, the 5D graviton field is felt as \\begin{itemize} \\item a 4D spin-2 graviton $h_{ij}$ (2 polarizations), \\item a 4D spin-1 gravi-vector (gravi-photon) $\\Sigma_i$ (2 polarizations), and \\item a 4D spin-0 gravi-scalar $\\beta$. \\end{itemize} The massive modes of the 5D graviton are represented via massive modes in all three of these fields on the brane. The standard 4D graviton corresponds to the massless zero-mode of $h_{ij}$.  In the general case of $d$ extra dimensions, the number of degrees of freedom in the graviton follows from the irreducible tensor representations of the isometry group as ${1\\over2}(d+1)(d+4)$.  \\newpage ",
        "conclusions": "\\label{conclusion}  Simple brane-world models provide a rich phenomenology for exploring some of the ideas that are emerging from M~theory. The higher-dimensional degrees of freedom for the gravitational field, and the confinement of standard model fields to the visible brane, lead to a complex but fascinating interplay between gravity, particle physics, and geometry, that enlarges and enriches general relativity in the direction of a quantum gravity theory.  This review has attempted to show some of the key features of brane-world gravity from the perspective of astrophysics and cosmology, emphasizing a geometric approach to dynamics and perturbations. It has focused mainly on 1-brane RS-type brane-worlds, but also considered the DGP brane-world models. The RS-type models have some attractive features: \\begin{itemize} \\item They provide a simple 5D phenomenological realization of the Horava--Witten supergravity solutions in the limit where the hidden brane is removed to infinity, and the moduli effects from the 6 further compact extra dimensions may be neglected. \\item They develop a new geometrical form of dimensional reduction based on a strongly curved (rather than flat) extra dimension. \\item They provide a realization to lowest order of the AdS/CFT correspondence. \\item They incorporate the self-gravity of the brane (via the brane tension). \\item They lead to cosmological models whose background dynamics are completely understood and reproduce general relativity results with suitable restrictions on parameters. \\end{itemize}  The review has highlighted both the successes and some remaining open problems of the RS models and their generalizations. The open problems stem from a common basic difficulty, i.e., understanding and solving for the gravitational interaction between the bulk and the brane (which is nonlocal from the brane viewpoint). The key open problems of relevance to astrophysics and cosmology are \\begin{itemize} \\item to find the simplest realistic solution (or approximation to it) for an astrophysical black hole on the brane, and settle the questions about its staticity, Hawking radiation, and horizon; and \\item to develop realistic approximation schemes (building on recent work~\\cite{sod_1, sod_2, sod_3, sod_4, sod_5, koy, rbbd_1, rbbd_2, kkt, elmw}) and manageable numerical codes (building on~\\cite{koy, rbbd_1, rbbd_2, kkt, elmw}) to solve for the cosmological perturbations on all scales, to compute the CMB anisotropies and large-scale structure, and to impose observational constraints from high-precision data. \\end{itemize}  The RS-type models are the simplest brane-worlds with curved extra dimension that allow for a meaningful approach to astrophysics and cosmology. One also needs to consider generalizations that attempt to make these models more realistic, or that explore other aspects of higher-dimensional gravity which are not probed by these simple models. Two important types of generalization are the following: \\begin{itemize} \\item \\emph{The inclusion of dynamical interaction between the brane(s) and a bulk scalar field,} so that the action is \\begin{eqnarray} S = {1 \\over 2\\kappa_5^2} \\int \\! d^5x\\sqrt{-{}^{(5)\\!}g}\\left[{}^{(5)\\!}R - \\kappa_5^2\\partial_A\\Phi\\partial^A\\Phi-2\\Lambda_5(\\Phi) \\right] + \\int_\\mathrm{brane(s)} \\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\!\\! d^4x\\sqrt{-g} \\left[\\!-\\lambda(\\Phi)+{K\\over \\kappa_5^2}+L_\\mathrm{matter}\\!\\right] \\nonumber \\\\ \\end{eqnarray}% (see~\\cite{mw, sca_1, sca_2, sca_3, sca_4, sca_5, sca_6, hs_1, hs_2, hs_3, hs_4, hs_5, hs_6, hs_7, hs_8, hs_9, hs_10, hs_11, hs_12, hs_13, hs_14, hs_15, hs_16, hs_17}). The scalar field could represent a bulk dilaton of the gravitational sector, or a modulus field encoding the dynamical influence on the effective 5D theory of an extra dimension other than the large fifth dimension~\\cite{ek5_1, ek5_2, ek5_3, rbbd_1, rbbd_2, mod_1, mod_2, mod_3, mod_4, mod_5}.  For two-brane models, the brane separation introduces a new scalar degree of freedom, the radion. For general potentials of the scalar field which provide radion stabilization, 4D Einstein gravity is recovered at low energies on either brane~\\cite{2b_1,2b_2,2b_3}. (By contrast, in the absence of a bulk scalar, low energy gravity is of Brans--Dicke type~\\cite{gt}.) In particular, such models will allow some fundamental problems to be addressed: \\begin{itemize} \\item The hierarchy problem of particle physics. \\item An extra-dimensional mechanism for initiating inflation (or the hot radiation era with super-Hubble correlations) via brane interaction (building on the initial work in~\\cite{kss_1, kss_2, ek_1, ek_2, ek_3, ek_4, ek_5, ek_6, ek_7, ek5_1, ek5_2, ek5_3, bub_1, bub_2, bub_3}). \\item An extra-dimensional explanation for the dark energy (and possibly also dark matter) puzzles: Could dark energy or late-time acceleration of the universe be a result of gravitational effects on the visible brane of the shadow brane, mediated by the bulk scalar field? \\end{itemize} \\item \\emph{The addition of stringy and quantum corrections to the Einstein--Hilbert action,} including the following: \\begin{itemize} \\item Higher-order curvature invariants, which arise in the AdS/CFT correspondence as next-to-leading order corrections in the CFT. The \\emph{Gauss--Bonnet} combination in particular has unique properties in 5D, giving field equations which are second-order in the bulk metric (and linear in the second derivatives), and being ghost-free. The action is \\begin{eqnarray} S&=&{1 \\over 2\\kappa_5^2} \\!\\int \\!\\! d^5x\\sqrt{-{}^{(5)\\!}g}\\left[{}^{(5)\\!}R-2\\Lambda_5+ \\alpha\\!\\left( {}^{(5)\\!}R^2\\!-4 \\:{}^{(5)\\!}R_{AB} \\:{}^{(5)\\!} R^{AB}\\!+{}^{(5)\\!}R_{ABCD} \\;{}^{(5)\\!}R^{ABCD} \\right)\\!\\right] \\nonumber \\\\ && +\\int_\\mathrm{brane} \\!\\!\\!\\!\\!\\!\\!\\!\\!\\! d^4x\\sqrt{-g} \\left[-\\lambda+{K\\over \\kappa_5^2}+L_\\mathrm{matter}\\right]\\!, \\end{eqnarray}% where $\\alpha$ is the Gauss--Bonnet coupling constant, related to the string scale. The cosmological dynamics of these brane-worlds is investigated in~\\cite{gbon_1, gbon_2, gbon_3, gbon_4, gbon_5, gbon_6, gbon_7, gbon_8, gbon_9, gbon_10, gbon_11, gbon_12, gbon_13, gbon_14, gbon_15, gbon_16}. In~\\cite{bmsv} it is shown that the black string solution of the form of Equation~(\\ref{bs1}) is ruled out by the Gauss--Bonnet term. In this sense, the Gauss--Bonnet correction removes an unstable and singular solution.  In the early universe, the Gauss--Bonnet corrections to the Friedmann equation have the dominant form \\begin{equation} H^2 \\propto \\rho^{2/3} \\end{equation} at the highest energies. If the Gauss--Bonnet term is a small correction to the Einstein--Hilbert term, as may be expected if it is the first of a series of higher-order corrections, then there will be a regime of RS-dominance as the energy drops, when $H^2 \\propto \\rho^2$. Finally at energies well below the brane tension, the general relativity behaviour is recovered. \\item Quantum field theory corrections arising from the coupling between brane matter and bulk gravitons, leading to an induced 4D Ricci term in the brane action. The original induced gravity brane-world is the DGP model~\\cite{dgp_1, dgp_2, dgp_3, dgp_4}, which we investigated in this review as an alternative to the RS-type models. Another viewpoint is to see the induced-gravity term in the action as a correction to the RS action: \\begin{equation} S = {1 \\over 2\\kappa_5^2} \\int d^5x \\sqrt{-{}^{(5)\\!}g}\\left[{}^{(5)\\!}R-2\\Lambda_5 \\right]+ \\int_\\mathrm{brane} \\!\\!\\!\\!\\!\\!\\!\\!\\!\\! d^4x \\sqrt{-g} \\left[\\beta{R}-\\lambda+{K\\over \\kappa_5^2}+L_\\mathrm{matter}\\right], \\end{equation} where $\\beta$ is a positive coupling constant.  The cosmological models have been analyzed in~\\cite{ind_1, ind_2, ind_3, ind_4, ind_5, ind_6, ind_7, ind_8, ind_9, ind_10, ind_11, ind_12, ind_13, ind_14, kmp}. (Brane-world black holes with induced gravity are investigated in~\\cite{kpz}.) Unlike RS-type models, DGP models lead to 5D behaviour on \\emph{large scales} rather than small scales. Then on an FRW brane, the late-universe 5D behaviour of gravity can naturally produce a late-time acceleration, even \\emph{without dark energy}, although the self-accelerating models suffer from a ghost. Nevertheless, the DGP model is a critical example of modified gravity models in cosmology that act as alternatives to dark energy.   \\end{itemize} \\end{itemize}  The RS and DGP models are 5-dimensional phenomenological models, and so a key issue is how to realize such models in 10-dimensional string theory. Some progress has been made. 6-dimensional cascading brane-worlds are extensions of the DGP model. 10-dimensional type IIB supergravity solutions have been found with the warped geometry that generalizes the RS geometry. These models have also been important for building inflationary models in string theory, based on the motion of D3 branes in the warped throat \\cite{Burgess:2001vr, Dvali:2001fw} (see the reviews~\\cite{Linde:2005dd, Baumann:2009ni} and references therein). The action for D3 branes is described by the Dirac-Born-Infeld action and this gives the possibility of generating a large non-Gaussianity in the Cosmic Microwave Background temperature anisotropies, which can be tested in the future experiments \\cite{Silverstein:2003hf, Kachru:2003sx} (see the reviews~\\cite{Chen:2010xk, Koyama:2010xj}).  These models reply on the effective 4-dimensional approach to deal with extra dimensions. For example, the stabilization mechanism, which is necessary to fix moduli fields in string theory, exploits non-perturbative effects and they are often added in the 4-dimensional effective theory. Then it is not clear whether the resultant 4-dimensional effective theory is consistent with the 10-dimensional equations of motion~\\cite{deAlwis:2003sn, deAlwis:2004qh, Kodama:2005cz, Koyama:2006ni}. Recently there has been a new development and it has become possible to calculate all significant contributions to the D3 brane potential in the single coherent framework of 10-dimensional supergravity~\\cite{Baumann:2007np, Baumann:2007ah, Baumann:2009qx, Baumann:2010sx}. This will provide us with a very interesting bridge between phenomenological brane-world models, where dynamics of higher-dimensional gravity is studied in detail, and string theory approaches, where 4D effective theory is intensively used. It is crucial to identify the higher-dimensional signature of the models in order to test a fundamental theory like string theory.  In summary, brane-world gravity opens up exciting prospects for subjecting M~theory ideas to the increasingly stringent tests provided by high-precision astronomical observations. At the same time, brane-world models provide a rich arena for probing the geometry and dynamics of the gravitational field and its interaction with matter.   \\newpage"
    },
    "1004/1004.3365_arXiv.txt": {
        "abstract": "We investigate observational constraints on  the generalized Chaplygin gas (GCG) model as the unification of dark matter and dark energy from  the latest observational data: the Union SNe Ia data, the  observational Hubble data, the SDSS baryon acoustic peak and the five-year WMAP shift parameter. It is obtained  that the best fit values of the GCG model parameters with their confidence level are $A_{s}=0.73^{+0.06}_{-0.06}$ ($1\\sigma$) $^{+0.09}_{-0.09}$ $(2\\sigma)$, $\\alpha=-0.09^{+0.15}_{-0.12}$ ($1\\sigma$) $^{+0.26}_{-0.19}$ $(2\\sigma)$. Furthermore in this model, we can see that the evolution of equation of state (EOS) for dark energy is similar to quiessence, and its current best-fit value is $w_{0de}=-0.96$ with the $1\\sigma$ confidence level $-0.91\\geq w_{0de}\\geq-1.00$. ",
        "introduction": "$}   {\\small {~~~}}~ The recently cosmic observations from the type Ia supernovae (SNe Ia) \\cite{SNe}, the cosmic microwave background (CMB) \\cite{CMB}, the clusters of galaxies \\cite{LSS} etc., all suggest that the expansion of present universe is speeding up rather than slowing down.  And it indicates that baryon matter component is about 5\\% for total energy density, and about 95\\% energy density in universe is invisible.  Considering the four-dimensional  standard cosmology, the accelerated expansion of the present universe is usually attributed to the fact that dark energy (DE) is an exotic component with negative pressure. And it is shown that DE takes up about two-thirds of the total energy density from cosmic observations. Many kinds of DE models have already been constructed such as $\\Lambda$CDM \\cite{LCDM}, quintessence \\cite{quintessence}, phantom \\cite{phantom}, quintom \\cite{quintom}, generalized Chaplygin gas (GCG) \\cite{GCG}, modified Chaplygin gas \\cite{MCG}, holographic dark energy \\cite{holographic}, agegraphic dark energy\\cite{agegraphic}, and so forth. Furthermore, model-independent method\\footnote{Using mathematical fundament, one expands  equation of state  of DE $w_{de}$ or deceleration parameter $q$ with respect to scale factor $a$ or redshit $z$. For example, $w_{de}(z)=w_{0}$=const \\cite{independent1}, $w_{de}(z)=w_{0}+w_{1}z$\\cite {independent2}, $w_{de}(z)=w_{0}+w_{1}\\ln(1+z)$ \\cite{independent3}, $w_{de}(z)=w_{0}+\\frac{w_{1}z}{1+z}$ \\cite{independent4}, $q(z)=q_{0}+q_{1}z$ \\cite{independent1}, $q(z)=q_{0}+ \\frac{q_{1}z}{1+z}$ \\cite{independent5}, where $w_{0}$, $w_{1}$, or $q_{0}$, $q_{1}$ are model parameters.} and modified gravity theories (such as scalar-tensor cosmology \\cite{scalar}, braneworld models \\cite{braneworld})  to interpret accelerating universe have also been discussed.  It is well known that the GCG model have been widely studied  for interpreting the accelerating universe \\cite{GCGpapers}. The most interesting property for this scenario is that, two unknown dark sections in universe--dark energy and dark matter  can be unified by using an exotic equation of state (EOS). In this paper, we use the latest observational data: the Union SNe Ia data \\cite{307Union}, the observational Hubble data (OHD) \\cite{OHD},  the baryon acoustic oscillation (BAO) peak from Sloan Digital Sky Survey (SDSS) \\cite{SDSS} and the five-year WMAP CMB shift parameter \\cite{5yWMAP} to constrain the GCG model. And we  discuss whether the parameter degeneration \\cite{GCG2}\\cite{GCG3} for the GCG model can be broken by the latest observed data,  since  it is always expected that the model degeneration problem can be solved  by the more accurate observational data.  The paper is organized as follows. In section 2, the GCG model as the unification of dark matter and dark energy is introduced briefly. Based on the   observational data, we constrain the GCG model parameter in section 3.  The evolutions of  EOS of DE and deceleration parameter for GCG model are presented in section 4. Section 5 is the conclusions. ",
        "conclusions": "$}  The constraints on the GCG model as the unification of dark matter and dark energy are studied in this paper by using the latest observational data: the Union SNe Ia data, the observational Hubble data, the SDSS baryon acoustic peak and the five-year WMAP shift parameter. We find that the model parameters $A_{s}$ and $\\alpha$ are degenerate, and their values  are constrained to $A_{s}=0.73^{+0.06}_{-0.06}$ ($1\\sigma$) $^{+0.09}_{-0.09}$ $(2\\sigma)$ and $\\alpha=-0.09^{+0.15}_{-0.12}$ ($1\\sigma$) $^{+0.26}_{-0.19}$ $(2\\sigma)$. This constraint on parameter $\\alpha$ is more stringent than the results in Refs. \\cite{GCG2}\\cite{GCG3}. Furthermore, it is shown that the evolution of EOS of dark energy for the GCG model is similar to quiessence, and the best fit value of current EOS of DE $w_{0de}=-0.96>-1$. And it indicates that the values of transition redshift and current deceleration parameter are $z_{T}=0.74^{+0.05}_{-0.05}$ $(1\\sigma)$, $q_{0}=-0.55^{+0.06}_{-0.05}$ $(1\\sigma)$.    \\textbf{\\ Acknowledgments } The research work is supported by  NSF (10573004) of PR China."
    },
    "1004/1004.4259_arXiv.txt": {
        "abstract": "We report on $G$-band emission observed by the Solar Optical Telescope onboard the {\\it Hinode} satellite in association with the X1.5-class flare on 2006 December~14. The $G$-band enhancements originate from the footpoints of flaring coronal magnetic loops, coinciding with non-thermal hard X-ray bremsstrahlung sources observed by the {\\it Reuven Ramaty High Energy Solar Spectroscopic Imager}. At the available 2 minute cadence, the $G$-band and hard X-ray intensities are furthermore well correlated in time. Assuming that the $G$-band enhancements are continuum emission from a blackbody, we derived the total radiative losses of the white-light flare (white-light power). If the $G$-band enhancements additionally have a contribution from lines, the derived values are overestimates. We compare the white-light power with the power in hard X-ray producing electrons using the thick target assumption. Independent of the cutoff energy of the accelerated electron spectrum, the white-light power and the power of accelerated electrons are roughly proportional. Using the observed upper limit of $\\sim$30 keV for the cutoff energy, the hard X-ray producing electrons provide at least a factor of 2 more power than needed to produce the white-light emission. For electrons above $40\\,{\\rm keV}$, the powers roughly match for all four of the time intervals available during the impulsive phase. Hence, the flare-accelerated electrons contain enough energy to produce the white-light flare emissions. The observed correlation in time, space, and power strongly suggests that electron acceleration and white-light production in solar flares are closely related. However, the results also call attention to the inconsistency in apparent source heights of the hard X-ray (chromosphere) and white-light (upper photosphere) sources. ",
        "introduction": "In association with solar flares, we sometimes observe enhancements of visible continuum, in which case the event is termed a ``white-light flare.'' Although white-light events had previously been mainly associated with energetic flares ({\\it GOES} X-class), there are now reports of continuum emission from events as weak as C-class flares \\citep{Matthews2003, Hudson2006, Wang2009, Jess2008} thanks to accurate photometry from space achieved by {\\it Yohkoh}, {\\it TRACE}, and {\\it Hinode}, and by improved ground-based instruments. However, white-light flares are still very infrequently observed and some energetic events do not show any enhancement in white light. The processes causing it remain unclear \\citep{Neidig1989}. Because there is a good correlation of light curves and sites of emission between optical continuum and hard X-rays \\citep[e.g.,][]{RustHegwer1975, Neidig1989, Hudson1992, Metcalf2003, Xu2006}, there is some consensus that the origin of white-light emission lies in the energy in accelerated particles, especially non-thermal electrons.  Using the thick-target model \\citep{Brown1971}, the energy in flare-accelerated electrons can be compared to the radiative losses in white light \\citep[e.g.,][]{Hudson1972}. If flare-accelerated electrons indeed produce the white-light emission, the energy content in electrons must be larger than the radiative losses in white light. Due to the (inferred) steep electron spectrum, the energy in electrons strongly depends on the cutoff of the electron spectrum at low energies. To match the energies, \\citet{Neidig1989} and \\citet{Ding2003} estimated the cutoff energy of electrons at more than $50\\,{\\rm keV}$, whereas \\citet{Fletcher2007} obtained values below $25\\,{\\rm keV}$ from a statistical analysis of {\\it Transition Region and Coronal Explorer (TRACE)} and {\\it Reuven Ramaty High Energy Solar Spectroscopic Imager (RHESSI)} observations. These differences might be due to the variation from flare to flare. In any case, a cutoff energy of $\\sim20\\,{\\rm keV}$ can supply the white-light power, but not $100\\,{\\rm keV}$.  The Solar Optical Telescope (SOT) of \\textit{Hinode} \\citep{Tsuneta2008, Suematsu2008, Shimizu2008, Ichimoto2008} makes observations in white light. Its broadband filter imager (BFI) take images in red ($668.40\\,{\\rm nm}$, width $0.4\\,{\\rm nm}$), green ($555.05\\,{\\rm nm}$, width $0.4\\,{\\rm nm}$) and blue ($450.45\\,{\\rm nm}$, width $0.4\\,{\\rm nm}$) continuum ranges. Radiation at these wavelengths comes from the photosphere and hence reflects the broadband continuum emission well. However, SOT normally obtains only infrequent images in these filters. More frequently, SOT takes images in the $G$-band ($430.50\\,{\\rm nm}$, width $0.83\\,{\\rm nm}$), formed mainly from CH~line opacity. \\cite{Carlsson2007} show contribution functions for these filters; the $G$-band has a photospheric and an upper-photospheric contribution. It therefore serves well to define the morphology of white-light flares and it was also used in the \\textit{Yohkoh} observations \\citep{Hudson1992,Matthews2003}. However, $G$-band emission could contain not only continuum emission, but also CH~line emission. If the $G$-band emission contains line emission, the radiative losses estimated from the $G$-band emission assuming blackbody radiation is overestimates of the true losses. However, in this paper, we treat the $G$-band emission mainly came from the continuum emission, and we therefore use $G$-band images as a proxy for the white-light images.  SOT observed white-light emission from three X-class flares in 2006 December \\citep{Wang2009}. $G$-band emission of the largest event (X3.4 flare on 2006 December~13) is reported by \\citet{Isobe2007} and \\citet{Jing2008}. \\citet{Isobe2007} concluded that the white-light emission could be produced by radiative back-warming resulting from particle-beam heating, and \\citet{Jing2008} noted that the white-light emissions appeared at the sites of the largest inferred reconnection rates. In this paper, we describe the white-light observations of the 2006 December~14 flare that was also observed by the {\\it RHESSI} \\citep{Lin2002}. We obtain X-ray energy spectra for each foot-point separately, using {\\it RHESSI} imaging spectroscopy, and compare the results with energy estimated from {\\it Hinode} $G$-band images (Section \\ref{observation}). In Section \\ref{relation}, energy estimates are discussed for different cutoff energies. ",
        "conclusions": "The observations discussed in this paper show that the solar flare white-light emission is closely related in time, space, and power to the acceleration of non-thermal electrons. To explain the observed correlation between white light and high energy hard X-ray emission in the simplest possible way, the two components should originate in the same source region. Continuum emission in the $G$-band emission comes from $0-100\\,{\\rm km}$ above the photosphere \\citep[see Figure~1 of ][]{Carlsson2007}, and hard X-ray emission in $50-100\\,{\\rm keV}$ originates in the chromosphere. Observationally, the emission site of $50-100\\,{\\rm keV}$ hard X-rays is estimated at $6.5 \\times 10^{3}\\,{\\rm km}$ height above the photosphere from (early) {\\it Yohkoh} observations \\citep{Matsushita1992} and around $600\\,{\\rm km}$ height by  \\textit{RHESSI} for a single event \\citep{Kontar2008}. This information is weak and we would like to see systematic \\textit{RHESSI} data analyses on this point, but the existing data suggest a difference of more than $500\\,{\\rm km}$ between the emission sites \\citep{Kontar2008, Carlsson2007}. Theoretically, a $50-100\\,{\\rm keV}$ electron should thermalize some $1000\\,{\\rm km}$ height above the photosphere; at this mid-chromospheric height, the density is about $10^{13.5}/{\\rm cm}^3$ \\citep{Neidig1989}. At these energies, however, the electrons cannot penetrate into the lower chromosphere, and thus they do not heat the photosphere. Electron energies more than $900\\,{\\rm keV}$ are necessary for penetration to the photosphere, even if the flare site has become ionized \\citep{Neidig1989}. However, the energy in $900\\,{\\rm keV}$ electrons is far too small (by about 4 orders of magnitude, assuming the power law seen at $40\\,{\\rm keV}$ can be extrapolated to $900\\,{\\rm keV}$) to produce the white-light emission. The data presented here therefore call attention to the need for a white-light emission model which can explain the good correlation with high-energy electron emission and difference of the emission height of white light and hard X-rays. Non thermal ionization levels enhance the continuum \\citep{Hudson1972} and also make back-warming possible \\citep[e.g.,][]{Metcalf2003}, but we do not have well-defined models for these processes in realistic physical conditions yet."
    },
    "1004/1004.4403_arXiv.txt": {
        "abstract": "Using the Hugenholtz-Van Hove theorem, we derive general expressions for the quadratic and quartic symmetry energies in terms of single-nucleon potentials in isospin asymmetric nuclear matter. These analytical relations are useful for gaining deeper insights into the microscopic origins of the uncertainties in our knowledge on nuclear symmetry energies especially at supra-saturation densities. As examples, the formalism is applied to two model single-nucleon potentials that are widely used in transport model simulations of heavy-ion reactions. ",
        "introduction": "One of the central issues currently under intense investigation in both nuclear physics and astrophysics is the Equation of State (EOS) of neutron-rich nuclear matter \\cite{JML04,AWS05,li1}. For cold nuclear matter of isospin asymmetry $\\delta=(\\rho_n-\\rho_p)/(\\rho_n+\\rho_p)$ at density $\\rho$, the energy per nucleon $E(\\rho,\\delta)$ can be expressed as an even series of $\\delta$ that respects the charge symmetry of strong interactions, namely, $E(\\rho,\\delta)=E_0(\\rho,0)+\\sum_{i=2,4,6...}E_{sym,i}(\\rho)\\delta^i$ where $E_{sym,i}(\\rho)$ is the so-called symmetry energy of the \\textit{i}th order \\cite{li1} and $E_0(\\rho,0)$ is the EOS of symmetric nuclear matter. The quadratic term $E_{sym,2}(\\rho)$ is most important and its value at normal nuclear matter density $\\rho_0$ is known to be around 30 MeV from analyzing nuclear masses within liquid-drop models. Essentially, all microscopic many-body calculations have indicated that the higher-order terms are usually negligible around $\\rho_0$, leading to the so-called empirical parabolic law of EOS even for $\\delta$ approaching unity for pure neutron matter. The $E_{sym,2}(\\rho)$ is then generally regarded as the symmetry energy. For instance, the value of the quartic term has been estimated to be less than 1 MeV at $\\rho_0$ \\cite{sie70,lee98}. However, the presence of higher-order terms at supra-saturation densities can significantly modify the proton fraction in neutron stars at $\\beta$-equilibrium and thus the cooling mechanism of proto-neutron stars \\cite{ste,zfs}. It was also found that a tiny quartic term can cause a big change in the calculated core-crust transition density in neutron stars \\cite{sjo,XCLM09}. Therefore, precise evaluations of the quartic symmetry energy in neutron-rich matter are useful. Although much information about the EOS of symmetric nuclear matter $E_0(\\rho,0)$ has been accumulated over the past four decades, our knowledge about the density dependence of $E_{sym,i}(\\rho)$ is unfortunately still very poor. It has been generally recognized that the $E_{sym,i}(\\rho)$, especially the quadratic and quartic terms, is critical for understanding not only the structure of rare isotopes and the reaction mechanism of heavy-ion collisions, but also many interesting issues in astrophysics \\cite{XCLM09,li0,bro,li2,dan,bar,Sum94,Bom01,LWC05,tsa,Cen09,Joe10,xia,wen}. Therefore, to determine the $E_{sym,i}(\\rho)$ in neutron-rich matter has recently become a major goal in both nuclear physics and astrophysics. While significant progress has been made recently in constraining the $E_{sym,2}(\\rho)$ especially around and below the saturation density, see, e.g., \\cite{LWC05,tsa,Cen09,Joe10}, much more work needs to be done to constrain more tightly the $E_{sym,i}(\\rho)$ at supra-saturation densities where model predictions are rather diverse \\cite{Das03,ulr,van,zuo,Fri05,she,Che05c,zhli,Sto03,pan,wir,kut}. As dedicated experiments using advanced new detectors have now been planned to investigate the high density behavior of $E_{sym,2}(\\rho)$ at many radioactive beam facilities around the world, it has become an urgent task to investigate theoretically more deeply the fundamental origin of the extremely uncertain high density behavior of $E_{sym,2}(\\rho)$. It is also of great interest to evaluate possible corrections due to the $E_{sym,4}(\\rho)$ term to the equation of state of asymmetric nuclear matter.  Among existing proposals for extracting information about $E_{sym,i}(\\rho)$ using terrestrial laboratory experiments, transport model simulations have shown that many observables in heavy-ion reactions are particularly useful for studying $E_{sym,i}(\\rho)$ in a broad density range. In these transport model simulations of heavy-ion reactions, the EOS enters the reaction dynamics and affects the final observables through the single-nucleon potential $U_{n/p}(\\rho,\\delta,k)$ where $k$ is the nucleon momentum. Except in situations where statistical equilibrium is established and thus many observables are directly related to the binding energy $E(\\rho,\\delta)$ after correcting for finite-size effects, what is being directly probed in heavy-ion reactions is the single-nucleon potential $U_{n/p}(\\rho,\\delta,k)$. The latter is, however, directly related to the symmetry energy $E_{sym,2}(\\rho)$ through the underlying nuclear effective interaction as first pointed out by Brueckner, Dabrowski and Haensel \\cite{bru64,Dab73} using K-matrices within the Brueckner theory. They showed that if one expands $U_{n/p}(\\rho,\\delta,k)$ to the leading order in $\\delta$ as in the well-known Lane potential \\cite{Lan62}, i.e., \\begin{equation}\\label{Lane} U_{n/p}(\\rho,\\delta,k)\\approx U_0(\\rho,k) \\pm U_{sym,1}(\\rho,k)\\delta \\end{equation} where $U_0(\\rho,k)$ and $U_{sym,1}(\\rho,k)$ are, respectively, the nucleon isoscalar and isovector (symmetry) potentials, the quadratic symmetry energy is then \\cite{bru64,Dab73} \\begin{equation}\\label{Esym} E_{sym,2}(\\rho)=\\frac{1}{3} t(k_F) + \\frac{1}{6} \\frac{\\partial U_0}{\\partial k}\\mid _{k_F}\\cdot k_F + \\frac{1}{2}U_{sym,1}(\\rho,k_F) \\end{equation} where $t(k_F)$ is the nucleon kinetic energy at the Fermi momentum $k_F=(3\\pi^2\\rho/2)^{1/3}$ in symmetric nuclear matter of density $\\rho$. The above equation indicates that the symmetry energy $E_{sym,2}(\\rho)$ depends only on the single-particle kinetic and potential energies at the Fermi momentum $k_F$. This is not surprising since the microscopic origin of the symmetry energy is the difference in the Fermi surfaces of neutrons and protons. The first term $E_{sym}^{kin}=\\frac{1}{3} t(k_F)=\\frac{\\hbar ^2}{6m} (\\frac{3\\pi^2}{2})^{\\frac{2}{3}} \\rho^{\\frac{2}{3}}$ is the trivial kinetic contribution due to the different Fermi momenta of neutrons and protons; the second term $\\frac{1}{6} \\frac{\\partial U_0}{\\partial k}\\mid _{k_F}\\cdot k_F$ is due to the momentum dependence of the isoscalar potential and also the fact that neutrons and protons have different Fermi momenta; while the term $\\frac{1}{2}U_{sym}(\\rho,k_F)$ is due to the explicit isospin dependence of the nuclear strong interaction. For the isoscalar potential $U_0(\\rho,k)$, reliable information about its density and momentum dependence has already been obtained from high energy heavy-ion collisions, see, e.g., ref. \\cite{dan}, albeit there are still some rooms for further improvements, particularly at high momenta/densities. On the contrary, the isovector potential $U_{sym,1}(\\rho,k)$ is still not very well determined, especially at high densities and momenta, and has been identified as the key quantity responsible for the uncertain high density behavior of the symmetry energy as stressed in ref.\\cite{li1}.  In the present work, we first show, using both the differential and integral formulations of the Hugenholtz-Van Hove (HVH) theorem \\cite{hug}, that the relation in Eq.(\\ref{Esym}) is valid in general. We then derive an expression for the quartic symmetry energy $E_{sym,4}(\\rho)$ in terms of the single-nucleon potential by keeping higher-order terms in the expansion of both the EOS and the single-nucleon potential. Applying the HVH formalism to two model single-nucleon potentials, namely, the Bombaci-Gale-Bertsch-Das Gupta (BGBD) potential \\cite{Bom01} and a modified Gogny Momentum-Dependent-Interaction (MDI)\\cite{Das03,dec}, which are among the most widely used ones in studying isospin physics based on transport model simulations of heavy-ion reactions \\cite{li1,bar}, we examine the relative contributions from the kinetic and various potential terms to $E_{sym,2}(\\rho)$ and $E_{sym,4}(\\rho)$. We put the emphasis on identifying those terms that dominate the high density behaviors of $E_{sym,2}(\\rho)$. Finally, we evaluate the relative importance of the $E_{sym,4}(\\rho)$ term by studying the $E_{sym,4}(\\rho)/ E_{sym,2}(\\rho)$ ratio as a function of density.  The paper is organized as follows. In Section~\\ref{symmetry}, based on the HVH theorem we derive general expressions for the higher-order symmetry energy terms $E_{sym,2}(\\rho)$ and $E_{sym,4}(\\rho)$ in terms of the single-nucleon isoscalar and isovector potentials. The derivation is carried out in Section~\\ref{hvh} using the differential form of the HVH theorem by starting from the neutron and proton chemical potentials and in Section~\\ref{fermi} using the integral form of the HVH theorem by starting from the total energy density of the system. Numerical results and discussions for both the BGBD and MDI interactions are given in Sections~\\ref{BGBD} and \\ref{MDI}, respectively. Finally, we give a summary in Section~\\ref{summary}. ",
        "conclusions": "As shown in the previous section, the symmetry energy can be explicitly separated into the kinetic energy term $T$ and the potential terms $U_0$ and $U_{sym,i}$ at the Fermi momentum $k_F$. To evaluate their relative contributions to the symmetry energies, especially for the second-order and fourth-order terms $E_{sym,2}(\\rho)$ and $E_{sym,4}(\\rho)$, we consider in this section two typical single-nucleon potentials that have been widely used in tansport model simulations of heavy-ion reactions.  \\subsection{The Bombaci-Gale-Bertsch-Das Gupta potential}\\label{BGBD}  As a first example, we use the phenomenological potential of Bombaci-Gale-Bertsch-Das Gupta \\cite{Bom01} \\begin{eqnarray} \\label{Bomba} &&U_\\tau(u,\\delta,k)= Au+Bu^\\sigma -\\frac{2}{3}(\\sigma-1)\\frac{B}{\\sigma+1} (\\frac{1}{2}+x_{3}) u^{\\sigma}\\delta^2 \\nonumber\\\\ &&\\pm \\left[{-\\frac{2}{3}A (\\frac{1}{2}+x_{0}) u - \\frac{4}{3}\\frac{B}{\\sigma+1}(\\frac{1}{2}+x_{3})u^{\\sigma}\\,}\\right]\\delta\\nonumber\\\\ && +\\frac{4}{5\\rho_0}\\left[{\\frac{1}{2} (3C-4z_1) \\mathcal{I_{\\tau}} + (C+2z_1)\\mathcal{{I_{\\tau^{\\prime}}}}}\\right]+ \\left({C \\pm \\frac{C-8z_1}{5}\\delta}\\right)u\\cdot g(k), \\end{eqnarray} where $u=\\rho/\\rho_0$ is the reduced density and $\\pm$ is for neutrons/protons. In the above, we have $\\mathcal{I}_\\tau=[2/(2\\pi)^3]\\int d^3k f_{\\tau}(k)g(k)$ with $g(k)= 1/[{1+({\\frac{k}{\\Lambda}})^2 }]$ being a momentum regulator and $f_{\\tau}(k)$ being the phase space distribution function. The parameter $\\Lambda$ has the value $\\Lambda=1.5k_F^0$, where $k_F^0$ is the nucleon Fermi wave number in symmetric nuclear matter at $\\rho_0$. With A=-144 MeV, B=203.3 MeV, C=-75 MeV and $\\sigma=7/6$, the BGBD potential reproduces all ground state properties including an incompressibility $K_0$=210 MeV for symmetric nuclear matter \\cite{Bom01}. The three parameters $x_0, x_3$ and $z_1$ can be adjusted to give different symmetry energy $E_{sym,2}(\\rho)$ and the neutron-proton effective mass splitting $m^*_n-m^*_p$ \\cite{Bom01,riz04,Li04}. For example, the parameter set $z_1=-36.75$ MeV, $x_0=-1.477$ and $x_3=-1.01$ leads to $m_n^*>m_p^*$ while the one with $z_1=50$ MeV, $x_0=1.589$ and $x_3=-0.195$ leads to $m_n^*<m_p^*$ at all non-zero densities and isospin asymmetries.  On expanding the BGBD potential in $\\delta$, the coefficients of the first few terms are \\begin{eqnarray} U_0(\\rho,k)&=& U_{n/p}|\\, _{\\delta=0} \\nonumber \\\\&=&   Au+Bu^\\sigma +\\frac{2C}{\\rho_0} \\frac{\\Lambda^2}{\\pi^2} \\left[k_F - \\Lambda \\textrm{tan}^{-1}\\left(\\frac{k_F}{\\lambda}\\right)\\right]+ C u\\cdot g(k), \\end{eqnarray} \\begin{eqnarray} U_{sym,1}(\\rho,k)&=&\\pm \\frac{1}{1!} \\frac{\\partial U_{n/p}}{\\partial \\delta}|\\, _{\\delta=0}\\nonumber\\\\ &=& \\left[{-\\frac{2}{3}A (\\frac{1}{2}+x_{0}) u - \\frac{4}{3}\\frac{B}{\\sigma+1}(\\frac{1}{2}+x_{3})u^{\\sigma}\\,}\\right]\\nonumber\\\\ &+& \\frac{C-8z_1}{5}u\\cdot g(k_F) + \\frac{C-8z_1}{5}u\\cdot g(k), \\end{eqnarray} \\begin{eqnarray} U_{sym,2}(\\rho,k) &=& \\frac{1}{2!} \\frac{\\partial^2 U_{n/p}}{\\partial \\delta^2}|\\, _{\\delta=0} \\nonumber \\\\&=& -\\frac{2}{3}(\\sigma-1)\\frac{B}{\\sigma+1} (\\frac{1}{2}+x_{3}) u^{\\sigma} - \\frac{C}{3}u \\frac{k_F^2}{\\Lambda^2} g(k_F)^2, \\end{eqnarray} \\begin{eqnarray} U_{sym,3}(\\rho,k) &=&  \\pm \\frac{1}{3!} \\frac{\\partial^3 U_{n/p}}{\\partial \\delta^3}|\\, _{\\delta=0} \\nonumber \\\\&=& \\frac{C-8z_1}{135}u \\frac{k_F^2}{\\Lambda^2}(\\frac{5k_F^2}{\\Lambda^2}+1) g(k_F)^3. \\end{eqnarray}  Thus, the second-order symmetry energy $E_{sym,2}(\\rho)$ is given by \\begin{eqnarray} E_{sym,2}(\\rho)  &=& \\frac{1}{3} t(k_F) + \\frac{1}{6} \\frac{\\partial U_0}{\\partial k}\\mid _{k_F}\\cdot k_F + \\frac{1}{2}U_{sym,1}(\\rho,k_F) \\nonumber\\\\ &=& \\frac{\\hbar^2}{6m} \\left(\\frac{3\\pi^2}{2}\\right)^{2/3} \\rho^{2/3} - \\frac{C}{3}u \\frac{k_F^2}{\\Lambda^2} g(k_F)^2 \\nonumber \\\\&+& \\left[{-\\frac{1}{3}A (\\frac{1}{2}+x_{0}) u - \\frac{2}{3}\\frac{B}{\\sigma+1}(\\frac{1}{2}+x_{3})u^{\\sigma}\\,}\\right] \\nonumber\\\\&+& \\frac{C-8z_1}{5}u\\cdot g(k_F), \\end{eqnarray} and the fourth-order symmetry energy $E_{sym,4}(\\rho)$ is \\begin{eqnarray} &&E_{sym,4}(\\rho) = \\frac{\\hbar^2}{162m} \\left(\\frac{3\\pi^2}{2}\\right)^{2/3} \\rho^{2/3} \\nonumber \\\\& + & \\left[ \\frac{5}{324}\\frac{\\partial U_0(\\rho,k)}{\\partial k}|_{k_F} k_F - \\frac{1}{108} \\frac{\\partial^2 U_0(\\rho,k)}{\\partial k^2}|_{k_F} k_F^2 +\\frac{1}{648} \\frac{\\partial^3 U_0(\\rho,k)}{\\partial k^3}|_{k_F} k_F^3 \\right. \\nonumber \\\\& - & \\left. \\frac{1}{36} \\frac{ \\partial U_{sym,1}(\\rho,k)}{\\partial k}|_{k_F} k_F + \\frac{1}{72} \\frac{ \\partial^2 U_{sym,1}(\\rho,k)}{\\partial k^2}|_{k_F} k_F^2 + \\frac{1}{12} \\frac{\\partial U_{sym,2}(\\rho,k)}{\\partial k}|_{k_F} k_F + \\frac{1}{4} U_{sym,3}(\\rho,k_F)\\right] \\nonumber \\\\& =  & \\frac{\\hbar ^2}{162m} (\\frac{3\\pi^2 \\rho}{2})^{\\frac{2}{3}}- \\frac{C}{81} u \\frac{k_F^2}{\\Lambda^2}(\\frac{10k_F^4}{\\Lambda^4}+\\frac{5k_F^2}{\\Lambda^2}+1) g(k_F)^4+ \\frac{C-8z_1}{135}u \\frac{k_F^2}{\\Lambda^2}(\\frac{5k_F^2}{\\Lambda^2}+1) g(k_F)^3.\\nonumber\\\\ \\end{eqnarray}  \\begin{figure}[!htb] \\centering \\includegraphics[width=12cm]{Graph1.EPS} \\caption{The kinetic energy part (T), the isoscalar potential part ($U_0$) and the isovector potential part ($U_{sym,1}$) of the symmetry energy $E_{sym,2}$ from the BGBD potential with $m_n^*>m_p^*$ (left) and for $m_n^*< m_p^*$ (right).} \\label{BGBDe2} \\end{figure}  In Fig.\\ \\ref{BGBDe2} we compare $E_{sym,2}(\\rho)$ and its three components in the two cases of $m_n^*>m_p^*$ and $m_n^*<m_p^*$. It is seen that the kinetic and isoscalar contributions are the same in both cases. However, they have significantly different isovector potentials $U_{sym,1}$, leading thus to different $E_{sym,2}(\\rho)$ especially at supra-saturation densities.  \\begin{figure}[!htb] \\centering \\includegraphics[width=12cm]{Graph2.EPS} \\caption{The kinetic and various potential contributions to the fourth-order symmetry energy $E_{sym,4}$ with the BGBD potential for $m_n^*>m_p^*$ (left) and for $m_n^*< m_p^*$ (right).} \\label{BGBDe4} \\end{figure}  Various contributions to the fourth-order symmetry energies $E_{sym,4}(\\rho)$ in the two cases are compared in Fig.\\ \\ref{BGBDe4}. Similar to $E_{sym,2}(\\rho)$, the contributions of the $T$ and $U_0$ terms to $E_{sym,4}(\\rho)$ are positive and they are the same in both cases. Interestingly, the $U_{sym,1}$ term also plays the most important role in determining the high-density behavior of $E_{sym,4}(\\rho)$. It is positive in the case of $m_n^*>m_p^*$ but negative in the case of $m_n^*<m_p^*$, resulting in very different behaviors of $E_{sym,4}(\\rho)$ at supra-saturation densities. Moreover, it is interesting to note that $E_{sym,4}(\\rho)$ receives no contribution from the $U_{sym,2}$ term. This is not surprising because the $U_{sym,2}$ term in the BGBD interaction is momentum independent and its contribution to $E_{sym,4}(\\rho)$ is actually $\\frac{1}{12} \\frac{\\partial U_{sym,2}(k)}{\\partial k}|_{k_F} k_F$=0. On the contrary, the $U_{sym,3}$ term still contributes to $E_{sym,4}(\\rho)$ via $\\frac{1}{4} U_{sym,3}(k_F)$ although it is momentum independent too. In the two cases considered here, the contributions from the $U_{sym,3}$ term also have opposite sign.  \\begin{figure}[!htb] \\centering \\includegraphics[width=12cm]{Graph3.EPS} \\caption{The ratio of $E_{sym,4}$ over $E_{sym,2}$ with the BGBD potential for $m_n^*>m_p^*$ (left) and for $m_n^*< m_p^*$ (right).} \\label{BGBDr} \\end{figure}  To compare the fourth-order term $E_{sym,4}(\\rho)$ with the second-order term $E_{sym,2}(\\rho)$ more clearly, we show in Fig.\\ \\ref{BGBDr} their ratio $E_{sym,4}(\\rho)/E_{sym,2}(\\rho)$ as a function of the reduced density $\\rho/\\rho_0$. Obviously, the relative value of $E_{sym,4}(\\rho)$ is generally small. However, it can reach up to about $\\pm 10\\%$ at high densities for both cases of $m_n^*> m_p^*$ and $m_n^*< m_p^*$ . It may thus lead to an appreciable modification in the proton fraction and therefore the properties of neutron stars at $\\beta$-equilibrium.  \\subsection{A modified Gogny Momentum-Dependent-Interaction}\\label{MDI}  In this subsection, we discuss the symmetry energy obtained from the MDI interaction \\cite{Das03}, which is derived from the Hartree-Fock approximation using a modified Gongy effective interaction \\cite{dec} \\begin{eqnarray}\\label{mdi} &&U(\\rho,\\delta,\\vec p,\\tau) = A_u(x)\\frac{\\rho_{\\tau'}}{\\rho_0} +A_l(x)\\frac{\\rho_{\\tau}}{\\rho_0}\\nonumber\\\\ &&+B(\\frac{\\rho}{\\rho_0})^{\\sigma}(1-x\\delta^2)-8\\tau x\\frac{B}{\\sigma+1}\\frac{\\rho^{\\sigma-1}}{\\rho_0^{\\sigma}}\\delta\\rho_{\\tau'} \\nonumber \\\\&& +\\frac{2C_{\\tau,\\tau}}{\\rho_0} \\int d^3p'\\frac{f_{\\tau}(\\vec r,\\vec p')}{1+(\\vec p-\\vec p')^2/\\Lambda^2}+\\frac{2C_{\\tau,\\tau'}}{\\rho_0} \\int d^3p'\\frac{f_{\\tau'}(\\vec r,\\vec p')}{1+(\\vec p-\\vec p')^2/\\Lambda^2}. \\end{eqnarray} In the above, $\\tau=1/2$ ($-1/2$) for neutrons (protons) and $\\tau\\neq\\tau'$; $\\sigma=4/3$ is the density-dependence parameter; $f_{\\tau}(\\vec r,\\vec p)$ is the phase space distribution function at coordinate $\\vec{r}$ and momentum $\\vec{p}$. The parameters $B, C_{\\tau,\\tau}, C_{\\tau,\\tau'}$ and $\\Lambda$ are obtained by fitting the nuclear matter saturation properties \\cite{Das03}. The momentum dependence of the symmetry potential stems from the different interaction strength parameters $C_{\\tau,\\tau'}$ and $C_{\\tau,\\tau}$ for a nucleon of isospin $\\tau$ interacting, respectively, with unlike and like nucleons in the background fields. More specifically, $C_{unlike}=-103.4$ MeV while $C_{like}=-11.7$ MeV. The quantities $A_{u}(x)=-95.98-x\\frac{2B}{\\sigma +1}$ and $A_{l}(x)=-120.57+x\\frac{2B}{\\sigma +1}$ are parameters. The parameters $B$ and $\\sigma$ in the MDI single-particle potential are related to the $t_0$ and $\\alpha$ in the Gogny effective interaction via $t_0 = \\frac{8}{3} \\frac{B}{\\sigma+1} \\frac{1}{\\rho_0^{\\sigma}}$ and $\\sigma = \\alpha + 1$ \\cite{dec}. The parameter $x$ is related to the spin(isospin)-dependence parameter $x_0$ via $x=(1+2x_0)/3$ \\cite{Xuli10a}. On expanding the single-nucleon potential in $\\delta$, the first four terms are \\begin{eqnarray} && U_0(\\rho,k) = U_{n/p}|\\, _{\\delta=0} \\nonumber \\\\&& = \\frac{(A_l + A_u)}{2} \\frac{ \\rho} {\\rho_0} + B(\\frac{\\rho}{\\rho_0})^{\\sigma} + \\frac{2(C_{\\tau,\\tau}+C_{\\tau,\\tau'})}{\\rho_0} \\frac{2}{h^{3}}\\pi \\Lambda ^{3} \\nonumber  \\\\&& \\times \\left[ \\frac{p_{F}^{2}+\\Lambda^{2}-p^{2}}{2p\\Lambda }\\ln \\frac{(p+p_{F})^{2} +\\Lambda ^{2}}{(p-p_{F})^{2}+\\Lambda ^{2}} + \\frac{2p_{F}}{\\Lambda }-2\\tan ^{-1}\\frac{p+p_{f}} {\\Lambda }+2\\tan ^{-1}\\frac{p-p_{f}}{\\Lambda }\\right],\\nonumber\\\\ \\end{eqnarray} \\begin{eqnarray} &&U_{sym,1}(\\rho,k)=  \\pm \\frac{1}{1!} \\frac{\\partial U_{n/p}}{\\partial \\delta}|\\, _{\\delta=0} \\nonumber \\\\&& = \\frac{(A_l - A_u)}{2} \\frac{ \\rho} {\\rho_0} - 2 x\\frac{B}{\\sigma+1}\\frac{\\rho^{ \\sigma}}{\\rho_0^{\\sigma}} + \\frac{2(C_{\\tau,\\tau} - C_{\\tau,\\tau'})}{\\rho_0}  \\frac{2 p_F^2 \\pi \\Lambda^2} {3h^3p} \\ln \\frac{(p+p_{F})^{2} +\\Lambda ^{2}}{(p-p_{F})^{2}+\\Lambda ^{2}},\\nonumber\\\\ \\end{eqnarray} \\begin{eqnarray} && U_{sym,2}(\\rho,k)= \\frac{1}{2!} \\frac{\\partial^2 U_{n/p}}{\\partial \\delta^2}|\\, _{\\delta=0} \\nonumber \\\\&& =  -Bx (\\frac{\\rho}{\\rho_0})^{\\sigma} + \\frac{2Bx}{1+\\sigma} (\\frac{\\rho}{\\rho_0})^{\\sigma} \\nonumber \\\\&& + \\frac{(C_{\\tau,\\tau}+C_{\\tau,\\tau'})}{3\\rho_0} \\frac{p_F^2 \\pi \\Lambda^2} {9h^3p} \\left[ \\frac{4 p \\,p_F(p^2-p_F^2+\\Lambda^2)}{[(p+p_F)^2+\\Lambda^2][(p-p_F)^2+\\Lambda^2]} - \\ln \\frac{(p+p_{F})^{2} +\\Lambda ^{2}}{(p-p_{F})^{2}+\\Lambda ^{2}} \\right],\\nonumber\\\\ \\end{eqnarray} \\begin{eqnarray} &&U_{sym,3}(\\rho,k)=  \\pm \\frac{1}{3!} \\frac{\\partial^3 U_{n/p}}{\\partial \\delta^3}|\\, _{\\delta=0} \\nonumber\\\\ && =   -  \\frac{(C_{\\tau,\\tau}-C_{\\tau,\\tau'})}{3\\rho_0} \\frac{4 p_F^2 \\pi \\Lambda^2} {81h^3p} \\nonumber \\\\&& \\times \\left[ \\frac{2p{p_F}\\left(2p^6 - 3{p_F}^6 + 5{p_F}^2{\\Lambda}^4 + 2{\\Lambda}^6 + p^4\\left(-7{p_F}^2 + 6{\\Lambda}^2 \\right) + p^2\\left(8{p_F}^4 - 2{p_F}^2{\\Lambda}^2+6{\\Lambda}^4 \\right) \\right)}{[(p+p_F)^2+\\Lambda^2]^2[(p-p_F)^2+\\Lambda^2]^2} \\right.\\nonumber \\\\&& \\left.-   \\ln \\frac{(p+p_{F})^{2} +\\Lambda ^{2}}{(p-p_{F})^{2}+\\Lambda ^{2}}\\right]. \\end{eqnarray}  According to Eq.(\\ref{Esym2}), the second-order symmetry energy $E_{sym,2}(\\rho)$ is \\begin{eqnarray} &&E_{sym,2}(\\rho) =  \\frac{1}{3} t(k_F) + \\frac{1}{6} \\frac{\\partial U_0}{\\partial k}\\mid _{k_F}\\cdot k_F + \\frac{1}{2}U_{sym,1}(\\rho,k_F) \\nonumber \\\\= && \\frac{\\hbar^2}{6m} \\left(\\frac{3\\pi^2}{2}\\right)^{2/3} \\rho^{2/3} \\nonumber\\\\ + && \\frac{(C_{\\tau,\\tau}+C_{\\tau,\\tau'})}{3\\rho_0} \\frac{\\pi\\Lambda^2}{h^3}\\left[4p_F-\\left(2p_F+\\frac{\\Lambda^2}{p_F}\\right)\\textrm{ln}\\left(\\frac{4p_F^2+\\Lambda^2}{\\Lambda^2}\\right)\\right] \\nonumber \\\\+ && \\frac{(A_l - A_u)}{4} \\frac{ \\rho} {\\rho_0} - x\\frac{B}{\\sigma+1}\\frac{\\rho^{ \\sigma}}{\\rho_0^{\\sigma}} \\nonumber \\\\+ && \\frac{(C_{\\tau,\\tau}-C_{\\tau,\\tau'})}{3\\rho_0} \\frac{\\pi\\Lambda^2}{h^3}2p_F\\textrm{ln}\\left(\\frac{4p_F^2+\\Lambda^2}{\\Lambda^2}\\right), \\end{eqnarray} and according to Eq.(\\ref{Esym4}) the fourth-order symmetry energy $E_{sym,4}(\\rho)$ is \\begin{eqnarray} &&E_{sym,4}(\\rho) = \\frac{\\hbar^2}{162m} \\left(\\frac{3\\pi^2}{2}\\right)^{2/3} \\rho^{2/3} \\nonumber\\\\ & + & \\left[ \\frac{5}{324}\\frac{\\partial U_0(\\rho,k)}{\\partial k}|_{k_F} k_F - \\frac{1}{108} \\frac{\\partial^2 U_0(\\rho,k)}{\\partial k^2}|_{k_F} k_F^2 +\\frac{1}{648} \\frac{\\partial^3 U_0(\\rho,k)}{\\partial k^3}|_{k_F} k_F^3 \\right. \\nonumber \\\\& - & \\left. \\frac{1}{36} \\frac{ \\partial U_{sym,1}(\\rho,k)}{\\partial k}|_{k_F} k_F + \\frac{1}{72} \\frac{ \\partial^2 U_{sym,1}(\\rho,k)}{\\partial k^2}|_{k_F} k_F^2 + \\frac{1}{12} \\frac{\\partial U_{sym,2}(\\rho,k)}{\\partial k}|_{k_F} k_F + \\frac{1}{4} U_{sym,3}(\\rho,k_F)\\right] \\nonumber \\\\&=& \\frac{\\hbar ^2}{162m} (\\frac{3\\pi^2 \\rho}{2})^{\\frac{2}{3}} - \\frac{C_{\\tau,\\tau}}{3^{5}\\rho _{0}\\rho }\\left( \\frac{4\\pi }{h^{3}}\\right) ^{2}\\Lambda ^{2}\\left[ 7\\Lambda ^{2}p_{f}^{2}\\ln \\frac{4p_{f}^{2}+\\Lambda ^{2}}{\\Lambda ^{2}}-\\frac{4(7\\Lambda ^{4}p_{f}^{4}+42\\Lambda ^{2}p_{f}^{6}+40p_{f}^{8}}{(4p_{f}^{2}+\\Lambda ^{2})^{2}}\\right] \\nonumber \\\\&-&\\frac{C_{\\tau,\\tau'}}{3^{5}\\rho _{0}\\rho }\\left( \\frac{4\\pi }{h^{3}}\\right) ^{2}\\Lambda ^{2}\\left[ (7\\Lambda ^{2}p_{f}^{2}+16p_{f}^{4})\\ln \\frac{% 4p_{f}^{2}+\\Lambda ^{2}}{\\Lambda ^{2}}-28p_{f}^{4}-\\frac{8p_{f}^{6}}{\\Lambda ^{2}}\\right].\\nonumber\\\\ \\end{eqnarray} As one expects, the above expressions are identical to those derived directly from the exact MDI EOS using \\cite{Chen09} \\begin{eqnarray} E_{\\mathrm{sym},2}(\\rho ) &=&\\frac{1}{2!}\\frac{\\partial ^{2}E(\\rho ,\\delta )} {\\partial \\delta ^{2}}|_{\\delta =0}\\label{Esyme2} \\nonumber\\\\ E_{\\mathrm{sym,4}}(\\rho ) &=&\\frac{1}{4!}\\frac{\\partial ^{4}E(\\rho ,\\delta ) }{\\partial \\delta ^{4}}|_{\\delta =0}  \\label{Esyme4}. \\end{eqnarray}  \\begin{figure}[htb] \\centering \\includegraphics[width=12cm]{Graph4.EPS} \\caption{The kinetic energy part (T), the isoscalar potential part ($U_0$) and the isovector potential part ($U_{sym,1}$) of the symmetry energy $E_{sym,2}$ from the MDI interaction with $x=1$, 0 and -1.} \\label{MDIe2} \\end{figure}  In Fig.\\ \\ref{MDIe2}, we show the kinetic (T), isoscalar ($U_0$) and isovector ($U_{sym,1}$) potential contributions to $E_{sym,2}$ for the three different spin (isospin)-dependence parameter $x=1$, 0, and -1. We notice that the kinetic ($T$) and the isoscalar potential ($U_0$) contributions are the same for the three different $x$ values. As pointed out in Ref.\\ \\cite{Das03}, it is the isovector potential $U_{sym,1}$ that is causing the different density dependence of $E_{sym,2}$. For instance, with $x=1$ the $U_{sym,1}$ term decreases very quickly with increasing density and thus results in a super-soft symmetry energy at supra-saturation densities. On the contrary, the symmetry energy $E_{sym,2}$ at supra-saturation densities is very stiff for both $x=0$ and $x=-$1 as the contribution of the $U_{sym,1}$ term becomes very positive with smaller values of $x$.  \\begin{figure}[htb] \\centering \\includegraphics[width=12cm]{Graph5.EPS} \\caption{The kinetic energy and potential contributions to the fourth-order symmetry energy $E_{sym,4}$ from the MDI interaction.} \\label{MDIe4} \\end{figure}  Unlike the second-order term $E_{sym,2}$, the fourth-order symmetry energy $E_{sym,4}$ is independent of the spin (isospin)-dependence parameter $x$. Shown in Fig.\\ \\ref{MDIe4} are the various contributions to the fourth-order symmetry energy $E_{sym,4}$. Comparing these with the results obtained using the BGBD in Fig.\\ \\ref{BGBDe4}, we find that the $T$ and $U_0$ terms from these two interactions are almost identical. However, there exists some differences for other terms. For the MDI interaction, the $U_{sym,2}$ term is negative and becomes very important for determining $E_{sym,4}$. One the contrary, the contributions from the $U_{sym,1}$ and $U_{sym,3}$ terms are positive and they are relatively small as compared to $U_{sym,2}$. Generally, the behavior of $E_{sym,4}$ from the MDI interaction is very similar to that from the BGBD interaction for the case of $m_n^*>m_p^*$.  \\begin{figure}[htb] \\centering \\includegraphics[width=12cm]{Graph6.EPS} \\caption{The ratio of $E_{sym,4}$ over $E_{sym,2}$ obtained from the MDI interaction as a function of reduced density $\\rho/\\rho_0$ for $x=1$, 0, and -1.} \\label{MDIr} \\end{figure}  To compare more directly $E_{sym,4}$ with $E_{sym,2}$, their ratio $E_{sym,4}/E_{sym,2}$ is plotted in Fig.\\ \\ref{MDIr} as a function of reduced density for $x=1$, 0, and -1. It is seen that with $x=1$ there is a sharp break in the curve around $3\\rho_0$. This is because the second-order symmetry energy $E_{sym,2}$ changes from positive to negative around $3\\rho_0$ in this case. However, this is not the case for both $x=-1$ and $x=0$ where $E_{sym,2}$ remains positive at all densities.  In all cases, $E_{sym,4}$ is very small compared to $E_{sym,2}$. For both BGBD and MDI interactions, the small values of $E_{sym,4}$ up to several times the normal density clearly shows that the parabolic approximation of the EOS is well justified for most purposes. However, cares have to be taken in evaluating the core-crust transition density where the energy curvatures are involved \\cite{XCLM09}."
    },
    "1004/1004.4129_arXiv.txt": {
        "abstract": "{} {We propose a method to be able to decide whether the planets of CoRoT-7 are moving on mutually inclined orbits in the order of  $i>10^{\\circ}$.} {The extrasolar system CoRoT-7 is very special with respect to the closeness of the planets to the host star, which results in a fast dynamical development. It would therefore be possible to determine the inclination of the innermost planet CoRoT-7b with respect to the observer after an observation of at least three years from space with the satellite CoRoT with sufficient precision. Different inclinations would cause different duration of the transit times of a planet in front of the star and would therefore give us a better knowledge of the architecture of this system. With the aid of numerical integrations and analytical estimations we checked how inclined orbits of additional planets would change the transit duration of CoRoT-7b} {After 3 years of observations when an additional planet would be on a inclined orbit with respect to CoRoT-7b ($I_{mutual} > 10^{\\circ}$) an increase of the order of minutes could be observed for the transit duration. } {} ",
        "introduction": "Most of the more than 440 extrasolar planetary systems (=EPS)\\footnote{see http://exoplanet.eu} are single planetary systems,  but as of October 2010 there are some 50 multiplanetary systems, where between 2 and 7 planets are known to orbit their host stars. A very special case is the newly discovered system HD 10180 with seven planets orbiting an early G-type star \\citep{Lov10}. Certainly the small number of multiple EPS is a biased sample because it is highly improbable that there is just one planet around a star. From the theory of formation one expects that several planets may originate from the disk of gas and dust around a young star, like it is the case for the solar system.   An interesting fact is that many of the EPS-planets seem to move on large eccentric orbits and consequently -- when we expect more planets orbiting the star -- strong perturbations will act on the planets and therefore the  orbital elements may change significantly. Unfortunately our present methods to detect planets are just snapshots of the dynamical life time such that we don't have access to the dynamical evolution of a system. This is especially true for the detections via RV measuremenent where it is impossible to determine the inclinations of the orbital planes with respect to the observer. Even via transit observations of one planet combined with the measurement of radial velocities it is impossible to determine mutual inclinations of the orbits of planets when only one -- clearly the innermost -- is transiting. The main problem to be solved would be to observe such an EPS for sufficiently long times to be able to see more than one snapshot in the dynamical evolution of the system. We are now in the same situation in astronomy as we had been during the last centuries when scientists wanted to determine the proper motion of stars or the orbital elements of double stars. Within the EPS most of the planets, especially the ones which we observe via transits, are relatively close to the host star and they have relatively large masses. So we might hope to see a signature of additional planets influencing the orbit of the inner planet which may result in a change of the transits times and/or the duration of the transit. Up to now it was impossible to make such precise measurements using ground-based observation but with the possibility of using satellites the situation changed. The CoRoT spacecraft\\footnote{The CoRoT space mission has been developed and is operated by CNES with the contribution of Austria, Belgium, Brazil. ESA, Germany and Spain} (launched 2006) and the NASA Kepler mission  (launched 2009) permit very accurate measurement of light curves of transiting planets in terms of photometric precision and timing of transits. The recent extension of the CoRoT mission for another three years is important with respect to the observation of Transit Time Variations (TTV) and also possible Transit Duration Variations (TDV) on a time line of several years. ",
        "conclusions": "In this investigation we estimated the TDV caused by inclined additional planets which reaches -- although relatively small -- values of some degrees and is therefore detectable with the CoRoT satellite. This possibility stems from the  fact that the CoRoT-7 system is a rapidly evolving extrasolar system with very close-in planets.  We did our study using numerical integrations for the long term development of the system, as well as an analytical approach for short time intervals in the order of years. It turned out the system is quite stable even in the 3 planet model. A quantitatively new result is the dependence on the difference in the ascending node on the short time development of \\ccc b. The small phase shift in the dynamical development is not important for the qualitative behaviour for long time, but essential for the short time evolution of the orbit of \\ccc b.  In our determination of the duration of the transit caused by the change of the inclination of \\ccc b we also discussed what kind of orbits of the outer planet(s) would lead to a transit of these planets.  Despite the constraints given by the incompleteness of the data derived from observations the main conclusion of our study is that after three years of observation of the EPS CoRoT-7 we would be able to determine  -- via transit observations from space of the CoRoT satellite and additional ground based RV -- whether the two (three planets) are on mutually inclined orbits or whether they have just small inclinations of the order of  $i<10^{\\circ}$. More work has to be done for a more detailed analysis with numerical studies but also with analytical approaches \\citep[e.g.][]{Bor03,Bor07} comparable to those which have been undertaken for the investigation of the change in occultations of eclipsing binaries caused by an additional star."
    },
    "1004/1004.2822_arXiv.txt": {
        "abstract": "{} { We are carrying out a physical and chemical study of the protostellar envelopes in a representative sample of IM Class 0 protostars. In our first paper we determined the physical structure (density-temperature radial profiles) of the protostellar envelopes. Here, we study the CO depletion and N$_2$H$^+$ deuteration.} {We observed the millimeter lines of C$^{18}$O, C$^{17}$O, N$_2$H$^+$ and N$_2$D$^+$ towards the protostars using the IRAM 30m telescope. Based on these observations, we derived the C$^{18}$O, N$_2$H$^+$ and N$_2$D$^+$ radial abundance profiles across their envelopes using a radiative transfer code. In addition, we modeled the chemistry of the protostellar envelopes.} {All the C$^{18}$O 1$\\rightarrow$0 maps are well fit when assuming that the C$^{18}$O abundance decreases inwards within the protostellar envelope until the gas and dust reach the CO evaporation temperature, $\\approx$20--25~K, where the CO is released back to the gas phase. The N$_2$H$^+$ deuterium fractionation in Class 0 IMs is [N$_2$D$^+$]/[N$_2$H$^+$]=0.005--0.014,  two orders of magnitude higher than the elemental [D/H] value in the interstellar medium, but a factor of 10 lower than in prestellar clumps. Chemical models account for the C$^{18}$O and N$_2$H$^+$ observations if we assume the CO abundance is a factor of $\\sim$2 lower than the canonical value in the inner envelope. This could be the consequence of the CO being converted into CH$_3$OH on the grain surfaces prior to the evaporation and/or the photodissociation of CO by the stellar UV radiation. The deuterium fractionation % is not fitted by chemical models. This discrepancy is very likely caused by the simplicity of our model that assumes spherical geometry and neglects important phenomena like the effect of bipolar outflows and UV radiation from the star. More important, the deuterium fractionation is dependent on the ortho-to-para H$_2$ ratio, which is not likely to reach the steady-state value in the dynamical time scales of these protostars.  } {} ",
        "introduction": "Intermediate-mass young stellar objects (IMs) share many characteristics with high-mass stars (clustering, PDRs) but their study presents an important advantage: many are located closer to the Sun (d $\\leq$ 1 kpc) and in less complex regions than massive star-forming regions. On the other hand, they are also important for understanding planet formation since Herbig Ae stars are the precursors of Vega-type systems. Despite this, IMs have been studied very little so far. A few works on HAEBE stars have been carried out at millimeter wavelengths (Fuente et al. 1998, 2002; Henning et al. 1998), but almost nothing has been done for their precursors, the Class 0 IM objects.  Chemistry has been successfully used to determine the physical structure and investigate the formation and evolution of low-mass YSOs. Chemical diagnostics have also been shown to be good indicators of the protostellar evolution in these objects (see e.g. Maret et al. 2004, J{\\o}rgensen et al. 2005). However, very few works deal with IMs. Fuente et al. (2005a) present a chemical study of the envelopes of the Class 0 IM protostar NGC~7129--FIRS~2 and the young Herbig Be star LkH$\\alpha$~234. They find that the changes in the physical conditions of the envelope during its evolution from the Class 0 to the Class I stage (the envelope is dispersed and warmed up) strongly influence the molecular chemistry. The Class 0 object NGC~7129--FIRS~2 presented evidence of H$^{13}$CO$^+$ depletion. Moreover, the deuterium fractionation, measured as the DCO$^+$/H$^{13}$CO$^+$ ratio, decreases by a factor of 4 from the Class 0 to the Herbig Be star, very likely owing to the increase in the kinetic temperature. Regarding the abundance of complex molecules, the beam-averaged abundances of CH$_3$OH and H$_2$CO increase from the Class 0 to the Herbig Be star. A hot core was also detected in NGC~7129--FIRS~2 (Fuente et al. 2005a,b). Although two objects are not enough to establish firm conclusions, these pioneering results suggest that chemistry is also a good indicator of the evolution of IMs.  We are carrying out a chemical study of a representative sample of IM Class 0 YSOs This is the first systematic chemical study of IM Class 0  objects that has been carried out so far. Some properties like the temperature of the protostellar envelope and the clustering degree depend on the final stellar mass, so the results for low-mass stars cannot be directly extrapolated to intermediate-mass objects. In the first paper (Crimier et al. 2010, hereafter C10), we determined the physical structure (density-temperature radial profiles) by modeling the dust continuum emission. We now present the observations of the millimeter lines of C$^{18}$O, C$^{17}$O, N$_2$H$^+$, and N$_2$D$^+$ in the same sample. Our goal is to investigate the CO depletion and N$_2$H$^+$ deuteration in these Class 0 YSOs. For comparison, we also include 2 Class I objects, LkH$\\alpha$~234, and S140. ",
        "conclusions": "We carried out a study of the CO depletion and N$_2$H$^+$ deuteration in a sample of representative IM Class 0 protostars.  Our results can be summarized as follows.  \\begin{itemize} \\item We observed the millimeter lines of C$^{18}$O, C$^{17}$O, N$_2$H$^+$, and N$_2$D$^+$ using the IRAM 30m telescope in a sample of 7 Class 0 and 2 Class I IM stars. We have found a clear evolutionary trend that differentiates Class 0 from Class I sources. While the emission of the N$_2$H$^+$ 1$\\rightarrow$0 peaks towards the star position in Class 0 protostars, it surrounds the FIR sources in the case of Class I stars. This occurs because the recently formed star has heated and disrupted the parent core in the case of Class I objects. The deuterium fractionation, R$_2$, is low, below a few 0.001 in all Class I sources. There is, however, a wide dispersal in the values of R$_2$ in Class 0 sources ranging from a few 0.001 to a few 0.01. This  at least two orders of magnitude greater than the elemental value in the interstellar medium, although a factor of 10-100 lower than in prestellar clumps.  It is impossible, however, to establish an evolutionary trend among Class 0 sources based on simple parameters such as the average CO depletion and the average N$_2$H$^+$ deuterium fractionation. This stems from the complexity of these regions (multiplicity) and the limited angular resolution of our observations, which prevents us from tracing the inner regions of the envelope. Interferometric observations are required to provide a more precise picture of the evolutionary stage of these objects.  \\item We used a radiative transfer code to derive the C$^{18}$O, N$_2$H$^+$, and N$_2$D$^+$ radial abundance profiles in 5 IM Class 0 stars. In particular, we fit the C$^{18}$O 1$\\rightarrow$0 maps by assuming that the C$^{18}$O abundance decreases inwards within the protostellar envelope until the gas and dust reach the CO evaporation temperature, T$_{ev}$. Our observational data are better fit with values of T$_{ev}$$\\sim$20--25~K, consistent with the binding energy of 1100~K, which is the measured in the laboratory for a CO-CO matrix.  \\item We determined the chemistry of the protostellar envelopes using the model by Caselli et al. (2002). A  spherical envelope and steady-state chemical model cannot account for the observations. We had to introduce modifications to better fit the C$^{18}$O and N$_2$H$^+$ maps. In particular the CO abundance in the inner envelope seems to be lower than the canonical value. This could be due to the conversion of CO into CH$_3$OH on the grain surfaces, the photodissociation of CO by the stellar UV radiation, or even geometrical effects. Likewise, we have problems fitting the low values of the deuterium fractionation ($\\sim$ a few 0.001) measured for some Class 0 IMs. Several explanations have been proposed to account for this discrepancy. \\end{itemize}"
    },
    "1004/1004.2363_arXiv.txt": {
        "abstract": "{% We discuss all possible sources of uncertainties in theoretical values of the photometric amplitudes and phases of B-type main sequence pulsators. These observables are of particular importance because they contain information about the mode geometry as well as about stellar physics. Here, we study effects of various parameters coming both from theory of linear nonadiabatic oscillations and from models of stellar atmospheres. In particular, we show effects of chemical composition, opacities and, for the first time, effects of the NLTE atmospheres.} ",
        "introduction": "To construct a seismic model of a star, knowledge about the geometry of observed modes is a precondition. In the case of B-type pulsators, mode identification cannot be done directly from oscillation spectra because they are sparse and lack equidistant patterns. An alternative way is to make use of the fact that information about the mode degree, $\\ell$, and the azimuthal order, $m$, is embedded in the photometric and spectroscopic variations of a pulsating star. Ones of the most popular tools to identify a pulsation mode are the amplitude ratios and phase differences in various photometric passbands. In the case of zero-rotation approximation, these observables are independent of the azimuthal order, $m$, and inclination angle, $i$.  The semi-analytical expression for the bolometric light variation was formulated by Dziembowski \\cite {dziemb}. Balona \\& Stobie \\cite{balst} and Stamford \\& Watson \\cite{stwa} expanded this expression for the light variation in photometric passbands. They showed that modes with different values of $\\ell$ are located in separated parts on the amplitude ratio $vs.$ phase difference diagrams based on multicolour photometry. Subsequently, this method was applied to various types of pulsating stars by Watson \\cite{watson}. Cugier, Dziembowski \\& Pamyatnykh \\cite{cdp} improved the method by including nonadiabatic effects in calculations for the $\\beta$ Cephei stars. Effects of rotation on photometric observables were studied by Daszy\\'nska-Daszkiewicz et al.\\,(\\cite{ref3}) for close frequency modes and by Townsend \\cite{townsend2003} and Daszy\\'nska-Daszkiewicz, Dziembowski \\& Pamyatnykh \\cite{dd2007} for long-period g-modes.  The goal of this paper is to examine all possible effects on theoretical values of the photometric amplitude ratios and phase differences for early B-type pulsators. As an example, we consider the main sequence models with a mass of 10 $M_\\odot$ and low degree modes, $\\ell$=0, 1, 2. All effects of rotation on pulsation are neglected. In Section 2, we recall basic formulas and describe our models. Effects of parameters coming from linear non\\-adiabatic theory of stellar pulsation are presented in Section 3. Effects of atmospheric parameters are discussed in Section 4. The last section contains Conclusions. ",
        "conclusions": "The photometric amplitude ratios and phase differences of the early B-type main sequence pulsators strongly depend on chemical composition and opacities. The effects of the atmospheric parameters are smaller but may become important when the nonadiabatic parameter $f$ is determined from observations instead taken from the pulsation theory. Here, we studied for the first time the NLTE effects on photometric observables of the $\\beta$ Cep star model."
    },
    "1004/1004.5063_arXiv.txt": {
        "abstract": "Nonbarred ringed galaxies are relatively normal galaxies showing bright rings of star formation in spite of lacking a strong bar.  This morphology is interesting because it is generally accepted that a typical galactic disk ring forms when material collects near a resonance, set up by the pattern speed of a bar or bar-like perturbation. Our goal in this paper is to examine whether the star formation properties of rings are related to the strength of a bar or, in the absence of a bar, to the non-axisymmetric gravity potential in general. For this purpose, we obtained H$\\alpha$ emission line images and calculated the line fluxes and star formation rates (SFRs) for 16 nonbarred SA galaxies and four weakly barred SAB galaxies with rings. For comparison, we combine our new observations with a re-analysis of previously published data on five SA, seven SAB, and 15 SB galaxies with rings, three of which are duplicates from our sample.  With these data, we examine what role a bar may play in the {\\it star formation process} in rings. Compared to barred ringed galaxies, we find that the inner ring SFRs and H$\\alpha$+[NII] equivalent widths in nonbarred ringed galaxies show a similar range and trend with absolute blue magnitude, revised Hubble type, and other parameters.  On the whole, the star formation properties of inner rings, excluding the distribution of HII regions, are independent of the ring shapes and the bar strength in our small samples. We confirm that the deprojected axis ratios of {\\it inner rings} correlate with maximum relative gravitational force $Q_g$; however, if we consider all rings, a better correlation is found when local bar forcing at the radius of the ring, $Q_r$, is used.  Individual cases are described and other correlations are discussed. By studying the physical properties of these galaxies, we hope to gain a better understanding of their placement in the scheme of the Hubble sequence and how they formed rings without the driving force of a bar. ",
        "introduction": "A nonbarred ringed galaxy is a normal galaxy that shows one or more large-scale ring-shaped patterns in the absence of a conspicuous bar. These galaxies are interesting because they pose a small dilemma for existing dynamical theories: based on numerical simulations, a bar is believed to be an {\\it essential element} in ring formation (e.g., Schwarz 1981, 1984a; see also reviews by Sellwood \\& Wilkinson 1993; see also review by Buta \\& Combes 1996). These simulations have indicated that rings form easily in galaxies through bar-driven gravity torques.  These torques can lead to the secular accumulation of interstellar gas into well-defined regions, usually the main low-order resonances associated with the bar pattern speed. Invariant manifolds of orbits near unstable Lagrangian points have also been proposed to explain specific characteristics of galaxy rings (Romero-G\\'{o}mez et al. 2006; Athanassoula et al. 2009).  Alternate theories have been proposed to explain the presence of rings in nonbarred galaxies. One such theory suggests that nonbarred ringed galaxies were once more strongly barred, but the bar has mostly or completely dissolved, leaving behind the rings that formed when the bar was stronger. Bar dissolution might occur if a bar drives enough gas to the center to build up a central mass concentration (CMC; e.g., Shen \\& Sellwood 2004), or if gas exerts sufficient torques on the bar (Bournaud et al. 2005). The idea of bar dissolution and reformation is an active topic at this time (Bournaud \\& Combes 2002; Bournaud et al. 2005), but has not been unambiguously confirmed by observations.  In 2002, Block et al. found some evidence for bar dissolution, but a similar analysis done by Buta et al. (2004) was not able to confirm the results.  Buta (1991) suggested that NGC 7702 could be an example of a ringed galaxy whose main bar has dissolved because of the presence of a nuclear bar, a possible relic of a past major bar episode (Friedli \\& Pfenniger 1991). Athanassoula (1996) examined the effects of bar dissolution on a simulated outer ring, and found that the ring survives for a long time after the bar dissolves.  It is possible that the rings of some nonbarred galaxies formed in response to a strong spiral density wave (e.g., the SA(r)bc spiral NGC 5364). Mark (1974) suggested that a ring could form at the inner Lindblad resonance (ILR) of the spiral wave pattern speed.  Buta \\& Combes (1996) argued that such rings would be infrequent due to the inefficiency of the ring formation process in the presence of what might be an unsteady spiral.  However, Rautiainen \\& Salo (2000) concluded that in models with a hot disk that never formed a bar (Toomre $Q$-parameter=2.5), a spiral potential can still effectively form a ring at the spiral ILR.  Another possible explanation for ring formation in nonbarred galaxies is that some of these rings are in fact still bar-driven features, but the bar is only detectable in the near-infrared where the galaxy is more transparent to dust. Most galaxies are selected for given properties from catalogs that only give blue-light classifications (e.g., de Vaucouleurs et al. 1991; Sandage \\& Tammann 1981). Blue light is a very good passband for detecting rings, dust, and young stars, but bars tend to be made of older stellar populations that are less prominent in such a waveband.  Infrared light penetrates the dust more effectively and is more sensitive to the stellar population typically found in bars. A recently identified example of a ringed galaxy classified as nonbarred in blue light but which shows a bar in the near-IR is the SA(rs)bc spiral NGC 3147 (Casasola et al. 2008). In related cases, a bar might be missed even if potentially detectable in blue light because it is viewed end-on at a high inclination angle (e.g., the Sb(r) galaxy NGC 7184; Sandage \\& Tammann 1981).  It is also possible that a ring of a nonbarred galaxy can form through the minor merger of a companion.  Sil'chenko \\& Moiseev (2006) suggested that the small rings seen in NGC 7217 and NGC 7742, both known to have counter-rotating components, are due to tidally induced distortion of the stellar disk due to the (now mostly-merged) companion. Some rings, as in Hoag's object and IC 2006, have been interpreted in terms of accretion of a gas-rich companion (e.g., Schweizer et al. 1987, 1989). These ideas have also been considered to explain the large star-forming UV ring seen in the nonbarred galaxy ESO 381$-$47 (Donovan et al. 2009).  In addition to the mode of ring formation, another intriguing aspect of some nonbarred ringed galaxies is the presence of counter-winding spiral patterns in which two sets of spiral structure appear to open opposite to one another.  The implication, assuming both features are in the same plane and are not counter-rotating, is that one set of spiral structure is trailing while the other is leading. The best-known example of counter-winding spiral structure is the SA(r)a galaxy NGC 4622 (Buta et al. 1992, 2003), and we have recently identified a new case in the SA(rs)bc spiral ESO 297$-$27 (Grouchy et al. 2008, hereafter Paper 1).  In NGC 4622, the inner ring is made partly of leading and trailing arm segments, while in ESO 297$-$27, the inner pseudoring is made mainly of a single inner spiral arm. Although it has been suggested in these papers that such patterns could indicate that an interaction has occurred, no definitive theoretical studies have yet been made that can explain the special characteristics of each case.  In this paper, we examine the star formation properties in both barred and nonbarred ringed galaxies with several objectives in mind. First, concentrated star formation is an aspect that both barred and nonbarred galaxy rings share (Sandage 1961; Kormendy 1979; Buta 1988 and references therein; Pogge 1989; Buta et al. 2004). This would seem to point to a galaxy-wide mechanism for collecting gas into rings. We wish to examine whether any aspect of the ring-shaped star formation, such as the star formation rate (SFR), might connect to the strength of the non-axisymmetric perturbation. We are also interested in the triggering mechanism of star formation in rings. If bars actually trigger the star formation in rings, we might expect that barred galaxies will on average show higher ring SFRs. Previous studies have suggested that this is not the case. Some nonbarred or weakly barred galaxies are exceptional sites of star formation (e.g., NGC 4736 and 7742), and these make it unclear whether a bar is an essential element to the {\\it star formation process} occurring in rings. The problem is akin to spiral structure: do density waves trigger star formation in galaxies, or do they act merely as pattern-organizing structures with little or no role in triggering the collapse of clouds?  Elmegreen \\& Elmegreen (1986) and McCall (1986) independently concluded that density waves are not likely to be the main triggering mechanism for star formation in galaxies.  Our second objective is to re-examine how intrinsic ring shape connects to perturbation strength and ring star formation. Buta (2002) had suggested that ring shape does not depend significantly on bar strength, defined as the maximum gravitational torque per unit mass per unit square of the circular speed.  This finding is at odds with numerical simulation studies (e.g., Salo et al. 1999), and we consider an alternative approach that uses the strength of the perturbation at the position of the ring as the controlling parameter for ring shape.  Our final goal is to bring more attention to the diversity of properties of the rings in nonbarred galaxies. This diversity is greater than what is seen in barred galaxy rings, whose statistical properties, such as relative sizes, shapes, and orientations relative to the bar have been studied previously (Buta 1995).  Our study is based on an analysis of new H$\\alpha$+[NII] images of 20 nonbarred ringed galaxies (Grouchy 2008) complemented by optical $BVI$ images that will be more fully presented in a separate paper. In order to cover a wide range of apparent bar strengths, our sample is combined with a re-analysis of previously published data for 20 strongly barred and 12 weakly barred galaxies initially analyzed by Crocker et al. (1996, hereafter CBB96).  The main finding of CBB96 was a correlation between the way HII regions are distributed around inner rings and the intrinsic axis ratio of the rings. Circular inner rings have HII regions distributed more uniformly in azimuth than do elliptical inner rings. In the latter, HII regions concentrate around the ring major axis. We also further explore this issue here.  \\enlargethispage{\\baselineskip} H$\\alpha$+[NII] images are used to derive SFRs for both the global HII region distribution and for the HII regions confined to specific rings, using the flux conversion technique of Kennicutt (1983, 1998a). Previous studies (Finn et al. 2004; Kennicutt et al. 1994) have found connections between a galaxy's type and its SFR (see Finn's Figure 9). This paper complements these previous studies with more comparisons of the general properties of nonbarred ringed galaxies (color, absolute magnitude, equivalent width, and type) to their SFRs.  This paper is arranged as follows. In Section 2, we describe the galaxies in our sample as well as the CBB96 sample.  We discuss the observations and the process of data analysis, including a re-analysis of the CBB96 data. We also describe the process of calculating the galaxy's SFR as well as our estimates in the accuracy of the calibration.  Section 3 focusses on the derivation of ring parameters and includes a description of the H$\\alpha$ distribution for each galaxy observed. In Section 4, we explain the process of deriving the non-axisymmetric perturbations.  In Section 5, we discuss the analysis of the galaxy and ring properties addressing possible correlations. The discussion of our findings and conclusions can be found in Sections 6 and 7, respectively. ",
        "conclusions": "  \\noindent 1. Rings in nonbarred galaxies are often well-defined, narrow zones of HII regions and H$\\alpha$ emission. This is consistent with the enhanced blue colors seen in broadband images, indicating that the rings are well-organized, active zones of star formation, a trait which is very much in common with the more abundant rings seen in barred galaxies (CBB96). The best examples we illustrate here are found in ESO 111$-$22, ESO 198$-$13, ESO 231$-$1,ESO 236$-$29, ESO 526$-$7, IC 1993, NGC 7020, NGC 7187, NGC 7217, and NGC 7742.  \\noindent 2. For the 20 ringed galaxies illustrated in this paper, azimuthally averaged H$\\alpha$ surface brightness profiles follow the approximate shape of $B$-band azimuthally averaged profiles. This is in contrast to the finding of Ryder \\& Dopita (1994) that H$\\alpha$ profiles have a longer radial scale length than optical $V$ or $I$-band profiles.  \\noindent 3. The organized nature of the H$\\alpha$ emission from rings allows us to define the rings well enough to selectively integrate the fluxes and estimate SFRs {\\it for the rings alone}.  Combining our sample of 20 nonbarred ringed galaxies with a larger sample of ringed and mostly barred galaxies from CBB96, we were able to investigate possible correlations between ring SFRs and bar strength. The analysis showed that for a typical ringed galaxy having an absolute blue magnitude of $\\approx$$-$20, inner ring SFRs show little or no dependence on the strength of the nonaxisymmetric perturbation. There are galaxies showing exceptional star-forming rings with little or no trace of a bar. The few outer rings in the sample are consistent with the results from the inner rings.  \\noindent 4. Our combined sample allowed us to further investigate the correlation, if any, between bar strength and intrinsic inner ring shape. A previous study (Buta 2002) had suggested that inner ring shape did not depend on bar strength as defined by $Q_g$. Galaxies having similar values of $Q_g$ can have very different values of $q_{dep}$ (see, for example, Buta et al. 2007). When our sample of nonbarred ringed galaxies is considered, some correlation with $Q_g$ is found. However, a better correlation is found when a newly defined parameter, the maximum relative torque at the position of the ring, $Q_r$, is used. This is in agreement with the results from numerical simulations, and implies that the metric properties of the rings are dictated by the local bar strength.  \\noindent 5. For barred galaxies having $Q_g$ $\\geq$ 0.15, a better correlation is found between $q_{dep}$ and $Q_r/Q_g$ than between $q_{dep}$ and $Q_r$. For barred galaxies, this suggests that the controlling parameter for inner ring shape is the location of the ring relative to the bar maximum. If the ring major axis radius is at $\\approx$1.4$r(Q_g)$, an inner ring can be nearly circular, while if the ring major axis radius is at $\\approx$1.1$r(Q_g)$, the ring is practically on top of the bar and is highly elongated. $Q_r/Q_g$ loses its usefulness when a bar is weak or absent.  \\noindent 6. The combined barred/nonbarred sample includes some intrinsically very large rings. A comparison of the deprojected linear diameters of inner rings in our sample with an earlier analysis of nearby galaxies by Buta \\& de Vaucouleurs (1982) shows that the average nonbarred ring diameters are comparable to those of the barred galaxies. In the Buta \\& de Vaucouleurs analysis, SA inner rings averaged a factor of 2 smaller than SB inner rings. The emphasis on large SA rings in our sample is not unexpected because of the way the sample was selected.  \\noindent 7. We verify previous analyses which have shown that nonbarred ringed galaxies are a diverse and relatively inhomogeneous class of objects whose ring formation mechanisms are likely varied. Evolved bar resonance rings, spiral density wave resonance rings, interaction-produced rings, and even the limitations of $B$-band images for recognizing bars could account for most of the features observed in our combined sample.  R.G. and R.B. have been supported by NSF grant AST 050-7140 to the University of Alabama. R.G. is now supported by an NSF International Research Fellowship (OISE-0852959).  E.L., H.S., and R.G. acknowledge the support of the Academy of Finland. This research has made use of the NED, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with NASA. We thank Drs$.$ Fran\\c{c}oise Combes, Lia Athanassoula, Patrick Treuthardt, and the anonymous referee for their helpful comments. We also thank NOAO and CTIO for use of the 1.5 m telescope.  \\clearpage"
    },
    "1004/1004.0700_arXiv.txt": {
        "abstract": "Low-ionization (\\MgII, \\FeII, \\FeIII) broad absorption line quasars (LoBALs) probe a relatively obscured quasar population, and could be at an early evolutionary stage for quasars.  We study the intrinsic fractions of LoBALs using the SDSS, 2MASS, and FIRST surveys.  We find that the LoBAL fractions of the near infra-red (NIR) and radio samples are approximately 5--7 times higher than those measured in the optical sample.  This suggests that the fractions measured in the NIR and radio bands are closer to the intrinsic fractions of the populations, and that the optical fractions are significantly biased due to obscuration effects, similar to high-ionization broad absorption line quasars (HiBALs).  We also find that the LoBAL fractions decrease with increasing radio luminosities, again, similar to HiBALs. In addition, we find tentative evidence for high fractions of LoBALs at high NIR luminosities, especially for FeLoBALs with a fraction of $\\sim$18 per cent at $M_{K_s} < -31$~mag.  This population of NIR luminous LoBALs may be at an early evolutionary stage of quasar evolution. We use a two-component model of LoBALs including a pure geometric component and a luminosity dependent component at high NIR luminosities, and obtain better fits than those from a pure geometric model. Therefore, the LoBAL population can be modelled as a hybrid of both the geometric and evolutionary models, where the geometric component constitutes 3.4$\\pm$0.3, 5.8$\\pm$0.4, and 1.5$\\pm$0.3 per cent of the quasar population for BI-LoBALs, AI-LoBALs, and FeLoBALs, respectively. Considering a population of obscured quasars that do not enter the SDSS survey, which could have a much higher LoBAL fraction, we expect that intrinsic fraction of LoBALs could be even higher. ",
        "introduction": "Broad absorption line quasars (BALQSOs) are a sub-sample of quasars exhibiting blue-shifted absorption troughs (e.g., Weymann et al.\\ 1991). In the traditional definition of Weymann et al.\\ (1991), absorption troughs must be at least 2000~\\kms\\ wide excluding the first 3000~\\kms\\ region blue-ward from the emission lines to classify quasars as BALQSOs. Less strict definitions have also been used, for example with a requirement of a trough to be at least 1000~\\kms\\ wide (e.g., Trump et al.\\ 2006). BALQSOs can also be further divided into a population containing absorption troughs from only high-ionization state species (e.g., \\CIV\\ and \\NV; HiBALs) and a population that exhibits absorption troughs in low-ionization species (e.g., \\MgII\\ and \\FeII; LoBALs). The majority of BALQSOs are HiBALs. In fact, all LoBALs also contain the high-ionization troughs in their spectra (e.g., Weymann et al.\\ 1991; Trump et al.\\ 2006). Besides the presence of low-ionization troughs, the optical continua of LoBALs are more reddened compared to HiBALs, suggesting stronger dust extinction (e.g, Sprayberry \\& Foltz 1992; Reichard et al.\\ 2003). In X-rays, LoBALs also have higher gas absorption column densities than HiBALs (e.g., Green et al.\\ 2001; Gallagher et al.\\ 2002). Therefore, LoBALs probe a relatively obscured quasar population. The origin of a small LoBAL fraction in quasars is unclear, and it has been attributed to geometric effects (e.g., Elvis et al.\\ 2000) or evolutionary effects (e.g., Voit et al.\\ 1993), like the BALQSO population in general. There are several tentative arguments supporting the view that LoBALs are young quasars at a stage of blowing out obscuring materials. First, several early studies of LoBAL fractions in the infra-red band showed larger fractions (e.g., Boroson \\& Mayers 1992) and associations with ultra-luminous infra-red galaxies (e.g., L{\\'{i}}pari et al.\\ 1994; Canalizo \\& Stockton 2000). Second, a few optical spectral analyses suggested that the covering fraction of the LoBAL wind is large (e.g., Voit et al.\\ 1993; Casebeer et al.\\ 2008). Third, some radio spectra of LoBALs resemble those of compact steep spectrum or gigahertz peaked spectrum sources, which are also candidates for young quasars (e.g., Montenegro-Montes et al.\\ 2008; Liu et al.\\ 2008). In particular, the LoBALs containing Fe absorption troughs (FeLoBAL) are viewed as the most promising candidates for young quasars (e.g., L{\\'{i}}pari et al.\\ 2009). The fractions of LoBALs are important constraints on the origin of the LoBAL populations.  Before studying the intrinsic fractions of LoBALs, it is important to compare the measurements of the intrinsic fractions of BALQSOs in the quasar population. Recently, a series of studies emerged on this topic. Dai, Shankar, and Sivakoff (2008a) studied the SDSS BALQSO (Trump et al.\\ 2006) fractions in the 2MASS bands (Skrutskie et al.\\ 2006), finding that the BALQSO fractions in the near-infrared (NIR) are twice those found in the optical band. In particular, we found the BALQSO fraction to be 20$\\pm$2 per cent for the traditional BALQSOs that satisfy the stricter Weymann et al. (1991) definition and 43$\\pm$2 per cent for the relaxed definition of Trump et al. (2006) that requires less broad absorption troughs. Dai et al.\\ (2008a) argued that the BALQSO fractions measured in the NIR bands are closer to the intrinsic fractions, based on the observations that significant obscuration is associated with BALQSOs in the optical bands (e.g., Reichard et al.\\ 2003), confirming the earlier estimate of Hewett \\& Foltz (2003). This result was confirmed by several studies, such as Ganguly \\& Brotherton (2008) using a different SDSS BALQSO catalogue, Maddox et al.\\ (2008) using the deeper UKIDSS survey, Shankar et al. (2008a) in the radio band, and Knigge et al.\\ (2008) by correcting the fraction in the optical bands directly. In particular, Ganguly \\& Brotherton (2008) extended the study to include narrow and associated absorbers and found the overall outflowing AGNs to be 60 per cent of the total quasar population. Recently, Allen et al.\\ (2011) claimed an even larger intrinsic fraction for the traditional BALQSOs of 41 per cent, by including the additional fraction of missing quasars that do not enter the SDSS survey. In particular, Allen et al.\\ (2011) found that the completeness for BALQSOs and non-BALQSOs in SDSS is very similar at $z < 2.1$ and $z > 3.6$, but can differ at other redshifts, e.g., for $z\\sim2.6$ and $z\\sim3.5$.  The larger fraction of BALQSOs makes the AGN wind a more promising candidate responsible for the feedback energy that is needed to explain the co-evolution between black holes and their host galaxies. Understanding evolutionary versus geometric models of AGNs can not only probe the AGN feedback, but also constrain the nature of the feedback, whether it is kinetic from winds (e.g., Granato et al.\\ 2004, Shankar et al.\\ 2006, 2008b) or thermal (e.g., Di Matteo et al.\\ 2005; Hopkins et al.\\ 2006). If BALQSOs provide the majority of the feedback energy, the feedback mechanism will be kinetic. Motivated by the results from BALQSOs, and the larger obscuration of LoBALs compared to HiBALs, we expect that the optical fractions for LoBALs are also biased low. This effect has already been noticed when only a few LoBALs were observed (Sprayberry \\& Foltz 1992); however, their study was limited by their small sample size. The large sample size enabled by SDSS warrants a new study on the intrinsic fractions of LoBALs.  The radio properties of BALQSOs provide additional constraints on the nature of these objects. In an early study, Stocke et al.\\ (1992) found no radio-loud BALQSOs within 68 BALQSOs. Later studies showed that radio emission is present in BALQSOs (Francis et al.\\ 1993; Brotherton et al.\\ 1998; Becker et al.\\ 2000); however, these BALQSOs are mostly radio-moderate. Matching the SDSS BALQSO catalogue in the FIRST survey (Becker et al.\\ 1995), Shankar et al.\\ (2008a) quantified the dependence of the BALQSO fraction on radio luminosities. We found that the BALQSO fraction drops at high radio luminosities confirming earlier claims of such an effect (Stocke et al.\\ 1992; Becker et al.\\ 2001). In addition, Shankar et al.\\ (2008a) also found that the BALQSO fraction at the low radio luminosity range is consistent with the NIR fraction of BALQSOs of Dai et al.\\ (2008a). This result further supports the view that the NIR BALQSO fraction is close to the intrinsic fraction (modulo corrections to the parent SDSS quasar selection), since there is also little absorption in the radio bands. The drop of the BALQSO fractions at high radio luminosities can be naturally explained under a geometric model of BALQSOs. If the radio emission has a preferred orientation, which is usually considered in the polar direction, the drop indicates that BALQSOs are less frequent in these viewing angles. Using a unification model between radio-loud and radio-quiet quasars (e.g., Urry \\& Padovani 1995), we were able to successfully reproduce the trend, thus explaining the majority of the BALQSOs with radio emission under a geometric model. Exceptions still exist, such as the polar radio-loud BALQSOs, which were identified based on the radio variability that implies too large brightness temperatures unless the radio emission is relativistic (e.g., Zhou et al.\\ 2006; Ghosh \\& Punsly 2007). Cold polar outflows are also present in blazars (e.g., Dai et al.\\ 2008b), and they could be related to polar BALQSOs as the outflow continuously extends close to the polar axis. However, these objects are rare, and their implication for the total BALQSO population is still uncertain.  In this paper, we study the intrinsic fraction of LoBALs by correlating SDSS quasars with detections in the NIR and radio bands. We also explore their radio properties and compare with BALQSOs to test whether LoBALs can be explained under a geometric or evolutionary model. We assume that $H_0 = 70~\\rm{km~s^{-1}~Mpc^{-1}}$, $\\Omega_{\\rm m} = 0.3$, and $\\Omega_{\\Lambda}= 0.7$ throughout the paper.  \\begin{figure*} \\includegraphics[width=17.5truecm]{f1.eps} \\caption{(top left) 2MASS $K_s$ magnitude versus SDSS $i$ magnitude for SDSS DR3 QSOs that are detected in all of the $J$, $H$, and $K_s$ bands in the redshift range of $0.5 \\le z \\le 2.15$. The subsamples of QSOs that are not LoBALs (non-LoBALs; black dots), QSOs that satisfied the traditional Weymann et al.\\ definition (BI-LoBALs; red triangles), and QSOs that satisfied the relaxed BAL definition of Trump et al.\\ (2006), but did not satisfy the traditional definition (AINB-LoBALs; green squares) are displayed separately. (top right) SDSS $i$ magnitude versus 1.4~GHz flux density. (bottom left) 2MASS $K_s$ magnitude versus 1.4~GHz flux density.  (bottom right) SDSS $r -$ 2MASS $K_s$ colour for the three samples, where the histograms for the two LoBAL samples are multiplied arbitrarily by 30 for clarity.  The LoBALs are significantly redder than the non-LoBAL population.  The inset in the bottom right panel shows the cumulative distributions of the colour for the three samples.  The K--S test results indicate that both of the LoBAL samples differ from the non-LoBAL sample with significances greater than 99.996 per cent, and that the two LoBAL samples are not significantly different from each other. \\label{fig:one}} \\end{figure*}  \\begin{figure*} \\includegraphics[width=17.5truecm]{f2.eps} \\caption{Same plots as Fig.~\\ref{fig:one}, but for FeLoBals in the redshift range of $1.19 \\le z \\le 2.24$.  The FeLoBALs are even more significantly redder than the non-LoBAL population.  The K--S test result indicates that the $r-K_s$ colours of the FeLoBAL population differ from those of the non-LoBAL population with a significance greater than $1-10^{-8}$. \\label{fig:two}} \\end{figure*} ",
        "conclusions": "\\label{sec:discuss} We find significantly high fractions of LoBALs in the quasar population compared to the values obtained using optical data only. For example, the BI-LoBAL and AI-LoBAL fractions in the optical data were measured as 0.55 and 1.31 per cent (T06), while our results are 5--7 times larger. Although the final intrinsic fractions depend on the choice of catalogues, the overall trend is found to be the same. For example, we perform a similar analysis to the BI-LoBALs from the G09 sample, and also find large intrinsic LoBAL fractions. Although there is a systematic offset between the BI-LoBAL fractions from the T06 and G09 samples, the overall trend for the observed fractions is the same. Our intrinsic fractions are obtained using two independent methods, one from the NIR and optical data and the other from the radio data, and the results are mutually consistent between the two methods. Combining the estimates from the two methods using the least square (minimum variance) method, we find that the intrinsic fractions for BI-LoBALs, AI-LoBALs, and FeLoBALs are $4.0\\pm0.5$, $7.1\\pm0.6$, and $2.1\\pm0.3$ per cent, respectively, using a pure geometric model.  For our hybrid model with both the geometric and evolutionary components, which fit the data better, the corresponding intrinsic fractions for the geometric component are $3.4\\pm0.3$, $5.8\\pm0.4$, and $1.5\\pm0.3$ per cent. The intrinsic fractions for the evolutionary component for LoBALs are functions of luminosities, and the total intrinsic fractions of LoBALs are the sums of the two components. The final combined estimates of the intrinsic LoBAL fractions are listed in Table~\\ref{tab:cfrac}. The results are not unexpected considering the large obscuration observed in LoBALs (e.g., Sprayberry \\& Foltz 1992). Although we find significantly larger intrinsic fractions of LoBALs, they still represent a small portion of the total population. Compared with the intrinsic fractions of $20\\pm2$ and $43\\pm2$ per cent for BI-BALQSOs and AI-BALQSOs (Dai et al.\\ 2008a), the LoBALs are about 20 per cent of BALQSOs.  Our method of calculating the intrinsic fractions of BALQSOs still depends on the completeness of optical quasar surveys.  The fraction of quasars that do not enter the optical surveys was not be accounted in this or our previous papers.  It was estimated that SDSS is about 90 per cent complete at $z < 2.2$ and $i < 19.1$~mag (Richards et al.\\ 2002).  The remaining 10 per cent of quasars, which are thought to be highly obscured, could potentially all be LoBALs or FeLoBALs.  Thus, in most optimistic estimates, the fractions of LoBALs or FeLoBALs can reach $\\sim$15 per cent for the geometric component.  However, the nature of the obscured quasars are still uncertain, and we are not sure whether the ultra-violet spectra of these quasars still show broad absorption lines, if the continua of these quasars are mostly obscured. Therefore, the fractions quoted in our paper represent conservative estimates based on observations.  The LoBAL fractions in the radio band are particularly interesting, since we find that the LoBAL fractions decrease with increasing radio luminosities.  This confirms the early result of Becker et al.\\ (2000) with about a dozen LoBALs. The trend is similar to that found in the total BALQSO population (Shankar et al.\\ 2008a). This trend found in both the total BALQSO population and LoBALs suggests that the majority of LoBALs and BALQSOs can be united under a similar physical scheme. In Shankar et al.\\ (2008a), we favoured a geometric model to interpret the trend. Applying the geometric model of Shankar et al.\\ (2008a) to LoBALs, we can successfully reproduce the data for BI-LoBALs, AI-LoBALs, and FeLoBALs. However, it is problematic under a pure evolutionary model to explain the radio-luminosity/LoBAL fraction trend as argued by Shankar et al.\\ (2008a).  The main question is why quasars spend the same fraction of time as BALQSOs in the radio-loud stage and radio-quiet stage. Therefore, we argue that the majority of BALQSOs and LoBALs can be understood in a geometric model. Polar BALQSOs/outflows (Zhou et al.\\ 2006; Ghosh \\& Punsly 2007; Dai et al.\\ 2008b) present a challenge to our results; however, these objects are rare and we are uncertain about their implications to the total quasar population. A modification of the geometric model to have both disc and polar outflows (e.g., Borguet \\& Hutsem{\\'e}kers 2010) may be needed to incorporate these objects. We also note that LoBALs dominate this population of polar BALQSO candidates (Ghosh \\& Punsly 2007).  There are other indications, such as the association with ULIRGs, radio spectra, and large covering fractions from spectral modeling, arguing that LoBALs belong to an earlier evolutionary stage of quasar population. However, most of these studies are based on a small sample size and may not extrapolate to the whole LoBAL population. Urrutia et al.\\ (2009) studied the fraction of LoBALs in the dust reddened quasars at high redshift, finding that all except one are LoBALs, supporting the young nature of LoBALs. However, the authors also noted that their selection method may be biased favouring LoBALs since they are associated with large dust reddening. In our study, we find LoBALs and BALQSOs are similar in most aspects, except that the fraction of LoBALs increases with increasing NIR luminosities. This is not consistent with BALQSOs in general, because the BALQSO fractions are mostly constant with increasing NIR luminosities for AI-BALQSOs (Dai et al.\\ 2008a). At the NIR luminous end, the observed LoBAL fraction is higher, although with large uncertainties, than the intrinsic fraction that we obtain for a pure geometric model.  It is possible that a portion of NIR luminous LoBALs are special compared to the rest of the population, and at an early evolutionary stage of the quasar cycle. This will reconcile some observations supporting the young quasar interpretation for LoBALs. We add an evolutionary LoBAL component to our model, and generally obtain a better fit to the data. However, the parameters of this evolutionary component are unconstrained from the data. The combination of a short early evolution model, with a large covering fraction close to 100 per cent, and a subsequent, longer geometric model, with covering fractions consistent with the intrinsic fraction measured in this paper and Dai et al.\\ (2008a), could simultaneously account for many of the differences and similarities between LoBALs and HiBALs. The longer-lived geometric model mainly sets the intrinsic fractions of BALQSOs of various species and the measured fraction as a function of their radio luminosities. If there is some small spherical outflow component at early times, this might also explain the predominance of LoBALs among the rare polar BALQSOs. Deeper IR and radio surveys are needed to increase the sample size and confirm this claim."
    },
    "1004/1004.2711.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract {We perform a sensitive (line confusion limited), single-side band spectral survey towards Orion KL with the IRAM 30m telescope, covering the following frequency ranges: 80-115.5 GHz, 130-178 GHz, and 197-281 GHz. We detect more than 14 400 spectral features of which 10 040 have been identified up to date and attributed to 43 different molecules, including 148 isotopologues and lines from vibrationally excited states. In this paper, we focus on the study of OCS, HCS$^+$, H$_2$CS, CS, CCS, C$_3$S, and their isotopologues. In addition, we map the OCS $J$=18-17 line and complete complementary observations of several OCS lines at selected positions around \\object{Orion IRc2} (the position selected for the survey). We report the first detection of OCS $\\nu_2$ = 1 and $\\nu_3$ = 1 vibrationally excited states in space and the first detection of C$_3$S in warm clouds. Most of CCS, and almost all C$_3$S, line emission arises from the hot core indicating an enhancement of their abundances in warm and dense gas. Column densities and isotopic ratios have been calculated using a large velocity gradient (LVG) excitation and radiative transfer code (for the low density gas components) and a local thermal equilibrium (LTE) code (appropriate for the warm and dense hot core component), which takes into account the different cloud components known to exist towards Orion KL, the \\textit{extended ridge}, \\textit{compact ridge}, \\textit{plateau}, and \\textit{hot core}. The vibrational temperature derived from OCS $\\nu_2$ = 1 and $\\nu_3$ = 1 levels is $\\simeq$210 K, similar to the gas kinetic temperature in the hot core. These OCS high energy levels are probably %mainly pumped by absorption of IR dust photons. We derive an upper limit to the OC$_3$S, H$_2$CCS, HNCS, HOCS$^+$, and NCS column densities. Finally, we discuss the D/H abundance ratio and infer the following isotopic abundances: $^{12}$C/$^{13}$C = 45$\\pm$20, $^{32}$S/$^{34}$S = 20$\\pm$6, $^{32}$S/$^{33}$S = 75$\\pm$29, and $^{16}$O/$^{18}$O = 250$\\pm$135.}  %  % context heading (optional) %  % {} leave it empty if necessary %   {To investigate the physical nature of the `nuc\\-leated instability' of %   proto giant planets, the stability of layers %   in static, radiative gas spheres is analysed on the basis of Baker's %   standard one-zone model.} %  % aims heading (mandatory) %   {It is shown that stability %   depends only upon the equations of state, the opacities and the local %   thermodynamic state in the layer. Stability and instability can %   therefore be expressed in the form of stability equations of state %   which are universal for a given composition.} %  % methods heading (mandatory) %   {The stability equations of state are %   calculated for solar composition and are displayed in the domain %   $-14 \\leq \\lg \\rho / \\mathrm{[g\\, cm^{-3}]} \\leq 0 $, %   $ 8.8 \\leq \\lg e / \\mathrm{[erg\\, g^{-1}]} \\leq 17.7$. These displays %   may be %   used to determine the one-zone stability of layers in stellar %   or planetary structure models by directly reading off the value of %   the stability equations for the thermodynamic state of these layers, %   specified %   by state quantities as density $\\rho$, temperature $T$ or %   specific internal energy $e$. %   Regions of instability in the $(\\rho,e)$-plane are described %   and related to the underlying microphysical processes.} %  % results heading (mandatory) %   {Vibrational instability is found to be a common phenomenon %   at temperatures lower than the second He ionisation %   zone. The $\\kappa$-mechanism is widespread under `cool' %   conditions.} %  % conclusions heading (optional), leave it empty if necessary %   {} ",
        "introduction": "\\label{sect_int}  The Orion KL (Kleinmann-Low) cloud is the closest ($\\simeq$ 414 pc, \\citealt{men07}) and most well studied high mass star-forming region in our Galaxy (see, e. g., \\citealt{gen89} for review). The prevailing chemistry of the cloud is particularly complex as a result of the interaction of the newly formed protostars, outflows, and their environment. The evaporation of dust mantles and the high gas temperatures produce a wide variety of molecules in the gas phase that are responsible for a spectacularly prolific and intense line spectrum \\citep{bla87, bro88, cha97}.  %\\bf{ %Early studies %soon realized that the large scale distribution of gas and %dust was heavily influenced by violent phenomena such %as the interaction of compact and large scale outflows %with the quiescent gas producing strong line and continuum %emission. Near- and mid-IR subarcsecond resolution imaging and (sub)millimeter interferometric observations have identified the main sources of luminosity, heating, and dynamics in the region. At first, IRc2 was believed to be the responsible for this complex environment. %of luminosity, heating and dynamics of the region. However, %Near- and mid-IR subarcsecond %resolution imaging and (sub)millimeter interferometric %observations have s our understanding %of the region. First, However, the 8-12 $\\mu$m emission peak of IRc2 is not coincident with the % Orion SiO maser %origin (related to the origin of the outflow(s) (and the Orion SiO maser origin), and its intrinsic IR luminosity (L$\\approx$1000 L$\\sun$) is only a fraction of the luminosity of the entire system (\\citealt{gez98}). In addition, %gezari et al. 1998 3.6-22 $\\mu$m images indicate that IRc2 is resolved into four non self-luminous components. Therefore, IRc2 is not presently the powerful engine of Orion KL and its nature remains unclear \\citep{dou93, shu04, gre04}. %(Dougados et al. 1993; Shuping et al. %2004; Greenhill et al. 2004). %A new step forward was %given by \\citet{men95} identified the very embedded radio continuum source I (a young star with a very high luminosity without an infrared counterpart, $\\simeq$10$^5$ L$_\\odot$, \\citealt{gez98, gre04}, located 0''.5 south of IRc2) as the source coinciding with the centroid of the SiO maser distribution \\citep{pla09, zap09a, god09b}. They also detected the radio continuum emission of IR source $n$, suggesting this source as another precursor of the large-scale phenomena. In addition, \\citet{beu04} detected a sub-millimeter source without IR and centimeter counterparts, SMA1, previously predicted by \\citet{dev02}, which may be the source driving the high velocity outflow \\citep{beu08}. % and suggested that it could also contribute to %the origin of some of the phenomena observed at larger %scales. Thus, the core of Orion KL contains the compact HII regions $I$ and $n$ (in addition to BN, which was resolved with high resolution at 7 mm by \\citealt{rod09}), which appear to be receding from a common point, an originally massive stellar system that disintegrated $\\simeq$500 years ago \\citep{gom05, zap09b}. %Thus, in addition to BN, the core of Orion KL %contains at least two more compact HII regions (I and %n) that seem to be running away from a common point, %suggesting that BN, I and n were originally part of a %common massive stellar system that disintegrated $\\simeq$500 %years ago \\citep{gom05}. %(G\u00a1\u00e4omez et al. 2005). Finally, submm aperture synthesis line surveys provided the spatial location and extent of many molecular species \\citep{bla96, wri96, liu02, beu05, god09b, pla09, zap09a}. %(Blake et al. 1996; %Wright et al. 1996; Liu et al. 2002; Beuther et al. 2005).  The chemical complexity of Orion KL has been demonstrated by several line surveys performed at different frequency ranges: 72.2-91.1 GHz by \\citet{joh84}; 215-247 GHz by \\citet{sut85}; 247-263 GHz by \\citet{bla86}; 200.7-202.3, 203.7-205.3 and 330-360 GHz by \\citet{jew89}; 70-115 GHz by \\citet{tur89}; 257-273 GHz by \\citet{gre91}; 150-160 GHz by \\citet{ziu93}; 325-360 GHz by \\citet{sch97}; 607-725 GHz by \\citet{sch01}; 138-150 GHz by \\citet{lee01}; 159.7- 164.7 GHz by \\citet{lee02}; 455-507 GHz by \\citet{whi03}; 795-903 GHz by \\citet{com05}; 44-188 $\\mu$m by \\citet{ler06}; 486-492, 541-577 GHz by \\citet{olo07} and \\citet{per07}; and 42.3-43.6 GHz by \\citet{god09a}.  In spite of this large amount of data, no line confusion limited survey has been carried out so far with a large single dish telescope. We performed such a line survey towards Orion IRc2 with the IRAM 30-m telescope at wide frequency ranges (a total frequency coverage of $\\simeq$ 168 GHz). Our main goal was to obtain a deep insight into the molecular content and chemistry of the Orion KL, an archetype high mass star-forming region (SFR), and to improve our knowledge of its prevailing physical conditions. It also allows us to search for new molecular species and new isotopologues, as well as the rotational emission of vibrationally excited states of molecules already known to exist in this source. Since the amount and complexity of the data is large, we divided our analysis into families of molecules so that model development and discussions could be more focused. In this paper, we concentrate on sulfur carbon chains, in particular carbonyl sulfide OCS (see previous studies by \\citealt{gol81}; \\citealt{eva91}; \\citealt{wri96}; \\citealt{cha97}), CS (previously analyzed by \\citealt{has84}; \\citealt{mur91}; \\citealt{zen95}; \\citealt{wri96}; \\citealt{joh03}), H$_2$CS (\\citealt{min91}; \\citealt{gar85}), HCS$^+$, CCS, CCCS, and their isotopologues.  Column density calculations, and therefore the estimation of isotopic abundance ratios and molecular excitation, have improved, with respect to previous works, due to the much larger number of available lines, their consistent calibration across the explored frequency range, the up-to-date information about the physical properties of the region and molecular constants, and the use of a LVG radiative transfer code to derive reliable physical and chemical parameters. Modeled brightness temperatures obtained from a fit to all observed lines have been convolved with the telescope beam profile, assuming a given size for each cloud component, to provide accurate source-averaged, and not beam-averaged, molecular column densities.  After presenting the line survey (Sects. \\ref{sect_obs} and \\ref{sect_sur}), this work concentrates on the detection of OCS, HCS$^+$, H$_2$CS, CS, CCS, and CCCS lines and their analysis, as well as providing upper limits to the abundance of several non-detected sulfur-carbon-chain bearing molecules such us OC$_3$S, H$_2$CCS, HNCS, HOCS$^+$, and NCS (Sects. \\ref{sect_res} to \\ref{sect_vib}). This is the first of a series of papers dedicated to the analysis of the millimeter emission from different molecular families towards Orion KL. ",
        "conclusions": "\\label{sect_dis}  The power of spectral line surveys at different mm and sub-mm wavelengths to search for new molecular species and derive the physical and chemical structure of molecular sources has been demonstrated (\\citealt{bla87}; \\citealt{sut95}; \\citealt{cer00}; \\citealt{sch01}; \\citealt{par07}). The main and final goal of our line survey is to provide a consistent set of molecular abundances derived from a systematic analysis of the molecular rotational transitions. Our line survey allows us to obtain with unprecedented sensitivity and completeness the census of the identified and unidentified molecules in Orion KL. These kinds of studies are necessary to understand the chemical evolution of this archetypal star-forming region. Moreover, that many rotational transitions of the same molecule have been observed in different frequency ranges (the 3 mm window illustrates more clearly the extended ridge component, whereas the 1.3 mm one identifies the warmest gas at the hot core and along the compact ridge), provide strong observational constraints on the source structure, gas temperature, gas density, and molecular column densities.  \\subsection{Molecular abundances} \\label{sect_dis_abu}  Molecular abundances were derived using the H$_2$ column density calculated by means of the C$^{18}$O column density provided in Sect. \\ref{sect_col_cs}, assuming that CO is a robust tracer of H$_2$ and therefore their abundance ratio is roughly constant, ranging from CO/H$_2$ $\\simeq$ 5$\\times$10$^{-5}$ (for the ridge components) to 2$\\times$10$^{-4}$ (for the hot core and the plateau). In spite of the large uncertainty in this calculation, we include it as a more intuitive result for the molecules described in the paper. We obtained N(H$_2$) = 7.5$\\times$10$^{22}$, 7.5$\\times$10$^{22}$, 2.1$\\times$10$^{23}$, and 4.2$\\times$10$^{23}$ cm$^{-2}$ for the extended ridge, compact ridge, plateau, and hot core, respectively. In addition, we assume that the H$_2$ column density spatially coincides with the emission from the species considered. Our estimated source average abundances for each Orion KL component are summarized in Table \\ref{tab_abun} (only available online), together with comparison values from other authors (\\citealt{sut95} and \\citealt{per07}). The differences between the abundances shown in Table \\ref{tab_abun} are mostly due to the different H$_2$ column density considered, to the assumed cloud component of the molecular emission and discrepancies in the sizes of these components.  \\subsection{Column density ratios} \\label{sect_dis_rat}  \\begin{table*} %t19 \\begin{center} \\caption{Column density ratios\\label{tab_rat}} \\resizebox{0.9\\textwidth}{!}{% \\begin{tabular}{l|lllll|lllll} \\hline \\hline Column & & This & work & & & $^{(1)}$ & $^{(1)}$ & $^{(2)}$ & Dark clouds & Hot core\\\\ Density ratio & ER & CR & P & HC & Total & A & B &  & (TMC1) & \\object{G327.3-0.6}\\\\ \\hline OCS/CS & 3 & 2 & 5 & 2 & 3 & 5 & 13 & 6(HC) & 0.5 & $<$4\\\\ CS/HCS$^+$ & 7700 & 100 & 50 & 420 & 180 & 1000 & 270 & ... & 10 & ...\\\\ CS/H$_2$CS & 14 & 50 & 2 & 12 & 7 & 25 & 8 & 4(CR) & 6 & $>5$\\\\ CS/CCS & 4000 & 1100 & 1200 & 420 & 530 & ... & ... & ... & 0.5 & 325\\\\ CS/C$_3$S & ... & ... & ... & 1050 & 1050 & ... & ... & ... & 4 & ...\\\\ CCS/C$_3$S & ... & ... & ... & 2.5 & 2.5 & ... & ... & ... & 8 & ...\\\\ CO/CS & 5000 & 500 & 10000 & 2500 & 2800 & 66700 & 133300 & 2400(P) & 20000 & 26200$^{(3)}$\\\\ HCO$^+$/HCS$^+$ & 270 & 4 & 9 & 27 & 13 & 12 & 10 & ... & 20 & ...\\\\ H$_2$CO/H$_2$CS & 12 & 8 & 18 & 11 & 12 & 6250 & 1250 & 15(CR) & 71 & $>3$\\\\ \\hline \\end{tabular} } \\end{center} $^{(1)}$: \\citet{nom04}\\\\ $^{(2)}$: \\citet{per07}\\\\ $^{(3)}$: assuming $^{16}$O/$^{17}$O = 2625\\\\ Note.-Derived column density ratios and comparison with other works and sources. Column 1 gives the considered ratio, Cols. from 2 to 6 show the results obtained in this work in the different spectral cloud components of Orion and the total value, Cols. 7 and 8 ratios obtained by \\citet{nom04} in their models of hot cores (in model B trapping of mantle molecules in water ice is assumed), Col. 9 gives values in Orion by \\citet{per07}, Col. 10 in dark clouds (TMC-1) (\\citealt{wal09} and references within), and Col. 11 provides these ratios for the molecular hot core G327.3-0.6, \\citet{gib00}.\\\\ \\end{table*}  To compare the chemistry of the different spectral cloud components related to sulfur-bearing carbon chains molecules, we derived the column density ratios showed in Table \\ref{tab_rat}. This table also shows the ratios found in chemical models of hot cores, other results found in the literature for Orion, and other sources (the dark cloud TMC-1 and the hot core G327.3-0.6). We found good agreement between our ratios and those derived by \\citet{per07}, both set of values corresponding to Orion KL. For the other molecular hot core, we noted a large difference in the ratio CO/CS. This discrepancy also occurs with the chemical models computed by \\citet{nom04}. We note that the chemical models cannot provide realistic values for the H$_2$CO/H$_2$CS column density ratio, as we have discussed in Sect. \\ref{sect_col_h2cs}. TMC-1 exhibits ratios very different by those of hot cores, as expected from their different chemical and physical conditions.  We find $N$(C$^{34}$S/OC$^{34}$S)$\\simeq$0.3, 0.6, 0.2, and 0.2 in the extended ridge, the compact ridge, the plateau, and the hot core, respectively. The chemical models for hot cores computed by \\citet{nom04} infer that $N$(CS)/$N$(OCS) = 0.2 (at 10$^{4}$ years). The $N$(CS)/$N$(CCS) abundance ratio is 300, 1143, 1000, and 280 for the extended ridge, the compact ridge, the plateau, and the hot core, respectively. For the hot core, we also derive $N$(CS)/$N$(C$_3$S)$_{Hot\\;\\;Core}$=700. Both CCS and C$_3$S have not been studied in the chemical models available for hot cores. As expected, these values are very different from those derived in the dark cloud TMC-1 for which $N$(CS)/$N$(CCS)=2.2 and $N$(CS)/$N$(C$_3$S)=7.8 \\citep{hir92}. However, we obtain $N$(C$_2$S)/$N$(C$_3$S) = 2.5, very similar to the 3.4 value derived by \\citet{hir92} in TMC-1 (cyanopolyyne peak) and the value of $\\simeq$ 3 found in the envelope of the C-rich star IRC+10216 by \\citet{cer87a}.  This is a surprising result because CCS is considered to be a typical molecule in cold dark clouds. Moreover, C$_3$S is found only in the hot core, which is indicative of an enhancement in the production of CCS and C$_3$S in the warm and dense gas. Although spectral confusion is large when observing weak lines such as those of C$_3$S, thanks to our survey, we detected 17 lines. They cover from the J=14-13 (E$_{up}$=29.1 K with v$_{LSR}$=4.2 kms$^{-1}$) up to J=47-46 (E$_{up}$=313 K with v$_{LSR}$=4.2 kms$^{-1}$), thus, we are fully confident in its detection. In addition, the observed velocities correspond definitively to the hot core.  Our results indicate that C$_3$S is efficiently formed in warm regions. That the C$_2$S/C$_3$S abundance ratio is similar to that of dark clouds or evolved stars may indicate that these species formed in the gas phase. Gas phase chemical models predict C$_2$S/C$_3$S $\\simeq$ 2 and 0.3 in TMC-1 and IRC10216, respectively (\\citealt{wal09}; \\citealt{cor09}). %In both models the effects of molecular anions have been included)  %The nature of the hot core has been widely discussed by several %authors. The amount of -CN bearing molecules and the lack of %many carbon-rich chains are proves that the hot core is dominated by a %a O-rich chemistry (REFERENCIAS). %There are clear indications of that the hot core is composed of %condensations with very different temperatrures (this work, %\\citealt{dev02}, \\citealt{wri96}). %However, there are controversies in different points: %Is the hot core is internally %heated or the radio source I (considered the main heating source %in the region \\citealt{men95}, \\citealt{gez98}, \\citealt{bla96}) %heats this molecular component?. \\citealt{bla96} have pointed out that %there is no %evidence of internal heating within the molecular hot core, %based on the distribution of the HC$_3$N J = 24-23 line in %the 1v7 vibrationally excited state. However, \\citealt{kau98} and %\\citealt{dev02} proposed that it is interanlly heated by young embedded %proto-stars. The work of \\citealt{dev02} was positive indicating that the %radio souce I may not dominate %the heating of the hot core, based on the distribution in the hot core of the J %= 10-9 rotational line of HC$_3$N in the vibrationally excited level %1v5. \\subsection{Orion KL cloud structure} \\label{sect_dis_str}  We have analyzed and discussed the emission lines of the studied molecules in terms of the four well-known Orion KL cloud components (hot core, extended ridge, compact ridge, and plateau). However, low angular resolution does not enable us to detect any possible %the vibrational temperature calculated here suggests variation in the excitation temperature across the hot core and the other Orion components. A more complex physical structure has been observed with sensitive interferometers \\citep{wri96, beu05, pla09, zap09a}. Further analysis of our survey indicates that at the position of IRc2, the lines of both SiS and the SiO maser emission show a velocity component at 15.5 km s$^{-1}$, an additional cloud component to those described above. Owing to the high energies involved in some emission lines, \\citet{sch01} and \\citet{com05} claimed that a hotter component exists at the hot core v$_{LSR}$ in their surveys at high frequency. In the same way, we detected the emission of vibrationally excited OCS and CS at the hot core LSR velocity. %that cloud %corresponds to the feature we have detected in SiS and SiO and in the %vibrational states of SiO, CS and OCS. In spite of the low angular resolution of our data, the amount of molecules, the large number of transitions, and the different vibrationally excited states found in the survey permit us to derive realistic source-averaged physical and chemical parameters."
    },
    "1004/1004.0857_arXiv.txt": {
        "abstract": "Force balance considerations put a limit on the rate of  AGN radiation momentum  output, $L/c$, capable of driving galactic  superwinds and reproducing the observed $\\mbh -\\sigma $ relation between black hole mass and spheroid velocity dispersion. We show that  black holes  cannot supply enough momentum in radiation to drive the gas out by pressure alone. Energy-driven winds give a  $\\mbh -\\sigma $ scaling favoured by a recent analysis but also fall short energetically once cooling is incorporated. We propose that outflow-triggering of star formation by enhancing the intercloud medium turbulent pressure and squeezing clouds can supply the necessary boost, and  suggest possible tests of this hypothesis. Our hypothesis simultaneously can account for the observed halo baryon fraction. ",
        "introduction": "\\begin{figure} \\epsscale{.8} \\plotone{rc1.eps} \\caption{Plot of $\\eta \\mbh c/(\\mg \\sigma c)$ (for $\\eta=0.1$)  versus $\\sigma$ from the sample of Gultekin et al 09. The horizontal line  is $0.1\\mbh=\\mg \\sigma /c.$ Black hole momentum falls short of the required momentum for most galaxies. } \\label{fig1} \\end{figure} There is a consensus that  the powerful  Active Galactic Nuclei (AGNs)  play a crucial role in shaping the general properties of  galaxies \\citep[e.g.,][]{SR,Croton,SN09} and clusters of galaxies \\citep[e.g.,][]{Vecc04,NSB,MN07}. AGNs are powered by accretion onto supermassive black holes believed to reside at the centers of most galaxies. An indication of the galaxy-black hole connection is the remarkable correlation between the black hole  mass, $\\mbh$, and the velocity dispersion, $\\sigma$, of the spheroidal galactic components \\citep[e.g.,][]{gult09}. Any successful model for galaxy formation must provide an explanation of this correlation. Self-regulated black hole growth offers a natural explanation for this relation \\citep{SR}. Both radiation pressure and mechanical outflows deposit momentum into the protogalactic gas.  If this results in a wind, force balance arguments (\\citep{Fabian99,King03,Murray,Thompson05} but see \\citep{Soker09}) lead to the  conclusion that winds driven by pressure of radiation from a central black hole can suppress the collapse of gas and hence regulate the growth of the black hole. However the available momentum is, we show, insufficient to give the required normalization of the $\\mbh-\\sigma$ scaling (Section 2).    The original self-regulation argument of \\cite{SR} relied on energy balance: AGN activity heats  the galactic  gas reservoir above the virial temperature, generating galactic winds and eventually terminating gas accretion onto the black hole. However, energy-driven winds suffer strong radiative cooling losses: while the  radiation  heats the gas nearby the black hole, the gas expands but cools rapidly, making the process inefficient (Section 3). Our preferred solution is to introduce positive AGN feedback via triggered star formation. We argue that this simultaneously resolves three problems: the required order-of-magnitude boost in the $\\mbh-\\sigma$ scaling (Section 4), the enhanced specific star formation rate in massive galaxies (addressed elsewhere by  \\cite{sadegh}, and the shortfall in the halo baryon fraction (Section 5). ",
        "conclusions": "The primary aim of our   paper is to highlight the scaling relation  normalisation problem for AGN feedback,  and to  propose a possible solution involving AGN-triggered star formation. Positive feedback may have important ramifications for star formation at high redshift, and  is inevitably followed by gas outflows driven by both AGN and supernovae, along with concomitant quenching of star formation.  From the data plotted in figure 1,  the momentum boosting by the starburst is a factor of a few. This will naturally yield the dispersion in the relations given the nature of the boost, e.g. by BH outflow  triggering of SNII.  The points that lie low  in the momentum condition had a larger boost, and this would  lead to a prediction that the residuals in  $\\mbh c$ vs $M_g\\sigma$ should anticorrelate with SNII tracers in chemical evolution, e.g. the  bulge $\\alpha/Fe$.   Small galaxies which formed stars before host galaxy AGN onset will survive. They should be seen as a bump in the galaxy luminosity function (GLF), analogously to what is seen in the MW \\citep{Koposov} and in the  K-band GLF \\citep{Smith}. These galaxies are distinguishable by being older and more metal-poor than their AGN-modulated successors which are primarily either low mass satellites or  massive early-type galaxies. For the MW, the failure of the \\cite{Koposov} model tuned to the numerous ultrafaint dwarfs to account for the admittedly sparse numbers of massive dwarfs is consistent with the lack of a large BH in our AGN feedback model. Feedback from IMBH can resolve this problem. We have suggested that the IMBH in $\\omega$Cen may be an example of a population of  halo IMBH that could have provided the additional feedback needed to both allow the LMC and similar dwarfs  to form and not  simultaneously overproduce the faint dwarfs. Such IMBH could easily, during an active accretion phase, have produced enough momentum to have swept the residual gas out of the outer halo.  Globular clusters are plausibly the most visible surviving component of the first generation of substructure. That they might have a direct connection to IMBH is weakly supported  by the possibility that  one of the most massive globular clusters, $\\omega$Cen,  might contain an IMBH. Another hint of a connection with globular clusters may be present in the apparent correlation  between black hole mass and mass of the host galaxy globular cluster system \\citep{Spitler}. A variation on this relation has recently been found that relates black hole mass to the number of globular clusters \\citep{Burkert2010}. Numerical simulations find that the SMBH-$\\sigma$ scaling relation can be preserved by hierarchical mergers of IMBH \\citep{Burkert}. This lends support to the possibility that globular clusters may serve as a proxy both for IMBH and for dwarf galaxies, and therefore provide a possible witness to the required baryonic cleansing role of satellites by IMBH  in our model."
    },
    "1004/1004.0299_arXiv.txt": {
        "abstract": "We present a sample of 68 low-$z$ \\ion{Mg}{2} low-ionization broad absorption-line (loBAL) quasars. The sample is uniformly selected from the Sloan Digital Sky Survey Data Release 5 according to the following criteria: (1) redshift $0.4<z\\leq0.8$, (2) median spectral $S/N>7~$pixel$^{-1}$, and (3) \\ion{Mg}{2} absorption-line width $\\Delta v_{c} \\geq 1600~$\\kms. The last criterion is a trade-off between the completeness and consistency with respect to the canonical definition of BAL quasars that have the `balnicity index' $BI>0$ in \\ion{C}{4} BAL. We adopted such a criterion to ensure that $\\sim 90\\%$ of our sample are classical BAL quasars and the completeness is $\\sim 80\\%$, based on extensive tests using high-$z$ quasar samples with measurements of both \\ion{C}{4} and \\ion{Mg}{2} BALs.  We found (1) \\ion{Mg}{2} BAL is more frequently detected in quasars with narrower H$\\beta$ emission-line, weaker [\\ion{O}{3}] emission-line, stronger optical \\ion{Fe}{2} multiplets and higher luminosity. In term of fundamental physical parameters of a black hole accretion system, loBAL fraction is significantly higher in quasars with a higher Eddington ratio than those with a lower Eddington ratio. The fraction is not dependent on the black hole mass in the range concerned. The overall fraction distribution is broad, suggesting a large range of covering factor of the absorption material. (2) [\\ion{O}{3}]-weak loBAL quasars averagely show undetected [\\ion{Ne}{5}] emission line and a very small line ratio of [\\ion{Ne}{5}] to [\\ion{O}{3}]. However, the line ratio in non-BAL quasars, which is much larger than that in [\\ion{O}{3}]-weak loBAL quasars, is independent of the strength of the [\\ion{O}{3}] line. (3) loBAL and non-loBAL quasars have similar colors in near-infrared to optical band but different colors in ultraviolet. (4) Quasars with \\ion{Mg}{2} absorption lines of intermediate width are indistinguishable from the non-loBAL quasars in optical emission line properties but their colors are similar to loBAL quasars, redder than non-BAL quasars. We also discuss the implication of these results. ",
        "introduction": "About 15\\% of quasars show broad absorption lines (BALs) of high ionization ions such as \\ion{N}{5}, \\ion{C}{4}, \\ion{Si}{4}, Ly$\\alpha$, \\ion{O}{6}, up to a velocity of $v\\sim 0.1~c$. BALs are detected occasionally (another $\\sim$15\\%) also in low ionization species such as \\ion{Mg}{2}, \\ion{Al}{3}. Two very different scenarios have been proposed to explain the BAL phenomenon. The first scenario, namely `unification model', suggests that BAL and non-BAL quasars are physically the same, and attributes their different appearance solely to different line of sight. According to the unification model, every quasar has a BAL region (BALR) with a covering factor of 10\\%-20\\%, and our line of sight passes through BALR only in BAL quasars, plausibly at low inclination angles (Tolea et al. 2002; Hewett \\& Foltz 2003; Reichard et al. 2003b; Trump et al. 2006; Gibson et al. 2009, hereafter G09). The second, so called evolutionary scenario, suggests that BAL quasars are in the early stage of quasar evolution with a gas/dust richer nuclear environment(Sanders et al. 1988; Hamann \\& Ferland 1993; Voit et al. 1993; Egami et al. 1996; Becker et al. 2000; Trump et al. 2006).  On the one hand, there are many pieces of observational evidence for the unification of BAL and non-BAL quasars, including the similarity of emission line spectrum between the two classes of quasars (Weymann et al. 1991), a small covering factor of BALR inferred from emission line profiles (Korista et al. 1993), spectropolarimetric observations of BAL quasars (e.g., Goodrich \\& Miller 1995; Cohen et al. 1995; Hines \\& Wills 1995; Ogle et al. 1999; Schmidt \\& Hines 1999), and the great similarity of the  spectral energy distribution (SED) between BAL and non-BAL quasars in the infrared to millimeter waveband (e.g., Willott et al. 2003; Gallagher et al. 2007). The notoriously weak X-ray emission from BAL quasars is often ascribed to strong absorption in the BAL direction, which is also supported by X-ray spectroscopy (e.g., Green et al. 1995; Brinkmann et al. 1999; Wang et al. 1999; Brandt et al. 2000; Gallagher et al. 2002, 2006; Fan et al. 2009). On the other hand, there are observations that cannot be understood in the simple unification scenario. First, radio morphology and radio variability study showed that BAL quasars are not observed at any preferred direction with respect to the radio axis (Jiang \\& Wang 2003; Brotherton et al. 2006; Zhou et al. 2006b; Ghosh \\& Punsly 2007; Wang et al. 2008a). Second, Boroson (2002) found that BAL quasars on average have higher Eddington ratios than non-BAL quasars in a small sample of BAL QSOs. A similar conclusion has been reached by Ganguly et al. (2007). It has also been suggested that BAL quasars are redder and more luminous than other quasars (Reichard et al. 2003b, Trump et al. 2006, cf., G09). The latter results indicate that BAL and non-BAL quasars can be unified, but the covering factor of BALR depends on their nuclear parameters. A wide range of covering factor of BALR has also been implied by comparison of the optical polarization between BAL and non-BAL quasars (Wang et al. 2005).  Significant differences between low-ionization BAL (loBAL) and high-ionization BAL (HiBAL)/non-BAL quasars are also seen in dust extinction and far-infrared emission. loBAL quasars show a redder spectrum than HiBAL and non-BAL quasars on average, consistent with a reddening of $E(B-V)\\sim$0.1 for a SMC-like dust extinction curve (Weymann et al. 1991; Richards et al. 2003). Dai et al. (2008) showed that BAL fraction among Two Micron All Sky Survey(2MASS) selected quasars are as high as $\\sim$ 44\\% (cf. Ganguly \\& Brotherton 2008). Surprisingly, when going down to a low flux limit in near infrared, Maddox \\& Hewett (2008) found a similar 30\\% fraction of BAL quasars. This indicates that BAL quasars on average are heavily reddened and thus many red BAL quasars have been overlooked in optical spectroscopic surveys. This can be interpreted physically in two very different scenarios: either BAL quasars are a distinct population with the nuclei preferring a gas and dust rich environment; or dust is preferably distributed in the outflow direction as suggested by the dusty disk wind models for BALR (Konigl \\& Kartje 1994), and overall covering factor is 30\\%. Isotropic properties, such as far-infrared emission, are of great importance to distinguish between the two. Boroson \\& Meyers (1992) found that the fraction of loBAL quasars in a small far infrared selected sample is much higher than that in optically selected samples. They also found that these quasars show weak [\\ion{O}{3}] and strong optical \\ion{Fe}{2} emission lines. A large sample of low-$z$ loBAL quasars are needed to confirm these findings.  Low-$z$ BAL quasars are of great interest also because a number of important spectral diagnostics can be accessed via the ground optical spectroscopic observations, such as narrow emission-lines (NELs), Balmer and \\ion{Fe}{2} broad emission lines (BELs). We can also inspect the properties of the host galaxies much easier at low-$z$. However, previous studies mainly focused on high-$z$ BAL or HiBAL quasars due to the rarity of loBAL quasars. With the advent of large area spectroscopic surveys, such as the Sloan Digital Sky Survey (SDSS; York et al. 2000), it is possible to perform a systematic study of low-$z$ BAL quasars based on a large sample.  In this paper, we present a sample of 68  \\ion{Mg}{2} BAL quasars at  $0.4<z\\leq0.8$ uniformly selected from the quasar sample of SDSS Data Release 5 (Schneider et al. 2007), and make a comparison study with non-BAL quasars. This paper is oraganized as follows. We analyze the SDSS spectra and compile the sample in \\S2. Various definition criteria have been adopted in previous studies, and different criteria result in different fractions of BAL quasars. We have made extensive tests for different criteria in order to obtain a more objective definition of loBAL quasars. The tests are described in detail in this section, together with the definition criterion that is used in our sample. In \\S3, we compare the properties between loBAL and non-BAL quasars. The implication of our findings is discussed in \\S4. Throughout  this paper, we assume a $\\Lambda$-dominated cosmology with $H_0=72$\\kms~Mpc$^{-1}$, $\\Omega_M=0.28$, and $\\Omega_\\Lambda=0.72$. ",
        "conclusions": "We have selected 68 \\ion{Mg}{2} BAL quasars from the SDSS spectroscopic quasar sample with a redshift of $0.4<z\\leq0.8$ and a median $S/N \\geq 7$ using criteria of a continuous absorption over a velocity interval greater than 1600 \\kms~for a depth of at least 10\\%. The BAL-selected criterion is a trade-off between the completeness and consistency with respect to the canonical definition of BAL quasars that have the 'balnicity index' $BI>0$ in \\ion{C}{4} BAL. We adopted such a criterion to ensure that $\\sim 90\\%$ of our sample are classical BAL quasars and the completeness is $\\sim 80\\%$, based on extensive tests using high-$z$ quasar samples with measurements of both \\ion{C}{4} and \\ion{Mg}{2} BALs. The low-$z$ sample is used to define the fraction of \\ion{Mg}{2} BAL quasars and its dependence on the continuum and emission line properties, the difference between \\ion{Mg}{2} and non-\\ion{Mg}{2} BAL quasars. We find that, (1) the fraction of \\ion{Mg}{2} BAL quasars in the optical survey is around 1.2\\%. The fraction does not include correction for internal dust extinction of BAL quasars and the color bias against reddened loBAL quasars. After correcting these factors, the true \\ion{Mg}{2} BAL fraction is likely in between 2\\% and 7\\%. (2) \\ion{Mg}{2} BAL quasars are more frequently found in quasars with low [\\ion{O}{3}] equivalent width and high continuum luminosity although they show a wide range of [\\ion{O}{3}] equivalent width. loBAL quasars display stronger narrow optical \\ion{Fe}{2} emission lines and UV \\ion{Fe}{2} emission, weaker even absent [\\ion{Ne}{5}] lines. (3) The fraction of quasars with \\ion{Mg}{2} BAL increases strongly with the Eddington ratio but does not correlate with the black hole mass. (4) There is an excess of intrinsic reddening in \\ion{Mg}{2} BAL quasars and quasars with intermediate width \\ion{Mg}{2} absorption lines with an average of 0.08 mag for SMC-like dust grain. In this section, we will discuss the implication of our results.%  It is generally believed that NELs are nearly isotropic in quasars because they are produced in an extended region. In contrast the BEL region and continuum are thought to be much more compact and can be blocked on some lines of sight by obscuration (e.g., Antonucci 1993). The excess extinction in \\ion{Mg}{2} BAL quasars with respect to non-\\ion{Mg}{2} BAL quasars will enhance the conclusion that [\\ion{O}{3}] equivalent width is lower in loBAL quasars. The conclusion will be further strengthened if we consider the anisotropic emission of optical continuum from the accretion disk because it is generally believed that BAL quasars are seen nearly edge-on. The large range of observed $EW_{[O\\;III]}$ among \\ion{Mg}{2} BAL quasars suggests that loBAL occurs in both [\\ion{O}{3}]-weak and [\\ion{O}{3}]-strong emission quasars, but the covering factor decreases strongly as the [\\ion{O}{3}] strength increases.  If [\\ion{O}{3}] is considered as an indicator of overall strength of NELs, the strength of other lines relative to [\\ion{O}{3}] should connect more to the physical state of narrow-line region (NLR) or ionizing continuum. The absence of [\\ion{Ne}{5}] in [\\ion{O}{3}]-weak \\ion{Mg}{2} BAL turns out to be a rather surprise. Previous studies have suggested that [\\ion{Ne}{5}] is produced in the high density, inner NLR (e.g., Heckman et al. 1981; De Robertis \\& Osterbrock 1984; Whittle 1985a, 1985b). Lack of [\\ion{Ne}{5}] emission indicates that there is no such region or the inner NLR does not expose to a hard ionizing continuum. Interaction of BAL outflow with inner NLR may destroy dense clouds in the inner NLR. However, the presence of strong narrow optical \\ion{Fe}{2} emission would suggest such dense inner NLR does exist but with low ionization parameters (see V{\\'e}ron-Cetty et al. 2004; also Wang, Dai \\& Zhou 2008). Then we look at the option that the NLR only sees a soft ionizing continuum. BAL quasars, loBAL quasars in particular, are weak in soft X-rays (Green et al. 1995; Brinkmann et al. 1999), which are required to produce  Ne$^{4+}$ (ionization potential 97 eV). If NLR sees a continuum similar to the observed one, the absence of [\\ion{Ne}{5}] can be naturally explained. Because the weakness of soft X-rays in BAL quasars is usually attributed to X-ray absorption rather than intrinsic weakness (Wang et al. 1999; Gallagher et al. 1999, 2002), this requires that the [\\ion{Ne}{5}] emission region is behind the X-ray absorber. There are two possibilities for this, the outflow has a large covering factor or [\\ion{Ne}{5}] emission region located coincidently  behind the outflow. Nagao et al. (2001) proposed that [\\ion{Ne}{5}] emission region is the inner region of the dust torus, and use the hypothesis to explain the unification of two type Seyfert galaxies. In their scenarios, [\\ion{Ne}{5}] emission region lies on the equatorial plane, which is coincident with the region shielded by disk wind.  In order to check this, we select 250 non-\\ion{Mg}{2} BAL quasars with $S/N >20$, $\\beta_{[3K,4K]} > -2.2$ and $EW_{[O\\;III]} < 20$, out of which 119 quasars do not show detectable [\\ion{Ne}{5}] emission line in the spectra. Fig.\\ \\ref {f13} shows the two composite spectra of these non-\\ion{Mg}{2} BAL quasars with/without [\\ion{Ne}{5}] emission, the emission parameters are shown in Table \\ref{tab3}. The composite spectrum of the 119 non-BAL quasars without [\\ion{Ne}{5}] is similar to that of \\ion{Mg}{2} BAL quasars, showing stronger optical narrow \\ion{Fe}{2} emission, weak [\\ion{O}{3}] strength and strong UV \\ion{Fe}{2} but with a blue continuum. These quasars are probably from the same parent population of loBAL quasars but our line of sight does not intersect the outflow. The  high fraction of non-BAL quasars without [\\ion{Ne}{5}] emission line indicates that the covering factor of loBALR is not large.  Several observed trends may be explained by the strong correlation between the frequency of loBAL and Eddington ratio, and the correlations of the Eddington ratio with the other parameters concerned. It was reported that narrow optical \\ion{Fe}{2} strength is fairly well correlated with the Eddington ratio for low redshift quasars (Dong et al. 2009b). Dietrich et al. (2002) demonstrated that [\\ion{O}{3}] strength is inversely correlated with the Eddington ratio for quasars. Weakness of \\ion{Mg}{2} in the red-side of \\ion{Mg}{2} line profile is difficult to be ascribed to the absorption, and can be understood in this context as well via a fairly strong anti-correlation between EW of \\ion{Mg}{2} and the Eddington ratio (Dong et al. 2009a). Thus, the Eddington ratio can be an underlying driver for the different covering factor of low ionization BALR. We note that even in the highest Eddington ratio bin, the fraction of loBAL quasars is only a factor of 2   of the value in the rest two low Eddington ratio bins. There is no correlation between the fraction of loBAL with black hole mass for this sub-sample. Narrow line Seyfert 1 galaxies (NLS1s) are believed to be the low mass counter-parts of high accretion rate quasars. Zhou et al. (2006a; 2006b) found that several NLS1s also show \\ion{Mg}{2} BALs. Twelve of our low-$z$ \\ion{Mg}{2} BAL quasars can be formally classified as NLS1s according to the formal criterion of $H\\beta <2000~km~s^{-1}$, which account for 1.6\\% of NLS1s in this redshift range. Therefore, NLS1s do not appear to show significantly different properties from other quasars with  similar optical luminosity. The black hole mass range of these loBAL-NLS1s is very narrow with [2.1, 7.8] $\\times 10^7$ M$_{\\sun}$. It is possible that black hole mass does not matter once it is above certain threshold. %   Finally, most quasars with \\ion{Mg}{2} BALs or intermediate width \\ion{Mg}{2} absorption lines show reddened colors. We already noticed that the two group quasars show significantly different properties of emission lines, and the BAL fraction among the latter group is less than 25\\% (refer \\S 2.2). Thus, it is unlikely due to mixing of un-identified BAL quasars. The ubiquity of dust in intermediate and LoBAL outflows may be naturally explained as both absorbers are large scale outflows (e.g., Dunn et al. 2010). The presence of dust will significantly boost the radiative force and thus it allow gas in a relative large distance from the nucleus to be accelerated by the quasar radiation. In other word, gas free from dust will not be accelerated to high velocity by radiation pressure. In  this case, dust reddening would be preferably observed in the outflow direction. As we have argued that [\\ion{Ne}{5}]-weak quasars may be \\ion{Mg}{2} BAL quasars seen from an off-BALR direction, and their color can be fairly blue, thus dust may not present in other direction. This is consistent with above argument. Certainly, critical test for this can be done with a comparative study of broad band infrared SED of BAL and non-BAL quasars.  Dust reddening is ubiquitous in broad ($\\Delta v_{c} \\geq 1000$ \\kms) \\ion{Mg}{2} absorbers, regardless of whether they are BAL or non-BAL quasars. The average excess reddening is E(B-V) $\\sim$ 0.08 mag for SMC-type dust for both groups. However, we think that the intermediate width \\ion{Mg}{2} absorption quasars might have somewhat lower reddening than \\ion{Mg}{2} BAL quasars. Because quasars with large extinction, more likely BAL quasars, are missed in the SDSS quasar sample due to color selection criteria of quasar target, this sample explores only the relative low extinction end of quasars.  Thus the similar color distribution for quasars with \\ion{Mg}{2} BAL and with intermediate width \\ion{Mg}{2} absorption lines may be caused by the color selection effect that introduces a truncation in the severely reddened quasars. Indeed, radio and infrared-selected \\ion{Mg}{2} BAL quasars show much larger extinction with E(B-V) up to 1.5 mag (Urrutia et al. 2008), while there is no good statistical work for intermediate width \\ion{Mg}{2}. So it is not conclusive whether \\ion{Mg}{2} BAL quasars and quasars with intermediate width \\ion{Mg}{2} absorption lines have similar dust extinctions."
    },
    "1004/1004.0250_arXiv.txt": {
        "abstract": "We study the effect of noise in the density field, such as would arise from a finite number density of tracers, on reconstruction of the acoustic peak within the context of Lagrangian perturbation theory. Reconstruction performs better when the density field is determined from denser tracers, but the gains saturate at $\\bar{n}\\sim 10^{-4}\\,(h\\,{\\rm Mpc}^{-1})^3$. For low density tracers it is best to use a large smoothing scale to define the shifts, but the optimum is very broad. ",
        "introduction": "Baryon acoustic oscillations (BAO) in the baryon-photon fluid provide a standard ruler to constrain the expansion of the Universe and have become an integral part of current and next-generation dark energy experiments \\cite{EisReview05}. These sound waves imprint an almost harmonic series of peaks in the power spectrum $P(k)$, corresponding to a feature in the correlation function $\\xi(r)$ at $\\sim$100 Mpc, with width $\\sim 10$\\% due to Silk damping \\cite{PeeYu70,SunZel70,DorZelSun78,Eis98,MeiWhiPea99,ESW}. Non-linear evolution leads to a damping of the oscillations on small scales \\cite{Bha96,MeiWhiPea99} (and a small shift in their positions \\cite{ESW07,CroSco08,Mat08a,Seo08,PadWhi09}), \\begin{equation} P_{\\rm obs}(k) = b^2 e^{-k^2\\Sigma^2/2} P_L(k) + \\cdots \\cdots \\label{eq:processed} \\end{equation} where we have assumed a scale-independent bias, $b$, and left all broad band and mode-coupling features implicit in the $\\cdots$. The damping of the linear power spectrum (or equivalently the smoothing of the correlation function) reduces the contrast of the feature and the precision with which the size of ruler may be measured and is given by the mean-squared Zel'dovich displacement of particles, \\begin{equation} \\Sigma^2 = \\frac{1}{3\\pi^2} \\int dp\\ P_L(p) \\qquad . \\label{eq:sigmal} \\end{equation}  In \\cite{ESSS07} a method was introduced for reducing the damping, sharpening the feature in configuration space or restoring the higher $k$ oscillations in Fourier space. This procedure was studied in \\cite{PadWhiCoh09,NohWhiPad09} using Lagrangian perturbation theory. In this brief note we generalize these treatements to show how the effects of noise in the density field, arising for example from the finite number density of tracers, affects reconstruction.  We shall concentrate on the broadening of the peak, and refer the reader to \\cite{PadWhiCoh09,NohWhiPad09} for details, discussion and notation. ",
        "conclusions": ""
    },
    "1004/1004.2255_arXiv.txt": {
        "abstract": "By comparing the outcome of $N$-body calculations that include primordial residual-gas expulsion with the observed properties of 20 Galactic globular clusters (GCs) for which the stellar mass function (MF) has been measured, we constrain the time-scale over which the gas of their embedded cluster counterparts must have been removed, the star formation efficiency the progenitor cloud must have had and the strength of the tidal-field the clusters must have formed in. The three parameters determine the expansion and mass-loss during residual-gas expulsion. After applying corrections for stellar and dynamical evolution we find birth cluster masses, sizes and densities for the GC sample and the same quantities for the progenitor gas clouds. The pre-cluster cloud core masses were between $10^5-10^7\\msun$ and half-mass radii were typically below $1$ pc and reach down to $0.2$ pc. We show that the low-mass present day MF (PDMF) slope, initial half-mass radius and initial density of clusters correlates with cluster metallicity, unmasking metallicity as an important parameter driving cluster formation and the gas expulsion process. This work predicts that PD low-concentration clusters should have a higher binary fraction than PD high-concentration clusters.  Since the oldest GCs are early residuals from the formation of the Milky Way (MW) and the derived initial conditions probe the environment in which the clusters formed, we use the results as a new tool to study the formation of the inner GC system of the Galaxy. We achieve time-resolved insight into the evolution of the pre-MW gas cloud on short time-scales (a few hundred Myr) via cluster metallicities. The results are shown to be consistent with a contracting and self-gravitating cloud in which fluctuations in the pre-MW potential grow with time. An initially relatively smooth tidal-field evolved into a grainy potential within a dynamical time-scale of the collapsing cloud. ",
        "introduction": "\\label{sec:intro} Globular clusters (GCs) and old low-mass stars have often been used as local probes of Galaxy formation, since they preserve information about ancient times. In the past, the formation of the Milky Way (MW) has been investigated by means of kinematic studies of stars and their abundances \\citep{CarolBeers07,bee02,bee01,els62}, as well as by horizontal branch morphology, metallicity and kinematic measurements of star clusters \\citep{bekki07,mvdb05,mg04,z93,sz78}. In terms of the origin of the GCs, this led to a picture of Galactic formation in which GCs are divided in the old halo (OH) and young halo (YH) clusters. The OH clusters are located inside a Galactocentric distance of $\\dgc\\approx8-10$ kpc. Many of them appear to have formed coevally with the collapse of the protogalaxy \\citep{els62,sw02,dea05,mf09}. The YH clusters have $\\dgc\\gtrsim8-10$ kpc and have been accreted over several Gyr \\citep*{sz78}. All this information was gained from knowledge about the present day (PD) parameters of GCs. Knowing about the initial conditions at star cluster birth, however, would provide a deeper insight into the early formation processes of the MW since it would allow to probe directly the environment in which the GCs formed \\citep*{g97,g98}.  Cluster masses at birth were larger than they are today since clusters suffer mass loss due to primordial residual-gas expulsion, stellar and dynamical evolution. And since the expulsion of gas leads to subsequent expansion, young and gas-embedded clusters were much more compact and denser than they are nowadays. Star clusters are the PD gravitationally bound remnants of these dense objects after the stars emerged from their natal cloud \\citep*{tut78,kah01,bg06,bk07}. The overall change in the potential due to the gas loss leads to cluster expansion and the loss of stars, with the \\textit{initial conditions} at the onset of this process deciding about cluster survival or destruction. The majority of the freshly hatched clusters are destroyed during this violent phase \\citep*{ll03,bg06,gg08,b08,gb08} and their member stars become stars of the field. In this view, the OH GCs are the massive remnants of an initial population of embedded clusters that rapidly formed the population II halo of the MW \\citep*{kah01,kb02,bkp08}. Some of them ended as bound clusters which, after stellar and two-body relaxation driven dynamical evolution, we can observe nowadays. So in order to understand physical properties of star clusters \\textit{today} it is essential to understand how star clusters \\textit{formed}, because the birth configuration determines the fate of a star cluster.  The time, $\\tau_M$, over which the natal gas is removed from the cluster determines crucially whether a star cluster survives gas expulsion or not. \\citet*{bkp08} provide an analytic formula to calculate $\\tau_M$, \\begin{equation} \\tau_M=7.1\\times10^{-8}\\frac{1-\\epsilon}{\\epsilon}\\frac{\\mcl}{\\msun}\\left(\\frac{r_h}{\\rm pc}\\right)^{-1}\\,\\rm Myr, \\label{eq:tau} \\end{equation} based on the amount of energy needed to be put into the gas to overcome its potential energy. The deeper the potential, i.e. the larger the progenitor cloud mass, $\\mcl$, the more difficult it is to remove the gas. A large star formation efficiency (SFE), $\\epsilon$, leads to more and also more massive stars with stronger winds and radiation and there is less gas to remove. Finally, the larger the half-mass radius, $r_h$, the faster the gas is expelled since the overall potential is shallower for fixed $\\mcl$ and $\\epsilon$.  The gas expulsion time, thus, depends on the mass and the size of the cluster. From theoretical considerations mass and radius are related. Dependent on the exact form of the mass-radius relation of young star clusters, the gas expulsion time-scale, $\\tau_M$, is affected more strongly or weakly. Virialised gas cores (radius $r_c$ and core mass $m_c$) are expected and observed to show a strong mass-radius relation, scaling as $r_c\\propto m_c^{1/2}$ \\citep*[i.e. constant surface density, e.g.][]{hp94}. Therefore the mass-radius relation of young clusters is expected to display a similar behaviour. If this relation is valid the gas removal time-scales would be essentially mass independent, $\\tau_M\\propto m_c^{1/12}$ \\citep*{pf09}. In contrast, observations of young clusters do not show any significant mass-radius relation, $r_{cl}\\propto m_{cl}^{0-0.1}$ \\citep*[e.g.][]{l04,k05}, and the influence on the gas expulsion time-scale is a stronger function of mass, $\\tau_M\\propto m_c^{1/2}$ \\citep*{pf09}.  Observationally determined SFEs lie between $\\epsilon=0.2$ and $0.4$ \\citep*{ll03}, so typically most of the mass remains in the gas to be expelled. $N$-body modelling showed that SFEs can be as low as $10-20$ per cent if the gas expulsion time-scales are sufficiently long and has to be at least $33$ per cent if the gas is lost instantaneously in order for a cluster to survive \\citep*{lmd84,gb01,bk03a,bk03b,bk07}. SFEs may vary locally and could be different between high- and low-mass cores, which is also a viable explanation to wipe out the mass-radius relations observed in gas cores \\citep*{az01}.  In Sec. \\ref{sec:varimf} we present the observational data to be compared to the results of $N$-body experiments and in Sec. \\ref{sec:models} we explain how we model the initial cluster size, mass and density. We derive the initial conditions and discuss correlations between and among initial and present day cluster parameters in Sec. \\ref{sec:initcond}. In Sec. \\ref{sec:galform} we connect the primordial conditions constrained in the former sections to develop a picture for the assembly of the old population of GCs located in the inner halo. Sec. \\ref{sec:sum} summarises our main results. ",
        "conclusions": ""
    },
    "1004/1004.2125_arXiv.txt": {
        "abstract": " ",
        "introduction": "Coherent oscillation of scalar field plays an important role to describe various phenomena in particle physics and particle cosmology. One of the most important examples is the so-called slow-roll model of the inflationary universe.~\\cite{SRINF,LindeBook} A scalar field, called as inflaton, is initially displaced from the potential minimum and its vacuum energy leads to the de-Sitter expansion of the universe.  This class of models elegantly solves the flatness and horizon problems of the standard Big Bang cosmology.  Moreover, it can give an origin of the density fluctuations which is strongly supported from the recent measurements of the cosmic microwave background radiation.~\\cite{Larson:2010gs} After the inflation ends, the inflaton starts to cause coherent oscillation around its potential minimum.  The energy of coherent oscillation is diluted due to the expansion of the universe as well as the energy transfer to particles though interaction of the oscillating field.  Produced particles are then thermalized and the hot universe can be realized.  The whole of these processes is called as the reheating.  In particular, the reheating of the slow-roll inflation gives an initial condition of the standard Big Bang cosmology. Therefore, the reheating process is crucial for understanding the very early universe.  In this paper, we focus on the first stage of the reheating, \\ie, particle production from coherent oscillation. This process has been widely discussed based on that coherent oscillation is considered as non-relativistic scalar particles,~\\cite{NR_particle} \\ and particles are produced through their decay and/or scattering processes.~\\cite{Albrecht:1982mp} \\ In this case the number of parent non-relativistic particles is given by the energy of coherent oscillation divided by the mass of the oscillating field, and the number of produced particles is determined by it.  On the other hand, it has been pointed out that, when the coupling with produced particles and also the amplitude of the oscillation become large, the non-perturbative effect becomes significant in the early stage of particle production~\\cite{Traschen:1990sw,Kofman:1994rk,Boyanovsky:1994me,Yoshimura:1995gc}. \\ This process is called as the preheating.~\\cite{Kofman:1994rk} \\ Especially, the explosive production of bosonic degrees of freedom can happen due to the broad parametric resonance effect. On the other hand, the fermion production at the preheating has been also investigated.~\\cite{Dolgov:1989us,Baacke:1998di,Greene:1998nh}  The purpose of this paper is to investigate the production of scalars and fermions from coherent oscillation, especially when the coupling constants of oscillating field are very small to avoid the effect of the broad parametric resonance.  For this purpose, we apply the method based on the Bogolyubov transformations.~\\cite{Bogolyubov:1958se,Birrell} In this case, the equation of motion for the mode functions of produced particles in the presence of coherent oscillation is solved and the growth of the mode functions are then interpreted as the production of particles.  First, we will present the analytical formulae for the distribution functions and the number densities of produced particles by using the perturbative expansion of the coupling constants. We will also discuss the conditions under which the perturbative results are justified.  Indeed, it will be shown that the leading-order results collapse in the end.  This is a signal that the non-perturbative effect becomes important even if the coupling constant is sufficiently small. Such a correction is crucial for describing the statistical properties of produced particles, namely the effects of the Bose condensation for the scalar production and the Pauli blocking for the fermion production.  In order to handle the annoying non-perturbative effects we will present the time averaging method, which is familiar in the nonlinear dynamical system.~\\cite{NAYFEH} It will be demonstrated that this method is powerful to extract the characteristic evolution of the occupation number for the growing mode, \\ie, the exponential growth for scalar production or the oscillation between 0 to 1 for fermion production. Furthermore, we will show that the results by the time averaging method obey the exact scaling property, which is obtained from the periodicity of the the equation of motion.~\\cite{Mostepanenko:1974} This gives a justification for the use of the time averaging method. Throughout the present analysis we neglect the expansion of the universe for simplicity.   The rest of this paper is organized as follows. In Sec.~\\ref{sec:Framework} we explain the model in this analysis. We perform in Sec.~\\ref{sec:Perturbative_Estimate} the perturbative estimation of the yields when the amplitude of the coherent oscillation is sufficiently small.  The importance of non-perturbative effects in particle production is addressed in Sec.~\\ref{sec:Time_Ave}.  We present the time averaging method to deal with such effects, and try to figure out the statistical properties of produced particles. Finally, the last section is devoted to conclusion. We also add Appendix~\\ref{sec:AP1} to explain the perturbative estimation of the number density. ",
        "conclusions": "\\label{sec:Conc} We have investigated the particle production from the coherent oscillation by using the method based on the Bogolyubov transformation.  For the case when the coupling constants of the oscillating field are very small, we have obtained the leading contributions to the distribution functions and the number densities of the produced particles.  When the amplitude of the oscillation is small ($\\Phi \\ll \\vev{\\phi}$), the leading contributions to the yields are found to be ${\\cal O}(g_{S,F}^2)$.  We have presented the exact expressions for the distribution functions of the produced $\\chi$ and $\\psi$ at the ${\\cal O}(g_{S,F}^2)$ order.  It has been shown that there exists the growing mode with $k = k_\\ast$ if $m_{\\chi, \\psi} < m_\\phi/2$, and its occupation number increases at the rate $t^2$ after sufficient numbers of the oscillation. The distribution function has a peak at $k \\simeq k_\\ast$ and the width of the peak decreases at the rate $1/t$.  As a result, the number density of produced particles is proportional to $t$.  The expression for the number density is found to be consistent with the one obtained by assuming that the coherent oscillation is a correction of non-relativistic scalar particles and the decay process is a main source of the particle production.  We have found that the above perturbative results fail to describe the exact ones for sufficient late times since the non-perturbative correction becomes significant even when the coupling constants are extremely small.  Indeed, the occupation number of $\\chi$ for the growing mode increases exponentially while that of $\\psi$ oscillates around $1/2$.  These distinctive features represent the statistical properties of the produced particles, \\ie, the effects of the Bose condensation for the scalar production or the Pauli blocking for the fermion production.  Due to these non-perturbative effects, the explosive production of $\\chi$ happens while the production of $\\psi$ becomes insignificant for late times.  To handle with these non-perturbative effects we have used the time averaging method, and have successfully described the evolution of the occupation number for the growing mode.  This method works well because the typical time scale of the evolution is much longer than the rapid $\\phi$ oscillation.  Furthermore, we have shown that the results obtained by the time averaging method satisfies the exact scaling properties in Ref.\\citen{Mostepanenko:1974}, which also gives the justification of the use of the time averaging method.   Throughout this analysis, we have neglected the back-reaction effect of the produced particles in the estimation of the yields.  When the occupation number of these particles is close to unity, such an effect should be taken into account.  In addition to this, the inclusion of the expansion of the universe is also necessary to reveal the reheating/preheating processes in the inflationary universe.  These issue will be discussed in elsewhere~\\cite{AN}."
    },
    "1004/1004.1060_arXiv.txt": {
        "abstract": "The Australia Telescope Compact Array (ATCA) has been used to measure positions with arcsecond accuracy for 379 masers at the 22-GHz transition of water.  The principal observation targets were 202 OH masers of the variety associated with star formation regions (SFR)s in the Southern Galactic plane.  At a second epoch, most of these targets were observed again, and new targets of methanol masers were added. Many of the water masers reported here are new discoveries and others had been reported, with position uncertainties exceeding 10 arcsec, from Parkes telescope single dish observations many years ago.  Variability in the masers is often acute, with very few features directly corresponding to those discovered two decades ago.  Within our current observations, less than a year apart, spectra are often dissimilar, but positions at the later epoch, even when measured for slightly different features, mostly correspond to the detected maser site measured earlier, to within the typical extent of the whole site, of a few arcseconds.  The precise water positions show that approximately 79 per cent (160 of 202) of the OH maser sites show coincident water maser emission, the best estimate yet obtained for this statistic;  however, there are many instances where additional water sites are present offset from the OH target, and consequently less than half of the water masers coincide with a 1665-MHz ground-state OH maser counterpart. Our less uniform  sample of methanol targets is not suitable for a full investigation of their association with water masers, but we are able to explore differences between the velocities of peak emission from the three species, and quantify the typically larger deviations shown by water maser peaks from systemic velocities.  Clusters of two or three distinct but nearby sites, each showing one or several of the principal molecular masing transitions, are found to be common.  We also report the detection of ultracompact \\HII regions towards some of the sites.  In combination with an investigation of correlations with IR sources from the GLIMPSE catalogue, these comparative studies allow further progress in the use of the maser properties to assign relative evolutionary stages in star formation to individual sites. ",
        "introduction": "Masers of OH (hydroxyl), water and methanol are key tools for investigating the formation of massive stars.  The masers reside in the dusty molecular envelope or torus of a massive star in its earliest stage of formation, and the masers are a sensitive probe for discovering stars in this embryonic state when the star is not visible because of obscuration from the dust.  Detailed studies of selected maser sites have been made comparing OH with water, and OH with methanol (Forster \\& Caswell 1989, 1999; Caswell 1997; Caswell Vaile \\& Forster 1995; hereafter FC89, FC99, C97, CVF95).  These suggest that the maser spots of all species are usually contained inside a region of diameter less than 30 mpc (= $0.93 ~ 10^{15}$ m = $0.93 ~  10^{17}$ cm  $\\approx 6200$ au $\\approx 0.1$ light year), corresponding to an angular diameter of about 1 arcsec at a typical distance of 6 kpc. Positional precision of about 1 arcsec is therefore desirable to establish whether masers of different species originate from the same site.  Within many star formation regions (SFRs), on a larger scale of 100 mpc or more, the FC89 and C97 studies revealed many instances where a number of maser sites are present in a small cluster, often with different combinations of maser species present.  There has been considerable speculation that the various combinations are indicative of massive young stars at a different evolutionary stage, or in a different mass range \\citep[e.g.][]{Breen09,Ellingsen07}. Water maser emission is especially puzzling.  It has commonly been thought to be most prolific at an early stage of stellar evolution,  but isolated water masers are only partly accounted for by very young sites preceding OH maser excitation.  Other offset positions indicate the additional occurrence of water masers towards lower mass stars, or in high velocity fragments that have travelled far from the initial site.  In the latter case, maser spots from individual fragments are sometimes ephemeral, and disappear on a timescale as short as months, but are often replaced by generally similar emission in the same region, although at a slightly different position and velocity.  To date, the most extensive unbiased survey for masers of the SFR variety completed in the Galactic plane is at the 1665-MHz transition of OH, presented as a catalogue reporting positions of arcsec accuracy for more than 200 masers (Caswell 1998;  hereafter C98). Sensitive observations of 6.6-GHz methanol at these positions have already been made, with a detection rate of 80 per cent after high precision positions are compared \\citep{C98,C09}.  Early 22-GHz water maser surveys with the Parkes radio telescope towards southern SFRs have been reported by \\citet{C+89} and references therein, with positions measured to about 10 arcsec accuracy. Until recent years, follow up observations to arcsec accuracy were restricted to northerly objects accessible to the VLA (FC89, FC99).  The availability of the 22-GHz frequency band at the ATCA has allowed us to search with high sensitivity and high positional precision for water masers towards the full sample of OH masers, as well as $\\sim$100 methanol maser sites.       \\section[]{Observations and data reduction}  Water maser observations were made with the ATCA on two separate occasions. During the first epoch of observations (2003 October 4 and 5), OH maser sites from C98 were targeted (see also Section 5.3) and during the second epoch (2004 July 25, 26 and 30) many water maser detections towards the OH masers were re-observed as well as a selected set of methanol masers \\citep[chiefly from][]{C09}. OH maser sources that were targeted in the first epoch but showed no detectable water maser emission were not observed at the second epoch, and only a few of the water maser sources that had previously been observed by FC89 with the VLA, and successfully confirmed during the first epoch, were reobserved.  Our first observations were made with an EW array yielding 10 baselines between 30 and 352 m (project c1190). The correlator sampled a single linear polarization, processed to give a 512-channel spectrum across a 32-MHz bandwidth. The observing strategy was to observe approximately 100 targets over each 12-h session. After initial calibration, the first 10 targets were observed for 1.5 min each, followed by a calibrator, and similarly for the remaining 90 targets. Then the cycle was repeated 2 more times, so as to provide for each source adequate uv-coverage, combined with a total integration time of 4.5 min.  Primary flux calibration is relative to PKS B1934-638 and in general is expected to be accurate to $\\sim$20 \\%. PKS B1921-293 was used for bandpass calibration.   The AIPS reduction package was used for  processing of the data collected in this first epoch, following the general procedure described in C97.  In the realignment of channels using the CVEL task, the adopted rest frequency was 22235.08 MHz, and the velocity scale was with respect to the local standard of rest (lsr).  The channel separation was 0.84 \\kms\\ which, with uniform weighting of the correlation function, yields a final velocity resolution of 1.0 \\kmsns.  The quite coarse resolution was a compromise chosen in order to allow a large velocity coverage of more than 400 \\kmsns.  With this coverage, it was possible to recognise any high velocity features indicative of close association with outflows (for which the water masers are renowned).  Total intensity maps were then produced of the channels with maser emission apparent in the scalar averaged spectrum, or in a vector averaged spectrum shifted to the location of the target OH or methanol maser emission.  The rms noise in an individual channel image was typically 150~mJy.  The synthesised beam has a halfpower width of approximately 8 arcsec in right ascension, but is larger in declination by a factor cosec(declination) as expected for an array aligned East-West with maximum baseline of 352 m.  Our second series of observations (project c1330) were made in similar fashion but with a different array configuration, H168. The correlator configuration, and therefore the spectral resolution and velocity coverage, was identical to that used in the 2003 observations. A few sample observations from this epoch were reduced, firstly with AIPS (as for the 2003 data), and secondly with  {\\tt miriad} software package \\citep{Sault}. There was excellent agreement between data reduced in the respective data reduction packages. The full data set from this epoch were reduced using {\\tt miriad}, applying the standard techniques for ATCA spectral line and continuum observations. Image cubes of the entire primary beam and velocity ranges were produced for each source. The flux densities of sources that were located away from the centre of the primary beam have been corrected to account for beam attenuation. Spectra for each source detected at this epoch were produced by integrating the emission in the ATCA image cubes for each source. The typical resultant rms noise in each spectrum was 40 - 50 mJy. For the H168 array used in 2004, the synthesised beam was typically 13x9 arcsec.%  The ATCA observations are most sensitive at the targeted positions, but provide useful measurements, albeit at lower sensitivity, of any other sources that happen to lie within the field of view of the primary beam; the full width to the first null is nearly 5 arcmin, and the HPBW is 2.29 arcmin = 137 arcsec.     \\begin{table*} \\caption{22-GHz water masers detected towards sites of OH and methanol masers. Column 1 shows the source name in Galactic coordinates, column 2 and 3 give the right ascension and declination, column 4, 5 and 6 give the velocity of the water maser peak, velocity range and peak flux density in the 2003 observations, while columns 7, 8 and 9 give the velocity of the water maser peak, velocity range and peak flux density in the 2004 observations. A `--' in either column 6 or 9 indicates that no observations were made of the given source during the 2003 or the 2004 epoch, respectively, while the presence of a number preceded by a `$<$' indicates that there was no emission detected above the quoted threshold. For some complicated sources a `t' is present in either column 6 or 9 and this indicates that the exact nature of the detection is discussed in Section~\\ref{sect:ind}. Associations are given in column 10, where the presence of an `o' denotes an OH maser, an `m' denotes a methanol maser, a `c' denotes the presence of a 22-GHz radio continuum source, a `g' the presence of a GLIMPSE point source and the presence of a `$\\gamma$' indicates that the water maser source is outside the range of the GLIMPSE survey region. A `$^{\\#}$' indicates that the proceeding associated source is strictly outside our association threshold but has been added through special circumstances. See Section 3 for a more extensive description.} \\begin{tabular}{lccrcrrcrl} \\hline \\multicolumn{1}{c}{\\bf Water maser} & {\\bf RA}  & {\\bf Dec} & {\\bf Vpeak} &{\\bf Vrange} & {\\bf Speak} & {\\bf Vpeak} &{\\bf Vrange} & {\\bf Speak} & {\\bf Associations}\\\\ \\multicolumn{1}{c}{\\bf ($l,b$)} & {\\bf (J2000)}  & {\\bf (J2000)}&  {\\bf (\\kmsns)}&  {\\bf (\\kmsns)} & {\\bf (Jy)} &{\\bf (\\kmsns)}&  {\\bf (\\kmsns)} & {\\bf (Jy)}     \\\\ \\multicolumn{1}{c}{\\bf (degrees)} & {\\bf (h m s)}&{\\bf ($^{o}$ $'$ $``$)}& {\\bf 2003} & {\\bf 2003} & {\\bf 2003}& {\\bf 2004} & {\\bf 2004} & {\\bf 2004} \\\\ \\hline \\\\ G\\,240.316+0.071 & 07 44 51.94 & --24 07 41.9 & 89 \t& 58,100\t\t& 10 \t\t&\t\t&\t\t\t&\t--\t& \to$\\gamma$\\\\ G\\,263.250+0.514 & 08 48 47.91 & --42 54 27.1 & 19 \t& 17,21\t\t& 3.0 \t\t& 20\t\t& 17,21\t\t&\t0.7\t& \tom$\\gamma$\\\\ G\\,284.350--0.418& 10 24 10.60 & --57 52 33.0 &--4\t &--42,72    \t& 60\t\t& 7\t\t&--78 ,71\t\t&\t102\t&\to$\\gamma$\t\\\\ G\\,285.260--0.067& 10 31 24.57 & --58 03 03.7 &--28  \t&--95,20    \t& 22\t\t&      \t--92\t&--96,20\t\t&\t30\t&\t$\\gamma$\t\\\\ G\\,285.263--0.050& 10 31 29.64 & --58 02 18.9 & 2    \t&--59,50    \t&1100\t& 3\t\t& --36,63\t\t& \t1651\t& \to$\\gamma$\t\\\\ G\\,287.371+0.644 & 10 48 04.25 & --58 27 00.7 & 1    \t&--14,6     \t\t&3.5\t\t&  --1\t& --13,1\t\t&\t4.9\t&\tom$\\gamma$    \\\\ G\\,290.374+1.661 & 11 12 17.98 & --58 46 21.6 & --12 \t&--20,--10   \t&3.5\t\t& --47\t&--47,--33\t\t&\t0.26\t&\tom$\\gamma$ \\\\ G\\,290.384+1.663 & 11 12 22.53 & --58 46 29.0 & --38 \t&--50,--35  \t& 6 \t\t&--37\t&--41,--35\t\t&\t4.5\t&$\\gamma$\t \\\\ G\\,291.270--0.719& 11 11 49.67 & --61 18 53.8 &--102 \t&--110,--23 \t&65 \t\t& --102\t&--109,--17 \t&\t53\t&   m$\\gamma$\\\\ G\\,291.274--0.709& 11 11 53.28 & --61 18 24.1 & --32 \t&--51,14  \t&60 \t\t&--32\t& --59,--14\t&\t64\t&   \tom$\\gamma$\t\\\\ G\\,291.284--0.716& 11 11 56.58 & --61 19 01.1 &--133 \t&--142,--120\t&930\t&--133\t&--139,--120\t&\t701\t&$\\gamma$\t\t      \\\\ G\\,291.578--0.434& 11 15 04.91 & --61 09 50.7 &\t\t&\t\t\t&$<$0.2\t&18\t\t&17,\t27\t\t&\t16 \t& $\\gamma$\\\\ % G\\,291.579--0.431& 11 15 05.67 & --61 09 41.0 & 12   \t&--29,41     \t\t&215\t& 13\t\t&--31,22\t\t&\t608\t& om$\\gamma$\\\\ G\\,291.581--0.435& 11 15 06.17 & --61 09 56.1 & 26   \t& 25,27     \t&4.0\t\t&\t\t&\t\t\t& $<$0.2\t& m$\\gamma$\\\\ G\\,291.610--0.529& 11 15 02.58 & --61 15 48.8 & 13   \t& --66,20   \t&18 \t\t&\t12\t& --66,28\t\t&\t39\t& oc$\\gamma$\\\\ G\\,291.627--0.529& 11 15 10.18 & --61 16 12.7 & 22   \t&8,23       \t\t&12 \t\t& 22\t\t&\t20,25\t&\t31  \t& $\\gamma$ \\\\ G\\,291.629--0.541& 11 15 08.88 & --61 16 54.8 & 11 \t&8,16       \t\t&70 \t\t&    \t10\t&--2,21\t\t&\t46\t& $\\gamma$ \\\\ G\\,294.511--1.622& 11 35 32.04 &--63 14 44.3 & --12 \t&--20,--6   \t&250\t& --12\t&\t--20,--4\t&\t112 & om$\\gamma$\\\\ G\\,294.976--1.733& 11 39 13.77 & --63 29 03.3 & \t  \t&           \t\t&   \t--\t&      1\t& --17,5\t\t& \t2.2\t\t& $\\gamma$\\\\ G\\,294.989--1.719& 11 39 22.56 & --63 28 25.1 &\t\t&           \t\t&   \t--\t&     --17  \t&  --18,--10\t& \t0.6\t&m$\\gamma$\\\\ G\\,297.660--0.974& 12 04 08.76 & --63 21 37.3 & 29\t&--80,38    \t&90 \t\t&\t26\t& --79,82\t\t&\t75\t& o\\\\ G\\,299.012+0.125 & 12 17 24.05 & --62 29 13.6 &\t\t&\t\t\t& $<$0.2\t&--26\t&--27,--24\t\t&\t0.44&    \\\\ G\\,299.013+0.128 & 12 17 24.58 & --62 29 04.8 & 19 \t&\t 18,30     \t&100\t&\t19\t&--1,\t50\t\t&\t83\t& omcg\\\\ G\\,300.491--0.190& 12 29 55.99 & --62 57 33.8 & \t24\t&\t23,25\t&2.3\t\t&25\t\t&22,25\t\t&\t1.6\t   &\\\\ G\\,300.504--0.176& 12 30 03.42 & --62 56 50.2 & 11   \t&--37,14    \t&180\t&11\t\t&--26,14\t\t&\t94\t& omg\\\\ G\\,300.968+1.143 & 12 34 52.51 & --61 39 57.9 &--61  \t&--86,--42  \t&70 \t\t&--58\t&--86,--3\t\t&\t26\t& $\\gamma$\\\\ G\\,300.971+1.143 & 12 34 54.34 & --61 39 57.1 &\t\t&\t\t\t&$<$3\t&--43\t&--46,--42\t\t&\t3.0 &$\\gamma$\\\\ % G\\,301.136--0.225& 12 35 34.76 & --63 02 28.5 &\t\t&\t\t  \t&t \t\t&\t--47\t&\t--48,--41\t&\t17 \\\\ G\\,301.136--0.226a& 12 35 34.93 & --63 02 34.5&--29\t&--56,--18  \t&80 \t\t&\t--29\t&--31,--28\t\t& \t78\t& g\\\\ G\\,301.136--0.226b& 12 35 34.84 & --63 02 31.1&\t\t&\t\t  \t&t \t\t&\t--45\t&\t--63,--43\t&\t92\t& (omc)\\\\ G\\,301.137--0.225& 12 35 35.21 & --63 02 30.6 &\t\t&\t\t  \t&t \t\t&\t--36\t&\t--40,--33\t&\t39\t& omc\\\\ G\\,305.191--0.006& 13 11 12.95 & --62 47 27.7&31    \t&28,36      \t&25  \t\t& \t32\t&29,35\t\t&\t4.8  \t& g   \\\\ G\\,305.198+0.007 & 13 11 16.38 & --62 46 37.4 &--39  \t&--42,--27  \t& 8 \t\t&     --35\t&--41,--24\t\t&\t3.5 \t& \\\\ G\\,305.208+0.207 & 13 11 13.37 & --62 34 40.0 &--39\t&--43,--37  \t&300\t&     \t--39\t&--44,--38\t\t& \t235\t& om \\\\ G\\,305.361+0.150 & 13 12 35.61 & --62 37 18.9 &--43  \t&--45,--28  \t&250\t&\t--36\t& --44,--29\t\t&\t126  & omg   \\\\ G\\,305.799--0.245& 13 16 42.92 & --62 58 31.7 &--26  \t&--45,35\t   \t&400\t&\t--34\t&\t--119,37\t&\t266\t& omg \\\\ G\\,306.318--0.331& 13 21 20.87 & --63 00 22.7 &--19\t&--24,--14  \t&2.1  \t\t& \t--18\t&\t--23,--15\t&\t0.7     \\\\ G\\,307.805--0.456& 13 34 27.32 & --62 55 12.4 &--13  \t&\t--15,--7   \t&1.4\t\t&\t--7\t&\t--11,--7\t&\t0.6\t& og\\\\ G\\,308.754+0.549 & 13 40 57.47 & --61 45 42.4 &--49  \t&--50,--48  \t& 3.4 \t\t&\t--49\t&\t--50,--48\t&\t0.7\t& omg\t\\\\ G\\,308.918+0.124 & 13 43 01.64 & --62 08 48.9 &--56\t&--66,--45  \t&0.7\t\t& --61\t&\t--62,--49\t&\t1.5\t& om$^{\\#}$\\\\ G\\,309.384--0.135& 13 47 24.01 & --62 18 11.4 &--50\t&--57,--41  \t&\t3.5\t&--50\t&\t--51,--48\t&\t2.2\t& omg\\\\ G\\,310.144+0.760 & 13 51 58.53 & --61 15 40.6 &--58  \t&--60,--55  \t&2.5\t\t&--58\t& --59,--57     \t&\t1.6\t& omg \\\\ G\\,310.146+0.760 & 13 51 59.61 & --61 15 39.7 &--63  \t&--64,--62  \t&8  \t\t&      \t\t&\t\t\t&$<$0.2 \\\\ G\\,311.643--0.380& 14 06 38.77 & --61 58 22.7 &35    \t&18,50      \t&320\t&36\t\t&8,56\t\t&\t167\t& omcg \\\\ G\\,312.106+0.278 & 14 08 46.06 & --61 12 30.1 &\t\t&\t\t\t&\t--\t&--54\t&--55,--53\t\t& 0.6\t\t& g\t\\\\ G\\,312.109+0.262 & 14 08 49.45 & --61 13 23.4 &\t\t&\t\t\t&\t--\t&--48\t&--48,--47\t\t&     0.43\t& mg\\\\ G\\,312.596+0.045 & 14 13 14.13 & --61 16 57.6 &\t--61\t& --62,--58\t&0.6\t\t&--59\t&   --64,--56\t&  1.5\t& m\\\\ G\\,312.599+0.046 & 14 13 15.19 & --61 16 51.9 &--75  \t&--102,--58 \t&2.7\t\t& --79\t& --96,--56\t\t&\t9\t& om\\\\ G\\,313.457+0.193 & 14 19 35.05 & --60 51 54.2 & \t--1\t& --2,0\t\t& 1.9\t\t&45 \t\t&   --1,46\t\t&1.0        \t& c \\\\ G\\,313.470+0.191 & 14 19 40.97 & --60 51 46.2 & --5 \t&--10,--2   \t&6  \t\t&--6\t\t&--10,--1\t\t& 4.5 \t& omg\\\\   \\hline \\label{tab:masers} \\end{tabular} \\end{table*}  \\begin{table*}\\addtocounter{table}{-1} \\caption{-- {\\emph {continued}}} \\begin{tabular}{lccrcrrcrl} \\hline \\multicolumn{1}{c}{\\bf Water maser} & {\\bf RA}  & {\\bf Dec} & {\\bf Vpeak} &{\\bf Vrange} & {\\bf Speak} & {\\bf Vpeak} &{\\bf Vrange} & {\\bf Speak} & {\\bf Associations}\\\\ \\multicolumn{1}{c}{\\bf ($l,b$)} & {\\bf (J2000)}  & {\\bf (J2000)}&  {\\bf (\\kmsns)}&  {\\bf (\\kmsns)} & {\\bf (Jy)} &{\\bf (\\kmsns)}&  {\\bf (\\kmsns)} & {\\bf (Jy)}     \\\\ \\multicolumn{1}{c}{\\bf (degrees)} & {\\bf (h m s)}&{\\bf ($^{o}$ $'$ $``$)}& {\\bf 2003} & {\\bf 2003} & {\\bf 2003}& {\\bf 2004} & {\\bf 2004} & {\\bf 2004} \\\\ \\hline \\\\ G\\,313.578+0.325 & 14 20 08.63 & --60 41 59.0 &--46\t&--58,--32   \t&55 \t\t&--47\t&--54,--32\t\t&\t46\t& omg \\\\ G\\,313.767--0.862& 14 25 01.71 & --61 44 57.2 & --54 \t&--56,--48  \t&160\t&--54\t&\t--57,--36\t&\t162 \t& omg\\\\ G\\,314.320+0.112 & 14 26 26.37 & --60 38 29.4 & --44 \t&--70,--40  \t&34 \t\t& --45\t&--70,--40\t\t&16\t\t& om \\\\ G\\,316.360--0.361& 14 43 11.07 & --60 17 10.3 & 3    \t&--4,21     \t\t& 30\t\t&3\t\t& --7,20\t\t&16 \\\\ G\\,316.361--0.363& 14 43 12.22 & --60 17 16.6 &--3\t& --4,--1\t\t& 6\t\t&--3 \t\t&--7,--1\t\t&2.6\t\t& \\\\ G\\,316.412--0.308& 14 43 23.22 & --60 12 58.8 & --22 \t&--30,6     \t\t&6  \t\t&--20\t&--24,5\t\t& 15\t\t& omcg \\\\ G\\,316.640--0.087& 14 44 18.39 & --59 55 10.7 & --19 \t& --104,--2 \t&12 \t\t& --15\t&--40,92\t\t& 12\t\t& omg\\\\ G\\,316.763--0.011& 14 44 56.32 & --59 47 59.8 & --47 \t&--51,--33  \t&60 \t\t& --48\t&--49,--33\t\t&4.1\t\t& og\\\\ G\\,316.812--0.057& 14 45 26.58 & --59 49 14.1 & --46 \t&--47,--36  \t&500\t& --46\t&--56,--11\t\t&408 \t& om\\\\ G\\,317.429--0.561& 14 51 37.72 & --60 00 18.2 & 16 \t& 12,25     \t&0.5\t\t&25\t\t& 24,25\t\t& 0.27\t&  oc \\\\ G\\,317.429--0.556& 14 51 36.75 & --60 00 03.4 &  26\t&\t25,36\t&0.5\t\t& 28\t \t& 27,29\t\t& 0.6\t\t& \t\\\\ G\\,318.044--1.404& 14 59 08.61 & --60 28 23.9 & 42   \t& 31,44     \t&3.5\t\t& 42\t\t& 32,42\t\t& 3.5\t\t& om$\\gamma$ \t\\\\ G\\,318.050+0.087 & 14 53 42.62 & --59 08 52.3 &--55  \t&--61,--39  \t&470\t&--48\t&--69,--38\t\t& 50\t\t& omg\t\\\\ G\\,318.948--0.196a& 15 00 55.18 & --58 58 51.6 &--41\t&--44,--28  \t&5  \t\t&\t--36\t&--44,--27\t\t&\t9\t& g\t\\\\ G\\,318.948--0.196b& 15 00 55.33 & --58 58 53.6&\t\t&\t\t\t& \tt\t&--38\t& --39,--21       \t&\t7\t& omg \\\\ G\\,319.399--0.012& 15 03 17.50 & --58 36 11.4  &--4.5 \t&--20,1     \t\t&10 \t\t& --5\t\t& --7,2\t\t&\t3.8\t& oc\t\\\\ G\\,319.836--0.196& 15 06 54.54 & --58 32 58.6  &--13\t&--23,0     \t\t&9  \t\t& --11\t& --19,0\t\t&\t2.8\t& omg\t\\\\ G\\,320.120--0.440& 15 09 43.83 & --58 37 06.3  &--46  \t&--70,--40  \t&1.2\t\t& --46\t&--158,30\t\t&\t0.9\t& o\t\\\\ G\\,320.221--0.281& 15 09 47.00 & --58 25 47.6 &\t\t&\t\t\t&$<$0.4\t&--73\t& --75,--72\t&\t0.7\t&\t\\\\ G\\,320.232--0.284& 15 09 51.92 & --58 25 38.0 &--63  \t&--70,--61  \t&9  \t\t& --67\t&--81,--72\t\t&\t2\t& om \\\\ % G\\,320.233--0.284& 15 09 52.50 & --58 25 35.7 &--60\t& --61,--58\t& 3.5\t\t&--60\t& --61,--54         &   \t3.4\t& c\t\\\\ G\\,320.255--0.305& 15 10 06.14 & --58 26 00.8 &\t\t&\t\t\t&\t--\t&--126 \t&  --144,--111\t&\t45\t& \t\\\\ G\\,320.285--0.308& 15 10 18.88 & --58 25 16.5 &\t\t&\t\t\t&\t--\t&--69\t&  --81,--56\t&\t15\t&g\t\\\\ G\\,321.028--0.484& 15 15 50.90 & --58 11 19.7 & --60 \t&--70,--55  \t&2.5\t\t& --58\t& --70,--52\t&\t5\t&\t\\\\ G\\,321.033--0.483 & 15 15 52.60 & --58 11 07.2 & \t--60\t&--68,--58\t\t&3.6\t\t& --61\t& --64,--49\t&\t0.25\t& m\t \\\\ G\\,321.148--0.529& 15 16 48.25 & --58 09 50.1 &--64  \t&--66,--62  \t&0.8\t\t& --97\t& --98,--61\t&\t1.6\t& omg\t\\\\ G\\,322.158+0.636 & 15 18 34.52 & --56 38 24.7 &--73  \t&--81,--61  \t&5  \t\t& --76\t& --86,--65\t&\t2.7\t& om \\\\ G\\,322.165+0.625 & 15 18 39.74 & --56 38 46.7 & --40 \t&--67,--36  \t&9  \t\t& --39\t& --67,--36\t&\t2.8\t&g\t\\\\ G\\,323.740--0.263& 15 31 45.48 & --56 30 49.6 & --50 \t&--72,--46  \t&140\t& --50\t& --88,--42\t&\t70\t& omg \\\\ G\\,324.201+0.122 & 15 32 52.76 & --55 56 04.9 & --87 \t&--100,--47 \t&50 \t\t& --87\t& --100,--48\t& \t14\t&o\t\\\\ G\\,324.716+0.342 & 15 34 57.41 & --55 27 22.3 & --55 \t&--72,--47  \t&10 \t\t&--58\t& --81,--30\t& 26\t\t& omg \\\\ G\\,326.662+0.521 & 15 45 02.73 & --54 09 03.3 &--39  \t&--50,--34  \t&256\t& \t\t&\t\t\t&\t--\t&m\t\\\\ G\\,326.665+0.553 & 15 44 55.82 & --54 07 25.6 &\t\t&\t\t\t&\tt\t&--42 \t&--127,--39\t&\t16\t&g\t\\\\ % G\\,326.670+0.554 & 15 44 57.03 & --54 07 10.6 &--42  \t&--44,--40  \t&26 \t\t& --40\t& --48,8\t\t&\t101\t& o\t\\\\ G\\,326.780--0.241& 15 48 55.10 & --54 40 38.6 &--64  \t&--66,--60  \t&36 \t\t&--66\t&--92,--52\t\t&\t18\t& og\t\\\\ G\\,326.859--0.676& 15 51 13.82 & --54 58 03.6 &\t\t&\t\t\t& --\t\t&--103 \t&  --104,--103\t& \t0.42\t& mg\t\\\\ G\\,327.119+0.511 & 15 47 32.56 & --53 52 39.3 &--87  \t&--90,--80  \t&25 \t\t& --88\t&--89,--57\t\t&\t19\t& omg\t\\\\ G\\,327.291--0.578& 15 53 07.65 & --54 37 07.2 &--56\t&--80,--39  \t&400\t& --63\t&--84,--36\t\t&\t668\t& omg\t\\\\ G\\,327.391+0.200 & 15 50 18.31 & --53 57 06.1 &\t\t&\t\t\t&\t--\t&--86\t&--92,--86\t\t& 0.40\t& mg \t\\\\ G\\,327.402+0.445 & 15 49 19.32 & --53 45 13.8 &--80  \t&--83,--68  \t&230\t& --81\t& --84,--69\t& 195\t& o$^{\\#}$mcg \\\\ G\\,327.581--0.077& 15 52 29.50 & --54 02 51.6 &\t\t&\t\t\t&\t--\t&--101  \t& --102,--95\t& 1.0\t         &   \t\\\\ G\\,327.594--0.095& 15 52 38.19 & --54 03 11.5 &\t\t&\t\t\t&\t--\t& --99 \t& --102,--91\t& 0.8\t\t&g\t\\\\ G\\,327.619--0.111& 15 52 50.31 & --54 03 00.0 &\t\t&\t\t\t&\t--\t&--85 \t& --85,--84\t& 0.20\t& mg \t\\\\ G\\,327.935--0.123& 15 54 33.09 & --53 51 29.1 &\t\t&\t\t\t&\t--\t&--98\t& --100,--76\t&2.1\t\t&\t\\\\ G\\,328.236--0.548& 15 57 58.21 & --53 59 25.4 &--38  \t&--39,--37 \t&30  \t\t&--38    \t& --40,10\t\t&20\t\t&omcg\t\\\\ G\\,328.254--0.532& 15 57 59.69 & --53 58 00.7 &--50 \t&--53,--48 \t&200 \t& --50\t& --51,--48\t& 155\t& omg\t\\\\ G\\,328.306+0.432 & 15 54 05.91 & --53 11 37.4 &--96  \t&--97,--87\t\t&40  \t\t& --93\t& --96,--87\t&\t141\t& c\t\\\\ G\\,328.808+0.633 & 15 55 48.23 & --52 43 05.2 &--46  \t&--47,--44 \t& 10 \t\t& --46\t& --48,--44\t&\t4.4\t& omcg\t\\\\ G\\,329.021--0.186& 16 00 24.32 & --53 12 16.9 &--42  \t&--43,--41 \t&2.4   \t\t&    --44\t& --44,--43\t&0.34\t&g \t\\\\ G\\,329.029--0.199& 16 00 30.22 & --53 12 34.3 &--38  \t&--40,--37 \t&1.6 \t\t&\t\t&\t\t\t&$<$0.2\t& og\t\\\\ G\\,329.030--0.205& 16 00 31.90 & --53 12 48.7 &--39  \t&--54,--34 \t&8   \t\t&     --46 \t& --52,--35\t& 6\t\t& om\t\\\\ G\\,329.031--0.198& 16 00 30.34 & --53 12 26.5 &--39  \t&--54,33   \t\t&5   \t\t&  --52\t& --65,33\t\t& 1.2\t\t& omg \\\\ G\\,329.066--0.307& 16 01 09.89 & --53 16 01.5 &--48  \t&--50,--47 \t&1.4 \t\t& --45\t& --46,--45\t& 0.6\t\t& omg\t\\\\ G\\,329.183--0.313& 16 01 46.90 & --53 11 41.7 &--51\t&--66,--36 \t&24  \t\t& --50\t&\t--60,--39\t& 34\t\t& omg\t\\\\ G\\,329.342+0.130 & 16 00 38.87 & --52 45 22.9 &\t--112\t&\t--115,--95\t\t& 1.1\t\t& \t--112 \t& --114,--100\t&\t2.2\t\\\\ G\\,329.404--0.459& 16 03 31.81 & --53 09 30.8 &--113 \t&--117,--106\t&3.1 \t\t&     --113 \t& --116,--111\t&   2.4\t& \t\\\\ G\\,329.405--0.459& 16 03 32.15 & --53 09 29.0 &--78  \t&--80,--60 \t&15  \t\t&\t--77\t&\t--79,--44\t& 10\t\t& om\t\\\\ G\\,329.407--0.459& 16 03 32.77 & --53 09 25.0 &--74  \t&--76,--72 \t& 80 \t\t&      \t\t&                      \t&  $<$0.2\t&  mg\\\\ G\\,329.421--0.167& 16 02 19.85 & --52 55 41.8 &\t\t&\t\t\t&$<$0.8\t& --77\t&  --78,--75        & \t2.4\t&\t\\\\ G\\,329.424--0.164& 16 02 20.03 & --52 55 25.9 &--78  \t&--83,--60 \t&0.8 \t\t&                 &      \t\t&$<$0.2\t&g\t\\\\   \\hline    \\end{tabular} \\end{table*}  \\begin{table*}\\addtocounter{table}{-1} \\caption{-- {\\emph {continued}}} \\begin{tabular}{lccrcrrcrl} \\hline \\multicolumn{1}{c}{\\bf Water maser} & {\\bf RA}  & {\\bf Dec} & {\\bf Vpeak} &{\\bf Vrange} & {\\bf Speak} & {\\bf Vpeak} &{\\bf Vrange} & {\\bf Speak} & {\\bf Associations}\\\\ \\multicolumn{1}{c}{\\bf ($l,b$)} & {\\bf (J2000)}  & {\\bf (J2000)}&  {\\bf (\\kmsns)}&  {\\bf (\\kmsns)} & {\\bf (Jy)} &{\\bf (\\kmsns)}&  {\\bf (\\kmsns)} & {\\bf (Jy)}     \\\\ \\multicolumn{1}{c}{\\bf (degrees)} & {\\bf (h m s)}&{\\bf ($^{o}$ $'$ $``$)}& {\\bf 2003} & {\\bf 2003} & {\\bf 2003}& {\\bf 2004} & {\\bf 2004} & {\\bf 2004} \\\\ \\hline \\\\ G\\,329.426--0.161& 16 02 19.71 & --52 55 12.8 &--73\t&--78,--71 \t&8   \t\t&   --73\t& --74,--72\t&\t10\t&\t\\\\ G\\,329.457+0.503 & 15 59 36.93 & --52 23 53.6 &\t\t&\t\t\t&\t--\t&--66\t& --68,--65\t&\t4.3\t&\t\\\\ G\\,329.622+0.138 & 16 02 00.28 & --52 33 57.7 &\t\t&\t\t\t&\t--\t&--82\t& --110,--66\t&\t30\t& mg\t\\\\ G\\,330.070+1.064 & 16 00 15.56 & --51 34 25.7 &\t\t&\t\t\t&\t--\t&--50\t&  --75,--45        &\t11\t& m$\\gamma$\t\\\\ G\\,330.879--0.367& 16 10 20.04 & --52 06 06.8 &--64  \t&--72,--28 \t& 90 \t\t& --60\t& --72,--25\t& 95\t\t& omc \\\\ G\\,330.954--0.182& 16 09 52.65 & --51 54 54.6 &--80  \t&--150,70  \t&240 \t& --91\t& --191,56\t& 323\t& ocg\t\\\\ G\\,331.132--0.244& 16 10 59.73 & --51 50 22.5 &--99\t&--102,--73\t&280 \t& --99\t& --118,--55\t& 47\t\t& omg\t\\\\ G\\,331.278--0.188& 16 11 26.51 & --51 41 55.8 &--86  \t&--104,--79\t&55  \t\t& --90\t& --118,--64\t& 42\t\t& omg\t\\\\ G\\,331.342--0.346& 16 12 26.49 & --51 46 14.9 &--60  \t&--62,--59 \t&1.8 \t\t& --62\t& --64,--60\t& 11\t\t& omg \\\\ G\\,331.418+0.252 & 16 10 10.56 & --51 16 52.2 &\t\t&\t\t\t&\t--\t&--71\t&--71,--70  \t& 0.6\t\t& g\t\\\\ G\\,331.442--0.187& 16 12 12.46 & --51 35 09.3 &--88\t&--93,--72 \t&70  \t\t& --88\t&--113,--80\t& 212\t&\tomcg \\\\ G\\,331.512--0.103& 16 12 10.01 & --51 28 36.7 &--89  \t&--162,--33\t&700 \t& --90\t& --159,--32\t& 534\t&\tocg \\\\ G\\,331.555--0.122& 16 12 27.20 & --51 27 42.6 &--99  \t&--170,--96\t&20  \t\t& --99\t& --166,--86\t& 9\t\t&\t\\\\ G\\,332.094--0.421& 16 16 16.68 & --51 18 26.2 &\t\t&\t\t\t&\t--\t& --59\t&   --59,--58\t&   0.29\t&\tm \\\\ G\\,332.296--0.094& 16 15 45.84 & --50 55 52.7 &\t\t&\t\t\t&\t--\t&--50\t& --71,--43\t& 6\t\t&\tm\\\\ G\\,332.349--0.433& 16 17 30.03 & --51 08 16.6 &\t\t&\t\t\t&\t--\t& --67\t& --68,--67\t& 2.9\t\t&\t\\\\ G\\,332.352--0.117& 16 16 07.10 & --50 54 31.0 &--45  \t&--48,--41 \t&2.0   \t\t& --60\t&--61,--60\t\t& 0.5\t\t&\tom \\\\ G\\,332.604--0.167& 16 17 29.45 & --50 46 11.7 &\t\t&\t\t\t&\t--\t&--46\t& --48,--45\t& 2.4\t\t&\tmg \\\\ G\\,332.725--0.621& 16 20 02.91 & --51 00 33.1\t&--42  \t&--43,--41 \t&0.6 \t\t& --58\t& --59,--56\t& 5\t\t&\tomg \\\\ G\\,332.826--0.549& 16 20 11.17 & --50 53 14.6 &--56\t&--72,--30 \t&45  \t\t& --59\t& --71,--35\t& 70\t\t&\tmc \\\\ G\\,332.964--0.679& 16 21 23.03 & --50 52 57.3 &\t\t&\t\t\t&\t--\t&--52\t&  --52,--50\t& 2.9\t\t&\tmg \\\\ G\\,333.030--0.063& 16 18 56.86 & --50 23 53.6 &\t\t&\t\t\t&\t--\t&--40\t& --153,--40\t& 3.4\t\t&\tmc\\\\ G\\,333.055--0.436& 16 20 42.47 & --50 38 46.4 &\t\t&\t\t\t&\t--\t&--47\t& --57,--48\t&0.39\t&\t\\\\ G\\,333.114--0.439& 16 20 59.30 & --50 36 21.9 &\t--62\t&\t--63,--60\t& 4.6  \t&--62\t& --62,--53\t& 2.0\t\t&\t\\\\ G\\,333.121--0.434& 16 20 59.70 & --50 35 50.8 &--57  \t&--59,12   \t\t&38 \t\t&  --47\t& --91,--33\t& 21\t\t&\tm\\\\ G\\,333.126--0.440& 16 21 02.69 & --50 35 54.1 &--50  \t&--70,--48 \t&19  \t\t&   --52\t& --72,--47\t& 12\t\t&\tm \\\\ G\\,333.128--0.440& 16 21 03.18 & --50 35 51.8 &\t\t&\t\t\t&$<$0.2  \t&--124\t& --125,--124\t& 0.79\t&\tm\\\\ G\\,333.130--0.425& 16 20 59.75 & --50 35 05.1 & --39 \t&--67,--38 \t&23  \t\t& --64\t& --65,--31\t&1.9\t\t&\t\\\\ G\\,333.132--0.560& 16 21 36.46 & --50 40 45.4 &\t\t&\t\t\t&\t--\t&--53\t&   --67,--46\t& 1.9\t\t&\t\\\\ G\\,333.219--0.062& 16 19 47.40 & --50 15 53.8 &\t \t&          \t\t&  $<$0.3\t& --13\t& --14,84\t\t&   0.5\t& g\t\\\\ G\\,333.234--0.060& 16 19 50.85 & --50 15 09.7 &--88  \t&--102,--83\t&140 \t&--88\t& --102,82\t& 117\t&\tog \\\\ G\\,333.315+0.106 & 16 19 28.75 & --50 04 39.7 &--48\t&--68,--41 \t&2.2 \t\t& --48\t& --60,--48\t& 6\t\t&\tomg \\\\ G\\,333.387+0.032 & 16 20 07.52 & --50 04 47.4 &--61\t&--63,--60 \t&0.4 \t\t&  --61\t& --61,--60\t& 0.14\t&\tomg \\\\ G\\,333.467--0.164& 16 21 20.20 & --50 09 46.1 &--44  \t&--46,--40 \t&3.5 \t\t& --42\t& --47,--40\t& 3.2\t\t&\tom\\\\ G\\,333.608--0.215& 16 22 11.08 & --50 05 56.3 &--51\t&--76,--45 \t&50  \t\t& --49\t& --83,--41\t& 24\t\t&\to\\\\ G\\,333.646+0.058 & 16 21 09.12 & --49 52 45.1 &\t\t&\t\t\t&\t--\t&--89\t&   --90,--84\t& 3.3\t\t&\tm \\\\ G\\,333.682--0.436& 16 23 29.67 & --50 12 07.4 &\t\t&\t\t\t& \t--\t&--3\t\t& --3,--2\t\t& 0.24\t&\tmg \\\\ G\\,333.930--0.134& 16 23 14.68 & --49 48 48.8 &\t\t&\t\t\t&\t--\t&--46\t& --50,--45\t&  0.18\t&\tm \\\\ G\\,334.635--0.015& 16 25 45.83 & --49 13 37.0 &\t\t&\t\t\t&\t--\t& --26\t& --29,--15\t& 49\t\t&\tmg \\\\% Only observed in 2004, at meth target G\\,334.935--0.098& 16 27 24.22 & --49 04 11.0 &\t\t&\t\t\t&\t--\t&--17\t& --18,--14\t&   1.0\t&\tmg\\\\ G\\,334.951--0.092& 16 27 26.96 & --49 03 14.7 &\t\t&\t\t\t&\t--\t&--21\t& --25,--20\t&  1.3\t&\t\\\\ G\\,335.059--0.428& 16 29 23.20& --49 12 31.3 & \t\t&\t\t\t&   t\t\t& --38\t& --39,--37\t& 1.7\t\t& \t\\\\ % G\\,335.060--0.428& 16 29 23.24 & --49 12 28.0 &--46  \t&--50,--37 \t&12  \t\t&--37\t& --44,15\t\t& 3.0\t\t&\tomg\\\\ G\\,335.070--0.423& 16 29 24.72 & --49 11 47.7 &--88\t& --109,--84\t&  1.0\t&--90\t& --105,--84\t& 5\t\t& g\t\\\\ G\\,335.585--0.285& 16 30 57.34 & --48 43 39.4 &--45  \t&--50,--40 \t&30  \t\t& --42\t& --49,--32\t&25\t\t&\tomg \\\\ G\\,335.586--0.290& 16 30 58.73 & --48 43 51.2 &--48  \t&--61,--33 \t& 5  \t\t& --56     \t& --57,--42\t& 20\t\t&\tomg \\\\ G\\,335.588--0.264& 16 30 52.52 & --48 42 39.5 &--51  \t&--56,--48 \t& 16  \t&      \t\t&                      \t&  $<$0.2\t&g\t\\\\ G\\,335.727+0.191 & 16 29 27.52 & --48 17 51.9 &\t\t&\t\t\t&\t--\t&--51\t& --51,--42\t& 13\t\t&\tm\\\\ G\\,335.787+0.177 & 16 29 46.18 & --48 15 49.1 &--55  \t&--56,--48 \t& 3 .0\t&   --49   \t& --59,--45\t&\t10\t&\t\\\\ G\\,335.789+0.174 & 16 29 47.33 & --48 15 50.8 &--46\t&--51,--45 \t& 3.0  \t&\t\t&\t\t\t& $<$0.2\t&\tomg \\\\ G\\,335.789+0.183 & 16 29 45.10 & --48 15 30.4 &--91  \t&--112,--89\t&4.2 \t\t&\t\t&\t\t\t& $<$0.2\t&\t\\\\ G\\,336.018--0.827& 16 35 09.35 & --48 46 47.7 &--54  \t&--59,--36 \t&120 \t&  --54\t& --59,--36\t&\t82\t&\tomc\\\\ G\\,336.352--0.149& 16 33 30.73 & --48 04 27.6 &\t--79\t& --81,--78\t& 0.4\t\t& --79\t& --81,--78\t&    0.43\t&\t\\\\ G\\,336.359--0.137& 16 33 29.37 & --48 03 41.5 &--67\t&--67,--66 \t&0.5 \t\t&\t\t&\t\t\t&$<$0.2\t&\tomc\\\\ G\\,336.433--0.262& 16 34 20.31 & --48 05 30.5 &\t\t&\t\t\t&\t--\t&--89\t& --90,--88\t& 0.6\t\t&\tmg \\\\ G\\,336.496--0.258& 16 34 34.52 & --48 02 34.3 &\t\t&\t\t\t&\t--\t&--25\t& --38,--14\t& 8\t\t&\t\\\\ G\\,336.830--0.375& 16 36 26.19 & --47 52 29.5 &\t\t&\t\t\t&\t--\t&--20\t& --45,--19\t& 0.28\t&\tmg \\\\ G\\,336.864+0.005 & 16 34 54.50 & --47 35 37.7 &--78\t&--80,--64 \t&3.0   \t& --66\t& --79,--65\t& 2.2\t\t&\tom\\\\ G\\,336.864--0.002& 16 34 56.02 & --47 35 55.3 &--73  \t&--76,--71 \t&4.5   \t& --73      \t&    --77,--71    \t& 1.5\t\t&\t\\\\ G\\,336.870--0.003& 16 34 57.91 & --47 35 42.7 &\t--77\t& --78,--55\t& 2.0\t\t&--77\t&   --78,--71\t& 3.7\t\t& \t\\\\ \\hline  \\end{tabular} \\end{table*}   \\begin{table*}\\addtocounter{table}{-1} \\caption{-- {\\emph {continued}}} \\begin{tabular}{lccrcrrcrl} \\hline \\multicolumn{1}{c}{\\bf Water maser} & {\\bf RA}  & {\\bf Dec} & {\\bf Vpeak} &{\\bf Vrange} & {\\bf Speak} & {\\bf Vpeak} &{\\bf Vrange} & {\\bf Speak} & {\\bf Associations}\\\\ \\multicolumn{1}{c}{\\bf ($l,b$)} & {\\bf (J2000)}  & {\\bf (J2000)}&  {\\bf (\\kmsns)}&  {\\bf (\\kmsns)} & {\\bf (Jy)} &{\\bf (\\kmsns)}&  {\\bf (\\kmsns)} & {\\bf (Jy)}     \\\\ \\multicolumn{1}{c}{\\bf (degrees)} & {\\bf (h m s)}&{\\bf ($^{o}$ $'$ $``$)}& {\\bf 2003} & {\\bf 2003} & {\\bf 2003}& {\\bf 2004} & {\\bf 2004} & {\\bf 2004} \\\\ \\hline \\\\ G\\,336.983--0.183& 16 36 12.38 & --47 37 59.1 &\t--76 \t&\t--77,--74 \t\t& 0.4\t\t& 45  \t& --78,45\t\t&  0.18\t& o$^{\\#}$mc\t\\\\ % G\\,336.991--0.024& 16 35 32.53 & --47 31 12.4 & --48 \t&--61,--44 \t& 4  \t\t& --49\t& --54,--44\t& 1.0\t\t&\tc\\\\ G\\,336.994--0.027& 16 35 34.01 & --47 31 12.2 &--121 \t&--137,--80\t&160 \t& --120\t& --177,--47\t& 158\t&\tomg \\\\ G\\,336.995--0.024& 16 35 33.53 & --47 31 00.6 &\t\t&\t\t\t&$<$0.2\t& --54\t& --54,--48\t& 1.1          &  g\t\\\\ G\\,337.258--0.101& 16 36 56.36 & --47 22 27.5 &--69  \t&--71,--52 \t&1.6 \t\t& --69\t& --69,--68\t& 0.33\t&\tomg \\\\  G\\,337.404--0.402& 16 38 50.57 & --47 28 00.8 &--40\t&--53,--37 \t&140 \t&--40\t& --49,--35\t&\t137\t&\tomcg \\\\ G\\,337.612--0.060& 16 38 09.46 & --47 05 00.3 &--52  \t&--101,--47\t&38  \t\t& --51\t& --99,--46\t& 17\t\t&\tomg \\\\ G\\,337.687+0.137 & 16 37 35.60 & --46 53 46.3 &\t\t&\t\t\t&\t--\t&--74\t& --151,--73\t&      0.6\t&\tmg \\\\ G\\,337.705--0.053& 16 38 29.72 & --47 00 35.7 &--48  \t&--147,5   \t\t&54  \t\t& --49\t& --159,--31\t&\t39\t&\tomcg \\\\ G\\,337.916--0.477& 16 41 10.49 & --47 08 02.9 &--46\t&--80,--28 \t&400 \t& --33\t& --65,--27\t& 321\t&\to\\\\ G\\,337.920--0.456& 16 41 06.14 & --47 07 02.3 &--40\t&--42,--39 \t&50  \t\t& --40\t& --69,--27\t& 22\t\t&\tom \\\\ G\\,337.994+0.133 & 16 38 48.63 & --46 40 15.7 &--114 \t&--117,--110\t&4.5\t\t& --113\t& --116,--111\t& 4.5\t\t&\t\\\\ G\\,337.998+0.137 & 16 38 48.53 & --46 39 56.4 &--39  \t&--50,--32 \t&30  \t\t& --38\t& --47,--30\t& 27\t\t&\tomg \\\\ G\\,338.069+0.011 & 16 39 37.98 & --46 41 44.9 &--37  \t&--40,--21 \t&1.3\t\t&  --28\t& --46,--22\t& 3.9\t\t&\tg \\\\ G\\,338.075+0.012 & 16 39 38.88 & --46 41 26.8 &--50  \t&--51,--50 \t&0.5\t\t&\t\t&\t\t\t& $<$0.2\t&\tomc\\\\ G\\,338.075+0.010 & 16 39 39.78 & --46 41 31.8 &--132 \t&--139,--21\t&1.2\t\t& --48\t& --51,--28\t& 1.2\t\t&\tm\\\\ G\\,338.077+0.019 & 16 39 37.69 & --46 41 05.6 &\t--40\t&\t--41,--38\t& 1.0 \t& --40\t&   --40,--39\t& 1.3\t\t&\tg \\\\ G\\,338.281+0.542 & 16 38 09.16 & --46 11 02.6 &--61 \t&--68,--59  \t&12  \t\t& --64\t& --67,--58\t& 4.6\t\t&\tomg \\\\ G\\,338.427+0.051 & 16 40 50.25 & --46 24 05.5 &\t\t&\t\t\t& --\t\t&--30\t&  --46,--29\t&   0.48\t&\t\t\\\\ G\\,338.430+0.053 & 16 40 50.50 & --46 23 53.0 &\t\t&\t\t\t&\t--\t&--44\t&   --51,--29\t& 3.2\t\t&\tg \\\\ G\\,338.433+0.057 & 16 40 50.08 & --46 23 33.8 &\t\t&\t\t\t&\t--\t&--29\t& --30,--29\t&  0.19\t&\tm\\\\ G\\,338.435+0.055 & 16 40 51.23 & --46 23 34.5 &\t\t&\t\t\t&\t--\t&--31\t& --31,--30\t&   0.58\t&\t\\\\ G\\,338.436+0.057 & 16 40 51.05 & --46 23 26.8 &\t\t&\t\t\t&\t--\t&--35\t&  --56,--34\t&    0.30\t&\t\\\\ G\\,338.440+0.064 & 16 40 49.87 & --46 22 59.5 &\t\t&\t\t\t&\t--\t&--81\t& --81,--80\t& 0.6\t\t&\tg \\\\ G\\,338.461--0.245& 16 42 15.57 & --46 34 18.6 &--53 \t&--119,--51 \t&9   \t\t& --52\t& --114,--49\t& 8\t\t& \tomg \\\\ G\\,338.462--0.259& 16 42 19.49 & --46 34 48.3 &\t\t&\t\t\t& $<$0.2\t&--54\t&--55,--53\t\t& \t0.35\t&\t\\\\ G\\,338.472+0.289 & 16 39 58.99 & --46 12 36.2 &--54 \t&--62,--22  \t&25  \t\t& --29\t& --62,--25\t& 10\t\t&\tomg \\\\ G\\,338.562+0.217 & 16 40 38.12 & --46 11 24.8 &\t\t&\t\t\t&\t--\t&--39\t& --40,--39\t&  0.22\t&\tmg \\\\ G\\,338.567+0.110 & 16 41 07.16 & --46 15 28.1 &\t\t&\t\t\t&\t--\t&--76\t& --92,--74\t& 3.2\t\t&\tm \\\\ G\\,338.682--0.084& 16 42 24.12 & --46 17 59.1 &--6  \t&--18,--5   \t&2.2 \t\t& --16\t& --18,--6\t\t& 0.7\t\t&\tocg \\\\ G\\,338.920+0.550 & 16 40 34.02 & --45 42 07.9& \t--110&\t--125,--65\t& 17\t\t& --68\t& --127,--64\t& 5.3\t\t&\tmg \\\\ G\\,338.925+0.556 & 16 40 33.63 & --45 41 37.9 &--63 \t&--85,--61  \t&160 \t& --62\t&--86,--6\t\t& 116\t&\tom\\\\ G\\,338.924--0.060& 16 43 13.25 & --46 06 04.2 &\t\t&\t\t\t&\t--\t&--64\t& --69,--63\t& 2.3\t\t&\t\\\\ G\\,339.582--0.127& 16 45 58.88 & --45 38 47.4 &\t\t&\t\t\t&\t--\t&--28\t& --32,--26\t& 0.27\t&\tmg \\\\ G\\,339.584--0.128& 16 45 59.63 & --45 38 44.3 &\t\t&\t\t\t&\t--\t&--40\t& --41,--26\t& 1.8\t\t&\tg \\\\ G\\,339.585--0.126& 16 45 59.18 & --45 38 36.7 &\t\t&\t\t\t&\t--\t&--41\t& --75,--39\t&  0.40\t&\tg \\\\ G\\,339.586--0.128& 16 45 59.92 & --45 38 39.7 &\t\t&\t\t\t&\t--\t&--106\t& --112,--105\t&  0.5\t&\tg \\\\ G\\,339.609--0.115& 16 46 01.57 & --45 37 07.3 &\t\t&\t\t\t& $<$0.2\t&--71\t&  --82,--67\t& 3.0\t\t&\t\\\\ G\\,339.622--0.121& 16 46 06.03 & --45 36 44.5 &--34 \t&--38,--32  \t&90  \t\t&--33\t& --36,--32\t& 43\t\t&\tom \\\\ G\\,339.762+0.055 & 16 45 51.56 & --45 23 31.0 &\t\t&\t\t\t&\t--\t&--57\t&  --58,--57\t&   0.25\t&\tmg \\\\ G\\,339.884--1.259& 16 52 04.71 & --46 08 33.6 &--32 \t&--50,--28  \t&50  \t\t& --51\t& --52,--24\t& 4.3\t\t&\tom$\\gamma$\\\\ G\\,340.054--0.243& 16 48 13.82 & --45 21 43.9 &--54 \t&--58,--44  \t&35  \t\t&      \t\t&\t\t\t&\t--\t&\tomg \\\\ G\\,340.785--0.096& 16 50 14.84 & --44 42 24.7 &--120\t&--121,--111\t&1.0   \t\t&\t\t&\t\t\t& $<$0.2\t&\tomg \\\\ G\\,341.218--0.212& 16 52 17.92 & --44 26 51.6 &--43 \t&--47,--38  \t&120 \t& --39\t& --49,--23\t& 33\t\t&\tomg \\\\ G\\,341.276+0.062 & 16 51 19.50 & --44 13 44.0 &--77 \t&--80,--58  \t&5   \t\t& --64\t& --83,--62\t& 10\t\t&\tomg \\\\ G\\,342.484+0.183 & 16 55 02.39 & --43 12 59.3 &\t\t&          \t\t&\t--\t&--43\t& --44,--36\t& 0.6\t\t&\tmg \\\\ G\\,343.126--0.065& 16 58 17.54 & --42 52 15.8&\t--16\t& --17,--15\t& 2.5\t\t&--21\t& --23,--20\t& 8\t\t&\t\\\\ G\\,343.127--0.063& 16 58 17.29 & --42 52 06.6 &--35 \t&--45,--20  \t&250         & --30\t& --46,--16\t& 208\t&\to\\\\ G\\,344.226--0.576& 17 04 09.36 & --42 18 58.0 &\t\t&\t\t\t&$<$0.2\t&--19\t&  --19,--18\t&   0.33\t&\t\\\\ G\\,344.228--0.569& 17 04 07.85 & --42 18 38.9 & --24\t&--28,31    \t&8   \t\t& --25\t& --52,--9\t\t&\t1.4\t&\tomg \\\\ G\\,344.421+0.046 & 17 02 08.53 & --41 46 56.4 &--27 \t&--32,--18  \t&1.8 \t\t& --26\t& --31,--24\t& 0.9\t\t&\t\\\\ G\\,344.582--0.024& 17 02 57.94 & --41 41 54.1 &--4  \t&--20,5     \t\t&250 \t& --4\t\t& --52,6\t\t&\t127\t&\tomcg \\\\ G\\,345.004--0.224& 17 05 11.12 & --41 29 04.1 &--23\t&--46,16   \t\t&3.0   \t\t& 15\t\t& --89,15\t\t&\t3.7\t&\tom$^{\\#}$cg \\\\ G\\,345.010+1.793 & 16 56 47.51 & --40 14 23.9 &--17\t& --19,10    \t&2.0   \t\t&\t\t&\t\t\t& $<$0.2\t&\tomc$\\gamma$\\\\ G\\,345.010+1.802 & 16 56 45.41 & --40 14 04.2 &--28\t&--33,--26   \t&7   \t\t& --25\t& --31,--18\t& 23\t\t&\t$\\gamma$\\\\ G\\,345.012+1.797 & 16 56 47.01 & --40 14 08.5 &--12\t&--13,7      \t& 50 \t\t& --12\t& --12,7\t\t& 29\t\t&\tm$\\gamma$\\\\ G\\,345.397--0.950& 17 09 33.08 & --41 36 20.4 &\t--21\t& --22,--20\t&3.1\t\t&--21\t&  --27,--20\t& 6\t\t&\t\\\\ G\\,345.402--0.948& 17 09 33.62 & --41 36 02.9 &--26\t&--32,--21   \t&8   \t\t& --23\t& --32,--21\t& 2.9\t\t&\t\\\\ G\\,345.405--0.947& 17 09 33.83 & --41 35 50.8 &\t\t&\t\t\t& $<$0.5\t&--28\t&  --28,--23\t&   0.5\t&\t\\\\ G\\,345.406--0.942& 17 09 32.71 & --41 35 37.6 &--15\t&--19,--12   \t&3.0 \t\t&  --18\t& --20,--17\t& 1.9\t\t&\t\\\\  \\hline \\end{tabular} \\end{table*}  \\begin{table*}\\addtocounter{table}{-1} \\caption{-- {\\emph {continued}}} \\begin{tabular}{lccrcrrcrl} \\hline \\multicolumn{1}{c}{\\bf Water maser} & {\\bf RA}  & {\\bf Dec} & {\\bf Vpeak} &{\\bf Vrange} & {\\bf Speak} & {\\bf Vpeak} &{\\bf Vrange} & {\\bf Speak} & {\\bf Associations}\\\\ \\multicolumn{1}{c}{\\bf ($l,b$)} & {\\bf (J2000)}  & {\\bf (J2000)}&  {\\bf (\\kmsns)}&  {\\bf (\\kmsns)} & {\\bf (Jy)} &{\\bf (\\kmsns)}&  {\\bf (\\kmsns)} & {\\bf (Jy)}     \\\\ \\multicolumn{1}{c}{\\bf (degrees)} & {\\bf (h m s)}&{\\bf ($^{o}$ $'$ $``$)}& {\\bf 2003} & {\\bf 2003} & {\\bf 2003}& {\\bf 2004} & {\\bf 2004} & {\\bf 2004} \\\\ \\hline\\\\ G\\,345.408--0.953& 17 09 35.85 & --41 35 56.5 &\t\t&\t\t\t&$<$0.5\t&--15\t& --16,36\t\t&    0.5\t&\to$^{\\#}$m$^{\\#}$c\\\\ G\\,345.412--0.955& 17 09 37.08 & --41 35 48.9 &\t\t&\t\t\t&$<$0.2\t&--55\t& --55,--54\t& 1.6\t\t&\t\\\\ G\\,345.425--0.951& 17 09 38.64 & --41 35 03.3 & --13\t&\t--15,--12\t&1.5\t\t&--13\t& --16,--13\t& 1.1\t\t&\tm \\\\ G\\,345.438--0.074& 17 05 56.75 & --41 02 54.9 &--11\t&--32,--9    \t& 40 \t\t& --12\t& --37,--8\t\t& 15\t\t&\to\\\\ G\\,345.482+0.309 & 17 04 28.41 & --40 46 52.1 &--52\t&--80,--24   \t&15  \t\t& --55\t& --82,--51\t& 1.8\t\t&\t \\\\  G\\,345.487+0.314 & 17 04 28.19 & --40 46 28.6 &--16\t&--24,2      \t&6   \t\t& --13\t& --39,--12\t& 0.7\t\t&\tm \\\\ G\\,345.493+1.469 & 16 59 41.47 & --40 03 46.2 &5   \t\t&3,6         \t\t&6   \t\t&\t\t&\t\t\t& $<$0.2\t&\to$\\gamma$\\\\ G\\,345.494+1.470 & 16 59 41.15 & --40 03 39.9 &\t1\t&\t--16,2\t&  1.0\t& 0\t \t& --18,--4\t\t& 1.8\t\t&\t$\\gamma$\\\\ G\\,345.495+1.473 & 16 59 40.78 & --40 03 28.9 &--60  \t&--62,--16 \t&4.0   \t\t& --9\t\t& --10,--8\t\t& 1.5\t\t&\t$\\gamma$\\\\ G\\,345.505+0.348 & 17 04 23.02 & --40 44 23.5 &--3   \t&--43,1    \t\t&22  \t\t&--4\t\t& --42,4\t\t& 4.5\t\t&\tom \\\\ G\\,345.505+0.343 & 17 04 24.35 & --40 44 32.1 &\t\t&\t\t\t& $<$0.2\t& --68\t& --70,--67\t& 0.8\t\t&\t\\\\ G\\,345.699--0.090& 17 06 50.72 & --40 50 58.9 &--10\t&--87,102    \t&200 \t& --5\t\t& --92,141\t& 216\t&\tog \\\\ G\\,346.480+0.132 & 17 08 22.67 & --40 05 26.9 &--10\t&--12,--8      \t&1.2 \t\t&\t\t&\t\t\t& $<$0.2\t&\tomg \\\\  G\\,346.522+0.085 & 17 08 42.35 & --40 05 08.1 & \t\t&\t\t\t&\t--\t&4\t\t& --2,14\t\t& 2.9\t\t&\tm \\\\ G\\,346.529+0.106 & 17 08 38.12 & --40 04 03.9 &\t\t&\t\t\t&\t--\t& 4\t\t& --1,5\t\t& 1.6\t\t&\t\\\\ G\\,347.588+0.213 & 17 11 27.61 & --39 09 08.1 &\t\t&\t\t\t&\t--\t&--93\t& --94,--93\t& 0.7\t\t&\t\\\\ G\\,347.623+0.148 & 17 11 50.09 & --39 09 45.8 &--118\t&--120,--117\t&0.5 \t\t& --118\t& --122,--116\t& 1.0\t\t&\t\\\\ G\\,347.628+0.149 & 17 11 50.85 & --39 09 29.6 &--125\t&--133,--122 \t&0.5 \t\t&\t\t&\t\t\t& $<$0.2\t&\tomg \\\\ G\\,347.632+0.210 & 17 11 36.30 & --39 07 06.5 &\t\t&\t\t\t&\t--\t& --88\t& --95,--24\t&\t6\t&\tmc\\\\ G\\,348.533--0.974& 17 19 16.20 & --39 04 30.6 &--56 \t&--59,--10  \t&3.2 \t\t& --31\t& --93,24\t\t& \t1.8\t&\t\\\\ G\\,348.534--0.983& 17 19 18.64 & --39 04 47.3 &--21 \t&--30,--11  \t&2.2 \t\t& --14\t& --109,--12\t& 0.9\t\t&\t\\\\ G\\,348.551--0.979& 17 19 20.61 & --39 03 49.2 &--30 \t&--32,--15  \t&1.4 \t\t& --18\t& --18,--17\t& \t0.28\t&\tmg \\\\ G\\,348.726--1.038& 17 20 06.42 & --38 57 13.2 &--11 \t&--64,42    \t&115 \t& --10\t& --77,61\t\t&\t162\t&\t$\\gamma$\\\\ G\\,348.885+0.096 & 17 15 50.25 & --38 10 12.3 &--81 \t&--84,--77  \t&8   \t\t& --80\t& --87,--77\t&\t8\t&\tomg \\\\ G\\,348.892--0.180& 17 17 00.26 & --38 19 27.9 & 7   \t&6,12       \t\t&2.0   \t\t&\t\t&\t\t\t&\t--\t&\tomg \\\\ G\\,349.052+0.002 & 17 16 43.25 & --38 05 18.2 &15   \t&14,17      \t&3.0 \t\t&      \t\t&                    \t&    $<$0.2&\t \\\\ G\\,349.067--0.018& 17 16 50.82 & --38 05 13.8 & 5   \t&3,7        \t\t&1.6 \t\t& 13\t\t& --14,19\t\t&\t1.0\t&\tomg \\\\ G\\,349.074--0.015& 17 16 51.23 & --38 04 49.3 &--27 \t&--32,--20  \t&1.6\t\t& --26\t& --32,--24\t& 1.2\t\t&\tg \\\\ G\\,349.068+0.110 & 17 16 19.24 & --38 00 48.3 &--25 \t&--32,--22  \t&3.0 \t\t& --21\t& --22,--20\t&\t1.4\t&\t\\\\ G\\,349.092+0.105 & 17 16 24.66 & --37 59 45.4 &--80 \t&--84,--74  \t&43  \t\t& --80\t&   --84,--72\t& 154\t&\tomg \\\\ G\\,350.015+0.433 & 17 17 45.43 & --37 03 12.9 &--35 \t&--50,--26   \t&4.7 \t\t&\t\t&\t\t\t&\t--\t&\tomg \\\\ G\\,350.098+0.099 & 17 19 21.82 & --37 10 41.2 &--44 \t&--45,--42  \t&1.8 \t\t&      \t\t&                      \t&  $<$0.2\t&\t \\\\ G\\,350.098+0.080 & 17 19 26.67 & --37 11 20.8 &--66 \t&--68,--65  \t&6 \t\t& --67\t& --69,--67\t& 2.6\t\t&\tg \\\\ G\\,350.100+0.081 & 17 19 26.74 & --37 11 09.8 &\t--68\t&\t--69,--67\t&4.0\t\t& --68\t& --69,--63\t&\t5\t&\t\\\\ G\\,350.105+0.084 & 17 19 26.78 & --37 10 51.5 &--71 \t&--102,--38 \t&6 \t\t& --71\t& --74,--29\t& 14\t\t&\tm\\\\ G\\,350.110+0.087 & 17 19 26.96 & --37 10 29.0 &--72 \t&--74,--71  \t&8   \t\t& --72\t& --73,--70\t& 14\t\t&\tg \\\\ G\\,350.112+0.089 & 17 19 26.96 & --37 10 19.9 &--172\t&--175,--110\t&2.0\t\t& --128\t& --164,--106\t& 2.9\t\t&\t\\\\ G\\,350.113+0.095 & 17 19 25.69 & --37 10 04.8 &--66 \t&--70,--64  \t&2.5 \t\t& --64\t& --79,--63\t&\t1.4\t&\tog \\\\ G\\,350.274+0.120 & 17 19 47.07 & --37 01 16.6 &\t\t&\t\t\t&\t--\t&--63\t&  --63,--61\t& 2.9\t\t&\t \\\\ G\\,350.299+0.122 & 17 19 50.89 & --37 00 00.6 &\t\t&\t\t\t& \t--\t&--68\t& --68,--67\t& 0.17\t&\tm\\\\ G\\,350.330+0.100 & 17 20 01.79 & --36 59 14.3 &--68 \t&--69,--60  \t&1.0\t\t& --62\t& --68,--48\t& 1.9\t\t&\to\\\\ G\\,350.341+0.140 & 17 19 53.60 & --36 57 19.0&\t\t&\t\t\t& --\t\t&--105\t& --105,--58\t&\t3.2\t&\tg \\\\ G\\,350.686--0.491& 17 23 28.71 & --37 01 47.9 &--23  \t&--22,--24  \t&0.6\t\t& --14\t& --14,--13\t& 1.0\t\t&\tomg \\\\ G\\,350.690--0.490& 17 23 29.19 & --37 01 35.6 &\t\t&\t\t\t& $<$0.6\t& --22  \t&   --24,--21\t& 3.4\t\t&\t \\\\ G\\,351.160+0.696 & 17 19 57.50 & --35 57 54.1 &--3   \t&--11.5,--0.5\t&9  \t\t&   --3\t& --10,2\t\t& 18\t\t&    \tomc\\\\ G\\,351.163+0.696 & 17 19 58.13 & --35 57 48.9 &\t\t&\t\t\t& t\t\t&--10\t& --13,--9\t\t& 10\t\t&\t\\\\ % G\\,351.240+0.668 & 17 20 17.98 & --35 54 57.6&\t\t&\t\t\t&--\t& --24  \t& --38,31           & 21\t\t&\t\\\\ G\\,351.243+0.671 & 17 20 17.76 & --35 54 42.8 &\t\t&\t\t \t&\t--\t& \t--77\t& --108,84\t& 453\t& m\t\\\\\t% G\\,351.246+0.668 & 17 20 19.01 &--35 54 38.2 &\t\t&\t\t\t& \t--\t&21 \t\t& 19,22\t\t& 1.3\t\t&c\t\\\\ G\\,351.417+0.646 & 17 20 53.29 & --35 46 58.3 &--10  \t&--58,50    \t&1400\t&\t\t&\t\t\t&\t--\t&om\t\\\\ G\\,351.582--0.353& 17 25 25.35 & --36 12 44.0 &--89  \t&--120,--87 \t&1600\t&\t\t&\t\t\t&\t--\t&omg \t\\\\ G\\,351.775--0.536& 17 26 42.50 & --36 09 15.9 &--2   \t&--32,21     \t\t&85 \t\t&\t\t&\t\t\t&\t--\t&om\t\\\\ G\\,352.098+0.160 & 17 24 45.25 &--35 29 48.8 &\t\t&\t\t\t&\t--\t&--75  \t& --75,--73\t&\t2.5\t& g\t\\\\ G\\,352.111+0.176 & 17 24 43.79 &--35 28 37.7 &\t\t&\t\t\t&\t--\t&--60  \t& --61,--60\t& \t1.3\t&mg\t\\\\ G\\,352.133--0.944& 17 29 22.46 &--36 04 59.9 &\t\t&\t\t\t&\t--\t& --11\t&  --17,--6\t\t& \t27\t&mg\t\\\\ G\\,352.162+0.199 & 17 24 46.36 & --35 25 19.9 &--45 \t&--46,--44   \t&1  \t\t& --45\t& --119,--43\t& 0.28\t&og\t\\\\ G\\,352.517--0.155& 17 27 11.34 & --35 19 32.0 &--49 \t&--53,--46   \t&12 \t\t&      \t\t&\t\t\t&\t--\t&om\t\\\\ G\\,352.525--0.158\t& 17 27 13.44\t&--35 19 15.7\t& --51\t&--53,--49\t&\t3.6\t&\t&\t&\t--\t&\tmg\t\\\\ % G\\,352.623--1.076& 17 31 14.93 & --35 44 47.5 &2\t \t&0,2         \t\t&28 \t\t& 2\t\t& --7,3\t\t&\t8\t&\t$\\gamma$\\\\ G\\,352.630--1.067& 17 31 13.94 & --35 44 08.8 &--2\t&--10,20     \t&35 \t\t&  0\t\t& --13,18\t\t&\t700\t&om$\\gamma$\\\\  \\hline  \\end{tabular} \\end{table*}  \\begin{table*}\\addtocounter{table}{-1} \\caption{-- {\\emph {continued}}} \\begin{tabular}{lccrcrrcrl} \\hline \\multicolumn{1}{c}{\\bf Water maser} & {\\bf RA}  & {\\bf Dec} & {\\bf Vpeak} &{\\bf Vrange} & {\\bf Speak} & {\\bf Vpeak} &{\\bf Vrange} & {\\bf Speak} & {\\bf Associations}\\\\ \\multicolumn{1}{c}{\\bf ($l,b$)} & {\\bf (J2000)}  & {\\bf (J2000)}&  {\\bf (\\kmsns)}&  {\\bf (\\kmsns)} & {\\bf (Jy)} &{\\bf (\\kmsns)}&  {\\bf (\\kmsns)} & {\\bf (Jy)}     \\\\ \\multicolumn{1}{c}{\\bf (degrees)} & {\\bf (h m s)}&{\\bf ($^{o}$ $'$ $``$)}& {\\bf 2003} & {\\bf 2003} & {\\bf 2003}& {\\bf 2004} & {\\bf 2004} & {\\bf 2004} \\\\ \\hline \\\\ G\\,353.273+0.641 & 17 26 01.57 & --34 15 14.7&\t\t&\t\t\t&\t--\t&--49 \t& --110,--5\t&\t366\t&m\t\\\\ G\\,353.408--0.350& 17 30 23.38 &--34 41 29.8 &\t\t&\t\t\t& $<$0.2\t&--14\t& --15,--13\t& 0.6\t\t& g\t\\\\ G\\,353.411--0.356& 17 30 25.31 & --34 41 35.0&\t\t&\t\t\t& $<$0.2\t&--20 \t& --21,--17\t& 6.4\t\t&\t\\\\ G\\,353.411--0.362& 17 30 26.78 & --34 41 46.6&\t\t&\t\t\t&$<$0.2\t&--7\t\t& --7,--5\t\t& 1.8\t\t& c\t\\\\ G\\,353.413--0.367& 17 30 28.47 & --34 41 49.1 &--17 \t&--29,--8    \t&2.0\t\t& --20\t& --26,--9\t\t&\t7\t& g\t\\\\ G\\,353.413--0.365& 17 30 27.84 & --34 41 44.6 &--20 \t&--24,--18    \t&4.5\t\t& --19\t& --25,--18\t&\t2.2\t& \t\\\\ G\\,353.414--0.363& 17 30 27.56 & --34 41 38.9 &--10 \t&--11,--9    \t&2  \t\t&    1\t\t& --22,7\t\t&\t1.7\t&\t\\\\ G\\,353.463+0.563 & 17 26 51.25 & --34 08 25.3 &--47 \t&--52,--46   \t& 1 \t\t&\t\t&\t\t\t&\t--\t& g\t\\\\ G\\,353.464+0.562 & 17 26 51.61 & --34 08 24.1 & --60\t& --61,--59         & 0.7         &                &                        &     --         & omg\\\\ G\\,354.594+0.469 & 17 30 14.47 & --33 15 03.2 &--22 \t&--25,--20   \t&9\t\t& \t\t&\t\t\t&$<$0.2\t&\t\\\\ G\\,354.615+0.472 & 17 30 17.22 & --33 13 54.4 &--38 \t&--40,--12   \t&1.6\t\t&\t\t&\t\t\t&$<$0.2\t& om\t\\\\ G\\,354.703+0.297 & 17 31 12.96 & --33 15 17.6&\t\t&\t\t\t& $<$0.4\t&105\t& 104,112\t\t&\t1.6\t&\t\\\\ G\\,354.712+0.293 & 17 31 15.38 & --33 14 57.8 &96   \t&95,105       \t&1.1\t\t&      \t\t&\t\t\t&$<$0.2\t& g\t\\\\ G\\,354.722+0.302& 17 31 14.87 & --33 14 08.3 &\t\t&\t\t\t&$<$0.2\t&99\t\t& 96,101\t\t&\t0.38\t& \t\\\\ G\\,355.343+0.147 & 17 33 29.00 & --32 48 00.1 & 17  \t&9,35        \t\t&34 \t\t&\t\t&\t\t\t&--\t\t& om\t\\\\ G\\,355.345+0.149 & 17 33 28.83 & --32 47 49.9 & 72\t& 8,109              & 2.2         &                &                          &  --           & m \\\\ G\\,357.965--0.164& 17 41 20.12 & --30 45 15.9 & --4\t& --5,--3\t\t& 53\t\t&--19 \t& --20,--19\t& 0.9\t\t&mg\t\\\\ G\\,357.967--0.163& 17 41 20.30 & --30 45 07.0 &0    \t&--80,100    \t&40 \t\t& --65\t&--81,87\t\t& 57\t\t&om\t\\\\ G\\,358.371--0.468& 17 43 32.01 & --30 34 11.4 &10   \t&--5,12     \t\t&2.7   \t\t& 2\t\t&2,31\t\t& 1.4\t\t&mg\t\\\\ G\\,358.386--0.483& 17 43 37.72 & --30 33 50.6 &--1  \t&--4,1      \t\t&4.5 \t\t& 0\t\t& --3,1\t\t& 3.6\t\t&omcg\t\\\\ G\\,359.137+0.032 & 17 43 25.61 & --29 39 18.6 &--1  \t&--117,26    \t&300 \t&\t\t&\t\t\t&\t--\t&omg\t\\\\ G\\,359.419--0.104& 17 44 38.27 & --29 29 12.0&\t\t&\t\t\t&$<$0.2\t&--26 \t& --59,2\t\t& 1.0\t\t&\t\\\\ G\\,359.436--0.102& 17 44 40.23 & --29 28 12.2 &--59 \t&--61,--50  \t&15  \t\t& --59\t& --61,--54\t& 9\t\t&mg\t\\\\ G\\,359.436--0.104& 17 44 40.66 & --29 28 16.1&--56\t& --57,--55\t& 1.5\t\t& --47\t& --48,--47\t& 0.37\t&om\t\\\\ G\\,359.441--0.111& 17 44 42.87 & --29 28 15.4 &--65 \t&--66,--64  \t&1.8 \t\t&      \t\t&                      \t&$<$0.2\t&\t\\\\ G\\,359.442--0.106& 17 44 42.00 & --29 28 00.6 &--53 \t&--54,--52  \t&9   \t\t& --53\t& --56,--52\t& 0.928\t& g \t\\\\ G\\,359.442--0.104& 17 44 41.47 & --29 27 58.8&\t\t&\t\t\t& t\t\t&--60 \t& --62,--58\t& 0.8\t\t&\t\\\\% text confused G\\,359.443--0.104& 17 44 41.82 & --29 27 56.5&\t\t&\t\t\t& t\t\t& --49\t& --52,--48 \t& 0.6\t\t& g\t\\\\ % G\\,359.615--0.243& 17 45 39.12 & --29 23 29.6 &22   \t&--16,73    \t&7   \t\t& 64\t\t& --15,68\t\t& 14\t\t&omg\t\\\\ G\\,359.969--0.457& 17 47 20.08 & --29 11 59.0 &11   \t&10,16      \t&27  \t\t&\t\t&\t\t\t& \t--\t&om\t\\\\ G\\,0.209--0.002  & 17 46 07.44 &--28 45 32.1 &\t\t&\t\t\t&\t--\t&39\t\t& 18,42\t\t&\t2.2\t&c\t\\\\ G\\,0.212--0.002  & 17 46 07.86 &--28 45 23.0 &\t\t&\t\t\t&\t--\t&56\t\t& 55,65\t\t&\t0.9\t&m\t\\\\ G\\,0.216--0.023  & 17 46 13.20 &--28 45 49.8 &\t\t&\t\t\t&\t--\t&--11\t& --16,8\t\t&\t1.7\t& g\t\\\\ G\\,0.308--0.177  & 17 47 02.38 &--28 45 55.2 &\t\t&\t\t\t&\t--\t&--26 \t&  --27,--26\t&\t1.8\t& g\t\\\\ G\\,0.316--0.201  & 17 47 09.29 &--28 46 15.5 &\t\t&\t\t\t&\t--\t&23\t\t&   14,31\t\t&\t22\t&mg\t\\\\ G\\,0.376+0.040   & 17 46 21.29 & --28 35 39.7 &40   \t\t&5,58       \t\t&55  \t\t&\t\t&\t\t\t&\t--\t&omg\t\\\\ G\\,0.497+0.188   & 17 46 03.96 & --28 24 50.7 &--8  \t\t&--72,25    \t&7   \t\t& 26\t\t& --50,34\t\t&\t2.5\t&omg\t\\\\ G\\,0.547--0.851  & 17 50 14.52 & --28 54 28.8 &20   \t\t&--60,110   \t&200 \t&\t\t&\t\t\t&\t--\t&omg\t\\\\ G\\,0.655--0.045 & 17 47 20.88 & --28 23 59.9  &  --52      & --55,--50           &      4.0   &       &      &    --&\\\\ % G\\,0.657--0.042 & 17 47 20.44 & --28 23 46.4  &   62        & 60,64                &    70      &        &       &   --& og\\\\% associtaions? strongest OH maser in sgrb2 south no meth, no ex oh G\\,0.665--0.032 & 17 47 19.07 & --28 23 03.5  & 4            &       0,7                  & 4.3          &               &         & -- &g\\\\ % G\\,0.668--0.035  & 17 47 20.26 & --28 23 02.6 &59   \t\t&50,90      \t&408 \t&\t\t&\t\t\t&\t--\t& og\t\\\\ G\\,0.677--0.028 & 17 47 19.64 & --28 22 18.9  & 36           & 14,132              & 330       &       &       &   --    & g\\\\ % G\\,2.143+0.009   & 17 50 36.13 & --27 05 46.4 &37   \t\t&35,70      \t&0.6 \t\t&\t\t&\t\t\t&\t--\t&omg\t\\\\ G\\,2.536+0.198   & 17 50 46.66 &--26 39 44.9 &\t\t&\t\t\t&\t--\t&25\t\t& --1,63\t\t&\t30\t&m\t\\\\ G\\,5.886--0.392  & 18 00 30.52 & --24 03 58.2 &11   \t\t&--10,23    \t&40  \t\t&\t\t&\t\t\t&\t--\t&\to$^{\\#}$m$^{\\#}$\\\\ G\\,5.897--0.445  & 18 00 43.90 & --24 04 58.1&\t\t&\t\t\t&\t--\t&19\t\t& 19,20\t&\t0.9\t&\t\\\\ G\\,5.901--0.430  & 18 00 40.95 &--24 04 19.6 &\t\t&\t\t\t&\t--\t&14\t\t&   --40,30\t&\t3.9\t&m\t\\\\ G\\,5.913--0.388  & 18 00 33.12 & --24 02 28.4 &--50 \t\t&--65,--47  \t&9.2 \t\t&      \t\t&                      \t&\t--\t& g\t\\\\ G\\,6.049--1.447  & 18 04 53.19 & --24 26 40.3 &20   \t\t&18,22      \t&1.0   \t\t&\t\t&\t\t\t&\t--\t&o$\\gamma$\t\\\\ G\\,6.534--0.105  & 18 00 49.44 &--23 21 40.2 &\t\t&\t\t\t&\t--\t&22\t\t&  6,23\t\t&\t0.5\t& g\t\\\\ G\\,6.611--0.082  & 18 00 54.15 &--23 17 00.8 & \t\t&\t\t\t&\t--\t&6\t\t& --34,11\t\t&\t7\t&mg\t\\\\ G\\,6.796--0.257  & 18 01 57.71 & --23 12 32.5 &14   \t\t&10,20      \t&10  \t\t& 1\t\t& --1,24\t\t&\t5\t&omg\t\\\\ G\\,8.139+0.226   & 18 03 00.83 &--21 48 09.5 &\t\t&\t\t\t&\t--\t&16\t\t& 1,24\t\t&\t18\t&mg\t\\\\ G\\,8.670--0.356  & 18 06 19.17 &--21 37 31.0 &34\t\t&30,44      \t&10  \t\t& 36\t\t& --9,43\t\t&\t3.4\t&omc\t\\\\ G\\,8.673--0.354  & 18 06 18.98 &--21 37 17.9 &\t--4\t&\t--23,--3\t&0.5\t\t&--8\t\t&--20,4\t\t&\t0.36\t&\t\\\\ G\\,9.620+0.194   & 18 06 14.97 & --20 31 37.5 &5    \t\t&--30,50    \t&25  \t\t& 6\t\t& --19,11\t\t&\t12\t&o\t\\\\ G\\,9.622+0.195   & 18 06 14.88 &--20 31 31.0 &\t21\t&\t0,22\t\t& 1.4\t\t& 22\t\t& 0,23\t\t&   1.3\t&om$^{\\#}$\t\\\\ G\\,9.986--0.028  & 18 07 50.20 &--20 18 55.9 &\t\t&\t\t\t&\t--\t&49\t\t& 44,60\t\t&\t4.0\t&mg\t\\\\ G\\,10.288--0.125 & 18 08 49.44 &--20 05 57.4 &\t\t&\t\t\t&\t--\t&9 \t\t& 9,11\t\t&\t0.5\t&m\t\\\\ \\hline  \\end{tabular} \\end{table*}  \\begin{table*}\\addtocounter{table}{-1} \\caption{-- {\\emph {continued}}} \\begin{tabular}{lccrcrrcrl} \\hline \\multicolumn{1}{c}{\\bf Water maser} & {\\bf RA}  & {\\bf Dec} & {\\bf Vpeak} &{\\bf Vrange} & {\\bf Speak} & {\\bf Vpeak} &{\\bf Vrange} & {\\bf Speak} & {\\bf Associations}\\\\ \\multicolumn{1}{c}{\\bf ($l,b$)} & {\\bf (J2000)}  & {\\bf (J2000)}&  {\\bf (\\kmsns)}&  {\\bf (\\kmsns)} & {\\bf (Jy)} &{\\bf (\\kmsns)}&  {\\bf (\\kmsns)} & {\\bf (Jy)}     \\\\ \\multicolumn{1}{c}{\\bf (degrees)} & {\\bf (h m s)}&{\\bf ($^{o}$ $'$ $``$)}& {\\bf 2003} & {\\bf 2003} & {\\bf 2003}& {\\bf 2004} & {\\bf 2004} & {\\bf 2004} \\\\ \\hline \\\\ G\\,10.307--0.270 & 18 09 24.29 & --20 09 08.8&\t\t&\t\t\t&\t--\t&33\t\t& 32,33\t\t&\t0.8\t&\t\\\\ G\\,10.323--0.160 & 18 09 01.57 &--20 05 07.6 &\t\t&\t\t\t&\t--\t&--3\t\t& --6,62\t\t&\t3.2\t&m\t\\\\ G\\,10.331--0.159 & 18 09 02.30 &--20 04 40.2 &\t\t&\t\t\t&\t--\t&11\t\t& 7,22\t\t& \t1.0\t&\t\\\\ G\\,10.342--0.143 & 18 09 00.11 & --20 03 35.8&\t\t&\t\t\t&\t--\t&8\t\t& --40,61\t\t&\t4.8\t&mg\t\\\\ G\\,10.445--0.018 & 18 08 45.01 & --19 54 35.1 &71  \t\t&58,85       \t&6   \t\t& 70\t\t& 69,81\t\t&\t2.7\t&\t\\\\ G\\,10.473+0.027  & 18 08 38.30 & --19 51 48.8 &60  \t\t&30,93       \t&45  \t\t& 62\t\t& 28,129\t\t& 169\t&omcg\t\\\\ G\\,10.480+0.034  & 18 08 37.69 & --19 51 12.4 &64  \t\t&63,65       \t&12  \t\t&\t\t&\t\t\t& $<$0.2\t&om\\\\ G\\,10.623--0.383 & 18 10 28.57 & --19 55 49.4 &2   \t\t&--11,5      \t&350 \t&\t\t&\t\t\t&\t--\t&og\t\\\\ G\\,10.625--0.335 & 18 10 18.07 & --19 54 20.5&\t\t&\t\t\t&\t--\t&\t--5 \t& --6,--5\t\t&\t0.40\t&\t\\\\ G\\,10.959+0.022  & 18 09 39.38 & --19 26 27.0&\t\t&\t\t\t&\t--\t&25\t\t& 7,26\t\t&\t2.9\t&\tmcg \\\\  G\\,11.034+0.062  & 18 09 39.75 & --19 21 21.1 &18  \t\t&16,19       \t&\t0.6 \t&\t\t&\t\t\t&\t--\t&\tom\\\\ G\\,11.498--1.486 & 18 16 22.32 & --19 41 26.1&\t\t&\t\t\t&\t--\t&17\t\t& --3,21\t\t&\t112\t&\tm$\\gamma$\\\\ G\\,11.903--0.142 & 18 12 11.41 & --18 41 33.0 &36  \t\t&35,38       \t&0.3 \t\t&      \t\t&\t\t\t& $<$0.2\t&\tom \\\\ G\\,12.203--0.107 & 18 12 40.25 & --18 24 47.4&\t33\t&\t32,35\t\t& 6&35\t\t& 32,37\t\t& 8\t\t&\tm\\\\ G\\,12.209--0.102 & 18 12 39.88 & --18 24 17.1 &21  \t\t&--12,42     \t&141  \t\t&   22\t&   --15,50\t\t&\t51\t&\tmcg \\\\ G\\,12.216--0.119 & 18 12 44.55 & --18 24 25.2 &26  \t\t&25,42       \t&6   \t\t& 26\t\t& 0,38\t\t& 10\t\t&\tog \\\\ G\\,12.681--0.183 & 18 13 54.82 & --18 01 47.0 &61\t\t&55,63       \t&1200 \t\t& 61\t\t& 17,62\t\t& 445\t&\tomg \\\\ % G\\,12.884+0.502  & 18 11 47.92 &--17 31 21.3 &\t42\t&\t41,43\t& 1.3\t\t&40\t\t& 34,41\t\t& 1.1\t\t&\t\\\\ G\\,12.889+0.489  & 18 11 51.54 & --17 31 27.9 &30  \t\t&28,32       \t&45  \t\t& 30\t\t& 28,39\t\t& 59\t\t&\tomg \\\\ G\\,12.901--0.242 & 18 14 34.50 & --17 51 51.6 &35  \t\t&23,37       \t&8\t\t& 36\t\t& 23,37\t\t& 12\t\t& g\t\\\\ G\\,12.908--0.260 & 18 14 39.45 & --17 52 01.4 &37  \t\t&    21,39      \t\t& 0.8\t\t&\t\t&\t\t\t&$<$0.2\t&\tom\\\\ G\\,15.016--0.679 & 18 20 23.11 & --16 12 38.8&\t\t&\t\t\t& t\t\t&19\t\t& 14,21\t\t&11\t\t&\t\\\\ % G\\,15.025--0.658 & 18 20 19.50 & --16 11 31.9&\t\t&            \t\t& t\t\t& 25\t\t&  14,31\t\t& 4.2\t\t&\t\\\\ G\\,15.026--0.654 & 18 20 18.72 & --16 11 23.3 &27  \t\t&27,31       \t&29   \t\t&  28\t\t& 26,31\t\t& 16\t\t&\t\\\\ G\\,15.028--0.673 & 18 20 23.04 & --16 11 48.4 &19  \t\t&17,28       \t&197 \t& 20\t\t& 14,31\t\t& 67\t\t&\t \\\\ G\\,15.032--0.667 & 18 20 22.35 & --16 11 27.2 &\t45\t&\t44,46\t& 1.8\t\t&40\t\t& --8,41\t\t& 1.1\t\t& g\t\\\\ G\\,15.032--0.670 & 18 20 23.01 & --16 11 30.1&\t\t&\t\t\t&t\t \t&26\t\t& 14,28            \t& 12\t\t&\t\\\\ G\\,15.034--0.667 & 18 20 22.58 & --16 11 18.4&\t\t&\t\t\t& t\t\t&17\t\t& 17,21\t\t& 3.8\t\t&\t\\\\ G\\,17.638+0.156  & 18 22 26.45 & --13 30 12.4 &27  \t\t&15,29       \t&230 \t& 27\t\t& 17,35\t\t& 245\t&\tom\\\\ \\hline \\end{tabular} \\end{table*}   \\section[]{Results}  The search for 22-GHz water masers carried out with the ATCA in 2003 October and 2004 July towards 202 OH maser sites and 104 methanol maser sites (with no reported OH maser emission) resulted in the detection of 379 distinct water maser sites (Table~\\ref{tab:masers}). Spectra of all detected sources are shown in Fig.~\\ref{fig:spectra}. For the majority of sources, the spectra are taken from the 2004 data, except where sources were either not observed or not detected at this epoch. For these latter cases we show the 2003 spectra and distinguish them from the 2004 spectra with a `2003' in the top left hand corner of each spectrum.  The 2003 spectra were obtained directly from the uv data with a phase shift to the source position, and amplitude correction for offset from the field centre.  For eight sources we show a spectrum from each epoch to either highlight that a weak source is a genuine detection (333.387+0.032 and 336.983-0.183) or to give an indication of the level of variability seen over the 10 month time-scale (284.350--0.418, 321.148--0.529, 327.291--0.578, 345.004--0.224, 15.026--0.654 and 15.028--0.673). A velocity range of 200 \\kms\\ is shown for the majority of sources, but there are several instances where we either decreased this value to clearly show individual features in spectra that are complex (or include multiple nearby sources) or increased it in order to display extremely high velocity features. A decreased velocity range of 100 \\kms\\ is shown for the following sources; 301.136--0.225, 301.136--0.226a, 301.136--0.226b, 301.137--0.225, 335.060--0.428/335.059--0.428, 336.991--0.024, 336.995--0.024, 359.441--0.111, 359.442--0.106, 359.442--0.104, 359.443--0.104. An increased range of 300 \\kms\\ was used for 320.120--0.440, 330.954--0.182, 333.219--0.062, 333.234--0.060, 345.699--0.090, 357.965--0.164, 357.967--0.163, 0.547--0.851, 0.668--0.035, 0.665--0.032, 0.655--0.045, 0.657--0.042 and 0.677--0.028.  A number of the sources that we detect have been observed previously and have been presented in the literature \\citep[e.g.][and references therein]{John72,C74,K76,GD77,B80,BS82,BE83,C+89,HC96} but the majority of these earlier observations (performed up to 20 years ago) were made with relatively poor positional accuracy.  Due to the intrinsically variable nature of water masers, many sources exhibit levels of variability so extreme that they display no common spectral features at epochs separated by many years. This, combined with the tendency of water masers to form in clusters, and the previously poor positional information, mean that it is almost impossible to accurately match up sources from the literature with our present data. We have therefore limited our references (in Section ~\\ref{sect:ind}) to previous detections of sources that were observed with high positional precision \\citep[e.g.][]{FC89,Breen,CP08}, or where there was little doubt that the sources were the same.   The majority of OH maser targets were observed in both 2003 and 2004, whereas the methanol maser targets were observed in 2004 only. Where appropriate data were available for both epochs, reported positions are the average of the two since, in general, it provides the most accurate positions for the sources. For sources north of declination --20\\degrees, we have used a weighting of 2:1 for the declinations in favour of the 2004 data to account for the three times more elongated beam of the 2003 observations (a consequence of the different array configurations). Sources that were observed at both epochs allowed a direct comparison of the positions for each of the sources and therefore afford verification of the positional uncertainties.  Additional to direct comparison of 2003 with 2004 data, an overall assessment of data quality and reliability was made in several other ways.  FC89 and FC99 used their VLA observations of a sample of more than 70 SFR targets with OH and water masers, to show that more than half were a simple association of water and OH masers coincident to within their combined relative errors (of typically 1 arcsec).  Subsequent observation of the more southerly OH masers in that sample with the ATCA (C98) showed that the most southerly ones (observed by the VLA inevitably at low elevation) had significantly larger position uncertainties, and corrections to the positions resulted in an increased number of close OH/water maser associations. Thus we may expect the majority of our sample to show a water maser position coincident with OH, and thus the OH position is an indirect check on the accuracy of the newly derived water positions.     A further assessment was made using the 35 masers north of declination --47 degrees which are present in the FC89 target list. The FC89 absolute positions for the more southerly targets, although of variable quality for the OH masers (where ionospheric effects at low elevation can be significant), appear to remain excellent for the water masers.  Thus we can directly compare our positions to those of FC89, to assess the errors in our current data. Furthermore, in some fields there is a strong ultracompact \\ionhy (\\UCHIIns) region that has been measured to subarcsecond accuracy, such as in the 6-GHz observations of C97 and C2001; where these are detectable in the current 22-GHz observations, they allow a further check on the positions, without the need for any assumptions concerning the true relative positions of the masers.    From these many comparisons, we are able to estimate our rms positional uncertainty as 2 arcsec.  The target OH and methanol masers have rms position uncertainties of 0.4 arcsec (C98, Caswell 2009). An additional positional uncertainty in characterising any water maser site by a single position arises because a single site sometimes consists of many separate spots with angular separations as extensive as 4 arcsec \\citep[e.g.][]{R88}, explicable by an outflow \\citep[e.g.][]{CP08}. We therefore regard our water maser sources to be associated with OH or methanol masers when they are separated by less than 3 arcsec, a threshold which captures most associations without diluting them with too many false, chance, coincidences;  see also further discussion in Section 5.2.  Where the water maser positions are derived from a single epoch, we relax this threshold to 4.5 arcsec. As the positions of the 22-GHz radio continuum have also been determined from a single epoch, a threshold of 4.5 arcsec is similarly adopted for continuum associations. Most of our proposed associations correspond to a much better accuracy (see Table 3) than our thresholds. There are, however, some more complex cases that have been judged on individual merit as discussed in detailed considerations summarized in Section~\\ref{sect:ind}.  For example, the required precision of agreement was relaxed for sources believed to be nearby, at a distance of less than 2 kpc. Comparison of our 379 water maser positions with the positions of OH and methanol masers shows that 128 are coincident with both species, 33 are coincident with OH only and 70 are coincident with the location of methanol masers (see Section 5.3 for more extensive discussion). Surprisingly, 148 sources have no association with other maser species and we describe these as `solitary'.  Details for the 379 sources that we detect are presented in Table~\\ref{tab:masers} and, following the usual practice, the Galactic longitude and latitude of each source, listed in the first column, is used as an identifying source name for each water maser.  These Galactic coordinates are derived from the more precise measurements of equatorial coordinates given in columns 2 and 3. The peak velocity and velocity range (w.r.t. lsr), followed by the peak flux density, are given in columns 4, 5 and 6 for the 2003 epoch and in columns 7, 8 and 9 for the 2004 epoch. The presence of a `--' in either column 6 or 9 indicates that no observations were made for that source during the 2003 or 2004 observations respectively and a `t' in either column indicates that there is a comment in the text of Section~\\ref{sect:ind} explaining the nature of the detection status at the indicated epoch. The presence of a number preceded by a `$<$' in either column 6 or 9 indicates that no emission above the quoted flux density was detected at that epoch. Column 10 gives a list of associations for each water maser source; here, the presence of an `o' denotes the presence of an associated OH maser, `m' the presence of an associated methanol maser, `c' the presence of associated 22-GHz continuum emission (in our observations) and `g' the presence of an associated GLIMPSE point source. A `$\\gamma$' in this column indicates that the source is outside the GLIMPSE survey region. A `$^{\\#}$' following an `o' indicates that the OH maser is strictly outside our association threshold but is associated with the methanol maser that falls within our threshold for a given source, meaning that either all three sources are coincident or the water maser is offset;  similarly for the case where a `$^{\\#}$' follows an `m'. In some cases we have the situation that both the OH and the methanol masers are strictly located outside the association threshold but we regard them as associated through special circumstances, in which case we have used the `$\\#$' after both the `o' and the `m'. In the case of 301.136-0.226 a second water maser site lies within the association threshold for the same methanol and OH masers and the association is shown in parentheses.  OH masers that were searched and resulted in no water maser detection are listed in Table~\\ref{tab:nowater}. The first column gives the name of the OH maser followed by its right ascension and declination. Column 4 gives the angular separation between the OH maser and the nearest methanol maser \\citep{C09} within 2.5 arcsec, and when a `--' is present, this signifies that there is no methanol maser within 2.5 arcsec of the OH maser.  An extensive list of the OH and methanol masers associated with our water maser sources, as well as water maser associations with 22-GHz continuum sources is given in Table~\\ref{tab:ass}.  All OH and methanol masers as well as 22-GHz radio continuum that fall within 5 arcsec of the water masers are presented. Column 1 in Table~\\ref{tab:ass} gives the water maser source name, and the name of the nearby OH and methanol masers are given in columns 2 and 4 respectively.  The source names, based on the precise positions of individual species, inevitably differ slightly in a few cases due to different small position errors. The angular separations between the water masers and OH masers are given in column 3, and between water and methanol masers in column 5. Columns 6, 7 and 8 give the peak velocity of the water (2004 values are given unless not available, in which case the 2003 value is used) and the coincident OH and methanol maser sources. Column 9 gives the Galactic coordinates of 22-GHz continuum sources that we detect, followed by the angular separation between the continuum source and the water maser in column 10. The discussion of individual sources in Section~\\ref{sect:ind} includes some comparisons between the positions of water maser sources and other masers, continuum and GLIMPSE sources.   A complete list of the continuum sources detected towards water maser sources in the 2004 observations is given in Table~\\ref{tab:cont}. Associations between the 29 water maser sources observed only during the 2003 observations and possibly associated 22-GHz continuum sources (see Section~\\ref{sect:cont}) have not been determined.          \\begin{table}  \\caption{OH masers with no associated water maser emission. Listed in column 1 is the OH maser source name followed in columns 2 and 3 by the right ascension and declination. Column 4 shows the angular separation between the listed OH maser and a nearby methanol maser; a -- in this column indicates that there is no known methanol maser sources within 2.5 arcsec.} \\begin{tabular}{lccrcrrcrl} \\hline \\multicolumn{1}{c}{\\bf OH maser}& {\\bf RA}  & {\\bf Dec} & {\\bf Methanol}\\\\ \\multicolumn{1}{c} {\\bf ($l,b$)} & {\\bf (J2000)}  & {\\bf (J2000)} & {\\bf maser}  \\\\ \\multicolumn{1}{c}{\\bf degrees}& {\\bf (h m s)}&{\\bf ($^{o}$ $'$ $``$)}& {\\bf sep. (arcsec)} \\\\ \\hline \\\\ G\\,232.621+0.996\t\t&07 32 09.82\t &--16 58 13.0\t& 0.7\\\\ G\\,300.969+1.147\t\t&12 34 53.24\t &--61 39 40.3\t&  0.5\\\\ G\\,305.200+0.019\t\t&13 11 16.90\t &--62 45 54.7\t&  0.5\\\\ G\\,305.202+0.208\t\t&13 11 10.61\t &--62 34 37.8\t&  1.3\\\\ G\\,306.322--0.334\t\t&13 21 23.00\t &--63 00 30.4\t&  0.9\\\\ G\\,309.921+0.479\t\t&13 50 41.73\t &--61 35 09.8\t&0.9 \\\\ G\\,313.705--0.190\t\t&14 22 34.72\t &--61 08 27.4\t& 0.8\\\\ G\\,316.359--0.362\t\t&14 43 11.00\t &--60 17 15.3\t& 2.5\\\\ G\\,321.030--0.485\t\t&15 15 51.67\t &--58 11 18.0\t& 0.9\\\\ G\\,323.459--0.079\t\t&15 29 19.36\t &--56 31 21.4\t& 1.4\\\\ G\\,328.307+0.430\t\t&15 54 06.48\t &--53 11 40.3\t& --\\\\ G\\,329.339+0.148\t\t&16 00 33.15\t &--52 44 39.8\t& 0.2 \\\\ G\\,331.542--0.066\t\t&16 12 09.05\t &--51 25 47.2\t&  0.5 \\\\ G\\,331.543--0.066\t\t&16 12 09.16\t &--51 25 45.3\t&0.2\\\\ G\\,331.556--0.121\t\t&16 12 27.19\t &--51 27 38.1\t&0.2\\\\ G\\,332.295+2.280\t\t&16 05 41.72\t &--49 11 30.5\t& 0.2 \\\\ G\\,332.824--0.548\t\t&16 20 10.23\t &--50 53 18.1\t& --\\\\ G\\,333.135--0.431\t\t&16 21 02.97\t &--50 35 10.1\t& 2.4\\\\ G\\,335.556--0.307\t\t&16 30 56.00\t &--48 45 51.0\t& 0.8\\\\ G\\,336.822+0.028\t\t&16 34 38.26\t &--47 36 33.0\t&  0.8 \\\\ G\\,336.941--0.156\t\t&16 35 55.22\t &--47 38 45.7\t& 0.4\\\\ G\\,338.875--0.084\t\t&16 43 08.23\t &--46 09 12.8\t&  0.2\\\\ G\\,339.053--0.315\t\t&16 44 49.16\t &--46 10 14.4\t& 2.2\\\\ G\\,339.282+0.136\t\t&16 43 43.12\t &--45 42 08.4\t& 0.4\\\\ G\\,339.682--1.207\t\t&16 51 06.21\t &--46 15 57.8\t& 0.4\\\\ G\\,343.930+0.125\t\t&17 00 10.92\t &--42 07 18.7\t& 0.6 \\\\ G\\,344.419+0.044\t\t&17 02 08.67\t &--41 47 08.6\t& 1.8 \\\\ G\\,345.498+1.467\t\t&16 59 42.81\t &--40 03 36.2\t&  0.4\\\\ G\\,347.870+0.014\t\t&17 13 08.80\t &--39 02 29.5\t& --\\\\ G\\,348.550--0.979\t\t&17 19 20.39\t &--39 03 51.8\t&0.3 \\\\ G\\,348.579--0.920\t\t&17 19 10.56\t &--39 00 24.5\t&  0.6\\\\ G\\,348.698--1.027\t\t&17 19 58.91\t &--38 58 14.1\t& --\\\\ G\\,348.703--1.043\t\t&17 20 03.96\t &--38 58 31.3\t&1.2\\\\ G\\,348.727--1.037\t\t&17 20 06.55\t &--38 57 08.2\t&  0.9\\\\ G\\,350.011--1.342\t\t&17 25 06.50\t &--38 04 00.7\t& 0.5\\\\ G\\,353.410--0.360\t\t&17 30 26.20\t &--34 41 45.5\t\t&0.3\\\\ G\\,354.724+0.300\t\t&17 31 15.52\t &--33 14 05.3\t&0.5\\\\ G\\,356.662--0.264\t\t&17 38 29.22\t &--31 54 40.6\t& 2.0 \\\\ G\\,3.910+0.001\t\t&17 54 38.77\t &--25 34 45.2\t&0.5\\\\ G\\,8.683--0.368\t\t&18 06 23.46\t &--21 37 10.2\t& 0.4\\\\ G\\,12.025-0.031\t\t&18 12 01.88\t &--18 31 55.6\t&0.3\\\\ G\\,15.034--0.677\t\t&18 20 24.75\t &--16 11 34.9\t& 0.6\\\\ \\hline \\end{tabular} \\label{tab:nowater} \\end{table}  \\begin{figure*} \\psfig{figure=fig1_page1.ps} \\caption{Spectra of the 22-GHz water masers detected in 2004 towards sites of OH and methanol masers.} \\label{fig:spectra} \\end{figure*}    \\begin{figure*}\\addtocounter{figure}{-1} \\psfig{figure=fig1_page2.ps} \\caption{--{\\emph {continuued}}} \\end{figure*}  \\begin{figure*}\\addtocounter{figure}{-1} \\psfig{figure=fig1_page3.ps} \\caption{--{\\emph {continuued}}} \\end{figure*}  \\begin{figure*}\\addtocounter{figure}{-1} \\psfig{figure=fig1_page4.ps} \\caption{--{\\emph {continuued}}} \\end{figure*}  \\begin{figure*}\\addtocounter{figure}{-1} \\psfig{figure=fig1_page5.ps} \\caption{--{\\emph {continuued}}} \\end{figure*}   \\begin{figure*}\\addtocounter{figure}{-1} \\psfig{figure=fig1_page6.ps} \\caption{--{\\emph {continuued}}} \\end{figure*}  \\begin{figure*}\\addtocounter{figure}{-1} \\psfig{figure=fig1_page7.ps} \\caption{--{\\emph {continuued}}} \\end{figure*}  \\begin{figure*}\\addtocounter{figure}{-1} \\psfig{figure=fig1_page8.ps} \\caption{--{\\emph {continuued}}} \\end{figure*}  \\begin{figure*}\\addtocounter{figure}{-1} \\psfig{figure=fig1_page9.ps} \\caption{--{\\emph {continuued}}} \\end{figure*}  \\begin{figure*}\\addtocounter{figure}{-1} \\psfig{figure=fig1_page10.ps} \\caption{--{\\emph {continuued}}} \\end{figure*}  \\begin{figure*}\\addtocounter{figure}{-1} \\psfig{figure=fig1_page11.ps} \\caption{--{\\emph {continuued}}} \\end{figure*}  \\begin{figure*}\\addtocounter{figure}{-1} \\psfig{figure=fig1_page12.ps} \\caption{--{\\emph {continuued}}} \\end{figure*}  \\begin{figure*}\\addtocounter{figure}{-1} \\psfig{figure=fig1_page13.ps} \\caption{--{\\emph {continuued}}} \\end{figure*}    \\begin{figure*}\\addtocounter{figure}{-1} \\psfig{figure=fig1_page14.ps} \\caption{--{\\emph {continuued}}} \\end{figure*}  \\begin{figure*}\\addtocounter{figure}{-1} \\psfig{figure=fig1_page15.ps} \\caption{--{\\emph {continuued}}} \\end{figure*}  \\begin{figure*}\\addtocounter{figure}{-1} \\psfig{figure=fig1_page16.ps} \\caption{--{\\emph {continuued}}} \\end{figure*}  \\begin{figure*}\\addtocounter{figure}{-1} \\psfig{figure=fig1_page17.ps} \\caption{--{\\emph {continuued}}} \\end{figure*}   \\clearpage   \\begin{figure*}\\addtocounter{figure}{-1} \\psfig{figure=fig1_page18.ps} \\caption{--{\\emph {continuued}}} \\end{figure*}      \\begin{table*} \\caption{Water maser sources with associated OH and methanol masers as well as 22-GHz continuum emission. Column 1 shows the water maser source name; column 2 gives the source name of the nearest OH maser within 5 arcsec (-- if none) of the detected water maser; column 3 gives the angular separation between the water and the OH masers; column 4 gives the source name of the nearest methanol maser within 5 arcsec (-- if none) of the detected water maser; column 5 gives the angular separation between the water maser and the methanol maser; columns 6, 7 and 8 give the water maser, OH maser and methanol maser peak velocities; column 9 gives detected \\UCHII regions with 5 arcsec of the detected water masers (-- if none); and column 10 gives the angular separation between the \\UCHII region and the detected water maser.} \\begin{tabular}{llclcllrlcl} \\hline \\multicolumn{1}{c}{\\bf Water} & {\\bf OH} & {\\bf Sep.} & {\\bf Methanol}  & {\\bf Sep.} & {\\bf Water}& {\\bf OH} & {\\bf Methanol} & {\\bf Continuum} & {\\bf Sep.} \\\\ \\multicolumn{1}{c}{\\bf ($l,b$)}  &{\\bf ($l,b$)} & &{\\bf ($l,b$)} & & {\\bf Vpeak} &  {\\bf Vpeak} &  {\\bf Vpeak}&{\\bf ($l,b$)} \\\\ \\multicolumn{1}{c}{\\bf (degrees)} & {\\bf (degrees)} & {\\bf (arcsec)}&{\\bf (degrees)} &{\\bf (arcsec)}& {\\bf (\\kms)} & {\\bf (\\kms)} &{\\bf (\\kms)} & {\\bf (degrees)} &{\\bf (arcsec)} \\\\ \\hline \\\\ G\\,240.316+0.071\t\t&\tG\\,240.316+0.071\t&\t0.7\t&\t--\t\t\t\t&            \t&   \t89\t&\t63\t&\t\t&\t\t--\t\t\t\t&\t\t&\t\\\\ G\\,263.250+0.514\t \t&\tG\\,263.250+0.514\t&\t2.1\t&\tG\\,263.250+0.514\t&\t1.5   \t& \t20\t&\t15.3\t&\t12.3\t&\t\t--\t\t\t\t&\t\t&\t\\\\ G\\,284.350--0.418\t \t&\tG\\,284.351--0.418\t&\t1.1\t&\t--\t\t\t\t&      \t\t&\t7\t&\t6\t&\t\t&\t\t--\t\t\t\t&\t\t&\t\\\\ G\\,285.263--0.050\t\t&\tG\\,285.263--0.050 \t&\t2.0\t&\t--\t\t\t\t&\t\t&\t3\t&\t6\t&\t\t&\t--\t\t\t\t\t&\t\t&\\\\ G\\,287.371+0.644\t \t&\tG\\,287.371+0.644 \t&\t1.6\t&\tG\\,287.371+0.644\t&\t1.5    &\t--11\t&\t--4\t&\t--1.8\t&\t\t--\t\t\t\t&\t\t&\\\\ G\\,290.374+1.661\t \t&\tG\\,290.374+1.661 \t&\t1.6\t&\tG\\,290.374+1.661 \t&\t1.0    &\t--12\t& --23.3\t& --24.2\t&\t\t--\t\t\t\t&\t\t&\\\\ G\\,291.270--0.719\t \t&\t--\t\t\t\t&\t \t&\tG\\,291.270--0.719\t&      2.5    &\t--102\t&\t\t& --26.5\t&\t\t--\t\t\t\t&\t\t&\\\\ G\\,291.274--0.709\t \t&\tG\\,291.274--0.709\t&\t1.3\t&\tG\\,291.274--0.709\t&\t0.7   \t&\t--32\t& --24.5\t& --29.6\t&\t\t--\t\t\t\t&\t\t&\\\\ G\\,291.579--0.431\t \t&\tG\\,291.579--0.431\t&\t0.6\t&   \tG\\,291.579--0.431  \t&\t0.7   \t&\t13\t&\t13\t&\t14.5\t&\t\t--\t\t\t\t&\t\t&\\\\ G\\,291.581--0.435\t \t&\t--\t\t\t\t&\t\t&\tG\\,291.582--0.435\t&      3.8    &\t26\t&\t\t&\t10.5\t&\t\t--\t\t\t\t&\t\t&\t\\\\ G\\,291.610--0.529\t \t&\tG\\,291.610--0.529 \t&\t0.7\t&\t--\t\t\t\t&            \t&\t12\t&\t18\t&\t\t&\tG\\,291.611--0.529\t\t&\t2.6 \t&\\\\ G\\,291.627--0.529\t \t&\t--\t\t\t\t&\t\t&\t--\t\t\t\t&      \t\t&\t\t&\t\t&\t\t&\tG\\,291.626--0.531 \t\t&\t4.8\t&\t\\\\ G\\,294.511--1.622\t \t&\tG\\,294.511--1.621 \t&\t2.1\t&\tG\\,294.511--1.621\t&\t1.8   \t&\t--12\t&--12.7\t&--12.3\t&\t\t--\t\t\t\t&\t\t&\\\\ G\\,294.989--1.719\t \t&\t--\t\t\t\t&\t\t&\tG\\,294.990--1.719\t&      2.5    & \t--17\t&\t\t&\t--12.3\t&\t\t--\t\t\t\t&\t\t&\\\\ G\\,297.660--0.974\t\t&\tG\\,297.660--0.973\t&\t2.0\t&\t--\t\t\t\t&            \t&\t26\t&\t27.6\t&\t\t&\t--\t\t\t\t\t&\t\t&\t\\\\ G\\,299.013+0.128\t \t&\tG\\,299.013+0.128\t&\t1.2\t&\tG\\,299.013+0.128\t&\t1.1    &\t19\t&\t20.3\t&\t18.4\t&\tG\\,299.012+0.128\t\t&\t3.3\t&\t\\\\ G\\,300.504--0.176\t \t&\tG\\,300.504--0.176 \t&\t0.6\t&\tG\\,300.504--0.176\t&\t1.8    &\t11\t&\t22.4\t&\t7.5\t&\t--\t\t\t     \t\t&\t\t&\\\\ G\\,301.136--0.226b\t \t&\tG\\,301.136--0.226\t&\t2.0\t&\tG\\,301.136--0.226\t&\t 2.5  \t&\t--44\t&  --40.2\t&\t--39.8&\tG\\,301.136--0.226     \t\t&\t1.0\t&\\\\ G\\,301.137--0.225\t \t&\tG\\,301.136--0.226\t&\t2.0\t&\tG\\,301.136--0.226\t&      2.0    &\t--35\t& --40.2\t&--39.8\t&\t--\t\t\t\t\t&\t\t&\\\\ G\\,305.208+0.207\t\t&      G\\,305.208+0.206\t&\t3.0\t&\tG\\,305.208+0.206  \t&\t2.7\t&\t--42\t&\t--38\t&\t--38.3\t&\t--\t\t\t\t\t&\t\t&\\\\ G\\,305.361+0.150\t \t&\tG\\,305.362+0.150 \t&\t2.1\t&\tG\\,305.362+0.150  \t&\t2.0\t&\t--36\t&\t--39.5\t&\t--36.5\t&\t--\t\t\t\t    \t&\t\t&\\\\ G\\,305.799--0.245\t \t&\tG\\,305.799--0.245\t&\t3.0\t&\tG\\,305.799--0.245  \t&    \t2.5    & \t--34\t&\t--36.7\t&\t--39.5\t&\t--\t\t\t&\\\\ G\\,307.805--0.456\t \t&\tG\\,307.805--0.456\t&\t1.5\t&\t--\t\t\t\t&\t\t& \t--7\t&\t--14.5\t&\t\t&\t\t--\t\t\t&\\\\ G\\,308.754+0.549\t \t&\tG\\,308.754+0.549\t&\t0.8\t&\tG\\,308.754+0.549\t&\t1.4    &\t--48\t&\t--43.5\t&\t--51.0\t&\t--\t\t\t&\\\\ G\\,308.918+0.124\t \t&\tG\\,308.918+0.123 \t&\t3.0\t&\tG\\,308.918+0.123\t&\t3.6    &\t--61\t&\t--54\t&\t--54.7\t&\t\t--\t\t\t&\\\\ G\\,309.384--0.135\t \t&\tG\\,309.384--0.135\t&\t1.3\t&\tG\\,309.384--0.135\t&\t0.6    &\t--50\t&\t--52\t&\t--49.6\t&\t\t--\t\t\t&\\\\ G\\,310.144+0.760\t \t&\tG\\,310.144+0.760\t&\t2.2\t&\tG\\,310.144+0.760   \t&\t1.0\t&\t--63\t&\t--57\t&\t--55.6\t&\t\t--\t\t\t&\\\\ G\\,311.643--0.380\t \t&\tG\\,311.643--0.380 \t&\t1.3\t&\tG\\,311.643--0.380   \t&\t0.4\t&\t36\t&\t38\t&\t32.5\t&\tG\\,311.643--0.380\t\t&\t1.3\t&\\\\ G\\,312.109+0.262\t \t&\t--\t\t\t\t&\t\t&\tG\\,312.108+0.262\t&     1.9     &\t--48\t&\t\t& \t--50.0\t&\t\t--\t\t\t&\\\\ G\\,312.596+0.045\t \t&\t--\t\t\t\t&\t\t&\tG\\,312.597+0.045\t&     1.6     &\t--59\t&\t\t&\t--60.0\t&   \t\t--\t\t\t&\\\\ G\\,312.599+0.046\t \t&\tG\\,312.598+0.045\t&\t2.1\t&\tG\\,312.598+0.045 \t&\t2.1   \t&\t--79\t&\t--65.2\t&\t--67.9\t&\t--\t\t\t&\\\\ G\\,313.457+0.193\t \t&\t--\t\t\t\t&\t\t&\t--\t\t\t        \t&            \t&    \t\t&\t\t&\t\t&\tG\\,313.458+0.193\t\t&     2.1\t&\\\\ G\\,313.470+0.191\t \t&\tG\\,313.469+0.190\t&\t0.9\t&\tG\\,313.469+0.190\t&\t1.1    &\t--15\t&\t--10\t&\t--9.4\t&--\\\\ G\\,313.578+0.325\t \t&\tG\\,313.577+0.325\t&      1.0\t&    \tG\\,313.577+0.325\t&\t1.9    &\t--47\t&\t--47\t&\t--47.9\t&\t--\t\\\\ G\\,313.767--0.862\t \t&\tG\\,313.767--0.863\t&      1.1\t&\tG\\,313.767--0.863\t&\t0.9    &\t--54\t&\t--53.5\t&\t--54.6\t&--\\\\ G\\,314.320+0.112\t \t&\tG\\,314.320+0.112 \t&\t2.2  \t&\tG\\,314.320+0.112\t&\t2.3    &\t--45\t&\t--45\t&\t--43.7\t&--\\\\ G\\,316.361--0.363\t \t&\t--\t\t\t\t&\t\t&\tG\\,316.359--0.362\t&     3.2     \t&\t--3\t&\t\t&\t3.5\t&--\\\\ G\\,316.412--0.308\t \t&\tG\\,316.412--0.308  \t&\t1.5\t&\tG\\,316.412--0.308\t&\t2.3   \t&\t--20\t&\t--2\t&\t--5.7\t&\tG\\,316.412--0.308  \t\t&\t0.7\t&\\\\ G\\,316.640--0.087\t \t&\tG\\,316.640--0.087  \t&\t0.7\t&\tG\\,316.640--0.087 \t&\t0.9   \t&\t--15\t&\t--22\t&\t--19.8\t&\t--\\\\ G\\,316.763--0.011\t \t&\tG\\,316.763--0.012 \t&\t1.0\t&\t--\t\t\t\t&            \t&\t--48\t&\t--40\t&\t\t&--\\\\ G\\,316.812--0.057\t \t&\tG\\,316.811--0.057\t&\t2.2\t&\tG\\,316.811--0.057\t&\t2.5   \t&\t--46\t&\t--43.5\t&\t--46.3\t&--\\\\ G\\,317.429--0.561\t \t&\tG\\,317.429--0.561 \t&\t2.1\t&\t--\t\t\t\t&            \t&\t25\t&\t25.5\t&\t\t&\tG\\,317.430--0.561\t\t&\t2.6\t&\\\\ G\\,318.044--1.404\t \t&\tG\\,318.044--1.405\t&\t2.0\t&\tG\\,318.043--1.404\t&\t1.6   \t&\t42\t&\t45\t&\t46.2\t&\t--\\\\ G\\,318.050+0.087\t \t&\tG\\,318.050+0.087\t&\t0.6\t&\tG\\,318.050+0.087\t&\t0.4    &\t--48\t&\t--53\t&\t--46.5\t&\t--\\\\ G\\,318.948--0.196b\t \t&\tG\\,318.948--0.196\t&\t0.8\t&\tG\\,318.948--0.196\t&\t0.9    &\t--38\t&\t--35.5\t&\t--34.7\t&\t--\\\\ G\\,319.399--0.012\t \t&\tG\\,319.398--0.012\t&\t1.1\t&\t--\t\t\t\t&            \t&\t--5\t&\t--1\t&\t\t&\tG\\,319.399--0.012\t\t&\t0.8\t&\\\\ G\\,319.836--0.196\t \t&\tG\\,319.836--0.196 \t&\t1.5\t&\tG\\,319.836--0.197\t&\t1.6    &\t--11\t&\t--10.5\t&\t--9.1\t&\t--\\\\ G\\,320.120--0.440\t \t&\tG\\,320.120--0.440 \t&\t0.5\t&\t--\t\t\t\t&            \t&\t--46\t&\t--55.5\t&\t\t&\t--\\\\ G\\,320.232--0.284\t\t&\tG\\,320.232--0.284 \t&\t0.4\t&\tG\\,320.231--0.284\t&\t0.6  \t&\t--67\t&\t--64\t&\t--66.5\t&\t--\\\\ G\\,320.233--0.284\t \t&\t--\t\t\t\t&\t\t&\t--\t\t\t\t&      \t\t&\t\t&\t\t&\t\t&\tG\\,320.234--0.283\t\t&\t3.5\t&\\\\ G\\,321.033--0.483\t \t&\t--\t\t\t\t&\t\t& \tG\\,321.033--0.483\t&\t0.5    &\t--61\t&\t\t&\t--61.6\t&\t--\\\\ G\\,321.148--0.529\t \t&\tG\\,321.148--0.529\t&\t1.1\t&\tG\\,321.148--0.529\t&\t1.1\t&\t--97\t&\t--63\t&\t--66.1\t&\t--\\\\ G\\,322.158+0.636\t \t& \tG\\,322.158+0.636\t&\t1.2\t&\tG\\,322.158+0.636 \t&\t1.2\t&\t--76\t&\t--61\t&\t--63.3\t&\t\t--\\\\ G\\,323.740--0.263\t \t&\tG\\,323.740--0.263\t&\t1.1\t&\tG\\,323.740--0.263   \t&\t0.6\t&\t--50\t&\t--39\t&\t--51.1\t&\t--\t&\\\\ \\hline  \\end{tabular} \\label{tab:ass} \\end{table*}  \\begin{table*}\\addtocounter{table}{-1} \\caption{-- {\\emph {continued}}} \\begin{tabular}{llclcllrlcl} \\hline \\multicolumn{1}{c}{\\bf Water} & {\\bf OH} & {\\bf Sep.} & {\\bf Methanol}  & {\\bf Sep.} & {\\bf Water}& {\\bf OH} & {\\bf Methanol} & {\\bf Continuum} & {\\bf Sep.} \\\\ \\multicolumn{1}{c}{\\bf ($l,b$)}  &{\\bf ($l,b$)} & &{\\bf ($l,b$)} & & {\\bf Vpeak} &  {\\bf Vpeak} &  {\\bf Vpeak}&{\\bf ($l,b$)} \\\\ \\multicolumn{1}{c}{\\bf (degrees)} & {\\bf (degrees)} & {\\bf (arcsec)}&{\\bf (degrees)} &{\\bf (arcsec)}& {\\bf (\\kms)} & {\\bf (\\kms)} &{\\bf (\\kms)} & {\\bf (degrees)} &{\\bf (arcsec)} \\\\ \\hline \\\\ G\\,324.201+0.122\t \t&\tG\\,324.200+0.121\t&\t2.9\t&\t--\t\t   \t\t&            \t&\t--87\t&\t--91.5\t&\t\t&\t--\t&\t\\\\ G\\,324.716+0.342\t \t&\tG\\,324.716+0.342\t&\t1.7\t&\tG\\,324.716+0.342 \t&\t1.5   \t&\t--58\t&\t--50\t&\t--46\t&\t--\t&\\\\ G\\,326.662+0.521\t \t&\t--\t\t\t\t&\t\t&\tG\\,326.662+0.521\t&      2.0    &\t --39\t&\t\t&\t--38.6\t& \t--\\\\ G\\,326.670+0.554\t \t&\tG\\,326.670+0.554\t&\t2.6\t&\t--\t\t\t\t&           \t&\t--40\t&\t--40.8\t&\t\t&\t--\\\\ G\\,326.780--0.241\t\t&\tG\\,326.780--0.241 \t&\t0.9\t&\t--\t\t\t\t&\t     \t&      \t--66\t&\t--65\t&\t\t&\t--\\\\ G\\,326.859--0.676\t\t&\t--\t\t\t\t&\t\t& \tG\\,326.859--0.677\t&\t3.4\t&\t--103\t&\t\t&\t--58.0\t&\t--\\\\ G\\,327.119+0.511\t\t&\tG\\,327.120+0.511\t&\t2.3\t&\tG\\,327.120+0.511 \t&\t1.8\t&\t--88\t&\t--80.5\t&\t--87.0\t&\t--\\\\ G\\,327.291--0.578\t \t&\tG\\,327.291--0.578  \t&\t1.2\t&\tG\\,327.291--0.578\t&\t0.8\t&\t--63\t&\t--50.5\t&\t--36.8\t&\t--\\\\ G\\,327.391+0.200\t \t& \t--\t\t\t\t&\t\t&    \tG\\,327.392+0.199\t&\t1.5    &\t--86\t&\t\t&\t--84.6\t&\t--\\\\ G\\,327.402+0.445\t\t&\tG\\,327.402+0.444\t&\t3.2\t&\tG\\,327.402+0.444 \t&\t1.6\t&\t--81\t&\t--77\t&\t--82.6\t&\tG\\,327.402+0.445\t\t&\t0.6\t&\\\\ G\\,327.619--0.111\t \t&\t--\t\t\t\t&\t\t&     G\\,327.618--0.111\t&\t0.9\t&\t--85\t&\t\t&\t--97.6\t&\t--\\\\ G\\,328.236--0.548\t \t&  \tG\\,328.237--0.547 \t&\t2.9\t&     G\\,328.237--0.547 \t&\t2.6\t&\t--38\t&\t--41\t&\t--44.5\t&\tG\\,328.236--0.547\t\t&\t2.1\t&\\\\ G\\,328.254--0.532\t \t&\tG\\,328.254--0.532\t&\t1.5\t&     G\\,328.254--0.532 \t&\t0.8\t&\t--50\t&\t--37\t&\t--37.5\t&\t--\\\\ G\\,328.306+0.432\t \t&\t--\t\t\t\t&\t\t&     --\t\t\t\t&            \t&\t\t&\t\t&\t\t&\tG\\,328.307+0.431\t\t&\t3.3\t&\\\\ G\\,328.808+0.633\t \t&\tG\\,328.809+0.633\t&\t2.9\t&     G\\,328.808+0.633\t&\t2.4\t&\t--46\t&\t--43.5\t&\t--43.8\t&\tG\\,328.808+0.633 \t\t&\t2.0\t&\t\\\\ G\\,329.029--0.199\t\t&\tG\\,329.029--0.200\t&\t1.9\t&     --\t      \t\t\t&           \t&\t--38\t&\t--38.5\t&\t\t&\t--\t\\\\ G\\,329.030--0.205\t \t&\tG\\,329.029--0.205   \t& \t1.5\t&\tG\\,329.029--0.205\t&\t1.3\t&\t--46\t&\t--38.5\t&\t--37.4\t&\t--\\\\ G\\,329.031--0.198\t \t&\tG\\,329.031--0.198   \t&\t1.3\t&\tG\\,329.031--0.198\t&\t0.8\t&\t--52\t&\t--45.5\t&\t--45.5\t&\t--\\\\ G\\,329.066--0.307\t \t&\tG\\,329.066--0.308   \t&\t1.5\t&\tG\\,329.066--0.308\t&\t1.2\t&\t--45\t&\t--43.5\t&\t--43.8\t&\t--\t\\\\ G\\,329.183--0.313\t \t&\tG\\,329.183--0.314   \t&\t2.3\t&\tG\\,329.183--0.314\t&\t1.9\t&\t--50\t&\t--53\t&\t--55.7\t&\t--\\\\ G\\,329.405--0.459\t \t&\tG\\,329.405--0.459\t&\t2.0\t&\tG\\,329.405--0.459\t&\t1.5\t&\t--77\t&\t--69.5\t&\t--70.5\t&--\t\t\\\\ G\\,329.407--0.459\t \t&\t--\t\t\t\t&\t\t&\tG\\,329.407--0.459\t&\t2.1   \t& \t--74\t&\t\t&\t--66.7\t&\t--\\\\ G\\,329.622+0.138\t \t&\t--\t\t\t\t&\t\t&\tG\\,329.622+0.138\t&\t1.7    &\t--82\t&\t\t&\t--84.8\t&\t--\\\\ G\\,330.070+1.064\t \t&\t--\t\t\t\t&\t\t&\tG\\,330.070+1.064\t&\t1.2    &\t--50\t&\t\t&\t--38.8\t&\t--\\\\ G\\,330.879--0.367\t\t&\tG\\,330.878--0.367a\t&\t0.9\t&\tG\\,330.878--0.367\t&\t2.6\t&\t--60\t&\t--61.8\t&\t--59.3\t&\tG\\,330.879--0.367\t\t&\t1.7\t&\\\\ &\tG\\,330.878--0.367b\t&\t1.2\t&\t--\t\t\t\t&           \t&\t\t&\t--65.6\t&\t\t&\t\\\\ G\\,330.954--0.182\t \t&\tG\\,330.954--0.182\t&\t1.3\t&\tG\\,330.953--0.182\t&\t3.9\t&\t--91\t&\t--85.5\t&\t--87.6\t&\tG\\,330.954--0.182\t\t&\t1.6\t&\t\\\\ G\\,331.132--0.244\t \t&\tG\\,331.132--0.244\t&\t0.2\t&\tG\\,331.132--0.244 \t&\t0.3\t&\t--99\t&\t--88.5\t&\t--84.3\t&\t--\\\\ G\\,331.278--0.188\t \t&\tG\\,331.278--0.188   \t&\t0.9\t&\tG\\,331.278--0.188\t&\t1.1\t&\t--90\t&\t--89.5\t&\t--78.2\t&\t--\t\\\\ G\\,331.342--0.346\t \t&\tG\\,331.342--0.346   \t&\t1.4\t&\tG\\,331.342--0.346\t&\t1.6\t&\t--62\t&\t--67\t&\t--67.4\t&\t--\\\\ G\\,331.442--0.187\t \t&\tG\\,331.442--0.186\t&\t0.5\t&\tG\\,331.442--0.187\t&\t0.9\t&\t--88\t&\t--83\t&\t--88.4\t&\tG\\,331.443--0.187\t\t&\t3.5\t&\\\\ G\\,331.512--0.103\t \t&\tG\\,331.512--0.103\t&\t1.4\t&\t--\t\t\t\t&\t      \t&\t--90\t&\t--88.2\t&\t\t&\tG\\,331.512--0.103\t\t&\t1.0\t&\\\\ G\\,331.555--0.122\t \t&\tG\\,331.556--0.121 \t&\t4.5\t&\tG\\,331.556--0.121\t&\t4.4\t&\t--99\t&\t--100\t& --103.4\t&\t--\\\\ G\\,332.094--0.421\t \t&\t--\t\t\t\t&\t\t&\tG\\,332.094--0.421\t&\t2.2    &\t--59\t&\t\t&\t--58.6\t&\t--\\\\ G\\,332.296--0.094\t \t&\t--\t\t\t\t&\t\t&\tG\\,332.295--0.094\t&\t4.4    &\t--50\t&\t\t&\t--47.0\t&\t--\\\\ G\\,332.352--0.117\t \t&\tG\\,332.352--0.117\t&\t0.8\t&\tG\\,332.352--0.117\t&\t0.2\t&\t--60\t&\t--44\t&\t--41.8\t&\t--\\\\ G\\,332.604--0.167\t \t&\t--\t\t\t\t&\t\t&\tG\\,332.604--0.167\t&\t1.6  \t&\t--46\t&\t\t&\t--50.9\t&\t--\\\\ G\\,332.725--0.621\t \t&\tG\\,332.726--0.621\t&\t1.4\t&\tG\\,332.726--0.621\t&\t1.0\t&\t--58\t&\t--48\t&\t--49.6\t&\t--\\\\ G\\,332.826--0.549\t \t&\t--\t\t\t\t&\t\t&\tG\\,332.826--0.549\t&\t3.0    &\t--59\t&\t\t&\t--61.7\t&\tG\\,332.826--0.549\t\t&\t1.1\t&\\\\ G\\,332.964--0.679\t \t&\t--\t\t\t\t&\t\t&\tG\\,332.963--0.679\t&\t1.6    &\t--52\t&\t\t&\t--45.8\t&\t--\\\\ G\\,333.030--0.063\t \t&\t--\t\t\t\t&\t\t&\tG\\,333.029--0.063\t&\t1.3    &\t--40\t&\t\t&\t--55.2\t&\tG\\,333.030--0.063\t\t&\t0.7\t&\\\\ G\\,333.121--0.434\t \t&\t--\t\t\t\t&\t\t&   \tG\\,333.121--0.434\t&\t1.2    &\t--47\t&\t\t&\t--49.3\t&\t--\\\\ G\\,333.126--0.440\t \t&\t--\t\t\t\t&\t\t&\tG\\,333.126--0.440\t&\t1.0    &\t--52\t&\t\t&\t--43.9\t&\t--\\\\ G\\,333.128--0.440\t \t&\t--\t\t\t\t&\t\t&  \tG\\,333.128--0.440\t&\t2.5    &\t--124\t&\t\t&\t--44.6\t&\t--\\\\ G\\,333.234--0.060\t \t&\tG\\,333.234--0.060\t&\t0.6\t&\t--\t\t\t\t&\t      \t&\t--88\t&\t--84\t&\t\t&\t--\\\\ G\\,333.315+0.106\t \t&    \tG\\,333.315+0.105 \t&\t2.8\t&\tG\\,333.315+0.105\t&\t3.0\t&\t--48\t&\t--47\t&\t--45\t&\t--\\\\ G\\,333.387+0.032\t \t&    \tG\\,333.387+0.032 \t&\t0.6\t&\tG\\,333.387+0.032\t&\t1.1\t&\t--61\t&\t--74\t&\t--73.9\t&\t--\\\\ G\\,333.467--0.164\t \t&    \tG\\,333.466--0.164\t&\t2.1\t&\tG\\,333.466--0.164\t&\t2.5\t&\t--42\t&\t--43.5\t&\t--42.5\t&\tG\\,333.466--0.163\t\t&\t4.7\t&\\\\ G\\,333.608--0.215\t \t&    \tG\\,333.608--0.215  \t&\t0.2\t&\t--\t\t\t\t&\t      \t&\t--49\t&\t--51\t&\t\t&--\\\\ G\\,333.646+0.058\t \t&\t--\t\t\t\t&\t\t&  \tG\\,333.646+0.058 \t&\t0.9\t&\t--89\t&\t\t&\t--87.3\t&\t\t--\\\\ G\\,333.682--0.436\t \t&\t--\t\t\t\t&\t\t&    \tG\\,333.683--0.437\t&\t1.6\t&\t--3\t&\t\t&\t--5.3\t&\t--\\\\ G\\,333.930--0.134\t \t&\t--\t\t\t\t&\t\t&    \tG\\,333.931--0.135\t&\t1.5\t&\t--46\t&\t\t&\t--36.7\t&\t--\\\\ G\\,334.635--0.015\t \t&\t--\t\t\t\t&\t\t&    \tG\\,334.635--0.015\t&\t1.0\t&\t--26\t&\t\t&\t--30\t&\t--\\\\ G\\,334.935--0.098\t \t&\t--\t\t\t\t&\t\t&    \tG\\,334.935--0.098\t&\t0.4\t&\t--17\t&\t\t&\t--19.5\t&\t--\\\\ G\\,335.060--0.428\t \t&\tG\\,335.060--0.427 \t&\t1.5\t&\tG\\,335.060--0.427\t&\t1.3\t&\t--37\t&\t--36\t&\t--47.0\t&\t--\\\\ G\\,335.585--0.285\t \t&    \tG\\,335.585--0.285  \t&\t0.5\t&\tG\\,335.585--0.285\t&\t0.7\t&\t--42\t&\t--48\t&\t--49.3\t&\t--\\\\ G\\,335.586--0.290\t \t&\tG\\,335.585--0.289\t&\t1.1\t&\tG\\,335.585--0.289\t&\t0.8\t&\t--56\t&\t--53.5\t&\t--51.4\t&\t--\\\\ &\t\t\t\t\t&\t\t&\tG\\,335.585--0.290\t&      2.3    &\t\t&\t\t&\t--47.3\t&\\\\ G\\,335.727+0.191\t \t&\t--\t\t\t\t&\t\t& \tG\\,335.726+0.191\t&\t2.0    &\t--51\t&\t\t&\t--44.4\t&--\t\\\\ G\\,335.789+0.174\t \t&\tG\\,335.789+0.174\t&\t0.5\t&\tG\\,335.789+0.174\t&\t0.9    &\t--46\t&\t--51.5\t&\t--47.6\t&--\t\\\\ \\hline    \\end{tabular} \\end{table*}  \\begin{table*}\\addtocounter{table}{-1} \\caption{-- {\\emph {continued}}} \\begin{tabular}{llclcllrlcl} \\hline \\multicolumn{1}{c}{\\bf Water} & {\\bf OH} & {\\bf Sep.} & {\\bf Methanol}  & {\\bf Sep.} & {\\bf Water}& {\\bf OH} & {\\bf Methanol} & {\\bf Continuum} & {\\bf Sep.} \\\\ \\multicolumn{1}{c}{\\bf ($l,b$)}  &{\\bf ($l,b$)} & &{\\bf ($l,b$)} & & {\\bf Vpeak} &  {\\bf Vpeak} &  {\\bf Vpeak}&{\\bf ($l,b$)} \\\\ \\multicolumn{1}{c}{\\bf (degrees)} & {\\bf (degrees)} & {\\bf (arcsec)}&{\\bf (degrees)} &{\\bf (arcsec)}& {\\bf (\\kms)} & {\\bf (\\kms)} &{\\bf (\\kms)} & {\\bf (degrees)} &{\\bf (arcsec)} \\\\ \\hline \\\\ G\\,336.018--0.827\t \t&    \tG\\,336.018--0.827\t&\t0.6\t&\tG\\,336.018--0.827\t&\t0.9\t&\t--54\t&\t--41.5\t&\t--53.4\t&\tG\\,336.018--0.828\t\t&\t0.5\t&\\\\ G\\,336.359--0.137\t \t&    \tG\\,336.358--0.137\t&\t3.0\t&\tG\\,336.358--0.137 \t&\t3.1\t&\t--67\t&\t--82\t&\t--73.6\t&\tG\\.336.360--0.137\t\t&\t3.8\t&\\\\ G\\,336.433--0.262\t \t&\t--\t\t\t\t&\t\t&    \tG\\,336.433--0.262\t& \t1.9    &\t--89\t&\t\t&\t--93.3\t&\t--\\\\ G\\,336.830--0.375\t \t&\t--\t\t\t\t&\t\t&    \tG\\,336.830--0.375\t&\t1.6    &\t--20\t&\t\t&\t--22.7\t&\t--\\\\ G\\,336.864+0.005\t \t&\tG\\,336.864+0.005\t&\t1.2\t&\tG\\,336.864+0.005\t&\t0.7\t&\t--66\t&\t--89\t&\t--76.1\t&\t--\\\\ G\\,336.983--0.183\t \t&\tG\\,336.984--0.183\t&\t4.2\t&\tG\\,336.983--0.183\t&\t3.0\t&\t45\t&\t--80.5&\t--80.8\t&\tG\\,336.984--0.184      \t&\t2.5\t&\\\\ G\\,336.991--0.024\t \t&\t--\t\t\t\t&\t\t&\t--\t\t\t \t&            \t&\t\t&\t\t&\t\t&\tG\\,336.990--0.025\t\t&\t2.3\t&\\\\ G\\,336.994--0.027\t \t&\tG\\,336.994--0.027\t&\t1.0\t&\tG\\,336.994--0.027  \t&\t0.6\t&\t--120\t&\t--123\t& --125.8\t& --\t\\\\ G\\,337.258--0.101\t \t&\tG\\,337.258--0.101\t&\t0.7\t&\tG\\,337.258--0.101\t&\t1.2\t&\t--69\t&\t--70\t&\t--69.3\t&\t--\\\\ G\\,337.404--0.402\t \t&\tG\\,337.405--0.402\t&      1.5\t&\tG\\,337.404--0.402\t&\t0.7\t&\t--40\t&\t--38\t&\t--39.7\t&\tG\\,337.404--0.403\t\t&\t2.0\t&\\\\ G\\,337.612--0.060\t \t&\tG\\,337.613--0.060\t&      0.6\t&\tG\\,337.613--0.060\t&\t0.9    & \t--51\t&\t--42\t&\t--42\t&\t--\\\\ G\\,337.687+0.137\t\t&\t--\t\t\t\t&\t\t&\tG\\,337.686+0.137\t&\t2.3\t&\t--74\t&\t\t&\t--74.9\t&\t--\\\\ G\\,337.705--0.053\t \t&\tG\\,337.705--0.053\t&      0.6\t&\tG\\,337.705--0.053\t&\t0.9\t&\t--49\t&\t--49\t&\t--54.6\t&\tG\\,337.706--0.054\t& 1.1\\\\ G\\,337.916--0.477\t \t&\tG\\,337.916--0.477\t&      0.6\t&\t--\t\t\t\t&\t      \t& \t--33\t&\t--51\t&\t\t&\t--\\\\ G\\,337.920--0.456\t \t&\tG\\,337.920--0.456\t&      0.7\t&\tG\\,337.920--0.456\t&\t0.9\t&\t--40\t&\t--39.5\t&\t--38.8\t&\t--\\\\ G\\,337.998+0.137\t \t&\tG\\,337.997+0.136\t&      1.3\t&\tG\\,337.997+0.136\t&\t1.1\t&\t--38\t&\t--35.5\t&\t--32.3\t&\t--\\\\ G\\,338.075+0.012\t \t&\tG\\,338.075+0.012\t&      2.1\t&\tG\\,338.075+0.012\t&\t2.3\t&\t--50\t&\t--47\t&\t--53.0&\tG\\,338.075+0.012\t\t&\t2.9\t&\\\\ G\\,338.075+0.010\t \t&\t--\t\t\t\t&\t\t&\tG\\,338.075+0.009\t&\t1.3   \t&\t--48\t&\t\t&\t--38.2\t&\t--\\\\ G\\,338.281+0.542\t \t&\tG\\,338.280+0.542\t&      1.2\t&\tG\\,338.280+0.542\t&\t1.1\t&\t--64\t&\t--61\t&\t--56.8\t&\t--\\\\ G\\,338.433+0.057\t \t&\t--\t\t\t\t&\t\t&\tG\\,338.432+0.058\t&    \t4.4    &\t--29\t&\t\t&\t--30.2\t&\t--\t\\\\ G\\,338.461--0.245\t \t&\tG\\,338.461--0.245\t&      0.4\t&\tG\\,338.461--0.245   \t&\t0.8\t&\t--52\t&\t--56\t&\t--50.4\t&\t--\\\\ G\\,338.472+0.289\t \t&\tG\\,338.472+0.289\t&      1.2\t&\tG\\,338.472+0.289\t&\t1.1\t&\t--29\t&\t--32\t&\t--30.5\t&\t--\\\\ G\\,338.562+0.217\t \t&\t--\t\t\t       \t&\t\t&\tG\\,338.561+0.218\t&\t2.0    &\t--39\t&\t\t&\t--40.8\t&\t--\\\\ G\\,338.567+0.110\t \t&\t--\t\t\t\t&\t\t&\tG\\,338.566+0.110\t&\t1.4    &\t--76\t&\t\t&\t--75\t&\t--\\\\ G\\,338.682--0.084\t \t&\t G\\,338.681--0.084\t&      1.4\t&      --\t\t\t\t&\t      \t&\t--16\t&\t--22\t&\t\t&\tG\\,338.681--0.085\t\t&\t1.5\t& \t\\\\ G\\,338.920+0.550\t \t&\t --\t\t\t       \t&\t\t&\tG\\,338.920+0.550\t&\t0.8  \t&\t--68\t&\t\t&\t--61.4\t&\t--\\\\ G\\,338.925+0.556\t \t&\t G\\,338.925+0.557\t&      0.9\t&\tG\\,338.925+0.557\t&\t1.3\t&\t--62\t&\t--61\t&\t--62.3\t&\t--\\\\ G\\,339.582--0.127\t \t&\t --\t\t\t\t&\t\t&\tG\\,339.582--0.127\t&\t0.6    &\t--28\t&\t\t&\t--31.3\t&\t--\\\\ G\\,339.622--0.121\t \t&\t G\\,339.622--0.121\t&      0.8\t&      G\\,339.622--0.121\t&\t1.3\t&\t--33\t&\t--37.3\t&\t--32.8\t&\t--\\\\ G\\,339.762+0.055\t \t&\t --\t\t\t       \t&\t\t&\tG\\,339.762+0.054\t&      1.6    &\t--57\t&\t\t&\t--51\t&\t--\\\\ G\\,339.884--1.259\t \t&\tG\\,339.884--1.259b\t&\t0.7\t&\tG\\,339.884--1.259 \t&\t0.8   \t&\t--51\t&\t--36\t&\t--38.7\t&\t--\\\\ &\tG\\,339.884--1.259a\t&\t1.2\t&\t--\t\t\t        \t&            \t&\t\t&\t--29\t&\t\t&\t\\\\ G\\,340.054--0.243\t \t&\tG\\,340.054--0.244\t&      1.4\t&\tG\\,340.054--0.244\t&\t0.8\t&\t--54\t&\t--53.6\t&\t--59.7\t&\t--\\\\ G\\,340.785--0.096\t \t&\tG\\,340.785--0.096   \t&      2.2   \t&\tG\\,340.785--0.096\t&\t1.6\t&\t--120\t&\t--102\t&  --105.1\t&\t--\\\\ G\\,341.218--0.212\t \t&\tG\\,341.218--0.212 \t&\t1.2\t&\tG\\,341.218--0.212\t&\t1.0\t&\t--39\t&\t--37.3\t&\t--37.9\t&\t--\\\\ G\\,341.276+0.062\t \t&\tG\\,341.276+0.062\t&    \t0.9\t&\tG\\,341.276+0.062\t&\t1.1  \t&\t--64\t&\t--73\t&\t--73.8\t&\t--\\\\ G\\,342.484+0.183\t \t&\t--\t\t\t\t&\t\t&\tG\\,342.484+0.183\t&      1.1    \t&\t--43\t&\t\t&\t--41.8\t&\t--\\\\ G\\,343.127--0.063\t \t&\tG\\,343.127--0.063\t&\t2.1\t&    \t--\t\t\t\t&            \t&\t--30\t&\t--31.5\t&\t\t&\t--\\\\ G\\,344.228--0.569\t \t&\tG\\,344.227--0.569\t&     1.4\t&\tG\\,344.227--0.569\t&\t1.0\t&\t--25\t&\t--30.5\t&\t--19.8\t&\t--\\\\ G\\,344.421+0.046\t \t&\t--\t\t\t\t&\t\t&\tG\\,344.421+0.045\t&\t3.4\t&\t--26\t&\t\t&\t--71.5\t&\t--\\\\ G\\,344.582--0.024\t \t&\tG\\,344.582--0.024\t&     2.2\t&\tG\\,344.581--0.024\t&\t2.5\t&\t--4\t&\t--2.3\t&\t1.4\t&\tG\\,344.582--0.024\t\t&\t1.7\t&\\\\ G\\,345.004--0.224\t \t&\tG\\,345.003--0.224 \t&     3.0\t&\tG\\,345.003--0.224   \t&      3.1    &\t\t&\t\t&\t--26.2\t&\tG\\,345.004--0.225\t\t&\t3.6\t&\\\\ &\t\t\t      \t\t&\t\t&\tG\\,345.003--0.223   \t&      3.3\t&\t15\t&\t--27\t&\t--22.5\t&\t\\\\ G\\,345.010+1.793\t \t&\tG\\,345.010+1.793\t&     1.5\t&\tG\\,345.010+1.792\t&      2.1 \t&\t--17\t&\t--22.5\t&\t--18\t&\tG\\,345.010+1.792\t\t&\t4.3\t&\\\\ G\\,345.012+1.797\t \t&\t--\t\t\t\t&\t\t&\tG\\,345.012+1.797\t&\t2.2\t&\t--12\t&\t\t&\t--12.7\t&\t--\\\\ G\\,345.408--0.953\t \t&\tG\\,345.407--0.952\t&     4.6\t&\tG\\,345.407--0.952\t&\t4.9\t&\t--15\t&\t--17.6\t&\t--14.4\t&\tG\\,345.408--0.952\t\t&\t3.2\t&\\\\ G\\,345.425--0.951\t \t&\t--\t\t\t\t&\t\t&\tG\\,345.424--0.951  \t&\t1.6\t&\t--13\t&\t\t&\t--13.5\t&\t--\t\\\\ G\\,345.438--0.074\t \t&\tG\\,345.437--0.074\t&     1.9\t&\t--\t\t\t\t&\t\t&\t--12\t&\t--24.3\t&\t\t&\t--\\\\ G\\,345.487+0.314\t \t&\t--\t\t\t\t&\t\t&\tG\\,345.487+0.314\t&\t0.6    &\t--13\t&\t\t&\t--22.6\t&\t--\\\\ G\\,345.493+1.469\t \t&\tG\\,345.494+1.469\t&     3.3\t&\t--\t\t\t\t&\t      \t&\t5\t&\t--12.7\t&\t\t&\t--\\\\ G\\,345.505+0.348\t \t&\tG\\,345.504+0.348\t&     1.8\t&\tG\\,345.505+0.348\t&\t2.2\t&\t--4\t&\t--19.5\t&\t--17.7\t&\t--\\\\ G\\,345.699--0.090\t \t&\tG\\,345.698--0.090\t&     1.2\t&\t--\t\t\t\t&\t      \t&\t--5\t&\t--6\t&\t\t&\t--\\\\ G\\,346.480+0.132\t \t&\tG\\,346.481+0.132 \t&     1.5\t&\tG\\,346.481+0.132\t&\t1.4\t&\t--10\t&\t--8\t&\t--5.5\t&\t--\\\\ G\\,346.522+0.085\t \t&\t--\t\t\t\t&\t\t&\tG\\,346.522+0.085\t&\t0.7    &\t4\t&\t\t&\t5.5\t&\t--\\\\ G\\,347.628+0.149\t \t&\tG\\,347.628+0.148\t&     2.0\t&\tG\\,347.628+0.149 \t&\t0.9\t&\t--125\t&\t--94.3\t&\t--96.6\t&\t--\\\\ G\\,347.632+0.210\t \t&\t--\t\t\t\t&\t\t&\tG\\,347.631+0.211\t&\t2.9\t&\t--88\t&\t\t&\t--91.9\t&\tG\\,347.632+0.210\t\t&\t0.9\t&\\\\ G\\,348.551--0.979\t \t&\tG\\,348.550--0.979\t&     3.7\t&\tG\\,348.550--0.979n\t&\t1.9\t&\t--18\t& --19.7\t& --20.0\t&\t--\\\\ &\t\t\t      \t\t&\t\t&\tG\\,348.550--0.979\t&      3.4 \t&\t\t&\t\t&\t--10\t&\t\\\\ G\\,348.726--1.038\t \t&\t--\t\t\t\t& \t\t&\tG\\,348.727--1.037 \t&\t4.4\t&\t--10\t&\t\t&\t--7.6\t&--\\\\ G\\,348.885+0.096\t \t&\tG\\,348.884+0.096\t&    1.2\t&\tG\\,348.884+0.096\t&\t1.4\t&\t--80\t& --73.2\t&\t--76.2\t&\t--\\\\  \\hline      \\end{tabular} \\end{table*}  \\begin{table*}\\addtocounter{table}{-1} \\caption{-- {\\emph {continued}}} \\begin{tabular}{llclcllrlcl} \\hline \\multicolumn{1}{c}{\\bf Water} & {\\bf OH} & {\\bf Sep.} & {\\bf Methanol}  & {\\bf Sep.} & {\\bf Water}& {\\bf OH} & {\\bf Methanol} & {\\bf Continuum} & {\\bf Sep.} \\\\ \\multicolumn{1}{c}{\\bf ($l,b$)}  &{\\bf ($l,b$)} & &{\\bf ($l,b$)} & & {\\bf Vpeak} &  {\\bf Vpeak} &  {\\bf Vpeak}&{\\bf ($l,b$)} \\\\ \\multicolumn{1}{c}{\\bf (degrees)} & {\\bf (degrees)} & {\\bf (arcsec)}&{\\bf (degrees)} &{\\bf (arcsec)}& {\\bf (\\kms)} & {\\bf (\\kms)} &{\\bf (\\kms)} & {\\bf (degrees)} &{\\bf (arcsec)} \\\\ \\hline \\\\ G\\,348.892--0.180\t \t&\tG\\,348.892--0.180  \t&     0.6\t&\tG\\,348.892--0.180\t&\t1.0\t&\t7\t&\t9.5\t&\t1.4\t&--\\\\ G\\,349.067--0.018\t \t&\tG\\,349.067--0.017\t&     1.1\t&\tG\\,349.067--0.017\t&\t1.1\t&\t13\t&\t15\t&\t6.9\t&--\\\\ G\\,349.092+0.105\t \t&\tG\\,349.092+0.106\t&     0.8\t&\tG\\,349.092+0.106\t&\t0.9\t&\t--80\t&\t--80\t& --80.4\t&\t--\\\\ &\t\t\t      \t\t&\t\t&\tG\\,349.092+0.105\t&      2.0 \t&\t\t&\t\t& --76.5\t&\\\\ G\\,350.015+0.433\t \t&\tG\\,350.015+0.433\t&     0.1\t&\tG\\,350.015+0.433\t&\t1.0\t&\t--35\t&\t--33\t& --31.7\t&--\t\\\\ G\\,350.113+0.095\t \t&\tG\\,350.113+0.095\t&     1.3\t&\t--\t\t\t\t&\t      \t&\t--64\t&\t--71\t&\t\t&--\t\\\\\tG\\,350.105+0.084\t \t&\t--\t\t\t\t&\t\t&\tG\\,350.104+0.084\t&\t2.0\t&\t--71\t&\t\t& --68.4\t&--\t\\\\ &\t\t\t\t\t&\t\t&\tG\\,350.105+0.083\t&\t3.3  \t&\t\t&\t\t& --74.0\t&\t\\\\ G\\,350.299+0.122\t \t&\t--\t\t\t\t&\t\t&\tG\\,350.299+0.122\t&\t0.7\t&\t--68\t&\t\t& --62.1\t&--\t\\\\ G\\,350.330+0.100\t \t&\tG\\,350.329+0.100\t&     2.5\t&\t--\t\t\t\t&\t\t&\t--62\t&\t--64\t&\t\t&\tG\\,350.331+0.099\t\t&\t3.9\t&\\\\ G\\,350.686--0.491\t \t&\tG\\,350.686--0.491\t&     0.4\t&\tG\\,350.686--0.491\t&\t1.3\t&\t--14\t& --14.5\t& --13.8\t&\t--\\\\ G\\,351.160+0.696\t \t&\tG\\,351.160+0.697\t&     2.5\t&\tG\\,351.160+0.697\t&\t1.3\t&\t--3\t&\t--8.5\t&\t--5.2\t&\tG\\,351.161+0.696\t\t&\t3.0\t&\\\\ G\\,351.243+0.671\t\t&\t--\t\t\t\t&\t\t& \tG\\,351.243+0.671\t&\t3.4\t& \t--77\t&\t\t&\t2.5\t&\\\\ G\\,351.246+0.668\t \t&\t--\t\t\t\t&\t\t&\t--\t\t\t\t&            \t&\t\t&\t\t&\t\t&\tG\\,351.247+0.667\t\t&\t3.6\t&\\\\ G\\,351.417+0.646\t \t&\tG\\,351.417+0.645\t&     3.7\t&\tG\\,351.417+0.646\t&      1.7\t&\t--10\t&\t--9.1\t& --11.2\t&\t--\\\\ &\t\t\t      \t\t&\t\t&\tG\\,351.417+0.645\t&      3.1    &\t       \t&\t\t& --10.4\t&\t\\\\ G\\,351.582--0.353\t \t&\tG\\,351.581--0.353\t&     1.6\t&\tG\\,351.581--0.353n\t&    \t2.1\t&\t--89 \t&  --97.6 \t&  --91.1 \t&\t--\\\\ &\t\t\t      \t\t&\t\t&\tG\\,351.581--0.353\t&\t3.4    &       \t&\t       \t&  --94.4 \t&\t\\\\ G\\,351.775--0.536\t \t&\tG\\,351.775--0.536\t&     1.8\t&\tG\\,351.775--0.536\t&\t1.9\t&\t--2     \t&   --2     \t&  1.3     \t&--\t\\\\ G\\,352.111+0.176\t \t&\t--\t\t\t      \t&\t\t&\tG\\,352.111+0.176\t&      2.8\t&\t --60  \t&         \t&  --54.8 \t&  --   \t\\\\ G\\,352.133--0.944\t \t&\t--\t\t\t\t&\t\t&\tG\\,352.133--0.944\t&    \t2.8\t&\t--11\t&        \t&  --16    \t&--\t\\\\ G\\,352.162+0.199\t \t&\tG\\,352.161+0.200\t&     1.0\t&\t--\t\t\t\t&\t\t&\t--45 \t&  --42.2 \t&         \t&--\t\\\\ G\\,352.517--0.155\t \t&\tG\\,352.517--0.155\t&     0.2\t&\tG\\,352.517--0.155\t&      0.3\t&\t--49 \t&  --50.6 \t& --51.2  \t&--\t\\\\ G\\,352.525--0.158\t\t&\t--\t\t\t\t&\t\t&       G\\,352.525--0.158    &      0.3    &       --51  &                &     --53          & -- \\\\ G\\,352.623--1.076\t \t&\t--\t\t\t      \t&\t\t&\tG\\,352.624--1.077   \t&   \t4.7\t&\t--6     \t&\t   \t&    5.8   \t&--\t\\\\ G\\,352.630--1.067\t \t&\tG\\,352.630--1.067\t&     0.5\t&\tG\\,352.630--1.067\t&      0.4\t&\t0\t&\t0     \t&    --2.8 \t&--\t\\\\ G\\,353.273+0.641\t \t&\t--\t\t\t\t&\t\t&\tG\\,353.273+0.641   \t&     0.3\t&\t--49\t&         \t&   --5.2  \t&--\t\\\\ G\\,353.411--0.362\t \t&\t--\t\t\t\t&\t\t&\t--\t\t\t\t&\t\t&          \t&        \t&         \t&\tG\\,353.411--0.362\t\t&\t1.7\t&\\\\ G\\,353.464+0.562\t \t&\tG\\,353.464+0.562\t&     0.9\t&\tG\\,353.464+0.562\t&\t1.9\t&\t--60  &    --45  \t&   --50.7\t&--\t\\\\ G\\,354.615+0.472\t \t&\tG\\,354.615+0.472\t&     1.9\t&\tG\\,354.615+0.472\t&\t1.8\t&\t--38\t&  --15.4 \t&  --24.6 \t&--\t\\\\ G\\,355.343+0.147\t \t&\tG\\,355.344+0.147\t&     2.0\t&\tG\\,355.344+0.147\t&\t1.7\t&\t17\t&  19      \t&  20      \t&--\t\\\\ &\t\t\t\t\t&\t\t&\tG\\,355.343+0.148\t&\t2.7    &    \t \t&\t     \t&  5.7   \t&\t\\\\ G\\,355.345+0.149\t\t&\t--\t\t\t\t&\t\t&      G\\,355.346+0.149     &      1.4     &    72       &                &      10      & --   \\\\ G\\,357.965--0.164\t \t&\t--\t\t\t      \t&\t\t&\tG\\,357.965--0.164    \t&     1.5\t&\t--19   \t&         \t&   --8.8  \t&\t--\\\\ G\\,357.967--0.163\t \t&\tG\\,357.968--0.163\t&     1.7\t&\tG\\,357.967--0.163\t&     0.5\t&\t--65   \t& --6.3    \t& --3.2    \t&\t--\\\\ G\\,358.371--0.468\t \t&\t--\t\t\t  \t&\t\t&\tG\\,358.371--0.468 \t&\t1.0\t&\t1       \t&        \t&     1     \t&--\t\\\\ G\\,358.386--0.483\t \t&\tG\\,358.387--0.482a\t&    \t2.3\t&\tG\\,358.386--0.483\t&\t1.5\t&     0       \t&\t--6.3 \t& --6.0   \t&\tG\\,358.387--0.483\t\t&\t3.4\t&\\\\ &\tG\\,358.387--0.482b \t&\t3.3\t&\t--\t\t\t \t&\t      \t&              \t&    --7.8 \t&           \t&\t\\\\ G\\,359.137+0.032\t \t&\tG\\,359.137+0.032\t&     1.3\t&\tG\\,359.138+0.031\t&\t1.5\t&\t --1\t&   --1     \t&  --3.9   \t&--\t\\\\ G\\,359.436--0.102\t \t&\t\t--\t\t\t&     \t\t&\tG\\,359.436--0.102\t&\t0.4\t&\t--59\t&    \t\t&  --53.6 \t&--\t\\\\ G\\,359.436--0.104\t \t&\tG\\,359.436--0.103\t&\t1.9\t&\tG\\,359.436--0.104\t&\t0.7\t&\t--47\t&       --52    \t&    --52  \t&--\t\\\\ G\\,359.615--0.243\t\t&\tG\\,359.615--0.243\t&     0.8\t&\tG\\,359.615--0.243\t&\t0.6\t&\t64\t&   22.5  \t&  22.5   \t&--\t\\\\ G\\,359.969--0.457\t \t&\tG\\,359.970--0.457\t&     1.2\t&\tG\\,359.970--0.457\t&\t1.3\t&     11     \t&   15.5  \t&  23.0   \t&--\t\\\\ G\\,0.209--0.002\t  \t&\t\t--\t\t\t&     \t\t&\t--\t\t\t \t&\t\t&           \t&         \t&         \t&\tG\\,0.209--0.002\t\t\t&\t2.4\t&\\\\ G\\,0.212--0.002\t   \t&\t\t--\t\t\t&     \t\t&    \tG\\,0.212--0.001  \t&\t3.7\t&\t56     &         \t&   49.2  \t&--\t\\\\ G\\,0.316--0.201\t   \t&\t\t--\t\t\t&     \t\t&      G\\,0.316--0.201        \t&\t0.7\t&\t23\t&\t\t&   21     \t&--\t\\\\ &\t\t\t\t\t&\t\t&\tG\\,0.315--0.201\t &     2.1    \t&\t\t&          \t&   18     \t&\t\\\\ G\\,0.376+0.040\t        \t\t&\tG\\,0.376+0.040  \t&\t1.3   \t&\tG\\,0.376+0.040\t \t&\t1.6\t&\t40   \t&      36\t&   37.1  \t&--\t\\\\ G\\,0.497+0.188\t        \t\t&\tG\\,0.496+0.188   \t&   \t2.1   \t&\tG\\,0.496+0.188\t \t&\t2.1\t&\t26\t&     --5.5 \t&   0.8    \t&--\t\\\\ G\\,0.547--0.851\t        \t&\tG\\,0.546--0.852   \t&  \t2.7   \t&\tG\\,0.546--0.852 \t &\t3.2\t&\t20\t&\t13.5\t&    13.8 \t&--\t\\\\ G\\,0.657--0.042\t\t&\tG\\,0.658--0.042\t&\t0.4\t&       G\\,0.657--0.041         &      4.8    &       62    &       52         &            & -- \\\\ G\\,0.668--0.035\t        \t&\tG\\,0.666--0.035   \t&  \t4.0   \t&\t--\t\t \t\t&\t\t&\t59\t&\t61\t&         \t&--\t\\\\ G\\,2.143+0.009\t     \t\t&    \tG\\,2.143+0.009     \t&\t0.8   \t&\tG\\,2.143+0.009    \t&\t0.1\t&\t37\t&\t59.8\t&    62.7 \t&   -- \t\\\\ G\\,2.536+0.198\t      \t\t&\t--\t\t\t\t&\t\t&\tG\\,2.536+0.198\t \t&\t2.5\t&     25\t&          \t&    3.2   \t&--\t\\\\ G\\,5.886--0.392\t\t& G\\,5.885--0.392\t\t&\t6.3\t&\tG\\,5.885--0.392\t&\t5.5\t&11\t\t&\t13.9\t&\t6.7\t& --\\\\ G\\,5.901--0.430\t      \t&\t--\t\t\t\t&\t\t& \tG\\,5.900--0.430 \t &\t1.7\t&\t14\t&\t\t&    10    \t& -- \t\\\\ G\\,6.049--1.447\t \t& \tG\\,6.048--1.447 \t&\t2.0\t&\t--\t\t\t\t&\t\t&\t20\t&     11.2 \t&         \t&\t--\\\\ G\\,6.611--0.082\t    \t&\t--\t\t\t\t&\t\t&      G\\,6.610--0.082 \t &\t2.1\t&\t6\t&\t\t&   0.7     \t&--\t\\\\ G\\,6.796--0.257\t    \t&     G\\,6.795--0.257 \t\t&\t2.1\t&\tG\\,6.795--0.257  \t&\t2.4\t&\t1\t&\t16.1\t&   26.6   \t&--\t\\\\ G\\,8.139+0.226\t        \t\t&     --\t\t\t\t&\t\t& \tG\\,8.139+0.226 \t&\t1.2\t&\t16\t&\t    \t&   20.0   \t&--\t\\\\  \\hline \\end{tabular} \\end{table*}  \\begin{table*}\\addtocounter{table}{-1} \\caption{-- {\\emph {continued}}} \\begin{tabular}{llclcllrlcl} \\hline \\multicolumn{1}{c}{\\bf Water} & {\\bf OH} & {\\bf Sep.} & {\\bf Methanol}  & {\\bf Sep.} & {\\bf Water}& {\\bf OH} & {\\bf Methanol} & {\\bf Continuum} & {\\bf Sep.} \\\\ \\multicolumn{1}{c}{\\bf ($l,b$)}  &{\\bf ($l,b$)} & &{\\bf ($l,b$)} & & {\\bf Vpeak} &  {\\bf Vpeak} &  {\\bf Vpeak}&{\\bf ($l,b$)} \\\\ \\multicolumn{1}{c}{\\bf (degrees)} & {\\bf (degrees)} & {\\bf (arcsec)}&{\\bf (degrees)} &{\\bf (arcsec)}& {\\bf (\\kms)} & {\\bf (\\kms)} &{\\bf (\\kms)} & {\\bf (degrees)} &{\\bf (arcsec)} \\\\ \\hline \\\\ G\\,8.670--0.356\t        \t&     G\\,8.669--0.356 \t\t&     2.9\t&\tG\\,8.669--0.356\t &\t2.7\t&\t36\t&    39.2  \t&  39.3    \t&\tG\\,8.670--0.356\t\t\t&\t1.1\t&\\\\ G\\,9.620+0.194\t        \t\t&     G\\,9.620+0.194 \t\t&\t1.6\t&\t--\t\t\t\t&\t\t&\t 6\t&\t22\t&    \t\t&--\t\\\\ G\\,9.622+0.195\t        \t\t&     G\\,9.621+0.196  \t&\t2.9\t&\tG\\,9.621+0.196\t \t&\t3.3\t&\t22\t&\t 1.4\t&    1.3\t&--\t\\\\ G\\,9.986--0.028\t        \t&     --\t\t\t\t&\t\t&\tG\\,9.986--0.028 \t &\t1.3\t&\t49     &     \t\t&    47.1   \t&--\t\\\\ G\\,10.288--0.125\t       \t&     --\t\t\t\t&\t\t&  \tG\\,10.287--0.125\t&\t1.9\t&\t9\t&\t     \t&    5        \t&--\t\\\\ G\\,10.323--0.160\t        \t&     --\t\t\t\t&\t\t& \tG\\,10.323--0.160\t&\t1.5\t&\t--3\t&          \t&   10      \t&--\t\\\\ G\\,10.342--0.143\t        \t&     --\t\t\t\t&\t\t& \tG\\,10.342--0.142\t&\t1.8\t&\t8\t&\t\t&   14.8   \t&--\t\\\\ G\\,10.445--0.018\t \t&     G\\,10.444--0.018 \t&\t3.3\t&\tG\\,10.444--0.018\t&\t3.7\t&\t70\t&    75.5 \t&    73.2   \t&--\t\\\\ G\\,10.473+0.027\t        \t&     G\\,10.473+0.027   \t&\t0.9\t&\tG\\,10.473+0.027\t&\t1.9\t&\t62    \t&    51.5 \t&    75      \t&\tG\\,10.473+0.027\t\t&\t1.8\t&\\\\ G\\,10.480+0.034\t        \t&     G\\,10.480+0.033 \t&\t4.5\t&\tG\\,10.480+0.033\t&\t4.5\t&\t64\t&    66    \t&     65     \t&--\t\\\\ G\\,10.623--0.383\t        \t&     G\\,10.623--0.383\t&\t0.6\t&\t--\t\t \t\t&\t\t&     2       \t&      --2    \t&         \t&--\t\\\\ G\\,10.959+0.022\t        \t&     --\t\t\t\t&\t\t& \tG\\,10.958+0.022\t&\t1.4\t&\t25\t&\t\t&     24.4\t&       G\\,10.959+0.022\t\t&\t0.7\t&\\\\ G\\,11.034+0.062\t\t&     G\\,11.034+0.062 \t&\t1.6\t&\tG\\,11.034+0.062\t&\t1.5\t&\t18\t&\t21.7\t&     20.6\t&   --    \\\\ G\\,11.498--1.486\t      \t&     --\t\t\t\t&\t\t&  \tG\\,11.497--1.485\t&\t2.9\t&\t17\t&\t\t&     6.7    \t&--\t\\\\ G\\,11.903--0.142\t        \t&     G\\,11.904--0.141\t&\t3.5\t&\tG\\,11.904--0.141\t&\t4.4\t&\t 36\t&\t40.5\t&     42.8\t&--\t\\\\  G\\,12.203--0.107\t      \t&     --\t\t\t\t&\t\t&   \tG\\,12.203--0.107\t&\t0.2\t&\t35     &          \t&    20.5  \t&--\t\\\\ G\\,12.209--0.102\t \t&     --\t\t\t\t&\t\t&      G\\,12.209--0.102\t&\t1.0\t&     22     \t&          \t&    19.8  \t&\tG\\,12.209--0.102\t\t&\t3.0\t&\\\\ G\\,12.216--0.119\t        \t&     G\\,12.216--0.119 \t&\t1.5\t&\t--     \t\t\t \t&\t\t&\t26     &    27.9  \t&          \t&\t--\\\\ G\\,12.681--0.183\t        \t&    G\\,12.680--0.183    \t&      1.0\t&\tG\\,12.681--0.182\t&\t1.1\t&\t61     &    64.5  \t&    57.6   \t&\t--\\\\ G\\,12.889+0.489\t        \t&    G\\,12.889+0.489  \t&\t3.0   & \tG\\,12.889+0.489\t&\t2.7\t&\t30\t&\t33\t&\t39.3\t&\t--\\\\ G\\,12.908--0.260\t        \t&    G\\,12.908--0.260 \t&     1.2\t&\tG\\,12.909--0.260\t&    1.8\t&\t37\t&    38\t&    39.9   \t& \t--\\\\ G\\,17.637+0.156\t\t&   G\\,17.638+0.157\t\t&\t2.1\t&\tG\\,17.638+0.157\t&\t2.2\t&\t27\t&    20     \t&    20.7    &    \t--\\\\ \\hline  \\end{tabular} \\end{table*}      \\begin{table*} \\caption{22-GHz continuum sources detected towards water maser sources (i.e. continuum sources within 4.5 arcsec of detected water masers). See Tables 1 and 3 for details of the water maser sources that these continuum sources are associated with. Columns one through to 5 show: the name of the continuum source in Galactic coordinates, the source right ascension, declination, peak flux density (mJy/beam) and integrated flux density (mJy). We present two additional continuum sources that fall just outside our association threshold and we distinguish these sources with a `$^{\\#}$' following the source name in column 1.} \\label{tab:cont} \\begin{tabular}{lllcc} \\hline \\multicolumn{1}{c}{\\bf Continuum} & {\\bf RA(2000)}  & {\\bf Dec(2000)} & {\\bf Ipeak} & {\\bf Total Flux} \\\\ \\multicolumn{1}{c}{\\bf ($l,b$)} & & &&{\\bf Density}\\\\ \\multicolumn{1}{c}{\\bf (degrees)}  &{\\bf (h m s)}&{\\bf ($^{o}$ $'$ $``$)}& {\\bf (mJy/beam)} & {\\bf (mJy)} \\\\ \\hline \\\\ G\\,291.611--0.529\t& 11 15 02.62 & --61 15 51.4 &  350\t\t\t& 446\\\\ G\\,291.626--0.531$^{\\#}$\t& 11 15 09.66 & --61 16 15.7 & 137\t\t\t& 357\\\\ G\\,299.012+0.128\t& 12 17 24.12 & --62 29 05.6 & 25\t\t\t& 38\\\\ G\\,301.136--0.226\t& 12 35 34.96 & --63 02 31.6 & 1017\t\t\t& 1116\\\\ G\\,311.643--0.380 \t& 14 06 38.73& --61 58 21.4 & 174\t\t\t& 181  \\\\ G\\,313.458+0.193\t& 14 19 35.04 & --60 51 52.1 & 205\t\t\t& 272      \\\\ G\\,316.412--0.308\t& 14 43 23.25 & --60 12 59.5 & 160\t\t\t& 167\\\\ G\\,317.430--0.561\t& 14 51 38.03 & --60 00 19.5 & 32\t\t\t& 35\\\\ G\\,319.399--0.012\t& 15 03 17.60 & --58 36 11.2 & 230\t\t\t& 298\\\\ G\\,320.234--0.283\t& 15 09 52.63 & --58 25 32.4 & 282\t\t\t& 293\\\\ G\\,327.402+0.445   & 15 49 19.35 & --53 45 13.3 & 86\t\t\t& 94\\\\ G\\,328.236--0.547\t& 15 57 58.15 & --53 59 23.4 & 26\t\t\t& 36\\\\ G\\,328.307+0.431\t& 15 54 06.24 & --53 11 38.8 & 3361\t\t\t& 3647  \\\\ G\\,328.808+0.633\t& 15 55 48.33 & --52 43 07.0 & 1208\t\t  \t& 1357\\\\ G\\,330.879--0.367\t& 16 10 19.95 & --52 06 05.3 & 324\t\t\t& 476\\\\ G\\,330.954--0.182\t& 16 09 52.51 & --51 54 54.1 & 2954\t\t\t& 3151\\\\ G\\,331.443--0.187\t& 16 12 12.83 & --51 35 10.2 & 40\t\t\t& 44  \\\\ G\\,331.512--0.103\t& 16 12 10.10 & --51 28 37.2 & 109\t\t\t& 113    \\\\ G\\,332.826--0.549\t& 16 20 11.12 & --50 53 13.7 & 2510\t\t\t& 2670   \\\\ G\\,333.030--0.063\t& 16 18 56.94 & --50 23 53.5 & 27\t\t\t& 25    \\\\ G\\,333.466--0.163$^{\\#}$\t& 16 21 19.71 & --50 09 45.6 & 141\t\t& 254     \\\\  G\\,336.018--0.828\t& 16 35 09.39 & --48 46 47.6 & 85.1\t\t\t& 80.8    \\\\ G\\,336.360--0.137\t& 16 33 29.64 & --48 03 38.8 &  189\t\t\t& 293\\\\ G\\,336.984--0.184\t& 16 36 12.60 & --47 37 57.8 & 51\t\t\t& 56.8   \\\\ G\\,336.990--0.025\t& 16 35 32.48 & --47 31 14.6 & 91\t\t\t& 102\\\\ G\\,337.404--0.403\t&16 38 50.59  & --47 28 02.8  & 117\t\t\t& 121\\\\ G\\,337.706--0.054\t& 16 38 29.83 & --47 00 35.7 & 244\t\t\t& 262\\\\ G\\,338.075+0.012\t& 16 39 39.15 & --46 41 26.2 & 818\t\t\t& 1144       \\\\ G\\,338.681--0.085\t& 16 42 24.19 & --46 18 00.4 & 74\t\t\t& 74   \\\\ G\\,344.582--0.024\t& 17 02 58.03 & --41 41 52.7 & 19\t\t\t& 21    \\\\ G\\,345.004--0.225\t& 17 05 11.36 & --41 29 06.5 & 353\t\t\t& 355  \\\\ G\\,345.010+1.792\t& 16 56 47.85 & --40 14 25.8 & 362\t\t\t& 367    \\\\ G\\,345.408--0.952\t& 17 09 35.62 & --41 35 54.6 & 493\t\t\t& 608\\\\ G\\,347.632+0.210\t& 17 11 36.22 & --39 07 06.2 & 46\t\t\t& 48\\\\ G\\,350.331+0.099\t& 17 20 02.09 & --36 59 12.8 & 55\t\t\t& 65\\\\ G\\,351.161+0.696\t& 17 19 57.65 & --35 57 51.8 & 238\t\t\t& 271\\\\ G\\,351.247+0.667\t& 17 20 19.29 & --35 54 39.4& 1280\t\t&      1567\\\\ G\\,353.411--0.362\t& 17 30 26.66 & --34 41 46.0 & 384\t\t\t& 788        \\\\ G\\,358.387--0.483\t& 17 43 37.96 & --30 33 49.2 & 109\t\t\t& 113     \\\\ G\\,0.209--0.002\t& 17 46 07.57 & --28 45 30.5 & 75\t\t\t& 95\\\\ G\\,8.670--0.356\t& 18 06 19.15 & --21 37 32.1 & 677\t\t\t& 689\\\\ G\\,10.473+0.027   \t& 18 08 38.41 &  --19 51 47.9 & 152\t\t\t& 151\\\\ G\\,10.959+0.022 \t& 18 09 39.43 &  --19 26 27.0 & 150\t\t\t&153\\\\ G\\,12.209--0.102\t& 18 12 39.85 &  --18 24 20.0 & 119\t\t\t& 140\\\\ \\hline    \\end{tabular} \\end{table*} ",
        "conclusions": "From a large sample of water masers measured with precise positions at two epochs, we conclude that spectra are highly variable but positions are generally persistent.  The occurrence of a water maser at nearly 80 per cent of the OH maser targets is comparable to that of methanol at OH sites.  This is despite the difference in favoured pumping schemes, where both OH and methanol depend on far IR radiation, whereas the favoured pumping scheme for water masers is collisional.  Our study of water masers at methanol maser sites is preliminary, but the common presence of water at methanol sites is confirmed.  We argue that there is indeed an important role for water masers in mapping the Galaxy and its velocity field.  The present contribution of a large number of water masers with accurate positions in the southern Galaxy has been an important step in advancing such a project, and reveals the value of conducting even larger future surveys with complete Galactic plane coverage."
    },
    "1004/1004.1583_arXiv.txt": {
        "abstract": "One of the important goals for future neutrino telescopes is to identify the flavors of astrophysical neutrinos and therefore determine the flavor ratio. The flavor ratio of astrophysical neutrinos observed on the Earth depends on both the initial flavor ratio at the source and flavor transitions taking place during propagations of these neutrinos. We propose a model independent parametrization for describing the above flavor transitions. A few flavor transition models are employed to test our parametrization. The observational test for flavor transition mechanisms through our parametrization is discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.3065_arXiv.txt": {
        "abstract": "We report a mass and rotational broadening ($v\\sin i$) for the pulsating white dwarf component of the WZ~Sge type Dwarf Nova GW~Lib based on high-resolution VLT spectroscopy that resolves the Mg~\\textsc{ii} 4481\\AA~absorption feature. Its gravitational redshift combined with white dwarf mass-radius models, provides us with a direct measurement of the white dwarf mass of $M_1 = 0.84 \\pm 0.02 M_\\odot$. The line is clearly resolved and if associated with rotational broadening gives $v \\sin i = 87.0 \\pm 3.4$\\kms, equivalent to a spin period of $97 \\pm 12$~s. ",
        "introduction": "The population of cataclysmic variables (CVs) provides an important sample of binary systems possessing homogeneous configurations of white dwarfs (WDs) accreting from near main-sequence donor stars.  A CV initially evolves towards shorter orbital periods, but near a period of $\\sim 75$min bounces back towards longer periods with its donor star turning into a degenerate brown dwarf. Recent CV searches, such as the significant harvest of CVs from SDSS \\citep{szkody09-1, gaensicke09-1} are finally unearthing large numbers of systems at short orbital periods as expected from binary evolution considerations.  Reliable estimates for binary parameters are needed to place individual systems on their evolutionary tracks. The faintness of the low mass donor stars in comparison to the WD and its accretion flow makes such parameter estimates difficult for short period systems (e.g WZ Sge; \\citealt{steeghsetal01-2}; \\citealt{steeghsetal07-1}), except in favourable cases where eclipse constraints can be exploited \\citep{littlefairetal06-2}. \\citet{patterson01-1} has identified an indirect method of inferring binary mass ratios using the tidally driven superhump modulations in the lightcurves of CVs. Although promising, this furthermore highlights the need for suitable calibrator systems for which binary parameters can be accurately determined independently.  GW Lib is a short period dwarf nova ($P_{\\mathrm{orb}}$=76.78 mins; \\citealt{thorstensenetal02-3}) displaying the typical characteristics of a CV near its period minimum. Crucially, it was the first CV discovered to contain a pulsating WD \\citep{vanzyletal04-1}, by analogy with field WD pulsators that lie in the instability strip for non-radial pulsations. So far, various methods to determine its system parameters have been exploited, including model fits to the WD absorption profiles, prominent in both the optical \\citep{thorstensenetal02-3} and the UV \\citep{szkodyetal02-4} providing estimates for the WD $T_{\\mathrm{eff}}$ and $\\log g$. Asteroseismological models, combined with a UV-flux limit, suggests a WD mass of $M_1 = 1.02 M_\\odot$ (\\citealt{townsleyetal04-1}). After its second recorded super-outburst in April 2007 \\citep{templeton07-1}, the superhump period suggest a mass ratio of $q=M_2/M_1=0.06$ \\citep{katoetal08-1} when combined with Patterson's empirical relation \\citep{pattersonetal05-3}.  In this letter we explore the possibilities of using the gravitational redshift of the Mg~\\textsc{ii} absorption line reported in \\citet{vanspaandonketal09-1} in combination with mass-radius relations to give an independent measurement of the   mass and spin of the WD in GW Lib.  \\section[]{Observations and reduction} \\label{sec:observations}  GW~Lib was observed within the 69.D-0591 program at the 8.4-m Very Large Telescope (\\textsc{vlt}), located at the Paranal Observatory in Chile as part of the European Southern Observatory, equipped with only the blue arm of the Ultraviolet and Visual Echelle Spectrograph (\\textsc{uves}: \\citealt{dekkeretal00-1}).  We retrieved a total of 45 science frames obtained during 2002 May 16-17 covering 2.4 and 2.8 binary orbits on the respective nights. The reduction of the raw frames was conducted using the most recent standard recipe pipeline release of the \\textsc{uves} Common Pipeline Library (\\textsc{cpl}) recipes. The resultant optimally-extracted spectra covered a wavelength range of $\\lambda$4020-5240\\AA\\ at a dispersion of 0.031\\AA~pixel$^{-1}$ and a spectral resolution of 0.10\\AA~(5.77\\kms) as measured from the skylines. The spectra were wavelength calibrated with one Thorium Argon arc per night. This calibration was tested against the two sky lines visible at $\\lambda$5197.92\\AA\\, and $\\lambda$5200.28\\AA. The corresponding science frames were corrected for any remaining shifts. The residuals were scattered around zero with a maximum amplitude of $0.01$~km~s$^{-1}$. Next, heliocentric velocity corrections were applied to the individual frames to deliver spectra in a common heliocentric rest frame.  No standard was observed in the correct settings on the nights and as no master response curve exists for the non-standard setup used, the frames were not flux calibrated. The exposure time was 500 seconds giving a signal to noise of $\\sim$6.5 per spectrum.  Details of the observations can be found in Table \\ref{tab:observations} while Figure \\ref{fig:spectrum} shows the average spectrum of GW~Lib on the 16th of May. Prominent features are the Balmer disc emission lines on top of broad absorption troughs from the WD, visible due to the low mass accretion rate in the system.  He~\\textsc{i} is also seen in emission as is He~\\textsc{ii} at 4685.75\\AA.  Our pre-outburst, intermediate resolution spectra of GW Lib showed Mg~\\textsc{ii} at 4481.21\\AA\\ in absorption (see figure 2, \\citealt{vanspaandonketal09-1}). The archival \\textsc{vlt/uves} data confirm the presence of this line and thanks to the superior spectral resolution combined with the low inclination of the system shows it to be unblended from the nearby He~\\textsc{i} emission at 4471\\AA. We measure an EW of $0.25\\pm0.01$\\AA\\ (similar to $0.24\\pm 0.03$\\AA\\, in 2004) and a Full Width at Half Maximum (\\textsc{fwhm}) of $1.35\\pm 0.04$\\AA.  \\section[]{Gravitational Redshift for Mg II} \\label{sec:massdetermination} In low accretion rate dwarf novae, the luminosity of the accretion disc is low enough to show the broad absorption features of the WD and can even show the narrow absorption features of metal lines due to freshly accreted gas. These lines open a window to probe the WD atmosphere directly and give independent measurements of stellar parameters. For lines formed near the primary, a gravitational red-shift is expected, introduced in the deep gravitational potential of the WD (\\citealt{eddington24-1}; \\citealt{greenstein+timble67-1}; \\citealt{sionetal94-1}). A measurement of the gravitational redshift in the rest frame of the binary could provide the WD mass directly when combined with mass-radius models (e.g. Eggleton's relation as quoted in \\citealt{verbunt+rappaport88-1}). This method has previously been used in CVs in the cases of U~Gem \\citep{long+gilliland99-1}, VW Hyi \\citep{smithetal06-1} and WZ~Sge \\citep{steeghsetal07-1}.  In the case of GW Lib,  the directly measured redshift of the magnesium line ($v_{\\mathrm{MgII}}$) needs to be rectified for several contributions in order to get the true gravitational redshift induced by the WD only ($v_{\\mathrm{grav}}(\\mathrm{WD})$). These corrections consist of the systemic velocity ($\\gamma$) of the binary system and the effects of the gravitational potential from the donor star ($v_{\\mathrm{grav}}(\\mathrm{donor})$). Hence the gravitational redshift due to the WD is given by:  \\[ v_{\\mathrm{grav}}(\\mathrm{WD}) = v_{\\mathrm{MgII}} - v_{\\mathrm{grav}}(\\mathrm{donor}) -\\gamma \\]  We will discuss the various contributions and measurements independently. Firstly, all 45 spectra of the two nights were individually continuum normalised and then binned into 20 equally spaced orbital phase bins to increase S/N.  \\subsection{$v_{\\mathrm{MgII}}$} To measure the redshift of the magnesium line to the best precision and minimise any orbital effect we compared several methods.  Firstly, we made a weighted orbital average and fitted the Mg~II absorption line with a weighted triple Gaussian as the Mg~\\textsc{ii} is a triplet line. The different components have rest wavelengths of $\\lambda$4481.126\\AA\\, with a  transition probability of $\\log(gf)= 0.7367$,  $\\lambda$4481.150\\AA\\, with $\\log (gf) = -0.5643$  and $\\lambda$4481.325\\AA\\, with $\\log (gf) = 0.5818$\\footnote{Data from The Atomic Line List Version 2.04: \\texttt{http://www.pa.uky.edu/\\~ {}peter/atomic/} }. The variables for the fit are the common offset (for all 3 lines) and the common \\textsc{fwhm} to provide a good fit to the blue wing of the absorption feature. The peak height is also a common variable but scaled according to the various transition probabilities. This gives a best fit with a mean offset of $35.2 \\pm 1.1$\\kms, a \\textsc{fwhm} of $1.32\\pm 0.08$\\AA\\, and a common absorption line depth of $2\\pm0.1$\\%.  Secondly, we fixed the \\textsc{fwhm} and peak of the Gaussians to these values and checked the individual spectra for orbital variablity. Unfortunately, the S/N is too low to give good individual fits (individual $1\\sigma$ errors on the offset are $\\sim 8$\\kms)  nor can we phase-lock this motion to the ephemeris of the system as determined in \\citet{vanspaandonketal09-1} due to the uncertainty in the period. The resultant radial velocity curve suggested motion, and a formal sine fit delivered a semi-amplitude of $K = 13 \\pm 2 $\\kms. To rectify for any orbital motion, we removed any measured shift compared to the mean from individual spectra.  An uniform orbital average was constructed to minimise any effects caused by varying S/N over the orbital period and finally we refit this last average with our triple Gaussian fit with as variables the common offset, the common \\textsc{fwhm} and the weighted peak as described before (see inset Figure \\ref{fig:spectrum}).   This gives a final measurement of $v_{\\mathrm{MgII}} = 35.8 \\pm 1.5$\\kms. The line has a  \\textsc{fwhm} $= 1.30 \\pm 0.05$\\AA. We have tested the accuracy of our measurement by following different recipes for combining and averaging the spectra but note that the uncertainty on the gravitational redshift is dominated by the S/N and resolution of the data and not the specific recipe used.  \\subsection{$ v_{\\mathrm{grav}}(\\mathrm{donor})$} The first correction is due to the influence of the gravitational potential of the donor on the magnesium line. However, the expected low mass of the donor star ($\\sim 0.05~M_\\odot$) gives only a small correction of $0.06\\pm 0.02$\\kms\\, near the WD surface, effectively negligible given the measurement uncertainty.   \\subsection{The systemic velocity $\\gamma$} \\label{subsec:gamma} From simultaneous double Gaussian fits to the double peaked H$\\beta$ and H$\\gamma$ disc lines, combined with radial velocity curve fits provide a systemic velocity of $-12.3 \\pm 1.2$\\kms\\, and a semi-amplitude of $K_{\\mathrm{disc}} = 36.4\\pm 1.8$\\kms.  These values are consistent with previously derived values \\citep{thorstensenetal02-3,vanspaandonketal09-1}.  As we measure $\\gamma$ from disc lines, we need to take into account that this  emission is red-shifted both by the WD and the donor star. As $v_{\\mathrm{grav}}$ is reciprocal to the distance, the amount we have to account for is minimal at the edge of the disc, $R_{\\mathrm{outer\\,disc}} = 2.2\\times 10^{8}$~m (Equation 2.61 of \\citealt{warner95-1}) giving a minimal correction to $v_{\\mathrm{grav}}$ of $1.7$\\kms. A more realistic value comes from the projected Keplerian velocity at the edge of the disc $K_{\\mathrm{disc}} \\sim 200$\\kms\\, (from the location of the disc ring in Doppler maps and the estimate for $i$, see \\citealt{vanspaandonketal09-1}) which gives a Kepler speed of $v_{\\mathrm{disc}} = K_{\\mathrm{disc}}/ \\sin i  \\sim 1070$~km~s$^{-1}$ and corresponds to a $v_{\\mathrm{grav}} = v^2/c \\sim 3.8$\\kms. Including the transverse Doppler redshift (at half the strenght of the gravitational redshift) this amounts to a total gravitational redshift at the location of the disc lines of $5.7\\pm 1.6$\\kms. The gravitational potential of the donor star has an effect of only $0.08\\pm0.02$\\kms\\, at a distance of $\\sim 3 \\times 10^{8}$~m. Hence the total $v_{\\mathrm{grav}}\\mathrm{(disc)}=5.8 \\pm  1.6$\\kms.  The $\\gamma$ from the disc lines combined with the above correction should be consistent with the $\\gamma$ suggested by the radial velocity curve from the donor star (corrected for the effects of the gravitational potential at its surface). From the Ca~\\textsc{ii} emission line in the \\textsc{i}-band we previously found $\\gamma = -13.1 \\pm 1.2$ \\citep{vanspaandonketal09-1}, the correction at the surface of the donor star is $0.34\\pm 0.15 $\\kms\\, from the donor star and $1.4\\pm0.2$\\kms\\, from the WD (again including the transverse Doppler shift). Thus we have $\\gamma_{\\mathrm{disc}}= (-12.3 \\pm 1.2)- ( 5.8\\pm 1.6) = -18.1 \\pm 2.0 $ and $\\gamma_{\\mathrm{donor}}= (-13.1 \\pm 1.2)- (1.7 \\pm 0.3) = -14.8\\pm 1.2$. Both estimates for $\\gamma$ are indeed consistent and with similar precision. Hence we have used $\\gamma_{\\mathrm{disc}}$ as it is derived from the same data set.   \\subsection{The implied WD mass} Combining all values we find for the final gravitational redshift: \\begin{eqnarray} v_{\\mathrm{grav}}\\mathrm{(WD)}&=& (35.8 \\pm 1.5) - (0.06 \\pm 0.02)  -(-18.1 \\pm 2.0) \\, \\mathrm{km\\,s}^{-1} \\nonumber\\\\ &=& 53.8 \\pm 2.5\\, \\mathrm{km\\,s}^{-1}\\nonumber \\end{eqnarray} When combined with theoretical and empirical models for the mass-radius relationship for WDs this gives a direct measurement of the WD mass. The models we used are Eggletons zero-temperature mass-radius relation as quoted by \\citet{verbunt+rappaport88-1} and several appropriate non-zero temperature models for GW Lib from \\citet{fontaineetal01-1}. We plot these models in Figure \\ref{fig:mass_radius_relation} together with the $M_1(R)$ line demanded by our measured $v_{\\mathrm{grav}}$. Accommodating the intersections between non-zero temperature models we find that $M_1 = 0.84\\pm0.02~ M_\\odot$.  \\subsection{System parameters and WD spin}  \\citet{katoetal08-1} reported the detection of superhump modulations in GW Lib. Combining this with the period from \\citet{thorstensenetal02-3} and the improved superhump excess - mass ratio relation given by \\citet{knigge06-1}, the implied mass ratio of the system is $q = 0.060 \\pm 0.008$.  We previously determined the projected radial velocity of the donor star, $K_2$ in \\citet{vanspaandonketal09-1}. Thus with an independent determination of $M_1$ from this paper, we can solve the system parameters for GW Lib under these constraints  and list these in Table \\ref{tab:systemparameters}.   If one assumes that the width of the Mg~\\textsc{ii} absorption line is dominated by rotational broadening, its \\textsc{fwhm} can constrain the spin period of the WD.  The \\textsc{fwhm} of the absorption line fit is measured to be $1.30 \\pm 0.05$\\AA\\, which corresponds to a  $v \\sin i $ of $87.0\\pm 3.4$\\kms\\, at this wavelength. Assuming $i$ is close to the value giving in Table \\ref{tab:systemparameters}, this translates into a rotation speed of $448\\pm 24$\\kms\\, at the surface of the WD. For a WD of mass $0.84\\pm 0.02~M_\\odot$ and a radius for $6.95\\pm 0.15\\times 10^{8}$~cm this results in a spin period of the WD of $97 \\pm 12$~seconds. ",
        "conclusions": "\\label{sec:discussion} We have measured the gravitational redshift of the Mg~\\textsc{ii} absorption line in high-resolution echelle spectra of GW~Lib during quiescence. Assuming an origin in the photosphere of the accreting WD, we combined this redshift with non-zero temperature mass-radius relations for WDs to derive a WD mass of $0.84\\pm 0.02M_\\odot$. Combining this independent measurement with other constrains confirms that GW~Lib is a low mass ratio system observed at very low inclination (Table \\ref{tab:systemparameters}). Because the Mg~\\textsc{ii} line was well-resolved in our data, we could also estimate the spin period of the WD to be $97 \\pm 12$~seconds if we assume the width of the line to be dominated by rotational broadening.  Our mass value is significantly below the $1.02~M_\\odot$ lower limit derived by \\citet{townsleyetal04-1}. However, their choice of preferred solution was largely driven by the WD size as implied by the measured UV flux. \\citet{szkodyetal02-4} fitted WD models to UV spectra of GW Lib. Two different single temperature solutions were explored;  $d=171$~pc ($M_1 = 0.6 M_\\odot / v_{\\mathrm{grav}} \\sim 29$\\kms) and $d=148$~pc ($M_1 = 0.8 M_\\odot / v_{\\mathrm{grav}} \\sim 49$\\kms) respectively. These solutions can be corrected to the latest parallax distance of $100^{+17}_{-13}$~pc (Thorstensen, private communication - improvement of the distance in \\citealt{thorstensenetal02-3}). At this distance, the observed UV flux implies a radius of $5\\times 10^{8}$~cm (for a WD temperature of $14\\,700$~K). However, such a small radius intercepts the mass-radius relations at $1.07~M_\\odot$, Figure \\ref{fig:mass_radius_relation} (\\textit{horizontal solid line}) and was indeed the reason \\citet{townsleyetal04-1} ruled out their lower mass solutions. Such a massive WD would result in $v_{\\mathrm{grav}} \\sim 95$\\kms, much larger than our measured redshift would suggest. We note that \\citet{szkodyetal02-4} fail to find a good single temperature fit and claim the best fit arises fitting a dual temperature model to the spectrum, allowing for a cooler zone that would place GW Lib within the ZZ Ceti instability strip for pulsating single WDs. However, the UV-flux versus implied WD radius then becomes a less clear-cut argument given the freedom of two temperatures and the surface coverage split between the two. We note that the reverse has been seen in U~Gem where the UV-flux places the WD at a lower mass than the gravitational redshift suggests \\citep{long+gilliland99-1, longetal06-1}.  As \\citet{townsleyetal04-1} warn, no WD rotation was included in their asteroseismology models and therefore their derived parameters ($M_1, M_{\\mathrm{acc}}, \\dot{M}$) may need to be revisited given our estimate of the spin period of the WD in GW Lib. A  $P_{spin}= 97 \\pm 12$~seconds should give rise to rotationally split modes. \\citet{vanzyletal04-1} conducted a thorough campaign with a baseline of over 4 years to find these modes in the power spectrum of GW Lib. The extensive search revealed a possible doublet around the 230s mode with a frequency difference between the components of $0.79$~$\\mu$Hz. If this doublet originates from the rotation splitting of an $l=1$ mode, the WD would have a spin period of $\\sim 7.3$~days which is much longer than expected for the WD in a CV.  Approached from the other side, if a similar mode is split as a result of a rotation period of $\\sim 100$~s, the frequency difference would be $\\sim 0.5$~mHz and should be visible in the power spectrum. Note that for the proposed spin period, the first order approximation for frequency splitting is no longer valid and the splitting would become asymmetric.  From both GW~Lib's and U~Gem's comparison between WD mass found by gravitational redshift and the constraints from the UV-flux we can conclude that the measurement of the gravitational redshift through WD absorption lines gives in principle a good measurement of the WD mass in CVs, but is not always fully consistent with other constraints. Nonetheless, it suggests photospheric absorption lines in the atmospheres of accreting WDs can be a viable tool for mass measurements, in particular for low mass accretion rate CVs where the light from the WD dominates and especially for low $i$ systems.  For GW Lib itself, we are still dependent on a number of assumptions in order to solve its system parameters. To properly align all methods, we need to revisit the mass from asteroseismology taking into account the effects of rotation and see if a WD model with a mass of $0.84 M_\\odot$ at the latest distance can be fitted to the UV data with reasonable temperatures. To be able to calculate a direct mass ratio, a determination of $K_1$ from the Mg~\\textsc{ii} should be possible through high S/N phase-resolved spectroscopy. If furthermore matched with an solid ephemeris using the emission from the donor star (giving $K_2$ and $\\gamma$), a fully consistent and accurate WD mass for this prototypical accreting WD pulsator is entirely feasible."
    },
    "1004/1004.3592_arXiv.txt": {
        "abstract": "Important clues to the chemical and dynamical history of elliptical galaxies are encoded in the abundances of heavy elements in the X-ray emitting plasma. We derive the hot ISM abundance pattern in inner ($0-2.3R_e$) and outer ($2.3-4.6R_e$) regions of NGC 4472 from analysis of {\\it Suzaku} spectra, supported by analysis of co-spatial {\\it XMM-Newton} spectra. The low background and relatively sharp spectral resolution of the {\\it Suzaku} XIS detectors, combined with the high luminosity and temperature in NGC 4472, enable us to derive a particularly extensive abundance pattern that encompasses O, Ne, Mg, Al, Si, S, Ar, Ca, Fe, and Ni in both regions. We apply simple chemical evolution models to these data, and conclude that the abundances are best explained by a combination of $\\alpha$-element enhanced stellar mass loss and direct injection of Type Ia supernova (SNIa) ejecta. We thus confirm the inference, based on optical data, that the stars in elliptical galaxies have supersolar $[\\alpha/Fe]$ ratios, but find that that the present-day SNIa rate is $\\sim 4-6$ times lower than the standard value. We find SNIa yield sets that reproduce Ca and Ar, or Ni, but not all three simultaneously. The low abundance of O relative to Ne and Mg implies that standard core collapse nucleosynthesis models overproduce O by $\\sim 2$. ",
        "introduction": "\\subsection{Context}  Galaxy formation encompasses the processes of star formation and largescale dynamics -- where dynamics is defined in its broadest sense to include the assembly of dark matter and baryons, as well as the exchange of mass and energy among the various galactic components and with the external environment. Star formation ultimately leads to the nuclear production of metals, while dynamical processes determine their destination. Therefore the abundance, abundance pattern, and location of metals in different galaxies in the {\\it local} universe provide fundamental insights into the construction and development of galaxies..  Chemical evolution modeling quantifies this connection. Heavy elements are both statically and explosively synthesized in evolving stars following episodes of star formation with abundances that reflects the total mass of stars formed and the initial mass function (IMF), as well as the production rate of Type Ia supernova (SNIa) progenitor binary stars. These metals are subsequently returned to the interstellar medium (ISM), and the enriched ISM may regenerate stars or escape into intergalactic space, depending on the level of supernova energy injection that accompanies the metal enrichment -- and how efficiently that energy is channeled into outflow. These processes must depend on initial conditions and environment in such a way as to produce the diverse and evolving universe of galaxies that we observe from the Local Group out to the distant universe at redshifts of 6 and greater. One can illuminate the history in the ensemble of galaxies in a cluster by measuring abundances in the intracluster medium (ICM; see, e.g., Loewenstein 2006), and in individual galaxies by measuring stellar and interstellar abundances.  It is evident that in giant elliptical galaxies, predominantly composed of stellar populations passively evolving since the early universe, the star-gas cycle was greatly accelerated and concentrated in time relative to the Milky Way and other late-type galaxies \\citep{pap06,vv07,jim07,per08}. Moreover, the prodigious metal content in the intergalactic medium of galaxy clusters dominated by ellipticals implies that strong galactic winds were driven in these galaxies. However, details such as the precise initial epoch and duration of star formation, the form of the IMF, the relative roles of SNIa and core collapse (SNII) supernovae in enriching and expelling ISM -- and how all these depend on mass and environment -- remain unclear.  Abundances in the stellar component and their correlation with structural parameters, provide key diagnostics of elliptical galaxy ages, star formation and accretion histories, and IMFs \\citep{tho05,pip09,cle09,cpm09,ts09,rec09}. Since stars dominate the baryon content of ellipticals, the level of stellar $\\alpha$-element enrichment is determined by the past integrated frequency of SNII, and hence the number of massive stars formed. For a given total mass in stars, this then constrains the IMF. The stellar $\\alpha/Fe$ ratio is then determined by the relative contributions of SNIa and SNII. Careful modeling of multiband photometry and very high signal-to-noise optical spectroscopy can, in principle, determine the abundances as well as the ages of ellipticals. However, since interpretation of colors and absorption line indices suffer from a degeneracy in their dependence on the SFH and overall metallicity -- even for an assumed IMF and distribution of abundance ratios \\citep{how05} -- optical studies of the composite stellar population cannot clearly and unambiguously measure these fundamental quantities. The additional consideration of Balmer emission line indices can assist in breaking some of the degeneracies \\citep{t00} but, e.g., does not easily distinguish a very old stellar population with a ``frosting'' of recent star formation episodes from a somewhat younger one \\citep{st07,ts09}. Since, optical line indices can only be interpreted within the context of a conjectured SFH and IMF, they are more robust in providing consistency tests of combinations of ages and abundance ratios ({\\it i.e.}, Mg/Fe) than they are in determining absolute abundances \\citep{how05}. Moreover, the effects of variations in individual elements is not easily disentangled \\citep{lee09}, and optical absorption features are only measured out to 1-2 (projected) effective radii ($R_e$) -- the abundances in, typically, half of the stellar mass are unknown.  Fortunately, mass lost by post-main-sequence stars in elliptical galaxies is heated to millions of degrees K as it is incorporated into the interstellar medium (ISM), where it is amenable to relatively robust and straightforward X-ray spectroscopic analysis \\citep{mb03}. Because the gas, with the possible exception of a few high oscillator strength lines, is optically thin the absolute abundances of a wide range of elements with prominent emission features that can include C, N, O, Ne, Mg, Al, Si, S, Ar, Ca, Fe, and Ni may be directly derived given sufficient spectral resolution, sensitivity, and bandpass. Because this approach relies on strong emission lines, X-ray derived abundances may be determined out to very large radii -- often extending to the edge of the optical galaxy and beyond \\citep{mat98}.  While hot ISM abundances may be more directly derived than stellar abundances, their interpretation requires careful deconstruction within the context of physical gasdynamical and chemical evolutionary models. Unlike the case of the ICM that dominates the baryon content in clusters, most of the metals produced by the evolving stellar population are expelled from galaxies or locked up in stars -- not accumulated in a reservoir of hot gas. Enrichment timescales for elements synthesized by SNIa, SNII, and intermediate mass stars are distinct, and coupled in a complex way to star formation, mass return, and outflow timescales. Nevertheless, one can employ reasonable assumptions, approximations, and simplifications to construct simple models that track the evolution of global abundances of these elements. Comparison of these models with observations constrain important features of the nature of the galaxy stellar population, as well as the rate and ultimate disposition of metals ejected as this population evolves. As such, they may serve as guides to subsequent more complex modeling.  In this paper, we introduce such models and apply them to the hot ISM abundances in NGC 4472 that we infer from analysis of {\\it Suzaku} and {\\it XMM-Newton} spectra. Because of its brightness and temperature structure, a particularly well-determined and wide-ranging X-ray abundance pattern is measurable in NGC 4472 -- one that is made more accessible utilizing the low internal background and sharp energy resolution of the {\\it Suzaku} XIS detectors.  \\subsection{Brief Survey of Previous X-ray Results}  High quality X-ray spectra for a large sample of elliptical galaxies was first obtained with the {\\it Advanced Satellite for Cosmology and Astrophysics} ({\\it ASCA}). Fe abundances in most of the X-ray brightest ellipticals were found to be within a factor of two of solar, with no strong evidence of non-solar ratios of $\\alpha$-elements to Fe \\citep{bf98,mom01}. X-ray spectra extracted from the {\\it Chandra} and {\\it XMM-Newton} CCD detectors take advantage of improved sensitivity and spectral resolution, and far superior angular resolution, to more cleanly derive abundance patterns and gradients (Humphrey \\& Buote 2006, and references therein). {\\it Chandra} results indicate that ISM Fe abundances are roughly solar and decline slowly with radius as far out as tens of kpc \\citep{hb06,ath07}. In addition, approximately solar Mg/Fe and Si/Fe ratios and subsolar O/Fe ratios are inferred.  Abundances derived from the {\\it XMM-Newton} reflection grating spectrometer (RGS) benefit from cleaner separation of individual features, so that effects of resonance scattering on optically thick lines, and multiphase gas on temperature-sensitive lines and line ratios, may be directly addressed. Analysis of {\\it XMM-Newton} reflection grating spectrometer (RGS) spectra of NGC 4636 by \\cite{xu02} yields subsolar Mg/Fe and O/Fe ratios, a roughly solar Ne/Fe ratio, and a supersolar N/Fe ratio. The difficulty in analyzing grating spectra of extended sources, and limitations in sensitivity and bandpass, restrict the applicability of RGS spectroscopy to X-ray luminous galaxies with the highest central X-ray surface brightnesses. \\cite{ji09} analyze {\\it Chandra} and {\\it XMM-Newton} EPIC and RGS data in 10 bright systems. They confirm the above abundance pattern with respect to O/Mg/Si/Fe, and extend to Ne and S that show no strong deviations from solar ratios, and Ni that is generally supersolar.  The lower internal background and sharper energy resolution of the {\\it Suzaku} XIS CCDs enable the derivation of abundance patterns to larger radii, and allow for a more accurate measurement of possible features at high energy originating from S, Ar, Ca, and Ni. The {\\it Suzaku} low energy sensitivity makes it suitable for measuring O as well. Results on NGC 1399 \\citep{mat07}, NGC 720 \\citep{taw08}, NGC 5044 \\citep{kom09}, NGC 507 \\citep{sat09}, and NGC 4636 \\citep{hay09} confirm the solar Si/Fe and subsolar O/Fe ratios, demonstrate that these persist to large radii, and extend this behavior to S/Fe. It is the case that this is also broadly true of Mg/Fe; however, Mg/Fe is somewhat lower in NGC 720 and NGC 1399. Ne/Fe seemingly shows more variation, coming in at solar in NGC 4636, subsolar in NGC 720 and supersolar in NGC 507 and NGC 5044. Ne does not seem to trace O as might be expected.  Optical spectra imply that the stellar ratio of $\\alpha$ elements to Fe, expressed logarithmically with respect to solar as $[\\alpha/Fe]$, is such that $[\\alpha/Fe]\\sim 0.3-0.5$ -- although there is little known about this ratio well beyond $R_e$ or about the relative abundances of different $\\alpha$ elements. If stellar mass loss dominates ISM enrichment, they should display the same $\\alpha$-element overabundance. Based on the the above summary, this is clearly not the case. Direct injection of SNIa will skew that abundance pattern towards lower $[\\alpha/Fe]$, but should do so in a way that does not effect relative abundances of elements with low SNIa nucleosynthetic yields such as O, Ne, and Mg. The ratios, with respect to Fe, of Si, S, Ar, and Ca may be less affected, depending on the level of SNIa enrichment since these elements are produced with moderate efficiency by SNIa. The solar ratio seen for most of the $\\alpha$ elements in the hot ISM indicates a significant role for SNIa enrichment; but, the fact that O is particularly underabundant while Mg and Ne are as abundant as Si and S does not comfortably fit into this scheme. We examine these issues more quantitatively below with respect to the ISM abundances in NGC 4472, where measurements of Ca and Ar in the ISM add a new twist to attempts to understand the enrichment of this galaxy. ",
        "conclusions": "We have derived the hot ISM abundance pattern in inner ($0-2.3R_e$) and outer ($2.3-4.6R_e$) regions of NGC 4472 from analysis of {\\it Suzaku} spectra and used them to study the chemical evolution in this elliptical galaxy. The low {\\it Suzaku} background and relatively sharp spectral resolution of the {\\it Suzaku} XIS detectors enabled us to extend the range of accurately measured abundances beyond what is feasible with {\\it Chandra} or {\\it XMM-Newton} to S, Ar, Ca, and possibly Al -- while we find general agreement with overlapping {\\it XMM-Newton} results. The abundances of Ne, Mg, Si, S, Fe, and Ni may be explained by enrichment via a combination of $\\alpha$-element enhanced stellar mass loss and direct injection of SNIa with W7 yields exploding at a current rate of 0.027 SNU; however, additional SNIa production of Ca and Ar is required. Ca and Ar are reproduced in models (with a SNIa rate of 0.035 SNU) where WDD1 yields are adopted, but at the price of disrupting the previous concordance for S and Ni.  Models with the standard SNIa rate of 0.16 SNU badly overpredict Fe, unless the enrichment from stellar mass loss is diluted. We introduced models where this is a consequence of outflow biased in favor of high metallicity, or an inertial effect resulting from the preexistence of a metal poor ISM. However, such models predict an underabundance of elements primarily synthesized in SNII (Ne and Mg for W7 and WDD1 SNIa yields, and Si, S, Ar, and Ca as well for W7) and offer a poorer match to the observed abundance pattern.  We confirm previous measurements of a low O abundance; the low O/Mg and O/Ne ratios imply that the SNII yields in standard models overproduce O by $\\sim 2$.  Analysis of additional {\\it Suzaku} and {\\it XMM-Newton} spectra of elliptical galaxies will be used to investigate the universality of these results, and probe for diversity in the star formation and SNIa histories of giant elliptical galaxies. Application of more rigorous chemical evolution models will sharpen these conclusions, and should yield new information on the formation and evolution of these systems and indirect constraints on the physics of Type Ia supernovae."
    },
    "1004/1004.1418_arXiv.txt": {
        "abstract": "{\\em Fermi} Gamma ray Space Telescope observations of the flat spectrum radio quasar 3C~454.3 show a spectral-index change $\\Delta \\Gamma \\cong 1.2\\pm 0.3$ at break energy $E_{br} \\approx 2.4\\pm0.3$ GeV. Such a sharp break is inconsistent with a cooling electron distribution and is poorly fit with a synchrotron self-Compton model. We show that a combination of two components, namely the Compton-scattered disk and broad-line region (BLR) radiations, explains this spectral break and gives a good fit to the quasi-simultaneous radio, optical/UV, X-ray, and $\\gamma$-ray spectral energy distribution observed in 2008 August. A sharp break can be produced independent of the emitting region's distance from the central black hole if the BLR has a gradient in density $\\propto R^{-2}$, consistent with a wind model for the BLR. ",
        "introduction": "\\label{intro}   The flat spectrum radio quasar (FSRQ) \\object{3C~454.3} (\\object{PKS~2251+158}) is representative of blazars with large, $\\gtrsim 10^{48}$ erg s$^{-1}$ apparent $\\gamma$-ray luminosities, and is one of the brightest sources seen with the recently launched {\\em Fermi} Gamma-Ray Space Telescope \\citep{tosti08_atel,escande09_atel}. It has also been detected with AGILE up to $\\approx 3$ GeV, including a giant outburst in 2007 prior to the {\\em Fermi} launch \\citep{vercellone09,vercellone10}.  Based on $> 100$ MeV data taken during the checkout and early science phase of the mission (2008 July 7-- October 6), this source was found to be variable on timescales down to $t_{v}\\sim$ few days \\citep{abdo09_3c454.3}.  A well-sampled data set using the ground-based SMARTS (Small and Moderate Aperture Research Telescope System) in the B, V, R, J, and K bands, and the {\\em Swift} X-ray telescope (XRT) and ultraviolet-optical telescope (UVOT) was collected simultaneously with {\\em Fermi} observations in the period 2008 August through December \\citep{bonning09}.  These multiwavelength observations revealed that the infrared, optical, and $\\g$-rays were have similar variability patterns.  The $\\gamma$ rays display similar relative flux changes compared to the longer optical wavelengths, but smaller relative flux changes as the short wavelength optical emission, which \\citet{bonning09} interpreted as being due to an underlying accretion disk.  The {\\em Swift} X-ray data did not strongly correlate with the other wavelengths, though \\citet{jorstad10} and \\citet{bonnoli10} find a correlation at different epochs.  The {\\em Fermi} observations of \\object{3C~454.3} revealed a broken power-law spectrum, with $\\Gamma\\cong 2.3\\pm 0.1$ (where the photon flux $\\Phi\\propto E^{-\\Gamma}$) and $\\Gamma\\cong 3.5\\pm 0.25$, below and above a break energy of $2.4\\pm0.3$ GeV, respectively \\citep{abdo09_3c454.3}.  This spectrum is not consistent with Compton scattering of a lower-energy photon source by jet electrons in the fast- or slow-cooling regimes, the standard interpretation for the high energy spectrum in leptonic models of blazars. \\citet{abdo09_3c454.3} also consider it unlikely that the spectral break is a result of $\\g\\g$ absorption with a local radiation field or the extragalactic background light (EBL), the former because this would require a lower jet bulk Lorentz factor than that inferred from superluminal radio observations of \\object{3C~454.3} \\citep[$\\Gamma_{bulk}\\approx15$;][]{jorstad05}, and the latter because the universe is transparent to 40~GeV photons at $z\\la1.0$ for all EBL models \\citep[e.g.,][]{finke10_EBLmodel,stecker06,gilmore09}. \\citet{abdo09_3c454.3} thus consider the most likely explanation to be that the spectral break is caused by an intrinsic break in the electron distribution.  However, as explained below, this is unlikely, given the broadband spectral energy distribution of this source.  In this letter, we show that a combination of two Compton scattering components can explain the break in the LAT spectrum of \\object{3C~454.3}. In particular, we consider a combination of components from Compton-scattered accretion disk and broad-line region (BLR) radiation, both of which should be strong seed photon sources in FSRQs.  The implications of this model for the BLR and for the variability of the emitted radiation are explored.  Finally, we conclude with a brief discussion, with implications for breaks in other sources.  We use parameters $H_0=71$ km s$^{-1}$ Mpc$^{-3}$, $\\Omega_m=0.27$, and $\\Omega_\\Lambda=0.73$ so that \\object{3C~454.3}, with a redshift of $z=0.859$, has a luminosity distance of $d_L=5.5$ Gpc. ",
        "conclusions": "Assuming the dual-component Compton-scattering scenario presented in this paper is correct, what can it tell us about the location of the $\\g$-ray emission region?  It implies that this region is $R_i < r \\ll \\G^4r_g$, so that, using the parameters from the modeling of 3C~454.3, the location of the jet falls over a large range from $0.1\\ \\pc < r \\ll 100\\ \\pc$.  However, the external energy density from disk and BLR radiations is $\\propto r^{-3}$, a steep decline, so that the blob must be within $\\approx 0.1$ pc for a significant scattering component. Recent modeling of 3C 454.3 \\citep{bonnoli10} also finds that the $\\gamma$-ray emitting region is within the BLR. By contrast, \\citet{sikora08} suggested that the $\\g$-rays observed by EGRET and AGILE were caused by Compton-scattering of radiation from hot dust by a blob located $\\sim 10$ pc from the central black hole, though detailed  modeling is still needed to determine if this model can explain the sharp spectral break in the {\\em Fermi} observations.  The blazar 3C 279 also shows a similar spectral break during the onset of an extended episode of optical polarization, while the jet emission region could still be within the BLR region \\citep{abdo10_3c279}.  Whether this model can predict similar spectral breaks in other FSRQs and low-synchrotron-peaked blazars, such as AO 0235+164 or PKS 1502+106 \\citep{abdo10_sed}, depends mainly on the properties of the BLR in these sources, in particular, whether $\\tau_{BLR} \\approx r_g/R_i$.  In specific blazars, especially at higher $z$, $\\g\\g$-absorption by scattered disk radiation \\citep{reimer07} could still be effective.  Detailed modeling of simultaneous SED data will test these scenarios, as it does for the case of \\object{3C~454.3}. If our model is correct, it provides evidence for a wind model of the BLR \\citep{murray95,chiang96,elvis00}."
    },
    "1004/1004.1903_arXiv.txt": {
        "abstract": "{ Double-lined, detached eclipsing binaries are our main source for accurate stellar masses and radii. In this paper we focus on the 1.15--1.70 \\Msun\\ interval where convective core overshoot is gradually ramped up in theoretical evolutionary models. } {We aim to determine absolute dimensions and abundances for the F-type detached eclipsing binary \\BK, and to perform a detailed comparison with results from recent stellar evo\\-lu\\-tio\\-nary models, including a sample of previously studied systems with accurate parameters.} {$uvby$ light curves and $uvby\\beta$ standard photometry were obtained with the Str\\\"omgren Automatic Telescope, ESO, La Silla, and high-resolution spectra were acquired with the FIES spectrograph at the Nordic Optical Telescope, La Palma.} { The $5\\fd49$ period orbit of \\BK\\ is slightly eccentric ($e$ = 0.053). The two components are quite different with masses and radii of ($1.414 \\pm 0.007$ \\Msun, $1.988 \\pm 0.008$ \\Rsun) and ($1.257 \\pm 0.005$ \\Msun, $1.474 \\pm 0.017$ \\Rsun), respectively. The measured rotational velocities are $16.6 \\pm 0.2$ (primary) and $13.4 \\pm 0.2$ (secondary) \\kms. For the secondary component this corresponds to (pseudo)synchronous rotation, whereas the primary component seems to rotate at a slightly lower rate. We derive an iron abundance of \\feh\\,$=-0.12\\pm0.07$ and similar abundances for Si, Ca, Sc, Ti, Cr and Ni. The stars have evolved to the upper half of the main-sequence band. Yonsei-Yale and Victoria-Regina evolutionary models for the observed metal abundance reproduce \\BK\\ at ages of 2.75 and 2.50 Gyr, respectively, but tend to predict a lower age for the more massive primary component than for the secondary. We find the same age trend for three other upper main-sequence systems in a sample of well studied eclipsing binaries with components in the 1.15--1.70 \\Msun\\ range. We also find that the Yonsei-Yale models systematically predict higher ages than the Victoria-Regina models. The sample includes BW\\,Aqr, and as a supplement we have determined a \\feh\\ abundance of $-0.07 \\pm 0.11$ for this late F-type binary. } {We propose to use \\BK, BW\\,Aqr, and other well-studied 1.15--1.70 \\Msun\\ eclipsing binaries to fine-tune convective core overshoot, diffusion, and possibly other ingredients of modern theoretical evolutionary models. } ",
        "introduction": "\\label{sec:intro}  Detached, double-lined eclipsing binaries (dEB) are our main source for stellar masses and radii, today accurate to 1\\% or better (Torres et al. \\cite{tag09}), and they also provide stringent tests of various aspects of stellar evolutionary models. For this purpose, well-established abundance information is needed, as demonstrated by e.g. Clausen et al. (\\cite{avw08}, hereafter CTB08). One of the troublesome ingredients in theoretical models for stars heavier than the Sun is the amount of and treatment of convective core overshoot, and in this paper we focus on that aspect.  The literature on the existence and calibration of core overshoot is extensive, and here we only draw attention to a few studies based on binary and cluster results. From a sample of 1.5--2.5 \\Msun\\ dEBs and turn-off stars in IC~4651 and NGC~2680, Andersen et al. (\\cite{anc90}) found strong evidence for convective overshoot in intermediate-mass stars. Clausen (\\cite{jvc91}) found indication for core overshoot for the 1.4+1.5 \\Msun\\ late-F type dEB \\object{BW\\,Aqr} and discussed \\BK\\ as well, and recently, Lacy et al. (\\cite{ltc08}) found that for the 1.5 \\Msun\\ F7~V system \\object{GX\\,Gem}, the lowest core overshoot parameter $\\alpha_{\\rm ov}$ consistent with observations is approximately 0.18 (in units of the pressure scale height). The question of mass-dependence of the degree of core overshoot has -- again based on dEB samples -- been addressed by e.g. Ribas et al. (\\cite{rjg00}) and Claret (\\cite{c07}), but they arrive at very different conclusions.  Most recent grids of stellar evolutionary calculations include core overshoot, but the recipes in terms of mass and abundance dependence are somewhat different. We refer to Sect.~\\ref{sec:models} for details.  Our main motivation to undertake a study of \\BK\\ has been that a) it has evolved to the upper part of the main-sequence band and is therefore well-suited for core overshoot tests, but published dimensions are not of sufficient quality, and b) few similar well-studied systems are known. Below we present absolute dimensions and abundances based on new $uvby$ light curves and high-resolution spectra and compare \\BK\\ and other similar systems with Yonsei-Yale and Victoria-Regina stellar evolutionary models.  In Appendix~\\ref{sec:bwaqr} we present a spectroscopic abundance analysis for BW\\,Aqr as supplement to the study of this binary by Clausen (\\cite{jvc91}). ",
        "conclusions": "\\label{sec:sum}  From state-of-the-art observations and analyses, precise (0.4--1.2\\%) absolute dimensions have been established for the components of the late F-type detached eclipsing binary \\BK\\ ($P = 5\\fd49$, $e$ = 0.0053); see Table~\\ref{tab:bkpeg_absdim}. A detailed spectroscopic analysis yields an iron abundance relative to the Sun of \\feh\\,$=-0.12\\pm0.07$ and similar relative abundances for Si, Ca, Sc, Ti, Cr, and Ni. The measured rotational velocities are $16.6 \\pm 0.2$ (primary) and $13.4 \\pm 0.2$ (secondary) \\kms. For the secondary component this corresponds to (pseudo)synchronous rotation, whereas the primary component seems to rotate at a slightly lower rate.  The 1.41 and 1.26 \\Msun\\ components of \\BK\\ have evolved to the upper half of the main-sequence band. Yonsei-Yale and Victoria-Regina solar scaled evolutionary models for the observed metal abundance reproduce \\BK\\ at ages of 2.75 and 2.50 Gyr, respectively, but tend to predict a lower age for the more massive primary component than for the secondary. If real, this might be due to less than perfect calibration of the amount of convective core overshoot of the models as function of mass (and metal abundance).  For this reason, we have performed model comparisons for a sample of eight additional well-studied binaries with component masses in the 1.15--1.70 \\Msun\\ interval where convective core overshoot is gradually ramped up in the models; see Table~\\ref{tab:systems}. We find that {\\it a}) the Yonsei-Yale models systematically predict higher ages than the Victoria-Regina models, and that {\\it b}) the three other most evolved systems in the sample share the age difference trend seen for \\BK.  We propose to use the sample to fine-tune the core overshoot treatment, as well as other model ingredients, and to clarify why the two model grids predict different ages. The sample should be expanded by a number of new F-type systems under study, binary cluster members, and the unique K0IV+F7V binary AI\\,Phe (1.23+1.19 \\Msun)."
    },
    "1004/1004.2900_arXiv.txt": {
        "abstract": "We present broadband (radio, optical, and X-ray) light curves and spectra of the afterglows of four long-duration gamma-ray bursts (GRBs\\,090323, 090328, 090902B, and 090926A) detected by the Gamma-Ray Burst Monitor (GBM) and Large Area Telescope (LAT) instruments on the \\fermi\\ satellite. With its wide spectral bandpass, extending to GeV energies, \\fermi\\ is sensitive to GRBs with very large isotropic energy releases (10$^{54}$ erg). Although rare, these events are particularly important for testing GRB central-engine models. When combined with spectroscopic redshifts, our afterglow data for these four events are able to constrain jet collimation angles, the density structure of the circumburst medium, and both the true radiated energy release and the kinetic energy of the outflows. In agreement with our earlier work, we find that the relativistic energy budget of at least one of these events (GRB\\,090926A) exceeds the canonical value of $10^{51}$\\,erg by an order of magnitude. Such energies pose a severe challenge for models in which the GRB is powered by a magnetar or neutrino-driven collapsar, but remain compatible with theoretical expectations for magneto-hydrodynamical collapsar models (e.g., the Blandford-Znajek mechanism).  Our jet opening angles ($\\theta$) are similar to those found for pre-\\fermi\\ GRBs, but the large initial Lorentz factors ($\\Gamma_{0}$) inferred from the detection of GeV photons imply $\\theta\\Gamma_{0} \\approx 70$--90, values which are above those predicted in magnetohydrodynamic models of jet acceleration. Finally, we find that these \\fermi-LAT events preferentially occur in a low-density circumburst environment, and we speculate that this might result from the lower mass-loss rates of their lower-metallicity progenitor stars.  Future studies of \\fermi-LAT afterglows in the radio with the order-of-magnitude improvement in sensitivity offered by the Extended Very Large Array should definitively establish the relativistic energy budgets of these events. ",
        "introduction": "\\label{sec:intro} Long-duration gamma-ray bursts (GRBs\\footnote{Throughout this work, we shall refer exclusively to long-duration GRBs (i.e., those apparently with massive-star progenitors) unless explicitly stated otherwise.}), like hydrogen-deficient Type Ib/c supernovae (SNe Ib/c), result from the gravitational collapse of the evolved core of a massive star.  The main characteristic that sets GRBs apart from other SNe is that a substantial fraction of the energy of the explosion is coupled to relativistic ejecta. A compact central engine is responsible for accelerating and collimating these jet-like outflows and driving the SN explosions \\citep{wb06,grf09,scp+10}.  The precise nature of the central engine which powers GRB-SNe, however, remains an open question.  Motivated by empirical constraints, all viable central-engine models for long-duration GRBs share some common characteristics (e.g., \\citealt{p05}).  They must produce a collimated outflow with an initial Lorentz factor ($\\Gamma_{0}$) of a few hundred on observed time scales of 10--100\\,s, with luminosities and kinetic energies of order 10$^{50}$\\,erg s$^{-1}$ and 10$^{51}$ erg, respectively. Leading models include the ``collapsar'' model in which a relativistic jet is produced from a rotating black hole/accretion disk system \\citep{w93,mw99}, and the ``magnetar'' model in which the rapid energy loss from a newly born millisecond neutron star (formed either from the gravitational collapse of a massive star or from accreting or coalescing white dwarfs) with a 10$^{15}$\\,G magnetic field drives a Poynting flux-dominated relativistic outflow \\citep{u92}.  These and other more exotic models for GRB central engines are highly constrained by their energetics. The prompt high-energy emission, when combined with a spectroscopically determined redshift (and hence distance measurement), yields the isotropic radiated gamma-ray energy ($E_{\\gamma,\\mathrm{iso}}$). Well-sampled afterglow observations allow both a measurement of the degree of collimation (and hence the true beaming-corrected energy release in the prompt emission, $E_\\gamma$) and the kinetic energy remaining in the shock that powers the broadband afterglow emission ($E_{\\mathrm{KE}}$). Such measurements, made nearly a decade ago, pointed to a total {\\it relativistic} energy yield ($E_{\\mathrm{rel}} \\approx E_\\gamma + E_{\\mathrm{KE}}$) of $\\sim 10^{51}$\\,erg \\citep{fks+01,pk01b,fw01,bfk03,bkf03}.  Since that time, there has been growing evidence for a considerable range in the relativistic energy scale $E_{\\mathrm{rel}}$, suggesting either a diversity in central engines or their properties. Most notably, a population of nearby (redshift $z \\lesssim 0.1$) subenergetic long-duration GRBs have been identified \\citep{bfk03,skb+04,skn+06}. They too are associated with SNe~Ib/c, but their relativistic energy release is a factor of 100 less than that of typical cosmological GRBs and their outflows are significantly less collimated (quasi-spherical). Since they can only be detected at low redshifts where the comparative volume for discovery is low, they are small in total number.  But their volumetric rate is inferred to be 10--100 times larger than that of the more distant long-duration GRBs \\citep{skn+06,cbv+06,lzv+07}.  More recently, evidence has been growing for a class of GRBs whose total relativistic energy release is at least an order of magnitude above the canonical value of $10^{51}$\\,erg (e.g., \\citealt{cfh+10} and references therein).  Unlike subluminous events, the total energy budget of these hyper-energetic events poses a significant challenge for some progenitor models.  In particular, models in which the GRB is powered by a magnetar or a neutrino-driven collapsar are strongly disfavored.  On the other hand, collapsars driven by magneto-hydrodynamical (MHD) processes, such as the Blandford-Znajek mechanism \\citep{bz77}, can naturally accomodate energy budgets as large as $10^{53}$\\,erg.  Unfortunately, it has been rather difficult to constrain the beaming-corrected energetics for the hundreds of GRBs detected by the \\swift\\ satellite \\citep{gcg+04}.  The reasons for this difficulty are now largely understood. First, the relatively narrow energy bandpass (15--150\\,keV) can miss entirely the peak of the gamma-ray spectrum, making estimates of $E_{\\gamma,\\mathrm{iso}}$ highly uncertain. Second, there has been a dearth of measurements of jet opening angles (e.g., \\citealt{p07,kb08,lrz+08,rlb+09}) and well-sampled multi-wavelength GRB afterglows (used to derived the afterglow kinetic energy $E_{\\mathrm{KE}}$). \\swift\\ GRBs are on average more than twice as distant \\citep{jlf+06} and therefore significantly fainter ($\\sim 1.5$\\,mag in the optical; \\citealt{bkf+05,kkz+07}) than GRBs in previous samples.  In large part this is due to selection effects: a combination of bandpass and sensitivity from \\swift\\ has preferentially selected the faint end of the luminosity function --- GRBs with low isotropic energy release but large opening angles \\citep{psf03}.  With its nearly seven decades in energy coverage (10\\,keV -- 100\\,GeV), \\fermi\\ can provide unparalleled constraints on this subsample of the most luminous events. In light of the empirical relation between the peak energy of the gamma-ray spectrum and the isotropic gamma-ray energy release (the $E_{\\mathrm{p}}$--$E_{\\gamma,\\mathrm{iso}}$, or ``Amati'' relation; \\citealt{a06}), MeV/GeV events detected by either the Gamma-Ray Burst Monitor (GBM, 8\\,keV -- 40\\,MeV; \\citealt{mlb+09}) or the Large Area Telescope (LAT, 20\\,MeV -- 300\\,GeV; \\citealt{aaa+09c}) onboard \\fermi\\ preferentially select a sample of GRBs with large isotropic energy release (Figure~\\ref{fig:egamma}). High-$E_{\\gamma,\\mathrm{iso}}$ events also have brighter X-ray and optical afterglows on average \\citep{nfp09}.  Follow-up afterglow observations can then determine whether these GRBs are highly beamed events ($\\theta \\lesssim 2^{\\circ}$) with a typical energy release or true hyper-energetic GRBs.  \\begin{figure}[t] \\epsscale{1.2} \\plotone{f1.eps} \\caption{Prompt isotropic gamma-ray energy release ($E_{\\gamma,\\mathrm{iso}}$) of GRBs.  With its soft, narrow bandpass (15--150\\,keV), \\swift\\ typically selects events with smaller isotropic energy release but larger opening angles than previous missions, which triggered predominantly in the MeV bandpass \\citep{psf03}.  GRBs detected at GeV energies with the \\fermi-LAT all fall at the brightest end of the isotropic energy distribution, and must therefore be highly collimated to achieve a canonical beaming-corrected energy release of $\\sim 10^{51}$\\,erg.  References: pre-\\swift: \\citet{a06}; \\swift: \\citet{bkb+07}; \\fermi-LAT: \\citet{gck+09}, this work.} \\label{fig:egamma} \\end{figure}  The \\fermi-LAT offers a further advantage over previous GRB missions sensitive only at MeV and keV energies by providing strict constraints on the initial Lorentz factor of the relativistic outflow.  To avoid $e^{+}-e^{-}$ pair production (and the accompanying thermal spectrum), the GRB jet must be moving towards the observer with ultra-relativistic speeds (the ``compactness'' problem; \\citealt{cr78}).  The higher the energy of the most energetic photon detected from a GRB, the more strict the lower limit on the outflow Lorentz factor will be. Combining the Lorentz factor limits for the most relativistic GRBs with inferred jet opening angles from broadband afterglow models can provide critical diagnostics of the jet acceleration mechanism.  \\begin{figure*}[t] \\plotone{f2.eps} \\caption{The broadband radio (blue), optical (red), and X-ray (black) light curve of \\grba.  The best-fit model is plotted in solid lines (see Table~\\ref{tab:0323mod} for parameters).  The identical model parameters for an isotropic explosion are plotted as the dashed lines.  The strength of the possible modulation of the radio afterglow caused by interstellar scintillation (e.g., \\citealt{fwk00}) is indicated by the light-blue shaded region. The model provides a reasonable fit in all bandpasses.  It is clear that any jet break must occur at $t > 10$\\,days, although the upper bound on the jet break time is only weakly constrained.} \\label{fig:0323} \\end{figure*}  Here we report on broadband (radio, optical, and X-ray) observations of four long-duration GRBs detected by the \\fermi-LAT at GeV energies: GRBs\\,090323, 090328, 090902B, and 090926A.  For each event we construct afterglow models to constrain the collimation and beaming-corrected energetics, and we compare these LAT events with previous GRBs detected at other energies (i.e., keV energies from \\swift, and MeV energies from pre-\\swift\\ satellites).  For three of these GRBs, we also present the optical spectra used to determine the afterglow redshift.  A more thorough analysis of the host-galaxy properties of these events will be presented in a forthcoming work.  Throughout this paper, we adopt a standard $\\Lambda$CDM cosmology with $H_{0}$ = 71\\,km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\mathrm{m}} = 0.27$, and $\\Omega_{\\Lambda} = 1 - \\Omega_{\\mathrm{m}} = 0.73$ \\citep{sbd+07}.  We define the flux-density power-law temporal and spectral decay indices $\\alpha$ and $\\beta$ as $f_{\\nu} \\propto t^{-\\alpha} \\nu^{-\\beta}$ (e.g., \\citealp{spn98}).  All quoted uncertainties are 1$\\sigma$ ($68\\%$) confidence intervals unless otherwise noted. ",
        "conclusions": "\\label{sec:conclusions} We have undertaken extensive broadband continuum (radio, optical, and X-ray) and spectroscopic observations of four long-duration GRBs (GRBs\\,090323, 090328, 090902B, and 090926A) detected by the LAT instrument on the \\fermi\\ satellite at GeV energies. This work was motivated by the realization that \\fermi\\ is especially sensitive to GRBs with large isotropic energy release, and hence provides an interesting sample of events to test GRB central-engine models and their relativistic outflows. Our afterglow models constrain the jet break times and the density of the circumburst medium, from which we derive the collimation angle and hence beaming-corrected energy release for each event. We find three GRBs with a total relativistic content about an order of magnitude in excess of the canonical 10$^{51}$ erg, with \\grbd\\ almost certainly in excess of $10^{52}$\\,erg. This analysis provides support for our earlier claim of a class of hyper-energetic GRBs. The discovery of more GRBs with total energy release $\\gtrsim 10^{52}$\\,erg is troubling for central engines in which the energy to drive the jet is derived either from a rotating magnetar or collapsars powered by neutrino annihilation.  For this reason, we are led to believe that, at least for hyper-energetic GRBs, the massive star progenitor collapses directly to a black hole and the rotational energy of this system is extracted via the Blandford-Znajek process.  Although we find relatively narrow opening angles for all four events ($\\theta \\lesssim 10^{\\circ}$), the extreme initial Lorentz factors inferred for these LAT events imply that the product $\\theta \\Gamma_{0}$ can be a factor of 5--10 larger than estimates of previous GRBs detected at MeV energies.  These values are inconsistent with recent simulations of low-magnetization MHD jets, suggesting that the outflow may be at least initially Poynting-flux dominated.  If this is indeed the case, it is unclear how the initial kinetic energy of the outflow is converted to prompt gamma-ray emission.  Interestingly, for the three events having sufficient radio coverage to derive a circumburst density, we find anomalously large values of $E_{\\mathrm{KE,iso}} / n$ (or, for a wind-like medium, $E_{\\mathrm{KE,iso}} / A_{*}$).  While the large $E_{\\mathrm{KE,iso}}$ values are simple to understand, the low circumburst densities require a more complex explanation.  One possibility is that the progenitor stars of LAT GRBs are somehow different from the progenitors of most previous GRBs detected at MeV or keV energies.  It is currently thought that GRB progenitors are distinguished from the progenitors of ordinary SNe Ib/c by their low metallicities (e.g., \\citealt{mw99,wh06,mkk+08}): the lower mass-loss rates allow the progenitors of GRBs to keep more angular momentum.  The increased rotation evacuates a cavity through which a relativistic jet can propagate.  If LAT events have larger initial Lorentz factors, it may be that they come from lower metallicity progenitors with minimal pre-explosion mass loss.  Observations of the host galaxies of these events, both through absorption and emission spectroscopy, may help shed light on this matter.  It is also possible that the low density preference is the result of other, more subtle, selection effects.  In particular, if the GeV emission arises in the external shock, the \\fermi-LAT could be biased towards events in low-density environments.  If the circumburst density is too high, the blastwave will decelerate at small radii (Equations~\\ref{eqn:gammaISM} and \\ref{eqn:gammawind}), where the outflow may be opaque to GeV photons.  More observations of LAT GRBs, particularly at very early times, would help to investigate this hypothesis.  We end by emphasizing the importance of afterglow observations of high-$E_{\\gamma,\\mathrm{iso}}$ events in the \\fermi\\ era to provide further confirmation of this picture. Such GRBs are either highly collimated outflows ($\\theta \\lesssim 2^{\\circ}$) with a typical energy release, or truly hyper-energetic events; both represent extreme tests of jet collimation and central-engine models, respectively. Current efforts suffer from delays in LAT localizations and limited ground-based afterglow follow-up efforts. The latter can be improved by focusing rare follow-up resources on \\fermi-LAT GRBs; as \\citet{nfp09} and \\citet{mkr+10} have shown, these events have brighter X-ray and optical afterglows on average, and are therefore accessible even for moderate-aperture optical facilities.  Targeting these bright afterglows will make it easier to measure the jet breaks, which have proven almost impossible to obtain in the \\swift\\ era. Finally, we note that one testable consequence of hyper-energetic GRBs is long-lived afterglow emission ($\\gtrsim 1$\\,yr). If the shock microphysics and the circumburst density do not undergo a drastic evolution, it should be possible to detect these afterglows with the upcoming generation of radio facilities and carry out calorimetry measurements."
    },
    "1004/1004.2077_arXiv.txt": {
        "abstract": "Type 2 AGNs with intrinsically  weak broad emission lines (BELs) would be  exceptions to  the  unified  model. After  examining  a number  of proposed   candidates  critically,   we  find   that  the   sample  is contaminated   significantly  by  objects   with  BELs   of  strengths indicating that  they actually contain intermediate-type  AGNs, plus a few Compton-thick sources as revealed by extremely low ratios of X-ray to nuclear IR luminosities.  We develop quantitative metrics that show two (NGC 3147  and NGC 4594) of the remaining  candidates to have BELs 2-3  orders   of  magnitude  weaker  than  those   of  typical  type-1 AGNs. Several  more galaxies remain as candidates  to have anomalously weak  BELs, but  this status  cannot  be confirmed  with the  existing information.  Although  the parent sample  is poorly defined,  the two confirmed objects are  well under 1\\% of its  total number of members, showing that  the absence of a  BEL is possible, but  very uncommon in AGN. We  evaluate these two  objects in detail  using multi-wavelength measurements  including new IR  data obtained  with {\\it  Spitzer} and ground-based  optical  spectropolarimeteric  observations.  They  have little   X-ray  extinction  with   $N_{\\rm  H}$   $<$  $\\sim$10$^{21}$ cm$^{-2}$.  Their IR spectra  show strong silicate emission (NGC 4594) or weak  aromatic features on a  generally power law  continuum with a suggestion of silicates  in emission (NGC 3147).  No  polarized BEL is detected in NGC  3147. These results indicate that  the two unobscured type-2 objects have circumnuclear tori that are approximately face-on. Combined  with their  X-ray and  optical/UV properties,  this behavior implies that  we have an unobscured  view of the nuclei  and thus that they have  {\\it intrinsically} weak BELs. We  compare their properties with those of  the other less-extreme candidates. We  then compare the distributions of bolometric luminosities  and accretion rates of these objects with theoretical models that predict weak BELs. ",
        "introduction": "Active galactic nuclei (AGNs) are optically classified into type 1 and type 2:  a type 1  object shows broad  emission lines (BELs)  that are absent in  type 2s.  The strict AGN  unified model \\citep{Antonucci93, Urry00} states  that the apparent large diversity  in AGN properties is   caused   almost entirely  by   different   viewing   angles  and   nuclear luminosities. An  AGN is identified  as type 2  when a dusty  torus is viewed  at  large inclination  (more  edge-on)  and  obscures the  BEL region, while a  torus in a type 1 AGN is  viewed at lower inclination and does not block the BEL region.  In  the  past few  decades,  the  unified model  has  achieved tremendous  success  in explaining  the  general  properties of  AGNs. However, there are  indications that its axioms may  not be universal; parameters  other  than orientation  may  influence  the observed  AGN properties.  For example,  the host galaxy may play  an important role in the type  1/type 2 division.  It has been  found that a significant fraction of type 2 AGNs are actually obscured by kpc-scale material in the    host   galaxy    instead   of    the   sub-kpc    dusty   torus \\citep[e.g.][]{Keel80, Maiolino95a,  Rigby06, Diamond-Stanic09a}.  The typical level of star formation in  a type 2 host galaxy is also found to be  higher than that in  a type 1  host \\citep{Maiolino95b, Shi07}, implying that  a more dusty type  2 host increases  the probability of obscuration  of   the  BEL  region.   In   radio-loud  AGNs,  infrared observations  indicate that  a significant  fraction of  type  2 radio galaxies  harbor  intrinsically weaker  nuclei  than  type 1  galaxies \\citep{Whysong04,  Shi05, Shi07,  Ogle06, Cleary07}.   There  are also claims  that  AGNs with  intrinsically  weak  BEL  emission may  exist \\citep[e.g.][]{Laor03, Hawkins04,  Diamond-Stanic09b, Hopkins09}.  For example,  spectropolarimetric observations  reveal  hidden BEL  (HBEL) regions  in  only  half  of  type 2  Seyferts  \\citep{Kay94,  Moran00, Tran01}.  The comparison between Seyfert 2 galaxies with and without HBELs  reveals  a significant  difference  in  the nuclear  luminosity between  the two  populations  \\citep{Tran01, Gu02}.   However, it  is unclear  if the  low nuclear  luminosity of  Seyfert 2  nuclei without HBELs is simply due to a  larger host galaxy contamination for a given nuclear  luminosity  or if  Seyfert  2  galaxies  without HBELs  harbor genuinely weak nuclei.  In  the standard  model, a  type 2  object shows  only  narrow optical emission lines in  total light and has strong  X-ray obscuration.  The HI  equivalent   column  density  can  be   inferred  through  fitting photoelectric absorption  models to the X-ray  spectrum.  As expected, the  HI column  densities  of 96\\%  of  type 2  AGNs  are larger  than 10$^{22}$  cm$^{-2}$, and they  are therefore  considered to  be X-ray obscured \\citep{Risaliti99}.   However, with the  increasing volume of X-ray spectral data,  a small sample of type 2 AGNs  has been found to be   relatively  unobscured  in   the  X-ray   \\citep{Pappa01,  Xia02, Panessa02,  Georgantopoulos03a,   Wolter05,  Gliozzi07,  Bianchi08a, Brightman08}.  In  the literature, they  are defined as type  2 AGNs with $N_{\\rm H}$ $<$ 10$^{22}$ cm$^{-2}$.  This sample of unusual AGNs may offer the best means to  refine the unified model and may identify exceptions  to the model  that would  modify our  overall view  of the subject.  Various explanations  have been proposed  to understand the  nature of this  unusual  type of  AGN:  (1)  Compton-thick  behavior: The  X-ray spectra of  Compton-thick objects in the  energy range $<$  10 keV are dominated by the  reflection of the nuclear emission  off the far side of the torus and the  HI column densities inferred from X-ray spectral fits are  similar to those  for unobscured objects; (2)  Dilution: The BELs  in these  objects are  overwhelmed  by host  galaxy light.   (3) Variability:  Some  AGNs  change  their optical  classification  on  a timescale of years to decades as a result of the location of obscuring material on pc scales.  Such objects might be classified as unobscured (in the X-ray) type 2 objects by chance, if the X-ray and optical data are  obtained at  different times.   (4) S/N  effect: The  BEL  can be hidden in a low-S/N optical spectrum while the Compton-thick signature can be  lost in a  low S/N X-ray  spectrum; (5) Exceptions to  the AGN unified model:  in contrast to  the above speculations that  still fit the  unified model,  these  X-ray-unobscured type  2  AGNs may  harbor genuinely  weak BEL  regions, and  thus  are not  the simple  obscured version of type 1 AGNs expected from the unified model.   In this paper,  we will first combine {\\it Spitzer}  IR data with data in the literature to update  the current list of X-ray-unobscured type 2 AGN candidates.  We then review the members of this list to identify objects  that can be  explained by  points 1  - 4  above and  that are therefore  not  true  unobscured  type-2 objects. Of 24 candidates from the literature, less than a third are possible true unobscured  type-2 AGNs. We evaluate these candidates, including using metrics for relative broad line strength, and find that the broad lines in two of them are at least two orders of magnitude weaker than in typical type-1 AGNs. Their IRS/Spitzer spectra indicate that we have  a pole-on view  of their circumnuclear tori,  consistent with the  X-ray  measurements  of  small  absorbing  columns  toward  their nuclei.  Thus,  the  properties  of  these two  AGNs  appear  to  be contradictory  between type-1  levels of  obscuration but  type-2 line properties. We conclude that AGN with intrinsically weak BELs are very rare, but that they do exist.  We   develop   our    arguments   as    follows.    In Section~\\ref{DATA},  we   describe  the  sample   selection  and  data analysis.  In  Section~\\ref{CONT_UNAGN}, we eliminate  contaminants to the sample of X-ray-unobscured type  2 AGNs, including type 1 AGNs and Compton-thick objects.   The IR observations of  the remaining objects are presented  in Section~\\ref{discussion} and placed  in context with observations  at other wavelengths.   Section~\\ref{conclusion} summarizes our conclusions. ",
        "conclusions": "We have presented a multi-wavelength study of unobscured type 2 AGNs. Our conclusions are the following:  (1) We  have found that the  original sample of this  proposed type of AGN  is contaminated by  many objects  with BELs,  as revealed  by our consistent optical  classification and new observations.  One of these objects, RXJ  1737.0+6601, has  a highly polarized  optical continuum. Additional contaminants  include several new  Compton-thick candidates with extremely low nuclear X-ray-to-IR ratios.  (2) We have identified two objects that appear to be true unobscured type-2 AGNs: NGC 3147 and NGC 4594. They have little  X-ray extinction  with $N_{\\rm H}$  $<$  $\\sim$10$^{21}$ cm$^{-2}$.  The  upper-limits  on the  BEL luminosities at a  given nuclear X-ray, IR or  radio luminosity are two orders of magnitude lower in relative BEL strength than the average of  typical type-1 AGNs. Several other galaxies remain as candidates to be weak-BEL AGN.  (3) From the small number of confirmed cases, unobscured Type 2 AGN do exist but they are very rare.  (4) The IR spectra of the unobscured  type 2 AGNs and of many of the candidates show silicate emission features. The presence of the silicate features demonstrates  the existence of dusty tori and that the tori are viewed approximately face-on.  (5) Thus, in contradiction to  the simple unified model, the X-ray and IR  properties indicate that  the nuclei  are viewed  directly without intervening obscuring  material, despite the  intrinsically weak broad emission lines.  (6) The  distributions of the bolometric luminosity  and the accretion rate  of these  objects contradict  some theoretical  studies  but are consistent with the work of \\citet{Nicastro00}."
    },
    "1004/1004.0843_arXiv.txt": {
        "abstract": "We investigate the evolution of Type Ib/c supernova (SN Ib/c) progenitors in close binary systems, using new evolutionary models that include the effects of rotation, with initial masses of 12 -- 25~\\Msun{} for the primary components, and of single helium stars with initial masses of 2.8 -- 20~\\Msun{}.  We find that, despite the impact of tidal interaction on the rotation of primary stars, the amount of angular momentum retained in the core at the presupernova stage in different binary model sequences converge to a value similar to those found in previous single star models. This amount is large enough to produce millisecond pulsars, but too small to produce magnetars or long gamma-ray bursts.  We employ the most up-to-date estimate for the Wolf-Rayet mass loss rate, and its implications for SN Ib/c progenitors are discussed in detail.  In terms of stellar structure, SN Ib/c progenitors in binary systems at solar metallicity are predicted to have a wide range of final masses up to about 7~\\Msun, with helium envelopes of $M_\\mathrm{He} \\simeq 0.16 - 1.5~\\mathrm{M_\\odot}$.  Our results indicate that, if the lack of helium lines in the spectra of SNe Ic were due to small amounts of helium (e.g. $M_\\mathrm{He} \\la 0.5$), the distribution of both initial and final masses of SN Ic progenitors should be bimodal. Furthermore, we find that a thin hydrogen layer ($0.001~\\mathrm{M_\\odot} \\la M_\\mathrm{H} \\la 0.01~\\mathrm{M_\\odot}$) is expected to be present in many SN Ib progenitors at the presupernova stage.  We show that the presence of hydrogen, together with a rather thick helium envelope, can lead to a significant expansion of some SN Ib/c progenitors by the time of supernova explosion.  This may have important consequences for the shock break-out and supernova light curve.  We also argue that some SN progenitors with thin hydrogen layers produced via Case AB/B transfer might be related to Type IIb supernova progenitors with relatively small radii of about $10~\\mathrm{R_\\odot}$. ",
        "introduction": "It is generally believed that Type Ib and Type Ic supernovae result from core collapse events of naked helium stars.  The helium stars are thought to be produced by the loss of the hydrogen envelope, via stellar winds mass loss from massive single stars or via mass transfers in close binary systems.  According to recent stellar models adopting the most up-to-date stellar winds mass loss rates \\citep{Meynet03, Meynet05, Eldridge06, Limongi06, Georgy09}, the final masses ($M_\\mathrm{f}$) of helium stars produced by mass-losing single stars appear to be too massive to produce typical SNe Ib/c (i.e., $M_\\mathrm{f} > 10$~\\Msun{} at solar metallicity).  Although the limiting mass for BH formation is not yet well determined, given their high binding energy, such massive progenitors of $M_\\mathrm{f} > 10$~\\Msun{} are likely to form black holes (BHs), producing faint supernovae or no supernova at all~\\citep[cf.][]{Fryer99}.  Although very bright SNe Ib/c like SN 1998bw could be produced from such  massive helium stars if, for example, powered by rapid rotation \\citep[e.g.,][]{Woosley93, Burrows07}, such events are shown to be rare \\citep[e.g.,][]{Podsiadlowski04, Guetta07}.  By contrast, helium stars with a wide range of  masses ($2.0~\\mathrm{M_\\odot} \\la M_\\mathrm{He} \\la 25~\\mathrm{M_\\odot}$) can be made from $12...60$~\\Msun{} primary components in close binary systems via the so-called Case A/B mass transfer.\\footnote{Case A, B or C mass transfer denotes mass transfer from the primary star during core hydrogen burning, helium core contraction/beginning of core helium burning, or    core helium burning and later stages, respectively. On the other hand, if mass transfer occurs during helium core contraction/beginning of core helium burning from a star that has already undergone Case A mass transfer, such a mass transfer phase is called Case AB. Case ABB or Case BB mass transfer denotes  mass transfer from the primary star during core helium burning and/or later stages, which has already undergone Case AB or Case B mass transfer, respectively.} Many of them may end their life as bright SNe Ib/c leaving neutron stars as remnants, if their final masses are less than about 7 -- 10~\\Msun{}.   Population studies indeed show that close binary stars can produce a sufficient number of SNe Ibc to explain their observed rate, without the need of invoking single star progenitors \\citep[e.g.,][]{Podsiadlowski92, deDonder98, Eldridge08}.  Therefore, it is most likely that the majority of typical SNe Ibc are produced in binary systems.  The observational evidence for the connection between SNe Ibc and long gamma-ray bursts (GRBs) has particularly motivated many observational studies to better understand SNe Ibc since the last decade \\citep[see][for a review]{WB06}. Theoretical stellar models of SNe Ibc progenitors are thus highly required nowadays.  The most comprehensive studies on the detailed characteristics of SNe Ibc progenitors in binary systems were conducted by \\citet{Woosley95} (hereafter, WLW95) using mass-losing pure helium star models, and by \\citet{Wellstein99} (hereafter, WL99) using self-consistent binary star models. Although more recent theoretical studies on  SNe Ibc progenitors in binary systems can be found in the literature, they have been focused on long GRB progenitors or stellar populations, rather than on the detailed nature of typical SNe Ibc progenitors \\citep[e.g.,][]{Brown00, Izzard04, Petrovic05a, Cantiello07, Heuvel07, Detmers08, Eldridge08}.  In this paper, we revisit the problem of SNe Ibc progenitors in close binary systems using both binary star and single helium star models up to the neon burning stage, with updated physics of two important ingredients. One is rotation, which was not considered in WLW95 and WL99, and the other is the mass loss rate of Wolf-Rayet (WR) stars (Sect.~\\ref{sect:review}).  This paper is organized as follows.  In Sect.~\\ref{sect:review}, we  briefly review recent developments of stellar evolution models regarding the effects of rotation and the WR star mass loss rate, arguing for the need of updated physics in binary star models.  Our adopted physical assumptions and numerical method are discussed in Sect.~\\ref{sect:method}.  In the following section (Sec.~\\ref{sect:rotation}), using our binary star evolution models including the effect of rotation and the transport of angular momentum due to hydrodynamic instabilities and magnetic torques, we explore the role of tidal interaction and mass transfer in the redistribution of angular momentum in primary stars.  In Sec.~\\ref{sect:sn}, the nature of SNe Ibc progenitors is investigated in terms of final masses, masses of helium and hydrogen layers, radii and mass loss rates at the presupernova stage, assuming these properties do not significantly change from neon burning to core collapse.  For this purpose, we also present mass-losing single helium star models as a complement to our binary star models, given that the parameter space explored with our binary model sequences is limited.  We conclude the paper by discussing observational implications of our results, in Sect.~\\ref{sect:discussion}. ",
        "conclusions": "\\label{sect:discussion}  We have presented new evolutionary models of massive close binary stars, considering tidal interaction, and transport of angular momentum and chemical species due to rotationally induced hydrodynamic instabilities and the Spruit-Tayler dynamo. We have investigated the redistribution of angular momentum in the primary star. Although mass transfer and tidal interaction can significantly affect the evolution of the rotation velocity of the primary star on the main sequence, the amount of angular momentum retained in the core in the late evolutionary stages is rather insensitive to the previous history of such binary interactions because of the self-regulating nature of the Spruit-Tayler dynamo.  We have also calculated non-rotating, mass-losing single helium star models and compared them with our primary star models in binary systems.  Our models adopt a much lower WR mass loss rate than in the previous studies by  WLW95 and WL99, and predict some new important properties of SN Ibc progenitors accordingly. The following discussions are based on the models with $f_\\mathrm{WR} = 5$, unless otherwise specified.   \\begin{enumerate}  \\item  The final masses of SN Ibc progenitors in binary systems at $Z\\simeq$~\\Zsun{} are not limited to $1.5~\\mathrm{M_\\odot} \\la M_\\mathrm{f} \\la 4~\\mathrm{M_\\odot}$ as predicted by WLW95 and WL99, but a more wide range $M_\\mathrm{f}$ is expected (i.e., $1.5~\\mathrm{M_\\odot} \\la M _\\mathrm{f} \\la 7.1~\\mathrm{M_\\odot}$ from $M_\\mathrm{init} \\simeq 12 ... 60$~\\Msun; see Fig.~\\ref{fig:mimf}).   \\item At $Z \\simeq$~\\Zsun, significant deficiency of helium ($M_\\mathrm{He} < 0.5$~\\Msun{}) is expected for $M_\\mathrm{f}/\\mathrm{M_\\odot} \\ga 5.5$ and for $1.5 \\la M_\\mathrm{f}/\\mathrm{M_\\odot} \\la 2.0$ (Fig.~\\ref{fig:dmhe}).  Rather large amounts of helium up to 1.5~\\Msun{} are expected for the other final mass range (i.e., $2.0 \\la M_\\mathrm{f}/\\mathrm{M_\\odot} \\la$ 5.5~\\Msun{}). At $Z \\simeq$~\\Zsmc, no such helium deficient SN progenitors are expected for the considered initial masses ($16 - 40$~\\Msun).  \\item A thin layer of hydrogen with $M_\\mathrm{H} = 10^{-4} - 10^{-2}$~\\Msun{} is predicted for SN Ibc progenitors with  $ 3.0~\\mathrm{M_\\odot} \\la M_\\mathrm{f} \\la 3.7~\\mathrm{M_\\odot}$ at $Z \\simeq $\\Zsun, and  $ 3.0~\\mathrm{M_\\odot} \\la M_\\mathrm{f} \\la 8 ~\\mathrm{M_\\odot}$ at $Z \\simeq$\\Zsmc, respectively (Fig.~\\ref{fig:dmh}; Table~\\ref{tab1}).  \\item Most SN Ibc progenitors with $M_\\mathrm{f} \\la 5 $~\\Msun{} rapidly expand during core carbon burning, resulting in  $ R =  \\sim 4.0 - 30~\\mathrm{R_\\odot}$ at the presuprnova stage.  This is much larger than found in WLW95 (Fig.~\\ref{fig:radius}). Compact progenitors of $R \\la ~\\mathrm{R_\\odot}$ are only expected for a relatively high mass ($M_\\mathrm{f} \\ga 5.5$~\\Msun{} at $Z\\simeq$~\\Zsun{} and $M_\\mathrm{f} \\ga 10$~\\Msun{} at $Z\\simeq$~\\Zsmc{} ; Fig.~\\ref{fig:radius}). \\end{enumerate}   The above results raise several important issues regarding observational consequences, as discussed below.  \\subsection{Implications for energetic explosions powered by rapid rotation}  Our binary star models show that the mass transfer during helium core contraction (Case AB or Case B) in a close massive binary system cannot remove the hydrogen envelope promptly enough to avoid the core braking due to the Spruit-Tayler dynamo during the helium core contraction phase. Comparison of our binary star models with the single star models by \\citet{Heger05} and \\citet{Yoon06} indicate that the amount of angular momentum retained in the core of the primary star at the presupernova stage should not be much different from those found in single star models if the Spruit-Tayler dynamo is adopted. I.e., a specific angular momentum of a few $10^{14}~\\mathrm{cm^2~s^{-1}}$ in the innermost $\\sim$1.4~\\Msun{} at the presupernova stage is expected in both single and binary stars.  This value is smaller by one or two orders of magnitude than what is necessary to make a long gamma-ray bursts by magnetar or collapsar formation, or very energetic supernovae (hypernovae) powered by rapid rotation and strong magnetic fields ~\\cite[e.g.][]{Burrows07}, although it may suffice to produce millisecond pulsars \\citep{Heger05}. Together with the work by \\citet{Petrovic05a},  our results thus indicate that binary interactions with Case AB/B mass transfers at $Z\\approx \\mathrm{Z_\\odot}$ may not particularly enhance the production of strongly rotation-powered events like long GRBs or hypernovae. This is consistent with the observational evidence that such events are rare compared to normal core collapse events \\citep[e.g.,][]{Podsiadlowski04, Guetta07}. This also confirms the theoretical consensus that other evolutionary paths are needed to produce long GRBs associated with SN Ibc, such as the quasi-chemically homogeneous evolution of a metal poor star \\citep{Yoon05, Yoon06, Woosley06, Cantiello07}, tidal spin-up of a WR star in a very close binary system with a neutron star or black hole companion \\citep[e.g.][]{Izzard04, Heuvel07}\\footnote{A recent study using detailed stellar evolution models by \\citet{Detmers08}, however, seriously questions this possibility.}, or binary evolution with Case C mass transfer with some specific conditions \\citep{Brown00, Podsiadlowski10}.  \\subsection{Progenitor size}  Larger radii of our SN Ibc progenitor models than those previously found should have consequences in shock break-outs and bolometric light curves.  For instance, a shock break-out from a larger envelope would be marked by a lower photosperic temperature . Detailed comparison of numerical calculations with observational data may thus give strong constraints on SNe Ibc progenitor properties \\citep[e.g.,][]{Calzavara04}.  Recent discovery of the X-ray outburst with SN 2008D by \\citet{Soderberg08} indeed suggests the usefulness of such a study for the probe of supernova progenitors \\citep[e.g.][]{Soderberg08, Chevalier08, Xu08, Modjaz08}, for which our new models would provide ideal input.  We will address this issue in a forthcoming paper.    \\subsection{Presence of helium and implications for SNe Ic}\\label{sect:snic}   \\begin{figure} \\epsscale{1.} \\plotone{fig14.eps} \\caption{ The predicted final masses of the primary stars in massive close binaries that undergo Case B mass transfer, as a function of the zero-age main sequence (ZAMS) mass, based on our binary and helium star models with $f_\\mathrm{WR} = 5$.  The absicssa is given in log scale.  The numbers right above the abscissa denote the initial masses of the helium stars that the primary stars of the corresponding ZAMS masses would produce.  The expected final outcomes of the primary stars according to different ZAMS masses are given by the labels right below the top: white dwarf (WD), type Ic supernova (Ic), type Ib supernova with a thin hydrogen layer (Ib with H) or type IIb supernova (IIb), and type Ib supernova without hydrogen (Ib without H). Here we assumed $M_\\mathrm{He} < 0.5~\\mathrm{M_\\odot}$ for SNe Ic progenitors.  Note that the each boundary would shift to a higher value of $M_\\mathrm{ZAMS}$ for close binary systems with Case A mass transfer. }\\label{fig:mimf} \\end{figure}  \\begin{figure} \\epsscale{1.} \\plotone{fig15.eps} \\caption{ Same as in Fig.~\\ref{fig:mimf}, but with $f_\\mathrm{WR}=10$.} \\label{fig:mimf2} \\end{figure}    Although the weak signature or no evidence of helium in SNe Ic spectra may indicate the deficiency of helium in their progenitors, it is not well known how much helium can be hidden in the supernova spectra.  It may also depend on the degree of mixing of nickel into helium rich layers \\citep{WE95}.  If we assume 0.5~\\Msun{}, for instance, as the maximum amount of helium allowed for hiding helium lines in SN spectra, our models indicate that most SNe Ic progenitors at solar metallicity should belong to two distinct classes in terms of both ZAMS and final masses, as summarized in Fig.~\\ref{fig:mimf} for the binary systems that undergo Case B mass transfer\\footnote{For the systems with Case A mass transfer, the each boundary should move to the right in the figure, but the parameter space explored with our model grid is not large enough to determine it quantitatively}. The same conclusion was also drawn by WL99 and \\citet{Pols02}. But Pols \\& Dewi considered different types of binary systems (see below), and the finding of \"two mass classes\" for the final masses  was not obvious in WL99 due to the very high WR rate adopted in their study, although it was clearly seen for the ZAMS masses.  If we assume that primary stars of $12.5 \\la M_\\mathrm{ZAMS}/\\mathrm{M_\\odot} \\la 13.5$ produce low-mass-class SNe Ic via Case BB mass transfer, about 62\\% of the SNe Ic from Case B systems should belong to the high mass class and the rest ($\\sim$38\\%) to the low mass class, at $Z \\approx \\mathrm{Z_\\odot}$ (see Fig.~\\ref{fig:mimf}).  It is important to note that the two classes are produced by different mechanisms.  The high mass class of SNe Ic progenitors (i.e., $M_\\mathrm{ZAMS} \\ga 33$~\\Msun{} in Fig~\\ref{fig:mimf}) is a consequence of WR winds mass loss, while the low mass class results from Case ABB/BB mass transfer  as discussed in Sect.~\\ref{sect:masses}.  The ZAMS mass range for the low mass class may not be much affected by metallicity, while it should be widened with increasing metallicity for the high mass class.  This leads to the conclusion that the low and high mass classes would dominate at low and high metallicities respectively. It should also be noted that the final mass range for the low mass class may not change much for different metallicities, while it may decrease with increasing metallicity for the high mass class, due to the increasing WR winds mass loss rates, as implied by the result of WL99.  In the massive close binary systems considered in this paper, the primary star masses become much smaller than those of the secondary stars when Case ABB or Case BB mass transfer begins. If the companion star mass were lower than the helium star in a close binary system, the mass transfer rate should become higher than in the systems of the present study.  For example, if a helium star is located in a very short period binary system ($P \\la 1~\\mathrm{day}$) with a less massive companion (e.g. a neutron star),  mass transfer from the helium star may occur rapidly enough to make a helium-deficient carbon star, even for $M_\\mathrm{He,i} \\approx$ 6.0~\\Msun{} as shown by \\citet{Pols02}, \\citet{Dewi02} and \\citet{Ivanova03}.  The final masses of such SN Ic progenitors may range from 1.5~\\Msun{} to 3.0~\\Msun{}.  This scenario was also suggested by \\citet{Nomoto94} to explain the fast light curve of Type Ic SN 1994I.  The ZAMS mass of such a SN Ic should be in the range of 12 -- 20~\\Msun{} \\citep{Pols02}. However, such close helium star plus neutron star systems are supposed to rarely form, and might not contribute much to the population of SNe Ic, compared to the systems considered in this study.  Our results should have several observational consequences.  As mentioned above,  the population of SNe Ic should be dominated by the high mass class for $Z \\ga \\mathrm{Z_\\odot}$. They would have higher ZAMS and final masses than those of typical SNe Ib progenitors \\citep[cf.][]{Kelly08, Anderson08}. Given that the parameter space for the high mass class SNe Ic may become larger with a higher WR mass loss rate, the number ratio of SNe Ic to SNe Ib, and that of high mass class SNe Ic to low mass class SNe Ic should increase with increasing metallicity \\citep[cf.][]{Prieto08, Anderson09, Boissier09}. The existence of the two mass classes of SNe Ic progenitors may be related to some aspects of the observational diversity of SNe Ic.  For example, SNe Ic of the low mass class is likely to be characterized by rather fast declining light curves and low luminosities \\citep[cf.][]{Iwamoto94, Richardson06, Young09},  implying that the observed population of SNe Ic is likely to be baised to high-mass-class SNe Ic.  We should also note the huge difference of the binding energy between the two classes.  The binding energy of the envelope above 1.4~\\Msun{} in the SN Ic progenitor star model of Seq.~3 ($M_\\mathrm{f} = 1.64$~\\Msun{}) is only about $10^49$~erg, while it should be one or two orders of magnitude higher for a SN Ic progenitor with $M_\\mathrm{f} \\ga 5.5$~\\Msun{} (see Table~\\ref{tab2}). As a consequence, the energetics of SNe Ic might be systematically different for the two different classes.   On the other hand, the assumption of $M_\\mathrm{He} < 0.5$~\\Msun{} for SN Ic progenitors leads to a ratio of type Ic to Ib supernova rate (Ic/Ib ratio) of about 0.4 from binary systems at solar metallicity \\footnote{ At SMC metallicity, the SN Ibc progenitor models of initial masses of 16 -- 40~\\Msun{} have $3.9\\la M_\\mathrm{f} \\la 12$~\\Msun{} with $M_\\mathrm{He} \\ga 1.2$~\\Msun{} (see Table~\\ref{tab1}).  This implies that the low-mass class SNe Ic would predominantly occur in binary systems at this metallicity.  If we assume stars with $12.5 \\la M_\\mathrm{init} \\la 13.5$ would produce low-mass class SNe Ic as in the case of solar metallicity,  the SN Ic/SN Ib ratio would be about 0.1 at SMC metallicity. The exploration of the exact mass range for the low-mass class SNe Ic is a time-consuming task, and we plan to investigate this in near future. }. This appears in contradiction with recent observations that indicate rather a high Ic/Ib ratio of about 2.0~\\citep[e.g.][]{Smartt09}.  This discrepancy would become even larger with $f_\\mathrm{WR} =10$, as implied by Fig.~\\ref{fig:mimf2}. This raises a question on the nature of SN Ic progenitors, and it should be kept in mind that we still do not fully understand what distinguishes SN Ic progenitors from those of SN Ib.  A recent work by \\citet{Dessart10} indicates that the mass fraction of helium in the outermost layers ($Y_\\mathrm{s}$), rather than the total mass of helium, may be more relevant for the presence of helium lines in supernova spectra. Specifically, it is shown that if helium is well mixed with CO material such that $Y_\\mathrm{s}$ becomes less than about 0.5, helium lines are not seen in early time spectra, despite rather a large total amount of helium ($M_\\mathrm{He} \\simeq 1.0$~\\Msun), if non-thermal excitation is absent. In our progenitor models, such a small $Y_\\mathrm{s}$ is realized only for $M_\\mathrm{f} \\ga$ 5.5~\\Msun{} at solar metallicity (with $f_\\mathrm{WR} = 5$).  This is not different from the above-discussed mass limit for having $M_\\mathrm{He} \\la 0.5$~\\Msun, implying that the initial mass range for SN Ic progenitors would not change much even if we adopted $Y_\\mathrm{s}$ as a criterion, at least for the high mass class. On the other hand, we have $Y_\\mathrm{s} = 0.98$ in the primary star of Seq.~3 at carbon exhaustion while the total mass of helium is less than 0.2~\\Msun. This imples that the initial mass range for the low mass class SN Ic progenitors might be affected if the condition of $Y_\\mathrm{s} < 0.5$ for SN Ic progenitors were applied.  But \\citet{Dessart10} did not yet calculate such a low mass SN progenitor model ($M_\\mathrm{f} \\la 2.0$~\\Msun), and their analyses were limited to early times of supernovae. It remains to be an important subject of future work to systematically investigate which types of supernova progenitors would lead to the presence or absence of helium lines in the supernova spectra at different epochs, including the effect of non-thermal excitation. Therefore, the above discussion on Type Ic progenitors based on the total amount of helium should only be considered indicative at this stage.   \\subsection{Presence of hydrogen}\\label{sect:dischydrogen}  It is interesting that, at $Z \\approx \\mathrm{Z_\\odot}$, the presence of a thin hydrogen layer is only expected for a limited range of the initial/final mass of SN Ib progenitors, as shown in Figs~\\ref{fig:dmh}, and~\\ref{fig:mimf}. The detection of hydrogen absorption lines at high velocity has been indeed reported in many SNe Ib \\citep[e.g.,][]{Deng00, Branch02, Elmhamdi06},  in favor of our model prediction for the presence of a thin hydrogen layer in SNe Ib progenitors.  This might provide a strong constraint for the progenitor masses of observed SNe Ib, in principle.  Note also that explosions of such helium stars with thin hydyrogen layers could be recognized as SN IIb rather than Ib, if hydrogen lines were detected short after supernova explosion, e.g., as in the case of SN 2008ax~\\citep{Chornock10} and as recently discussed by \\citet{Baron10}   and \\citet{Dessart10}. The radii of these progenitor models range from $\\sim 10^{11}~\\mathrm{cm}$ to $\\sim 10^{12}~\\mathrm{cm}$. They may corredpond to the ''compact'' category of SN IIb progenitors, which is discussed in \\citet{Chevalier10}. The relatively low ejecta masses of such SNe IIb are consitent with our model predictions.  On the other hand, Case~C mass transfer can also leave helium cores covered  with small amounts of hydrogen envelope. As the life time of such stars made via Case C mass transfer should be rather short, they can retain much more hydrogen ($M_\\mathrm{H} > 0.1$~\\Msun), than what is predicted from our binary models with Case AB/B mass transfer. Such a star may eventually explode as a SN IIb like SN 1993J \\citep[e.g.][]{Podsiadlowski93, Maund04}, with a much extended envelope ($\\sim 10^{13} - 10^{14}~\\mathrm{cm}$).  Therefore, the two categories of SN IIb progenitors according to their sizes, which has been recently suggested by \\citet{Chevalier10}, may be understood within the framework of binary evolution; SNe IIb of the compact type  may be produced via Case AB/B mass transfer (especially at $Z \\la \\mathrm{Z_\\odot}$), and SNe IIb of the extended type via Case C mass transfer."
    },
    "1004/1004.2846_arXiv.txt": {
        "abstract": "\\noindent An algorithm for creating synthetic telescope images of Smoothed Particle Hydrodynamics (SPH) density fields is presented, which utilises the adaptive nature of the SPH formalism in full.  The imaging process uses Monte Carlo Radiative Transfer (MCRT) methods to model the scattering and absorption of photon packets in the density field, which then exit the system and are captured on a pixelated image plane, creating a 2D image (or a 3D datacube, if the photons are also binned by their wavelength).  The algorithm is implemented on the density field directly: no gridding of the field is required, allowing the density field to be described to an identical level of accuracy as the simulations that generated it.  Some applications of the method to star and planet formation simulations are presented to illustrate the advantages of this new technique, and suggestions as to how this framework could support a Radiative Equilibrium algorithm are also given as an indication for future development. ",
        "introduction": "\\noindent The physics of dust is of paramount importance when modelling star and planet formation: dust signatures are seen in both molecular clouds and circumstellar discs, and are a crucial component in their observed properties over a wide range of wavelengths.  Indeed, the existence of dust is essential if planets are to be formed inside these discs.  The presence of dust has important effects on the radiative signatures that can be detected by astronomers: dust will scatter and polarise at short wavelengths, as well as reprocessing this radiation to longer wavelengths.  The nature of the scattering and polarisation will depend on the geometry of both the system and the dust in the system, as well as the physical properties of the dust itself (composition, morphology, grain size, etc).  In recent years, telescopes/networks such as HST, SCUBA, MERLIN and Spitzer have efficiently probed stellar systems from the IR to the radio, with future ground and space-based missions such as Herschel, ALMA, SCUBA II, JWST, SPICA and e-MERLIN improving the quality of this data.  Astronomers will soon be faced with a wealth of new, high-fidelity astronomical data on stellar objects across a wide spectral range, allowing detailed studies of the evolution of dusty disc systems, in particular the interplay between disc dust and disc gas.  For numerical simulations to inform these observations, the simulation of radiative transfer (RT) in these systems must be pursued, so that imaging of numerical simulations can provide theoretical insights to observations.  Monte Carlo Radiative Transfer (MCRT) has become a popular tool in simulating and imaging astrophysical systems \\citep{Whitney_and_Hartmann_92,Wood_et_al_96,MCFOST}.  The technique involves the tracking of \\emph{photon packets} in the medium, allowing them to scatter and absorb in the medium (as well as attaining a non-zero polarisation).  The photons are tracked until they escape the medium, and can then be captured on an image plane.  Traditionally, the method requires the density field to be defined to as high an accuracy as possible - this is usually achieved by gridding the density field in 3 dimensions.  Smoothed Particle Hydrodynamics (SPH) is a Lagrangian method which represents a fluid by a particle distribution \\citep{Lucy,Gingold_Monaghan}.  Each particle is assigned a mass, position, internal energy and velocity.  From this, state variables such as density can then be calculated at any position in the system by interpolation - see reviews by \\citet{Monaghan_92, Monaghan_05}. It has been successful in modelling astrophysical systems of various scales and geometries, from protostellar systems to galactic systems and beyond.  Its key advantage is the ability to follow the change of density through many orders of magnitude adaptively.  Gravity can be included in SPH calculations, and optimised using traditional hierarchical tree methods \\citep{Hernquist_and_Katz_89}. The formalism can also be modified to simulate magneto-hydrodynamics (MHD) \\citep{Hosking_Whitworth_04,Price_Monaghan_SPMHD2,Price_Monaghan_SPMHD3,Price_and_Bate_07}.  Radiative physics is crucial to the simulation of any astrophysical fluid.  Radiative transfer in SPH has a long history, with efforts ranging from simple parametrisations for the radiative cooling time (e.g. \\citealt{Ken_1}), optical-depth dependent radiative cooling (e.g. \\citealt{Stam_2007}), through to flux-limited diffusion models (e.g. \\citealt{WB_1,Bastien_diffusion,Viau_et_al_06,WB_2, Mayer_et_al_07, intro_hybrid}).  While these algorithms are extremely well-suited to calculating gas temperatures during runtime, they lack the precision and insight offered by radiative transfer techniques which do not average over frequency, such as MCRT.  MCRT techniques can be applied to the SPH density field, and have been on several occasions.  Early attempts began by binning the particle distribution onto a grid - however, the choice of geometry strongly influences the final gridded field, and often adaptive meshes are required to correctly represent the matter distribution \\citep{Oxley_Woolfson_2003,Kurosawa_et_al_04,Stam_MCRT,Bisbas_et_al_09}.  Later efforts have utilised ray-tracing directly in SPH fields \\citep{Kessel_Deynet_and_Burkert_00,SPHRAY,TRAPHIC}, allowing the full power of the SPH formalism to be applied to the calculation of optical depths. However, these methods currently only calculate optical depths along the ray for the purpose of photoionisation, etc, and do not account for detailed scattering and polarisation, which are important in imaging small-scale systems such as protoplanetary discs.  This paper introduces an algorithm for imaging SPH simulations of star and planet formation directly using MCRT, without requiring gridding, and including scattering and polarisation.  It utilises the same techniques that a traditional gridded MCRT code uses, with the advantage that it can trace rays in the density field with the same adaptive capability as SPH, such that it can model radiative effects with at least the same resolving power.  The paper is organised as follows: section \\ref{sec:method} will outline the algorithm, sections \\ref{sec:applications} and 4 will describe some applications of the technique to imaging numerical simulations, and section \\ref{sec:conclusions} will summarise the results of this work. ",
        "conclusions": "\\label{sec:conclusions}  \\noindent This paper has outlined a means for applying Monte Carlo Radiative Transfer (MCRT) techniques directly to Smoothed Particle Hydrodynamics (SPH) density fields of star and planet formation.  In doing so, it gives a means for synthetic telescope images to be made, allowing the flexibility of SPH simulations to be retained in calculating optical depths, scattering and polarisation.  This work has shown two applications of the method to SPH simulations of protostellar discs: non-trivial features in the imaging of these environments are well-traced, and the resulting images give new opportunities for theoretical input to observed discs.  It can be shown what resolutions and sensitivities are required to observe these features.  ALMA appears to be able to resolve and image perturbed structures like tidal arms in discs undergoing close encounters.  But, while it has sufficient resolution, it cannot image unperturbed spiral structure (or the shadows it casts) without an improvement in sensitivity by at least a factor of 100.  Its inherently graphical nature allows the code to be greatly optimised, whether by standard parallelisation methods or by the use of Graphical Processing Units (GPUs) to handle the ray intersections.  While used for imaging in this work, the method is not restricted to this alone.  As the algorithm only modifies how the optical depth is calculated (in order to work in SPH fields), it can be applied to any circumstances traditional gridded MCRT has been applied to.  This includes radiative equilibrium simulations where the temperature structure is calculated from stellar emission, or moving beyond continuum emission to perform line radiative transfer calculations.  In effect this places SPH on a par with grids or meshes for MCRT techniques."
    },
    "1004/1004.4889_arXiv.txt": {
        "abstract": "\\baselineskip 11pt We present an analytical model which reproduces measured galaxy number counts from surveys in the wavelength range of 500 $\\mu$m to 2 mm. The model involves a single high-redshift galaxy population with a Schechter luminosity function which has been gravitationally lensed by galaxy clusters in the mass range $10^{13}$ to $10^{15}\\Msun$. This simple model reproduces both the low flux and the high flux end of the number counts reported by the BLAST, SCUBA, AzTEC and the SPT surveys. In particular, our model accounts for the most luminous galaxies detected by SPT as the result of high magnifications by galaxy clusters (magnification factors of 10-30). This interpretation implies that submillimeter and millimeter surveys of this population may prove to be a useful addition to ongoing cluster detection surveys. The model also implies that the bulk of submillimeter galaxies detected at wavelengths larger than 500 $\\mu$m lie at redshifts greater than 2. ",
        "introduction": "\\label{sec:intro}  Over the last decade, submillimeter (submm) surveys have yielded significant advances in our understanding of the galaxy population responsible for the high-redshift component of the cosmic infrared background (CIB) \\citep{Fixetal96,SmaIviBla97,Hugetal98,Baretal98,Dweetal98,Fixetal98,Greetal04,Popetal06,Copetal06,Devetal09}. With typical far-infrared (FIR) luminosities $> 10^{12} \\Lsun$, submm galaxies are presumed to be the high-redshift counterparts to (ultra) luminous infrared galaxies (LIRGs, ULIRGs). The high luminosity of these galaxies is the result of star formation rates of 100--1000 $\\Msun$\\,yr$^{-1}$. Approximately half of these galaxies are located at $1.9 \\lesssim z \\lesssim 2.9$ \\citep{Chaetal05,Areetal07}, dominating the total star formation rate at this epoch \\citep{Peretal05, Micetal09}.  One way to express the results of submm surveys is through number counts of galaxies as a function of flux for each observed wavelength.  The shape of these counts has been interpreted as arising from different populations of galaxies whose characteristics evolve over cosmic time \\citep{Lagetal03,Lagetal04,Peaetal09,LeBetal09}. These empirical models have successfully reproduced the counts. However, they may be masking a simpler explanation for the departure of the counts from a Schechter distribution at the high-flux end: magnification due to high redshift galaxy clusters and groups \\citep[e.g.][]{Bla96,Peretal02,Negetal07}.  Millimeter wavelength surveys have also aimed at detecting galaxy clusters via the Sunyaev Zel'dovich (SZ) effect \\citep{Hinetal08, Caretal09}; the first results, including CMB power spectra and cluster catalogs, have been released recently \\citep{Fowetal10, Staetal09, Vanetal10}.  The number of detected clusters remains relatively low, primarily due to the low value of $\\sigma _8$.  However, other effects could be reducing the sensitivity of the surveys \\citep[e.g.][]{LimJaiDev09, Limetalinprep}.  The South Pole Telescope (SPT) has measured number counts of dusty galaxies at wavelengths $\\lambda=1.4$ mm and $2.0$ mm over an area of $87$ deg$^2$ \\citep{Vieetal09}. The observed numbers at the bright end % are higher than expected: these galaxies are either at high redshifts and intrinsically exceptionally luminous, or have been magnified by gravitational lensing, or are simply at much lower redshifts than the bulk of the population of submm galaxies. The latter possibility is disfavored by the lack of detected counterparts in other surveys that probed the low redshift population \\citep{Vieetal09}. The possibility that these galaxies are at high redshifts and intrinsically bright would require them to be far more luminous than an underlying Schechter-like luminosity function would permit. Thus the favored explanation is that they have typical luminosities for high-$z$ galaxies, but have been magnified by foreground galaxies or clusters. In fact, lensing of high-redshift background submm galaxies has been observed in a number of systems \\citep{SmaIviBla97,Smaetal02,Wiletal08,Rexetal09,Gonetal09,Swietal10}.  In this {\\it Letter}, we explore the possibility that the existing observed galaxy number counts over a wide range of wavelengths can be reproduced by a {\\it single} population of galaxies at high-redshift.  Foreground galaxy groups and clusters gravitationally lenses the background submm population \\citep{LimJaiDev09} and leads to significant enhancements of the high-flux end to the galaxy counts. In \\S~\\ref{sec:lensing} we describe the lensing magnification formalism, which we then apply to a high-$z$ galaxy population and present results in \\S~\\ref{sec:results}. We discuss implications for high-$z$ galaxies and the cluster searches in \\S~\\ref{sec:discussion}.  \\begin{figure} \\resizebox{88mm}{!}{\\includegraphics[angle=0]{dNdS_zs_SPT_2000_dustnoiras.eps}} \\caption{Intrinsic and lensed number counts $dn/dS$ for a Schechter function describing galaxies at different redshifts. Also shown are the observed counts for SPT {\\it dusty} submm galaxies at $\\lambda=2.0$~mm, after removal of low-redshift galaxies with IRAS counterparts. } \\label{fig:dndS_SPT_lens} \\end{figure}  Throughout, we use a fiducial cosmology for a flat universe with parameter values based on the results of the Wilkinson Microwave Anisotropy Probe fifth year data release \\citep[WMAP5,][]{Kometal09}. % The cosmological parameters (and their values) are the normalization of the initial curvature spectrum $\\delta_\\zeta (=2.41\\times 10^{-4})$ at $k=0.02$ Mpc$^{-1}$ (corresponding to $\\sigma_8=0.8$), its tilt $n (=0.96)$, the baryon density relative to critical $\\Omega_bh^2 (=0.023)$, the matter density $\\Omega_{\\rm m} h^2 (=0.13)$, and two dark energy parameters: its density $\\Omega_{\\rm DE} (=0.74)$ and equation of state $w(=-1)$, which we assume to be constant. Since lensing effects depend on the abundance of dark matter halos, which is exponentially sensitive to linear matter perturbations, we also consider changes in $\\sigma_8$ consistent with the WMAP5 errors of $\\Delta \\sigma_8 \\approx 0.03$. Our central value and uncertainty for $\\sigma_8$ is  consistent with the WMAP7 analysis \\citep{Kometal10}. ",
        "conclusions": "\\label{sec:discussion}  We have considered the possibility that bright millimeter and submillimeter galaxy counts arise largely from a galaxy population at high redshifts that is lensed by intervening galaxy groups and clusters.  Our model predictions match the counts from $500~\\mu$m to $2~$mm from the BLAST, SCUBA, AzTEC and SPT surveys (see Fig.~\\ref{fig:dndS_all}). We find that the high flux SPT number counts can be explained by highly magnified galaxies from this high-$z$ population (once the known, low-$z$ counterparts detected in IRAS are removed). This high-$z$ galaxy population is described by a Schechter luminosity function with $L^* \\sim 2.5 \\times 10^{12} L_\\odot$ and a source redshift $z_s=3.0$.  Our model predictions fit the data for  $500~\\mu$m~$< \\lambda < 2~$mm by varying $S^*$ with wavelength within the range of typical submillimeter galaxy SEDs (see Fig.~\\ref{fig:Sstar}).  Our model has some simplifying assumptions, such as the fixed source redshift $z_s=3$, and there are significant measurement and theoretical uncertainties. A complete analysis would require a more detailed treatment of the lensing and inclusion of the measured galaxy clustering. In addition, it has been established that at the shorter wavelengths probed by BLAST, an increasing fraction of sources lie at low-$z$ -- hence we can expect some smooth variation in the fraction of low-$z$ galaxies with observed wavelength.  Nevertheless our results in Fig.~\\ref{fig:dndS_all} imply that current number counts do not require an additional galaxy population to explain the high flux measurements. Such a population has been invoked in theoretical models \\citep[e.g.][]{Lagetal03} and suggested as a possible explanation (as well as lensing of the high-$z$ population) for the recent SPT measurements \\citep{Vieetal09}. Indeed, adding a significant fraction of a second population could cause the predictions to exceed the measured counts once magnification effects for the high-$z$ population are included. It would also be difficult to explain how the scaling with wavelength of the high flux number counts is the same as the lower flux counts if they came from different galaxy populations.  Since lensing does not depend on frequency, our model naturally follows this common scaling.  Current SZ surveys with sensitivities of $3-7 \\times 10^{14} h^{-1}\\Msun$ cannot detect most halos that produce this lensing contribution. Conversely, looking for extremely bright objects in millimeter and submillimeter wavelengths provides a way to find high-$z$ lensing halos associated with galaxy groups and clusters. Note that the number of halos that can be found in this way is only a small fraction of all  halos within a given mass range. As discussed above, our model most likely underestimates the contribution from halos with $M\\simlt 10^{13}\\Msun$. It is therefore of great interest to investigate the lenses corresponding to the bright sources in current data.  Follow up observations with optical telescopes should be able to identify lensing group/cluster candidates up to $z\\simeq 1$. These clusters and groups host Brightest Cluster Galaxies with luminosities $L\\sim 10^{11}-10^{12} L_\\odot$,  based on the low-$z$ results of \\cite{Johetal07} from the  Sloan Digital Sky Survey. Multi-band optical imaging with limiting magnitude of $23-24$ (for the $r$ band) would enable identification of these groups and clusters. Conversely, targeted observations of the high magnification regions of known strong lensing clusters could provide detections of faint submm galaxies (which would lie below the detection threshold without the magnification boost). Similar to the use of clusters as gravitational telescopes in optical imaging, this may also help resolve submm galaxies. Planned observations with AzTEC and the Large Millimeter Telescope have considered such an approach (David Hughes, private communication)."
    },
    "1004/1004.1422.txt": {
        "abstract": "We present high spatial resolution $Chandra$ X-ray images of the NGC 2237 young stellar cluster on the periphery of the Rosette Nebula.  We detect 168 X-ray sources, 80\\% of which have stellar counterparts in USNO, 2MASS, and deep FLAMINGOS images.  These constitute the first census of the cluster members with $0.2 \\la M \\la 2$~M$_\\odot$. Star locations in near-infrared color-magnitude diagrams indicate a cluster age around 2~Myr with a visual extinction of $1\\la A_V \\la 3$ at 1.4 kpc, the distance of the Rosette Nebula's main cluster NGC 2244.  We derive the K-band luminosity function and the X-ray luminosity function of the cluster, which indicate a population $\\sim$400--600 stars. The X-ray-selected sample shows a $K$-excess disk frequency of 13\\%. The young Class II counterparts are aligned in an arc $\\sim 3$~pc long suggestive of a triggered formation process induced by the O stars in NGC~2244. The diskless Class III sources are more dispersed. Several X-ray emitting stars are located inside the molecular cloud and around gaseous pillars projecting from the cloud. These stars, together with a previously unreported optical outflow originating inside the cloud, indicate that star formation is continuing at a low level and the cluster is still growing.  This X-ray view of young stars on the western side of the Rosette Nebula complements our earlier studies of the central cluster NGC~2244 and the embedded clusters on the eastern side of the Nebula.  The large scale distribution of the clusters and molecular material is consistent with a scenario in which the rich central NGC 2244 cluster formed first, and its expanding HII region triggered the formation of the now-unobscured satellite clusters RMC XA and NGC 2237.  A large swept-up shell material around the HII region is now in a second phase of collect-and-collapse fragmentation, leading to the recent formation of subclusters.  Other clusters deeper in the molecular cloud appear unaffected by the Rosette Nebula expansion. ",
        "introduction": "The triggered formation of the lower mass stars in the vicinity of massive stars is a complex process that is only now being characterized in detail \\citep[see][for a recent review]{Briceno07}. In their immediate neighborhood, massive stars suppress further star formation by quickly ionizing and dispersing surrounding molecular material \\citep{Herbig62}. At greater distances, OB stars are more constructive to star formation activity; the shocks driven by ionization or stellar winds are crucial in triggering the collapse of molecular cores \\citep{Whitworth94,Lefloch94}.  Triggered star formation events by massive stars have been observed at different spatial scales, for example, small bright-rimmed clouds on the periphery of HII regions \\citep{Sugitani95, Getman07,Ogura07}, an embedded cluster in a molecular cloud core on the edge of the dispersed Cep~OB3b \\citep{Getman06}, multiple generations of star formation in W5 \\citep{Koenig08}, a broad ridge of young stars along the southwestern boundary of M~17 \\citep{Jiang02,Broos07}, and a rich secondary cluster on the edge of Sharpless~219 \\citep{Deharveng06}.  The Rosette star forming complex has been considered an excellent candidate for triggered star formation \\citep{Cox90, PL97} following the framework developed by \\citet{Elmegreen77}.  The massive young cluster NGC 2244 \\citep[$d\\sim 1.4$ kpc, $t\\sim 2$ Myr;][]{Hensberge00} powers a visually spectacular expanding HII region known as the Rosette Nebula.  The ionized nebula is clearly interacting with the adjacent Rosette Molecular Cloud (RMC) to the east of the Nebula which has a collection of embedded young stellar clusters, each with a few hundred pre-main sequence stars (Table 6 in Wang et al. 2009; see also Poulton et al. 2008).  However, recent investigations of the ages and disk fractions of these embedded stellar populations at near-infrared (NIR), mid-infrared, and X-ray wavelengths do not obviously support a sequential, triggered origin \\citep{RomanZuniga08,Poulton08,Wang08a}. Instead, Wang et al.\\ suggest that the previously unnoticed, more evolved, and less-obscured RMC~XA cluster lying between NGC~2244 and the molecular cloud may have been triggered by the expanding Rosette HII region $\\sim 2$~Myr in the past.  To the west of NGC~2244 cluster lies the NGC~2237 cluster\\footnote{The designations of clusters and nebulae in the Rosette region have been confused since the 19th century.  NGC~2237 refers both to the stellar cluster and its associated nebulosity on the west side of the Rosette nebula \\citep{Sulentic73}.  The star cluster is listed as Ocl 512 in the catalogue of \\citet{Alter70} and included in the study of the Monoceros star clusters by \\citet{Perez91}.  NGC~2238 and 2246 are old designations for bright emission regions around the Rosette Nebula; these names are rarely used today.  The well-studied rich star cluster exciting the nebula is NGC~2244; NGC~2239 is an obsolete designation for this central cluster.  The listing in the SIMBAD database for NGC~2237 was incorrect until recently.  Note that the open cluster NGC 2239 studied in \\citet{Bonatto09} in fact refers to NGC 2237.  A thorough review of the Rosette star formation complex is provided by \\citet{RomanZuniga08b}.}.  NGC~2237 was not studied until the recent 2MASS and FLAMINGOS $JHK$ surveys \\citep{Li05, RomanZuniga08, Bonatto09}. Li (2005) estimates that the cluster has a diameter of 3.2~pc with a population of 232 stars. \\citet{Bonatto09} studied NGC 2237 (called NGC 2239) with field-star-decontaminated 2MASS photometry, and suggested that it may be a young ($5\\pm 4$~Myr) cluster located in the background of NGC 2244 ($d=3.9$~kpc). In the Digital Sky Survey optical plate (Figure~\\ref{fig:dss}), the NGC~2237 region lies on the boundary between the HII region and the cold neutral molecular material \\citep[Core E;][]{Blitz80,RomanZuniga08b} on the western side of the nebula. Numerous dusty elephant trunks and pillars point towards the O stars in NGC~2244. Its structure contributes to the petal-like morphology of the Rosette Nebula and resembles the pillared boundaries of other HII regions such as M~16 and IC~1396 where the molecular cloud is photoionized and ablated by the O star ultraviolet light and winds.  However, the stellar content of NGC~2237 has not been well-characterized and most of its members have not been individually identified.  The difficulty is that the near-infrared surveys are dominated by older field stars unrelated to the Rosette star forming complex so that only the statistical enhancement in stellar surface density, and the fraction of members with $K$-band excesses, have been studied.  X-ray surveys are more effective at uncovering the full pre-main sequence (PMS) population of young clusters.  Low mass young stars emit X-rays from violent magnetic reconnection flares, orders of magnitude more luminous than seen in older Galactic stars \\citep[see review by][]{Feigelson07}. Hence, an X-ray study can individually identify young stars including the diskless Class III sources missed in the infrared-excess samples.  X-ray studies of NGC 2244 \\citep[][henceforth Paper I]{Wang08a} and the embedded clusters in the RMC \\citep[][Paper II]{Wang08b} on the eastern side of the Rosette Nebula produced a richer census of young stars than available from infrared-excess studies.  A handful of X-ray stars in NGC~2237 have been identified with the {\\it ROSAT} satellite \\citep[see][]{BC02}, but observations with the {\\it Chandra X-ray Observatory} with its high spatial resolution and sensitivity, are much more effective.  In this paper (the third in a series on the Rosette Complex), we focus on the X-ray point sources identified in the image of the NGC~2237 region obtained with the {\\em Chandra X-ray Observatory}.  We identify over 150 individual X-ray emitting young stars and investigate their spatial distribution and stellar properties to study the star formation activity on the western side of the Rosette Nebula. Proximity to dense molecular structures and correlated age-spatial gradients in the stellar distributions are particularly valuable in evaluating the timescale and efficiency of the triggering process that may have occurred. We combine these results on NGC~2237 with earlier work on the central NGC~2244 cluster and the embedded RMC clusters to elucidate a large-scale view of the star formation history in the Rosette complex.  A distance of $d=1.4$~kpc to NGC 2237 is adopted throughout the paper and the uncertainty in distance \\citep[see][]{Bonatto09} will be discussed. ",
        "conclusions": "}  We present the first high spatial resolution X-ray image of the NGC 2237 cluster in the Rosette Nebula obtained via a single {\\it Chandra} observation, which is part of our survey of the star formation activities in the Rosette complex.  Prior to this study, only 36 $K$-band excess cluster members were identified \\citep{RomanZuniga08}. The effectiveness of the X-ray survey technique is clearly demonstrated.  Our main findings are:  1.  In this 20 ks observation, 168 X-ray point sources are detected with a limiting X-ray sensitivity of $L_t\\sim 10^{30}$ ergs s$^{-1}$, nearly complete to the solar mass range.  We associate 134 of the 168 X-ray sources with cluster stars and infer about 24 other embedded sources, giving a total of $\\sim 160$ stellar members identified here in NGC~2237.  2. The locations of most ACIS sources in the NIR color-magnitude diagram are consistent with a small population of PMS low mass stars ($M \\la 2 M_{\\odot}$) in the NGC 2237 region with a visual extinction of $A_V \\sim 2$ at 1.4 kpc, assuming a similar age to NGC 2244 (2 Myr old). Our X-ray sample provides the first probe of the low mass population in this satellite cluster to NGC 2244.  We derive an overall $K$-excess disk frequency of 13\\% for stars with mass $M \\ga 1M_{\\odot}$ using the X-ray-selected sample, consistent with the reported disk fraction from NIR study.  3.  We examine the X-ray properties of a number of B stars in the field. Several X-ray emitting stars are located at the peripheries of pillar objects. A previously unknown optical outflow feature is seen originating inside the optically dark region, which may be a newly identified HH object; this supports the ongoing formation of stars in this region.  4. The X-ray selected population provides a reliable probe of the cluster center and structure, which agree well with previous NIR studies. Similar to RMC XA, collapse of pre-existing massive molecular clumps accompanying the formation of the NGC 2244 cluster may have formed NGC 2237 as a satellite cluster.  There is tantalizing evidence suggestive of a triggered formation process in the optical dark pillars by the NGC 2244 O stars, including discovery of young Class II sources aligned in an arc and embedded hard X-ray sources across the optical Photon Dominated Region.  5.  We derive the KLF and the XLF for the NGC 2237 cluster.  A significant excess of stars with apparent $K_s\\sim 11-12$ is seen in the KLF, likely associated with the presence of a few foreground stars. The slope of the NGC 2237 XLF matches the ONC XLF well, and the relative scaling implies a smaller cluster population $\\sim$400--600 for NGC 2237, which is consistent with the estimates from NIR studies. We discuss the uncertainty in the distance to NGC 2237 and prefer the near distance.  6.  The large-scale star formation in the Rosette Complex is reviewed in the context of the ``collect-and-collapse'' scenario, given the locations of clusters and the morphology of the ISM.  The formation of NGC 2244, RMC XA, and NGC 2237 represents the earliest star formation episode in the complex.  The IRAS 100$\\mu$m image and CO emission map show sites of currently active star formation--four potential clusters associated with the fragments of the large shell of swept-up neutral material from the expanding HII region.  In the RMC, a temporal sequence of star formation across the complex is present \\citep[see ][]{Li05c,Wang08b,RomanZuniga08}.  The cause of the sequence appears to be related to the molecular cloud, perhaps primordial, not limited to the originally proposed triggering scenario in \\citet{Elmegreen77} where one cluster provokes the formation of the next one.  We thank the anonymous referee for his/her careful reading and comments that significantly improved the clarity of this paper.  We thank Travis Rector and Mark Heyer for kindly providing the KPNO MOSAIC images of the Rosette Nebula and the CO emission maps of the Rosette Complex, respectively.  This work was supported by the Chandra ACIS Team (G. Garmire, PI) through NASA contract NAS8-38252.  E.D.F., P.S.B., and L.K.T. also received support from NASA grant NNX09AC74G. FLAMINGOS was designed and constructed by the IR instrumentation group (PI: R. Elston) at the University of Florida, Department of Astronomy with support from NSF grant AST97-31180 and Kitt Peak National Observatory.  The data were collected under the NOAO Survey Program, ``Towards a Complete Near-Infrared Spectroscopic Survey of Giant Molecular Clouds'' (PI: E. Lada) and supported by NSF grants AST97-3367 and AST02-02976 to the University of Florida. E.A.L. also acknowledges support from NASA LTSA NNG05D66G. This publication makes use of data products from the Two Micron All Sky Survey (a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by NASA and NSF), and the SIMBAD database and the VizieR catalogue access tool (operated by the CDS).  %------------------------------------------------------------ %"
    },
    "1004/1004.1712_arXiv.txt": {
        "abstract": "Through analysis of archival images and photometry from the {\\it Spitzer} GLIMPSE and MIPSGAL surveys combined with 2MASS and \\msx\\ data, we have identified 488 candidate young stellar objects (YSOs) in the giant molecular cloud M17 SWex, which extends ${\\sim}50$~pc southwest from the prominent Galactic \\hii region M17. Our sample includes ${>}200$ YSOs with masses ${>}3$~\\Msun\\ that will become B-type stars on the main sequence. Extrapolating over the stellar initial mass function (IMF), we find that M17 SWex contains ${>}1.3\\times10^4$ young stars, representing a proto-OB association. The YSO mass function is significantly steeper than the Salpeter IMF, and early O stars are conspicuously absent from M17 SWex. Assuming M17 SWex will form an OB association with a Salpeter IMF, % these results reveal the combined effects of (1) more rapid circumstellar disk evolution in more massive YSOs and (2) delayed onset of massive star formation. % ",
        "introduction": "\\citet{EL76} discovered an extended molecular cloud associated with the well-known Galactic \\hii region M17. The size ($70~{\\rm pc}\\times 15$~pc at $d=2.1$~kpc), mass \\citep[${>}2\\times 10^5$~\\Msun;][]{ELD79}, and fragmentary morphology of this cloud were reminiscent of the largest Galactic OB associations. \\citet{EL76} suggested that an extended OB association will eventually form near M17, and they noted an apparent progression of ages among the known OB populations, \\hii regions, and molecular cloud cores across the complex. % The M17 complex became an observational touchstone for the theory of sequential triggered massive star formation by propagating ionization fronts \\citep{EL77}.  While the M17 \\hii region itself and the interface between the ionization front and the massive molecular core known as M17 SW have remained the focus of intense study \\citep[][and references therein]{CH08}, interest in the extended molecular cloud, which we name M17 SWex (since M17 SW is ambiguous), has gradually waned. M17 SWex lacks very massive stars; a handful of low-luminosity compact and ultracompact \\hii regions have been identified, % but they are insufficiently powerful to trigger massive star formation throughout the cloud \\citep{ELD79,JF82}.  M17 SWex presents one of the most striking infrared dark cloud (IRDC) morphologies in the Galaxy, revealed in great detail by the {\\it Spitzer Space Telescope} as part of the Galactic Legacy Infrared Mid-Plane Survey Extraordinaire \\citep[GLIMPSE;][]{GLIMPSE} and the Multiband Imaging Photometer for {\\it Spitzer} Galactic Plane Survey \\citep[MIPSGAL;][]{MIPSGAL}. IRDCs have been the subject of much recent work because they are % laboratories for studying the initial conditions for massive (OB stars, $m\\ge 8$~\\Msun) star formation \\citep{RJS06,BS07,BT09}.  In this Letter, we use archival GLIMPSE and MIPSGAL data to find and characterize individual young stellar objects (YSOs) throughout the ${\\sim}1\\degree$ extent of M17 SWex. We model the mass function of the YSO population and find that it bears the imprints of both mass-dependent circumstellar disk evolution and delayed massive star formation in M17 SWex. ",
        "conclusions": "\\subsection{YSO Mass Spectrum Parameterization}  Assuming a Salpeter--Kroupa IMF applies to M17 SWex, % the steep YMF slope can be understood in terms of evolution and selection effects. Our sample % is populated by ongoing star formation and % includes {\\it only} YSOs with circumstellar material (disks). We parameterize these effects using the linear form of the mass function, the {\\it mass spectrum,} \\begin{equation} \\phi(m)={dN}/{dm}\\propto m^{-\\alpha},~~\\alpha=\\Gamma+1 \\end{equation} \\citep{BCM10}. The YSO mass spectrum $\\psi(m)$ is populated over time as \\begin{equation}\\label{prime} \\psi^{\\prime}(m)=\\int_{\\tau_0(m)}^{\\tau}\\frac{\\partial\\phi}{\\partial{t}}dt, \\end{equation} where $\\tau$ is the time since star formation began in the cloud, $\\partial\\phi(m)/\\partial{t}$ is the star formation rate (SFR) in mass interval $(m,~m+dm)$, and $\\tau_0(m)$ is the time at which star formation began {\\it for each mass interval}. Assuming for simplicity that the SFR is not time-dependent, $\\partial\\phi(m)/\\partial{t}\\to\\phi(m)/\\tau$ and the integral in Equation~\\ref{prime} becomes \\begin{equation}\\tag{3a}\\label{primea} \\psi^{\\prime}(m)=\\phi(m)[1-\\tau_0(m)/\\tau]=\\phi(m)f_0(m), \\end{equation} where $f_0(m)$ is the fraction of stars that have {\\it already} formed versus all stars that will {\\it eventually} form in each mass interval. In the limiting cases of $\\tau_0(m)=0$ (first stars to form in cloud) or $\\tau_0(m)\\ll\\tau$ (timescale for differential onset of star formation small compared to age of population), $f_0(m)=1$ and $\\psi^{\\prime}(m)=\\phi(m)$.  The YSO sample is further biased by disk evolution. Taking $\\tau_d(m)$ to be the disk lifetime as a function of mass, \\begin{equation}\\label{disks} \\psi(m)=\\psi^{\\prime}(m)\\frac{\\tau_d(m)}{\\tau-\\tau_0(m)}=\\psi^{\\prime}(m)f_d(m), \\end{equation} where $f_d(m)$ is the (mass-dependent) disk fraction. For $\\tau_d(m)\\ge\\tau-\\tau_0(m)$, $f_d(m)=1$.  Combining Equations \\ref{primea} and \\ref{disks}, \\begin{equation}\\label{combined} \\psi(m)=\\phi(m)f_0(m)f_d(m). % \\end{equation} Since the observed YMF can be fit with a power law (Fig.~\\ref{YMF}), we parametrize $f_0(m)$ and $f_d(m)$ as $m^{-\\omega}$ and $m^{-\\delta}$, respectively. Hence the observed deviation of the YMF slope from the standard IMF is \\begin{equation}\\label{powers} \\Delta\\alpha=\\Delta\\Gamma=\\Gamma_{\\rm YMF}-\\Gamma_{\\rm IMF}=\\omega+\\delta=2.2\\pm 0.6 \\end{equation} for $m\\ge M_c$ (Fig.~\\ref{YMF}).  \\subsection{Implications of a Steepened YSO Mass Function: Disk Evolution, Delayed Massive Star Formation, or Both?}  We have shown that a steep YMF could be produced by delayed formation of massive stars or by shorter disk lifetimes for more massive stars. Both interpretations have physical basis. Can either of these effects alone explain the observed YMF?  \\begin{figure} \\epsscale{1.0} \\plotone{f3.eps} \\caption{YMF of Fig.~\\ref{YMF}, binned up by a factor of 2 and subdivided by evolutionary stage. \\label{YMF_stages}} \\end{figure} {\\it Disk evolution.} The YMF is broken down by evolutionary stage in Fig.~\\ref{YMF_stages}, and it is apparent that the steep slope is driven by Stage II sources. % We thus consider the case of disk evolution alone: $\\omega\\to0$, $\\delta\\to2.2$ (Equation~\\ref{powers}). \\citet{JH07} measured the disk fractions in the $\\sigma$ Orionis cluster (age ${\\sim}3$~Myr) to be ${\\sim}35\\%$ for $m\\le1$~\\Msun\\ and ${\\sim}10\\%$ for $m>2$~\\Msun, giving $\\delta<1.8$. This does not rule out disk evolution as the primary driver of the YMF slope, given the uncertainties, but the disk lifetimes themselves, which set the duration of the Stage II phase, become problematically short. Adopting $\\delta = 1.8$ and $\\tau_d(1$~\\Msun$)=2$~Myr \\citep{HLL01}, we would predict $\\tau_d(4$~\\Msun$)=0.16$~Myr. This is comparable to the typical duration of the % envelope-accretion phase in low-mass stars \\citep[${\\sim}0.1$~Myr;][]{KH95}. % The Stage 0/I objects in our sample span a wide range in mass (Fig.~\\ref{YMF_stages}), yet the typical accretion timescale is similarly $\\tau_A\\sim0.1$~Myr ($\\tau_A=M_{\\star}/\\dot{M}_{\\rm env}$, where $\\dot{M}_{\\rm env}$ is the accretion rate in the RW06 models; see Table~\\ref{table} and P09). We expect the {\\it maximum} age of a 4~\\Msun\\ Stage II source in our sample to be $(N_{\\rm II}/N_{\\rm 0I}+1)\\tau_A\\sim0.7$~Myr ($N_{\\rm II}/N_{\\rm 0I}\\approx 6$ at $m=4$~\\Msun; Fig.~\\ref{YMF_stages}). This longer, more realistic disk lifetime implies $\\delta\\sim 0.9$.  {\\it Delayed Massive Star Formation.} Since disk evolution alone cannot explain the observed YMF, delayed massive star formation must {\\it also} contribute ($\\omega\\ne 0$). While diskless low- and intermediate-mass stars are lost among the overwhelming field star population in the IR images, % massive stars cannot easily hide within a dense molecular cloud. The \\citet{Kroupa} IMF, scaled to match the YMF peak (Fig.~\\ref{YMF}), predicts ${>}10$ O stars, including at least one early O star ($m>50$~\\Msun); such stars ionize dusty compact and ultracompact \\hii regions. The few \\hii regions in M17 SWex (visible as compact, bright extended sources in Fig.~\\ref{image}) are insufficiently luminous to contain O stars \\citep{ELD79,JF82}. We note also that the ratio $N_{\\rm II}/N_{\\rm 0I}$ steadily decreases with $m$ for $m\\ga 4$~\\Msun\\ and actually inverts for $m>6$~\\Msun\\ (Fig.~\\ref{YMF_stages}), suggesting that the more massive YSOs are preferentially younger. This is not a firm result, however, because it is based on ${<}20$ sources, and the trend may instead reflect instability of disks around massive YSOs.  While the data presented here do not support more than simple parameterizations, we can constrain the delay timescale for the most massive stars: $\\tau_0(m>20$~\\Msun$)>\\tau$. If the YMF break at $M_c\\sim 4$~\\Msun\\ (Fig.~\\ref{YMF}) is real, not an artifact of incompleteness, then $\\psi(m\\le M_c)=\\phi(m\\le M_c)$, hence $f_d(m\\le M_c)=1$ (Equation~\\ref{combined}), and $\\tau=\\tau_d(M_c)\\sim0.7$~Myr (Equation~\\ref{disks}) is age of the oldest YSOs in the cloud.  \\subsection{Present-Day Star Formation Rate}  Integrating the scaled \\citet{Kroupa} IMF (Fig.~\\ref{YMF}) over $m\\ge 0.1$~\\Msun\\ yields a total population of $1.3\\times 10^4$ YSOs in M17 SWex, with a total stellar mass of $8\\times 10^3$~\\Msun. \\citet{EL76} predicted that M17 SWex would eventually form an OB association; these numbers, lower limits due to possible incompleteness, show that the M17 proto-OB association is already forming. Adopting $\\tau=0.7$~Myr, the present-day SFR in M17 SWex is 0.011~\\Msun~yr$^{-1}$, comparable to the time-averaged SFR of M17 itself (P09), albeit distributed over a much larger volume.  \\subsection{Sequential Star Formation in the M17 Complex}  \\citet{ELD79} compared the morphology of the M17 complex to the dust lanes and ``beads-on-strings'' \\hii regions of extragalactic spiral structure. P09 discovered an extended, diffuse \\hii region northeast of M17, called M17 EB, ionized by a group of optically revealed O stars in a 2--5~Myr old cluster or association. When M17 SWex is included, the entire M17 complex presents a clear sequence of star formation extending ${>}100$~pc % parallel to the Galactic midplane and spanning several Myr in age. As \\citet{EL76} noted, Galactic spiral density waves propagate in the direction of decreasing age. % The difference between the Galactic rotation speed and the spiral pattern speed at the location of M17 is $221~\\kms-130~\\kms=91~\\kms$ \\citep{BB93,MM04}. The timescale for a spiral shock to cross the entire M17 complex % is ${\\sim}1$~Myr, the same order of magnitude as the observed age spread.  The rapid passage of the M17 complex through the Sagittarius spiral arm served as the ``global'' trigger ultimately responsible for the formation of this large OB association. Has sequential ``local'' triggering driven by OB stars \\citep{EL77} played an important role? % While there is strong circumstantial evidence for locally triggered star formation on the periphery of the M17 and M17 EB \\hii regions (P09 and references therein), it remains unclear whether any of the {\\it major} OB clusters were similarly triggered. Global triggering appears to be a far more likely explanation for the formation of the proto-OB association in M17 SWex. The ${\\la}0.7$~Myr spread in ages among the YSO sample is small compared to the 2--3~Myr timescale for sequential triggering by OB stars \\citep{EL77}. % The spatial distribution of both the stars and the dust (Fig.~\\ref{image}) suggest a disturbed molecular cloud that is experiencing global collapse and fragmentation."
    },
    "1004/1004.3785_arXiv.txt": {
        "abstract": "Using cosmological simulations with a dynamic range in excess of $10^7$, we study the transport of gas mass and angular momentum through the circumnuclear region of a disk galaxy containing a supermassive black hole (SMBH). The simulations follow fueling over relatively quiescent phases of the galaxy's evolution (no mergers) and without feedback from active galactic nuclei (AGNs), as part of the first stage of using state-of-the-art, high-resolution cosmological simulations to model galaxy and black hole co-evolution. We present results from simulations at different redshifts ($z=6$, $4$, and $3$) and three different black hole masses ($3\\times10^7$, $9\\times10^7$, and $3\\times10^8 \\dim{M}_{\\sun}$; at $z=4$), as well as a simulation including a prescription that approximates optically thick cooling in the densest regions. The interior gas mass throughout the circumnuclear disk shows transient and chaotic behavior as a function of time. The Fourier transform of the interior gas mass follows a power law with slope $-1$ throughout the region, indicating that, in the absence of the effects of galaxy mergers and AGN feedback, mass fluctuations are stochastic with no preferred timescale for accretion over the duration of each simulation ($\\sim 1$-$2 \\dim{Myr}$). The angular momentum of the gas disk changes direction relative to the disk on kiloparsec scales over timescales less than $1 \\dim{Myr}$, reflecting the chaotic and transient gas dynamics of the circumnuclear region. Infalling clumps of gas, which are driven inward as a result of the dynamical state of the circumnuclear disk, may play an important role in determining the spin evolution of an SMBH, as has been suggested in stochastic accretion scenarios. ",
        "introduction": "Supermassive black holes (SMBHs) with masses from $\\sim 10^6$ to more than $\\sim 10^9 \\dim{M}_{\\sun}$ are found in the centers of most galaxies \\citep[e.g.,][]{KormRich, Magorrian98}. Observations of correlations between the masses of SMBHs and several properties of their host galaxies indicate that the growth of SMBHs is closely tied to the evolution of their hosts \\citep[e.g.,][]{Magorrian98, FM00, Geb00, Trem02}. Scenarios including feedback from active galactic nuclei (AGNs) have been explored in order to understand both the growth of SMBHs and the nature of their relationship with their host galaxies \\citep[e.g.,][]{SilkRees98, KaufHae00, WyitheLoeb03, DiMatteo05, Springeletal05a, Springeletal05b, Crotonetal06, DiMatteo08, Hopkins07, JohanssonNB09}. AGNs are thought to be powered by accreting SMBHs \\citep[e.g.,][]{LyndenBell69}, so that observations of the various phases of AGN activity may provide clues about the connection between SMBH growth and galaxy evolution, as models explored by, e.g., \\citet {Hopkins05} and \\citet{Sijackietal07} have shown. It is therefore essential to characterize accretion onto SMBHs in order to understand AGNs and their relevance for galaxy evolution.  Substantial SMBH growth requires a mechanism for replenishing fuel on the subparsec scales feeding the SMBH accretion disk. Large scale tidal torques caused by galaxy mergers are an effective mechanism for transporting gas from super-galactic and galactic scales down to scales of several parsecs, where other mechanisms may become more viable for funneling material the rest of the distance toward the SMBH \\citep[e.g.,][]{Hernquist89, BarnesHernquist92}. Secular evolution may also play a role as disk instabilities lead to the formation of bars and spiral waves which can transport material \\citep{Robertsetal79, barswinbars, RegTeu04}. There are several observations of inflowing gas in the circumnuclear regions of low-luminosity AGNs, apparently driven by global instabilities arising from secular evolution \\citep[see, e.g.,][]{NUGA4, NUGA7, StorchiBergmann07, NUGA8, Riffeletal08}. Mergers are thought to be effective for building SMBHs in more massive early-type galaxies, which appear to form earlier \\citep[see, e.g.,][and references therein]{Hopkinsetal08a, Hopkinsetal08b}, ultimately building the SMBH-bulge relationships, whereas secular evolution may play a larger role in the growth of late-type galaxies, particularly at lower redshift \\citep[see, e.g.,][and references therein]{KormKenn04}.  Whichever is the dominant mechanism for fueling SMBHs, the specific details of how fueling occurs in the circumnuclear regions of galaxies are still not understood. Fueling may occur continuously, over the course of large-scale dynamical instabilities in the galaxy, or it may be an intermittent process, dependent entirely on the dynamics on small scales. Cold gas, stochastically accreted from the circumnuclear regions of galaxies, may be able to sustain the fueling of low-luminosity AGNs \\citep{HopkinsHernquist06}. Intermittent accretion episodes, consisting of in-falling clouds of gas with randomly oriented angular momentum vectors, may contribute to the spin-down of SMBHs, thus lowering their radiative efficiency and allowing them to grow faster \\citep[see][]{KingPringle06,KingPringle07,Kingetal08}. \\citet{Wangetal09} have used a Soltan-type argument \\citep{Soltan82} to show that the average radiative efficiency of SMBHs has decreased over time, as might be expected if SMBHs grow predominately through episodic accretion events that act to lower their spin. Using cosmological simulations which followed the buildup of large SMBHs at $z=6$, \\citet{Sijackietal09} showed that black hole mergers can contribute to the spin-down of SMBHs as well, resulting in low radiative efficiencies and rapid growth. A decrease in radiative efficiency must be reconciled with observations of an increase in the fraction of radio-loud quasars with decreasing redshift \\citep{Jiangetal07}, which would suggest SMBHs with larger spins (and hence larger efficiencies) at low-$z$. Therefore, understanding the behavior of angular momentum in the circumnuclear region is an important part of modeling accretion, and subsequently AGN feedback.  Simulations that follow the growth of SMBHs over cosmological times, or over the duration of a galaxy merger, often cannot follow the circumnuclear regions of galaxies with high enough resolution to describe the accretion flow in detail. Such simulations must make approximations for the accretion rates using the properties of the galaxies on scales that are resolved. A common technique is to assume that the unresolved disk is fed by Bondi-type accretion \\citep{HoyleLyttleton39, BondiHoyle44, Bondi52}, as in for example the smoothed particle hydrodynamics simulations of \\citet{Springeletal05a} who estimate the accretion rate based on the properties of the gas on a scale of $\\sim100 \\dim{pc}$. The assumption of Bondi accretion appears to be reasonable for following evolution over cosmological times, where the average accretion rate onto the black hole cannot be more than the average accretion rate through any radius, since the fuel supply is limited by large scales. Previous cosmological simulations have been successful in reproducing observed population demographics and trends, such as the black hole mass function and galaxy colors and morphologies \\citep[as in, e.g.,][]{DiMatteo08, Croftetal09, McCarthyetal09}. However, for detailed studies of the growth and evolution of individual SMBHs or of the coupling between AGN feedback and the black hole accretion rate, more accurate descriptions of the accretion rate and its dependence on the small-scale features of the host galaxy are desirable.  Small-scale simulations have addressed gas dynamics in subgalactic-scale disks with high resolution \\citep{Fukuda00, Wada01, WadaNorman01,Escala07,WadaNorman07,KawakatuWada08}, finding the development of a turbulent, multi-phase interstellar medium. The approximation of Bondi accretion in such an environment is not necessarily invalid. \\citet{Krumholzetal05} have shown that modified forms of the Bondi prescription can describe accretion in turbulent environments. A simulation must be equipped to model the properties of the turbulence in order to employ the modified Bondi prescription. If the gas contains a significant amount of angular momentum, the unmodified Bondi prescription gives inaccurate estimates of the accretion rate as well \\citep{ProgaBegelman03, Krumholzetal06}. However, even modified Bondi prescriptions become inapplicable in the case of the self-gravitating, rotationally supported disks that are likely to form as large amounts of gas are driven inward in high-redshift galaxies. As the use of adaptive techniques increasingly improves the resolution of cosmological simulations, accretion onto black holes can no longer be described by approximate prescriptions such as Bondi accretion.  In the present paper, we use cosmological adaptive mesh refinement simulations with a large dynamic range to study the transport of gas and angular momentum through the circumnuclear disk of an SMBH host galaxy over time. The goal is to provide a description of accretion that can be compared to prescriptions typically applied in larger scale simulations. It will be shown that in the limiting case of relative quiescence (e.g., in between merger events and without AGN feedback) accretion is a stochastic rather than a continuous process. The method behind the simulations used here is described in detail in \\citet{Levineetal08}, hereafter Paper I, and in \\citet{MyThesis}. In Section \\ref{sec:simtr}, we briefly summarize the details of the simulations. The results of the simulations are given in Sections \\ref{sec:mass}-\\ref{sec:mom}, including an analysis of the mass accretion rate and the angular momentum of the gas in the circumnuclear region of the galaxy. Finally, we summarize and discuss the results in Section \\ref{sec:disctr}. ",
        "conclusions": "\\label{sec:disctr}  We have studied the transport of gas in the circumnuclear region of a disk galaxy within a cosmological simulation at different redshifts, for different SMBH masses, and for a run including an approximation of optically thick cooling. The mass accretion rate is not steady, but fluctuates randomly and substantially, the accretion rate through the circumnuclear disk being almost as often negative as positive.  This result is consistent with that of \\citet{HopkinsQuataert09}, who find a large amount of variation in the instantaneous accretion rates in their galaxy simulations on scales $< 1 \\dim{kpc}$. We find that there is no preferred timescale for accretion over the course of the $\\sim$million years spanned by each of the simulations. The flat slope of the Fourier transform of the accretion rate characterizes the fluctuations as a function of timescale, revealing the stochastic nature of accretion in the simulations. This complex and chaotic behavior is apparent at each of the three redshifts explored in detail here: $z=3$, $4$, and $6$, as well as in $z=4$ simulations containing black holes with larger masses, and the $z=4$ run including optically thick cooling.  The dynamic and turbulent nature of the simulated circumnuclear disk is not well modeled by Bondi accretion, even when using a modified prescription that takes into account the vorticity and turbulent properties of the gas. Bondi prescriptions do not take into account the effect of the disk's self-gravity.  Additionally, the modified Bondi accretion rates, estimated from the properties of the disk in the highest-resolution part of the simulation (where the density rises steeply with decreasing radius), are highly super-Eddington on the smallest scales (assuming a radiative efficiency of $0.1$). As the Eddington rate provides an upper limit for the accretion rate in the presence of radiative feedback from the SMBH, highly super-Eddington rates are not likely to endure in nature.  The Eddington rate (or a few times the Eddington rate) is set as an upper limit to the accretion rate in many of the simulations that employ Bondi type prescriptions to model accretion \\citep[e.g.,][]{Springeletal05a, DiMatteo05, Lietal07b, DiMatteo08}. On scales corresponding to the spatial resolution of those simulations ($\\gtrsim 100 \\dim{pc}$), the accretion rates predicted by the modified Bondi prescription in the present simulations are closer to the Eddington limit than they are at smaller scales. However, black holes in other simulations spend the majority of time accreting well below this limit, at substantially sub-Eddington rates. The difference between the rates determined here and those of other simulations is even larger when comparing with prescriptions that multiply the Bondi accretion rate by a large factor to approximate the effects of the small scale physics, since no such factor has been included in our calculations \\citep[see discussion of the parameter $\\alpha$][]{BoothSchaye09, JohanssonBN09, JohanssonNB09}. However, it is not straightforward to make a direct comparison between accretion rates in the present simulations and the rates in other simulations that contain prescriptions for AGN feedback. As shown in galaxy merger simulations \\citep[e.g.,][]{Springeletal05a, Debuhretal09}, the presence of AGN feedback regulates black hole growth by forcing gas out of the circumnuclear region. Future work incorporating AGN feedback into the present simulations will address this issue.  The absence of AGN feedback in our simulations may result in an excessively dense circumnuclear gas disk, which corresponds to large, super-Eddington accretion rates within the Bondi prescription. Allowing the black hole particle to grow in situ in our simulations, or using a more physically motivated recipe for star formation in the zoom-in simulation would further change the distribution of mass in the circumnuclear gas disk. The circumnuclear disk is self-gravitating, almost certainly behaving outside the Bondi regime, even with the modifications of \\citet{Krumholzetal06}. This result reinforces the need for high-resolution simulations incorporating detailed gas dynamics and radiative processes, in order to truly model the complicated dynamics on small scales, which ultimately govern accretion onto SMBHs.  The chaotic behavior of the disk may be consistent with models suggesting that stochastic accretion of molecular clouds from the circumnuclear region of the galaxy can power low-luminosity AGNs \\citep{HopkinsHernquist06}. This manner of accretion may determine the spin of the black hole. We find that the angular momentum vector on scales $\\lesssim 100 \\dim{pc}$ can vary substantially from the direction of angular momentum on kiloparsec scales. In each simulation, the axis of the disk is closely aligned with the axis on kiloparsec scales at the beginning of the zoom-in simulation, but slowly over time shifts direction as the circumnuclear disk dynamically evolves.  The angular momentum can be misaligned with the rest of the circumnuclear disk, resembling the stochastic fueling scenarios mentioned in Section \\ref{sec:mom}. As is observed in several of the simulations presented here, clumps of gas develop in the circumnuclear disk, which change the angular momentum of the gas that ultimately feeds the accretion disk. The change can occur as clumps fall into the center or as larger clumps develop within the disk and change the potential. However, as cautioned above, the simulation does not yet include feedback from an accreting black hole, which will heat the gas and drive out some of the material in the circumnuclear disk. Future work including more realistic physics in the simulations will allow us to constrain better the effects of accreting gas on the properties of the SMBH and the feedback it produces."
    },
    "1004/1004.0609_arXiv.txt": {
        "abstract": "{ Rotationally-split modes can provide valuable information about the internal rotation profile of stars. This has been used for years to infer the internal rotation behavior of the Sun. The present work discusses the potential additional information that rotationally splitting asymmetries may provide when studying the internal rotation profile of stars. We present here some preliminary results of a method, currently under development, which intends: 1) to understand the variation of the rotational splitting asymmetries in terms of physical processes acting on the angular momentum distribution in the stellar interior, and 2) how this information can be used to better constrain the internal rotation profile of the stars. The accomplishment of these two objectives should allow us to better use asteroseismology as a test-bench of the different theories describing the angular momentum distribution and evolution in the stellar interiors.} ",
        "introduction": "The study of internal rotation of stars is one of the main issues in stellar physics. Rotation is present in almost all the stars, and it interacts with other physical processes acting in the stellar interior. In particular, understand the transport of angular momentum in the interior of stars is crucial to correctly and precisely describe the evolution. Turbulence, meridional circulation, mixing of elements responsible of different $\\mu$ gradients during evolution, dynamo effects due to the presence of magnetic fields, etc., are some of the physical phenomena and processes affected by rotation (see e.g. Goupil et al. 2005, Goupil 2009, and Goupil \\& Talon 2009, for a review on this topic).  Nowadays, it is possible to probe the internal structure of stars thanks to asteroseismology. In what regards rotation, progress has been made, during the last decades, on the knowledge about the rotation-pulsation interaction. Up to now, this problem has been tackled using the perturbation techniques to compute the stellar oscillations. These methods (see e.g. Dziembowski \\& Goode 1992; Soufi et al. 1998; Su\\'arez, Goupil \\& Morel 2006) are only valid for slow-to-moderate rotators (see a review on the effects of rotation on stellar p modes by Goupil 2009). For faster rotators, the distortion of the stellar structure due to the centrifugal is too large and invalids the perturbation techniques. Non-perturbative approaches must thus be considered (Ligni\\`eres et al. 2006; Reese et al. 2006).  The present work consider slow-to-moderate rotators, i.e. those for which the parameters  $\\epsilon=\\Omega/(G\\,M/R^3)^{1/2}$ and $\\mu=\\Omega/\\nu_{n,\\ell}$ are small, i.e. the stellar structure is not significantly deformed by the centrifugal force, and oscillation frequencies are much larger than the angular rotation rate, respectively. The failure of the perturbative approach comes first for high radial-order frequency modes. We restrict thus this study to rotating stars showing oscillations in a relatively low-order frequency domain (low g- and p modes, or mixed modes), like some \\dss, \\gds, solar-like stars and some massive stars like \\bceph.  The heterogeneity of internal structures (chan\\-ging with evolution and different from one star to another) and processes therein, are strong arguments to assume non-uniform rotation. It is thus necessary to take into account possible variations (in the radial or angular directions) of the angular momentum transport, and thus in the shape of the internal rotation profile. To do so, we use the oscillation code {\\sc filou} (Su\\'arez 2002; Su\\'arez \\& Goupil 2008) which corrects the oscillation frequencies for up to second-order effects of rotation (including near degeneracy effects) in presence of radial differential rotation. The study of radial differential rotation has also been used for the asteroseismic studies, e.g. Casas~(2006, 2009), Fox Machado ~(2006), Bruntt et al.~(2007a, 2007b); Su\\'arez et al.~(2005a), Su\\'arez et al.~(2006b); Su\\'arez et al.~(2007), for \\dss, Rodr\\'{\\i}guez et al.~(2006a, 2006b), Uytterhoeven et al.~(2008), Moya et al.~(2005), Su\\'arez et al.~(2005b), for \\gds, or Su\\'arez et al.~(2009) for \\bceph\\ stars.  This technique is applied in the present work to analyze the asymmetries of the mode splittings due to rotation. In particular, we are interested in examining the behavior of the rotational splitting and their asymmetries for different mode types (g and p), in presence of radial differential rotation, i.e. to understand physically how variations of the internal rotation profile affect the splitting asymmetries. ",
        "conclusions": ""
    },
    "1004/1004.2867_arXiv.txt": {
        "abstract": "When the electrons stored in the ring of the European Synchrotron Radiation Facility (ESRF, Grenoble) scatter on a laser beam (Compton scattering in flight) the lower energy of the scattered electron spectra, the Compton Edge (CE), is given by the two body photon-electron relativistic kinematics and depends on the velocity of light. A precision measurement of the position of this CE as a function of the daily variations of the direction of the electron beam in an absolute reference frame provides a one-way test of Relativistic Kinematics and the isotropy of the velocity of light. The results of GRAAL-ESRF measurements improve the previously existing one-way limits, thus showing the efficiency of this method and the interest of further studies in this direction. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.2056_arXiv.txt": {
        "abstract": "We discuss a model of Poynting-dominated gamma-ray bursts from the collapse of very massive first generation (pop. III) stars. From redshifts of order 20, the resulting relativistic jets would radiate in the hard X-ray range around 50 keV and above, followed after roughly a day by an external shock component peaking around a few keV. On the same timescales an inverse Compton component around 75 GeV may be expected, as well as a possible infra-red flash. The fluences of these components would be above the threshold for detectors such as Swift and  Fermi, providing potentially valuable information on the formation and properties of what may be the first luminous objects and their black holes in the high redshift Universe. ",
        "introduction": "\\label{sec:intro}  Population III stars are widely considered to consist mainly of `very massive stars' (VMS) in the the range of hundreds of solar  masses \\citep{ohkubo+06,yoshida+06}.  The VMS are expected to be very fast rotating, close to the break-up speed, and accretion leads to a mass upper limit which may be around $10^3\\msun$.  Those in the $140\\msun \\siml M_\\ast \\siml 260\\msun$ range are expected to be subject to pair instability and explode as supernovae without leaving any compact remnant behind, while those above $\\sim 260\\msun$ are expected to undergo a core collapse leading directly to a central black hole \\citep{heger-woosley02}, whose mass would itself be hundreds of stellar masses. Accretion onto such massive black holes could lead to a scaled-up collapsar gamma-ray burst \\citep{heger+03,komissarov-barkov09}.  % In this paper we discuss a specific scenario for pop. III VMS collapsars at redshifts of order $z\\sim 20$, resulting in Poynting dominated relativistic jets which produce GRBs with characteristic radiation properties extending from soft X-rays to multi-GeV energies. ",
        "conclusions": "\\label{sec:disc}  In this paper we have explored in some detail the spectral and temporal properties of high redshift population III GRBs within the context of a Poynting dominated relativistic jet model. The core collapse of a pop. III very massive star of $250\\siml (M_\\ast /\\msun) \\siml 10^3$ will result in an intermediate mass black hole and a temporary accretion torus, and for fast rotating objects the extraction of rotational energy from the black hole could power a Poynting dominated outflow. The largest uncertainty, in such a model, is the value of the Poynting luminosity $L$ and its time-dependence. The default assumption made here is that, for a constant magnetic $\\alpha$ viscosity in the torus, $L$ is approximately constant for $1/\\alpha$ times the free-fall time from the boundary of the star.  This is of course uncertain because the efficiency of field build-up is unknown, as is whether the magnetic stresses get up to a significant fraction of the gas pressure $P$. The  quantity that determines the jet luminosity is the field strength around the hole, which depends on the peak density (and peak pressure $P$) near the inner part of the disk. In the standard alpha model the density in the disk goes as $r^{-3/2}$.  However, for a radiation dominated gas (which is more compressible, even if the radiation is trapped on the relevant timescale) it would in principle be possible for the density law to be closer to $r^{-3}$ (and this could happen if the effective $\\alpha$ were to decrease towards small $r$). Thus, the disk could have a higher peak density (and steeper profile) whatever the initial stellar profile was  -- and the jet could have a much higher luminosity for a shorter period. Our timescale estimates are also subject to uncertainty. We took nominally the star to be rotating at half the break-up speed, but the stellar angular velocity could be non-uniform. If the outer regions were rotating more slowly than we assumed, the the disk would obviously be smaller, but the times scale would still be the free-fall time, so the accretion time could be up to a factor $\\alpha=10^{-1}\\alpha_{-1}$ shorter. Thus, there are large uncertainties in the timescales and luminosities we have estimated, which could be off in either direction. Nonetheless, the simplifying assumption made here may be appropriate considering the preliminary state of knowledge about population III stars.  The spectrum of the `prompt' emission within the first day is shown to extend from soft X-rays to the multi-GeV range, with a characteristic time evolution. As a rough estimate, the luminosity of the annihilation photosphere (mainly in the form of photons) can be taken to be of the order of the remaining kinetic luminosity $L/2$, with an efficiency of conversion into annihilation photons of $\\eta_\\gamma \\sim 1/2$. The annihilation component of eq.(7) peaks at 50 keV and extends as $N(E)\\propto E^{-2}$ up to $\\siml 3 k m_e c^2$ in the comoving frame, or $\\sim 1$ MeV in the observer frame.  The the spectral energy flux  per decade $E^2(dN/dE)$ is the total energy times $1/(\\ln(E_{max}/E_{min})=1/\\ln(20)\\sim 0.31$. The radiation falling in the BAT band, 50-150 keV has $1/\\ln(150/50)\\sim 0.9$, so we can take the spectral efficiency in the BAT range, say 50-150 keV, as $\\eta_{BAT}\\sim 0.5\\times 0.3\\times 0.9 \\sim 0.13 \\sim 10^{-1}$. The X-ray flux might be even larger than this, adding a roughly comparable contribution from the external shock synchrotron component of eq. (9).  With this efficiency, the predicted X-ray flux would be detectable in X-rays and hard X-rays by instruments such as the BAT detector on Swift, as estimated in \\S \\ref{sec:pop3col}. This would be detectable also in the GBM detector on Fermi, whose sensitivity is slightly better than Swift's. The GeV range Large Area detector (LAT) on Fermi has a fluence sensitivity for times $t\\simg 3\\times 10^4\\s$ of $\\sim 3\\times 10^{-9}t^{1/2}$ even with significantly less than 10\\% of the luminosity in the GeV band at $t\\sim 10^5\\s$, this component would be detectable by the LAT. To detect them, however, it may be necessary to adjust the flux trigger algorithms to respond to a low level, very extended increase in the flux.  The spectral signature would have an initial hard, $\\sim 50\\keV$ X-ray rise from the annihilation photons (equation [\\ref{eq:Ea}]) lasting for about a day, with a possible extension  out to $\\sim 25\\GeV$ from up-scattering in the pair photosphere (equation [\\ref{eq:Ephotsc}]). There could be a cascade component from external photons leading almost simultaneously to soft X-rays in the few keV range, which is subject to considerable uncertainties. These would be followed, after a delay of hours up to a day, by an external shock synchrotron component in the few keV range (equation [\\ref{eq:Esyfor}]). An inverse Compton component at energies in the 70 GeV range (equation [\\ref{eq:Essc}]) may also be expected, lagging by about ten minutes after the keV range external shock synchrotron component. If the jets acquire a non-negligible baryon load at some stage before the external shock, a reverse shock may result in an infrared flash of $\\simg 13$th magnitude. An afterglow similar to that of lower redshift GRBs would follow over the next days, gradually shifting into the optical, infrared and radio frequency bands.  The detection of such very high redshift GRBs would be of great value, as it might be the first and perhaps the only way to trace the formation of the first generation of stellar objects in the Universe. It could give important information about the redshift at which the initial objects form, the rate at which they form, and the input of radiation into the Universe at those early epochs and its contribution to the reionization of the intergalactic medium."
    },
    "1004/1004.1396_arXiv.txt": {
        "abstract": "Transmission spectroscopy at UV wavelengths is a rich and largely unexplored source of information about the upper atmospheres of extrasolar planets.  So far, UV transit observations have led to the detection of atomic hydrogen, oxygen and ionized carbon in the upper atmosphere of HD209458b.  The interpretation of these observations is controversial -- it is not clear if the absorption arises from an escaping atmosphere interacting with the stellar radiation and stellar wind, or the thermosphere inside the Roche lobe.  In this paper, we introduce an empirical model that can be used to analyze UV transit depths of extrasolar planets.  We use this model to interpret the transits of HD209458b in the H~Lyman~$\\alpha$ and the 1304~\\AA~O~I triplet emission lines.  The results indicate that the mean temperature of the thermosphere is $T =$~8,000--11,000~K and that the H$_2$/H dissociation front is located at pressures between $p =$~0.1--1~$\\mu$bar, which correspond to an altitude of $z \\approx$~1.1~R$_p$.  The upper boundary of the model thermosphere is located at altitudes between $z =$~2.7--3~R$_p$, above which the atmosphere is mostly ionized.  We find that the H~I transit depth reflects the optical depth of the thermosphere in the wings of the H~Lyman~$\\alpha$ line but that the atmosphere also overflows the Roche lobe.  By assuming a solar mixing ratio of oxygen, we obtain an O~I transit depth that is statistically consistent with the observations.  An O~I transit depth comparable to or slightly larger than the H~I transit depth is possible if the atmosphere is undergoing fast hydrodynamic escape, the O/H ratio is supersolar, or if a significant quantity of neutral oxygen is found outside the Roche lobe.  We find that the observations can be explained solely by absorption in the upper atmosphere and extended clouds of ENAs or atoms strongly perturbed by radiation pressure are not required.  Due to the large uncertainty in the data, repeated observations are necessary to better constrain the O~I transit depths and thus the composition of the thermosphere. ",
        "introduction": "\\label{sc:intro}  Visible and near-IR spectroscopy of transiting extrasolar planets has led to the detection of Na~I, H$_2$O, CH$_4$, CO, and CO$_2$ \\citep{charbonneau02,tinetti07,swain08,swain09} in the atmospheres of extrasolar giant planets (EGPs).  Discoveries such as these have helped to shift the emphasis in the study of extrasolar planets from mere detection towards characterization of the dynamics and composition of their atmospheres.  Modeling and observations of a large variety of objects allow us to constrain theories of the formation and evolution of different planets and their atmospheres.  In this respect, studies of mass loss from the atmospheres of close-in EGPs and terrestrial planets are of particular interest.  In addition to being interesting in their own right, these studies are valuable for constraining models of planetary atmospheres during the early history of the solar system.  Approximately 25\\% of the currently known extrasolar planets orbit within 0.1 AU from their host stars\\footnote{See the Extrasolar Planets Encyclopaedia, maintained by J. Schneider, at http://www.exoplanet.eu/.}.  At such small orbital distances many of these planets are subject to intense stellar irradiation.  In particular, the outermost layers of the atmosphere are strongly ionized and heated by the absorption of FUV and EUV radiation, which leads to the formation of the thermosphere \\citep[e.g.,][]{yelle04,garciamunoz07,koskinen07}.  Modeling indicates that the thermospheres of close-in extrasolar giant planets are hot, greatly extended, and their composition is dominated by atomic constituents.  The same should apply to close-in super-Earths orbiting M dwarf stars.   These are the closest types of object to an Earth-like habitable planet whose atmospheres we may be able to probe with space-born or even ground-based telescopes in the near future \\citep[e.g.,][]{palle09,charbonneau09}.  As a consequence, the UV transit depths of extrasolar planets should be much larger than the visible and IR transit depths because the observations probe the upper atmosphere that strongly absorbs UV radiation in electronic resonance lines.  Transit observations of extrasolar planets in the FUV were pioneered by \\citet{vidalmadjar03} (hereafter VM03) and \\citet{vidalmadjar04} (hereafter VM04) who detected atomic hydrogen (H~I), oxygen (O~I), and ionized carbon (C~II) in the upper atmosphere of the transiting planet HD209458b.  All of these detections were obtained by using the Space Telescope Imaging Spectrograph (STIS) onboard the Hubble Space Telescope (HST).  VM03 used the G140M grating to observe three transits of HD209458b in the wavelength range covering the stellar H~Lyman~$\\alpha$ (H~Ly$\\alpha$) emission line.  The spectral resolution of these observations was $\\sim$0.08~\\AA, which allowed for the details of the line profile to be resolved.  VM04 used the G140L grating with a low spectral resolution of $\\sim$2.5~\\AA~to observe four transits in the wavelength range of [1180,1710]~\\AA~and detected absorption in the O~I and C~II stellar emission lines.  VM03 deduced a H~I transit depth of 15~$\\pm$~4~\\% from the ratio of the flux in two wavelength regions around the core of the H~Ly$\\alpha$ line to the flux in the wings of the line during transit.  Based on this observation, they concluded that H~I overflows the Roche lobe with a mass loss rate of $dM/dt >$~10$^{7}$~kg~s$^{-1}$.  They also suggested that the planet is followed by a cometary tail that is shaped by stellar radiation pressure acting on the escaping hydrogen.  Later, \\citet{benjaffel07} (hereafter BJ07) and \\citet{benjaffel08} (hereafter BJ08) presented a thorough and convincing reanalysis of the same data.  Based on this analysis, BJ08 argued that there is no evidence for a cometary tail in the transit light curve or the detailed transit depth measurements.  Further, he showed that the observed absorption may arise from atomic hydrogen below the Roche lobe and that the H~I absorption line profile is broadened by thermal and natural broadening in the thermosphere of the planet.  However, his results also imply that hydrogen overflows the Roche lobe, and thus the atmosphere is still evaporating \\citep{vidalmadjar08}.  VM04 used the low resolution G140L data to obtain O~I and C~II transit depths of 13~$\\pm$~4.5 and 7.5~$\\pm$~3.5~\\%, respectively.  They also calculated a full-width H~I transit depth of 5~$\\pm$~2~\\% from the unresolved H~Ly$\\alpha$ line and claimed that this transit depth is consistent with the much stronger absorption observed within a limited section of the line profile by VM03.  Further, they argued that the large O~I and C~II transit depths are possible because these species overflow the Roche lobe and the absorption lines are broadened by the velocity dispersion of the escaping gas.  Thus the observations were interpreted as proof that the atmosphere is undergoing fast hydrodynamic escape.  \\citet{benjaffel10} (hereafter BJ10) published a thorough reanalysis of the G140L data.  They obtained revised full-width H~I, O~I, and C~II transit depths of 6.6~$\\pm$~2.3~\\%, 10.5~$\\pm$~4.4~\\%, and 7.4~$\\pm$~ 4.7~\\%, respectively.  We note that these depths are only 2, 1.93, and 1.15$\\sigma$, respectively, away from the FUV continuum transit depth of $\\sim$2~\\%.  The more detailed H~I transit depth measurements reported by BJ08 provide stronger constraints for the H~Ly$\\alpha$ line but no such constraints are available for the O~I and C~II transits.  In order to explain the large transit depths in these lines, BJ10 argued that oxygen and ionized carbon are preferentially heated to a temperature more than ten times higher than the temperature of hydrogen within a layer in the atmosphere located between $\\sim$2.25~R$_p$ and the boundary of the Roche lobe at 2.9~R$_p$.  Most analyses of the HD209458b UV absorption signatures to date have either been limited to first order deductions, such as the effective size of the absorbing obstacle (VM03,VM04) or been based on complicated first principle models for the atmosphere \\citep[][]{tian05,garciamunoz07,murrayclay09,benjaffel10}.  We believe that there is an important role for an intermediate class of models that satisfy some basic physical constraints, but parameterize aspects of the atmosphere that are difficult or impossible to predict accurately.  For example, it is well established that the thermosphere of HD209458b should be composed primarily of H and H$^+$, but the location of the transition from H$_2$ to H is uncertain with different physical models making vastly different predictions \\citep{liang03,yelle04,garciamunoz07}.  Moreover, the boundary between the atmosphere and interplanetary space is dependent upon the unknown strength of the stellar wind and the planetary and interplanetary magnetic field and has yet to be modeled in a realistic fashion.  Finally, although it seems well established that the temperature is of the order 10,000~K, the precise value depends on the heating efficiencies, which have yet to be calculated, and radiative cooling by minor species, which is not included in any of the models.  Because of these uncertainties it is important to analyze the data in a way that makes clear what aspects of the atmosphere are constrained by the observations and which are not.  In this paper we introduce a generic methodology that can be used to interpret UV transit light curves in stellar emission lines.  We demonstrate this methodology by using a simple empirical model of the thermosphere to analyze the H~I and O~I transit depths of HD209458b summarized by BJ10.  In Section~\\ref{sc:methods} we introduce a model for calculating transit light curves for planets with extended atmospheres and discuss the basic features of the model thermosphere.  In Section~\\ref{sc:results} we discuss the H~I transit depth measurements in detail and confirm that they \\textit{can} be explained by absorption by atomic hydrogen below the Roche lobe.  Nevertheless, we also show that the core of the H~Ly$\\alpha$ absorption line is optically thick up to the Roche lobe and that the atmosphere is evaporating.  Further, we demonstrate that the disagreement between BJ08 and VM03 is due to differences in the treatment of the data and different definitions of the transit depth.  In the rest of Section~\\ref{sc:results}, we discuss the transits in the O~I triplet lines and show that the empirical model thermosphere with a solar abundance of oxygen can be used to obtain transit depths that are statistically consistent with the observations.  We also address the feasibility of the idea that energetic oxygen atoms are present in the thermosphere of HD209458b and present alternative ways to explain O~I transit depths that are comparable to or larger than the full-width H~I depth.  We conclude Section~\\ref{sc:results} by discussing a variety of different models for the thermosphere of HD209458b and use them to calculate transit depths.  In Section~\\ref{sc:discussion} we discuss the feasibility of our assumptions and in Section~\\ref{sc:conclusion} we summarize our findings and conclusions. ",
        "conclusions": "\\label{sc:conclusion}  Transmission spectroscopy in UV wavelength bands is a rich source of information about the upper atmospheres of extrasolar planets.  In this paper we have introduced a generic method that can be used to interpret and analyze UV transit light curves.  This method is based on tracing the emitted flux from the stellar atmosphere through the atmosphere of the planet and the interstellar medium to the observing instrument at Earth.  We have demonstrated the method in practice by applying it to the existing UV transit depth measurements of HD209458b in the H~Ly$\\alpha$ and O~I triplet lines (VM03,VM04,BJ10).  In order to interpret the measured transit depths we used a simple empirical model of the thermosphere based on the generic features of more complex models \\citep[e.g.,][]{yelle04,garciamunoz07,koskinen10a} to simulate absorption by the occulting atmosphere.  The H~I transit depth is sensitive to interstellar absorption and the full-width transit depth reflects the optical depth of the atmosphere in the wings of the H~Ly$\\alpha$ line.  We found that it is easy to generate models that explain the observed absorption without the need to introduce external ENAs or hydrogen atoms accelerated by radiation pressure outside the Roche lobe.  According to the best-fitting models (see models M1 and M7 in Table~\\ref{table:models}), the mean temperature of the thermosphere is $T =$~8000--11,000~K and the H$_2$/H dissociation front is located at $p_b =$~0.1--1~$\\mu$bar ($z_b \\approx$~1.1~R$_p$).  The upper boundary of the model is located at $Z_e \\sim$~2.9~R$_p$, which is near the boundary of the Roche lobe.  Below $Z_e$ the density profiles are approximately hydrostatic and above $Z_e$ the atmosphere is mostly ionized.  By using the M1 model, we obtain a full-width H~I transit depth of 6.6~\\% and the model transmission of the H~Ly$\\alpha$ line matches the data points reported by BJ08.  The apparent disagreement between BJ08 and the earlier analysis by VM03 arises from differences in the treatment of the data and different definitions of the transit depth.  In particular, there is no definite observational evidence for a cometary tail following the planet in the transit light curve and the absorption is not significantly stronger in the blue side of the H~Ly$\\alpha$ line.  However, we agree with VM08 that the core of the absorption line is optically thick up to the Roche lobe and thus that the atmosphere is evaporating.  We estimate a mass loss rate of $dM/dt \\approx$10$^{7}$--10$^{8}$~kg~s$^{-1}$ based on the density profile of the M1 model and a range of possible radial velocities.  It is possible that ENAs are present outside the Roche lobe as suggested by E10, but the optical depth of the ENA clouds has to be consistent with the underlying models of the thermosphere.  We do not agree with E10 that the optical depth of the thermosphere below $z =$~2.8~R$_p$ is small.  We do believe, however, that one-dimensional hydrodynamic models are not adequate in modeling the distribution of ionized gases outside the Roche lobe.  Multidimensional plasma models are required to study the interaction of the ionosphere with the stellar wind and magnetospheric plasma self-consistently.  The mean O~I transit depth of 10.5~\\% is only 1.93$\\sigma$ away from the FUV continuum transit depth of $\\sim$2~\\%.  By using the M1 model, we obtain an O~I transit depth of 4.3~\\%, which is marginally consistent with the observations.  Full-width O~I transit depths that are comparable or slightly higher than the H~I depth are possible if the abundance of oxygen in the thermosphere is supersolar, the atmosphere is escaping with supersonic velocities or if large external clouds of neutral oxygen are present above the Roche lobe.  More precise measurements are required to constrain the O~I transit depth further.  Out of the currently available instruments, the STIS G140M grating can be used to obtain repeated measurements with higher S/N and spectral resolution."
    },
    "1004/1004.1991_arXiv.txt": {
        "abstract": "{To determine the physical parameters of a transiting planet and its host star from photometric and spectroscopic analysis, it is essential to independently measure the stellar mass. This is often achieved by the use of evolutionary tracks and isochrones, but the mass result is only as reliable as the models used.} {The recent paper by Torres et al (2009) showed that accurate values for stellar masses and radii could be obtained from a calibration using $T_{eff}$, log $g$ and $[Fe/H]$. We investigate whether a similarly good calibration can be obtained by substituting log $\\rho$ - the fundamental parameter measured for the host star of a transiting planet - for log $g$, and apply this to star-exoplanet systems.} {We perform a polynomial fit to stellar binary data provided in Torres et al (2009) to obtain the stellar mass and radius as functions of $T_{eff}$, log $\\rho$ and $[Fe/H]$, with uncertainties on the fit produced from a Monte Carlo analysis. We apply the resulting equations to measurements for seventeen SuperWASP host stars, and also demonstrate the application of the calibration in a Markov Chain Monte Carlo analysis to obtain accurate system parameters where spectroscopic estimates of effective stellar temperature and metallicity are available.} {We show that the calibration using log $\\rho$ produces accurate values for the stellar masses and radii; we obtain masses and radii of the SuperWASP stars in good agreement with isochrone analysis results. We ascertain that the mass calibration is robust against uncertainties resulting from poor photometry, although a good estimate of stellar radius requires good-quality transit light curve to determine the duration of ingress and egress.} {} ",
        "introduction": "There are currently over 400 known exoplanets, of which more than 60 transit their host stars\\footnotemark \\footnotetext[1]{www.exoplanet.eu}. This important transiting subset are the only planets for which the orbital inclination, and hence precise stellar and planetary parameters, may be determined. The fundamental parameters found for the host star and transiting planet are stellar density (see below) and planetary surface gravity \\citep{southworth04}. To convert these into values for the radii of both, it is necessary to find the stellar mass. This is often arrived at iteratively via deriving a stellar density from the lightcurve analysis and a stellar effective temperature from spectroscopy and using these with model evolutionary tracks and isochrones of appropriate metallicity to find a stellar mass and age \\citep{sozzetti07}. Further photometric and spectroscopic analysis may then be performed to arrive at final values for the masses and radii of the star and planet, see e.g. \\citet{hebb09}. The resulting values for masses and radii are therefore only as reliable as the evolutionary models used. A recent study by \\citet{southworth09} highlighted the fact that discrepancies between different sets of evolutionary models represent the dominant source of systematic uncertainty in planetary parameters. For example, they find that the spread of mass values obtained for HD 209458 using different models is around $4\\%$.  Here we develop a new one-step approach to determining the masses of exoplanet host stars from their effective temperatures, metallicities and photometric bulk densities. We base our method on the recent study by \\citet{torres09} of the masses and radii of a large sample of well-characterised main-sequence stars belonging to non-interacting, eclipsing spectroscopic binaries. \\citet{torres09} showed that accurate stellar masses and radii could be obtained using a calibration of stellar surface gravity, effective temperature and metallicity. They used a set of well-determined measurements of log $g$, $T_{eff}$, $[Fe/H]$, $M$ and $R$ from binary stars to obtain coefficients that allow mass and radius to be calculated directly for any normal star, without isochrone fitting.  Recently, the use of log $\\rho$ in place of log $g$ in the determination of star-planet system parameters has become widespread, see for example \\citet{sozzetti07}, \\citet{winn08b}, \\citet{sozzetti09} and \\citet{fernandez09}. Where high quality photometric data can be obtained of the transit event, the stellar parameters can be obtained more precisely using the stellar density value derived from the lightcurve than using the stellar surface gravity value from spectral analysis \\citep{sozzetti07}.  In Section 2 we review the methodology for determining exoplanet host-star densities from the transit geometry. We re-determine the mass and radius calibrations of \\citet{torres09} using their data, and obtain comparably tight mass and radius calibrations using log $\\rho$ in place of log $g$. In Section 3 we apply the method to the host stars of several transiting planets for which isochrone mass determinations have been published recently. In Section 4 we show how the method can be incorporated directly in a Markov-chain Monte Carlo (MCMC) analysis, to give the stellar mass as a derived parameter. ",
        "conclusions": "We have presented a new calibration for stellar masses and radii based on stellar effective temperature, metallicity and stellar density. We have shown that the resulting equations provide a good fit to data for 38 stars from \\citet{torres09}, and also to values for masses and radii of exoplanet host stars obtained from isochrone analyses. We have demonstrated that accurate stellar masses may be obtained for such exoplanet host stars via a Markov-chain Monte Carlo analysis of photometric and spectroscopic data, using spectroscopically determined temperatures and metallicities as input.  Even where poor photometry yields an uncertain estimate of stellar density, the mass estimate from the calibration is encouragingly robust. However, the stellar radius depends strongly on the stellar density estimate which in turn requires good knowlege of the impact parameter. Thus in establishing planet radii there is no substitute for good quality photometry, though the Main Sequence prior can provide a useful additional constraint if the star can be shown via independent means to be unevolved."
    },
    "1004/1004.0503_arXiv.txt": {
        "abstract": "We have examined trapping of two-armed nearly vertical oscillations in polytropic disks. Two-armed nearly vertical oscillations are interesting in the sense that they are trapped in an inner region of disks with proper frequencies, if the inner edge of disks is a boundary that reflects oscillations. The frequencies of the trapped oscillations cover the frequency range of kHz QPOs to low frequency QPOs in LMXBs, depending on the modes of oscillations. Low frequency trapped oscillations are particularly interesting since their trapped region is wide. These low frequency oscillations are, however, present only when $\\Gamma(\\equiv 1+1/N)$ is close to but smaller than 4/3 (when spin parameter $a_*$ is zero), where $N$ is the polytropic index. The above critical value 4/3 slightly increases as $a_*$ increases. ",
        "introduction": "One of the possible candidates of high-frequency quasi-periodic oscillations (HFQPOs) observed in low-mass X-ray binaries is resonantly excited p- and/or g-mode oscillations in deformed disks (Kato 2004, 2008a,b; Ferreira and Ogilvie 2008; Oktariani et al. 2010). The disk deformation required in this model is a warp or an eccentric deformation in equatorial plane. Recent numerical MHD simulations (Henisey et al. 2009) seem to have preliminary confirmed the presence of the excitation mechanism.  The high frequency QPOs observed in black-hole candidates appear in pairs with frequency ratio of 3 : 2. This characteristic of high-frequency QPOs in black-hole candidates seems to be described by the above model, if we assume that geometrically thin disks are surrounded by hot tori and that the QPOs photons are disk photons Comptonized in the tori (Kato and Fukue 2006). The twin kHz QPOs observed in neutron stars, on the other hand, have always not the 3 : 2 frequency ratio. They change their frequencies with time with correlation. If we want to describe such characteristics of kHz QPOs by the above model, a warp or eccentric disk deformation must have a time-dependent precession. In the case of neutron stars, distinct from the case of black holes, such precession of the deformation might be expected, since the central stars have surfaces and this might become causes of time-dependent precession through magnetic and radiative couplings between the central sources and the disks. It is not clear, however, whether possible time-dependent precession of deformation has a time scale consistent with the time variations of the kHz QPOs.  In this context, it is worthwhile examining whether there are other kinds of disk deformations that can become a possible source of resonant excitation of high-frequency oscillations in disks. This problem has been examined by Kato (2009), and it was suggested that two-armed nearly vertical oscillations\\footnote{ There will be no commonly accepted classification and terminology concerning disk oscillations. Here, we classify oscillations into  four types, i.e., p-mode, g-mode, c-mode and vertical p-mode oscillations (see, for example, Kato 2001; Kato et al. 2008). The oscillations considered here are the vertical p-mode oscillations. } can excite high-frequency p- and/or g-mode oscillations. Hence, it will be interesting to examine whether there are two-armed nearly vertical oscillation modes that can become a global deformation of disks.  More interestingly, the two-armed vertical disk oscillations themselves may be origins of the varous types of QPOs observed in neutron-star LMXBs. This is because two-armed nearly vertical osillations occur in the inner region of disks and cover a wide range of frequency by difference of modes.  Based on these considerations, we examine in this paper basic properties of the two-armed vertical disk oscillations. That is, we examine their characteristics of trapping, eigen-frequency and its dependence on polytropic index specifying the vertical disk structure, using the mathematical formulations already prepared by Silbergleit et al. (2001) to study the trapping of the corrugation waves (c-mode oscillations). ",
        "conclusions": ""
    },
    "1004/1004.2466.txt": {
        "abstract": "Interferometric observations of the W33A massive star-formation region, performed with the Submillimeter Array (SMA) and the Very Large Array (VLA) at resolutions from 5\\arcsec ~ (0.1 pc) to 0.5\\arcsec ~  (0.01 pc) are presented.  Our three main findings are: (1) parsec-scale, filamentary structures of cold molecular gas are detected. Two filaments at different velocities intersect in the zone where the star formation is occurring. This is consistent with triggering of the star-formation activity by the convergence of such filaments, as predicted by numerical simulations of star formation initiated by converging flows. (2) The two dusty cores (MM1 and MM2) at the intersection of the filaments  are found to be at different evolutionary stages, and each of them is resolved into multiple condensations. MM1 and MM2 have markedly different temperatures, continuum spectral indices, molecular-line spectra, and masses of both stars and gas. (3) The dynamics of the ``hot-core'' MM1 indicates the presence of a rotating disk in its center (MM1-Main) around a faint free--free source. The stellar mass is estimated to be $\\sim 10~\\Msun$. A massive molecular outflow is observed along the rotation axis of the disk. ",
        "introduction": "\\label{intro}  Stars form by accretion of gas in dense molecular-cloud cores. However, the differences, if any, in the details of the formation process of massive stars (those with roughly $M_\\star > 8 ~ M_\\odot$) compared to low-mass stars are not well understood. Recent reviews on the topic are those by \\cite{Beu07}, and \\cite{ZY07}.  We are carrying out a program aimed at studying how the formation of massive stars in clusters proceeds in the presence of different levels of ionization, from the onset of detectable free--free emission to the presence of several bright ultracompact (UC) \\HII regions. In this paper we present our first results on the massive star-formation region W33A (also known as G12.91-0.26), at a kinematic distance of 3.8 kpc \\citep{Jaffe82}. W33A is part of the W33 giant \\HII region complex \\citep{W58}. It was recognized as a region with very high far infrared luminosity ($\\approx 1 \\times 10^5$ $L_\\odot$ ), but very faint radio-continuum emission by \\cite{Stier84}. \\cite{vdT00} modeled the large-scale (arcminute) cloud as a spherical envelope with a power-law density gradient, based on single-dish mm/submm observations. Those authors also presented mm interferometric observations at several-arcsecond resolution that resolved the central region into two dusty cores separated by $\\sim 20,000$ AU. The brightest mm core contains faint ($\\sim 1$ mJy at cm wavelengths) radio-continuum emission \\citep{RH96} resolved at 7 mm into possibly three sources separated by less than $1\\arcsec$ ($\\approx 4000$ AU) from each other \\citep{vdTM05}. These radio sources were interpreted by \\cite{vdTM05} as the gravitationally trapped \\HII regions set forth by \\cite{Keto03}. However, the earlier detection by \\cite{Bunn95} of near-infrared recombination line (Br$\\alpha$)  emission with FWHM $=155~\\kms$ suggests that at least some of the radio free--free emission is produced by a fast ionized outflow. More recently, \\cite{Davies10} reported spectroastrometry observations of Br$\\gamma$ emission toward W33A. The Br$\\gamma$ emission appears to be produced by at least two physical components: broad line wings extending to a few hundreds of kilometers per second from the systemic velocity appear to trace a bipolar jet on scales of a few AU, while the narrow-line emission may be attributed to a dense \\HII region \\citep{Davies10}. Being a bright mid- and far-infrared source, W33A has also been target of interferometry experiments at these wavelengths, which reveal density gradients and non-spherical geometry in the warm dust within the inner few hundred AU \\citep{dW07,dW09}.  Here we report on millimeter and centimeter interferometric observations performed with the Submillimeter Array (SMA) and the Very Large Array (VLA) at angular resolutions from $\\sim 5\\arcsec$ to $0.5\\arcsec$. We find a massive star-forming cluster embedded in a parsec-scale filamentary structure of cold molecular gas. The dense gas is hierarchically fragmented into two main dusty cores, each of them resolved into more peaks at our highest angular resolution. The main cores appear to be at different evolutionary stages, as evidenced from their differing spectra, masses, temperatures, and continuum spectral indices. The warmer core harbors faint free--free emission centered on a rotating disk traced by warm molecular gas. The disk powers a massive molecular outflow, indicating active accretion. In Section 2 of this paper, we describe the observational setup. In Section 3 we list our results, in Section 4 we present a discussion of our findings, and in Section 5 we give our conclusions. ",
        "conclusions": "\\label{conclu}  We present for the first time resolved observations in both mm continuum and molecular-line emission for the massive star formation region W33A, characterized by a very high luminosity ($L\\sim10^5~L_\\odot$) and very low radio-continuum emission ($\\sim 1$ mJy). Both of the previously known mm cores (MM1 and MM2) are resolved into multiple peaks, and appear to be at very different evolutionary stages, as indicated by their molecular spectra, masses, temperatures, and continuum spectral indices. The brightest core (MM1-Main at the center of MM1) is centered on a very faint free--free source and the gas dynamics up to a few thousand AU of it indicates the presence of a circumstellar disk rotating around a stellar mass of $M_\\star \\sim 10~\\Msun$. MM1-Main also drives a powerful, high-velocity molecular outflow perpendicular to the disk. MM2, the coldest and most massive core, is not detected in hot-core lines but appears to drive a more modest outflow. Both MM1 and MM2 are located at the intersection of parsec-scale filamentary structures with line-of-sight velocity offset by $\\approx 2.6~\\kms$. Analysis of the position--position--velocity structure of these filaments and a comparison with recent numerical simulations suggests that star formation in W33A was triggered by the convergence of filaments of cold molecular gas."
    },
    "1004/1004.0818_arXiv.txt": {
        "abstract": "The non-Gaussian distribution of primordial perturbations has the potential to reveal the physical processes at work in the very early Universe. Local models provide a well-defined class of non-Gaussian distributions that arise naturally from the non-linear evolution of density perturbations on super-Hubble scales starting from Gaussian field fluctuations during inflation. I describe the $\\delta N$ formalism used to calculate the primordial density perturbation on large scales and then review several models for the origin of local primordial non-Gaussianity, including the cuvaton, modulated reheating and ekpyrotic scenarios. I include an appendix with a table of sign conventions used in specific papers. ",
        "introduction": "The common presumption that primordial density perturbations have a Gaussian distribution is a powerful simplifying assumption that allows one to specify all the properties of the distribution once the two-point correlation function is known in real space, or equivalently the power spectrum in Fourier space. In particular the three-point and connected higher moments of the distribution vanish. On the other hand the statement that a distribution is non-Gaussian opens up an infinite array of possibilities. This has led to an assortment of empirical tests for non-Gaussianity of the primordial perturbations. By contrast there are relatively few non-Gaussian distributions that are motivated by theoretical models for the origin of structure in the early universe. While one can argue that a Gaussian distribution could describe density perturbations arising from a wide range of possible sources, any detection of deviations from a Gaussian distribution predicted by a specific theoretical model would be strong evidence in support of that model.  Vacuum fluctuations in light, weakly-coupled scalar fields during a period of inflation (defined here as accelerated expansion, $\\ddot{a}>0$) in the very early universe provide a natural origin for an almost Gaussian distribution of field perturbations on large scales. Fluctuations in a free quantum field on small scales with comoving wavenumber, $k$, are swept up to scales much larger than the comoving Hubble scale, $H^{-1}/a=1/\\dot{a}$, which shrinks during inflation. On super-Hubble scales ($k<aH$) damping drives the fluctuations into a squeezed state and they can effectively be treated as a classical random distribution \\cite{Polarski:1995jg}. Models of slow-roll inflation in the early universe are driven by canonical fields whose self-interactions are small relative to the Hubble scale and hence their distribution on scales close to the Hubble scale remains well described by a Gaussian distribution with vanishing third- and higher-order connected functions \\cite{Maldacena:2002vr,Seery:2005gb}.  As inflation continues to stretch the perturbations up to scales far larger than the Hubble scale, spatial gradients are expected to become negligible and the we may treat the evolution locally as effectively ``separate universes'' \\cite{Salopek:1990jq,Wands:2000dp} whose initial conditions are set by the Gaussian distribution of the scalar field values during inflation. This motivates the study of local models of primordial non-Gaussianity where the distribution of primordial density perturbations, $\\zeta$, can be described by a local function of one or more Gaussian random fields, $\\zeta(\\delta\\varphi^I)$. A linear function of Gaussian fields is itself Gaussian, so a non-Gaussian distribution implies non-linearity and the study of non-linear cosmological perturbations. This local model for non-Gaussianity turns out to be a very good description of non-Gaussianity in some simple physical models for the origin of structure in the very early universe.  Conventionally the primordial perturbations are characterised by the metric potential in the matter-dominated era, $\\Phi=(3/5)\\zeta$. In Fourier space we define the power spectrum and bispectrum as \\bea \\langle\\Phi_{\\bki}\\Phi_{\\bkii} \\rangle = (2\\pi)^3 P_{\\Phi}(k_1) \\delta^{3}(\\bki+\\bkii) \\,,\\\\ \\label{bispectrum} \\langle \\Phi_{\\bki} \\Phi_{\\bkii} \\Phi_{\\bkiii} \\rangle = (2\\pi)^3 B_{\\Phi}(k_1,k_2,k_3) \\delta^{3}(\\bki+\\bkii+\\bkiii) \\,, \\eea and the amplitude of the bispectrum relative to the power spectrum is then given by the dimensionless parameter \\be \\label{deffnl} \\fnl(k_1,k_2,k_3) \\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_3) P_{\\Phi}(k_1) \\right]} \\,. \\ee  In a perturbative expansion, the second-order expression for the metric potential as a local function of a single Gaussian field, $\\phi$, is given by~\\cite{Komatsu:2001rj} \\be \\label{originalPhi} \\Phi(\\bx) = \\phi(\\bx) + \\fnl \\left( \\phi^2(\\bx) - \\langle \\phi^2 \\rangle \\right) +\\ldots \\,. \\ee In this special case $\\fnl$ is, by construction, a constant parameter independent of spatial position or scale. However in more general local models the ratio of the bispectrum to the power spectrum given by Eq.~(\\ref{deffnl}) may be taken as a definition of $\\fnl$ which is then scale and shape-dependent. Note that following Komatsu and Spergel~\\cite{Komatsu:2001rj} we adopt a sign convention for the metric potential $\\Phi$, and hence $\\fnl$, which is the opposite of that used by, for example, Mukhanov et al~\\cite{Mukhanov:1990me} and Maldacena \\cite{Maldacena:2002vr}. Specific sign conventions are summarised in a table in an Appendix.  In section 2 I review the $\\delta N$-formalism that is commonly used to calculate the primordial density perturbation on large scales and its higher-order correlations, including the use of Feynman-type diagrams and how the $\\delta N$ formalism extends to the description of primordial isocurvature density perturbations. In section 3 I review several examples of models for the origin of structure in the very early universe and the local non-Gaussianity they give rise to. I conclude in section 4 with brief review of current observational constraints. ",
        "conclusions": "The distribution of primordial density perturbations contains more information than just the power-spectrum. Among a plethora of different possible forms of non-Gaussianity the local models provide a clearly defined class which arise naturally from the evolution of the density perturbation on super-Hubble scales from initially Gaussian vacuum fluctuations during inflation. In particular the $(2+n)$-point function of a non-Gaussian distribution generated by a local function of a single Gaussian field can be described at leading order by the two-point function and $n$ non-linearity parameters, $\\fnl$, $\\gnl$, etc.  Non-Gaussianities have the potential to reveal the physical interactions at work in the very early Universe. For adiabatic perturbations during inflation the non-linearity parameters are related to the scale-dependence of the power spectrum \\cite{Maldacena:2002vr,Creminelli:2004yq} and hence must be small. Thus any detection of local non-Gaussianity of the primordial density perturbation would be evidence of non-adiabatic perturbations on super-Hubble scales, e.g., the presence of multiple light fields during slow-roll inflation.  Current observational limits on $\\fnl$ in the single-field local model come from combining WMAP cosmic microwave background bounds ($-10<\\fnl<74$ \\cite{Komatsu:2010fb}) with large-scale structure ($-29<\\fnl<70$ \\cite{Slosar:2008hx}), yielding $-5<\\fnl<70$ at 95\\% confidence limit~\\cite{Komatsu:2010fb}. Recent work also places a bound $-3.80<\\gnl/10^6<3.88$ from the WMAP 5-year data \\cite{Smidt:2010sv}. Care should be taken when applying these bounds to non-local or even multi-variate local models. CMB bounds are obtained using matched filtering techniques to construct an optimal estimator for $\\fnl$ defined by Eq.~(\\ref{originalPhi}). It estimates the amplitude of a second-order primordial perturbation which has a specified correlation with the first-order perturbation. On the other hand bounds from large scale structure arise from the form of the (general relativistic) Poisson equation \\cite{Wands:2009ex} which implies that a local form for the primordial Newtonian potential requires a non-local form for the comoving density contrast, leading to a distinctive scale-dependent bias \\cite{Dalal:2007cu}.  There is great potential for future discovery coming from future CMB missions such as ESA's Planck satellite which has the potential to bound $|\\fnl|<5$ and large-scale structure surveys, including 21cm line radio surveys \\cite{Komatsu:2009kd}. There is also plenty of scope for further theoretical developments given the data that already exists to test non-Gaussian models. New theoretical templates will need to be derived for optimised constraints to be placed on alternative theoretical models of non-Gaussianity. At some point primordial non-Gaussianity will be discovered as it inevitably arises from the non-linear evolution of density perturbations and the last-scattering of the CMB photons \\cite{Bartolo:2004if,Pitrou:2010sn}, and it will then be a question of using all the available data to disentangle the different contributions."
    },
    "1004/1004.2506_arXiv.txt": {
        "abstract": "We investigate statistics of the decay process in the equal-mass three-body problem with randomized initial conditions. Contrary to earlier expectations of similarity with ``radioactive decay'', the lifetime distributions obtained in our numerical experiments turn out to be heavy-tailed, i.e. the tails are not exponential, but algebraic. The computed power-law index for the differential distribution is within the narrow range, approximately from $-1.7$ to $-1.4$, depending on the virial coefficient. Possible applications of our results to studies of the dynamics of triple stars known to be at the edge of disruption are considered. ",
        "introduction": "Disruption of a three-body gravitational system is an enigmatic dynamical process, statistics of which is mostly unexplored, at long timescales especially. \\cite{V88} supposed that the lifetime distribution for a three-body system is an exponentially decaying function, in analogy with ``radioactive decay''. Recent statistical numerical studies by \\cite{MT07} of this process in the equal-mass problem revealed new important data on the statistics of the disruption times and raised new questions on the nature of this process. \\cite{MT07} have explored a statistics of the disruption times $T_\\mathrm{d}$ in the equal-mass three-body problem. (The disruption time $T_\\mathrm{d}$ is the system lifetime as a bound system.) They computed the disruption times for equal-mass systems with randomized initial conditions, and fitted the found disruption time distributions by exponential functions.  However, exponential decay, in the asymptotically long time scale, is in conflict with the theoretical result of \\cite{A83} that the average lifetime of a general isolated three-body system is infinite. What is more, recently it was shown by \\cite{S09} on the basis of the Kepler map theory, applied to a hierarchical three-body problem, that at the edge of disruption the orbital periods and the size of the orbit of the escaping body behave as L\\'evy flights. Due to them, the time decay of the survival probability (the integral distribution of lifetimes) is heavy-tailed with the power-law index equal to $-2/3$. Combining the Kepler map theory and earlier theoretical findings of \\cite{H93}, \\cite{S09} made a conclusion that the $T_\\mathrm{d}^{-2/3}$ law is expected to be quite universal.  Here we explore the problem of the lifetime statistics in the three-body problem by means of numerical simulations. We pay particular attention to analysis of the tails of the lifetime distributions, and find the algebraic behaviour. We compare our results with previous numerical work and discuss why the algebraic tails had not been identified earlier. Possible applications of our results to studies of the dynamics of triple stars known to be at the edge of disruption are considered. ",
        "conclusions": "We have investigated statistics of the disruption times $T_\\mathrm{d}$ in the equal-mass three-body problem. The statistics has been explored in massive numerical simulations with randomized initial conditions. The distributions of the disruption times have turned out to be heavy-tailed with the power-law index being within narrow range near the value of $-5/3$ (for the differential distribution). The observed range is from $-1.7$ to $-1.4$; the value of the slope index slightly depends on the virial coefficient.  Our result is in conflict with earlier expectations and analogies with ``radioactive decay'', which implied exponential tails, and also in conflict with numerical results of \\cite{MT07}, where exponential decay was found. We explain the divergence with the latter results by the fact that \\cite{MT07} fitted the distribution as a whole (and what is more, taken at insufficiently long timescale), while the behaviour of the tails should be treated separately from the initial exponential drop.  On the other hand, our finding of the algebraic decay is in agreement with the theoretical result of \\cite{A83}, stating that the average lifetime of a general isolated three-body system is infinite. Moreover, it is in agreement with the Kepler map theory \\citep{S09}, predicting just the same value of the slope index ($\\sim -5/3$) that we have observed in the majority of our numerical experiments. This makes evident the existence of L\\'evy flights in the process of decay.  At intermediate values of the virial coefficient, $k \\sim 0.5$, we have observed the slope index $\\sim -1.5$ (for the differential distribution). This is typical for free diffusion in the central part of a chaotic layer until the finite width of the layer becomes important. We believe that the free diffusion predominance is merely an effect of insufficient times of integration. If one extends the integration time to sufficient limits, the asymptotic behaviour with the slope index $\\sim -5/3$, characteristic of the L\\'evy flights in the given problem, should become apparent.  Several actual multiple stars are known nowadays to be near the edge of the stability region in the phase space of the motion, in particular, HD~40887, HD~76644 (ADS~7114$=$ $\\iota$~Uma), HD~217675 ($o$~And), HD~222326 (ADS~16904) \\citep{Or05}. Numerical integrations, based on initial data from further high-precision astrometric observations of these systems, would allow one to reveal long-term evolution features, in particular, the theoretically predicted L\\'evy flights in the orbital periods and the geometrical sizes of the systems."
    },
    "1004/1004.2730_arXiv.txt": {
        "abstract": "Orbital parameters of binary radio pulsars reveal the history of the pulsars' formation and evolution including dynamic interactions with other objects. Advanced technology has enabled us to determine these orbital parameters accurately in most of the cases. Determination of post-Keplerian parameters of double neutron star binaries (especially of the double pulsar system PSR J0737-3039) provide clean tests of the general theory of relativity and in the future may lead us to constrain the dense matter Equations of State. For binary pulsars with main-sequence or white dwarf companions, knowledge about the values of the orbital parameters as well as of the spin periods and the masses of the pulsars and the companions might be useful to understand the evolutionary history of the systems. As accreting neutron star binaries lead to orbit circularization due to the tidal coupling during the accretion phase, their descendants $i.e.$ binary millisecond pulsars are expected to be in circular orbits. On the other hand, dense stellar environments inside globular clusters cause different types of interactions of single stars with pulsar binaries like ``fly-by\", ``exchange\", and ``merger\". All these interactions impart high eccentricities to the pulsar binaries. So it is quite common to get eccentric millisecond pulsar binaries in globular clusters and we find that ``fly-by\" causes intermediate values of eccentricities while ``exchange\" or ``merger\" causes high values of eccentricities. We also show that ``ionization\" is not much effective in the present stage of globular clusters. Even in the absence of such kinds of stellar interactions, a millisecond pulsar can have an eccentric orbit as a result of Kozai resonance if the pulsar binary is a member of a hierarchical triple system. PSR J1903+0327 is the only one eccentric millisecond pulsar binary in the galactic disk where stellar interactions are negligible. The possibility of this system to be a member of a hierarchical triple system or past association of a globular cluster have been studied and found to be less likely. ",
        "introduction": "\\label{sec:intro}  Radio pulsars are rotating, magnetized neutron stars (NSs) emitting radio beams from magnetic poles. Their masses are in the range of $1-2~M_{\\odot}$, radii $\\sim10$ km, densities $\\sim10^{14}~{\\rm g~cm^{3}}$ and surface values of the dipolar magnetic field are around $10^{8}-10^{14}$ gauss. Presently there are 1864 radio pulsars with their spin periods covering a wide range 1.39 ms $-$ 11.78 s \\footnote{http://www.atnf.csiro.au/research/pulsar/psrcat/}. This number is increasing quite rapidly after the first serendipitous discovery in 1967 by Jocelyn Bell, a graduate student in Cambridge \\cite{joc68}.  Pulsars are believed to be born with spin period $\\sim 0.1$ s which increases very slowly with time (at the rate of $10^{-20}-10^{-14}~{\\rm s \\, s^{-1}}$) as their rotational energy is emitted in the form of electromagnetic energy. Millisecond spin periods are explained as a result of spin-up process, which occurs when a neutron star has a normal star as its binary companion and accretes matter from the companion (at its red giant phase) causing transfer of angular momentum. Depending upon the evolution of the companion, the resultant millisecond pulsar can be either isolated or a member of a binary system. In this article we concentrate on binary pulsars and discuss how the knowledge of their orbital parameters help us to understand the structure, formation, evolution and dynamic history of pulsars. As an example, binary millisecond pulsars are expected to be in circular orbits as tidal coupling during the accretion phase lead to orbit circularization. But if there is no history of accretion phase, the orbit of the binary pulsar can be eccentric because of the asymmetric kick imparted to the pulsar during its birth in a supernova. A millisecond pulsar binary can also have an eccentric orbit if it experiences interactions with other near-by stars (which is possible in dense stellar environments like globular clusters) or if it is a member of a hierarchical triple.  It is noteworthy to mention here that in addition to radio wavelengths, pulsars emit in x-rays and/or gamma rays too. But radio pulsars are the most useful astronomical objects firstly because the underlying neutron stars are more stable than in the case of X-ray pulsars which have episodically varying accretion torques; secondly high sensitivities of radio telescopes and data analysis software make it possible to determine their spin periods very accurately. Moreover, for radio pulsars in binary systems ($\\sim 8\\%$ of presently known radio pulsars are in binary systems), it is possible to determine the mass of the pulsar and its companion, size and shape of the orbit $etc$ by modelling Keplerian and post-Keplerian parameters. Such accurate determination of those parameters are difficult for X-ray pulsars at present. That is why we concentrate only on binary radio pulsars when we try to use the knowledge about their orbital parameters to understand the properties of pulsars. The first binary pulsar PSR B1913+16 was discovered in 1974 at Arecibo observatory by R. Hulse and J. H. Taylor \\cite{hultal74} which is a 59 ms pulsar having orbital period 0.323 days and orbital eccentricities 0.671. Interestingly, its $\\sim 1.06~M_{\\odot}$ mass companion star is most probably another neutron star $i.e.$ the first discovered binary pulsar is also the first discovered double neutron star (candidate) system. Afterwords nine other double neutron star (DNS) systems have been discovered. Among these ten DNSs, seven systems are confirmed to be DNSs and the rest are most likely to be DNSs \\cite{dunc08}. These DNSs are useful tools to test the validity of the general theory of relativity \\cite{ingrid03} and even in principle to constrain the dense matter Equations of State \\cite{lat06}. In Section \\ref{sec:dns} we discuss how the study of DNSs can lead to a better understanding of the structure of neutron stars ($i.e$ the Equation of State), the physics of the pulsar magnetosphere and the formation scenario for the special case of the double pulsar system.  Most of the binary pulsars have either main sequence (MS) or white dwarf (WD) companions. Knowledge about the values of the orbital parameters as well as of the spin periods and the masses of these pulsars (and their companions) might be useful to understand the evolutionary history of these systems. The possible causes of eccentricities of millisecond pulsar binaries have been studied in this article. When these systems are in dense stellar environments as in globular clusters, the eccentricities can arise as a result of different types of interactions of single stars with pulsar binaries like ``fly-by\", ``exchange\" and ``merger\" (Section \\ref{sec:gc}). The eccentricity of the only eccentric millisecond pulsar PSR J$1903+0327$ in the galactic disk is difficult to explain as the hierarchical triple scenario faces difficulties to satisfy observed parameters of this system and the globular cluster origin also appears to be unlikely (section \\ref{sec:triple}). All the findings of the present work have been summarized in section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  In this article we discuss how orbital parameters of binary radio pulsars reveal the history of their formation and evolution including dynamic interactions with other objects, the physics of the matter at extreme densities as inside the pulsar as well as the properties of the pulsar magnetosphere. Firstly we try to understand the importance of studying DNSs. If at least one of the NS of a DNS is a radio pulsar, we can determine the moment of inertia of the pulsar through accurate radio timing analysis by measuring the extra advancement of the periastron of the orbit due to the spin-orbit coupling over the standard post-Newtonian one. Then by matching this observed value of the moment of inertia with the theoretically predicted value, one can constrain the dense matter EoS. One necessary condition for this to be done is that the pulsar must be rapidly spinning and its companion NS must be much slower; the double pulsar system PSR J0737-3039(A/B) seems to be the most potential candidate for this purpose till date as we know that PSR J0737-3039A has $P_{s} = 22.7 $ ms whereas PSR J0737-3039B has $P_{s} = 2.77 $ s. Moreover the masses of both the NSs are well determined. We also checked how good the system really is for this purpose by studying the effect of various parameters on the spin-orbit coupling term as a larger value of this term will facilitate the moment of inertia measurement. We found that this term has larger effect if the following conditions are satisfied - (i) short spin period  (ii) short orbital period, (iii) high eccentricity.  Unfortunately, PSR J0737-3039 lacks conditions (iii) and even (i) to some extent.  So it will be nice if we can find better systems satisfying all the above conditions. Globular clusters are potential sources for favorable DNS systems in the above sense as DNSs can be formed by exchange collisions due to the high stellar density and the abundance of millisecond pulsars in globular clusters. On the contrary, only one DNS has been found in a GC so far which might be a result of selection effects as GCs are usually at larger distances in comparison to the disk pulsars. The resulting binaries are more probable to be highly eccentric with a millisecond pulsar facilitating  moment of inertia estimation. We have also discussed how the study of the eclipse of the radio beam of  PSR J0737-3039A by the magnetosphere of  PSR J0737-3039B can be used as a tool to understand the magnetospheric properties of a pulsar in a better extent. Some suggestive works about the formation of PSR J0737-3039A/B have been mentioned too.  We explain the presently observed orbital eccentricity and period data of GC binary pulsars by stellar interaction scenarios of fly-bys, exchanges and mergers with field stars. Binaries with $e > 0.1$ are most probably the result of exchange or merger events whereas binaries with $0.01 > e > 0.00001 $ are products of fly-by of single stars. A number of wide orbit intermediate eccentricity pulsars seen in the galactic disk are absent in the GC sample because they have been kicked up to relatively high eccentricities by passing stars in the dense stellar environments in GCs. In some GCs such as Ter 5, the stellar densities are so high, and the velocity dispersion so modest that the interaction timescale for exchange and fly-by interactions is relatively short. In such GCs a typical binary system may undergo multiple interactions. If the original binary contains a spun-up millisecond pulsar in a relatively ``soft\" binary, then the exchange interaction may even produce a single millisecond pulsar in the cluster. This may explain the higher incidence of isolated millisecond pulsars in GCs compared to that in the galactic disk \\cite{cam05}; in GCs 52 out of 59 isolated pulsars are millisecond pulsars whereas in the disk 23 out of 1565 isolated pulsars are millisecond pulsars. In addition, exchange interactions, as we have seen, can lead to highly eccentric orbits and the system can be ejected from the cluster core. If the last encounter took place not too long ago, the system can be at a relatively large offset from the cluster core, albeit being still spatially co-located with the GC. We have found that the eccentricity of PSR B1638+36B (in M13) is difficult to explain by any type of interactions - fly-by, exchange or merger. So we propose that either its true eccentricity is  lower than the presently known upper limit or it is a member of a hierarchical triple. We have also considered the effects of collision induced ionization on the present day distribution of orbital parameters of radio pulsars in GCs. In the galactic disk we find 12 pulsars with $ 100 < P_{orb} < 1000 $ days (although few of them have more massive companions than the pulsar binaries in the GCs), while there are only 3 pulsars in this range of orbital period. Although it is tempting to speculate that these systems are missing from the present day GCs because they have been ionized in the past, we find that the ionization probability becomes substantial in this orbital period range only for very small values of companion masses. Observational selection effects can be the cause of this discrepancy between the number of large orbital period pulsars in the disk in comparison to that in GCs. Many galactic disk binary pulsars are seen in the region of $e-P_{orb}$ plane predicted by Phinney due to the fluctuation dissipation of convective eddies and the resultant orbital eccentricities that are induced. These pulsars are missing from the GC sample. There is no reason not to expect these systems to form in the GCs (although due to the lower metallicity of the stellar companions in GCs, they are expected to be in slightly different region in the $e-P_{orb}$ plane). This can be explained by the substantial probability of them being knocked out of their original phase space due to stellar interaction in GCs. Indeed there are some wide binaries in the present day GCs with moderately high eccentricities (e.g. $0.01< e <0.1$ and $60 < P_{orb} < 256 \\; \\rm d$) which could have arisen out of fly-bys or exchanges from progenitor binaries with ``relic\" eccentricities $e \\sim 10^{-4}$ (merger in unlikely because of moderately low values of $m_c$s). The circularity of group III pulsars in GCs are thought to be due to either their special formation channels and young age or due to their locations at the outskirts of the GCs.  PSR J1903$+$0327 is the only one eccentric millisecond pulsar in the galactic disk discovered so far. Its millisecond spin-period suggests the spin-up scenario through accretion in the past. As such recycled pulsars are expected to be in circular orbits due to tidal-coupling induced circularization, this system has created lots of interest among pulsar researchers. As it is well known that in globular clusters, stellar interactions can impart eccentricity to an initially circular orbit, the possibility of past GC association of this system have been studied and almost rejected. Moreover, the possibility of this system being a member of a hierarchical triple system also seems to be less unlikely. So this system remains an open question for us.  We hope in the future with more advanced technologies like SKA, many more new interesting types of pulsars will be discovered and that the exciting and challenging situation will ultimately lead to a better understanding of the physics of pulsars."
    },
    "1004/1004.0384_arXiv.txt": {
        "abstract": "The new Wide Field Camera 3/IR observations on the Hubble Ultra-Deep Field started investigating the properties of galaxies during the reionization epoch. To interpret these observations, we present a novel approach inspired by the conditional luminosity function method. We calibrate our model to observations at $z=6$ and assume a non-evolving galaxy luminosity versus halo mass relation. We first compare model predictions against the luminosity function measured at $z=5$ and $z=4$. We then predict the luminosity function at $z\\geqslant 7$ under the sole assumption of evolution in the underlying dark-matter halo mass function. Our model is consistent with the observed $z \\gtrsim 7$ galaxy number counts in the HUDF survey and suggests a possible steepening of the faint-end slope of the luminosity function: $\\alpha(z \\gtrsim 8) \\lesssim -1.9$ compared to $\\alpha=-1.74$ at $z=6$. Although we currently see only the brightest galaxies, a hidden population of lower luminosity objects ($L/L_{*} \\gtrsim 10^{-4}$) might provide $\\gtrsim 75\\%$ of the total reionizing flux. Assuming escape fraction $f_{esc} \\sim 0.2$, clumping factor $C\\sim 5$, top heavy-IMF and low metallicity, galaxies below the detection limit produce complete reionization at $z\\gtrsim 8$. For solar metallicity and normal stellar IMF, reionization finishes at $z\\gtrsim 6$, but a smaller $C/f_{esc}$ is required for an optical depth consistent with the WMAP measurement. Our model highlights that the star formation rate in sub-$L_*$ galaxies has a quasi-linear relation to dark-matter halo mass, suggesting that radiative and mechanical feedback were less effective at $z \\geq 6$ than today. ",
        "introduction": "The new Wide Field Camera 3/IR (WFC3) Hubble Ultra Deep Field (HUDF09) observations opened a new window on high-redshift galaxy formation \\citep{oesch09_size,oesch09_zdrop,bouwens09_ydrop,bouwens09_slope,bunker09,mclure09,fink09}. Yet the sample of $z \\gtrsim 6.5$ galaxies is too small ($16$ z-dropouts in \\citealt{oesch09_zdrop}, and $5$ Y-dropouts in \\citealt{bouwens09_ydrop}) for a precise determination of the galaxy luminosity function (LF), especially after taking into account the systematic uncertainty introduced by cosmic variance \\citep{trenti08}.  Measuring the galaxy LF is important to assess their contribution to cosmic reionization, which started at $z\\gtrsim 10$, as inferred from the Thomson scattering optical depth $\\tau_e$ in the CMB background \\citep{komatsu09}. The nature of the reionizing sources is currently debated. Are normal galaxies the agents of reionization, or are other sources responsible, such as Population III stars or Mini-QSOs \\citep{madau04,sokasian04,shull08}? Within uncertainties, galaxies detected at $z\\sim 6$ barely keep the Universe reionized \\citep{stiavelli04b,bunker04}. The LF evolution established from $z\\sim 4$ to $z \\sim 6$ \\citep{bouwens07} seems to continue into the dark ages, with progressively fewer bright sources \\citep{bolton07,bouwens08,bouwens09_ydrop,oesch09}.  The exploration of the link between LF and underlying dark-matter halo mass function (MF) helps us understand the processes regulating star formation. This has been studied via the conditional luminosity function (CLF) method locally and at high redshift \\citep{vale04,cooray05a,cooray05b,cooray06,stark07,bouwens08,lee09}. Key results are: (1) significant redshift evolution of galaxy luminosity versus halo mass, $L(M_h)$, \\citep{cooray05b,lee09}; (2) only a fraction $\\epsilon_{DC}\\sim 20-30\\%$ of halos appears to host Lyman Break galaxies (LBG) \\citep{stark07,lee09}; (3) the predicted LF at $z \\gtrsim 6$ deviates significantly from Schechter form, missing the sharp drop in density of bright ($M_{AB} \\lesssim -20$) galaxies \\citep{bouwens08}. These findings suggest limitations of the current models extrapolated to the highest redshift. In fact, because of the young age of the Universe during the reionization epoch ($\\Delta z =1$ corresponds to $\\lesssim 170~\\mathrm{Myr}$ at $z \\gtrsim 6$), it becomes difficult to justify rapid evolution of $L(M_h)$, unless the IMF changes. A low $\\epsilon_{DC}$ also appears problematic: the halo MF evolves rapidly at $z\\gtrsim 6$: the number density of $M_h > 10^{11} M_{\\sun}$ halos (hosting $\\sim L_*$ galaxies) increases by a factor three from $z=7$ to $z=6$. Hence, $\\epsilon_{DC} \\lesssim 0.3$ implies that the majority of recently formed halos at $z \\geqslant 6$ did not experience significant star formation. Finally, the absence of a well-defined knee in the predicted $z\\gtrsim 6$ LF differs from the rarity of observed bright galaxies (see \\citealt{bouwens09_ydrop}).  To overcome these limitations, we present a novel implementation of the CLF model, tailored for application at $z\\gtrsim 5$. Instead of a duty cycle, we adopt another simple assumption: only halos formed within a given time interval host a detectable LBG (Section~\\ref{sec:ml}). Section~\\ref{sec:clf} contains the predictions for the $z \\gtrsim 7$ LF, compared to WFC3-HUDF09 observations. Section~\\ref{sec:reion} discusses the contribution of galaxies to reionization. ",
        "conclusions": "\\label{sec:conc}  In this Letter we construct a model for the evolution of the galaxy luminosity function at $z\\gtrsim 4.5$ based on a modification of the CLF method. We derive the relation between galaxy luminosity and dark-matter halo mass at $z=6$, assuming a one-to-one correspondence between observed galaxies and halos that formed in a period $\\Delta t = 200~\\rm Myr$. Using $L(M_h)$ fixed at $z=6$, we derive the expected LFs between $z=4$ and $z=9$, assuming only evolution of the underlying dark-matter MF. The $z=5$ LF is consistent with observations, but our model is less accurate at lower redshift because it underestimates the faint-end slope. At $z \\gtrsim 6$, we predict a moderate decrease of $\\phi_*$, a possible steepening of the faint-end slope, and continued evolution of $L_*$ toward lower values (Figure~\\ref{fig:lf} and Table~\\ref{tab:lf}). At all epochs, our predicted LF is well fitted by a Schechter function with a prominent ``knee''.  The predicted number counts for the HUDF09-WFC3 field are a good match to the dropouts observed at $z \\sim 7$ and $z \\sim 8$ (Table~\\ref{tab:hudfcounts}). Overall our ICLF model is consistent with the observed galaxy LF from $z \\sim 5$ to $z \\sim8$ with no evolution in $L(M_h)$. DM halos assembly can explain LF evolution at $z\\geq 5$ without invoking a change in the properties of LBG star formation. This is in agreement with the constant specific star formation rate inferred at $z\\gtrsim 5$ \\citep{gonzalez09}.  Our model provides evidence for a reduced impact of feedback in low-mass $z\\gtrsim 5$ halos. In fact, we derive a star formation efficiency weakly dependent on halo mass ($\\eta \\propto M_h^{0.3}$), compared to the strong quenching of star formation derived at $z=0$ ($\\eta \\propto M_h^3$; \\citealt{cooray05a}), providing a testable prediction for cosmological simulations. The strong suppression of star formation in $M_h\\lesssim 10^{11}~M_{\\sun}$ halos suggested by \\citet{bouche10} and \\citet{maiolino08} contrasts with $\\alpha(z=6) \\sim 1.7$ measured for halos with $M_h \\gtrsim 2 \\times 10^{10}~M_{\\sun}$ (Section~\\ref{sec:ml}). Such strong feedback would also imply that galaxies appear incapable of sustaining reionization. In fact, with a steep LF, sources below the HUDF-WFC3 detection limit may contribute $\\gtrsim 75\\%$ of the ionizing flux, sufficient for full reionization if $C/f_{esc} \\lesssim 25$. A metal-poor and top-heavy IMF, or smaller $C/f_{esc}$, are required to complete reionization at $z\\gtrsim 8$ for consistency with $(\\tau_e)_{WMAP}$. While our extrapolation is physically motivated to $T_{vir} \\geqslant 2\\times10^4~\\mathrm{K}$, it extends for $8$ magnitudes. Deeper observations are thus crucial to verify that the LF faint-end remains steep."
    },
    "1004/1004.5568_arXiv.txt": {
        "abstract": "{ {We present a comprehensive review of MHD wave behaviour in the neighbourhood of coronal null points: locations where the magnetic field, and hence the local Alfv\\'en speed, is zero.} {The behaviour of all three MHD wave modes, i.e. the Alfv\\'en wave and the fast and slow magnetoacoustic waves, has been investigated in the neighbourhood of 2D, 2.5D and {{(to a certain extent)}} 3D magnetic null points, for a variety of assumptions, configurations and geometries.} {In general, it is found that the fast magnetoacoustic wave behaviour is dictated by the Alfv\\'en-speed profile.  In a $\\beta=0$ plasma, the fast wave is focused towards the null point by a refraction effect and all the wave energy, and thus current density,  accumulates close to the null point. Thus, {\\emph{null points will be locations for preferential heating by fast waves}}.} {Independently, the Alfv\\'en wave is found to propagate along magnetic fieldlines and is confined to the fieldlines it is generated on. As the wave approaches the null point, it spreads out due to the diverging fieldlines. Eventually, the Alfv\\'en wave accumulates along the separatrices (in 2D) or along the spine or fan-plane (in 3D). Hence, {\\emph{Alfv\\'en wave energy will be preferentially dissipated at these locations}}.} {It is clear that the magnetic field plays a fundamental role in the propagation and properties of MHD waves in the neighbourhood of coronal null points. This topic is a fundamental plasma process and results so far have also lead to critical insights into reconnection, mode-coupling, quasi-periodic pulsations and phase-mixing.} } ",
        "introduction": "Magnetohydrodynamic (MHD) wave motions (e.g.  Roberts \\citeyear{Bernie};  Nakariakov \\& Verwichte \\citeyear{NV2005}; De Moortel \\citeyear{DeMoortel2005}) are  ubiquitous throughout the solar corona (Tomczyk et al. \\citeyear{Tomczyk}). Several different types of MHD wave motions have now been observed by various solar instruments:  slow magnetoacoustic  waves have been seen in {\\emph{SOHO}} data (e.g. {{Ofman {\\it{et al.}} \\citeyear{Ofman1997}; DeForest \\& Gurman \\citeyear{plumes}}};  Berghmans \\& Clette \\citeyear{Berghmans1999}; Kliem {\\it{et al.}} \\citeyear{Kliem}; Wang {\\it{et al.}} \\citeyear{Wang2002}) and {\\emph{TRACE}} data (De Moortel {\\it{et al.}} \\citeyear{DeMoortel2000}). Fast magnetoacoustic waves have been seen with {\\emph{TRACE}} (Aschwanden {\\it{et al.}} \\citeyear{Aschwandenetal1999}, \\citeyear{Aschwandenetal2002}; Nakariakov {\\it{et al.}} \\citeyear{Nakariakov1999}; Wang \\& Solanki \\citeyear{Wang2004}) and {\\emph{Hinode}} (Ofman \\& Wang \\citeyear{OW2008}). Non-thermal line narrowing / broadening due to Alfv\\'en waves has been reported by Harrison {\\it{et al.}} (\\citeyear{Harrison2002}) / {Erd{\\'e}lyi} {\\it{et al.}} (\\citeyear{E1998}), {{Banerjee {\\it{et al.}} (\\citeyear{Banerjee1998}) }}and  O'Shea {\\it{et al.}} (\\citeyear{Oshea}). Alfv\\'en waves have possibly been observed in the corona  (Okamoto {\\it{et al.}} \\citeyear{Okamoto}; Tomczyk {\\it{et al.}} \\citeyear{Tomczyk}) and chromosphere (De Pontieu {\\it{et al.}} \\citeyear{Bart2007}; Jess {\\it{et al.}} \\citeyear{Jess2009}), although these claims are subject to intense discussion (Erd{\\'e}lyi \\& Fedun \\citeyear{RF2007}; Van Doorsselaere {\\it{et al.}} \\citeyear{Tom2008}).    It is clear that the  coronal magnetic field  plays a fundamental role in the  propagation and properties of MHD waves, and to begin to understand this  inhomogeneous, magnetised  environment, it is useful to look at the topology (structure) of the magnetic field itself.  Potential-field extrapolations of the coronal magnetic field can be made from photospheric magnetograms, and such extrapolations show the existence of important features of the topology: {\\it{null points}} - locations where the magnetic field, and hence the Alfv\\'en speed, is zero, and {\\it{separatrices}} - topological features that separate regions of different magnetic flux connectivity. A comprehensive review can be found in Longcope (\\citeyear{L2005}).   This paper will provide a comprehensive literature review of the nature of MHD wave propagation in the neighbourhood of coronal null points. This topic exists at the overlap of two important areas of solar physics: {{MHD wave and magnetic null-point theories}}. A brief introduction to both of these areas is provided in $\\S\\ref{section:MHDwaves}$ and $\\S\\ref{section:toplogy}$, including a mathematical description of magnetic null points. $\\S\\ref{section:polar}$ reviews the early work that considers a 2D null point in a cylindrically symmetric geometry and describes the system in terms of normal modes. $\\S\\ref{section:cartesian}$ reviews work performed in a 2D cartesian geometry that focuses on externally driven perturbations, and $\\S\\ref{section:nonlinear}$ describes the extension of these investigations into the nonlinear regime. $\\S\\ref{guide_field}$ details the effects of threading the 2D X-point with an orthogonal  weak-guiding field. $\\S\\ref{section:threedimensionalnullpoints}$ details the behaviour of MHD wave propagation in the neighbourhood of 3D null points, and the conclusions and summary are given in $\\S\\ref{section:conclusions}$.     \\subsection{MHD Equations}\\label{section:basic_MHD_equations}    The viscous, resistive, compressible  MHD equations {{utilised in this paper are}}: \\begin{eqnarray} \\rho \\left[ {\\partial {\\bf{v}}\\over \\partial t} + \\left( {\\bf{v}}\\cdot\\nabla \\right) {\\bf{v}} \\right] &=& - \\nabla p + {\\frac{1}{\\mu}}\\left(   { \\nabla \\times {\\bf{B}}} \\right)\\times {\\bf{B}} +  \\nu  \\nabla \\cdot {\\underline{\\underline{\\pi}}}  \\;   \\; ,\\nonumber \\\\ {\\partial {\\bf{B}}\\over \\partial t}  &=& \\nabla \\times \\left ({\\bf{v}}\\times {\\bf{B}}\\right ) + \\eta \\nabla ^2  {\\bf{B}}\\;\\; ,\\nonumber \\\\ {\\partial \\rho\\over \\partial t} + \\nabla \\cdot \\left (\\rho {\\bf{v}}\\right ) &=& 0\\; \\;, \\nonumber \\\\ \\rho \\left[{\\partial {\\epsilon}\\over \\partial t}  + \\left( {\\bf{v}}\\cdot\\nabla \\right) {\\epsilon}\\right] &=& - p \\nabla \\cdot {\\bf{v}} + {{\\frac{1}{\\sigma}}} \\left| {\\bf{j}} \\right| ^2 +  \\nu \\varepsilon_{ij} \\pi_{ij}   \\;\\; \\label{MHDequations}  , \\end{eqnarray} where $\\rho$ is the mass density, ${\\bf{v}}$ is the plasma velocity, ${\\bf{B}}$ the magnetic induction (usually called the magnetic field), $p$ is the plasma pressure,  $ \\mu = 4 \\pi \\times 10^{-7} \\/\\mathrm{Hm^{-1}}$  is the magnetic permeability,{{ $\\nu$ is the coefficient of classical viscosity, $\\pi_{ij} = \\varepsilon_{ij} - \\delta_{ij}\\nabla \\cdot {\\bf{v}}$ is the stress tensor, $\\varepsilon_{ij} = \\left( {\\partial v_i/ \\partial x_j} + {\\partial v_j/ \\partial x_i} \\right)/2$ is the rate-of-strain tensor,}} $\\sigma$ is the electrical conductivity,  $\\eta=1/ {\\mu \\sigma} $ is the magnetic diffusivity, $\\epsilon= {p / \\rho \\left( \\gamma -1 \\right)}$ is the specific internal energy density, where $\\gamma={5 / 3}$ is the ratio of specific heats and ${\\bf{j}} = {{\\nabla \\times {\\bf{B}}} / \\mu}$ is the electric current density. $\\nu$ and $\\eta$ are assumed to be constants.     {{Note that the classical viscous term used in equations (\\ref{MHDequations}) is in fact not the most appropriate for the solar corona since,  in the presence of strong magnetic fields, the viscosity takes the form of a non-isotropic tensor. However, only the papers of Craig \\& Litvinenko (\\citeyear{CL2007}) and Craig (\\citeyear{Craig2008}) mentioned in this review  will invoke the non-isotropic viscous tensor and so, for brevity, we do not provide a full description here. The mathematical details of the  non-isotropic viscous tensor can be found in Braginskii (\\citeyear{Braginskii1965}) and, for example, Van der Linden {\\it{et al.}} (\\citeyear{Linden}), Ofman {\\it{et al.}} (\\citeyear{Ofman1994}) and  Erd{\\'e}lyi \\& Goossens (\\citeyear{EG1994}; \\citeyear{EG1995}).}}          \\subsection{MHD waves}\\label{section:MHDwaves}   {{A wave is a disturbance that propagates through space and time, usually with the transference of energy. Such a disturbance, either continuous or transient, propagates by virtue of the elastic nature of the medium.}} In MHD, the magnetic tension provides an elastic restoring force, such that we would expect waves to propagate along uniform magnetic fieldlines with a characteristic speed: \\begin{eqnarray*} v_A= \\frac{ | {\\bf{B}} |}{ \\sqrt{ \\mu \\rho}} \\;\\;, \\end{eqnarray*} where  $v_A$ is called the {\\emph{Alfv\\'en speed}}. Transverse waves travelling at this speed along magnetic fieldlines are called {\\emph{Alfv\\'en waves}}.  If we consider a compressible medium, then we can define the {\\emph{sound speed}} as: \\begin{eqnarray*} c_s  = \\sqrt{\\frac{\\gamma p}{\\rho}} \\;\\;. \\end{eqnarray*} When assuming  a compressible medium, the Alfv\\'en wave still remains, but the sound and Alfv\\'en speed can now couple together to give {\\emph{magnetoacoustic}} waves. Two combinations arise: the higher frequency mode is known as the {\\emph{fast magnetoacoustic wave}} and the lower frequency wave is known as the {\\emph{slow magnetoacoustic wave}}. These three wave types, the Alfv\\'en wave and the fast and slow magnetoacoustic waves, make up the three MHD waves considered in this review paper.   The fundamental properties and nature of linear MHD waves in uniform magnetic fields have been reported in detail by several authors, for example in an unbounded homogeneous medium (Cowling \\citeyear{Cowling1976}), and in a bounded inhomogeneous slab / cylindrical density profile embedded in a uniform magnetic field (Roberts \\citeyear{Roberts1981a}) / (Edwin \\& Roberts \\citeyear{ER1983}; Cally \\citeyear{Cally1986}; Roberts \\& Nakariakov \\citeyear{RN2003}).   Finally, in MHD it is important to consider the ratio of magnetic pressure to thermal pressure. This ratio is called the plasma$-\\beta$ and is given by: \\begin{eqnarray} \\beta=  \\frac{2 \\mu p}{| {\\bf{B}} |^2} = {\\frac{2}{\\gamma}}{\\frac{c_s^2}{v_A^2}}\\;\\;.\\label{plasmabetaequation} \\end{eqnarray} The properties of the fast and slow  magnetoacoustic waves have a strong dependence on the magnitude of the plasma$-\\beta$, namely because it is directly proportional to the square of the ratio of the sound speed to the Alfv\\'en speed. Thus, in a regime where $\\beta \\ll 1$, magnetic pressure and magnetic tension dominate the propagation and vice versa. Table 1 lists the main properties of the three wave types depending upon their environment. Note that the Alfv\\'en wave behaviour is independent of the plasma-$\\beta$, as it is a purely magnetic wave (in the linear regime).  The plasma$-\\beta$ parameter varies greatly with height across the layers of the solar atmosphere (see Gary \\citeyear{Gary2001} for well-constrained values). However, magnetic pressure generally dominates thermal pressure in the solar corona, and thus it is usual to assume plasma$-\\beta \\ll 1$ when modelling a coronal environment. Hence, when we talk about coronal null points, we are talking about null points in a low$-$ or zero$-\\beta$ environment, although there are some caveats to this ($\\S\\ref{section:mode conversion}$).    \\begin{table}[h] \\begin{tabular}{|c|c|c|} \\hline & plasma-$\\beta\\gg1$ (high-$\\beta$) & plasma$-\\beta\\ll 1$ (low-$\\beta$)\\\\ \\hline\\hline Alfv\\'en wave & \\multicolumn{2}{c}{Transverse wave propagating at speed $v_A$} \\begin{tabular}{c} \\end{tabular} \\vline\\\\ \\hline Fast MA wave & \\begin{tabular}{c} \\vspace{0.1cm}Behaves like sound wave\\\\ (speed $c_s$) \\end{tabular} & \\begin{tabular}{c} Propagates roughly isotropically \\\\ Propagates across magnetic fieldlines \\\\ (speed $v_A$) \\\\ \\end{tabular}\\\\ \\hline Slow MA wave & \\begin{tabular}{c} Guided along ${\\bf{B}}$  \\\\ (speed $v_A$) \\end{tabular} & \\begin{tabular}{c} Guided along ${\\bf{B}}$\\\\ Longitudinal wave propagating \\\\ at speed $c_s$\\\\ \\end{tabular}\\\\ \\hline \\end{tabular} \\caption{Properties of MHD waves depending on the plasma$-\\beta$.} \\end{table}       \\subsection{Magnetic Topology}\\label{section:toplogy}    The magnetic field plays an essential role in understanding the myriad of phenomena in the solar corona. A realistic magnetic field can have many different components, and we can use {\\emph{topology}} nomenclature to reduce a complicated set of fieldlines to something more understandable. In 2D, a general magnetic configuration contains {\\emph{separatrix curves}} (separatrices) which split the magnetic plane into topologically distinct regions, in the sense that within a specific region all the fieldlines start at a particular source and end at a particular sink. There is a second important topological aspect: {\\emph{magnetic null points}} (or neutral points) are single-point locations where the magnetic field vanishes (${\\bf{B}}={\\bf{0}}$). There are two types of magnetic null point: {\\emph{X-type null points}}, commonly called X-points, which  occur at the intersection of separatrix curves, and {\\emph{O-type null points}}, or O-points, located at the center of magnetic islands. Magnetic topologies that contain null points are common in the presence of multiple magnetic sources.   A magnetic fieldline that joins two null points (itself a special type of separatrix) is called a {\\emph{separator}}. Thus, instead of showing all the magnetic field lines in a region, we can just show the important aspects of the topology; such a picture of the magnetic structure is called the {\\emph{magnetic skeleton}} of the field. In 3D, we have similar properties, now with {\\emph{separatrix surfaces}} separating the volume into topologically distinct regions, and these surfaces intercept at a separator.      \\subsubsection{Mathematical description of null points}\\label{2Dnulls}      \\begin{figure}[t] \\includegraphics[width=6.0in]{XXXXFigure1.eps} \\caption{$(a)$ X-type null point for $\\alpha^2=1$. This potential neutral point has separatrices (red lines) intersecting at an angle of $\\pi/2$. $(b)$ O-type null point, with $\\alpha =-1$. A blue star denotes the null point. $(c)$ Single potential magnetic null point configuration created by interaction of two dipoles. Here, $A_z= y/ \\left[ \\left( x+\\lambda \\right)^2 +y^2\\right] + y/ \\left[ \\left( x-\\lambda\\right)^2 + y^2 \\right]$, for $\\lambda=0.5$. Red lines denote the separatrices.} \\label{figure1} \\end{figure}    Let us first consider null points in 2D {{(e.g. Dungey \\citeyear{Dungey1953}; \\citeyear{Dungey1958}). Following  $\\S1.3.1$ of Priest \\& Forbes ({\\citeyear{magneticreconnection2000}})}}, we assume a magnetic field of the form: \\begin{eqnarray*} {\\bf{B}}= \\left[ B_x(x,y), B_y(x,y),0\\right]\\;\\;. \\end{eqnarray*} A null point occurs at the point $(x_0,y_0)$ if: \\begin{eqnarray*} B_x(x_0,y_0) =0\\;\\;{\\rm{and}}\\;\\;B_y(x_0,y_0) =0\\;. \\end{eqnarray*} Expanding $B_x$ and $B_y$ in a Taylor series about $(x_0,y_0)$ gives the linear approximation: \\begin{eqnarray} B_x &=& \\left.{\\frac{\\partial B_x}{\\partial x}}\\right|_{(x_0,y_0)} \\left(x-x_0\\right) + \\left.{\\frac{\\partial B_x}{\\partial y}}\\right|_{(x_0,y_0)} \\left(y-y_0\\right)\\nonumber \\\\ &=& a(x-x_0)+b(y-y_0)\\;\\;,\\label{a1}\\\\ B_y &=& \\left.{\\frac{\\partial B_y}{\\partial x}}\\right|_{(x_0,y_0)} \\left(x-x_0\\right) + \\left.{\\frac{\\partial B_y}{\\partial y}}\\right|_{(x_0,y_0)} \\left(y-y_0\\right) \\nonumber \\\\ &=& c(x-x_0)-a(y-y_0)\\;\\;,\\label{a2} \\end{eqnarray} where the coefficients $a,b,c$ are arbitrary.  Let us now introduce the vector potential (also called the flux function), ${\\bf{A}}$, such that ${\\bf{B}} = \\nabla \\times {\\bf{A}}$, and in 2D we have ${\\bf{A}} = (0,0,A_z)$. Thus, we have: \\begin{eqnarray} {\\bf{B}}= \\left(\\frac{\\partial A_z}{\\partial y}, -\\frac{\\partial A_z}{\\partial x},0\\right)\\;\\;.\\label{fluxfunction} \\end{eqnarray}  Integrating equations (\\ref{a1}) and (\\ref{a2}) gives the corresponding vector potential as: \\begin{eqnarray*} A_z= a(x-x_0)(y-y_0)+ \\frac{b}{2} (y-y_0)^2 - \\frac{c}{2}(x-x_0)^2\\;\\;, \\end{eqnarray*} where we have chosen the arbitrary constant of integration such that $A_z$ vanishes at  $(x_0,y_0)$. Further simplification is possible by rotating the $xy$-axes through an angle $\\theta$ to give new $x'$, $y'$-axes, and choosing the angle $\\theta$ such that $\\tan{2\\theta}= -2a/(b+c)$. This simplification gives the corresponding vector potential as: \\begin{eqnarray} A_z= {\\frac{B}{2L}}\\left[ \\left(y'-y'_0\\right)^2 - \\alpha^2 \\left(x'-x'_0\\right)^2\\right]\\;\\;,\\label{ff_simpleXpoint} \\end{eqnarray} where \\begin{eqnarray*} {\\frac{B}{L}}= \\frac{2a^2+b^2-c^2}{\\sqrt{4a^2+(b+c)^2}}\\;\\;,\\;\\; \\alpha ^2 = {\\frac{ 4a^2}{ 2a^2+b^2-c^2} -1} \\;\\;. \\end{eqnarray*}  Here, $B$ is the characteristic strength of the magnetic field and $L$ is the characteristic length-scale over which the field varies. The corresponding field components are: \\begin{eqnarray} B_x =\\frac{B}{L}(y'-y'_0)  \\;\\; {\\rm{and}} \\;\\; B_y =\\frac{B}{L}\\alpha^2 (x'-x'_0) \\label{BX_and_BY}\\;\\;. \\end{eqnarray}  Magnetic field lines are defined by $A_z$ equal to a constant. For $\\alpha^2>0$, the fieldlines are hyperbolic, giving an X-type null point. The separatrices are given by $y'-y'_0 = \\pm \\alpha (x'-x'_0)$ and are inclined at an angle $\\pm \\arctan{\\alpha}$ to the $x'$-axis.  The magnetic fieldlines for $\\alpha^2=1$, $x'_0=y'_0=0$ can be seen in  Figure \\ref{figure1}a.    \\begin{figure*} \\includegraphics[width=6.0in]{XXXXFigure2.eps} \\caption{Magnetic fields containing two null points. $(a)$ Magnetic configuration containing both X-type and O-type null points. Here $A_z= \\lambda^2 x- y^2 - (x-\\lambda)^3/3$ (${\\bf{B}}=\\left[-2y, \\left(x-\\lambda\\right)^2 - \\lambda^2\\right]$) for $\\lambda=0.5$. Red lines/blue star denotes the separatrices/O-type null point. $(b)$ Potential magnetic configuration containing two X-type null points connected by a separator. Here, $A_z= -x^2y+y^3/3+ \\lambda^2 y$, where $\\lambda=1$.  $(b)$ Potential magnetic configuration containing two X-type null points not connected by a separator. Here, $A_z= -xy^2+x^3/3-\\lambda^2 x$, where $\\lambda=1$. Red lines denote the separatrices.} \\label{figure2} \\end{figure*}    The value of $\\alpha$ (and thus the angle between the separatrices) is related to the current density. The current density is given by: \\begin{eqnarray*} {\\bf{j}}= {\\frac{1}{\\mu}}\\left( {\\nabla \\times {\\bf{B}}}\\right) = -{\\frac{1}{\\mu}} \\nabla^2 A_z  {\\hat{\\bf{z}}} =  -{\\frac{B}{\\mu L}} \\left( 1-\\alpha^2 \\right) {\\hat{\\bf{z}}}  \\;\\;. \\end{eqnarray*} Thus, a null point is potential if $\\alpha=\\pm1$ (i.e. an X-type configuration of rectangular hyperbola) and the angle between the separatrices is $\\pi/2$. Note that for an O-point, $\\alpha$ is imaginary and so the current density is always non-zero. Thus, O-type neutral points can never be potential.  An O-type null point magnetic configuration can be seen in Figure \\ref{figure1}b, for $\\alpha^2=-1$, $x'_0=y'_0=0$. However, note that the simple 2D magnetic field configuration of equation (\\ref{BX_and_BY}) is only valid close to the null point: as $x'$ and/or $y'$ get very large, ${\\bf{B}}$ becomes unphysically large. Figure \\ref{figure1}c denotes a more realistic single magnetic null point configuration created by the interaction of two dipoles. This configuration comprises of four separatrices and an X-point, and as $x'$ and/or $y'$ get large, the field strength becomes smaller (i.e. a more physical field).   Magnetic configurations can also contain multiple null points, and it can be argued that null points appear in pairs; a double null point may arise as a local  bifurcation of a single 2D null point (see e.g. Galsgaard {\\it{et al.}} \\citeyear{KRRN1996}; Brown \\& Priest \\citeyear{BP1998}).    Figure \\ref{figure2}a.  shows a magnetic configuration containing both X-type and O-type null points. Figures \\ref{figure2}b and \\ref{figure2}c present magnetic configuration containing two X-type null points connected by a separator and not connected by a separator, respectively.         \\begin{figure} \\hspace{-1.0cm} \\includegraphics[width=6.0in]{XXXXFigure3.eps} \\caption{$(a)$ Proper radial null point, described by ${\\bf{B}}=(x,y,-2z)$, {\\it{i.e.}} $\\epsilon=1$. $(b)$ Improper radial null point, described by ${\\bf{B}}=\\left(x,\\epsilon y,- \\left[\\epsilon+1 \\right]z\\right)$, for $\\epsilon={1}/{2}$.  Note for $\\epsilon={1}/{2}$, the field lines rapidly curve such that they run parallel to the $x-$axis along $y=0$. In both figures, the  $z-$axis indicates the {\\emph{spine}} and the $xy-$plane at $z=0$ denotes the {\\emph{fan-plane}}. The red fieldlines have been tracked from the $z=1$ plane, the blue from $z=-1$.} \\label{figure3} \\end{figure}    \\subsubsection{Three-dimensional magnetic null points}\\label{section:threeDnulls}     Magnetic null points also exist in three dimensions, but  occur in a different form to those described in \\S\\ref{2Dnulls}. In 3D, potential null points are of the form: \\begin{eqnarray} {\\bf{B}} = \\frac{B}{L} \\left(x,\\epsilon y,-\\left[ \\epsilon +1 \\right] z \\right)   \\label{Bfield}  \\;\\;, \\end{eqnarray} where the parameter $\\epsilon$ is related to the predominate direction of alignment of the fieldlines in the fan plane.  Parnell {\\it{et al.}} (\\citeyear{Parnell1996}) investigated and classified the different types of linear magnetic null points that can exist (our $\\epsilon$ parameter is called $p$ in their work). Topologically, this 3D null consists of two key parts: the $z-$axis represents a special, isolated fieldline called the {\\emph{spine}} which approaches the null from above and below (as found by Priest \\& Titov \\citeyear{PriestTitov1996}) and the $xy-$plane through $z=0$ is known as the {\\emph{fan-plane}} and consists of a surface of fieldlines spreading out radially from the null. Figure \\ref{figure3} shows two examples of 3D null points: $\\epsilon=1$ (Figure \\ref{figure3}a) and  $\\epsilon=1/2$ (Figure \\ref{figure3}b). Titov \\& Hornig (\\citeyear{TH2000}) have investigated the steady state structures of 3D magnetic null points.  Equation (\\ref{Bfield}) is the general expression for the linear field about a potential magnetic null point (see Parnell {\\it{et al.}} \\citeyear{Parnell1996}: ${\\S}{IV}$).  For $\\epsilon \\ge 0$, 3D nulls are described as {\\it{positive}} nulls, {\\it{i.e.}} the spine points into the null and the field lines in the fan are directed away. In addition, all potential nulls are designated  {\\it{radial}}, {\\it{i.e.}} there is no spiral motions in the  fan-plane.  In this review paper, we only consider positive, potential null points, and thus there are three cases to consider: \\begin{itemize} \\item{$\\epsilon  = 1$: describes a {\\it{proper null}} (Figure \\ref{figure3}a). This magnetic null has cylindrical symmetry about the spine axis (so is actually only a 2.5D null point).} \\item{$\\epsilon > 0,\\: \\epsilon \\neq 1$: describes an {\\it{improper null}}  (Figure \\ref{figure3}b). Field lines rapidly curve such that they run parallel to the $x-$axis if $0<\\epsilon <1$ and parallel to the $y-$axis if $\\epsilon >1$.} \\item{$\\epsilon=0$: equation (\\ref{Bfield}) reduces to a simple 2D X-point potential field in the $xz-$plane and forms a null line along the $y-$axis through $x=z=0$.} \\end{itemize}      \\subsection{Statistics of coronal null points}\\label{statistics}    We have provided a mathematical description of null points, but how common are null points in the corona? Null points are an inevitable consequence of the distributed isolated magnetic flux sources at the photospheric surface. {{Using photospheric magnetograms to provide the field distribution on the lower boundary,  both potential and non-potential (nonlinear force-free) field extrapolations suggest}} that there are always likely to be null points in the corona. The number of such singular points will depend upon the magnetic complexity of the photospheric flux distribution. Detailed investigations of the coronal magnetic field, using such potential field calculations, can be found in {Beveridge} {\\it{et al.}} (\\citeyear{Beveridge2002}) and {Brown \\& Priest (\\citeyear{BrownPriest2001})}. The properties of coronal null points have also been considered through theoretical considerations (e.g. Parnell {\\it{et al.}} \\citeyear{Parnell1996}; Brown \\& Priest \\citeyear{BrownPriest2001};  {Beveridge} {\\it{et al.}} \\citeyear{Beveridge2002}; Parnell \\& Galsgaard \\citeyear{PG2004}; Parnell {\\it{et al.}} \\citeyear{Parnell2008}).        The statistics of coronal null points has been investigated  using two methodologies: direct measurement from potential field extrapolations (e.g. Close {\\it{et al.}} \\citeyear{Close2004}; R\\'egnier {\\it{et al.}} \\citeyear{Stephane2008}) and, secondly, as an estimate from the Fourier spectrum of magnetograms (Longcope \\& Parnell \\citeyear{LP2009}). Close {\\it{et al.}} (\\citeyear{Close2004}) calculated a potential field extrapolation from a high resolution MDI magnetogram and found $1.7\\times10^{-3}$  magnetic null points per square megameter. R\\'egnier {\\it{et al.}} (\\citeyear{Stephane2008}) performed a similar investigation using a magnetogram from the Narrowband Filter Imager onboard {\\emph{Hinode}} and found  $6.7\\times10^{-3}$ ${\\rm{Mm^{-2}}}$. Longcope \\& Parnell (\\citeyear{LP2009}) investigated 562 MDI magnetograms using the Fourier spectrum of magnetograms and found $3.1\\times10^{-3}\\pm 3.0\\times10^{-4}$  coronal  null points per square megameter (at altitudes greater than 1.5 Mm). Alternatively, we can estimate the total number of coronal null points by multiplying these results by the surface area of the Sun (i.e. to provide a crude estimate, where we assume the Sun is free of active regions and coronal holes). This corresponds to approximately $10,000$  (Close {\\it{et al.}} \\citeyear{Close2004}),  $19,000$  (Longcope \\& Parnell \\citeyear{LP2009}) or $40,000$  (R\\'egnier {\\it{et al.}} \\citeyear{Stephane2008}) coronal null points.   {{More recently,  Cook {\\it{et al.}} (\\citeyear{Cook2009}) investigated the solar cycle variation of coronal null points using a potential field source surface model in spherical geometry, and find that there is no significant variation in the number of nulls found from the rising to the declining phase (indicating that null points are present throughout both phases of the solar cycle).}}     {{Several investigations also consider specific examples of null points in the corona. For example, Aulanier {\\it{et al.}} (\\citeyear{Aulanier}) investigated a class M3 flare that occured on 14 July 1998 above a $\\delta-$spot. Using potential field extrapolations, the authors recreated the pre-flare magnetic topology  from  Kitt Peak line-of-sight magnetograms and revealled a single coronal null point located above the $\\delta-$spot. Secondly, Ugarte-Urra {\\it{et al.}} (\\citeyear{Ugarte}) investigated the magnetic topology of  26 CME events by performing potential field extrapolations from corresponding MDI magnetograms, and find that magnetic null points are present in a large number of the pre-CME topologies.}}     Finally, we note that a null point plays a key role in the {\\emph{magnetic breakout model}} (e.g. Antiochos \\citeyear{Antiochos1998}; Antiochos {\\it{et al.}} \\citeyear{Antiochos1999}; MacNeice {\\it{et al.}} \\citeyear{Macneice2004}; Lynch {\\it{et al.}} \\citeyear{lynch2004}; Choe {\\it{et al.}} \\citeyear{choe2005}). The equilibrium set-up of the  magnetic breakout model  model consists of a quadrupolar photospheric flux distribution coupled with an overlying field, and such a set-up contains a coronal null point. Such a null point is an ideal candidate for the study of MHD wave behaviour about coronal null points (i.e. the investigations detailed in this review paper). However, it is important to stress  that the null point in the breakout model is not the only candidate - the ideas and investigations detailed below apply equally well around coronal null points found elsewhere (i.e. quiet Sun and inside active regions). Thus,  the coronal null points we are describing in this paper are not solely those involved in the magnetic breakout model.      \\subsection{Why is this area of study interesting or important?}\\label{section:motivation}   {{The motivation for investigating the behaviour of MHD wave propagation in the neighbourhood of magnetic null points can be summarised as follows: \\begin{itemize} \\item{MHD wave propagation in inhomogeneous media is a fundamental plasma process, and the study of MHD wave behaviour in the neighbourhood of magnetic null points directly contributes to this area.} \\item{We now know that MHD wave perturbations are omnipresent in the corona. We also know that null points are an inevitable consequence of the distributed isolated magnetic flux sources at the photospheric surface (where the number of such singular points will depend upon the magnetic complexity of the photospheric flux distribution). Thus, these two areas of scientific study; MHD waves and magnetic topology, {\\emph{will}} encounter each other at some point,  i.e. MHD waves will propagate into the neighbourhood of coronal null points (e.g. blast waves from a flare will at some point encounter a null point). Thus, the study of MHD waves around null points is itself a fundamental coronal process.} \\item{The study of MHD wave behaviour in the neighbourhood of magnetic null points is also interesting in its own right and, as we shall see, often  provides critical insights into other areas of plasma physics, including: mode-conversion ($\\S\\ref{section:mode conversion}$), reconnection ($\\S\\ref{section:nonlinear}$), quasi-periodic pulsations ($\\S\\ref{QPPs_section}$) and  phase-mixing ($\\S\\ref{PM_section}$).} \\end{itemize}   }} ",
        "conclusions": "\\label{section:conclusions}  The behaviour of all three MHD wave types; Alfv\\'en, fast and slow wave, has been investigated in the neighbourhood of  2D, 2.5D and{{ (to a certain extent) }}3D magnetic null points, in a variety of geometries and under a variety of assumptions. The main conclusions may be summarised as follows: \\begin{itemize} \\item{The linear, fast magnetoacoustic wave behaviour is dictated by the equilibrium fast wave speed profile (i.e. ${\\sqrt{v_A^2+c_S^2}}$), which in low-$\\beta$ plasmas can be thought of as the equilibrium Alfv\\'en-speed profile. The fast wave is guided towards the null point by a refraction effect and wraps around it. The fast wave slows as it approaches the null, leading to a decrease in length scales and thus an increase in current density close to the null point. In a $\\beta=0$ plasma, the fast wave cannot cross the null point and the build-up of current is exponential, indicating that dissipation will occur on a timescale related to $\\log {\\eta}$. Thus, linear fast wave dissipation is very efficient, and {\\emph{null points will be locations for preferential heating}}. For $\\beta \\neq 0$, the fast wave can cross the null point, due to the finite sound speed there, and wave energy can now escape the null point. In this case, there exists two competing phenomena and the dominate effect is determined by the value of the plasma-$\\beta$.} \\item{The linear Alfv\\'en wave  propagates along the equilibrium fieldlines and a fluid element is confined to the fieldline it starts on. Since the propagation follows the fieldlines, the Alfv\\'en wave spreads out as it approaches the diverging null point. In 2D, all the Alfv\\'en wave energy accumulates along the separatrices and the current build-up is exponential in time. In 3D, for an Alfv\\'en wave generated along the fan-plane, the wave accumulates along the spine and for an Alfv\\'en wave generated across the spine,  the value of $\\epsilon$ determines where the wave accumulation will occur:  fan-plane ($\\epsilon=1$), along the $x-$axis ($0<\\epsilon <1$) or along the $y-$axis ($\\epsilon>1$). Hence, all the {\\emph{Alfv\\'en wave energy will be dissipated along the separatrices/ separatrix surfaces}} and these will be the locations for preferential heating.} \\item{The behaviour of the slow wave in the neighbourhood of null points has received the least attention in the literature. The linear slow wave is found to be wave-guided and accumulates along the separatrices. A low-$\\beta$ fast wave can generate/convert into both a high-$\\beta$ fast and high-$\\beta$ slow wave as it crosses the $v_A=c_S$ mode-conversion layer. Such a layer is a natural consequence for a null point emersed in a $\\beta \\neq 0$ plasma. In fact, the value of $\\beta$ grows as $r^{-2}$ close to  a null point.} \\item{The addition of a weak guiding field leads to linear coupling between the fast and Alfv\\'en waves in a low-$\\beta$ plasma, and thus the propagation of either mode can generate the other. Such a configuration is, of course, no longer a null point, but rather an X-line. However, the nature of mode-coupling for 3D null points is, at this time, uncertain. Fast waves have been shown to be generated by the propagation of the Alfv\\'en wave, but it is unclear if this is due to the fieldline geometry, nonlinear coupling or both waves being simultaneous generated by a common driver.} \\item{Results in 2D show that in the  nonlinear regime, the fast magnetoacoustic wave can deform the equilibrium X-point configuration, leading to a cycle of horizontal and vertical current sheets and associated changes in connectivity. Thus, the system exhibits {\\emph{oscillatory reconnection}}.} \\item{It is clear that the equilibrium  magnetic field  plays a fundamental role in the  propagation and properties of MHD waves. In general, an arbitrary  disturbance/perturbation will generate all three wave modes and current accumulation could occur at all the null points, and/or along the spine, fan and separators. Thus, the results described in this review all highlight the importance of understanding the magnetic topology in determining  the locations of wave heating.} \\end{itemize}             However, several big questions still remain in this area: \\begin{itemize} \\item{The nature of the coupling of the three modes in 3D needs to be addressed, and the importance of coupling  due to the magnetic geometry verses nonlinear coupling should be investigated.} \\item{The theory of nonlinear fast waves driving oscillatory reconnection should be extended to study more general disturbances, and to investigate how robust the initial findings of McLaughlin {\\it{et al.}} (\\citeyear{MDHB2009}) are.} \\item{The key results for the linear fast and Alfv\\'en wave make clear predictions as to where preferential heating can occur. It would be interesting to see the theoretical models developed with forward modelling (see e.g. Kilmchuk \\& Cargill \\citeyear{KC2001}; {De Moortel} \\& {Bradshaw} \\citeyear{ineke2008}) to provide tell-tale observational signatures, and for these synthetic results to be compared with observational data.} \\end{itemize}        {{In conclusion, we have seen that the study of MHD wave behaviour in the neighbourhood of magnetic null points is  a fundamental plasma process, and can provide critical insights into other areas of plasma behaviour including: mode-conversion ($\\S\\ref{section:mode conversion}$), oscillatory reconnection ($\\S\\ref{section:nonlinear}$), quasi-periodic pulsations ($\\S\\ref{QPPs_section}$) and  phase-mixing ($\\S\\ref{PM_section}$).}}     {{We now know that the corona is full of MHD wave perturbations   (Tomczyk et al. \\citeyear{Tomczyk}). We also know that null points are an inevitable consequence of the distributed isolated magnetic flux sources at the photospheric surface, and potential and non-potential field extrapolations suggest that there are always likely to be null points in the corona (see $\\S\\ref{statistics}$). Thus, these two areas of scientific study (MHD wave behaviour and magnetic topology) will inevitably encounter each other at some point, i.e. MHD waves {\\emph{will}} propagate in the neighbourhood of  coronal null points. Thus, MHD wave propagation about magnetic null points  is itself - theoretically -  a fundamental coronal process. }}       {{However, there is as yet no clear observational evidence for MHD wave behaviour in the neighbourhood of coronal null points. In the lead author's opinion, the successful detection of MHD oscillations around coronal null points will require input from two areas: high-spatial/temporal resolution imaging data as well as potential/non-potential extrapolations from co-temporal magnetograms. Two of the instruments onboard the recently launched Solar Dynamics Observatory (SDO) may satisfy these requirements: the {\\emph{Atmospheric Imaging Assembly}} (which will provide high-quality imaging data) and the {\\emph{Helioseismic and Magnetic Imager}} (which will provide vector magnetograms). Thus, the first detection of MHD waves in the neighbourhood of coronal null points may be reported in the near future.}}"
    },
    "1004/1004.2038_arXiv.txt": {
        "abstract": "We present new $U$-band photometry of the magnetic Helium-strong star \\sorie, obtained over 2004--2009 using the SMARTS 0.9-m telescope at Cerro Tololo Inter-American Observatory. When combined with historical measurements, these data constrain the evolution of the star's 1\\fd19 rotation period over the past three decades. We are able to rule out a constant period at the $\\pnull = 0.05\\%$ level, and instead find that the data are well described ($\\pnull = 99.3\\%$) by a period increasing linearly at a rate of 77\\,ms per year. This corresponds to a characteristic spin-down time of 1.34\\,Myr, in good agreement with theoretical predictions based on magnetohydrodynamical simulations of angular momentum loss from magnetic massive stars. We therefore conclude that the observations are consistent with \\sorie\\ undergoing rotational braking due to its magnetized line-driven wind. ",
        "introduction": "\\label{sec:intro}  The helium-strong star \\sorie\\ (HD~37479; B2Vpe; $V=6.66$) has long been known to harbor a circumstellar magnetosphere in which plasma is trapped and forced into co-rotation by the star's strong ($\\sim 10\\,{\\rm kG}$) dipolar magnetic field \\citep[see, e.g.,][]{LanBor1978,GroHun1982}. This magnetosphere is largely responsible for the star's distinctive eclipse-like dimmings, which occur when plasma clouds transit across the stellar disk twice every 1\\fd19 rotation cycle \\citep{TowOwo2005}. Some fraction of the star's photometric variations likely also arise from its photospheric abundance inhomogeneities, as in other chemically peculiar stars \\citep[e.g,][]{Mik2009}; but for \\sorie\\ the magnetospheric contribution to the variations is dominant \\citep{Tow2008}.  This paper presents new $U$-band photometry of the star's primary light minimum\\footnote{The term `primary' stems from early mis-identifications of the star as an eclipsing binary system \\citep[e.g.,][]{Hes1976}; here, it simply indicates the deeper of the star's two light minima.}, obtained over four seasons spanning 2004--2009 using the SMARTS 0.9-m telescope at Cerro Tololo Inter-American Observatory (CTIO). When combined with historical measurements by \\citet{Hes1977}, these new data allow a precise measurement of the star's rotation period, and its evolution, over the past three decades.  A description of the observations, both archival and new, is provided in the following section. In \\S\\ref{sec:analysis}, we discuss a procedure for accurately measuring the times \\tmin\\ of primary light minimum, and then use these measurements to construct a standard observed-minus-corrected (\\OC) diagram for the star, allowing us to assess the evolution of the star's rotation period. We discuss and summarize our findings in \\S\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  The Helium-strong star HD~37776 was found by \\citet{Mik2008} to exhibit a progressive lengthening in its $1\\fd5387$ rotation period, with a characteristic spin-down time $\\tspin \\equiv P/\\dot{P} = 0.25\\,{\\rm Myr}$. For \\sorie\\ the absence of photometric data in the 1980s and 1990s means that we cannot empirically differentiate between steady spin-down and a sequence of abrupt braking episodes. However, the steady scenario is lent strong support by magnetohydrodynamical (MHD) simulations of angular momentum loss in magnetically channelled line-driven winds \\citep{udD2009}, which indicate that the lengthening of rotation periods should be a smooth process. Therefore, the use of a quadratic ephemeris (cf.~eqn.~\\ref{eqn:ephem}) appears justified, and we derive a characteristic spin-down time $\\tspin = 1.34^{+0.10}_{-0.09}\\,{\\rm Myr}$. This value coincides very well with the $\\tspin = 1.4\\,{\\rm Myr}$ predicted specifically for \\sorie\\ by \\citet{udD2009}, from their MHD-calibrated scaling law for spin-down times. (Such a close agreement is partly fortuitous, given the uncertainties in stellar and wind parameters). Assuming that \\tspin\\ has remained constant over the lifetime of the star implies that it can be no older than $1.16^{+0.09}_{-0.08}\\,{\\rm Myr}$ (otherwise, it would at some stage have been rotating faster than the critical rate $P_{\\rm crit} \\sim 0\\fd5$). This upper limit on the age fits within the lower portion of the $0.5-8\\,{\\rm Myr}$ age range estimated for the $\\sigma$ Orionis cluster \\citep[][and references therein]{Cab2007}.  In summary, then, we conclude that the observations are consistent with \\sorie\\ undergoing rotational braking due to its magnetized line-driven wind. This result is significant: although magnetic rotational braking is inferred from population studies of low-mass stars \\citep[e.g.,][]{DonLan2009}, direct measurement of spin-down in an individual (non-degenerate) object is noteworthy, and has been achieved so far for only handful of magnetic B and A stars \\citep[see][]{Mik2009}."
    },
    "1004/1004.4102_arXiv.txt": {
        "abstract": "Following an extremely interesting idea \\cite{R1}, published long ago, the work function associated with the emission of ultra-relativistic electrons from magnetically deformed metallic crystal of astrophysical relevance is obtained using relativistic version of Thomas-Fermi type model. In the present scenario, surprisingly, the work function becomes anisotropic; the longitudinal part is an increasing function of magnetic field strength, whereas the transverse part diverges. ",
        "introduction": "In condensed matter physics the work function is the minimum energy needed to remove an electron from a solid to a point immediately outside the solid surface or equivalently, the energy needed to move an electron from the Fermi level into vacuum. Here the word immediately means that the final electron position is far from the surface on the atomic scale but still close to the solid on the macroscopic scale. It is also known that the work function is a characteristic property for any solid face of a substance with a conduction band, which may be empty or partly filled. For a metal, the Fermi level is inside the conduction band, indicating that the band is partly filled. For an insulator, however, the Fermi level lies within the band gap, indicating an empty conduction band; in this case, the minimum energy to remove an electron is about the sum of half the band gap, and the work function.  In the free electron model the valence electrons move freely inside the metal but find a confining potential step, say $C$ at the boundary of the metal. In the system's ground state, energy levels below the Fermi energy are occupied, and those above the Fermi Level are empty. The energy required to liberate an electron in the Fermi Level is the work function and is given by $W_f=C-\\mu_e$, where $\\mu_e$ is the Fermi energy. Therefore, the work function of a metal is usually defined as the smallest energy needed to extract an electron at zero temperature. Formally, this definition is made for an infinitely large crystal plane in which one takes an electron from infinitely deep inside the crystal and brings it through the surface, infinitely far away into the vacuum. The exact definition of work function, which is valid in the atomic scale is the work done in bringing an electron from far below the surface, compared to atomic dimensions but not far compared to crystal dimensions.  The work function of metals varies from one crystal plane to another and also varies slightly with temperature. For a metal, the work function has a simple interpretation. At absolute zero, the energy of the most energetic electrons in a metal is referred to as the Fermi energy; the work function of a metal is then equal to the energy required to raise an electron with the Fermi energy to the energy level corresponding to an electron at rest in vacuum. The work function of a semiconductor or an insulator, however, has the same interpretation, but in these materials the Fermi level is in general not occupied by electrons and thus has a more abstract meaning (see \\cite{sn,fr,tb,sk,ej1,jb} as some of the fundamental papers on work function).  The work function is associated with three types of electron emission processes: photo-emission, thermionic emission and field emission. All of them have a large number of applications in various branches of science and technology; starting from electronic valves, a very old type electronic device, to modern opto-electronic devices used to convert optical signals to electrical signals etc. The other important applications are in photo-multiplier tube (PMT), CCD, etc. In the recent years, it is found that the field emission process has a lot of important applications in modern nano technology.  Similar to the down to earth application of work function in condensed matter physics and its various scientific and technological applications, the work function also plays vital role in many astrophysical processes, e.g., in the formation of magneto-sphere by the emission of electrons from the polar region induced by strong electrostatic field in strongly magnetized neutron stars. The present investigation is mainly associated with the emission of high energy electrons from the dense metallic iron crystal present at the crustal region of strongly magnetized neutron stars. The study of the formation of plasma in a pulsar magnetosphere is a quite old but still an unresolved astrophysical issue. In the formation of magneto-spheric plasma, it is generally assumed that there must be an initial high energy electron flux from the magnetized neutron stars. At the poles of a neutron star the emitted charged particles flow only along the magnetic field lines. Further a rotating magnetized neutron star generates extremely high electro-static potential difference at the poles. The flow of high energy electrons along the direction of magnetic field lines and their penetration through the light cylinder is pictured with the current carrying conductors. Naturally, if the conductor is broken near the pulsar surface the entire potential difference will be developed across a thin gap, called polar gap. This is of course based on the assumption that above a critical height from the polar gap, because of high electrical conductivity of the plasma, the electric field $E_{\\vert\\vert}$, parallel with the magnetic field near the poles is quenched. Further, a steady acceleration of electrons originating at the polar region of neutron stars, travelling along the field lines, will produce magnetically convertible curvature $\\gamma$-rays. If these curvature $\\gamma$-ray photons have energies $>2m_ec^2$ (with $m_e$ the electron rest mass and $c$ the velocity of light), then pairs of $e^--e^+$ will be produced in enormous amount with very high efficiency near the polar gap. These produced $e^--e^+$ pairs form what is known as the magneto-spheric plasma \\cite{R2,R3,R4,R5,R6,R7,R8}. The emission of electrons from the polar region of neutron stars is mainly dominated by the cold emission or the field emission \\cite{R12}, driven by electro-static force at the poles, produced by the strong magnetic field of rotating neutron stars. The work function plays a major role in the emission processes. Moreover, it is interesting to note that electrons are only emitted from the polar region of magnetized neutron stars, but never from some region far from the poles. A large number of theoretical and numerical techniques were developed to obtain analytical expression and also the numerical values of work functions of various materials for which the experimental values are also known. Self-consistent jellium-background model, embedding atom-jellium model, density functional theory (DFT) etc., were used to obtain numerical values for work functions \\cite{ek,wc,mvvv}. In a recent article the work function has been identified with the exchange energy of electrons within the material \\cite{hph,pl,rs}.  The process of extracting electrons from the outer crust region of strongly magnetized neutron stars, including the most exotic stellar objects, the magnetars, it requires a more or less exact description of the structure of matter in that region. From the knowledge of structural deformation of atoms in strong magnetic field; we expect that the departure from spherical nature to a cigar shape allows us to assume that because of high density the electron distribution around the iron nuclei at the outer crust region may be replaced by Wigner-Seitz type cells of approximately cylindrical in structure \\cite{R9,R10,R11}. We further assume that the electron gas inside the cells are strongly degenerate and at zero temperature (In presence of strong quantizing magnetic field the atoms get deformed from their usual spherical shape and also get contracted, then it is quite natural that the electron density becomes so high inside the cells that their chemical potential $\\mu_e \\gg T$, the temperature of the degenerate electron gas. Then the temperature of the system may be assumed to be very close to zero). It is well known that the presence of extraordinarily large magnetic field not only distorts the crystalline structure of dense metallic iron, also significantly modifies the electrical properties of such matter. As for example, the electrical conductivity, which is otherwise isotropic, becomes highly anisotropic in presence of strong quantizing magnetic field \\cite{pbh,po,py}. In presence of strong magnetic field, iron crystal is highly conducting in the direction parallel to the magnetic field, whereas flow of current in the perpendicular direction is severely inhibited.  The aim of this article (i) is to show that the work function associated with the emission of electrons along the direction of magnetic field at the polar region of strongly magnetized neutron stars increases with the strength of magnetic field and (ii) it will be shown that the same quantity associated with the emission of electrons in the direction transverse to the magnetic field direction is infinitely large. Therefore the main result of this article is to show that in presence of a strong quantizing magnetic field, the work function becomes anisotropic (see also \\cite{pl,rs} and \\cite{md1,md2}). We have noticed that the result does not depend on the dimension of the crustal matter. The scenario is very much analogous to the charge transport mechanism within the metal in presence of strong magnetic field. The electron emission process may be assumed to be a kind of charge transport process from inside the metal to outside. We shall show in this article that analogous to the internal charge transport, this is also anisotropic in presence of strong magnetic field and if the magnetic field strength is high enough, the emission process in the transverse direction is totally forbidden. To the best of our knowledge, using this simple idea for cold cathode emission \\cite{R1}, the study of anisotropic nature of work function in presence of strong quantizing magnetic field, relevant for strongly magnetized neutron star crustal region, has not been studied earlier. Of course, a completely different approach, popularly known as density functional theory (DFT) has been used by Lai et. al. to study properties of dense neutron star crustal matter. They have also obtained the work function as a function of magnetic field \\cite{md1,md2,ad} (see also \\cite{lk}). Surprisingly, our simple minded model gives almost same kind of magnetic field dependence for the work function as obtained by Lai et. al.  We have organized this article in the following manner: in section 2, we shall obtain an approximate analytical expression for work function associated with the emission of electrons along the field direction. In section 3, we shall obtain the same quantity corresponding to the emission transverse to the direction of magnetic field, and show that this particular component is infinitely large. Whereas in the last section we shall give conclusions and future prospects of this work. ",
        "conclusions": "In conclusion, we mention that the main purpose of this article is to show that (i) in presence of strong quantizing magnetic field the work function becomes anisotropic, (ii) the transverse part is infinitely large, (iii) the longitudinal part is finite but increases with the strength of magnetic field and finally, (iv) low field values of work functions are more or less consistent with the tabulated values. In the present work, with such simple minded model, to the best of our knowledge, the anisotropic nature of work function in presence of strong quantizing magnetic field is predicted for the first time. Further, the diverging character of work function associated with the electron emission in the transverse direction is also obtained for the first time, and so far our knowledge is concerned, it has not been reported in any published work. We have also noticed that in the low magnetic field limit (within the limitation of this model) the numerical values of work function are of the same orders of magnitude with the known laboratory data. However, we are not able to compare our results with the variation from one metal to another. The anisotropic nature of work function is apparently coming from the deformed cylindrical nature of electron distribution around iron nuclei caused by strong magnetic field at the outer crust region of a strongly magnetized neutron star. Now because of Ohmic decay, when the field strength will become low enough $(<B_c^{(e)})$, we do expect that the spherically symmetric nature of the electron distribution around the nuclei will be restored. As a consequence, the electron emission processes will no longer be affected by strong magnetic field through work function of the metallic outer crust. If the diverging nature of work function associated with the emission of electrons through the curved faces of the so called metallic atoms (they are in reality the electron distribution around the iron nuclei), is solely because of cylindrical deformation, which in the present model caused by the presence of ultra-strong magnetic field, we therefore expect that for the emission of electrons through the curved surface of any object having cylindrical structure, the transverse component of work function will be much larger compared to the longitudinal part. As a consequence the electron emission current along the axial direction will be large enough compared to the transverse emission current \\cite{pl}. The work function associated with the emission of electrons along the direction of magnetic field may be used to obtain electron field emission current at the polar region of strongly magnetized neutron stars, which will help us to understand the formation of magneto-spheric plasma."
    },
    "1004/1004.3208_arXiv.txt": {
        "abstract": "{Gamma Ray Burst models predict the broadband spectral evolution and the temporal evolution of the energy flux. In contrast, standard data analysis tools and data repositories provide count-rate data, or use single flux conversion factors for all of the data, neglecting spectral evolution.} {To produce \\Swift\\ BAT and XRT light curves in flux units, where the spectral evolution is accounted for.} {We have developed software to use the hardness ratio information to track spectral evolution of GRBs, and thus to convert the count-rate light curves from the BAT and XRT instruments on \\Swift\\ into accurate, evolution-aware flux light curves.} {The \\Swift\\ Burst Analyser website\\thanks{http://www.swift.ac.uk/burst\\_analyser} contains BAT, XRT and combined BAT-XRT flux light curves in three energy regimes for all GRBs observed by the \\Swift\\ satellite. These light curves are automatically built and updated when data become available, are presented in graphical and plain-text format, and are available for download and use in research.} {} ",
        "introduction": "\\label{sec:intro}  Gamma Ray Bursts are the most powerful explosions known, and the \\Swift\\ satellite (Gehrels \\etal2004) has revolutionised our understanding of them, both in filling some gaps in our knowledge and raising new questions and challenges to theory. See Zhang (2007) for a recent review of GRBs and the advances made by \\Swift.  One of the difficulties inherent in confronting theory with the wealth of data that \\Swift\\ has produced is that models predict how the flux and spectrum of a GRB or its afterglow will evolve, whereas the data are in units of count rate over some bandpass; in the presence of spectral evolution the count-rate cannot be seen as a proxy for the flux.  It is thus desirable to create GRB flux light curves\\footnote{We create light curves in units both of flux and flux density; however for concision, we use the collective phrase, `flux light curves' to refer to both types.} which employ a time-dependent flux conversion factor to account for spectral evolution. Also, since the bandpasses of the \\Swift\\ Burst Alert Telescope (BAT; Barthelmy \\etal 2005; bandpass: 15--350 keV) and X-ray Telescope (XRT; Burrows \\etal2005; bandpass: 0.3--10 keV), are close to each other, it is often informative to consider the two instruments' data together, extrapolated to a single bandpass, be it the XRT band (e.g. O'Brien \\etal2006; Willingale \\etal2007) or part of the BAT bandpass (e.g. 15--25 keV, Sakamoto \\etal2007).  We have therefore created the \\Swift\\ Burst Analyser. In this first paper relating to the facility, we present an online repository of BAT and XRT unabsorbed flux light curves in three energy regimes: 0.3--10 keV flux, 15--50 keV flux, and the flux density at 10 keV.  The Burst Analyser provides BAT and XRT flux light curves separately and combined; an example is shown in Fig.~\\ref{fig:example}. It also includes a time evolution history for each instrument of the counts-to-flux conversion factor, and of the spectral photon index, \\G\\ (i.e.\\ for a power-law spectrum with the number of photons at energy $E$ is given by $N(E) dE \\propto E^{-\\Gamma}$). In addition, we provide BAT flux light curves where spectral evolution is not included, for comparison with the non-evolving XRT data already available online\\footnote{http://www.swift.ac.uk/xrt\\_curves\\label{foot:lc}} (Evans \\etal2007, 2009). One example of the advantage of considering spectral evolution is shown in Fig.~\\ref{fig:paper2}; for GRB 060729 the rapid decay phase is steeper when viewed in flux space with spectral evolution accounted for, which allows us to see the turn-on of the afterglow. (Note that Grupe \\etal2007, 2010 also accounted for spectral evolution in their analysis of this GRB. The evolution-induced dip feature shown in Fig.~\\ref{fig:paper2} is less prominant in the 0.3--10 keV band which is why it is not seen in their analysis.) We consider the physical interpretation of GRB light curves enabled by our new method in Paper II (Evans \\etal in prep.).  All of the data created by the Burst Analyser are available from:  \\begin{center}{\\bf http://www.swift.ac.uk/burst\\_analyser}\\end{center}  \\noindent This includes graphical plots and the data in plain-text format.  In this paper we introduce the \\Swift\\ Burst Analyser. In Section~\\ref{sec:create} we explain how the light curves are created and spectral evolution accounted for. Section~\\ref{sec:caveats} highlights the limitations of our method and some recommended checks users should apply. We also give details of when the light curves are created and how they can be accessed (Section~\\ref{sec:avail}), and the usage policy (Section~\\ref{sec:usage}).  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90,width=8.2cm]{fig1.eps}} \\caption{An example of a flux density light curve from the Burst Analyser. The spectrally-evolving BAT-XRT flux light curve of GRB 061121, is shown as unabsorbed flux density at 10 keV; the lower panel shows the evolution of the photon index of the power-law spectrum. The last 3 BAT data points suffer from a poorly-constrained spectrum; see Section~\\ref{sec:errors} for details.} \\label{fig:example} \\end{figure}   \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics[angle=-90,width=8.2cm]{fig2.eps}} \\caption{The flux density light curve of GRB 060729 from the Burst Analyser. Accounting for spectral evolution shows the flux decline during the steep decay to be more rapid than in count space, and reveals the turn-on of the afterglow.} \\label{fig:paper2} \\end{figure} ",
        "conclusions": ""
    },
    "1004/1004.3722_arXiv.txt": {
        "abstract": "\\noindent We propose a six dimensional Universal Extra Dimensions (UED) model compactified on a real projective plane $RP^2$, a two-sphere with its antipodal points being identified. We utilize the Randjbar-Daemi-Salam-Strathdee spontaneous sphere compactification with a monopole configuration of an extra $U(1)_X$ gauge field that leads to a spontaneous radius stabilization. Unlike the sphere and the {\\color{black} $S^2/Z_2$ orbifold} compactifications, the massless $U(1)_X$ gauge boson is safely projected out. We show how a compactification on a non-orientable manifold results in a chiral four dimensional gauge theory by utilizing 6D chiral gauge and Yukawa interactions. The resultant Kaluza-Klein mass spectra are distinct from the ordinary UED models compactified on torus. We briefly comment on the anomaly cancellation and also on a possible dark matter candidate in our model. ",
        "introduction": "\\label{Introduction_and_Conclusions} Universal Extra Dimensions (UED) allow fairly lower Kaluza-Klein (KK) scale even after taking into account the precision electroweak constraints~\\cite{Appelquist:2000nn,Appelquist:2002wb,Gogoladze:2006br}, comparing to a model where some of the Standard Model (SM) fields are localized in the extra dimensions, and therefore the UED can be directly tested at the CERN Large Hadron Collider (LHC) which is just being started. Furthermore the Lightest KK Particle (LKP) in the UED model can serve as a natural dark matter candidate~\\cite{Cheng:2002iz,Servant:2002aq}. A six dimensional (6D) UED model is particularly attractive since three generations of SM fermion are naturally required~\\cite{Dobrescu:2001ae,Borghini:2001sa} for the cancellation of $SU(2)_W$ global anomaly, which is analogous to the Witten anomaly.\\footnote{\\label{anomaly_footnote} This three-generation argument of global anomaly is based on the computation with $S^6$~\\cite{Bershadsky:1997sb}. Strictly speaking, when extra dimensions are compactified on $B$, analysis must be redone on $S^4\\times B$. We thank S.\\ Yamaguchi for pointing out this issue. }   In 1953, Pauli initiated the compactification on two-sphere $S^2$~\\cite{O'Raifeartaigh:1998pk,Straumann:2000zc} and DeWitt further developed it in 1963 (see~\\cite{Pons:2006vz} for a historical review). After a half century, it has been still attracting interest as a beautiful tool to incorporate the spontaneous compactification (see~\\cite{RandjbarDaemi:1982hi} and citations thereof), the gauge-Higgs unification~\\cite{Manton:1979kb,Hosotani:1983vn,Hatanaka:1998yp,Dvali:2001qr,Lim:2006bx,Nomura:2008sx}, supergravity models~\\cite{Salam:1984cj,Fujii:1986nh} (see also~\\cite{Salvio:2007mb} and references therein), etc. Recently, a UED model is constructed on the {\\color{black} $S^2/Z_2$ orbifold}~\\cite{Maru:2009wu}, in which there exists a stable particle due to the ``$Z_2'$ KK parity,'' potentially being a possible dark matter candidate.  In this paper we first consider a UED model compactified on the two-sphere, based on the spontaneous compactification mechanism by Randjbar-Daemi, Salam and Strathdee (RSS), which utilizes a monopole configuration of an extra $U(1)_X$ gauge field in the two-sphere~\\cite{RandjbarDaemi:1982hi}. A remarkable feature of the RSS sphere compactification is that, in contrast to the ordinary UED models on torus, the compactification radius is spontaneously fixed by the monopole configuration.\\footnote{ See Refs.~\\cite{Okada:1984cv,Okada:1984sf} for cosmology of the RSS model. } We show that our model can supply a possible dark matter candidate, due to what we call the KK angular momentum. Actually the stable particle is exactly the same one as in~\\cite{Maru:2009wu} and the so-called $Z_2'$ parity conservation is nothing but the result of the KK angular momentum conservation.\\footnote{ We also note that our two-sphere compactification can serve chiral fermions~\\cite{RandjbarDaemi:1982hi} so that the $Z_2$ orbifold projection in~\\cite{Maru:2009wu}, which is $(\\theta,\\phi)\\sim(\\pi-\\theta,-\\phi)$ in the standard polar coordinates, is rather redundant. }   However, the two-sphere compactification suffers from the existence of massless $U(1)_X$ gauge boson, which obviously contradicts existing experimental observation~\\cite{Smith:1999cr} (see also~\\cite{Adelberger:2009zz} for a review).\\footnote{\\label{anomaly_giving_superheavy} In the {\\color{black} $S^2/Z_2$ orbifold} model~\\cite{Maru:2009wu}, there exists exactly the same problem and it is assumed that {\\color{black} the SM-singlet $\\mc N_-$ is not put and hence} the $U(1)_X$ is broken by an anomaly~\\cite{Scrucca:2003ra}. It is somewhat puzzling how a classical gauge configuration is allowed when $U(1)_X$ is broken by an anomaly at an UltraViolet (UV) scale $\\Lambda$. That is, if we have a gauge boson mass term like $\\Lambda^2X_MX^M$~\\cite{Scrucca:2003ra}, it will spoil the monopole vacuum expectation value (vev)~\\cite{RandjbarDaemi:1982hi} due to the extra contribution to the energy momentum tensor in the Einstein equation. } As a remedy, we propose a compactification on a real projective plane $RP^2$ that is a two-sphere with its antipodal points being identified: $(\\theta,\\phi)\\sim(\\pi-\\theta,\\phi+\\pi)$.\\footnote{ A compactification on a flat projective plane is considered in a different context in~\\cite{Hebecker:2003we}, where all the curvature is put into conical singularities. {\\color{black} This singular space is also utilized to construct another UED model~\\cite{Cacciapaglia:2009pa}.} } By this identification, the unwanted zero mode of $U(1)_X$ gauge field is projected out, as well as the first KK photon that is a possible dark matter candidate in the $S^2$ and $S^2/Z_2$ UED models. In our $RP^2$ model, $U(1)_X$ gauge interaction is made anomaly free and the first KK mode of $U(1)_X$ gauge field becomes a possible dark matter candidate.\\footnote{ In 5D UED model, radiative corrections tend to make gauge bosons lighter than fermions~\\cite{Cheng:2002iz}. If this turns out to be the case in our model, the $U(1)_X$ KK gauge boson becomes the dark matter. Generically one expects that there are also tree level threshold corrections at UV cut-off scale. In that case, the first KK neutrino might become the dark matter. } We also discuss possible cosmological constraint from the decay of the second lightest KK particle into the dark matter.  The $RP^2$ is a non-orientable manifold and so is the whole space $M^4\\times RP^2$, where $M^4$ is the four dimensional (4D) Minkowski space. It is standard lore that a quantum field theory must be either $P$ or $CP$ symmetric on a non-orientable manifold and hence cannot account for the SM~\\cite{Zeldovitch_Novikov,Ellis:1970ey}. This argument is based on the fact that a non-orientable manifold necessarily has a closed loop that gives one a parity transformation $P$ when one circulates around it. Therefore, the Lagrangian must be symmetric under either $P$ or $CP$ transformation for the theory to be single-valued on the space. In four dimensions, this is obviously in contradiction to the experimentally observed $P$ and $CP$ violation.   The key observation is that the higher dimensional $CP$ transformation is different from the four dimensional one. We can identify the field after the circulation of the closed loop by an outer automorphism~\\cite{Hebecker:2001jb} with the 6D-$CP$ transformation. {\\color{black} Even if a} theory is vector-like (non-chiral) under the 6D-$CP$ transformation, it can be chiral under the 4D-$P$ and $CP$ transformations after the compactification. {\\color{black} The (4D) $CP$ violation in Yukawa sector results} from the fact that we can identify SM spinors by 6D-$CP$ transformation and Higgs by 6D-$P$ transformation.   We obtain the KK mass spectrum that is distinctive from those in the 5D and 6D UED models with orbifolded toroidal compacitifications: $S^1/Z_2$ and $T^2/(Z_2\\times Z_2')$, respectively. From the conservation of the KK angular momentum, the single production of a KK mode from SM matters cannot occur, in contrast to the ordinary UED where only KK parity is conserved and hence one can produce a single KK-even mode from a scattering of zero modes. Our model contains a possible dark matter candidate, the first KK mode of the $U(1)_X$ gauge field, as well as the more conventional one, the fist KK neutrino, depending on the  input at the UV cutoff. ",
        "conclusions": "\\label{Introduction_and_Conclusions} Universal Extra Dimensions (UED) allow fairly lower Kaluza-Klein (KK) scale even after taking into account the precision electroweak constraints~\\cite{Appelquist:2000nn,Appelquist:2002wb,Gogoladze:2006br}, comparing to a model where some of the Standard Model (SM) fields are localized in the extra dimensions, and therefore the UED can be directly tested at the CERN Large Hadron Collider (LHC) which is just being started. Furthermore the Lightest KK Particle (LKP) in the UED model can serve as a natural dark matter candidate~\\cite{Cheng:2002iz,Servant:2002aq}. A six dimensional (6D) UED model is particularly attractive since three generations of SM fermion are naturally required~\\cite{Dobrescu:2001ae,Borghini:2001sa} for the cancellation of $SU(2)_W$ global anomaly, which is analogous to the Witten anomaly.\\footnote{\\label{anomaly_footnote} This three-generation argument of global anomaly is based on the computation with $S^6$~\\cite{Bershadsky:1997sb}. Strictly speaking, when extra dimensions are compactified on $B$, analysis must be redone on $S^4\\times B$. We thank S.\\ Yamaguchi for pointing out this issue. }   In 1953, Pauli initiated the compactification on two-sphere $S^2$~\\cite{O'Raifeartaigh:1998pk,Straumann:2000zc} and DeWitt further developed it in 1963 (see~\\cite{Pons:2006vz} for a historical review). After a half century, it has been still attracting interest as a beautiful tool to incorporate the spontaneous compactification (see~\\cite{RandjbarDaemi:1982hi} and citations thereof), the gauge-Higgs unification~\\cite{Manton:1979kb,Hosotani:1983vn,Hatanaka:1998yp,Dvali:2001qr,Lim:2006bx,Nomura:2008sx}, supergravity models~\\cite{Salam:1984cj,Fujii:1986nh} (see also~\\cite{Salvio:2007mb} and references therein), etc. Recently, a UED model is constructed on the {\\color{black} $S^2/Z_2$ orbifold}~\\cite{Maru:2009wu}, in which there exists a stable particle due to the ``$Z_2'$ KK parity,'' potentially being a possible dark matter candidate.  In this paper we first consider a UED model compactified on the two-sphere, based on the spontaneous compactification mechanism by Randjbar-Daemi, Salam and Strathdee (RSS), which utilizes a monopole configuration of an extra $U(1)_X$ gauge field in the two-sphere~\\cite{RandjbarDaemi:1982hi}. A remarkable feature of the RSS sphere compactification is that, in contrast to the ordinary UED models on torus, the compactification radius is spontaneously fixed by the monopole configuration.\\footnote{ See Refs.~\\cite{Okada:1984cv,Okada:1984sf} for cosmology of the RSS model. } We show that our model can supply a possible dark matter candidate, due to what we call the KK angular momentum. Actually the stable particle is exactly the same one as in~\\cite{Maru:2009wu} and the so-called $Z_2'$ parity conservation is nothing but the result of the KK angular momentum conservation.\\footnote{ We also note that our two-sphere compactification can serve chiral fermions~\\cite{RandjbarDaemi:1982hi} so that the $Z_2$ orbifold projection in~\\cite{Maru:2009wu}, which is $(\\theta,\\phi)\\sim(\\pi-\\theta,-\\phi)$ in the standard polar coordinates, is rather redundant. }   However, the two-sphere compactification suffers from the existence of massless $U(1)_X$ gauge boson, which obviously contradicts existing experimental observation~\\cite{Smith:1999cr} (see also~\\cite{Adelberger:2009zz} for a review).\\footnote{\\label{anomaly_giving_superheavy} In the {\\color{black} $S^2/Z_2$ orbifold} model~\\cite{Maru:2009wu}, there exists exactly the same problem and it is assumed that {\\color{black} the SM-singlet $\\mc N_-$ is not put and hence} the $U(1)_X$ is broken by an anomaly~\\cite{Scrucca:2003ra}. It is somewhat puzzling how a classical gauge configuration is allowed when $U(1)_X$ is broken by an anomaly at an UltraViolet (UV) scale $\\Lambda$. That is, if we have a gauge boson mass term like $\\Lambda^2X_MX^M$~\\cite{Scrucca:2003ra}, it will spoil the monopole vacuum expectation value (vev)~\\cite{RandjbarDaemi:1982hi} due to the extra contribution to the energy momentum tensor in the Einstein equation. } As a remedy, we propose a compactification on a real projective plane $RP^2$ that is a two-sphere with its antipodal points being identified: $(\\theta,\\phi)\\sim(\\pi-\\theta,\\phi+\\pi)$.\\footnote{ A compactification on a flat projective plane is considered in a different context in~\\cite{Hebecker:2003we}, where all the curvature is put into conical singularities. {\\color{black} This singular space is also utilized to construct another UED model~\\cite{Cacciapaglia:2009pa}.} } By this identification, the unwanted zero mode of $U(1)_X$ gauge field is projected out, as well as the first KK photon that is a possible dark matter candidate in the $S^2$ and $S^2/Z_2$ UED models. In our $RP^2$ model, $U(1)_X$ gauge interaction is made anomaly free and the first KK mode of $U(1)_X$ gauge field becomes a possible dark matter candidate.\\footnote{ In 5D UED model, radiative corrections tend to make gauge bosons lighter than fermions~\\cite{Cheng:2002iz}. If this turns out to be the case in our model, the $U(1)_X$ KK gauge boson becomes the dark matter. Generically one expects that there are also tree level threshold corrections at UV cut-off scale. In that case, the first KK neutrino might become the dark matter. } We also discuss possible cosmological constraint from the decay of the second lightest KK particle into the dark matter.  The $RP^2$ is a non-orientable manifold and so is the whole space $M^4\\times RP^2$, where $M^4$ is the four dimensional (4D) Minkowski space. It is standard lore that a quantum field theory must be either $P$ or $CP$ symmetric on a non-orientable manifold and hence cannot account for the SM~\\cite{Zeldovitch_Novikov,Ellis:1970ey}. This argument is based on the fact that a non-orientable manifold necessarily has a closed loop that gives one a parity transformation $P$ when one circulates around it. Therefore, the Lagrangian must be symmetric under either $P$ or $CP$ transformation for the theory to be single-valued on the space. In four dimensions, this is obviously in contradiction to the experimentally observed $P$ and $CP$ violation.   The key observation is that the higher dimensional $CP$ transformation is different from the four dimensional one. We can identify the field after the circulation of the closed loop by an outer automorphism~\\cite{Hebecker:2001jb} with the 6D-$CP$ transformation. {\\color{black} Even if a} theory is vector-like (non-chiral) under the 6D-$CP$ transformation, it can be chiral under the 4D-$P$ and $CP$ transformations after the compactification. {\\color{black} The (4D) $CP$ violation in Yukawa sector results} from the fact that we can identify SM spinors by 6D-$CP$ transformation and Higgs by 6D-$P$ transformation.   We obtain the KK mass spectrum that is distinctive from those in the 5D and 6D UED models with orbifolded toroidal compacitifications: $S^1/Z_2$ and $T^2/(Z_2\\times Z_2')$, respectively. From the conservation of the KK angular momentum, the single production of a KK mode from SM matters cannot occur, in contrast to the ordinary UED where only KK parity is conserved and hence one can produce a single KK-even mode from a scattering of zero modes. Our model contains a possible dark matter candidate, the first KK mode of the $U(1)_X$ gauge field, as well as the more conventional one, the fist KK neutrino, depending on the  input at the UV cutoff."
    },
    "1004/1004.2042_arXiv.txt": {
        "abstract": "Short Gamma-Ray Bursts (SGRBs) are expected to form from the coalescence of compact binaries, either of primordial origin or from dynamical interactions in globular clusters. In this paper, we investigate the possibility that the offset and afterglow brightness of a SGRB can help revealing the origin of its progenitor binary. We find that a SGRB is likely to result from the primordial channel if it is observed within 10 kpc from the center of a massive galaxy and shows a detectable afterglow. The same conclusion holds if it is 100 kpc away from a small, isolated galaxy and shows a weak afterglow. On the other hand, a dynamical origin is suggested for those SGRBs with observable afterglow either at a large separation from a massive, isolated galaxy or with an offset of 10-100 kpc from a small, isolated galaxy. We discuss the possibility that SGRBs from the dynamical channel are hosted in intra-cluster globular clusters and find that GRB 061201 may fall within this scenario. ",
        "introduction": "The afterglows of several Short Gamma Ray Bursts (SGRBs) have recently been localized on the sky with Swift (Gehrels et al. 2004), allowing for a determination of their redshift and host galaxy (see e.g. Berger 2009; Fong, Berger \\& Fox 2010).  SGRBs are been observed in all kinds of galaxies (from star-bursts to ellipticals and also associated with galaxy clusters) with a wide range of offsets.  SGRBs are currently believed to result from the coalescence of a compact binary, either a double neutron star or a neutron star (NS) and a black hole (BH) binary (Nakar 2007). Compact binaries are known to form in different environments along two main channels: (i) from the evolution of massive stars in primordial binaries rising in the galactic field (Narayan, Paczynski \\& Piran 1992), and (ii) through three (or a few)-body dynamical interactions among stars and compact remnants in globular clusters (GCs) (Grindlay, Portegies Zwart \\& McMillan 2006). Salvaterra et al. (2008) have shown that both formation channels are needed in order to reproduce the {\\it Swift} redshift distribution of SGRBs, with the dynamical channel in GCs contributing mainly at $z\\simlt 0.3$.  In this paper, we investigate the nature of the observed offsets in relation to the galaxy type and burst environment in order to discriminate between the two channels and highlight the origin of SGRBs.  A wide range of offsets is expected in both channels.  In the primordial scenario, a large separation from the host galaxy can originate from the natal kick rising at the time of formation of the compact object (Belczynski et al. 2006). This mechanism is not present in case of dynamical origin: few body interactions that are at the origin of dynamical double neutron stars do not release binaries with large recoil velocities (Devecchi et al. in preparation). So, the offset in this latter case has to be ascribed to the underlying GC spatial distribution. For isolated galaxies this reflects the GC distribution that is known to decline more gently compared to stars (see Brodie \\& Strader 2006 and references therein).  In galaxy clusters, GCs may have a wider spread, since there are hints (both theoretical and observational) that a population of intra-cluster GCs (ICGCs) can exist. SGRBs inside ICGCs can explain large potential offsets in galaxy groups and clusters, besides natal kicks.   We outline here three different cases for the origin of the offsets: (i) primordial SGRB kicked from a galaxy, isolated (Belczynski et al. 2006) or in a cluster (Niino \\& Totani 2008), by a natal kick; (ii) dynamical SGRB in a GC bound to an isolated or cluster galaxy, and (iii) dynamical SGRB in a ICGC. ",
        "conclusions": "In this paper, we showed how the combination of the observed spatial offset, host type/mass and afterglow brightness can shed light on the pathway of formation of SGRB binary progenitors, resulting either from primordial stellar binaries, or from dynamical interactions in GCs. The main results of our study are summarized in Fig.~3. Key results of our analysis are:  \\begin{itemize} \\item{If a SGRB is observed with a large separation from a massive, isolated galaxy we expect that it belongs to the dynamical channel and, residing in a globular cluster, it may show an observable X-ray and optical afterglow;} \\item{A SGRB exploding with a large offset from a small, isolated host galaxy is likely to ensue from the primordial channel. In this case the progenitor binary is ejected in the intergalactic medium following a large natal kick.  Nor or weak afterglow is expected given the very low density of the medium;} \\item{A SGRB observed within 10 kpc from the center of a massive galaxy (either isolated or in a cluster) should arise from the primordial channel with a detectable afterglow;} \\item{A SGRB associated to a galaxy cluster but not to any of its specific members may result from both formation scenarios. A possible discriminant could be the detection of an optical afterglow that is expected to be brighter in the case of the denser intra-cluster globular cluster natal environment;} \\item{A SGRB with an offset of 10-100 kpc from a small, isolated galaxy is more probably originated from the dynamical channel and this may be confirmed from the identification of the afterglow.} \\end{itemize}   We now apply our results to a few observed SGRBs. Troja et al. (2008) suggested that SGRBs showing an extended-duration soft emission component in their prompt emission preferentially have small projected physical offsets.  Recent \\textit{HST} observations of a sample of SGRB host galaxies show that the distribution of offsets has a median of $\\sim 5$ kpc, about 5 times larger than for long GRBs, with no evidence of differences between SGRBs with and without extended emission (Fong et al. 2010).  In the case of a host galaxy at low redshift, deep optical/NIR imaging can allow to resolve the galaxy surface brightness profile; the detection of an optical afterglow can then clearly pinpoint the SGRB position with respect to the host galaxy. This is the case of GRB\\,071227 ($z = 0.38$) and GRB\\,060505 ($z = 0.09$). GRB\\,071227 was firmly classified as a SGRB (Sato et al. 2007, Golenetskii et al. 2007; Onda et al.  2008) and occurred on the plane of a large ($r \\sim 15$ kpc) spiral galaxy at a relatively large offset ($\\sim 15$ kpc; D'Avanzo et al 2009; Fong et al 2010). This could equally favor both the primordial or dynamical formation channel for its progenitor (Fig.~3). We note however that its location on the galactic plane and within the light of the host likely favors a primordial origin. The case of GRB\\,060505 is even more interesting, given that the classification of this GRB is still debated. The duration ($T_{90} \\sim 4$ s) and the spectral lag point towards a long GRB classification (McBreen et al. 2008; Xu et al 2009), the non-detection of an associated supernova down to deep limits favors for a SGRB (Ofek et al. 2007), while detailed study of its host galaxy led to different interpretations of its progenitor (Levesque \\& Kewley 2007; Thoene et al. 2008).  Nevertheless, we note that GRB\\,060505 is an outlier of the $E_{\\rm p,i}-E_{\\rm iso}$ correlation like all SGRBs of known redshift and peak energy $E_{\\rm p,i}$ (Amati et al.  2007). In the scenario of a double compact object merger progenitor for GRB\\,060505, the position of the afterglow at an offset of 6.5 kpc from a large spiral galaxy (Levesque \\& Kewley 2007; Thoene et al. 2008), makes a primordial origin highly probable (Fig.~3). A primordial origin is also a valuable hypothesis for the progenitor of the farthest short-duration ($T_{90} \\sim 1.3$ s) GRB, occurred at $z \\sim 2.6$ (GRB\\,090426; Levesque et al. 2010; Antonelli et al. 2009). The position of the optical afterglow inside the host galaxy, the intrinsic absorption measured in the X-ray spectrum, and the redshift could be indicative of a ``primordial'' binary system that merged in a relatively short time ($10^7-10^8$ yr). On the other hand, a core-collapse origin for this burst cannot be excluded in light of its consistency with the $E_{\\rm p,i}-E_{\\rm iso}$ correlation that holds for long GRBs (Antonelli et al.  2009), making the classification of GRB\\,090426 not straightforward.  As discussed above, the measure of the redshift of a SGRB is a key issue to accurately evaluate its offset and to relate it with the size of its host galaxy. However, some conclusions can also be drawn for those SGRBs with no measured redshift. The detection of an optical afterglow with sub-arsec precision\\footnote{We note that most of the largest offsets inferred for SGRBs are based on X-ray positions, while only an optical (sub-arcsec) localization could give a firm GRB-galaxy association.} coincident with the profile of a galaxy (see, e.g. GRB\\,070707 and GRB\\,051227; Piranomonte et al. 2008, D'Avanzo et al. 2009) strongly hints for an association between the two objects. For these bursts, offsets are consistent with zero, and the SGRB positions follow the host light. From a statistical point of view the majority of these systems probably originate from the primordial channel as a population of SGRBs exploding inside GCs should result in larger offsets.     A few interesting counter examples do exist.  GRB~061201 and GRB~070809 are two SGRBs for which the detection of an optical afterglow has allowed to determine their position with high accuracy. Deep, follow-up observations fail to find a host galaxy coincident with the optical afterglow. The field of GRB~061201 has been studied by Stratta et al. (2007) with VLT and Fong et al. (2010) with HST.  No host is revealed down to R(AB)$=26.1$ (Stratta et al. 2007).  Two galaxies are in the field of GRB~061201: a spiral galaxy at $z=0.111$ (Stratta et al. 2007) at a projected distance of $\\sim 32.5$ kpc and a fainter object with undetermined redshift at 1.8$^{\\prime\\prime}$ (Fong et al. 2010).  Although the ejection hypothesis from one of these host candidates can not be excluded, we note that the burst should have likely exploded in a gas poor environment far away from the host. This is in contrast with the relative bright afterglow observed. GRB~061201 is also know to be 0.9 Mpc away from the galaxy cluster ACO S 995 at $z=0.0865$. While the ejection hypothesis from the cluster center appears unlikely due to the high kick velocity required, we note that a sizable fraction ($\\sim 30$\\%, see Fig.~2) of the ICGC population is still present at $R>0.9$ Mpc. We suggest that GRB~061201 originates from the dynamical channel inside a ICGC of ACO S 995. The denser environment inside the host GC may be responsible for the brightness of its optical afterglow.  A second interesting case is GRB~070809 that is likely a SGRB (see Barthelmy et al. 2007). Similar to GRB~061201, this burst shows an optical afterglow but no underlying host galaxy to g(AB)=26.3 (Perley et al. 2008).  A possible host candidate has been identified in a small spiral galaxy ($\\sim 2\\times 10^{10}\\;\\msun$) at $z=0.2187$ (Perley et al. 2008)\\footnote{Note the existence of another, very faint host candidate with undetermined redshift $\\sim 2.3^{\\prime\\prime}$ away from the optical afterglow position (Perley et al. 2008).}. At this redshift, the projected distance is $\\sim 20$ kpc. Accordingly to Fig.~3 the probability of detecting a SGRB with this offset from a small galaxy is rather low ($\\sim 10$\\%) for the primordial channel. A dynamical origin is preferred in this case. As for GRB~061201, the optical detection of the afterglow supports this interpretation. We also notice that the dynamical channel is expected to contribute most to the SGRB population at the redshift of the putative host of GRB~070809 (Salvaterra et al. 2008).  These two examples show how the results presented in this paper can provide a powerful tool to discriminate the origin of SGRBs."
    },
    "1004/1004.0047_arXiv.txt": {
        "abstract": "{} {Employing photometric rotation periods for solar-type stars in NGC 1039 [M\\,34], a young, nearby open cluster, we use its mass-dependent rotation period distribution to derive the cluster's age in a distance independent way, i.e., the so-called gyrochronology method.} {We present an analysis of 55 new rotation periods, using light curves derived from differential photometry, for solar type stars in the open cluster NGC 1039 [M\\,34]. We also exploit the results of a recently-completed, standardized, homogeneous BVIc CCD survey of the cluster, performed by the Indiana Group of the WIYN open cluster survey, in order to establish photometric cluster membership and assign B-V colours to each photometric variable.  We describe a methodology for establishing the gyrochronology age for an ensemble of solar-type stars. Empirical relations between rotation period, photometric colour and stellar age (gyrochronology) are used to determine the age of M\\,34. Based on its position in a colour-period diagram, each M\\,34 member is designated as being either a solid-body rotator ({\\em interface or I-star}), a differentially rotating star ({\\em convective or C-star}) or an object which is in some transitory state in between the two ({\\em gap or g-star}). Fitting the period and photometric colour of each I-sequence star in the cluster, we derive the cluster's mean gyrochronology age.} {Of the photometric variable stars in the cluster field, for which we derive a period, 47 out of 55 of them lie along the loci of the cluster main sequence in V/B-V and V/V-I space. We are further able to confirm kinematic membership of the cluster for half of the periodic variables [21/55], employing results from an on-going radial velocity survey of the cluster. For each cluster member identified as an I-sequence object in the colour-period diagram, we derive its individual gyrochronology age, where the mean gyro age of M\\,34 is found to be $193 \\pm  9$ Myr.} {Using differential photometry, members of a young open cluster can be easily identified in a colour-magnitude diagram from their periodic photometric variability alone. Such periodicity can be used to establish a period-colour distribution for the cluster, which for M\\,34, we have used to derive its gyrochronology age of $193 \\pm  9$ Myr. Formally, our gyro age of M\\,34 is consistent (within the errors) with that derived using several {\\em distance-dependent}, photometric isochrone methods ($250 \\pm 67$ Myr). } ",
        "introduction": "Over the past century, as instrumentation and detector technology have advanced, the methods by which stellar age is determined have also evolved. While we still rely heavily upon the use of theoretical model isochrones, fitting empirical luminosity-temperature data (on an H-R diagram - e.g., Sandage 1958; Demarque \\& Larson 1964; Demarque \\& Gisler 1975, Mengel et al. 1979; and more latterly, Meynet, Mermilliod \\& Maeder 1993; D'Antona \\& Mazzitelli 1997; Naylor 2009), several alternative techniques for determining stellar age now exist which exploit a more diverse assembly of fundamental stellar properties, such as magnetic activity, elemental abundances, white dwarf cooling time-scales and angular momentum content.   Inspired by the pioneering efforts of Wilson (1963), observations of solar type stars have shown that there is a clear age dependence in the strength of magnetic activity indicators (eg., \\caiihk and X-rays), whose emissions are presumably due to solar-like magnetic field activity (eg., Wilson 1964; Wilson \\& Skumanich 1964; Wilson 1966; Skumanich 1972; Barry et al. 1981; Noyes et al. 1984; Maggio et al. 1987; Soderblom, Duncan \\& Johnson 1991; Henry et al. 1996; Mamajek \\& Hillenbrand 2008). While there is not a one-to-one relationship between activity and age for any given star, current data samples unequivocally show there to be a strong correlation between decreasing magnetic activity and increasing age.   However, even such a well-observed phenomenon does not deliver unequivocal results, as illustrated for instance by the factor of two range in ages derived by Giampapa et al. (2006) for photometric and kinematic members of the 4Gyr-old open cluster M\\,67. Moreover, stellar activity cycles, akin to those occurring in the Sun, can also lead to difficulties in establishing ages for individual stars. For instance, Soderblom et al. (1991) show that errors on chromospheric ages are roughly $50\\%$. Presumably, this error could be driven down statistically in a cluster by measuring many stars, and by removing close binaries or other contaminants. Soderblom, Jones \\& Fischer (2001) have provided chromospheric H$\\alpha$ measurements for a significant number of M\\,34 stars, but to our knowledge, a chromospheric age for M\\,34 has not yet been published.   More recently, age determinations for young open clusters have been advanced based on an analysis of their light element abundance distributions, specifically using lithium as a tracer. This technique, the so-called lithium depletion boundary method (Basri, Marcy \\& Graham 1996; Rebolo et al. 1996; Stauffer, Schultz \\& Kirkpatrick 1998; Barrado y Navascu\\'{e}s, Stauffer \\& Patten 1999), does however have both merits and limitations. Line strengths and elemental abundance ratios are fairly straight forward to measure, and are indeed fundamental stellar properties, perhaps more so than compared to magnetic activity indicators. Unfortunately, the strength of this technique lies in detecting lithium in M-stars and cooler, which for any group of stars at a given distance are among the faintest. Even for nearby open clusters, detecting the lithium boundary thus requires use of precious 8-10m class telescope time. Perhaps the ultimate limitation for this technique is that it is only discernibly sensitive to stars of \\lsi 250 Myr (\\eg Stauffer et al. 1998), severely limiting its widespread applicability.   The use of heavy-element abundances to establish stellar age has also been attempted, however its implementation is as yet limited. Cayrel et al. (2001) advocate using the detection of the singly-ionized uranium 238 line at 3859.6\\r{A} as an age proxy, although as they themselves note, the accuracy of this technique is currently limited by deficiencies in the nuclear data - \\eg poorly-known oscillator strengths. Even with the advent of well-calibrated atomic data for uranium species, dating of open clusters using detection of stellar uranium will only be feasible for clusters older than the Hyades [600Myr], due to the prohibitively long half-lives of the two dominant naturally occurring uranium species ($\\tau_{1/2}= 7.04\\times10^{8}$ \\& $4.47\\times10^{9}$ yrs for U$^{235}$ \\& U$^{238}$ respectively - CRC 2005).   Spectroscopic and photometric observations of white dwarfs, in the Galactic field and in globular and open clusters, can also yield an age estimate of the their stellar content. The detailed physics of how white dwarfs cool as they age is now well developed (Mestel 1952; Cox \\& Giuli 1965; Beaudet, Petrosian \\& Salpeter 1967), and their model-dependent cooling time-scales are well described using theoretical stellar models (Koester \\& Sch\\\"{o}nberner 1986; D'Antona \\& Mazzitelli 1989; Iben \\& Laughlin 1989; Wood 1992; Fontaine, Brassard \\& Bergeron 2001). This technique has been historically rather difficult to implement in open clusters. First, even for the nearby clusters, their white dwarf members are the faintest objects comprising the mass function, typically being fainter than V=$20$th magnitude (\\eg Rubin et al. 2008), which is challenging for all but the largest aperture terrestrial telescopes.  Second, white dwarfs are intrinsically blue objects (typically B-V$_{o}$~$<0.2$), which until recently was problematic for capturing photons efficiently using astronomical, red sensitive spectrographs and detectors. Third, and most importantly, to properly derive the age of the cluster one must observe and characterize objects extending right down to the bottom of the white dwarf sequence; \\ie those objects which are the dimmest, the coolest and hence the {\\em oldest} white dwarfs in the cluster. Consequently, in light of such procedural difficulties, relatively few open clusters have white dwarf cooling-time age estimates available in the literature (\\eg Richer et al. 1998; Bedin et al. 2005; Kalirai et al. 2003, 2005; Rubin et al. 2008).   In terms of rotation, Kawaler (1989) uses theoretical stellar spin-down models (from Kawaler 1988) to analytically derive a relationship uniquely connecting a star's B-V photometric colour and age to its period. Using a photometric colour, rotation period dataset for G \\& K-stars in the Hyades open cluster, he shows the Hyades rotation age to be $4.9\\pm1.1 \\times 10^{8}$ years. More recently, Barnes (2003, 2007 - hereafter B03, B07 respectively) exploit the morphology of photometric colour-rotation period distributions in co\\\"{e}val samples of stars (\\eg open clusters and binary systems) to establish stellar age. In essence, this so-called {\\em gyrochronology} technique allows one to construct rotational isochrones, in order to trace the boundaries of age dependent colour-period distributions. Crucially however, these boundary definitions further suggest the identification and coupling-state of each star's basic internal stratification in terms of their radiative zone and convection envelope.   Solar-type stars lying close to or on the $I$, or interface sequence, are probably rotating as solid bodies or close to it, and their spin-down evolution is Skumanich-like (Skumanich 1972), i.e., directly proportional to the square root of stellar age (spin-down~$\\propto ~t^{1/2}$). The most rapidly rotating solar-type stars at a given mass lie on or close to the $C$, or convective sequence. These stars are thought to be rotationally {\\em de-coupled}; that is to say, it is likely that {\\em only} their outer convective envelope is spinning down, with an exponential time dependency (spin-down~$\\propto ~e^{f(t)}$). Early incarnations of rotational evolution models incorporating {\\em de-coupled} stellar structure were referred to as core-envelope (de)coupling models (\\cf Endal \\& Sofia 1981; Stauffer et al. 1984; Soderblom et al. 1993a, Jianke \\& Collier Cameron 1993). Those stars which are in a transitory state between the $C$ and $I$-sequences, representing a scenario whereby the outer convective envelope is {\\em re-coupling}, probably magneto-hydrodynamically, to the inner radiative zone, constitute the so-called {\\em gap} or $g-$stars in the gyrochronology paradigm.   Establishing stellar age for open clusters using gyrochronology models of course has its limitations. For instance for open clusters, enough stars must be photometrically monitored in order to derive rotation periods, so that a clear, well-defined distribution of interface, gap, or convective sequence objects is apparent. Such an observing programme typically requires extensive allocations of telescope time and considerable human effort. Moreover, differential rotation of solar-type stars and multiple star-spot groups on their surfaces can also introduce ambiguities into gyrochronology analyzes, since it can act to smear-out the distribution of rotation periods of a given mass.  For small samples of rotation periods in open clusters, the gyrochronology method is not applicable in multiple systems where other effects might interfere with, or even overwhelm, the regular wind-related loss of angular momentum. For instance, close binaries ($a$ \\lsi 0.1 AU) experience a tidal torque which acts to synchronize their rotation with the orbital motion of the system, in addition to any magnetic torque they experience due to stellar winds (Zahn 1989; 1994); thus, as they evolve along the main-sequence, components in the closest binaries tend to rotate faster (on average) than single stars of similar mass and age. Observations of RS CVn and BY Dra systems have shown Zahn's theoretical framework to be correct (Hall 1976; Bopp \\& Fekel 1977; Fekel, Moffet \\& Henry 1986). Binary/multiple systems must therefore be avoided, where the data allow, in establishing gyrochronology ages of stars (see Meibom, Mathieu \\& Stassun 2006 for a more thorough discussion).  A recent International Astronomical Union symposium (IAUS 258), dedicated to the subject of {\\em The Age of Stars}, now provides our community with a thorough overview of an historical and modern approach to understanding the how and why of stellar ages. A careful perusal of the symposium proceedings allows, for both the novice and expert alike, a detailed insight into the theoretical framework unpinning stellar age determinations as well as the observational data calibrating and constraining such models. Of special interest to this manuscript were the oral presentations by Barnes (2009), Jeffries (2009) and Meibom (2009).   The genesis of this gyrochronology project for M\\,34, a $\\simeq$ 200\\,Myr open cluster, occurred nearly ten years ago. At that epoch, SAB and GWL obtained a differential photometry dataset at Lowell Observatory over 17-nights for the Western half of the cluster, with the goal of deriving rotation periods for photometric variables in the field. The project lay dormant for some years until we noticed that Irwin et al. (2006) published new rotation period data for a sample of solar-type stars in the cluster. Unfortunately in terms of rotation, there are several fundamental problems with the Irwin et al. dataset which hamper our ability to pursue a rigorous gyrochronology analysis using their data (see \\S~\\ref{IRWINcomp} for details). A natural progression and indeed amelioration of their period distribution analysis is to derive the gyrochronology properties of this cluster, now incorporating rotation periods derived from the differential photometry that we had obtained during the 1998 Lowell campaign. In this manuscript therefore, we now report an analysis of our differential photometry dataset for the cluster.  New BVI photometry is reported in \\S~\\ref{indiana} for each photometric variable discovered in our Lowell observing campaign. Rotation periods derived from light-curve analysis of the photometric variables in our field of view are described in \\S~\\ref{newProt}, which includes a comparison with the Irwin et al. study. Using a clean sample of rotation periods for photometric and/or kinematic members of the M\\,34 cluster, we describe in \\S~\\ref{period-distribution} how colour-period data are used in undertaking a gyrochronology analysis, which results in the derivation of the rotation age of the cluster. Also included is a commentary on how error in the gyrochronology age can be ascribed to disparate physical properties of the colour-period dataset, such as cluster non-membership, binarity and differential rotation. ",
        "conclusions": "\\label{discussion}  We have presented the results of a 17-night differential photometry campaign over the Western half of the central region of the open cluster M\\,34. For all photometrically variable stars, we construct differential photometry light-curves, from which we derive periodicity. It is assumed that the photometric variability of objects consistent with cluster membership is due to the presence of magnetic-field induced starspots on their surfaces, rotating with the star; the derived periods are therefore representative of the angular velocity of the star.   In order to assess cluster membership, we exploit an extensive standardized CCD survey of the cluster, which shows that the majority of photometric variable stars [47/55] lie along the loci of the cluster main sequence in V/B-V and V/V-I space. Moreover, we are able to confirm kinematic membership of the cluster for 21 stars from an on-going radial velocity survey of the cluster (Meibom et al. in prep), 5 of which show radial velocity variations (multiple systems). Four (4) of the photometric variables are kinematic non-members. Stars which are either photometric or kinematic non-members were excluded from the gyrochronology analysis.   We note the existence of another photometric period dataset for solar-type stars in the M\\,34 cluster (Irwin et al. 2006). We present an analysis comparing periods for stars in common between our Lowell campaign and their study. We detail several concerns that we have with the integrity of their dataset, notably their short observing window ($<10$ nights) and their half-night observing cadence. After having verified that several Irwin et al. periods are most likely measurement aliases, and that our Lowell campaign periods agree those derived from an independent, separate, long time base-line differential photometry survey, we preferentially employ our period results in pursuing a gyrochronology age analysis for this cluster.   Rotation periods for {\\em bona fide} cluster members, in concert with photometric B-V colours, are used to create colour-period distributions, from which we outline a specific methodology for performing a gyrochronology analysis. In order to calculate the actual gyrochronology age of the cluster, we assign each M\\,34 member its CgI status, as judged from its position in the colour-period diagram. I-sequence stars in M\\,34 are classified specifically as those objects in the colour-period diagram lying above the I-sequence locus of the younger M\\,35 cluster, yielding a mean gyro age for all I-sequence stars in M\\,34 of $193 \\pm  9$ Myr.   There are two existing sets of age determination for the M\\,34 cluster, using the traditional isochrone fitting method and the white dwarf cooling time-scale. Together with gyrochronology, all three methods suffer from various forms of dependency on stellar models under-pinning their theoretical frameworks. Furthermore, both the isochrone fitting and white dwarf methods suffer from a dependence on (or an assumption of) cluster distance, while gyrochronology is free from distance as an input parameter.  The isochronal age of M\\,34 is not that well defined (see \\S~\\ref{extantM34} \\& Table~\\ref{table-age}), with a broad range in determined values. Unfortunately, the lack of observed giants in the cluster, and its broad {\\em turn-off} locus in the V/B-V colour magnitude diagram precludes a precise assessment of its isochronal age. Based on existing studies, the mean isochronal age of the cluster is $250 \\pm 67$ Myr.  The white dwarf cooling time-scale age for M\\,34 is also problematic because its result is based upon only five (5) objects whose membership of the cluster is not yet confirmed. Furthermore, the terminus of the white dwarf sequence has not yet been observed. In simple terms, the faintest, coolest, and therefore oldest white dwarfs have yet to be characterized, and the cluster's cooling time-scale age of $64.0 \\pm 12.9$ Myr (Rubin et al. 2008) must be considered as a lower limit.   We conclude by noting that a study of M\\,34 by Meibom et al., including both a multi-month differential photometry campaign and also a multi-year radial velocity membership and binarity campaign, similar to their prior work in M\\,35, is in preparation. Their survey will yield several hundred new photometric periods of M\\,34 solar-type stars, the I- and C-sequences of which should be far more clearly defined than they are from our Lowell campaign results. Their dataset should allow them to better define the gyrochronology age of M\\,34, and we await their results eagerly."
    },
    "1004/1004.3933_arXiv.txt": {
        "abstract": "{Luminous infrared galaxies (LIRGs) are an important class of objects in the low$-z$ universe bridging the gap between normal spirals and the strongly interacting and starbursting ultraluminous infrared galaxies (ULIRGs). Since a large fraction of the stars in the Universe have been formed in these objects, LIRGs are also relevant in a high-$z$ context. Studies of the two-dimensional physical properties of LIRGs are still lacking.} {We aim to understand the nature and origin of the ionization mechanisms operating in the extra-nuclear regions of LIRGs as a function of the interaction phase and infrared luminosity.} {This study uses optical integral field spectroscopy (IFS) data obtained with VIMOS. Our analysis is based on over 25\\,300 spectra of 32 LIRGs covering all types of morphologies (isolated galaxies, interacting pairs, and  advanced mergers), and the  entire $10^{11} - 10^{12} L_{\\sun}$ infrared luminosity range. } {We found strong evidence for shock ionization, with a clear trend  with  the dynamical status of the system.  Specifically, we quantified the variation with interaction phase of several line ratios indicative of the excitation degree. While the \\nha{} ratio does not show any significant change, the \\sha{} and \\oha{} ratios are higher for more advanced interaction stages. Velocity dispersions are higher than in normal spirals and increase with the interaction class (medians of 37, 46, and 51~km~s$^{-1}$ for class 0, 1, and 2, respectively). We constrained the main mechanisms causing the ionization in the extra-nuclear regions (typically for distances ranging from $\\sim$0.2-2.1~kpc  to $\\sim$0.9-13.2~kpc) using diagnostic diagrams. Isolated systems are mainly consistent with ionization caused by young stars. Large fractions of the extra-nuclear regions in interacting pairs and more advanced mergers are consistent with ionization caused by shocks of $v_s \\lsim 200$~km~s$^{-1}$. This is supported by the relation between the excitation degree and the velocity dispersion of the ionized gas, which we interpret as evidence for shock ionization in interacting galaxies and advanced mergers but not in isolated galaxies. This relation does not show any dependence with the infrared luminosity (i.e. the level of star formation). All this indicates that tidal forces play a key role in the origin of the ionizing  shocks in the extra-nuclear regions. We also showed for the first time what appears to be a common  $\\log$(\\oha) - $\\log (\\sigma)$ relation for the extranuclear ionized gas in interacting (U)LIRGs  (i.e. covering the entire $10^{11.0}-10^{12.3}$~L$_{\\odot}$ luminosity range). This preliminary result needs to be investigated further with a larger sample of ULIRGs. } {} ",
        "introduction": "Luminous and ultraluminous infrared galaxies (LIRGs and ULIRGs) are defined as those objects with an infrared luminosity of $L_{IR} = L (8-1000 \\mu \\mathrm{m}) =  10^{11} - 10^{12} L_{\\sun}$ and $L_{IR} \\ga 10^{12} L_{\\sun}$, respectively \\citep[see][for a review]{san96,lon06}. They are systems which contain large amounts of gas and dust \\citep[e.g.][]{eva02} and which are undergoing an intense star-formation episode in their (circum)nuclear regions \\citep[e.g.][]{sco00,alo06}. This activity is the main cause of their huge luminosity in about $\\sim$80\\% of these systems, although some contribution from an AGN is present and even dominant in some cases \\citep[e.g.][]{gen98,ris06,far07,nar08}.  These systems usually present some degree of interaction whose importance increases with luminosity. While the majority  of local LIRGs can be classified as isolated spirals or interacting pairs \\citep[e.g.][]{arr04,alo06,san04}, most of the ULIRGs show signs of a clear merging process \\citep[e.g.][]{cle96,bor00,cui01,bus02,vei02}. While (U)LIRGs are an oddity in the local Universe, recent mid-infrared and submillimeter surveys show how they present a strong evolution with redshift, increasing their number by two orders of magnitude at $z \\sim 0.8 - 1.2$ \\citep{elb02}. Indeed they are the dominant population of the infrared selected galaxies at high redshift, making a significant contribution to the star-formation rate density at $0.5 < z < 2$ \\citep{per05,lef05}.  The study of the ionization properties of the gas in these objects is relevant for two main reasons. On the one hand, the ionization is important to investigate the nature (i.e. starburst, AGN) of the dominant source that causes the huge luminosity in the infrared. In the optical, this has mainly been done via long-slit observations of the nuclear regions of large samples of (U)LIRGs \\citep[e.g.][and references therein]{kim95,vei99}. These studies established trends with the luminosity and interaction stage, and found an increase in the frequency of AGN-dominated systems with luminosity. These results have been recently revisited using the new optical classifications provided by the use of \\emph{Sloan Digital Sky Survey} (SDSS) data \\citep{yua10}. They show that most of the (U)LIRGs previously classified as \\emph{Low-ionization nuclear emission-line region}  (LINER), now are classified as starburst-AGN composite galaxies. The presence of an obscured AGN has been also revealed by the detection of ionization cones with integral field spectroscopy (IFS) data \\citep[e.g. Arp~299,][]{gar06}.  On the other hand, the ionization structure helps to understand how the interaction/merger process as well as the release of energy and material from the central source and/or starbursts are affecting the extended structure of the galaxies in general and its interstellar medium in particular. In that sense, the presence of Super Galactic Winds (SGWs) in (U)LIRGs has been suggested using emission \\citep{hec90,leh96} and absorption \\citep{hec00,rup02,rup05b,rup05a} lines.  Tidally induced forces associated with the interaction process itself have been also suggested as the cause for the ionization of the gas \\citep{mcd03,col05}. Given the complex structure of these systems, where the selection of a preferential direction is specially difficult, these studies would benefit from IFS data thanks to which it is possible to obtain homogeneous two-dimensional spectral information.  Using this technique \\citet[][hereafter MAC06]{mon06} have studied a sample of  six ULIRGs (nine galaxies),  and  found that wide areas of the extra-nuclear extended regions presented line ratios typical of LINERs according to the diagnostic diagrams of \\citet{vei87}. In addition, it was shown that the velocity dispersion is positively correlated with the degree of ionization supporting the idea that shocks are the main cause of the ionization in these areas. However, these results were based on a relatively small sample, which covered a restricted range in luminosity  ($\\log(L_{IR}/L_\\odot) = 12.03 - 12.40$) and interaction phase.  In this paper we extend that study to a larger sample of  32 systems, which cover the entire $\\log(L_{IR}/L_\\odot) = 11.00 - 12.00 $ luminosity range (i.e. the LIRGs range), and the different interaction types (i.e., isolated galaxies, interaction pairs, and mergers remnants).  The present study is part of a wider project devoted to the study of the internal structure and kinematics of a representative sample of low-redshift LIRGs and ULIRGs using optical and near-IR IFS facilities \\citep{arr08}. Specifically we used the INTEGRAL+WYFFOS facility \\citep{arr98,bin94} and the \\emph{Potsdam Multi-Aperture Spectrograph}, PMAS \\citep{rot05} in the Northern Hemisphere, and VIMOS \\citep{lef03} and SINFONI \\citep{eis03}  in the southern one. The corresponding catalogs for the PMAS, INTEGRAL and VIMOS samples can be found in \\citet{alo09}, \\citet{gar09} and Rodr\\'{\\i}guez-Zaur\\'{\\i}n et al.  (in prep.), respectively.  The paper is structured as follows: in Sect. \\ref{secdos} we describe the sample used in this work as well as the characteristics of the instrumental configuration and technical details regarding data reduction and analysis; Sect. \\ref{sectres} quantifies how the ionization degree varies with interaction stage and constrains the possible  mechanisms that cause the ionization of the gas. Finally, a comparison with the previous results for ULIRGs  and a discussion about  the origin of the ionization produced by shocks in terms of the star formation and the interaction process are presented.  Throughout the paper, a cosmology with  70~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_M =  0.3$ and $\\Omega_\\Lambda = 0.7$  is assumed. ",
        "conclusions": "The two-dimensional ionization structure of the extended (few to several kpc) ionized gas in a representative sample of 32 low-$z$ LIRGs (i.e. $\\log(L_{ir}/L_\\odot)=11.00-12.00$ luminosity range) was investigated with the VIMOS integral field spectrograph. The sample covers isolated galaxies, as well as interacting galaxies and systems in an advanced stage of the merger. This paper investigates the nature and origin of the main ionization mechanisms operating in the extra-nuclear regions of these systems based on several thousands of independent measurements of the emission line ratios (up to twenty 24\\,000 for \\nha) and velocity dispersions.  The present study is part of a larger project devoted to the study of the two-dimensional structure for the stars and ionized gas as well as its kinematics and ionization conditions in representative samples of LIRGs and ULIRGs using optical IFS. The main results of this study can be summarized as follows:  1. The distribution of the \\nha\\  line ratio does not show any significant variations with the interaction class, with most regions presenting a ratio typical of \\ion{H}{ii} regions. The \\sha\\ and \\oha\\ line ratios do however show a change in their distribution with an extension towards higher excitation (i.e. LINER-like excitation) for galaxies classified as interacting pairs and advanced mergers. This change is more pronounced for the \\oha\\ ratio.  2. There is an anti-correlation between the ionization degree and the \\ha\\ surface brightness independently of the interaction type and similar to what occurs in our Galaxy or in the so-called DIG in other spiral galaxies. Most of the observed line ratios are similar to those found in \\ion{H}{ii} regions in our Galaxy, but there is a relatively large percentage of line ratios similar to those for DIGs.  3. The nature of the ionization sources was investigated comparing the measured \\sha\\ vs. \\nha{} and \\oha{} vs. \\nha{} line ratios with the predictions of ionization due to stars (\\ion{H}{ii} regions), TML, shocks (DIGs) and power-law (AGN) spectra. Turbulent Mixing Layers do not seem to play a major role in the ionization of the extra-nuclear regions. Line ratios in LIRGs classified as isolated can mostly be explained as caused by ionization due to young stars. On the other hand, the ionization in a large fraction of the regions in systems with some degree  of interaction cannot be due to stars but is better explained by high velocity shocks. This is particularly evident when using the best shock tracer, i.e. the \\oha\\ vs. \\nha\\ diagram. Independently of the ionization mechanisms, only models with metallicity between solar and twice solar are able to explain the observed line ratios.  4. Local velocity dispersions  increase with the interaction degree, with medians of 37, 46, and 51~km~$^{-1}$ for class 0, 1, and 2 respectively, and are higher than those typically found in normal spirals \\citep[$\\sim20-30$~km~s$^{-1}$,][]{epi10}. This indicates that the ionized ISM in LIRGs is dynamically hotter than the quiescent ISM of normal galaxies due to strong shocks produced by tidal forces and stellar winds associated to the nuclear starburst.  5. There is a positive relation between the degree of excitation (as traced by the emission line ratios) and the velocity dispersion of the ionized gas in LIRGs classified as interacting systems, and mergers, while this relation is not observed in isolated systems. This relation is better seen when using the \\oha\\ and \\sha\\ line ratios, and supports the scenario where the relevance of shocks as ionizing sources in the extranuclear extended regions of LIRGs increases when there is some degree of interaction.  6. The relation between the degree of excitation and the velocity dispersion of the ionized gas in interacting and merging LIRGs does not clearly improve with the infrared luminosity (i.e. star formation rate) of the systems. Thus the interaction process itself rather than superwinds caused by the star formation seems to be the main origin of the shocks in the extended extra-nuclear regions, assuming that the release of energy into the ISM is proportional to the SFR. This result is still compatible with stellar superwinds in the internal regions of these systems, and/or along certain preferential directions associated with AGN-related outflows.  7. A comparison between the sub-sample of interacting/merging LIRGs and a small sample of ULIRGs suggests the existence of a common positive $\\log$(\\oha) - $\\log(\\sigma)$ relation. If confirmed, these results will provide further evidence for the tidal origin of shocks in these galaxies over the entire LIRG and ULIRG luminosity range. A study with a larger sample of ULIRGs is under way to confirm the result."
    },
    "1004/1004.3102_arXiv.txt": {
        "abstract": "Results   from  the   first  fully   general   relativistic  numerical simulations  in  axisymmetry  of  a  system formed  by  a  black  hole surrounded by  a self-gravitating torus in  equilibrium are presented, aiming to assess the influence  of the torus self-gravity on the onset of the  runaway instability. We  consider several models  with varying torus-to-black  hole  mass  ratio  and angular  momentum  distribution orbiting in equilibrium around a non-rotating black hole. The tori are perturbed  to induce  the mass  transfer towards  the black  hole. Our numerical simulations show that  all models exhibit a persistent phase of  axisymmetric  oscillations  around  their equilibria  for  several dynamical   timescales   without  the   appearance   of  the   runaway instability, indicating  that the self-gravity  of the torus  does not play a critical role favoring the onset of the instability, at least during the first few dynamical timescales. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.1331_arXiv.txt": {
        "abstract": "We present deep Keck spectroscopy for 17 morphologically-selected field spheroidals in the redshift range $1.05<z<1.60$ in order to investigate the continuity in physical properties between the claimed massive compact red galaxies (``nuggets'') at $z\\simeq2$  and well-established data for massive spheroidal galaxies below $z\\simeq1$.  By combining Keck-based stellar velocity dispersions with \\hst-based sizes, we find that the most massive systems (M$_{\\rm dyn}>10^{11}$ M$_\\odot$) grew in size over $0<z<1.6$ as (1+z)$^{-0.75\\pm0.10}$ (i.e., $\\times 2$ since $z = 1.5$) whereas intermediate mass systems ($10^{11} {\\rm M}_\\odot > {\\rm M}_{\\rm dyn} > 10^{10} {\\rm M}_\\odot$) did not grow significantly. These trends are consistent with a picture in which more massive spheroidals formed at higher redshift via ``wetter'' mergers involving greater dissipation. To examine growth under the favored ``dry'' merger hypothesis, we also examine size growth at a fixed velocity dispersion. This test, uniquely possible with our dynamical data, allows us to consider the effects of ``progenitor bias.'' Above our completeness limit ($\\sigma > 200$ km~s$^{-1}$), we find size growth consistent with that inferred for the mass-selected sample, thus ruling out strong progenitor bias. To maintain continuity in the growth of massive galaxies over the past 10 Gyr, our new results imply that size evolution over $1.3<z<2.3$, a period of 1.9~Gyr, must have been even more dramatic than hitherto claimed if the red sources at $z>2$ are truly massive and compact. ",
        "introduction": "The observation that many red galaxies with large stellar masses at $z\\simeq2$ are 3--5 times more compact than equivalent ellipticals in the local Universe \\citep[e.g.,][]{Daddi2005,Trujillo2007,vanDokkum2008,Buitrago2008,Damjanov2009} has been a source of much puzzlement. How can an early galaxy grow primarily in physical size without accreting significant stellar mass as required if these objects are the precursors of the most massive ellipticals observed today? Furthermore, studies of the fundamental plane and other stellar population indicators do not permit substantial recent star formation since $z\\sim2$ in massive galaxies, thus precluding growth by accretion of young stars or via gas-rich (``wet\") mergers \\citep[e.g.,][hereafter T05]{Treu2005}. Some have questioned the reliability of the observations, suggesting an underestimate of physical sizes or an overestimate of stellar masses (\\citealt{Hopkins2009}; however, see \\citealt{Cassata2009} for a contrasting view). Others have proposed size expansion driven by self-similar dissipationless ``dry'' mergers, or mass accretion from minor mergers \\citep[][and references therein]{Khochfar2006,Naab2009,Hopkins2010}.  To verify the compact nature of distant sources and to track their evolution in size and mass, it is preferable to use dynamical masses \\mdyn\\, from absorption line spectra, which do not suffer from uncertainties associated with the assumed initial mass function and stellar mass estimates derived from broad-band photometry \\citep[e.g.,][]{Muzzin2009}. \\mdyn\\ measurements are available for relatively large samples out to $z\\sim1$ (T05; \\citealt{vanderWel2008}, hereafter vdW08), suggesting a small but detectable difference in average size at fixed mass when compared to the local universe. But beyond $z\\simeq1$, there is little high-quality dynamical data for field spheroidals. \\citet{vanDokkum2009} undertook an heroic observation of a single $z>2$ source with a stellar mass $\\simeq 2 \\times 10^{11}M_\\odot$ and an effective radius $r_e=0.8$ kpc typical of compact galaxies at $z\\simeq2.3$. The spectrum has a claimed stellar velocity dispersion of $\\sigma=510^{+165}_{-95}$ km~s$^{-1}$, suggesting a remarkably dense system. van Dokkum et al.~postulate the initial dissipative collapse at $z\\simeq3$ of a high mass ``core'' but are unable to account for its subsequent evolution onto the $z\\simeq1$ scaling relations. The quantitative effect of minor mergers on the physical size of a galaxy involves many variables, and it is unclear whether such dramatic size evolution is possible while maintaining the tightness of the fundamental plane and its projections \\citep{Nipoti2009}.  Interpretation of the observed trends at fixed \\mdyn\\, is further complicated by the so-called ``progenitor bias'' \\citep{vanderWel2009}: if galaxies grow by dry mergers, the main progenitor of a present-day massive galaxy did not have the same mass at $z\\sim2$. Similarly, if galaxies become recognizable as spheroidals only above a certain threshold in stellar velocity dispersion $\\sigma_{\\rm ET}$ that depends on redshift, it is clear that the addition of a new -- and less dense -- population could mimic a false evolutionary trend. This bias can be reduced by considering galaxy sizes at fixed $\\sigma$.  Foremost, $\\sigma$ changes very little under a variety of growth mechanisms \\citep[e.g.,][]{Hopkins2010} and it is therefore a better ``label'' than \\mdyn~to track the assembly history. Secondly, $\\sigma$ is closely correlated with stellar age \\citep{vanderWel2009} and therefore offers the most direct way to track the evolving population.  Given there is no clear consensus in understanding the continuity between the galaxy population at $z<1$ and that at $z>2$, we have embarked on a campaign to measure $\\sigma$ and \\mdyn for a large sample of field spheroidals at $1<z<1.7$. This has recently become practical using multi-object optical spectrographs equipped with deep depletion red-sensitive CCDs. Our goal is to extend the earlier work at $z<1$ (T05, vdW08) to within $\\simeq$1 Gyr of the sample of ultracompact galaxies at $z\\simeq2.3$. In this first analysis, we present new results spanning the redshift range $1.05<z<1.60$.  We adopt a $\\Lambda$CDM cosmology with $(\\Omega_m, \\Omega_v, h) = (0.3, 0.7, 0.7)$; all magnitudes are in the AB system. A Chabrier IMF is assumed where necessary. ",
        "conclusions": "Our Keck spectra have shown the utility of securing individual spectroscopic and photometric measures for a representative sample of $z>1$ massive spheroidals. By probing to $z\\simeq1.6$, we are sampling velocity dispersions, sizes and dynamical masses within 1.2 Gyr of the puzzling population of compact red galaxies at $z \\simeq 2.3$.  Importantly, the size evolution we infer over $0<z<1.6$ at fixed dynamical mass is modest: $\\simeq\\times2$ for the most massive ($\\log M_{\\rm dyn}/M_\\odot>11$) examples but much smaller for lower mass systems. If the compact red galaxies at $z\\simeq2-2.3$ are their precursors, they must have grown dramatically in size over a very short time interval."
    },
    "1004/1004.3965_arXiv.txt": {
        "abstract": "We report unusual near- and mid-infrared photometric properties of \\object{G\\,196--3\\,B}, the young substellar companion at 16\\arcsec~from the active M2.5-type star \\object{G\\,196--3\\,A}, using data taken with the IRAC and MIPS instruments onboard {\\sl Spitzer}. \\object{G\\,196--3\\,B} shows markedly redder colors at all wavelengths from 1.6 up to 24 $\\micron$ than expected for its spectral type, which is determined at L3 from optical and near-infrared spectra. We discuss various physical scenarios to account for its reddish nature, and conclude that a low-gravity atmosphere with enshrouded upper atmospheric layers and/or a warm dusty disk/envelope provides the most likely explanations, the two of them consistent with an age in the interval 20--300 Myr. We also present new and accurate separate proper motion measurements for \\object{G\\,196--3\\,A} and B confirming that both objects are gravitationally linked and share the same motion within a few mas\\,yr$^{-1}$. After integration of the combined spectrophotometric spectral energy distributions, we obtain that the difference in the bolometric magnitudes of \\object{G\\,196--3\\,A} and B is 6.15\\,$\\pm$\\,0.10 mag. Kinematic consideration of the Galactic space motions of the system for distances in the interval 15--30\\,pc suggests that the pair is a likely member of the Local Association, and that it lay near the past positions of young star clusters like $\\alpha$ Persei less than 85~Myr ago, where the binary might have originated. At these young ages, the mass of \\object{G\\,196--3\\,B} would be in the range 12--25\\,M$_{\\rm Jup}$, close to the frontier between planets and brown dwarfs. ",
        "introduction": "The characterization of brown dwarfs and planetary-mass objects of known distance, metallicity, and age can provide critical tests for evolutionary models as well as empirical references for understanding the substellar population in the field, including planets orbiting stars. Confirmed members of nearby open star clusters and substellar companions to stars and massive brown dwarfs can become good targets to establish the properties of benchmark objects. Among them, those located at the nearest distances are preferred since brown dwarfs and planets are intrinsically faint and they evolve towards lower luminosities with age (e.g., D'Antona \\& Mazzitelli \\cite{dantona94}; Burrows et al$.$ \\cite{burrows97}; Baraffe et al$.$ \\cite{baraffe98}). Tens of brown dwarfs are known to orbit stars and more massive brown dwarfs, including the discovery of a 5 Jupiter-mass planet around a brown dwarf in the young TW Hya association (Chauvin et al$.$ \\cite{chauvin05}). They show L and T spectral types characterized by surface temperatures below $\\sim$2200 K.  One ultra-cool substellar companion to a nearby low-mass star is \\object{G\\,196--3\\,B}, found by direct imaging and proper motion studies at a separation of $\\sim$16\\arcsec~from the active M2.5 star \\object{G\\,196--3\\,A} (Rebolo et al$.$ \\cite{rebolo98}). In the discovery paper, the authors discussed that this system is young with an age in the interval 20--300 Myr. The young limit is imposed by the lack of the lithium absorption feature at 670.8 nm (Christian \\& Mathioudakis \\cite{christian02}) in the optical spectrum of the primary star, implying that \\object{G\\,196--3\\,A} has efficiently depleted this element by nuclear reactions. This happens at ages $\\ge$20 Myr for the mass and temperature of this particular star. The intense emission lines like H$\\alpha$ and other activity properties, particularly the X-ray and UV emission of the M2.5 star, are quite similar to $\\alpha$~Persei and Pleiades stars of the same temperature, suggesting that \\object{G\\,196--3} could have a likely age of about 100 Myr. Gizis et al$.$ \\cite{gizis02} imposed a conservative upper limit of 640 Myr to the age of \\object{G\\,196--3\\,A} based on the intensity of the most relevant chromospheric emission lines. More recently, Shkolnik et al$.$ \\cite{shkolnik09} have reviewed the age of \\object{G\\,196--3\\,A} and have set it at 25--300\\,Myr.  Since its discovery and due to its relative brightness ($J$ = 14.8 mag), \\object{G\\,196--3\\,B} has been observed spectroscopically at optical and near-infrared wavelengths by several groups. Therefore, its spectroscopic properties are well defined: \\object{G\\,196--3\\,B} is a moderate rotator ($v$\\,sin\\,$i$ = 10 km\\,s$^{-1}$), with no H$\\alpha$ emission (Mohanty \\& Basri \\cite{mohanty03}). Its rotation rate is quite similar to the rotational velocity of the primary star ($v$\\,sin\\,$i$ = 15 km\\,s$^{-1}$, Gizis et al$.$ \\cite{gizis02}). The lithium feature is seen with a relatively strong intensity (a few to several Angstroms, Rebolo et al$.$ \\cite{rebolo98}; Mart\\'\\i n et al$.$ \\cite{martin99}; Kirkpatrick et al$.$ \\cite{kirk01}), strongly supporting its substellar nature. In addition, \\object{G\\,196--3\\,B} displays all spectroscopic hallmarks of a low gravity atmosphere (see next Sections), thus confirming a young age. More recently, Cruz et al$.$ \\cite{cruz09} has compared the optical spectrum of \\object{G\\,196--3\\,B} to other field dwarfs of similar temperature and spectroscopic properties, and have assigned a spectral type of L3, slightly cooler than the L1 given in Mart\\'\\i n et al$.$ \\cite{martin99} and the L2 measured by Kirkpatrick et al$.$ \\cite{kirk01}. As estimated in Rebolo et al$.$ \\cite{rebolo98}, the mass of \\object{G\\,196--3\\,B} is 25$^{+15}_{-10}$ times that of Jupiter.  In this paper we focus on the description and interpretation of the photometric properties of \\object{G\\,196--3\\,B} from the visible to the mid-infrared wavelengths (0.6--24 $\\mu$m). In addition, we present new, accurate measurements of the proper motion of each member of the pair used to determine the Galactic kinematics of the system. ",
        "conclusions": "We presented near-infrared, low-resolution spectrum (obtained with TNG/NICS instrument and Amici prism) and mid-infrared photometry ({\\sl Spitzer}/IRAC and MIPS) of the brown dwarf \\object{G\\,196--3\\,B}, which is a wide companion of the young, active M2.5-type star \\object{G\\,196--3\\,A} (Rebolo et al$.$ \\cite{rebolo98}). Using the near-infrared spectrum (0.8--2.4 $\\mu$m), we confirmed an L3 spectral classification for the brown dwarf in agreement with the L3 typing derived from optical data by Cruz et al$.$ \\cite{cruz09}; we also confirmed previously reported detection of spectroscopic features (``triangular'' shape of the $H$-band and strong VO$+$H$_2$O absorption) consistent with a low-gravity atmosphere and a young age. New optical and near-infrared images obtained on different occasions between 1998 and 2010, in addition to data available from public archives, allowed us to determine an accurate proper motion for each member of the binary and a projected separation and position angle of 15\\farcs99\\,$\\pm$\\,0\\farcs06 and 209.3\\,$\\pm$\\,0.3 deg, which have not changed in the last 12 yr (i.e., orbital motion below our astrometric detectability). This provides strong support for the gravitational link of the pair. We found that the kinematics of \\object{G\\,196--3} (adopting a probable distance interval of 15--30 pc) is consistent with the binary being a likely member of the Local Association, thus supporting the young age (20--300\\,Myr) assumed for the system. A simple exercise of linearly back tracing the Galactic orbit of \\object{G\\,196--3} located the system relatively close to the past positions of the open clusters Collinder~65 (20--30 Myr) and $\\alpha$~Persei (50--85 Myr), where \\object{G\\,196--3} may have originated. This implies that the age of the system could likely be 20--85 Myr (i.e., younger than the Pleiades star cluster). At this young age, \\object{G\\,196--3\\,B} would have a mass of 0.012--0.25~$M_{\\odot}$ and would be burning deuterium.  From the {\\sl Spitzer}/IRAC and MIPS photometry, we concluded that \\object{G\\,196--3\\,A} spectral energy distribution (SED) covering 0.44--24 $\\mu$m is purely photospheric, and that \\object{G\\,196--3\\,B} shows infrared colors significantly redder than expected for its L3 spectral classification and much more similar to those of L7--L8 sources. The integration of the observed SEDs yielded a bolometric magnitude difference $\\Delta m_{\\rm bol}$\\,=\\,6.15\\,$\\pm$\\,0.10 mag between the two objects. The SED of \\object{G\\,196--3\\,B} appears overluminous longwards of 1.6 $\\mu$m when compared to the average SED of field L2--L3 dwarfs. We discussed that a low-gravity atmosphere with enshrouded upper atmospheric layers and/or a warm dusty disk/envelope provides the most likely explanations, the two of them consistent with an age in the interval 20--85 Myr. \\object{G\\,196--3\\,B} can be used as a benchmark object to understand L-type sources with similar photometric and spectroscopic properties in the field."
    },
    "1004/1004.3362_arXiv.txt": {
        "abstract": "{The identification of long-gamma-ray-bursts (LGRBs) is still uncertain, although the collapsar engine of fast-rotating massive stars is gaining a strong consensus.} {We propose that low-metallicity Be and Oe stars, which are massive fast rotators, as potential LGRBs progenitors.} {We checked this hypothesis by 1) testing the global specific angular momentum of Oe/Be stars in the ZAMS with the SMC metallicity, 2) comparing the ZAMS ($\\Omega/\\Omega_{\\rm c},M/M_{\\odot}$) parameters of these stars with the area predicted theoretically for progenitors with metallicity $Z\\!=\\!0.002$, and 3) calculating the expected rate of LGRBs/year/galaxy and comparing them with the observed ones. To this end, we determined the ZAMS linear and angular rotational velocities for SMC Be and Oe stars using the observed \\vsini~ parameters, corrected from the underestimation induced by the gravitational darkening effect.} {The angular velocities of SMC Oe/Be stars are on average $\\langle\\Omega/\\Omega_{\\rm c}\\rangle\\!=\\!0.95$ in the ZAMS. These velocities are in the area theoretically predicted for the LGRBs progenitors. We estimated the yearly rate per galaxy of LGRBs and the number of LGRBs produced in the local Universe up to z=0.2. We have considered that the mass range of LGRB progenitors corresponds to stars hotter than spectral types B0-B1 and used individual beaming angles from 5 to 15\\degr. We thus obtain $R^{\\rm pred}_{\\rm LGRB}\\sim10^{-7}$ to $\\sim10^{-6}$ LGRBs/year/galaxy, which represents on average 2 to 14 LGRB predicted events in the local Universe during the past 11 years. The predicted rates could widely surpass the observed ones [(0.2-3)$\\times10^{-7}$ LGRBs/year/galaxy; 8 LGRBs observed in the local Universe during the last 11 years] if the stellar counts were made from the spectral type B1-B2, in accordance with the expected apparent spectral types of the appropriate massive fast rotators.} {We conclude that the massive Be/Oe stars with SMC metallicity could be LGRBs progenitors. Nevertheless, other SMC O/B stars without emission lines, which have high enough specific angular momentum, can enhance the predicted $R_{\\rm LGRB}$ rate.} ",
        "introduction": "Since their discovery, data on gamma ray bursts (GRBs) considerably increased, which today allow us to aim at preliminary statistical conclusions. The properties of GRBs and their possible connection to SN were reviewed by \\citet{woosley2006} and \\citet{fryer2007}. Thus, there are at least 2 classes of GRBs according to the duration of the phenomenon: 1) short bursts that last less than 1 s and 2) long bursts, which are longer than 1-3 s hereafter LGRBs. Although the short GRBs might correspond to a violent merging of two compact objects, evidence is growing that suggests that the disruption of a massive star can be behind an LGRB \\citep{fryer2007}. A possible, but rare, third class of GRBs can be explained by a binary-star scenario \\citep{tutu2007}. However,  proper identification of both the progenitor and the final nature of the phenomenon are still fairly uncertain. Nevertheless, a kind of consensus is gaining in the community that the collapsar engine during a supernova SNIb,c explosion, as proposed by \\citet{woosley1993}, might lead to allow the explanation of LGRBs. According to this hypothesis, the explosion follows a massive star collapse to a black hole. The circumstellar disk accretes onto the black hole and a bi-polar jet is formed \\citep{hirschi2005}. The infalling material must have enough angular momentum to remain in a disc before accretion.\\par  New observational findings and the latest theoretical developments are steadily giving best founded constraints to understand the LGRBs phenomenon. The observations carried out by \\citet{iwamoto1998,iwamoto2000} support the idea that massive fast-rotating stars are at the origin of the LGRBs. \\citet{thone2008} find that the LGRB \\object{GRB060505} is hosted in a low-metallicity galaxy, which is characterized by a high star-formation rate. The event seems to come from a young environment (6 Myrs) and from an object about 32 M$_{\\odot}$. The latest models have tried to reproduce the evolution of stars until the end of their lives. In particular, for massive stars, they have been worked out until the GRBs phase. In the search for SNIb,c progenitors, WR stars with He-rich envelopes are recognized as possibly behind the GRBs events. In fact, \\citet{hammer2006} find that the LGRBs occur in areas of galaxies with WR stars.\\par Rotation has been recognized as a key point for understanding the appearance of GRBs \\citep{woosley1993,hirschi2005,yoon2006}. Accordingly, to keep a large amount of angular momentum up to the last evolutionary phases before the collapse, GRBs progenitors should be massive objects with low initial metallicities and possibly display anisotropic winds \\citep{meynet2007}. Thus, from \\citet{yoon2006} it seems that WR stars with metallicities $Z\\!\\lesssim0.002$ can be progenitors of GRBs. In the same sense, \\citet{modjaz2008} find that the SNIc-GRB association occurs when metallicities are lower than $0.6\\times{Z}_{\\odot}$. Because stars with low metallicities can on average rotate faster, \\citet{hirschi2005} and \\citet{yoon2006} foresee that the number of GRBs must grow as the redshift increases, which seems to be confirmed observationally \\citep{fryer2007}. It might still be that the first massive stars, which were very metal-poor stars, were fast rotators. This is an additional reason for the frequency of GRBs being higher as the redshift increases. In this sense, \\citet{kewley2007} find a link between the location of cosmological LGRBs and the very low-metallicity galaxies.\\par Thanks to fast rotation and the concomitant efficient mixing of chemical elements, massive stars can undergo quasi-chemically homogeneous evolution to end up as helium WR stars satisfying the requirements for the collapsar scenario \\citep{yoon2006,vmarle2008}. \\citet{yoon2006} have calculated the quasi-chemically homogeneous evolution of magnetized massive stars. They produced diagrams of LGRBs progenitors as a function of their ZAMS rotational velocities and masses and of different initial metallicities. From these diagrams, it emerges that there must be an upper limit to the initial metallicity of LGRBs progenitors, which approximatively corresponds to that of the Small Magellanic Clouds (SMC). According to \\citet[][and references therein]{maeder1999}, the SMC average metallicity is $Z\\!\\sim0.002$. The diagram of \\citep{yoon2006} also indicates that, for metallicities $Z\\!\\lesssim0.002$, the WR phenomenon can appear in stars having lower masses than those in the Milky Way (hereafter MW). In this context, \\citet{martins2009} have observed several SMC WR stars, whose evolutionary status and chemical properties can be understood if they are fast rotators.\\par According to the listed observational and theoretically inferred requirements, stars might be potential LGRBs progenitors if they\\par  \\begin{itemize} \\item are massive enough at the end of their evolution to form a black hole; \\item have lost their hydrogen envelope and have a fast-rotating core; \\item are formed in low-metallicity environments, where there must also be high star-formation rates to ensure having enough massive stars. \\end{itemize}  Since the lower the metallicity the faster the rotation and the lower the mass of stars that can undergo the WR phase, according to \\citet{yoon2006}, the most massive B-type stars, as well as O-type stars, could become WR stars if they rotated fast enough. The required rotational velocities must be close to those what are typical of Oe/Be stars. Then, we simply ask whether the more massive stars  displaying the Be phenomenon today in the SMC, can be potential progenitors of LGRBs. This possibility has already been raised by \\citet{woohe2006}. We recall that the ``Be phenomenon'' appears in the main sequence evolutionary phase of fast-rotating O- and B-type stars. They are characterized by emission lines produced in a decretion disk formed by continuous and episodic matter ejections from the central star. These objects are the fastest known rotators in the main sequence, and their rotation can be closer to the critical one when the metallicity is lower \\citep{dds2003,meilland2007,fremat2005,vinicius2006,marta2007}.\\par Our research on whether the SMC Oe/Be stars can be progenitors of LGRBs is based on three tests: 1) determination of the global specific stellar angular momentum (Sect.~\\ref{iozrv}), 2) inference of the ZAMS angular velocity ratios of the studied SMC Oe/Be stars and comparison with the model predicted requirements to be LGRB-progenitor (Sect.~\\ref{lgrbp}), and 3) estimation of the LGRBs rate based on the SMC Oe/Be- and fast-rotator population, and its comparison with the observed ones (Sect.~\\ref{prates}). A general discussion is given in Sect.~\\ref{opprd}.\\par ",
        "conclusions": "While the first models focused primarily on fast-rotating WR stars as probable progenitors LGRBs, the latest developments indicate that fast-rotating, low-metallicity massive B- and O-type stars could also be behind the LGRBs. Such massive B- and O-type stars can actually be the Be and Oe stars of low metallicity found in the SMC. To test this assumption, we used fundamental parameters derived for Oe/Be stars in the SMC observed with VLT-FLAMES instruments. We first compared the average \\vsini~determined for several mass subsamples of SMC Oe/Be stars with the theoretically predicted rotational velocities. We found that today these objects on average have a surface angular velocity ratio \\omc$\\simeq0.99$. Then, we determined the ZAMS rotational velocities again for Be stars in the SMC that we had studied in previous works and determined the ZAMS rotational velocities for a new subsample composed by SMC Oe stars. The average angular velocity rate of all these objects in the ZAMS is $\\langle$\\omc$\\rangle\\!=\\!0.95$. The ZAMS \\omc~rates thus determined were compared with those of fast-rotating massive stars foreseen by model calculations of LGRBs progenitors having an initial metallicity similar to the average one in the SMC. This comparison indicates that massive Oe/Be stars in the SMC could certainly have the properties required to be taken as LGRBs progenitors.  A lower limit of 17\\% of galaxies in the near Universe are of irregular Magellanic-type that, by definition, all have metallicities $Z\\lesssim0.002-0.004$. All can then host LGRBs progenitors. From the discussion it appears that not only can massive Oe/Be stars be LGRBs progenitors, but a larger population of O/B stars without emission, which are less massive and with apparent lower surface rotation velocities, can also partake of this quality. The total number of LGRBs estimated from the Oe/Be star populations based on spotted beaming angles and counts from B0-B1 to O8 stars is on average $R_{\\rm LGRB}^{\\rm pred}\\!\\sim\\!10^{-7}$ to  $\\sim\\!10^{-6}$ LGRBs/yr/galaxy, which represent $N_{\\rm LGRB}^{\\rm pred}\\!\\sim\\!2$ to 14 in 11 years. These predictions  increase beyond the observed ones, if the stellar counting starts at spectral types B1-B2, in accordance with the apparent spectral types of massive fast rotators. Massive fast rotators without emission lines can still enhance these predictions by more than 30\\%. On account of all possible uncertainties that may affect the calculated and the observed rates of LGRBs in the local Universe, we can consider that both overlap correctly. Young star formation regions in the SMC like NGC346 could then have hosted or will host one LGRB in the next Myrs.\\par"
    },
    "1004/1004.3538_arXiv.txt": {
        "abstract": "We present two independent, homogeneous, global analyses of the transit \\lcs, radial velocities and spectroscopy of \\kepivb, \\kepvb, \\kepvib, \\kepviib\\ and \\kepviiib, with numerous differences over the previous methods.  These include: i) improved decorrelated parameter fitting set used, ii) new limb darkening coefficients, iii) time stamps modified to BJD for consistency with RV data, iv) two different methods for compensating for the long integration time of Kepler LC data, v) best-fit secondary eclipse depths and excluded upper limits, vi) fitted mid-transit times, durations, depths and baseline fluxes for individual transits.  We make several determinations not found in the discovery papers: i) We detect a secondary eclipse for \\kepviib\\ of depth $(47 \\pm 14)$\\,ppm and statistical significance 3.5-$\\sigma$. We conclude reflected light is a much more plausible origin than thermal emission and determine a geometric albedo of $A_g = (0.38 \\pm 0.12)$. ii) We find that an eccentric orbit model for the Neptune-mass planet \\kepivb\\ is detected at the 2-$\\sigma$ level with $e = (0.25 \\pm 0.12)$. If confirmed, this would place \\kepivb\\ in a similar category as GJ 436b and HAT-P-11b as an eccentric, Neptune-mass planet. iii) We find weak evidence for a secondary eclipse in Kepler-5b of 2-$\\sigma$ significance and depth $(26 \\pm 17)$\\,ppm. The most plausible explanation is reflected light caused by a planet of geometric albedo $A_g = (0.15 \\pm 0.10)$. iv) A 2.6-$\\sigma$ peak in \\kepvib\\ TTV periodogram is detected and is not easily explained as an aliased frequency. We find that mean-motion resonant perturbers, non-resonant perturbers and a companion extrasolar moon all provide inadequate explanations for this signal and the most likely source is stellar rotation. v) We find different impact parameters relative to the discovery papers in most cases, but generally self-consistent when compared to the two methods employed here. vi) We constrain the presence of mean motion resonant planets for all five planets through an analysis of the mid-transit times. vii) We constrain the presence of extrasolar moons for all five planets. viii) We constrain the presence of Trojans for all five planets. ",
        "introduction": "\\label{sec:introduction}   The \\emph{Kepler Mission} was successfully launched on March 7th 2009 and began science operations on May 12th of the same year. Designed to detect Earth-like transits around Sun-like stars, the required photometric precision is at the level of $20$\\,ppm over 6.5 hours integration on 12$^{\\mathrm{th}}$ magnitude stars and early results indicate this impressive precision is being reached \\citep{bor09}. In 2010, the first five transiting exoplanets (TEPs) to be discovered by the \\emph{Kepler Mission} were announced by the Kepler Science Team \\citep{bor10a}, known as \\kepivb, \\kepvb, \\kepvib, \\kepviib\\ and \\kepviiib\\ (\\citet{bor10b}; \\citet{koc10}; \\citet{dun10}; \\citet{lat10}; \\citet{jen10a}). These expanded the sample of known transiting exoplanets to about 75 at the time of announcement.  The main objective of the \\emph{Kepler Mission} is to discover Earth-like planets, but the instrument naturally offers a vast array of other science opportunities including detection of gas giants, searches for thermal emission \\citep{deming:2005} and/or reflection from exoplanets, detection of orbital phase curves \\citep{knutson:2007} and ellipsoidal variations \\citep{wel10}, asteroseismology \\citep{christ:2010} and transit timing \\citep{agol:2005}, to name a few. Confirmation and follow-up of exoplanet transits is known to be a resource intensive activity and since the detection of new planets is Kepler's primary objective, it is logical for many of these other scientific tasks to be conducted by the astronomical community as a whole.  Independent and detailed investigations of the Kepler photometry provides an ``acid-test'' of the methods employed by the Kepler Science Team. Indeed, the distinct analysis of any scientific measurement has always been a fundamental corner stone of the scientific method. In this paper, we present two independent analyses of the discovery photometry for the first five Kepler planets. We aim to not only test the accuracy of the methods used in the discovery papers, but also test our own methods by performing two separate studies. Both methods will be using the same original data, as published in the discovery papers.  Some additional data-processing tasks are run through the Kepler reduced data, which we were not used in the original analyses presented in the discovery papers, and are discussed in \\S3. In this section, we also discuss the generation of new limb darkening coefficients and methods for compensating for the long integration time of the Kepler long-cadence photometry. We also perform individual transit fits for all available Kepler transits in order to search for transit timing variations (TTV), transit duration variations (TDV) and other possible changes. These will be used to provide a search for perturbing planets and companion exomoons. ",
        "conclusions": "\\label{sec:discussion}  Due to the large number of results presented in this paper, we will here summarize the key findings not available in the original discovery papers:  \\begin{itemize}  \\item[{\\tiny$\\blacksquare$}] Secondary eclipse of \\kepviib\\ is detected to 3.5-$\\sigma$ confidence with depth $(F_{pd}/F_\\star) = 47 \\pm 14$\\,ppm, indicative of a geometric albedo of $A_g = (0.38 \\pm 0.12)$.  \\item[{\\tiny$\\blacksquare$}] A marginally significant orbital eccentricity is detected for the Neptune-mass planet \\kepivb\\ to 2-$\\sigma$ confidence, with an eccentricity of $e = (0.25 \\pm 0.12)$.  \\item[{\\tiny$\\blacksquare$}] A marginally significant (1.8-$\\sigma$) secondary eclipse is detected for \\kepvb\\ with a depth of $(26 \\pm 17)$\\,ppm, for which the most plausible explanation is reflected light due to a geometric albedo of $A_g = (0.15 \\pm 0.10)$.  \\item[{\\tiny$\\blacksquare$}] A 2.6-$\\sigma$ significance peak in the TTV periodogram of \\kepvib\\ is detected, which is not easily explained as an alias frequency. Perturbing planets and exomoons are unlikely to be responsible either and currently our favored hypothesis is one of stellar rotation.  \\item[{\\tiny$\\blacksquare$}] We derive significantly different impact parameters for all of the Kepler planets except \\kepviib.  \\end{itemize}  \\subsection{Comparison of three hot-Neptunes}  The eccentricity of \\kepivb\\ is marginally detected to 2-$\\sigma$ confidence. If confirmed, this would mean that three of the four hot-Neptunes discovered through the transit method have eccentricities around 0.15-0.20, the other two being GJ 436b and HAT-P-11b. HAT-P-26b is the fourth transiting hot-Neptune, recently discovered by \\citet{hartman:2010} and this system also possesses a marginal eccentricity of $e = 0.124 \\pm 0.060$. With HAT-P-26b possessing a FAP of 12\\% and Kepler-4b being 11.8\\%, it would seem likely that at least one of these two systems is genuinely eccentric.  In all cases, the  tidal circularization time is expected to be much lower than the age of the system for Jovian-like tidal dissipation values, which raises the question as to why these planets still retain eccentric orbits. Generally two hypotheses have been put forward: i) an unseen perturbing planet pumps the eccentricity of the hot Neptune ii) hot Neptunes have larger $Q_P$ values than expected.  At the time when only one example of an eccentric Neptune was known, GJ 436b, \\citet{ribas:2008} made the reasonable deduction that is was more likely an unseen perturber was present. However, given that three such planets are now known, Occam's razor seems to favor the alternative hypothesis. For all three hot-Neptunes to have perturbing planets which all remain hidden from detection and yet produce nearly identical eccentricity pumping levels appears to be a less likely scenario than that of a common solution due to intrinsically different tidal dissipation values and formation history. For \\kepivb, the observed eccentricity indicates $Q_P \\geq 10^6$ if no eccentricity pumping is occurring.  We also point out that the physical properties of Kepler-4b differ greatly between the eccentric and circular fits (see Table~\\ref{tab:neptunes}), due to the large impact of eccentricity on the lightcurve derived stellar density \\citep{kip10a}. Further RV measurements are likely required to resolve the eccentricity of this system, since an occultation is generally not expected to be observable given the very low $R_P/R_*$.   \\begin{table*} \\caption{ \\emph{Comparison of two known hot-Neptunes with Kepler-4b. GJ 436b values taken from \\citet{tor08} except eccentricity which comes from \\citet{gillon:2007}; HAT-P-11b values taken from \\citet{bakos:2009}. The classification of Kepler-4b depends greatly upon whether the system is confirmed as maintaining an eccentric orbit or not.} } \\centering % \\begin{tabular}{l l l l l l l} % \\hline\\hline % Planet & Mass/$M_J$ & Radius/$R_J$ & Density/gcm$^{-3}$ & log$(g/$cgs) & e & $T_{eq}$/K \\\\ \\hline GJ 436b & $0.0729_{-0.0025}^{+0.0025}$ & $0.376_{-0.009}^{+0.008}$ & $1.69_{-0.12}^{+0.14}$ & $3.107 \\pm 0.040$ & $0.16 \\pm 0.02$ & $650 \\pm 60$ \\\\ HAT-P-11b & $0.081 \\pm 0.009$ & $0.422 \\pm 0.014$ & $1.33 \\pm 0.20$ & $3.05 \\pm 0.06$ & $0.198 \\pm 0.046$ & $878 \\pm 15$ \\\\ Kepler-4b (circ) & $0.081 \\pm 0.031$ & $0.460_{-0.084}^{+0.272}$ & $0.86_{-0.63}^{+0.97}$ & $2.90_{-0.38}^{+0.25}$ & $0^{\\mathrm{*}}$ & $1777_{-132}^{+308}$ \\\\ Kepler-4b (ecc) & $0.096 \\pm 0.023$ & $0.79_{-0.109}^{+0.145}$ & $0.24_{-0.13}^{+0.46}$ & $2.58_{-0.20}^{+0.28}$ & $0.25 \\pm 0.12$ & $2215_{-339}^{+233}$ \\\\ \\hline Neptune & $0.05395$ & $0.3464$ & $1.638$ & $3.047$ & $0.011$ & 9.6 \\\\ Uranus &  $0.0457$ & $0.3575$ & $1.27$ & $2.939$ & $0.044$ & 12.1 \\\\ [1ex] Saturn & $0.299$ & $0.843$ & $0.687$ & $3.019$ & $0.056$ & 26.0 \\\\ [1ex] \\hline\\hline % \\end{tabular} \\label{tab:neptunes} % \\end{table*}  \\subsection{Secondary eclipses}  We detect a secondary eclipse for the bloated \\kepviib\\ of depth $(47 \\pm 14)$\\,ppm and 3.5-$\\sigma$ confidence. This is above our formal detection threshold of 3-$\\sigma$ thus represents an unambiguous detection. Thermal emission is an unlikely source for the eclipse since the depth indicates a brightness temperature of $2570_{-85}^{+108}$\\,K, which is far in excess of the equilibrium temperature of the planet of $(1554 \\pm 32)$\\,K. A geometric albedo of $A_g = (0.38 \\pm 0.12)$ offers a more plausible explanation for this eclipse depth. However, we do note that a significant nightside flux is apparently present which is not consistent with the reflection hypothesis. We believe that this nightside flux is likely an artifact of the Kepler pipeline whose effects may mimic a long-cut filter. Although \\kepviib\\ exhibits similarities to HD 209458b in terms of its very low density, the albedo clearly marks the planet at distinct given that $A_g < 0.17$ for HD 209458b \\citep{row08}. The study of \\citet{burrows:08} seems to indicate that the albedo requires the presence of reflective clouds, possibly composed of iron and/or silicates, as seen in L-dwarf atmospheres. It is interesting to note that \\kepvii\\ is $1.3 \\pm 0.2$ times more metal-rich than HD 209458, and the planet has a higher equilibrium temperature of $(1565 \\pm 30)$\\,K than that of HD 209458b with $(1130 \\pm 50)$\\,K.  We also note that, to our knowledge and if confirmed, the above measurement is the first determination of a transiting planet's albedo at visible wavelengths, and thus the first detection the associated reflected light. This compliments the recent polarimetric detection of reflected light of HD 189733b was made by \\citet{ber08}, and recently confirmed in \\citet{ber11}. Visible secondary eclipses of other planets have been made for HAT-P-7b \\citep{bor09} and CoRoT-1b \\citep{sne09} but in both cases the eclipse can be explained through a combination of reflected light plus thermal emission or even simply pure thermal emission for models of large day-night contrasts, which is in fact expected for these extremely hot-Jupiters \\citep{fortney:2008}. In contrast, thermal emission cannot explain the Kepler-7b secondary eclipse unless the planet has an internal heat source with a luminosity equivalent to that of a M6.5 dwarf star, which is highly improbable.  We also find weak significance of a secondary eclipse for \\kepvb\\ of $\\sim$2-$\\sigma$ confidence. Further transits will either confirm or reject this hypothesis but we note that the obtained depth is quite reasonable for this planet. Whilst thermal emission again seems unlikely based upon the required temperature of 2500\\,K versus the equilibrium temperature of 1800\\,K, a geometric albedo of $A_g = (0.15 \\pm 0.10)$ offers a satisfactory explanation for the eclipse.  For \\kepvib\\ and \\kepviib, we find no evidence of a secondary eclipse, but the measurements do constrain the geometric albedos to $A_g \\leq 0.32$ and $A_g \\leq 0.63$ at 3-$\\sigma$ confidence. For \\kepivb, the measurement place no constraints on the geometric albedo.  \\subsection{Differences with the discovery papers}  For almost all cases, we find significantly different impact parameters from the discovery papers, but self-consistent between methods A and B. There are also several examples of other parameters being different. These can be easily reviewed by consulting Tables \\ref{tab:kep4tab}, \\ref{tab:kep5tab}, \\ref{tab:kep6tab}, \\ref{tab:kep7tab} and \\ref{tab:kep8tab}.  For example, for \\kepivb, we find a very low stellar density implying a larger stellar radius and thus larger planetary radius.  % Other examples of differing parameters include the eccentricity and secondary eclipse parameters already discussed. %  \\subsection{On the error budget of planetary parameters}   The exquisite \\lcs\\ measured by Kepler lead to very low uncertainties on certain parameters, even in the long-cadence mode.  One example is the fitted ratio-of-radii $p$, where the error is about 0.5\\% (e.g., $p=\\kepviieccLCrprstar$ for \\kepviib). The median error for the first 70 known transiting exoplanets (TEPs) on $\\rpl/\\rstar$ is 0.9\\%, about double that of \\kepviib. Limiting factors on the precision of this purely geometrical factor are the limb-darkening values for the \\kept\\ bandpass and the slight degeneracy with other parameters that are affected by the long-cadence binning, even if re-sampled models are fitted (such as the $b$ impact parameter). Using short-cadence data, deriving precise limb-darkening coefficients, and accumulating many transits will significantly improve errors in $\\rpl/\\rstar$.  The typical errors on the period for the first five \\kept\\ planets are $0.00004$\\,days, about 6 times greater than the median error of the period of the first 70 published TEPs, or about 20-times the error quoted in e.g.~\\citet{hartman:09} for HAT-P-12b, an example of a TEP discovery using photometry data spanning 2 years. This error on the \\kept\\ planet periods, however, comes from a dataset with short time-span (44\\,days), while the ground-based results come from multi-year campaigns. The median error on the ephemeris ($T_c$) for the current 5 \\kept\\ planets is $0.000145$\\,days, about half that of the known TEPs ($0.0003$\\,days). Precision of the ephemeris data will greatly improve during the lifetime of the \\kept\\ mission just by accumulating more data, and switching to short-cadence mode.  At first glance perhaps somewhat surprisingly, the errors on the \\kept\\ planetary masses and radii are not different from the rest of the transiting exoplanet population. The typical error on planetary masses for 70 transiting exoplanets is $0.06$\\,\\mjup, whereas the error for radii is $0.05\\,\\rjup$. In our analysis the errors for planetary masses and radii for Kepler planets have roughly the same values.  There are multiple reasons behind this. Planetary mass and radius scale with the respective parameters of the host star. In our analysis, \\mstar\\ and \\rstar\\ are determined from stellar isochrones, using \\teff, \\feh\\ and \\arstar. Each of these quantities have significant errors. The stellar \\teff\\ and \\feh\\ come from SME analysis of the spectra, which, in turn, due to the faintness of the \\kept\\ targets, are of low S/N, and have larger than usual errors. The other input parameter, \\arstar, is related to the \\lc\\ parameters \\zrstar, $b$, and the RV parameters $k$ and $h$. While \\zrstar\\ is quite precise for the \\kept\\ transits (median error 0.02 as compared to 0.15 for the known TEPs), the $b$ impact parameter from the current LC \\kept\\ data-set have errors comparable to ground-based transits. The main limiting factor, however, are the $k$ and $h$ Lagrangian orbital parameters from the RV data. Due to the faintness of the targets, the orbital fits have considerable errors in $k$ and $h$. As a result, the error on \\arstar\\ for the current \\kept\\ data in this work (median error 0.26) is similar to that of the known TEPs (median error 0.2). This effect is well visible on the isochrone figures (Fig.~\\ref{fig:kep4iso}, \\ref{fig:kep5iso}, \\ref{fig:kep6iso}, \\ref{fig:kep7iso}, \\ref{fig:kep8iso}), where the 1-$\\sigma$ and 2-$\\sigma$ confidence ellipses are shown for all \\kept\\ planets for both the circular and eccentric solutions, with the latter clearly covering a much larger area on the isochrones. Finally, the error in the planetary mass also scales with the $K$ RV semi-amplitude, which has typical errors for the \\kept\\ planets similar to the median error of the known planets ($\\Delta K \\approx 4$\\,\\ms). Given the faintness of the \\kept\\ targets this is an impressive result.  A breakthrough in both accuracy and precision of \\mpl\\ and \\rpl\\ can be expected from exploiting the extraordinary precision of \\kept, and nailing down the stellar parameters via asteroseismology (see \\citet{mulet:2009}, \\citet{christ:2010}). Alternatively (or simultaneously), if parallaxes for the \\kept\\ targets become available (possibly from the \\kept\\ data), then analyses similar to that performed for HAT-P-11b \\citep{bakos:2009} can be performed, where the ``luminosity-indicator'' in the isochrone fitting will be the absolute magnitude of the stars, as opposed to the $\\arstar$ constraint. A final possibility is a dynamical measurement of the stellar mass by measuring the transit timing variations of two TEPs within the same system \\citep{agol:2005}.    \\subsection{Transit timing with LC data}  After the discovery of an exoplanet in Kepler's field, the subsequent photometry is switched to short-cadence (SC) mode. It was generally expected that this would be the only way to produce meaningful TTV studies. However, we have found that by carefully correcting for the long integration time, the long-cadence (LC) photometry yields transit times consistent with a linear ephemeris and of r.m.s.~scatters of 10-20 seconds are possible.  Despite the evident timing precision possible with the LC photometry, some issues require further investigation. The effect of phasing, as defined in \\S2.3.4 remains somewhat unclear but in numerous instances we find that peaks in the TTV and TDV periodograms occur near to the phasing period.  Subsequent aliasing and frequency mixing also seems to lead to numerous false positives in the periodograms and these issues may be diminished by using SC data, but are unlikely to completely disappear.  Out of all of the studied planets, we found only one peak in the timing periodograms which does not appear spurious. \\kepvib\\ exhibits a 2.6-$\\sigma$ peak of period $(17.27 \\pm 0.84)$\\,days and amplitude $(19.7 \\pm 5.0)$\\,seconds. We investigated the plausibility of this signal being a perturbing planet or exomoon but found no convincing supporting evidence, although such hypotheses could not excluded. Our favored hypothesis is that of stellar rotation inducing such a signal, as reported by \\citet{alo08} for CoRoT-2b. The host star has a rotational period of $23.5_{-5.9}^{+11.7}$\\,days and thus it is consistent with the TTV period. A bisector analysis, as performed by \\citet{alo08}, is not possible for this data due to the long cadence. However, the SC data will be able to either confirm or reject this hypothesis.  Finally, we note that the uncertainties in all of the parameters from individual fits show evidence for being overestimates, based upon the observed scatter of these parameters between various transits. Both methods employed in this analysis found very similar errors and thus the reason for consistent overestimation of the errors remains unclear. Both methods employed Markov Chain Monte Carlo (MCMC) techniques for obtaining these errors, which may be linked to the overestimation. If the MCMC trials are unable to sample the true global minimum, perhaps due to slightly erroneous limb darkening coefficients or insufficient numerical resolution in the integration time compensation procedure, the errors are expected to be overestimated. Nevertheless, in searching for evidence of signals, the accurate estimation of error bars is just as important as the best-fit value. We have proposed a F-test periodogram which is insensitive to absolute errors, only their relative weightings, and also penalizes overly-complex models. Further detailed investigation with future Kepler timings will be possible and should shed light on this issue.  \\subsection{Constraints on planets, moons \\& Trojans}  We find no convincing evidence for perturbing planets, moons or Trojans in any of the systems studied here. We are able to exclude the presence of mean-motion resonant (MMR) planets for \\kepvb, \\kepvib\\ and \\kepviiib\\ of masses $\\geq$ $0.79 \\mearth$, $0.38 \\mearth$ and $0.50 \\mearth$ respectively, to 3-$\\sigma$ confidence. The corresponding resonant configurations are period ratios of 3:2, 4:3 and 4:3. We do not yet possess sufficient temporal baseline to investigate MMR planets for \\kepivb\\ and \\kepviib.  Extrasolar moons at a maximum orbital separation are excluded by the TTV measurements for \\kepivb\\ through to \\kepviiib\\ of masses $\\geq$ $11.0 \\mearth$, $10.6 \\mearth$, $4.8 \\mearth$, $2.5 \\mearth$ and $2.1 \\mearth$ respectively, to 3-$\\sigma$ confidence. Extrasolar moons at a minimum orbital separation are excluded by the TDV measurements of masses $\\geq$ $13.3 \\mearth$, $17.3 \\mearth$, $8.2 \\mearth$, $5.9 \\mearth$ and $21.5 \\mearth$ respectively, to 3-$\\sigma$ confidence.  For \\kepivb, \\kepviib\\ and \\kepviiib, we do not yet possess sufficient temporal baseline to search for Trojans through TTV measurements. However, for \\kepvb\\ and \\kepvib\\ we are able to exclude Trojans of angular displacement $\\sim 10^{\\circ}$ from L4/L5 of cumulative mass $\\geq$ $3.14 \\mearth$ and $0.67 \\mearth$ respectively, to 3-$\\sigma$ confidence. For all five planets, we can inspect the photometry at $\\pm P/6$ from the transit center for signs of a photometric dip due the occultation of the host star by the Trojans. This search yielded no detections but does exclude Trojans of effective radii $\\geq$ $1.22 R_{\\oplus}$, $1.13 R_{\\oplus}$, $0.87 R_{\\oplus}$, $1.11 R_{\\oplus}$ and $1.58 R_{\\oplus}$ for the five planets sequentially, to 3-$\\sigma$ confidence."
    },
    "1004/1004.3823_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{s.intro}  After three decades of development our understanding of the space weathering phenomenon may be converging on its cause and effects. Originally proposed \\citep{bib.cha73} as a solution to the mismatch between the color of the most common meteorites, ordinary chondrites, and their likely source region, inner main belt S-complex asteroids, the space weathering hypothesis is that the surface colors of asteroids change with exposure to the space environment \\citep{bib.hap73,bib.hap68}.  The idea has held up to the test of time and in recent years has led to measurements of the rate of color change on S-complex asteroids yielding empirical models that predict the rate of reddening of their surfaces \\citep{bib.jed04, bib.nes05,bib.wil08}.  This work extends our space weathering model to the youngest dated asteroids in the main belt that are less than $10^6$ years old and, for the first time, measures the surface gardening rate on the asteroids due to impact-generated regolith cycling.  \\citet{bib.wil08}'s space weathering model applies exclusively to S-complex asteroids.  Other types of asteroids may also undergo space weathering and we expect that they would obey a model of similar functional form with different parameters.  For instance, if C-complex asteroids undergo space weathering they must have a different color range because they occupy a different region of principal component space and the weathering rate and even the sense of coloration may be different \\citep{bib.nes05}.  The gardening rate would probably also be different due to different material density and strength.  In this work we study the space weathering of S-complex asteroids without particular concern for the agent causing the weathering though the likely culprit is energy deposition due to particle bombardment \\eg micrometeorites, solar protons and cosmic rays.  \\citet{bib.mar06} claim that Sun-related effects are dominant but they assumed that micrometeorite impacts are independent of heliocentric distance. Since neither their flux nor their velocity are distance independent \\citep{bib.cin92} it is unclear whether it is justified to downgrade their contribution to the space weathering agent inventory. Laboratory-based studies suggest that the asteroidal space weathering mechanism, long confused with the lunar processes producing agglutinates, involves the coating of near surface semi-transparent nanophase grains by a vapor or sputter deposited film of metallic iron \\citep[\\eg][]{bib.sas01,bib.pie00}.  While measuring the color of asteroid surfaces is relatively straightforward and asteroid taxonomy based on colors has been a mainstay of planetary science for decades the determination of asteroid surface ages is a relatively recent innovation.  The surface ages of asteroid family members may be determined from dynamical simulations of the evolution of the family's orbit element distribution\\citep[\\eg][]{bib.vok06a, bib.vok06b, bib.nes05, bib.nes02, bib.mar95}.  The dynamical methods include family size frequency distribution (SFD) modeling, global main belt SFD modeling, modeling of family spreading via thermal forces and backward numerical integration of orbits.  The combination of asteroid surface colors with their ages allowed the first astronomical determination of the space weathering rate \\citep{bib.jed04,bib.par08}.  Space weathering rate measurements utilize the fact that asteroid family members are fragments of a single collisionally disrupted parent body that formed at the same time and assume that fresh unweathered material from the collision debris cloud coats all members of a family with a homogeneous regolith layer.  Therefore, all family members should start out with similar color and weather in tandem.  We ignore grain size color effects for a number of reasons.  First, we believe that unless grain size is correlated with age its effect will average out over the members of a family.  Second, we will see below that km scale asteroids, in contrast to moon sized bodies with highly comminuted regoliths, are gravitationally sorted rubble piles whose surfaces are dominated by boulders.  Fines are largely absent, perhaps sequestered beneath the bouldered surface, so asteroid regolith appears to vary with size.  This would diminish grain size observational effects in the sub-milligee environment of small asteroids. Finally, \\citet{bib.nes05} and \\citet{bib.jed04} did not find any correlation between color and size, just color and age.  \\citet{bib.jed04}'s space weathering rate measurements encapsulated `the color' of an asteroid's surface as its first principal component color \\begin{equation} PC_1 = 0.396\\;(u-g) + 0.553\\;(g-r) + 0.567\\;(g-i) + 0.465\\;(g-z) \\label{eq.pc1} \\end{equation} \\noindent that correlates strongly with the average slope of the spectrum.  The principal component color provides the linear combination of filter magnitudes having the greatest variability over the sample of asteroids.  In this case $PC_1$ is specific to the asteroids in the Sloan Digital Sky Survey (SDSS) 2$^{nd}$ Moving Object Catalog.  While most of the color variability in the SDSS asteroid sample is captured in $PC_1$, the second principal component, \\begin{equation} PC_2 = -0.819\\;(u-g) + 0.017\\;(g-r) + 0.090\\;(g-i) + 0.567\\;(g-z), \\label{eq.pc2} \\end{equation} \\citep{bib.nes05} corresponds to the asteroid spectrum's curvature and is used in this analysis to identify asteroid taxonomy.  Motivated by the fact that uniform surface irradiation should produce an exponentially decaying amount of unweathered surface \\citet{bib.jed04} showed that the color of S-complex asteroids as a function of time $t$ can be expressed as \\begin{equation} PC_1(t) = PC_1(0) + \\Delta PC_1[1 - e^{-(t/\\tau)^\\alpha}] \\label{eq.jedpc1} \\end{equation} \\noindent where $PC_1(0)$ is the unweathered color of fresh surface material, $\\Delta PC_1$ is the magnitude of the weathering color change after a long period of time, $\\tau$ is the characteristic time for space weathering and $\\alpha$ is a generalizing exponent. \\citet{bib.wil08} improved upon the earlier work by refining the color of the $\\lae 5$~My old Iannini family, eliminating the Eos family that is no longer considered to be in the S-complex, and carefully refitting the data to the functional form given above to determine $PC_1(0) = 0.31 \\pm 0.04$, $\\Delta PC_1 = 0.31 \\pm 0.07$, $\\tau = 570 \\pm 220$ My and $\\alpha = 0.53 \\pm 0.19$.  The characteristic time scale of $570 \\pm 220$~My for color change in the main asteroid belt \\citep{bib.wil08} is in agreement with pulsed laser experiments \\citep{bib.sas01} on silicate pellets intended to simulate micrometeorite bombardment.  They suggested a characteristic time for weathering of 100 My at $1$~AU, equivalent to about 700 My in the main belt assuming that the Sun is the source of the weathering agent and a $r^{-2}$ dependence on its strength.  However, \\citep{bib.loe09} argue that \\citet{bib.sas01} overestimated the space weathering characteristic timescale due to an incorrect flux calculation leading.  Basically, the time scale extrapolation from lab results is complicated because it is unclear how to evaluate the contribution from several different factors.  For instance, in the main belt micrometeorite impact speeds will be lower ($\\sim 5$ km/s) than at 1 AU ($\\sim 20$ km/sec) and their impact energy will be $\\sim$16x weaker.  \\citet{bib.pie00} estimated characteristic aging times of $100-800$ My for lunar surfaces by comparing ages from craters dated radiometrically or by cosmic ray exposure ages to spectral differences.  Correcting the rate to the center of the main belt suggests weathering times on the order of $\\sim$600-4800~My.  Even though the lunar surface is not identical to S-complex asteroid surfaces the time scales are similar.  Finally, we note that craters on (243)~Ida are bluer than their surrounding background terrain \\citep{bib.vev96}.  The craters correspond to freshly exposed and unweathered regolith \\citep{bib.vev96} while other parts of the asteroid's surface indicates an age of about 1 Gy \\citep{bib.gre96}.  The wide range in diameters (a proxy for crater age) of blueish crater suggests that the space weathering time must be long.  On the other hand there are claims of faster surface weathering time scales such as \\cite{bib.tak08}'s upper limit of 450 kyr based on a shallow 1 $\\mu$m absorption band observed on (1270)~Datura.  However, we expect that the space weathering phenomenon is a relatively subtle effect easily masked by stochastic variations between asteroids due to \\ie mineralogical and morphological differences and/or collisional and cratering events.  The space weathering effect can only be identified in specific regions on an asteroid as observed with {\\it in situ} spacecraft measurements on (951)~Gaspra \\citep{bib.hel94}, (243)~Ida \\citep{bib.vev96}, and (25143)~Itokawa \\citep{bib.ish07}] or as an ensemble effect on a statistically large sample of asteroids.  It is therefore difficult to make a general conclusion based on a single asteroid such as (1270) Datura.  Indeed, our measurement of that asteroid's $PC_1 = 0.41 \\pm 0.02$ is redder by $2.2 \\sigma$ than its predicted color from \\citet{bib.wil08}.  Takato's measurement of a shallow 1 $\\mu$m absorption band which would tend to redden the overall spectrum is therefore generally consistent with our measurement.  We expect that the `redness' of individual small rubble pile asteroids in a sample is affected more by random surface variations than larger regolith-rich asteroids with finer surface materials.  In this case (1270) Datura is redder than the Datura cluster average so its particular value can be misleading.  The four (1270) Datura members in this work have mean $PC_1 = 0.305$ with a RMS of $0.278$.  The large standard deviation is not due to measurement error --- it is a result of intrinsic color differences between the members of the (1270) Datura family and is typical of other families.  Therefore we would discount the resulting 450 kyr weathering time.  Another short time scale measurement was reported from recent lab experiments on olivine powder by \\citet{bib.loe09} (and references therein) simulating solar wind effects at 1 AU by bombardment with 4 keV protons.  They conclude that spectral reddening caused by the solar wind should saturate in $\\sim$ 5 kyr.  \\citep{bib.ver09} combined spectral data from a sample of four members of young clusters with archival meteorite spectra and found that space weathering is substantially complete in $<$ 1 My but then continues at a slower pace up to several Gy.  They attribute the first stage to the solar wind (\\citet{bib.str05} lab simulation) and the second to micrometeorite bombardment.  This scenario depends critically on the starting color that they derive from meteorite data.  It is not clear how to reconcile the information from these various lab results that differ in time scale by several orders of magnitude with our observations.  \\citet{bib.wil08}'s weathering model was derived from families that were several~My (\\eg Karin, Iannini) to several Gy (\\eg Eunomia, Maria) old.  The family age estimates have typical errors of $\\sim 40$\\% resulting from fundamental limitations in the dating techniques \\citep{bib.nes07}.  Although there were nine known S-complex families with ages from tens of Mys to a few Gys at that time there were only two younger than 10 My.  Refining the space weathering rate at even younger ages requires a large sample of young asteroid families, which are typically produced by the catastrophic disruption of small asteroids, have only a small number of detectable members, and are therefore difficult to identify.  But in recent years the number of catalogued asteroid orbits has reached hundreds of thousands and includes many asteroids smaller than 1~km.  This large sample of asteroids allows the identification of rare small clusters\\footnote{As pairs of asteroids are (tautologically) composed of two asteroids, the established families have dozens to thousands of members, and known `clusters' have between three and seven members we define a cluster as a small family with three to about ten asteroids.  Pairs have distinct formation mechanisms from families and clusters.  Families form from larger parent bodies than clusters and are therefore older on average.  The youngest known family, Iannini, is $\\sim$3 My old while the clusters are less than 1 My old.} originating from collisions less than 1 My ago.  \\citet{bib.nes06a,bib.nes06b} found four such clusters with a total of 16 members.  The clusters are named after their largest known members: (1270)~Datura, (14627)~Emilkowalski, (21509)~Lucascavin and (16598) 1992 YC2.  The progression to ever smaller clusters reached its logical limit with \\citet{bib.vok08}'s discovery of 60 pairs of asteroids in extremely similar orbits --- much more alike than would be expected based on the density of proper elements for other asteroids with similar orbit elements.  The pairs are thought to have formed less than 500 kyr ago based on the dynamical evolution time of the pair member's orbits.  Their list was updated and restricted to only 36 pairs \\citep{bib.pra09b}.  Pairs belonging to known young families such as (1270) Datura or (832) Karin were excluded because of the possibility that the catastrophic impact and subsequent family formation process may create paired asteroids.  Such cases were eliminated to focus on pairs that were created in isolation and presumably by the same method.  The formation method of pairs may involve a critical distinction from that of clusters or families. Whereas the latter form in catastrophic collisions, pairs have alternative possible formation methods.  Some of these methods may be gentle processes not involving resetting of the surface as discussed in \\S\\ref{ss.comparemodel}.  Shortly after the publication of \\citet{bib.wil08} we obtained observations of some members of the sub-My old families and found that they were significantly redder than predicted.  Did this imply a problem with our space weathering model or is the space weathering phenomenon more complicated than suggested by the simple model of eq.~\\ref{eq.jedpc1}?  We were also concerned with the functional form of the space weathering model of \\citet{bib.wil08} \\ie eq. \\ref{eq.jedpc1}.  If space weathering is an isolated process the exponent $\\alpha$ should be unity but \\citet{bib.jed04} fit $\\alpha = 0.53 \\pm 0.19$.  Their argument was that $\\alpha$ is a generalizing factor that accounts for both space weathering and regolith gardening which acts to counteract the surface aging by slowing turning over the asteroid's surface.  To investigate these questions we collected color and spectral data from several sources for members of the young families and pairs.  We obtained spectra of members of the sub-My asteroid clusters \\citep{bib.nes06a,bib.nes06b} and were provided spectra for some of the objects observed by \\citet{bib.mot08}.  We also located archived photometry for 19 pair members in the Sloan Digital Sky Survey Data Release 7 Moving Object Catalog 4 (SDSS DR7 MOC4) \\citep{bib.par08}. We then developed a new asteroid surface color-age model that explicitly separates the effects of weathering and gardening and eliminates the need for the unexplained generalizing factor ($\\alpha$).  The new model allowed us to measure the characteristic timescales for both weathering and gardening on the S-complex asteroids.  Finally, we independently calculated the gardening timescale on main belt asteroids from their size distribution, impact rates and estimates of crater and ejecta blanket size. ",
        "conclusions": "\\subsection{Sub-My cluster member taxonomy} \\label{ss.submyrClusterTaxonomy}  We observed the brightest member of each of the four sub-My clusters to identify their taxonomic type.  The spectra or multiband photometry for each object are shown in Figure~\\ref{f.emillucdatvisir}. (1270) Datura, (16598) 1992 YC2 and (21509)~Lucascavin all show classic S-complex characteristics in the visible --- a 0.75~$\\mu$m peak and a 1.0~$\\mu$m absorption band.  Datura also shows an inflection near 0.55 $\\mu$m indicating a fresher surface relative to older S asteroids with smoother spectra \\citep{bib.ish07, bib.hir06}.  (21509)~Lucascavin also shows the 2.0~$\\mu$m band typical of pyroxene. Using the techniques described in \\S\\ref{s.classid} we identify (1270)~Datura as an Sk, close to \\citet{bib.mot08}'s identification of Sl, (21509)~Lucascavin also as an Sk, and (16598) 1992 YC2 as member of the S-complex.  Our visible and IR spectra of (14627)~Emilkowalski does not show the 1.0 and 2.0~$\\mu$m absorption bands typical of the S-complex and we classify it as T type.  As our space weathering model only applies to asteroids within the S-complex we ignore (14627)~Emilkowalski for the remainder of this work.  \\begin{figure}[!ht]\\small \\centerline{\\includegraphics[width=5in,angle=90]{emillucdatvisir.ps}} \\caption\\small{Visible and near IR spectra of (14627)~Emilkowalski and (21509)~Lucascavin obtained with UH2.2~m/SNIFS and IRTF/SpeX, (1270)~Datura visible spectrum obtained with UH2.2~m/SNIFS, and SDSS (16598) 1992 YC2 photometry.  The spectrum of (14627)~Emilkowalski is normalized to 1.0 at 0.55~$\\mu$m and the others are offset vertically.  The gap at $\\sim 1.9 \\mu$m results from removing a sky absorption band.  The (1270)~Datura visible spectrum, (14627)~Emilkowalski and (21509)~Lucascavin IR spectra are all smoothed fits to the data (causing a spurious mismatch to the visible spectrum), the others are binned.} \\label{f.emillucdatvisir} \\end{figure}  It is unsurprising that (1270)~Datura and (21509)~Lucascavin were identified in the S-complex as both clusters are located in the inner main belt which is dominated by S-complex asteroids.  Similarly, (14627)~Emilkowalski and (16598) 1992 YC2 are located in the middle of the main belt where X and S-complex types are common.  Having identified three of the sub-My clusters within the S-complex for which we intend to examine the effects of space weathering and gardening we obtained data for eight of the cluster members as provided in Table~\\ref{t.Sub-My-cluster-pc1}.  \\begin{table}[!ht]\\small \\begin{center} \\title{Table~\\ref{t.Sub-My-cluster-pc1}. S-complex sub-My cluster members's derived spectral and age data} \\begin{tabular}{cccccc} & & \\\\ \\tableline\\tableline Cluster    & Asteroid   & Source   & Slope                & $PC_1$ & Age \\\\ &            &          & (Reflectance/$\\mu$m) &          & (kyr) \\\\ \\tableline &(1270) Datura& MDN     & $0.355$              & $0.391$  & $530 \\pm 20$ \\\\ &(1270) Datura&this work& $0.382$              & $0.414$  & \" \\\\ Datura     &(203370) 2001 WY$_{35}$& MDN&$-0.155$         &$-0.053$  & \" \\\\ & (60151) 1999 UZ$_6$   & MDN&$0.606$          & $0.609$  & \" \\\\ & (90265) 2003 CL$_5$   & MDN&$0.205$          & $0.260$  & \" \\\\ \\tableline\\tableline &            &          &                      mean & $0.305 \\pm 0.278$ \\\\ &            &          &        \\\\ \\tableline & (21509) Lucascavin&this work& $0.476$        & $0.496$  & $300-800$ \\\\ Lucascavin & (209570) 2004 XL$_{40}$&MDN                  & $0.314$  & $0.355$  & \" \\\\ & (180255) 2003 VM$_{9}$& this work            & $0.327$  & $0.366$  & \" \\\\ \\tableline\\tableline &            &          &                 mean & $0.406 \\pm 0.078$ \\\\ &            &          &        \\\\ \\tableline 1992 YC2   & (16598) 1992 YC2&MDN  & $0.083 \\pm 0.02$     & $0.155 \\pm 0.027$ & $135-220$ \\\\ \\tableline\\tableline &            &          &        \\\\ &            &          & Sample Mean & $0.36 \\pm 0.07$ & $511 \\pm 10$ \\\\ \\tableline\\tableline \\end{tabular} \\end{center} \\caption\\small{Derived color data and age estimates \\citep{bib.nes06a,bib.nes06b} for members of three S-complex sub-My clusters.  $PC_1$ was calculated using eq.~\\ref{eq.slopepc1}. Errors on the slope and $PC_1$ are not included as the error on the family's color is dominated by the distribution of $PC_1$ values within a cluster.  The errors on cluster means are standard deviations except for (16598) 1992 YC2 for which the measurement error is provided since it was the only object observed in the cluster.  The Sample Mean includes all eight cluster members with the error on the mean.} \\label{t.Sub-My-cluster-pc1} \\end{table}    \\subsection{Taxonomy and orbit distribution of asteroid pair members} \\label{ss.pairMemberTaxonomy}  \\begin{table}[!ht] \\myfontsize % \\begin{center} \\title{Table~\\ref{t.Asteroid Pairs}. Asteroid Pairs} \\begin{tabular}{rlcccrlccccc} & & \\\\ \\tableline\\tableline \\multicolumn{2}{c}{$Ast_1$}   & SMASS  &Tholen&Quality& \\multicolumn{2}{c}{$Ast_2$}& a      & e        & i       & $H_1$ & $H_2$ \\\\ &        &        &      &     &     &       & (AU)    &          &(degrees)&         &         \\\\ \\tableline 1986& JN$_{1}$   & X,  Xe & EMP, EMP & good & 2001& XO$_{105}$&$1.946$&$0.0601$& $23.710$& $13.5$ & $17.4$ \\\\ 2000& WX$_{167}$ & Xe, T  & EMP, T   & fair & 2007& UV    & $1.909$ & $0.0613$ & $23.096$& $16.2$ & $17.1$ \\\\ 2001& MD$_{30}$  & Xe, X  & EMP, EMP & fair & 2004& TV$_{14}$&$1.938$&$0.0886$ & $19.987$& $14.9$ & $17.2$ \\\\ {\\bf 2000} & {\\bf NZ}$_{{\\bf 10}}$ & {\\bf L, Sl} & {\\bf S, S} & {\\bf good} & {\\bf 2002} & {\\bf AL$_{{\\bf 80}}$} & {\\bf 2.287} & {\\bf 0.1801} & {\\bf 4.097} & {\\bf 14.1} & {\\bf 16.2} \\\\ \\smallskip {\\bf 2002} & {\\bf AL}$_{{\\bf 80}}$ & {\\bf Sl, S} & {\\bf S, S} & {\\bf good} & {\\bf 2000} & {\\bf NZ$_{{\\bf 10}}$} & {\\bf '' }   & {\\bf '' }    & {\\bf '' }   & {\\bf 16.2} & {\\bf 14.1} \\\\ 1999& KF    & L, Sl & S, S     & good & 2008& GR$_{90}$   & $2.327$ & $0.2339$ & $1.777$ & $15.0$ & $17.2$ \\\\ 2002& GP$_{75}$& L, S  & S, S  & good & 2001& UR$_{224}$  & $2.340$ & $0.1727$ & $3.865$ & $15.7$ & $17.2$ \\\\ 2006& AL$_{54}$& L, Sl & S, S  & good & 2000& CR$_{49}$   & $2.272$ & $0.1763$ & $4.591$ & $16.8$ & $14.3$ \\\\ \\smallskip 1962& RD    & Sl, Ld & S, S    & good & 1999& RP$_{27}$   & $2.198$ & $0.1775$ & $1.129$ & $13.1$ & $15.3$ \\\\ 1997& CT$_{16}$  & Sl, Sa & S, S     & good & 2002& RZ$_{46}$& $2.186$ & $0.1672$ & $4.599$ & $15.4$ & $16.4$ \\\\ 2000& RV$_{55}$  & Sl, Sa & S, S     & good & 2006& TE$_{23}$& $2.657$ & $0.1026$ & $2.245$ & $14.9$ & $16.8$ \\\\ 2004& RJ$_{294}$ & S,  Sr & S, S     & good & 2004& GH$_{33}$& $2.268$ & $0.0981$ & $4.238$ & $18.2$ & $16.7$ \\\\ \\smallskip 2003& SC$_{7}$   & Sk, K  & S, S   & good & 1998& RB$_{75}$  & $2.264$ & $0.1114$ & $7.263$ & $16.6$ & $14.6$ \\\\ 2000& GQ$_{113}$ & Sq, Sk & S, S   & good & 2002& TO$_{134}$ & $2.324$ & $0.1319$ & $5.515$ & $14.4$ & $16.3$ \\\\ 1983& WM    & Sr, Sa & S, S   & good & 1999& RC$_{118}$ & $2.320$ & $0.0790$ & $5.726$ & $13.7$ & $14.6$ \\\\ 2003& YK$_{39}$  & Sr, Q  & S, Q   & good & 1998& FL$_{116}$ & $2.187$ & $0.0845$ & $3.736$ & $18.3$ & $15.0$ \\\\ \\smallskip 1999& TE$_{221}$ & Q,  V  & Q, V   & fair & 2001& HZ$_{32}$  & $2.308$ & $0.1540$ & $5.642$ & $16.5$ & $15.0$ \\\\ 2000& LU$_{15}$  & V,  Q  & V, Q   & good & 1992& WJ$_{35}$  & $2.313$ & $0.0701$ & $5.742$ & $16.1$ & $13.7$ \\\\ 2001& XH$_{209}$ & A,  Sa & A, S   & good & 2004& PH         & $2.401$ & $0.2150$ & $3.638$ & $15.6$ & $16.4$ \\\\ \\tableline\\tableline \\end{tabular} \\end{center} \\caption{Asteroid pair member data for objects appearing in SDSS DR7 MOC4.  We determined SMASS taxonomy for $Ast_1$ using $ugriz$ photometry from SDSS DR7 MOC4 as described in the text.  The first and second ranked SMASS classes (\\S \\ref{s.classid}) are provided along with the corresponding Tholen classes and an assessment of that match's quality (degree of identification certainty). Semi-major axis (a), eccentricity (e), inclination (i) and absolute magnitudes ($H_1$,$H_2$) of both asteroids are from \\citet{bib.vok08,bib.pra09b}.  Members shown in bold constitute the only complete pair.} \\label{t.Asteroid Pairs} \\end{table}  None of the members of 36 non-family asteroid pairs (\\citet{bib.pra09b, bib.vok08}, \\S\\ref{s.intro} describes paired asteroids' discovery based on similar orbits.) are available in either the SMASSI or SMASSII \\citep{bib.bus02b} spectra databases or the Eight Color Asteroid Survey (ECAS) \\citep{bib.tho89}.  This is unsurprising considering that the pair members are considerably smaller than the typical asteroid in those surveys.  However, the 19 pair members identified in Table~\\ref{t.Asteroid Pairs} were found in the SDSS MOC4 \\citep{bib.par08} from which we obtained the five-filter solar-corrected $ugriz$ photometry in Table~\\ref{t.ugriz}.  Most of the pair members are located in the inner main belt in a region dominated by S-complex asteroids.  Their SDSS photometry indicate that they belong to various taxonomic types typical of the inner belt including L, S and V classes.  Our formal identification of the pair member's taxonomy using the methods described in \\S\\ref{s.classid} are provided in Table~\\ref{t.Asteroid Pairs} which shows that we have identified 14 of the 19 pair members with the S-complex in which we also include L-class and Q-class.  We will examine the colors of these asteroids in the context of our space weathering model later in this section.  We also identified three X-class, one V-class, and one A-class asteroid.  The taxonomic variety of the pair members is also represented in Figures \\ref{f.pairspc1pc2} and \\ref{f.orbitaldistn} which shows that their $PC_1$ color distribution is narrower than the full span of SMASS classes while the distribution of their $PC_2$ values extends beyond the SMASS class range.  Several pair members have $PC_2<-0.4$. Since $PC_2$ corresponds roughly to a spectrum's curvature this indicates an unusually convex spectrum.  A couple pair members lie close to the V-class region while only one lies in the C-complex.  We believe that the dearth of C-complex objects is an observational artifact because pair members tend to be small and would be difficult to detect with the low albedo of C-complex members in the outer belt.  The assignation of 1999~TE$_{221}$ to the Q class is important to our space weathering analysis since it has been suggested \\citep{bib.bin04} that Q-class objects are actually very young, essentially unweathered, S-complex asteroids.  Thus, we assume that it is a particularly young member of the S-complex with a deep $1 \\mu$m band as has been predicted for young S-complex asteroids.  However, in osculating element space it is located close to 2000~LU$_{15}$, another member of one of our asteroid pairs from table \\ref{t.Asteroid Pairs} that we assigned to the V-class.  Both the asteroids lie close to the edge of the Vesta family region as shown in fig. \\ref{f.pairspc1pc2} (again, in osculating elements). This opens the possibility that 1999~TE$_{221}$ could be a Vestoid with a slightly shallower $1 \\mu$m absorption band.  To confirm our Q-class identification for 1999~TE$_{221}$ and as an additional check on our type-identifications we examined whether the pair members are of taxonomic types typical of their orbit element phase space region.  To do so we identified each pair member's five nearest orbit element neighbors (using the $D$-criterion of \\citet{bib.nes05}) in the set of 1175 objects from SMASS that also have osculating orbital elements in \\citet{bib.ast08}.  We found that 17 out of 19 pair members match their nearest neighbor's complex suggesting that our identification methods identify the right complex $\\sim 90$\\% of the time and that the pair members are representative of the composition of the main belt region in which they are located.  As mentioned earlier, this test was particularly important for the cases of 1999~TE$_{221}$ and 2000~LU$_{15}$ as both lie on the periphery of the Vesta family region.  We identify the $1^{st}/2^{nd}$ most likely types for these two objects as Q/V and V/Q respectively. 1999~TE$_{221}$'s five nearest neighbors include four in the S-complex with one being a Sq and none in the V-class.  On the other hand, 2000~LU$_{15}$ has three V-class neighbors.  This supports our ability to reliably distinguish Q from V.  Remember that we place the Q-class within the S-complex and, since 1999~TE$_{221}$ is by far the bluest member of the S-complex pair members, its inclusion in our analysis could have a substantial impact on the mean $PC_1$ of the pair members and on our measurement of the space weathering rate of S-complex asteroids.  The effect of including or excluding 1999~TE$_{221}$ in our analysis is described later.  Our interest in and utilization of the asteroid pairs for the purpose of measuring young asteroid surface ages assumes that the pair members are genetically related and fissioned by some as yet undefined process $<0.5$ My ago \\citep{bib.vok08}.  If the members of a pair are genetically related asteroids then our expectation is that they will display nearly identical spectra.  Only one complete pair (2000~NZ$_{10}$ and 2002~AL$_{80}$, see Table \\ref{t.Asteroid Pairs}) was identified among the 19 pair members available in the SDSS MOC4. Figure~\\ref{f.matchpair} shows that the colors of the two objects match and therefore supports a genetic origin of the pair.  Using the taxonomic identification methods of \\S\\ref{s.classid} we find that 2000 NZ$_{10}$ is SMASS L-class and 2002 AL$_{80}$ is Sl-class - adjacent classes in $PC_2$ vs. $PC_1$ color space as shown in Figure~\\ref{f.pairspc1pc2}.  Thus, this one line of evidence suggests that asteroid pairs are genetically related.  \\begin{figure}[!ht]\\small \\centerline{\\includegraphics[width=4.5 in,angle=90]{matchpair.ps}} \\caption\\small{SDSS filter photometry for both members of the only complete dynamical pair in the MOC4.  The central wavelength for each data point corresponds to the band centers for the $ugriz$ filters except for a small horizontal offset for clarity while their width represents the band pass.  The data is normalized such that a straight line between the $g$ and $r$ data points passes through unity at 0.55~$\\mu$m.} \\label{f.matchpair} \\end{figure}  Having established the taxonomic composition of the asteroid pairs and their likely genetic relationship we would like to examine their taxonomic-orbit distribution --- does the pair member taxonomy match that of their neighbors in orbit element space?  The answer to this question could shed light on the relative internal strengths of the different types or provide information on the mechanism for asteroid pair creation.  \\ie if C-complex asteroids split into pairs more frequently it could imply that they are weaker than other types or that the pair formation mechanism acts more efficiently on them. Unfortunately, answering this question is beyond the scope of this work and we leave it to the future.  Instead, we make a couple simple observations on the pair's orbit element distribution.  \\begin{figure}[!h]\\small \\centerline{\\includegraphics[width=5.0in,angle=90]{orbitaldistn.ps}} \\caption\\small{Osculating $\\sin$(inclination) versus semi-major axis for 36 pair members (large colored points) superimposed on the proper element distribution for main belt asteroids identified in the SDSS MOC4 (black dots).  For semi-major axis $<2.1$~AU we show osculating elements for main belt asteroids from Astorb \\citep{bib.orb08}.  The 18 distinct pair members with SDSS photometry were identified as the following types: violet squares are X-complex, blue triangles are S-complex, green diamonds are L-complex, the orange $\\times$ is V-class, the fuchsia $\\times$ is Q-class, and the fuchsia $+$ is A-class.  The 18 red asterisks represent pairs for which neither member is present in the SDSS MOC4.  The $\\nu_6$ and 3:1 resonances are shown for orientation along with the Hungaria family region.} \\label{f.orbitaldistn} \\end{figure}  The axis-inclination structure of the main belt and the pairs is shown in Figure~\\ref{f.orbitaldistn} revealing that the 36 pairs are distributed in two clumps; a high inclination clump inside 2.0~AU within the Hungaria family region, a group dynamically protected from perturbations by Mars via their high inclination, and a clump on the inner edge of the main belt with 2.1~AU$\\lae a \\lae$~2.4~AU. There are also two outliers in the middle belt with semi-major axes in the range 2.65~AU$<a<$~2.75~AU.  That most of the pairs are located in the innermost main belt is almost certainly an observational selection effect --- asteroid pairs are composed of small asteroids that are only visible when they are located close to the Earth \\ie in the inner main belt.  The Hungaria clump includes six pairs of which three have SDSS photometry that we identify as (SMASS) X-complex asteroids --- consistent with \\citet{bib.gra82}'s claim that roughly 70\\% of the asteroids in the Hungaria region are Tholen E or R class (the Tholen E class is contained within the SMASS X-complex).  Thus, our identification of three SDSS X-complex members in the region is unsurprising and provides further support for our taxonomic classification techniques.   \\subsection{Space weathering on sub-My clusters and asteroid pairs} \\label{ss.comparemodel}  Figure~\\ref{f.spaceWeathering} combines our previous work \\citep{bib.wil08} with the new color-age data in this work for S-complex asteroids in three sub-My clusters and eleven sub-My asteroid pairs.  The mean $PC_1$ for the three sub-My clusters is $0.36 \\pm 0.07$ --- within 1$\\sigma$ of \\citet{bib.wil08}'s predicted color of $PC_1 = 0.31 \\pm 0.04$ for the clusters's mean age of $440 \\pm 60$ kyr.  The good agreement with the prediction lends confidence to the space weathering model which now agrees with the cluster color-age data over five orders of magnitude in age in the decades from $10^5 - 10^{10}$ years.   \\begin{figure}[!ht]\\small \\centerline{\\includegraphics[width=5in,angle=90]{submyr.measured.ps}} \\caption\\small{$PC_1$ color and dynamically determined ages for S-complex asteroid families adapted from \\citet{bib.wil08}.  (The corresponding spectral slope is shown on the right.)  The dashed curve represents their space weathering model (eq.~\\ref{eq.jedpc1}) extrapolated to the sub-My region.  Three S-complex sub-My clusters ((1270)~Datura, (21509)~Lucascavin, (16598) 1992 YC2) are shown individually and with their mean value indicated by the red `Sub-My cluster average' data point.  The solid curve represents the dual weathering/gardening model fit to the family data including the sub-My cluster point.  The `SDSS pairs' point represents the average color of 12 unique S-complex sub-My pairs found in SDSS DR7 MOC4. Errors are standard errors on the mean except for (16598) 1992 YC2, a single object, for which we provide the measurement error.} \\label{f.spaceWeathering} \\end{figure}  However, the weighted mean $PC_1 = 0.49\\pm0.03$ color (the error is weighted error on the mean) for the S-complex young pairs is over $5\\sigma$ redder than predicted by \\citet{bib.wil08}.  Excluding the Q-type pair member 1999~TE$_{221}$ discussed in \\S\\ref{ss.pairMemberTaxonomy} increases the mean $PC_1$ to $0.54 \\pm 0.04$ and the discrepancy to over $7\\sigma$.  The very young asteroid pairs clearly do not follow the space weathering function proposed by \\citet{bib.wil08}.  The disparity may be explained in a number of ways: 1) the asteroid pairs do not represent a recent breakup of a parent body and are not genetically related or 2) the asteroid pairs are the result of a recent breakup but with only partial resurfacing which did not `reset' the space weathering clock or 3) the space weathering model of \\citet{bib.wil08} is either too simplistic or wrong.  We examine each of these scenarios in turn: \\begin{enumerate} \\item We consider it unlikely that the asteroid pairs are not genetically related for two reasons.  First, the pairs were specifically selected \\citep{bib.vok08} because they are statistically likely to be related asteroids.  Furthermore, the colors of the pair for which both members exist in the SDSS MOC4 agree extremely well (see Fig. \\ref{f.matchpair}).  While asteroids that inhabit the same region of the main belt often have similar colors the scale of agreement in both the orbit and colors argues persuasively for a genetic link between the pair members. \\item Some pair formation scenarios may not be as violent as the formation of large asteroid families through the catastrophic disruption of a parent body and may not reset the entire surface to zero age.  \\citet{bib.vok08} cite three possible methods of forming pairs: catastrophic collision followed by fragment reaccumulation into binary orbits \\citep[\\eg][]{bib.dur04,bib.nes06c}, YORP induced rubble pile spin-up leading to calving of the secondary \\citep[\\eg][]{bib.wal08}, or YORP induced angular acceleration of contact binaries leading to their separation \\citep[\\eg][]{bib.mer02,bib.pra09a,bib.dur04}.  The first scenario will `reset' the age of the entire surface of both the primary and secondary.  The third mechanism might leave large portions of the surface undisturbed since the only portion that is necessarily exposed is the binary contact region.  It is unclear how much of the surface would be affected by the second scenario.  Therefore there is a clear distinction between the possible formation mechanisms of pairs and larger groupings.  Asteroid families and clusters with more than two members form only by the first catastrophic collisional scenario while pairs can also form by the two variations on YORP spin-up.  This yields at least one pair formation scenario that could leave some of the members' surface undisturbed.  In this case the globally averaged surface color would correspond to a misleading age somewhere between fresh surface and the age of the parent body's surface.  Taking the color of the pairs at face value and interpreting it in the framework of the \\citet{bib.wil08} space weathering model indicates an average surface age of $\\sim 440$~My --- over 1000$\\times$ older than their dynamical age.  If we assume that the original parent body's surface was reddened to saturation prior to separation, and taking the dynamical ages of $<500$~kyr for the pairs at face value such that freshly exposed surface is essentially unweathered, then $\\sim 64$\\% of the surface must be disturbed in the pair separation process.  Considering that it is unlikely that the surface was fully weathered prior to separation allows us to set an upper limit on the fraction of disturbed surface at $\\la 64$\\%.  (Excluding the Q-type object 1999~TE$_{221}$ only changes the upper limit to $\\la 48$\\%.)  A gentle binary separation due to slow YORP spinup may be consistent with this scenario. Indeed, \\citet{bib.pra09a} also provide evidence from reconstruction of the initial configuration of the 6070-54827 pair and rotation rate observations that the non-family pair formation process is a gentle event. If we envision a bi-lobed asteroid gradually accelerating in angular velocity and finally fissioning at the neck that joins the lobes then it seems reasonable that the portion of disturbed surface would be $\\la64$\\%.  In the binary/pair formation mechanism proposed by \\citet{bib.wal08,bib.wal06} an asteroid's polar surface material migrates to the equator as the object's rotation rate increases and eventually flies into orbit around the parent body where it reaccumulates into a satellite.  The primary and satellite eventually separate and their orbits evolve dynamically.  Our impression is that this model would generate fresh (blue) surface on both the primary and satellite in conflict with our observation of reddish surfaces on the asteroid pairs.  However, it is not difficult to envision a slow migration process that allows material to weather on the primary's surface before being shed and accumulating into a secondary object. \\item It is possible that the space weathering model of \\citet{bib.wil08} that built upon the earlier work of \\citet{bib.nes05} and \\citet{bib.jed04} is simply wrong; that the apparent change in color of S-complex asteroids with age is a statistical fluke or due to some other underlying effect (though obvious possibilities were considered in detail in the early works). However, \\citet{bib.par08} confirm the weathering effect in an independent updated analysis of the SDSS DR7 MOC4 data.  We consider it more likely that the space weathering mechanism is more complicated than simply affecting the average spectral slope (or $PC_1$) as a function of time.  For instance, it is well known that space weathering affects not only the slope of the spectrum but also the depths of the 1~$\\mu$m and ultraviolet absorption bands and the surface albedo.  We consider it not only possible but likely that these effects occur during space weathering at different rates.  The apparent redness of the young asteroid pair members relative to the expectation of the simplistic space weathering model could indicate that a `fast' space weathering process takes place in $\\la 10^5$ years.  However, this scenario requires a rather contrived sequence of events.  Consider the three color change processes: decreasing depth of the ultraviolet band shortward of 0.4~$\\mu$m, decreasing depth of the 1~$\\mu$m absorption band and continuum reddening between these two bands.  The first process is the only one that produces a bluer color.  Therefore, accounting for the anomalous redness of the pair members requires the unlikely scenario that one of the latter two processes dominate on short time scales which is then belatedly overwhelmed by the first process which is then finally dominated by the third. \\end{enumerate} Given the disagreement between the asteroid pair color-age and the \\citet{bib.wil08} space weathering model, and considering our enumerated arguments above, we continue with our analysis under the assumption that we can ignore the colors of the asteroid pairs in this new determination of the space weathering and gardening rates.  First, we fit all the S-complex color data including the sub-Myr clusters but not the asteroid pairs to the `old' space weathering function of eq. \\ref{eq.jedpc1}.  Considering the good agreement between the predicted \\citep{bib.wil08} and observed colors of the sub-Myr cluster members it is unsurprising that the fit including the new data matches the previous fit in all four parameters to within 1-$\\sigma$ with $PC_1(0)=0.34 \\pm 0.02$, $\\Delta PC_1=0.28 \\pm 0.05$, $\\tau = 700 \\pm 270$~My, $\\alpha = 0.58 \\pm 0.17$.  However, as we observed in \\citet{bib.wil08}, fitting the color-age data to the form of eq. \\ref{eq.jedpc1} suffers from multiple and wide minima in the fit-parameter space.  Furthermore, the function does not explicitly separate the weathering and gardening effects.  On the other hand, we found that fitting the same color-age data to our new function that incorporates both space weathering and gardening (eq. \\ref{eq.garden}) is better behaved because the solution space does not show multiple local $\\chi^2$ minima.  The best fit (lowest $\\chi^2$) including the sub-My clusters (but, again, not the asteroid pair colors) yields $PC_1(0)=0.37 \\pm 0.01$, $\\Delta PC_1=0.33 \\pm 0.06$, $\\tau_w = 960 \\pm 160$~My, $\\tau_g = 2000 \\pm 290$~My as shown by the solid curve in Figure~ \\ref{f.spaceWeathering}.  Note that it is not correct to compare the new $\\tau_w$ or $\\tau_g$ to the old $\\tau$ value because of the complicating and non-physical use of $\\alpha$ in the generalizing exponent in the old functional form. In essense, the `old' value was an effective weathering time that combined the effects of both regolith gardening and weathering.  Since the effective weathering time can be shown to be equivalent to the time corresponding to the inflection point on the new weathering function, Figure~\\ref{f.spaceWeathering} shows that the two models are in good agreement.  The old and new $\\Delta PC_1$ are not strictly comparable either. Formerly in the weathering only case the entire surface would eventually reach the color $PC_1 = PC_0 + \\Delta PC$.  The weathering/gardening case will produce a lower equilibrium value with $PC_1 < PC_0 + \\Delta PC$.  The new space weathering-only time frame of $\\sim$1~Gy is consistent with the `slow weathering' measured in lab-based measurements as discussed in the introductions. The `slow' school includes our result of $\\sim$1~Gy based on space observations of S-complex families, \\citet{bib.sas01}'s equivalent value of 700~My at 2.6~AU based on pulsed-laser irradiated silicate pellets, and \\citet{bib.pie00}'s estimate of $100-800$ My for lunar surfaces based on crater counts and spectra and color-terrain correlations on (243) Ida \\citep{bib.vev96}. But these slow space weathering results disagree with the `fast weathering' school that includes lab-based ion irradiation results from \\citet{bib.ver09} ($\\lae$~1~My) and \\citet{bib.loe09}'s ($\\sim$ 5 kyr), and \\cite{bib.tak08}'s $<$450~kyr time scale based on a shallow 1 $\\mu$m absorption band observed on (1270) Datura.  We do not have an explanation for the discrepancy between the fast and slow weathering results.   \\subsection{Gardening Time} \\label{s.gdntime}  In \\S\\ref{s.gardenimpact} we calculated the gardening time for asteroid regolith as a function of target diameter as shown in Figure. \\ref{f.gardentaulog}.  Over a broad range of asteroid sizes the gardening time scale is $200-300$ My years.  This result is in dramatic disagreement with a similar calculation by \\citet{bib.mel09} that was based on a collisional cratering model by \\citet{bib.gil02} and yielded a time scale of 2000 years for resurfacing Trojan asteroids.  It is difficult to reconcile the two results that differ by five orders of magnitude.  Two factors that we are aware of that explain some of the discrepancy are 1) that they used a Trojan impact probability double that of the main belt and 2) the slope of the Trojan size-frequency distribution was assumed to be $-3$ whereas the data we used had slope of $\\sim -2.20$ appropriate to the main belt \\citep{bib.bot05b}.  But these two differences only account for a fraction of the difference between our results.  Figure \\ref{f.gardentaulog} also shows that our measured gardening rate from the S-complex asteroid color-age relationship is an order of magnitude different from our calculated rate based on impacts.  Is it possible to reconcile this difference?  There is considerable uncertainty in the various terms involved in calculating the impact gardening time from eq. \\ref{eq.tauofR} but we have done our best to select the best contemporary values in each case.  While the gray region on the figure represents the formal 1-$\\sigma$ error on the calculation based on the reported errors on each input quantity there is considerable unreported systematic error in the calculation associated with its sensitivity to the input parameters and functions.  In particular, we examined the sensitivity of the calculation to the: \\begin{itemize} \\item asteroid size frequency distribution, $N(D)$ \\item specific shattering energy function, $Q_S^*$, that determines the largest non-disruptive impactor (we used the \\citet{bib.mel97} function as a comparison), \\item impact speed and, \\item $D_{min}$, the smallest impactor that creates ejecta \\end{itemize}  In each case we varied the input parameter or function over a range of of 2 to 4 in each direction.  In most cases the gardening rate for large target objects is only slightly affected.  The rate for objects $\\sim$1~km diameter changed by a factor of about two or slightly more. The two most important factors in determining the gardening rate are 1) $D_{min}$ and 2) the amount of area covered by crater ejecta. Increasing $D_{min}$ or decreasing the ejecta area coverage both work to increase the gardening timescale.  Leaving all other parameters at their nominal values, increasing $D_{min}$ to $\\sim$30~m is sufficient to increase the gardening time to our measured value of $\\sim$2~Gy.  A similar result is obtained by decreasing the diameter of the ejecta field by a factor of two (the affected area by a factor of four).  In a related observation, \\citet{bib.cha05} and others note the unexpected paucity of craters $<$200 m diameter on (433)~Eros.  Taken at face value the observation would imply a deficit of impactors $\\la 20$~m in diameter but \\citet{bib.obr09} has shown that no such deficit can exist because it would in turn generate a `wave' in the observed size-frequency distribution for larger asteroids that is not observed.  Thus, we have identified two independent problems --- the mismatched measured color-age and impact-calculated gardening times and the crater deficiency on (433)~Eros --- that can both be resolved if impactors $\\la 20$ m do not leave a crater record.  While \\citet{bib.ric04,bib.ric05} have proposed that seismic shaking can erase the small craters the mechanism can only explain the crater shortage not the gardening time mismatch.  This is because any process that increases the gardening rate (\\ie seismic shaking) shortens the calculated impact gardening time and worsens the mismatch with measured color-age gardening rate.  If we believe that there is no deficit of impactors of $\\la 20$~m diameter and that there is a deficit of craters and regolith gardening caused by impactors in the same size range then it must be the case that the impacts take place without creating craters in the `normal' manner.  Perhaps the surfaces of small asteroids `absorb' the small impactors in an inelastic collision without creating a crater an order of magnitude larger or ejecting significant material.  Like throwing a small stone into a bee's empty honeycomb."
    },
    "1004/1004.2708_arXiv.txt": {
        "abstract": "Different theoretical methodologies lead to order-of-magnitude variations in predicted galaxy-galaxy merger rates. We examine how this arises, and quantify the dominant uncertainties. Modeling of dark matter and galaxy inspiral/merger times contribute factor $\\sim2$ uncertainties. Different estimates of the halo-halo merger rate, the subhalo ``destruction'' rate, and the halo merger rate with some dynamical friction time delay for galaxy-galaxy mergers, agree to within this factor $\\sim2$, provided proper care is taken to define mergers consistently. There are some caveats: if halo/subhalo masses are not appropriately defined the major merger rate can be dramatically suppressed, and in models with ``orphan'' galaxies and under-resolved subhalos the merger timescale can be severely over-estimated. The dominant differences in galaxy-galaxy merger rates between models owe to the treatment of the baryonic physics. Cosmological hydrodynamic simulations without strong feedback and some older semi-analytic models, with known discrepancies in mass functions, can be biased by large factors ($\\sim5$) in predicted merger rates. However, provided that models yield a reasonable match to the total galaxy mass function, the differences in properties of {\\em central} galaxies are sufficiently small to alone contribute small (factor $\\sim1.5$) additional systematics to merger rate predictions. But variations in the baryonic physics of {\\em satellite} galaxies in models can also have a dramatic effect on merger rates. The well-known problem of satellite ``over-quenching'' in most current semi-analytic models (SAMs) -- whereby SAM satellite populations are too efficiently stripped of their gas -- could lead to order of magnitude under-estimates of merger rates for low-mass, gas-rich galaxies. Models in which the masses of satellites are fixed by observations (or SAMs adjusted to resolve this ``over-quenching'') tend to predict higher merger rates, but with factor $\\sim2$ uncertainties stemming from the uncertainty in those observations. The choice of mass used to defined ``major'' and ``minor'' mergers also matters: stellar-stellar major mergers can be more or less abundant than halo-halo major mergers by an order of magnitude. At low masses, most true major mergers (mass ratio defined in terms of their baryonic or dynamical mass) will appear to be minor mergers in their stellar mass ratio - observations and models using just stellar criteria could underestimate major merger rates by factors $\\sim3-5$. We discuss the uncertainties in relating any merger rate to spheroid formation (in observations or theory): in order to achieve better than factor $\\sim3$ accuracy, it is necessary to account for the distribution of merger orbital parameters, gas fractions, and the full efficiency of merger-induced effects as a function of mass ratio. ",
        "introduction": "\\label{sec:intro}  In the now established $\\Lambda$CDM cosmology, structure grows hierarchically \\citep[e.g.][]{whiterees78}, making mergers an inescapable element in galaxy formation. The galaxy-galaxy merger rate, as a function of properties such as galaxy mass, redshift, and mass ratio, is a quantity of fundamental interest. It is critical for informing models of the growth and assembly of galaxies, the distribution of bulge and spheroid mass, properties of disks such as their thickness, morphologies, and flaring, and the fueling and growth of the most luminous starbursts and infrared systems as well as massive BHs and quasars.  The last few years have seen the emergence of a ``concordance'' precision $\\Lambda$CDM cosmology \\citep[see e.g.][and references therein]{komatsu:wmap5}, with remarkable convergence between different probes of structure formation. Meanwhile, cosmological dark-matter simulations have developed the ability to track large populations of dark matter halos and substructure within halos over cosmological volumes and timescales. Different calculations now yield results for e.g.\\ the dark matter halo mass function that agree at the $\\sim5-10\\%$ level \\citep[see e.g.][and references therein]{reed:halo.mfs}.  Despite these advances, however, there has yet to be a similar convergence in theoretical predictions of the {\\it galaxy} merger rate. Depending on the models used to convert dark matter halo growth histories into galaxy growth histories, different theoretical predictors of the merger rate in the same cosmology can differ by more than an order of magnitude \\citep[see e.g.\\ the comparison in][]{jogee:merger.density.08.conf.proc, bertone:merger.rate.vs.stellar.wind.model,lopezsanjuan:merger.fraction.to.z1}. And the differences cannot be trivially attributed to e.g.\\ the known fact that models predict somewhat different absolute galaxy masses and abundances -- in terms of mergers {\\em per galaxy}, there is similar variation.  Some of this may owe to the fact that dark matter merger rates are -- unlike dark matter mass functions -- derivative (instantaneous) quantities and so are more sensitive to choices of definition and methodology in simulations. But much of the difference owes to the fact that the mapping from dark matter merger rates to galaxy merger rates is non-trivial, and requires some relation between the history of a given dark matter halo and a galaxy inside that halo. In models that attempt to predict galaxy formation in an {\\em a priori} sense, from physical prescriptions for gas cooling, star formation, and feedback from stars and accreting BHs, the resulting predicted halo to galaxy ``map'' and corresponding merger rates can be very sensitive to the input prescriptions. It has been shown that even in otherwise identical models, small changes in the treatment of gas cooling or stellar wind physics can lead to order-of-magnitude changes in the predicted major merger rates: compare the merger rates in \\citet{bower:sam} and \\citet{font:durham.sam.update} (our \\S~\\ref{sec:compare.hods:satellites:resolution}), or those in \\citet{delucia:sam} and \\citet{bertone:stellar.wind.recycling.sam} \\citep[presented in][]{bertone:merger.rate.vs.stellar.wind.model}.  These quantitative differences in the number of mergers per galaxy have led to qualitatively different conclusions regarding galaxy formation. Most models, especially those based on empirical halo occupation constraints or direct cosmological hydrodynamic simulation, have found that there are too many mergers in low mass systems to explain the survival and prevalence of disks in the simplest scenario \\citep{granato:sam,somerville:sam,somerville:new.sam, koda:disk.survival.prescriptions,stewart:merger.rates,khochfarsilk:new.sam.dry.mergers, hopkins:disk.survival.cosmo,stewart:disk.survival.vs.mergerrates, sommerlarsen03:disk.sne.fb,abadi03:disk.structure,governato:disk.formation, robertson:cosmological.disk.formation, okamoto:feedback.vs.disk.morphology,scannapieco:fb.disk.sims, kazantzidis:mw.merger.hist.sim,kazantzidis:thin.disk.thickening, purcell:minor.merger.thindisk.destruction}.  More directly, {\\em observed} merger rates find that $\\sim5-10\\%$ of low or intermediate-mass galaxies ($M_{\\ast}\\lesssim10^{10}\\,\\msun$ at redshifts $z\\sim0.2-1.2$) are in mergers -- strongly morphologically disturbed or in ``major'' similar-mass pairs about to merge (at small scales and small relative velocities) \\citep{bridge:merger.fractions,bridge:merger.fraction.new,kartaltepe:pair.fractions, conselice:mgr.pairs.updated, jogee:merger.density.08,lotz:merger.fraction,lin:mergers.by.type, robaina:mgrs.make.es.sufficient}. This is equal to or higher than the predicted rates from the models above. Even without correcting for the short expected duty cycles/lifetimes of this merger phase \\citep[see e.g.][]{ lotz:merger.selection,conselice:mgr.pairs.updated}, these fractions are {\\em already} higher than the observed fraction of bulge-dominated galaxies at these masses \\citep{allen:bulge-disk,% benson:disk.spheroid.mfs, dominquez:bulge.disk.sample,weinzirl:b.t.dist}. Likewise the number density of IR luminous systems apparently in mergers at high redshifts is so high as to yield tension with the relic bulge mass in the Universe \\citep[see e.g.][]{younger:highz.smgs, genel:smg.numden.vs.mergers,tacconi:smg.mgr.lifetime.to.quiescent, hopkins:merger.lfs,cimatti:highz.compact.gal.smgs}. Typically, some additional physics (often some form of strong stellar feedback) must be invoked to preserve high gas fractions and suppress the efficiency of bulge formation. Indeed, this is {\\em always} the case in hydrodynamic simulations of galaxy formation. This has led to the development of various feedback models \\citep[e.g.][and references therein]{sommerlarsen03:disk.sne.fb,abadi03:disk.structure, governato:disk.formation,okamoto:feedback.vs.disk.morphology,scannapieco:fb.disk.sims} and emphasis on gas-richness as a stabilizing factor \\citep{robertson:disk.formation,hopkins:disk.survival, stewart:disk.survival.vs.mergerrates,moster:2010.thin.disk.heating.vs.gas} that allows disks to survive such mergers. However, some semi-analytic calculations have claimed the opposite: that there are not enough mergers at low masses to explain even the small bulges observed at these masses \\citep{parry:sam.merger.vs.morph}.   On the other hand, at $\\gtrsim L_{\\ast}$, most empirical and semi-analytic models predict similar numbers of mergers and agree that massive bulges are primarily formed in major mergers \\citep[see e.g.][and references therein]{hopkins:merger.rates, parry:sam.merger.vs.morph,cattaneo:new.sam.w.mergerfx}. But some simulations have claimed that there are more minor and fewer major mergers at these high masses \\citep{naab:etg.formation}.  As observations of e.g.\\ pair and morphologically disturbed fractions improve, it is rapidly becoming possible to empirically constrain merger rates and fractions at a level better than the apparent order-of-magnitude scatter in predictions \\citep[see e.g.][]{kartaltepe:pair.fractions,lin:mergers.by.type,domingue:2mass.merger.mf, bundy:merger.fraction.new,conselice:mgr.pairs.updated, robaina:2009.sf.fraction.from.mergers, robaina:mgrs.make.es.sufficient,comerford:2009.bh.binaries, lotz:merger.fraction,bridge:merger.fractions,lopezsanjuan:mgr.rate.pairs}. However, the conversion of such observations to merger rates still requires knowledge of appropriate timescales and contamination from projection effects and/or minor mergers -- all of which mean that interpretation of the observations is somewhat dependent on how well or poorly theoretical predictions agree on basic quantities. Moreover, an observed pair or morphologically disturbed fraction is fundamentally {\\em not} invertible at much better than the factor of $\\sim2$ level, as disturbances and dynamics are non-unique and depend on other quantities such as galaxy gas content; as such, they must be forward-modeled.   Fundamentally, the galaxy-galaxy merger rate depends on two basic quantities: the dark matter halo merger rate and the manner in which galaxies populate halos. However, models arrive at these quantities by various methods. In this paper, we attempt to examine in detail different theoretical models and identify the sources of apparently large variations in galaxy merger rates between the models. We are interested both in defining the ``error budget'' of theoretical predictions, as well as in identifying the definitions, physics, and prescriptions that lead to substantial differences.  In the follow sections we show that combining empirical halo occupation constraints with dark matter merger histories from high-resolution simulations leads to a relatively small (factor $\\sim2$) systematic uncertainty in merger rates (\\S~\\ref{sec:empirical}). We then compare various different substructure based calculations for subhalo-subhalo mergers and/or merger time delays (\\S~\\ref{sec:compare.dark}). We highlight several caveats and potential problems in adopting these merger rates: without proper care, it is possible to obtain apparently order-of-magnitude different dark matter merger rates from the same simulation or analytic calculation.  We reveal that most of the differences between model predictions owe to the treatment of baryons in galaxy formation (\\S~\\ref{sec:compare.hods}). Specifically, we show that cosmological hydrodynamic merger simulations lacking feedback do not reproduce the observed $M_{\\rm gal}(M_{\\rm halo})$ relation (not just the normalization of this relation, but also its shape) and so do not map robustly between halo-halo and galaxy-galaxy mergers (\\S~\\ref{sec:compare.hods:central}). Semi-analytic models, on the other hand, perform well for central galaxies, but commonly have well-known problems reproducing observed properties of satellite galaxies; this accounts for most of the discrepancies both between different semi-analytic models and between those models and empirical approaches/calculations (\\S~\\ref{sec:compare.hods:satellites}). In \\S~\\ref{sec:definitions}, we discuss the importance of consistent and appropriate mass ratio definitions, and the variations between different definitions, as well as the range in the effects of a given merger assuming changes in galaxy gas fractions, orbital parameters, and structural properties. We summarize our conclusions and discuss future tests and improvements in \\S~\\ref{sec:discuss}.  Except where otherwise specified, we adopt a ``concordance'' cosmology with $(\\Omega_{\\rm M},\\,\\Omega_{\\Lambda},\\,h,\\,\\sigma_{8},\\,n_{s})$= $(0.3,\\,0.7,\\,0.7,\\,0.9,\\,1.0)$ and a \\citet{chabrier:imf} IMF, and appropriately normalize all observations and models shown. The choice of IMF systematically shifts the normalization of stellar masses herein, but does not otherwise change our comparisons.  Throughout, we use the notation $M_{\\rm gal}$ to denote the baryonic (stellar+cold gas) mass of galaxies; the stellar, cold gas, and dark matter halo masses are denoted $M_{\\ast}$, $M_{\\rm gas}$, and $M_{\\rm halo}$, respectively. When we refer to merger mass ratios, we use the same subscripts to denote the relevant masses used to define a mass ratio (e.g.\\ $\\mu_{\\rm gal} = M_{\\rm gal,\\, 2}/M_{\\rm gal,\\,1}$), always taken such that $0<\\mu<1$ ($M_{\\rm gal,\\,1}>M_{\\rm gal,\\,2}$).   \\breaker ",
        "conclusions": "\\label{sec:discuss}  A simple comparison of different predictions of the galaxy-galaxy merger rate from e.g.\\ different semi-analytic models and simulations demonstrate that the predictions vary by an order of magnitude (Figure~\\ref{fig:mgr.rate.vs.sams}). We have attempted to survey the sources of uncertainty and systematic differences between various theoretical attempts to predict the galaxy merger rate, to identify what drives these differences and address how progress can be made.    \\subsection{Dark Matter and Dynamics} \\label{sec:discuss:budget:dark}  All models depend similarly on the ``background'' dark matter merger rate. However, with proper caution in adopting definitions, and modern convergence in high-resolution cosmological simulations, the uncertainties in this rate can be reduced to the factor $\\sim2-3$ level. In other words, if how galaxies populate halos is appropriately fixed, there is relatively little uncertainty in the merger rate owing to the dynamics of the dark matter (Figure~\\ref{fig:mgr.rate.vs.model}). The relevant quantities considered here include:  {\\bf Cosmology:} Nominally, changing the cosmology within the statistical uncertainty in modern observational constraints \\citep{komatsu:wmap5} leads to factor $\\sim1.5$ differences in the merger rate. However, most of this owes to changes in the mass function which can be normalized out -- if e.g.\\ the galaxy population is normalized so as to match the observed stellar mass function and large-scale bias, then the resulting differences in the merger rate are negligible compared to the other sources of uncertainty below (Figure~\\ref{fig:mgr.rate.vs.model}).  {\\bf Halo-Halo Merger Rate Determinations:} Again, nominally different halo-halo merger rate determinations differ by factors of $\\sim2$. Some of this stems from e.g.\\ the inherent ambiguity in masses and extents of halos. However, much owes to the definitions adopted when attempting to fit/quantify the instantaneous merger rate ``function.'' If these definitions are properly accounted for, or if halo merger trees are used directly to track galaxies (rather than defining halo merger rates at all), the real uncertainties are small, a factor $\\sim1.5$. Effects of baryons contribute similarly small ($\\sim10-20\\%$) uncertainties in the overall halo-halo merger rate (Figure~\\ref{fig:model.halorates}). However, care is needed with definitions -- defining halo mergers simply in terms of ``instantaneous'' mass ratios, and/or applying cuts which can be useful for constructing {\\em galaxy} merger trees, or using timestep-sensitive mass ratio definitions can lead to an apparent (artificial) suppression of major halo-halo mergers by an order-of-magnitude or more (Figures~\\ref{fig:model.halorates.durham}-\\ref{fig:model.halorates.timestepping}). This requires careful consideration in simulation-based semi-analytic models.  {\\bf Subhalo versus Halo Merger Rates:} Defining the merger rate not when halos first merge into e.g.\\ a larger friends-of-friends group, but when subhalos merge/are destroyed into the central (primary) group subhalo, may be more representative of galaxy-galaxy mergers. Doing so, however, introduces additional uncertainties owing to e.g.\\ how subhalos are identified. It is also important, in this case, that the subhalo ``mass'' (for merger mass ratio purposes) be defined as the ``infall'' or maximum pre-accretion/pre-stripping mass (i.e.\\ maximum mass before the system became a subhalo), otherwise the mass of all systems $\\rightarrow0$ at merger, by definition. Different calculations from various simulations, with different subhalo identification/destruction criteria, yield rates converged to within a factor $\\sim2$ (Figure~\\ref{fig:model.subhalorates}). Using hydrodynamic simulations to ``tag'' mergers (using the galaxies therein as the ideal subhalo/halo ``tracers'') yields consistent results.  {\\bf Merger Timescales:} In models without subhalo resolution or with Press-Schechter merger trees, galaxy-galaxy mergers are often assumed to be ``delayed'' with respect to the halo-halo merger by a timescale given by e.g.\\ the dynamical friction time (this approximates the subhalo evolution). We show that calibrations of such timescales, from high-resolution galaxy merger simulations, are relatively well-converged (Figure~\\ref{fig:tmerger.compare}). Adopting one of these calibrations, we show that the implied merger rate agrees well with that obtained from full tracking of subhalos, with similar factor $\\sim2$ uncertainties (Figure~\\ref{fig:model.subhalorates.tdf}). For {\\em major} mergers, both methods agree well with the simple halo-halo merger rate -- this is because the timescale for such a merger to complete ($\\ll t_{\\rm Hubble}$) is short compared to the time between such mergers ($\\sim t_{\\rm Hubble}$).  However, we show that adopting an artificially long merger timescale, or adopting e.g.\\ the simple Chandrasekhar timescale with a normalization higher than the specific numerical calibrations here, can suppress high-redshift mergers by factors $\\sim5-10$. We also caution that these calibrations are designed for the time from halo-halo merger -- they are {\\em not} constructed for application in ``hybrid'' models, where subhalos are followed in simulations down to some resolution limit, and then a merger time is applied to the residual ``orphan'' based on its instantaneous radius and mass. We show that in such cases, if the system is lost to resolution at large radius $>0.2\\,R_{\\rm vir}$, the calibrated formulae are not self-consistent, and can lead to over-estimates of the total merger timescale by factors of $\\sim2-8$ (Figure~\\ref{fig:tmerger.orphan.issues}).   \\subsection{Baryonic Physics} \\label{sec:discuss:budget:baryons}  Controlling for some of the caveats and definitions above, the dominant uncertainties in predicted merger rates owe to variations in how galaxies populate halos. The differences can be identified with the shape of the assumed halo occupation distribution (HOD) -- essentially, the distribution of galaxy masses in given (sub)halo masses, $M_{\\rm gal}(M_{\\rm halo})$. Some model for this is necessary to translate halo-halo mergers (or halo-subhalo mergers) into galaxy-galaxy mergers. We therefore consider how these physics are modeled and how they lead to variations in the merger rate, in three classes of models with very different approaches towards modeling the galaxy-halo connection:   {\\bf Semi-Empirical (HOD-based) Models:} In semi-empirical models the HOD is adopted explicitly from observational constraints. As such, it is subject to the attendant uncertainties and limited to the dynamic range observed. At low redshifts and galaxy masses within factors of several around $\\sim L_{\\ast}$, the constraints are tight and different methods agree well -- resulting uncertainties in the galaxy-galaxy merger rate (holding halo-halo merger rates {\\em fixed}) are a factor $\\sim1.5$ (Figures~\\ref{fig:mgr.rate.vs.model} \\&\\ \\ref{fig:model.HODs}). The uncertainties grow to a factor $\\sim2$ at the lowest and highest masses, and $z\\sim1-2$. Above $z\\sim2$, the uncertainties grow very rapidly: there are not sufficient observational constraints on the HOD to make strong statements about galaxy merger rates. Predicted merger rates and pair counts from such models agree well with direct observations over a stellar mass range $\\sim10^{9.5}-10^{11.5}\\,\\msun$ and redshifts $z=0-2$ \\citep[see Figures above and][]{stewart:merger.rates, hopkins:merger.rates}. At this level, it appears, there is no tension between $\\Lambda$CDM merger rates and merger/pair counts. To the extent that other models disagree significantly with the observed merger fractions, it should owe to issues in the model baryonic physics leading to a $M_{\\rm gal}(M_{\\rm halo})$ distribution different from that observed.   {\\bf Cosmological Hydrodynamic Simulations:} In hydrodynamic simulations, the $M_{\\rm gal}(M_{\\rm halo})$ distribution is predicted in an {\\em a priori} manner based on the cooling, star formation, and feedback models implemented in the simulation. Unfortunately, it is not yet possible to run large-volume cosmological hydrodynamic simulations (needed to quantify the galaxy-galaxy merger rate) with the spatial and mass resolution and detailed prescriptions for feedback from stars and black holes that it is becoming clear are necessary to form ``realistic'' galaxies. Simulations available that do not include feedback yield poor agreement with the observed HOD and stellar mass function, predicting a relationship closer to the efficient star formation limit ($M_{\\rm gal}=f_{b}\\,M_{\\rm halo}$). Relative to what the merger rates would be, taking the same dynamics and merger locations but re-populating galaxies with masses chosen to fit the observed HOD (stellar mass and clustering), these simulations tend to over-predict merger rates at low masses and high redshifts, and under-predict rates at high masses and low redshifts by factors $\\sim3-5$, as well as over-predicting the relative importance of minor versus major mergers at all masses (Figure~\\ref{fig:model.HODs}). Because it is the {\\em shape} of the $M_{\\rm gal}-M_{\\rm halo}$ relation that is most important, simply re-normalizing predicted masses by a uniform factor will not correct for these effects. Re-normalizing all masses to their ``correct'' masses given some observed HOD is an improvement, but care is still needed, since the incorrect masses and morphologies will affect quantities such as the dynamical friction time.  {\\bf Semi-Analytic Models:} In semi-analytic models, the difficulties and expense of simulations are replaced by use of analytic prescriptions, given some background dark matter population, to predict ultimate galaxy properties. Such models are adjusted to give good agreement with the galaxy stellar mass function (and clustering) at $z=0$; as such, some agreement with the $M_{\\rm gal}(M_{\\rm halo})$ distribution of {\\em central} galaxies (which dominate the stellar mass function at all masses) is implicitly guaranteed. Indeed, we find that adopting the predicted SAM HODs for central galaxies, instead of the empirically determined HOD, yields no systematic difference and scatter within the same factor $\\sim2$ allowed by different observational constraints (Figures~\\ref{fig:model.HODs}). At high redshifts, the uncertainties in both grow. Some SAMs yield growing discrepancies relative to semi-empirical models, directly related to issues such as e.g.\\ the known tendency of SAMs to over-predict the abundance of low-mass galaxies at high redshift, but these are still within the factor $\\sim3-5$ level at $z<3$ (Figure~\\ref{fig:model.HODs.z}).  However, the HOD of {\\em satellite} galaxies in semi-analytic models is extremely sensitive to prescriptions for cooling, stellar feedback, and halo mixing/stripping in satellites. Moreover, satellite masses are not strongly constrained by the stellar mass function, so there is no implicit guarantee/check that these are correct, and more detailed comparison with e.g.\\ observed group catalogs and small-scale clustering must be used to calibrate the models.  In detail, it is well-known that most SAMs have difficulty reproducing the observed properties of satellite galaxy populations: they tend to predict satellite galaxies that are under-massive and over-quenched, relative to observations \\citep[see Figures~\\ref{fig:model.HODs.sat} \\&\\ \\ref{fig:sat.overquenching} and e.g.][and references therein]{ weinmann:obs.hod,weinmann:group.cat.vs.sam, wang:sat.shutdown.slower.vs.sam,kimm:passive.sats.vs.centrals}. This can occur even in models where the small-scale clustering of satellites is over-predicted, relative to observations \\citep[see e.g.][who find both effects, although the clustering discrepancy is probably due to the large value of $\\sigma_{8}$ used in the simulation]{guo:2010.millenium.2.update}. In most of the models, satellites lose their entire halo gas reservoir at the moment they become such -- even in major mergers (and even when the moment of ``becoming'' a satellite occurs at several times the primary virial radius). Moreover, the combination of simple stellar wind feedback and cooling models often leads to the satellites {\\em also} losing almost all of their cold/disk gas reservoir.  As such, at low masses where gas fractions are important, initially major and even equal-mass mergers can easily become minor mergers in the model, potentially suppressing the predicted merger rate by a large factor. Models which yield more ``over-quenched'' satellite populations tend to yield correspondingly smaller merger rates (Figures~\\ref{fig:model.HODs.sat} \\&\\ \\ref{fig:font.vs.bower}). Correcting for these differences, for example by enforcing that satellite galaxies obey a similar HOD to central galaxies, or by adjusting the prescriptions for cooling onto satellite galaxies such that they better reproduce the observed color and star formation rate distributions of satellites, leads to larger merger rates that converge with the merger rates predicted from semi-empirical models (Figure~\\ref{fig:font.vs.bower}). Whether or not this is necessary for matching e.g.\\ close pair counts and small-scale clustering is unclear \\citep[see e.g.\\, the comparison in][]{kitzbichler:mgr.rate.pair.calibration, guo:2010.millenium.2.update}, but it demonstrates an important uncertainty in predictions of merger rates.   \\subsection{Impact of Mergers on Galaxies} \\label{sec:discuss:budget:impact}  Of course, even with perfect knowledge of the merger rate and distribution of mass ratios under some definition, it is not trivial to relate this to either observable or physical quantities such as merger fractions or the amount of bulge mass formed in mergers. We consider how two basic aspects of this relate to uncertainties in merger rates and their consequences:  {\\bf Mass Ratio Definitions:} As outlined in \\citet{stewart:massratio.defn.conf.proc}, a halo-halo major merger is not necessarily a galaxy-galaxy major merger, and vice versa. But there are also several means of defining galaxy-galaxy mass ratio, including e.g.\\ the stellar-stellar mass ratio $\\mu_{\\ast}$ and baryon (stellar mass plus cold gas within the galaxy disk)-baryon mass ratio $\\mu_{\\rm gal}$. At high (low) masses, the merger rate in terms of stellar-stellar mass ratio is enhanced (suppressed) by an order of magnitude relative to the halo-halo merger rate (Figure~\\ref{fig:merger.fx.cautions}). In comparing e.g.\\ model predictions to stellar mass ratio-selected pair samples, clearly the stellar mass ratio is most applicable. However, in comparing to luminosity-ratio selected pair samples, or to morphological studies, where the total baryonic mass is what matters, the merger rate in terms of the baryon-baryon mass ratio is more relevant. At high masses, this behaves similarly to the merger rate in terms of the stellar-stellar mass ratio. At low masses, however, galaxies are increasingly gas rich -- thus many mergers with minor stellar-stellar mass ratios in fact have major baryon-baryon mass ratios, and the baryon-baryon major merger rate is a factor $\\sim3$ higher than the stellar-stellar merger rate. This may, in fact, explain at least part of the well-established fact that morphology-inferred merger rates at low masses tend to be systematically higher than pair count-inferred merger rates \\citep[see e.g.][and references therein]{lopezsanjuan:merger.fraction.to.z1, lopezsanjuan:mgr.rate.pairs}.  Although often a superior definition, even the baryonic mass ratio misses the tightly-bound dark matter that is important for the dynamics of final merger, which will not be stripped because it is within the baryonic radii. We therefore propose a new ``dynamical'' mass and mass ratio which approximates the most important quantity in high-resolution simulations, and includes both baryons and some dark matter. Merger rates in terms of this quantity behave similarly to baryonic major mergers, but with somewhat ($\\sim20-50\\%$) higher (lower) rates at low (high) masses (Figure~\\ref{fig:merger.fx.cautions}).  If one wishes to infer the amount of bulge or other ``damage'' done by major mergers but uses {\\em either} the stellar-stellar mass ratio or halo-halo mass ratio, the estimate can easily be systematically incorrect at the factor $\\sim3-10$ level. This is critical for observational studies seeking to infer the amount of bulge formed by major mergers, using a stellar-mass-ratio selected sample \\citep[see e.g.][]{bundy:merger.fraction.new}; at low masses, there might be $\\sim3$ times as many ``damaging'' mergers as given by this statistic. Also, many analytic models simply use the stellar (or even halo) mass ratio as a criterion for determining the impact of a merger -- the systematic errors introduced by this assumption can be {\\em larger} than those from halo mis-identification, complete ignorance of subhalos, assignment of incorrect merger timescales, or discrepancies between model and observed stellar mass functions. Yet despite the considerable literature on those problems, there has been relatively little focus on the adoption of better mass ratio proxies.   {\\bf Other Parameters and Merger `Impact':} Even when controlling for the above uncertainties, we show that at perfectly well-defined merger rates and merger mass ratios, there is a large variation in the physical effects of mergers, owing to other parameters such as the merger dynamics and orbit, merger gas fractions, and initial structural properties of the merging systems (Figure~\\ref{fig:merger.fx.cautions}). At the same merger mass ratio, a prograde orbit can build twice as much bulge as a retrograde orbit, and will lead to much more dramatic tidal features and distortions observable in morphology-selected samples. Also, the merger timescale will be significantly different at moderate/small radii $<100\\,$kpc, so pair-selected samples will also see biased distributions. A gas-poor major merger will typically violently relax the entire stellar disk and funnel the gas entirely into a starburst that builds central bulge mass, but a very gas-rich merger will experience very inefficient angular momentum transfer, suppressing the burst mass by a factor $\\sim (1-f_{\\rm gas})$, a factor $\\sim3-5$ at low masses and high redshifts. More subtle properties lead to scatter and offsets at a smaller, but non-negligible level. Many of these are discussed in detail in \\citet{hopkins:disk.survival}.  From a purely empirical perspective, without knowledge of these properties, it is difficult to assess the impact of mergers (in terms of e.g.\\ the bulge mass formed) at better than the factor $\\sim2-3$ level. From the theoretical perspective, neglecting quantities such as orbital parameters and, especially, gas fractions, in forward-modeling merger remnants again will introduce systematic uncertainties that are larger than any of the uncertainties from modeling the dark matter distribution, subhalos, merger timescales, and the like. Many of these are outlined in \\citet{hopkins:disk.survival.cosmo}; at high masses, the systematic uncertainties are less severe, but at low masses, where systems tend to be gas-rich, they can be factors $\\sim5-10$ in the total bulge mass formed.  Another important point is that mergers are continuous -- there is no special division at the traditional 1:3 distinction between ``major'' and ``minor'' mergers. In many models, and in many observational assessments of merger ``effects,'' it is assumed that major mergers are completely destructive, while minor mergers do no damage. This simple assumption introduces systematic factor $\\sim2$ errors in the total merger ``damage budget'' and total bulge formation efficiency \\citep[see][]{hopkins:merger.rates}; the effect in terms of skewing which mergers do ``more'' or ``less'' for bulge formation can obviously be severe.      \\subsection{Observational Constraints and Outlook} \\label{sec:discuss:obs}  Testing these models and determining the true merger rate in a robust manner will of course ultimately depend on observations. Continued observations of the merger rate, bearing in mind the caveats above, are of obvious importance. In improving such estimates, calibration of specific samples to high-resolution $N$-body simulations, {\\em specifically with mock observations matched to the exact selection and methodology adopted}, will be critical \\citep[see e.g.][]{lotz:merger.selection}. And even with such a calibration, the discussion above makes it clear why the relation between observed merger rates and bulge buildup (depending on a number of secondary properties not directly observed), in detail, must be {\\em forward-modeled}.  Tighter observational constraints on the halo occupation distribution -- from e.g.\\ group catalogs, kinematics, weak lensing, and clustering -- in particular at low masses and at high redshifts, will directly improve the semi-empirical models, and put strong constraints on the a priori galaxy formation models in the areas that have greatest effect on predicted merger rates. As we have shown, constraints on satellite populations specifically will be valuable. However, the satellite populations of particular interest are not the historically well-studied extreme cases of e.g.\\ dwarfs in the local group or low-mass galaxies in Virgo and massive clusters (from which much of our intuition regarding dynamical friction, stripping, and satellite gas exhaustion comes). The case of interest for merger rates is that of major mergers (i.e.\\ near equal-mass galaxies) in field or loose group environments (i.e.\\ systems analogous to the local group, but where either the Milky Way or Andromeda is the ``satellite'' of interest). Also, at low stellar masses and high redshifts, large uncertainties remain in galaxy gas masses, and these matter as much or more relative to the stellar mass in assessing the ``impact'' of mergers."
    },
    "1004/1004.5402_arXiv.txt": {
        "abstract": "We present the first ultraviolet (UV) and multi-epoch optical spectroscopy of 30~Dor~016, a massive O2-type star on the periphery of 30~Doradus in the Large Magellanic Cloud.  The UV data were obtained with the Cosmic Origins Spectrograph on the {\\em Hubble Space Telescope} as part of the Servicing Mission Observatory Verification program after Servicing Mission~4, and reveal \\#016 to have one of the fastest stellar winds known.  From analysis of the C~\\4 \\lam\\lam1548-51 doublet we find a terminal velocity, $v_\\infty$\\,$=$\\,3450\\,$\\pm$\\,50\\,\\kms.  Optical spectroscopy is from the VLT-FLAMES Tarantula Survey, from which we rule out a massive companion (with 2\\,d\\,$<$\\,$P$\\,$<$\\,1\\,yr) to a confidence of 98\\%.  The radial velocity of \\#016 is offset from the systemic value by $-$85\\,\\kms, suggesting that the star has traveled the 120\\,pc from the core of 30 Doradus as a runaway, ejected via dynamical interactions. ",
        "introduction": "30~Doradus in the Large Magellanic Cloud (LMC) is the richest H\\,\\2 region in the Local Group, providing an excellent template with which to study regions of intense star formation, and both stellar and cluster evolution.  It harbors a significant fraction of the most massive and luminous stars known, with a rich population of the earliest O-type stars (e.g., Melnick 1985; Walborn \\& Blades 1997), particularly in R136, the dense cluster at its core (Massey \\& Hunter 1998).  The O3 spectral class was introduced by Walborn (1971) to accommodate stars in which He~\\1 \\lam4471 was absent in moderate-resolution photographic spectra, compared to the very weak absorption seen in O4-type spectra.  The classification scheme was extended further by Walborn et al. (2002) to include the new types of O2 and O3.5 to delineate the behavior of the N~\\3, N~\\4, and N~\\5 features seen in the earliest types. Even in the age of large multi-object surveys only a few tens of stars are known with O2--O3.5 types.  Although very rare, their influence is far-reaching as they are expected to evolve rapidly into nitrogen rich Wolf--Rayet stars (WN types), plausible progenitors of supernovae and, potentially, gamma-ray bursts (Smartt 2009).  Observations with the 2-degree Field (2dF) instrument at the Anglo-Australian Telescope revealed a new O2-type star on the western fringes of 30~Doradus (Figure~\\ref{fig1}), with a radial velocity of $\\sim$85\\,\\kms\\/ lower than the systemic velocity of nearby massive stars (e.g., Bosch, Terlevich \\& Terlevich 2009).  New multi-epoch spectroscopy of this star, 30~Dor~016 in the VLT-FLAMES Tarantula Survey (Evans et al. 2010), now enables us to rule out the presence of a close massive companion to a high level of confidence, suggesting that the star might have been ejected from the denser central region. Here, we combine these observations with new ultraviolet (UV) spectroscopy of \\#016, some of the first data taken with the Cosmic Origins Spectrograph (COS) on the {\\em Hubble Space Telescope (HST)}, which reveal the star to have one of the highest wind terminal velocities seen to date in any massive star.  \\begin{figure} \\begin{center} \\includegraphics[width=8cm]{fig1} \\caption{Digital Sky Survey ``IR'' ($\\sim${\\em I}-band) image showing the location of 30~Dor~016 (encircled) relative to the 30~Dor complex.} \\label{fig1} \\end{center} \\end{figure} ",
        "conclusions": ""
    },
    "1004/1004.3542_arXiv.txt": {
        "abstract": "We develop a pseudo power spectrum technique for measuring the lensing power spectrum from weak lensing surveys in both the full sky and flat sky limits. The power spectrum approaches have a number of advantages over the traditional correlation function approach.  We test the pseudo spectrum method by using numerical simulations with square-shape boundary that include masked regions with complex configuration due to bright stars and saturated spikes.  Even when 25\\% of total area of the survey is masked, the method recovers the $E$-mode power spectrum at a sub-percent precision over a wide range of multipoles $100\\simlt\\ell\\simlt 10^4$. The systematic error is smaller than the statistical errors expected for a 2000 square degree survey. The residual $B$-mode spectrum is well suppressed in the amplitudes at less than a percent level relative to the $E$-mode.  We also find that the correlated errors of binned power spectra caused by the survey geometry effects are not significant.  Our method is applicable to the current and upcoming wide-field lensing surveys. ",
        "introduction": "\\label{sec:intro}  Large scale structure deflects light rays as they propagate from distant galaxies to us, thus distorting the shapes of these galaxies \\citep[e.g.,][for thorough reviews]{BS01,HoekstraJain08}. This weak lensing or cosmic shear signals measures a combination of the total matter distribution projected along the line-of-sight and the angular diameter distance. Since the first measurements of weak lensing only a decade ago \\citep{VW00,Wittman00,Bacon00,Kaiser00}, there have been major improvements as surveys continue to grow in size and depth \\citep[e.g.][for the latest results]{Fuetal08,Ichiki09,Schrabback10}. Gravitational lensing is one of the most promising methods of constraining cosmology including the nature of dark energy \\citep[e.g.][]{TakadaBridle07}. There are various on-going and planned surveys aimed at studying dark energy through the high-precision weak lensing measurements: the CFHT Legacy Survey\\footnote{http://www.cfht.hawaii.edu/Science/CFHLS/}, the Hyper Suprime-Cam Weak Lensing Survey \\citep{Miyazaki06}, the Dark Energy Survey (DES)\\footnote{http://www.darkenergysurvey.org}, and ultimately Large Synoptic Survey Telescope (LSST)\\footnote{http://www.lsst.org}, Euclid \\citep{Refregier10}, and Joint Dark Energy Mission (JDEM).  How should we analyze these new lensing data set?  Most researchers use the two-point correlation function to characterize the cosmic shear signals. The correlation function method can be easily applied to complex survey geometries involving partial sky coverage and masked regions. However, the errors in the measurement are highly correlated between different bins \\citep[see][for the detailed studies]{Schneider02,Joachimietal08}. Even if the shear field follows Gaussian statistics, which is a good approximation in the linear regime, there are large correlations between different angular-scale bins.  These correlations are even larger, and more model-dependent on the smaller scales that contain most of the current observational information. Since the accurate estimate of the covariances is essential for robust cosmological constraints, a large number of numerical simulations are necessary \\citep{Semboloni07}.  The power spectrum, the Fourier- or Harmonic-transformed counterpart of the two-point correlation function, is an alternative means of measuring the cosmic shear correlations.  Whilst the correlation function and power spectrum are mathematically equivalent, the power spectrum measurement of cosmic shear has been used less \\citep[see the COMBO 17-survey by][]{Brown03}. The power spectrum approach has a number of advantages: its theoretical interpretation is simpler and there are weaker correlations between band powers at different multipoles. For example, the different bins are independent for the Gaussian field or on large angular scales. Even for small angular scales affected by nonlinear structure formation, the power spectrum covariances are relatively well understood through both analytical models and simulations of nonlinear structure formation \\citep{HuWhite01,CoorayHu01,TakadaJain09,Sato09,Pielorzetal09}.  The disadvantage is the presence of finite sky coverage and masked regions, which breaks the orthogonality of Fourier/Harmonic components.  One needs to properly deal with the survey geometry effect to estimate unbiased power spectrum.  The purpose of this paper is to eliminate this disadvantage.  We employ the {\\em pseudo} power spectrum technique, which is well developed in the CMB studies \\citep[e.g.][]{Wandelt01} \\footnote{See \\citet{Seljak98} and \\citet{HuWhite01} for the maximum likelihood method of shear power spectrum estimation.}. For the first time, we apply the method to recover the lensing power spectrum from the shear field taking into account incomplete survey geometry.  To assess the performance of this method, we make simulated shear maps including a realistic configuration of masked regions due to bright stars and saturated spikes. Furthermore, we develop the method for both the full-sky and flat-sky approaches.  The full-sky approach is adequate for reconstructing large angular-scale modes that are relevant for the curvature of the sky.  On the other hand, the flat-sky approach should serve as a practically useful approximation of sub-degree scale modes, which carry most of useful cosmological information in the shear power spectrum.  We find that the pseudo power spectrum method allows for an unbiased estimate of the underlying $E$-mode power spectrum over a range of angular scales we study. We also show that the residual $B$-mode power spectrum, which is leaked from $E$-mode power due to an imperfect reconstruction, can be well suppressed. Our method can be applied to the existing data and forthcoming weak lensing surveys.   The paper is organized as follows: Section \\ref{sec:method} describes the pseudo spectrum method to deconvolve shear power spectra with inhomogeneous survey mask.  Section \\ref{sec:simulations} describes the simulation maps we use to test the deconvolution method. We employ two different simulation maps: one is Gaussian shear maps and the other is the ray-tracing simulations of shear maps including the non-Gaussian effects due to nonlinear structure formation.  Section \\ref{sec:results} shows the results of both the full-sky and flat-sky approaches.  Section~\\ref{sec:summary} is devoted to the summary and conclusions. ",
        "conclusions": "\\label{sec:summary}  We develop a pseudo-spectrum method for reconstructing the cosmic shear power spectra from actual lensing data. The observed shear field is limited to be a finite patch of sky and furthermore roughly 25\\% of the survey area is masked due to bright stars.  We apply for the first time the pseudo-spectrum technique developed in CMB studies to the lensing field and show that our method successfully recovers the shear spectra over a wide range of multipoles from $100$ to $10^4$ in both full- and flat-sky approaches.  We test the flat-sky method using ray-tracing simulations assuming a square-shaped survey region.  The 25\\% fraction of the total area is masked by realistic configurations for a ground-based survey: circular shapes for bright stars; rectangular shapes for bright star spikes; zero padding in one direction for bad pixels. We show that both the full- and flat-sky methods reconstruct the input $E$-mode power spectrum in sub-percent accuracy and the residuals are much smaller than the statistical errors of a power spectrum measurement for a survey of 2000 square degree coverage. The residual $B$-mode power spectrum from the $E/B$-mode mixing due to the imperfect correction of survey geometry is also suppressed below a percent of the $E$-mode power spectrum.  Our method offers a new means of measuring the cosmic shear correlations and separating the $E/B$ modes from an actual survey data.  Although the pseudo-spectrum technique is promising, our method still yields sub-percent residual of $B$-mode in the reconstructed power spectra. To further suppress the $B$-mode spectrum, one has to remove ``ambiguous mode'' that is inevitably generated in a finite patch of sky \\citep{Bunn03}. The ambiguous mode satisfies both the $E$-mode (rotation-free) and the $B$-mode conditions (divergence-free) and thus contaminates E/B-mode spectrum reconstructed using the simple pseudo-spectrum method that we adopt. To eliminate such contamination, \\citet{Bunn03} introduces pure E/B modes that is orthogonal to the ambiguous modes. Pure pseudo $C_l$ estimator and its optimization technique of sky apodization have been developed \\citep{Smith06,SmithZal07,Grain09,Kim10}. This technique can be straightforwardly applied to the lensing case.  In this paper we assume that masked regions are uncorrelated with the cosmological shear field.  In reality the regions with large shear is preferentially masked: in a region of massive clusters, we cannot obtain a fair sample of background galaxies in the central region due to the dense concentration of member galaxies, where the shearing effect on background galaxies are greater. Masking such a crowded region may bias the power spectrum measurement.  This contaminating effect can be estimated by combining ray-tracing simulations with halo catalogs in the underlying N-body simulations. This is our future project, and will be presented elsewhere.  \\bigskip  We deeply appreciate Masanori~Sato for kindly providing ray-tracing simulation data. We also thank an anonymous referee for careful reading and providing useful comments.  C.H. acknowledges support from a Japan Society for Promotion of Science (JSPS) fellowship.  This work is in part supported in part by JSPS Core-to-Core Program ``International Research Network for Dark Energy'', by Grant-in-Aid for Scientific Research from the JSPS Promotion of Science (18072001,21740202), by Grant-in-Aid for Scientific Research on Priority Areas No. 467 ``Probing the Dark Energy through an Extremely Wide \\& Deep Survey with Subaru Telescope'', and by World Premier International Research Center Initiative (WPI Initiative), MEXT, Japan."
    },
    "1004/1004.3297_arXiv.txt": {
        "abstract": "\\noindent Supersymmetric models based on anomaly-mediated SUSY breaking (AMSB) generally give rise to a neutral wino as a WIMP cold dark matter (CDM) candidate, whose thermal abundance is well below measured values. Here, we investigate four scenarios to reconcile AMSB dark matter with the measured abundance: 1. non-thermal wino production due to decays of scalar fields ({\\it e.g.} moduli), 2. non-thermal wino production due to decays of gravitinos, 3. non-thermal wino production due to heavy axino decays, and 4. the case of an axino LSP, where the bulk of CDM is made up of axions and thermally produced axinos. In cases 1 and 2, we expect wino CDM to constitute the entire measured DM abundance, and we investigate wino-like WIMP direct and indirect detection rates. Wino direct detection rates can be large, and more importantly, are bounded from below, so that ton-scale noble liquid detectors should access all of parameter space for $m_{\\tz_1}\\alt 500$ GeV. Indirect wino detection rates via neutrino telescopes and space-based cosmic ray detectors can also be large. In case 3, the DM would consist of an axion plus wino admixture, whose exact proportions are very model dependent. In this case, it is possible that both an axion and a wino-like WIMP could be detected experimentally. In case 4., we calculate the re-heat temperature of the universe after inflation. In this case, no direct or indirect WIMP signals should be seen, although direct detection of relic axions may be possible. For each DM scenario, we show results for the minimal AMSB model, as well as for the hypercharged and gaugino AMSB models. \\vspace*{0.8cm} ",
        "introduction": "\\label{sec:intro}  Supersymmetric (SUSY) models of particle physics are very attractive in that they stabilize the gauge hierarchy problem, and provide an avenue for the incorporation of gravity via local SUSY, or supergravity\\cite{wss}. They also receive some indirect experimental support via the unification of gauge couplings under Minimal Supersymmetric Standard Model (MSSM) renormalization group evolution (RGE)\\cite{drw}, and they provide several different candidates (neutralinos, gravitinos, axions/axinos, $\\cdots$) which can serve as cold dark matter (CDM) in the universe. If evidence for SUSY is found at LHC, then a paramount question will be: what is the mechanism of SUSY breaking, and how is it communicated to the visible sector? Some of the possibilities proposed in the literature include: gravity-mediation (SUGRA) with a gravitino mass $m_{3/2}\\sim 1$ TeV\\cite{sugra}, gauge-mediation (GMSB) with $m_{3/2}\\ll 1$ TeV\\cite{gmsb}, and anomaly mediation (AMSB) with $m_{3/2}\\sim 100$ TeV\\cite{rs,glmr,recent}.  Anomaly-mediated supersymmetry breaking (AMSB) models have received much attention in the literature due to their attractive properties\\cite{rs,glmr}: 1. the soft supersymmetry (SUSY) breaking terms are completely calculable in terms of just one free parameter (the gravitino mass, $m_{3/2}$), 2. the soft terms are real and flavor invariant, thus solving the SUSY flavor and $CP$ problems and 3. the soft terms are actually renormalization group invariant\\cite{jj}, and can be calculated at any convenient scale choice. In order to realize the AMSB set-up, it was proposed that the hidden sector be ``sequestered'' on a separate brane from the observable sector in an extra-dimensional universe, so that tree-level supergravity breaking terms do not dominate the soft term contributions. Such a set-up can be realized in brane-worlds, where SUSY breaking takes place on one brane, with the visible sector residing on a separate brane.  A further attractive feature of AMSB models arises due to the scale of their gravitino mass. SUGRA-type models with $m_{3/2}\\sim 1$ TeV suffer from the cosmological gravitino problem. There are two parts to the gravitino problem\\cite{gprob}. 1. If the re-heat temperature after inflation $T_R\\agt 10^{10}$ GeV, then the high rate of thermal gravitino production leads to an overabundance of neutralino dark matter\\cite{moroi}. 2. Even for lower values of $T_R\\sim 10^5-10^{10}$ GeV, thermal production of $\\tG$ followed by late decays to $particle+sparticle$ pairs injects high energy particles into the cosmic soup during or after BBN, thus disrupting one of the pillars of Big-Bang theory\\cite{moroi}. If $m_{3/2}\\agt$ 5 TeV, then the lifetime $\\tau_{\\tG}$ drops below $0.1-1$ sec,  and gravitino decay occurs before or at the onset of BBN. In AMSB models where $m_{3/2}\\sim 100$ TeV, the gravitino is much too short-lived to be afflicted by the BBN bounds.  In spite of their attractive features, AMSB models suffer from the well-known problem that slepton mass-squared parameters are found to be negative, giving rise to tachyonic states. The original ``solution'' to this problem was to posit that scalars acquire as well a universal mass $m_0$, which when added to the AMSB SSB terms, renders them positive\\cite{rs,glmr}. The derived form of soft SUSY breaking terms, supplemented by a universal scalar mass $m_0$ and implemented at the GUT scale, constitutes what is usually called the minimal AMSB, or mAMSB model. In mAMSB and the additional models described below, it is assumed that electroweak symmetry is broken radiatively due to the large top quark mass, so that the magnitude of the $\\mu$ parameter is determined to gain the correct value of $M_Z$, and the bilinear soft term $B$ is traded for the ratio of Higgs field vevs, $\\tan\\beta$.  An alternative set-up for AMSB has been advocated in Ref. \\cite{dvw}, known as hypercharged anomaly-mediation (HCAMSB). It is a string-motivated scenario which uses a similar setup as the one envisioned for AMSB. In HCAMSB, the MSSM resides on a D-brane, and the hypercharge gaugino mass is generated in a geometrically separated hidden sector. An additional contribution to the $U(1)_Y$ gaugino mass $M_1$ is generated, and its magnitude is parametrized by an additional parameter $\\alpha$. The large value of $M_1$ feeds into slepton mass evolution through the MSSM RGE, and acts to lift the weak-scale slepton soft masses beyond tachyonic values. Thus, the HCAMSB model naturally solves the tachyonic slepton mass problem which is endemic to pure AMSB scenarios.   A third scenario has recently been proposed in Ref. \\cite{inoamsb}, under the name gaugino AMSB, or inoAMSB. The inoAMSB model is suggested by recent work on the phenomenology of flux compactified type IIB string theory\\cite{shanta}, which reduces to $N=1$ supergravity below the compactification scale. The essential features of this scenario are that the  gaugino masses are of the anomaly-mediated SUSY breaking (AMSB) form, while scalar and trilinear soft SUSY breaking terms are highly suppressed: they are taken as $m_0=A_0\\simeq 0$ at energy scale $Q\\sim M_{GUT}$, at first approximation. The normally large value of $M_1$ as generated in AMSB models feeds into the scalar soft term evolution, lifting slepton soft masses to generate an allowable sparticle mass spectrum, while at the same time avoiding tachyonic sleptons or charged LSPs (lightest SUSY particles). Charged LSPs are common in models with negligible soft scalar masses, such as no-scale\\cite{noscale} or gaugino mediation models\\cite{inoMSB}. Since scalar and trilinear soft terms are highly suppressed, the SUSY induced flavor and $CP$-violating processes are also suppressed in inoAMSB.  All three of these models-- mAMSB, HCAMSB and inoAMSB-- share the common feature that the lightest MSSM particle is a neutral wino, while the lightest chargino is wino-like with a mass $m_{\\tw_1}\\sim m_{\\tz_1}$. The $\\tw_1$-$\\tz_1$ mass gap is of order $\\sim 200$ MeV\\cite{matchev}, so that dominantly $\\tw_1^\\pm\\to\\tz_1\\pi^\\pm$, with the decay-produced pion(s) being very soft. The small mass gap makes the $\\tw_1$ rather long lived ($\\tau_{\\tw_1}\\sim 10^{-9}$ sec), and it may yield observable highly ionizing tracks (HITs) of order $cm$ in length at LHC detectors\\cite{amsb_coll}.  An important consequence of wino-like neutralinos is that the thermal abundance of neutralino cold dark matter falls generally an order of magnitude or so below the measured abundance: \\be \\Omega_{CDM}h^2=0.1123\\pm 0.0035\\ \\ \\ 68\\%\\ CL \\ee according to the WMAP7 data analysis\\cite{wmap7}. This latter fact has led many to consider AMSB-like models as perhaps less interesting than SUGRA-type models, wherein the bino-like or mixed bino-higgsino neutralino can more easily yield the measured relic abundance.  In this paper, we address the question of the dark matter abundance in AMSB models. While the calculated thermal abundance of wino-like neutralinos is found to be below measured values (for $m_{\\tz_1}\\alt 800$ GeV), we find that there exists a variety of attractive methods to augment the wino abundance, thus bringing the calculated abundance into accord with experiment. These include: \\begin{enumerate} \\item Decay of scalar ({\\it e.g.} moduli) fields into sparticles, ultimately terminating in $\\tz_1$ production\\cite{mr}. In this case, the LSP is expected to be a relic wino-like neutralino, which would constitute {\\it all } of the CDM. \\item Thermal production\\cite{bbs,ps} of gravitinos $\\tG$ and also possibly gravitino production via moduli\\cite{kyy} or inflaton\\cite{ety} decay, followed by $\\tG\\to particle+\\ sparticle\\to \\tz_1 + X$ (where $X=$ assorted SM debris). Here also, the LSP would be a relic wino-like neutralino, which would constitute all of the CDM. \\item Thermal production of heavy axinos\\cite{bs} followed by $\\ta\\to particle +\\ sparticle\\to\\tz_1 +X$\\cite{ckls}. Here, the LSP is again a relic wino-like neutralino, but the CDM would consist of a wino-like WIMP plus axion mixture. \\item A scenario where $m_{\\ta}<m_{\\tz_1}$, so the $\\ta$ is instead the stable LSP\\cite{ckkr}. In this case, a combination of thermally produced axinos plus vacuum mis-alignment produced axions would constitute the CDM\\cite{bs_ou,bbs_ou}. \\end{enumerate}   In Sec. \\ref{sec:models}, we present some details of the three AMSB models which we investigate. In Sec. \\ref{sec:dm}, we present four methods of reconciling the AMSB CDM relic abundance with the measured value. Given cases 1 and 2, and possibly 3, where we expect all (or some fraction) of the measured abundance to consist of relic winos, in Sec. \\ref{sec:winodet} we present rates for direct and indirect detection of wino-like neutralinos. Unlike SUGRA models, the wino CDM direct detection rate is bounded from below. We find the current experiments like Xenon-100 should be able to explore the parameter space of AMSB-like models with a wino-like neutralino up to $m_{\\tz_1}\\alt 200$ GeV. Next generation detectors such as ton-scale noble liquids or SuperCDMS should be able to push to $m_{\\tz_1}\\sim 500$ GeV. This would correspond to a reach in $m_{\\tg}\\sim 3850$ GeV, {\\it i.e.} well beyond the projected reach of LHC. We also find excellent prospects for indirect detection of wino-like CDM via detection of wino annihilation into $\\gamma$s, $e^+$s, $\\bar{p}$s or $\\bar{D}$s in  the galactic halo. In fact, Kane {\\it et al.} have already proposed wino CDM as an explanation for the recent anomalies seen by Pamela, ATIC, Fermi and others\\cite{kane}. Neutrino telescopes such as IceCube will also have a reach for wino-like neutralinos, especially for large $\\tan\\beta$.  In case 4, we would expect the axino $\\ta$ to be the LSP, and so here no direct or indirect WIMP detection signals are expected. However, it may be the case that large amounts of {\\it axions} $a$ are produced in the early universe, in which case direct detection of axions may be possible at ADMX\\cite{admx}. We present parameter expectations in Sec. \\ref{ssec:ata} for the scenario of mixed axion/axino CDM to occur. In Sec. \\ref{sec:conclude}, we present a summary and our conclusions. ",
        "conclusions": "\\label{sec:conclude}  In this paper, we have investigated aspects of cold dark matter in three models of anomaly mediation: mAMSB, HCAMSB and inoAMSB. Typically, each gives rise to a wino-like lightest neutralino, unless very high values of $m_0$ (for mAMSB) or $\\alpha$ (for HCAMSB) are used, in which case the $\\tz_1$ becomes a mixed wino-higgsino state. In this class of models with a wino-like $\\tz_1$, the thermal abundance of neutralino CDM is well below measured values, unless $m_{\\tz_1}\\agt 1300$ GeV. We discuss four ways to reconcile the predicted abundance of CDM with experiment: \\begin{enumerate} \\item enhanced neutralino production via scalar field ({\\it e.g.} moduli) decay, \\item enhanced neutralino production via gravitino decay, where gravitinos may arise thermally, or by moduli or inflaton decay, \\item enhanced neutralino production via heavy axino decay, and \\item neutralino decay to axinos, where the bulk of CDM comes from a mixture of vacuum mis-alignment produced axions and thermally produced axinos. \\end{enumerate} Cases 1 and 2 should lead to a situation where all of CDM is comprised of wino-like WIMPs; they will be very hard, perhaps impossible, to tell apart. Case 3 would contain a mixture of axion and wino-like WIMP CDM. It is a scenario where it is possible that both a WIMP and an axion could be detected. Case 4 predicts no direct or indirect detection of WIMPs, but a possible detection of relic axions. It is important to note that more than one of these mechanisms may occur at once: for instance, we may gain additional neutralino production in the early universe from moduli, gravitino and axino decay all together.  In Sec. \\ref{sec:winodet}, we presented rates for direct and indirect detection of relic wino-like WIMPs. The SI direct detection cross sections are bounded from below. Ultimately, ton-scale noble liquid or SuperCDMS experiments should probe out to $m_{\\tz_1}\\sim 500$ GeV, which would exceed the 100 fb$^{-1}$ reach of LHC; a non-observation of signal would put enormous stress on AMSB-like models as new physics. We also evaluated SD direct detection: current experiments have little reach for AMSB-like models, although IceCube DeepCore and possibly COUPP upgrades may probe more deeply.  We also presented indirect WIMP detection rates for all three AMSB models. The IceCube experiment has some reach for WIMPs from AMSB models, especially at high $\\tan\\beta$ or when the $\\tz_1$ picks up a higgsino component. We noted an interesting inverse resonance effect in the muon flux detection rate, caused by transition from solar core annihilations to $VV$ states, to annihilations to mainly $b\\bar{b}$ states. The detection of $\\gamma$s, $e^+$s, $\\bar{p}$s and $\\bar{D}$s are all elevated in AMSB-like models compared to mSUGRA, due to the high rate for $\\tz_1\\tz_1\\to VV$ annihilation in the galactic halo. The results do depend on the assumed halo profile, especially for $\\gamma$-ray detection in the direction of the galactic core. Generally, if a signal is seen in the $e^+$ channel, then one ought to be seen in the $\\bar{p}$ channel, and ultimately in the $\\gamma$, $\\bar{D}$ (if/when GAPS flies) or direct detection channel. In addition, a sparticle production signal should ultimately be seen at LHC, at least for $m_{\\tg}\\alt 2400$ GeV, once 100 fb$^{-1}$ of integrated luminosity is accrued.  As a final remark, we note here that the dark matter detection signals all provide complementary information to that which will be provided by the CERN LHC. At LHC, each model-- mAMSB, HCAMSB and inoAMSB-- will provide a rich assortment of gluino and squark cascade decay signals which will include multi-jet plus multi-lepton plus missing $E_T$ events. In all cases, the wino-like lightest neutralino state will be signaled by the well-known presence of highly ionizing tracks (HITs) from quasi-stable charginos with track length of order {\\it cm}s, before they decay to soft pions plus a $\\tz_1$. It is noted in Ref's \\cite{hcamsb} and \\cite{inoamsb} that the three models should be distinguishable at LHC by the differing opposite-sign/same flavor dilepton invariant mass distributions. In the case of mAMSB, with $m_{\\tell_L}\\simeq m_{\\tell_R}$, we expect a single mass edge from $\\tz_2\\to\\ell\\tell_{L,R}\\to\\ell^+\\ell^-\\tz_1$ decay. In HCAMSB, the sleptons are rather heavy, and instead $\\tz_2\\to\\tz_1 Z$ occurs at a large rate, leading to a bump in $m(\\ell^+\\ell^-)\\sim M_Z$, upon a continuum distribution. In inoAMSB, with $m_{\\tz_2}>m_{\\tell_{L,R}}$, but with $\\tell_L$ and $\\tell_R$ split in mass (due to different $U(1)_Y$ quantum numbers), a characteristic {\\it double mass edge} is expected in the $m(\\ell^+\\ell^-)$ invariant mass distribution."
    },
    "1004/1004.0449.txt": {
        "abstract": "{The CoRoT 5-month long observation runs give us the opportunity to analyze a large variety of red-giant stars and to derive fundamental stellar parameters from their asteroseismic properties.}% % {We perform an analysis of more than 4\\,600 CoRoT light curves to extract as much information as possible. We take into account the characteristics of the star sample and of the method in order to provide asteroseismic results as unbiased as possible. We also study and compare the properties of red giants of two opposite regions of the Galaxy.}% % {%Compared to previous works, we have obtained precise asteroseismic results with the analysis of the time series with the envelope autocorrelation function. With this method, we manage We analyze the time series with the envelope autocorrelation function in order to extract precise asteroseismic parameters with reliable error bars. We examine first the mean large frequency separation of solar-like oscillations and the frequency of maximum seismic amplitude, then the parameters of the excess power envelope. With the additional information of the effective temperature, we derive the stellar mass and radius.}% % {We have identified more than 1\\,800 red giants among the 4\\,600 light curves and have obtained accurate distributions of the stellar parameters for about 930 targets. We were able to reliably measure the mass and radius of several hundred red giants. We have derived precise information on the stellar population distribution and on the red clump. Comparison between the stars observed in two different fields shows that the stellar asteroseismic properties are globally similar, but with different characteristics for red-clump stars.}% % {This study shows the efficiency of statistical asteroseismology: validating scaling relations allows us to infer fundamental stellar parameters, to derive precise information on the red-giant evolution and interior structure and to analyze and compare stellar populations from different fields.} ",
        "introduction": "The high-precision, continuous, long  photometric time series recorded by the CoRoT satellite allow us to study a large number of red giants. In a first analysis of CoRoT red giants, \\cite{2009Natur.459..398D} reported the presence of radial and non-radial oscillations in more than 300 giants. \\cite{2009A&A...506..465H}, after a careful analysis of about 1000 time series, have demonstrated a tight relation between the large separation and the frequency of maximum oscillation amplitude.  \\cite{2009A&A...503L..21M} have identified the signature of the red clump, in agreement with   synthetic populations. \\cite{2009ASPC..404..307K} have exploited the possibility of measuring stellar mass and radius from the asteroseismic measurements, even when the stellar luminosity and effective temperature are not accurately known.  In this paper, we specifically focus on the statistical analysis of a large set of stars in two different fields observed with CoRoT \\citep{2009A&A...506..411A}. One is located towards the Galactic center (LRc01), the other in the opposite direction  (LRa01). We first derive precise asteroseismic parameters, and then stellar parameters. We also examine how these parameters vary with the  frequency $\\numax$ of the maximum amplitude. The new analysis that we present in this paper has been made possible by the use of the autocorrelation method \\citep{2009A&A...508..877M}, which is significatively different from what has been used in other works \\citep{2010A&A...511A..46M, 2009A&A...506..465H, 2009CoAst.160...74H}. It does not rely on the identification of the excess oscillation power, but on the direct measurement of the acoustic radius $\\tau$ of a star. This acoustic radius is  related to the large separation commonly used in asteroseismology ($\\dnu = 1/2\\tau$). The chronometer is provided by the autocorrelation of the asteroseismic time series, which is sensitive to the travel time of a pressure wave crossing the stellar diameter twice, i.e. 4 times the acoustic radius. %Autocorrelation of the time series has often been used for various purposes in helioseismology \\citep[e.g.][]{1999A&A...343..608F}. The calculation of this autocorrelation as the Fourier spectrum of the Fourier spectrum with the use of narrow window for a local analysis in frequency has been suggested by \\cite{2006MNRAS.369.1491R}. \\cite{2009A&A...508..877M} have formalized and quantified the performance of the method based on the envelope autocorrelation function (EACF).  With this method, and the subsequent automated pipeline, we search for the signature of the mean large separation of a solar-like oscillating signal in the autocorrelation of the time series. \\cite{2009A&A...508..877M} have shown how to deal with the noise contribution entering the autocorrelation function, so that they were able to determine the reliability of the large separations obtained with this method. Basically, they have scaled the autocorrelation function according to the noise contribution. With this scaling, they have shown how to define the threshold level above which solar-like oscillations are detected and a reliable large separation can be derived.  An appreciable advantage of the method consists in the fact that the large separation is determined first, without any assumptions or any fit of the background. As a consequence, the method directly focusses on the key parameters of asteroseismic observations: the mean value $\\dnumoy$ of the large separation and the frequency $\\numax$ where the oscillation signal is maximum. Since it does not rely on the detection of an energy excess, it can operate at low signal-to-noise ratio, as shown by \\cite{2009A&A...506...33M}. The value of the frequency $\\numax$, inferred first from the maximum autocorrelation signal, is then refined from the maximum excess power observed  in a smoothed Fourier spectrum corrected from the background component. The different steps of the pipeline for the automated analysis of the time series are presented in \\citet{ma10}.  The method has been tested on CoRoT main-sequence stars \\citep{2009A&A...507L..13B, 2009A&A...506...51B, 2009A&A...506...41G, Ballot2010, Deheuvels2010, Mathur2010} and has proven its ability to derive reliable results rapidly for low signal-to-noise ratio light curves, when other methods fail or give questionable results \\citep{2009A&A...506...33M, Gaulme2010}. %It has been able to provide the measurement of the large separation of a K1V target hosting an exoplanet despite a time series where the mean seismic excess power represents only 2.5\\,\\% of the power of the background \\citep{Gaulme2010}. The method also allowed the correct identification of the degree of the eigenmodes of the first CoRoT target HD49933 \\citep{2005A&A...431L..13M, 2008A&A...488..705A, 2009A&A...508..877M}. The EACF method and its automated pipeline have been tested on the CoRoT red giants presented by  \\cite{2009Natur.459..398D}  and \\cite{2009A&A...506..465H}, and also on the Kepler red giants \\citep{Stello2010, Bedding2010}.  The paper is organized as follows. In Sect.~\\ref{methode}, we present the analysis of the CoRoT red giants using the EACF and define the way the various seismic parameters are derived. We also determine the frequency interval where we can extract unbiased global information. Measurements of the asteroseismic parameters $\\dnumoy$ and $\\numax$ are presented in Sect.~\\ref{scaling} and compared to previous studies. We also present the variation $\\deltanunu$ performed with the EACF. Section~\\ref{power} deals with the parameters related to the envelope of the excess power observed in the Fourier spectra, for which we propose scaling laws. From the asteroseismic parameters $\\numax$ and $\\dnumoy$, we determine the red-giant mass and radius in Sect.~\\ref{masserayon}. Compared to \\cite{2009ASPC..404..307K}, we benefit from the stellar effective temperatures obtained from independent photometric measurements, so that we do not need to refer to stellar modeling for deriving the fundamental parameters. We then specifically address the properties of the red clump in Sect.~\\ref{redclump}, so that we can carry out a quantitative comparison with the synthetic population performed by \\cite{2009A&A...503L..21M}. The difference between the red-giant populations observed in 2 different fields of view is also presented in Sect.~\\ref{redclump}. Section~\\ref{conclusion} is devoted to discussions and conclusions.  %__________________________________________________________________  \\begin{table} \\caption{Red-giant targets}\\label{runs} \\begin{tabular}{lccccc} \\hline run  & $T$ (d)&$\\nzero$& $\\nun$ & $\\nde$& $\\ntr$ \\\\ \\hline %LRc1 & 150 d&      & 1228   & 875   & 520\\\\ %LRa1 & 125 d& 1380 & 1271   &  498   & 219\\\\ %LRc1 & 150 d& 9938 & 1228   & 875   & 520\\\\ LRc01 & 142 d& 9938 & 3388   & 1399    & 710    \\\\ LRa01 & 128 d& 2826 & 1271   &  428    & 219    \\\\ \\hline total&  --  &12764 & 4659   & 1827    & 929    \\\\ \\hline \\multicolumn{2}{l}{ratio (\\%)} &  --   &100\\,\\% & 39\\,\\%  & 20\\,\\% \\\\ \\hline \\end{tabular}  Among $\\nzero$ targets a priori identified as red giants in the input catalog of each field, $\\nun$ light curves were available and analyzed in each run. Among these, $\\nde$ targets show solar-like oscillation patterns for which we can derive precise values of $\\dnumoy$ and $\\numax$. Envelope parameters can be precisely determined for $\\ntr$ targets. \\end{table} ",
        "conclusions": "}  We have shown the possibility of extracting statistical information from the high-precision photometric time series of a large sample of red giants observed with CoRoT and analyzed with an automated asteroseismic pipeline. We summarize here the main results as well as open issues:  - Out of more than 4\\,600 time series, we have identified more than 1\\,800 red giants showing solar-like oscillations. We have extracted a full set of precise asteroseismic parameters for more than 900 targets.  - Thanks to the detection method, we are able to observe precise large separations as small as 0.75\\muHz.  We obtain reliable information for the seismic parameters  $\\dnumoy$ and  $\\numax$ for $\\numax$ in the range [3.5, 100\\muHz]. We have shown that the detection and measurement method does not introduce any bias for $\\numax$ above 6\\muHz. This allows us to study in detail the red clump in the range [30, 40\\muHz].  - We have proposed scaling relations for the parameters defining the envelope where the asteroseismic power is observed in excess. We note that the relation defining the full-width at half-maximum $\\nuenv$ of the envelope cannot be extended to solar-like stars. The scaling relation between $\\nuenv$ and $\\numax$ is definitely not linear for giants: $\\nuenv \\propto \\numax^{0.90}$. The maximum amplitude scales as $\\numax^{-0.85}$ or $(L/M)^{0.89}$. Deriving bolometric amplitudes will require more work: examination of the equipartition of energy between the modes and stellar atmosphere modeling.  - When supplemented by the effective temperature, the asteroseismic parameters $\\dnumoy$ and $\\numax$ give access to the stellar mass and radius. Red-giant masses derived from asteroseismology are degenerate, but it is still possible to estimate their value with a typical uncertainty of about 20\\,\\%. We have established a tight link between the maximum amplitude frequency $\\numax$ and the red-giant radius from an unbiased analysis in the range [7 - 30\\,$R_\\odot$] which encompasses the red-clump stars. This relation scales as $R\\ind{RG} \\propto \\numax^{-0.48}$.  - From this result, and taking into account the scaling law $\\dnumoy \\propto \\numax^{0.75}$, we have shown that the ratio $\\numax / \\nuc$ is constant for giants. A similar analysis performed on main-sequence stars and sub-giants reaches the same result: $\\numax / \\nuc$ is also nearly constant.  - As a by-product, we have shown that scaling laws are slightly but undoubtedly different between giants, sub-giants and dwarfs. For red-giant stars only, the fact that the temperature is nearly a degenerate parameter plays a significant role. As a consequence, global fits encompassing \\emph{all} stars with solar-like oscillations may be not precise, since they do not account for the different physical conditions between main-sequence and giant stars.  - The comparison of data from 2 runs pointing in different directions at different galactic latitudes has shown that the stellar properties are similar; the dispersion around the global fits are too weak to be noticed. The main difference between the 2 runs consists in different stellar populations. The distributions of the asteroseismic parameters are globally similar, except for the location of the red clump. %Red giants with a radius larger than 10\\,$R_\\odot$ are more common in LRc01 towards the inner region of the Galaxy.  - We have obtained precise information for the red-clump stars. Statistical asteroseismology makes it possible to identify the expected secondary clump and to measure the distribution of the fundamental parameters of the red-clump stars. We have shown that the relative importance of the two components of the clump is linked to the stellar population. The precise determination of the red-clump parameters will benefit from the asteroseismic analysis and the modeling of individual members of the clump. \\\\  These points demonstrate the huge potential of asteroseismology for stellar physics.  %______________________________________________________________"
    },
    "1004/1004.2919_arXiv.txt": {
        "abstract": "We report the {\\it Swift} discovery of the nearby long, soft gamma-ray burst GRB\\,100316D, and the subsequent unveiling of its low redshift host galaxy and associated supernova. We derive the redshift of the event to be $z = 0.0591 \\pm 0.0001$ and provide accurate astrometry for the GRB-SN. We study the extremely unusual prompt emission with time-resolved $\\gamma$-ray to X-ray spectroscopy, and find that the spectrum is best modelled with a thermal component in addition to a synchrotron emission component with a low peak energy. The X-ray light curve has a remarkably shallow decay out to at least 800~s. The host is a bright, blue galaxy with a highly disturbed morphology and we use Gemini South, VLT and HST observations to measure some of the basic host galaxy properties. We compare and contrast the X-ray emission and host galaxy of GRB\\,100316D to a subsample of GRB-SNe. GRB\\,100316D is unlike the majority of GRB-SNe in its X-ray evolution, but resembles rather GRB\\,060218, and we find that these two events have remarkably similar high energy prompt emission properties. Comparison of the host galaxies of GRB-SNe demonstrates, however, that there is a great diversity in the environments in which GRB-SNe can be found. GRB\\,100316D is an important addition to the currently sparse sample of spectroscopically confirmed GRB-SNe, from which a better understanding of long GRB progenitors and the GRB--SN connection can be gleaned. ",
        "introduction": "\\label{sec:intro} The connection between long-duration gamma-ray bursts (GRBs) and Type Ic core-collapse supernovae (SNe) has long been established, beginning with the association of GRB\\,980425 with SN\\,1998bw \\citep{Galama,Pian98}. Subsequent associations between nearby GRBs and spectroscopically-confirmed SNe include GRB\\,030329/SN\\,2003dh \\citep{Hjorth,Stanek}, GRB\\,031203/SN\\,2003lw \\citep{Malesani} and most recently GRB\\,060218/SN\\,2006aj \\citep{Campana,Pian2}. These are spectroscopically confirmed examples of nearby GRB-SN associations (out to $z = 0.17$). The characteristic `bump' of a supernova has been noted in the photometric data for a dozen or more GRBs out to $z\\sim1$ \\citep[e.g.][]{Zeh,Ferrero,Dellavalle,Woosley,Tanvir}, while the majority of GRBs lie at higher redshifts where such signatures would be impossible to detect \\citep[$\\langle z \\rangle = 2.2$][]{Jakobsson,Fynbo2}. The GRBs with SNe are therefore rare, but provide a crucial insight into the GRB--SN connection and the progenitors of long GRBs.  The GRBs with spectroscopically confirmed SNe are generally underluminous and subenergetic in comparison to a typical long GRB, though GRB\\,030329 is a notable exception to this. The prompt emission of these bursts has a lower energy spectral peak than a typical GRB \\citep{Kaneko} and they are suggested to have less relativistic outflows, or be viewed more off-axis. It has been suggested that the observed nearby ($z<0.1$), underluminous/subenergetic GRBs, including GRB\\,100316D, may be drawn from a different population to the cosmological GRB sample, motivated by estimates of the local GRB rate several times greater than the rate of cosmological GRBs \\citep[e.g.][]{Cobb,Chapman,Liang}. Radio observations of GRBs suggest that mildly relativistic ejecta and often high luminosities are the properties that distinguish the supernovae associated with GRBs from the non-relativistic ordinary Type Ib/c supernovae \\citep{Kulkarniradio,Soderbergradio}.  GRB\\,060218 was a landmark long-duration, low-luminosity, soft spectral peak event, detected by the {\\it Swift} satellite \\citep{Gehrels} with unprecedented multiwavelength coverage of the prompt emission. The prompt spectrum was found to comprise both the non-thermal synchrotron emission ascribed to most GRBs and thought to originate in the collision of fast-moving shells within the GRB jet \\citep[e.g.][]{Rees}, and a thermal component. The presence of this thermal component, together with its evolution, led to the suggestion that we were observing the shock breakout of the supernova for the first time \\citep{Campana,Waxman}. However, the non-thermal emission did not differ greatly from that of the X-ray flash class of soft GRBs and an outflow speed close to the speed of light could be inferred \\citep{Toma}, alternatively suggesting this to be an extension of the typical GRB population and not requiring significantly slower ejecta or any special (off-axis) geometry \\citep{Ghisellini}. \\cite{Toma} speculate that the long-duration, low-luminosity events such as this one may point to a different central engine for these events compared with more typical GRBs: a neutron star rather than a black hole \\citep[see also][]{Mazzali,Fan}.  Also of relevance to this discussion are the supernovae for which no GRB was detected, yet for which mildly relativistic ejecta provide a good explanation for the radio data: SN\\,2007gr \\citep{Paragi} and SN\\,2009bb \\citep{Soderberg}. As GRB jets are thought to be highly collimated, it is expected that many will go undetected if the emission is directed away from our line of sight \\citep[e.g.][]{Mazzali1}, and a fraction will also lie below current detector sensitivity limits. \\cite{Soderberg} estimate the fraction of Type Ib/c SNe with central engines at about one per cent, consistent with the inferred rate of nearby GRBs. In contrast, there are nearby GRBs for which an accompanying supernova could be expected but none has been detected to extremely deep limits. GRBs 060505 and 060614 \\citep{Fynbo,Dellavalle,Gal-Yam} are the two best examples of this, highlighting the need for a greater understanding of the relationship between GRBs and SNe.  We present the discovery of a new GRB-SN, GRB\\,100316D \\citep{Stamatikos} associated with SN\\,2010bh \\citep{Wiersema2,Bufano,Chornock}. This is an unusually long-duration, soft-spectrum GRB positioned on a nearby host galaxy. Many of the properties of this GRB appear unlike those of typical GRBs, resembling rather GRB\\,060218. In this paper we report accurate astrometry for the GRB-SN and the redshift of the underlying host galaxy. We examine the GRB prompt and afterglow emission observed with all 3 instruments on-board {\\it Swift}, and the broad characteristics of the host galaxy as observed with the Gemini South telescope, the Very Large Telescope (VLT) and the Hubble Space Telescope (HST). We compare these properties with the GRB sample as a whole and with a subset of GRB-SNe, to understand the origins of the GRB emission and further our knowledge of the GRB--SN connection. ",
        "conclusions": "\\label{sec:discuss}  GRB\\,100316D is an atypical gamma-ray burst both in its temporal and spectral behaviour. The very soft spectral peak and extended and slowly decaying flux emission are highly unusual among the prompt emission of GRBs \\citep[e.g.][]{Sakamoto}. The estimated total isotropic energy ($E_{\\rm iso} \\ge$~(5.9$\\pm$0.5)~$\\times$10$^{49}$ erg), when considered with the 90\\% confidence range of spectral peak values of 10--42~keV we measure in Section \\ref{sec:spec}, places GRB\\,100316D on the low $E_{\\rm iso}$--low $E_{\\rm pk}$ tail of the Amati correlation which relates $E_{\\rm pk}$ and $E_{\\rm iso}$ for the majority of typical GRBs \\citep{Amati1}. The notable outlier to this relation is GRB\\,980425 (SN\\,1998bw) \\citep[e.g.][]{Kaneko,Amati}. We also compare this GRB with the relation reported by \\cite{Gehrels2} between XRT flux at 11~h (here, (3.2$\\pm$0.8)$\\times$10$^{-13}$ erg cm$^{-2}$ s$^{-1}$) and BAT fluence and find, using the lower limit to the fluence, that it lies along the correlation with values that are typical for the sample of {\\it Swift} GRBs reported in \\cite{Evans}.  GRB\\,100316D is among the longest duration {\\it Swift} GRBs ($T_{\\rm 90}\\ge1300$~s), after GRB\\,090417B ($T_{\\rm 90}\\sim2130$~s) and GRB\\,060218 ($T_{\\rm 90}\\sim2100$~s). GRBs with durations of 400~s or longer, as quantified using the $T_{\\rm 90}$ parameter, account for only 1.4\\% of the {\\it Swift} sample observed to date. Owing to their long prompt emission durations, these GRBs are more likely to have exceptional broadband coverage, lending themselves to prompt emission studies. We show the X-ray light curve and hardness ratio of GRB\\,100316D in relation to the $T_{\\rm 90}>400$~s XRT-observed {\\it Swift} events in Figure \\ref{XRTlcs_long} (omitting GRB\\,060218 as it can be seen in Figure \\ref{XRTlcs_GRBSNe}; data are taken from the {\\it Swift} XRT GRB light curve repository\\footnote{\\url{http://www.swift.ac.uk/xrt_curves/}}, \\citealt{Evans1}).  This category comprises a number of different light curve types, from those which are long due to triggering on a precursor or the presence of a broad flare to the very gradually decaying prompt emission of GRB\\,100316D. The {\\it Swift} subsample shown here indicates two GRBs, in addition to GRB\\,060218, whose emission resembles GRB\\,100316D both in temporal decay and in spectral hardness: GRBs 070616 \\citep{Starling070616} and 090419 \\citep{Stratta}. We note that GRB\\,090417B is also spectrally hard at early times, but this may be attributed to strong X-ray flaring (Holland et al. in preparation).  GRB\\,070616 was particularly well sampled, with $\\gamma$-ray through optical coverage at early times clearly showing the movement of the spectral peak energy of the Band function \\citep{Band} and the occurrence of additional spectral softening \\citep{Starling070616}. However, the spectrum was much harder than that of 100316D - observed by both {\\it Swift} BAT and {\\it Suzaku} WAM, and neither the redshift nor the presence of any SN could be established as the source lay close to a bright star.  We now discuss the properties of GRB\\,100316D in the context of GRB-SNe, paying particular attention to GRB\\,060218.  \\subsection{Comparison to GRBs with accompanying supernovae: high energy emission} \\label{sec:cfGRBSNe} In Figure \\ref{XRTlcs_GRBSNe} we show the X-ray light curves and hardness ratios of a subset of the GRBs with associated SNe (either spectroscopically or photometrically confirmed). On this plot we also include GRB\\,051109B which was tentatively associated with a very low redshift \\citep[$z = 0.08$,][]{Perley} host galaxy and has no reported SN that we are aware of, and the well-studied nearby GRBs 060505 and 060614 for which SN searches to very deep limits registered no detections \\citep{Fynbo,Dellavalle,Gal-Yam}. The prompt emission shows a broader range of observed fluxes than the late-time decays, perhaps suggesting they are unrelated or that an extra component is present in the early emission which contributes differently in each GRB. The majority of the GRB-SNe (and GRBs without SNe shown here) light curves decay steeply over the first 1000~s, with GRBs 100316D and 060218 being the only exceptions. The hardness ratios reveal that GRB\\,100316D and GRB\\,060218 have a similar spectral hardness evolution, clearly different from those of most other GRB-SNe which remain at a stable and softer spectral shape throughout the prompt and late-time emissions. GRBs 060614 (without a SN) and 090618 (photometrically discovered SN) also transition from hard to soft, but do this earlier, over the first 200~s during their steep decays. We note, however, that X-ray data during the early emission for a number of the GRB-SNe, including the spectroscopically confirmed GRB-SNe 980425, 030329, 031203 and 050525A, are either not available or inadequate for spectral studies.  \\begin{figure*} \\begin{center} \\includegraphics[width=12cm, angle=-90]{fig7.ps} \\caption{X-ray light curves against time since trigger in flux density at 1.7 keV for a selection of long GRBs with associated supernovae (either spectroscopically confirmed - generally coloured in shades of red with GRB\\,100316D in black, or photometrically determined - shades of blue and green), two nearby long GRBs with no observed supernovae to deep limits and a further long GRB likely at $z\\le0.1$ with no reported supernova (grey shades). The lower panel shows the hardness ratios for the {\\it Swift} XRT-observed GRBs, derived from the 1.5--10 keV and 0.3--1.5 keV count rate light curves in the same colour scheme.\\newline References for the data sources: 100316D$^2$, 980425 \\citep{Pian}, 030329 \\citep{Willingale}, 031203 \\citep{Watson}, 050525A \\citep[][and $^2$]{Blustin}, 060218$^2$, 081007$^2$, 970228 \\citep{Costa,DePasquale}, 011121 \\citep{Piro}, 050824$^2$, 070419A$^2$, 080319B$^2$, 090618$^2$, 051109B$^2$, 060505$^2$, 060614$^2$.} \\label{XRTlcs_GRBSNe} \\end{center} \\end{figure*}  \\subsection{Comparison to GRBs with accompanying supernovae: host galaxies} \\label{sec:cfGRBSNehosts} Nearby GRBs show a large variety in their host galaxy properties: the prototype GRB-SN, GRB\\,980425 and SN\\,1998bw, occured in a dwarf spiral, in a small H\\,{\\sc II} region right next to a very large star formation complex; whereas GRB\\,060218 and SN\\,2006aj, occured in a very faint, blue compact dwarf galaxy \\citep{Wiersema1}. This diversity is illustrated in Figure \\ref{GRBSNehosts}.  Assuming that the properties of source A are at least roughly representative of the host galaxy as a whole, we can compare the spectroscopic properties with those of the other nearby GRB-SN hosts. Electron temperature based oxygen abundances have been derived for several of these, see Table \\ref{table:grbsnhosts}. The case of GRB\\,980425 is of particular interest: besides many similarities between SN\\,1998bw and SN\\,2010bh there are also several similarities in host galaxy properties. The $T_e$ metallicity measured at the location of the GRB and at the location of a nearby, bright, WR star region \\citep{Hammer} are clearly similar to the metallicity measured for the H\\,{\\sc II} region A in the host of GRB\\,100316D. In addition, the size and brightness (and possibly morphology) of these two hosts are comparable, though an important difference between the host of 980425 and 100316D is evident from Figure \\ref{GRBSNehosts}: the morphology of the host of 100316D is highly disturbed, possibly indicative of a recent merger.  The host of GRB\\,980425 has been studied using integral field spectroscopy \\citep{Christensen}, showing that the H\\,{\\sc II} region in which the GRB exploded is similar in mass, SFR, reddening and line equivalent width to other H\\,{\\sc II} regions in this host, with the exception of a bright WR-star rich H\\,{\\sc II} region 800 pc away from the GRB site. In terms of metallicity, the WR region and GRB site have a somewhat lower metallicity than other H\\,{\\sc II} regions in this host \\citep{Christensen}. If this situation is comparable to the host of GRB\\,100316D, we may expect that the properties measured above of source A are indeed representative of most of the host galaxy properties.  \\begin{figure*} \\begin{center} \\includegraphics[width=16cm, angle=0]{fig8.ps} \\caption{A mosaic of the host galaxies of spectroscopically confirmed supernovae, associated with GRBs, and imaged by the Hubble Space Telescope. The centre image shows the host of GRB\\,100316D, while other examples are shown in smaller insets. The physical scale across each image is approximately 7~kpc, and the positions of source A and of the SN are marked with crosshairs. Note, in the case of GRB\\,060218 and GRB\\,100316D there is still some contamination from the supernova at the time of the image.} \\label{GRBSNehosts} \\end{center} \\end{figure*}  \\begin{table*} \\caption{Host galaxies of nearby GRB-SNe. The absolute magnitudes $M_B$ were taken from the compilation of Levesque et al.~(2010).\\label{table:grbsnhosts}} \\begin{tabular}{@{}lllll@{}} GRB     & $T_e$ oxygen abundance             &  Host type      & Absolute magnitude & References\\\\ & (12 + log(O/H))                    &                 & $M_B$              &           \\\\ \\hline 980425  & 8.25 (GRB site)                    & Dwarf spiral    &  $-$17.6             &  Sollerman et al. (2005);         \\\\ & 8.39 (nearby WR region)            &                 &                    &   Hammer et al. (2006)                             \\\\ 020903  & 7.97                               & Irr             &  $-$18.8             & Hammer et al. (2006)          \\\\ 030329  & 7.72                               & Irr             &  $-$16.5             &  Levesque et al. (2010)  \\\\ 031203  & $8.02 \\pm 0.15$ (integrated)       & Irr             &  $-$21.0             & Prochaska et al. (2004) \\\\ 060218  & $7.54^{+0.16}_{-0.1}$  (integrated)& Irr             & $-$15.9              & Wiersema et al. (2007) \\\\ 100316D & $8.23 \\pm 0.15$ (source A)         & Spiral? Irr?      &   $-$18.8                 & This work \\\\ \\hline \\end{tabular} \\end{table*}  \\subsection{Comparison to GRB\\,060218} \\label{sec:cf060218} In Section \\ref{sec:cfGRBSNe} we found that among the GRB-SNe, there is a clear commonality in the X-ray behaviour of 100316D and 060218. Comparing the prompt emission spectra in Section \\ref{sec:spec}, we again find parallels in that both bursts seem to require a similar thermal component in addition to the typical GRB synchrotron emission for which the synchrotron peak is observed to move to lower energies with time. The time-evolving prompt spectral parameters are compared in Figure \\ref{Epkevolcf060218}. Both events have low isotropic equivalent energies, of order 4$\\times$10$^{49}$ erg, and very little flux above 50 keV.  These two events lie in apparently quite different environments: 060218 in a faint, compact dwarf \\citep{Wiersema1,Modjaz} and 100316D in a luminous, disturbed possibly spiral host galaxy. The metallicities are also different, with the host of 100316D likely being more metal-rich (though the metallicity is below Solar). A prompt optical component, as observed in 060218, could not be detected in 100316D, but conditions for its detection were far less optimal in this case due to the superposition on a brighter host galaxy and the possibility of a higher extinction along the line-of-sight. The higher intrinsic X-ray column density we measure for 100316D compared with similar modelling of 060218 does not necessarily imply a higher optical extinction \\citep{Watson2,Campana2}.  In summary, both 100316D and its predecessor 060218 are nearby ($z=0.059, 0.033$), long-duration ($T_{90} \\ge$1300~s,~2100~s), initially relatively constant in X-ray flux (Figure \\ref{XRTlcs_long}), spectrally soft (very few or no counts above 100 keV; low $E_{\\rm pk}$, Figure \\ref{Epkevolcf060218}), subenergetic (both have $E_{\\rm iso}\\sim 4\\times 10^{49}$ erg) GRBs with an associated SN. These two events show similar prompt emission properties and stand out among the GRB-SNe subsample considered here for their unusual X-ray evolution. However, their host galaxies are rather different in morphology and metallicity, with the host of 100316D more closely resembling the host of 980425.  The thermal X-ray component, with a luminosity, temperature and radius similar to that observed in 060218, dictates that the shock break-out of the supernova must be considered. The optical/UV thermal component observed in 060218 is not seen here: if the extinction is similar we would expect a shock breakout similar to 060218 to be two or more times fainter in 100316D and with the relatively brighter host galaxy it is perhaps no surprise that early optical emission is not detected with UVOT. The presence of a thermal component in the prompt $\\gamma$-ray emission of BATSE (Burst And Transient Source Experiment) GRBs, in addition to non-thermal emission, was proposed by \\cite{Ryde}. The discovery of thermal components in the {\\it Swift} XRT X-ray spectra shown here for 100316D and in \\cite{Campana} for 060218 will therefore be important to consider in future studies of GRB prompt emission. The dominant component of the high energy emission in 100316D remains the synchrotron-like non-thermal spectrum common to all types of GRB (with and without SNe) thought to originate in internal shocks in a relativistic jet. The long duration of the early X-ray emission is curious, and exceedingly rare, perhaps suggesting a greater reservoir of material is available to feed the central engine and prolong its activity.  The discovery of GRB\\,100316D and its associated supernova SN\\,2010bh and the analyses of its high energy emission and host galaxy properties presented here illustrates the diversity in GRB-SNe characteristics that must be understood if we are to fully appreciate the relationship between GRBs and core-collapse SNe."
    },
    "1004/1004.2258_arXiv.txt": {
        "abstract": "Recent X-ray observations of galaxy clusters suggest that cluster populations are bimodally distributed according to central gas entropy and are separated into two distinct classes: cool core (CC) and non-cool core (NCC) clusters. While it is widely accepted that AGN feedback plays a key role in offsetting radiative losses and maintaining many clusters in the CC state, the origin of NCC clusters is much less clear. At the same time, a handful of extremely powerful AGN outbursts have recently been detected in clusters, with a total energy $\\sim 10^{61}-10^{62}$ erg. Using two dimensional hydrodynamic simulations, we show that if a large fraction of this energy is deposited near the centers of CC clusters, which is likely common due to dense cores, these AGN outbursts can completely remove CCs, transforming them to NCC clusters. Our model also has interesting implications for cluster abundance profiles, which usually show a central peak in CC systems. Our calculations indicate that during the CC to NCC transformation, AGN outbursts efficiently mix metals in cluster central regions, and may even remove central abundance peaks if they are not broad enough. For CC clusters with broad central abundance peaks, AGN outbursts decrease peak abundances, but can not effectively destroy the peaks. Our model may simultaneously explain the contradictory (possibly bimodal) results of abundance profiles in NCC clusters, some of which are nearly flat, while others have strong central peaks similar to those in CC clusters. A statistical analysis of the sizes of central abundance peaks and their redshift evolution may shed interesting insights on the origin of both types of NCC clusters and the evolution history of thermodynamics and AGN activity in clusters. ",
        "introduction": "\\label{section:intro}  The hot gas in galaxy clusters emits prolifically in X-rays and has been extensively studied by X-ray telescopes {\\it Chandra} and {\\it XMM-Newton}, which reveal a striking bimodality in the properties of cluster cores. According to their core gas entropy, clusters are separated into two distinct classes: cool core (CC) and non-cool core (NCC) clusters. While the former have a low central gas entropy peaked at $S\\equiv k_{\\rm B}T/ n_{\\rm{e}}^{2/3} \\sim 15 \\, {\\rm keV \\, cm^{2}}$, the latter usually have high-entropy cores peaked at $S \\sim 150 \\, {\\rm keV \\, cm^{2}}$, and there is a distinct less-populated gap between $\\sim 30$ and $\\sim 50\\, {\\rm keV \\, cm^{2}}$ \\citep{cavagnolo09}. This bimodality also appears in cluster temperature profiles, which decrease significantly toward the center in CC clusters, but remain relatively flat in NCC clusters \\citep{sanderson06}. Observations also suggest that AGN activity and star formation in cluster central dominant galaxies are much more pronounced in CC clusters \\citep{burns90,cavagnolo08, rafferty08}.  The origin of this dichotomy has been studied by many authors. Galaxy clusters may naturally reach the CC state due to the interplay between radiative cooling and heating by active galactic nuclei (AGNs). It is widely thought that AGN outbursts play a key role in heating the intracluster medium (ICM), thus preventing cooling catastrophe (see \\citealt{mcnamara07} for a recent review). It was also shown that AGN may operate as a self-regulating feedback mechanism, which is essential in suppressing global thermal instability and thus in maintaining the ICM in the CC state \\citep{guo08b}. In contrast, the origin of NCC clusters is much less clear; three competing explanations have been proposed: \\begin{enumerate} \\item Mergers: Mergers are shown to significantly disturb cluster CCs and mix the ICM \\citep{ricker01,gomez02,ritchie02}. However, recent simulations by \\citet{poole08} find that CC systems are remarkably robust and can be disrupted only in direct head-on or multiple collisions; even so, the resulting warm core state is only transient. Cosmological simulations by \\citet{burns08} suggest that NCC clusters may form when they experience major mergers early in their evolution which destroy embryonic CCs. Note that clusters in these simulations suffer from the overcooling problem since they do not incorporate mechanisms (such as AGN feedback) to stop cooling catastrophe. Furthermore, the relatively low numerical resolution in these cosmological simulations ($15.6 \\, h^{-1} {\\rm kpc}$) may preclude firm conclusions about core structure and evolution.  \\item  Pre-heating: \\citet{mccarthy08} suggested that early pre-heating prior to cluster collapse could explain the formation of  CC/NCC systems, which receive lower/higher levels of pre-heating. A possible concern in such scenarios is whether one can pre-heat the ICM to a high adiabat and yet retain sufficient low entropy gas in lower mass halos to obtain a realistic galaxy population. Furthermore, observations by \\citet{rossetti10} find that most NCC clusters host regions with low-entropy, but relatively high metallicity gas, which are probably a signature of recent CC to NCC transformation, apparently inconsistent with both the scenarios proposed by \\citet{mccarthy08} and \\citet{burns08}.  \\item AGN outbursts: Strong AGN outbursts with $E_{\\rm agn}$ $\\sim 10^{61}$-$10^{62}$ erg have been detected in clusters, e.g., Hydra A \\citep{nulsen05a}, Hercules A \\citep{nulsen05} and MS0735.6+7421 (\\citealt{mcnamara05}; \\citealt{mcnamara09}). Recently in \\citet[hereafter GO09]{guo09}, we show that such strong AGN outbursts could bring a CC cluster to the NCC state, which can be stably maintained by conductive heating from the cluster outskirts. AGN outbursts may also drag magnetic field lines radially, thus enhancing thermal conduction, which could significantly heat CCs in high-temeprature clusters. \\end{enumerate}  The origin of NCC clusters may not necessarily explain the origin of the CC/NCC bimodality. To ensure the pronounced bimodality seen in observations, clusters need to stay in the NCC state for a duration at least comparable to the cooling time. In fact, many NCC clusters have a short cooling time ($\\sim 1$ Gyr; \\citealt{sanderson06}), and the NCC peak in the cluster central entropy distribution may not exist unless heating largely offsets radiative cooling for a relatively long time. While clusters may be maintained in the CC state by episodic AGN feedback, thermal conduction can stably keep clusters in the NCC state \\citep{guo08b}. This has been confirmed by numerical simulations in GO09, which further predicts that the CC/NCC dichotomy is more pronounced in higher-temperature clusters, due to the fact that the heating and stablizing effects of conduction decline with temperature.  In GO09, we adopted a simplified `effervescent model' for AGN heating \\citep{begelman01}, and performed one-dimensional (1D) calculations. X-ray cavities in this model are in pressure equilibrium with and rise buoyantly in the ICM, resulting in a very gentle AGN heating. Furthermore, the 1D model assumes spherical symmetry. However, real AGN outbursts produce jets and X-ray cavities in two opposite directions, which are by no means spherically symmetric. The formation and evolution of X-ray cavities are probably much more dynamic than assumed in the `effervescent model'. Shock waves have been detected in many clusters with X-ray cavities, and are usually thought to be a natural result of AGN energy released in the ICM. In this paper, we study the evolution of X-ray cavities formed as AGN energy is injected into the ICM, using two-dimensional (2D) hydrodynamical simulations. Our primary goal is to investigate if strong AGN outbursts can transform a CC cluster to the NCC state, and if so, how this transformation happens.  The metallicity distribution of the ICM, extensively measured with X-ray observations, contains important clues about the physics of galaxy clusters. The mass and distribution of metals in the ICM constrain the integrated history of past star formation (metals are released via supernovae explosions and winds) and the ICM enrichment processes (see \\citealt{bohringer10} for a recent review). The spatial abundance distribution is also significantly affected by transport processes in the ICM, e.g., turbulent mixing triggered by central AGN outbursts \\citep{rebusco05, rebusco06, roediger07}. Of great interest is the bimodality in central abundance profiles between CC and NCC clusters. CC clusters usually have a peak in the iron abundance profile at the cluster center, while many NCC clusters show a relatively flat radial abundance profile \\citep{degrandi01, degrandi04,leccardi10}. Nevertheless, recent observations by \\citet{leccardi08} and \\citet{sanderson09} suggest that some NCC clusters also have central iron abundance peaks. The differences in abundance profiles between CC and NCC systems have not been studied with detailed computational models. If NCC clusters are transformed from CC systems (i.e., not primordial), the same process transforming the CC to NCC systems may also explain the differences seen in their abundance profiles. In this paper we follow the evolution of iron abundance in our simulations and investigate if strong AGN outbursts can significantly change abundance profiles as they transform a CC cluster to the NCC state. We show that strong AGN outbursts efficiently mix metals in cluster central regions, but the final abundance profile in the NCC state strongly depends on the size of the initial abundance peak at the CC cluster center. Our model may naturally explain the the range of abundance profiles observed in NCC clusters. CC systems with broad/narrow central abundance peaks can be transformed by powerful AGN outbursts into NCC systems with/without central abundance peaks. A statistical analysis of the sizes of abundance peaks and their redshift evolution could test the validity and applicability of our model and may shed interesting insights on the origin of the CC/NCC bimodality.  The rest of the paper is organized as follows. In Section~\\ref{section2}, we present our model, including basic equations and numerical setup. Our results are presented in Section~\\ref{section:results}. We summarize our main results in Section~\\ref{section:conclusion} with a discussion of the implications. The cosmological parameters used throughout this paper are: $\\Omega_{m}=0.3$, $\\Omega_{\\Lambda}=0.7$, $h=0.7$. We have rescaled observational results if the original paper used a different cosmology. ",
        "conclusions": "\\label{section:conclusion}  X-ray observations of galaxy clusters indicate a striking bimodality in the properties of cluster cores, which separate clusters into two distinct classes: CC and NCC clusters. The origin of this bimodality remains unclear. At the same time, a handful of extremely powerful, core-changing AGN outbursts have recently been detected in clusters with a total energy $\\sim 10^{61}-10^{62}$ erg \\citep{nulsen05a, nulsen05, mcnamara05}. By conducting a suite of two-dimensional axisymmetric hydrodynamical simulations, we inquire if such strong AGN outbursts can transform a CC cluster to the NCC state, and if so, how this transformation happens. Cluster abundance profiles have been well observed by X-ray observations and contain important information on physical processes in clusters. Consequently, we also follow the evolution of iron abundance profiles and study how they are affected by AGN outbursts during the CC to NCC transformation.  We assume that AGN energy is injected into the ICM mainly in the form of CRs, and follow the co-evolution of these CRs  with the cluster gas in our simulations. We find that when a large fraction of this energy is injected near the cluster center, strong AGN outbursts with energy $\\sim 10^{61}-10^{62}$ erg (see Table 1) can completely remove cool cores, transforming the CC cluster (A1795 as our fiducial cluster) to the NCC state. In view of the high central gas density in CC clusters, deposition of AGN energy within the cluster core may be common. The deposition of CRs produces shocks in the ICM, which propagate outward and heat the core gas to high entropies ($\\sim 100$ kev cm$^{2}$) very quickly (within $\\sim 10$-$50$ Myr in our main run D1). The CRs also drive a strong gas outflow, producing a large, low-density cavity near the cluster center (with a radius of $\\sim 80-100$ kpc), which breaks up at time $t\\sim 200$ Myr. During the break-up of the cavity, high-entropy thermal gas flows back to the cluster center and is efficiently mixed. The cluster relaxes to the NCC state with a central entropy $\\sim 100$ kev cm$^{2}$ by time $t\\sim 200$-$300$ Myr.  Recent X-ray observations reveal several extremely powerful AGN outbursts with energy $\\sim 10^{61}-10^{62}$ erg, including Hydra A \\citep{nulsen05a}, Hercules A \\citep{nulsen05} and MS0735.6+7421 (\\citealt{mcnamara05}; \\citealt{mcnamara09}). MS0735.6+7421 hosts the most powerful AGN outburst currently known with an estimated AGN energy $\\sim 10^{62}$ erg, but most of the energy seems to be deposited beyond $40$ kpc (see \\citealt{guo10} for a detailed modeling of this outburst), and the CC in this cluster, though significantly heated, is not completely removed. In contrary, the temperature profile in Hydra A is quite flat, suggesting that the CC in this cluster may have been removed by the observed powerful AGN outburst. Furthermore, the observed ``Swiss-cheese-like\" topology associated with three pairs of X-ray cavities in Hydra A \\citep{wise07} is also well reproduced in our run D1-A, where CRs are injected into the ICM continuously by a (jet) source moving out from the central AGN. Hercules A hosts a young ($t\\sim 59$ Myr) AGN outburst ($\\sim 3 \\times 10^{61}$ erg), which has  significantly disturbed the cluster core \\citep{nulsen05}. While Hercules A is not fully relaxed, there is already some evidence suggesting that it may be a NCC cluster \\citep{white97}. Deep X-ray observations are needed to reveal its detailed thermal structure.  During the transformation of CC to NCC systems, AGN outbursts efficiently mix metals in cluster central regions (within $\\sim 50-60$ kpc $\\sim 0.03 r_{200}$ for A1795), where abundance peaks are usually seen in CC systems. The maximum abundance in the central peak thus decreases, but the peak is not removed if its spatial size ($r_{\\rm metal}$) is large. However, when $r_{\\rm metal}$ is comparable to or smaller than the characteristic radius within which AGN outbursts can efficiently mix metals (as in the cluster A2199), the central peak is effectively removed and the resulting NCC abundance profile is relatively flat. Our model naturally explains these two distinct types of NCC clusters seen in observations -- those with and without central abundance peaks -- both of which can develop following powerful AGN outbursts in CC clusters with different initial spatial abundance peak sizes.  In contrast, the merger scenario for creating NCC clusters may not explain those NCC systems without strong central abundance peaks. Simulations by \\citet{poole08} indicate that merging clusters which initially host central abundance peaks do not yield merger remnants with flat metallicity profiles (but also see \\citealt{vazza10}), though the dependence of their results on the size of the initial abundance peak remains to be explored. A statistical analysis of the peak sizes and their redshift evolution combined with a study of the possibly redshift-dependent relative fraction of both types of NCC clusters may shed interesting insights on the evolution history of the ICM thermal state and AGN activity in galaxy clusters."
    },
    "1004/1004.0893_arXiv.txt": {
        "abstract": "The methods of effective field theory are used to study generic theories of inflation with a single inflaton field and to perform a general analysis of the associated non-Gaussianities. We investigate the amplitudes and shapes of  the various generic three-point correlators, the   bispectra, which may be generated by different classes of single-field inflationary models.  Besides the well-known results for the DBI-like models and the ghost inflationary theories, we point out that  curvature-related interactions may give rise to large non-Gaussianities  in the form of bispectra characterized by a flat shape which, quite interestingly, is independently produced by several interaction terms. In a subsequent work, we will perform a similar general analysis for the non-Gaussianities generated by the generic four-point correlator,  the  trispectrum. ",
        "introduction": "The inflationary paradigm is central in modern cosmology as it naturally provides an explanation for  long standing issues as the flatness and horizon problems. The simplest standard, single-field slow-roll models of inflation already nicely account for the scale invariant primordial power spectrum of the Cosmic Microwave Background (CMB) anisotropies \\cite{smoot92,bennett96,gorski96,wmap3,wmap5,kom}. On the other hand, in order to get a deeper understanding of the inflaton dynamics, one needs to consider observables sensitive to deviations from Gaussianity, which might be studied in the form of the three-point correlator of cosmological perturbations, the so-called bispectrum \\cite{bisp,bispectrum,BabichCZ,chen-bis}, the connected four-point correlator, the so-called trispectrum \\cite{trispectrum}, loop corrections to the power spectrum \\cite{loop} and so on. In recent years, a lot of progress has been done in this direction resulting in the characterization of a wide variety of inflationary models through their higher order correlation functions~\\cite{reviewk}.  Complementing this theoretical effort, new experiments provide a better sensitivity to  deviations from Gaussian statistics (see, e.g.~\\cite{Liguori1,VerdeM}). The very recent launch of the Planck satellite~\\cite{Pl,Mand} (and the continued analysis of WMAP data~\\cite{kom}) presents us with the exciting opportunity to actually test this zoo of possibilities and makes more urgent the classification and refinement of the different predictions at the level of bispectrum and  trispectrum of curvature perturbations (see, e.g.,~\\cite{koyamarev,ourrev,chenrev} for comprehensive and updated reviews).   In this context, the effective theory approach proposed in \\cite{luty,eft08} and further analyzed in \\cite{w-e} represents a powerful tool. Indeed, the effective Lagrangian approach provides an efficient way to  generically study large classes of  inflationary models. Furthermore, this formalism  provides a  clear-cut dictionary between deviations from the standard scenario (almost Gaussian perturbations) and higher-order invariant operators in the action; it sheds new light on effects due to symmetries in the action (notably, a reduced speed of sound often automatically results in an enhanced non-Gaussianity); it allows some general and immediate calculational advantages, by exploiting the so called decoupling regime; finally, one can classify almost any specific inflationary model switching on some particular operators in the Lagrangian, thus allowing a more unifying approach. The effective field theory formalism naturally comprises Lagrangians with higher order terms. These terms call for an UV completion of the underlying theory. In this perspective one looks with particular interest at string theory inspired models, such as DBI infaltion \\cite{DBI}, and possible explanations in this context for k-inflation models \\cite{mukh1,mukh2}.  They are  retrievable in the effective theory  approach as the effective field theory that lives in a particular region of the parameters space.  In the spirit of effective theory, one should also include all possible curvature-generated terms ($\\delta R^\\mu_{\\,\\,\\,\\,\\mu}{}^2$,...) in writing the most generic effective theory  Lagrangian  of inflation up to the desired order in perturbation theory. In this paper we perform a thorough analysis of the amplitude and shapes of the three-point correlators which arise from the various classes of inflationary models driven by a single scalar field described within the generic effective theory approach. Besides reproducing the known results exisiting in the literature, that is to say that large non-Gaussianities may be generated if the sound speed is smaller than unity, we also show that single-field models of inflation may produce large non-Gaussianities in the form of three-point correlators whose amplitude is peaked in the so-called flat configuration, and precisely corresponding to a shape of the bispectrum where two of the momenta are roughly one half of the third one $k_1 =2 k_2=2 k_3$. Such bispectra emerge from curvature-generated interactions that have not been discussed so far in the literature. Let us recall that a shape of this type, peaked on flat triangles, was first found in \\cite{ssz05} where it was obtained as the result of a linear combination of two operators with proper weights \\footnote{It is important to note that these weights are not fixed, indeed there is an order one interval in \\cite{ssz05} that produces shapes compatible with the flat template.}. We stress here that our results further expand the space of operators that generates this kind of shape and allow one to easily obtain it from different, independent operators. Before the results of \\cite{ssz05}, the only scenario compatible with such a flat shape (also called folded or squashed) for the primordial bispectrum was that of models of inflation which relax the assumption of a regular Bunch-Davies vacuum, as studied in~\\cite{chen-bis,HT,Pier}.  In fact, up to \\cite{ssz05}, it was standard procedure to associate the primordial bispectrum evaluated in the Bunch-Davies vacuum  to the so-called equilateral or local shapes~\\cite{BabichCZ}.\\footnote{We refer to~\\cite{FShellard,Liguori2} for a detailed analysis of possible shapes of the primordial bispectra and to the discussion in Ref.~\\cite{ssz05} for an approach that includes the flat shape. (see also Ref.~\\cite{chen-bis}).} Here  we highlight that such a flat configuration for the primordial non-Gaussianities might be generated in single-field models of inflation in a more general context than previously thought. It is worth pointing out another aspect in which our results differ from previous ones. We are able to provide a solution for the equations of motion of the primordial perturbations that can interpolate between known cases, like DBI or Ghost inflation. We are able to study in  full generality these intermediate situations also when interaction terms are included. In a follow-up paper \\cite{future}, we will perform a similar analysis for the trispectrum.  The  paper is organized as follows. In section 2 we give a brief account on how to write the Lagrangian we employ in the following sections, following Ref.~\\cite{eft08}. The reader who is familiar with this approach might wish  to skip this part and jump  directly to the third-order Eq.~(\\ref{action}). We also comment on the inflationary models one can span within this approach and on the claim of its generality. The calculational algorithm used in this section can be extended to higher orders as well. In section 3 we derive the equation of motion for the scalar field which in the effective theory approach is used to describe the cosmological perturbations \\cite{tilted} and we provide an analytical solution to this general equation. In section 4 we present the calculation of the amplitude of the bispectrum  for each interaction term in six different configurations capturing the simplest single field slow roll model, DBI inflation, ghost inflation and intermediate cases. We focus on  regions in the parameter space where non-Gaussianity is  large. In section 5 we study  the shape of the bispectra arising from the  individual interaction terms identifying, among other things, interesting shapes from the  curvature-related contributions. We then summarize our findings and comment on further work in section 6. ",
        "conclusions": "The purpose of this paper was to employ a powerful tool such as effective field theory to obtain as general as possible a knowledge of primordial non-Gaussianities generated in a very general set-up of single field models of inflation. We have performed a study of the corresponding bispectra, in terms of both amplitudes and shapes.  We extended the results existing in the literature in different ways. First, we have improved the treatment of the wavefunction describing the behaviour of  the cosmological  perturbations at second order in perturbation theory. Secondly, we computed the amplitude and the shape of bispectra for theories which interpolate between the most studied ones, {\\it i.e.}, the DBI and the ghost models. Third, we have pointed out the importance of the curvature-induced operators. Their study has revealed that large non-Gaussianities may be generated with a flat shape which is quite uncommon for single-field models of inflation so far analyzed. It would be interesting to identify in  which (class of) theories of inflation such operators arise. A natural way to extend our findings is to study with the same philosophy the four-point correlator, the trispectrum, and to identify, for instance, those classes of inflationary models where the  bispectra are suppressed, but large trispectra are generated. These issues are currently under investigation \\cite{future}."
    },
    "1004/1004.2896_arXiv.txt": {
        "abstract": "By performing a series of two-dimensional, special relativistic magnetohydrodynamic (MHD) simulations, we study signatures of gravitational waves (GWs) in the magnetohydrodynamically-driven core-collapse supernovae. In order to extract the gravitational waveforms, we present a stress formula including contributions both from magnetic fields and special relativistic corrections.  By changing the precollapse magnetic fields and initial angular momentum distributions parametrically, we compute twelve models. As for the microphysics, a realistic equation of state is employed and the neutrino cooling is taken into account via a multiflavor neutrino leakage scheme. With these computations, we find that the total GW amplitudes show a monotonic increase after bounce for models with a strong precollapse magnetic field ($10^{12}$G) also with a rapid rotation imposed. We show that this trend stems both from the kinetic contribution of MHD outflows with large radial velocities and also from the magnetic contribution dominated by the toroidal magnetic fields that predominantly trigger MHD explosions. For models with weaker initial magnetic fields, the total GW amplitudes after bounce stay almost zero, because the contribution from the magnetic fields cancels with the one from the hydrodynamic counterpart. These features can be clearly understood with a careful analysis on the explosion dynamics. We point out that the GW signals with the increasing trend, possibly visible to the next-generation detectors for a Galactic supernova, would be associated with MHD explosions with the explosion energies exceeding $10^{51}$ erg. ",
        "introduction": "Successful detection of neutrinos from SN1987A paved the way for {\\it Neutrino Astronomy} \\citep{hirata,bionta}, alternative to conventional astronomy by electromagnetic waves. Core-collapse supernovae are now expected to be opening yet another astronomy, {\\it Gravitational-Wave Astronomy}. Currently long-baseline laser interferometers LIGO \\citep{firstligonew}, VIRGO$^1$\\footnotetext[1]{http://www.ego-gw.it/}, GEO600$^2$\\footnotetext[2]{http://geo600.aei.mpg.de/}, TAMA300 \\citep{tamanew}, and AIGO$^3$\\footnotetext[3]{http://www.gravity.uwa.edu.au/} with their international network of the observatories, are beginning to take data at sensitivities where astrophysical events are predicted (see, e.g., \\citet{hough} for a recent review). For these detectors, core-collapse supernovae have been proposed as one of the most plausible sources of gravitational waves (GWs) (see, e.g., \\citet{kota06,ott_rev} for recent reviews).  Although the explosion mechanism of core-collapse supernovae has not been completely clarified yet, current multi-dimensional simulations based on refined numerical models show several promising scenarios. Among the candidates is the neutrino heating mechanism aided by convection and standing accretion shock instability (SASI) (e.g., \\citet{marek,bruenn,sche04,fog,suwa}), the acoustic mechanism \\citep{burr06,burrows2}, and the magnetohydrodynamic (MHD) mechanism (e.g., \\citet{arde00,kota04b,kotake05,ober06b,burr07,taki09} and references therein). For the former two to be the case, the explosion geometry is expected to be unipolar and bipolar, and for the MHD mechanism to be bipolar.  Since the GW signatures imprint a live information of the asphericity at the moment of explosion, they are expected to provide us an important hint to solve the supernova mechanism. So far, most of the theoretical predictions of GWs have focused on the bounce signals in the context of rotational core-collapse (e.g., \\citet{mm,zweg,kotakegw,shibaseki,ott,ott_prl,ott_2007,dimm02,dimmelprl,dimm08,simon}). For the bounce signals having a strong and characteristic signature, the iron core must rotate enough rapidly. The waveforms are categorized into the following three types, namely types I, II, and III. Type II and III waveforms are shown less likely to appear than type I, because a combination of general relativity (GR) and electron capture near core bounce suppresses multiple bounce in the type II waveforms \\citep{dimmelprl,ott_prl,ott_2007}. In general, a realistic nuclear equation of state (EOS) is stiff enough to forbid the type III waveforms. After bounce, asymmetries due to convection \\citep{burohey,muyan97,fryersingle,mueller04}, SASI \\citep{marek_gw,kotake07,kotake09,kotake_ray,murphy}, and g-mode oscillations of protoneutron stars \\citep{ott_new}, are expected to account for sizable GW signals.  In general, detection of these GW signals in the postbounce phase (except for the g-mode oscillation) is far more difficult than the bounce signals, because they do not possess a clear signature like bounce signals, but change stochastically with time as a result of chaotically growing convection as well as SASI in the non-linear hydrodynamics (\\citet{kotake09,kotake11,marek_gw,murphy}).  Rapid rotation, necessary for the strong bounce signals, is likely to obtain $\\sim$ 1\\% of massive star population (e.g., \\citet{woos_blom}). However this can be really the case for progenitors of rapidly rotating metal-poor stars, which experience the so-called chemically homogeneous evolution \\citep{woos06,yoon}. The high angular momentum of the core as well as a strong precollapse magnetic field is preconditioned for the MHD mechanism, because the MHD mechanism relies on the extraction of rotational free energy of the collapsing core via magnetic fields. The energetic MHD explosions are receiving great attention recently as a possible relevance to magnetars and collapsars (e.g., \\citet{harikae_a,harikae_b} for collective references), which are presumably linked to the formation of long-duration gamma-ray bursts (GRBs) (e.g., \\citet{mesz06}).  Among the previous studies mentioned above, only a small portion of papers has been spent on determining the GW signals in the MHD mechanism \\citep{yama04,kota04a,ober06b,cerd07,shib06,scheid}. This may be because the MHD effects on the dynamics as well as their influence over the GW signals can be visible only for cores with precollapse magnetic fields over $B_{0} \\gtrsim 10^{12}$ G \\citep{ober06b,kota04a}. Considering that the typical magnetic-field strength of GRB progenitors is at most $\\sim 10^{11-12}$ G \\citep{woos06}, this is already an extreme situation. Interestingly in a more extremely case of $B_0 \\sim 10^{13}$ G, a secularly growing feature in the waveforms was observed \\citep{ober06a,shib06,scheid}. Moreover \\citet{ober06a} called a waveform as type IV in which quasi-periodic large-scale oscillations of GWs near bounce are replaced by higher frequency irregular oscillations. Some of these MHD simulations follow adiabatic core-collapse, in which a polytropic EOS is employed to mimic supernova microphysics. At this level of approximation, the bounce shock generally does not stall and a prompt explosion occurs within a few ten milliseconds after bounce. Therefore a main focus in these previous studies has been rather limited to the early postbounce phase ($\\lesssim$ several 10 ms). However, for models with weaker precollapse magnetic fields akin to the current GRB progenitors, the prompt shocks stall firstly in the core like a conventional supernova model with more sophisticated neutrino treatment (e.g., \\citet{burr07}). In such a case, the onset of MHD explosions, depending on the initial rotation rates, can be delayed till $\\sim 100$ ms after bounce \\citep{burr07,taki09}. There remains a room to study GW signatures in such a case, which we hope to study in this work.  In this study, we choose to take precollapse magnetic fields less than $10^{12}$ G based on a recent GRB-oriented progenitor models. By this choice, it generally takes much longer time after bounce than the adiabatic MHD models to amplify magnetic fields enough strong to overwhelm the ram pressure of the accreting matter, leading to the magnetohydrodynamically-driven (MHD, in short) explosions. Even if the speed of jets in MHD explosions is only mildly relativistic, Newtonian simulations are not numerically stable because the Alfv\\'{e}n velocity ($\\propto B/\\sqrt{\\rho}$) could exceed the speed of light unphysically especially when the strongly magnetized jets (i.e., large $B$) propagate to a stellar envelope with decreasing density ($\\rho$). To follow a long-term postbounce evolution numerically stably, we perform special relativistic MHD (SRMHD) simulations \\citep{taki09}, in which a realistic EOS is employed and the neutrino cooling is taken into account via a multiflavor neutrino leakage scheme. Note in our previous study of GWs in magneto-rotational core-collapse \\citep{kota04a} that we were unable to study properties of the GWs long in the postbounce phase because the employed Newtonian simulations quite often crashed especially in the case of strong MHD explosions. To include GR effects in this study, we follow a prescription in \\citet{ober06a} which is reported to capture basic features of full GR simulations quite well. By changing precollapse magnetic fields as well as initial angular momentum distributions parametrically, we compute twelve models. By doing so, we hope to study the properties of GWs in MHD explosions systematically and also address their detectability.  The paper opens up with descriptions of the initial models and numerical methods employed in this work (section \\ref{sec2}). Formalism for calculating the gravitational waveforms in SRMHD is summarized in section \\ref{sec3}. The main results are given in Section \\ref{sec4}. We summarize our results and discuss their implications in Section \\ref{sec5}. ",
        "conclusions": "}  By performing a series of two-dimensional SRMHD simulations, we studied signatures of GWs in the MHD-driven core-collapse supernovae. In order to extract the gravitational waveforms, we presented a stress formula including contributions both from magnetic fields and special relativistic corrections. By changing the precollapse magnetic fields and initial angular momentum distributions parametrically, we computed twelve models.  As for the microphysics, a realistic equation of state was employed and the neutrino cooling was taken into account via a multiflavor neutrino leakage scheme. With these computations, we found that the total GW amplitudes show a monotonic increase after bounce for models with a strong precollapse magnetic field ($10^{12}$G) also with a rapid rotation imposed. We pointed out that this trend stems both from the kinetic contribution of MHD outflows with large radial velocities and also from the magnetic contribution dominated by the toroidal magnetic fields that predominantly trigger MHD explosions. For models with weaker initial magnetic fields, the total GW amplitudes after bounce stay almost zero, because the contribution from the magnetic fields cancels with the one from the hydrodynamic counterpart. These features can be clearly understood with a careful analysis on the explosion dynamics. It was pointed out that the GW signals with an increasing trend, possibly visible to the next-generation detectors for a Galactic supernova, would be associated with MHD explosions exceeding $10^{51}$ erg.   \\begin{figure} \\begin{center} \\begin{tabular}{cc} \\resizebox{55mm}{!}{\\includegraphics{fig8left.eps}} & \\resizebox{55mm}{!}{\\includegraphics{fig8right.eps}} \\\\ \\end{tabular} \\caption{Snapshots at 33 ms after bounce for model B12X20$\\beta0.1$ in which the leakage scheme (left panel) or the $Y_e$ prescription (right panel) is employed, respectively. In each panel, density (logarithmic, left-half) and entropy (right-half) distributions are shown. The side length of each plot is 600x600(km).} \\label{fig11} \\end{center} \\end{figure}  \\begin{figure} \\begin{center} \\begin{tabular}{ccc} \\resizebox{50mm}{!}{\\includegraphics{fig9leftyep.eps}} & \\resizebox{50mm}{!}{\\includegraphics{fig9middleyep.eps}} & \\resizebox{50mm}{!}{\\includegraphics{fig9rightyep.eps}}\\\\ \\resizebox{50mm}{!}{\\includegraphics{fig9leftleak.eps}} & \\resizebox{50mm}{!}{\\includegraphics{fig9middleleak.eps}} & \\resizebox{50mm}{!}{\\includegraphics{fig9rightleak.eps}}\\\\ \\end{tabular} \\caption{Left and middle panels show magnetic pressure (red line) vs. ram pressure (blue line) for model B12X20$\\beta 0.1$ around 32 ms after bounce with the $Y_e$ prescription (top panels) or the leakage scheme (bottom panels) along the polar axis (left panel) or the equatorial plane (middle panel). Matter pressure is shown by green line as a reference. The right panels show velocity profiles along the pole near after the stall of the bounce shock (red lines). Both in the two different deleptonization schemes, the MHD-driven explosions are indeed obtained.} \\label{fig12} \\end{center} \\end{figure}  \\begin{figure}[htb] \\begin{center} \\begin{tabular}{c} \\resizebox{66mm}{!}{\\includegraphics{fig10gw.eps}} \\end{tabular} \\caption{Gravitational waveforms for model B12X1$\\beta1$ near bounce with different grid points. Three lines starting from 0 (at $t - t_{\\rm b}$ = -2 ms with $t_{\\rm b}$ being the epoch of bounce), correspond to models with different angular grid points (30 (green), 60(red), 120(blue)) while fixing the radial grid points to be 300. The bottom three lines are set to start from -100 in the GW amplitudes (just for convenience), and they correspond to models with different radial grid points (250 (pink), 300 (orange), 600 (brown)) while fixing the lateral grid points to be 60. Note that the fiducial set employed in this work is 300(r)x60($\\theta$).} \\label{fig14} \\end{center} \\end{figure}   Although the presented simulations have utilized the leakage scheme to approximate the deleptonization, it would be more accurate (especially before bounce) to employ a formula developed by \\citet{matthias05}, which was designed to fit 1D Boltzmann results. Figure \\ref{fig11} shows snapshots at around 33 ms after bounce for model B12X20$\\beta0.1$ in which the leakage scheme (left panel) or the $Y_e$ prescription (right panel) is employed, respectively. Note that ``G15'' is taken in our simulation until bounce among the parameter sets in \\citet{matthias05}\\footnote{The inner-core mass at bounce for the employed parameter set is 0.6 $M_{\\odot}$ for our non-rotating 25 $M_{\\odot}$ progenitor. This value is higher than that obtained in GR simulations (0.45-0.55 $M_{\\odot}$) in \\citet{dimm08}. This may be because the pseudo GR potential employed in this work underestimates the GR gravity, which could potentially lead to a large inner-core mass.} and is switched to the leakage scheme at the postbounce phase. As is shown, the shock revival also occurs for the model with the $Y_e$ prescription (right panel). Figure \\ref{fig12} depicts the magnetic pressure (red line) vs. ram pressure (blue line) along the polar axis (left panel) or the equatorial plane (middle) near the rebirth of the stalled shock for the model with the $Y_e$ prescription (top panels) or the leakage scheme (bottom panels). For the equator, the magnetic pressure is much less than the ram pressure (middle panel), while the magnetic pressure amplified by the field wrapping along the pole becomes as high as the ram pressure of the infalling material at the shock front, leading to the MHD-driven shock formation (see right panels). Regardless of the two different deleptoniaztion schemes, these important features associated with the MHD explosions are shown to be quite similar.  Now we mention a comparison between the obtained results and relevant MHD simulations. Model R4E1CF in \\citet{scheid} whose precollapse rotational parameter is $\\beta = 0.5$ \\% with a uniform rotation imposed, and whose initial poloidal magnetic field is set to $10^{12}$ G, is close to our model B12X20$\\beta0.1$. From their Figure 23, the jet propagates to $\\sim$ 300 km along the rotational axis at around $\\sim 18$ ms after bounce. In our counterpart model, the MHD-driven shock revives after around 30 ms after bounce, and it reaches to $300$ km at around $\\sim$ 10 ms, which is equivalent to $\\sim$ 40 ms after bounce. Considering that our model ($\\beta = 0.1$\\%) is a slower rotator than model R4E1CF ($\\beta = 0.5$\\%), the delay of the shock revival for our model seems reasonable. Model s20A1B5-D3M12 in \\citet{pablo} whose precollapse angular velocity is 4 rad/s (the rotational parameter should be close to $\\beta= 0.1$ \\%) with uniform rotation and whose initial poloidal magnetic field is $10^{12}$ G, is close to our model B12X20$\\beta0.1$. From their Figure 10, the MHD jet propagates to $\\sim$ 240 km at 51 ms after bounce. The employed EOS is the same as ours (the Shen EOS), while the deleptonization scheme taken in their study was the $Y_e$ formula \\citep{matthias05}. As mentioned above, the dynamics is rather close to our corresponding model. Among the computed models in \\citet{burr07}, model M15B12DP2A1H which has a precollapse angular velocity of $\\pi$ rad/s (the rotational parameter should be close to $\\beta = 0.06$\\%) with initial dipolar magnetic field of $10^{12}$ G, is close to our model B12X20$\\beta0.1$. The interval before the launch of the MHD shock for their model is 80 ms after bounce (e.g., their Table 1) is much later than our model (28 ms after bounce). This may be due to the larger initial angular momentum ($\\beta = 0.1 \\%$) assumed in our study.  Model A3B3G5-D3M13 in Obergaulinger et al. (2005) which has a rotational parameter of $\\beta=0.9$\\% with a differential rotation imposed (the radial cut off is 500 km) and the initial poloidal magnetic field is $10^{13}$ G, is closer to our model B12X1$\\beta0.1$. The MHD jet reaches to 500 km at around 7 ms after bounce, which is also the case of our counterpart model. As discussed above, our results are compatible to the ones obtained in the relevant foregoing results.   The major limitation of this study is the assumption of axisymmetry. Recently it was reported in three-dimensional (3D) MHD core-collapse simulations \\citep{simon,scheid} that the fast growth of the spiral SASI hinders the efficient amplification of the toroidal fields, which could suppress the formation of jets rather easily realized in 2D simulations. As a sequel of this study, we plan to investigate the 3D effects in SRMHD. Regarding a resolution dependence of our results, Figure \\ref{fig14} indicates that our standard resolution is adequate to follow the evolution of the computed models. However, it is not sufficient at all to capture the magneto-rotational instability (MRI, e.g., \\citet{balb98}). At least $10-100$ times finer mesh points are required for resolving the MRI \\citep{MRI}, which may require some adaptive-mesh-refinement treatment, a very important component that remains to be improved. If the MRI could be resolved, the increasing-type waveform could emerge also for models with weakly initial magnetic fields because a more efficient field amplification could be captured. Although the general relativistic effects were treated only by a very approximative way, we think that the general relativity should not drastically change our results qualitatively, because the central protoneutron stars will not collapse to a black hole during our simulation time as inferred from a simple argument of the compactness of the inner-core. As for the microphysics, the neutrino heating is not included in this study. However the inclusion of the neutrino heating may play a minor role in the waveforms, since the timescales before the neutrino-driven explosions set in (e.g., \\cite{marek,bruenn,suwa} and references therein) are much longer than the MHD explosions.  As one possible extension of this study, we think it interesting to study GWs from anisotropic neutrino radiation in the MHD models. Extrapolating the outcome obtained in previous studies (e.g., \\citet{mueller04,kotake_ray}), we anticipate that the total amplitudes become larger when we include the neutrino GWs. This is because the neutrino GWs may make a positive contribution since the neutrino emissions from the oblately deformed protoneutron stars become stronger toward the polar direction \\citep{jamu,kota03a,walder,ott_angle}. If this is really the case, the total GW amplitudes especially for the increasing type should be much larger, possibly making its detection more promising. Furthermore, the neutrino signals from MHD explosions, may have a sharp directional dependence through neutrino oscillations, reflecting the aspherical propagation of the shock to the stellar envelope \\citep{kawa}. Taking a correlation analysis between the neutrino and GW signals could help to reveal the hidden nature of the central engines. In fact, several observational proposals have been made in this direction recently \\citep{van,aso,leonor}. The MHD-driven core-collapse supernovae, albeit rather minor among typical type II supernova explosions, seem to still contain a number of rich research subjects."
    },
    "1004/1004.0555_arXiv.txt": {
        "abstract": "In order to study the origin of high-frequency quasi-periodic oscillations observed in X-ray binaries, \\citet{Kat04} suggested a resonant excitation mechanism of disk oscillations in deformed disks. In this paper, we investigate numerically, following his formulation, whether trapped g-mode oscillations in a warped disk, where the warp amplitude varies with radius, can be excited by this mechanism. For simplicity, we adopt Newtonian hydrodynamic equations with relativistic expressions for the characteristic frequencies of disks. We also assume that the accretion disk is isothermal. We find that the fundamental modes of trapped g-mode oscillations with eigenfrequencies close to the maximum of epycyclic frequency are excited. The intermediate oscillations found are isolated in a narrow region around the resonance radius. After varying some parameters, we find that the growth rate increases as the warp amplitude or the black hole spin parameter increases, while it decreases as the sound speed increases. ",
        "introduction": "\\label{sec:intro} The Rossi X-ray Timing Explorer (RXTE) satellite has detected  high-frequency quasi-periodic oscillations (QPOs) in X-ray binaries. They are kilohertz (kHz) and hectohertz (hHz) QPOs in neutron-star low-mass X-ray binaries, and high-frequency (HF) QPOs ($\\geq 100$ Hz) in black-hole X-ray binaries. HF QPOs occur at fixed frequencies and the appearance is correlated to the state of the sources, that is, they appear only in high-luminosity states where $L > 0.1 L_{E}$, with $L_{E}$ being the Eddington luminosity. They are suggested to be phenomena originating in a strong gravitational field, which occur in the innermost region of the relativistic disk. HF QPOs are regarded as being a powerful tool to explore the mass and spin of the central black hole, and also to explore the physical states of the innermost region of the accretion disk.  Various models have been proposed to explain HF QPOs. For example, \\citet{Abr01} and \\citet{Klu01} pointed out the importance of resonant processes as being the cause of HF QPOs. In their model, HF QPOs are the result of resonant couplings between the vertical and horizontal epicyclic oscillations at a particular radius. There is, however, uncertainty concerning the excitation process of the oscillation system as a whole.  \\authorcite{Kat04}\\ (\\yearcite{Kat04, Kat08a, Kat08b}) proposed a model where HF QPOs are regarded to be disk oscillations resonantly excited in deformed disks. The deformation considered is a warp or an eccentric disk deformation in the equatorial plane. An outline of the model is as follows. A non-linear coupling between a disk oscillation (hereafter, an original oscillation) and a deformed part of the disk (warp or eccentric deformation) causes some forced disk oscillations (hereafter, intermediate oscillations). The intermediate oscillations make a resonant coupling with the disk, and then feedback to the original oscillation. Since the nonlinear feedback process involves a resonance, the original oscillation is amplified or damped. There are two kinds of resonances, i.e., horizontal resonance (Lindblad resonance) and vertical resonance. When the deformation is a warp, the resonance that can excite p- and g-mode oscillations is the horizontal resonance (\\authorcite{Kat04}\\ \\yearcite{Kat04, Kat08a, Kat08b}). Hence, in this paper we consider the horizontal resonance alone. In the Keplerian disks this resonant process works only when the disk is general relativistic. That is, a non-monotonic radial distribution of epicyclic frequency is necessary for the appearance of resonance and for trapping of oscillations. Recently, \\citet{Fer08} generalized \\authorcite{Kat04}'s idea on the excitation of trapped g-mode oscillations and made detailed numerical calculations of the growth rates of the oscillations in the case of rotating black holes. Their results indicate that the coupling mechanism can provide an efficient excitation of trapped g-mode oscillations, provided that the global deformations reach the inner part of the disk with non-negligible amplitude.  In this paper, we consider nearly the same problem examined by \\citet{Fer08}, following the formulation by \\authorcite{Kat08a} \\ (\\yearcite{Kat08a,Kat08b}). The purpose of this paper is to make comparison of results obtained from two different approaches, \\authorcite{Fer08}'s fully numerical approach and \\authorcite{Kat08a}'s analytical one. We especially consider the case where the disk is warped. In section 2, we summarize the basic equations for original oscillations and intermediate oscillations, and an analytical expression for growth rate, which we use in later sections. In section 3, we present our numerical results, and section 4 is devoted to conclusions and a discussion. A detailed formulation of the nonlinear coupling terms is provided in Appendix. ",
        "conclusions": "\\subsection{Original Oscillations}  We assume an unperturbed disk to be steady and axisymmetric with a flow $\\textit{\\textbf{u}}_{0}$. Then, by using a displacement vector $\\boldsymbol{\\xi}$, a weakly nonlinear hydrodynamical equation for adiabatic perturbations is written as (\\cite{Ly67})  \\begin{equation} \\rho_{0}\\frac{\\partial^{2}\\boldsymbol{\\xi}}{\\partial t^{2}} + 2\\rho_{0} \\left( \\boldsymbol{u_{0} \\cdot \\nabla} \\right)\\frac{\\partial \\boldsymbol{\\xi}}{\\partial t}+ \\boldsymbol{L(\\xi )} = \\rho_{0}\\boldsymbol{ C(\\xi ,\\xi )}, \\label{eq:Kato1} \\end{equation} where $\\textit{\\textbf{L}}(\\boldsymbol{\\xi}$) is a linear Hermitian operator with respect to $\\boldsymbol{\\xi}$, given by \\begin{eqnarray} \\boldsymbol{L(\\xi )} &= &\\rho_{0}\\left( \\boldsymbol{u_{0} \\cdot \\nabla} \\right)\\left( \\boldsymbol{u_{0} \\cdot \\nabla} \\right) \\boldsymbol{\\xi} + \\rho_{0}\\left( \\boldsymbol{\\xi \\cdot \\nabla} \\right)(\\boldsymbol{\\nabla} \\psi_{0})+ \\boldsymbol{\\nabla} \\left[(1-\\Gamma_{1}) p_{0}\\mathrm{div} \\boldsymbol{\\xi} \\right] \\nonumber \\\\ && - p_{0} \\boldsymbol{\\nabla} (\\mathrm{div} \\boldsymbol{\\xi} ) - \\boldsymbol{\\nabla} \\left[ \\left( \\boldsymbol{\\xi \\cdot \\nabla} \\right) p_{0} \\right] + \\left( \\boldsymbol{\\xi \\cdot \\nabla} \\right) (\\boldsymbol{\\nabla} p_{0}), \\label{eq: Kato2} \\end{eqnarray} where $\\rho_{0}$(\\textit{\\textbf{r}}) and $p_{0}$(\\textit{\\textbf{r}}) are the density and pressure in the unperturbed state, respectively, $\\Gamma_{1}$ is the barotropic index specifying the linear part of the relation between Lagrangian variations $\\delta p$ and $\\delta \\rho $, and $\\psi_{0}$ is a general-relativistic gravitational potential. The right-hand side of equation (\\ref{eq:Kato1}) represents the weakly nonlinear terms. For $\\Gamma_{1} = 1$ (isothermal perturbations), it is given by (\\cite{Kat04}) \\begin{equation} \\rho_{0}\\boldsymbol{C(\\xi ,\\xi )} = - \\frac{1}{2} \\rho_{0} \\xi_{i} \\xi_{j} \\frac{\\partial^{2}}{\\partial r_{i} \\partial r_{j}}(\\boldsymbol{\\nabla} \\psi_{0}) - \\frac{\\partial}{\\partial r_{i}} \\left( p_{0} \\frac{\\partial \\xi_{i}}{\\partial r_{j}} \\boldsymbol{\\nabla} \\xi_{j} \\right). \\label{eq:Kato3} \\end{equation}  We write the displacement vector $\\boldsymbol{\\xi}$ in the form of a normal mode, \\begin{equation} \\boldsymbol{\\xi}(r,t) = \\hat{\\boldsymbol{\\xi}}(r,z) \\exp \\left[ i(\\omega t - m\\varphi ) \\right] , \\label{eq:disp} \\end{equation} where $\\omega$ is the frequency and $m$ is the azimuthal wavenumber. In the case where the disk is isothermal in the vertical direction and the perturbations have a short radial wavelength compared with the radial characteristic length of the unperturbed disks, we can separate $\\boldsymbol{\\xi}$ into $\\textit{r}$-, $\\varphi$-, and $\\textit{z}$-components as (\\cite{Oka87})  \\begin{equation} \\hat{\\xi}_{r} (r,z) = \\breve{\\xi}_{r,n} (r) \\mathcal{H}_{n} (z/H), \\label{eq: Kato20} \\end{equation} \\begin{equation} \\hat{\\xi}_{\\varphi} (r,z) = \\breve{\\xi}_{\\varphi ,n} (r) \\mathcal{H}_{n} (z/H), \\label{eq: Kato21} \\end{equation} \\begin{equation} \\hat{\\xi}_{z} (r,z) = \\breve{\\xi}_{z,n} (r) \\mathcal{H}_{n-1} (z/H), \\label{eq: Kato22} \\end{equation} where $n$ is a non-negative integer characterizing the number of node(s) of the oscillation in the vertical direction and $\\mathcal{H}_{n}$ is the Hermite Polynomial of argument $z/H$, with $H$ being the vertical scale-height of the disk, which is given by $H = c_{s}/\\Omega_{\\perp} $, where $c_{s}$ is the sound speed and $\\Omega_{\\perp}$ is the vertical epicyclic frequency.  We express the $r$-, $\\varphi$-, and $z$- components of the homogeneous parts of equation (\\ref{eq:Kato1}) as  \\begin{equation} \\left[ -(\\omega - m\\Omega)^{2} + \\kappa^{2} -4\\Omega^{2} -c_{s}^{2} \\frac{d^{2}}{dr^{2}} \\right] \\breve{\\xi}_{r,n} - i2\\Omega (\\omega - m\\Omega) \\breve{\\xi}_{\\varphi ,n} + \\Omega_{\\perp}^{2} H \\frac{d\\breve{\\xi}_{z,n}}{dr} = 0 , \\label{eq: Kato23} \\end{equation} \\begin{equation} -(\\omega - m\\Omega)^{2} \\breve{\\xi}_{\\varphi ,n}+ i2\\Omega (\\omega - m\\Omega) \\breve{\\xi}_{r,n} = 0 , \\label{eq: Kato24} \\end{equation} \\begin{equation} \\left[ -(\\omega - m\\Omega)^{2} + n\\Omega_{\\perp}^{2} \\right] \\breve{\\xi}_{z,n} -  n\\Omega_{\\perp}^{2} H \\frac{d\\breve{\\xi}_{r,n}}{dr} = 0 . \\label{eq: Kato25} \\end{equation} Note that these equations are the same as equations (23) - (25) of \\citet{Kat08b}. From these equations we have the basic equation describing the original oscillation (\\cite{Oka87}): \\begin{equation} \\frac{d^{2}\\breve{\\xi}_{r,n}}{dr^{2}} = \\frac{\\left[ (\\omega -m\\Omega)^{2} - \\kappa^{2} \\right] \\left[ n\\Omega _{\\perp}^{2} - (\\omega -m\\Omega)^{2} \\right]}{\\Omega _{\\perp}^{2}H^{2} (\\omega -m\\Omega)^{2}} \\breve{\\xi}_{r,n}, \\label{eq:origosc} \\end{equation} where $\\Omega $ is the (Keplerian) angular frequency of disk rotation, and $\\kappa$ is the horizontal epicyclic frequency, respectively. In the general relativistic potential around a rotating black hole, these frequencies are given by (e.g., \\cite{Kat90}) \\begin{equation} \\Omega = (r^{3/2}+a)^{-1}, \\label{eq: Fer15} \\end{equation} \\begin{equation} \\kappa = \\Omega \\sqrt{1-\\frac{6}{r}+\\frac{8a}{r^{3/2}}-\\frac{3a^{2}}{r^{2}}}, \\label{eq: Fer16} \\end{equation} \\begin{equation} \\Omega_{\\perp} = \\Omega \\sqrt{1-\\frac{4a}{r^{3/2}}+\\frac{3a^{2}}{r^{2}}}, \\label{eq: Fer17} \\end{equation} where $a \\ (-1\\leq a\\leq 1)$ is the central star spin parameter, and $r$ is in units of the gravitational radius $R_{g}$, where $R_{g} = GM/c^{2}$ with $M$ being the mass of the central star.  \\subsection{Warp and Intermediate Oscillations}  As mentioned in section 1, our resonant excitation model for the high-frequency QPOs assumes the disk to be warped or to have an eccentric deformation. In this paper, we consider the case where the disk is warped. The Lagrangian displacement associated with the deformation, $\\boldsymbol{\\xi}^{W} (\\boldsymbol{r},t)$, is given by \\begin{equation} \\boldsymbol{\\xi}^{W} (\\boldsymbol{r},t) = \\exp (-i m_{W}\\varphi)\\hat{\\boldsymbol{\\xi}}^{W} (r,z), \\label{eq:Kato5} \\end{equation} where $m_{W} = 1$. \\begin{figure} \\centering \\FigureFile(6cm,6cm){warpgrafba5.ps}  \\caption{Warp function $W(r)$ for $c_{s}= 10^{-3}$, $W_{0}= 0.01$, and $a=0.5$. } \\label{fig:warpsol5} \\end{figure} The warp solution is described by $u_{r}^{W}=-W(r)\\Omega z$, $u_{\\varphi}^{W}= iW(r)z d(r\\Omega)/dr$, and $u_{z}^{W}= W(r)\\Omega r$, where $u_{r}^{W}$, $u_{\\varphi}^{W}$, and $u_{z}^{W}$ are the Eulerian perturbation of the velocity field and $W(r)$ is the solution of  \\begin{equation} \\frac{d}{dr} \\left(\\frac{\\Omega^{2}}{\\kappa^{2}-\\Omega^{2}}\\frac{dW}{dr} \\right) + \\frac{1}{r}\\frac{dW}{dr} = \\frac{\\Omega^{2}-\\Omega_{\\bot}^{2}}{c_{s}^{2}} W \\label{eq:solwarp} \\end{equation} (\\cite{Pap95}; see also \\cite{Fer08}). As the boundary condition, we specify the value of warp amplitude at the inner radius,$W_{0}$, and set $dW/dr = 0$ there, following \\citet{Fer08}. Here and hereafter, we take the inner radius to be the radius of the Innermost Stable Circular Orbit (ISCO), which is the radius where the horizontal epicyclic frequency, $\\kappa$, goes to zero. Figure \\ref{fig:warpsol5} shows a typical warp function $W(r)$, where $c_{s}= 10^{-3}$, $W_{0}= 0.01$, and $a=0.5$.  Using the warp function $W(r)$, we can write $\\hat{\\boldsymbol{\\xi}}^{W}$ as \\begin{equation} \\hat{\\xi}_{r}^{W} = - i W(r) z = -i W(r) H \\mathcal{H}_{1} (z/H), \\label{eq: warp1} \\end{equation} \\begin{equation} \\hat{\\xi}_{\\varphi}^{W} = - W(r) z = -W(r) H \\mathcal{H}_{1} (z/H), \\label{eq: warp2} \\end{equation} \\begin{equation} \\hat{\\xi}_{z}^{W} = i W(r) r = i W(r) r \\mathcal{H}_{0} (z/H). \\label{eq: warp3} \\end{equation}  Nonlinear coupling between the original oscillation, $\\boldsymbol{\\xi}$, characterized by $(\\omega,m$) and the deformation $\\boldsymbol{\\xi}^{W}$ characterized by $(0,1)$ introduces intermediate oscillations, $\\boldsymbol{\\xi}^{\\mathrm{int}}$, which can be expressed as \\begin{equation} \\boldsymbol{\\xi}^{\\mathrm{int}}_{\\pm} (\\boldsymbol{r},t) = \\exp \\left[ i(\\omega t - \\tilde{m} \\varphi) \\right] \\hat{\\boldsymbol{\\xi}}^{\\mathrm{int}}_{\\pm} (r,z), \\label{eq: Kato7} \\end{equation} where $\\tilde{m}=m+1$ or $m-1$. Here $\\hat{\\boldsymbol{\\xi}}^{\\mathrm{int}}_{+}$ represents the intermediate oscillations with $\\tilde{m} = m + 1$ resulting from the coupling between $\\hat{\\boldsymbol{\\xi}}$ and $\\hat{\\boldsymbol{\\xi}}^{W}$, while $\\hat{\\boldsymbol{\\xi}}^{\\mathrm{int}}_{-}$ represents those with $\\tilde{m} = m - 1$ resulting from the coupling between $\\hat{\\boldsymbol{\\xi}}$ and $\\hat{\\boldsymbol{\\xi}}^{W\\ast}$, where the asterix represents the complex conjugate. To consider these various coupling cases separately, we write $\\hat{\\boldsymbol{\\xi}}^{\\mathrm{int}}_{\\pm}$ in the following form: \\begin{equation} \\hat{\\xi}^{\\mathrm{int}}_{r,\\pm} (r,z) = \\breve{\\xi}^{\\mathrm{int}}_{r,\\pm ,\\tilde{n}} (r) \\mathcal{H}_{\\tilde{n}} (z/H), \\label{eq: Kato26} \\end{equation} \\begin{equation} \\hat{\\xi}^{\\mathrm{int}}_{\\varphi ,\\pm} (r,z) = \\breve{\\xi}^{\\mathrm{int}}_{\\varphi ,\\pm ,\\tilde{n}} (r) \\mathcal{H}_{\\tilde{n}} (z/H), \\label{eq: Kato27} \\end{equation} \\begin{equation} \\hat{\\xi}^{\\mathrm{int}}_{z,\\pm} (r,z) = \\breve{\\xi}^{\\mathrm{int}}_{z,\\pm ,\\tilde{n}} (r) \\mathcal{H}_{\\tilde{n}-1} (z/H), \\label{eq: Kato28} \\end{equation} where $\\tilde{n} = n + 1$ or $n-1$ for the coupling with the warp.  The nonlinear coupling terms are separated into terms proportional to $\\exp [i(\\omega t -\\tilde{m}\\varphi)]$ and $\\mathcal{H}_{\\tilde{n}}(z/H)$. In the case of coupling through $\\boldsymbol{\\xi}^{W}$, we write the coupling terms as \\begin{equation} \\frac{1}{2} \\rho_{0} \\left[ \\boldsymbol{C}(\\boldsymbol{\\xi}, \\boldsymbol{\\xi}^{W}) + \\boldsymbol{C}(\\boldsymbol{\\xi}^{W}, \\boldsymbol{\\xi}) \\right]_{r} =  \\rho_{0} \\sum_{\\tilde{n}} \\breve{A}_{r,+,\\tilde{n}} (r) \\exp \\left[ i(\\omega t - \\tilde{m}\\varphi) \\right] \\mathcal{H}_{\\tilde{n}} (z/H) + \\cdots , \\label{eq: Kato29} \\end{equation} \\begin{equation} \\frac{1}{2} \\rho_{0} \\left[ \\boldsymbol{C}(\\boldsymbol{\\xi}, \\boldsymbol{\\xi}^{W}) + \\boldsymbol{C}(\\boldsymbol{\\xi}^{W}, \\boldsymbol{\\xi}) \\right]_{\\varphi} =  \\rho_{0} \\sum_{\\tilde{n}} \\breve{A}_{\\varphi,+,\\tilde{n}} (r) \\exp \\left[ i(\\omega t - \\tilde{m}\\varphi) \\right] \\mathcal{H}_{\\tilde{n}} (z/H) + \\cdots , \\label{eq: Kato30} \\end{equation} \\begin{equation} \\frac{1}{2} \\rho_{0} \\left[ \\boldsymbol{C}(\\boldsymbol{\\xi}, \\boldsymbol{\\xi}^{W}) + \\boldsymbol{C}(\\boldsymbol{\\xi}^{W}, \\boldsymbol{\\xi}) \\right]_{z} =  \\rho_{0} \\sum_{\\tilde{n}} \\breve{A}_{z,+,\\tilde{n}} (r) \\exp \\left[ i(\\omega t - \\tilde{m}\\varphi) \\right] \\mathcal{H}_{\\tilde{n}-1} (z/H) + \\cdots , \\label{eq: Kato31} \\end{equation} where $\\tilde{m} = m+1$ and $+\\cdots$ denotes terms orthogonal to both $\\mathcal{H}_{\\tilde{n}}$ and $\\mathcal{H}_{\\tilde{n}-1}$. The subscript $+$ of $\\breve{A}$'s denotes that they are related to the $\\varphi$-dependence of $\\exp [ -i(m+1)\\varphi]$. Note that these equations are the same as equations (29) -- (31) of \\citet{Kat08b}. In the case of coupling through $\\boldsymbol{\\xi}^{W\\ast}$, the nonlinear coupling terms, $(1/2) \\rho_{0} [ \\boldsymbol{C}(\\boldsymbol{\\xi},\\boldsymbol{\\xi}^{W\\ast})+\\boldsymbol{C}(\\boldsymbol{\\xi}^{W\\ast},\\boldsymbol{\\xi})]$, can be expressed in similar forms, by introducing $\\breve{A}_{r,-,\\tilde{n}}$,$ \\breve{A}_{\\varphi ,-,\\tilde{n}}$, and $\\breve{A}_{z,-,\\tilde{n}}$ instead. Then, the basic equations describing intermediate oscillations are written as \\begin{eqnarray} &\\left[  -(\\omega - m\\Omega)^{2} + \\kappa^{2} -4\\Omega^{2} -c_{s}^{2} \\frac{d^{2}}{dr^{2}} \\right] & \\breve{\\xi}_{r,\\pm, \\tilde{n}}^{\\mathrm{int}} - i2\\Omega (\\omega - \\tilde{m}\\Omega) \\breve{\\xi}_{\\varphi ,\\pm, \\tilde{n}}^{\\mathrm{int}} \\nonumber \\\\ && +  \\Omega_{\\perp}^{2} H \\frac{d\\breve{\\xi}_{z,\\pm, \\tilde{n}}^{\\mathrm{int}}}{dr} = \\breve{A}_{r,\\pm,\\tilde{n}} , \\label{eq: Kato32} \\end{eqnarray} \\begin{equation} -(\\omega - \\tilde{m}\\Omega)^{2} \\breve{\\xi}_{\\varphi ,\\pm, \\tilde{n}}^{\\mathrm{int}}+ i2\\Omega (\\omega - \\tilde{m}\\Omega) \\breve{\\xi}_{r,\\pm, \\tilde{n}}^{\\mathrm{int}} = \\breve{A}_{\\varphi ,\\pm,\\tilde{n}} , \\label{eq: Kato33} \\end{equation} \\begin{equation} \\left[ -(\\omega - \\tilde{m}\\Omega)^{2} + \\tilde{n}\\Omega_{\\perp}^{2} \\right] \\breve{\\xi}_{z,\\pm, \\tilde{n}}^{\\mathrm{int}} - \\tilde{n}\\Omega_{\\perp}^{2} H \\frac{d\\breve{\\xi}_{r,\\pm, \\tilde{n}}^{\\mathrm{int}}}{dr} = \\breve{A}_{z,\\pm,\\tilde{n}}, \\label{eq: Kato34} \\end{equation} where both cases of ${\\tilde m}=m+1$ and ${\\tilde m}=m-1$ are written together.  We now eliminate $\\breve{\\xi}_{\\varphi,\\pm,{\\tilde n}}^{\\rm int}$ and $\\breve{\\xi}_{z,\\pm,{\\tilde n}}^{\\rm int}$ from the set of equations (\\ref{eq: Kato32}) -- (\\ref{eq: Kato34}) to derive an equation with respect to $\\breve{\\xi}_{r,\\pm,{\\tilde n}}^{\\rm int}$. By using the approximation that the radial wavelength of oscillations, $\\lambda$, is shorter than $r$, i.e., $\\lambda\\ll r$, which has been used in deriving equations (\\ref{eq: Kato23}) -- (\\ref{eq: Kato25}), we have \\begin{eqnarray} &&\\left[ - (\\omega -\\tilde{m} \\Omega )^{2} + \\kappa^{2} \\right] \\breve{\\xi}^{\\mathrm{int}}_{r,\\pm ,\\tilde{n}} - \\frac{c_{s}^{2}(\\omega-{\\tilde m}\\Omega)^2}{(\\omega-{\\tilde m}\\Omega)^2-{\\tilde n}\\Omega_\\bot^2} \\frac{d^{2}\\breve{\\xi}^{\\mathrm{int}}_{r,\\pm ,\\tilde{n}}}{dr^{2}} \\nonumber  \\\\ &&    = \\breve{A}_{r,\\pm ,\\tilde{n}} - i \\frac{2\\Omega}{\\omega - \\tilde{m} \\Omega} \\breve{A}_{\\varphi,\\pm ,\\tilde{n}} +\\frac{\\Omega_\\bot^2H}{(\\omega-{\\tilde m}\\Omega)^2-{\\tilde n}\\Omega_\\bot^2}\\frac{d{\\breve A}_{z,\\pm,{\\tilde n}}}{dr}. \\label{eq:Kato35} \\end{eqnarray} The horizontal resonance occurs at the radius where $-(\\omega-{\\tilde m}\\Omega)^2+\\kappa^2=0$ holds. After obtaining $\\breve{\\xi}^{\\mathrm{int}}_{r,\\pm ,\\tilde{n}}$ by solving this equation, we can obtain $\\breve{\\xi}^{\\mathrm{int}}_{\\varphi,\\pm ,\\tilde{n}}$ and $\\breve{\\xi}^{\\mathrm{int}}_{z,\\pm ,\\tilde{n}}$ from equations (\\ref{eq: Kato33}) and (\\ref{eq: Kato34}), respectively.  It is noted that in the coupling through ${\\tilde n}=n-1$ in warped disks ($n=1$), which is the case to be considered below, we have $\\breve{A}_{z,\\pm ,\\tilde{n}}=0$ (see the next subsection) and $\\breve{\\xi}^{\\mathrm{int}}_{z,\\pm ,\\tilde{n}}=0$.  \\subsection{Expressions for Coupling Terms}  Detailed expressions for the coupling terms, i.e., the right-hand side of equations (\\ref{eq: Kato32}) -- (\\ref{eq: Kato34}), are very complicated (see Appendix for details). Their expressions, however, can be simplified by using the approximations of $H\\ll \\lambda \\ll r$. In the case of ${\\tilde n}=n-1$, after a lengthy manipulation we have \\begin{eqnarray} \\breve{A}_{r,\\pm ,n-1}= \\left[ - \\frac{d\\Omega_{\\bot}^{2}}{dr} \\breve{\\xi}_{z,n} - \\Omega_{\\bot}^{2} H  \\frac{d}{dr} \\left(  \\breve{\\xi}_{r,n} \\frac{d}{dr} \\right) + (n-1) \\Omega_{\\bot}^{2} \\breve{\\xi}_{z,n} \\frac{d}{dr} \\right] \\left( \\begin{array} {cc} \\breve{\\xi}_{z}^{W} \\\\ \\breve{\\xi}_{z}^{W\\ast} \\end{array} \\right), \\label{eq:Ar+} \\end{eqnarray} \\begin{eqnarray} \\breve{A}_{\\varphi,\\pm ,n-1}=  \\left[ \\pm i n \\Omega_{\\bot}^{2} \\frac{H}{r} \\left( \\frac{d\\breve{\\xi}_{r,n}}{dr} + \\breve{\\xi}_{r,n}  \\frac{d}{dr} \\right) \\mp i \\frac{n-1}{r} \\Omega_{\\bot}^{2} \\breve{\\xi}_{z,n}  \\right] \\left( \\begin{array} {cc} \\breve{\\xi}_{z}^{W} \\\\ \\breve{\\xi}_{z}^{W\\ast} \\end{array} \\right) , \\label{eq:Aphi+} \\end{eqnarray} \\begin{eqnarray} \\breve{A}_{z,\\pm ,n-1} = \\mathrm{small} . \\label{eq:Az+} \\end{eqnarray} Here, $\\breve{A}_{z,\\pm,n-1}$ is small in the sense that $H (d\\breve{A}_{z,\\pm,n-1}) / dr$ is negligible compared with $\\breve{A}_{r,\\pm,n-1}$ in magnitude.  Similarly, in the case of ${\\tilde n}=n+1$, we have \\begin{eqnarray} \\breve{A}_{r,\\pm ,n+1} = \\left[ \\Omega_{\\bot}^{2} H \\frac{d \\breve{\\xi}_{r,n}}{dr} \\frac{d}{dr} \\mp \\frac{i}{r} \\Omega_{\\bot}^{2} H \\frac{d \\breve{\\xi}_{\\varphi ,n}}{dr}\\right] \\left( \\begin{array} {cc} \\breve{\\xi}_{z}^{W} \\\\ \\breve{\\xi}_{z}^{W\\ast} \\end{array} \\right) , \\label{eq:Ar-} \\end{eqnarray} \\begin{equation} \\breve{A}_{\\varphi,\\pm ,n+1} = \\mathrm{smaller \\  than\\ } \\breve{A}_{r, \\pm, n+1} \\mathrm{\\ in \\ magnitude} \\label{eq:Aphi-} \\end{equation}  \\begin{eqnarray} \\breve{A}_{z,\\pm ,n+1}= \\left[ - \\frac{d\\Omega_{\\bot}^{2}}{dr} \\breve{\\xi}_{r,n} + n \\Omega_{\\bot}^{2} \\breve{\\xi}_{r,n} \\frac{d}{dr} \\mp i n \\frac{\\Omega_{\\bot}^{2}}{r} \\breve{\\xi}_{\\varphi ,n} \\right] \\left( \\begin{array} {cc} \\breve{\\xi}_{z}^{W} \\\\ \\breve{\\xi}_{z}^{W\\ast} \\end{array} \\right). \\label{eq:Az-} \\end{eqnarray}  It is noted that in the above expressions, ${\\breve \\xi}_z^{\\rm W}$ is adopted for $\\breve {\\mbox{\\boldmath $A$}}_{+,{\\tilde n}}$, while ${\\breve \\xi}_z^{\\rm W *}$ is adopted in the case of $\\breve {\\mbox{\\boldmath $A$}}_{-,{\\tilde n}}$, where the asterisk denotes the complex conjugate.  \\subsection{Growth Rate}  By the feedback process through the intermediate oscillations, the original oscillation is amplified or damped, so the frequency $\\omega$ can no longer be real. \\citet{Kat08b} showed that the growth rate $-\\omega_{\\rm i}$, where $\\omega_{\\rm i}$ is the imaginary part of $\\omega$, can be written as  \\begin{equation} -\\omega_{{\\rm i},\\pm } = \\frac{W_{\\pm }}{2E}, \\label{eq: Kato12} \\end{equation} where $E$ is the wave energy of the original oscillation, $W_{\\pm}$ is the rate at which work is done on the original oscillation by the nonlinear resonant process, and $\\pm$ denotes the cases of the coupling through $\\breve{\\mbox{\\boldmath $\\xi $}}_{+}^{\\rm int}$ and $\\breve{\\mbox{\\boldmath $\\xi $}}_{-}^{\\rm int}$, respectively. Here, $W_{\\pm}$ are written as follows,  \\begin{eqnarray} W_{+}&=\\frac{\\omega_{0}}{2} \\Im \\int \\frac{1}{2} \\rho_{0} \\hat{\\boldsymbol{\\xi}}^{\\ast}& \\left[ \\boldsymbol{C}(\\hat{\\boldsymbol{\\xi}}_{+}^{\\rm int},\\hat{\\boldsymbol{\\xi}}^{W\\ast}) + \\boldsymbol{C}(\\hat{\\boldsymbol{\\xi}}^{W\\ast}, \\hat{\\boldsymbol{\\xi}}_{+}^{\\rm int}) \\right] dV, \\nonumber\\\\ &= \\frac{\\omega_{0}}{2}\\Im\\int\\rho_{00}(r)&[(2\\pi)^{3/2}\\tilde{n}!rH] \\nonumber \\\\ && \\times \\left[\\breve{\\xi}_{r,+,\\tilde{n}}^{\\rm int}\\breve{A}_{r,+,\\tilde{n}}^{\\ast}+ \\breve{\\xi}_{\\varphi ,+,\\tilde{n}}^{\\rm int}\\breve{A}_{\\varphi ,+,\\tilde{n}}^{\\ast} + \\frac{1}{\\tilde{n}} \\breve{\\xi}^{\\rm int}_{z,+,\\tilde{n}} \\breve{A}_{z ,+,\\tilde{n}}^{\\ast} \\right]dr, \\label{eq: Kato1336} \\end{eqnarray}  \\begin{eqnarray} W_{-}&= \\frac{\\omega_{0}}{2} \\Im \\int \\frac{1}{2} \\rho_{0} \\hat{\\boldsymbol{\\xi}}^{\\ast} & \\left[ \\boldsymbol{C}(\\hat{\\boldsymbol{\\xi}}_{-}^{\\rm int},\\hat{\\boldsymbol{\\xi}}^{W}) + \\boldsymbol{C}(\\hat{\\boldsymbol{\\xi}}^{W}, \\hat{\\boldsymbol{\\xi}}_{-}^{\\rm int}) \\right]dV \\nonumber\\\\ & =  \\frac{\\omega_{0}}{2}\\Im\\int\\rho_{00}(r) & [(2\\pi)^{3/2}\\tilde{n}!rH]\\nonumber \\\\ &&  \\times \\left[\\breve{\\xi}_{r,-,\\tilde{n}}^{\\rm int}\\breve{A}_{r,-,\\tilde{n}}^{\\ast}+\\breve{\\xi}_{\\varphi ,-,\\tilde{n}}^{\\rm int}\\breve{A}_{\\varphi ,-,\\tilde{n}}^{\\ast} + \\frac{1}{\\tilde{n}} \\breve{\\xi}^{\\rm int}_{z,-,\\tilde{n}} \\breve{A}_{z ,-,\\tilde{n}}^{\\ast}\\right]dr, \\label{eq: Kato1437} \\end{eqnarray} where $\\omega_{0}$ is the frequency of the original oscillation before the mode couplings, and the wave energy, $E$, is given by \\begin{eqnarray} E &=& \\frac{1}{2} \\omega_{0} \\int \\rho_{0} \\hat{\\xi}^{\\ast} \\left[ \\omega - i(\\boldsymbol{u \\cdot \\nabla}) \\right] \\hat{\\xi}dV \\nonumber \\\\ && = \\frac{(2\\pi)^{3/2}}{2} \\omega_{0}^{2} (r^{4} H \\rho_{00})_{c} E_{n}, \\label{eq: Kato1538} \\end{eqnarray} where $E_{n}$ is a dimensionless quantity given by \\begin{equation} E_{n} = \\int \\frac{ rH \\rho_{00}}{(rH \\rho_{00})_{c} } \\frac{\\omega_{0} -m\\Omega }{\\omega_{0} } \\left( n! \\frac{\\mid \\breve{\\xi}_{r,n}\\mid^{2}}{r_{c}^{3}} + (n-1)! \\frac{\\mid \\breve{\\xi}_{z,n}\\mid^{2}}{r_{c}^{3}} \\right) dr \\label{eq: Kato39} \\end{equation} with subscript c denoting the values at the resonant radius, where $-(\\omega_{0}-\\tilde{m}\\Omega)^{2} + \\kappa^{2} = 0$.    We have numerically studied the excitation of trapped g-mode oscillations in warped disks around black holes, based on a scenario proposed by \\authorcite{Kat04}\\ (\\yearcite{Kat04, Kat08a, Kat08b}). We first obtained trapped g-mode oscillations with eigenfrequencies close to the maximum of the horizontal epicyclic frequency. Then, we examined whether these modes are excited via the resonant coupling with the warp of the disk. We have found that the fundamental modes of the trapped g-modes are excited except in the case of non-rotating black holes. We have performed calculations of the growth rate of the fundamental mode for some parameter values such as the warp amplitude in the inner radius, the isothermal sound speed in the disk and the black hole spin parameter. We have found that the growth rate increases as the warp amplitude and spin parameter increases or the sound speed decreases. The dependence of the growth rate on these parameters qualitatively agrees with that obtained by \\citet{Fer08} for the fundamental mode, but quantitatively there are significant differences between the two. We suspect that these differences are due to differences of the model considered.  Our results presented in this paper show that the excitation of trapped g-mode oscillations by global disk deformation is rather sensitive to the disk structure, because the resonant radius given by $\\omega_{0}=\\Omega-\\kappa$ and a characteristic radius of the oscillation where ${\\breve \\xi}_{r,n}$ or $\\partial{\\breve \\xi}_{r,n}/\\partial r$ has a peak are close. We also suppose that the damping (or amplification) of the intermediate oscillations at the corotation resonance is not an essential ingredient for the excitation of disk oscillations in deformed disks. In the next paragraphs, we discuss the excitation mechanism from a viewpoint of energy flow among oscillations.  Energy flow between original and intermediate oscillations occurs at two aspects of the resonance. The first occurs at the stage where the oscillation resulting from nonlinear coupling between the original oscillation and the warp acts as a forcing source on the intermediate oscillation. The second occurs at the feedback process from the intermediate oscillation to the original oscillation where the intermediate oscillation acts as a source to the original oscillation after coupling with the warp. The net energy input rate to the original oscillation through these coupling processes is given by equation (\\ref{eq: Kato12}). The results derived from this equation show that in a simplified situation the signs of wave energy of original and intermediate oscillations must be opposite for excitation of the original oscillation (e.g., \\cite{Kat08a}). This suggests that a physical cause of resonant excitation of disk oscillations in deformed disks is an interaction of two waves with different signs of the wave energy at a resonant radius (\\authorcite{Kat08a}\\ \\yearcite{Kat08a} \\yearcite{Kat08b}: \\cite{Fer08}). In the present case, the original oscillation is axisymmetric and has positive energy, i.e., $E>0$. Hence, a necessary condition for wave amplification to occur is a negative wave energy of the intermediate oscillation. The major process of excitation is thus a positive energy flow from the intermediate oscillation with negative energy to the original oscillation with positive energy at the resonant region. By this energy flow, both original and intermediate oscillations grow.  The main role of the warp will be a kind of catalyzer for energy transport. In equation (\\ref{eq: Kato12}), energy exchange between the oscillation and the disk in the resonant region will be involved, but it will be subsidiary in the present excitation mechanism. If the intermediate oscillation has an interaction with environment and its negative wave energy is transferred to the environment, the oscillation has a tendency to be damped. However, a negative energy will be transported to the intermediate oscillation from the original one to replenish the loss. Then, the original oscillation will be excited more strongly than that in the case of absence of such negative energy loss from the intermediate oscillation. This is the case considered by \\citet{Fer08}.  Here, a brief comment is necessary on the procedure adopted in this paper to evaluate the growth rate. Our procedure is based on the fact that the growth rate can be calculated by using $W_\\pm$ given by equations (\\ref{eq: Kato1336}) and (\\ref{eq: Kato1437}). With these equations, we need not to solve an inhomogeneous wave equation for the original oscillations. In other words, the inhomogeneous wave equation to be solved by taking into account the resonant effects is only that of intermediate oscillations, as we have done in this paper. This is justified by comparing terms in equations (\\ref{eq: Kato1336}) [or equation (\\ref{eq: Kato1437})] as follows.  Let us, for example, consider $\\Im ({\\breve \\xi}_r^{\\rm int} {\\breve A}_r^*)$ as a typical term in the integrand of equation (\\ref{eq: Kato1336}) (the subscripts $+$ and ${\\tilde n}$ are omitted for simplicity). The term consists of ${\\breve \\xi}_{r,{\\rm r}}^{\\rm int} {\\breve A}_{r, {\\rm i}}^* + {\\breve \\xi}_{r,{\\rm i}}^{\\rm int} {\\breve A}_{r, {\\rm r}}^*$. If the effect of resonant processes on the original oscillation is neglected, ${\\breve A}_r^*$ is purely imaginary, i.e.,  ${\\breve A}_{r, {\\rm r}}^*=0$ [see equations (\\ref{eq:Ar+}) or (\\ref{eq:Ar-})]. Then,  $\\Im ({\\breve \\xi}_r^{\\rm int} {\\breve A}_r^*)={\\breve \\xi}_{r,{\\rm r}}^{\\rm int} {\\breve A}_{r, {\\rm i}}^*$. This is nothing but the case we have treated in this paper. If the effect of resonant processes on the original oscillation is taken into account, however, a real part of ${\\breve A}_r^*$ appears, i.e., ${\\breve \\xi}_{r,{\\rm i}}^{\\rm int}{\\breve A}_{r,{\\rm r}}^*\\not= 0$. Then, the problem to be considered here is whether the term ${\\breve \\xi}_{r,{\\rm i}}^{\\rm int}{\\breve A}_{r,{\\rm r}}^*$ is of importance in evaluating $W_+$. The term is, however, found to be smaller than ${\\breve \\xi}_{r,{\\rm r}}^{\\rm int}{\\breve A}_{r,{\\rm i}}^*$ in magnitude by the following reasons.  ${\\breve A}_{r,{\\rm r}}^*$ is non-zero, but smaller than ${\\breve A}_{r,{\\rm i}}^*$ in magnitude, since it is a small correction term of ${\\breve A}_r^*$ introduced by the resonant processes. Hence, the term ${\\breve \\xi}_{r,{\\rm i}}^{\\rm int} {\\breve A}_{r, {\\rm r}}^*$ is negligible compared with ${\\breve \\xi}_{r,{\\rm r}}^{\\rm int} {\\breve A}_{r, {\\rm i}}^*$, unless  ${\\breve \\xi}_{r,{\\rm i}}^{\\rm int}$ is much larger than ${\\breve \\xi}_{r,{\\rm r}}^{\\rm int}$ in magnitude. The results in this paper show that ${\\breve \\xi}_{r,{\\rm i}}^{\\rm int}$ and ${\\breve \\xi}_{r,{\\rm r}}^{\\rm int}$ are of the same order in the resonant region (see figures~\\ref{fig:intosc}). Therefore, ${\\breve \\xi}_{r,{\\rm i}}^{\\rm int} {\\breve A}_{r, {\\rm r}}^*$ can be safely neglected compared with ${\\breve \\xi}_{r,{\\rm r}}^{\\rm int} {\\breve A}_{r, {\\rm i}}^*$. In conclusion, when we calculate the growth rate by using $W_{\\pm}$, the inhomogeneous wave equation that should be considered is only that for the intermediate oscillations, as we have done in this paper.  The results of this paper as well as those of \\citet{Fer08} confirm that the resonant excitation process of disk oscillations in deformed disks work to excite trapped oscillations. The excitation process is thus one of the most promising causes of the observed high frequency QPOs. The trapped g-mode oscillations themselves, however, cannot account for the observed characteristic that high frequency QPOs often appear in pairs with frequency ratio close to 3:2. Furthermore, in the case of kHz QPOs in neutron-star low-mass X-ray sources, the frequencies of QPOs vary with time. To describe such observational characteristics by the present disk oscillation models, it might be necessary to relax the restriction of trapping, i.e., consideration of non-trapped oscillations might be necessary, as \\authorcite{Kat04} (\\yearcite{Kat04}, \\yearcite{Kat08a}) did.  \\bigskip  We thank the anonymous referee for invaluable comments, which greatly helped us improve the paper. FO thanks Masayuki Fujimoto for helpful discussions. She also acknowledges the scholarship from Ministry of Education, Culture, Sports, Science and Technology. ATO is grateful for the finantial support via research grants from Hokkai-Gakuen Educational Foundation and Japan Society for the Promotion of Science (20540236).  \\bigskip"
    },
    "1004/1004.1637.txt": {
        "abstract": "The strong dependence of the large-scale dark matter halo bias on the (local) non-Gaussianity parameter, $f_\\mathrm{NL}$, offers a promising avenue towards constraining primordial non-Gaussianity with large-scale structure surveys.  In this paper, we present the first detection of the dependence of the non-Gaussian halo bias on halo formation history using  $N$-body simulations.  We also present an analytic derivation of the expected signal based on the extended Press-Schechter  formalism. In excellent agreement with our  analytic prediction, we find that the halo formation history-dependent contribution to the non-Gaussian halo bias (which we call non-Gaussian halo assembly bias) can be factorized in a form approximately independent of redshift and halo mass. The correction to the non-Gaussian halo bias due to the halo formation history can be as large as 100\\%, with a suppression of the signal for recently formed halos and enhancement for old halos. This could in principle be a problem for realistic galaxy surveys if observational selection effects were to pick galaxies occupying only recently formed halos.   Current semi-analytic galaxy formation models, for example, imply an enhancement in the expected signal of $\\sim 23\\%$ and $\\sim 48\\%$ for galaxies at $z=1$ selected by stellar mass and star formation rate, respectively. ",
        "introduction": "Placing constraints on  deviations from Gaussian primordial fluctuations offers the possibility to test  inflationary models \\cite{BKMR04,KomatsuWhitepaper} and probes aspects of inflation (namely the interactions of the inflaton) that are difficult to probe by other means.  In this paper we focus on the so-called local non-Gaussianity, which  describes inflation-motivated departures from Gaussian initial conditions  and is   parameterized by  \\cite{salopek/bond:1990,  gangui/etal:1994, VWHK00, komatsu/spergel:2001}: \\begin{equation} \\label{eq:fnlnongauss} \\Phi = \\phi + f_\\mathrm{NL}  (\\phi^2 - \\left<\\phi^2\\right>)\\,. \\end{equation}  Here $\\phi$ denotes a Gaussian field and  $\\Phi$ denotes Bardeen's gauge-invariant potential, which on sub-Hubble scales reduces to the usual Newtonian peculiar gravitational potential, up to a minus sign.  The parameter  $f_\\mathrm{NL}$ is the amplitude of the non-Gaussian correction; since $\\phi \\sim 10^{-5}$ and current observational limits restrict $|f_\\mathrm{NL}| < 100$ \\cite{slosar/etal:2008,komatsu/etal:2010}, we are considering corrections of order $10^{-3}$.  Recently, Refs.~\\cite{dalal/etal:2008b, matarrese/verde:2008} showed that primordial non-Gaussianity affects the clustering of dark matter halos (i.e., density extrema), inducing a scale-dependent bias for halos on large scales. The strong scale-dependence of halo bias ($\\propto 1/k^2$) predicted for non-Gaussianity of the local type \\cite{dalal/etal:2008b} can provide constraints on $f_\\mathrm{NL}$ competitive with those available from  the Cosmic Microwave Background  \\cite{dalal/etal:2008b,carbone/verde/matarrese:2008, verde/matarrese:2009}.  Analytic estimates of the amplitude of the scale-dependent bias show good agreement with results from $N$-body simulations \\cite{dalal/etal:2008b,grossi/etal:2009, pillepich/porciani/hahn:2010, desjacques/etal:2009}.  Slosar et al.~(2008) \\cite{slosar/etal:2008} argue that the amplitude of the non-Gaussian halo bias should depend on the halo merger history.  Motivated by the idea that quasar activity is triggered by recent mergers, they estimate the amplitude of the non-Gaussian halo bias for recent mergers.  In this paper we extend their reasoning to a more general dependence on the halo merger history through the halo formation redshift $z_f$.  We compare this analysis with the dependence of the non-Gaussian halo bias on halo merger history detected in the $N$-body simulations of Grossi et al.~(2009) \\cite{grossi/etal:2009}.  By comparison with the halo merger history dependence of the halo occupation distribution of certain galaxies in semi-analytic models of galaxy formation, we estimate the possible impact of these results on predictions for the amplitude of the non-Gaussian bias in upcoming large scale structure surveys.  This paper is organized as follows. In Section 2 we first revisit the extended Press Schechter non-Gaussian halo merger bias derivation of Ref.~\\cite{slosar/etal:2008} and then generalize it to arbitrary halo formation redshifts.  In Section 3 we detect the effect in $N$-body simulations and show the agreement with the analytic description, including the simple halo mass and redshift dependence predicted by the model.  We explore the consequences for  practical determination of $f_\\mathrm{NL}$ in Section 4 and we conclude in Section 5. The appendix presents our methodology for fitting our simulation results for the amplitude of the non-Gaussian halo bias mode by mode, i.e. without computing a binned power spectrum. ",
        "conclusions": "We have demonstrated that the impact of assembly bias on the amplitude of the non-Gaussian halo bias can be quite strong.  We have expanded the arguments in Slosar et al. (2008) \\cite{slosar/etal:2008} using extended Press-Schechter theory to express non-Gaussian assembly bias in terms of halo formation redshift $z_f$ for arbitrary $f \\geq 0.5$, where $z_f$ is the redshift at which a halo has accreted a fraction $f$ of its final mass.  This theory predicts that halo subsamples containing a fraction $x$ of the earliest (latest) forming halos (compared with other halos with the same mass) have a non-Gaussian halo bias that differs from the full parent halo sample by a fractional correction dependent only on $x$; when using this variable, the non-Gaussian assembly bias correction is independent of halo mass and redshift.  The $N$-body simulations of Grossi et al. (2009) \\cite{grossi/etal:2009} are in good agreement with these ePS predictions.  The implications of these results for galaxy redshift surveys are extremely uncertain.  If the commonly adopted assumption that the probability of a halo hosting a particular type of galaxy only depends on the halo mass, then there will be no non-Gaussian halo assembly bias contribution to the galaxy sample's non-Gaussian bias.  However, in principle, galaxy formation depends on the entire history of host dark matter halos to some degree, and semi-analytic models of galaxy formation attempt to account for this dependence.  In Section \\ref{implications}, we found that a relatively mild preference for early-forming halos for both stellar mass and star formation rate selected $z=1$ samples translates into an increase in the expected non-Gaussian galaxy bias of $\\sim 23-48\\%$ compared with the average signal expected from the samples' Gaussian bias values.  This result is particularly counter-intuitive for star-forming galaxies, since star formation is triggered by galaxy mergers in these models.  One should bear in mind that a galaxy merger does not necessarily correspond to a major merger of the host dark matter halo, and that it is reasonable to expect some time-lag between the dark matter halo merger and the merger of the galaxies populating them.  Furthermore, we caution that we have not been extensive in our exploration of galaxy sample selection space, or precise enough to make predictions for upcoming experiments.  There may be certain populations for which this effect may be much larger or much smaller.  On the other hand, it is possible that in an analysis aimed at constraining $f_\\mathrm{NL}$, one may be able to weight galaxies by some color or spectral property in order to enhance the non-Gaussian signal in the survey. More work is needed to further quantify the impact of such an approach on the recovered constraints  on $f_\\mathrm{NL}$ from realistic surveys."
    },
    "1004/1004.0249_arXiv.txt": {
        "abstract": "We present a new mechanism for the ejection of a common envelope in a massive binary, where the energy source is nuclear energy rather than orbital energy. This can occur during the slow merger of a massive primary with a secondary of $1-3\\,$M$_\\odot$ when the primary has already completed helium core burning. We show that, in the final merging phase, hydrogen-rich material from the secondary can be injected into the helium-burning shell of the primary. This leads to a nuclear runaway and the explosive ejection of both the hydrogen and the helium layer, producing a close binary containing a CO star and a low-mass companion. We argue that this presents a viable scenario to produce short-period black-hole binaries and long-duration gamma-ray bursts (LGRBs). We estimate a LGRB rate of $\\sim 10^{-6}$\\,yr$^{-1}$ at solar metallicity, which implies that this may account for a significant fraction of all LGRBs, and that this rate should be higher at lower metallicity. ",
        "introduction": "Roughly half of the known short-period black-hole binaries have low-mass companions (Lee, Brown \\& Wijers 2002). However, it has been realized for more than a decade now that such systems are difficult to form (see the discussion in Podsiadlowski, Rappaport \\& Han 2003 [PRH] and further references therein). The problem is that, in order to produce a short-period system (i.e., $\\la 15$ hrs) with a low-mass donor star, the progenitor system has to pass through a common-envelope (CE) phase where the binary's period is reduced from a typical period of several years to less than a few days. In the standard model for CE evolution (Paczy\\'nski 1976), this requires that the orbital energy released in the spiral-in process is sufficient to eject the massive envelope of the primary. For a low-mass companion (i.e., $\\la 2~M_\\odot$), this is energetically difficult, if not impossible, in particular considering the large binding energies of the massive envelope of the primaries (Dewi \\& Tauris 2001; PRH).\\footnote{In a recent study by Yungelson et al.\\ (2006), as well as in some other studies, this problem did not seem to arise. However, these authors used a prescription for the binding energy of the envelope that {\\em underestimates} the binding energy by up to a factor of 10 compared to the actual binding energies calculated from realistic models for the structure of massive supergiants (Dewi \\& Tauris 2001; PRH).}  This problem has led to the suggestion of a number of more exotic formation scenarios for low-mass black-hole binaries, e.g., involving a triple scenario (Eggleton \\& Verbunt 1986) or the formation of the companion from a disrupted envelope (Podsiadlowski, Cannon \\& Rees 1995; PRH). Alternatively, the low-mass black-hole binaries could descend from intermediate-mass systems (Justham, Rappaport \\& Podsiadlowski 2006), as is the case for the majority of low-mass neutron-star binaries (Pfahl, Rappaport \\& Podsiadlowski 2003). In this paper, we present a new mechanism for the ejection of the common envelope: ``explosive common-envelope ejection'', involving nuclear rather than orbital energy, which can be highly efficient in ejecting a massive envelope even if the companion is a relatively low-mass star.  A particularly interesting black-hole binary is the well studied system GRO J1655--40 (Nova Scorpii 1994). It has an orbital period of 2.6\\,d and a black hole with a mass $\\sim 5.4\\,M_\\odot$ (Beer \\& Podsiadlowski 2002). Israelian et al.\\ (1999) claimed that the secondary in this system has been highly enriched with the products of explosive nucleosynthesis (e.g., Mg, Si, S, Ti) produced in the supernova explosion that produced the black hole.\\footnote{In this context, we refer the reader to two related studies: one by Foellmi, Dall \\& Depagne (2007) challenging the original claim of these overabundances, and one by Gonz\\'alez Hern\\'andez, Rebolo \\& Israelian (2008) re-affirming them.}  Based on the actual abundance ratios, Podsiadlowski et al.\\ (2002) found some tentative evidence that these are better explained by an energetic supernova explosion, a hypernova, with a typical ejecta energy of $\\ga 10^{52}$\\,ergs (i.e., 10 times the energy in a `typical' supernova). This may suggest that the formation of the black hole in this system could have been accompanied by a long-duration gamma-ray burst (LGRB).  However, in this case it is puzzling why the black-hole progenitor would have been rapidly rotating as required in the collapsar model for LGRBs (Woosley 1993; MacFadyen \\& Woosley 1999) since the core of the primary should have been spun down rather than have been spun up during its evolution (Heger, Woosley \\& Spruit 2005).  As we will show later in this paper, in the case of explosive CE ejection, this problem does not arise, possibly linking the formation of compact black-hole binaries having low-mass secondaries to LGRBs.\\footnote{Brown et al.\\ (2000) were the first to suggest that the formation of the black hole in GRO J1655-40 was associated with an LGRB and proposed a general link between low-mass black-hole binaries and LGRBs (see, in particular, Brown, Lee \\& Moreno M\\'endez 2007). Similar to the present study, they suggest late Case C mass transfer for the progenitor systems, but argue for the spin-up of the cores rather than spin-down during the CE phase, invoking tidal locking, an assumption that remains to be proven.}   This process of explosive CE was discovered in a systematic study of the slow merger of massive stars (Ivanova 2002; Ivanova \\& Podsiadlowski 2003; Ivanova \\& Podsiadlowski 2010), where it was found that, in some cases in the late stage of the spiral-in process, hydrogen-rich material could be mixed into the helium-burning shell leading to a thermonuclear runaway which released enough energy to eject both the hydrogen and the helium envelope (see Figure~1 for a schematic representation), possibly explaining why to date all supernovae associated with LGRBs appear to be Type Ic supernovae (i.e., supernovae without hydrogen and helium in their spectrum; cf. Podsiadlowski et al.\\ 2004).\\footnote{We note that Siess \\& Livio (1999a,b) were the first to point out the potential importance of nuclear energy on the CE ejection process in their modelling of the dissolution of planets/brown dwarfs in red-giant stars.}  In this paper, we will first discuss the numerical method in Section~2, and the physics of the process of explosive CE ejection in Section~3. In Section~4 we apply it to the formation of low-mass black-hole binaries and LGRBs, and end with a broader discussion in Section~5. ",
        "conclusions": "In this paper, we have presented a new mechanism for the ejection of a common envelope where the energy source is not orbital energy but nuclear energy. This provides a new channel to produce plausible progenitors of short-period black-hole binaries and long-duration GRBs.  We have also demonstrated that this scenario may be able to explain both the origin and the main properties of short-period black-hole binaries, such as the period distribution, although the distribution of spectral types is still not fully satisfactory.   Two of the most attractive features of the explosive CE ejection scenario are that (1) it leads to the ejection of both the hydrogen and the helium layer, explaining why all LGRB supernovae to-date have been classified as Type Ic supernovae, and (2) this ejection occurs late in the evolution of the star; hence the progenitor will only experience a short Wolf-Rayet phase, in which it will not be spun down significantly by wind mass loss.  This may also explain why extended Wolf-Rayet wind bubbles are not being found around LGRBs; e.g., in the case of GRB 021004, van Marle, Langer \\& Garcia-Segura (2005) found that the Wolf-Rayet phase had to last less than $10^4$\\,yr; this is consistent with the explosive CE-ejection scenario, since it always requires late case C mass transfer. This also predicts that such Wolf-Rayet bubbles are terminated by a dense shell from the ejected common envelope.  Our rate estimate for this channel ($\\sim 10^{-6}\\,$yr$^{-1}$) implies that this can produce a significant fraction of all LGRBs; this rate should be higher at lower metallicity, because case C mass transfer is expected to be more common at lower metallicity.  Unlike some of the single-star progenitor models for LGRBs (Yoon \\& Langer 2005; Woosley \\& Heger 2006; Yoon, Langer \\& Norman 2006), LGRBs may occur even at solar metallicity, but they are expected to be more common at low metallicity. Indeed, there is some evidence now that LGRBs can also occur in super-solar host galaxies (see, e.g., Levesque et al.\\ 2010). The different metallicity biases may provide a possible way of distinguishing between these two different scenarios.  Explosive CE ejection also operates for lower-mass primaries that are expected to produce neutron stars rather than black holes.  It is tempting to associate these with rapidly rotating neutron stars and possibly magnetars, for which SN 2006aj, which was associated with the X-ray flash GRB 060218, may provide an observed example in nature (Mazzali et al.\\ 2006).  \\bigskip  \\noindent{\\bf ACKNOWLEDGEMENTS} \\medskip  \\noindent NI gratefully acknowledges support from the NSERC of Canada and from the Canada Research Chairs Program.  SJ is partially supported by the National Science Foundation of China under grant numbers 10903001 and 10950110322, and by the Chinese Postdoc Fund (award number 20090450005)."
    },
    "1004/1004.0280_arXiv.txt": {
        "abstract": "We extend the concept of galaxy environment from the local galaxy number density to the gravitational potential and its functions like the shear tensor. For this purpose we examine whether or not one can make an accurate estimation of the gravitational potential from an observational sample which is finite in volume, biased due to galaxy biasing, and subject to redshift space distortion. Dark halos in a $\\Lambda$CDM simulation are used in this test. We find that one needs to stay away from the sample boundaries by more than 30$h^{-1}$Mpc to reduce the error within 20\\% of the root mean square values of the potential or the shear tensor. The error due to the galaxy biasing can be significantly reduced by using the galaxy mass density field instead of the galaxy number density field. The error caused by the redshift space distortion can be effectively removed by correcting galaxy positions for the peculiar velocity effects. We inspect the dependence of dark matter halo properties on four environmental parameters; local density, gravitational potential, and the ellipticity and prolateness of the shear tensor. We find the local density has the strongest correlation with halo properties. This is evidence that the internal physical properties of dark halos are mainly controlled by small-scale physics. In high density regions dark halos are on average more massive and spherical, and have higher spin parameter and velocity dispersion. In high density regions dark halos are on average more massive and spherical, and have higher spin parameter and velocity dispersion. We also study the relation between the environmental parameters and the subtypes of dark halos. The spin parameter of satellite halos depends only weakly on the local density for all mass ranges studied while that of isloated or central halos depends more sensitively on the local density. The gravitational potential and the shear tensor have weaker correlations with halo properties, but have environmental information independent of the local density. ",
        "introduction": "One of recent developments in the study of galaxy formation is quantitative understanding of the dependence of galaxy properties on environment (Park et al. 2007; Hwang \\& Park 2009 among others). It has already been noticed since the 1930's that galaxy luminosity and morphology depend on the local density: the high density regions preferentially harbor more luminous and early morphological type of galaxies (Hubble \\& Humason 1931). In the beginning of the studies on the environmental effects on galaxy properties (on the simulation side, see Lemson \\& Kauffmann 1999; Gao et al. 2005; Jing et al. 2007; Hahn et al. 2007; Maccio et al. 2007; on the observation side, see Blanton \\& Berlind 2007; Park et al. 2007; Cervantes-Sodi et al. 2008; Skibba et al. 2009; Blanton \\& Moustakas 2009), the environment was distinguished according to the large-scale structure where the galaxies under study are located. For example, comparative studies for galaxies located within massive galaxy clusters, groups, or voids were carried out (Oemler 1974; Dressler 1980; Postman \\& Geller 1984; Rojas et al. 2004 among many others).  Another trend of the same kind of study used continuous parameters that measure the local galaxy number density. Various kinds of smoothing kernel were used to estimate the local number density from galaxy positions in redshift space. The most popular one is the truncated cylindrical cone which is motivated from the fact that massive clusters appear as Fingers of God in redshift space (Hogg et al. 2004; Kauffmann et al. 2004; Kuehn \\& Ryden 2005; Reid \\& Spergel 2009). The Gaussian and spline filters are also often used. These filters can have a fixed size or can vary in size to include a fixed number of galaxies (Park et al. 1994; Monaco \\& Efstathiou 1999; Park et al. 2007; Kim et al. 2008).  Even though these two approaches look quite different, they essentially use the local galaxy number density to distinguish among different environments. Extension of the concept of environment beyond the `local number density' has been started by several authors. Park, Gott \\& Choi (2008) and Park \\& Choi (2009) used the mass density due to the galaxy plus dark halo systems as a new environmental parameter. They also divided the `local' density into the large-scale background mass density and the small-scale density attributed to the nearest neighbor galaxy. The mass density is estimated from galaxy luminosity and mass-to-light ratios. It turned out that the environment set up by the nearest neighbor was critically important in determining galaxy properties. Lee et al. (2009) used galaxy luminosity density and local color (difference between luminosity densities in two bands) as environmental parameters. Lee \\& Lee (2008) inspected the relation between the ellipticity of the tidal shear and galaxy morphology.  It is expected that some galaxy properties depend on the local galaxy number/mass density sensitively. Color and recent star formation activity may be such properties. However, the root-mean-square (RMS) displacement  of mass is about $10 h^{-1}$Mpc till the present epoch in the $\\Lambda$CDM model best fit to the Wilkinson Microwave Anisotropy Probe (WMAP) 3-year data (Park \\& Kim 2009), and therefore the local density at the present location of galaxies cannot fully represent the environment where galaxies formed and evolved. This is true particularly for intermediate and high density regions since their sizes are typically only a few Mpcs. Therefore, it may be useful to consider environmental parameters other than local density to understand the environmental effects on galaxy formation.  A theoretically motivated environmental parameter is the gravitational potential. Dark matter and baryons are expected to fall into the deep gravitational potential well to form massive objects. Since the gravitational potential field picks up the fluctuation power at scales much larger than those of the density field, the correlation between them will not be perfect. And it will be interesting to see how galaxy properties are related with the `local' gravitational potential.  Furthermore, it is expected that galaxy angular momentum is generated from the large-scale gravitational shear force. According to tidal torque theory (Hoyle 1951; Peebles 1969; Doroshkevich 1970; White 1984; Lee \\& Pen 2000; Vitvitska et al. 2002; Porciani et al 2002), the origin of galactic angular momentum is originated by the tidal torque operating on primordial gas lump that will form a galaxy. The torque is given by $\\tau = T \\times I$, where $T$ is a shear tensor generated by external material and $I$ is a moment of inertial tensor of material being torqued. Navarro et al. (2004) presented supporting evidence for the theory with direction of galaxy rotation axis. An accurate estimation of the tidal shear tensor will enable one to verify if the tidal torque theory is really responsible for galaxy spin (Porciani et al. 2002; Lee \\& Pen 2002).  In this paper we will study how accurately one can estimate the gravitational potential and its functions from a simulated sample of galaxies. The error sources in this estimation are divided into three categories:  1. finite volume of the survey,  2. galaxy biasing, and  3. redshift space distortion.  \\noindent We then inspect the dependence of dark matter halo properties on various environmental parameters including the `local' gravitational potential. It is hoped that a generalization of environmental parameter beyond the local density allow us to better understand galaxy formation and evolution. ",
        "conclusions": "In this paper we demonstrated the gravitational potential and its functions can be reasonably accurately estimated from an observational sample that covers only a finite volume of the universe and suffers from tracer biasing and redshift space distortion. We found that the error in the gravitational potential and shear tensor decreases rapidly as one moves inside the survey boundaries. In the case of shear tensor the error becomes less than 20\\% of its RMS value in the regions separated from the survey boundaries by more than about 30$h^{-1}$Mpc. This requires the sample size to be much larger than 60$h^{-1}$Mpc for an environment study with accurate estimation of the potential and its functions. Our study also shows that the effects of halo biasing on the gravitational potential estimation can be greatly reduced by weighting dark halos (or galaxies) by their mass as was done by Park et al. (2008) and Park \\& Choi (2009) in their studies of small and large-scale density dependence of galaxy properties. Accuracy in the estimation starts to fall down rapidly when the halo mass cut is larger than $10^{13}h^{-1}$M$_{\\odot}$i. This means that the mass density and potential fields estimated from the distribution of the Luminous Red Galaxies (LRGs) will be quite inaccurate because the mean separation of a volume-limited sample of the SDSS LRGs is about 20$h^{-1}$Mpc (Gott et al. 2009), which corresponds to the halo mass cut of about $2\\times 10^13 h^{-1} M_{\\odot}$.  The error due to the redshift space distortion effects can be also reduced by using dark halo mass density in the estimation of potential. But even more reduction can be achieved by correcting the observed (i.e. redshift space) distribution of dark halos for the peculiar velocity. It was sufficient to use the peculiar velocity linearly estimated from the redshift space distribution of dark halos. After making the peculiar velocity correction and using halo mass weight the error due to the redshift space distortion becomes tiny.  We showed there exists large dispersion in the gravitational potential and the shear at fixed local density. It demonstrates the potential has large-scale information independent of local density. We inspected the dependence of dark matter halo properties on local density, gravitational potential, shear ellipticity and prolateness. Among these environmental parameters the local density shows the strongest correlation with the internal physical parameters of dark halos. When halo mass is fixed, the spin and shape parameters are nearly independent of the potential and shear tensor but depend sensitively on the background density in the case of massive halos, in particular.  In the following paper we will analyze the Main Galaxy sample of the SDSS DR7 catalog to examine the dependence of various galaxy properties on these environmental parameters. This will extend our understanding on the environmental effects on galaxy formation and evolution."
    },
    "1004/1004.2766_arXiv.txt": {
        "abstract": "We simulate time-dependent particle acceleration in the blast wave of a young supernova remnant (SNR), using a Monte Carlo approach for the diffusion and acceleration of the particles, coupled to an MHD code. We calculate the distribution function of the cosmic rays concurrently with the hydrodynamic evolution of the SNR, and compare the results with those obtained using simple steady-state models. The surrounding medium into which the supernova remnant evolves turns out to be of great influence on the maximum energy to which particles are accelerated. In particular, a shock going through a $\\rho \\propto r^{-2}$ density profile causes acceleration to typically much higher energies than a shock going through a medium with a homogeneous density profile. We find systematic differences between steady-state analytical models and our time-dependent calculation in terms of spectral slope, maximum energy, and the shape of the cut-off of the particle spectrum at the highest energies. We also find that, provided that the magnetic field at the reverse shock is sufficiently strong to confine particles, cosmic rays can be easily re-accelerated at the reverse shock. ",
        "introduction": "\\label{sec:intro} Cosmic rays with energies up to at least $\\sim 10^{15}$~eV are thought to originate in supernova remnants (SNRs). They are high-energy particles that have a simple power-law energy spectrum that extends over five decades in energy. The need for an efficient acceleration mechanism in SNRs has motivated the development of the theory of diffusive shock acceleration (DSA), according to which particles are accelerated at shock fronts \\citep[see~][for a comprehensive review]{2001MalkovDrury}.  Young supernova remnants are ideal locations for studying this process, because of the high shock velocity, and the presence of a few nearby young remnants for which detailed observations are available \\citep[e.g.~][]{2004Hwangetal, 2005Bambaetal, 2009Aceroetal}. The presence of high-energy electrons spiralling in a $\\sim 10\\ \\umu$G magnetic field has been established from the emission of synchrotron radiation from radio wavelengths to X-rays. Both the magnetic field strength and the typical particle energy are inferred from synchrotron theory and can be used to compare with theoretical predictions \\citep{1994Achterbergetal,2003VinkLaming, 2005Voelketal, 2005Vink}. Even though synchrotron emission only indicates the presence of relativistic electrons, the presence of energetic protons is suggested by the observations of TeV gamma rays \\citep[e.g.~][]{2009Aharonianetal}, and by indications of magnetic field amplification beyond that what is expected from simple shock compression \\citep{2005Warrenetal, 2008CassamChenaietal}. In addition, there are indications that the SNR blast waves are not simple hydrodynamic blast waves in a single-component gas. They behave in a way that indicates that a significant fraction of the pre-shock energy density resides in cosmic rays \\citep{2000Decourchelleetal,2009Helderetal}.  Various groups work on trying to get an integral picture of the interaction between SNR shocks, magnetic fields, and the associated particle acceleration \\citep[e.g.~][]{1999BerezhkoEllison,2005KangJones,2005AmatoBlasi, 2006Vladimirovetal,2006BerezhkoVoelk,2007ZirakashviliAharonian,2010Ferrandetal}. The difficulty is the fact that the process is inherently non-linear. The spectral slope of the particles is determined mainly by the difference between the plasma velocities on the two sides of the shock. This difference depends on the compression ratio, which in turn is determined by the effective equation of state of the gas-cosmic ray mixture: the presence of cosmic rays tends to soften the equation of state around the shock, which in turn increases the total compression. In addition, the gradient in cosmic ray pressure slows down the incoming flow in the cosmic-ray precursor to the shock. High-energy particles with a larger scattering mean free path probe a larger region around the shock and `feel' a higher compression ratio. For these reasons the spectrum flattens at the high-energy end in fully non-linear models. An additional nonlinearity arises when the magnetic fields are amplified by the streaming of cosmic rays.  Cosmic rays isotropise by scattering off Alfv\\'en waves. In gyro-resonant scattering the scattering rate depends on the cosmic ray energy through the slope of the spectrum of the scattering waves. It is often assumed that the diffusion rate is described by Bohm diffusion, where the diffusion coefficient scales as $\\kappa_{\\rm B} \\propto E$, with $E$ the energy of the particle. These waves are self-generated by the streaming of cosmic rays away from the shock, through a resonant instability \\citep[e.~g.][]{1975Skilling}, and / or through a non-resonant instability \\citep{2001BellLucek, 2004Bell, 2009LuoMelrose}.   Various authors focus on the feedback of cosmic rays onto the hydrodynamics near the shock \\citep{1997Malkov, 2007Blasietal, 2009Patnaudeetal,2010Ferrandetal}. The distribution function of the particles is calculated, from which the cosmic ray pressure can be determined. The pressure term alters the equation of state and feeds back on the hydrodynamics. Alternatively, a standard power law distribution is assumed, simplifying and speeding up the process, making it fast enough for application on larger scales \\citep{2007Ensslinetal}. The disadvantage of this approach is that -- generally speaking-- the time-dependence or the influence of a complicated magnetic field geometry can not be taken into account.    Our work focuses on the time- and position dependence of cosmic ray spectra in the supernova remnant, where we use the test-particle approximation, neglecting feedback of the particles onto the plasma. The high-energy particles isotropise on both sides of the shock, due to scattering off waves, and are accelerated by repeated crossing cycles at the shock. The acceleration of relativistic particles can be described by a set of stochastic differential equations (SDEs) \\citep{1992AchterbergKruells,1994KruellsAchterberg} and this has been applied succesfully by a number of authors  \\citep{1999MarcowithKirk, 2004vanderSwaluwAchterberg, 2010MarcowithCasse} in various hydrodynamic codes.  We use the adaptive mesh refined magneto-hydrodynamics (MHD) code: AMRVAC \\citep{2003Keppensetal, 2007HolstKeppens} as the framework for our particle acceleration method. Not only do we calculate the acceleration/deceleration of test particles due to compression/expansion of the flow, we also model the dependence of diffusion and radiative losses, while keeping computational costs down with the adaptive mesh strategy. Our approach has the advantage that it is able to also tackle a more complicated circumstellar density profile than other models. The disadvantage is (for now) having to neglect the nonlinear feedback of the cosmic rays onto the plasma.  We describe the different models we use to calculate the particle spectrum in supernova remnants in \\S~\\ref{sec:1Dmodels}. In \\S~\\ref{sec:theory} we will discuss the theory of diffusive shock acceleration and the evolution of the supernova remnant, and derive some analytical estimates for the expected cosmic ray spectrum. The method and set-up of the simulations will be described in \\S~\\ref{sec:method}. In \\S~\\ref{sec:testmodels} we will describe some test models and results obtained with this method. We will subsequently present the results for the particle spectra from the SNR models in \\S~\\ref{sec:snr} and conclude with a discussion and summary in \\S~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  In this paper, we have calculated diffusive shock acceleration through the first-order Fermi mechanism, using stochastic differential equations. We treat the cosmic rays as test-particles and follow their acceleration and propagation along with the evolution of a supernova remnant. We have extended the model as described by \\citet{1992AchterbergKruells} \\citep[and used by e.g.~][]{2004vanderSwaluwAchterberg} to account for spherical geometry, something that is relevant when using this model to simulate cosmic ray acceleration in supernova remnants. The model is set up generically so it can use local magnetic field strengths to calculate the diffusion coefficient of the test-particles. However, in this paper we employ a constant magnetic field strength in the ISM/CSM. The ejecta magnetic field is parametrised in such a way that it either decays as expected for an expanding field that is frozen into the plasma, or we fix the magnetic field at the same strength as used for the ISM/CSM. This is because the numerical method that we use is not applicable in cases where strong gradients of the diffusion coefficient that may result from gradients in the magnetic field are present. The unique set-up of the calculation of the acceleration of the cosmic rays concurrent with the evolution of the SNR allows us to model the time- and location-dependent spectrum more accurately than other models. The disadvantage of our method is that it does not (yet) include feedback of the cosmic ray pressure onto the local plasma. Our calculations show the following results.  \\subsubsection*{Energy spectrum and spectral slope:}  With our method we can accurately model the spectrum of particles, where the slope is close to $q=2$ for an adiabatic index of $5/3$, and to $q=1.5$ for an adiabatic index of $4/3$. The slope of the spectrum depends on the location: near the shock the slope of the spectrum is closest to the analytically predicted value for steady and planar shocks, while the slope of the cumulative spectrum (all particles in the source) is steeper. The cumulative spectrum is steeper than that at the shock because it also contains particle populations from downstream of the shock, for which the particles were in the acceleration process for a shorter time and hence have lower cut-off energies. We also find that in spherical geometry the overall spectrum is slightly steeper when compared to a simulation in slab geometry. The main reason for this is the inclusion of adiabatic expansion, which causes the acceleration process to be slightly less effective: particles lose a fraction of their energy in the downstream part of the shock crossing cycle, reducing the mean energy gain per cycle. The analytical solution for the steady state parallel geometry where adiabatic losses are excluded therefore does not strictly apply in spherical geometry. The detailed shape of the spectrum also depends on the injection rate of particles as a function of time, which differs in the ISM and the CSM models that we employ.  \\subsubsection*{Maximum particle energy:}  There are additional differences between the results of our approach and those obtained using steady-state analytical models. The cut-off energy $E_{\\rm max}$ is different from that obtained analytically from the balance between acceleration and losses. For protons the $E_{\\rm max}$ is determined by the age of the source. The CSM simulations show a higher $E_{\\rm max}$ than the analytical estimate, whereas for the ISM models the trend is the other way around. We attribute this to the difference in cosmic ray age distribution, where, since we assume the injection rate to be proportional to the amount of swept-up mass, the fraction of `older' particles, which are accelerated for a longer time, is relatively high in the CSM model and relatively low for the ISM model. For Bohm diffusion, the high-energy end of the spectrum shows a distinct cut-off, either caused by the finite age of the SNR or, in the case of electrons, by synchrotron losses.  \\subsubsection*{Shape of the cut-off in the energy spectrum:}  The shape of the cut-off region for the cumulative spectrum follows quite nicely the quasi-exponential drop that is often assumed. However, if one looks at the spectrum in the close vicinity of the shock, the proton spectrum falls off slightly sharper than for the overall proton spectrum, while for electrons in the loss-limited regime the cut-off of the spectrum at the shock is more gradual than the cut-off in the overall spectrum.  \\subsubsection*{CSM versus ISM:}  The strong dependence of the acceleration rate on the shock velocity causes the cosmic ray distribution to become sensitive to the surrounding medium. The density profile of the environment into which the supernova explodes determines the shock velocity and its evolution. The shock velocity (together with the magnetic field) therefore determines the acceleration rate and the maximum particle energy that can be attained. Because in our models the average shock velocity is much higher for a SNR expanding into a CSM, and because the average cosmic ray age in the remnant is larger in the CSM case, $E_{\\rm max}$ is larger in our CSM models than in the ISM models.  This suggests that core-collapse SNe in dense environments, such as expected around a red supergiant, may be the most efficient particle accelerators and therefore the dominant contributors to cosmic rays up to the ``knee''-energy. In the absence of a significant magnetic field in the ejecta, the particles are mostly located between the blast wave and the contact discontinuity. The distance between those is also sensitive to the velocity-evolution of the SNR and therefore the surrounding medium, and determines how long electrons are subjected to synchrotron losses.  \\subsubsection*{Re-acceleration at the reverse shock:}  If the magnetic field is sufficiently amplified at the reverse shock, cosmic rays that are advected away from the blast wave can be re-accelerated at the reverse shock. This has important consequences for the maximum energy and the distribution of the cosmic rays. The maximum attainable energy for protons becomes significantly higher if re-acceleration at the reverse shock takes place, whereas for electrons the additional synchrotron losses in the now strongly magnetised SNR interior can have the opposite effect. In reality it is conceivable that due to localized magnetic field amplification, the net effect is a higher maximum energy for electrons, too.   \\bigskip Overall, we conclude that a time-dependent calculation of diffusive shock acceleration in SNRs shows significant differences compared with steady-state plane-parallel analytical models. The environment of the SNR has a large impact on the maximum attainable energy of the cosmic rays. The age distribution of the cosmic rays determines whether a time-dependent approach yields higher or lower maximum attainable energies."
    },
    "1004/1004.0139_arXiv.txt": {
        "abstract": "{% Time-resolved spectroscopic observations of rapidly oscillating Ap (roAp) stars show a complex picture of propagating magneto-acoustic pulsation waves, with amplitude and phase strongly changing as a function of atmospheric height. We have recently conducted numerical, non-linear MHD simulations to get an insight into the complex atmospheric dynamics of magnetic pulsators. Here we use the resulting time-dependent atmospheric structure and velocity field to predict line profile variations for roAp stars. These calculations use realistic atmospheric structure, account for vertical chemical stratification and treat the line formation in pulsating stellar atmosphere without relying on the simplistic single-layer approximation universally adopted for non-radial pulsators. The new theoretical calculations provide an essential tool for interpreting the puzzling complexity of the spectroscopic pulsations in roAp stars. } ",
        "introduction": "Rapidly oscillating Ap (roAp) stars is a group of cool magnetic Ap stars pulsating in high-overtone, non-radial modes with periods around 10~min. Excitation of these pulsations and the physics of their propagation in the stellar envelopes and atmospheres is closely connected to the presence of global magnetic fields of several kG strength (e.g., Balmforth et al. 2001; Saio 2005).  Recent time-resolved spectroscopic observations of roAp stars (Kochukhov \\& Ryabchikova 2001; Mkrtichian, Hatzes \\& Kanaan 2003; Ryabchikova et al. 2007) showed a remarkably complex and diverse pulsational variability of spectral lines of different chemical elements. In particular, one often finds a factor of 100 difference in amplitude and phase jumps of up to $\\pi$ radian between the lines which should originate in very similar regions of normal stellar atmosphere. This unique behaviour is understood to be a result of vertical chemical stratification (e.g., Kochukhov et al. 2006), combined with a rapid intrinsic height variation of the magneto-acoustic pulsation waves propagating in stellar atmosphere.  The spatial filtering effect of chemical inhomogeneities opens interesting prospects for horizontal and vertical resolution of pulsation modes in roAp stars, as can be done for no other type of pulsating stars except for the Sun. The horizontal mapping of roAp pulsations has already been performed with the help of an extended Doppler imaging technique (Kochukhov 2004). However, the vertical resolution of pulsation modes turns out to be considerably more challenging because one has to abandon the standard single-layer approximation universally adopted in detailed line profile modelling of non-radially pulsating stars (e.g., Briquet \\& Aerts 2003; Schrijvers et al. 1997).  In an effort to get a better insight into the complex atmospheric dynamics of these stars, we are conducting numerical, non-linear magneto-hydrodynamic (MHD) simulations of pulsational wave propagation (Khomenko \\& Kochukhov 2009). Here we use the resulting time-dependent atmospheric structure and velocity field to predict line profile variations for roAp stars accounting for vertical chemical stratification and using realistic line formation calculations. This theoretical modelling represents a key step towards understanding the puzzling complexity of the spectroscopic pulsations in roAp stars and eventually resolving 3-D structure of their pulsation modes.  \\begin{figure*}[!t] \\centering \\includegraphics[width=\\textwidth]{kochukhov2_f1.eps} \\caption{ Initial model atmosphere and chemical stratification adopted for the MHD calculations and line profile synthesis. {\\bf a)} Temperature (solid line) and density as a function of height and optical depth for an unperturbed atmospheric model with $T_{\\rm eff} = 7750$~K and $\\log g = 4.0$. {\\bf b)} Depth-dependence of the sound (solid line) and Alfv\\'en (dotted lines) speeds for several values of the magnetic field strength. {\\bf c)} Vertical stratifications of Ca, Fe and three different distributions of Nd employed in the spectrum synthesis.} \\label{atmos} \\end{figure*}  \\begin{figure*}[!t] \\centering \\includegraphics[width=140mm]{kochukhov2_f2.eps} \\caption{ Height-time variation of the vertical and horizontal pulsation velocities (top panels), temperature and density (bottom panels). These results illustrate MHD simulations of the $P=10$~min pulsation at the intermediate magnetic co-latitude of the $B_{\\rm p}=3$~kG dipolar magnetic field. The local field strength is $1.9$~kG and the field inclination is 44.8\\degr\\ with respect to the surface normal. The horizontal dotted lines indicate the surfaces of constant optical depth $-6\\le\\log\\tau_{5000}\\le1$. } \\label{mhd} \\end{figure*} ",
        "conclusions": "\\begin{itemize} \\item We have presented the first calculations of the pulsational LPV in roAp stars based on detailed MHD models and realistic treatment of the spectral line formation in magnetic, chemically stratified atmosphere. \\item In agreement with observations, our calculations show a large increase of the pulsational amplitude and a gradual delay in phase from Fe and Ca to rare-earth spectral lines. \\item Even for the simple $\\ell=1$, $m=0$ pulsation mode considered in our modelling the behaviour of line profiles, as well as pulsational amplitude and phase, depends strongly on the local field strength and inclination. \\item Pulsational changes of pressure and temperature contribute non-negligibly to the profile variations of rare-earth lines and can dominate variations of Fe and Ca lines. \\end{itemize}"
    },
    "1004/1004.4006_arXiv.txt": {
        "abstract": "We report the discovery of H{\\sc i} 21-cm absorption towards the well-studied GHz Peaked-Spectrum source CTA~21 (4C~16.09) using the Arecibo Telescope on 2009 September 20 and 21. Recently, the frequency band between 700 and 800 MHz was temporarily opened up to radio astronomy when US TV stations were mandated to switch from analog to digital transmissions, with new frequency allocations.  The redshifted H{\\sc i} frequency  for CTA~21 falls within this band. CTA~21 has a complex radio structure on a range of scales.  The innermost prominent components are separated by $\\sim$12 mas while weak diffuse emission extends for up to $\\sim$300 mas.  The H{\\sc i} absorption profile that we find has two main components, one narrow, the other wider and blue-shifted. The total H{\\sc i} column density is 7.9$\\times$10$^{20}$ cm$^{-2}$, assuming a covering factor of unity and a spin temperature of 100 K.  This H{\\sc i} absorption confirms the recently determined optical redshift of this faint galaxy of z$\\sim$0.907.  We discuss this new detection in the light of H{\\sc i} absorption studies towards compact radio sources, and also the possibility that CTA~21 may be exhibiting multiple cycles of nuclear activity.  This new detection in CTA~21 is consistent with a strong trend for detection of H{\\sc i} absorption in radio galaxies with evidence of episodic nuclear/jet activity. ",
        "introduction": "In the widely accepted model of active galactic nuclei (AGN), the source of energy is the gravitational potential energy of the material being accreted by a supermassive black hole via an accretion disk (e.g. Krolik 1999). This nuclear region is surrounded by a torus consisting of ionized, atomic and molecular gas components (e.g. Urry \\& Padovani 1995 and references therein). Studying the kinematics and distribution of the gaseous components in the circumnuclear region is important for understanding a number of aspects, such as the fueling of the AGN activity, the anisotropy of the radiation field and thereby testing the unified schemes for active galaxies, interaction of the jets with the external gas clouds and probing star formation in the central regions of an AGN. At radio wavelengths, the ionized component may be investigated via polarization observations of compact radio structure in the nuclear regions of AGNs (e.g. Saikia \\& Salter 1988; Mantovani et al. 1994; Udomprasert et al.  1997; Junor et al.  1999; Saikia \\& Gupta 2003; Rossetti et al.  2009). These include the compact radio cores, the nuclear radio jets, and the compact steep spectrum (CSS) and Gigahertz peaked spectrum (GPS) sources.  CSS sources are defined as having a projected linear size $<$15 kpc (H$_o$ = 71 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm m}$ = 0.27, and $\\Omega_{\\Lambda}$ = 0.73), and  a steep high-frequency radio spectrum ($\\alpha > 0.5$, where S$(\\nu)\\propto\\nu^{-\\alpha}$). GPS sources have spectra which turn over at, or above, 1 GHz, and are more compact than CSS objects  whose spectra could turn over below 1~GHz.  It is believed that GPS sources evolve into CSS objects, which later evolve to form the larger radio galaxies and quasars (Fanti et al.  1995; Readhead et al.  1996; O'Dea 1998; Snellen et al. 2000).  An important way of investigating atomic gas on subgalactic scales is via H{\\sc i} absorption towards the compact components of CSS and GPS sources or the radio nuclei of larger objects (e.g. van Gorkom et al. 1989; Conway \\& Blanco 1995; Peck et al. 2000; Pihlstr\\\"om et al. 2003; Vermeulen et al. 2003; Gupta et al. 2006; Morganti et al. 2009 and references therein).  These studies have shown that approximately 50\\% of GPS objects exhibit H{\\sc i} absorption compared with about 35\\% for  CSS sources (cf. Gupta et al.  2006). The H{\\sc i} column density also exhibits an anticorrelation with source size (Pihlstr\\\"om et al. 2003; Vermeulen et al. 2003; Gupta et al. 2006). The H{\\sc i} spectra exhibit a variety of line profiles with substantial red and blue shifts from the systemic velocities (Vermeulen et al. 2003; Gupta et al.  2006), suggesting that the atomic gas possesses complex motions, and may be out-flowing or in-falling, interacting with the jets, or rotating around the nucleus.  Another interesting class of AGN are those which exhibit signs of episodic activity. In radio-loud AGN these can be seen as two or more pairs of radio lobes on opposite sides of the active nucleus, or as young radio lobes embedded in diffuse emission from earlier cycles of activity. The objects with two or more pairs of distinct lobes have been designated double-double radio galaxies or DDRGs (e.g. Schoenmakers et al. 2000; Saikia et al. 2006).  Saikia et al. (2007) reported the detection of H{\\sc i} absorption towards the central region of the DDRG J1247+6723, and suggested from available information that there might be a strong correlation between the detection of H{\\sc i} in absorption and the occurrence of rejuvenated radio activity. The detection of H{\\sc i} absorption in the rejuvenated radio galaxy 4C~29.30 is also consistent with this trend (Chandola et al. 2010).  In order to extend H{\\sc i}-absorption investigations to a larger number of GPS and CSS objects, especially those that are of lower luminosity or more distant, and also towards the cores of larger sources and the central regions of rejuvenated radio sources, we have been observing these sources both with the Arecibo 305-m telescope and the Giant Metrewave Radio Telescope (GMRT). The results obtained in the first phase of this study were reported by Gupta et al. (2006). Recently, the spectral region between 700 and 800 MHz became temporarily available to radio astronomers due to the U.S. television switch from analog to digital transmissions, with new frequency allocations. We have used this opportunity to begin a search for highly redshifted H{\\sc i} and OH absorption within this band against the continuum emission from  CSS/GPS radio sources of appropriate redshift.  The first source observed, the well-known galaxy CTA~21 (4C~16.09), shows strong H{\\sc i} absorption.  We summarize the properties of CTA~21 in Section 2, the observations and results obtained with the Arecibo telescope in Section 3, and provide a discussion and concluding remarks in Section 4. ",
        "conclusions": "\\label{discussion} With a total column density of {\\it N}(H{\\sc i})$= 7.92 \\times 10^{20}$~cm$^{-2}$, CTA~21 is a source with a high H{\\sc i} column density.  In the compilation of H{\\sc i} absorption observations towards CSS and GPS sources by Gupta et al. (2006), there are 96 sources in their `full sample', of which only five (J0111+3906, J1357+4354, J1415+1320, J1819-6345, J1945+7055) have measured column densities which are higher. It is also the highest redshift CSS or GPS source for which H{\\sc i} absorption has been detected. With an overall linear size of $\\sim$0.3 kpc, it is close to the upper envelope of the {\\it N}(H{\\sc i}) vs projected linear size diagram (Pihlstr\\\"om et al.  2003; Vermeulen et al. 2003; Gupta et al. 2006), and consistent with the overall trend. The redshift of the strongest absorption component in Table~1, $z \\sim 0.9057$, agrees well with the optically-determined value of 0.907 (Labiano et al. 2007).  The highest-resolution VLBI images at $\\sim$5 and 15 GHz (Jones 1984; Kellermann et al. 1998), with angular resolutions of a few milliarcsec, do not show evidence for a core component. A comparison of these two images suggests that the prominent components have steep spectral indices. The estimated upper limit on the core flux density of $\\sim$10 mJy at 15 GHz implies that the fraction of emission from the core is less than $\\sim$1.5\\% at this frequency, which corresponds to an emitted frequency of $\\sim$29 GHz.  Although the {\\it N}(H{\\sc i}) vs core fraction diagram to test the unified scheme for active galaxies and probe the geometry of the H{\\sc i} disk has a large scatter (e.g. Gupta \\& Saikia 2006), CTA~21 appears broadly consistent with this.  The non-detection of a radio core is also consistent with its identification as a radio galaxy.  The existence of a compact double-lobed structure of size  $\\sim$12 mas seen in the highest-resolution VLBI images (Jones 1984; Kellermann et al. 1998), plus the more extended diffuse emission discussed above, raises the possibility that CTA~21 may be undergoing repeated cycles of activity. For example, the image of Kellermann et al.  (1998) shows evidence of weaker emission on opposite sides separated by $\\sim$40 mas, while Dallacasa et al. (1995) find evidence of diffuse emission extending up to $\\sim$300 mas towards the south.  The different resolutions of these images make it difficult to reliably estimate spectral indices for the different components.  However, on the basis of the compact double-lobed structure and diffuse extended emission we classify CTA~21 as a candidate rejuvenated radio galaxy. Rejuvenated radio sources cover a large range of linear sizes (see Saikia \\& Jamrozy 2009, for a review), and it has been suggested that jet activity in compact radio sources may be intermittent on time scales of $\\sim$10$^4$$-$10$^5$ yr (Reynolds \\& Begelman 1997). Saikia et al. (2007) and Chandola et al. (2010) explored a possible relationship between rejuvenation of radio or jet activity and the detection of H{\\sc i} in absorption. Unfortunately the number of sources is still small because most rejuvenated radio sources have weak radio emission in the central or nuclear region. The well-known examples of rejuvenated radio sources where H{\\sc i} absorption has been detected are the giant radio galaxy 3C236 which has a projected linear size of $\\sim$4250 kpc and exhibits evidence of star formation (Conway \\& Schilizzi 2000), the giant radio galaxy J1247+6723 with a GPS core (Saikia et al. 2007), the misaligned DDRG 3C293 (Beswick et al. 2004) which also exhibits fast outflowing gas blue-shifted by up to $\\sim$1000 km s$^{-1}$ (Emonts et al. 2005), the large southern radio galaxy Centaurus A (Sarma et al. 2002; Morganti et al. 2008), and the rejuvenated radio galaxy 4C~29.30 (Chandola et al.  2010).  The archetypal 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 significant nuclear H{\\sc i} absorption (Conway 1999).  While the sample size needs to be increased, the detection of absorbing H{\\sc i} gas in  rejuvenated galaxies appears to be even more frequent than for CSS and GPS objects.  Considering the GPS objects listed by Gupta et al. (2006), these have the highest H{\\sc i} detection rate of $\\sim$50\\%, and a median column density of $\\sim$3$\\times$10$^{20}$ cm$^{-2}$. In comparison, the rejuvenated radio galaxies discussed here have column densities in the range of $\\sim$8$-$50$\\times$10$^{20}$ cm$^{-2}$, and  tend to exhibit complex multi-component absorption profiles. The estimated column density of CTA~21 would be consistent with the range for the other rejuvenated radio galaxies."
    },
    "1004/1004.3760_arXiv.txt": {
        "abstract": "We present new dynamical models of the merger remnant NGC~7252 which include star formation simulated according to various phenomenological rules.  By using interactive software to match our model with the observed morphology and gas velocity field, we obtain a consistent dynamical model for NGC~7252.  In our models, this proto-elliptical galaxy formed by the merger of two similar gas-rich disk galaxies which fell together with an initial pericentric separation of $\\sim2$ disk scale lengths approximately $620$~Myr ago.  Results from two different star formation rules--- density-dependent and shock-induced--- show significant differences in star formation during and after the first passage.  Shock-induced star formation yields a prompt and wide-spread starburst at the time of first passage, while density-dependent star formation predicts a more slowly rising and centrally concentrated starburst.  A comparison of the distributions and ages of observed clusters with results of our simulations favors shock-induced mechanism of star formation in NGC~7252.  We also present simulated color images of our model of NGC~7252, constructed by incorporating population synthesis with radiative transfer and dust attenuation.  Overall the predicted magnitudes and colors of the models are consistent with observations, although the simulated tails are fainter and redder than observed.  We suggest that a lack of star formation in the tails, reflected by the redder colors, is due to an incomplete description of star formation in our models rather than insufficient gas in the tails. ",
        "introduction": "NGC~7252 is a well-studied merger remnant \\citep{too77,s83}, which provides strong observational evidence supporting the idea that at least some disk galaxy mergers produce elliptical galaxies.  Its luminosity profile follows a $r^{1/4}$ law \\citep{s82}, as do the profiles of elliptical galaxies \\citep{deV, kor}.  The core contains a central counter-rotating disk of ionized gas within $8\\arcsec$ ($\\simeq2.5$~kpc\\footnote{In order to compare with previous observational results, throughout this paper we adopt $H_{0}=100h=75$~km~s$^{-1}$~Mpc$^{-1}$.  This places NGC~7252 at a distance of 64.4~Mpc and a projected scale of $1\\arcsec=$~312~pc.}) of the center, providing a kinematic signature of a major merger event in the recent past \\citep{s82}.  Molecular gas is also closely associated with the central disk \\citep{dup,wang}.  The two tidal tails are rich in H\\,{\\small I}, suggesting that the progenitors were late-type spirals \\citep[][hereafter HGvGS94]{hb94}.  About $500$ candidate young clusters are found in the main body in the {\\em Hubble Space Telescope} images \\citep{whit,m97}, indicating that clusters can form in large numbers in the violent star formation that accompanies a galaxy merger, and that cluster formation can increase the specific globular cluster frequency of the merger remnant.  Moreover, since the clusters are single burst stellar populations, the gravitationally bound clusters surviving several $10^{8}$ years provide a record of the merger's star formation history.  Ages and metallicities of these clusters can be used to reconstruct that history; \\citet[hereafter SS98]{ss} found five young globular clusters with ages of $400$--$600$~Myr, which presumably formed shortly after the first close encounter of the two gas-rich disk galaxies that later merged to produce the present-day remnant.  Metallicities of these clusters support the idea that elliptical galaxies with bimodal globular cluster systems can form through major mergers.  In addition to the extensive observations of NGC~7252, many numerical models have been made to investigate the history of the system.  \\citet[hereafter HM95]{hm} revised encounter parameters from previous studies \\citep{br, m93} with detailed H\\,{\\small I} observations from the Very Large Array (VLA, HGvGS94).  Their dynamical simulation determined the age of NGC~7252 to be about $773$ ($\\sim580h^{-1}$)~Myr after periapse of the initial orbit and predicted the future infall rate of the tidal tail materials.  \\citet{mdh} further constrained the dark halos in the NGC~7252 progenitor galaxies.  These models matched the morphology and kinematics of the system very well, and have provided robust results for the dynamical history.   However, gas dynamics and star formation have not yet been included in any dynamical modeling of NGC~7252.  \\citet{ag} attempted to reconstruct NGC~7252's star formation history and spectro-photometric properties using an evolutionary spectral synthesis model, but were forced to postulate that the interaction-induced star formation began well before the first passage.  In this paper we use simulations including gas dynamics and star formation to construct possible star formation histories for NGC~7252.  We follow \\citet{b04} in examining how two different star formation rules--- \\dd and $shock$-$induced$--- affect the star formation history.  In addition to matching the observed morphology and H\\,{\\small I} kinematics, we compare our star formation models with the ages of young globular clusters observed by SS98.  We examine the different star formation models and try to constrain the triggers of star formation in merging galaxies.  Moreover, we also compare the modeled colors of NGC~7252 with the observed values.  Our goal is to match the overall morphology, kinematics and photometric properties while insuring that the initial disks have reasonable parameters for late-type spirals.  This represents a new and significant step beyond pure dynamical modeling.  This paper is organized as follows: in Section~2 we describe the methodology and the simulation technique, while in Section~3 we describe the matching process to the H\\,{\\small I} observation as well as our best-fit model.  In Section~4 we present the predicted star formation history and a comparison with the ages of SS98 clusters.   We discuss the simulated images of NGC~7252 and a comparison of the photometric properties in Section 5, and finally give conclusions in Section 6.  For the purpose of a consistent discussion throughout the paper, we introduce physical scales in Section 2 using the scaling factors of the best-fit model, which will be described in Section~3. ",
        "conclusions": "We describe new simulations of the galactic merger NGC~7252 with gas dynamics and star formation included.  In our models NGC~7252 is formed by the merger of two similar gas-rich disk galaxies which fell together with an initial pericentric separation of $\\sim1.9$~disk scale lengths about 620~Myr ago.  Starting with plausible models for the pre-encounter disks, our simulation produce fairly well matches with the observed morphology and kinematics.  We emphasize that rather than a perfect reproduction, our goal is to present a plausible dynamical model of NGC~7252, which provides insight into the star formation history and matches the photometric properties; this nonetheless represents a new and significant step beyond pure dynamical modeling.  Following the study of NGC~4676 by \\citet{b04} and \\citet{c07}, we consider two star formation mechanism in our simulation and use the ages of young globular clusters from \\citet{ss} to constrain the trigger of star formation.  \\citet{b04} suggested that \\dd star formation rule gives an incomplete description of large scale star formation in interacting galaxies.  In our simulations of NGC~7252, \\dd star formation happens mostly at gas rich regions concentrated in galactic centers throughout the simulation, whereas more products of past star formation in \\si model are distributed at wider projected distances and in tails.  The total mass born in bursts are on the same order in each model, yet the median projected radius of the burst populations is 0.6~kpc in \\dd model and 2.1~kpc in \\si model.  This supports the suggestion in \\citet{b04} that \\si star formation may provide a better match to the extended distributions of young clusters in NGC~7252 and other merger remnants.  When compared with observations, the star formation histories of both models successfully reproduce the range of ages spanned by the observed clusters.  Those that have ages of $\\sim570$--$600$~Myr, in particular, are likely to form during the starburst at first passage, and this seems consistent with the prediction of the \\si simulation.  Within the projected annulus where the observed clusters lie, the age distribution of stellar particles formed in the \\si simulation shows a peak around such age, while in the \\dd simulation the age distribution is flat and declining.  Again, almost all interaction-induced star formation occurred in the very central regions for the \\dd model, and the \\si model is more likely to explain the concentration of ages of clusters observed in SS98.  We also describe a method to incorporate simple stellar population synthesis models into our dynamical simulations in order to compare the predicted photometric properties with the observed values.  Beyond the attempt of matching the morphology as in many previous simulations of interacting galaxies, we also try to match the photometric properties.  While the overall magnitudes and colors are consistent with the observed values, the tails seem to be fainter and redder than observed.  We argue that this lack of star formation, reflected by the redder colors, is due to an incomplete description of star formation, rather than insufficient material in the tails.  An additional mode of star formation which is independent of large-scale conditions could help to solve this problem.  Finally we discuss some future directions for making our star-forming models more realistic, which could help improve the modeled photometric properties.  First, as discussed in Section \\ref{mphot}, our models do not incorporate the gradual mass return from evolving stellar populations.  In order to incorporate population synthesis models to produce luminosities, we have to make adjustments which are not entirely straightforward.  Moreover, recycled gas might help boost the star formation rate throughout the last stages of the merger, improving the luminosities and colors of our models of NGC~7252.  Another concern regarding the colors of the merger remnant tails comes from the colors of the progenitor disks.  Disk galaxies generally become moderately but systematically bluer outward from the bulge, and this effect is caused by the age and metallicity gradients in the stellar populations as well as internal dust extinction \\citep[e.g.][]{wk87, BdJ00}.  In order to improve our modeling of the progenitor galaxies, we could include radial gradients of age and metallicity in the initial stellar disks.  This could influence the photometric properties of the remnant tails, since they mostly come from the outer parts of the progenitor disks.  With regard to our models of NGC~7252, it might also help to improve the luminosities and colors of the tails.  However, unless star formation in the outer regions and tails can be sustained throughout the merger process, the outer populations will fade fast after the encounter, reducing the effect of age gradients on tail properties."
    },
    "1004/1004.4971_arXiv.txt": {
        "abstract": "We study the line widths in the [\\ion{O}{3}]$\\lambda$5007 and H$\\alpha$ lines for two groups of planetary nebulae in the Milky Way bulge based upon spectroscopy obtained at the \\facility{Observatorio Astron\\'omico Nacional in the Sierra San Pedro M\\'artir (OAN-SPM)} using the Manchester Echelle Spectrograph.  The first sample includes objects early in their evolution, having high H$\\beta$ luminosities, but [\\ion{O}{3}]$\\lambda 5007/\\mathrm H\\beta < 3$.  The second sample comprises objects late in their evolution, with \\ion{He}{2}~$\\lambda 4686/\\mathrm H\\beta > 0.5$.  These planetary nebulae represent evolutionary phases preceeding and following those of the objects studied by \\citet{richeretal2008}.  Our sample of planetary nebulae with weak [\\ion{O}{3}]$\\lambda$5007 has a line width distribution similar to that of the expansion velocities of the envelopes of AGB stars, and shifted to systematically lower values as compared to the less evolved objects studied by \\citet{richeretal2008}.  The sample with strong \\ion{He}{2}~$\\lambda 4686$ has a line width distribution indistinguishable from that of the more evolved objects from \\citet{richeretal2008}, but a distribution in angular size that is systematically larger and so they are clearly more evolved.  These data and those of \\citet{richeretal2008} form a homogeneous sample from a single Galactic population of planetary nebulae, from the earliest evolutionary stages until the cessation of nuclear burning in the central star.  They confirm the long-standing predictions of hydrodynamical models of planetary nebulae, where the kinematics of the nebular shell are driven by the evolution of the central star. ",
        "introduction": "Hydrodynamical models of planetary nebulae have long predicted a particular kinematic evolution for the nebular shells, driven primarily by the evolution of the central stars \\citep[e.g.,][]{kwoketal1978, kahnwest1985, schmidtvoigtkoppen1987a, schmidtvoigtkoppen1987b, breitschwerdtkahn1990, kahnbreitschwerdt1990, martenschonberner1991, mellema1994, villaveretal2002, perinottoetal2004, schonberneretal2007}.  Initially, the central stars are cool and their winds relatively slow.  This wind interacts with the wind that the precursor asymptotic giant branch (AGB) star emitted in a momentum-conserving mode \\citep{kwok1982}.  However, the central star's temperature and wind velocity increase rapidly, with the consequences that an ionization front is driven through the AGB envelope and the interaction between the two winds switches to an energy driven mode, and a hot bubble is created behind the shocked wind.  The ionization front first accelerates the AGB envelope, now seen as the rim of the planetary nebula.  In time, once the internal pressure of the hot bubble exceeds that of the nebular shell, it further accelerates the nebular shell.  Theoretically, the latest phases of evolution are less clear, though the central star will cease nuclear energy generation, fade rapidly, and emit an ever-weaker wind, in principle allowing the inner part of the nebular envelope to backfill into the central cavity \\citep[e.g.,][]{garciaseguraetal2006}.  Although the many extant observations of the kinematics of planetary nebulae all show expansion of the nebular shell, there are few systematic observations of how these shells acquire their motion and how it evolves with time.  \\citet{dopitaetal1985, dopitaetal1988} were the first to provide observational support for the early acceleration of the nebular shell from studies of planetary nebulae in the Magellanic Clouds.  Studies of Milky Way planetary nebulae provided much less convincing results \\citep[e.g.,][]{chuetal1984, heap1993, medinaetal2006}.  Recently, \\citet{richeretal2008} demonstrated the acceleration of nebular shells in bright planetary nebulae in the Milky Way bulge during the early evolution of their central stars and were able to associate the acceleration seen in different evolutionary stages with the phases of acceleration expected from theoretical models.  %  Here, we undertake a study that complements \\citet{richeretal2008}, selecting objects earlier and later in their evolution than they did. The addition of these objects % allows us to study the evolution of the kinematics of the nebular shell from the earliest stages of the planetary nebula phase until the cessation of nuclear burning in the central stars. In section \\ref{sec_observations}, we present our new data and their analysis.  In section \\ref{sec_results}, our results and their implications are outlined, the principal ones being the similarity of the line widths in H$\\alpha$ and [\\ion{O}{3}] $\\lambda 5007$, that our sample of least evolved objects has a line width distribution shifted to the lowest values while the sample of most evolved objects has a size distribution with the largest sizes, and that the evolutionary state correlates with the H$\\beta$ luminosity.  % In section \\ref{sec_conclusions} we summarize our conclusions. ",
        "conclusions": "We have obtained kinematic data for two samples of planetary nebulae in the Milky Way bulge, selected so as to include objects very early and late in their evolution (\\S \\ref{sec_sample_def}).  % We measure line widths for H$\\alpha$ and [\\ion{O}{3}] $\\lambda 5007$ in most cases.  We combine these data sets with that studied by \\citet{richeretal2008}.  We define four evolutionary groups, based upon the temperature of the central star, which allow us to study the kinematics of the nebular shell from the earliest phases until the central star ceases nuclear burning.  Generally, we find a near equality of the line widths for the H$\\alpha$ and [\\ion{O}{3}] $\\lambda 5007$ lines in any given object. Ionization stratification likely accounts for the deviations: The [\\ion{O}{3}] $\\lambda 5007$ line widths are systematically smaller than the H$\\alpha$ line widths for the smallest and largest H$\\alpha$ line widths, corresponding to the earliest and latest evolutionary phases, respectively.  We see clear evolution of the kinematics of the nebular shell.  The least evolved objects, our planetary nebulae with weak [\\ion{O}{3}] $\\lambda 5007$, have cool central stars  and the nebular envelopes have a line width distribution similar to that of the envelope expansion velocities of AGB stars, indicating that the ionization front has not yet been able to substantially accelerate the nebular shell.  In subsequent phases \\citep{richeretal2008}, the nebular shell is first accelerated by the passage of an ionization front and then further accelerated once the central star's wind produces a hot bubble.  The line width distributions for the planetary nebulae in these three evolutionary phases are statistically distinct.  The most evolved objects, with high \\ion{He}{2}~$\\lambda 4686$ ratios, have a similar line width distribution to evolved [\\ion{O}{3}] $\\lambda 5007$-bright objects, suggesting that no further acceleration occurs as the central stars reach their highest temperatures, their nuclear reactions cease, and their winds decline.  %  This kinematic evolution of the nebular shell has long been predicted by hydrodynamical models \\citep[e.g.,][]{kahnwest1985, martenschonberner1991, mellema1994, villaveretal2002, perinottoetal2004}.  % Our results, together with those of \\citet{richeretal2008}, based upon a large sample of Galactic planetary nebulae from a single stellar population, clearly confirm these predictions.  At least until the point at which nuclear reactions cease in the central stars, their ionizing fluxes and winds continuously accelerate the nebular envelopes.  What happens thereafter is not yet completely clear, and would require samples of planetary nebulae chosen specifically to contain hot central stars of low luminosity.  %"
    },
    "1004/1004.0696_arXiv.txt": {
        "abstract": "A significant fraction of the mature FGK stars have cool dusty disks at least an orders of magnitudes brighter than the solar system's outer zodiacal light. Since such dusts must be continually replenished, they are generally assumed to be the collisional fragments of residual planetesimals analogous to the Kuiper Belt objects. At least 10\\% of solar type stars also bear gas giant planets. The fraction of stars with known gas giants or detectable debris disks (or both) appears to increase with the stellar mass. Here, we examine the dynamical evolution of systems of long-period gas giant planets and residual planetesimals as their host stars evolve off the main sequence, lose mass, and form planetary nebula around remnant white dwarf cores. The orbits of distant gas giant planets and super-km-size planetesimals expand adiabatically. During the most intense AGB mass loss phase, sub-meter-size particles migrate toward their host stars due to the strong hydrodynamical drag by the intense stellar wind.  Along their migration paths, gas giant planets capture and sweep up sub-km-size planetesimals onto their mean-motion resonances. These planetesimals also acquire modest eccentricities which are determined by the mass of the perturbing planets, the rate and speed of stellar mass loss. The swept-up planetesimals undergo disruptive collisions which lead to the production of grains with an extended size range. The radiation drag on these particles is ineffective against the planets' resonant barrier and they form 30-to-150-AU-sizes rings which can effective reprocess the stellar irradiation in the form of FIR continuum. We identify the recently discovered dust ring around the white dwarf WD 2226-210 at the center of the Helix nebula as a prototype of such disks and suggest such rings may be common. ",
        "introduction": "The discovery of over 300 planets shows that at least 10 \\% of nearby solar type stars have Jupiter-mass planets around them (Cumming {\\it et al.} 2008). According to widely adopted sequential accretion hypothesis, these planets formed through condensation of heavy elements, emergence of planetesimals which coagulated into protoplanetary embryos and cores prior to efficient gas accretion. Through a series of population-synthesis simulations based on this scenario (Ida \\& Lin 2004a,b, 2005, 2008a,b), we were able to reproduce the observed mass-period ($M_p-P$) distribution of the known gas giant planets (Schlaufman {\\it et al.} 2008) around a modest fraction ($\\eta_J$) of FGK main sequence dwarf stars.  Mass accretion rate $\\dot M$ from protostellar disks onto stars more massive than the Sun is observed to increase with square of the stellar mass {\\it i.e} $M_\\ast^2$ (Natta {\\it et al.} 2006, Garcia-Lopez {\\it et al.} 2006, cf Clarke \\&  Pringle 2006). If the initial mass of these disks also increases with  $M_\\ast$, the greater availability of heavy elements would provide more rapid growth and larger of protoplanetary embryos around intermediate mass stars (Ida \\& Lin 2004b, 2005, Kennedy {\\it et al.} 2007). Furthermore, the extended local surface density and pressure maxima at the ionization radius (where the mid-plane temperature $T_c \\simeq 1000$ K), snow line (Kretke {\\it et al.}  2009), and outer edge of the dead zone provide more efficient barrier against hydrodynamic drag for the sub-km grains and type-I migration for the critical-mass ($\\sim 10 M_\\oplus$) cores.  The enhance retention efficiency of these building blocks enhance emergence probability of gas giants (Laughlin {\\it et al.} 2004, Ida \\& Lin 2004b).  On the observational side, high-precision radial-velocity surveys of intermediate-mass main sequence stars have been handicapped by their hot and active atmosphere as well as their fast spin (Griffin {\\it et al.} 2000). However, these problems are significantly reduced when these stars evolve off their main sequence. Several planets have been found around G sub-giants and giants (Sato {\\it et al.}  2005) which have more expanded envelopes and relatively cool atmospheres. Preliminary surveys indicate that the fraction of these relatively massive G giant stars with gas giant planets ($\\eta_J \\sim 1/3$) may be more than double that around the solar-type G dwarf main sequence stars (Sato {\\it et al.} 2007, Johnson {\\it et al.} 2007, Dollinger {\\it et al.}  2007, Hatzes 2008, Johnaon 1008).  Another area where observations provide valuable clues on the planet formation process is the search for residual building-block planetesimals around their host stars. Since 1995 a population of Kuiper Belt objects (KBO's) has been found in the solar system and beyond the orbit of Neptune (Jewitt \\& Luu 1995, Brown {\\it et al.} 2005). These objects are thought to be the parent bodies of a diffuse dusty ring which emits cold zodiacal light with far infrared (FIR) radiation.  A series of recent systematic 70 $\\mu$m (FIR) observational surveys with the Spitzer Infrared Space Telescope indicate that 1) none of the nearby M stars show detectable excess (Gautier {\\it et al.} 2007), 2) 15 \\% of 274 FGK stars show excess (Bryden {\\it et al.} 2006, Beichman {\\it et al.} 2006, Trilling {\\it et al.} 2008), 3) 30 \\% of 160 A stars also show excess (Rieke {\\it et al.} 2005; Su \\& Rieke 2007). The threshold level for a marginal detection by the Spitzer Telescope is at least an order of magnitude more intense that of the solar system zodiacal light.  Since both the fraction of stars with gas giants ($\\eta_J$) and that with detectable debris disks ($\\eta_d$) increases with the stellar mass ($M_\\ast$), we can extrapolate that a major fraction of intermediate-mass stars may have both gas giants and persistent, rich debris disks.  This inferred association of gas giant planets with debris disks is highlighted by the recent discovery (with adaptive imaging technique) of three gas giants and a debris disk outside them around an A5V star HR 8799 (Marois {\\it et al.} 2008). The mass and periods of these three planets are (10, 10, 7) $M_J$ and (24, 38, 68) AU.  Around this $M_\\ast= 1.5 M_\\odot$ star, the ice line is located at $\\sim$10 AU. In \\S2, we provide a more detailed discussion on the association of debris disks and planet formation in the Solar System and around nearby stars.  With the evidence of existence of many detectable debris disks and gas giant planets around nearby, mature, intermediate-mass stars, we now consider their dynamical evolution during the post main sequence evolution of their host stars. The goal of our study here is to identify whether the inconspicuous planet building blocks may reveal their presence. The analysis to be presented below is particularly relevant for the AGB mass loss phase where the mass loss rates range up to $\\sim 10^{-4} M_\\odot$ yr$^{-1}$ and where most of the stars may reach this mass loss rate near the end of the AGB phase (Willson 2000). The typical outflow speed is $\\sim 10$ km s$^{-1}$ and $\\dot M_\\ast$ may vary on the time scale of $10^{3-4}$ yr (Weidemann 1987, Blocker 1997). After losing mass on the AGB, these stars rapidly evolve across the HR diagram to form planetary nebula around emerging white dwarfs. During the post AGB phase, although high-speed mass loss continue to occur, the mass flux is much reduced (cf Kwok 1982).  In this paper, we mostly consider stars with a main-sequence masses in the range of $2 M_\\odot < M_\\ast < 4 M_\\odot$ (A-F stars). We choose these stars not only because many, perhaps most of them, have gas giants and debris disks (similar to those around the less massive Sun and HR 8799), but they are also thought to be the progenitors of some well known planetary nebulae such as the Helix nebula where the mass of the remnant white dwarf WD 2226-210 is $\\sim 0.58 M_\\odot$ (Su {\\it et al.} 2007). This rapid and large-magnitude change in the gravitational potential can lead to great orbital expansions for their planets. This evolutionary tendency may lead to the possibility of resolvable HST imaging of long-period planets which are otherwise undetectable.  Evidence for the existence of residual planetesimals comes from the recent discovery of a dusty ring around a young white dwarf WD 2226-210 at the center of the Helix nebula (Su {\\it et al.}  2007). This ring extends between about 35 and 150 AU from a central white dwarf WD 2226-210 (well interior to the complex, helix structure), and have a total mass of about 0.13 $M_ \\varoplus $. We propose that this ring is the byproduct of planets and planetesimals' orbital evolution during the epoch when the central star rapidly lost most of its mass. Based on the extrapolation of the system around HR 8799, our basic conjectures are:  \\noindent 1) gas giant planets and debris disks are common around A main sequence stars;  \\noindent 2) at the epoch of their main sequence turnoff, these relatively massive stars have debris disks with more than 10$M_\\oplus$ in total mass and most of it is contained in km-size residual planetesimals  \\noindent 3) gas giant planets' orbits expanded adiabatically by nearly an order of magnitude while there host stars evolve through AGB mass loss phase with rapid stellar mass loss;  \\noindent 4) the KBO-equivalent planetesimals are captured onto the main motion resonance of gas giants; and  \\noindent 5) through collisions, these parent bodies generate dust which reprocess the radiation of their central star.  These basic assumptions are based on their solar system and HR 8799 analogues. A detailed and systematic analysis of these processes is useful for understanding the long-term evolution of this class of planetary systems. In \\S3, we first construct a working model for a solar-system like configuration. The orbital evolution of the planets and planetesimals is computed and analyzed in \\S4. Here we take into account of the central star's mass loss, the hydrodynamic drag by the expanding envelope and the planet's dynamical perturbation on the residual planetesimals. In \\S5, we show that collisions between the resonant planetesimals lead to their disruptions and fragmentation. We derive the condition for the retention of dust grains into rings outside the orbits of some gas giant planets. We model the structure of the disk around WD 2226-210 in \\S6.  Finally, we summarize our results and discuss their implication, especially in the context of possible detectable gas giant planets around young white dwarfs in \\S7. ",
        "conclusions": "In this paper we construct a dynamic model for the post main sequence evolution of planetary systems. With this model, we find that as a consequence of mass loss from the central star, km-size planetesimals can be captured by mean-motion resonance and form a disk beyond the planets. During the subsequent orbital expansion, the eccentricity of the resonant planetesimals is excited which enhances their collision frequency. Highly destructive impacts among these planetesimals can continuously provide a ring of dust around the mature star, which can be seen in the FIR and MIR wave band. This phenomenon can also be extended to other A-F stars, with the inference that a large fraction of stars may have planet systems and substantial debris disks during their main sequence epoch. Stars with higher main sequence mass, greater mass loss rate and lower velocity of star wind tend to have higher ability to capture the planetesimals.  We applied this model to account for the dust ring around WD 2226-210 at the center of the Helix nebula. It can also be used to study other young white dwarfs. For example, we anticipate this process will occur when the Sun evolve to its AGB mass loss phase, albeit among of fragmentary dust would be limited and the intensity of the reprocessed radiation may be weak. Since this process depends sensitively on the initial heavy element content and the presence of sufficiently massive planets, we expect there to be large dispersion in the signatures of dust debris around young white dwarfs. Nevertheless, we expect to see more dusty rings around planetary nebulae with A star progenitors because the fractions of stars with planets and debris disks is observed to increase with stellar mass.  The present work is develpoed for the $\\rho_{wind} \\propto 1/r^2$ (or $v_{wind}$~=~constant) portion of a stellar wind and thus not applicable to the region very close to the star. In addition, apart from the planetesimals that are captured by the planet, the fate of other planetesimals also need to be discussed. These may also play an important role in the evolution of the white dwarf.  Some observational tests of these models should be possible with present technology or will be possible soon. As the planetesimals are collected by the giant planet, it is most likely that the gas giant planet survives the AGB period, and the orbit changes to several times the original one. If such giant planet around white dwarf can be detected, it would provide support to our model. Although there are several failed attempts in these searches (Debes {\\it et al.} 2005a, b, 2007), the upper mass limits set for the unseen planet is far larger than that needed in our model.  We thank Sverre Aarseth for provide the Hermit integrator, Y.-H. Chu and an anonymous referee for useful correspondence and helpful suggestions. This work is supported by NSFC(10233020), NCET (04-0468), NASA (NNX07A-L13G, NNX07AI88G, NNX08AM84G), JPL (1270927), and NSF(AST-0908807)."
    },
    "1004/1004.0005_arXiv.txt": {
        "abstract": "We study the formation of disc galaxies in a fully cosmological framework using adaptive mesh refinement simulations. We perform an extensive parameter study of the main subgrid processes that control how gas is converted into stars and the coupled effect of supernovae feedback. We argue that previous attempts to form disc galaxies have been unsuccessful because of the universal adoption of strong feedback combined with high star formation efficiencies. Unless extreme amounts of energy are injected into the interstellar medium during supernovae events, these star formation parameters result in bulge-dominated S0/Sa galaxies as star formation is too efficient at $z\\sim 3$. We show that a low efficiency of star formation more closely models the sub-parsec physical processes, especially at high redshift. We highlight the successful formation of extended disc galaxies with scale lengths $r_{\\rm d}=4-5\\kpc$, flat rotation curves and bulge-to-disc ratios of B/D$\\,\\sim1/4$. Not only do we resolve the formation of a Milky Way-like spiral galaxy, we also observe the secular evolution of the disc as it forms a pseudo-bulge. The disc properties agree well with observations and are compatible with the photometric and baryonic Tully-Fisher relations, the $\\Sigma_{\\rm SFR}-\\Sigma_{\\rm gas}$ (Kennicutt-Schmidt) relation and the observed angular momentum content of spiral galaxies. We conclude that the underlying small-scale star formation physics plays a greater role than previously considered in simulations of galaxy formation. ",
        "introduction": "\\label{sect:intro} The prevailing picture of galaxy formation emerged more than 30 yr ago \\citep[][]{WhiteRees78,FallEfstathiou80}. Within the framework of the broadly accepted $\\Lambda$ Cold Dark Matter ($\\Lambda$CDM) scenario \\citep{Komatsu2009}, gravity assembles structures in a bottom-up fashion. Haloes of dark matter acquire angular momentum via tidal torques \\citep{Peebles69,FallEfstathiou80} from interacting structures, and as gas cools and condenses into their central parts, star-forming galaxies form. A realistic angular momentum content can be accounted for if most of the angular momentum is retained in the assembly process. In this picture, the host halo is responsible for the final galaxy characteristics \\citep[e.g.][]{MoMaoWhite98}. While several aspects of the theory of galaxy formation are still being developed, e.g. the underlying physics of the missing satellite problem \\citep{Klypin1999,moore99} and the role of cold stream accretion \\citep[][]{Keres05,Keres09,Dekel09}, the model has proven successful for understanding global properties of galaxy assembly.  Given the complexity and non-linearity of the involved processes, computer simulations have become the ideal tool for studying the formation of structure. The formation of a late-type spiral galaxy, such as our own Milky Way, has been studied numerically in fully $\\Lambda$CDM cosmological context by many authors \\citep[e.g.][]{Abadi03a,SommerLarsen03,Governato04,Robertson04,Okamoto05,Governato07,croft09,Scannapieco09,Piontek09b,Agertz09b}. To date, no attempt has yielded a realistic candidate. The dominant reason for this is the so called \"angular momentum problem\" which leads to small, centrally concentrated discs dominated by large bulges \\citep{NavarroBenz91,NavarroWhite1994}. Merging substructures lose angular momentum to the outer halo via dynamical friction, forcing the associated baryons to end up in the central parts of the proto-galaxy as a spheroid rather than a disc. This poses a problem for the theoretical understanding of extended late-type galaxies. This might in part stem from numerical issues:  the commonly used Smoothed Particle Hydrodynamics (SPH) \\citep{GingoldMonaghan77,Lucy77} technique is known to incorrectly treat boundaries, hence poorly treating multiphase fluids \\citep[e.g.][]{Agertz07,Read2010}. This can lead to artificial angular momentum transfer at the interface between cold disc and a hot halo \\citep{Okamoto05}.  Many proposed solutions exists to the angular momentum problem, all amounting to the same process: keep the gas from cooling and forming stars too efficiently in the merging dark matter satellites at high redshift. One natural source is the cosmological UV background, being responsible for reionization at $z\\gtrsim6$ which heats the gas, preventing it to cool efficiently into star-forming dwarf galaxies \\citep{ThoulWeinberg96,Quinn96,gnedin00,hoeft06}. However, the impact on objects larger than $v_{\\rm circ}\\sim 10\\kms$ is unclear due to e.g. self-shielding and efficient collisional cooling \\citep{Dijkstra04}.  Gas in low-mass haloes can also be blown out by supernova driven winds \\citep{DekelSilk86,Efstathiou00}, hence lowering the resulting star formation efficiency (SFE), enriching the intergalactic medium (IGM) in the process. \\cite{MacLowFerrara99} demonstrated that while dwarf galaxies of mass $10^6-10^9\\Msol$ efficiently can expel metals in supernovae-driven winds, virtually no mass is lost for systems of mass $\\gtrsim10^7\\Msol$ \\citep[see also][]{dubois08}. The inefficiency in driving winds from dwarfs was also reported by \\cite{Marcolini06} who attributed this to the extended dark matter halo and efficient metal cooling. In this scenario, mass-loss and IGM enrichment will occur due to tidal and ram-pressure stripping \\citep[e.g.][]{mori00}. Phenomenological models of e.g. momentum driven winds have proven successful in reproducing the high-$z$ IGM \\citep{oppenheimerdave06} but it is uncertain how it regulates star formation and in what manner the expelled gas is re-accreted at later times \\citep{oppenheimer10}.  Various recipes of supernovae feedback have been developed for numerical simulations \\citep[e.g.][]{NavarroWhite93,Kay02,Scannapieco06}, and the methods have proven successful in removing low angular momentum material from central parts of galaxies \\citep[e.g.][]{SommerLarsen03,Okamoto05,Governato07}, yielding more extended galaxies in comparison to models without feedback. However, it is unclear to what extent this way of reducing star formation can account for disc-dominated spiral galaxies like the Milky Way. Recently \\cite{Scannapieco09} demonstrated numerically, in a fully cosmological setting, how a set of 8 Milky Way sized haloes failed to form significant discs. While half of the sample were early type galaxies resulting from late time mergers, the other half of the sample had less than 20 per cent of their stellar mass in discs. This can be a result of the inability of the adopted feedback to remove or redistribute low angular momentum material, but is also a strong indication that something else might regulate star formation at high redshift. On the same topic, \\cite{Sawala10} argues that modern simulations of dwarf galaxy formation \\citep{Valcke08,Stinson09,Governato10} all yield much larger stellar masses than expected from observations as well as gas-to-star conversion efficiencies almost an order of magnitude too large. \\cite{Dutton09} found that, for SNe feedback to yield realistic galaxies, it must be very efficient, converting 25 per cent of the SN energy into outflows. If too strong feedback is employed, the discs can be destroyed by internal processes as too much material is ejected into the halo, preventing efficient disc reformation from cold gas, and possibly violating the upper bounds of halo gas found in X-ray surveys [see \\cite{Bregman07} and references within]. In light of these studies, it is unclear if supernovae feedback is the sole agent in regulating star formation. Note that SNe explosions can regulate star formation in galaxies without expelling gas, being a driver of galactic turbulence \\citep{maclow:review04}.  Fundamentally, star formation is regulated by the availability of H$_2$. The observed Kennicutt-Schmidt (from now on \\emph{K-S}) relation \\citep{kennicutt98}, that relates $\\Sigma_{\\rm SFR}$ to $\\Sigma_{\\rm gas}$, varies strongly among individual spiral galaxies and can not be fit with a single power law \\citep{bigiel2008}. $\\Sigma_{\\rm SFR}$ behaves very differently for $\\Sigma_{\\rm gas}$ greater or smaller than $\\approx9\\Msol\\,{\\rm pc}^{-2}$, marking the transition from atomic to fully molecular star-forming gas \\citep{Leroy08}, and is dependent on gas metallicity, dust content, turbulence, small scale clumpiness and local dissociating UV field \\citep{mckeeostriker07}. The inclusion of these processes and its impact on global star formation in discs has recently been studied both numerically \\citep{RobertsonKravtsov08,Gnedin09,Pelupessy09} as well as analytically \\citep[e.g.][]{Krumholz09}. A natural outcome of this treatment is an order of magnitude lower amplitude of the \\emph{K-S} relation at high redshifts ($z\\sim3$) \\citep{GnedinKravtsov2010}. This agrees well with the observation of damped Ly$\\alpha$ systems \\citep[DLA:][]{WolfeChen06} as well as Lyman Break Galaxies \\citep{Rafelski09}. This indicates that star formation can be made inefficient at high redshift, leaving gas for late-time star formation in a disc like environment, but not necessarily by expelling gas in supernova-driven winds. In addition, \\cite{Murray10} argues that the disruption time-scale of giant molecular clouds (GMCs) due to jets, H{\\,\\small II} gas pressure, and radiation pressure also serves to regulate the SFE in galaxies. The disruption occurs well before the most massive stars exit the main sequence, meaning that supernovae in principle have little effect on GMC lifetimes.  In this paper we investigate to what extent supernovae feedback and the underlying small scale star-forming physics can affect the formation and evolution of realistic spiral galaxies in a fully cosmological setting. The former effect is studied via well tested numerical implementations of SNII, SNIa feedback coupled to metal enrichment, as well as stellar mass-loss. The latter influence is achieved by considering different normalizations of the Schmidt-law star formation efficiency. We conduct a comprehensive analysis of the resulting $z=0$ discs and compare them to observational relations.  The paper is organized as follows. In Section\\,\\ref{sect:num}, we describe the numerical method used in this work, including the adopted feedback and star formation prescriptions. In Section\\,\\ref{sect:IC}, we present the cosmological initial conditions and discuss the free parameters of this work. Section\\,\\ref{sect:discs} outlines the disc analysis and summarizes the final properties of the simulation suite. In Section\\,\\ref{sect:effect} and Section\\,\\ref{sect:FB}, we present a detailed analysis of the impact of small-scale SFE and supernova feedback respectively. In Section\\,\\ref{sect:observations} we compare our simulations to modern observations. Finally, Section\\,\\ref{sect:discussion} summarizes and discusses our conclusions. ",
        "conclusions": "\\label{sect:discussion} In this paper we have presented a set of AMR simulations studying the assembly of large Milky Way-like disc galaxies. The self-consistent formation of a late-type disc galaxy has remained elusive in the field of numerical galaxy formation, mainly due to the strong loss of angular momentum in the galaxy assembly process. A popular solution to this problem is to regulate star formation at high redshift via supernova explosions that drive galactic winds, transporting material out of star-forming regions hence lowering the local star formation rate.  We have investigated the plausibility of this mechanism in comparison to a small scale ($\\sim 100\\,\\pc$) physical approach where star formation is made inefficient by modifying the Schmidt-law star formation normalization. In a very crude way, this mimics unresolved physics such as H$_2$ formation, small scale turbulence and radiative effects. We find that the Schmidt-law efficiency of star formation is far more successful way of regulating star formation towards realistic galaxies than what can be achieved via supernova feedback. Our most successful models reproduce Milky Way galaxies with flat rotation curves, where the small bulge component is formed via secular processes. The main conclusions of this work can be summarized as follows. \\begin{enumerate} \\item Disk characteristics such as $\\Sigma_*(r)$, $\\Sigma_{\\rm gas}(r)$, $v_{\\rm rot}(r)$ and B/D strongly depend on the choice of star formation efficiency per free-fall time, $\\epsilon_{\\rm ff}$. The parameter will essentially set the mode of global star formation, hence governing the final spiral Hubble type, where low efficiencies of $\\epsilon_{\\rm ff}\\sim1$ per cent render discs of Sb or Sbc type, while $\\epsilon_{\\rm ff}=5$ per cent moves the discs closer to Sa/S0 types. Simulations at low efficiencies agree well with observational constraints on disc characteristics \\citep{Courteau97,oleggnedin07}, as well as the angular momentum content of disc galaxies \\citep{navarrosteinmetz00}, the Tully-Fisher relationship \\citep{Pizagno07} and the $\\Sigma_{\\rm SFR}$-$\\Sigma_{\\rm gas}$ relation \\citep{kennicutt98,bigiel2008}. The origin of the successful Milky Way-like galaxy formation is a well motivated suppression of star formation at $z\\sim 3$, the epoch at which the violent assembly process would form a slowly rotating bulge in case of efficient star formation.  \\item Supernova feedback does not regulate star formation efficiently at low input energies. Only when the injected energy per supernova event is five times the canonical value, i.e. $5\\times 10^{51}\\,{\\rm erg}$, do we find lower and more realistic B/D ratios in the simulations tuned to the standard \\cite{kennicutt98} star formation law, leading to a flatter rotational velocity profile, hence resembling the galaxies formed without strong feedback but with a low Schmidt-law efficiency. This comes at the cost of a significantly distorted gas disc at $z=0$, as well as a less bright stellar disc as gas is expelled into the halo, leaving less fuel for star formation at late times. In essence, we find that changes in $\\epsilon_{\\rm ff}$ can play a much greater role in shaping a spiral galaxy than gas redistribution via supernovae-driven winds.  It is plausible that at very high resolution, or using a drastically different recipe of supernovae feedback, lower values of $E_{\\rm SNII}$ may be successful in regulating the SFE. If so, it will still need to mimic the low efficiency on scales of a few $100\\pc$ which, as argued in this work, can be absorbed by the $\\epsilon_{\\rm ff}$-term.  \\item If the star formation efficiency parameter is tuned to match the standard $z=0$ \\emph{K-S} data \\citep{kennicutt98}, i.e. requiring on the order of $\\epsilon_{\\rm ff}\\ge 5\\,{\\rm per cent}$ \\citep[e.g.][]{Stinson06}, star formation is likely to be overestimated at high redshift ($z=3$) where the amplitudes of $\\Sigma_{\\rm SFR}$ are an order of magnitude lower \\citep{WolfeChen06,GnedinKravtsov2010}. All efficiencies studied in this work ($\\epsilon_{\\rm ff}=1-5$ per cent) are compatible with modern data of the THINGS survey \\citep{bigiel2008} but only when $\\epsilon_{\\rm ff}\\sim 1$ per cent can the constraints from $z=3$ data be met and late-type, disc dominated systems form. As the true SFE varies in space and time, being dependent on small scale physics governing H$_2$ formation \\citep[see e.g.][]{Gnedin09}, present day simulations based on single valued efficiency parameter have little predictive power. \\end{enumerate}  We argue \\citep[see also][]{Gnedin09} that the results presented in this paper indicate that other processes in the ISM in addition to, or in conjunction with, supernova feedback are important in explaining the evolution of the galaxy population, as well as regulating observed disc sizes. Some form of outflow process must be responsible for enriching the IGM \\citep{oppenheimerdave06}, which together with an inefficient star formation might explain the faint end of the stellar mass function \\citep{SomervillePrimack99,Keres09}. The same argument can be used for the mass-metallicity relationship \\citep{Brooks07}, although \\cite{Tassis08} demonstrated that it could be reproduced without supernova-driven outflows. Galaxies of masses considered in this work are situated at the knee of the stellar mass function, where the observed and simulated functions \\citep[even without feedback; see][]{Keres09} are in closest agreement. This circumstance might explain why even our simulations without feedback resulted in realistic discs. At this galaxy mass, supernova driven winds cannot escape the deep potential well, and are impeded by the hot halo. On the other hand, AGN feedback, which recently has been introduced into galaxy formation simulations \\citep{DiMatteo05}, is probably not relevant for the Milky Way since the black hole might not be massive enough for efficient AGN radio-heating. At higher masses, and/or at high redshift, the inclusion of AGN is probably necessary to correctly reproduce the observed abundances and stellar masses. This is the greatest uncertainty of our work, which we leave for a future study.  The way in which galaxies populate dark matter haloes is an important topic, see e.g. \\cite{Dutton2010} and references within for a compilation of recent observational data and theoretical work. Recently, \\cite{Guo2010} [see also \\cite{Moster2010} and \\cite{Behroozi2010}] matched dark matter halo mass function from cosmological $N-$body simulations to the stellar mass function of the galaxies from the SDSS \\citep{LiWhite2009}. This analysis yields the required galaxy formation efficiency, $\\eta=(M_{*}/M_{\\rm halo})(\\Omega_{\\rm m}/\\Omega_{\\rm b}$), i.e. what fraction of the universal baryons that must have condensed into stars at a given halo mass. In our \"best-case\" model (n01e1ML, see Table\\,\\ref{table:simsummary1}), the total stellar and dark matter halo virial mass is $\\sim10^{11}\\,\\Msol$ and $\\sim 10^{12}\\,\\Msol$ respectively. This results in a stellar fraction of 10 per cent, which corresponds to almost 60 per cent of the cosmic baryon fraction. The rest of the baryons reside in the stellar halo, gaseous disc and ionized gas halo. At this halo mass, abundance matching requires that the stellar disc accounts for only $\\sim 20$ per cent of the cosmic baryon fraction, i.e. a factor of three lower. Similar discrepancies exist in all modern work of numerical galaxy formation \\citep{Abadi03b,Okamoto05,Governato07,Scannapieco09,Piontek09b}, and its origin is not yet know, although AGN is a compelling mechanism at the high mass end, as discussed above. This issue is the topic of a follow-up paper in preparation.  \\cite{Behroozi2010} performed a comprehensive analysis of abundance matching, accounting for systematic errors in e.g. the stellar mass estimates, the halo mass function, cosmology etc.  Our simulated galaxy formation efficiencies would be in  $\\sim 2\\sigma$ agreement with their result (see their Fig. 11). We note that our own Galaxy and M31 also might be strong outliers in this analysis, considering the inferred $\\eta$ from mass modelling \\citep{Klypin2002,Seigar2008} as well as via recent MW halo mass estimates \\citep{Xue2008}.  Abundance matching is insensitive to the actual Hubble types of the galaxies, forcing all galaxies of a specific mass to be linked to only one halo mass. At  the stellar mass scale of the Milky Way ($5-7\\times 10^{10}\\,\\Msol$), only $\\sim 25$ per cent of galaxies are of Sb/Sbc type \\citep{NairAbraham2010}. It is plausible that the more active merger histories associated with ellipticals and early type disc galaxies have led to a stronger mass expulsion, via e.g. AGN, in comparison to the more disc dominated counterparts. In this scenario, late-type discs are expected to be outliers in the galaxy formation efficiency vs. stellar mass relation, considering the strong bias towards early type systems. A detailed sub-division into Hubble types has not yet been performed when matching galaxies to haloes, although color separations into red and blue systems have been made in studies using weak-lensing \\citep{Mandelbaum2006} and satellite kinematics \\citep{more2010}. These studies indicate a different galaxy formation efficiency for galaxies similar in mass to the Milky Way; a late-type galaxy is associated with a halo of $\\sim 0.5$ dex lower halo mass compared to an equally massive early type \\cite[see e.g. fig. 11 in][]{more2010}.  Understanding, from a numerical perspective, the spread of baryon fractions across dark matter haloes of different masses, accretion histories and environments is a complicated problem, and will require a large sample of high-resolution simulations, which we leave for a future investigation."
    },
    "1004/1004.4626_arXiv.txt": {
        "abstract": "Using estimates of dark halo masses from satellite kinematics, weak gravitational lensing, and halo abundance matching, combined with the Tully-Fisher and Faber-Jackson relations, we derive the mean relation between the optical, $\\Vopt$, and virial, $V_{200}$, circular velocities of early- and late-type galaxies at redshift $z\\simeq 0$.  For late-type galaxies $\\Vopt\\simeq V_{200}$ over the velocity range $\\Vopt = 90-260 \\kms$, and is consistent with $\\Vopt = \\Vmaxh$ (the maximum circular velocity of NFW dark matter haloes in the concordance \\LCDM cosmology). However, for early-type galaxies $\\Vopt \\ne V_{200}$, with the exception of early-type galaxies with $\\Vopt \\simeq 350 \\kms$.  This is inconsistent with early-type galaxies being, in general, globally isothermal.  For low mass ($\\Vopt \\lta 250 \\kms$) early-types $\\Vopt > \\Vmaxh$, indicating that baryons have modified the potential well, while high mass ($\\Vopt \\gta 400 \\kms$) early-types have $\\Vopt < \\Vmaxh$. Folding in measurements of the black hole mass - velocity dispersion relation, our results imply that the supermassive black hole - halo mass relation has a logarithmic slope which varies from $\\simeq 1.4$ at halo masses of $\\simeq 10^{12}h^{-1}\\Msun$ to $\\simeq 0.65$ at halo masses of $10^{13.5}h^{-1}\\Msun$.  The values of $\\Vopt/V_{200}$ we infer for the Milky Way and M31 are lower than the values currently favored by direct observations and dynamical models. This offset is due to the fact that the Milky Way and M31 have higher $\\Vopt$ and lower $V_{200}$ compared to typical late-type galaxies of the same stellar masses. We show that current high resolution cosmological hydrodynamical simulations are unable to form galaxies which simultaneously reproduce both the $\\Vopt/V_{200}$ ratio and the $\\Vopt-\\Mstar$ (Tully-Fisher/Faber-Jackson) relation. ",
        "introduction": "\\label{sec:intro}  It is theoretically expected and observationally established that galaxies are surrounded by extended haloes of dark matter (White \\& Rees 1978; Blumenthal \\etal 1984; van Albada \\& Sancisi 1986; Zaritsky \\& White 1994; Brainerd \\etal 1996; Prada \\etal 2003).  In recent years much progress has been made in understanding the relation between the masses of dark matter haloes and the properties of the galaxies that reside in them.  The majority of work in the literature has focused on the relation between halo mass and galaxy luminosity or stellar mass (e.g.  Yang \\etal 2003; Kravtsov \\etal 2004; van den Bosch \\etal 2004; Hoekstra \\etal 2005; Mandelbaum \\etal 2006; Conroy \\etal 2007; More \\etal 2009).  An alternative approach is to link galaxies to haloes via kinematics, for example, by measuring the relation between the circular velocity ($\\Vcirc(r)= \\sqrt{G M (<r)/r}$, for a spherical system) within the optical part (e.g. the half light radius or 2.2 disk scale lengths) of galaxies, $\\Vopt$, to the circular velocity at the virial radius of the dark matter halo, $V_{200}$.  This approach has the advantage that it is free from uncertainties in luminosities or stellar masses due to uncertainties in extinction from dust, stellar populations or the stellar initial mass function (IMF). In addition, because it is a {\\it dynamical} link, it gives a direct measurement of the slope of the global mass density profile within the virial radius. For example, for a singular isothermal density profile, $\\rho(r) \\propto r^{-2}$, and thus $V_{\\rm c} =const.$ which implies $\\Vopt/V_{200}=1$.  By measuring a departure from this prediction, one can rule out galaxies being globally isothermal. This is a relevant issue because it is common practice to infer halo masses from observed $\\Vopt$ by assuming $\\Vopt=V_{200}$, even though this assumption lacks both theoretical and observational support. For example, in \\LCDM cosmologies dark matter haloes are not isothermal, so a detection of non-isothermality would be useful from a \\LCDM haloes perspective.  The $\\Vopt/V_{200}$ ratio also contains information on the relative importance of baryons vs dark matter in the optical regions of galaxies. For galaxy mass \\LCDM haloes, the ratio between the maximum circular velocity of the dark matter halo, $\\Vmaxh$, and the virial velocity of the dark matter halo $V_{\\rm max,h}/V_{200}\\simeq 1.1-1.2$ (Bullock \\etal 2001), and thus if $\\Vopt/V_{200}$ is observed to be significantly larger than this, it indicates that baryons have modified the potential well.  Using weak lensing measurements of halo masses combined with the Tully-Fisher (1977, hereafter TF) relation, Seljak (2002) inferred $\\Vopt/V_{200}=1.8$ (with a 95\\% C.I. lower limit of 1.4), for late-type $L_*$ galaxies. Such high values of $\\Vopt/V_{200}$ are naturally explained by the combined effects of baryons adding to the optical circular velocity directly, and indirectly by inducing contraction of the dark matter halo (Blumenthal \\etal 1986).  However, a high value of $\\Vopt/V_{200}$ introduces problems in reconciling the TF relation with the halo mass function and galaxy luminosity function, LF. Semi-analytic galaxy formation models that are able to simultaneously reproduce the zero point of the TF relation and the galaxy LF assume that $\\Vopt = V_{200}$ (e.g.  Somerville \\& Primack 1999), or $\\Vopt=V_{\\rm max,h}$ (e.g.  Croton \\etal 2006). However, models that take into account the contribution of the baryons to the rotation curve and the effect of halo contraction have been unable to match both the LF and TF relation (e.g. Benson \\etal 2003; Benson \\& Bower 2010).  Dutton \\etal (2007) showed that the $\\Vopt/V_{200}$ ratio can place strong constraints on dark halo structure. In particular they argued that models with adiabatic contraction, standard $\\Lambda$CDM halo concentrations, and standard stellar IMFs are unable to {\\it simultaneously} reproduce the zero points of the TF and size-luminosity relations, and the requirement that $\\Vopt\\simeq \\Vmaxh$ (in order to reproduce the LF).  However, if halo contraction was somehow avoided, or even if haloes could {\\it expand} in response to galaxy formation, then models could be constructed that reproduced all of the observational constraints.  In this paper we derive the relation between $\\Vopt$ and $V_{200}$ for both early- (red, bulge dominated) and late-type (blue disk dominated) galaxies at redshifts $z\\simeq 0$. For late-types, we use $\\Vopt=V_{2.2}$, where $V_{2.2}$ is the rotation velocity at 2.2 disk scale lengths. For early-types we use $\\Vopt=1.65\\sigma(R_{50})$, where $\\sigma(R_{50})$ is the velocity dispersion within the projected half light radius of the galaxy, and the factor of 1.65 is from Padmanabhan \\etal (2004, see also \\S\\ref{sec:fj}). We determine the mean relation between virial velocity and stellar mass using published measurements of halo masses from satellite kinematics, weak lensing, and halo abundance matching studies. By comparing these relations with the observed TF and Faber Jackson (1976, hereafter FJ) relations we derive the mean relation between $\\Vopt$ and $V_{200}$. In a future paper (Dutton \\etal in prep) we will combine these results with other scaling relations of early- and late-type galaxies to place constraints on the structure of their dark matter haloes.  This paper is organized as follows.  In \\S2 we describe the observations of halo masses and galaxy scaling relations.  In \\S3 we derive and discuss the relation between $\\Vopt$ and $V_{200}$.  In \\S4 we discuss the implications of our results for the slope of the relation between super-massive black hole mass and halo mass.  In \\S5 we compare our results with predictions for $\\Lambda$CDM dark matter haloes. In \\S6 we compare our results with halo masses derived for the Milky Way and M31 galaxies, as well as galaxies formed in cosmological hydrodynamical simulations. In \\S7 we give a summary.  Unless otherwise specified, throughout this paper we adopt a Hubble parameter, $H_0=100\\, h \\,\\rm km \\,s^{-1}\\, Mpc^{-1}$, i.e. stellar masses are expressed in $h^{-2}\\Msun$ units, halo masses are expressed in $h^{-1}\\Msun$ units, and halo sizes are expressed in $h^{-1} \\kpc$ units.  \\begin{table*} \\caption{Cosmological and stellar mass parameters adopted by studies of the halo masses - stellar mass relation. Notes--- Col. (2) Method used to derive the halo mass - stellar mass relation: GC--Group Catalog; AM--Abundance Matching; WL--Weak Gravitational Lensing; SK--Satellite Kinematics. Col. (3) Mean density of dark matter haloes in units of the critical density of the universe ($\\Delta_{\\rm vir}=\\bar{\\rho}/\\rho_{\\rm crit}$). Col. (4) Redshift zero cosmological matter density. Col. (5) Redshift zero cosmological dark energy density. Col. (6) Redshift zero Hubble parameter in units of $\\rm km\\,s^{-1}\\, Mpc^{-2}$. Col. (7) Reference for stellar masses and stellar initial mass function: B03--Bell \\etal (2003); BR07--Blanton \\& Roweis (2007); G06--Gallazzi \\etal (2006); K03--Kauffmann \\etal (2003); P07--Panter \\etal (2007); diet-Sal.-- Salpeter (1955) $-0.15$dex; Kroupa--Kroupa (2001); Chab.--Chabrier (2003); Col (8) Offset we apply to stellar mass, in dex. } \\begin{tabular}{lcccccccc} \\hline \\hline Reference  &  Method    &  $\\Delta_{\\rm vir}$ & $\\Omega_{\\rm m}$ & $\\Omega_{\\Lambda}$ & $H_0$ & Stellar Masses & Mass Offset & Morphology\\\\ (1) &(2) & (3) & (4) & (5) & (6) & (7) & (8) & (9)\\\\ \\hline Yang \\etal 2007       & GC & \\pt71.4\\pt& 0.238   & 0.762   & 73 & B03, diet-Sal. & $-0.1$ & All\\\\ Moster \\etal 2010     & AM & \\pt95.62  & 0.26\\po & 0.74\\po & 72 & P07, Chab. & $\\phantom{+}0.0$ & All\\\\ Guo \\etal 2010        & AM & 200.0\\pt& 0.25\\po & 0.75\\po & 73 & BR07, Chab. & $+0.1$ & All\\\\ Behroozi \\etal 2010   & AM & \\pt91.8\\pt& 0.27\\po & 0.73\\po & 70 & BR07, Chab. & $+0.1$ & All\\\\ Mandelbaum \\etal 2006 & WL & \\pt54.0\\pt& 0.30\\po & 0.70\\po & 70 & K03, Kroupa & $\\phantom{+}0.0$ & Elliptical/Disk\\\\ Mandelbaum \\etal 2008 & WL & \\pt54.0\\pt& 0.30\\po & 0.70\\po & 70 & G06, Chab. & $\\phantom{+}0.0$ & Elliptical \\\\ Schulz \\etal 2010     & WL & \\pt60.0\\pt& 0.30\\po & 0.70\\po & 70 & K03, Kroupa & $\\phantom{+}0.0$ & Elliptical\\\\ Conroy \\etal 2007     & SK & 200.0& 0.30\\po & 0.70\\po & 100h& B03, diet-Sal. & $-0.1$ & All/Red/Blue\\\\ More \\etal 2010       & SK & \\pt47.6\\pt& 0.238   & 0.762   & 73 & B03, diet-Sal. & $-0.1$ & All/Red/Blue\\\\ Klypin \\etal 2010     & SK & \\pt91.8\\pt& 0.27\\po & 0.73\\po & 70 & BR07, Chab. & $+0.1$  & All\\\\ \\hline \\hline \\end{tabular} \\\\ \\label{tab:mm} \\end{table*} ",
        "conclusions": "\\label{sec:sum} We combine measurements of halo virial masses from weak lensing, satellite kinematics, and halo abundance matching with the Tully-Fisher (1977) and Faber-Jackson (1976) relations to place constraints on the {\\it average} relation between the optical ($\\Vopt$) and virial ($V_{200}$) circular velocities of early- ($\\Vopt\\equiv V_{50}=1.65 \\sigma(R_{50})$) and late-type ($\\Vopt\\equiv V_{2.2} = V_{\\rm rot}(2.2 R_{\\rm d})$) galaxies at redshift $z\\simeq 0$.  We summarize our results as follows:  \\begin{itemize}  \\item The stellar mass to halo virial mass fractions of late-type (blue/disk dominated) galaxies increase from $\\simeq 0.015$ at a stellar mass of $\\Mstar = 10^{9.4} h^{-2}\\Msun$, to $\\simeq 0.043$ at a stellar mass of $\\Mstar \\simeq 10^{11.0} h^{-2} \\Msun$ (assuming $h=0.7$). These correspond to integrated star formation efficiencies, $\\epsilon_{\\rm SF} = \\Mstar/(f_{\\rm bar}M_{200})$, of $\\epsilon_{\\rm SF}\\simeq 9\\%$ and $\\epsilon_{\\rm SF}\\simeq 26\\%$, respectively (assuming a cosmic baryon fraction, $f_{\\rm bar}=0.165$). After accounting for cold gas, the galaxy formation efficiencies, $\\epsilon_{\\rm GF} = (\\Mstar+M_{\\rm gas})/(f_{\\rm bar} M_{200})$, are $\\epsilon_{\\rm GF}\\simeq 17\\%$ and $\\epsilon_{\\rm GF}\\simeq 33\\%$, respectively.  \\item The central galaxy formation efficiencies of early-type (red/bulge dominated) galaxies reach a peak of $\\simeq 12\\%$ at $\\Mstar=10^{10.5} h^{-2} \\Msun$.  The efficiency drops to $\\epsilon_{\\rm GF}\\simeq 2.8\\%$ at $\\Mstar=10^{11.4} h^{-2}\\Msun$.  \\item For late-type galaxies $V_{2.2}/V_{200}$ is roughly constant, and close to unity, over the range of stellar masses or rotation velocities that we probe: ($ 10^{9.35} \\lta \\Mstar \\lta 10^{11} h^{-2} \\Msun$, or $ 90 \\lta V_{2.2} \\lta 260 \\kms$).  \\item For late-type galaxies, the maximum circular velocity of $\\Lambda$CDM haloes, $\\Vmaxh$, is consistent with being equal to the maximum observed rotation velocity, $V_{\\rm max}$. This result is not reproduced by current cosmological simulations of disk galaxy formation (e.g. Governato \\etal 2007; Piontek \\& Steinmetz 2009).  \\item Cosmological simulations have been unable to form galaxies that fall on both the $V_{200}-\\Mstar$ and $\\Vopt-\\Mstar$ relations. These relations provide a strong and simple test for galaxy formation models.  \\item For early-type galaxies $V_{50}/V_{200}$ is not a constant, which rules out most early-type galaxies from being globally isothermal (i.e. constant circular velocity from the optical half light radius to the virial radius of the halo).  \\item Early-type galaxies have $V_{50}/V_{200} \\simeq 1$ only for $\\Vopt\\simeq350\\kms$ (i.e. $\\sigma \\simeq 215 \\kms$). This is consistent with the strong plus weak lensing result of Gavazzi \\etal (2007) which inferred that massive early-types with $\\sigma\\simeq 240\\kms$ have close to isothermal global density profiles.  \\item There is some evidence that for early-type galaxies in lower mass haloes ($V_{200} \\lta 250\\kms$) the $V_{50}/V_{200}$ ratio is higher than the $\\Vmaxh/V_{200}$ ratio, which would indicate that baryons must have modified the potential well, either through their own gravity, or by adiabatic contraction of the halo.  \\item For early-type galaxies in higher mass haloes ($V_{200} \\gta 500\\kms$), $V_{50}/V_{200} < 1 $, indicating that the half light radii of the galaxies are much smaller than the NFW scale radius.  \\item The mass dependence of the $V_{50}/V_{200}$ ratio for early-types implies that the slope of the black hole mass - halo mass relation varies with mass.  Specifically it varies from $\\simeq 0.65$ at high halo masses ($\\simeq 10^{13.5}\\Msun$), to $\\simeq 1.4$ at low halo masses ($\\simeq 10^{12}\\Msun$). The shallow slope at high mass is at odds with that predicted by the self-regulating feedback model of Wyithe \\& Loeb (2003) and the cosmological simulation of Booth \\& Schaye (2010) which both predict a slope of $\\simeq 1.5$.  \\item For the Milky Way (MW) ($\\Vopt=220 \\kms$) and Andromeda (M31) ($\\Vopt=260 \\kms$) galaxies our results predict that \\begin{equation} \\Vopt/V_{200} = 1.11\\, (+0.22,-0.20, 2\\sigma) \\end{equation} This value is significantly lower than those suggested by the dynamical models for the MW and M31 of Klypin, Zhao \\& Somerville (2002) which correspond to $\\Vopt/V_{200}=1.66$ for MW and $\\Vopt/V_{200}=1.62$ for M31.  Other measurements of the virial mass of the MW using kinematics of halo stars also infer relatively high values: $\\Vopt/V_{200}= 1.48^{+0.25}_{-0.26}$ (90\\% C.I.)  (Smith \\etal 2007); $\\Vopt/V_{200}=1.67^{+0.31}_{-0.24}$ ($2\\sigma$) (Xue \\etal 2008).  If both our relation and the dynamical models are correct this discrepancy would imply that the MW and M31 do not live in typical dark matter haloes for their optical rotation velocity. \\end{itemize}"
    },
    "1004/1004.3971_arXiv.txt": {
        "abstract": "The uncertainty in the absolute value of the fluorescence yield is still one of the main contributions to the total error in the reconstruction of the primary energy of ultra-energetic air showers using the fluorescence technique. A significant number of experimental values of the fluorescence yield have been published in the last years, however reported results are given very often in different units (photons/MeV or photons/m) and for different wavelength intervals. In this work we present a comparison of available results normalized to its value in photons/MeV for the 337~nm band at 800~hPa and 293~K. The conversion of photons/m to photons/MeV requires an accurate determination of the energy deposited by the electrons in the field of view of the experimental setup. We have calculated the energy deposition for each experiment by means of a detailed Monte Carlo simulation including when possible the geometrical details of the particular setup. Our predictions on deposited energy, as well as on some geometrical factors, have been compared with those reported by the authors of the corresponding experiments and possible corrections to the fluorescence yields are proposed. ",
        "introduction": "\\label{sec:intro}  Charged particles of a cosmic ray shower, mostly electrons, passing through the atmosphere lose energy by inelastic collisions with air molecules. A small fraction of the total deposited energy is emitted by molecular nitrogen as UV fluorescence radiation in the $\\sim 300 - 400$~nm spectral range. The detection of this radiation provides a precise determination of the longitudinal profile allowing the reconstruction of the cosmic ray properties, in particular its energy. This technique~\\cite{Bunner_PhD} was first successfully used by the Fly's Eye experiment~\\cite{FlysEye} and later by HiRes~\\cite{HiRes}. Fluorescence telescopes are presently being employed by the Pierre Auger Observatory~\\cite{Auger} and the Telescope Array experiment~\\cite{TelescopeArray}. On the other hand, the JEM-EUSO project~\\cite{JEM_EUSO} and the S-EUSO free-flying satellite mission~\\cite{S_EUSO} are being designed for the detection of fluorescence traces of air showers viewing downward from the top of the atmosphere.  The ratio between number of fluorescence photons and deposited energy, i.e., the fluorescence yield $Y$, is a key calibration parameter determining the energy scale of the fluorescence telescopes. A number of absolute measurements of the fluorescence yield have been carried out in laboratory experiments in the last years.  At atmospheric pressure, fluorescence emission in the spectral range of interest basically comes from the Second Positive (2P) system of N$_2$ and, to a much less extent, the First Negative (1N) system of N$^+_2$. Each one results from the transition between two electronic states and consists of a set of molecular bands corresponding to different combinations of vibrational levels $v$-$v'$ of the respective upper and lower states. Excitation cross section of the 1N system decreases with energy in the $10^2-10^6$~eV range showing a very smooth growing behavior at larger energies. That of the 2P system peaks at about 15~eV decreasing strongly with an $E^{-2}$ dependence. As a consequence, high-energy electrons themselves are very inefficient for the generation of air fluorescence. In fact, as is well known, fluorescence emission along the track of an energetic primary electron is mainly induced by low-energy secondary electrons ejected in successive ionization processes. The efficiency for fluorescence emission varies with the energy spectrum of these secondaries\\footnote{For instance, electrons with energies below the threshold for fluorescence excitation promptly lose all their energy, but induce no fluorescence.}, which in principle depends on the primary energy. However, a theoretical calculation using a detailed Monte Carlo simulation~\\cite{Blanco,Arqueros_Astropart_Phys,Arqueros_NJP} has demonstrated that the fluorescence yield is nearly constant for energies above a few kiloelectronvolts. This $E$-independent behavior of the fluorescence yield has been checked experimentally~\\cite{FLASH_thick,MACFLY,5th_FW_AIRFLY_E}.  A number of experimental values of the absolute fluorescence yield are available in the literature. Essentially, in these experiments a beam of electrons crosses a collision chamber filled with air at known conditions. However different techniques can be used. Many authors~\\cite{MACFLY,Kakimoto,Nagano_03,Nagano_04,Lefeuvre,AirLight} have used electrons from a source of $^{90}$Sr with energy around 1~MeV. Other absolute measurements have been performed with higher energy electrons from accelerators. The MACFLY Collaboration~\\cite{MACFLY} used the CERN/SPS-X5 test beam facility which delivers a pulsed electron beam of about 10$^4$ electrons per spill (4.8~s duration) every 16.8~s. Measurements at 20 and 50~GeV were carried out using this facility. The FLASH Collaboration~\\cite{FLASH_06,FLASH_08} used the Final Focus Test Beam facility at the Stanford Linear Accelerator Center which provided 28.5~GeV electrons in 3~ps pulses of about 10$^8$ electrons at a rate of 10~Hz. The AIRFLY Collaboration is carrying out absolute measurements of the air-fluorescence yield using different accelerators. Preliminary results obtained with 350~MeV electrons at the Beam Test Facility of the Instituto Nazionale di Fisica Nucleare have been presented in~\\cite{5th_FW_AIRFLY}. New measurements using a 120~GeV proton beam from the Meson Beam Test Facility at Fermilab were presented at the 6th AFW~\\cite{6th_FW} and very likely final results will be published soon.  In this work we will compare the absolute values reported by the above authors. In many occasions comparison cannot be done directly. For instance, some authors measure single intense fluorescence bands while others detect the integrated fluorescence in a wide spectral range. Comparison between these experimental results can be carried out if the relative intensities of the bands along the fluorescence spectrum are known. We will normalize all the results to the most intense fluorescence band at 337~nm using accurate experimental data in full agreement with theoretical intensity ratios proposed in~\\cite{Blanco,Arqueros_Astropart_Phys,Arqueros_NJP}.  On the other hand, some authors report their measurements in units of photons per electron and meter while others give directly the $Y$ parameter in photons per unit deposited energy (e.g., photons/MeV). The conversion factor between both magnitudes is the energy deposited per electron and unit path length. As it has been pointed out~\\cite{Blanco,Arqueros_Astropart_Phys,Arqueros_NJP}, an accurate determination of the deposited energy requires discounting that corresponding to secondary electrons outside the field of view of the optical system used to detect the emitted fluorescence. To evaluate this effect we have carried out a MC simulation of the various experimental setups. When possible a dedicated simulation including geometrical details has been performed. Also the effect of the spatial distribution of deposited energy on the optical efficiency of the experimental setup has been evaluated in some cases. As a result of this study absolute values of the fluorescence yield normalized to that of the 337~nm band at 293~K and 800~hPa in dry air will be shown with a discussion on possible corrections to be applied. A previous comparison at 1013~hPa for some of the above measurements has been published recently~\\cite{Arqueros_NJP}. In this work we will update those results and extend the comparison to other experimental results. ",
        "conclusions": "\\label{sec:conclusions}  A direct comparison of the absolute measurements of the air fluorescence yield available in the literature is not possible since very often the reported data correspond to different wavelength intervals and the magnitude is expressed in different units. The conversion factors for such comparison depend on the particular features of the experimental setup. In this work we have described a procedure for the normalization of the various experimental results to a common unit and wavelength interval.  In the first place, fluorescence yields have been reduced to that corresponding to the most intense band at 337~nm using accurate experimental data on intensity ratios in full agreement with theoretical relationships. Absolute values can be expressed in photons/MeV units as far as the energy deposited in the field of view of the setup is known. In this work we have used an upgraded simulation algorithm to calculate the energy deposition of electrons in a wide energy range. A generic algorithm has provided results as a function of the pressure and the overall size of the collision chamber. In addition, for several experiments, we have performed a detailed simulation (in full agreement with the predictions of the generic one) including in some cases the calculation of geometrical factors. The results of these simulations have been compared with those reported by the authors of the corresponding experiments and eventual disagreements have been discussed.  As a result of this work a table of the absolute fluorescence yields normalized to 800~hPa and 293~K have been presented. According to our simulation, a non-negligible correction has to be applied to those results obtained under the assumption that the energy lost by the primary electrons is fully deposited in the field of view of the optical system (i.e., Kakimoto~\\etal~\\cite{Kakimoto}, Nagano~\\etal~\\cite{Nagano_03,Nagano_04}) and Lefeuvre~\\etal~\\cite{Lefeuvre}). In particular the Nagano's fluorescence yield should be increased by about 7\\%. Possible implications of these corrections to the $Y$ values in the determination of the end of the energy spectrum of cosmic rays have been discussed recently in~\\cite{Nagano_NJP}.  On the other hand, our calculations of energy deposition are in general in reasonable agreement with those from experiments that carried out a detailed simulation of their setup. Even though, in some cases we have found non-negligible discrepancies and correction factors have been proposed. Experiments claiming high accuracy should be very careful in the evaluation of the error contribution from the deposited energy.  From the above table, and assuming the corrections proposed in this work, most fluorescence yield measurements at those reference conditions are contained within the 6~-~7~ph/MeV interval. Notice, however that the fluorescence yield reported by MACFLY, as well the preliminary result of AIRFLY, are smaller by more than 10\\%."
    },
    "1004/1004.2092_arXiv.txt": {
        "abstract": "So far, there have been no theories or observational data that deny the presence of interaction between dark energy and dark matter. We extend naturally the holographic dark energy (HDE) model, proposed by Granda and Oliveros, in which the dark energy density includes not only the square of the Hubble scale, but also the time derivative of the Hubble scale to the case with interaction and the analytic forms for the cosmic parameters are obtained under the specific boundary conditions. The various behaviors concerning the cosmic expansion depend on the introduced numerical parameters which are also constrained. The more general interacting model inherits the features of the previous ones of HDE, keeping the consistency of the theory. ",
        "introduction": "The observation of the type Ia Supernovae \\cite{ref1} has shown that our universe is undergoing an era of accelerated expansion. Data from other observations such as Cosmic Microwave Background Radiation \\cite{ref2} and SDSS \\cite{ref3} also give support to the phenomenon. These data reveal that the present universe is dominated by 70\\% exotic component, dubbed dark energy, which has negative pressure and is pushing the universe into accelerated expansion. Therefore, in order to explain this bizarre phenomenon, various models of dark energy have been put forward, ranging from the simplest one -- a cosmological constant -- to the scalar field theories of dark energy and modified gravity theories as well. The cosmological constant corresponds to the vacuum energy with constant energy and pressure, and an equation of state $w=-1$. However, it is confronted with two fundamental problems: the fine-tuning problem and the cosmic coincidence problem. To alleviate the drawbacks, various scalar field theories of dark energy emerge continually \\cite{ref4}. On the other hand, the new geometric structures of space-time are also taken to realize the accelerated expansion of the universe \\cite{ref5}.  Recently, another model inspired by the holographic principle has been put forward to explain the dark energy \\cite{ref6,ref7}. The holographic principle, enlightened by the quantum properties of black holes, is one of the most important results in recent investigations of quantum gravity \\cite{ref8,ref9}. It was first put forward by 't Hooft in the context of black hole physics \\cite{ref10} and later extended to string theory by Susskind \\cite{ref11}. According to the holographic principle, the entropy of a system scales not with its volume, but with its surface area \\cite{ref12,ref13}. Equally the degrees of freedom of a piece of spatial region reside not in the bulk but only at the boundary of the region and the number of degrees of freedom per Planck area is not greater than unity. From the work of Cohen and collaborators \\cite{ref14}, it is proposed that in quantum field theory a short distance cut-off is related to a long distance cut-off in virtue of the limit set by formation of a black hole. In other words, when $\\rho_{de}$ is taken as the quantum zero-point energy density caused by a short distance cut-off, the total energy in the region of size $L$ is not more than the mass of a black hole of the same size, where the inequality $L^3\\rho_{de} \\leqslant LM_p^2$ emerges. Saturating it is equivalent to taking the largest $L$ allowed, thus \\begin{equation}\\label{eq1} \\rho_{de}=3c^2M_p^2L^{-2}, \\end{equation} where $3c^2$ is a numerical constant introduced for convenience and $M_p$ is the reduced Planck mass, $M_p^2=(8\\pi G)^{-1}$.  In the context of cosmology, the existence of an unknown vacuum energy is suggested and when $L$ is taken as the size of the current universe, for which the most common choice is the Hubble scale i.e., $L \\sim H^{-1}$, the vacuum energy density is proportional to the square of the Hubble scale $\\rho_{de} \\propto H^2$, according to the holographic principle \\cite{ref14,ref15}. (In this Letter we use terms like the vacuum energy and dark energy interchangeably.) Through the Friedmann equation \\begin{equation}\\label{eq2} 3M_p^2 H^2=\\rho_{de}+\\rho_m, \\end{equation} it turns out that the dark matter density $\\rho_m$ has the same behavior as $\\rho_{de}$ with the expression $\\rho_m=3(1-c^2)M_p^2 H^2$. That is to say, the dark energy is also pressureless, namely, $p_{de}=w\\rho_{de}=0$, meaning the equation of state parameter $w=0$. Obviously the case violates the condition of the cosmic accelerated expansion $w<-1/3$, although having solved the fine-tuning problem in principle \\cite{ref16}. A better choice for the infrared cut-off $L$ is the particle horizon \\cite{ref9,ref13}, which is defined as $R_{ph}=a\\int_0^t\\frac{dt'}{a}=a\\int_0^a\\frac{da'}{Ha'^2}$. But as is shown in Li's work \\cite{ref6}, the equation of state of dark energy in this situation is larger than $-1/3$, out of the mechanism of cosmic acceleration likewise. Furthermore, Li takes the future event horizon, whose definition is $R_{eh}=a\\int_t^\\infty\\frac{dt'}{a}=a\\int_a^\\infty\\frac{da'}{Ha'^2}$, as the infrared cut-off instead, for the sake of ameliorating the holographic model. The aspect of amelioration is the satisfaction of acceleration and that $c=1$ makes the universe evolve into the de Sitter space-time while $c<1$ the phantom region \\cite{ref17}. The concomitant defect is encounter of the causality problem pointed out by Cai \\cite{ref18}. It is argued that according to the definition of the future event horizon, the history of dark energy depends on the future evolution of the scale factor $a(t)$, as violates causality. Moreover, for a spatially flat Friedmann-Robertson-Walker universe, only the cosmic expansion is accelerating can the future event horizon exist. So it is contrived that in order to interpret the cosmic accelerated expansion, the holographic dark energy model based on the future event horizon has presumed the acceleration. To avoid these problems originated by the introduced future event horizon, Cai also proposed a dark energy model, dubbed ``agegraphic dark energy'', characterized by the age of the universe, which was chosen as the length measure instead of the horizon distance of the universe. Correlative studies of agegraphic dark energy can be perused in \\cite{ref19,ref20} and references therein. We give no unnecessary details here.  Recently, inspired by the holographic dark energy model, Gao et al. \\cite{ref21} put forward the so-called Ricci dark energy (RDE) model. The idea is taking the average radius of Ricci scalar curvature $|R|^{-1/2}$ as the infrared cut-off. So from Eq. (\\ref{eq1}) the dark energy density is proportional to the Ricci scalar curvature \\begin{equation} \\rho_{de}=3\\alpha M_p^2\\left(\\dot{H}+2H^2\\right)=-\\frac{\\alpha}{2}M_p^2 R, \\end{equation} where $\\alpha$ is a constant to be determined and a spatially flat universe is presumed. With some values of cosmological parameters which are consistent with current observations \\cite{ref22,ref23}, the equation of state of Ricci dark energy can evolve across the cosmological constant boundary $w=-1$ \\cite{ref24}. It means this model, which differs slightly for different values of cosmological parameters, can be classified as a quintom one \\cite{ref25}. The Ricci dark energy model set up without the future event horizon is naturally free of the causality problem. In this case the dark energy is determined by the local Ricci scalar curvature rather than the event horizon of a global concept. Also, the fine-tuning problem is avoided because the dark energy is associated with the space-time scalar curvature, but not with Planck or other high energy physical scales and the coincidence problem is solved as well. In \\cite{ref26}, Granda and Oliveros extended the Ricci dark energy to a more general form in which the energy density is \\begin{equation}\\label{eq4} \\rho_{de}=3M_p^2(\\alpha H^2+\\beta\\dot{H}), \\end{equation} where $\\alpha$ and $\\beta$ are constants to be determined. So it can be realized that when $\\alpha=2\\beta$ the model reduces to Ricci dark energy. In the same way the model is phenomenologically viable, fitting with the current observational data, as well the causality and coincidence problems are solved. Correlative works of this kind are shown in \\cite{ref27} from various perspectives of scalar field, spatial curvature, braneworld cosmology, observational data and so forth.  On the other hand, in the forementioned models, dark energy and dark matter evolve separately, keeping to different conservation equations of energy, with the standard evolution of dark matter $\\rho_m \\propto a^{-3}$. But so far, neither theories nor observational data have denied the interaction between them. For the sake of generality, the interaction term is naturally considered. In the context of dark energy, the interaction has been introduced to study some issues, e.g., raising accelerated expansion in an interacting dark energy model in which the Hubble scale is treated as the infrared cut-off \\cite{ref28}, avoiding the big rip singularity \\cite{ref29} and so forth \\cite{ref30}.  Benefiting from investigations done already, we extend naturally the more general holographic Ricci dark energy model shown in \\cite{ref26} to the form with interaction between the two major cosmic components, expecting to enrich the theoretical studies of series of holographic dark energy models. In the next section we expatiate upon the interacting model and obtain analytic expressions for the cosmic evolution. With the help of Suwa's constraints on the interacting Ricci dark energy (IRDE) model \\cite{ref31}, we also constrain parameters in our model by analogism. The last section is for discussion. ",
        "conclusions": "Hereinafter, we review the consistency of formulae of the holographic dark energy density in various models. First of all, from the definition (\\ref{eq9}) the dark energy density can be expressed generally $\\rho_{de}=3M_p^2 \\Omega_{de}H^2$. As is seen, if $\\Omega_{de}$ is a constant, the energy density will correspond to the holographic model with the Hubble scale as the infrared cut-off, namely, $\\rho_{de} \\propto H^2$. It is understood easily. For $\\rho_{de}=3c^2M_p^2 H^2$, \\begin{equation} r\\equiv\\frac{\\rho_m}{\\rho_{de}}=\\frac{1-c^2}{c^2}=\\mbox{const.}, \\end{equation} which makes $\\Omega_{de}=\\frac{1}{1+r}=c^2$, so $\\Omega_{de}$ must be a constant in this situation. Whereas, a variable $\\Omega_{de}$, which is time dependent, will lead to other holographic dark energy models ever put forth. Concretely speaking, for the models with the particle horizon and the future event horizon, $L=R_h$ is presumed, where $R_h$ represents $R_{ph}$ or $R_{eh}$, \\begin{equation} \\Omega_{de}(t)=\\frac{\\rho_{de}}{3M_p^2 H^2}=\\frac{3c^2M_p^2 R_h^{-2}}{3M_p^2 H^2}=\\frac{c^2}{H^2 R_h^2}, \\end{equation} so $\\Omega_{de}$ is time dependent. Also in the general model (\\ref{eq4}), into which Eq. (\\ref{eq8}) is  substituted for $\\dot{H}$, the dark energy density reads \\begin{equation}\\label{eq24} \\rho_{de}=3M_p^2\\left[\\alpha-\\frac{3\\beta}{2}\\left(1+\\frac{w}{1+r}\\right)\\right]H^2, \\end{equation} where $r$ is in a state of dependence on time, and so is $w$. It means that terms in the square bracket of Eq. (\\ref{eq24}) is a function of time, representing $\\Omega_{de}(t)$.  Furthermore, we could see it from a different point of view. In \\cite{ref28}, it is interestingly argued that for the Hubble scale case with interaction considered, if $c^2$ is time dependent, i.e., $c^2=c^2(t)$, the equation of state of dark energy will be \\begin{equation} w=-\\left(1+\\frac{1}{r}\\right)\\left[\\frac{\\Gamma}{3H}+\\frac{(c^2)^\\cdot}{3Hc^2}\\right]. \\end{equation} Thus both a varying $c^2$, which leads to a varying $r$, further a varying $\\Omega_{de}$, and interaction can make $w$ more negative up to the cosmic accelerated expansion. In other words, the infrared cut-off may well remain $L=H^{-1}$, not necessarily changed. And the variation of $c^2(t)$ in $\\rho_{de}=3c^2(t)M_p^2 H^2$ denotes the change of the degree of saturation of the holographic bound $\\rho_{de} \\leqslant M_p^2 L^{-2}$. It is equivalent to the above case $\\rho_{de}=3M_p^2 \\Omega_{de}(t)H^2$.  So far as it goes, these holographic dark energy models seem harmonious. However, as a matter of fact, although the infrared cut-off $L$ has been changed continually to solve problems such as fine-tuning, cosmic coincidence, accelerated expansion and causality step by step, the underlying physical mechanism of HDE remains vague yet. Therefore, we can but say that the series of HDE models are provided with consistency in sense of formalism and the internal relationship among them is claimed.  On the other hand, from analysis of curves above compared with that in \\cite{ref35}, note that the increasing $b^2$ leads the range of $w$ and $\\Omega_{de}$ to narrow down by reason of decay of the dark energy into dark matter, we have introduced up front. Also the larger $b^2$ is, the lower $w$ becomes in the early universe of high redshift, which means the behavior of dark energy is more different from that of dark matter, and the earlier the cosmic accelerated expansion occurs. But for smaller $b^2$, the cosmic expansion will gain larger acceleration in the far future. The conclusion here is coincident with that from Wang et al. \\cite{ref35}.  In essence the dark energy problem should be an issue of quantum gravity, nevertheless there is no mature theory of quantum gravity with so many things unknown and uncertain at present. As regards the reason why the holographic dark energy density form containing a term $\\dot{H}$ is motivated, Gao et al. bring forth their viewpoint from two perspectives of construction of a $K$-essence scalar field model and quantum fluctuation \\cite{ref21}. Although the more general interacting holographic dark energy model acts phenomenologically viable, we sometimes avoid arguing round and round the subject phenomenologically and look forward to more penetrating insight into the holographic theory of dark energy."
    },
    "1004/1004.3094_arXiv.txt": {
        "abstract": "Lensing flux-ratio anomalies have been frequently observed and taken as evidence for the presence of abundant dark matter substructures in lensing galaxies, as predicted by the cold dark matter (CDM) model of cosmogony. In previous work, we examined the cusp-caustic relations of the multiple images of background quasars lensed by galaxy-scale dark matter haloes, using a suite of high-resolution $N$-body simulations (the Aquarius simulations). In this work, we extend our previous calculations to incorporate both the baryonic and diffuse dark components in lensing haloes. We include in each lensing simulation: (1) a satellite galaxy population derived from a semi-analytic model applied to the Aquarius haloes, (2) an empirical Milky-Way globular cluster population and (3) satellite streams (diffuse dark component) identified in the simulations. Accounting for these extra components, we confirm our earlier conclusion that the abundance of intrinsic substructures (dark or bright, bound or diffuse) is not sufficient to explain the observed frequency of cusp-caustic violations in the CLASS survey. We conclude that the observed effect could be the result of the small number statistics of CLASS, or intergalactic haloes along the line of sight acting as additional sources of lensing flux anomalies. Another possibility is that this discrepancy signals a failure of the CDM model. ",
        "introduction": "The cold dark matter (CDM) cosmogony predicts that cosmic structures form hierarchically through a succession of mergers and accretions. The mass of the Milky Way's own dark halo is currently constrained to $\\sim1-5\\times10^{12}\\,M_{\\sun}$ (e.g. \\citealt{Guo09arxiv}). Numerical simulations predict a wealth of dark matter substructures (self-bound subhaloes) surviving in haloes of this mass (e.g. \\citealt{Gao2004a}; \\citealt{Gao2004b, DKM08Nature, volker08Aq}). These substructures have a power-law mass function; scaling this to a satellite galaxy luminosity function (LF) by adopting a fixed mass-to-light ratio consistent with the brightest objects overpredicts the number of satellite galaxies of the Milky Way (MW) and M31 by a factor of several hundred $-$ the so-called `missing satellite' problem (e.g. \\citealt{Klypin1999apj, Moore1999apj}; for a recent review see \\citealt{Kravtsov2010Review}). For many years, theoretical models of galaxy formation have predicted that the star formation efficiency in low-mass haloes can be strongly suppressed through a combination of photoionization (\\citealt{Efstathiou1992}) and supernova feedback (\\citealt{WhiteRees1978}). This suppression renders many such haloes permanently `dark' and may solve the apparent discrepancy (e.g. \\citealt{Kauffmann1993, Bullock2000, Gnedin2000, Benson02sats, Okamoto09Frenk, Maccio2009ExplainLF}; \\citealt{LiHelmiLucia2010}; \\citealt{Stringer09arxiv}). Recent discoveries of low luminosity satellites in the MW and M31 and corrections for incompleteness (e.g. \\citealt{Koposov2008MWLF, Tollerud2008}) imply reasonable agreement between observations and theoretical predictions for MW-like haloes (\\citealt{Benson02sats}). However, because they rely on uncertain baryonic physics and currently limited data, these comparisons so far only provide an indirect and uncertain test of the abundance and mass function of CDM substructures. Since gravitational lensing traces mass, regardless of its luminosity, it can provide a powerful alternative tool to investigate otherwise `hidden' subhaloes and directly test the CDM model (e.g. \\citealt{MKA09SubImageAnomaly, VKBTG2009}).   Flux-ratio anomalies observed in multiple images of lensed quasars are often cited as evidence for substructures in lensing galaxies (e.g. \\citealt{MS1998mn, MM2001, MZ2002, DK2002, Chiba2002, KD2004, Metcalf2004, Chiba2005, Sugai2007, McKean2007, More2009}). Some studies, e.g. \\citet{DK2002}, \\citet{BS2004aa} using CDM simulations have concluded that the substructure population expected for a galaxy-scale lens can reproduce the observed flux anomaly cases. However, other studies, e.g. \\citet{MaoJing04apj}, \\citet{AB06mn}, \\citet{Maccio2006b}, \\citet{Maccio2006} and our own previous work (\\citealt{Dandan09AquI}) have argued that the simulated CDM substructure population is \\textit{not} sufficient to explain the observed frequency of lensing flux-ratio anomalies. This negative conclusion does not yet rule out the predictions of CDM, as the observed sample of `anomalous' lenses is small. Intergalactic haloes along the line of sight may also perturb the lensing potential and cause the observed lensing anomalies (\\citealt{Chen2003, Wambsganss2005, Metcalf2005a, Metcalf2005b}; \\citealt{Miranda2007, PunchweinHilbert2009}).    One extreme case of flux-ratio anomalies is the cusp-caustic violation, which has been observed in five lens systems of a cusp geometry: B0712+472 (\\citealt{Jackson2000B0712,Jackson98B0712}), B1422+231 (\\citealt{Impey1996B1422,Patnaik2001B1422}), B2045+265 (\\citealt{Fassnacht1999B2045}), RXJ1131-1231 (\\citealt{Sluse2003aaJ1131}) and RXJ0911+0551 (\\citealt{Bade1997aa317, Burud1998apjl}). In the first three cases (those detected in radio observations) these have been attributed to substructure lensing, while the remaining two are thought to be due to microlensing (\\citealt{Anguita2008Microlensing, Morgan2006RXJ1131}).  In \\citet{Dandan09AquI} (which we summarise below) we found that even with the well-resolved subhalo population from the Aquarius simulations, the observed cusp lenses still violate the cusp-caustic relation more frequently than expected for a concordance CDM model. However, our earlier work did not consider the effect of baryons concentrated at the centre of subhaloes, which may modify their density profiles and so alter their lensing properties. In this work, we carry out a more thorough analysis of the lensing effect of baryons in substructures, proceeding as follows. In Section 2 we use a semi-analytic model of galaxy formation to predict the baryon content of subhaloes in each Aquarius simulation: we model the density distribution of baryons in each subhalo and re-run our lensing simulations to study cusp-caustic violations. In Section 3, to estimate the effects of possible baryonic concentrations below the resolution limit of the simulation, we further include an empirical Galactic globular cluster population (Harris 1996) in our model. Finally, as typical surface density fluctuations of only a few per cent would be sufficient to cause flux-ratio anomalies (\\citealt{MS1998mn, Li2006SPH}), we consider irregularities in the surface density of the main halo due to the debris of tidally disrupted satellites (the diffuse component) in Section 4. Section 5 provides a short summary and a brief discussion of our results. ",
        "conclusions": "We briefly discuss a few issues which may affect our conclusions. Firstly, the observed sample is small, consisting of only five cusp lenses, among which are the three systems detected in the radio band by the CLASS survey (\\citealt{Myers2003, Browne2003}). This survey shows an anomalously high incidence of secondary lenses (bright satellites) (\\citealt{BMK2008,Jackson2009}).  A larger sample of galaxy-scale lenses from future instruments will provide better statistical constraints on substructure abundance in galaxy-sized CDM haloes \\citep{VK2009}. New lenses are likely to be discovered in optical surveys (such as PanSTARRS and LSST), although radio and mid-infrared follow-up will be necessary to distinguish anomalies attributable to substructure from microlensing events \\citep[e.g.][]{WambsganssPaczynskiSchneider1990, WittMaoSchechter1995, Pooley2007, Eigenbrod2008}. High-resolution radio surveys (to be carried out by LOFAR and SKA, for example) will be a productive way to discover more of these systems.  Secondly, most lensing galaxies are massive elliptical galaxies due to their large lensing cross-sections (\\citealt{TOG}). Compared with their spiral and irregular counterparts, massive ellipticals are known to host more GCs (e.g. \\citealt{FMM1982, Harris1991, Harris1993, West1993}). In this work, we consider the lensing effect of GCs by imposing the observed Milky-Way GC population. If the lens galaxies are much more massive than the Aquarius haloes, this assumption may underestimate the appropriate GC abundance. However, even with an increased number of GCs typical of massive ellipticals, we would not expect major changes to our conclusions.  Thirdly, the semi-analytic galaxies that we have included in the lensing simulation are all hosted by dark matter subhaloes above the resolution limit of the Aquarius simulations. However, semi-analytic models also follow the attached stellar components even after their host dark matter (sub)haloes are tidally stripped below 20 particles ($\\sim 10^5 h^{-1} M_{\\odot}$). The baryonic masses of these `sub-resolution' stellar components are in the same range as those of satellite galaxies, and each is associated with a bound dark structure no more massive than $10^5 h^{-1} M_{\\odot}$. In {\\sc galform} their survival depends on an estimate of their merger timescales, which is subject to the treatment of dynamical friction. The radial distribution of the centres (most-bound particles) of these objects is more concentrated than that of their counterparts with resolved dark haloes. Nevertheless, including these sub-resolution `orphan galaxies' only marginally increases the cusp-violation probability under the assumption of singular isothermal spheres for their density profiles.  Finally, when including semi-analytic galaxies to account for the lensing effect of baryonic substructures, we have assumed that the overall galactic potential and the dynamics of subhaloes would not be significantly modified if baryons were added to our dark matter-only simulations. However, a concentration of baryons at the centre of a subhalo may increase its resilience to tidal stripping. The abundance of subhaloes in the inner region of the main halo may increase as a result of this effect. Our model does not account for this `baryon-enhanced' survival of a small number of massive objects, although we do not consider this omission to be significant. There are likely to be few such objects, as a subhalo must typically lose more than $\\sim90\\%$ of its mass before the structure of its central (stellar) core influences further stripping (see e.g. \\citealt{Jorge08tidal}). This naturally limits the number of `marginally destroyed' cases that would change in a fully self-consistent model. Another effect of including baryons would be to increase the concentration of the main halo in our simulation, resulting in a stronger tidal field. However, studies of disc-subhalo interactions using the Aquarius haloes (Lowing et al. in prep) support our assumption that the subhalo population (in the regime relevant to lensing) is only marginally affected by the presence of a realistic stellar disc. Disc shocking (e.g. \\citealt{Kazantzidis2009, OnghiaSpringle2009DiskShocking}) may also reduce the inner substructure abundance, which would further aggravate the problem of explaining the observed lensing anomalies. \\\\    In \\citet{Dandan09AquI}, the dark matter-only Aquarius simulations were used to study the relation between dark matter substructure abundance and lensing flux-ratio anomalies, in particular the cusp-caustic violations observed for multiply imaged quasars. We found that the dark substructures intrinsic to a typical galaxy-scale lensing halo were not sufficient to explain the observed frequency of cusp-caustic violations. In this work, we have considered whether this expectation changes when satellite galaxies in subhaloes are taken into account, or when the lensing halo is assumed to contain a MW-like globular cluster population that is not associated with surviving dark matter substructures. We have also considered streams of dark matter identified in the Aquarius simulations, in order to estimate the lensing effect of irregularities in the halo itself. We conclude that the abundance of intrinsic substructures, dark or bright, bound or diffuse, cannot fully account for the observed cusp-violation frequency. Taken at face value, this lack of substructure suggests a serious problem for the CDM model. Warm dark matter models, which could reduce the satellite abundance and may help to bring the dwarf galaxy LF into agreement with observations without invoking strong feedback or photoionization effects, (e.g. \\citealt{SavalaWhite2010}) would only make this problem worse. However, it is possible that the observed frequency of flux anomalies are strongly biased by the small number statistics of CLASS.  Previous studies have shown that intergalactic haloes ($M \\lesssim 10^{10} M_{\\odot}$) projected along the line of sight can cause surface density fluctuations at the level of 1-10 per cent, making them a probable source of lensing flux anomalies \\citep{Chen2003, Wambsganss2005, Metcalf2005a, Metcalf2005b,Miranda2007}. In future work, we will examine the contribution of this large-scale structure to observations of the cusp-caustic violation rate, using high-resolution cosmological volume simulations, which self-consistently model the halo mass function and clustering along the line of sight \\citep[e.g.][]{PunchweinHilbert2009}."
    },
    "1004/1004.3741_arXiv.txt": {
        "abstract": "We explore how finite integration times or equivalently temporal binning induces morphological distortions to the transit lightcurve.  These distortions, if uncorrected for, lead to the retrieval of erroneous system parameters and may even lead to some planetary candidates being rejected as ostensibly unphysical.  We provide analytic expressions for estimating the disturbance to the various lightcurve parameters as a function of the integration time. These effects are particularly crucial in light of the long-cadence photometry often used for discovering new exoplanets by for example \\emph{CoRoT} and the \\emph{Kepler Mission} (8.5 and 30 minutes).  One of the dominant effects of long integration times is a systematic underestimation of the lightcurve derived stellar density, which has significant ramifications for transit surveys. We present a discussion of numerical integration techniques to compensate for the effects and produce expressions to quickly estimate the errors of such techniques, as a function of integration time and numerical resolution. This allows for an economic choice of resolution before attempting fits of long-cadence lightcurves. ",
        "introduction": "Transiting extrasolar planets have become a powerful tool since the first discovery by \\citet{cha00} and \\citet{hen00}.  The detection of a transit allows for a measurement of the planet-star ratio-of-radii, which may be used to infer the planetary radius once the stellar radius has been determined.  Multi-wavelength measurements may be used to retrieve a spectrum of the light being absorbed by the planetary atmosphere and thus a deduction of the molecular constituents (e.g. Tinetti et al. 2007).  Timing of the eclipse minima and duration permits for the detection of other bodies in the system (Agol et al. 2005; Holman et al. 2005) and companion moons (Kipping 2009a,b).  In recent years, various authors have published increasingly ingenious methods of characterizing exoplanets through transits, including thermal emission detections (Deming et al. 2006), spin-orbit alignment (Winn et al. 2005), oblateness measurements (Seager \\& Hui 2002; Carter et al. 2010) and many more.  As we look for these small whispers of the exoplanet's nature, we are in fact examining the morphology of the transit lightcurve in increasing detail.  It is therefore paramount that the methods for modeling these lightcurves are in the most precise form possible.  In this paper, we explore the effects of finite integration time on the transit lightcurve, or equivalently the act of temporal binning.  Long-cadence (LC) data smears out the transit lightcurve signal into a broader shape which will lead to an erroneous retrieval of the system parameters, if unaccounted for.  We will first outline the nature of the morphological distortions in \\S2 and then derive analytic expressions for the approximate systematic error in the retrieved parameters, as a function of integration time.  It is shown that the effect can be significant enough to even lead to some planetary candidates being rejected as being unphysical, unless the effect is correctly accounted for.  Obviously, this is critical for missions like \\emph{CoRoT} and \\emph{Kepler} which employ LC photometry to detect new transiting planets.  In \\S3, we will discuss numerical techniques to correctly model the transit lightcurve and derive the errors of these methods as a function of integration time and numerical resolution. These expressions allow for a simple evaluation of the expected consequences of finite integration time and thus will guide observers in choosing how to deal with this issue. ",
        "conclusions": "We have explored how long-cadence data, with particular focus on \\emph{Kepler}, causes severe systematic errors in the retrieved physical parameters, unless accounted for.  The effect is valid for any finite exposure time but increases with longer cadences. Long-cadence data smears out the lightcurve morphology, which acts to stretch out the ingress/egress duration and suppress limb darkening in the lightcurve trough.  These two effects act to increase and decrease the retrieved impact parameter respectively.  Critically, overestimating the impact parameter is shown to lead to severe underestimations of the stellar density which could lead to planetary candidates being rejected on the basis of being unphysical.  Numerical integration techniques permit for improved modeling of the transit lightcurve. We discuss two particular methods, the composite Simpson's method and resampling. We provide expressions for estimating the errors of these techniques and find that both methods produce an error which scales as $N^{-2}$ where $N$ is the numerical resolution of the techniques. Out of these two discussed methods, the resampling approach yields a greater efficiency. %"
    },
    "1004/1004.3788_arXiv.txt": {
        "abstract": "The protostellar mass function (PMF) is the Present-Day Mass Function of the protostars in a region of star formation. It is determined by the initial mass function weighted by the accretion time. The PMF thus depends on the accretion history of protostars and in principle provides a powerful tool for observationally distinguishing different protostellar accretion models. We consider three basic models here: the Isothermal Sphere model \\citep{shu77}, the Turbulent Core model \\citep{mck03}, and an approximate representation of the Competitive Accretion model \\citep{bon97, bon01a}. We also consider modified versions of these accretion models, in which the accretion rate tapers off linearly in time. Finally, we allow for an overall acceleration in the rate of star formation. At present, it is not possible to directly determine the PMF since protostellar masses are not currently measurable. We carry out an approximate comparison of predicted PMFs with observation by using the theory to infer the conditions in the ambient medium in several star-forming regions. Tapered and accelerating models generally agree better with observed star-formation times than models without tapering or acceleration, but uncertainties in the accretion models and in the observations do not allow one to rule out any of the proposed models at present. The PMF is essential for the calculation of the Protostellar Luminosity Function, however, and this enables stronger conclusions to be drawn \\citep{off10}. ",
        "introduction": "The Initial Mass Function (IMF) is of central importance in star formation since the properties of a star and its effects on the surrounding medium are determined primarily by its initial mass. One of the main approaches for inferring the IMF is to apply evolutionary models to observations of the Present-Day Mass Function (PDMF) of a group of stars. The IMF is created during the process of star formation, when the mass of each protostar grows by accretion until it reaches its final value. Observations of a region of star formation can, in principle, allow one to infer the PDMF of the protostars there; we refer to this as the Protostellar Mass Function (PMF). The PMF depends on both the IMF, which determines the relative number of stars when the star formation is complete, and on the process of star formation, which determines how long each protostar spends at a given mass. Because of this latter property, the PMF is potentially a powerful tool for inferring the nature of the star formation process. For example, inside-out collapse of an isothermal sphere \\citep{shu77} leads to protostellar lifetimes that are proportional to the mass of the final star, whereas models based on competitive accretion \\citep{zin82,bon97} have protostellar lifetimes that are independent of the mass of the final star; the inside-out collapse of a turbulent core \\citep{mck02,mck03} has an intermediate behavior. As we shall see below, it follows that the PMFs predicted by these models are quite different.  There are of course other approaches for observationally distinguishing among the different theories of star formation. One is to study the relation between the mass distribution of density concentrations in molecular clouds (the Core Mass Function, or CMF) and the stellar IMF, which are observed to be similar \\citep{mck07}. This similarity is the basis for recent theories of the IMF, which are predicated upon the assumption that stellar masses are determined by the production of gravitationally bound cores in turbulent molecular clouds \\citep{pad02,pad07,hen08,hen09}. Such a direct connection between the CMF and IMF would appear to be a natural prediction of theories based on the inside-out collapse of gravitationally bound cores \\citep{shu77,mck02}, but inconsistent with the theory of competitive accretion (e.g., \\citealp{bon97}).  Indeed, in inside-out collapse theories, the PMF and CMF are closely related, since protostars are forming in a significant fraction of cores. \\citet{cla07} have pointed out that the mass distribution of cores depends on their lifetime, just as we shall see below that the PMF depends on the accretion time, and that this dependence would make the slope of the IMF significantly steeper than that of the CMF. However, in the Turbulent Core model \\citep{mck03} the star formation time depends only weakly on the mass of the core, and \\citep{hen09} have argued that as a result the slopes of the CMF and IMF would agree within the observational errors. It is clear that additional approaches for testing theories of star formation are needed.  The principal difficulty with the PMF method developed here is that at present, it is difficult to infer the masses of individual protostars, both because the evolutionary tracks of protostars are uncertain (e.g., \\citealp{cha07}) and because of the effects of accretion on the spectrum of the protostar, which are difficult to quantify. These difficulties will presumably be overcome in the future. Currently, the best way to determine the PMF appears to be through observations of the Protostellar Luminosity Function (PLF), which wlll be discussed in Offner \\& McKee (2010; hereafter Paper II).  In this work, we define a protostar as an embedded source that is still experiencing signficant accretion and thus is more than a few percent from its initial stellar mass. Observationally, young stellar objects are characterized on the basis of the slopes of their spectral energy distributions into four classes, 0-III \\citep{ada87,andre94}. However, since geometric effects due to disk and outflow orientation influence the radiation reprocessing of the embedded source, it is difficult to directly map classes to physical stages. For the latter, we use the classification proposed by \\citet{cra08}: Stage 0 objects have protostellar masses that are less than or equal to their envelope, while Stage I objects have protostellar masses larger than the envelope. Once the envelope mass falls below 0.1 $\\msun$, the object is considered to have completed its main accretion phase and entered the Stage II phase. Although these stages do not directly correlate with class definitions, in practice, Class 0 and Class I sources approximately correspond to Stage 0 and Stage I (see \\S 6 for additional discussion).   In our analysis, we shall make two main assumptions: First, we shall generally adopt a \\citet{cha05} IMF with an imposed upper mass cutoff at $\\mup$. This IMF has a log-normal form below $1 M_\\odot$ and a power-law form above. For star formation regions large enough to fully sample the IMF, $\\mup\\sim 150 M_\\odot$ \\citep{fig05}. However, as we show below, regions of low-mass star formation often have maximum stellar masses of only a few $M_\\odot$, even though in some cases one would expect stars more massive than that according to the \\citet{cha05} IMF. It does not appear that this is a selection effect, since the data we compare with \\citep{eva09} represent a complete survey of nearby star-forming regions. If the deficit of high-mass stars is not a selection effect or a rare statistical fluctuation, then it must be the result of a physical inability to form high-mass stars in some regions---e.g., \\citet{kru08}. In any case, our assumption is that the IMF above $1 M_\\odot$ in these regions can be approximated by a truncated power-law im which the upper limit on the final masses of the protostars is set by the inferred upper limit on the masses of the more numerous newly formed stars (primarily Class II sources) in the sample. In the applications to observations below, we shall generally adopt an upper mass limit of $3 M_\\odot$. Second, we shall assume that the accretion rate onto the protostar can be expressed as a simple function of the instantaneous protostellar mass, the final protostellar mass (i.e., the initial stellar mass), and the time. In particular, we ignore complications associated with an initial Larson-Penston accretion phase, when the accretion rate can be much larger than the value expected on the basis of dimensional analysis (e.g., \\citealp{mck07}; however, in one of the most complete simulations of the formation of a star to date, \\citet{mac09} found only a small enhancement in the accretion rate at  early times); Furthermore, we average over any temporal fluctuations in the accretion rate, such as occur in FU Ori outbursts \\citep{har96}. In principle, the accretion rate can depend on the location of the protostar in its natal cluster; in our treatment, any such dependence is encoded in the dependence on the final mass of the protostar.  The protostellar mass and luminosity functions were first considered by \\citet{fle94a,fle94b}. In their two companion papers, they derived the time-dependent mass and luminosity functions for young embedded stellar clusters, including the luminosity contributions from protostars through main-sequence stars. Our work differs from theirs in several important respects: (1) We consider a variety of different theories for star formation, allowing for the accretion rate to depend on the protostellar mass, $m$, and the final stellar mass, $m_f$, whereas they assumed that the accretion rate for all the stars was constant. Non-constant accretion allows us to test different theories of star formation. In particular, as noted by \\citet{shu87}, the constant accretion-rate model, which was developed for low-mass star formation, is unlikely to apply to high-mass star formation. (2) As we mention above, observations in the solar neighborhood suggest an upper cutoff in the mass distribution at a few solar masses; we explicitly allow for this in our adopted IMF, whereas Fletcher \\& Stahler did not. (3) They treated the protostellar, pre-main sequence and main sequence stages, whereas we consider only the protostellar case. (4) They focused on the case in which the star formation rate is constant for a given time period, whereas we consider both the steady-state case and the case of accelerating star formation \\citep{pal99,pal00}. (5) Finally, we note a difference in approach: they determined the probability that a star of a given mass would cease accreting and become a pre-main sequence star, whereas we label each protostar by its instantaneous mass and its final mass. As we shall see below, this makes it straightforward for us to consider the case of tapered accretion, in which the accretion rate declines prior to the end of the protostellar stage.   We begin with the derivation of the PMF in terms of the IMF and the accretion history of the protostars in \\S 2. We consider four accretion histories in \\S 3: the collapse of an isothermal sphere \\citep{shu77}, the Turbulent Core model \\citep{mck02,mck03}, a blend of these two models, and the Competitive Accretion model \\citep{bon97}. We also consider the effects of tapering the accretion rate as the mass approaches the final mass, as suggested by \\citet{mye98}. In \\S 4, we evaluate the PMF for these accretion histories, both analytically for a \\cite{sal55} IMF and numerically for a \\citet{cha05} IMF. \\citet{pal00} have suggested that nearby low-mass star-forming regions have star formation rates that are accelerating in time, and we show how this affects the PMF in \\S 5. We make a brief comparison with observation in \\S 6 (a more extensive comparison will be given in Paper II) and summarize our conclusions in \\S 7. ",
        "conclusions": "We have analyzed the Protostellar Mass Function (PMF), which is the Present-Day Mass Function (PDMF) of a cluster of protostars. The PMF builds on the IMF, which measures the final mass distribution of a cluster of stars after completion of the formation process. The PMF depends on the mass dependence of the formation time of the stars, $t_f$, and therefore can in principle provide an observational diagnostic of theories of star formation. Direct observation of the PMF must await further observational and theoretical advances, since it is not currently possible to infer the mass for protostars because they are generally extremely obscured and the luminosity is dominated by accretion. Nonetheless, it is possible to compare theoretical models with measurements of the mean lifetime of protostars in different molecular clouds as we have done here. Furthermore, the PMF is the basis for the Protostellar Luminosity Function, which is directly accessible to observation (Paper II).   We have made two key assumptions in inferring the PMF. First, since no stars above 3 $\\msun$ are seen in the low-mass clusters we have analyzed \\citep{eva09, enoch09}, we have assumed that none of the smaller number of protostars in these clusters will grow to more than $3 \\msun$; we therefore imposed an upper cutoff of $3 M_\\odot$ on the IMF in our analysis. Given the number of Class II sources, which have very nearly reached their final masses, we note that 12 Class II sources with $m>3\\msun$ would be expected in these clusters based on a \\citet{cha05} IMF. These clusters were selected on the basis of their proximity so as to allow a thorough study of the Class 0, I, and II sources, so it is unlikely that this anomalous IMF is due to a selection effect. The second key assumption we have made is that the accretion rate is a simple function of only the current protostellar mass, the final mass and the time. We thus do not allow for variations in the accretion rate due to a brief high-accretion Larson-Penston phase or to temporal fluctuations in the accretion rate (although insofar as such fluctuations are random and there is a statistically large sample of protostars, they should not significantly affect the PMF).    We consider four accretion rate histories: the classical Isothermal Sphere accretion \\citep{shu77}, the Turbulent Core model \\citep{mck02,mck03}, a blend of the two (Two-Component Turbulent Core, 2C Turbulent Core), and an analytic approximation for the Competitive Accretion model \\citep{bon97,bon01a}. There are substantial uncertainties in the accretion rates for each model: In all cases, one must allow for the effect of protostellar outflows, which can reduce the accretion rate by a factor of a few \\citep{mat00}. For the first three, there is a countervailing correction needed to allow for an initial infall velocity. Our approximation for the Competitive Accretion model captures many of its essential features, but since the model itself is based primarily on numerical simulations, there is no fully analytic form for it.  In comparing the models with observation, we assume that the star formation is steady or accelerating; since the Competitive Accretion model has been developed for the evolution of individual star clusters, the comparison with observation is valid for this model only if a number of clusters are sampled, either because a forming cluster is comprised of a number of sub-clusters or because data from different clusters are averaged together.  The mean protostellar mass (Fig. \\ref{meanmassplot}) and the ratio of the median mass to the mean mass (Fig. \\ref{medmassplot}) depend sensitively on the accretion history. The Turbulent Core and Competitive Accretion Models have accretion rates that increase with mass and therefore with time ($\\dot m\\propto m^j$, with $j=\\frac 12,\\, \\frac 23$ for the two models respectively). As a result, protostars of a given final mass, $\\mf$, spend a smaller fraction of their lives at high mass than in the Isothermal Sphere model. Furthermore, these two models have accretion rates that increase with $\\mf$, so that it takes less time to form a high-mass star than in the Isothermal Sphere model. Both effects are stronger for the Competive Accretion model. As a result, the mean protostellar mass increases systematically from the Competitive Accretion model to the Turbulent Core model to the Two-Component Turbulent Core model to the Isothermal Sphere model. The ratio of the median to mean protostellar mass follows the same ordering, and the same effect shows up in the plots of the PMF in Figure \\ref{pmf}.   A common feature of all the accretion models is that the accretion rate remains constant or (usually) increases until the time at which the protostar reaches its final mass, when it abruptly ceases. In reality, as pointed out by \\citet{mye98}, the accretion will turn off gradually. To allow for this, we have inserted  a factor $[1-(t/t_f)^n]$ into the accretion rates; we refer to this as tapered accretion. In practice, we focused on the case $n=1$, which gives an accretion time $t_f$ twice as long as would be expected in the absence of tapering. This has the effect of increasing the temperature, column density or density of the model needed to match a given observed formation time. For example, in the Isothermal Sphere model, the accretion rate is proportional to $T^{3/2}$, so the temperature needed to match the observations of a given formation time is $2^{2/3}$ times greater for a tapered model than for an untapered one. Not only is tapering physically plausible, it also generally results in models that are in better agreement with observation. As shown in Figure \\ref{allpmf}, tapering moves the peak of the PMF to higher masses since stars spend a larger fraction of their lives at high mass when the accretion slows down at the end of the accretion process.  The rate of star formation should accelerate in time in a contracting gas cloud, and \\citet{pal00} found direct evidence for such acceleration in a number of nearby star-forming clusters. We generalized our analysis of the PMF to the case in which the star formation rate is time dependent in \\S \\ref{sec:accel}. For simplicity, we have assumed that the acceleration applied only to the rate at which stars formed, not to the accretion rates of individual stars. This is a reasonably good approximation for the cases we analyzed, which have star-formation times that are significantly smaller than the time scale for acceleration. Moreover, the approximation is even better for for the Isothermal Sphere model, since the accretion rate depends on the temperature and radiative losses maintain an approximately constant temperature.  On the other hand,  the time scale for acceleration is often comparable to or less than the mean lifetime of Class II sources, so the ratio of the number of protostars to the number of Class II sources is larger than in the non-accelerating case. As a result, as shown in the Appendix, the ``observed\" star-formation time, which is given by equation (\\ref{eq:tftii}), exceeds the actual star-formation time. We find that acceleration does not have a substantial effect on the PMF. Rather, its primary effect is to reduce the inferred time scale for the formation of individual stars, thereby increasing the inferred temperature, column density or mean density, depending on the accretion model.  In the absence of any direct information on protostellar masses, we were able to carry out only a very crude comparison with observation: Using the observed star-formation time scales in two different clusters, we computed the implied temperature (Isothermal Sphere model), surface density (Turbulent Core model), and mean density (Competitive Accretion model), and then compared with the observed values of these parameters. We found that the tapered accretion and accelerating star formation models were somewhat better than untapered, non-acceleration models, but we could not draw any firm conclusions due to uncertainties in both the observations and in the models, which have accretion rates that are probably uncertain by a factor of 2. In addition, the molecular clouds have an unknown internal structure and the IMF can have significant statistical and perhaps physical fluctuations from one cloud to another. In Paper II we shall show that the Protostellar Luminosity Function is a more powerful diagnostic for inferring the accretion mechanism."
    },
    "1004/1004.2351_arXiv.txt": {
        "abstract": "Using archival RXTE data, we show that the ultracompact X-ray binary in NGC~1851 exhibits large amplitude X-ray flux varations of more than a factor of 10 on timescales of days to weeks and undergoes sustained periods of months where the time-averaged luminosty varies by factors of two.  Variations of this magnitude and timescale have not been reported previously in other ultracompact X-ray binaries.  Mass transfer in ultracompact binaries is thought to be driven by gravitational radiation and the predicted transfer rates are so high that the disks of ultracompact binaries with orbits as short as that of this object should not be susceptible to ionization instabilities. Therefore the variability characteristics we observe were unexpected, and need to be understood.  We briefly discuss a few alternatives for producing the observed variations in light of the fact that the viscous timescale of the disk is of order a week, comparable to the shorter time scale variation that is observed but much less than the longer term variation.  We also discuss the implications for interpretation of observations of extragalactic binaries if the type of variability seen in the source in NGC~1851 is typical. ",
        "introduction": "Ultracompact X-ray binaries are semi-detached binary star systems in which a neutron star or a black hole accretes material from a Roche-lobe overflowing white dwarf.  They are of astrophysical interest for a variety of reasons.  It is expected that some ultracompact X-ray binaries should be detectable as gravitational wave sources with missions like LISA (see e.g. Benacquisita 1999). Additionally, they present an opportunity to study plasma astrophysics in hydrogen-free gas, where the mass-to-charge ratio is different than in most other astrophysical systems.  The most recent compilation of ultracompact X-ray binaries in the Galaxy contained 27 candidates (in 't Zand, Jonker \\& Markwardt 2007 -- IZJM07).  The most secure identifications of ultracompact X-ray binaries come through direct measurements of their orbital periods, ranging in the observed sample from 11 to 50 minutes (IZJM; Galloway et al. 2002; Middleditch et al. 1981; Markwardt et al. 2002, 2003; Stella et al. 1987; Homer et al. 1996; White \\& Swank 1982; Dieball et al. 2005; and Zurek et al. 2009 for the period measurement of the object discussed in this paper).  At the time of publication of IZJM07, only 7 had well-estimated orbital periods.  The remainder were classified as ultracompact X-ray binaries on the bases of some combination of tentative orbital period measurements, deep optical spectra lacking hydrogen emission lines, high ratios of X-ray to optical flux, or persistent emission at low fractions of the Eddington rate (IZJM07 and references within).  The persistent ultracompact X-ray binaries typically have X-ray luminosities from about $10^{36}-10^{37}$ ergs/sec.  Understanding the formation mechanisms for ultracompact X-ray binaries is also of great interest.  The shortest period ultracompact X-ray binaries (i.e. those with orbital periods less than 30 minutes), for example, are predominantly in globular clusters (see e.g. the tabulation in IZJM07), and thus their production may be dominated by mechanisms which are inherently stellar dynamical, such as direct collisions between neutron stars and red giants (e.g. Verbunt 1987). However, theoretical work does suggest that it is possible to form white dwarf-neutron star binaries with orbital periods less than 30 minutes by going through an intermediate phase where the donor star is a helium star (e.g. Savonije, de Kool \\& van den Heuvel 1986).  At least the shortest period ultracompact binaries are relatively easy to detect, as they are bright ($L_X \\gtsim 10^{36}$ ergs/sec), persistent X-ray emitters. These thus represent samples of double degenerate stars with relatively well understood selection effects, and which can be seen nearly throughout the Galaxy.  Other classes of binary compact objects which are of great astrophysical interest include double white dwarfs, which may be the progenitors of Type Ia supernovae (Iben \\& Tutukov 1984) and double neutron stars, which are showing increasingly strong evidence for being the progenitors of short-hard gamma-ray bursts (e.g. Bloom et al. 2006).  Understanding the evolution of double compact binaries is likely to require using multiple source classes to gather constraints. Double neutron stars, for example, are both rare, and very difficult to detect.  Mass transferring double white dwarf systems can be studied in great deal once they are detected (e.g. Roelofs et al. 2007), but their sample sizes are only slightly larger than those for the ultracompact X-ray binaries, and the relative robustness to selection effects in surveys of neutron star ultracompact X-ray binaries and mass transferring double white dwarf systems is not yet clear.  In the absence of orbital period estimates, candidate ultracompact X-ray binaries can be found through measurements of large X-ray to optical flux ratios (e.g. Juett et al. 2001).  This signature is expected for short orbital period systems since the optical emission from most bright X-ray binaries is dominated by reprocessing of X-rays in the outer accretion disk, leading to a correlation between X-ray to optical flux ratio and orbital period (van Paradijs \\& McClintock 1994).  Additional candidate UCXBs can be identified from strong upper limits on hydrogen emission lines in optical spectra (e.g. Nelemans et al. 2004), or by detection of persistent bright hard X-ray emission combined with low rates of Type I X-ray bursts (IZJM07).  Bildsten \\& Deloye (2004) have suggested that ultracompact X-ray binaries may be the dominant population in elliptical galaxies.  The typical X-ray luminosity functions seen in elliptical galaxies (e.g. Kundu et al. 2002; Kim \\& Fabbiano 2004; Gilfanov 2004), and the similarities between the luminosity functions of globular cluster X-ray sources and non-cluster X-ray sources in the elliptical galaxies (Kundu et al. 2002) are a major part of the evidence they suggest for this idea.  The elliptical galaxy X-ray observations used in this analysis typically consist of snapshots of about half a day.  Strong variability can cause the observed luminosity functions in snapshots to be different from the luminosity function which would be obtained when averaging over long durations.  In this paper, we present the X-ray light curve of 4U~0513-40 from RXTE pointed observations.  This X-ray binary, in the Galactic globular cluster NGC~1851, has a 17-minute orbital period -- short enough that irradiated disc models (Lasota et al. 2008) predict it should not be subject to the standard ionisation instability (i.e. it should always be sufficiently ionised to exist in the high viscosity state), and hence should persistently be a bright X-ray source.  It {\\it is} persistent in the sense that it is always detectable above $L_X$ of about $10^{36}$ ergs/second.  However, we also present evidence for variability of a factor of $\\sim10$ in X-ray luminosity on timescales of $\\sim$ weeks, and a factor of more than 20 overall. Such variability is unusual for ultracompact X-ray binaries, and contrary to the expectations from ionisation instability theory for such source.  4U~0513-40 also shows variations at the factor of 2 level in the time averaged luminosity over years, with no strong indications of any periodicities in the long term light curve. ",
        "conclusions": "The ultracompact X-ray binary 4U~0513-40 in the globular cluster NGC~1851 shows large amplitude variability on a variety of timescales. Of particular interest are the variations of factors of $\\sim10$ on timescales of several weeks and the variations in the luminosity when time averaged over $\\sim$ 1 year.  While it may be possible to explain the $\\sim$ weeks timescale variability through tidal instabilities in the accretion disc, the latter almost certainly must be explained in terms of modulation of the actual accretion rate.  Given that this system's evolution is thought to be driven primarily by gravitational radiation, and that the mean luminosity agrees well with that scenario, the finding of significant deviations from a constant mass transfer rate is interesting, and needs to be explained."
    },
    "1004/1004.0448_arXiv.txt": {
        "abstract": "We demonstrate that it is possible to calculate not only the mean of an underlying population but also its dispersion, given only a single observation and physically reasonable constraints (i.e., that the quantities under consideration are non-negative and bounded). We suggest that this counter-intuitive conclusion is in fact at the heart of most modeling of astronomical data. ",
        "introduction": "In their discussion of optimizing search strategies for the simultaneous discovery and characterization of variable stars using predetermined but highly non-uniform sampling, Madore \\& Freedman (2005) investigated the small-sample limits. Their emphasis was on situations where only a handful of observations (2-20 say) might be available from which it would be necessary to determine periods, measure mean magnitudes and derive full amplitudes. In this paper we discuss how it is possible use a single observation to derive confidence intervals on the mean (i.e., limits on variability, or estimates of the population variance, or its full amplitude, etc.).  Shortly after the original manuscript was completed it was pointed out that similar arguments have been made much earlier by Gott (1993). Gott was interested in the estimation of ages and lifetimes of events based single instances; and he presented confidence intervals for the longevity of an object given a single observation of its present age. He assumed a flat prior (of finite duration), and took the stance that there cannot be anything special about any given time that a random observation is made of an object that has a finite lifetime. We return to a contextual discussion of this example, and a generalization of it to modeling, at the end of this paper. ",
        "conclusions": "Our interest in this formalism started with a desire to characterize luminosity variations of astronomical objects (in the first instance through their first two moments, chosen to be their means and amplitudes) having ahighly restricted number of observations. Here we have taken that thought experiment to its limit of a singular observation. For an underlying distribution of finite extent (or duration) the dispersion and the full amplitude are equivalent, differing only by a constant scale factor in any given case. Accordingly, our derivation of the variance of a population based on a single observation is conceptually equivalent to Gott's (1993) derivation of future longevity (that is, Gott's longevity is a one-directional semi-amplitude) for an object based on a ``single'' observation of its present age.  When dealing with physical observations certain assumptions are tacitly taken for granted. The obvious assumptions are, that those quantities are non-negative, and that the underlying population from which they are drawn does not have an infinite range (in time, mass, energy, size, {\\it etc.})  However, those same assumptions carry additionally useful (prior) quantitative information that can be used to constrain limits on observations as they are obtained. This paper had the intent of making those assumptions explicit and then formalizing the mathematical consequences expressed as confidence intervals on the underlying population as derived from one observation  While it was our hope that this intentionally short contribution might stimulate others into finding applications not obvious to the author, the referee requested that examples be given of how this formalism might be applied to astronomy. In keeping with the spirit of this paper we offer a modest example from the past, and predict that there will be future examples; all of this based on a sample of N = 1.  Laplace developed his ``rule of succession'' when confronted with a question as to the mathematical (not the physical) probability that the Sun will rise tomorrow given its past (statistical) performance. After observing N events, Laplace derived that the probability of the next occurrence was $(N+1)/(N+2)$. This would suggest that on the first day (N = 0) the probability of the Sun rising was actually 50:50. On the second day (N = 1) the probability would have gone up to 66\\%, and so on. Gott (1993) asked a similar question, not just about the next occurrence of something that has a past persistence, but about the sum of all future occurrences. How long will a thing last, given that we know how old it is now? Gott was reformulating Laplace's question to be, if we have a single measurement (N = 1) of the age of something, what can one say about its total lifetime (its full amplitude), or rather its future longevity (a semi-amplitude). Of course, any such prediction is best described in terms of probabilities, and so rather than predicting a firm lifetime based on a precise age, Gott predicted a forward-looking probability distribution (expressed as a variance) based on a backward-looking age.  One could argue that Gott actually required two observations: one of the time at which something began and another of the time at which the prediction was being made. This may be seen as quibbling but it is, in fact, equivalent to our physical prior, state above, that none of the quantities to which this method applies can drop below zero or become infinite in amplitude (mass, luminosity, time, etc).  Grounded with post-dictions on the longevity of the Soviet Union and the Berlin Wall, Gott went on to predict the longevity of a variety of things astronomically big and small: from the expected demise of {\\it Nature} magazine itself (somewhere between 3.15 and 4,800 years), to the probability that we will end as a civilization (in 5,100 to 8 million years with 95\\% probability), or colonize the Galaxy (the odds are against it). Each of these predictions were based on a single observation, N = 1 and a uniform prior.   It is quite clear that the uniform, non-negative prior distribution of data points discussed above is at one extreme (of simplicity, or of ignorance.)  However, this extreme is also rather inclusive. If the true range of the underlying distribution is smaller, more centrally peaked, or more skewed toward the upper bound of the distribution function, than a uniform distribution, then their confidence intervals will also be smaller than those calculated for a uniform prior; under those conditions the uniform prior is likely to provide a conservative {\\it upper} bound on the uncertainty.  \\subsection{Is This Just Another Name for Modeling?}   Finally, we suggest that aspects of the scientific enterprise as a whole, as practiced by many astronomers in interpreting observational data, might simply be a generalization of the the Sigma One methodology discussed here. Seen in that retrospective light, Sigma One becomes a fairly benign and low-level form of what would otherwise be called ``modeling''.  Consider the observation of a color-magnitude diagram for a composite stellar population, in a nearby galaxy, say. Based on that single observation one could ask what the magnitude and color of any given star might be on the next exposure (whenever that may be.) Depending on the amount of prior knowledge about the underlying distribution function for that star one could make a prediction. Indeed we do this all the time. It is known that intrinsic variables (Cepheids, Miras, RR Lyrae stars, etc.)  occupy fairly well-delineated regions of the color-magnitude diagram. Armed with known amplitudes and timescales one could invoke those distribution functions with their specific means and variances to predict the expected variance in those selected stars. Stars in regions not known to be variable on those same time-scales would have different priors used to predict their means and variances.  But all of this could also be recast into a very different form of the underlying distribution function, in the case where extremely long (astronomically long) timescales are being considered.  The predictive prior for a single observation would then become stellar evolution theory itself. That is, given a star observed (once) today at a given place in the color-magnitude diagram, what is its color and magnitude distribution function integrated over its projected future existence? And then how might the ensemble change with time?  We apparently have no problem in undertaking population synthesis modeling, for example, taking a single integrated spectrum and/or a single color-magnitude diagram and extrapolating it to encompass the entire life history (backward and forward in time) of a given star and/or its associated contemporary population (an entire galaxy).  So our point here is that if we are comfortable extracting very complex ``moments'' (in time and composition, etc.)  from single (but admittedly very rich) observations by invoking very complicated priors (i.e., models), then it should come as no surprise that it is possible to extract more than just one moment (i.e., a mean and a variance at least) from a single data point by assuming very simple priors (i.e., models), in the form of well known, commonly invoked, but simple, distribution functions."
    },
    "1004/1004.2167_arXiv.txt": {
        "abstract": "We explore the emissions by accelerated electrons in shocked shells driven by jets in active galactic nuclei (AGNs). Focusing on powerful sources which host luminous quasars, the synchrotron radiation and inverse Compton (IC) scattering of various  photons that are mainly produced in the  core are considered as radiation processes. We show that the radiative output is dominated by the IC emission for compact sources ($\\lesssim 30{\\rm kpc}$), whereas the synchrotron radiation is more important for larger sources. It is predicted that, for powerful sources ($L_{\\rm j} \\sim 10^{47}{\\rm ergs~s^{-1}}$), ${\\rm GeV}-{\\rm TeV}$ gamma-rays produced via the IC  emissions can be detected by the Fermi satellite and modern Cherenkov telescopes such as MAGIC, HESS and VERITAS if the source is compact. ",
        "introduction": "Relativistic jets in  radio-loud active galactic nuclei (AGNs) dissipate their kinetic energy via interactions with surrounding  interstellar medium (ISM) or intracluster medium (ICM), and inflate a bubble composed of decelerated jet matter, which is often referred to as cocoon. % Initially, the cocoon is highly overpressured against the ambient ISM/ICM  and a strong  shock is driven into the ambient matter. Then a thin shell is formed around the cocoon by the compressed ambient medium. As in other astrophysical shocks, % the shells are expected to be a promising site for particle accelerations, since the shocks are driven into tenuous plasmas. In the present study, we  explore  the evolution of the non-thermal emissions by the accelerated electrons in the shocked shells. We properly take into account the Comptonization of photons of various origins which were not considered in the previous studies.\\cite{FKY07} Focusing on the powerful sources which host luminous quasar in its core, we show, in particular, that the  energy of accelerated electrons is efficiently converted through the IC scattering to high energy $\\gamma$-rays of up to $\\sim 10~{\\rm TeV}$ if the source is relatively compact. ",
        "conclusions": "\\label{summ}  We have explored the temporal evolution of the emissions by accelerated electrons in the shocked shell produced by AGN jets. Below we summarize our main findings in this study.   \\noindent (i) When the source is young and small ($R \\lesssim R_{\\rm IC/syn} \\sim  27 L_{\\rm IR, 46}^{1/2} B_{-5}^{-2}{\\rm kpc}$), the dominant radiative process is the IC scattering of IR photons emitted from the dust torus. For larger sources, on the other hand, the synchrotron emissions dominate over the IC emissions, since the energy density of photons becomes smaller than that of magnetic fields ($U_{B} > U_{\\rm ph} \\propto R^{-2}$). Through the entire evolution, the spectrum is rather broad and flat, and the peak luminosity is approximately given by $\\nu L_{\\rm \\nu, peak} \\sim  3.0 \\times 10^{40} \\epsilon_{-2}L_{45}~{\\rm ergs~s^{-1}}$, since it is roughly equal to the energy injection rate, which is in turn determined by the jet power $L_{\\rm j}$ and acceleration efficiency $\\epsilon_e$.  \\noindent (ii) The spectra of the IC emissions extend up to $\\sim 10~{\\rm TeV}$ gamma-ray energies for a wide range of source size ($R\\sim 1-100~{\\rm kpc}$) and jet power ($L_{\\rm j}\\sim 10^{45}-10^{47}~{\\rm ergs~s^{-1}}$). For most powerful nearby sources  ($L_{\\rm j} \\sim 10^{47}~{\\rm ergs~s^{-1}}$, $D \\lesssim 100~{\\rm Mpc}$), ${\\rm GeV}-{\\rm TeV}$  gamma-rays produced via the IC emissions can be detected by Fermi/LAT as well as by the modern Cherenkov telescopes such as MAGIC, HESS and VERITAS if the source is compact ($R \\lesssim R_{\\rm IC/syn}$). The observation of these emissions  enable us to probe the acceleration efficiency  $\\epsilon_e$ of which little has been known so far."
    },
    "1004/1004.2217_arXiv.txt": {
        "abstract": "{Methanol has a rich rotational spectrum providing a large number of transitions at sub-millimetre wavelengths from a range of energy levels in one single telescope setting, thus making it a good tracer of physical conditions in star-forming regions. Furthermore, it is formed exclusively on grain surfaces and is therefore a clean tracer of surface chemistry.} {Determining the physical and chemical structure of low-mass, young stellar objects, in particular the abundance structure of CH$_3$OH, to investigate where and how CH$_3$OH forms and how it is eventually released back to the gas phase.} {Observations of the Serpens Molecular Core have been performed at the James Clerk Maxwell Telescope using the array receiver, Harp-B. Maps over a 4\\farcm5$\\times$5\\farcm4 region were made in a frequency window around 338 GHz, covering the 7$_K$--6$_K$ transitions of methanol. Data are compared with physical models of each source based on existing sub-millimetre continuum data.} {Methanol emission is extended over each source, following the column density of H$_2$ but showing up also particularly strongly around outflows. The rotational temperature is low, 15--20~K, and does not vary with position within each source. None of the Serpens Class 0 sources show the high-$K$ lines seen in several other Class 0 sources. The abundance is typically 10$^{-9}$ -- 10$^{-8}$ with respect to H$_2$ in the outer envelope, whereas ``jumps'' by factors of up to 10$^{2}$--10$^{3}$ inside the region where the dust temperature exceeds 100 K are not excluded. A factor of up to $\\sim$ 10$^3$ enhancement is seen in outflow gas, consistent with previous studies. In one object, SMM4, the ice abundance has been measured to be $\\sim$~3~$\\times$ 10$^{-5}$ with respect to H$_2$ in the outer envelope, i.e., a factor of 10$^3$ larger than the gas-phase abundance. Comparison with C$^{18}$O $J$=3--2 emission shows that strong CO depletion leads to a high gas-phase abundance of CH$_3$OH not just for the Serpens sources, but also for a larger sample of deeply embedded protostars.} {The observations illustrate the large-scale, low-level desorption of CH$_3$OH from dust grains, extending out to and beyond 7500~AU from each source, a scenario which is consistent with non-thermal \\mbox{(photo-)desorption} from the ice. The observations also illustrate the usefulness of CH$_3$OH as a tracer of energetic input in the form of outflows, where methanol is sputtered from the grain surfaces. Finally, the observations provide further evidence of CH$_3$OH formation through CO hydrogenation proceeding on grain surfaces in low-mass envelopes. } ",
        "introduction": "A long-standing goal of astrochemistry has been to determine the physical and chemical conditions prevailing in star-forming regions \\citep[e.g.,][]{vandishoeck98}. In this respect, different molecules act as tracers of different physical components, all depending on their formation history, their abundances, their chemical properties, etc. To effectively trace physical conditions such as density and temperature over the large range of values found in star-forming regions over the time-scale of star-formation, it is of great importance to have as many independent tracers as possible. Methanol (CH$_3$OH), with its rich rotational spectrum, is an excellent candidate tracing both temperature, density, grain surface formation and energy injection simultaneously during all phases of the early stages of stellar evolution.  CH$_3$OH is a slightly asymmetric top molecule with numerous rotational transitions observable at millimetre- and sub-millimetre wavelengths. Because of the large number of transitions observable in a single frequency window it is possible to obtain a coherent data set very efficiently, making methanol a very suitable tracer of physical conditions, in particular in low-mass star forming regions where the emission is optically thin. Moreover, since the molecule is a slightly asymmetric top molecule, it traces very efficiently both density and temperature \\citep[e.g.,][]{maret05, jorgensen05, leurini07}.  Methanol forms exclusively on ice-covered dust grain surfaces primarily through hydrogenation of CO \\citep{watanabe02, fuchs09}. Observations of interstellar ices show that methanol is indeed a prominent ice component, with abundances of up to almost 30\\% with respect to solid-state H$_2$O or a few $\\times$ 10$^{-5}$ with respect to gas-phase H$_2$ \\citep[e.g.,][]{dartois99, gibb04, pontoppidan04}. In contrast, pure gas phase reactions produce negligible CH$_3$OH abundances of less than 10$^{-10}$ \\citep{garrod06}. The question that naturally arises is how methanol desorbs from the surface of a dust grain to be observed in the gas phase. Is it through thermal heating of the entire grain, or is it through non-thermal desorption, where cosmic rays, UV-photons or exothermic reactions provide local heating of the grain? The former mechanism is at play close to young stellar objects (YSOs), in the inner-most part of the molecular envelope where the gas temperature exceeds 100~K \\citep{vandishoeck95, ceccarelli00, vandertak00, schoier02, maret05, jorgensen05} and in outflows where hot gas sputters the icy mantles \\citep[e.g.][]{bachiller95, bachiller97}. This mechanism allows for an effective methanol enrichment of the environment and abundances are typically in the range of 10$^{-7}$ to 10$^{-6}$. The non-thermal mechanism dominates in cold, dark clouds and the outer parts of molecular envelopes \\citep{hasegawa93, herbst06, garrod07, oberg09a}. Here, reported abundances typically have values of 10$^{-10}$ to 10$^{-8}$. Hence, methanol also acts as a tracer of energetic processes in star-forming regions.  So far most studies have concentrated on spectra at a single position or at most a few around them. Recently, large-scale mapping of weak molecular lines has become very efficient with the advent of array receivers such as the 16 pixel Harp-B receiver on the James Clerk Maxwell Telescope (JCMT). This allows for a direct study of the entire protostellar system (primarily envelope and outflow) on scales of several arcminutes at 15\\arcsec\\ resolution and for determination of density, temperature and energy input into the system. These observations will eventually be compared directly to observations of another important grain surface product, H$_2$O, to be done with the Herschel Space Observatory. By mapping the entire region, comparison with the different Herschel-beams (9\\arcsec--40\\arcsec) will be straight-forward.  The Serpens molecular core (also known as cluster A) is located at a distance of 230$\\pm$20 pc, following the discussion in \\citet{eiroa08}. The Serpens molecular core consists of several deeply embedded sources, of which at least four are identified as containing protostars \\citep{wolf-chase98, hogerheijde99}, SMM1, SMM3, SMM4 and S68N. Large-scale continuum-emission studies have been performed to quantify the spectral energy distribution of all sources in order to classify their evolutionary stage as well as the dust properties of the molecular envelopes surrounding each source \\citep[e.g.,][]{casali93, hurt96, testi98, davis99, larsson00, williams00}, most recently with the Spitzer Space Telescope as part of the Cores to Disks legacy program \\citep[c2d;][]{harvey07, evans09}. These studies show that three of the sources (SMM3, SMM4 and S68N) have relatively low bolometric luminosities of $\\sim$~5~$L_\\odot$ each, whereas SMM1 has a higher luminosity of $\\sim$~30~$L_\\odot$ \\citep[e.g.,][]{hogerheijde99, larsson00}. The mass of each system (envelope and star) is in all cases less than 10~$M_\\odot$. Recent interferometer observations by \\citet{choi09} show that SMM1 is a binary system with a projected separation of $\\sim$~500~AU. The binary SMM1b appears less embedded than the primary (SMM1a). This discovery has been refuted by \\citet{enoch09} and \\citet{vankempen09} who both resolve the disk around SMM1. Through detailed SED modelling, \\citet{enoch09} finds a very high disk mass of $\\sim$ 1 $M_\\odot$ and that the inner parts of the envelope have been cleared out to distances of 500 AU.  The region has been studied extensively at millimetre (mm) and sub-millimetre (sub-mm) wavelengths in numerous molecular transitions \\citep[e.g.,][]{mcmullin94, white95, wolf-chase98, hogerheijde99, mcmullin00}, however it was not included in the molecular surveys of isolated Class 0 and I objects in Perseus and Ophiucus \\citep{jorgensen04, jorgensen05, maret05}. The previous molecular studies conclude that SMM1, SMM3, SMM4 and S68N are all very similar to other young, low-mass stars with similar luminosities, such as IRAS16293-2422 and NGC1333 IRAS4A and 4B in terms of abundances of simple molecules that may be formed directly in the gas phase, e.g., HCO$^+$, CS, HCN \\citep{mcmullin94, hogerheijde99, mcmullin00}. In Serpens, little has been done to quantify excitation and gas-phase abundances of molecules predominantly formed on dust grain surfaces, like CH$_3$OH, even though several lines have been detected by \\citet{mcmullin94, mcmullin00} and \\citet{hogerheijde99}.  More direct observations of grain surface products have been made by \\citet{pontoppidan04}, who mapped infrared absorption by molecules in the ice over a region extending 40\\arcsec\\ south of SMM4, but still located well within the molecular envelope. The primary ice constituents were found to be H$_2$O, CO (0.4--0.9 with respect to H$_2$O), CH$_3$OH (0.28 with respect to H$_2$O) and CO$_2$ \\citep[0.3--0.5 with respect to H$_2$O;][]{pontoppidan08}. The solid-state CH$_3$OH abundance is one of the highest reported to date, both when compared to H$_2$O but also with respect to gas-phase H$_2$ (3$\\times$10$^{-5}$). At distances greater than 12\\,000~AU the CH$_3$OH-ice abundance drops beneath the detection limit, corresponding to 3$\\times$10$^{-6}$ with respect to H$_2$.  Besides the protostellar objects themselves, the region is permeated by large-scale outflows extending several arcminutes from the different sources with CO $J$=2--1 velocities ranging from $\\pm$10--15 km\\,s$^{-1}$ with respect to $\\varv_{\\rm LSR}=8-8.5$km\\,s$^{-1}$ \\citep[][Graves et al. subm.]{davis99}. \\citet{garay02} observed the outflows from SMM4 and S68N in CH$_3$OH 3$_K$--2$_K$ emission and inferred CH$_3$OH column densities of 1--2$\\times$10$^{15}$~cm$^{-2}$, corresponding to molecular abundance enhancements of $\\sim$~50--330, depending on outflow position, consistent with studies of other outflows \\citep[e.g.,][]{bachiller95, bachiller98}.  Here, the first map of rotationally excited methanol in the Serpens Molecular Core is presented of transitions which cover the energy range of $E_{\\rm up} \\sim 65-115$~K. The paper is structured as follows. In Sect. \\ref{sect:obs} the observations are presented, and in Sect. \\ref{sect:res} the observational results are provided along with radiative transfer modelling. Section \\ref{sect:dis} presents a discussion of the results, with a particular focus on the formation, desorption and excitation processes. Section \\ref{sect:conc} concludes the paper. ",
        "conclusions": "\\label{sect:dis}  \\begin{figure*} \\center \\includegraphics[width=1.\\columnwidth]{14182fg9a} \\includegraphics[width=0.5\\columnwidth]{14182fg9b} \\caption{Density and temperature profiles of the SMM1 and SMM4 envelopes as obtained from {\\sc Dusty} modelling. The radial density profile is displayed in blue (left axis) and kinetic dust temperature in red (right axis). The profiles are shown for the SMM1 envelope (full line) and the SMM4 envelope (dashed line). The CO freeze-out zone, characterised by temperatures lower than 25~K and densities greater than 10$^5$ cm$^{-3}$, is shown in dark gray, whereas the methanol evaporation zone ($T$ $>$ 80~K) is shown in light gray. To the right is a cartoon illustrating the differences between the envelopes surrounding SMM1 and SMM4. Both illustrations are to scale and the dark blue regions correspond to the CO freeze-out zones while red indicates CH$_3$OH evaporation zones. Note the much smaller CO freeze-out zone for SMM1.} \\label{fig:model} \\end{figure*}  The inferred abundances are all larger than what can be produced by pure gas-phase reactions \\citep[e.g.][]{garrod06}. Thus the observed CH$_3$OH gas must be formed on grains and subsequently have desorbed. In the following, possible scenarios for methanol desorption and excitation are discussed. Finally speculations on the nature of the individual young stellar objects in Serpens are made.  \\subsection{CH$_3$OH grain surface formation}  Methanol is formed on the surfaces of interstellar dust grains through hydrogenation of CO, a mechanism which has been studied in detail in the laboratory and in theoretical models of surface chemistry \\citep{hiraoka02, watanabe02, fuchs09}. To form large amounts of CH$_3$OH, it is imperative that CO freezes effectively out onto (water-ice covered) dust grains. Toward several pre-stellar cores \\citep[e.g., B68,][]{bergin02} and Class 0 objects \\citep[e.g.,][]{jorgensen05b} CO has been observed to freeze out very efficiently at temperatures lower than $\\sim$~20~K and densities greater than $\\sim$ 10$^5$ cm$^{-3}$. This ``catastrophic'' CO freeze-out has been observed directly through observations of CO ice abundances which show an increase in the amount of solid CO with respect to H$_2$O ice in the densest regions \\citep{pontoppidan06}. Through use of {\\sc Dusty} modelling discussed above it is possible to estimate the extent of the CO freeze-out zone in each of the Serpens Class 0 objects quantitatively, which is illustrated in Fig. \\ref{fig:model}. Here the predicted density and temperature profiles are shown as a function of distance from the protostar itself. The density profiles of all four envelopes are very similar, however SMM1 stands out in terms of temperature profile. Due to the higher luminosity (30~$L_\\odot$) of SMM1 with respect to the other source luminosities ($\\sim$~5~$L_\\odot$), the envelope temperature is also higher throughout.  The zone over which CO freezes out in SMM1 is significantly smaller than that of the other sources, SMM3, SMM4 and S68N. In the case of SMM1, the present freeze-out zone extends from $\\sim$ 1000--5000 AU whereas the other envelopes have freeze-out zones starting at a few hundred AU and extend out to the same distance as for SMM1 (see Fig. \\ref{fig:model}). This indicates that at present CO is not freezing out efficiently in the SMM1 envelope, so any efficient methanol formation has probably ceased at this point in time. The colder envelopes surrounding SMM3, SMM4 and S68N could still be forming CH$_3$OH.  In terms of observing methanol directly in the ice itself, \\citet{pontoppidan04} studied a small region extending south of SMM4 from 4000 AU to 12\\,000 AU by observing the 3.54~$\\mu$m CH$_3$OH features in absorption against background and embedded stars. The extent is illustrated by the white line in Fig. \\ref{fig:find}. In this region the CH$_3$OH-ice abundance is constant at 28\\% with respect to water ice or 3$\\times 10^{-5}$ with respect to gas-phase H$_2$. This corroborates the interpretation presented here, that the gas phase abundance of CH$_3$OH is low and constant out to 12\\,000 AU in the SMM4 envelope. Beyond this line the CH$_3$OH-ice abundance drops by at least an order of magnitude and \\citeauthor{pontoppidan04} could only determine upper limits. It is interesting to note that the location where the ice abundance drops is where one of the outflow knots starts (SMM4-S; see Fig. \\ref{fig:find}). Thus the reason for the drop is a combination of the envelope being more tenuous ($n_{\\rm H}<$ 10$^4$ cm$^{-3}$) far from the protostar, so that CO does not freeze out very efficiently, while at the same time whatever methanol is in the ice is sputtered into the gas phase by the outflow.  \\citet{cuppen09} have recently studied the formation of CH$_3$OH on ice surfaces using a Monte Carlo method, in which the gas-grain chemistry based on laboratory data is simulated microscopically over long time-scales. The limiting factors in producing methanol is the availability of both CO and H on the grain surfaces, and so results show that CO hydrogenates efficiently to form CH$_3$OH, especially at temperatures lower than 12 K where atomic hydrogen can be retained efficiently. After 10$^{5}$ years CH$_3$OH may form up to 100 individual mono-layers on the grain, comparable to that of water ice. Hence, the ice abundance of methanol in the outer parts of the envelope will depend strongly on the temperature. Because the envelopes of SMM3, SMM4 and S68N are significantly colder than that of SMM1, they would still be actively producing methanol, which could explain the higher gas phase abundances (see Table \\ref{tab:abun}).  Observations show a different behaviour of C$^{18}$O at the center position of S68N compared with the center position of SMM1. Figure \\ref{fig:co_ch3oh} presents C$^{18}$O $J$=3--2 emission in colour and methanol emission from the 7$_0$--6$_0$ A$^+$ line overlaid as contours. The morphology of the C$^{18}$O emission shows a peak toward SMM1, as expected for a warm envelope, and a ring of emission toward S68N. The C$^{18}$O abundances towards the two central positions are measured to be 3 and 1$\\times$10$^{-8}$ respectively, based on integrated C$^{18}$O intensities of 4.3 and 2.4 K\\,km\\,s$^{-1}$ and the gas column densities derived from dust continuum emission (Table \\ref{tab:abun}). This is to be compared to a standard C$^{18}$O abundance of 1.8$\\times$10$^{-7}$ for a normal abundance of CO/H$_2$ = 10$^{-4}$ and an $^{16}$O/$^{18}$O ratio of 550 \\citep{wilson94}. Thus, CO is depleted by a factor of $\\sim$ 6 and 18 for the two sources as averaged over the entire envelope. For these two sources there is a clear anti-correlation between CO and CH$_3$OH gas phase abundances.   \\begin{figure} \\center\\includegraphics[width=0.8\\columnwidth]{14182fg10} \\caption{Integrated C$^{18}$O, $J$=3--2 emission ($\\int T_{\\rm mb} {\\rm d}\\varv$) is shown in colour, with integrated CH$_3$OH 7$_0$--6$_0$ A$^+$ emission overlaid as contours. The image is centered on Serpens SMM1 with S68N located to the north. Outflow directions are indicated by dashed (red-shifted) and dotted (blue-shifted) lines.} \\label{fig:co_ch3oh} \\end{figure}   \\begin{figure} \\center\\includegraphics[width=\\columnwidth]{14182fg11} \\caption{CH$_3$OH gas abundance versus CO gas abundance for a selection of Class 0 and I sources based on data and analysis from \\citet{schoier02, maret05, jorgensen05} and this work. The Serpens sources are in red and IRAS16293-2422 and NGC1333-IRAS2A are in blue. Sources for which only upper methanol limits exist are marked by arrows. The two dashed lines indicate the 1$\\sigma$ confidence limits of the power-law fit.} \\label{fig:c18o_ch3oh} \\end{figure}   To examine whether this anti-correlation is unique to Serpens or a general feature of embedded sources, the outer-envelope CH$_3$OH abundances of the sample of \\citet{maret05} and \\citet{jorgensen05} was coupled with CO abundances for the same sources from \\citet{jorgensen02} (see Fig. \\ref{fig:c18o_ch3oh}). The CO gas phase abundance averaged over the extent of the envelope is taken as a tracer of CO depletion in the sense that the total gas and ice CO abundance is assumed to be constant. For sources where the presence of a jump zone can reproduce the observed CH$_3$OH emission, only the outer abundance is plotted here. A typical uncertainty of 30\\% is assumed both for the CO and CH$_3$OH abundances. Except for three sources (L1157, L1448-I2 and L483) there is a correlation between the two abundances as illustrated by the value of the Pearson correlation coefficient of $r$ = 0.70. The Serpens sources are characterized by the combination of low CO-abundance and high CH$_3$OH abundance. This result implies that the CH$_3$OH gas phase abundance is directly related to the current production of CH$_3$OH in the outer, cold parts of the envelope, and that any difference from source to source is due to a difference in the amount of CO frozen out onto the grains. Thus, not all solid CH$_3$OH is formed during the cold pre-stellar core phase, consistent with the lack of detected CH$_3$OH ice toward background stars behind quiescent dense clouds. This also implies that the lack of CH$_3$OH emission at normal cloud positions between the sources is because the density is lower, and hence the timescale for CO to freeze out is much higher, i.e., the absolute CH$_3$OH abundance will drop.   \\subsection{CH$_3$OH desorption mechanism}  Once methanol has formed on a grain surface it can desorb according to two different mechanisms, thermal and non-thermal desorption. In the first mechanism the entire grain is heated thermally (macroscopic grain heating), and the icy mantle evaporates entirely, releasing all adsorbed species into the gas phase. The ice mantles typically evaporate at grain temperatures of 30--100 K, depending on species. The other mechanism is non-thermal desorption, in which the ice mantle is ``heated'' on a microscopic (local) scale, either due to absorption of a single UV-photon \\citep{oberg09b}, impact of a cosmic ray particle \\citep{leger85, hasegawa93, herbst06}, or the binding energy being released from the formation of a new molecule \\citep{garrod07} or sputtering in outflows \\citep{jimenezserra08}.  \\subsubsection{Thermal desorption}  In the case of methanol, the thermal evaporation temperature has been determined experimentally to 80--100 K \\citep{brown07, green09} as is also indicated in Fig. \\ref{fig:model}. The lower temperature corresponds to evaporation of a pure CH$_3$OH ice and the higher to a mix of CH$_3$OH and H$_2$O. Because CO freezes out on top of a water ice, CH$_3$OH is expected to be present in the same layer. Thus the desorption temperature should be close to that of a pure CH$_3$OH ice ($\\sim$ 85 K), which is slightly lower, but not much, than that of CH$_3$OH mixed with H$_2$O. This is also found to be the case experimentally \\citep{bisschop07}. Thus, for the thermal-desorption mechanism to be active, it is necessary that all grains are heated to greater than 80 K, i.e., close to the protostar itself. From the physical structure models above it is predicted that the radius at which $T_{\\rm dust} > 80$ K is of the order of 50--100 AU or 0\\farcs2--0\\farcs4 for $d$=230 pc, depending on source, leading to a beam-dilution factor of 1500--5000. Methanol emission originating from very close to the protostar itself has previously been observed for two low-mass sources, IRAS16293-2422 \\citep[$d$=125 pc;][]{vandishoeck95, schoier02} and NGC1333 IRAS2A \\citep[$d$=250 pc;][]{maret05}. Even though the regions have not been spatially resolved with single-dish data, it has been possible to infer the existence of ``hot cores'' based on observations of high-$K$ lines ($K >$ 3) and measured rotational temperatures of $>$ 80 K. This indicates that the emission arises in a compact, warm and dense region, consistent with it originating close to the protostar itself and also consistent with high-spatial resolution interferometer data \\citep[e.g.,][]{jorgensen05c, jorgensen07}. For the sources presented here, it has only been possible to provide an upper limit on the inner abundance, which in most cases is a few $\\times$10$^{-7}$ with respect to H$_2$. One notable exception is SMM1, where the upper limit on the inner abundance is close to the outer abundance, a result that will be discussed further below (Sect. \\ref{sect:hotcore}).  \\subsubsection{Non-thermal desorption}  Since the temperature in the outer envelope is significantly lower than 100 K, the primary desorption mechanism must be non-thermal. This is supported by the fact that the abundance only follows the column density, and does not appear to depend on temperature. No attempt is made to distinguish between the above mentioned mechanisms here. If desorption is induced by secondary UV photons from cosmic ray ionization of H$_2$ the gas phase abundance is expected to be 10$^{-4}$--10$^{-3}$ times the ice abundance \\citep{oberg09a, oberg09c}. In the case of SMM4 where the CH$_3$OH-ice abundance is $\\sim$ 3$\\times$10$^{-5}$ with respect to H$_2$ \\citep{pontoppidan04} the UV-induced desorption alone can account for a gas phase abundance of 10$^{-9}$--10$^{-8}$. This is remarkably close to the observed gas phase abundance of $\\sim$ 2$\\times$10$^{-8}$.  \\citet{garrod07} modelled the release of CH$_3$OH from the grain surface into the gas phase by examining whether the release of binding energy would be enough to evaporate the molecule into the gas phase. They found that the fraction was probably in the range of 1--10\\%, but could not pin it down any further. In the case of SMM4, results presented here combined with ice observations indicate that much less than 1--10\\% is desorbed, since the gas/ice abundance is $\\sim$ 10$^{-4}$. Results presented in \\citet{hasegawa93} show that direct cosmic ray desorption is not an efficient desorption mechanism for tightly bound species like H$_2$O and CH$_3$OH when compared to the accretion timescale. In fact, they find that the difference between the two rates is of the order of 10$^4$. Previous studies show that direct cosmic ray induced desorption is most efficient for volatile species \\citep[see, e.g.,][]{shen04,roberts07} and for species adsorbed on very small grains \\citep[$r \\leq 0.05~\\mu$m;][]{herbst06}.  \\subsubsection{Shock-induced desorption}\\label{sect:shock_des}  In outflows the main desorption mechanism is sputtering of the grain mantle, i.e., impacts of high-temperature gas molecules and atoms, primarily H$_2$ and He. This can in principle be compared to the direct cosmic ray desorption mechanism, except that the flux of H$_2$ and He is orders of magnitude higher and the impact energy is much lower. The efficiency of this desorption mechanism is clearly seen from the fifth column of Table \\ref{tab:abun_out}, where the CH$_3$OH abundance in the outflow positions is compared to the ambient envelope abundance. In particular, the SMM4-S methanol outflow abundance is $\\sim$ 10$^{-5}$ with respect to H$_2$, i.e., only a factor of 3 lower than the measured CH$_3$OH-ice abundance in a region located $\\sim$ 15$''$ (4000 AU) to the north \\citep[see Fig. \\ref{fig:find} and][]{pontoppidan04}. If the envelope ice-abundance in this region were also 3$\\times$10$^{-5}$ prior to the impact of the shock, this would indicate that one third or more of the CH$_3$OH ice is sputtered from the grain mantle. However, the measured gas-phase abundance as determined from $^{12}$CO emission is an upper limit as briefly discussed above. The true value of the abundance probably lies in the range of 10$^{-8}$ to 10$^{-6}$ where the lower limit is obtained from C$^{18}$O emission.    \\subsection{The nature of the embedded YSOs in Serpens}\\label{sect:hotcore}  Two low-mass YSOs have been confirmed observationally to have the same chemical characteristics as high-mass hot cores, IRAS16293-2422 and NGC1333-IRAS2A. These two sources are at the same evolutionary stage as SMM1 based on classical indicators, e.g., $T_{\\rm bol}$. The characteristics are (1) a large number of detected, saturated, complex organic molecules and (2) a high rotational temperature for methanol ($>$ 80 K). At the same time, neither of these two envelopes show signs of extended methanol emission \\citep[][Kristensen et al. in prep.]{vandishoeck95}. IRAS16293-2422 is a close binary system surrounded by a circum-binary envelope. The inner parts of this envelope are passively heated to temperatures $>$ 80 K, which, along with small-scale shocks impinging on the inner wall of the envelope \\citep{chandler05}, give rise to the observed hot-core signature. NGC1333-IRAS2 is also a binary (possibly triple) system, but the projected separation is greater ($\\sim$ 30\\arcsec). It is believed that the same mechanisms are at play in NGC1333-IRAS2A causing the hot-core signatures \\citep{maret05, jorgensen05}. In the following the Serpens sources, and in particular SMM1, will be compared to these two low-mass hot cores, and differences will be discussed.  There are several possible explanations for the lack of ``hot core'' characteristics in the Serpens sources which may be categorized in the following manner: \\begin{enumerate} \\item Physical: No gas is present close to the source, or no hot gas is present. \\item Chemical: Hot gas is present, but the methanol abundance is very low. \\item Observational: Warm methanol is present close to the source, but either the extent is very small, or emission is optically thick. \\end{enumerate}  \\citet{choi09} recently reported that SMM1 is a close binary with a projected separation of $\\sim$ 500 AU. The primary would be the sub-mm source associated with SMM1 (SMM1a), while the binary (SMM1b) would be less embedded. SMM1b can be associated with a source observed by {\\it Spitzer} at wavelengths shorter than 24 $\\mu$m, where SMM1a is not detected. If the projected distance is similar to the actual separation between the two sources, then this could have cleared the inner part of the envelope of gas, explaining the absence of hot gas close to the (sub-mm) source. However, both IRAS16293-2422 and NGC1333-IRAS2A have been shown to have inner holes or cavities similar in size to that of SMM1 \\citep{schoier02, jorgensen05b}, therefore this cannot be the entire explanation.  Recent millimeter interferometry of SMM1 indicates that the region close to the protostar has been cleared of gas. Both \\citet{enoch09} and \\citet{vankempen09} report the detection of a resolved disk surrounding SMM1. The disk is unusually massive ($\\sim$ 1 $M_\\odot$) and has a modelled radius of up to 300 AU \\citep{enoch09}, while the inner 500 AU of the envelope have been cleared of gas. In such a dense disk, only the upper-most layers will be warm or hot, depending on the distance to the protostar, as the dust extinction grows rapidly towards the midplane of the disk. Thus, the column density of hot methanol will be very low. Through SED modeling of the disk, \\citet{enoch09} estimate that the inclination is 15\\degr, i.e., nearly face-on. This is in conflict with observations of the 3.6-cm radio jet, which indicate that the outflow is moving very close to the plane of the sky \\citep{rodriguez89, moscadelli06}. Moreover, the $^{12}$CO line profiles are very symmetric along the large-scale outflow, i.e., the molecular outflows are also moving close to the plane of the sky. This would indicate that the disk is seen close to edge-on. If so, then the beam-dilution will be very high implying that any hot part will not be detectable in our single-dish beam. The disk surrounding NGC1333-IRAS2A is less massive compared to SMM1 \\citep[$\\sim$ 0.056 $M_\\odot$;][]{jorgensen09} and is viewed closer to face-on resulting in less extinction and lower beam dilution. This implies that the high-$K$ emission observed in IRAS16293-2422 and NGC1333-IRAS2A is originating in the disk close to the protostar. The actual heating mechanism (passive heating or small-scale shocks) cannot be distinguished with the current observations.  If the inner part of the envelope has not been cleared out of warm gas or if the disk is seen face-on, then it is possible that the abundance of methanol is low. As has been shown by current observations, the CH$_3$OH gas-phase abundance is comparatively high in the outer parts of the envelopes, so to decrease the abundance in the inner part of the system, destruction of CH$_3$OH must be present. This destruction mechanism must be very efficient if it is to destroy all CH$_3$OH in the hot-core parts of the envelope, where the abundance is expected to rise to $>$ 10$^{-6}$ w.r.t. H$_2$. CH$_3$OH can be destroyed in the gas phase through direct reactions with other species, however the rate coefficients are typically low. An alternative mechanism of both the gas and the ice destruction is UV-dissociation. However, this destruction mechanism must also be at play in IRAS16293-2422 and NGC1333-IRAS2A and there is currently no reason why UV-photodissociation would be more efficient in Serpens sources than in the other two nor that the UV-field is enhanced in SMM1.  Finally, a low-mass hot core may be present in all of the Serpens sources, but not observable. This can be due to beam dilution or due to optical depth effects. The region over which methanol desorbs from the grain surface is expected to be $\\sim$ 100 AU at most, and so beam dilution would be of the order of 10$^3$. However, the same beam dilution would apply to at least NGC1333-IRAS2A, which is located at a similar distance of 250 pc, and can thus not be used as an argument. The line optical depth has been calculated in the {\\sc Ratran} simulations, and is typically of the order of 0.1 or less for the transitions observed here, even for lines arising in the inner-most part of the envelope. The dust opacity at 338 GHz is less than 0.08 at all times, as estimated from the {\\sc Dusty} modelling.  Of the three explanations presented above the first is the more plausible if the distance is indeed 230 pc as assumed. Recent VLBA observations indicate that the distance may be closer to 415 pc \\citep{dzib10}, in which case the third explanation is more plausible. The other two reasons can be disproved through comparison with IRAS16293-2422 and NGC1333-IRAS2A. However, only SMM1 has been suggested to have a massive disk, and it is also the source that resembles the two low-mass hot cores the most in terms of luminosity. The other sources (SMM3, SMM4 and S68N) have considerably lower luminosities by factors 4--10. Thus it may be that the hot-core regions around these sources are indeed much smaller and beam-diluted ($\\sim$ 5$\\times$ 10$^3$), comparable to several of the low-mass sources in the sample of \\citet{maret05} and \\citet{jorgensen05}. With the current observations only upper limits have been determined of the molecular abundance in the inner envelope of the lower-luminosity sources."
    },
    "1004/1004.0738_arXiv.txt": {
        "abstract": "{ The PLATO satellite mission project is a next generation ESA Cosmic Vision satellite project dedicated to the detection of exo-planets and to asteroseismology of their host-stars using ultra-high precision photometry. The main goal of the PLATO mission is to provide a full statistical analysis of exo-planetary systems around stars that are bright and close enough for detailed follow-up studies. Many aspects concerning the design trade-off of a space-based instrument and its performance can best be tackled through realistic simulations of the expected observations. The complex interplay of various noise sources in the course of the observations made such simulations an indispensable part of the assessment study of the PLATO Payload Consortium. We created an end-to-end CCD simulation software-tool, dubbed PLATOSim, which simulates photometric time-series of CCD images by including realistic models of the CCD and its electronics, the telescope optics, the stellar field, the pointing uncertainty of the satellite (or Attitude Control System [ACS] jitter), and all important natural noise sources. The main questions that were addressed with this simulator were the noise properties of different photometric algorithms, the selection of the optical design, the allowable jitter amplitude, and the expected noise budget of light-curves as a function of the stellar magnitude for different parameter conditions. The results of our simulations showed that the proposed multi-telescope concept of PLATO can fulfil the defined scientific goal of measuring more than 20000 cool dwarfs brighter than m$_V$=11 with a precision better than 27 ppm/h which is essential for the study of earth-like exo-planetary systems using the transit method. } ",
        "introduction": "The main objective of the PLATO mission is to study exo-planetary systems around bright stars using ultra-high precision photometry to facilitate ground-based follow-up studies. Asteroseismology of the host-stars will be the key to pinpoint characteristics such as the age of a system or the mass and density of the detected planets. A detailed description of the PLATO mission can be found in Catala (2009) and Claudi (2010).  The motivation to go to space to acquire photometry of stars for asteroseismology or detecting exo-planets is mainly the lack of atmospheric disturbances and the diurnal cycle. PLATO will be situated in a Lissajous orbit around the Lagrange point L2 which will lead to a minimization of external influences from Earth. A single field in the sky can be monitored for months to years with a very high duty cycle and with the same instrument, which leads to a very homogeneous data set, much lower photometric noise levels, and the avoidance of daily aliasing in power spectra. Nevertheless, still many noise sources remain and must be quantified in detailed studies to estimate their impact on the quality of the observations. The main unavoidable natural noise source is photon shot noise, whose relative effect on data can be minimized by collecting more photons. The instrumental noise sources which are due to the technical construction of the spacecraft and due to the limits of the electronics must be minimized to ensure that photon noise remains the major contributor of the overall noise budget. Due to the complex interplay of the noise sources, numerical simulations are ideally suited to study the overall performance of space-based observations.  PLATOSim is a software tool dedicated to the realistic simulation of CCD photometry of stars in the optical range acquired by a space-based instrument. It is designed in a way to account for the most important natural and instrumental noise sources that have an impact on the performance of the instrument. Due to the need of such a simulator for the assessment study for PLATO Mission, we created a homogeneous software tool based on pre-existing codes that had been developed for the Eddington and MONS space missions (Arentoft et al. 2004; De Ridder et al. 2006).\\\\ PLATOSim has been coded in C++ and features a graphical user interface (see Fig.~\\ref{fig:screenshot}) based on Qt4. The contributions of the different noise sources are described through mathematical models and use a large number of parametrised inputs and numerical pre-computed input data like PSFs, jitter time-series, and sky background brightness.  Several improvements have been applied to the simulator compared to the pre-existing version. These concern mainly a significant increase in computational speed, a more realistic treatment of jitter, and the inclusion of two different photometric algorithms and statistical tools to estimate the noise properties of the data. The graphical interface permits in an easy way to test several different instrument set-ups by modifying the parameters for the simulations. \\begin{figure} \\includegraphics[width=82mm,clip,angle=0]{screenshot1.eps} \\caption{Screenshot of PLATOSim. On the left panels, parameters for the simulations can be set. The right panel shows a simulated image which contains several visible noise features such as saturation smearing (bright vertical spikes of stars) and frame-transfer trails (vertical lines).} \\label{fig:screenshot} \\end{figure} ",
        "conclusions": ""
    },
    "1004/1004.4563_arXiv.txt": {
        "abstract": "We study time evolutions of superfluid neutron stars, focussing on the nature of the oscillation spectrum, the effect of mutual friction force on the oscillations and the hydrodynamical spin-up phase of pulsar glitches. We linearise the dynamical equations of a Newtonian two-fluid model for rapidly rotating backgrounds. In the axisymmetric equilibrium configurations, the two fluid components corotate and are in $\\beta$-equilibrium. We use analytical equations of state that generate stratified and non-stratified stellar models, which enable us to study the coupling between the dynamical degrees of freedom of the system. By means of time evolutions of the linearised dynamical equations, we determine the spectrum of axisymmetric and non-axisymmetric oscillation modes, accounting for the contribution of the gravitational potential perturbations, i.e. without adopting the Cowling approximation. We study the mutual friction damping of the superfluid oscillations and consider the effects of the non-dissipative part of the mutual friction force on the mode frequencies. We also provide technical details and relevant tests for the hydrodynamical model of pulsar glitches discussed by \\cite*{2010MNRAS.tmp..554S}. In particular, we describe the method used to generate the initial data that mimic the pre-glitch state, and derive the equations that are used to extract the gravitational-wave signal. ",
        "introduction": "\\label{sec:Int}   Mature neutron stars are expected to have superfluid and superconducting components in their interior. Shortly after a neutron star's birth the temperature decreases below $T\\simeq10^{9}~\\rm{K}$, at which point superfluid neutrons should be present both in the inner crust and the outer core, while the core protons should form a superconductor. At all relevant temperatures, the electrons form a ``normal'' fluid that is tightly locked to the protons due to the electromagnetic interaction. This suggests that the dynamics of mature neutron stars depends on the detailed interaction between coupled superfluids-superconductors~\\citep*{2010arXiv1001.4046G}, i.e. represents a complex physics problem. The situation is not expected to simplify if one also accounts for the inner neutron star core, at several times the nuclear saturation density, where exotic states like hyperon superfluid mixtures or deconfined quark condensates may be present.  Although it is generally appreciated that neutron stars have this very complicated structure, the evidence for the presence of the different superfluid phases remain indirect. The strongest support comes from observed pulsar glitches, rapid spin-up events seen in a number of young pulsars (and also some magnetars) during their magnetic slow-down phase. The typical glitch size is very small, representing a relative change ($\\Delta \\Omega$) in the observed rotation rate ($\\Omega$) in the range $10^{-9} < \\Delta \\Omega / \\Omega < 10^{-5}$. The currently accepted model for these events relies on the transfer of angular momentum between a (faster spinning) superfluid neutron component and the star's (slower spinning) elastic crust (to which the magnetic field is anchored).  The exchange is thought to be mediated by neutron vortices (by means of which the superfluid mimics bulk rotation) and the associated mutual friction~\\citep{ALS84}.  A challenge for future observations is to probe the detailed physics of a neutron star's interior. In this context, asteroseismology associated with either gravitational or electromagnetic signals seems particularly promising. In fact, the quasiperiodic oscillations seen in the tails of giant magnetar flares may have provided us with the first opportunity to test our theoretical models against observational data~\\citep[see for instance][and references therein]{2007Ap&SS.308..625W}. The observed variability likely originates from crustal oscillations and depends on the detailed crust dynamics and the interaction with the neutron star's magnetic field. These observations have led to a resurgence of interest in neutron-star seismology and a renewed assault on the problem of magnetic star oscillations, a seriously challenging problem from the theory point-of-view~\\citep*[see][for a discussion of the literature]{2009MNRAS.396.1441C}.  In the context of the present paper, the potential relevance of the neutron superfluid that penetrates the neutron star crust is particularly relevant~\\citep*{2009MNRAS.396..894A,2009CQGra..26o5016S}.  The prospect of detecting gravitational waves from oscillating neutron stars is also exciting, especially since the associated signals will allow us to probe the high-density region and hence the supranuclear equation of state (EoS)~\\citep*{1998MNRAS.299.1059A,benhar-2004-70,2007MNRAS.374..256S, 2009arXiv0912.0384A}.  In order to faciliate future observations and the decoding of collected data, we need to improve our models considerably. The superfluid aspects are particularly interesting in this respect, since the oscillation spectrum of a superfluid star is more complex than that of a single fluid model. In superfluid regions fluid elements can execute both co- and counter-moving motion, leading to the existence of unique ``superfluid'' oscillation modes. Our understanding of the nature of the additional degree(s) of freedom and the effect on observables must be improved by detailed modelling, ultimately in the context of general relativistic multi-fluid dynamics.  The present work presents recent progress towards this goal.  We study the oscillations of superfluid neutron stars by evolving in time the linearized two-fluid equations in Newtonian gravity. We improve on the analysis of~\\cite{2009MNRAS.396..951P} by including the perturbations of the gravitational potential. We also account for the mutual friction force associated with vortices, and implement quadrupole extraction of the gravitational-wave signal associated with the fluid motion.  We provide the detailed analysis (and relevant code tests) for the configurations that we recently used to study the hydrodynamics of pulsar glitches~\\citep{2010MNRAS.tmp..554S}.  We consider two simple analytical EoS and construct two distinct sequences of rapidly rotating stars, the main difference being the presence or absence of composition gradients.  Such gradients impact on the superfluid dynamics, as the co- and counter-moving degrees of freedom are coupled in stratified models.  From time-evolutions of the relevant perturbation equations, with the gravitational potential perturbation included, we determine the axi- and non-axisymmetric oscillation modes for models that rotate up to the mass shedding limit. Finally, we account for the (standard form of the) mutual friction force. This adds two coupling terms to the equations of motion.  One component is dissipative and damps an oscillation mode, while the other modifies the frequencies of the superfluid modes. We study both these effects and infer an analytical relation for the associated frequency change of the non-axisymmetric superfluid fundamental and inertial modes. ",
        "conclusions": "} \\label{sec:concl}  We have studied the dynamics of superfluid rotating neutron stars, focussing on the nature of the oscillation spectrum, the effects of the mutual friction force on the oscillations and the hydrodynamic spin-up phase of pulsar ``glitches''. Adopting the Newtonian two-fluid model, we evolved in time the perturbed dynamical equations on axisymmetric equilibrium configurations. This approach allows us to derive the spectrum of axisymmetric and non-axisymmetric oscillation modes of stellar models that rotate up to the mass shedding limit. In this work, we have improved on previous studies by including the gravitational perturbation and the mutual friction force. The spectrum is then determined with a better accuracy, as we no longer use the Cowling approximation~\\citep{2009MNRAS.396..951P}. From the computational point of view, we have to solve the perturbed Poisson equation together with the linearised momentum and mass conservation equations. We have numerically evolved the hyperbolic equations with a Mac-Cormack algorithm, while the elliptic equation for the gravitational potential is solved at each time step with a spectral method.  In our current model the rotating background models are pure fluid, i.e. without an elastic crust region, neutrons and protons corotate and are in $\\beta$-equilibrium. In superfluid stars, the co- and counter-phase motion of the two fluid constituents can be coupled by composition gradients and this influences the dynamics. In order to consider this effect we have studied two simple polytropic equations of state that generate distinct sequences of stratified and non-stratified rotating stars. These background models are simplistic, and we must improve on this aspect if we want to decode the complexity of astrophysical observations. Certainly, we must add an elastic crust to the model and relax the co-rotation assumption between the two fluids. If we want to use more realistic equations of state we also need to translate the model to General Relativity. We are currently working on all these issues.  In neutron stars, the mutual friction force may have both dissipative and non-dissipative effects. The dissipative part of the force, which is dominant in the weak drag regime, mainly damps an oscillation mode. Meanwhile, the non-dissipative term dominates in the strong drag regime, essentially modifying the oscillation spectrum.  We have studied the two drag regimes and showed that our numerical code effectively reproduces the mutual friction damping of the two-fluid relative motion. For non-stratified stars, the co- and counter-moving degrees of freedom are uncoupled and only the superfluid modes are damped. When the stellar model is stratified, the damping affects also the ordinary modes. The accuracy of our numerical code has also been tested in~\\cite{2010MNRAS.tmp..554S}, where determined the glitch spin-up time and compared it to a simple analytic formula. However, we are not yet able to extract (with useful precision) the mutual friction damping time of individual oscillation modes from our numerical evolutions. More work is needed to establish to what extent one should expect to do this within our computational framework. For the strong drag regime, we have studied the effect of the mutual friction and composition variation on the rotational splitting of the superfluid $l=m=2$ f-mode and on the frequencies of the $l=m=2$ superfluid r-mode. The main effect is a change of propagation direction of the modes with respect to the background rotation. A mode that is pro-grade (retro-grade) in the weak drag regime may become retro-grade (pro-grade) in the strong drag regime. We have determined the numerical frequencies of the f- and r-modes for the rotating sequence of non-stratified stellar models and provided simple empirical expressions based on the numerical data. For constant mutual friction parameters, the non-axisymmetric splitting of the superfluid f-mode and the r-mode frequencies depends on the inverse of the proton fraction.  Finally, we provided relevant technical details for the hydrodynamical models for pulsar glitches discussed by~\\cite{2010MNRAS.tmp..554S}. The initial conditions for the glitch evolutions describe two fluids that rotate with a small velocity lag. These configurations were been determined using a perturbative approach first introduced by~\\cite{2004MNRAS.347..575Y}. We extended this method to implement different EoS and consider non-corotating initial configurations that conserve the mass of each fluid constituent. Moreover, we derived the detailed quadrupole gravitational extraction formulae for $l=2$ oscillation modes of a superfluid star. We determined the perturbative expressions for the momentum and stress formulae that can be used to improve the numerical extraction of the gravitational-wave signal (reducing the order of the time derivative of the standard quadrupole formula). We determined the gravitational-wave strain for the two independent initial glitch configurations that are obtained with the~\\cite{2004MNRAS.347..575Y} approach. For a given background rotation, these results can be used to estimate the gravitational signal for any glitch size. Furthermore, we have showed the effect of the Cowling approximation on the glitch gravitational-wave strain and the oscillation spectrum.  With the progress described in this paper, our programme of studying superfluid neutron star dynamics by time-evolutions of the linearised equations has reached the point where we need to add key physics to the model.  The natural step would be to account for the elastic neutron star crust with the expected interpenetrating neutron superfluid. This requires us to change the computational framework somewhat, as it is natural to discuss the elasticity in term of Lagrangian perturbation theory. Moreover, we need to address various issues associated with vortex pinning by the crust nuclei. This problem requires additional force contributions at the level of individual vortices, and we need to develop a suitable smooth-averaged hydrodynamics description if we want to make progress. We are currently working on both these issues.  It would also be relevant to extend our models to general relativity.  This is essential if we want to be able to use realistic supranuclear equations of state.  As long as we make use of the relativistic analogue of the Cowling approximation this generalisation should be straightforward, but if we want to account for the dynamics of spacetime the problem becomes much more involved.  If we want to consider realistically ``layered'' neutron stars we also need to improve our understanding of the different phase-transitions, e.g. in the vicinity of the critical density/temperature for the onset of superfluidity, and how these regions affect the large scale dynamics. We face a number of challenging questions, but there is no reason why we should not be able to resolve the relevant issues and progress towards the construction of realistic dynamical neutron star models."
    },
    "1004/1004.1493_arXiv.txt": {
        "abstract": "Constantly accumulating observational data continue to confirm that about 70 \\% of the energy density today consists of dark energy responsible for the accelerated expansion of the Universe. We present recent observational bounds on dark energy constrained by the type Ia supernovae, cosmic microwave background, and baryon acoustic oscillations. We review a number of theoretical approaches that have been adopted so far to explain the origin of dark energy. This includes the cosmological constant, modified matter models (such as quintessence, k-essence, coupled dark energy, unified models of dark energy and dark matter), modified gravity models (such as $f(R)$ gravity, scalar-tensor theories, braneworlds), and inhomogeneous models. We also discuss observational and experimental constraints on those models and clarify which models are favored or ruled out in current observations. ",
        "introduction": "The discovery of the late-time cosmic acceleration reported in 1998 \\cite{Riess,Perlmutter} based on the type Ia Supernovae (SN Ia) observations opened up a new field of research in cosmology. The source for this acceleration, dubbed dark energy \\cite{Huterer}, has been still a mystery in spite of tremendous efforts to understand its origin over the last decade \\cite{Sahnireview,Carroll,Peebles,Paddy,SahniLecture,review,Durrerreview,Caldwell09,Tsujibook}. Dark energy is distinguished from ordinary matter in that it has a negative pressure whose equation of state $w_{\\rm DE}$ is close to $-1$. Independent observational data such as SN Ia \\cite{SNLS,Gold1,Gold2,Essence1,Essence2,Kowalski}, Cosmic Microwave Background (CMB) \\cite{WMAP1,WMAP3,WMAP5,WMAP7}, and Baryon acoustic oscillations (BAO) \\cite{Eisenstein,Percival1,Percival2} have continued to confirm that about 70 \\% of the energy density of the present Universe consists of dark energy.  The simplest candidate for dark energy is the so-called cosmological constant $\\Lambda$ whose equation of state is $w_{\\rm DE}=-1$. If the cosmological constant originates from a vacuum energy of particle physics, its energy scale is significantly larger than the dark energy density today \\cite{Weinberg} ($\\rho_{\\rm DE}^{(0)} \\simeq 10^{-47}$~GeV$^4$). Hence we need to find a mechanism to obtain the tiny value of $\\Lambda$ consistent with observations. A lot of efforts have been made in this direction under the framework of particle physics. For example, the recent development of string theory shows that it is possible to construct de Sitter vacua by compactifying extra dimensions in the presence of fluxes with an account of non-perturbative corrections \\cite{KKLT}.  The first step toward understanding the property of dark energy is to clarify whether it is a simple cosmological constant or it originates from other sources that dynamically change in time. The dynamical dark energy models can be distinguished from the cosmological constant by considering the evolution of $w_{\\rm DE}$. The scalar field models of DE such as quintessence \\cite{quin1,Ford,quin2,quin3,quin4,Ferreira1,Ferreira2,CLW,quin5,Zlatev,Paul99} and k-essence \\cite{kes1,kes2,kes3} predict a wide variety of variations of $w_{\\rm DE}$, but still the current observational data are not sufficient to provide some preference of such models over the $\\Lambda$-Cold-Dark-Matter ($\\Lambda$CDM) model. Moreover, the field potentials need to be sufficiently flat such that the field evolves slowly enough to drive the present cosmic acceleration. This demands that the field mass is extremely small ($m_{\\phi}\\simeq10^{-33}$\\,eV) relative to typical mass scales appearing in particle physics \\cite{Carrollqui,Kolda}. However it is not entirely hopeless to construct viable scalar-field dark energy models in the framework of particle physics. We note that there is another class of modified matter models based on perfect fluids--so-called (generalized) the Chaplygin gas model \\cite{Kamen,Bento}. If these models are responsible for explaining the origin of dark matter as well as dark energy, then they are severely constrained from the matter power spectrum in galaxy clustering \\cite{Waga}.  There exists another class of dynamical dark energy models that modify General Relativity. The models that belong to this class are $f(R)$ gravity \\cite{fR1,fR1d,fR2,fR3,Nojiri03} ($f$ is a function of the Ricci scalar $R$), scalar-tensor theories \\cite{st1,st2,st3,st4,st5}, and Dvali, Gabadadze and Porrati (DGP) braneworld model \\cite{DGP}. The attractive feature of these models is that the cosmic acceleration can be realized without recourse to a dark energy component. If we modify gravity from General Relativity, however, there are stringent constraints coming from local gravity tests as well as a number of observational constraints such as large-scale structure (LSS) and CMB. Hence the restriction on modified gravity models is in general very tight compared to modified matter models. We shall construct viable modified gravity models and discuss their observational and experimental signatures.  In addition to the above mentioned models, there are attempts to explain the cosmic acceleration without dark energy. One example is the void model in which an apparent accelerated expansion is induced by a large spatial inhomogeneity \\cite{Tomita1,Tomita2,Celerier,Iguchi,Alnes}. Another example is the so-called backreaction model in which the backreaction of spatial inhomogeneities on the Friedmann-Lema\\^{i}tre-Robertson-Walker (FLRW) background is responsible for the real acceleration \\cite{Rasanen,Kolb1,Kolb2}. We shall discuss these models as well.  This review is organized as follows. In Sec.~\\ref{obsec} we provide recent observational constraints on dark energy obtained by SN Ia, CMB, and BAO data. In Sec.~\\ref{cossec} we review theoretical attempts to explain the origin of the cosmological constant consistent with the low energy scale of dark energy. In Sec.~\\ref{mattersec} we discuss modified gravity models of dark energy--including quintessence, k-essence, coupled dark energy, and unified models of dark energy and dark matter. In Sec.~\\ref{mosec} we review modified gravity models and provide a number of ways to distinguish those models observationally from the $\\Lambda$CDM model. Sec.~\\ref{withoutsec} is devoted to the discussion about the cosmic acceleration without using dark energy. We conclude in Sec.~\\ref{consec}.  We use units such that $c=\\hbar=1$, where $c$ is the speed of light and $\\hbar$ is reduced Planck's constant. The gravitational constant $G$ is related to the Planck mass $m_{{\\rm pl}}=1.2211\\times10^{19}$\\,GeV via $G=1/m_{{\\rm pl}}^{2}$ and the reduced Planck mass $M_{{\\rm pl}}=2.4357\\times10^{18}$\\,GeV via $\\kappa^{2}\\equiv8\\pi G=1/M_{{\\rm pl}}^{2}$, respectively. We write the Hubble constant today as $H_0=100~h$\\,km\\,sec$^{-1}$\\,Mpc$^{-1}$, where $h$ describes the uncertainty on the value $H_0$. We use the metric signature $(-,+,+,+)$. ",
        "conclusions": "\\label{consec}  We summarize the results presented in this review. \\begin{itemize} \\item The cosmological constant ($w_{\\rm DE}=-1$) is favored by a number of observations, but theoretically it is still challenging to explain why its energy scale is very small. \\item Quintessence leads to the variation of the field equation of state in the region $w_{\\phi}>-1$, but the current observations are not sufficient to distinguish between quintessence potentials. \\item In k-essence it is possible to realize the cosmic acceleration by a field kinetic energy, while avoiding the instability problem associated with a phantom field. The k-essence models that aim to solve the coincidence problem inevitably leads to the superluminal propagation of the sound speed. \\item In coupled dark energy models there is an upper bound on the strength of the coupling from the observations of CMB, large-scale structure and SN Ia. \\item The generalized Chaplygin gas model allows the unified description of dark energy and dark matter, but it needs to be very close to the $\\Lambda$CDM model to explain the observed matter power spectrum. There is a class of viable unified models of dark energy and dark matter using a purely k-essence field. \\item In $f(R)$ gravity and scalar-tensor theories it is possible to construct viable models that satisfy both cosmological and local gravity constraints. These models leave several interesting observational signatures such as the modifications to the matter power spectrum and to the weak lensing spectrum. \\item The dark energy models based on the Gauss-Bonnet term are in conflict with a number of observations and experiments in general and hence they are excluded as an alternative to the $\\Lambda$CDM model. \\item The DGP model allows the self-acceleration of the Universe, but it is effectively ruled out from observational constraints and the ghost problem. However, some of the extension of works such as Galileon gravity allow the possibility for avoiding the ghost problem, while satisfying cosmological and local gravity constraints. \\item The models based on the inhomogeneities in the distribution of matter allow the possibility for explaining the apparent accelerated expansion of the Universe. The void model can be consistent with the SN Ia data, but it is still challenging to satisfy all other constraints coming from the CMB and the kinematic Sunyaev-Zeldovich effect. \\end{itemize}  When the author submitted a review article \\cite{review} on dark energy to International Journal of Modern Physics D in March 2006, we wrote in concluding section  that ``over 900 papers with the words `dark energy' in the title have appeared on the archives since 1998, and nearly 800 with the words `cosmological constant' have appeared''. Now in April 2010, I need to change the sentence to ``over 2250 papers with the words `dark energy' in the title have appeared on the archives since 1998, and nearly 1750 with the words `cosmological constant' have appeared''. This means that over 4000 papers about dark energy and cosmological constant have been already written, with more than 2300 papers over the past 4 years. Many cosmologists, astrophysicists, and particle physicists have extensively worked on this new field of research after the first discovery of the cosmic acceleration in 1998. We hope that the future progress of both theory and observations will provide some exciting clue to reveal the origin of dark energy."
    },
    "1004/1004.1170_arXiv.txt": {
        "abstract": "We measure the mass and size of cloud fragments in several molecular clouds continuously over a wide range of spatial scales ($0.05 \\lesssim r / {\\rm pc} \\lesssim 3$). Based on the recently developed ``dendrogram-technique'', this characterizes dense cores as well as the enveloping clouds. ``Larson's 3$^{\\rm rd}$ Law'' of constant column density, $m(r) \\propto r^2$, is not well suited to describe the derived mass-size data. Solar neighborhood clouds not forming massive stars ($\\lesssim 10 \\, M_{\\sun}$; Pipe Nebula, Taurus, Perseus, and Ophiuchus) obey $$ m(r) \\le 870 \\, M_{\\sun} \\, (r / {\\rm pc})^{1.33} \\, . $$ In contrast to this, clouds forming massive stars (Orion~A, G10.15$-$0.34,  G11.11$-$0.12) do exceed the aforementioned relation. Thus, this limiting mass-size relation may approximate a threshold for the formation of massive stars. Across all clouds, cluster-forming cloud fragments are found to be---at given radius---more massive than fragments devoid of clusters. The cluster-bearing fragments are found to roughly obey a mass-size law $m \\propto r^{1.27}$ (where the exponent is highly uncertain in any given cloud, but is certainly smaller than 1.5). ",
        "introduction": "Most of our present understanding of star formation processes is based on detailed studies of solar neighborhood molecular clouds (closer $\\sim 500 ~ \\rm pc$). To this end past research has, e.g., studied the masses and sizes of dense cores in molecular clouds ($\\lesssim 0.1 ~ \\rm pc$ size) such as Perseus, Taurus, Ophiuchus, Orion, and the Pipe Nebula (e.g.\\ \\citealt{motte1998:ophiuchus}, \\citealt{johnstone2000:rho_ophiuchi}, \\citealt{hatchell2005:perseus}, \\citealt{enoch2007:cloud_comparison}). Further research studied clumps in these clouds (some $0.1 ~ \\rm pc$) and the clouds ($\\gtrsim 10 ~ \\rm pc$) containing the cores (e.g., \\citealt{williams1994:clumpfind}, \\citealt{cambresy1999:extinction}, \\citealt{kirk2006:scuba-perseus}; see \\citealt{williams2000:pp_iv} for definitions of cores, clumps, and clouds). This research does, however, not probe the \\emph{relation} between the properties of cores, clumps, and clouds: traditionally, every domain is characterized and analyzed separately. As a result, even in the solar neighborhood, it is still not known how the core densities (and thus star-formation properties) relate to the state of the surrounding cloud.  To be precise, we do in principle know a bit about the relation between cloud structure at large and small scale. For structure within molecular clouds, \\citet{larson1981:linewidth_size} concluded (in his Eq.\\ 5) that the mass contained within the radius $r$ obeys a power-law, \\begin{equation} m(r) = 460 \\, M_{\\sun} \\, (r / {\\rm pc})^{1.9} \\, . \\label{eq:mass-size-larson} \\end{equation} Most subsequent work refers to this relation as ``Larson's 3$^{\\rm rd}$ law'', and replaces the original result with $m(r) \\propto r^2$ (e.g., \\citealt{mckee2007:review}). This ``law of constant column density'' (with respect to scale, $r$) is now considered one of the fundamental properties of molecular cloud structure (e.g., \\citealt{ballesteros-paredes2007:ppv}, \\citealt{mckee2007:review}, \\citealt{bergin2007:dense-core-review}). We do, however, not know whether this relation is still consistent with up-to-date column density maps of molecular clouds.\\medskip  Part I of the present series \\citep{kauffmann2010:mass-size-i} describes a new technique to extract mass-size relations from cloud maps. It is based on ``dendrograms'', a tree-based segmentation of cloud structure \\citep{rosolowsky2008:dendrograms}. Here, we employ this technique to study the molecular clouds in Perseus, Taurus, Ophiuchus, Orion, and the Pipe Nebula. To illustrate the properties of more massive clouds, we also include data for two more distant clouds of high density (farther than $2 ~ \\rm kpc$; G10.15$-$0.34 and G11.11$-$0.12).\\medskip  The present paper summarizes the analysis method in Sec.\\ \\ref{sec:method}. The main quantitative analysis of the maps is presented in Sec.\\ \\ref{sec:sample-analysis}. Section \\ref{sec:interpretation} systematizes the results and interprets them in the context of our present knowledge of star formation regions. This discussion is supported by model calculations in Appendices \\ref{sec-app:mass-size-models} and \\ref{sec-app:polytropes}. We conclude with a summary in Sec.\\ \\ref{sec:summary}. ",
        "conclusions": "} It is thought since long that clouds need to achieve a high column density in order to produce dense cores and stars. \\citet{johnstone2004:extinction_threshold_ophiuchus}, in particular, introduced the concept of extinction (or column density) thresholds for dense core formation (also see: \\citealt{onishi1998:c180-taurus}, \\citealt{hatchell2005:perseus}, \\citealt{enoch2007:cloud_comparison}). Similarly, \\citet{lombardi2006:pipe} and \\citet{lada2009:california-cloud} argue that a low fraction of mass at high column density yields low star formation activity (also see \\citealt{kainulainen2009:column-density-pdf}). This is in line with the Kennicutt-Schmidt law between star formation rate and mass surface density, $\\Sigma_{\\rm SFR} \\propto \\Sigma_{\\rm gas}^p$ (e.g., \\citealt{kennicutt1998:ks-law}). Since $p \\sim 1$ for star-forming clouds \\citep{evans2009:c2d-summary}, this relation predicts an increase of star formation activity with increasing column density.  The analysis in Sec.\\ \\ref{sec:sample-analysis} shows that these laws do also manifest in our data. Sections \\ref{sec:mass-size-trends} and \\ref{sec:clustered-isolated} suggest that the ability to form clusters and massive stars increases with increasing mass, when considering a given radius. Since $\\langle N_{\\rm H_2} \\rangle \\propto m / r^2$, one could also say that cloud fragments do appear to only form massive stars and clusters when they have a high mean column density---just as suggested by the aforementioned laws.\\medskip  Mass-size studies do, however, also provide information not available from the aforementioned plain column density studies. First, note that ``extinction threshold'' studies (e.g., \\citealt{johnstone2004:extinction_threshold_ophiuchus}) do only consider a single spatial scale, i.e.\\ the beam used to construct the extinction map. This is a major difference to mass-size studies, where many spatial scales are considered. Second, observe that studies of column density distributions (PDFs; e.g., \\citealt{lombardi2006:pipe}) \\emph{do} consider all scales of a map, but do not register which part of the signal shown on the histogram originates is which part of the cloud. (Depending on the analysis, this is not necessarily a problem.) Mass-size studies, in contrast, treat individual cloud fragments separately.  Mass-size data sets give a new twist to discussions of extinction thresholds. To see this, consider the solar neighborhood clouds examined here. These are all similar at large spatial scale (Eq.\\ \\ref{eq:mass-reference_large-radii}). Still, at smaller scale, they differ significantly in mass and star formation activity (Sec.\\ \\ref{sec:clustered-isolated}). The processes determining the star formation activity must thus operate on spatial scales smaller than the entire cloud. In this sense, \\emph{the star formation activity depends} (at least in our sample clouds) \\emph{on a cloud's ability to create, from a given mass reservoir, a small number of fragments that dominate the mass reservoir and concentrate it into a small volume}. The presence of high column densities are then a consequence of the cloud structure, but not the governing reason for the formation of dense cores and stars.  \\subsection{Slopes and Intercepts: Constraints on Density and Physical Cloud Models\\label{sec:density-pressure}} The observed power-law-like mass-size relations, $$ m(r) = m_0 \\, (r / {\\rm pc})^b \\, , $$ are characterized by slopes, $b$, and intercepts, $m_0$. As we show here, slopes and intercepts can be used to gauge densities and their gradients. At the same time, it is possible to constrain the absolute level of pressure, and the nature of its origin.  Consider, for example, an infinite equilibrium sphere with power-law density profile that is supported by isothermal pressure from gas at temperature $T_{\\rm g}$. Then, \\begin{equation} m(r) = 2.6 \\, M_{\\sun} \\, \\left( \\frac{T_{\\rm g}}{10 ~ \\rm K} \\right) \\, \\left( \\frac{r}{0.1 ~ \\rm pc} \\right) \\label{eq:mass-size-sis} \\end{equation} (\\citealt{kauffmann2008:mambo-spitzer}, Eq.\\ [13], in their case $\\epsilon \\to \\pi/2$). Thus, if this model holds, the intercept encodes information on the gas temperature supporting the cloud. Conversely, the slope predicted by the model, $b = 1$, can be used to validate the model; if observations yield $b \\neq 1$, then the model does not apply.  It has to be kept in mind that we consider mass-size laws derived from two-dimensional maps. These are related to, but not identical with, mass-size laws obtained from three-dimensional density maps. This is illustrated by the experiments conducted by \\citet{shetty2009:ppp-ppv} who use the fragment identification technique also used by us. Their analysis is based on three-dimensional numerical simulations of turbulent clouds. As part of their experiments, they derive mass-size slopes from their data. For their particular set of simulations, the power-law slopes derived in this fashion are similar to the number of dimensions used for mass measurements (i.e., 3 when based on density, and 2 when based on column density). This underlines that the number of dimensions considered has to be kept in mind.  \\subsubsection{Density Laws} The above discussion can be extended to include many more models of cloud structure. We do this in two Appendices. The results of this analysis are summarized here.  First, let us examine the connection between mass-size laws and cloud density profiles. Consider a sphere with about constant density for radii smaller than some flattening radius, $s_0$, a power-law drop at intermediate radii, $n(s) \\propto s^{-k}$ (where $s$ is the distance from the center), and vanishing density beyond some outer truncation radius, $R$. Such profiles provide a good match to observations of dense cores \\citep{tafalla2002:depletion}. They are a good approximation to the structure of isothermal equilibrium spheres \\citep{dapp2009:density-profile}. As we show in Appendix \\ref{sec-app:density-spheres}, for apertures of radius $r$, this yields mass-size relations of the form \\begin{equation} m(r) \\propto n_{\\rm c} \\, r^{3 - k} \\quad {\\rm for} \\quad s_0 \\ll r \\ll R \\, , \\end{equation} where $n_{\\rm c}$ is the density for $s = 0$, and mass-size slopes \\begin{equation} {\\rm d} \\, \\ln(m) / {\\rm d} \\, \\ln(r) = 3 - k \\quad {\\rm for} \\quad s_0 \\ll r \\ll R \\, . \\end{equation} Both relations apply only if $k < 3$. As the equations show, the intercept contains information on the central density, and the mass-size slope depends on the slope of the density law, $k$. Both does, of course, only hold at intermediate radii, $s_0 \\ll r \\ll R$. For reference, we note that the column density obeys $N \\propto r^{1 -   k}$.  A generalized version of power-law spheres are tri-axial ellipsoids. Appendix \\ref{sec-app:density-triaxial} considers the case in which $n(s) \\propto (s / s_0)^{-k}$ along any main axis, but with $s_0$ depending on the direction chosen (Eq.\\ \\ref{eq-app:density-law-ellipsoidal}). Detailed analysis shows that such ellipsoids follow the same mass-size relations as spheres, when $r = (A / \\pi)^{1/2}$. Thus, the laws listed above apply.  In a next step, one may wish to consider models of cylindrical clouds of length $\\ell$. Here, we adopt density drops $n(s) \\propto s^{-k}$ perpendicular to the cylinder axis for intermediate values of $s$. At intermediate radii $s_0 \\ll r \\ll R$, such clouds obey \\begin{equation} m(r) \\propto n_{\\rm c} \\, r^{4 - 2 k} \\quad \\Rightarrow \\quad {\\rm d} \\, \\ln(m) / {\\rm d} \\, \\ln(r) = 4 - 2 k \\label{eq:slope-cylinder_perp} \\end{equation} if their major axis is perpendicular to the line of sight, and \\begin{equation} m(r) \\propto n_{\\rm c} \\, r^{2 - k} \\quad \\Rightarrow \\quad {\\rm d} \\, \\ln(m) / {\\rm d} \\, \\ln(r) = 2 - k \\, , \\label{eq:slope-cylinder_par} \\end{equation} if the axes are aligned. (Intermediate angles are not considered here.) Meaningful relations are only obtained for $k < 2$. In both relations, the radius is defined as $r = (A / \\pi)^{1/2}$.  Thus, just as one may naively expect, the mass-size slope gauges the slope of the density profile. Further, the intercepts of mass-size relations constrain the absolute density of cloud fragments. There is, however, one less obvious fact that calls for attention: the exact relations between mass-size slopes, intercepts, and density law depends on the cloud model and viewing angle. It is therefore not possible to derive the true density profile without further information on the cloud geometry. Such information may, e.g., be derived by studying the elongation of cloud fragments. Also, the above power-law relations do only apply at intermediate radii, $s_0 \\ll r \\ll R$. This domain might not exist in actual observed clouds. Then, the central density plateau and the finite size have to be considered. These give \\begin{equation} m(r) \\left\\{ \\begin{array}{llll} \\propto n_{\\rm c} \\, r^2 & {\\rm for} & r \\ll s_0 &{\\rm and}\\\\ \\approx M & {\\rm for} & r \\gtrsim R & .\\\\ \\end{array} \\right. \\end{equation}  \\subsubsection{Polytropic Equilibria} The density slopes themselves depend on the processes shaping the model cloud. As a first example, here we consider static equilibrium models in which pressure gradients are in balance with self-gravity. We assume a polytrophic equation of state, $P \\propto n^{\\gamma_P}$, in which pressure and density are related by the polytrophic exponent, $\\gamma_P$.  In Appendix \\ref{sec-app:polytropes} we show that \\begin{equation} k = \\frac{2}{2 - \\gamma_P} \\label{eq:slope-gammap} \\end{equation} for polytropic equilibrium spheres (if $\\gamma_P < 4/3$) and cylinders (if $\\gamma_P < 1$). The density and mass-size slopes are, thus, related to the polytrophic exponent. Isothermal pressure, for which $\\gamma_P = 1$, implies $k = 2$ in spheres, for example. Then, ${\\rm d} \\, \\ln(m) / {\\rm d} \\, \\ln(r) = 1$ in spherical model clouds; laws too complex to be considered here apply to cylinders. As seen in Figs.\\ \\ref{fig:global-slopes_sf} and \\ref{fig:slope-comparison}, such a model can explain some, but not most slope measurements.  Polytropic exponents $\\gamma_P = 1/2$ are sometimes suggested to describe ``turbulent'' pressure within clouds, as e.g.\\ arising from Alfv\\'en waves \\citep{mckee1995:alfven_waves}. In this case, $k = 4/3$, and so ${\\rm d} \\, \\ln(m) / {\\rm d} \\, \\ln(r)$ assumes values of $5/3 \\approx 1.67$ (spheres and ellipsoids), $4/3 \\approx 1.33$ (perpendicularly viewed cylinder), and $2/3 \\approx 0.67$ (end-on cylinder) for the different models. Among these, spheres, ellipsoids, and cylinders viewed from the side provide an acceptable match to the observed mass-size slopes $b > 1$. Cylinders viewed along their major axis yield too shallow mass-size laws (and such a viewing direction is highly unlikely).  For a given physical model, the intercept can be used to gauge a cloud's stability against collapse, respectively suggest the level of supporting pressure. Here, we limit ourselves to the isothermal case, i.e.\\ $\\gamma_P = 1$. Stability considerations (e.g., of Bonnor-Ebert-type; \\citealt{ebert1955:be-spheres}, \\citealt{bonnor1956:be-spheres}) imply \\begin{equation} M \\le M_{\\rm cr} \\approx 2 \\, \\frac{\\sigma^2(v) \\, R}{G} \\label{eq:stability} \\end{equation} for the total mass, where $\\sigma(v)$ is the characteristic one-dimensional velocity dispersion (Eqs.\\ \\ref{eq-app:mass-limit-spheres} and \\ref{eq-app:mass-limit-cylinders}). For spheres, $R$ is the radius, while $R \\to \\ell$ in cylinders. If $\\sigma(v)$ is known (e.g., for thermal pressure), $M > M_{\\rm cr}$ implies collapse of the object considered. Conversely, depending on the situation, $\\sigma(v)$ can be inferred by requiring that $M = M_{\\rm cr}$. Required values of $\\sigma(v)$ significantly in excess of the thermal velocity dispersion of the mean free particle might, e.g., suggest the presence of significant non-thermal pressure. If we only require that pressure balances gravity, and drop the constraint that the object is to be stable against perturbations, the above law yields Eq.\\ (\\ref{eq:mass-size-sis}).  As particular example, consider B68. It is thought that this dense core has a structure very similar to a Bonnor-Ebert sphere \\citep{alves2001:b68}. Thus, one would expect the mass and size of B68 to obey Eq.\\ (\\ref{eq:mass-size-sis}), when considering large enough radii. This is indeed the case, as seen in Fig.\\ \\ref{fig:cloud-sample}.  \\subsection{Synoptic and physical Density Slopes} For intuitive communication, it may be helpful to report synoptic density slopes, \\begin{equation} \\left[ - \\frac{{\\rm d} \\, \\ln(n)}{{\\rm d} \\, \\ln(s)} \\right]_{\\rm syn} = 3 - \\frac{{\\rm d} \\, \\ln(m)}{{\\rm d} \\, \\ln(r)} \\, , \\end{equation} i.e.\\ the density slope a sphere of the same mass-size slope would have when observed at intermediate radii. The synoptic slopes give a good first idea of the true density slopes. First, recall that the model mass-size laws do not sensitively depend on the assumption of exact spheres; the same relation holds for ellipsoids. Also, in the observed range $1 \\lesssim {\\rm d} \\, \\ln(m) / {\\rm d} \\, \\ln(r) < 2$, spheres (or ellipsoids) and perpendicularly viewed cylinders (the end-on view is statistically insignificant) imply similar slopes; for these geometries, the synoptic slopes exceed the true ones by less than 0.5, assuming intermediate radii. Thus, we derive $$ \\left[ - {\\rm d} \\, \\ln(n) / {\\rm d} \\, \\ln(s) \\right]_{\\rm syn} = 1 ~ {\\rm to} ~ 2 $$ for mass-size slopes $1 \\lesssim {\\rm d} \\, \\ln(m) / {\\rm d} \\, \\ln(r) < 2$.  A limited comparison of these density-size slopes with previous results is possible. \\citet{tafalla2002:depletion}, e.g., study the dust emission of five starless dense cores, and derive density-size slopes of $2.0 ~ {\\rm to} ~ 2.5$ for four of them. This is a typical result for cloud fragments of $\\lesssim 0.1 ~ \\rm pc$ size \\citep{bergin2007:dense-core-review}, a spatial domain not too well covered by our data.  Concerning the analysis method and spatial range considered, the extinction study by \\citet{stuewe1990:structural-analysis} might provide the best match to our work. Based on star counts, he derives $- {\\rm d} \\, \\ln(n) / {\\rm d} \\, \\ln(s) > 1.0 \\pm 0.4$ on scales of up to $\\sim 1 ~ \\rm pc$. This is consistent with our results, also given the differences in map construction (he uses optical star counts).     Analysis of a limited sample of solar neighborhood cloud complexes (Taurus, Ophiuchus, Perseus, Pipe Nebula, Orion~A), as well as more distant clouds (G10 \\& G11), yields first some basic constraints on mass-size cloud structure. These are as follows. \\begin{enumerate} \\item On large spatial scales, $\\ge 1 ~ \\rm pc$, the most massive fragments in solar neighborhood clouds---with the possible exception of Orion~A---obey $$ m(r) = 400 \\, M_{\\sun} \\, (r / {\\rm pc})^{1.7} $$ (Eq.\\ \\ref{eq:mass-reference_large-radii}) with deviations $< 40 \\%$. This relation resembles the original mass-size law derived by \\citet{larson1981:linewidth_size}, $m(r) = 460 \\, M_{\\sun} \\, (r / {\\rm pc})^{1.9}$. The more distant clouds in the sample, however, deviate from this relation by up to an order of magnitude in mass. \\item No single mass-size law can be used to describe all fragments in all clouds. In particular, ``Larson's 3$^{\\rm rd}$ Law'' of constant column density, $m(r) \\propto r^2$, provides a bad global description; today's data are too complex to warrant the use of such relations. To give examples, power-law slopes vary with radius within a given cloud (Fig.\\ \\ref{fig:slope-comparison}), and clouds can differ massively in mass at given radius (Fig.\\ \\ref{fig:mass-size-comparison}).  In practice, different definitions of mass-size laws are used by different researchers. Also, different definitions may serve different purposes. This must be taken into account when comparing different mass-size laws. \\end{enumerate} Most importantly, the mass-size data can be used to learn about the formation of stars in molecular clouds. We derive the following constraints. \\begin{enumerate} \\setcounter{enumi}{2} \\item Sample clouds not forming massive stars ($\\gtrsim 10 \\, M_{\\sun}$) adhere to a limiting mass size relation, $$ m(r) \\le 870 \\, M_{\\sun} \\, (r / {\\rm pc})^{1.33} $$ (Eq.\\ \\ref{eq:mass-reference_limit-low-mass}), while our sample clouds forming such stars violate this law (Fig.\\ \\ref{fig:mass-size-comparison}). This suggests that the above relation describes the typical mass-size range of molecular clouds not forming high-mass stars. Also, the observations advocate that this boundary constitutes a mass limit for massive star formation. However, such conclusions are based on a small sample and are thus preliminary. \\item Across all clouds studied here, cloud fragments forming clusters are more massive than fragments not doing so (Figs.\\ \\ref{fig:cloud-sample} and \\ref{fig:mostmassive}; Sec.\\ \\ref{sec:clusters-dominate-host}). At given size, cluster-forming fragments dominate the mass reservoir of their host cloud. \\item The mass-size trend of cluster-forming fragments can e.g.\\ be captured by global mass-size slopes (i.e., from $0.05 ~ \\rm pc$ to $5.0 ~ \\rm pc$ radius; Sec.\\ \\ref{sec:global-slopes_sample}). Our cluster-forming sample clouds are consistent with a common slope $\\sim 1.27$. The uncertainties are, unfortunately, significant for a given cloud; slopes may well differ between clouds. In the case of Orion~A, e.g., the slope might be as low as $\\sim 0.7$. However, in any event slopes smaller 1.5 do hold. \\end{enumerate} Theoretical discussions show that mass-size laws of the form $m(r) = m_0 \\, (r / {\\rm pc})^b$ can be related to physical cloud models characterized by power-law density gradients, $n(s) \\propto s^{-k}$, or polytropic equations of state, $P \\propto n^{\\gamma_P}$ (Sec.\\ \\ref{sec:density-pressure}). Provided certain idealizations apply, $b$, $k$, and $\\gamma_P$ are directly related to another. This analysis suggests the definition of a synoptic density slope, $ \\left[ - {\\rm d} \\, \\ln(n) / {\\rm d} \\, \\ln(s) \\right]_{\\rm syn} = 3 - {\\rm d} \\, \\ln(m) / {\\rm d} \\, \\ln(r) $ (i.e., assuming the fragment considered was a sphere). This slope provides a first rough estimate of the true density law. Our data gives synoptic density slopes in the range $1 ~ {\\rm to} ~ 2$."
    },
    "1004/1004.2490_arXiv.txt": {
        "abstract": "The next generation of telescopes aim to directly observe the first generation of galaxies that initiated the reionization process in our Universe. The Ly$\\alpha$ emission line is robustly predicted to be the most prominent intrinsic spectral feature of these galaxies, making it an ideal target to search for and study high redshift galaxies. Unfortunately the large Gunn-Peterson optical depth of the surrounding neutral intergalactic medium (IGM) is thought to render this line extremely difficult to detect prior to reionization. In this paper we demonstrate that the radiative transfer effects in the interstellar medium (ISM), which cause Ly$\\alpha$ flux to emerge from galaxies at frequencies where the Gunn-Peterson optical depth is reduced, can substantially enhance the prospects for detection of the Ly$\\alpha$ line at high redshift. In particular, scattering off outflows of interstellar \\ion{H}{I} gas can modify the Ly$\\alpha$ spectral line shape such that $\\gsim 5\\%$ of the emitted Ly$\\alpha$ radiation is transmitted directly to the observer, {\\it even through a fully neutral IGM}. It may therefore be possible to directly observe `strong' Ly$\\alpha$ emission lines (EW$\\gsim 50$ \\AA\\hs rest frame) from the highest redshift galaxies that reside in the smallest \\ion{H}{II} `bubbles' early in the reionization era with JWST. In addition, we show that outflows can boost the fraction of Ly$\\alpha$ radiation that is transmitted through the IGM during the latter stages of reionization, and even post-reionization. Coupled with the fact that the first generation of galaxies are thought to have very large intrinsic equivalent Ly$\\alpha$ equivalent widths, our results suggest that the search for galaxies in their redshifted Ly$\\alpha$ emission line can be competitive with the drop-out technique out to the highest redshifts that can be probed in the JWST era. ",
        "introduction": "\\label{sec:intro}  One of the main science drivers of the next generation of telescopes is to detect the first generation of galaxies\\footnote{The definition of `the first generation of galaxies' is somewhat arbitrary. We take it to mean any star forming galaxy during the earliest stages of the reionization that might be detected (in  either their continuum or in on of their lines) by the next generation of telescopes. Although the analysis presented in this paper also applies to the first stars that likely formed one--by--one in minihalos, the overall fluxes from these sources is much too faint to be detected in the near future (regardless of radiative transfer effects that may boost the detectability of Ly$\\alpha$ emission from such sources).} that formed in our Universe. These galaxies contained hotter and more compact stars \\citep[e.g.][]{TS00,Bromm01}, whose initial mass function (IMF) was likely top-heavy \\citep[][]{La98,Bromm02}. Both the top-heavy IMF and low (or zero) gas metallicity enhanced the number of ionizing photons that the first galaxies emitted compared to that of local galaxies, at a fixed star formation rate \\citep{TS00,Bromm01,S02,S03}. This enhancement in ionizing luminosity results in larger \\ion{H}{II} regions in the interstellar medium (ISM) of these galaxies. As a result, one of the key predicted properties of the first galaxies are prominent nebular emission lines, dominated in flux by hydrogen Ly$\\alpha$ ($\\lambda=1216$\\hs\\AA, see e.g. Johnson et al. 2009). The first generation of galaxies are therefore likely to have been strong Ly$\\alpha$ emitters, with equivalent widths possibly as high as EW$\\sim 1500$ \\AA \\hs\\citep[][also see Partridge \\& Peebles 1967, Meier 1976]{S02,S03,J09b}.   In this paper we investigate the prospects for detecting this Ly$\\alpha$ emission. The first galaxies were surrounded by a mostly neutral intergalactic medium (IGM), which is extremely optically thick to Ly$\\alpha$ radiation. \\citet{LR99} showed that scattering of Ly$\\alpha$ photons through a neutral IGM causes galaxies to be surrounded by diffuse Ly$\\alpha$ halos \\citep[also see][]{Ko04,Ko06}. While the total flux in these halos can be substantial, their large angular size results in surface brightness levels that are beyond reach of even future telescopes such as the {\\it James Webb Space Telescope} (JWST, see \\S~\\ref{sec:RL}).  If Ly$\\alpha$ radiation from the first galaxies is to be detected, it will therefore be via Ly$\\alpha$ photons that were transmitted directly to the observer. In a neutral IGM, there are two mechanisms by which this can be achieved:  \\noindent ({\\it i}) The first mechanism is due to ionizing radiation that escapes from the galaxy, which creates a surrounding \\ion{H}{II} `bubble', that can strongly boost the detectability of Ly$\\alpha$ emission\\footnote{Indeed, the fact that the reionization process likely affects the observed number and distribution of high--redshift Ly$\\alpha$ emitting galaxies is exactly the reason why Ly$\\alpha$ emitting galaxies are thought to probe this epoch \\citep[][]{HS99,MR04,F06,Ka06,LF,McQ,Iliev08,Mesinger,Dayal10}.} \\citep[e.g.][]{Haiman02,Cen05}. This is because Hubble expansion redshifts Ly$\\alpha$ photons while they propagate freely through the ionized gas. As as result, a fraction of photons enter the neutral intergalactic medium on the red side of the line center, where the IGM optical depth optical depth can be smaller by orders of magnitude. For example, the Gunn-Peterson optical depth at redshift $z$ is given by \\citep[e.g.][]{BL01}  \\begin{equation} \\tau_{\\rm GP,0}\\approx 7.30\\times 10^5 x_{\\rm HI}\\Big{(} \\frac{1+z}{10}\\Big{)}^{3/2}, \\label{eq:taugp} \\end{equation} where $x_{\\rm HI}$ denotes the neutral volume fraction of hydrogen in the IGM. The Gunn-Peterson optical depth (Eq~\\ref{eq:taugp}) reduces to \\citep{M98,DW06} \\begin{equation} \\tau_{\\rm GP}(\\Delta v)\\approx 2.3\\Big{(} \\frac{\\Delta v}{600\\hs{\\rm km\\hs s}^{-1}}\\Big{)}^{-1}\\Big{(} \\frac{1+z}{10}\\Big{)}^{3/2} \\label{eq:redgp} \\end{equation} for photons that enter the neutral IGM with a redshift of $\\Delta v$  from line center. Eq~\\ref{eq:redgp} illustrates that photons that enter the neutral IGM significantly redward of the Ly$\\alpha$ resonance are scattered with only a weak optical depth. However given the high density of the surrounding IGM it is generally assumed that the first galaxies would not have had large enough \\ion{H}{II} `bubbles' to prevent complete damping of the line.  \\noindent ({\\it ii}) The second mechanism for enhancing the transmission of Ly$\\alpha$ photons from high redshift galaxies is due to  radiative transfer effects within the ISM of galaxies (in particular outflows of interstellar \\ion{H}{I} gas), which can shift Ly$\\alpha$ photons to the red side of the line before it reaches the IGM. Observed interstellar metal absorption lines (\\ion{Si}{II}, \\ion{O}{I}, \\ion{C}{II}, \\ion{Fe}{II} and \\ion{Al}{II}) in Lyman Break Galaxies (LBGs) are typically strongly redshifted relative to the galaxies' systemic velocity (with a median off-set of $\\sim 160$ km s$^{-1}$), while the Ly$\\alpha$ emission line is strongly redshifted \\citep[with a median velocity offset of $\\sim 450$ km s$^{-1}$ ][also see Shapley et al. 2003]{Steidel10}. This suggest that large scale outflows are ubiquitous in LBGs \\citep{Shapley03,Steidel10}. Furthermore, scattering of Ly$\\alpha$ photons by \\ion{H}{I} in outflows can successfully explain observed Ly$\\alpha$ line shapes in Ly$\\alpha$ emitting galaxies at $z=3-6$ \\citep[][]{V06,V08,V10}.  In this paper we explore the outflow mechanism in more detail. We will show that these `local' (i.e. inherent to the galaxy itself) processes in the ISM can cause as much as $\\gsim 5\\%$ of the emitted Ly$\\alpha$ radiation to be directly transmitted to the observer {\\it even through a fully neutral IGM}. This result is important for the study of high redshift galaxies, because it suggest that detecting the Ly$\\alpha$ emission line of the first generation of galaxies may well be within reach of the next generation of telescopes including JWST.  The outline of this paper is as follows: we describe our models and present our results in \\S~\\ref{sec:results}. We discuss our models and the implications of this work in \\S~\\ref{sec:disc}, before we conclude in \\S~\\ref{sec:conc}. The cosmological parameter values used throughout our discussion are $(\\Omega_m,\\Omega_{\\Lambda},\\Omega_b,h)=(0.27,0.73,0.046,0.70)$ \\citep{Komatsu08}. \\begin{figure*} \\vbox{\\centerline{\\epsfig{file=fig1.ps,angle=270,width=14.0truecm}}} \\caption[]{The observed properties of Ly$\\alpha$ halos (a.k.a Loeb-Rybicki halos) that surround galaxies embedded within a fully neutral comoving IGM. The `noisy' features in these plots are due to the finite (here $N_{\\rm phot}=10^6$) number of photons used in our Monte-Carlo calculations. {\\it Left panel:} the {\\it histogram} shows the integrated (over the entire area on the sky) emerging spectrum, under the assumption that all Ly$\\alpha$ photons were emitted at line center, and assuming an emitted Ly$\\alpha$ luminosity of $L_{\\alpha}=10^{43}$ erg s$^{-1}$. The lower horizontal axis shows the dimensionless frequency parameter $x$, while the upper horizontal axis shows the frequency parameter that was employed by Loeb \\& Rybicki (1999). The peak flux density occurs at $x\\sim -300$ which in our model corresponds to a redshift of $\\sim 660$ km s$^{-1}$. The FWHM of the spectrum is $\\sim 1500$ km s$^{-1}$. The vertical axes are expressed in arbitrary units on the left, and physical units on the right. The peak integrated flux density reaches 30 nJy. For comparison NIRSpec aboard JWST is expected to reach a sensitivity of $\\sim$ hundreds of nJy in $10^4$ s (S/N=10) for R=2700, or line fluxes of $\\gsim 10^{-19} $ erg s$^{-1}$ cm$^{-2}$. {\\it Right panel:} the integrated (over frequency) surface brightness $S$ as a function of impact parameter ($\\theta$) from the galaxy. We find that $S< 10^{-21}$ erg s$^{-1}$ cm$^{-2}$ at all $\\theta$. This plot illustrates that it will be extremely difficult to directly detect these Ly$\\alpha$ halos (see text). } \\label{fig:rl} \\end{figure*} ",
        "conclusions": "\\label{sec:conc} \\begin{figure*} \\vbox{\\centerline{\\epsfig{file=fig6.ps,angle=90,width=14.0truecm}}} \\caption[]{Schematic explanation for why outflows promote the detectability of Ly$\\alpha$ emission from galaxies surrounded by significant amounts of neutral intergalactic gas: In the {\\it top panel} a galaxy is surrounded by a large bubble of ionized gas, which in turn is surrounded by neutral intergalactic gas. Ly$\\alpha$ emission from this galaxy redshifts away from resonance as it propagates freely through the \\ion{H}{II} bubble (as indicated by the line color). Once the Ly$\\alpha$ photons reach the neutral IGM, they have redshifted far from resonance where the Gunn-Peterson optical depth is reduced tremendously (see Eq~\\ref{eq:redgp}). Because of the reduced GP optical depth, some fraction of the emitted Ly$\\alpha$ is transmitted to the observer without scattering in the IGM. In this drawing, the thickness of the line represents the specific intensity of the Ly$\\alpha$ radiation field. The {\\it bottom panel} shows that outflows surrounding star forming regions (represented by the expanding ring. The far side is receding form the observer and has a larger redshift, which is represented by the color) can Doppler boost Ly$\\alpha$ photons to frequencies redward of the Ly$\\alpha$ resonance. In this scenario, a fraction of Ly$\\alpha$ can propagate directly to the observer {\\it without the \\ion{H}{II} bubble}.} \\label{fig:scheme} \\end{figure*} The next generation of telescopes aim to directly observe the first generation of galaxies that initiated the reionization of our Universe. The Ly$\\alpha$ emission line is robustly predicted to be the most prominent intrinsic spectral feature of these galaxies. In this paper we investigated the prospects for detecting this Ly$\\alpha$ emission, taking account of radiative transfer effects that are likely to occur in the interstellar medium (ISM) of these galaxies.  Observed interstellar metal absorption lines (\\ion{Si}{II}, \\ion{O}{I}, \\ion{C}{II}, \\ion{Fe}{II} and \\ion{Al}{II}) in Lyman Break Galaxies (LBGs) are typically strongly redshifted relative to the galaxies' systemic velocity (with a median off-set of $\\sim 160$ km s$^{-1}$), while the Ly$\\alpha$ emission line is strongly redshifted \\citep[with a median velocity offset of $\\sim 450$ km s$^{-1}$ ][also see Shapley et al. 2003]{Steidel10}. This suggest that large scale outflows are ubiquitous in LBGs \\citep{Shapley03,Steidel10}. Furthermore, scattering of Ly$\\alpha$ photons by \\ion{H}{I} in outflows has successfully explained the observed Ly$\\alpha$ line shapes in Ly$\\alpha$ emitting galaxies at $z=3-6$ \\citep[e.g.][]{V06,V08,V10}. Scattering off outflows of interstellar \\ion{H}{I} gas can shift Ly$\\alpha$ photons to the red side of the line before it reaches the IGM (Fig~\\ref{fig:out}). At these frequencies the Gunn-Peterson optical depth may be reduced to order unity as a result.  In this paper we investigated the detectability of Ly$\\alpha$ radiation under the assumption that the outflows observed at low redshift also occur in the highest redshift galaxies. We found that outflows may cause as much as $\\gsim 5\\%$ of the emitted Ly$\\alpha$ radiation to be transmitted directly to the observer, even through a fully neutral IGM (\\S~\\ref{sec:outflow}). Since the intrinsic (restframe) equivalent width of the Ly$\\alpha$ line can be as high as EW$_{\\rm int}=1500$ \\AA for the first generation of galaxies, the observed EW$=f_{\\rm trans}$EW$_{\\rm int}=75(f_{\\rm trans}/0.05)({\\rm EW}_{\\rm int}/1500\\hs{\\rm \\AA})$ \\AA. For comparison, only $4\\%$ of the $z=3$ LBG population have larger EWs \\citep{Shapley03}. We showed that for $f_{\\rm trans} \\gsim 3\\%$ it may be easier to detect Ly$\\alpha$ line emission with NIRSPEC on the {\\it James Webb Space Telescope} (JWST), than continuum radiation with JWST's NIRCAM in the same integration time. We also note that the next generation of ground-based 30-m telescopes with diffraction limited AO are expected to be more (less) sensitive than JWST at $\\lambda \\gsim 1.3 \\mu$m for $R\\gg 100$ ($R \\lsim 100$) spectroscopy (see Mountain et al. 2009, their Fig~2). That is, ground based high-resolution spectroscopic searches for high redshift galaxies can detect fainter galaxies at a fixed integration time when $f_{\\rm trans} \\gsim 5\\%$. Such searches can therefore be competitive with searches that employ the drop-out technique. Irrespective of the survey strategy that is used to search for the highest redshift galaxies, the prospect that Ly$\\alpha$ can provide galaxies with spectroscopic redshifts is promising, and important as no robust predictions exists for the detectability of other emission lines (\\S~\\ref{sec:lines}).  This paper has focused on the first generation of galaxies that were surrounded by a neutral IGM, but our work also applies more broadly. For example we argued in \\S~\\ref{sec:EoR} that \\ion{H}{I} outflows promote the detectability of the Ly$\\alpha$ emission line during later stages of reionization when much of the absorption is resonant absorption in a highly ionized HII region. As a result outflows can reduce the minimum \\ion{H}{II} bubble size that is required to render LAEs `visible'. This is illustrated schematically in Figure~\\ref{fig:scheme}. Similarly, \\ion{H}{I} outflows also promote the detectability of the Ly$\\alpha$ emission line after reionization has been completed\\footnote{Recently, \\citet{Z10b} showed that scattering in the `local' IGM immediately surrounding Ly$\\alpha$ sources can introduce a unique anisotropy in the two-point correlation function of LAEs at $z=5.7$. \\citet{Z10b} argued that this scattering--induced signature is reduced when the intrinsic (i.e. prior to scattering) Ly$\\alpha$ line width is enhanced. Winds are therefore also expected to reduce this clustering signature. That is, the clustering of LAEs post-reionization can provide constraints on the importance of winds in shaping the Ly$\\alpha$ line shape.} (\\S~\\ref{sec:EoR}).  In summary, radiative transfer effects in the ISM of high redshift galaxies have been shown to broaden, and--in the case of outflows--redshift the emergent Ly$\\alpha$ flux to a level that allows $5\\%$ or more of the photons to escape absorption by the neutral IGM. Coupled with the large intrinsic Ly$\\alpha$ line EW of the first generation of galaxies, we have shown that searches for galaxies in their redshifted Ly$\\alpha$ emission line can be competitive with the drop-out technique out to the highest redshifts that can be probed observationally in the JWST era.  { \\bf Acknowledgments} MD thanks the School of Physics at the University of Melbourne, where most of this work was done, for their kind hospitality. MD is supported by Harvard University funds. We thank Zolt\\'{a}n Haiman \\& Zheng Zheng for helpful comments on earlier versions of this paper."
    },
    "1004/1004.2345_arXiv.txt": {
        "abstract": "\\fe\\ K line fluorescence is commonly observed in the X-ray spectra of many X-ray binaries and represents a fundamental tool to investigate the material surrounding the X-ray source.  In this paper we present a comprehensive survey of 41 X-ray binaries (10 HMXBs and 31 LMXBs) with \\chandra\\, with specific emphasis on the Fe K region and the narrow \\fe\\ \\ka\\  line, at the highest resolution possible. We find that: {\\it a}) The \\fe\\ \\ka\\ line is always centered at $\\lambda=1.9387\\pm 0.0016$ \\AA, compatible with Fe \\textsc{i} up to Fe \\textsc{x}; we detect no shifts to higher ionization states nor any difference between HMXBs and LMXBs.  {\\it b}) The line is very narrow, with $FWHM\\leq 5$ m\\AA, normally not resolved by \\chandra\\ which means that the reprocessing material is not rotating at high speeds.  {\\it c}) \\fe\\ \\ka\\ fluorescence is present in all the HMXB in the survey. In contrast, such emissions are astonishingly rare ($\\sim 10$ \\% ) among low mass X-ray binaries (LMXB) where only a few out of a large number showed \\fe\\ K fluorescence. However, the line and edge properties of these few are very similar to their high mass cousins.  {\\it d}) The lack of Fe line emission is always accompanied by the lack of any detectable K edge.  {\\it e}) We obtain the empirical curve of growth of the equivalent width of the \\fe\\ \\ka\\ line versus the density column of the reprocessing material, i.e. $EW_{\\rm K\\alpha}$ vs $N_{\\rm H}$, and show that it is consistent with a reprocessing region spherically distributed around the compact object.  {\\it f}) We show that fluorescence in X-ray binaries follows the X-ray Baldwin effect as previously only found in the X-ray spectra of active galactic nuclei. We interpret this finding as evidence of decreasing neutral Fe abundance with increasing X-ray illumination and use it to explain some spectral states of Cyg X-1 and as a possible cause of the lack of narrow Fe line emission in LMXBs.  {\\it g}) Finally, we study anomalous morphologies such as Compton shoulders and asymmetric line profiles associated with the line fluorescence. Specifically, we present the first evidence of a Compton shoulder in the HMXB X1908+075.  Also the \\fe\\ \\ka\\ lines of 4U1700$-$37 and LMC X-4 present asymmetric wings suggesting the presence of highly structured stellar winds in these systems. ",
        "introduction": "\\fe\\ K fluorescence lines constitute a fundamental tool to probe the physical characteristics of the material in the close vicinity of X-ray sources \\citep{george91}.  In the spectra of accreting X-ray Binaries (XRBs) these lines are very prominent due to of their intrinsic X-ray brightness and ubiquitous stellar material.  High mass X-ray binaries (HMXB) are composed of a compact object, either a neutron star (NS) or a stellar size black hole (BH), accreting from the powerful wind of a massive OB type star. The compact object in these cases is deeply embedded into the stellar wind of the donor providing an excellent source of illumination. In contrast, in low mass (LMXBs) and intermediate mass X-ray binaries (IMXBs), the donor stars are not significant sources of stellar winds. However they usually feature strong and matter rich accretion disks around the orbiting compact object and also have associated outflow processes which, in turn, tend to be not so important in HMXB.   Fluorescence excitation occurs whenever there is a low ionization gas illuminated by X-rays. In XRBs, the strong point like source of X-rays, allows to observe strong fluorescence emission.  The X-ray source, powered via accretion, irradiates the circumstellar material, either the wind or the accretion disk.  Whenever this material is more neutral than Li-like, the Fe atoms present in the stellar wind absorb a significant fraction of continuum photons blueward of the K edge (at $\\sim 1.74$ \\AA) thereby removing K shell electrons.  The vacancy thus produced will be occupied by electrons from the upper levels producing K$\\alpha$ (L $\\rightarrow$ K) and K$\\beta$ (M $\\rightarrow$ K) fluorescence emission lines at $\\sim$ 1.94 \\AA\\ and 1.75 \\AA\\ respectively (Fig. \\ref{fig:fecomplex}). K-shell fluorescence emission is highly inefficient for electron numbers $Z\\leq 16$.  The \\emph{Auger} effect dominates at lower $Z$ although K shell emission can be observed under very specific circumstances~\\citep{schulz02}. However, the fluorescence yield increases monotonically with $Z$. In the case of \\fe, fluorescence yields are already quite competitive \\citep[0.37, ][]{palmeri03}. Furthermore, the Fe is abundant and appears in an unconfused part of the spectrum. Therefore \\fe\\ K fluorescence is observed in a wide range of objects.  The \\fe\\ \\ka\\ line can present a composite structure. A broad line component, with {\\it FWHM} of the order of keV and a narrow line component with {\\it FWHM} much lower, of the order of some eV \\citep{miller02, hanke09}.  However, as has been shown in \\cite{hanke09}, the detection of the broad component by \\chandra\\ is very difficult and requires simultaneous {\\it RXTE} coverage. Furthermore, Nowak et al. (in prep.), observing Cyg X-1 simultaneously with several X-ray telescopes, have shown that while {\\it Suzaku} and {\\it RXTE} show clearly the broad component, \\chandra\\ does not. On the other hand, {\\it Suzaku} agrees on the narrow component detected by \\chandra. In the present survey, we will focus specifically on the narrow component, which is best studied at high resolution.  \\chandra\\ high energy transmission gratings (HETG, \\cite{canizares05}) are very well suited for this purpose. While {\\it RGS} instrument on board {\\it XMM-Newton} has the required spectral resolution, it lacks of effective area shortward of 6 \\AA.   The photons emitted during fluorescence must further travel through the stellar wind to reach the interstellar medium. In some cases, these photons can be Compton downscattered to lower energies and a 'red shoulder' can be resolved in the Fe line (\\cite{watanabe03}).  In such a case, the Compton shoulder can be used as a further probe of the wind material.  \\cite{gott95} established a comprehensive catalog of Fe line sources using \\emph{EXOSAT} GSPC.  These authors were able to detect iron line emission in 51 sources out of which 32 were identified as X-ray binaries (XRB). From these, 20 ($\\sim 63\\%$) were LMXB and 12 ($\\sim 37\\%$) HMXB. On average, the former showed a broad ($\\sim 1$ keV) line centered at $6.628\\pm 0.012$ keV, while the latter tended to show narrower ($\\sim 0.5$ keV) lines centered at $6.533\\pm 0.003$ keV. \\emph{EXOSAT} GSPC had a spectral resolution of about this amount and, therefore,  a width of $\\sim 0.5$ keV represents an upper limit.  More recently, \\cite{asai00} have performed a study of the Fe K line in a sample of 20 LMXB using ASCA GIS and SIS data. These authors were able to detect significant Fe line emission in roughly half of the sources. This line tended to be centered around 6.6 keV but showed large scatter with extreme values going from $\\sim 6.55$ to 6.7 keV. In general, the FWHM is not resolved but for those sources where the width could be measured was $\\sim 0.5$ keV.  In this paper we study in a homogeneous way and at the highest spectral resolution, the narrow component of the Fe line for the whole sample of HMXB and LMXB currently public within the \\chandra\\ archive. Specific studies of some individual sources within our sample have been published elsewhere (e.g. \\cite{watanabe06} for Vela X-1). However, we have reprocessed the entire sample to guarantee its homogeneity and have reanalyzed it, focused specifically on the narrow component.  \\begin{figure} \\includegraphics[angle=-90,width=\\columnwidth]{2733_m4.ps} \\caption{{\\it Chandra HETG} spectra of the HMXB GX301$-$2, included in this survey, ObsID 2733, showing all the relevant features discussed in the present work: \\fe\\ \\ka\\ and  \\fe\\ \\kb\\ fluorescence lines, Fe K edge, Compton shoulder and a hot line.} \\label{fig:fecomplex} \\end{figure} ",
        "conclusions": "We have reprocessed and analyzed the HETG spectra of all X-ray binaries publicly available at the \\chandra\\ archive with specific focus on the \\fe\\ K line region. The following conclusions can be drawn from this analysis:   Fe $K\\alpha$ fluorescence emission seem to be ubiquitous in HMXBs. This emission varies throughout time for a specific source. In particular, Cyg X-1 shows a rather weak emission, only in two out of nine \\chandra\\ observations analyzed here. This emission is detected in two different spectral states: during a low hard state (ObsID 3815) and an intermediate state (ObsID 2415).  It vanishes for brighter (softer) states. A possible explanation could reside in the X-ray Baldwin effect studied before. Only in low luminosity states there remains a significant fraction of near neutral Fe, although much less than in other binaries of this survey. This remaining neutral Fe, however, becomes ionized in brighter states with the disappearance of the fluorescence line. However, the lack of detection during other intermediate (ObsID 107) and low hard states (ObsID 1511) defies this explanation and means that this can not be the only mechanism at work in this system.  In contrast, such emissions are found to be very rare amongst LMXBs. Only four sources, out of 31 analyzed in this work, display the narrow component of Fe $K\\alpha$ in emission: 4U1822-37, GX1+4, Her X-1, and Cir X-1. This lack of narrow iron line is always accompanied by the lack of any detectable K edge. This finding, strongly suggests, that the neutral absorption column in LMXB is dominated by the interstellar medium while in HMXB the local absorption is very significant.   This finding is in contrast with the previous work by \\cite{gott95}, based on spectra of lower resolution, where the majority of sources showing Fe line emission were LMXB. Our findings, however, do not contradict the ASCA study of Asai et al 2000 which claim the detection of hot lines in the majority of LMXB. We point out though, that these hot lines are not as frequently observed in \\chandra\\ HETG spectra.  The curve of growth ($EW(K \\alpha)$ vs $N_{\\rm H}$) is fully consistent with a spherical distribution of reprocessing matter, formed by cold near neutral Fe, around the X-ray source. This reprocessing material must be close to the X-ray source, as the line and continuum variations are closely correlated. Those few LMXB displaying Fe $K\\alpha$ in emission follow the same correlations found for the HMXB.  Since the nature of the winds is very different in LMXB and HMXB this means that the circumsource material has lost all 'memory' of the donor. This is consistent with a reprocessing site location very close to the compact object.  We observe a moderate anticorrelation between $EW$ and the $L_{\\rm X}$ of the source, on average. Some sources follow individually this trend, over timescales from days to months, while other do not. This 'X-ray Baldwin effect' is reported here for the first time, for the XRBs as a class. The immediate interpretation is the increase in the degree of ionization of Fe with the increasing $L_{\\rm X}$ of the illuminating source which produces a concurrent decrease in the efficiency of the fluorescence process.  We observe a Compton shoulder in the supergiant HMXB X1908$+$075 formed by single Compton scattering of primary \\ka\\ photons. Together with the hypergiant system GX 301-2, where such a shoulder has first been reported~\\citep{watanabe03}, they form the very small group of galactic sources with such a feature detected. Other sources (LMC X-4, 4U1700$-$37) show asymmetric wings, with the blue wings rising sharply and red wings declining progressively. This effect can be explained as 'hot Compton shoulders'.  Fe line fluorescence is produced in the stellar wind of massive stars. Systems with a substantial wind component will show this line. Therefore it can be naturally observed in all HMXB. On the other hand, donor stars in LMXB are not significant sources of stellar winds and other sites must be invoked for the origin of this line. Relativistic accretion disks, as observed in some Seyfert galaxies, might be suitable candidates for those few LMXB with positive detections. However, the lack of the narrow Fe line in the spectra of the vast majority of LMXB might be an indication that such accretion disks are too hot or not illuminated."
    },
    "1004/1004.0560_arXiv.txt": {
        "abstract": "We present the {\\it Spitzer}/Infrared Spectrograph spectrum of the main-sequence star HD165014, which is a warm ($\\gtrsim 200$~K) debris disk candidate discovered by the {\\it AKARI} All-Sky Survey. The star possesses extremely large excess emission at wavelengths longer than 5$~\\micron$. The detected flux densities at 10 and 20$~\\micron$ are $\\sim 10$ and $\\sim 30$ times larger than the predicted photospheric emission, respectively. The excess emission is attributable to the presence of circumstellar warm dust. The dust temperature is estimated as 300--750~K, corresponding to the distance of 0.7--4.4~AU from the central star. Significant fine-structured features are seen in the spectrum and the peak positions are in good agreement with those of crystalline enstatite. Features of crystalline forsterite are not significantly seen. HD165014 is the first debris disk sample that has enstatite as a dominant form of crystalline silicate rather than forsterite. Possible formation of enstatite dust from differentiated parent bodies is suggested according to the solar system analog. The detection of an enstatite-rich debris disk in the current study suggests the presence of large bodies and a variety of silicate dust processing in warm debris disks. ",
        "introduction": "Debris disks were discovered in main-sequence stars by infrared excess over the photospheric emission in observations by {\\it IRAS} in the 1980s. Since debris disks are thought to be ``extra-solar zodiacal light'' formed via dust production from asteroids and/or comets \\citep[e.g.][]{backman93,lecavelier96}, it is interesting to examine mineralogical characteristics of debris dust and to explore connection between the debris dust and the dust in the solar system. Debris disks are also important as a probe of planetesimals, building blocks of planets, in extra-solar systems.  Mid-infrared (MIR) spectroscopy is a strong tool to investigate the properties of debris dust. Recent MIR observations with InfraRed Spectrograph \\citep[IRS;][]{houck04} on board {\\it Spitzer} \\citep{werner04} have revealed the presence of abundant crystalline forsterite (Mg$_2$SiO$_4$) and silica (SiO$_2$) in several debris disks \\citep[e.g.][]{chen06,lisse09}.  In this Letter, we present an MIR low resolution spectrum of the main-sequence star HD165014 obtained with {\\it Spitzer}/IRS. HD165014 is a debris disk candidate with large MIR excess selected from a search for debris disks \\citep{fujiwara09c} in the {\\it AKARI}/Infrared Camera (IRC) All-Sky Survey data \\citep{ishihara09}. We report the detection of abundant crystalline enstatite (MgSiO$_3$) dust compared to forsterite toward the star and discuss the origin of the debris dust around HD165014. ",
        "conclusions": "As mentioned above, it is evident that crystalline enstatite is abundant in the debris disk around HD165014 and forsterite features are not clearly seen. \\cite{chen06} conduct a comprehensive MIR spectroscopic survey of debris disks. Among their 59 debris disk samples, five objects (HR3927, $\\eta$~Crv, HD113766, HR7012, and $\\eta$~Tel) are found to possess spectral features that are well-modeled by $\\micron$-sized amorphous and crystalline silicates. Forsterite features are dominant compared to enstatite toward four samples (HR3927, $\\eta$~Crv, HD113766, and $\\eta$~Tel). HR7012, which shows a very strong peak at 9.1--$9.2~\\micron$ in its spectrum, was initially modeled by enstatite-rich silicate by \\cite{chen06}. However, \\cite{lisse09} conclude that the observed 9.1--$9.2~\\micron$ feature originates from abundant silica grains and that the feature strengths of enstatite and forsterite are almost comparable. Although there are a few additional debris disk stars with significant dust features in the $N$-band \\citep[$\\beta$~Pic, HIP8920, HD145263;][]{knacke93,song05,honda04}, no sample with high abundance of enstatite is reported to date. HD165014 is the only known debris disk that shows strong enstatite features compared to forsterite.  In the solar system, a number of achondrite meteorites, which are primarily composed of enstatite, named Aubrites, have been discovered. E-type asteroids are thought to have enstatite achondrite surfaces and to be parent bodies of aubrites \\citep[e.g.][]{zellner77} based on their reflection spectra. E-type asteroids form a large proportion of asteroids inward of the main belt known as Hungaria asteroids and are believed to be fragments of differentiated larger bodies that were heated at least to 1700~K \\citep{keil89}. Considering that enstatite is dominant around HD165014 rather than forsterite, the observed debris dust may originate from an analog of E-type asteroids in the solar system, being harmonic with the idea that the debris disk is formed by means of collisions between asteroids \\citep{backman93}. It is also known that Mercury's surface is rich in enstatite \\citep{sprague98}. If a Mercury-like planet exists around HD165014 and a mechanism to scatter the surface material of the planet, for example infall of small bodies, occurs, an enstatite-rich debris disk might be formed. It should also be noted that Mercury's high density is interpreted as a result of stripping of the surface crustal material \\citep{benz88}. The possible link between the debris material around HD165014 and Mercury needs to be further explored.  Although a number of MIR spectra of protoplanetary disks associated with younger stars, Herbig Ae/Be and T Tau stars, have been collected so far, objects that show enstatite features are rare. HD179218 is a Herbig Ae/Be star, which has the most enstatite-rich spectrum known to date \\citep{bouwman01,schutz05}. We show the $N$-band spectrum of HD179218 obtained with TIMMI2 \\citep{vanboekel05} in the bottom panel of Figure~\\ref{bestfit}. While crystalline enstatite features are clearly seen in the spectrum, features at 9.85 and $11.20~\\micron$ attributable to forsterite are also seen with comparable strengths as enstatite, suggesting the presence of abundant crystalline forsterite as well as enstatite. Annealing experiments of a magnesium silicate smoke made by \\cite{rietmeijer86} show that the initially formed forsterite and silica react with each other and form enstatite by the following reaction \\begin{eqnarray} {\\rm Mg}_2{\\rm SiO}_4 + {\\rm SiO}_2 \\longrightarrow 2{\\rm MgSiO}_3. \\end{eqnarray} \\cite{bouwman01} suggest that the presence of enstatite around HD179218 might be due to the high luminosity ($300L_\\odot$), which gives rise to rapid dispersal of the gas, resulting in a high efficiency of the reaction. \\cite{vanboekel05} suggest that enstatite might be produced by means of chemical equilibrium processes in high temperature environments, i.e., inner regions of the protoplanetary disk. \\cite{sato06} show that crystalline enstatite can be produced by simultaneous evaporation of SiO grains and Mg vapor in a plasma in laboratory experiments. It is now widely accepted that the turbulence in disks is attributed to the magnetorotational instability \\citep[MRI;][]{balbus91}, suggesting that at least part of the disk is sufficiently ionized \\citep[e.g.][]{inutsuka05}. Therefore the formation process of crystalline enstatite suggested by \\cite{sato06} may work efficiently in protoplanetary disks. It is possible that crystalline enstatite grains produced in the protoplanetary disk were once stored in small bodies such as comets and have been released recently in the HD165014 system. Indeed crystalline enstatite is detected toward some comets in the solar system by astronomical MIR observations \\citep[e.g.][]{lisse06} as well as by the STARDUST sample return mission \\citep[e.g.][]{zolensky06}. However, it remains an open question why crystalline forsterite is depleted around HD165014. Further experimental and theoretical studies are required to discuss the origin of abundant crystalline enstatite in the debris disk around HD165014 in detail."
    },
    "1004/1004.3343.txt": {
        "abstract": "A group of Mira variables in the solar neighborhood show unusual spatial motion in the Galaxy. To study this motion in much larger scale in the Galaxy, we newly surveyed 134 evolved stars off the Galactic plane by SiO maser lines, obtaining accurate radial velocities of 84 detected stars. Together with the past data of SiO maser sources, we analyzed the radial velocity data of a large sample of sources distributing in a distance range of about 0.3 -- 6 kpc in the first Galactic quadrant. % (but excluding near the galactic plane). At the Galactic longitudes between 20 and $40^{\\circ}$, we found a group of stars with large negative radial velocities, which deviate by more than 100 km s$^{-1}$ from the Galactic rotation. We show that these deviant motions of maser stars are created by periodic gravitational perturbation of the Bulge bar, and that the effect appears most strongly at radii between corotation and outer Lindblad resonances. The resonance effect can explain the displacement of positions from the Galactic plane as well. ",
        "introduction": "Stellar OH and SiO maser sources are powerful probes of Galactic structure and stellar evolution \\citep{hab06,deg08}. Radial velocity databases of these sources are useful to investigate dynamical motions of stars in the Galactic disk and Bulge \\citep{izu95,sev01,deg04b,fuj06}. Because of a recent progress of studies of tidal streams surrounding our Galaxy and in the solar neighborhood \\citep{bel07,gri09}, one of urgent issues in this field is to separate a Bulge-bar resonance stream from tidal streams of relic dwarf galaxies in the radial velocity data.  It has been known that the Galactic disk has two components, thick and thin; the former has a thickness of 1 -- 2 kpc involving metal-poor stars and  kinematically peculiar stars, while the latter has a thickness of about 300 pc involving young new populations. A hypothesis that the thick disk is a relic of past merging processes has been proposed for the origin of thick disk \\citep{hel99,nav04,hel06}.  Moving groups in the solar neighborhood have also been known, and they are considered to be fossils keeping dynamical information of their birth. Two famous examples are the Arcturus and Hercules groups of stars \\citep{egg96}; the former is a group of metal poor stars with a coherent spatial motion with a $\\sim 100$ km s$^{-1}$ lag to the Galactic rotation, and the latter is a stellar group with a heterogeneous mixture of metal abundances and with a smaller rotational lag. Spatial motions of these moving groups are well investigated optically based on the Hipparcos (proper motions) and the RAVE (the radial velocities) databases. However, these investigations have a limitation of distance up to about 1 kpc due to lack of proper motion data. Because of heterogenous metal abundances of member stars, the origin of the Hercules stream is attributed to a resonance of the bar-like Bulge \\citep{ben07}. \\citet{fea00} investigated an outward motion of short-period Mira variables near the Sun, and attributed it to the resonance effect of the Bulge bar.  In this paper, we reinvestigate the radial velocity data of SiO maser sources toward the region of $l=20$ -- 60$^{\\circ}$, and $-30<b<60^{\\circ}$ (excluding the galactic plane, $|b|<3^{\\circ}$). In this region, \\citet{deg07} found a group of SiO maser stars with large negative velocities ($v_{\\rm LSR} \\sim -70$ km s$^{-1}$), which may be attributed to a resonance effect of the Bulge bar or a relic streaming motion. However, the previous off-plane SiO search was aimed to find distant debris stars such as those in the Sgr dwarf streams, the search %in the different directions was somewhat shallow in depth except toward Sgr dwarf streams. Therefore, we have made a new sensitive observation by SiO maser lines toward the thick disk, and have added more data in radial velocity database. %The new data clearly show %a concentration of negative velocity stars in a limited range of Galactic longitudes %($\\Delta l \\sim 20^{\\circ}$), but spreading widely in Galactic latitudes ($\\Delta b \\sim 60^{\\circ}$). %Though the kinematics of these stars has a similarity with that of the Hercules moving group of stars %or the short-period mira group proposed by \\citet{fea00} %in the solar neighborhood, they appear at much larger distances and more deviant %in radial velocities, which are explored in the present survey. Here, we analyze kinematics of this SiO maser star stream, and test if %the net motion  of this star group is outward from the Galactic center. We conclude that this stream is originated by the gravitational perturbation of the Bulge bar.  % Apex    Antapex %\\cite{fea96} %\\cite{fam05} %\\cite{ben07} % \\cite{hel06} %\\cite{deg07} %\\cite{nav04} % \\cite{ecu07} %Massive stars at the low-mass limit ($M \\simeq 10--15 M_{\\odot}$) spend %\\newpage  % Hercules stream (U_LSR, V_LSR)=(-42, -51, -6) km/s   U toward Galactic center  V toward l=90 deg % (l,b)_Her=(50.5,5.2)    V_tot=66+-32 km/s % opposite direction (230, -5)  --> 07 04 50.47  -17 27 20.0 % Galactic rotational velocity  --> positive  not observed %  6% of K and M giants around the sun belong to this stream. ",
        "conclusions": "We detected 84 out of 134 infrared objects off the Galactic plane by the SiO $J=1$--0 $v=1$ or 2 lines. Some of these objects exhibit large negative radial velocities particularly at $l=20$ -- $40^{\\circ}$, where the Galactic rotation should give positive ones. Their distribution is scattered in the latitude range $\\Delta b \\sim 60^{\\circ}$. This negative velocity group of stars spreads between 0.3 kpc to 6 kpc in distance. It is possible to interpret that the brightest part of this deviant group is the Hercules stream of stars found in the solar neighborhood, and slightly distant part of this group as a part of outward flow found in short-period mira, both of which have been explained by the resonance effect of the Bulge bar. Though our simple calculation of the velocity field based on weak bar theory cannot fit the velocities of the nearest group of selected stars, it successfully explains the large negative-velocity stars located between the outer Lindblad and corotation resonances. We have also shown that the resonant coupling due to the periodic perturbation of the Bulge bar can create the star motion perpendicular to the Galactic plane near the corotation resonance. These facts strongly suggest that the deviant group of stars is produced by the gravitational perturbation of the Bulge bar.  \\  % Authors thank ....  for reading the manuscripts. We thank Dr. Tsuyoshi Sakamoto for reading the manuscripts and useful comments. This research made use of the SIMBAD and VizieR databases operated at CDS, Strasbourg, France, and as well as use of data products from Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and National Science foundation. %, and from the Midcourse Space %Experiment at NASA/ IPAC Infrared Science Archive, which is operated by the %Jet Propulsion Laboratory, California Institute of Technology, %under contract with the National Aeronautics and Space Administration.  %%%%%%%%%%%%%%%%%%%%%%% figure 11 %%%%%%%%%%%%%%%%%%%%%%%% % \\setcounter{figure}{11} % \\clearpage \\begin{figure*} \\begin{center} \\FigureFile(170mm,70mm){fig11.eps} %%% \\FigureFile(width,height){filename} \\end{center} \\caption{H$_2$O maser spectra for the detected sources. }\\label{fig: fig11} \\end{figure*}   %%%%%%%%%%  Appendix %%%%%%%%%%%%%"
    },
    "1004/1004.5173_arXiv.txt": {
        "abstract": "We have derived the temporal power spectra of the horizontal velocity of the solar photosphere. The data sets for 14 quiet regions observed with the \\textit{G}-band filter of \\textit{Hinode}/SOT are analyzed to measure the temporal fluctuation of the horizontal velocity by using the local correlation tracking (LCT) method. Among the high resolution ($\\sim$0\\farcs2) and seeing-free data sets of \\textit{Hinode}/SOT, we selected the observations whose duration is longer than 70 minutes and cadence is about 30 s. The so-called $k$-$\\omega$ diagrams of the photospheric horizontal velocity are derived for the first time to investigate the temporal evolution of convection. The power spectra derived from $k$-$\\omega$ diagrams typically have a double power law shape bent over at a frequency of 4.7 mHz. The power law index in the high frequency range is -2.4 while the power law index in the low frequency range is -0.6. The root mean square of the horizontal speed is about 1.1 km s$^{-1}$ when we use a tracer size of 0\\farcs4 in LCT method. Autocorrelation functions of intensity fluctuation, horizontal velocity, and its spatial derivatives are also derived in order to measure the correlation time of the stochastic photospheric motion. Since one of possible energy sources of the coronal heating is the photospheric convection, the power spectra derived in the present study will be of high value to quantitatively justify various coronal heating models. ",
        "introduction": "At least three types of convection are known to ubiquitously exist at the solar surface: granules, mesogranules, and supergranules. Their properties such as size, lifetimes, shapes have been studied by many authors (see reviews by, e.g., Leighton 1963, Spruit et al. 1990, Nordlund et al. 2009 and references therein). Through the magnetic field, the convective energy is transported upward to supply the sufficient energy to maintain the 1 MK corona. Therefore, the dynamics of the convective motions plays a key role in the so-called coronal heating problem, one of the most important issues in the solar physics. However, the temporal evolution of convection motion have not been sufficiently elucidated.   The vertical motion in convection is easily observed by using Doppler analysis. This motion generates compressible waves such as magnetohydrodynamic slow mode waves or fast mode waves. Although a lot of energy can be transported upward by the waves caused by convection, compressible waves are not considered to contribute to coronal heating, because the wave energy will be significantly reduced before reaching the corona by shock dissipation or reflection in the chromosphere \\citep{holl81}. On the other hand, the horizontal motion in convection plays an important role in the coronal heating. The interaction between magnetic field and the convection generates Alfv\\'{e}n waves \\citep{uchi74} while the stochastic photospheric motion braids the field lines to store the energy in the corona as electric current \\citep{park83}. Due to the lack of observations, it has been poorly understood what mechanisms contribute to coronal heating.  The local correlation tracking (LCT) method commonly is used to derive the horizontal velocity field \\citep{nove88,berg98}. Since LCT method uses apparent motion of granules to derive the velocity, it is better to use high spatial resolution and seeing-free data sets. From the horizontal velocity derived by LCT method, mesogranules and supergranules can be observed (e.g. Kitai et al. 1997, Ueno \\& Kitai 1998, Shine et al. 2000). Although various studies reveal the frequency distribution \\citep{titl89,berg98} and the spatial power spectra \\citep{rieu00, rieu08} of horizontal velocity fields, few studies explicitly mention the temporal power spectra \\citep{tarb90}. In a previous study, we derived the photospheric horizontal velocity in 14 different quiet regions to justify the Alfv\\'{e}n wave model for coronal heating \\citep{mats10}. In this study, we will show the temporal evolution of the horizontal motion of the photosphere in more detail. When discerning between various coronal heating models (see reviews by, e.g.,  Mandrini et al. 2000, Aschwanden et al. 2001, Klimchuk 2006, and references therein), the power spectra derived in the present study will be of high value. ",
        "conclusions": "We derived the apparent velocity in \\textit{G}-band data from \\textit{Hinode}/SOT by using the LCT method. Figure \\ref{kwdiagram} shows the so-called $k$-$\\omega$ diagram of the horizontal velocity normalized by the maximum power. The time series of LCT flow maps were Fourier transformed to estimate the power spectra. The dotted line in the figure represents $\\omega / k=5~\\mathrm{km~s}^{-1}$. Since we use the subsonic filtering method, spectral power above the dotted line is significantly reduced. The region in the lower frequency and the smaller wave number has greater power.  Integrating $k$-$\\omega$ diagram over wave number space, we can estimate the frequency power spectrum density, $P_{\\nu}$. The power spectrum density is defined as, \\begin{equation} <V_{\\perp}^2> = \\int _{\\nu_{min}} ^{\\nu_{max}}~P_{\\nu}~d\\nu \\label{eq_power}, \\end{equation} where $V_{\\perp}$ is the LCT velocity, $\\nu_{min}$ is the lowest frequency determined by total duration, and $\\nu_{max}$ is the highest frequency coming from the observational sampling time. The symbol, $<>$, denotes the temporal average. Figure \\ref{tspectrum} shows an example of the power spectrum density. Power spectra can be fitted to double power-law function, $P_\\nu \\propto \\nu^{\\alpha_L}$ when $\\nu < \\nu_b$ (so-called break frequency) and $P_\\nu \\propto \\nu^{\\alpha_H}$ when $\\nu \\ge \\nu_b$. Typically, $\\sqrt{<V_{\\perp}>},\\nu_b, \\alpha_L, $ and $ \\alpha_H$ of LCT velocity are 1.1 km s$^{-1}$, 4.7 mHz, -0.6, and -2.4, respectively in this work. By using the same analysis, temporal power spectra of divergence and rotation of velocity fields are also estimated (the dashed line and the dotted line in figure \\ref{tspectrum}). The total power, $\\alpha_L$, and $\\alpha_H$ of the divergence is 8.4 $\\times$ 10$^{-3}$ s$^{-1}$, -0.2, -1.5 respectively and the divergence will contribute to produce the acoustic mode waves. The total power, $\\alpha_L$, and $\\alpha_H$ of the rotation is 6.8 $\\times$ 10$^{-3}$ s$^{-1}$, -0.2, -1.3 respectively and the rotation is related to the generation of torsional Alfv\\'{e}n waves. The power spectra of the LCT velocity has a harder power law index than that of its derivatives ($\\nabla \\cdot V, \\nabla \\times V$).   The resulting power spectra are quite essential for evaluating models of coronal heating by Alfv\\'{e}n waves \\citep{holl82,kudo99,suzu05,suzu06,mats10}. In their numerical simulations, nonlinear Alfv\\'{e}n waves are driven by photospheric convection to heat the corona. \\cite{mats10} include the observed temporal power spectra shown here to generate Alfv\\'{e}n waves and succeed in producing sufficient energy flux to heat the corona.   Next, in order to derive the``correlation time\" of the stochastic convection motion, we derived the autocorrelation function (ACF), \\begin{eqnarray} R(\\tau) = {\\left< f(t) f(t+\\tau) \\right> \\over <f(t)^2>}\\label{eq_acf} \\end{eqnarray} , where $f$ is one of physical variables associated with turbulent motion and $<>$ denotes temporal averaging process. The correlation time is determined as the e-folding time of the ACF. Correlation time often scales with coronal heating rate in various heating models that use stochastic convection motion \\citep{stur81, park88, gals96}. Figure \\ref{crtime} shows the ACF of the \\textit{G}-band intensity fluctuation (solid line), horizontal velocity (dashed line), divergence (dotted line), and rotation (dash-dotted line). For \\textit{G}-band intensity, the correlation time is about 250 sec, which is comparable with the value in \\cite{titl89} and references therein. The horizontal velocity has shorter correlation time ($\\sim$ 100 sec) than the \\textit{G}-band intensity fluctuation. \\cite{rieu00} also derived the ACF of the horizontal velocity and found a relatively larger correlation time of about 35-45 min, because they were looking at larger spatial/temporal scale phenomena. ACFs of rotation and divergence fall rapidly and their correlation times are about 50 sec. The correlation time in the coronal heating models is usually defined by the velocity correlation time (e.g. Sturrock \\& Uchida 1981). However, the lifetime of granules, which corresponds to the intensity correlation time, is usually used as the correlation time in the previous studies. Since the velocity correlation time turned out to be shorter than the intensity correlation time, the coronal heating rate will be greatly reduced.  Let us explain why the correlation time of intensity fluctuation ($\\tau _I$) is longer than that of velocity ($\\tau _V$) qualitatively. The power spectrum of intensity ($P_I$) can be regarded as the power spectrum of spatial scale. From the dimensional analysis, the power spectrum of velocity ($P_V$) is proportional to $P_I \\nu ^2$ so that $P_I$ has a spectrum that is at most a power 2 softer than $P_V$. Since Fourier transform of ACF corresponds to power spectrum, softening of the power spectrum means hardening of the ACF decay or long correlation time. Therefore it is reasonable that $\\tau_I$ is longer than $\\tau_V$. From the similar dimensional analysis, the shorter correlation time of divergence or rotation can be explained.  If we integrate $k$-$\\omega$ diagram over frequency space, the spatial power spectrum density can be estimated (Figure \\ref{turb}). The power law index for the high $k$ range ($k \\sim 10^{-3}$) becomes -5/3, which represents the Kolmogorov type turbulence. Our results are complementary to the results of \\cite{rieu08} since they investigate large spatial structures. \\cite{espa93} showed the same power law index in the power spectra of vertical photospheric velocity. These results indicate that the granules are Kolmogorov turbulent eddies.  We also analyze motion of test particles in the LCT flow field. From the spatial deviation of test particles ($\\delta$) and the elapsed time ($\\tau$), diffusion coefficient ($D_{ph} \\equiv \\delta ^2 / 2 \\tau$) can be derived, if we assume motion of each test particle can be approximated by random walk. For our data set, the diffusion coefficient $D_{ph}$ is around 500 km$^2$ s$^{-1}$, which is comparable with that of \\cite{tarb90}. When we apply this value to the electric current cascading coronal heating model \\citep{ball86}, heating rate that scales with the diffusion coefficient of the photospheric horizontal motion becomes less than half of the required rate. Therefore when we consider coronal heating in the quiet sun, current cascading may not contribute significantly to the heating rate.  Since the LCT method tracks the apparent motion of the photospheric convection, the velocity derived here would contain some artifacts from the apparent velocity even though we have carefully removed the effect of acoustic waves by subsonic filtering method. The decrease in density or increase in temperature causes the destruction of CH molecules, resulting in \\textit{G}-band intensity changes \\citep{steiner01}. These intensity changes can affect the LCT velocity. It is important to compare the result of LCT measurements with the recent realistic 3D simulations \\citep{stei98,rieu01,geor06,geor07}.  In conclusion, $k$-$\\omega$ diagrams of LCT velocity are derived for the first time to investigate the temporal evolution of the photospheric convection. Integrating the power spectra over wave number space generally reveals a double power law spectral shape whose break frequency is about 4.7 mHz. The power law index in the low frequency range is -0.6 while the power law index in the high frequency range is -2.4. The correlation time of \\textit{G}-band intensity and horizontal velocity is about 200 sec and 100 sec respectively while that of divergence and rotation is 50 sec. The diffusion coefficient of photospheric convection is around 500 km$^2$ s$^{-1}$, which is not sufficient to heat the corona above quiet region through current cascading. By using the power spectra derived in the present study, some of the coronal heating models will be justified (e.g. Matsumoto \\& Shibata 2010)."
    },
    "1004/1004.0756_arXiv.txt": {
        "abstract": "{ We present chemical evolution models  of the Galactic disk with different $Z$-dependent yields. We find that a moderate mass loss rate for massive stars of solar metallicity produces an excellent fit to the observed C/H and C/O gradients of the Galactic disk. The best model also fits: the H, He, C, and O abundances derived from recombination lines of M17, the protosolar abundances, and the C/O-O/H, C/Fe-Fe/H, and O/Fe-Fe/H relations derived from solar vicinity stars. The agreement of the model with the protosolar abundances implies that the Sun originated at a galactocentric distance similar to the one it has. Our model for $r=3$ kpc implies that a fraction of the stars in the direction of the bulge formed in the inner disc. We obtain a good agreement between our model and the C/O versus O/H relationship derived from extragalactic H~{\\sc ii} regions in spiral galaxies. }   \\addkeyword{galaxies: abundances} \\addkeyword{galaxy: bulge} \\addkeyword{galaxies: evolution} \\addkeyword{H~II regions: M17, Orion nebula } \\addkeyword{ISM: abundances} \\addkeyword{Sun: abundances}  \\begin{document} ",
        "introduction": "\\label{sec:intro}  The comparison of detailed Galactic chemical evolution models, GCE models, with accurate abundance determinations of stars and gaseous nebulae provides a powerful tool to test the chemical evolution models and the accuracy of observational abundance determinations of stars of different ages and of H{\\sc~ii} regions located at different galactocentric distances.  In this paper we will compare our models with stellar and H{\\sc~ii} regions abundances to test if the H{\\sc~ii} region abundances derived from recombination lines agree with the stellar abundances, in particular with the protosolar abundances that correspond to those present in the interstellar medium 4.5 Gyr ago.  Also our GCE models can be used to constrain the C yields for massive stars, the C yield is not well known and we will vary it to obtain the best fit between our GCE models and the observational data.  Carigi and Peimbert (2008, hereinafter Paper I) presented chemical evolution models of the Galactic disk for two sets of stellar yields that provided good fits to: a) the O/H and C/H gradients (slope and absolute value) derived from H{\\sc~ii} regions based on recombination lines (Esteban et al. 2005) and including the dust contribution \\citep{est98}, and b) the $\\Delta Y/\\Delta Z$ value derived from the Galactic H{\\sc~ii} region M17, and the primordial helium abundance, $Y_p$ obtained from metal poor extragalactic H{\\sc~ii} regions (Peimbert et al. 2007). In Paper I, based on our GCE models and combined with the constraints available, we were not able to discriminate between the stellar evolution models assuming high wind yields for massive stars, HWY,  and those assuming low wind yields for massive stars, LWY.  Previous works have focused on the test of stellar yields using GCE models constrained by chemical gradients obtained by different methods, gradients that in general show similar slopes but a considerable spread in the absolute O/H ratios, e.g. \\citet{pra94,car96,chi03a,rom10}. To test the stellar yields it is necessary to have good absolute abundance values of stars and H{\\sc~ii} regions.  To determine the O/H abundances of H{\\sc~ii} regions in irregular and spiral galaxies many methods have been used in the literature. Most of them have been based on fitting photoionized models to observations or by determining the electron temperature, $T$, from the 4363/5007 [O III] ratio directly from observations. A comparison of  many of the different methods used has been made by \\citet{kew08}. They find that the O/H differences derived by different methods between two given H{\\sc~ii} regions amount to $0.10 - 0.15$ dex.  Alternatively for all the methods the absolute difference for a given H{\\sc~ii} region is considerably larger reaching values of 0.7 dex for extreme cases (see Figure 2 in their paper and the associated discussion).  Most of the differences among the various calibrations are due to the temperature distribution inside the nebulae. In this paper we will use only abundances of H{\\sc~ii} regions based on recombination lines of H, He, C, and O, these lines depend weakly on the electron temperature, they are roughly proportional to $1/T$, therefore the relative abundances among these four elements are practically independent of the electron temperature.   There are two frequently used methods to derive C and O gaseous abundances from H{\\sc~ii} regions: a) the most popular one based on collisionally excited lines (or forbidden lines) and the  $T$(4363/5007) temperatures, the FL method, and b) the one based on C and O recombination lines, the RL method. The RL method produces gaseous O and C abundances higher by about 0.15 to 0.35 dex than the FL method.  The RL method is almost independent of the electron temperature, while the FL method is strongly dependent on the electron temperature. It is possible to increase the FL abundances under the assumption of temperature inhomogeneities to reach agreement with the RL values. The temperature distribution can be characterized by the average temperature, $T_0$, and the mean square temperature variation, $t^2$, \\citep[e.g.][]{pei67,pei02}. The $t^2$ values needed to reach agreement between the RL and the FL abundances are in the 0.02 to 0.05 range, while the photoionization models predict typically $t^2$ values in the 0.003 to 0.01 range, this discrepancy needs to be sorted out \\citep[e.g.][and references therein]{pei11}.  Paper I is controversial because the C/H and O/H gaseous abundances of the H{\\sc~ii} regions have been derived from recombination lines (that is equivalent to the use of  $t^2 \\neq 0.000$ and forbidden C and O lines) and the assumption that 20\\% of the O atoms and 25 \\% of the C atoms are trapped in dust grains (0.08 dex and 0.10 dex respectively). These assumptions increase the O/H ratio by about 0.25 to 0.45 dex relative to the gaseous abundances derived from $T$(4363/5007),  the forbidden O and C lines, and the assumption that $t^2 = 0.00$.  Due to the controversial nature of the  H{\\sc~ii} region abundances used in Paper I and that we were not able to discriminate between the two sets of stellar yields adopted we decided to test our GCE models further by including additional observational constrains: a) the Asplund et al. (2009) protosolar abundances that provide us with the O/H, C/H, Fe/H, and $\\Delta Y/\\Delta O$ in the interstellar medium 4.5 Gyr ago, b) the C/H, O/H, and Fe/H by \\citet{ben06} for young F and G stars of the solar vicinity, c) the O/H, C/H, and He/H derived from B stars by \\citet{prz08}, and d) throughout this paper for all the H{\\sc~ii} regions we will use abundances derived from recombination lines  and to obtain the total abundances we will increase the gaseous abundances by 0.10 dex in C and 0.12 dex in O to take into account the fraction of atoms trapped in dust grains, with the exception of the metal poor irregular galaxies for which we will use an 0.10 dex depletion for O \\citep{est98,mes09,pea10}.   We also decided to compare our best models with the C/O versus O/H results derived by  \\citet{est02,est09} from bright H{\\sc~ii} regions in nearby spiral galaxies based on recombination lines and to make a preliminary discussion of a comparison between our models for $r$ = 3kpc and the stars in the direction of the galactic bulge obtained by \\citet{ben10a} and \\citet{zoc08}.  The symbols  $C$, $O$, $X$, $Y$, and $Z$ represent carbon, oxygen, hydrogen, helium, and heavy element abundances by unit mass respectively; while C/H, O/H, Fe/H,  C/O, C/Fe, and O/Fe represent the abundance ratios by number.  In Section~\\ref{sec:models} we discuss the general properties of the chemical evolution models, we discuss infall models for the Galaxy with two sets of stellar yields, the HWY and the LWY. In Section~\\ref{sec:gradient} we show the prediction of the current abundance gradients for the interstellar medium (ISM) of the Galactic disk and compare them with Galactic H{\\sc~ii} region abundances derived from recombination lines that include the dust correction, in addition for the solar vicinity we present the chemical history of the ISM and compare it with the chemical abundances of stars of different ages; we define the solar vicinity as a cylinder perpendicular to the galactic plane, centered in the Sun, with a radius of 0.5 kpc, that extends into the halo to include the stars in the cylinder.  In Section~\\ref{sec:M17} we compare our chemical evolution models with the protosolar chemical abundances and with those of the H{\\sc~ii} region M17. Based on the comparison between the observations and the models in Section~\\ref{sec:intermediate} we present a new Galactic chemical evolution model, with intermediate mass loss due to interstellar winds, IWY, that produces considerably better adjustments with the observations. In Section~\\ref{sec:other} we compare the IWY model with additional observations, those provided by extragalactic H{\\sc~ii} regions in spiral galaxies, and those provided by stars in the direction of the Galactic bulge.  The conclusions are presented in Section~\\ref{sec:conclusions}.  A preliminary account of some of the results included in this paper was presented elsewhere \\citet{pei10}. ",
        "conclusions": "\\label{sec:conclusions}  We have made models with three different sets of yields that differ only on their $Z$ dependence at solar metallicities for massive stars: HWY, IWY, and LWY.  The HWY and the LWY have been used before by us, while in this paper we introduce the IWY, that are given by (HWY + LWY)/2. We find that the IWY galactic chemical evolution models produce better fits to the observational data than either the HWY or the LWY galactic chemical evolution models.  We present a Galactic chemical evolution model based on the IWY for the disk of the Galaxy that is able to fit: a) the C/O vs O/H, C/Fe vs Fe/H, O/Fe vs Fe/H, and Fe/H vs $t$ relations derived from halo and disk stars of different ages in the solar vicinity, b) the O/H, C/H, and C/O abundance gradients (slopes and absolute values) derived from Galactic H~{\\sc{ii}} regions, c) the He/H, C/H, O/H, Fe/H protosolar abundances, and d) the He/H and O/H values of the galactic H~{\\sc{ii}} region M17.  We find that in general about half of the freshly made helium is produced by massive stars and half by LIMS, and that a similar situation prevails for carbon, while  most of the oxygen is produced by massive stars. The agreement of the He/O and C/O ratios between the model and the protosolar abundances implies that the Sun formed from a well mixed ISM.  We note that the agreement of our model with the protosolar abundances and the Sun-formation time supports the idea that the Sun originated at a galactocentric distance similar to that of the solar vicinity.  We show that chemical evolution models for the Galactic disk are able to reproduce the observed $\\Delta Y$ and $\\Delta O$ protosolar values and the $\\Delta Y$ and $\\Delta O$  values derived for M17 based on H, He and O recombination lines, but not the M17 $\\Delta Y$ and $\\Delta O$ values derived from $T$(4363/5007) and O collisionally excited lines under the assumption of $t^2 = 0.00$. This result provides a consistency check in favor of the presence of large temperature variations in H~{\\sc{ii}} regions and on the method based on the H, He, C and O recombination lines to derive abundances in H~{\\sc{ii}} regions.  We obtain that the IWY chemical evolution model of the Galactic disk for the present time, in the galactocentric range $ 6 < r$(kpc) $<$ 11, produces a reasonable fit to the O/H vs C/O relationship derived from H~{\\sc{ii}} regions of nearby spiral galaxies. The yields predict an increase of the C/O ratio with O/H starting from 12+log(O/H)$\\sim$8.4 that is observed in our Galaxy and in nearby galaxies. The O/H vs C/O relationship might imply  that spiral galaxies have a similar IMF, no selective outflows, and probably a formation scenario similar to that of our galaxy.  We find a remarkable parallelism of the O/H gradients for M33, M101, and the Galaxy when they are plotted with respect to $r/R_0$ ($R_0$ is the galactic photometric radius to 25 mag per square second), suggesting some common mechanisms in the formation and evolution of spiral galaxies. The O/H ratio at a given $r/R_0$ differs by a constant among M33, M101 and the MW.  The MW shows a 0.24 dex higher O/H ratio than the average of M33 and M101 at a given $r/R_0$.  We also find that the results for our model at $r=3$ kpc can explain: a) the C/O-O/H and O/Fe-Fe/H, relations, and b) partially the Fe/H-time relation and the Fe distribution function derived from stellar observations in the direction of the Galactic bulge. We find that stars belonging to the thin and thick discs make a significant contribution to these relations.  Future work to advance in this subject requires: a) to advance in the study of the galactic bulge to be able to quantify the stellar contributions due to the inner disk and the true bulge, b) to increase the H~{\\sc{ii}} regions $r/R_0$ coverage, mainly in the shorter galactocentric distance, with high accuracy C/H and O/H abundance determinations; c) to increase the sample of spiral galaxies of different types with O/H gradients of high accuracy, and d) to make specific models for each galaxy.  All in order to sort out possible differences in the galactic formation and evolution of other galaxies relative to that of the Galaxy."
    },
    "1004/1004.3409.txt": {
        "abstract": "We made deep near-infrared ($JHK$s) imaging polarimetry toward the Serpens cloud core, which is a nearby, active cluster forming region. The polarization vector maps show that the near-infrared reflection light in this region mainly originates from SVS2 and SVS20, and enable us to detect 24 small infrared reflection nebulae associated with YSOs. Polarization measurements of near-infrared point sources indicate an hourglass-shaped magnetic field, of which symmetry axis is nearly perpendicular to the elongation of the C$^{18}$O ($J=1-0$)  or submillimeter continuum emission. The bright part of C$^{18}$O ($J=1-0$),  submillimeter continuum cores as well as many class 0/I objects are located just toward the constriction region of the hourglass-shaped magnetic field. Applying the Chandrasekhar \\& Fermi method and taking into account the recent study on the signal integration effect for the dispersion component of the magnetic field, the magnetic field strength was estimated to be $\\sim$100 $\\mu$G, suggesting that the ambient region of the Serpens cloud core is moderately magnetically supercritical. These suggest that the Serpens cloud core first contracted along the magnetic field to be an elongated cloud, which is perpendicular to the magnetic field, and that then the central part contracted cross the magnetic field due to the high density in the central region of the cloud core, where star formation is actively continuing.  Comparison of this magnetic field with the previous observations of molecular gas and large-scale outflows suggests a possibility that the cloud dynamics is controlled by the magnetic field, protostellar outflows and gravitational inflows. Furthermore, the outflow energy injection rate appears to be larger than the dissipation rate of the turbulent energy in this cloud, indicating that the outflows are the main source of turbulence and that the magnetic field plays an important role both in allowing the outflow energy to escape from the central region of the cloud core and enabling the gravitational inflows from the ambient region to the central region. These characteristics appear to be in good agreement with the outflow-driven turbulence model and imply the importance of the magnetic field to continuous star formation in the center region of the cluster forming region. ",
        "introduction": "Stars are formed by gravitation in molecular clouds having both turbulence and  magnetic fields in the Galaxy,  and most of stars are thought to be formed in clusters \\citep[e.g.,][]{lad03, al07}. A mass spectrum of prestellar condensations  is reported to have the power similar to that of the stellar IMF both in dust continuum observations \\citep[][and references therein]{re06} and molecular-line observations \\citep[e.g.,][]{ik07}, and theoretical studies of turbulent molecular clouds \\citep[][and subsequent works]{kl98} suggest that these condensations were formed through turbulent shock. One of  the most promising sources of ordinary turbulence is outflows from protostars, which are ubiquitous in star forming regions and are believed  to be formed through the mediation of magnetic field. Magnetic fields are also considered to play an important role in dynamical evolution of molecular clouds and control of star formation, i.e., formation of molecular cloud cores and their collapse \\citep[e.g.,][]{mck07}.  Recently, \\cite{li06} and \\cite{nakam07} presented realistic 3D MHD simulations of cluster formation, taking into account the effect of protostellar outflows as well as initial turbulence and a magnetic field. In their simulations, they indicated that the initial turbulence is quickly replaced by turbulence generated by protostellar outflows, keeping the quasi-equilibrium state with a slow rate  of star formation, and that magnetic fields are dynamically important if their initial strengths are not far below the critical value for static cloud support because of  the amplification by the outflow-driven turbulent motions. The magnetic field is expected to influence the directions of outflow ejection and propagation and the transmission of outflow energy and momentum to the ambient medium. However, the magnetic field structures have not always been observationally clear in/around cluster forming regions, particularly around nearby cluster forming regions because of the lack of deep, wide-field near-infrared (NIR) polarimetry data.  The Serpens cloud core is one of the nearby\\footnote{We assume a distance of $\\sim$260 pc for the Serpens cloud, following the most of the recent papers on the Serpens cloud and based on the discussion of \\cite{st03} on the center distance of the Aquila Rift system.}, active  low-mass star forming regions at the northern part of the Serpens cloud and many observational works have been done \\citep[][and references therein]{ei08}. Recent mid-IR studies \\citep[e.g.,][]{ka04, ha07, wi07} revealed that a lot of embedded young stellar objects (YSOs), including Class 0/I objects, are located toward an aggregate of (sub)millimeter dust continuum cores \\citep[e.g.,][]{da99, ka04, en07}, which consists of two sub-clumps \\citep[NW and SE sub-clumps;][]{ol02} in the central region and is enveloped by ambient molecular gas \\citep[e.g. $^{13}$CO, and C$^{18}$O;][]{mc00, ol02}.  Many outflow activities that are related to star formation have been  taking place in the Serpens cloud core. CO high velocity flows are reported to be widely spread over the cloud core \\citep[e.g., ][]{wh95, da99, nar02}. Compact molecular outflows of higher density tracers and H$_2$ jet-like knots are associated with the submillimeter cores \\citep[e.g.,][]{cu96, her97, wo98, ho99, wil00}. The direction of these compact outflows was reported to be PA$\\sim155\\degr$ on an average with deviation of  a few $10\\degr$ \\citep[see Table 5 of][]{ol02}, which is nearly parallel to the alignment direction, from NW to SE, of the two sub-clumps \\citep[e.g.,][]{da99, ka04, en07}, or to the cloud elongation in $^{13}$CO, C$^{18}$O, and other higher density tracers \\citep[e.g. ][]{mc00, ol02}. In addition, \\cite{da99} and \\cite{zi99} found, through optical narrow-band imaging, that many HH objects emanate from the two sub-clumps to the ambient region, penetrating the dense part of the central region.  The Serpens Reflection Nebula (SRN) illuminated by SVS 2 \\citep{st76} has been extensively studied by polarimetric measurements both in  optical and near-infrared (NIR)  wavelengths  \\citep{ki83, wa87, go88, so97, hu97}. NIR polarimetric measurements \\citep{so97, hu97} probed also some other obscured reflection nebulae around SVS 2 in detail. \\cite{go88} suggested the magnetic field of a NW-SE direction based on the elongation of the reflection nebulae around several YSOs in the central region of Serpens cloud core. In contrast, the NIR polarization measurement of a background star candidate suggested rather different direction of magnetic field because its polarization angle was nearly perpendicular to the NW-SE direction  \\citep{so97}. However, this measurement was only for one background candidate, which in fact has a possibility of being a YSO in the Serpens cloud core and its polarization originating from the YSO itself. Therefore, it is vitally important to measure more background stars to resolve this discrepancy and to know the magnetic field structure toward the Serpens cloud core.  We conducted deep, $JHK$s imaging polarimetry of the Serpens cloud core to reveal the magnetic field structure in this region. We also searched for more NIR reflection nebulae associated with YSOs. Here, we present the results of our imaging polarimetry in the Serpens cloud core by comparing the data from the previous observations and discuss the role of the magnetic field in this region. ",
        "conclusions": "\\subsection{Shape of the magnetic field}  We have modeled the shape of the magnetic field with the polarization vectors measured at $H$ for point sources, following \\cite{gi06} and \\cite{kan09}. The magnetic field was fitted with a parabolic function of $x=g+gCy^2$, with a counterclockwise tilted $y$-axis (the parabolic magnetic field axis of symmetry) by $\\theta_{\\rm PA}$ and a symmetric center $(x, y)_{\\rm center}$, where the $y$ is the distance from the horizontal axis ($x=0$) and the $x$ is the distance from the parabolic magnetic field axis of symmetry. The value of $\\tan^{-1}(dy/dx)+90\\degr$ corresponds to the position angle of the polarization ($\\theta$). Only the point sources, except YSOs, having $P/{\\mathit \\Delta}P > 3$ and $P < 6.2 (H-Ks)$, were used for the fitting. The error of the polarization angle (${\\mathit \\Delta}\\theta$) was used to compute a weight for the datum, $1/({\\mathit \\Delta}\\theta)^2$.  In Figure \\ref{f7}, the best-fit magnetic field is shown as well as the measured polarization vectors for 149 sources. The position angle of the parabolic magnetic field axis of symmetry is $\\sim70\\degr$, and the coefficient $C$ of $y^2$ is $\\sim7.1\\times 10^{-6}$ pixel$^{-2}$. The root mean square  (r.m.s.) of the residuals is $\\sim22\\degr$.  We executed one-parameter fitting of the magnetic field in local areas, in order to more accurately calculate the r.m.s.  of the residuals, with the same $\\theta_{\\rm PA}$ and $(x, y)_{\\rm center}$ obtained in the global fitting above. We selected three corners and one more area of the image where the source density is relatively high and/or the magnetic field seems to be rather ordered (areas outlined by dashed boxes in Figure \\ref{f7}). Toward the SE corner of the image ( $x<400$ and $y<400$ in Figure \\ref{f7}; 30 sources), the coefficient $C$ of $y^2$ was determined to be $(7.99\\pm 0.76)\\times10^{-6}$ pixel$^{-2}$, similar to that of the global fitting, and the r.m.s. of the residual was calculated to be $12.9\\pm 0.9\\degr$, and toward the SW corner ($x>500$ and $y<230$; 20 sources), $C=(7.52\\pm 1.00)\\times 10^{-6}$ pixel$^{-2}$ and r.m.s. = $27.0\\pm2.0\\degr$ were obtained. Removing the dispersion due to the measurement uncertainties of the polarization angles $4.2\\pm3.0\\degr$ and $3.2\\pm 2.2\\degr$, we obtained the dispersions from the best-fit model, $12.2\\pm1.4\\degr$and $26.8\\pm2.0\\degr$ for the SE and SW corners, respectively. Toward the NE corner  ($x<300$ and $y>800$; 18 sources), $C=(3.36\\pm 0.92)\\times 10^{-6}$ pixel$^{-2}$ and r.m.s. = $14.8\\pm 1.6\\degr$ were evaluated, and the intrinsic dispersion from our model of $13.7\\pm2.0\\degr$ was obtained with the measurement uncertainty of $5.6\\pm2.4\\degr$. This smaller $C$ indicates that the curvature of the magnetic field here is rather looser than that expected from the global fitting, i.e., slightly bended to the direction parallel to the symmetry axis of the magnetic field. Toward the area next to the NE corner ($400<x<700$ and $y>800$; 17 sources), $C=(6.85\\pm 0.55)\\times 10^{-6}$ pixel$^{-2}$ and r.m.s. = $13.3\\pm1.4\\degr$ were evaluated, and the intrinsic dispersion of $12.6\\pm 1.8\\degr$ was obtained with the measurement uncertainty of $4.2\\pm 2.9\\degr$.  \\subsection{Comparison of the magnetic field with the submillimeter and millimeter data}  \\subsubsection{850 $\\micron$ continuum}  We compare our $H$-band measured polarization vectors and the modeled magnetic field with the 850 $\\micron$ dust continuum map of \\cite{da99} in Figure \\ref{f8}. Note that the green lines of this figure do not present lines of magnetic force, just the direction of the magnetic field.  The high intensity ridge of  the 850 $\\micron$ continuum is elongated along the NW-SE direction, having two sub-clumps (NW and SE sub-clumps), both of which consist of several dense cores (e.g., SMM 1--11, S68Nb--d, and PS2 in Figure \\ref{f8}). This distribution of the 850 $\\micron$ continuum is very similar to that of the bright parts of the $^{13}$CO($J=1-0$) and C$^{18}$O($J=1-0$) emission \\citep{mc00, ol02}, although the global distribution of the $^{13}$CO($J=1-0$) emission is not always elongated, but rather roundly extended \\citep{ol02}. It is evident that the symmetric axis ($y'$-axis) of the best-fit magnetic field with a parabolic function is nearly perpendicular to the elongation direction of these continuum and molecular line emissions. The horizontal axis ($x'$-axis) of the parabolic magnetic field is situated nearly along the  850 $\\micron$ continuum ridge, although there are some deviations of the continuum emission from the horizontal axis. The symmetric axis of the parabolic magnetic field runs through the northern part of the SE sub-clump, not through the middle point of the two sub-clumps, which looks like the center of gravity of the Serpens cloud core when we glance at the 850 $\\micron$ continuum map.  \\cite{da99} suggested the presence of three extended cavity-like structures to the east of SMM 3 (hereafter CLS 1), south-west of SMM 2 (hereafter CLS 2), and north-west of SMM 4 (hereafter CLS 3), which consist of three pairs of filaments that protrude the 850 $\\micron$ continuum ridge. They mentioned that these cavity structures (CLS 1--3) are probably shaped by outflows rather than by global cloud collapse along, say, magnetic field lines.  As is in Figure \\ref{f8}, the filaments to the north-east of SMM 3 and east of SMM 2 form CLS 1, those to the south-east of SMM2/PS2 and south of SMM11 form CLS 2, and those to west of SMM3 and east of SMM4 form CLS 3. It appears that the two filaments of CLS 1 jut almost along the magnetic field from the SE sub-clump and that the symmetry axis ($y'$-axis) of the magnetic field go through the inside of CLS 1 as well as CLS 3.  \\subsubsection{CO emission}  Here we compare our best-fit magnetic field with the $^{12}$CO $J=2-1$, $^{12}$CO $J=1-0$, $^{13}$CO $J=1-0$, and C$^{18}$O $J=1-0$ observations \\citep{wh95, da99, mc00, nar02, ol02}.  \\subsubsubsection{$^{12}$CO $J=2-1$}  As was mentioned above, the bright parts of the $^{13}$CO $J=1-0$ and C$^{18}$O $J=1-0$ emission maps are elongated and confined in the ridge, while the global distributions of $^{12}$CO $J=2-1$ and $^{13}$CO $J=1-0$, i.e., the low density molecular gas, are extended \\citep[e.g.,][]{wh95,da99,ol02}.  Figure \\ref{f9} presents our best-fit magnetic field superposed on the CO $J=2-1$ contour map and 850 $\\micron$ image of \\cite{da99}, where the CO $J=2-1$ map is considered to show the ambient molecular gas of the Serpens cloud core, but not the dense cores. \\cite{da99} mentioned that toward the two filaments of CLS 1 and one CLS 2 filament to the south-east of SMM2/PS2 the CO $J=2-1$ emission and 850 $\\micron$ continuum distributions coincide well. As mentioned above, the two filaments seem to run almost along the magnetic field, indicating that the CO $J=2-1$ filaments are also related with the magnetic field. For the CLS 2 filament to the south-east of SMM2/PS2, the same situation as the CLS 1 filaments may be also seen. Two other CO $J=2-1$ filaments/extensions to the north-west of SMM 9 and west of SMM 1 are also noticeable in Figure \\ref{f9}. Although considerable parts of these two filaments are out of our polarimetry image, the extrapolation of our best-fit magnetic field cloud predict that these two filaments run along the magnetic lines.  \\subsubsubsection{C$^{18}$O $J=1-0$ and $^{13}$CO $J=1-0$}  \\cite{mc00} showed that a velocity gradient running from a LSR velocity centroid of 9 km s$^{-1}$ at the north-west end of the C$^{18}$O $J=1-0$ emission to 7.5 km s$^{-1}$ at the south-east end \\citep[Figure 2 of][]{mc00}, i.e., along the elongation direction of C$^{18}$O. On the other hand, \\cite{ol02} suggested that the Serpens cloud exhibits a velocity gradient roughly from east to west, based on their model fitting of velocity gradients in C$^{18}$O $J=1-0$, $^{13}$CO $J=1-0$, C$^{34}$S $J=1-0$, adopting their map center, which is the middle point of the two sub-clumps, as the reference position for analysis. However, according to their channel and centroid velocity maps \\citep[Figures 7 and 8 of][]{ol02}, the bright parts of C$^{18}$O $J=1-0$ and $^{13}$CO $J=1-0$ are similar to that of \\cite{mc00}, and a steep velocity gradient from NW to SE almost along the normal line of the symmetry axis ($y'$-axis) of the magnetic field can be seen at just south of their reference position in $^{13}$CO, although at the reference position a velocity gradient from West to East is seen. It is surprising that the normal line of the steep velocity gradient almost coincides with the symmetry axis ($y'$-axis) of the magnetic field.  In summary, the direction of velocity gradient is nearly along the elongation of the Serpens cloud core and is nearly perpendicular to the symmetry axis of the magnetic field with a coincidence of the normal line of the steep velocity gradient and the axis of the magnetic field. It could be possible that this normal line of the velocity gradient is an axis of the global rotation of the Serpens cloud core if the real center of gravity of the Serpens cloud core is located on the symmetry axis of the magnetic field.  It is interesting to examine the presence of C$^{18}$O $J=1-0$ and $^{13}$CO $J=1-0$ features that coincide with the filaments of the CO $J=2-1$ emission and 850 $\\micron$ emission. In the C$^{18}$O $1-0$ integrated emission maps of \\cite{wh95}, \\cite{mc00} and \\cite{ol02}, a feature to the north-east of SMM 3 could coincide with one of the CSL 1 filament, but one to the east of SMM 2 is not clear. In the channel map of $^{13}$CO $J=1-0$ \\citep[Figure 7 of][]{ol02}, a filament feature to the east of SMM 2 is clearly visible in the blue-shifted emission at the panel of $V_{\\rm LSR}$=5--7.3 km s$^{-1}$. This filament looks likely to coincide with the CSL 1 filament to the east of SMM 2, but we can clearly recognize that it is located just outside this CSL 1 filament, i.e., between this CSL 1 filament and the CSL 2 filament to the south-east of SMM 2/PS2. At the same panel, a feature to the north and north-east of SMM 3 or near SMM 8 is also visible. This feature appears to be just outside the CLS 1 filament to the north-east of SMM 3. At the panel of $V_{\\rm LSR}$=8.4--12.4 km s$^{-1}$, a red-shifted feature that protrudes from the SE sub-clump is visible, but it is located toward the inside region of CLS 1. The presence of this red-shifted feature and the blue-shifted features are probably consistent with red-shifted velocity region that jut from the SE sub-clump and with blue-shifted regions toward both sides of this red-shifted region, respectively, in the $^{13}$CO $J=1-0$ centroid velocity map of \\cite{ol02}.  \\subsubsubsection{CO outflows}  \\cite{da99} presented the integrated intensity contours of CO $J=2-1$ blue- and red-shifted outflows \\citep[Figures 4 and 8 of][]{da99}. These figures imply that the 850 $\\micron$ filaments that coincide the CO $J=2-1$ filaments are shaped by outflows. On the basis of a fact that these filaments run along the magnetic field, the outflows that protruded from the ridge to its ambient are most likely to be guided by the magnetic field or to drag the magnetic field. The outflows may be guided by the magnetic field since the magnetic field seems to be strong enough to be ordered at least over our polarimetric imaging area.  The CLS 1 filaments are associated with red-shifted outflows, but no red-shifted CO $J=2-1$ emission is visible at the root of CLS 1. However, CO $J=1-0$ obervations \\citep{nar02} showed U-shaped, red-shifted high velocity flow at the root of CLS 1. This CO $J=1-0$ feature and our best-fit magnetic field support the idea of \\cite{da99} that the CLS 1 filaments of the CO $J=2-1$ and 850 $\\micron$ emission illustrate the action of a wide-angled wind powered by a source within the SMM 2/3/4 cluster, which has swept up gas and dust into a warm, compressed shell, although there is a possibility that the wind is powered by multiple sources within the cluster.  \\subsection{Magnetic field strength}  We try to make an evaluation of the magnetic field strength toward four areas where we calculated the angular dispersions (residuals) for our best-fit magnetic field, using the Chandrasekhar \\& Fermi (CF) method \\citep{cha53}. On the basis  of the conclusions of recent MHD studies that the introduction of a correction factor is needed for evaluating the plane-of-sky component of the magnetic field \\citep{os01,pa01,he01,ku03}, \\cite{hou04b} mentioned that a correction factor of $\\sim$0.5 is appropriate in most cases when the magnetic field is not too weak. Since the magnetic field seems to be ordered over the Serpens cloud core, the magnetic field is expected to be strong. Therefore we first adopt a correction factor of 0.5 to evaluate the magnetic field strength. We need the mass density and velocity dispersion of the matter coupled to the magnetic field to evaluate the magnetic field strength. Here, we use those estimated from the C$^{18}$O observation \\citep{ol02}.  Toward the four areas, the H$_2$ column densities from C$^{18}$O could be estimated to be $\\sim 6\\times10^{22}$ cm$^{-2}$ from Figure 11 of \\cite{ol02}. Adopting the approximate C$^{18}$O extent of $\\sim$12\\arcmin \\citep[$\\sim$0.9 pc at d$\\sim$260 pc; Figure 2 of ][]{ol02} as the depth of these area, we obtain the H$_2$ densities of $\\sim 2.1 \\times10^{4}$ cm$^{-3}$. From Figure 10 of \\cite{ol02}, the C$^{18}$O velocity widths could be estimated to be $\\sim$1.6--1.8 km s$^{-1}$ toward the SE and NE corners, and $\\sim$1.8--2.0 km s$^{-1}$ toward the area next to the NE corner. Toward the SW corner with a complex distribution of velocity width, the velocity width may be $\\sim$1--2 km s$^{-1}$. Using a mean molecular mass, $\\mu$, of 2.3 and these values to derive the velocity dispersions, we roughly evaluated the magnetic field strength of the plane-of-the-sky of $B_{\\parallel} \\sim$160--180 $\\mu$G toward the SE corner, $\\sim$150--160 $\\mu$G toward the NE corner, and $\\sim$180--200 $\\mu$G toward the area next to the NE corner. Although $\\sim$50--90 $\\mu$G can be evaluated toward the SW corner, this value might be more uncertain than those toward the other areas due to the larger uncertainty of the velocity width. The magnetic field strength evaluated here is higher than those measured around dark cloud complexes and prestellar cores, a few 10 $\\mu$G  \\citep[e.g.,][respectively]{alv08, kan09}, but smaller than those around HII regions, a few mG \\citep[e.g.,][]{hou04b} and of a protostellar envelope, a few mG \\citep[][]{gi06}.  Recently, \\cite{hou09} showed  how the signal integration through the thickness of the cloud and the area of the telescope beam affects on the measured angular dispersion and apply their results to OMC-1. Based on their estimated number (N=21) of the independent turbulent cells contained within the column probed by the telescope beam, they found that a correction factor of $1/\\sqrt{N} \\sim0.2$ is applicable to OMC-1. In our case, although the area of the telescope beam is negligibly small due to the point sources, the thickness of the cloud should be taken into account and the correction factor should be somewhat smaller than $\\sim0.5$. If we assume that the effect of the cloud thickness is similar to that of OMC-1, we obtain $N\\sim11$, suggesting a factor of $\\sim0.3$. Adopting this factor of $\\sim0.3$, the above estimated values are reduced by a factor of $\\sim0.6$ and $B_{\\parallel}\\sim100$ may be appropriate for the ambient region of the Serpens cloud core, except the SW corner.  Here, we roughly derive the mass to magnetic flux ratio $M_{\\rm cloud}/\\Psi$ using our estimated value of  $B \\sim$100 $\\mu$G, and compare it with the critical value for a magnetic stability of the cloud, ($M_{\\rm cloud}/\\Psi$)$_{\\rm critical}=(4\\pi^2 G)^{-1/2}$ \\citep{nak78}. With a formula  $M_{\\rm cloud}/\\Psi = (\\pi R^2 \\mu m_{\\rm H} N)/(\\pi R^2 B)=\\mu m_{\\rm H}N/B$ and the H$_2$ column density $N\\sim 6 \\times10^{22}$ cm$^{-2}$ where we estimated $B$, we derive $M_{\\rm cloud}/\\Psi \\sim3.8 \\times(M_{\\rm cloud}/\\Psi)_{\\rm critical}$, where $R$ is a radius of the cloud and $m_{\\rm H} $ is the mass of a hydrogen atom. Although this derived value is slightly larger than the critical value, $M_{\\rm cloud}/\\Psi$ could be much larger in the inner region of the cloud core because the column density of the inner region is much higher than those where we estimated $B$, but the magnetic field may be slightly larger than that we estimated in the ambient region, judged from the slowly curved shape of the magnetic field. We note that the adopted strength of the magnetic field is that estimated for the projection of the magnetic field in the plane of the sky, suggesting a slightly smaller $M_{\\rm cloud}/\\Psi$ than the estimated one. These imply that  the ambient region is marginally supercritical, while the inner region is supercritical. This situation is considered to be quite consistent with the hourglass shape of the magnetic field and with the cluster formation within the sub-clumps.  It is interesting to examine whether the magnetic field can maintain the outflow collimation along the magnetic field in the ambient region of the sub-clumps, i.e., whether the magnetic field can guide the outflows. The magnetic pressure, $P_{B}=B^2/8\\pi$, is calculated to be $\\sim$ 4 $\\times $10$^{-10}$  dyn, adopting $B\\sim100 \\mu$G. Assuming the average density and velocity width due to turbulence for the outflow to be $3 \\times 10^3$ cm$^{-3}$, which would be consistent with the optically thin condition of the high velocity gas \\citep{wh95}, and 3 km s$^{-1}$, which is larger than the C$^{18}$O velocity width by a factor of $\\sim$1.5--2.0, we obtain the turbulent pressure, $P_{\\rm turb}=\\rho \\sigma_{\\rm turb}^2$, of $\\sim 2 \\times $10$^{-10}$  dyn. Taking into account the fact that the adopted strength of the magnetic field is that estimated for the projection and that $P_{B}$ is proportional to $B^2$, these estimates imply that the magnetic field can guide the outflows in the ambient region of the Serpen cloud core.   \\subsection{Comparison with outflow-driven turbulence model for cluster formation}  From our analysis, the magnetic field seems to be important in considering the cloud stability that is related to star formation or cluster formation and the feed back from the star formation activity, such as outflows.  The hourglass-shaped magnetic field suggests that the Serpens cloud core first contracted along the straight magnetic field to be a filament or elongated cloud, which is perpendicular to the magnetic field, and that then the central part contracted cross the magnetic field due to the high density in the central region of the cloud core. This situation is very similar to the contraction of the low-mass core that is penetrated by the uniform magnetic field \\cite[e,g.,][]{gi06, kan09}. In addition, there might exist the cloud rotation, of which axis agrees with that of the hourglass-shaped magnetic field. It was reported that many small-scale outflows spread to or penetrate the NW and SW sub-clupms\\citep[e.g.,][]{her97,ho99,da99,zi99}, and the ambient, larger-scale outflows (filaments) seem to run along the magnetic field as shown above \\citep{da99,nar02}. Moreover, it is possible that the blue-shifted $^{13}$CO ($1-0$) features just outside CLS 1, which correspond to the red-shifted CO ($2-1$) outflows, are inflows from the ambient to the central part of the SE sub-clump. Considering these altogether, we may have to take into account the magnetic field, outflows, inflows, cloud rotation, and contraction as well as the turbulence of the molecular gas in the cluster formation process of the Serpens cloud core (see Figure \\ref{f10}).  The structures mentioned above seem to be in good agreement with the outflow-driven turbulence modelof \\cite{li06} and \\cite{nakam07} who performed 3D MHD simulation of cluster formation takinginto accout the effect of protostellar outflows. They demonstrated that protostellar outflows can generate supersonic turbulence in pc-scale cluster forming clumps like the Serpens cloud core. One of the important characteristics of outflow-driven turbulence is that gravitational infall motions almost balance the outward motions driven by outflows, creating very complicated density and velocity structure \\citep[see e.g., Figure 4 of][]{nakam07}. The resulting quasi-equilibrium state can be maintained through active star formation in the central dense region. In the presence of relatively strong magnetic field, both outflow and inflow motions in the less dense envelope tend to be guided by large scale ordered magnetic field lines. As a result, filamentary strucutures that are roughly converging toward the central dense region appear in the envelope, whereas the density structure tends to be more complicated in the central dense region where self-gravity and turbulence may dominate over the magnetic field. Infall motions detected by $^{13}$CO ($1-0$) in the Serpens core may correspond to such filamentary structures created by gravitational infall.  To clarify how the outflows and magnetic field affect the dynamical state of the cloud, we assess the force balance in the cloud, following  \\cite{mau09}. To prevent the global gravitational contraction, the following pressure gradient is needed to achieve the hydrostatic equilibrium: \\begin{equation} {dP_{\\rm grav} \\over dr} \\simeq -G{M(r) \\rho (r) \\over r^2} \\left(1-\\alpha^{-2}\\right) \\ , \\end{equation} where $M(r)$ is the mass contained within the radius $r$ and we assume that the cloud is spherical. The effect of magnetic field is taken into account by the factor $(1-\\alpha^{-2})$ and $\\alpha$ is the mass-to-magnetic flux ratio normalized to the critical value  and is approximated as \\begin{equation} \\alpha \\simeq {2\\pi G^{1/2} M/\\pi r^2 \\over B} \\end{equation} \\citep[e.g.,][]{nak98}.   Assuming the density profile of $\\rho \\propto r^{-2}$, the pressure needed to support the cloud against the gravity is estimated to be \\begin{equation} P_{\\rm grav} \\simeq {GM(R)^2 \\over 8\\pi R^4} \\left(1-\\alpha^{-2}\\right) \\ . \\end{equation} The force needed to balance the gravitational force is thus evaluated to be \\begin{equation} F_{\\rm grav} \\simeq 4 \\pi R^2 P_{\\rm grav} (R) = {GM(R)^2 \\over 2 R^2} \\left(1-\\alpha^{-2}\\right) \\ . \\end{equation} Adopting $M(R)= 210 M_\\odot$, $R=0.46$ pc \\citep{ol02}, $B=100 \\mu$G and $\\alpha = 3.8$, $F_{\\rm grav}$ can be estimated to be $ \\sim 4.3 \\times 10^{-4} M_\\odot$ km s$^{-1}$ yr$^{-1}$. The moderately strong magnetic field of $\\alpha = 3.8$ can reduce the gravitational force by $\\sim$7\\%. We note that we rescaled the cloud mass and radius derived from \\cite{ol02} by assumingthe distance to the cloud of 260 pc. Hereafter, we also use other values rescaled for this distance.  On the basis of the CO ($J=2-1$) observations, \\cite{da99} detected many powerful CO outflows in this cloud, and derived the physical properties of the outflows. From their analysis, we can evaluate the total force exerted by the outflows in this region as \\begin{equation} F_{\\rm outflow} \\simeq {p_{\\rm outflow} \\over t_{\\rm dyn}} \\sim {8.7\\textendash17.5M_\\odot {\\rm km \\ s}^{-1} \\over 2.5 \\times 10^4 {\\rm yr}} \\sim (3.4\\textendash7.0) \\times 10^{-4} M_\\odot \\ {\\rm km \\ s^{-1} \\ yr^{-1}}  \\ \\end{equation} where $p_{\\rm outflow}$ is the total outflow momentum, and $t_{\\rm dyn}$ is the representative dynamical time of the outflows. The force due to the outflows, $F_{\\rm outflow}$, is comparable to or somewhat larger than the force needed to stop the global gravitational collapse, $F_{\\rm grav}$, suggesting that the outflows play a crucial role in the cloud dynamics. This result, however, apparently contradicts that of \\cite{ol02} who suggested that the cloud may be undergoing a global contraction, although the further justification is needed to confirm their interpretation. This apparent inconsistency may come from our assumption of the spherical cloud. Since the relatively strong magnetic field associated with the cloud can guide the large scale outflow motions along the global magnetic field as discussed in the previous subsection, the force exerted by the outflows is expected to be weak along the cross-field direction. As a result, the cloud may be able to contract along the cross-field direction. For the Serpens core, both the magnetic field and the outflows are likely to control the cloud dynamics.   The outflows are also expected to be the major source for generating supersonic turbulence in the Serpens core. From the physical quantities of the outflows measured by \\cite{da99}, we can evaluate the total energy injection rate due to the outflows in this region as \\begin{equation} {dE_{\\rm outflow} \\over dt} \\simeq {E_{\\rm outflow} \\over t_{\\rm dyn}} \\sim {(12.7\\textendash48.3) \\ {\\rm J} \\over 2.5\\times 10^4 \\ {\\rm yr}} \\sim (0.5\\textendash2) L_\\odot   \\ . \\end{equation} where $E_{\\rm outflow}$ is the total outflow energy. The energy dissipation rate of supersonic turbulence is obtained by \\citet{mac99} as \\begin{equation} {dE_{\\rm turb} \\over dt} = f {1/2 M \\Delta V^2 \\over \\lambda_d / \\Delta V} \\end{equation} where $f (=0.34)$ is the non-dimensional constant determined from the numerical simulations, and $M$ is the cloud mass, and $\\Delta V$ is the 1D FWHM velocity width. The driving scale of the turbulence $\\lambda_d$ is estimated to be $\\lambda_d \\sim 0.4$ pc for the outflow-driven turbulence \\citep{ma07, nakam07}. The energy dissipation rate of the turbulence can be estimated to be $\\sim 0.12 L_\\sun$, where the FWHM velocity width of about 2 km s$^{-1}$ is adopted \\citep{ol02}. This energy dissipation rate is somewhat smaller than the outflow energy input rate. In the Serpens cloud core, the relatively strong magnetic field tends to guide the outflows and therefore the significant amount of the outflow energy might escape away from the cloud along the magnetic field, as inferred from the magnetic field and outflow structures discussed above. In any case, the outflows seem to have sufficient energy to power supersonic turbulence in this region and the magnetic field seems to play an important role in the escape of the outflow energy from the cloud. These characteristics appear to be in agreement with the the outflow-driven turbulence model for cluster formation, and imply the importance of the magnetic field for the continuous star formation  in the central region of the Serpens cloud core under the condition where the outflow energy injection rate is high. The Serpens cloud core is expected to be one of the good examples of the outflow-driven turbulence model for cluster formation.   \\subsection{Summary}  We have conducted deep and wide ($\\sim$7\\farcm7 $\\times$ 7\\farcm7) $JHK_{\\rm s}$ imaging polarimetry of the Serpens cloud core.    The main findings are as follows:  1. The central part of the infrared reflection nebula is illuminated mainly by two sources; the north by SVS 2 (SRN)  and the south by SVS 20 with two centrosymmetric patterns. The characteristics of the nebula are consistent with those reported in the previous infrared polarimetric works. Detailed inspection enabled us to find  24 YSOs associated with IR nebulae, in addition to SVS 2 and SVS 20.  2. Polarization of NIR point sources was measured and those sources, except YSOs, have an upper limit of polarization degree similar to that of the nearby star forming regions. It is consistent with the dichroic origin, i.e., the polarization vectors of the near-IR point sources could indicate the direction of the averaged local magnetic field.  3. The polarization vectors suggest a clear hourglass shape. We have made a model fitting of this shape with a parabolic function and found that the symmetry axis ($\\theta_{\\rm PA} \\sim$70\\degr)  of the hourglass magnetic field is nearly perpendicular to the elongation ($\\sim150 \\degr$) of the bright parts of C$^{18}$O ($J=1-0$) or submillimeter continuum emissions, i.e., the alignment direction of NW and SE sub-clumps. The submillimeter continuum filaments and CO outflow lobes, which protrude from these sub-clumps, seems to run along the best-fit magnetic field in the ambient region and some $^{13}$CO velocity features also seem to be along the magnetic field.  4. The evaluation of the magnetic field strength has been done with the CF method toward the ambient area of the Serpens cloud core, taking into account the recent study on the signal integration effect for the dispersion component of the magnetic field. The mass to magnetic flux ratio was estimated with the evaluated magnetic field strength of $\\sim100\\mu$G and the parameters of the previous C$^{18}$O ($J=1-0$) observations, and found to be slightly larger than the critical value of magnetic instability in the  the ambient area. This suggests a possibility that the central region is magnetically unstable, which is consistent with the fact that star formation is actively taking place in the central region. We estimated the magnetic pressure and the turbulent pressure of the outflow using the evaluated magnetic field strength and possible turbulent parameters, and found that  the magnetic pressure could be high enough to  guide the outflows in the ambient region.  5. The bright part of C$^{18}$O ($J=1-0$),  submillimeter continuum cores as well as many class 0/I objects are located just toward the constriction region of the hourglass-shaped magnetic field. These suggest that the Serpens cloud core first contracted along the magnetic field to be an elongated cloud and that then the central part contracted cross the magnetic field due to the high density in the central region of the cloud core.  6.  Comparisons of the best-fit magnetic field with the previous observations of molecular gas and large-scale outflows suggest  a possibility that the cloud dynamics is controlled by the magnetic field, protostellar outflows and gravitational inflows. In addition, the outflow energy injection rate appears to be  the same as or larger than the dissipation rate of the turbulent energy in this cloud, indicating that the outflows are the main source of turbulence and that the magnetic field plays an important role both in allowing the outflow energy to escape from the central region of the cloud core and enabling the gravitational inflows from the ambient region to the central region. These characteristics appear to be in good agreement with the outflow-driven turbulence model for cluster formation and imply the importance of the magnetic field to continuous star formation in the center region."
    },
    "1004/1004.0882.txt": {
        "abstract": "\\noindent The quantum theory of cosmological perturbations in single field inflation is formulated in terms of a path integral. Starting from a canonical formulation, we show how the free propagators can be obtained from the well known gauge-invariant quadratic action for scalar and tensor perturbations, and determine the interactions to arbitrary order. This approach does not require the explicit solution of the energy and momentum constraints, a novel feature which simplifies the determination of the interaction vertices. The constraints and the necessary imposition of gauge conditions is reflected in the appearance of various commuting and anti-commuting auxiliary fields in the action. These auxiliary fields are not propagating physical degrees of freedom but need to be included in internal lines and loops in a diagrammatic expansion. To illustrate the formalism we discuss the tree-level 3-point and 4-point functions of the inflaton perturbations, reproducing the results already obtained by the methods used in the current literature. Loop calculations are left for future work. ",
        "introduction": "Perhaps the most remarkable aspect of inflation~\\cite{Guth:1980zm} is its ability to imprint fluctuations on cosmic scales through a confluence of quantum mechanics and general relativity. This connection, first realized more than 25 years ago \\cite{Mukhanov:1981xt, Hawking:1982cz, Starobinsky:1982ee, Guth:1982ec, Bardeen:1983qw, Sasaki:1986hm}~\\footnote{The first rigorous and quantitatively accurate treatments of inflationary perturbations were \\cite{Mukhanov:1981xt} and \\cite{Sasaki:1986hm}.}, has given inflationary theory the impetus which positioned it as the leading paradigm for approaching the physics of the early universe. Since then, the primordial fluctuations have been measured in the CMB~\\cite{Bennett:1996ce} with ever increasing accuracy and resolution~\\cite{Larson:2010gs, Komatsu:2010fb} and will be scrutinized even further in the near future~\\cite{:2006uk}. It is not surprising then that over these past decades a lot of effort has been devoted to fleshing out the predictions inflation makes for these fluctuations in a variety of theoretical settings. Since these fluctuations are initially small, of order $10^{-5}$ at $z \\simeq 1090$, the linearized theory of perturbations has been developed to a significant degree and has been used, rather successfully, to compare theory to observation. Over the past few years, efforts have intensified to explore inflationary perturbations beyond linear order, mostly in the context of the related non-Gaussianity which has become a significant subfield of cosmological research~\\cite{Turner:2008zza}. The amount of work that has been done on the subject is by now rather voluminous with many authors examining various aspects. A definitive calculation was performed by Maldacena~\\cite{Maldacena:2002vr} showing that single field inflation leads to small primordial non-Gaussianities (see also \\cite{Acquaviva:2002ud}) with more complicated single- and multi-field models providing more possibilities for larger non-Gaussianity, see {\\it e.g.} \\cite{Bernardeau:2002jf, Zaldarriaga:2003my, Rigopoulos:2005ae, Rigopoulos:2005us, Vernizzi:2006ve, Alishahiha:2004eh, Enqvist:2004ey, Lyth:2005fi, Malik:2006pm, Sasaki:2006kq, Langlois:2008qf, Barnaby:2008fk, Chen:2008wn, Byrnes:2008wi, Sasaki:2008uc,Hotchkiss:2009pj, Chambers:2009ki}. On the observational side, a major effort is under way to develop observational measures of non-gaussianity \\cite{Komatsu:2010fb, Creminelli:2005hu, Liguori:2010hx, Jeong:2009wi, Bucher:2009nm, Yadav:2007yy, Komatsu:2008hk, Smith:2009jr, Fergusson:2008ra, Seljak:2008xr, Slosar:2008hx}.  The foray into non-linear corrections and the associated non-Gaussianity has been motivated by a number of reasons. If one neglects the running of the spectral indices, at linear order inflationary theories predict four numbers: the amplitude and spectral index of scalar and tensor perturbations, and a variety of different models can coincide on these predictions. However, non-Gaussianity can be rather discriminatory for different models due to its much richer, and more complicated, structure. As an example, the detection of a significant three-point function would immediately rule out single-field inflation, as well as some simple multi-field generalizations, may favor alternative models, or provide evidence for the processes involved in heating up the universe after Inflation. Entwined with considerations of testing inflationary theory against observations and non-Gaussianity are considerations of theoretical understanding: calculating and controlling higher order quantum loop corrections to inflationary predictions, backreaction issues~\\cite{Vilenkin:1982wt} and various divergences which appear at higher orders of perturbation theory -- see for example~\\cite{Tsamis:1996qm, Weinberg:2005vy, Tsamis:1993ub, Weinberg:2006ac, Sloth:2006az, Sloth:2006nu, Seery:2007we, Seery:2007wf, Bilandzic:2007nb, Janssen:2008dw, Janssen:2009nz, Campo:2009fx, Urakawa:2009my, Senatore:2009cf}.   So far, all the work on inflationary non-linear corrections has been done using the operator language in the interaction picture. However, in many branches of modern physics it is the path integral formulation for quantum mechanical systems which has proved quite useful. For example, it has been of paramount importance in understanding gauge theories and their experimental consequences for particle physics, the theoretical description of condensed matter systems and is the preferred language in which modern theories of fundamental physics are formulated and quantized. In this spirit, and hoping to throw more light on the understanding of inflation, here we develop a path integral formulation for inflationary perturbations to arbitrary order in interaction terms. An early path-integral formulation of linear perturbation theory can be found in \\cite{Anderegg:1994xq}. A novel feature of our approach is that there is no need to explicitly solve the energy and momentum constraints to a particular perturbative order as is usually done when working in a particular gauge. This allows us to obtain the interaction terms, and hence the vertices, to arbitrary order in a closed form.  We start from the more fundamental canonical path integral and show how to obtain the configuration phase space path integral, making in the process a connection with the well known gauge invariant linear perturbation theory. With a single inflaton there is only one scalar -- expressed in terms of the the Sasaki-Mukhanov variable -- and two tensor degrees of freedom which propagate, as expected, while the vectors completely drop out at quadratic order. However, various (real and commuting) auxiliary fields which do not appear in the in-state, the external lines in the ``in-in diagrammatic'' expansion, must be included in computations of N-point functions with $N\\geq 4$ or calculations involving loops. These fields arise because the constraints are not explicitly solved. Furthermore, anti-commuting ghosts arise from the path integral measure which must also be taken into account at a certain loop order. Here we focus on standard potential-energy-dominated single-field inflation, but the generalization to more fields and more complex theories is in principle straightforward as long as a canonical formulation is available.  The paper is organized as follows: In section \\ref{Quantization of Inflationary perturbations} we derive the action for perturbations in a form that includes interactions to all orders in a closed expression, discuss the role of constraints and gauge conditions and formulate the transition amplitude between two quantum states separated by (background) time t in terms of a path integral. In the process the action is written in a form that makes contact with known results from gauge-invariant linear perturbation theory \\cite{Mukhanov:1990me} and simplifies the computation of the propagators. In section \\ref{The functional in-in formalism and expectation values} we discuss the ``in-in'' generating functional  which provides the appropriate diagrammatic rules for the computation of the  N-point functions relevant to cosmology. In section \\ref{The 3-point and 4-point functions of Inflaton perturbations} we apply the above formalism to give expressions for the tree level 3-point and 4-point functions in terms of propagators, reproducing the results already obtained using an operator formalism and a different methodology. We close in section \\ref{Discussion} with a discussion of our results. ",
        "conclusions": "\\label{Discussion}  In this paper we developed a path integral formulation for inflationary perturbations in single field inflation. Using a phase space formulation of the system as a starting point, we were able to bring the free action into a particularly simple form which allows for a straightforward calculation of the various propagators, recovering in the process standard results of linear gauge-invariant cosmological perturbation theory. The interaction terms can then be obtained without the necessity of first solving the energy and momentum constraints as has been done so far in the literature. A notable feature is the appearance of two types of auxiliary fields in the path integral beyond the physical propagating degrees of freedom: a set of commuting non-dynamical fields related to the existence of the constraints, as well as a set of anticommuting Faddeev-Popov ghost fields induced by the imposition of gauge conditions. The resulting path integral is independent of the choice of gauge, provided the asymptotic states are defined in terms of the linearized fields and the corresponding linearized gauge transformations.  We then briefly described how to obtain N-point expectation values and commented on the meaning of two different gauges and their relation to cosmological observables in this formalism. They are the uniform field gauge and the tensor gauge which have been widely used in past considerations of non-Gaussianity but in the context of an interaction picture operator approach. Correlators can be expressed in a systematic expansion in diagrams in the in--in formalism. We found that, when quantum loops are taken into account, anticommuting ghost fields must be included in the computation. Furthermore, internal lines in diagrams should also include the commuting auxiliary fields and we demonstrated their role in the computation of 4-point functions, in which the leading order contributions indeed come from diagrams with internal lines involving auxiliary fields. This way of obtaining the effective 4-point interactions seems simpler than having to go through the solution of the constraints at second order and the substitution of these solutions back in the action~\\cite{Seery:2006vu,Jarnhus:2007ia}.  So far, we have only considered standard single field inflation but the generalization to multi-field and more complicated models should be straightforward as long as the canonical formulation is known. We have also explicitly considered only tree diagrams in the tensor gauge, noting that the standard non-linear redefinitions (gauge transformations) can be used to obtain results for the curvature perturbation on comoving slices. Of course we could have obtained this result by working directly in the uniform field gauge, at the price of more complicated interaction terms. Relating results in the two gauges would be particularly interesting if quantum loop corrections are taken into account and questions of renormalization arise. This is a regime where the role of the anticommuting ghosts would become crucial. Currently such questions are only addressed using the (essentially classical) $\\Delta N$ formalism on long wavelengths. Other issues involve the appearance of infrared divergences, backreaction and a more rigorous definition of stochastic inflation. A path integral formulation might prove very useful for such considerations which will be the focus of future work."
    },
    "1004/1004.4783_arXiv.txt": {
        "abstract": "The collision of  two white dwarfs is a quite  frequent event in dense stellar systems,  like globular clusters and galactic  nuclei. In this paper we  present the  results of  a set of  simulations of  the close encounters and  collisions of two  white dwarfs. We use  an up-to-date smoothed particle  hydrodynamics code that  incorporates very detailed input physics  and an improved treatment of  the artificial viscosity. Our  simulations have  been done  using  a large  number of  particles ($\\sim  4\\times10^5$) and  covering  a wide  range  of velocities  and initial distances of the colliding white dwarfs.  We discuss in detail when  the initial eccentric  binary white  dwarf survives  the closest approach,  when a  lateral collision  in which  several  mass transfer episodes occur is  the outcome of the newly  formed binary system, and which range of input parameters  leads to a direct collision, in which only  one   mass  transfer  episode  occurs.   We   also  discuss  the characteristics of the final  configuration and we assess the possible observational  signatures  of  the  merger,  such  as  the  associated gravitational waveforms  and the fallback luminosities.   We find that the overall  evolution of the  system and the main  characteristics of the final object agree with  those found in previous studies.  We also find  that the  fallback luminosities  are close  to  $10^{48}$ erg/s. Finally,  we find  as well  that  in the  case of  lateral and  direct collisions   the   gravitational   waveforms  are   characterized   by large-amplitude peaks  which are followed by a  ring-down phase, while in  the case  in which  the binary  white dwarf  survives  the closest approach,  the  gravitational pattern  shows  a distinctive  behavior, typical of eccentric systems. ",
        "introduction": "In recent  years, the study  of stellar collisions has  attracted much interest from  the astronomical community  working on the  dynamics of dense  stellar  systems,  like  the  cores of  globular  clusters  and galactic nuclei (Shara  2002). One of the reasons for  this is that in these systems,  stellar collisions are  rather frequent (Hills  \\& Day 1976).  In fact, it has been predicted that up to 10\\% of the stars in the core  of typical globular  clusters have undergone a  collision at some point during the lifetime of the cluster (Davies 2002).  The most  probable collisions are those  in which at least  one of the colliding stars has the largest possible cross section --- a red giant --- and those in which at least one of the stars is most common (Shara \\&  Regev 1986).   This later  type of  collisions  obviously includes those in  which a  main sequence star  is involved.   However, because white dwarfs  are the most common  end point of  stellar evolution and because  both globular clusters  and galactic  nuclei are  rather old, these   stellar  systems   contain  many   collapsed   and  degenerate objects.  Therefore, we  expect that  collisions in  which one  of the colliding stars is a white dwarf should be rather common.  Collisions  between  two  main  sequence  stars  are  supposed  to  be responsible for the observed population of blue stragglers in globular clusters (Sills  \\& Bailyn  1999; Sills et  al.  2005).   Although the collision  between  a main  sequence  star and  a  red  giant is  more probable than  that of two main  sequence stars because  of the larger geometrical cross section of the  red giant star, they probably do not produce an  interesting astrophysical object.  The reason  for this is the low density of the envelopes  of red giants. In most cases, during the encounter the red giant is deprived of part of its envelope and is able to recover its appearance (Freitag \\& Benz 2005).  The collisions of a white dwarf  and a red giant or a main  sequence star are also of large interest,  since they  may be responsible  for the  formation of some  interesting astrophysical  objects.  Unfortunately,  due  to the very   different  dynamical   scales   involved,  the   hydrodynamical simulation of  these events is difficult and  realistic simulations of their outcome are  still lacking. However, in the case  in which a red giant and a  white dwarf collide it is thought  that the most probable outcome  is the  ejection of  the envelope  of the  red giant  and the formation of a double white dwarf binary system (Tuchman 1985), whilst in the case in which a main sequence star and a white dwarf collide it has been shown (Shara \\& Regev 1986) that only a small fraction of the disrupted main  sequence star remains  bound to the white  dwarf. More recent  simulations (Ruffert  1992) predict  the formation  of  a disk around  the  white  dwarf.   However,  we  emphasize  that  all  these simulations used rudimentary input  physics and, thus, the outcomes of these simulations are dubious.  The collision of two white  dwarfs deserves study for various reasons. In particular, the collision of two white dwarfs can produce a Type Ia supernova. Although the most standard  scenario for a Type Ia outburst --- the  so-called  single-degenerate scenario  ---  involves a  white dwarf accreting from a non-degenerate companion, the double-degenerate scenario (Webbink 1984; Iben \\& Tutukov 1984), in which the merging of two  carbon-oxygen white  dwarfs with  a  total mass  larger than  the Chandrasekhar limit occurs, has been one of the most favored scenarios leading to  Type Ia supernovae.  In  fact, it has  been predicted that the white  dwarf merger  rate leading to  super-Chandrasekhar remnants will  be  increased  by   an  order  of  magnitude  through  dynamical interactions  (Shara \\&  Hurley  2002). Therefore,  collisions of  two white  dwarfs of  sufficiently large  masses could  explain supernovae occurring in the  nuclei of galaxies.  Moreover, it  has been recently suggested that such  a process would lead to both  a Type Ia supernova explosion  and  to the  formation  of  a  magnetar (King,  Pringle  \\& Wickramasinghe   2001).   This   scenario  would   explain   the  main characteristics  of  soft  gamma-ray  repeaters  and  anomalous  X-ray pulsars  like 1E2259+586.   Also, dynamical  interactions  in globular clusters can  form double  white dwarfs with  non-zero eccentricities, which would be powerful sources of gravitational radiation (Willems et al. 2007). Moreover, the initial  stages of the coalescence of a white dwarf binary system  could be one of the  most interesting sources for the detection of gravitational  waves using space-borne detectors like LISA ({\\tt http://lisa.jpl.nasa.gov}).  Thus, a close encounter of two white dwarfs would be a potentially observable source of gravitational waves and,  hence, characterizing the gravitational  waveforms is also of interest. Finally, the  temperatures achieved in a direct collision are substantially high  and, consequently, we expect that  some of the nuclearly processed material could  be ejected, leading to a pollution of the  environment where  it occurs, either  a globular cluster  or a galactic nucleus.  Despite of its  potential interest, there are very  few simulations of the collision of  two white dwarfs, the only  exception being those of Benz et  al. (1989), Rosswog et  al. (2009) and Raskin  et al. (2009). All  three sets  of simulations  used Smoothed  Particle Hydrodynamics (SPH) to  model the collisions.   However, the simulations of  Benz et al.  (1989) were done using  a small number of particles whereas those of Rosswog  et al.  (2009)  and Raskin et  al.  (2009) have  studied a limited range of impact parameters. To be specific, the simulations of Benz  et al.   (1989) used  $5\\times 10^3$  particles, while  those of Rosswog et  al. (2009)  and Raskin et  al. (2009)  used, respectively, $2.5\\times  10^6$ and $8\\times  10^5$ SPH  particles. The  very recent simulations of  Rosswog et al.  (2009)  and Raskin et  al. (2009) were aimed  to  produce a  thermonuclear  explosion  and,  thus, they  only studied  direct  collisions, while  the  simulations  of  Benz et  al. (1989) covered a broader range of initial conditions. Additionally, in the early work of Benz et al.  (1989) the classical expression for the artificial viscosity (Monaghan \\& Gingold 1983) was used, while in the very  recent calculations  of Rosswog  et  al.  (2009)  and Raskin  et al. (2009) more elaborated  prescriptions for the artificial viscosity were employed.  In  sharp  contrast,  the  coalescence  of  binary  white  dwarfs  was extensively  studied in  the  past and  also  has been  the object  of several  recent  studies.   For  instance,  the  pioneering  works  of Mochkovitch \\& Livio  (1989, 1990) used an approximate  method --- the so-called  Self-Consistent-Field method (Clement  1974) ---  while the full  SPH  simulations of  Benz,  Thielemann  \\&  Hills (1989),  Benz, Cameron \\&  Bowers (1989), Benz,  Hills \\& Thielemann (1989),  Benz et al.   (1990),  Rasio \\&  Shapiro  (1995)  and  Segretain, Chabrier  \\& Mochkovitch (1997)  studied the problem using  reduced resolutions and the  classical expression  for the  artificial viscosity  (Monaghan \\& Gingold 1983).  Later, Guerrero et  al.  (2004) opened the way to more realistic simulations, using an  increased number of SPH particles and an improved prescription for the artificial viscosity.  More recently, the  simulations of  Yoon et  al.  (2007) and  of Lor\\'en--Aguilar  et al. (2005, 2009)  were carried out using modern  prescriptions for the artificial viscosity and even larger numbers of particles. All in all, it is noticeable the lack of SPH simulations of white dwarf collisions and  close encounters  when compared  to the  available  literature on white dwarf mergers.  In  the present  paper  we study  the  collision of  two white  dwarfs employing   an  enhanced  spatial   resolution  ($4\\times   10^5$  SPH particles) and  an improved formulation for  the artificial viscosity. We pay  special attention to  discern the range of  initial conditions that produce the  tidal disruption of the less  massive white dwarf or those  for which  the initial  eccentric binary  survives  the closest approach.  The  number of  particles used in  our simulations  is much larger than those  used in the simulations of Benz  et al.  (1989) and in line with  those used in modern simulations  (Rosswog et al.  2009; Raskin  et al.  2009).   However, our  calculations encompass  a broad range of initial conditions of the colliding white dwarfs, in contrast to most modern simulations, in which  only a few cases were studied in detail.  The paper  is organized as follows.  In \\S  2 we describe our input physics and the  method of calculation, paying special attention to describe with some detail our  SPH code.  It follows \\S 3, which is devoted  to discuss  the  initial conditions  adopted  in the  present study,  while in  \\S 4  we describe  the results  of  our simulations. Finally  in  \\S  4  we  summarize  our  main  findings  and  draw  our conclusions. ",
        "conclusions": "In this paper  we have studied the collisions  and close encounters of two  white  dwarfs,  using  a  state  of  the  art  Smoothed  Particle Hydrodynamics code.   Collisions between two  white dwarfs are  not as frequent as  binary mergers.  However, as discussed  in Timmes (2009), Rosswog  et al.  (2009)  and Raskin  et al.  (2009), they  most likely occur in globular clusters and  the central regions of galaxies, where the stellar  densities are very high.  The  collision time, $\\tau_{\\rm coll}$,  adopting a Maxwellian  velocity distribution  with dispersion $\\sigma$, and  assuming a closest approach distance  $r_{\\min} < 2R_*$ is (Binney \\& Tremaine 1987):  \\begin{equation} \\frac{1}{\\tau_{\\rm coll}} = 16\\sqrt{\\pi}n_{\\rm WD}\\sigma R_*^2 \\left(1+\\frac{v_{\\rm esc}^2}{4\\sigma^2}\\right) \\end{equation}  \\noindent  where  $n_{\\rm WD}\\simeq  10^4$  pc$^{-3}$  is the  typical number density  of white  dwarfs in a  globular cluster,  $v_{\\rm esc} \\simeq  4000$~km/s is  the  white dwarf  escape  velocity and  $\\sigma \\simeq  5$~km/s is  the  relative velocity  dispersion  of both  white dwarfs  in a  the globular  cluster,  which is  entirely dominated  by gravitational  focusing. Consequently,  the rate  of collisions  for a typical globular cluster is given by:  \\begin{equation} r_{\\rm GC} \\sim \\frac{1}{2}\\frac{n_{\\rm WD}}{\\tau_{\\rm coll}} \\; \\frac{4}{3} \\pi r_{\\rm c}^3 \\end{equation}  \\noindent  where $r_{\\rm  c}\\sim 1.5$  pc is  the core  radius  of the globular  cluster (Peterson  \\& King  1975).  Adopting  the previously mentioned typical values, we  obtain $r_{\\rm GC}\\sim 8\\times 10^{-10}$ yr$^{-1}$.  Taking into account  that the density of globular clusters is $n_{\\rm GC}  = 4.2$ Mpc$^{-3}$ (Brodie \\&  Strader 2006) as Rosswog et al.  (2009)  did, we obtain an overall  rate of interactions $R\\sim 3\\times   10^{-10}$  Mpc$^{-3}$   yr$^{-1}$.   Thus,   although  these interactions are not  very frequent, they are not  unlikely and, thus, there is a possibility of detecting them.  Motivated by this we have  studied the collisions and close encounters of  two otherwise  typical white  dwarfs  of masses  $0.8$ and  $0.6\\, M_{\\sun}$, respectively,  for a broad range of  initial conditions and employing a large number of SPH particles. Our initial conditions have been chosen  in such a  way that a  close encounter or a  collision is always guaranteed, and are summarized in Table \\ref{outcome}.  We have found  that the outcome  of the  interactions can  be either  a direct collision,  a  lateral   collision,  in  which  several  mass-transfer episodes may occur,  and finally the survival of  the eccentric binary system  of two  white  dwarfs.   We have  characterized  the range  of initial velocities  and distances --- or, alternatively,  the range of energies  and angular  momenta ---  which lead  to each  one  of these outcomes. We find that when  the distance between the two white dwarfs at the closest  approach is smaller than $0.009\\pm  0.002 \\, R_{\\sun}$ the  final outcome  is  a  direct collision,  when  it ranges  between $0.009\\pm  0.002\\,  R_{\\sun}$   and  $0.033\\pm  0.004\\,R_{\\sun}$,  the outcome  is a lateral  collision and  otherwise the  double degenerate binary system survives.  In all the cases in which a collision is the result of the interaction we  obtain that  little  mass  is ejected  during  the entire  merging process and that  a central white dwarf surrounded  by a debris region is  formed.  If  the collision  is  direct this  region has  spherical symmetry,  whilst  if the  collision  is  lateral  we obtain  a  heavy rotationally-supported  keplerian  disk.    In  both  cases  the  peak temperatures  achieved  during   the  interaction  exceed  the  carbon ignition temperature  and some nucleosynthesis  occurs. However, since these high temperatures  are not achieved during long  periods of time the   abundances  of   heavy  nuclei   are  not   large   (see  tables \\ref{tab-chem-disk} and \\ref{tab-chem-corona}).  Naturally, the extent to  which nuclear  burning proceeds  depends  on the  strength of  the interaction, and  hence the  production of heavy  nuclei is  larger in direct  collisions. Most of  the nuclear  reactions occur  when matter from the  less massive white  dwarf is shocked  on the surface  of the most massive one.  Consequently, we find that the maximum temperatures of the  merged system occur  on a hot  corona around the  most massive white dwarf.  We have  also paid special attention  to a specific case  of a lateral collision in which several (up to 7) mass-transfer episodes occur.  We have found  that mass-transfer only occurs during  the periastron, and that at  each passage the  distance between the two  interacting white dwarfs decreases, that  the mass lost by the  less massive white dwarf increases,  and  that  there  is substantial  synchronization  of  the system.  We  have  also computed  possible  observational  signatures of  these events. Specifically, we have calculated the emission of gravitational waves and the fallback luminosities in the aftermath of the merger. We have  shown  that  it is  very  unlikely  that  LISA will  detect  the gravitational waves radiated during  these interactions.  Only at very late  phases,  when  the  orbits  are circularized,  the  emission  of gravitational waves from those systems in which an eccentric binary is initially formed is  possible. However, at these very  late stages all the information about the close encounter is completely lost.  We have also  computed  the  fallback  luminosities which  result  from  those interactions in  which a merger  occurs.  We have found  that although these  luminosities are somewhat  smaller than  those obtained  in the merger of  a binary  white dwarf system  they are still  rather large, allowing the future detection of these events.  Finally, we would like to emphasize  that our main aim was to study to the post-capture scenario for a  fixed pair of masses of the colliding white  dwarfs. Our  study differs  from  those of  Rosswog (2009)  and Raskin  et al. (2009)  in the  adopted masses  of the  colliding white dwarfs and  in the  initial conditions. This  is a consequence  of the different  motivations of  all three  works.  Whereas  the  studies of Rosswog (2009) and  Raskin et al. (2009) were  aimed at producing Type Ia supernova  outbursts, our  work was focused  at studying  the tidal disruption  of typical white  dwarfs in  globular clusters.   For this reason both Rosswog (2009) and and Raskin et al. (2009) studied direct collisions adopting  different masses  for the colliding  white dwarfs than those adopted in this study.  Additionally, Rosswog et al. (2009) placed the colliding white dwarfs  in a parabolic orbit and, moreover, the total mass of the system was larger than Chandrasekhar's mass.  In this sense,  all three studies are complementary  but, obviously, more studies are needed to explore the full range of possibilities."
    },
    "1004/1004.3992_arXiv.txt": {
        "abstract": "The spontaneous  onset of magnetic reconnection in thin collisionless current sheets is shown to result from a thermal-anisotropy driven magnetic Weibel-mode, generating seed-magnetic field {\\sf X}-points in the centre of the current layer. ",
        "introduction": "The idea of magnetic reconnection as the main plasma process that converts stored magnetic energy into kinetic energy originates from the intuitive geometric picture of annihilating antiparallel magnetic field lines when approaching each other  \\citep[see, e.g.,][]{sweet1957,parker1958,dungey1961}.  Observations in space  have unambiguously confirmed the presence of reconnection under collisionless conditions \\citep[see, e.g.,][]{fujimoto1997,oieroset2001,nagai2001} when the fluid theoretical approaches break down. However, no convincing theoretical argument for the spontaneous occurrence of reconnection has so far been given. In collisionless numerical simulations reconnection is artificially ignited \\citep[cf., e.g.,][]{zeiler2000}, mostly by {\\it ad hoc} imposing a seed {\\sf X}-point in the current sheet separating the anti-parallel fields. The ongoing search for the mechanism of spontaneous onset of collisionless reconnection points to the `missing microphysics' in thin current sheets.  In the present Letter we show that instability of the inner current layer gives rise to the self-consistent generation of local magnetic fields ${\\bf B}=(B_x,0,B_z)$ transverse to the current layer.  Such local fields are equivalent to the generation of microscopic seed {\\sf X}-points in the current sheet centre and are capable of spontaneously igniting reconnection as is known from two-dimensional PIC particle simulations. Since in an ideal current sheet ions and electrons become non-magnetic on their respective inertial scales $\\lambda_{\\,i,e}=c/\\omega_{\\,i,e}$, where $\\omega_{\\,i,e}=e\\sqrt{N/\\epsilon_0m_{\\,i,e}}$ are the plasma frequencies of ions and electrons, (classical) collisionless convective transport of magnetic fields into the current layer takes place up to a vertical distance $z\\sim\\lambda_{\\,e}$ from the centre of the current sheet. The region between $\\lambda_{\\,e}\\lesssim z\\lesssim\\lambda_{\\,i}$ is known as the `Hall-current' \\citep{sonnerup1979} or (mistakenly, as there is no diffusion present)  `ion-diffusion' region. Being a by-product of thinning of the current layer, the Hall currents are -- presumably -- not involved in the reconnection process proper.\\footnote{The question of the role of Hall currents in reconnection is not resolved yet. Classically they are unimportant for the reconnection process. It is, however, not certain whether or not on the microscopic scales non-classical (quantum-Hall) effects are induced by the environmental conditions \\citep[einselection effects, see][]{zurek2003} in which Hall-electrons would more directly be involved.} They close along the magnetic field by electrons that are accelerated in the oblique lower-hybrid-drift/modified-two-stream instability driven by magnetised Hall-electrons on the non-magnetic ion background thereby coupling the reconnection site to the auroral ionosphere \\citep{treumann2009}. %  When speaking of a current sheet, we refer to ideal current sheets separating strictly  antiparallel fields. The observational paradigm of a reconnecting current sheet is the magnetospheric tail-current sheet. This current sheet is not ideal in the above sense as it is embedded into a quasi-dipolar field which still might preserve a weak rudimentary (normal) magnetic field component $B_z$pointing northward. This $B_z$ component re-magnetises the central-sheet electrons and affects the evolution of (collisionless) tearing modes \\citep{galeev1975}. Nevertheless below, when using numbers, we will for reasons of resolution refer to conditions in the magnetotail even though our theory might better apply to the magnetopause, interplanetary space or astrophysics.  In principle, observation of the electron-inertial (`electron-diffusion') region is difficult because of its narrow width. Unambiguous observations do not yet exist. At the magnetopause, in particular, very narrow electron layers have sometimes been reported assuming that they relate to the electron-inertial region during reconnection \\citep[for a recent discussion of the experimental prospects of resolving the electron-inertial region cf.,][]{scudder2008}. ",
        "conclusions": ""
    },
    "1004/1004.3489_arXiv.txt": {
        "abstract": "We studied 16 sunspots with different sizes and shapes using the observations with the Hinode Solar Optical Telescope. The ratio of G-band and CaII H images reveal rich structures both within the umbra and penumbra of most spots. The striking features are the compact blob at the foot point of the umbra side of the penumbral fibrils with disk center-limb side asymmetry. In this paper, we present properties of these features using the spectropolarimetry and images in G-band, CaII and blue filters. We discuss the results using the contemporary models of the sunspots. ",
        "introduction": "\\noindent {\\it \"Clearly, observations must point the way, because it is evident now that the sunspot is too complicated a structure to be the product of a single overwhelming theoretical effect. The sunspot results from the conjunction of several effects, and we can spend a lot of time guessing what those effects might be without getting anywhere\".~~~~~~~~~~~~Eugene Parker (Personal communication 2010)}  \\vspace{0.1 cm}  We analyzed the sunspots observations obtained with Solar Optical Telescope (SOT) onboard Hinode to look for the signatures that are relevant to identify theoretical effect of different models. A sunspot may be made up of individual flux tubes pressed together or a monolithic block of magnetized plasma immersed in the solar convective zone \\citep{par79a,mey77}. High spatial resolution observations of sunspot formation show that individual magnetic knots appear on the solar surface that get collected to form larger bodies perhaps by mutual hydrodynamic attraction \\citep{par78}. The question still remains that after the  knots are collected together to form a sunspot, do they still retain their individuality? In other words, are there gaps between the individual flux tubes that contain plasma devoid of magnetic field? Existence of field free plasma below the visible layer of the sunspots distinguish the two types of sunspots models \\citep{del01}. The physical processes leading to the stability and decay of the sunspot would depend largely on its subphotospheric structure. The time series of high spatial resolution Hinode observations of sunspots, without the interruption of atmospheric seeing, provides opportunity to study the finer structures and their temporal evolution governed by the subphotospheric dynamics. The G-band images serve as a flux tube diagnostic tool due to their association with magnetic concentrations, with efficient heat transport leading to the weakening of molecular band and shifting of optical depth \\citep{ish07,san01}. The Ca II H images provides information on the layer slightly higher than the height of G-band formation, the combined understanding of which would provide new understanding of sunspot magnetic structure \\citep{car07}. ",
        "conclusions": "Figure 1 (a) shows the example of the ratio of G-band and CaII H images of a sunspot observed on 02 May 2007. The important features of the ratio images are: (1) The inner penumbra of the ratio images are brighter compared to the outer part where the CaII H signal is higher compared to the G-band. There is a sharp boundary at the middle of penumbra dividing bright inner and dark outer area. (2) The foot points of the penumbra, which are the locations of the peripheral umbral dots, are brighter compared to the inner umbral dots. The disk-centerward penumbral boundary has brighter points than their limb side counterparts. (3) The umbral dots are not uniformly distributed in the umbra. (4) The umbral dots, closer to the penumbra, are associated with elongated structures directed inwards. (5) Most umbral dots are not circular but elongated elliptical shape. The light bridges consists of compact bright points that have elongated structures. Figure 1 (b) shows the G-band flux and magnetic filling factor as a function of continuum intensity in BFI red filter through at 6685 \\AA\\/ and 6302.5 \\AA\\/. The solid curve represent the quadratic function of continuum intensity that follows the trend of both G-band flux and filling factor.  We interpret these observations using the theoretical models for sunspot structure below and above the visible layer with forest of field free gaps between the flux tubes \\citep{par79b,spr06}. Figure 2 describes the physical processes leading to the observed sunspot bright point. The figure shows seven flux tubes in a bundle embedded in field free plasma expanding above photosphere as the gas pressure drops. The field free plasma penetrates in the gap between the flux tubes as shown by gray arrows that transport heat from non magnetic convective zone. The penetrating field free plasma between the flux tubes oscillates, represented by hatched columns with double arrow, due to longitudinal overstability \\citep{par79b} and heats the upper layers represented by shaded columns, dissociate the molecules in them and produce umbral dots by altering the optical depth to expose the deeper layers. The inclinaed flux tubes result viewing of larger part of hot flux tube wall \\citep{rut01}. The oscillating plasma columns themselves do not appear at the visible surface as blocked by the overlying magnetic canopy. The penumbral flux tubes that are at the periphery of the sunspot get extra heating by direct contact with the convective zone hot plasma represented by hatched side arrows leading to the brighter penumbral foot points. The penumbral flux tubes are highly inclined which contribute to the G-band brightness with limb-disk side asymmetry. As the matter in the flux tube gets heated it expands leading to downward drop in pressure resulting directional flow indicated by arrows in flux tube 1 and 7 \\citep{deg93}. The umbra-ward motion of penumbral grains may also be due to this effect. In some locations of the umbra the flux tubes are so dense that there is not much gap between them for the oscillating field free plasma represented by shaded flux bundle between the 4 and 5 flux tube. At these locations there are few umbral dots leading to non uniform distribution in the umbra. In the penumbra, the field free plasma would heat localized regions and produce the bright points represented by shaded and hatched region near flux tube 1 and 7 and produce hotter flux tube plasma in the outer penumbra where the plasma is tenuous due to the expanding flux tubes."
    },
    "1004/1004.5099_arXiv.txt": {
        "abstract": "GRB 090926A was detected by both the GBM and LAT instruments on-board the \\textit{Fermi} Gamma-Ray Space Telescope.  \\textit{Swift} follow-up observations began $\\sim$13 hours after the initial trigger.  The optical afterglow was detected for nearly 23 days post trigger, placing it in the long lived category.  The afterglow is of particular interest due to its brightness at late times, as well as the presence of optical flares at T0+10$^5$ s and later, which may indicate late-time central engine activity.  The LAT has detected a total of 16 GRBs; 9 of these bursts, including GRB 090926A, also have been observed by \\textit{Swift}.  Of the 9 \\textit{Swift} observed LAT bursts, 6 were detected by UVOT, with 5 of the bursts having bright, long-lived optical afterglows.  In comparison, \\textit{Swift} has been operating for 5 years and has detected nearly 500 bursts, but has only seen $\\sim$30\\% of bursts with optical afterglows that live longer than 10$^5$ s.  We have calculated the predicted gamma-ray fluence, as would have been seen by the BAT on-board \\textit{Swift}, of the LAT bursts to determine whether this high percentage of long-lived optical afterglows is unique, when compared to BAT-triggered bursts.  We find that, with the exception of the short burst GRB 090510A, the predicted BAT fluences indicate the LAT bursts are more energetic than 88\\% of all Swift bursts, and also have brighter than average X-ray and optical afterglows. ",
        "introduction": "The \\textit{Fermi} Gamma-ray Space Telescope has opened a new era of gamma-ray burst (GRB) observations.  With the on-board Gamma-ray Burst Monitor (GBM) and Large Area Telescope (LAT) instruments (Atwood et al. 2009; Meegan et al. 2009), the GRB prompt emission can be probed at higher energies than ever before.  Used in conjunction with \\textit{Swift} (Gehrels et. al. 2004), GRB afterglows can be studied across a nearly continuous band from GeV energies to optical wavelengths.  As of April 1, 2010 the \\textit{Fermi} GBM has triggered on nearly 450 GRBs, 16 of which have also been seen by the LAT.  Of the 16 LAT-detected GRBs, one was simultaneously localized by the \\textit{Swift} Burst Alert Telescope (BAT; Barthelmy et al. 2005), and 8 others had \\textit{Swift} follow-up observations at late times.  The \\textit{Swift} X-ray Telescope (XRT; Burrows et al. 2005a) detected the X-ray afterglow from 7 of the 9 LAT bursts; 6 of the 7 with X-ray afterglows detected by the \\textit{Swift} UV/Optical Telescope (UVOT; Roming et al. 2005).  All but one of the UVOT afterglows stand out due to their brightness and length of detectability.  During the first 5 years of operation, \\textit{Swift} detected nearly 500 GRBs, but only a small percentage ($\\lesssim$30\\%) had bright, long-lived optical afterglows that extended beyond 10$^5$ s.  Comparing the \\textit{Swift} GRBs to the LAT GRBs, two-thirds of the LAT bursts with follow-up observations have optical afterglows that rival the brightest and longest lived of the \\textit{Swift} sample.  Such a high percentage raises the question as to whether the LAT bursts differ from the \\textit{Swift} sample.  Two possibilities are that the LAT is observing GRBs that exhibit extended energy injection, resulting in bright optical afterglows at late times, or the LAT bursts could simply be brighter, at all wavelengths, than the `average' BAT-triggered burst allowing for later detections of the afterglow (cf. Gehrels et~al.).  GRB 090926A is a LAT-detected burst with a bright, long-lived UVOT afterglow.  In this paper, we present the multiwavelength study of GRB 090926A, examining the X-ray and UV/optical wavelengths as observed by \\textit{Swift}.  In an attempt to understand the high percentage of LAT-detected bursts with optical afterglows, we also use the \\textit{Fermi} observations of the prompt emission to calculate the expected fluence as would have been observed by the BAT.  We perform this same calculation for the 6 other LAT bursts also detected by XRT and compare them to a sample of BAT-triggered bursts.  We use the power-law representation of flux density, $f{_\\nu}(t) \\propto t^{\\alpha} \\nu^{\\beta}$, where $\\alpha$ and $\\beta$ are the temporal and spectral indices, respectively.  Errors are reported at 1$\\sigma$, unless otherwise specified. ",
        "conclusions": "We have presented the \\textit{Swift} and \\textit{Fermi} observations of GRB 090926A, a recent LAT-detected GRB with a bright, long lived optical afterglow detected by UVOT.  We have compared this burst, to other LAT-detected and \\textit{Swift} BAT bursts in an attempt to show whether the GRBs detected by the LAT are simply brighter than the average BAT-triggered GRB or whether they represent a new type of GRB that commonly exhibit bright, long duration optical afterglows due to some form of energy injection.  We find that the LAT-detected bursts are generally brighter than their BAT-triggered counterparts.  We find that their fluence is consistently higher than the `average' BAT burst, and that their X-ray and UV/optical afterglows are brighter than $\\sim$80\\% of BAT GRBs.  Although we are working with a small sample of LAT bursts, and therefore suffer the consequences of small number statistics, our preliminary results indicate that LAT bursts exhibit bright late time X-ray and UV/optical afterglows because they are brighter at all wavelengths than the `average' burst, assuming the higher than average fluence can be extrapolated down to X-ray and UV/optical wavelengths.  This seems to be the most likely explanation, given the known correlation between prompt emission and afterglow emission brightness (Gehrels et~al. 2008).  We cannot say definitively, however, that this is the reason for the bright afterglows at late times, due to the presence of flares, which indicate possible late time central engine activity that could cause a rebrightening.  Without coverage of the early afterglow, it is impossible to say how the afterglow arrived at the state in which we observe it $\\sim$70 ks after the trigger.  If we simply extrapolate the optical light curve of GRB 090926A backward, we find that they could have peaked as high as \\textit{v} = 10 mag within the first hundred seconds after the trigger.  Extrapolating the LAT spectrum of GRB 090926A to the \\textit{v} band yields a peak magnitude of \\textit{v} = 4, or if we assume a cooling break at GeV energies, the spectral index changes to $\\beta \\approx -0.76$, yielding a magnitude of \\textit{v} = 15, consistent with our extrapolation backwards and the idea that LAT bursts are uniquely bright at all wavelengths.  However, if the early afterglow was fainter than \\textit{v} $\\approx$ 15 mag, then some sort of sustained energy injection would be required to keep the flux elevated at a level where we could then observe the bright afterglow at 70 ks after the trigger.  Such an energy injection would test our current theoretical understanding of GRB optical afterglows.  Our ability to determine the true nature of LAT-detected burst is contingent on our ability to follow-up LAT-detected GRBs at earlier times than has been achieved with the current sample."
    },
    "1004/1004.0958_arXiv.txt": {
        "abstract": "We report the discovery of three nearby old halo white dwarf candidates in the Sloan Digital Sky Survey (SDSS), including two stars in a common proper motion binary system. These candidates are selected from our 2800 square degree proper motion survey on the Bok and U.S. Naval Observatory Flagstaff Station 1.3m telescopes, and they display proper motions of $0.4-0.5\\arcsec$ yr$^{-1}$. Follow-up MMT spectroscopy and near-infrared photometry demonstrate that all three objects are hydrogen-dominated atmosphere white dwarfs with $T_{\\rm eff} \\approx 3700 - 4100$ K. For average mass white dwarfs, these temperature estimates correspond to cooling ages of $9-10$ Gyr, distances of $70-80$ pc, and tangential velocities of $140-200$ km s$^{-1}$. Based on the $UVW$ space velocities, we conclude that they most likely belong to the halo. Furthermore, the combined main-sequence and white dwarf cooling ages are 10-11 Gyr. Along with SDSS J1102+4113, they are the oldest field white dwarfs currently known. These three stars represent only a small fraction of the halo white dwarf candidates in our proper motion survey, and they demonstrate that deep imaging surveys like the Pan-STARRS and Large Synoptic Survey Telescope should find many old thick disk and halo white dwarfs that can be used to constrain the age of the Galactic thick disk and halo. ",
        "introduction": "White dwarf (WD) cosmochronology provides an independent and accurate age dating method for different Galactic populations \\citep{winget87,liebert88}. Using 43 cool WDs in the Solar Neighborhood, \\citet{leggett98} derived a disk age of 8 $\\pm$ 1.5 Gyr. \\citet{kilic06} and \\citet{harris06} significantly improved the field WD sample by using SDSS and USNO-B astrometry to select high proper motion candidates. However, their survey suffered from the magnitude limit of the Palomar Observatory Sky Survey plates and they were unable to find many thick disk or halo WD candidates.  Substantial investment of the $Hubble~Space~Telescope$ time on two globular clusters, M4 and NGC 6397, revealed clean WD cooling sequences. \\citet{hansen04,hansen07} and \\citet{bedin09} use these data to derive cooling ages of $\\approx12$ Gyr for the two clusters. The coolest WDs in these clusters are about 650 $\\pm$ 230 K cooler than the coolest WDs in the disk \\citep{kowalski07a}. These studies demonstrate that the Galactic halo is older than the disk by $\\geq2$ Gyr \\citep{hansen02,fontaine01,kowalski07a}. Even though the WDs in globular clusters provide reliable age estimates, these clusters may not represent the full age range of the Galactic halo. The required exposure times to reach the WD terminus in globular clusters limit these studies to the nearest few clusters. In addition, only two-filter ($V$ and $I$) photometry is used to model the absolute magnitude and color distribution of the oldest WDs to derive ages. The far closer and brighter WDs of the local halo field are an enticing alternative as well as complementary targets, with the additional potential to constrain the age range of the Galactic halo. Accurate ages for field WDs can be obtained through optical and near-infrared photometry and trigonometric parallax measurements. Nearby WDs can also be used to understand the model uncertainties and put the Globular cluster ages on a more secure footing.  The quest for field halo WDs has been hampered by the lack of proper motion surveys that go deep enough to find the cool halo WDs. The initial claims for a significant population of halo WDs in the field \\citep{oppenheimer01a} and in the Hubble Deep Field \\citep{ibata00,mendez00} were later rejected by detailed model atmosphere analysis \\citep[see][and references therein]{bergeron05} and additional proper motion measurements \\citep{kilic04,kilic05}. To date, the coolest known probable halo WDs are WD 0346+246 \\citep{hambly97,bergeron01} and SDSS J1102+4113, with $T_{\\rm eff} \\approx 3800$ K \\citep{hall08}. There are also about a dozen ultracool WDs detected in the SDSS \\citep{gates04,harris08} that may be thick disk or halo WDs, but current WD atmosphere models have problems in reproducing their intriguing spectral energy distributions (SEDs). Therefore, their temperatures and ages remain uncertain.  Here we report the identification of three old halo WD candidates discovered as part of our Bok and USNO proper motion survey. The details of this survey and our follow-up observations are discussed in Section 2, whereas our model fits and analysis are discussed in Section 3. ",
        "conclusions": "J2137+1050 and J2145+1106 are cool WDs with hydrogen-dominated atmospheres. Our effective temperature estimates of 3730-3780 K make J2137+1050 and J2145+1106N the coolest WDs known in the Solar neighborhood. Our best-fit models imply total ages of $\\approx 10-11$ Gyr, distances of 70-80 pc, and Galactic space velocities that are inconsistent with thick disk population within $2\\sigma$. We conclude that these targets most likely belong to the halo. However, trigonometric parallax observations are required in order to constrain the distances, masses, and ages of our targets accurately. Such observations are currently underway at the MDM 2.4m telescope.  Like WD 0346+246 and SDSS J1102+4113 \\citep{bergeron01,hall08}, our three halo WD candidates have hydrogen-rich atmospheres. The oldest WDs are likely to accrete from the interstellar medium within their $\\sim$10 Gyr lifetimes and end up as hydrogen-rich WDs even if they start with a pure helium atmosphere. However, the current sample of halo WD candidates is not large enough to conclude that most or all of the oldest WDs are hydrogen-rich. Observations of larger samples of field WDs will be necessary to check whether all WDs turn into hydrogen-rich atmosphere WDs or not \\citep[see the discussion in][]{kowalski06b}.  The three targets that we present here make up only a small fraction of the halo WD candidates in our proper motion survey. Follow-up observations of these targets will be necessary to confirm many more halo WD candidates that can be used to study the age and age dispersion of the Galactic thick disk and halo. Already we can see, however, that these halo (or possibly thick disk) WDs indicate a gap of 1--2 Gyr between the star formation in the halo and the star formation in the disk at the solar annulus. Our observations further demonstrate that deep, wide-field proper motion surveys ought to find many old halo WDs. Using the \\citet{liebert88} WD luminosity function for the Galactic thin disk and a single burst 12 Gyr old population with 10\\% and 0.4\\% local normalization for the thick disk and halo, we estimate that there are 3200 thick disk and 140 halo WDs per 1000 square degree (for a Galactic latitude of 45$^\\circ$) down to a limiting magnitude of $V=21.5$ mag (our survey limit). Pushing the limiting magnitude down to $V=24$ mag and assuming 50\\% sky coverage, we estimate that future surveys like the Pan-STARRS and LSST will image $\\sim$ 1.3 million thick disk and $\\sim$ 80,000 halo WDs. These surveys will be invaluable resources for halo WD studies."
    },
    "1004/1004.0675_arXiv.txt": {
        "abstract": "We present the McMaster Unbiased Galaxy Simulations (MUGS), the first 9 galaxies of an unbiased selection ranging in total mass from 5$\\times10^{11}$ M$_\\odot$ to 2$\\times10^{12}$ M$_\\odot$ simulated using n-body smoothed particle hydrodynamics (SPH) at high resolution.  The simulations include a treatment of low temperature metal cooling, UV background radiation, star formation, and physically motivated stellar feedback.  Mock images of the simulations show that the simulations lie within the observed range of relations such as that between color and magnitude and that between brightness and circular velocity (Tully-Fisher).  The greatest discrepancy between the simulated galaxies and observed galaxies is the high concentration of material at the center of the galaxies as represented by the centrally peaked rotation curves and the high bulge-to-total ratios of the simulations determined both kinematically and photometrically.  This central concentration represents the excess of low angular momentum material that long has plagued morphological studies of simulated galaxies and suggests that higher resolutions and a more accurate description of feedback will be required to simulate more realistic galaxies.  Even with the excess central mass concentrations, the simulations suggest the important role merger history and halo spin play in the formation of disks. ",
        "introduction": "\\label{intro} Forming a galaxy like our own Milky Way remains a challenge for the currently accepted $\\Lambda$ Cold Dark Matter ($\\Lambda$CDM) cosmogony.  The Milky Way is comprised of three distinct, stellar components: a flattened, rotating \\emph{disk}; a compact, central and spheroidal \\emph{bulge}; and a diffuse, spherical \\emph{halo} of stars.  Any consistent cosmogony needs to explain the origin and evolution of each of these components.  $\\Lambda$CDM posits that the energy budget of the Universe is currently dominated by vacuum energy ($\\Lambda$), and the majority of the mass is invisible (dark) and only interacts with baryons via gravity.  Thus, in the early Universe, thermal baryonic pressure did not support the dark matter, and because it is non-relativistic (cold) the dark matter first collapsed into small structures.  Subsequently, the small structures merged hierarchically to form larger structures like the Milky Way.  The $\\Lambda$CDM paradigm provides a simple explanation for the formation of the stellar halo:  stars formed early in small satellite galaxies, which got tidally stripped as their orbits brought them inside the tidal radius of the main galaxy.  While early observations indicated that stars in the Milky Way's halo might have condensed out of a monolithically collapsing gas cloud \\citep{eggen62}, later observations found instead that formation through mergers like those proposed in the CDM paradigm are more likely \\citep{searle78}.  Today, digital surveys of the sky reveal structures in the Milky Way's stellar halo such as streams and the remnant cores of dwarf galaxies \\citep{Majewski1993,Belokurov2006}.  These are the exact signatures left by tidally disrupted satellites in simulations \\citep{bullock05,abadi06}.  Unlike halos, disks are not so neatly explained by $\\Lambda$CDM.  Although conservation of angular momentum naturally creates rotating disks, the hierarchical buildup of structure impairs their formation, since disks form most efficiently in a quiet environment where gas cools and collapses smoothly.  In $\\Lambda$CDM-inspired simulations of substructure mergers, satellites orbiting disks tidally heat stars turning thin disks into thick disks \\citep{Toth1992,Quinn1993,Velazquez1999,Kazantzidis2008}.  Simulations of larger galactic mergers transform disks into centrally concentrated, spheroidal systems as the disks experience significant angular momentum loss \\citep{Barnes1996,Cox2006}.  However, the observed distribution of galaxy morphologies can be reproduced with simple, analytic models.  In these models uniformly rotating spheres collapse into centrifugally supported disks \\citep{Fall1980,Dalcanton97,mo98,vdB2001}.  The spheres start with angular momentum and mass profiles predicted using simulations of CDM structure formation.  The contradiction between mergers and disk formation may indicate that halo spin, not merger history, plays the dominant role in determining the morphology of galaxies.   The bulge is a spheroid of stars at the center of a galaxy.  The spheroidal shape indicates that they formed as the result of mergers.  However, other evidence reveals that some bulges may have a secular origin.  \\citet{Kormendy1993} suggested that observations of rapidly rotating bulges indicates the existence of ``pseudo-bulges'', and recent simulations show that the central regions of isolated disks can buckle and cause stars to evolve through ``peanut'' shaped orbits into spheroidal distributions \\citep{Debattista2004}.  Whether disks or spheroids form affects many galactic properties like their kinematics, color, light distribution, star formation history and metallicity in addition to morphology.  Disk kinematics are dominated by ordered circular velocity, while in spheroidal components random velocity dominates over circular velocities.  Because star formation usually happens in galaxy disks where gas densities are sufficiently high, galaxies with more prominent disks have more recent star formation and display bluer colors, while spheroids tend to be more metal-rich and red.  The prominence of morphological components also has an impact on the radial distribution of galactic light profiles.  Galaxies with more prominent spherical components exhibit more centrally concentrated light profiles. \\\\ Models of galaxy formation require high resolution, hydrodynamic numerical simulations.  Analytic modeling can evolve a $\\Lambda$CDM-motivated Gaussian density field into a spectrum of mass structures and populate those structures with stars such that they match the observered luminosity function of galaxies \\citep{Cole2000, Benson2003, Somerville2008}.  However, because of the non-linear interactions of processes such as gas cooling, merging, tidal stripping, star formation, stellar feedback and active galactic nuclei, it is difficult for these models to predict galaxy morphologies, though \\citet{Benson2009} and \\citet{Dutton2009} represent recent attempts.  As an alternative, numerical simulations allow us to study how halos, disks, and bulges were created and evolve.  Simulations that include gas are able to follow more physical processes than simulations that only track the gravitational interaction of dark matter.  Gas can be modeled using smoothed particle hydrodynamics (hereafter SPH), which partitions the gas in the Universe into particles and  with a Lagrangian approach follows the motions of those particles.  SPH is effective because it concentrates computational resources on the high density regions of a simulation, where galaxies form.  Several studies of galaxies in a cosmological context have generated individual objects that are similar to the observed local galaxies \\citep{Governato2004, Robertson2004, Okamoto2005, Brook2006, Governato2007, Scannapieco2008, Scannapieco2009, Ceverino2009, SanchezBlazquez2009, Governato2009, Piontek2009, Martinez-Serrano2009}. These require high resolution and large computational resources in order for several important properties of the simulations to converge, such as the galactic structure, motions and star formation history.  As a consequence of the high computational cost, the initial conditions have been carefully chosen to maximize the chance that the simulation will produce the desired type of object (usually a disk galaxy).  The previous cosmological simulations have shown that the collapse of gas is more complicated than smooth collapse into centrifugally supported disks.  They show that a significant fraction of gas flows into galaxies along cold filaments \\citep{Keres05,Brooks2009}.  This raises the question: are centrifugal forces all that should support disks?  Observations suggest that disks are supported by an equipartition of energy between thermal, magnetic, and cosmic ray pressure support \\citep{Cox2005}.  \\subsection{Stellar Feedback} One way to introduce pressure support into simulations is by harnessing the energy massive stars release in stellar winds and supernovae explosions. Recent simulations continue to suffer from the overcooling that has long plagued morphological studies of simulated galaxies.  \\citet{Navarro1991} described how angular momentum transferred from gas in the disk to the dark matter causes excess central mass concentrations, which leads to massive bulges that do not compare well to observations of disk galaxies whose bulges are fainter and less massive than their disks \\citep{Allen2006}.  \\citet{Navarro1991} instead proposed that stellar feedback could eliminate a significant amount of low angular momentum gas.  They implemented a method to kinematically excite gas particles around recently formed stars, but this showed little improvement in eliminating low angular momentum gas \\citep{Navarro2000}.  Stellar feedback plays a larger role in the development of satellite galaxies that merge with the parent galaxy.  The effects of feedback in satellites can contribute to the final morphology of the main galaxy.  If a satellite brings in more gas, it will contribute to make a larger disk; more stars will produce a larger spheroid.  In the main galaxies, stellar feedback may also determine the rate at which gas loses angular momentum and migrates through the disk into the dense, star forming center.  Due to insufficient resolution, there is currently no satisfactory treatment for stellar feedback.  Multiphase gas particles attempt to capture the phenomenology of the ISM inside individual particles \\citep{SH03}, but it is difficult to determine the appropriate pressure for each particle to exert on the others and sometimes a stiff equation of state needs to be enforced \\citep{Springel05}.  If the gas dynamics are separated into hot and cold gas \\citep{Ritchie2001}, it is ambiguous when and how much gas should move from one phase to the other.  Disabling the gas cooling reproduces stellar feedback \\citep{Gerrit97, TC2000, Stinson2006}, but resolution is often insufficient to turn off cooling in the proper amount of gas for the right length of time, and SPH does not allow single particles to create their own outflow.  Driven winds reproduce galactic outflows \\citep{Navarro2000, Springel03, Oppenheimer2006, Okamoto2009}, but it is yet to be determined whether wind particles should be allowed to interact with surrounding gas and whether they provide sufficient pressure support to forming disks.  Many potential avenues need to be followed to see how each feedback recipe affects galaxy formation differently.  To date, no stellar feedback model has been effective at reducing the central mass concentration in high mass galaxies, though strong adiabatic feedback combined with altering the star formation threshold to a higher density in very high resolution simulations of low mass galaxies has had the most success \\citep{Mashchenko2008,Governato2010}.  Computational resources can now support surveys that include a range of galaxies simulated at moderate (1 million dark matter particles inside $r_{vir}$) resolution \\citep{Scannapieco2009,Okamoto2009,Piontek2009} with different stellar feedback recipes.  Each survey uses SPH, and each of the previous surveys have used \\textsc{gadget}.  \\citet{Scannapieco2009} separated hot and cold gas and used supernova feedback as a conduit from the cold to the hot phase.  \\citet{Okamoto2009} used multiphase particles from \\citet{Springel05} combined with driven winds of different strengths.  \\citet{Scannapieco2009} found lower disk fractions than late type spirals in their simulations.  \\citet{Piontek2009} tested several stellar feedback recipes and did not find striking success with any of them.  Observational samples of galaxies from galaxy redshift surveys, like the 2dFGRS (Colless et al. 2001) and SDSS (York et al. 2000), now contain millions of objects, allowing a much more complete view of not only the properties of typical galaxies, but also a quantification of how galaxies are distributed within the multivariate parameter space of galaxy properties. In contrast, while many researchers are performing sophisticated galaxy formation simulations, each can only produce a handful of galaxies. Evaluating the success of these simulations requires a larger sample of simulated galaxies. When only a small number of simulations exist, it is easy to find a good observational match for any one simulation; however, when a \\emph{sample} of simulated galaxies exist that predict a mean and spread for any galactic property, these predictions can be directly confronted with the large observational samples.  In order to address these problems, we have begun the McMaster Unbiased Galaxy Simulations (MUGS) project. The goal of MUGS is to generate a large sample of sophisticated galaxy formation simulations that sample potential sites of L* galaxy formation in an unbiased manner for direct comparison to the large observational samples now available. In this paper, we describe the methodology used for MUGS and present an overview of the first 9 simulations, particularly focusing on the relative formation of disks and bulges.  MUGS provides an extended look at galaxies simulated using similar physics to \\citet{Governato2007} and \\citet{Governato2009}.  Namely, supernovae are modeled with adiabatic ``blastwave'' feedback described in \\citet{Stinson2006}.  In \\S \\ref{sec:sims}, we describe how we created the initial conditions for MUGS.  \\S \\ref{sec:code} details the algorithm that evolves the simulations. \\S \\ref{sec:results} examines the properties of the galaxies including their brightness, color, mass-to-light ratios, density profiles, bulge-to-total ratio, star formation history, and metallicity. ",
        "conclusions": "We presented 9 galaxies from the McMaster Unbiased Galaxy Simulations simulated using N-body gravity and SPH.  The galaxies are selected from a mass range around the mass of the Milky Way and from isolated environments, but their selection was otherwise unbiased for factors such as accretion history and angular momentum.  The galaxies were examined using the radiative transfer program \\textsc{sunrise} to enable comparisons between the simulated galaxies and real galaxies in the observed plane.  The simulated galaxies have colors and magnitudes that compare well with a sample of inclination corrected isolated galaxies from SDSS, and in particular separate into the well-known red sequence and blue cloud.  However, both simulated populations tend too much towards the ``green valley'', indicating that they contain more old stars than blue cloud galaxies and more young stars than galaxies along the red sequence.  The surface brightness profiles of the simulated galaxies can all be fit with exponential disks combined with a central de Vaucouleurs $r^{1/4}$ law similar to real galaxies.  However, the proportion of bulge to total light (B/T) is higher than what is typically observed.  The B/T ratios are also high when the stars are decomposed into the spheroid and disk based on their kinematics.  There are no galaxies with a B/T fraction less than 0.5 whereas observed samples find many galaxies with B/T $<$ 0.5.  This result is similar to that found in many previous simulations.  We note that many of the recently formed stars that are classified as part of the spheroid form with orbits in the disk plane in the central regions of the disk.  We also note that most of the stars that comprise the spheroid formed \\emph{in situ}, but we leave the question of how the stars may migrate from the disk to a spherical distribution for future work.  As to the question of why galaxies form with more or less spheroid, there seems to be a modest trend with accretion history.  We find that the largest disks (g1536 and g15784) form with a quiet merger history in which they had their last major merger prior to $z$=3, while the largest bulges all resulted from recent last major mergers.  There also appears to be a dependence on halo spin as g25271 has a relatively quiet merger history, but low $\\lambda'$ and results in a significant bulge and red color.  Since all the galaxies used in MUGS fall in a limited mass range, these conclusions do not include the impact of mass.  We compare the brightness of the final galaxies with their mass at two different radii in the final output.  First, we compare our galaxies with the observed Tully-Fisher relationship that probes the amount of light a galaxy produces with the mass contained in its inner regions.  Since the final mass concentration of all the galaxies is too high, we use a rotation velocity from 3.5 disk scale lengths away from the galaxy center.  Previous studies have shown that this is the radius at which rotation velocities converge.  The galaxies are still slightly \\emph{fainter} than the observed sample based on these inner velocities.  Second, we compare the brightness of our galaxies with observations of the whole halo mass derived using a number of different methods.  In each case, the simulated galaxies are \\emph{brighter} than  comparable observed galaxies at the same mass.  This indicates that too many stars form in the simulation.  The high central velocities used in the Tully-Fisher relationship indicate that mass gets too concentrated at the centers of the halos and while the galaxies are fainter than the observed TF relationship, the lack of resolution makes it difficult to determine whether the amount of stars formed is too many or too few.  We are thus left with the challenge of creating more realistic simulations in order to obtain more accurate insights into how the important physical processes involved in galaxy formation result in the observed population of galaxies.  Fortunately, there has been much recent work that guides the way forward.  It has been shown repeatedly that higher resolution makes better disks \\citep{Governato2004,Governato2007}.  More recently, it has been shown that high resolution combined with clustered star formation can remove central density concentrations from dwarf galaxies \\citep{Mashchenko2008,Governato2010}.  While these simulations open many possibilities for expanding our understanding of how galaxies form, they also show that there is much work left to be done before we can claim to have simulated a sample of galaxies that compares well with real ones."
    },
    "1004/1004.4506_arXiv.txt": {
        "abstract": "{It has recently been suggested that the compilation of the calibrated time-ordered-data (TOD) of the Wilkinson Microwave Anisotropy Probe (WMAP) observations of the cosmic microwave background (CMB) into full-year or multi-year maps may have been carried out with a small timing interpolation error. \\protect\\postrefereechanges{A large fraction} of the previously estimated WMAP CMB quadrupole signal would be an artefact of incorrect Doppler dipole subtraction if this hypothesis were correct.  } {Since observations of bright foreground objects constitute part of the TOD, these can be used to test the hypothesis. } {Scans of an object in different directions should be shifted by the would-be timing error, causing a blurring effect.  For each of several different timing offsets, three half-years of the calibrated, filtered WMAP TOD are compiled individually for the four W band differencing assemblies (DA's), with no masking of bright objects, giving 12 maps for each timing offset. Percentiles of the temperature-fluctuation distribution in each map at HEALPix resolution $N_{\\mathrm{side}}=2048$ are used to determine the dependence of all-sky image sharpness on the timing offset. The Q and V bands are also considered.} {In the W band, which is the band with the shortest exposure times, the 99.999\\% percentile, i.e. the temperature fluctuation in the $\\approx$ 503-rd brightest pixel, is the least noisy percentile as a function of timing offset. Using this statistic, the hypothesis that a $-25.6$~ms offset relative to the timing adopted by the WMAP collaboration gives a focus at least as sharp as the uncorrected timing is rejected at $4.6\\sigma$ significance, assuming Gaussian errors and statistical independence between the maps of the 12 DA/observing period combinations.  The Q and V band maps also reject the $-25.6$~ms offset hypothesis at high statistical significance. } {The requirement that the correct choice of timing offset must maximise image sharpness implies that the hypothesis of a timing error in the WMAP collaboration's compilation of the WMAP calibrated, filtered TOD is rejected at high statistical significance in each of the Q, V and W wavebands. However, the hypothesis that a timing error was applied during {{\\em calibration\\/}} of the raw TOD, \\protect\\postrefereechanges{leading to a dipole-induced difference} signal, is not excluded by this method. } ",
        "introduction": "\\label{s-intro} The nature of the large-scale cosmic microwave background (CMB) signal in the Wilkinson Microwave Anisotropy Probe (WMAP) observations \\nocite{WMAPbasic}({Bennett} {et~al.} 2003b) is of fundamental importance to observational cosmology. A lack of structure on the largest scales in maps of CMB temperature fluctuations would be a hint that the comoving spatial section of the Universe is multiply connected \\nocite{deSitt17,Fried23,Fried24,Lemaitre31ell,Rob35}({de Sitter} 1917; {Friedmann} 1923, 1924; {Lema{\\^i}tre} 1931; {Robertson} 1935) and that its size may have been detected \\nocite{Star93,Stevens93}({Starobinsky} 1993; {Stevens} {et~al.} 1993). The WMAP observations have confirmed that the large-scale CMB signal analysed either as a quadrupole signal or as a two-point autocorrelation function signal is weak \\nocite{WMAPSpergel,Copi07,Copi09}({Spergel} {et~al.} 2003; {Copi} {et~al.} 2007, 2009).  Several analyses indicate that either a $T^3$ comoving space \\nocite{WMAPSpergel,Aurich07align,Aurich08a,Aurich08b,Aurich09a}({Spergel} {et~al.} 2003; {Aurich} {et~al.} 2007; {Aurich} 2008; {Aurich} {et~al.} 2008, 2010) or a Poincar\\'e dodecahedral space ($S^3/I^*$) model \\nocite{LumNat03,Aurich2005a,Aurich2005b,Gundermann2005,Caillerie07,RBSG08,RBG08}({Luminet} {et~al.} 2003; {Aurich} {et~al.} 2005a, 2005b; {Gundermann} 2005; {Caillerie} {et~al.} 2007; {Roukema} {et~al.} 2008a, 2008b) fit the WMAP data better than an infinite, simply connected flat model, while other analyses disagree \\nocite{KeyCSS06,NJ07}({Key} {et~al.} 2007; {Niarchou} \\& {Jaffe} 2007). Highly significant evidence for either an infinite or a finite model of comoving space has not yet been obtained.  \\nocite{LL08hotsources}{Liu} \\& {Li} \\nocite{LL08hotsources,LL09Nobs,LL09lowquad,LL10toffset}({Liu} \\& {Li} 2008; {Li} {et~al.} 2009; {Liu} \\& {Li} 2009; {Liu} {et~al.} 2010) have suggested that several systematic errors may have been made in the data analysis pipelines used by the WMAP collaboration in the compilation of their original time-ordered-data (TOD) into single-year or \\postrefereechanges{multi-year} maps. \\nocite{LL08hotsources}{Liu} \\& {Li} (2008) recommended that new maps should be made using the full TOD in order to avoid some of these suspected errors.  This requires moderately heavy computational resources (RAM and CPU) on typical present-day desktop computers.  \\nocite{Aurich09a}{Aurich} {et~al.} (2010) sidestepped this by using a phenomenological correction for the hot pixel effect \\nocite{LL08hotsources}({Liu} \\& {Li} 2008). They found that their correction applied to the 5-year WMAP W band data removes the anti-correlation in temperature fluctuations at angular separations of nearly 180{\\ddeg} and improves the fit of a $T^3$ model to the data.  \\fartificialquadrupole  \\nocite{LL09lowquad}{Liu} \\& {Li} (2009) carried out a full analysis of the 5-year WMAP TOD. \\postrefereechanges{Their analysis pipeline appeared to give test results largely compatible with those of the WMAP collaboration.  However, the CMB quadrupole amplitude was found to be much weaker, and sub-degree power was also found to be a little weaker.} Later, \\nocite{LL10toffset}{Liu} {et~al.} (2010) traced the difference between their analysis and that of the WMAP collaboration to a difference in interpolating the timing of individual observations from the times recorded in the TOD files. The TOD FITS format files contain observational starting times in both the ``Meta Data Table'' and the ``Science Data Table''. The full set of observing times of individual observations are not recorded in either data table; they must be interpolated from the smaller set of observing times that are recorded. The starting times in the Meta Data Table are 25.6~ms greater than the corresponding times in the Science Data Table. A computer script for checking this is given in Appendix~\\ref{a-meta_vs_science}. The durations of individual observations in the Q, V and W bands are 102.4~ms, 76.8~ms, and 51.2~ms, respectively \\nocite{WMAPExplSupp100405}(Section 3.2, {Limon} {et~al.} 2010). The 21 April 2010 version of the WMAP Explanatory Supplement \\nocite{WMAPExplSupp100405}(Section 3.1,  {Limon} {et~al.} 2010) discusses the relation between the start times in the two tables, but does not refer to the 25.6~ms offset, i.e. half of a W band observing interval or a quarter of a Q band observing interval. \\postrefereechanges{\\nocite{Bennett03MAP}{Bennett} {et~al.} (2003a) state (\\SSS{6.1.2}) that ``a relative accuracy of 1.7~ms can be achieved between the star tracker(s), gyro and the instrument'', i.e. the combined random and systematic error should be much smaller than 25.6~ms.} Given that the timing offset is numerically implicit in the WMAP TOD files, it could, strictly speaking, be referred to as the WMAP collaboration's ``implicitly claimed but ignored'' timing offset. However, for simplicity, the timing offset will be referred to here as \\nocite{LL10toffset}{Liu} {et~al.} (2010)'s hypothesis.    Using their Eqs~(1) and (2), \\nocite{LL10toffset}{Liu} {et~al.} (2010) showed that the error in estimating the Doppler dipole of the spacecraft velocity implied by the would-be timing error can be used to produce a map \\nocite{LL10toffset}(Fig.~2 left,  {Liu} {et~al.} 2010) that matches the direction and amplitude of the CMB quadrupole signal as estimated by the WMAP collaboration \\nocite{LL10toffset}(Fig.~2 right,  {Liu} {et~al.} 2010) remarkably well, despite using only the history of spacecraft attitude quaternions (representing three-dimensional directions) and no sky observations. This is easy to check using the WMAP calibrated, filtered, 3-year TOD and \\nocite{LL10toffset}{Liu} {et~al.} (2010)'s software. Figure~\\ref{f-yr1V1mr80} shows a \\protect\\postrefereechanges{dipole-induced difference} map\\footnote{\\protect\\postrefereechanges{\\nocite{LL10toffset}{Liu} {et~al.} (2010) refer to this as a ``differential dipole'' map. However, this term is frequently used to refer to monopole-subtracted dipole maps. Moreover, the difference between two (fixed) dipole signals in different directions is itself also a dipole, not a quadrupole. To avoid misinterpretations, the term ``dipole-induced difference'' is suggested here.}} which is consistent with \\nocite{LL10toffset}{Liu} {et~al.} (2010)'s \\protect\\postrefereechanges{dipole-induced difference} map shown in their Fig.~2 (left), for a half-observing-interval offset in the V band, i.e. $-38.4$~ms. Using a velocity dipole model and a $-25.6$~ms offset, \\nocite{MSS10}{Moss} {et~al.} (2010) find a similar \\postrefereechanges{but weaker effect of reduction in the quadrupole amplitude.} \\postrefereechanges{Earlier, \\nocite{Jarosik06Beam}{Jarosik} {et~al.} (2007) (\\SSS{2.4.1}) noted that an error during the calibration process could introduce an artificial, dipole-induced quadrupole.}  Interpolation of the observing times and sky positions from the WMAP TOD files cannot be avoided. Given that the WMAP Explanatory Supplement does not presently account for the 25.6~ms offset between the two data tables in each TOD file, and given the potential importance of a 70--80\\% overestimate of the CMB quadrupole amplitude, it is useful to see if a relatively simple, robust method of analysing the TOD can determine which of the two timing interpolation methods, that of the WMAP collaboration or that proposed by \\nocite{LL10toffset}{Liu} {et~al.} (2010), is correct.  The WMAP TOD include many observations of bright foreground objects, including Solar System planets and objects in the Galactic Plane.  The complex scanning pattern of WMAP implies that individual objects are likely to be scanned in different directions over a series of successive observations.  The greatest differences in scanning directions are likely to occur over observing intervals of weeks or more. An error in the assumed sky direction due to a timing error should cause observations of a compact source made at different times to be slightly offset from one another in the compiled map, i.e. it should cause an overall shift in sky position and a blurring effect. For a large enough timing offset, double imaging of compact sources should occur (public communication, Crawford 2010\\footnote{\\url{http://cosmocoffee.info/viewtopic.php?t=1537}}; see also \\nocite{MSS10}{Moss} {et~al.} 2010). The W band observations, with the highest angular resolution, have an estimated FWHM beam size of about 12--12.6{\\arcm} \\nocite{WMAP1beam}(Table 5,  {Page} {et~al.} 2003). An error of 25.6~ms corresponds\\footnote{One spin during 2.2 minutes, i.e. \\protect\\postrefereechanges{129.3~s, \\protect\\nocite{WMAP03syserr}(e.g. Sect~3.4.1,  {Hinshaw} {et~al.} 2003)} corresponds to $360\\ddeg \\sin(70.5\\ddeg) = 339\\ddeg$ in great circle degrees, so 25.6~ms corresponds to \\postrefereechanges{4.0{\\arcm}}. } to an error of about \\postrefereechanges{4.0{\\arcm}}. Hence, a slight blurring of the images of compact sources rather than double imaging should occur. \\postrefereechanges{If the WMAP collaboration's timing offset were correct, then this blurring could explain the slight loss of sub-degree power found by \\nocite{LL10toffset}{Liu} {et~al.} (2010) as an artefact for the latter's choice of timing offset.} \\postrefereechanges{The blurring} could be detectable in beam shape analysis, but this may not be easy.  \\nocite{WMAP1beam}{Page} {et~al.} (2003) model the beam shape based on the physical structure of the WMAP receivers, but estimate centroids iteratively. It is not obvious that this iterative procedure would have detected an artificial blurring effect of just a few arcminutes.  \\fGCmfive  \\fGCzero \\fGCpzerofive  \\fGCpfive  Nevertheless, the blurring effect should be statistically detectable in maps compiled from the TOD. The effect should be strong if the timing error is exaggerated beyond that proposed by \\nocite{LL10toffset}{Liu} {et~al.} (2010). By considering the timing offset as a free parameter, it should be possible to find the value that minimises blurring. This should correspond to the correct choice of timing interpolation. In principle, since a timing error also induces an error in absolute sky position, this could also be used to determine the correct method. However, it is conceivable that a positional offset may have been at least partially removed at some stage in the pointing calibration, making it hard to detect, as suggested by \\nocite{MSS10}{Moss} {et~al.} (2010). On the other hand, a blurring effect is unlikely to have been removed at any step in the production of WMAP maps. Hence, the analysis here uses the blurring effect. In \\SSS\\ref{s-method}, an empirical method of estimating the blurring using all-sky maps is presented. Images of some Galactic Centre objects as a function of timing offset and the results of the statistical analysis are presented in \\SSS\\ref{s-results}. Conclusions are given in \\SSS\\ref{s-conclu}. ",
        "conclusions": "\\label{s-conclu}  The difference in image sharpness for \\nocite{LL10toffset}{Liu} {et~al.} (2010)'s and the WMAP collaboration's respective timing interpolation choices are not obvious in the images of the Galactic Centre shown in Figs~\\ref{f-GC00} and \\ref{f-GCp05}. However, use of the all-sky images gives estimates of image sharpness that depend strongly on the fractional time offset $\\delta t$.  These calculations do not require any fitting of source or beam profiles, nor any selection of one or more preferred sources, nor any assumptions regarding pointing accuracy.   The data analysis pipeline is performed using \\nocite{LL10toffset}{Liu} {et~al.} (2010)'s script, with only minor modifications, on slightly under a quarter of the full set of WMAP calibrated, filtered TOD. \\nocite{LL10toffset}{Liu} {et~al.} (2010)'s timing error hypothesis is rejected at very high significance in the Q, V and W wavebands separately.  The lack of an explanation in the WMAP Explanatory Supplement \\nocite{WMAPExplSupp100405}(Section 3.1,  {Limon} {et~al.} 2010) for the 25.6~ms offset between the start times in the Meta Data Set and the Science Data Set in the TOD files is an unfortunate gap in the documentation, but it clearly did not lead to an error of this magnitude in the WMAP collaboration's compilation of the calibrated, filtered TOD.  A possible explanation for the offset could be that it is a convention inserted by software into the Meta Data Set that was forgotten about, rather than a physical time offset. In principle, it would be good if the WMAP collaboration could find the line(s) in their software where the 25.6~ms offset is inserted, in order to confirm that its origin is fully understood, and that it has no unintended consequences.  Nevertheless, the fact that \\nocite{LL09lowquad}{Liu} \\& {Li} (2009) accidentally applied a time offset that approximately cancels the $\\delta t = 0.5$ CMB quadrupole in amplitude and direction, and that this can be modelled as a \\protect\\postrefereechanges{Doppler-dipole-induced difference} signal \\nocite{LL10toffset}(Eq.~(2), {Liu} {et~al.} 2010, ``deviation of differential dipole''), remain interesting coincidences. The relations between the quadrupole, octupole and ecliptic \\nocite{Copi06ecliptic,Copi10}(e.g.,  {Copi} {et~al.} 2006; {Sarkar} {et~al.} 2010) have long been known. What remains unestablished is whether these are coincidences or artefacts. \\nocite{MSS10}{Moss} {et~al.} (2010) discuss the possibility of relationships between the various coincidences and suggest that another effect might couple with the WMAP spacecraft velocity dipole in order to produce an artificial component of the quadrupole.  One interesting possibility for further work would be to consider a variation on the hypothesis proposed by \\nocite{LL10toffset}{Liu} {et~al.} (2010): could it be possible that the would-be timing error was introduced when {\\em calibrating the uncalibrated\\/} TOD? \\postrefereechanges{Although the velocity dipole is only removed from the data during the mapmaking step, it is first used for calibrating the raw data into a calibrated signal in mK. So, a} small timing error during the calibration step could correspond to artificially adding a \\protect\\postrefereechanges{dipole-induced difference} signal.  \\postrefereechanges{During the following step of compiling the calibrated TOD into maps, the dipole would then be subtracted at the correct positions, but without compensating for the incorrect calibration.} \\postrefereechanges{As noted above, \\nocite{Jarosik06Beam}{Jarosik} {et~al.} (2007) (\\SSS{2.4.1}) have earlier considered the possibility that a calibration error (not necessarily induced by a timing error) could introduce an artificial, dipole-induced quadrupole.}  The calibration method for the 7-year WMAP analyses are stated \\nocite{WMAP7JarosikBasic}({Jarosik} {et~al.} 2010) to be those used for the 5-year analyses, i.e. as described in \\SSS{4} of \\nocite{WMAP5Hinshaw}{Hinshaw} {et~al.} (2009), retaining the estimate of about 0.2\\% absolute calibration error. Since the dipole amplitude is about 3~mK, a 0.2\\% error corresponds to about 6~$\\mu$K. The 7-year quadrupole estimated for $\\delta t =0.5$ is \\postrefereechanges{$\\sqrt{(3/\\pi)C_2} \\approx 14 \\mu$K} \\nocite{WMAP7JarosikBasic}(Section 4.1.1, {Jarosik} {et~al.} 2010), only a factor of two higher than this calibration uncertainty (though the model-dependent systematic error is high for the infinite flat model). \\nocite{LL09lowquad}{Liu} \\& {Li} (\\SSS{4}, Fig.~6, 2009) estimate the quadrupole-like difference between their Q1 DA $\\delta t=0$ map and the WMAP collaboration's $\\delta t=0.5$ map to have an r.m.s. of 6.6~$\\mu$K.  Moreover, the spacecraft's velocity is assumed by \\nocite{WMAP5Hinshaw}{Hinshaw} {et~al.} (2009) and \\nocite{WMAP7JarosikBasic}{Jarosik} {et~al.} (2010) to be known exactly, and fits are made over intervals ``typically between 1 and 24 hours''.  It is not obvious that a timing error causing the spacecraft's velocity vector to be slightly misestimated, which in turn causes a quadrupole-like \\protect\\postrefereechanges{dipole-induced difference} signal over a year's observations, of amplitude $\\sim 5$--$10 \\mu$K, would have been detected or removed during the calibration process. \\postrefereechanges{Indeed, \\nocite{KCover09}{Cover} (2009) claims the presence of considerable systematic uncertainty in the calibration of the WMAP TOD.}  The calibration only applies to (differential) radio flux density estimates, not to positional data, so it would not affect \\postrefereechanges{the latter}, i.e. the spacecraft attitude quaternions. The Meta Data Tables containing the quaternions appear to be identical between the ``uncalibrated'' and ``calibrated, filtered'' WMAP 3-year TOD files, as can be expected.  Hence, an incorrect calibration that \\protect\\postrefereechanges{leads to a dipole-induced difference} signal would not cause an arcminute-scale blurring effect nor a positional error in point sources. It would not be detectable by the method used in this paper. Since a large fraction of the brightest point sources at WMAP frequencies are variable, it may not be easy to use point sources to calibrate the uncalibrated data independently of the dipole calibration. \\nocite{LL09lowquad}{Liu} \\& {Li} (2009) and \\nocite{LL10toffset}{Liu} {et~al.} (2010) appear to use the word ``raw'' to refer to the calibrated (filtered for 3-year analyses, unfiltered for 5-year analyses) TOD and do not appear to have carried out the calibration step.  Independently of these considerations, the correct way to compile the {\\em calibrated\\/} WMAP TOD into sky maps must clearly be the one that optimises the sharpness of the images of bright sources. Whether or not the calibration step itself was carried out with a small timing error---\\protect\\postrefereechanges{leading to an artificial dipole-induced difference} signal---remains to be investigated. \\nocite{LL10toffset}{Liu} {et~al.} (2010)'s calculations suggest that the latter is an interesting possibility."
    },
    "1004/1004.3783_arXiv.txt": {
        "abstract": "The light echo systems of historical supernovae in the Milky Way and local group galaxies provide an unprecedented opportunity to reveal the effects of asymmetry on observables, particularly optical spectra. Scattering dust at different locations on the light echo ellipsoid witnesses the supernova from different perspectives and the light consequently scattered towards Earth preserves the shape of line profile variations introduced by asymmetries in the supernova photosphere.  However, the interpretation of supernova light echo spectra to date has not involved a detailed consideration of the effects of outburst duration and geometrical scattering modifications due to finite scattering dust filament dimension, inclination, and image point-spread function and spectrograph slit width.  In this paper, we explore the implications of these factors and present a framework for future resolved supernova light echo spectra interpretation, and test it against Cas~A and SN~1987A light echo spectra.  We conclude that the full modeling of the dimensions and orientation of the scattering dust using the observed light echoes at two or more epochs is critical for the correct interpretation of light echo spectra.  Indeed, without doing so one might falsely conclude that differences exist when none are actually present. ",
        "introduction": "\\label{sec:intro}  The first light echoes (LEs) were discovered around Nova Persei 1901 \\citep{Ritchey01b, Ritchey01a, Ritchey02, Kapteyn02}. Since then, LEs (whereby we mean a simple scattering echo rather than fluorescence or dust re-radiation) have been seen in the Galactic Nova Sagittarii 1936 \\citep{Swope40} and the eruptive variable V838 Monocerotis \\citep[e.g.,][]{Bond03}. LEs have also been observed from extragalactic supernovae (SNe), with SN~1987A \\citep{Kunkel87} being the most famous case \\citep{Crotts88, Suntzeff88, Bond90, Xu95, Sugerman05a, Sugerman05b, Newman06}, but also including SNe~1991T \\citep{Schmidt94, Sparks99}, 1993J \\citep{Sugerman02, Liu03}, 1995E \\citep{Quinn06}, 1998bu \\citep{Cappellaro01}, 2002hh \\citep{Meikle06, Welch07}, 2003gd \\citep{Sugerman05, VanDyk06}, and 2006X \\citep{Wang08}.  The suggestion that historical SNe might be studied by their scattered LEs was first made by \\citet{Zwicky40} and attempted by \\citet{vandenBergh65, vandenBergh66}. The first LEs of centuries-old SNe were discovered in the Large Magellanic Cloud (LMC) \\citep{Rest05b}. They found three LE complexes associated with three small and therefore relatively young SN remnants (SNRs) in the LMC which were subsequently dated as being between 400--1000 years old using the LMC distance and LE feature positions and apparent motions \\citep{Rest05b}. These findings provided the extraordinary opportunity to study the spectrum of the SN light that reached Earth hundreds of years ago and had never been recorded visually or by modern scientific instrumentation. A spectrum of one of the LEs associated with SNR~0509-675 revealed that the reflected LE light came from a high-luminosity SN~Ia \\citep{Rest08a}, similar to SN~1991T \\citep{Filippenko92, Phillips92} and SN~1999aa \\citep{Garavini04} --- the first time that an ancient SNe was classified using its LE spectrum.  Analysis of {\\it ASCA} X-ray data of SNR 0509-675 led \\citet{Hughes95} to suggest that this SNR likely came from a SN~Ia. Recent analysis of its {\\it Chandra} X-ray spectra by \\citet{Badenes08} also supported the classification of this object as a high-luminosity SN~Ia.  The interpretation of LEs, however, is fraught with subtleties. \\citet{Krause05}, for example, identified moving Cas~A features (called ``infrared echoes'' or ``IR echoes'') using mid-IR imaging from the {\\it Spitzer Space Telescope}.  IR echoes are the result of dust absorbing the outburst light, warming and re-radiating at longer wavelengths.  The main scientific conclusion of \\citet{Krause05} was that most of the IR echoes were caused by a series of recent X-ray outbursts from the compact object in the Cas~A SNR based on their apparent motions. However, the analysis was flawed because it did not account for the fact that the apparent motion strongly depends on the inclination of the scattering dust filament \\citep{Dwek08,Rest11_casamotion}. \\citet{Dwek08} also showed that X-ray flares are not energetic enough to be the source of these IR echoes, but that instead the LEs must have been generated by an intense and short burst of EUV-UV radiation associated with the break out of the shock through the surface of the Cas~A progenitor star. Therefore, the most likely source of all LEs associated with Cas~A, both infrared and scattered, is the Cas~A SN explosion itself.  The need to properly analyze the growing collection of LE spectra is what motivated the present work.  The first scattered LEs of Galactic SNe associated with Tycho's SN and the Cas~A SN were discovered by \\citet{Rest07, Rest08b}. Contemporaneously, \\citet{Krause08a} obtained a spectrum of a scattered optical LE spatially coincident with one of the Cas~A IR echoes, and identified the Cas~A SN to be of Type~IIb.  Because \\citet{Rest08b} discovered several LEs with multiple position angles relative to the SNRs, there is a new opportunity for observational SN research --- the ability to measure the spectrum of the same SN from several different directions. At any given instant after the SN light has reached an observer by the direct path, an ellipsoidal surface exists where scattered light from the SN may reach the observer, whose arrival to the observer is delayed by the additional path length. Interstellar dust concentrations must lie at one or more locations on the ellipsoid for the scattered LE to be detectable. Since each LE is at a unique position on the ellipsoid, each LE has a unique line of sight to the SN.  Therefore, the scattered spectral light can provide spectral information on the asymmetry of the SN photosphere. In the limiting case of dust on the ellipsoid opposite the observer, the spectrum of the scattered LE carries the signature of conditions on the SN hemisphere usually hidden from the observer. As the time since the explosion increases, the scattering ellipsoid surface expands and may intersect additional dust concentrations. We present a framework on how to interpret observed LE spectra depending on the scattering dust properties, seeing, and spectrograph slit width and position.  In this paper, we examine the pitfalls of previous analyses, suggest improvements for future studies, and present the application of our methods to data.  In Section~\\ref{sec:case}, we describe how the properties of the scattering dust can cause significant differences to observed LE spectra, and show a specific example of how this has been overlooked in the past.  In Section~\\ref{sec:analysis}, we quantify how observed LE spectra depend on potential differences in dust characteristics (thickness and inclination), as well as observational characteristics (seeing and spectroscopic slit size and orientation). We then apply this framework to observations of Cas~A (Section~\\ref{sec:casa}) and SN~1987A (Section~\\ref{sec:87A}).  In Section~\\ref{sec:discussion}, we discuss how differences in the dust and observing conditions can effect an analysis of LE spectra and suggest future applications of LE observations.  We conclude in Section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In recent years, LEs of SNe have offered the rare opportunity to spectroscopically classify the SN type hundreds of years after the light of the explosion first reached the Earth \\citep[e.g.,][]{Rest08a}.  Recently, we have obtained the first spectroscopic observations of a historical SN from different lines of sight using LEs \\citep{Rest10_casaspec}.  These observations provide the unique opportunity to examine a SN from different directions, with the light for each LE coming from slightly different hemispheres of the SN.  In this paper, we have demonstrated that to properly separate differences in the observed LE spectra caused by intrinsic asymmetries in the SN from those caused by scattering dust, seeing, and spectroscopic slit position, it is necessary to properly model the latter.  Throughout this paper, we have shown how to both determine the properties of the scattering dust and their influence, as well as that of the PSF and spectroscopic slit position, on the observed LE spectra. We found that the dust width and inclination are the dominant factors, especially when the dust filament is thin.  These two factors can be degenerate with respect to the LE profile shape; therefore, it is essential to determine the dust filament inclination independently using the LE proper motion from images taken at two (or more) epochs. The slit width and its misalignment with the major axis of the LE are is also a significant factor.  The image PSF smears out the observed LE and thus also impacts the observed LE spectrum.  Taking all of these factors into account, we found that the observed LE spectrum is not the light-curve weighted integration of the spectra at all epochs, but rather the integration of the spectra weighted by an effective light curve.  This effective light curve is the original light curve convolved with a window function, and it is different for every LE location.  In general, the most significant difference between the effective light curves and the original light curve was that later epochs tend to be truncated in the effective light curves.  We used fits to the LE profiles of Cas~A and SN~1987A to test the consistency and veracity of our model.  We fit the model to the observed LE profiles, where the three free parameters were the LE profile peak height, LE position, and -- physically most important -- the dust filament width.  The model reproduced the observed LE profiles very well, and we found that it was possible to determine all factors influencing the observed LE spectrum {\\it a priori} just from images, without using the observed spectrum.  We constructed window functions and effective light curves for the Cas~A LEs using SN~1993J as a template. We found a much better match of the SN~1993J spectra template to the observed Cas~A LE spectrum if we used the effective light curve instead of the real light curve for weighting.  The most significant difference was in the late-phase emission lines [\\ion{O}{1}] $\\lambda\\lambda 6300$, 6363, [\\ion{Ca}{2}] $\\lambda\\lambda 7291$, 7324, and the \\ion{Ca}{2} NIR triplet, consistent with our expectation from the model. Although we do not have access to the \\citet{Krause08a} images, we find that their Cas~A spectrum is consistent with an integrated spectrum weighted with a light curve truncated at 80~days. Our expectation is that the majority of differences between the Cas~A LE spectrum and the full light-curve weighted SN~1993J spectrum are due to scattering dust properties.  We also constructed a window function and effective light curve for a SN~1987A LE which constitutes a special test case since the light curve and spectral data of the SN are available.  With this data, we constructed a LE spectrum from the original SN and compared it to the observed LE spectrum.  We found that the constructed LE spectrum correctly predicted the line strength of the H$\\alpha$ line, while a simple integration weighted by the light curve was a poor fit.  We presented several additional investigations that can be attempted with LEs in the near future.  Two of these studies, temporally resolved spectra and spatially resolved light curves, require thin dust filaments.  Until now, the assumption of thick scattering dust prevented such studies."
    },
    "1004/1004.3056_arXiv.txt": {
        "abstract": "{ The extraction of a `haze' from the WMAP microwave skymaps is based on subtraction of known foregrounds, {\\em viz.} free-free (bremsstrahlung), thermal dust and synchrotron, each traced by other skymaps. While the 408~MHz all-sky survey is used for the synchrotron template, the WMAP bands are at tens of GHz where the spatial distribution of the radiating cosmic ray electrons ought to be quite different because of the energy-dependence of their diffusion in the Galaxy. The systematic uncertainty this introduces in the residual skymap is comparable to the claimed haze and can, for certain source distributions, have a very similar spectrum and latitudinal profile and even a somewhat similar morphology. Hence caution must be exercised in interpreting the `haze' as a physical signature of, {\\it e.g.,} dark matter annihilation in the Galactic centre.  } ",
        "introduction": "The measurement of anisotropies in the cosmic microwave background (CMB) by COBE and WMAP has ushered in an exciting new era in cosmology. The study of the cosmic signal requires careful subtraction of galactic foreground emissions and this will become even more crucial for studies of the `B-mode' polarisation signal by PLANCK and the proposed CMBPol satellites \\cite{Dunkley:2008am}. It is interesting in this context that the subtraction of all known foregrounds, {\\it i.e.,} free-free (bremsstrahlung), thermal dust and synchrotron (as well as the CMB), from the WMAP skymaps leaves an anomalous emission --- the ``WMAP haze'' \\cite{Finkbeiner:2003im}. This has a roughly spherical morphology localised around the centre of the Galaxy, and a harder spectrum \\cite{Dobler:2007wv} than synchrotron radiation by relativistic cosmic ray (CR) electrons from standard astrophysical sources like supernova remnants (SNRs). An independent analysis has confirmed the existence of the haze \\cite{Bottino:2009uc}, but others do not find the evidence to be significant \\cite{Dickinson:2009yg,Cumberbatch:2009ji}.  It was believed initially that the haze is free-free emission from ionised gas too hot to be traced by recombination line maps but too cold to be visible in X-rays \\cite{Finkbeiner:2003im}. However it was suggested later that it is in fact synchrotron emission from a new population of relativistic electrons,\\footnote{Here and in the following we use ``electrons'' when referring to both electrons and positrons.} produced by dark matter annihilation \\cite{Finkbeiner:2004us}. It is indeed thus possible to explain the haze \\cite{Hooper:2007kb} although other authors argue that the annihilation cross-section needs to be significantly boosted over the usual estimate for thermal relic dark matter \\cite{Cumberbatch:2009ji}. There have also been attempts to fit both the morphology and spectrum of the haze by ascribing it to electrons emitted by pulsars with a hard spectrum \\cite{Kaplinghat:2009ix, Harding:2009ye}; however the expected haze is then less spherical since most pulsars are in the galactic disk. This is also true of SNRs which have in fact recently been invoked \\cite{Blasi:2009hv, Ahlers:2009ae} as sources of positrons with a hard spectrum to explain the rise in the cosmic ray positron fraction at high energies measured by PAMELA \\cite{Adriani:2008zr}.  The presence of an additional population of relativistic electrons in the galactic centre appears to be supported by a recent analysis \\cite{Dobler:2009xz} of the $\\gamma$-ray sky as observed by Fermi-LAT \\cite{Atwood:2009ez}. It is argued that an excess over known components is also present in $\\gamma$-rays, most likely due to inverse-Compton scattering (ICS) by relativistic electrons, and that the underlying electron distribution is compatible with the WMAP haze \\cite{Dobler:2009xz}. While a signature in ICS is naturally expected if there is indeed an additional population of electrons with a hard spectrum, it was pointed out \\cite{Linden:2010ea} that some template maps applied in this analysis \\cite{Dobler:2009xz} are in fact inappropriate and underestimate both the $\\pi^0$ decay and ICS contributions to the $\\gamma$-ray emission, in particular in the galactic centre region. The analysis in Ref.~\\cite{Dobler:2009xz} using the Fermi diffuse model that is believed to be a better tracer of $\\pi^0$ decay and ICS, however, again shows a residual.  The Fermi itself collaboration has not claimed any excess in the galactic centre region over the standard diffuse $\\gamma$-ray background \\cite{Abdo:2010nz,Casandjian:2009wq}.  A crucial ingredient of both studies \\cite{Finkbeiner:2003im,Bottino:2009uc} that identify a microwave haze is the extrapolation of the morphology of the synchrotron radiation template from 408 MHz to the WMAP bands at 23 (K), 33 (Ka), 41 (Q), 61 (V) and 94 (W) GHz, {\\em i.e.} over two orders of magnitude in frequency. In fact the spatial distribution of the radiating CR electrons is likely to differ significantly given their energy dependent diffusive transport in the Galaxy. Instead of attempting such a bold extrapolation, other studies, including the analysis by the WMAP collaboration \\cite{Gold:2010fm}, employ the K-Ka difference map as a tracer of synchrotron emission (despite some contamination by free-free emission and an anomalous component which has been interpreted (see, {\\it e.g.,}\\cite{deOliveiraCosta:2003az}) as spinning dust \\cite{Draine:1998gq}). However although both maps are dominated by synchrotron radiation, such a template could also contain any unidentified radiation, such as a possible haze, and therefore cannot exclude it.  CR transport in the Galaxy is dominated by diffusion through interstellar magnetic fields with an {\\em energy-dependent} diffusion coefficient $D(E) = D_0 E^{\\delta}$ where $\\delta = 0.3 \\mathellipsis 0.7$ \\cite{Strong:2007nh}. Taking the energy loss rate $b (E) = \\mathrm{d}E/ \\mathrm{d}t = b_0 E^2$ as is appropriate for synchrotron and ICS, the diffusion length $\\ell$ is \\begin{equation*} \\ell(E) \\approx 5 \\left(\\frac{E}{\\text{GeV}}\\right)^{(\\delta-1)/2}\\,\\text{kpc}\\,, \\end{equation*} for the standard values $D_0 = 10^{28}\\,\\text{cm}^{2}\\text{s}^{-1}$ and $b_0 = 10^{-16}\\,\\text{s}^{-1}$ \\cite{Strong:2007nh}. Therefore, the distance that GeV energy electrons can diffuse is comparable to the Kpc scale on which the source distribution varies; moreover it changes by a factor of 2.4 (1.5) for $\\delta = 0.3$ (0.7) in the energy range $\\sim4-50$ GeV (corresponding to peak synchrotron frequencies between 408 MHz and 50 GHz for a magnetic field of $6\\,\\mu\\text{G}$).  As a consequence the $\\sim 50 \\, \\text{GeV}$ electrons will trace the source distribution much better than the $\\sim 4 \\, \\text{GeV}$ electrons which diffuse further away from the sources and wash out their distribution.  The synchrotron map at 408 MHz {\\em cannot} therefore be a good tracer of synchrotron radiation at much higher, in particular WMAP, frequencies. Relying on such a crude extrapolation of the morphology of synchrotron emission can thus potentially introduce unphysical residuals. We estimate these by simulating synchrotron skymaps at 408 MHz and the WMAP frequencies and feeding these into the template subtraction process \\cite{Dobler:2007wv}. We show that this leads to residuals of the same order as the claimed haze, which can in fact be matched in spectrum and latitudinal profile for a particular source distribution in the galactic disk. We conclude therefore that the WMAP haze might be an artifact of inappropriate template subtraction rather than evidence of an exotic origin, e.g. dark matter annihilation. ",
        "conclusions": "We have investigated systematic effects in WMAP foreground subtraction stemming from the na\\\"ive extrapolation of the 408 MHz map. To this end we have considered two illustrative cosmic ray diffusion models assuming different source distributions, the first one based on a pulsar survey, and the second one exponential in galacto-centric radius. Both models are able to reproduce the synchrotron radiation at 408 MHz, the locally measured electron flux and are furthermore consistent with nuclear cosmic ray fluxes and secondary-to-primary ratios. When our `foreground' 408 MHz map is subtracted from the 23 GHz map, we find a residual whose size and morphology depends on the source and diffusion model adopted. Thus the energy-dependent diffusion of relativistic electrons makes the 408~MHz skymap a {\\em bad} tracer of synchrotron radiation at microwave frequencies, as had been suspected earlier \\cite{Bennett:2003ca}. Such a template subtraction produces a residual of the same overall intensity as the haze and can for particular source distributions give the same latitudinal profile.  For the Lorimer source distribution, the residual is of \\emph{opposite} sign to the ``haze'' and can therefore certainly not explain the ``haze'' as a residual of the template subtraction. However since it is of comparable magnitude and its morphology is strikingly similar, it is important to keep this issue in mind when interpreting the ``haze'' as an excess over standard synchrotron emission from electrons injected by SNRs. We emphasise that the significant uncertainty thus introduced has a considerable effect on the parameter space available for possible explanations of the ``haze\", {\\it e.g.,} dark matter annihilation or pulsars.  The residual obtained from the exponential source distribution does not perfectly reproduce the morphology found in Ref. \\cite{Dobler:2007wv} (although it is {\\em not} disk-like but rather clustered around the galactic centre). However, a quantitative assessment of the discrepancy is not straightforward, mainly because Ref. \\cite{Dobler:2007wv} does not provide any objective measure, {\\it e.g.,} the ellipticity of equal intensity contours. On the other hand, even the numerical {\\tt GALPROP} model we employed for our analysis is very likely too simple to fully capture the complexity of synchrotron emission in the Galaxy. For instances, not only the source density but also the galactic magnetic field is supposed to be correlated with the galactic spiral arms, which will break the symmetry in $r$ (and hence in $\\ell$) and can therefore considerably modify the morphology. Furthermore, much of the `diffuse' synchrotron emission from the disk may originate in the shells of old supernova remnants which have grown very large in their radiative phase \\cite{Sarkar:1980}. Exactly the same argument concerning the energy-dependent diffusion length that we applied to the cosmic ray source distribution can be applied to such localised structures too. Therefore the 408 MHz survey skymap is not expected to trace the emission from the latter at higher frequencies either. One can easily imagine that such localised structures (of which Loop I is a nearby example) might at least in part modify the morphology of the residual and bring the simulated map into agreement with the one determined from the subtraction of real templates.  \\subsection*{Note added}  As we were about to submit this manuscript, a related study appeared \\cite{McQuinn:2010ju}. Although we agree on the importance of diffusion-loss steepened electron spectra for producing the haze there is a major difference between our approaches --- while the authors of Ref. \\cite{McQuinn:2010ju} consider the haze to be {\\em physical}, we argue that it might in fact be an artifact of the foreground subtraction. Our models are also more constrained insofar as we reproduce the observed radio emission at 408 MHz and match the direct measurements of the electron spectrum at our position. Furthermore, we allow for spatial dependence of the $\\vec{B}$ field, and convection and reacceleration of cosmic ray electrons, which are all essential in order explain all these datasets simultaneously.   \\subsection{Acknowledgements}  PM acknowledges support by the EU Marie Curie Network ``UniverseNet'' (HPRN-CT-2006-035863) and a STFC Postgraduate Studentship.   \\appendix"
    },
    "1004/1004.2050_arXiv.txt": {
        "abstract": "We investigate the effect of three important processes by which AGN-blown bubbles transport material: drift, wake transport and entrainment. The first of these, drift, occurs because a buoyant bubble pushes aside the adjacent material, giving rise to a net upward displacement of the fluid behind the bubble. For a spherical bubble, the mass of upwardly displaced material is roughly equal to half the mass displaced by the bubble, and should be $\\sim 10^{7-9}\\,{\\rm M_\\odot}$ depending on the local ICM and bubble parameters. We show that in classical cool core clusters, the upward displacement by drift may be a key process in explaining the presence of filaments behind bubbles. A bubble also carries a parcel of material in a region at its rear, known as the wake. The mass of the wake is comparable to the drift mass and increases the average density of the bubble, trapping it closer to the cluster centre and reducing the amount of heating it can do during its ascent. Moreover, material dropping out of the wake will also contribute to the trailing filaments. Mass transport by the bubble wake can effectively prevent the build-up of cool material in the central galaxy, even if AGN heating does not balance ICM cooling. Finally, we consider entrainment, the process by which ambient material is incorporated into the bubble. Studies of observed bubbles show that they subtend an opening angle much larger than predicted by simple adiabatic expansion. We show that bubbles that entrain ambient material as they rise, {\\em will} expand faster than the adiabatic prediction; however, the entrainment rate required to explain the observed opening angle is large enough that the density contrast between the bubble and its surroundings would disappear rapidly. We therefore conclude that entrainment is unlikely to be a dominant mass transport process. Additionally, this also suggests that the bubble surface is much more stable against instabilities that promote entrainment than expect for pure hydrodynamic bubbles. ",
        "introduction": "The hot, gaseous atmospheres of galaxy clusters often show depressions in the X-ray surface brightness \\citep[see, for example][]{mc02,birzan,dunn05,rafferty06}. These depressions are indicative of empty cavities, or bubbles, embedded in the hot gas \\citep[e.g.][]{mcnuls}. The presence of bubbles is generally taken to be a signature of AGN feedback. In this model, a fraction of the material cooling from the gaseous atmosphere is accreted by a supermassive black hole located in the central galaxy. This releases vast amounts of energy often in the form of outflows which couple to the hot atmosphere \\citep[e.g.][]{calori,babul}.  AGN feedback is widely considered to have important consequences for the evolution of single galaxies, as well as galaxy groups and clusters. For example, feedback is thought to be key in determining the upper mass cut-off of the galaxy mass function \\citep[][]{benson, croton05, bower06}, balancing the radiative losses in cool-core galaxy clusters, as well as providing the non-gravitational `pre-heating' that may be important for non cool-core clusters \\citep[e.g.][]{bab,mac04,mc08}. Feedback also seems to play a role in determining the relationship between the temperature and X-ray luminosity of hot atmospheres across a range of halo masses \\citep[][]{bab,mac04,puchwein, bower08, dave, pope09}. As a direct consequence, feedback also regulates supermassive black hole growth \\citep[e.g.][] {sr,chur06} and, therefore, its relation to the properties of the host galaxy.  Considerable effort has been invested in understanding the impact of supermassive black holes. One of the key challenges is understanding how the energy from the AGN affects and couples to the broader environment. Most studies that have sought to describe this relationship have focused on bulk motion, shock waves and pressure/gravity waves induced by AGN outbursts and the subsequent dissipation of the associated energy \\citep[see][for a review]{mcnuls}. In this article we focus on the transport of material out of the cluster centre by AGN-blown bubbles. As an example, the filaments of cool material observed behind AGN-blown bubbles \\citep[e.g.][]{consel,crawf05,hatch} are a strong indication of bubble-induced mass transport within several 10s of kiloparsecs of the cluster centre. The most obvious example is the Perseus cluster in which the filaments contain some $ 10^{8}$ solar masses of cool material \\citep[e.g.][]{salome,salome08}.  More generally, some of the transport mechanisms we discuss in this paper may carry matter out to $\\sim$100 kiloparsecs from the cluster centre.  Mass transport by bubbles not only reduces the amount of material available for forming new stars in the central galaxy, allowing the normally adopted stringent requirement that AGN heating perfectly balance cooling to be some relaxed, but also affects the bubble dynamics and energetics. Hence, a better understanding of the main mass transport mechanisms is essential for describing the energy balance in the ICM. Here, we focus on three important processes by which bubbles can transport material: drift, wake transport and the entrainment of ambient material into the bubble.  The aim of this article is to quantitatively describe these three main mechanisms and also the corresponding implications.  The discussion is largely analytical, intended to facilitate a better understanding of the underappreciated aspects of these mechanisms and aid in the interpretation of both observations and numerical simulations.  The article is arranged in the following way.  Section 2 outlines the basic bubble model that serves as a backdrop for subsequent discussions.  Sections 3, 4 and 5 focus on each of the three mechanisms --- drift, wake transport and entrainment.  In each section, we discuss a process as well as associated, potentially observable consequences, like trailing optical filaments.  In Section 6, we consider how mass transport modifies the ongoing debate about whether or not AGN heating need balance cooling precisely.  We summarize our key findings in Section 7. ",
        "conclusions": "Having examined the three main mass transport processes associated with a buoyant bubble individually, we now consider the three processes jointly and ask:\\hfill\\break (1) how much material in total can actually be transported per bubble?\\hfill\\break (2) how frequently must bubbles be inflated in order for the mass outflow to balance the cooling inflow?  Regarding the first question, we have already established that the bubbles cannot be carrying very much entrained material.  The drift can carry upward an amount of ambient material comparable to 50\\% of the mass initially displaced by the bubble, or approximately $\\sim 10^{8}\\,{\\rm M_\\odot}$.  However, most of this material will not be displaced upwards by more than a bubble diameter and will likely fall back towards the cluster centre.  In addition, for clusters with a flat entropy core, the overall fluid displacement might be negligible. This then means that the drift, though likely responsible for observable features like the cool optical filaments, is not an important mechanism for affecting mass outflow.  This leaves the wake.  As noted previously, fluidised bed experiments \\citep[e.g.][]{yang} show that wake typically contains a approximately a quarter of the mass initially displaced by the bubble.  Assuming that laboratory value of $q$ can be applied to the ICM, AGN-blown bubble wakes could easily transport $q M_{\\rm dis,0} \\sim 10^{8}\\,{\\rm M_\\odot}$, see equation (\\ref{eq:disp}).  Moreover, the mass in the wake can be carried out to large distances. As indicated by equation (\\ref{eq:6a}), the maximum bubble height is a strong function of $q$. For $q \\sim 0.24$ (and $\\eta_{0} = 100$), depending on the temperature profile, the bubble can rise to at least 5$z_{\\rm 0}$, where $z_{0}$ is the height at which the bubble was inflated.  For $z_{0}\\sim 10$ kpc, the bubble will rise to 50 kpc.  Let us consider an AGN that inflates a sequence of bubbles.  We can write the time-averaged mass outflow rate from the cluster centre due to a train of bubbles as \\begin{equation}\\label{eq:beg} \\dot{M}_{\\rm out}\\approx qM_{\\rm dis,0}/\\tau, \\end{equation} where $\\tau$ is the average time between bubble inflation.  Since this outflow counteracts the inflow due to cooling, cumulatively bubble mass transport may allow for the relaxing of the stringent (but commonly adopted) requirement that AGN heating must balance radiative cooling.  In fact, there is some debate as to whether AGN heating in fact balances cooling, especially in the most luminous clusters. Estimates of the time-averaged heating rate by \\cite{best08}, for example, indicates that AGN heating {\\em does not} balance cooling in the most luminous clusters.\\footnote{We note that \\cite{dunn08} claim a close balance between heating and cooling but in their comparison, they refer to the heating rate due to an individual bubble. Averaging over the time between successive bubbles will lower the power estimate, leading to a discrepancy between heating and cooling in agreement with \\cite{best08}.}  As we show below, mass outflow can play an important role in limiting the mass build-up at the cluster centre in such instances.  To illustrate this, let us consider the following: in the absence of heating, the mass inflow rate due to cooling is related to the radiative losses of the ICM according to \\begin{equation} \\dot{M}_{\\rm cool,0} =  \\frac{(\\Gamma_{\\rm ICM}-1)}{\\Gamma_{\\rm ICM}} \\frac{\\mu m_{\\rm p}}{k_{\\rm B}T} L_{\\rm X}, \\end{equation} where $T$ is the average temperature of the ICM.  Studies of the X-ray spectra of the ICM in cluster cores indicate that the actual rate at which the gas is cooling is $\\dot{M}_{\\rm cool}\\simlt 0.2 \\dot{M}_{\\rm cool,0}$ \\cite[e.g.][]{peterson01,kaas01,peterson03}. If the mass inflow rate is balanced by the mass transport rate by the bubbles: $\\dot{M}_{\\rm cool}\\approx \\dot{M}_{\\rm out}$.  The actual gas cooling rate, $\\dot{M}_{\\rm cool}$, is the result of a difference between cooling and heating: \\begin{equation}\\label{eq:end} \\dot{M}_{\\rm cool} = \\frac{(\\Gamma_{\\rm ICM}-1)}{\\Gamma_{\\rm ICM}} \\frac{\\mu m_{\\rm p}}{k_{\\rm B}T} \\left(L_{\\rm X}-\\dot{E}_{\\rm heat}\\right), \\end{equation} where $\\dot{E}_{\\rm heat} = \\alpha_{\\rm heat} E_{\\rm bub,0}/\\tau$. Here, $E_{\\rm bub,0}$ is the enthalpy of an AGN-inflated bubble (see equation \\ref{eq:1}) and $\\alpha_{\\rm heat}$ allows for the possibility that the bubble enthalpy may under-represent the total energy injected by the AGN \\citep[e.g.][]{nusser,binney07}, especially if the bubbles are inflated supersonically, which would engender shocks that enhance the heating.  Combining equations (\\ref{eq:beg})---(\\ref{eq:end}) and taking $\\Gamma_{\\rm ICM} = 5/3$, $\\Gamma_{\\rm b} = 4/3$ and $T/T_0 \\sim 3$, we find that the build-up of excessive amounts of cold gas in the central galaxy can be prevented if \\begin{equation} \\dot{E}_{\\rm heat} \\approx \\frac{\\alpha_{\\rm heat} }{\\left[ 2q+\\alpha_{\\rm heat}\\right]}L_{\\rm X} \\end{equation} and \\begin{equation} \\dot{M}_{\\rm out} \\approx \\frac{2q }{\\left[2q+\\alpha_{\\rm heat}\\right]}\\dot{M}_{\\rm cool,0}. \\end{equation} If the AGN does no heating (i.e.~$\\alpha_{\\rm heat}=0$), mass transport by the wake can, by itself, prevent rate of cold gas in the central cluster galaxy as long as the bubble recurrence timescale is \\begin{equation}% \\tau \\approx 1 \\times 10^{7} \\bigg(\\frac{E_{\\rm bub,0}}{10^{59}\\,{\\rm erg}}\\bigg)\\bigg(\\frac{L_{\\rm X}}{10^{44}\\,{\\rm erg\\,s^{-1}}}\\bigg)^{-1}\\,{\\rm yrs}, \\end{equation} where $10^{59}\\,{\\rm erg}$ is the typical enthalpy of bubbles in a cluster with luminosity $L_{\\rm X} \\sim 10^{44}\\,{\\rm erg s^{-1}}$ \\citep[e.g.][]{birzan,dunn05} and we have taken $q\\approx 0.24$.  Of course, the corresponding $\\dot{M}_{\\rm cool}$ will exceed the limits derived from the X-ray spectra.  If, in addition to mass transport, the bubbles also heat the ICM such that the amount of energy deposited in the ICM is twice the enthalpy of the bubble (i.e.~$\\alpha_{\\rm heat}=2$), the rate of cold gas build-up can be prevented if the bubbles are inflated on a timescale \\begin{equation}\\label{eq:tcritx} \\tau \\approx 5 \\times 10^{7} \\bigg(\\frac{E_{\\rm bub,0}}{10^{59}\\,{\\rm erg}}\\bigg)\\bigg(\\frac{L_{\\rm X}}{10^{44}\\,{\\rm erg\\,s^{-1}}}\\bigg)^{-1}\\,{\\rm yrs}. \\end{equation} It is worth noting that in this particular case, the time-averaged heating rate does not balance the cooling rate, $\\dot{E}_{\\rm heat}\\approx 0.8 L_{\\rm X}$, and the mass cooling rate is the maximum allowed, $\\dot{M}_{\\rm cool} \\approx 0.2\\dot{M}_{\\rm cool,0}$, but mass build-up in the central cluster galaxy is successfully avoided by mass transport due to the bubbles.  The values of $\\tau$ obtained from equation (\\ref{eq:tcritx}) are reasonably consistent with indirect observational estimates, though there are large uncertainties. For example, \\cite{best08} found radio AGN duty cycles of $\\sim 30\\%$ in Brightest Cluster Galaxies (BCGs), which corresponds to an  typical `observed' time between outbursts of $\\tau_{\\rm obs} \\sim 3 t_{\\rm on}$, where $t_{\\rm on}$ is the duration of an AGN outburst. Therefore, since $t_{\\rm on}$ seems to vary between $10^{7}-10^{8}\\,{\\rm yrs}$ \\citep[e.g.][]{birzan}, it seems possible that AGN-blown bubbles can significantly reduce the rate at which cold gas collects in the galaxy.  Finally, we note that throughout this work we have only considered bubbles that remain intact. Bubbles which deform into vortex rings behave somewhat differently. In addition, the fluid circulation associated with the vortex ring exerts a force that can lift up additional material behind the bubble, \\citep[see][]{pav}. However, with the possible exception of M87 \\citep[][]{m87rad} structures resembling vortex rings have not been observed, so this additional mechanism of mass transport might be unimportant except in rare cases."
    },
    "1004/1004.2861_arXiv.txt": {
        "abstract": "{Mid-infrared (mid-IR) emission lines of molecular hydrogen (H$_2$) are useful probes to determine the mass of warm gas present in the surface layers of circumstellar disks. In the past years, numerous observations of Herbig Ae/Be stars (HAeBes) have been performed, but only two detections of H$_2$ mid-IR emission toward HD~97048 and AB~Aur have been reported.} { We aim at tracing the warm gas in the circumstellar environment of five additional HAeBes with gas-rich environments and/or physical characteristics close to those of AB~Aur and/or HD~97048, to discuss whether the detections toward these two objects are suggestive of peculiar conditions for the observed gas.} { We search for the H$_2$ S(1) emission line at 17.035 $\\mu$m using high-resolution mid-IR spectra obtained with VLT/\\visir, and complemented by CH molecule observations with VLT/\\uves. We gather the H$_2$ measurements from the literature to put the new results in context and search for a correlation with some disk properties.} { None of the five \\visir\\ targets shows evidence for H$_2$ emission at 17.035 $\\mu$m. From the 3$\\sigma$ upper limits on the integrated line fluxes we constrain the amount of optically thin warm ($>150~K$) gas to be less than $\\sim 1.4~ M_{\\rm Jup}$ in the disk surface layers. There are now 20 HAeBes observed with \\visir\\ and \\texes\\ instruments to search for warm H$_2$, but only two detections (HD~97048 and AB~Aur) were made so far. We find that the two stars with detected warm H$_2$ show at the same time high 30/13\\,$\\mu$m flux ratios and large PAH line fluxes at 8.6 and 11.3$\\,\\mu$m compared to the bulk of observed HAeBes and have emission CO lines detected at $4.7\\,\\mu$m.  We detect the CH 4300.3\\,$\\AA$ absorption line toward both HD~97048 and AB~Aur with \\uves. The CH to H$_2$ abundance ratios that this would imply if it were to arise from the same component as well as the radial velocity of the CH lines both suggest that CH arises from a surrounding envelope, while the detected H$_2$ would reside in the disk.} {The two detections of the S(1) line in the disks of HD~97048 and AB~Aur suggest either peculiar physical conditions or a particular stage of evolution. New instruments such as \\textit{Herschel / PACS} should bring significant new data for the constraints of thermodynamics in young disks by observing the gas and the dust simultaneously.} ",
        "introduction": "Planets are supposed to form in circumstellar disks composed of gas and dust around stars in their pre-main sequence phase. At this evolutionary stage, the disk mass is essentially dominated by the gas (99\\%), especially molecular hydrogen (H$_2$). Although H$_2$ is the principal gaseous constituent in disks, it is very challenging to detect. Molecular hydrogen is an homonuclear molecule which means that its fundamental transitions are quadrupolar in nature. Hence their Einstein spontaneous emission coefficients are small and produce only weak lines. For circumstellar disks another challenge is that the weak H$_2$ lines must be detected on top of the strong dust continuum emission. In the mid-infrared, high spectral resolution instruments like \\visir\\ at the VLT are required to disentangle these weak lines from the continuum.  However, one should bear in mind a fundamental issue concerning the structure of a gas-rich optically thick disk. Molecular lines are produced in the hot upper surfaces of the disk where the molecular gas and accompanying dust are optically thin. Therefore, molecular line emission is not sensitive to and does not probe the mid-plane interior layers of the disk because it is optically thick. Because the amount of molecular gas in the optically thin surface layers is small, the expected H$_2$ line fluxes are very weak. As an example of this, \\cite{Carmona08} calculated the expected H$_2$ line fluxes from typical Herbig Ae disks assuming a two-layer \\cite{Chiang97} disk model, $T_{\\rm gas}=T_{\\rm dust}$, a gas-to-dust-ratio of 100, LTE emission and a distance of 140 pc. They found that the expected line H$_2$ fluxes are much fainter than the detection limits of current instrumentation. Indeed, numerous non-detections of H$_2$ mid-IR pure rotational lines in the circumstellar environment of young stars have been reported in the past few years \\citep{Bitner08, Carmona08, klr08b, klr09a}. Nevertheless, lines can reach a detectable level if the gas-to-dust to ratio is allowed to be higher than 100 and/or the $T_{\\rm gas}>T_{\\rm dust}$ in the surface layers of the disk.        \\begin{table*} \\begin{center} \\caption{Astrophysical parameters of the sample stars.} \\begin{tabular}{lcccccccccccccccccc} \\hline \\hline Star        & Sp.    & $T_{eff}$     & \\Av         &  \\vrad $^{(a)}$  & $d$        & Disk         & Evidence for  \\\\ & Type   & (K)          & (mag)        & (\\kms)       & (pc)       & resolved     &  gas-rich \\\\ &         &              &              &              &            &              &  environment \\\\ \\hline HD~142527   & F6~IIIe & 6300$^{(1)}$  & 1.49$^{(1)}$ & -3.5$^{(2)}$  & 140$^{(3)}$ & yes$^{(4)}$   &   \\\\ HD~169142   & A8~Ve   & 8130$^{(5)}$  & 0.37$^{(6)}$ & -3.0$^{(7)}$  & 145$^{(5)}$ & yes$^{(8)}$   & yes$^{(9)}$   \\\\ HD~150193A  & A1~Ve   & 9300$^{(1)}$  & 1.61$^{(1)}$ & -6.0$^{(10)}$  & 150$^{(1)}$ & yes$^{(11)}$  &    \\\\ HD~163296   & A1~Ve   & 9300$^{(1)}$  & 0.25$^{(1)}$ & -4.0$^{(2)}$  & 122$^{(1)}$ & yes$^{(12)}$  & yes$^{(13)}$    \\\\ HD~100546   & B9~Vne  & 10470$^{(1)}$ & 0.25$^{(1)}$ & +17$^{(13)}$   & 103$^{(1)}$ & yes$^{(14)}$ & yes$^{(13)}$    \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item $^{(a)}$ radial velocity of the star in the heliocentric rest frame.  \\item References: $^{(1)}$ \\cite{VdA98b}; $^{(2)}$ SIMBAD database; $^{(3)}$ \\cite{DeZeeuw_99}; $^{(4)}$ \\cite{Fukagawa06}; $^{(5)}$ \\cite{Acke05}; $^{(6)}$ \\cite{Malfait98}; $^{(7)}$ \\cite{Dunkin97}; $^{(8)}$\\cite{Kuhn01}; $^{(9)}$ \\cite{Panic08}; $^{(10)}$ \\cite{Reipurth96}; $^{(11)}$ \\cite{Fukagawa03}; $^{(12)}$ \\cite{GRADY00}; $^{(13)}$ \\cite{LECAV03}; $^{(14)}$ \\cite{Augereau01}. \\\\ \\end{list} \\label{param} \\end{center} \\end{table*}     \\begin{table*} \\begin{center} \\caption{Summary of the observations.} \\begin{tabular}{lcccccccccccccccccc} \\hline \\hline Star      &  $t_{exp}$ & Airmass   & Optical   & Standard  & Airmass   & Optical   & Asteroid  & Airmass   & Optical  \\\\ &  (s)      &           & Seeing    &  Star     &           & Seeing    &           &           & Seeing     \\\\ &           &           & ('')      &           &           &  ('')     &           &           &  ('')    \\\\ \\hline HD~142527   & 3600      & 1.10-1.28 & 0.71-1.08 & HD~211416 & 1.23-1.25 & 0.79-1.02 & HEBE      & 1.19-1.20 & 0.97-1.06 \\\\ HD~169142   & 3600      & 1.17-1.50 & 0.71-1.01 & HD~211416 & 1.23-1.25 & 0.79-1.02 & HEBE      & 1.19-1.20 & 0.97-1.06 \\\\ HD~150193A  & 3600      & 1.31-1.88 & 0.85-1.40 & HD~211416 & 1.24-1.25 & 0.81-0.99 & PSYCHE    & 1.13-1.16 & 0.75-0.83 \\\\ HD~163296   & 3600      & 1.60-2.70 & 0.79-1.19 & HD~211416 & 1.24-1.25 & 0.81-0.99 & IRIS      & 1.00-1.01 & 0.99-1.28 \\\\ HD~100546   & 3600      & 1.44-1.52 & 0.73-1.09 & HD~89388  & 1.31-1.32 & 0.61-0.73 & HEBE      & 1.19-1.20 & 0.97-1.06 \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item  The airmass and seeing intervals are given from the beginning to the end of the observations. \\\\ \\end{list}  \\label{log_obs} \\end{center} \\end{table*}      Molecular hydrogen mid-IR lines have been detected in two Herbig Ae/Be stars, namely HD~97048 and AB Aur, respectively observed with VLT/\\visir\\ \\citep{klr07} and \\texes\\ \\citep{Bitner07}. These detections imply particular physical conditions for the gas and dust, such as $T_{\\rm gas}>T_{\\rm dust}$, as mentioned above. These conditions may be created by gas heated by X-rays or UV photons in the surface layers of the disks. We note that these two detections indicate that the gas has not completely dissipated in the inner part of these disks in a lifetime of about 3 Myrs (ages of the stars), while photoevaporation of the gas is expected to clear up this inner region within roughly the same time \\citep[$<$3 Myrs; e.g.][]{Takeuchi05, Alexander08}. Indeed, \\cite{klr09a} have shown that the emitting H$_2$ around HD~97048 was more likely distributed in an extended region within the inner disk, between 5~AU and 35~AU of the disk, and \\cite{Bitner07} concluded from their observation of the disk of AB~Aur that the H$_2$ emission arised around 18~AU from the central star.  We present new observations with the \\visir\\ high-resolution spectroscopic mode, to search for the mid-IR H$_2$ emission line at 17.035 $\\mu$m, the most intense pure rotational line observable from the ground, in a sample of well studied Herbig Ae/Be stars known to harbor extended gas-rich circumstellar disks. The main goal of these observations is to enlarge the global sample of HAeBes observed at 17.035 $\\mu$m, and to better constrain the particular physical conditions observed in the disks of HD~97048 and AB~Aur. In order to better understand the detections of H$_2$ in the disks of HD~97048 and AB~Aur, we also analyze VLT/\\uves\\ spectra of these two stars to observe spectral lines of the CH and CH$^+$ molecules which are linked to the formation and excitation of H$_2$. In Sects.~\\ref{results} and \\ref{discuss} we present the analysis of the \\visir\\ spectra of the five HAeBes of our sample and relate it to the context of the global search for molecular hydrogen at mid-IR wavelengths. We perfom a statistical analysis of the whole sample of HAeBes where H$_2$ emission has been searched for in the mid-IR, and explore the possible link between the H$_2$ and the dusty disk properties. Finally, we discuss the possible origin of the mid-IR H$_2$ emission in the disks of HD~97048 and AB~Aur. ",
        "conclusions": "We reported here on a search for the H$_2$ S(1) emission line at 17.035 $\\mu$m in the circumstellar environments of five well known HAeBes with the high resolution spectroscopic mode of \\visir. No source shows evidence for H$_2$ emission at 17.035 $\\mu$m. From the 3$\\sigma$ upper limits on the integrated line fluxes, we derived limits on column densities and masses of warm gas as a function of the temperature. The present work brings to 18 the number of non-detections of the H$_2$ S(1) line in a global sample of 20 Herbig Ae/Be stars observed with \\visir\\ and \\texes\\ \\citep{klr07, Bitner07, Carmona08, klr08b, Bitner08}.  The detections of H$_2$ emission at 17.035 $\\mu$m by \\cite{Bitner07} and \\cite{klr07} show that at least a few circumstellar disks have sufficiently high H$_2$ mid-infrared emission to be observed from the ground. The most likely explanation for this is that the optically thin surface layers of the disk has $T_{gas} > T_{dust}$ and that the gas-to-dust ratio is higher than the canonical ratio of 100 \\citep{Carmona08, klr07}.  Indeed, in the surface layers of the disk, low densities or dust settling and coagulation may conduct to a spatial decoupling between the gas and the dust. Photoelectric heating can thus play a significant role in the gas heating process and the physical conditions may rapidly differ from the LTE ones, i.e., $T_{gas} > T_{dust}$ \\citep{Kamp04, Jonkheid07}. On the other hand, UV and X-ray heating can be responsible for the excitation of the observed gas and can heat the gas to temperatures significantly hotter than the dust \\citep{Nomura05, Glassgold07, Ercolano08}. In any case, one would need to observe the lower H$_2$ $J$-levels (i.e. $J=0$ and $J=1$) to definitively constrain the kinetic temperature of the gas, and better understand the excitation mechanisms responsible of the mid-IR emission.  As a second step, we also performed a statistical analysis of the whole sample of Herbig Ae/Be stars observed at 17.035 $\\mu$m with \\visir\\ as well as with \\texes. This analysis allowed us to identify a population of stars, including HD~97048 and marginally AB~Aur, with properties that depart from the bulk of our sample. This raises numerous questions about the origin of the detected gas and the status of HD~97048 and AB~Aur. From our \\uves\\ observations we clearly demonstrated that the observed mid-IR H$_2$ emission does not come from the envelope, but from the disk.  Are HD~97048 and AB~Aur peculiar stars (and why)? Due to the similarities of these two stars ($T_{eff}$, age, mass, disk size...), one would expect that the physical conditions of their circumstellar gas are typical of a particular (short) stage of evolution of the disks. However, to confirm this assumption, we would need to observe other similar HAeBes. Unfortunately, no other nearby Herbig Ae/Be star observable with the existing instruments presents the same observational properties. A global diagnostic of the gaseous content of the disks is thus now required. To better constrain the physical and chemical properties of the gas, multi-wavelengths observations and a deep modeling would be very useful. In this context, the {\\it Herschel} satellite will be very helpful because it will allow us to constrain the thermodynamics in young disks by observing the gas and the dust simultaneously."
    },
    "1004/1004.5393_arXiv.txt": {
        "abstract": "We investigate scaling relations of bulges using bulge-disk decompositions at 3.6~$\\mu$m and present bulge classifications for 173 E-Sd galaxies within 20~Mpc. Pseudobulges and classical bulges are identified using S\\'ersic index, HST morphology, and star formation activity (traced by 8~$\\mu$ emission). In the near-IR pseudobulges have $n_b<2$ and classical bulges have $n_b>2$, as found in the optical. S\\'ersic index and morphology are essentially equivalent properties for bulge classification purposes.  We confirm, using a much more robust sample, that the S\\'ersic index of pseudobulges is uncorrelated with other bulge structural properties, unlike for classical bulges and elliptical galaxies. Also, the half-light radius of pseudobulges is not correlated with any other bulge property.  We also find a new correlation between surface brightness and pseudobulge luminosity; pseudobulges become more luminous as they become more dense. Classical bulges follow the well known scaling relations between surface brightness, luminosity and half-light radius that are established by elliptical galaxies.  We show that those pseudobulges (as indicated by S\\'ersic index and nuclear morphology) that have low specific star formation rates are very similar to models of galaxies in which both a pseudobulge and classical bulge exist. Therefore, pseudobulge identification that relies only on structural indicators is incomplete.  Our results, especially those on scaling relations, imply that pseudobulges are very different types of objects than elliptical galaxies. ",
        "introduction": "\\label{sec:intro}  A preponderance of observational evidence suggests that there are at least two types of bulges. Reviews of evidence for a dichotomy of bulges can be found in \\cite{kk04}, \\cite{athan05}, \\cite{kormendyfisher2005}, and for a review of more recent literature see \\cite{combes2009}. The dichotomy in bulge properties can be summarized as follows: many bulges have properties as described by \\cite{renzini99} of ``little elliptical galaxies surrounded by a disk'', while other bulges are similar in many properties to disk galaxies. We call those bulges that are similar to E-type galaxies ``classical bulges'', and those that are more like disk-galaxies are referred to as ``pseudobulges.''   Many bulges in the nearby Universe are filled with young stars \\citep{peletier1996, gadotti2001,carollo2002,macarthur2003}, and many bulges are gas rich \\citep{regan2001bima, helfer2003, sheth2005,jogee2005}. Also, \\cite{peletier1996} find that the ages and stellar populations of bulges vary greatly from galaxy-to-galaxy; however, the stellar population of the outer disk is an excellent predictor of that of the bulge. \\cite{fisher2006} shows that bulge morphology is also a good predictor of star formation in the bulge. He finds that pseudobulges are actively forming stars, and have ISM properties that are like that of their outer disks. \\cite{fdf2009} study star formation in bulges and their outer disks using Spitzer and GALEX data.  They show that pseudobulges are presently increasing the bulge-to-total ratio of stellar light (here after $B/T$) via internal star formation. If pseudobulges have similar star formation histories as their outer disks, then the mean historic star formation rate (over the past $\\sim$10~Gyr) ought to be 1-3 times that of present day star formation rate \\citep{kennicutt94}.  The results of \\cite{fdf2009} thus indicate that present day star formation rates are high enough to account for the entire stellar mass of most pseudobulges.  Also, \\cite{fdf2009} find positive correlations between bulge mass and star formation rate density suggesting that long term {\\em in situ} growth may have formed pseudobulges. Furthermore, those pseudobulges with highest star formation rate density are exclusively found in the most massive disks, suggesting that pseudobulge growth is connected to outer disk properties.  They also find that classical bulges, in contrast, are not forming significant amounts of stars today. Furthermore, data from \\cite{regan2001bima,helfer2003} show that classical bulges are gas poor when compared to their outer disk.  The shape of stellar density profiles is also thought to participate in the bulge dichotomy.  \\cite{andredak95} and, more recently, \\citealp{scarlata2004} show that the distribution of surface brightness profile shapes of bulges is bimodal.  Also, \\cite{carollo1999} finds that many of the the 'amorphous nuclei' of intermediate-type disk galaxies are better fit by exponential profiles, rather than the traditional $r^{1/4}$ profile. This leads \\cite{kk04}, among others, to suggest that surface brightness is tied to the pseudobulge -- classical bulge dichotomy. \\cite{fisherdrory2008} directly test this hypothesis. They identify pseudobulges with bulge morphology.  Using $V$-band bulge-disk decompositions they find that that 90\\% of pseudobulges have $n<2$ and all classical bulges have $n>2$.  \\cite{fisherdrory2008} also find that the S\\'ersic index of pseudobulges does not correlate with bulge luminosity, half-light radius or bulge-to-total ratio as it does for classical bulges and elliptical galaxies.  It is well known that elliptical galaxies follow a ``fundamental plane'' in parameter space that relates size, surface density, and velocity dispersion \\citep[e.g.][]{djorgovski1987,dressler1987,faber1989}.  To lowest order, these correlations are a consequence of the virial theorem; small deviations from virial predictions in slopes of these correlations represent variation in mass-to-light ratios and non-homologous density profiles. Disks correlate differently than ellipticals in fundamental plane parameter space \\citep{bbf92,kfcb}. Yet, the location of pseudobulges compared to that of classical bulges, elliptical galaxies, and disks in structural parameter correlations is poorly understood. Thus, knowing where pseudobulges lie with respect to fundamental plane correlations would help in interpreting their structural properties.  \\cite{graham2001},\\cite{macarthur2003}, and \\cite{dejong2004} each investigate the scaling relations of bulges. All find parameter correlations with high scatter.  The \\cite{sersic1968} function is highly degenerate \\citep{graham1997}, and the dynamic range available to fit bulges is limited. Indeed, \\cite{balcells2003} and \\cite{gadotti2008} show that high spatial resolution is necessary to accurately determine fit parameters when using the S\\'ersic function. Accordingly, \\cite{balcells2007} and \\cite{fisherdrory2008} use composite HST and ground-based data to calculate bulge-disk decompositions. The central isophotes are calculated with HST data, providing adequate spatial resolution, and outer isophotes are calculated with ground based data. However, this method is more labor intensive, and the the resulting data set is smaller.  They find a high probability of correlation among many structural parameters of bulges. Nonetheless, the scatter remains high.  \\cite{carollo1999} find that exponential bulges are systematically lower in effective surface brightness than those better fit by an $r^{1/4}$-profile. \\cite{falcon2002} find that bulges which deviate from the edge-on projection of the fundamental plane are found in late type (Sbc) galaxies, and it is generally assumed that pseudobulges are more common in late-type galaxies \\citep{kk04}.  \\cite{gadotti2009} finds a large population of pseudobulges that are much fainter in surface brightness for a given half-light radius than predicted by a correlation fit to elliptical galaxies.  There is a great variety of properties which participate in the dichotomy of bulges and motivate the association of pseudobulges with disk-like objects. \\citet{kk04} suggest that kinematics dominated by rotation; flattening similar to that of their outer disk; nuclear bar, nuclear ring and/or nuclear spiral; near-exponential surface brightness profiles are all features of pseudobulges and not classical bulges. \\cite{fisherdrory2008} use morphology to identify bulge-type and find that morphology correlates with S\\'ersic index. \\cite{gadotti2009} uses the position in $\\mu_e-r_e$ parameter space to study the distribution of bulge properties from SDSS data.  Yet, there is no reason to think that all pseudobulges must have a small S\\'ersic index, nor is it necessary that no pseudobulges overlap in $\\mu_e-r_e$ parameter space with classical bulges. In this paper, we will identify pseudobulges using multiple methods including morphology, star formation, and structural properties. We will present a quantitative prescription for identifying pseudobulges using these multiple methods, and apply that method to all non-edge on bulge-disk galaxies within 20~Mpc that have been observed by the Spitzer Space Telescope (SST). We will also use this more robust method for identifying bulge types to study the behavior of pseudobulges in photometric projections of the fundamental plane. ",
        "conclusions": "In this paper we develop a set of pseudobulge diagnostics that incorporate morphology, S\\'ersic index, and specific star formation rate. To accomplish this we carry out bulge-disk decompositions on 173 E-Sd galaxies in the Spitzer~IRAC 3.6~$\\mu$m band. We also measure the 3.6-8.0~$\\mu$m color as a rough estimate of specific star formation for all S0-Sd galaxies in the sample. We include in the appendix a list of notes on bulge diagnosis for all S0-Sd galaxies in our sample.  We find that S\\'ersic index and morphology are essentially interchangeable at 3.6~$\\mu$m as a means of identifying pseudobulges and classical bulges, confirming the results of \\cite{fisherdrory2008} with near-IR data. Pseudobulges have $n_b<2$ and classical bulges have $n_b>2$. Furthermore, with a much more robust sample we confirm that the S\\'ersic index of pseudobulges is uncorrelated with other bulge structural properties.  A careful investigation of van~Albada's (\\citeyear{vanalbada82}) result shows that simulated elliptical galaxies formed via more violent collapses and lumpier initial conditions create resulting surface brightness profiles with larger S\\'ersic index. Also, \\cite{eliche2006} studies minor mergers onto disks with collisionless $n$-body simulations and finds that mergers in general drive the S\\'ersic index of the bulge component up. Also, \\cite{hopkins2009_sersic_merging} finds that more merging generates larger S\\'ersic index. All three of these results seem fairly straight-forward and quite intuitive. Surface brightness profiles with larger S\\'ersic index are characterized by having more light in the tail of the distribution. Also, violent relaxation is a mechanism that can easily convert ordered-rotational orbits to radial orbits, which in turn puts more stellar light at larger radius. Thus more merging, and hence more violent relaxation, puts more stellar mass at large radius which in turn drives up S\\'ersic index. Thus at the minimum the now quite robust result that pseudobulges are marked by having very low S\\'ersic index seems to indicate that they have a more passive history than classical bulges and elliptical galaxies.  We also study a group of pseudobulges (as indicated by S\\'ersic index and nuclear morphology) who are located in the same location in $M_{3.6}-M_{8.0}$~vs.~$M_{3.6}$ parameter space. We call these bulges ``inactive pseudobulges''. We make a selection of $M_{3.6}<-18 AB~mag$ and $M_{3.6}-M_{8.0}<0$; note a similar (but much smaller) group of bulges was found by \\cite{fdf2009}).  We show that those inactive pseudobulges are very similar to models of galaxies in which both a pseudobulge and classical bulge exist in roughly equal parts and the classical bulges has a low S\\'ersic index. Additionally, star formation in inactive pseudobulges must be suppressed.  Therefore, pseudobulge identification that relies only on structural indicators is unable to detect composite systems. Also pseudobulge identification that relies only on stellar populations or star formation rates will underestimate the number of pseudobulges significantly.  We also show that pseudobulges and classical bulges differ in fundamental plane scaling relations. We find that the half-light radius of pseudobulges does not correlate with any other bulge parameter. Across 9~magnitudes in luminosity we find that the radial size of pseudobulges changes only slightly. Thus it appears that the only parameter in bulge-disk decompositions which is linked to pseudobulge half-light radius is the scale length of the outer disk. We find a positive correlation among pseudobulges between mean surface brightness and bulge luminosity. Pseudobulges that are more luminous are more dense. This relation is opposite that of classical bulges and elliptical galaxies, and implies that pseudobulges are very different types of objects than elliptical galaxies."
    },
    "1004/1004.2785_arXiv.txt": {
        "abstract": "The analysis of a sample of 52 clusters with precise and hypothesis-parsimonious measurements of mass, derived from caustics based on about 208 member velocities per cluster on average, shows that low mass clusters and groups are not simple scaled-down versions of their massive cousins in terms of stellar content: lighter clusters have more stars per unit cluster mass. The same analysis also shows that the stellar content of clusters and groups displays an intrinsic spread at a given cluster mass, i.e. clusters are not similar each other in the amount of stars they contain, not even at a fixed cluster mass.  The stellar mass fraction depends on halo mass with (logarithmic) slope  $-0.55\\pm0.08$ and with $0.15\\pm0.02$ dex of intrinsic scatter at a fixed cluster mass.  These results are confirmed by adopting masses derived from velocity dispersion. The intrinsic scatter at a fixed cluster mass we determine for gas mass fractions taken from literature is smaller, $0.06\\pm0.01$ dex. The intrinsic scatter in both the stellar and gas mass fractions is a distinctive signature that, when taken individually, the regions from which clusters and groups collected matter, a few tens of Mpc, are yet not representative, in terms of gas and baryon content, of the mean matter content of the Universe. The observed stellar mass fraction values are in marked disagreement with gasdynamics simulations with cooling and star formation of clusters and groups. Instead, amplitude and cluster mass dependency of observed stellar mass fractions are those requested not to need any AGN feedback to describe gas and stellar mass fractions and X-ray scale relations in simple semi-analytic cluster models. By adding stellar and gas masses and accounting for the intrinsic variance of both quantities, we found the the baryon fraction is fairly constant for clusters and groups with masses between $10^{13.7}$ and  $10^{15.0}$ solar masses and it is offset from the WMAP-derived value by about 6 sigmas. The offset is unlikely to be due to an underestimate of the stellar mass fraction, and could be related to the possible non universality of the baryon fraction, pointed out by our measurements of the intrinsic scatter. Our analysis is the first that does not assume that clusters are identically equal at a given halo mass and it is also more accurate in many aspects. The data and code used for the stochastic computation are distributed with the paper. ",
        "introduction": "Knowledge of the baryon content of clusters and groups is a key ingredient in our understanding of the physics of these objects and in their use as cosmological probes. In fact, clusters have accreted matter from a region of some tens of Mpc, large enough that their content should be representative of the mean matter content of the Universe (White et al. 1993). If this is the case, by measuring the baryon fraction in clusters, $f_b$, and coupling it with an estimate of $\\Omega_b$, for example from primordial nucleosynthesis arguments or from CMB anisotropies, gives $\\Omega_m=\\Omega_b \\  f_b$ (e.g. White et al. 1993; Evrard et al. 1997). Second, the study of how baryons are distributed in gas and stars and the way this splitting depends on halo (cluster or group) mass, should provide clues to the role played by the various physical mechanisms potentially active in clusters and groups.  However, the baryon fraction is far from being fully understood: WMAP-derived value of the baryon fraction is larger than all values found in X-ray analysis (i.e. Vikhlinin et al. 2006) even accounting for baryons in stars (e.g. Gonzalez, Zaritsky, \\& Zabludoff 2007), and gas depletion (e.g. Nagai et al. 2007). X-ray scaling relations (e.g. halo mass vs Temperature or X-ray luminosity) predicted on the assumption that the thermal energy of the gas comes solely from the gravitational collapse are notoriously in disagreement with observed scalings (e.g. Vikhlinin et al. 2006). Observed and predicted scalings may be bring in agreement by allowing star formation, and, eventually a further feedback (e.g. Kravsov et al. 2005, Nagai, Kravtsov, \\& Vikhlinin 2007; Bode et al. 2009; Fabjan et al. 2009). In particular, whether a further feedback, i.e. in addition to the stellar one, is needed, is largely unknown because of the uncertainty of the observed stellar mass content of clusters (e.g. Bode et al. 2009). More generally, recent works on the subject achieve to reproduce X-ray derived quantities (e.g. baryon fraction or mass-temperature scaling relations) by basically adding to the cluster model a further degree of freedom associated with star formation (e.g. Nagai, Kravtsov, \\& Vikhlinin 2007; Bode et al. 2009; Fabjan et al. 2009), without adding the corresponding observational constraint, i.e. requiring that the stellar mass produced in the model fit the data. We emphasise that gas properties strongly depend on the amount of stellar mass allowed in the model (e.g. Nagai \\& Kravtsov 2005; Kravtsov et al. 2005; Nagai, Kravtsov \\& Vikhlinin 2007) and a constraint on the stellar component has a direct and important consequence on the gas component of the model.   Several observational determinations of the stellar mass fraction suffer by important limitations: published works studied clusters with {\\it unmeasured}, or very poorly measured, masses and {\\it unmeasured} reference radii, while these quantities are requested to be known for the determination of the stellar mass fraction, as discussed in later sections. It is clear, therefore, that an observational measurement of the stellar mass fraction of clusters with known masses and reference radii is valuable.  The caustic method (Diaferio \\& Geller 1997; Diaferio 1999) offers a robust path to estimating cluster mass and reference radii. It relies on the identification in projected phase-space (i.e. in the plane of line-of-sight velocities and projected cluster-centric radii, $v,R$) of the envelope defining sharp density contrasts (i.e. caustics) between the cluster and the field region. The amplitude of such an envelope is a measure of the mass inside $R$. As opposed to masses derived in other ways (e.g. from X-ray, from velocity dispersion, from the virial theorem, from the Jeans method, etc.) caustic masses do not require that the cluster is in dynamical equilibrium (see Rines \\& Diaferio 2006 for a discussion). There is a good agreement between caustic and lensing masses for the very few clusters where both measurements are available (Diaferio, Geller, \\& Rines 2005). On larger cluster samples, caustic masses also shows a good agreement with virial masses (e.g. Rines \\& Diaferio 2006, Andreon \\& Hurn 2010) and with the extrapolation to larger radii of dynamical masses derived through a Jean analysis (Biviano \\& Girardi 2003). Both virial and Jean masses require, however, assume that the cluster is in dynamical equilibrium.    This paper addresses: a) the determination of the stellar mass fraction in a sample of clusters and groups with well determined masses and reference radii derived by the caustic method, using, on average, 208 members per cluster; and b) the determination of the average baryon content of clusters and groups.   Throughout this paper we assume $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. Magnitudes are quoted in their native system (quasi-AB for SDSS magnitudes). ",
        "conclusions": "We analysed a sample of 52 clusters with precise and hypothesis-parsimonious measurements of mass, derived from caustics based on about 208 member velocities per cluster on average, and with measured $r_{200}$ values. We found that low mass clusters and groups are not simple scaled-down version of their massive cousins in terms of stellar content: lighter clusters have more stars per unit cluster mass. The same analysis also shows that the stellar content of clusters displays an intrinsic spread at a given cluster mass, i.e. clusters are not similar each other in the amount of stars they contain, not even at a fixed cluster mass. The amplitude of the spread in stellar mass, at a fixed cluster mass, is $0.15\\pm0.02$ dex. The stellar mass fraction depends on halo mass with (logarithmic) slope  $-0.55\\pm0.08$. These results are confirmed by adopting masses derived from velocity dispersion.  The intrinsic scatter at a fixed cluster mass we determine for gas mass fractions taken from literature, is smaller, $0.06\\pm0.01$ dex. The intrinsic spread is not restricted to low mass systems only, but extend to massive systems. Since the studied systems look relaxed in X-ray images, the found spread is not due the presence in the sample of clusters manifestly out of equilibrium (e.g. merging). The non-zero intrinsic scatter of the gas mass fraction decreases the efficiency of $f_{gas}$ for cosmological studies, and asks to inquire about whether studied samples are representative, in terms of $f_{gas}$, to the population of clusters in the Universe.  The intrinsic scatter in both the stellar and gas mass fraction is a distinctive signature that when taken individually the regions in which clusters and groups collected matter are yet not representative, in terms of stellar and gas content and therefore in the baryon content, of the mean matter content of the Universe.  The observed stellar mass fraction values are in marked disagreement with gasdynamics simulation with cooling and star formation of clusters and groups. Instead, amplitude and cluster mass dependency of observed stellar mass fraction are those requested not to need any AGN feedback to describe X-ray scale relations and gas and stellar mass fractions in simple semi-analytic cluster models.     By adding the stellar and gas masses, or, more precisely speaking, by fitting both them and accounting for the intrinsic variance of both quantities, we found that the baryon fraction is fairly constant for clusters and groups with $13.7<\\log M200 < 15.0$ solar masses and it is offset from the WMAP-derived value by about $6$ sigmas. The offset is unlikely to be due to an underestimate of the stellar mass fraction and could be related to the possible non-universality of the baryon fraction, pointed out by our measurements of the intrinsic scatter.  Our analysis is the first that does not assume that clusters are identically equal at a given halo mass and it is also more accurate in many aspects than previous works. The data and code used for the stochastic computation are distributed with the paper."
    },
    "1004/1004.2266_arXiv.txt": {
        "abstract": "We revisit the issue of relaxation to thermal equilibrium in the so-called ``sheet model'', i.e., particles in one dimension interacting by attractive forces independent of their separation. We show that this relaxation may be very clearly detected and characterized by following the evolution of order parameters defined by appropriately normalized moments of the phase space distribution which probe its {\\it entanglement} in space and velocity coordinates. For a class of quasi-stationary states which result from the violent relaxation of rectangular waterbag initial conditions, characterized by their virial ratio $R_0$, we show that relaxation occurs on a time scale which (i) scales approximately linearly in the particle number $N$, and  (ii) shows also a strong dependence on $R_0$, with quasi-stationary states from colder initial conditions relaxing much more rapidly. The temporal evolution of the order parameter may be well described by a stretched exponential function. We study finally  the correlation of the relaxation times with the amplitude of fluctuations in the relaxing quasi-stationary states, as well as the relation between temporal and ensemble averages. ",
        "introduction": "\\label{introduction}  The so-called ``sheet model'' is an interesting toy model for the study of self-gravitating systems, or more generally of systems with long-range interactions. It is simply the one dimensional (1D) generalisation of Newtonian gravity, consisting of particles interacting by attractive forces independent of their separation (or, equivalently, infinite parallel planes embedded in three dimensions interacting via Newtonian  gravity). Because the particle trajectories are exactly integrable between crossings, it has the nice feature that its numerical integration can be performed with an accuracy limited only by machine precision. It has been the subject of (mostly numerical) study in the literature for several decades (see, e.g., \\cite{yawn+miller_2003} for a review of the literature on the model) following earlier analytical studies \\cite{camm, oort_1932}. A fundamental question about this system --- and more generally for any system with long-range interactions --- is whether they relax to the statistical equilibrium calculated in the microcanonical or canonical ensemble. For this model the latter were first calculated exactly,  for any particle number $N$,  by Rybicki \\cite{rybicki}.  The literature on this model --- which we will discuss in  greater detail in our conclusions section below ---  is marked by differing results (or, rather, interpretation of results) from different groups, and even some controversy.  Work by two groups in the eighties (see, e.g. \\cite{reidl+miller_1987} for a summary) led to the conclusion that relaxation could not be observed, except perhaps in some special cases.  Studies by two other groups over a decade ago  \\cite{tsuchiya+gouda+konishi_1996, milanovic+posch+thirring} found results indicating relaxation, and  \\cite{tsuchiya+gouda+konishi_1996} gave a  determination of the $N$ dependence of the characteristic time. However doubts about the interpretation of these latter results as establishing relaxation to equilibrium were raised by further analysis \\cite{tsuchiya+gouda+konishi_1997, yawn+miller_erg_1997}. In more recent work \\cite{yawn+miller_2mass_1997, yawn+miller_2003} clear evidence for relaxation in a version of the model in which there are different particle masses has been found, but the dependence on $N$ has not been determined\\footnote{Other variants of the model have also been studied in \\cite{miller+youngkins_concentric_1998,valageasOSC_1,valageasOSC_2}.} The mechanism of relaxation (if it indeed takes place) in these models remains, as in other long-range interacting systems, very poorly understood, and a basic subject of research in the statistical mechanics of long-range interacting systems (for recent reviews see e.g. \\cite{campa_etal_LRreview_2009, bouchet+mukhamel_LRreview_2010}). In this article we report an essentially numerical study of relaxation in the single mass sheet model. We introduce a simple but, as we will see, very useful tool for the characterisation of the long-time evolution and relaxation of the system. This tool allows us to resolve some outstanding issues about the relaxation in this system, and, in particular, to establish more definitively both that relaxation does indeed occur and the scaling with particle number of the time characterizing it. We consider a broader range of initial conditions, which allows us to establish also dependences of relaxation on these.  We also study the fluctuations --- both in time and over realizations of the initial conditions ---  about the average macroscopic evolution of the system,  showing phenomenologically  the correlation of their amplitude with the lifetimes of the intermediate ``quasi-stationary\" states.  We will discuss in greater detail in our conclusions the relation of our results to those in the previous literature, but it is useful at the outset to say a little more about the more general context of this study. In recent years there has been considerable interest in the statistical mechanics of long range interactions (see, e.g., \\cite{dauxoisetal, Assisi, campa_etal_LRreview_2009}), stimulated by the need to understand the physics of various laboratory systems with interactions of this kind, as well as by the more classical case of self-gravitating matter relevant in astrophysics and  cosmology. In this context one toy model in particular, and various variants of it, has been much studied:   the Hamiltonian Mean Field (HMF) model (see e.g. \\cite{yamaguchi_etal_04, campa_etal_2007,chavanis+campa_HMF_2010} and references therein), which is a model of particles on a circle interacting by a cosine potential.  Its study has shown that it shares many of the qualitative features well documented in the most studied of realistic long-range interacting systems --- self-gravitating systems in astrophysics --- and believed to be generic in such systems. Starting from generic initial conditions, the system evolves rapidly (by ``violent relaxation'') to a virialized macroscopically stationary state. These states  --- commonly referred to in the more recent literature as ``quasi-stationary states'' (QSS) --- are out of equilibrium states, which are described theoretically in the framework of Vlasov equation (more usually referred to as the ``collisionless Boltzmann equations'' in the astrophysical literature). On much longer time scales an evolution towards the true thermal equilibrium (i.e. that determined by the maximization of the Boltzmann entropy in a mean field approximation) is postulated.  For realistic systems --- such as Newtonian gravity in three dimensions ----  it is very difficult numerically to simulate the evolution on sufficiently long time scales to probe the relaxation. Studies in the literature (see e.g. \\cite{theis+spurzem_1999, diemandetal_2body, knebe, elzant_2006, levin_etal_2008}) provide some results but give still a very limited characterization and understanding of it.  The HMF model has the particular feature that the potential energy of any particle may be expressed as a function of its (angular) position and the mean potential energy due to all particles --- it is for this reason that it is ``mean-field'' ---   so that the calculation of the forces in a system with $N$  particles requires only of order $N$ operations (rather than $N^2$ in a typical long-range interacting system).  Further the force is continuous at zero separation, so that the difficulties associated in the case of gravity with the regulation of the potential at small scales are avoided. This allows the regime of relaxation to be accessed numerically even for quite large particle numbers.  The study of \\cite{yamaguchi_etal_04} found a scaling of the relaxation time in proportion to $N^{1,7}$ (but see also e.g. \\cite{campa_etal_2007} which finds indications of longer lifetimes for other initial conditions).  It can clearly be of interest to study different toy models, to determine in particular features which are indeed generic. The ``sheet model\" is probably the oldest toy model of long range interactions --- it was first explored in astrophysics as a toy model for self-gravitating systems  in three dimensions --- and is also, arguably, closer to reality than the HMF which is constrained on a circle. It has, further, as mentioned above the nice feature that it numerical integration can be performed up to machine precision. Despite this, the results concerning its dynamics and relaxation are less clearly determined than for the HMF, and the literature on the subject has, as we have discussed above,  been marked by some controversy and results showing that the model has, apparently, some very peculiar behaviours --- rapid relaxation to equilibrium for some classes of initial states \\cite{luwel+severne+rousseeuw_1983, reidl+miller_1991}, persistent phase space structures impeding relaxation to QSS \\cite{rouet+feix, mineau+rouet+feix_1990}, macroscopically chaotic behaviour in the long time evolution \\cite{tsuchiya+gouda_2000} --- which indicate that it might not be a very useful toy model (in that its behaviours are perhaps non-generic). In this article our main conclusion is that, 1) by using appropriate diagnostics of the macroscopic evolution and 2) by  extending simulations to sufficiently large $N$ and/or averaging over sufficiently large numbers of realization,  one finds behaviour in this toy model very similar in crucial respects to that in the HMF:  to a very good first approximation a generic initial configuration relaxes to a long-lived QSS, and then relaxes to its statistical equilibrium at sufficiently long time. This latter phase can be characterized apparently by a single time-scale, with the evolution of the order parameter during relaxation well fit by a simple function (in our case a better fit is obtained using a simple stretched exponential, rather than a hyperbolic tangent in the HMF as in \\cite{yamaguchi_etal_04}). On the other hand the $N$ dependence of this time scale,  linearly proportional to the number of particles $N$, is different to that found in  \\cite{yamaguchi_etal_04} for the HMF. This latter result, however, applies to spatially homogeneous states which, in the HMF, can occur due to the periodicity of the system.  Relaxation which is slower than linear in $N$ is expected in this case, as shown using  analysis of kinetic equations (see, e.g. contributions of P.H. Chavanis, and of  F. Bouchet and J. Barr\\'e in \\cite{Assisi}).  The article is organised as follows. In the next section we recall the basic definitions of the model, and relevant results on its statistical equilibrium.  We then explain the choice of the macroscopic parameters  (``order parameters'') we choose to monitor the evolution of  the system. In the following section we first describe our numerical simulations and the initial conditions we study, and then give our results. In presenting them we give first results for single realizations, and then use temporal averages and finally ensemble averages to derive the scaling with $N$ of the relaxation time. This is followed by further study of the fluctuations about the average behaviours of the order parameters.  Considering both temporal fluctuations and those in the ensemble, which we show to be very consistent with one another, we observe the correlation between their amplitude in the QSS and the observed relaxation time. In the conclusion sections we return to a more detailed discussion of the previous literature, presenting further results which allow one to understand the reasons for the divergence in conclusions in certain cases. ",
        "conclusions": "\\subsection{Summary}  Our primary aim in this paper has been to establish and characterize more fully than in the previous literature the relaxation to thermodynamic equilibrium of one of the simplest toy models for long-range interacting systems: equal mass self-gravitating  particles in one dimension (or infinite sheets in three dimensions). Compared to the much studied HMF model, notably, the basic properties of this model have remained somewhat unclear, and indeed marked by some controversy in the literature. The novelty of our work compared to previous studies is not just that we do more and larger simulations from a broader range of initial conditions, but that we have identified a tool which is very useful to characterize the evolution of the system: the measurement of appropriately normalized moments of the distribution function which characterize the ``entanglement\" of the one particle distribution function in configuration and velocity space.  This is particularly appropriate as a measure simply because the thermal equilibrium has the property that such entanglement is absent while, we have shown, in {\\it any other} stationary solution of the Vlasov-Poisson equations such entanglement is present. We note that this result, which we showed to be valid for any interaction in one dimension (but, as noted, excluding periodic systems like the HMF),  can be generalized easily to three dimensions if we restrict to stationary solutions which have radial symmetry in space and velocity. This suggests that these ``order parameters\" may also be useful indicators of relaxation in much more general, and perhaps, as we discuss below,  even  useful tools for understanding the mechanisms of such relaxation.  With the aid of these macroscopic measures, we have shown in our numerical study,  of a range of  simple ``waterbag\" and cold initial conditions,  that the system manifests the behaviour thought to be generic in long-range systems: there are essentially two phases in the evolution with two completely different time-scales. An initially non-stationary state evolves first,  on timescales characterized by the ``dynamical time\"  $t_{\\rm dyn}$ (roughly the crossing time of a particle in the mean-field due to the others),  to a QSS, an out of equilibrium state, which then evolves on a much longer time scale, dependent on the number of particles,  to thermal equilibrium. In other words it is reasonable to suppose that the system is ergodic (and mixing) on these very long time scales, but not so on the shorter time scales. Further we can identify clearly that the QSS resulting from different initial conditions (i.e. different values of $R_0$)  are very different macroscopically, characterized by very different phase space entanglement.  Focussing on the the $N$ dependence of the relaxation, averaging over a very large number of realizations to average out the fluctuations, we have concluded that the characteristic time scale for relaxation behaves, to a very good approximation, as \\be t_{\\rm relax}  \\sim  f_{QSS} \\, N \\, t_{\\rm dyn} \\label{relaxation-scaling} \\ee where $f_{QSS}$ is a numerical factor which depends on the initial condition, which we have denoted in this way as we expect that this dependence is not strictly on the initial condition but on the QSS which results from it.  We have seen that this prefactor increases as $R_0$ does, by about a factor of ten between $R_0=0$ and $R_0=0.5$, and approximately a further factor of ten  for $R_0=1$. We have noted that the overall normalization of $f_{QSS}$ is rather arbitrary, as it depends greatly on the exact criterion used to define the relaxation time-scale.  Given that the evolution towards zero of $\\phi_{11}$, which is what we have used to determine this time scale, is in fact well fit by the simple functional behaviour as a function of the time on a logarithmic scale, the normalisation of $f_{QSS}$ can differ by two orders of magnitude  by a trivial change in its definition. More specifically we have seen, that in the case where we have accumulated the greatest statistics allowing us to constrain the temporal evolution, a very good fit to our order parameter $\\phi_{11}$ is obtained to a stretched exponential form.  Although the relaxation of this system, and in general long-range interacting systems, is not well understood, we can say that this finding of a linear scaling --- besides the fact that it is, as we will discuss below,  in line with less complete previous analyses --- is not a surprising result: such a scaling can be anticipated both on the basis of simple naive estimates of the effects of collisionality, as well from general considerations based on kinetic theory.  A simple ``derivation\" of this scaling, along the lines of those applied originally by Chandrasekhar (see \\cite{chandra43} or \\cite{binney}) to obtain such an estimate for 3D self-gravitating systems, may be given as follows. Relaxation is in principle due to the discretisation, in $N$ particles, of a continuous mass distribution. Let us suppose that this latter field density varies spatially on a scale, $\\ell$.  The typical fractional  change in the velocity $v$  of a test particle due to its interaction with one (localized) particle, rather than the continuous mass distribution,  can be estimated as  $\\sim g \\ell/m v^2$. As it crosses the system (in a time  $\\sim t_{dyn}$) such a particle will interact with all $N$  particles. Assuming the contribution from each particle adds incoherently, one estimates \\be \\frac{\\delta v^2}{v^2} \\sim N \\left(\\frac{gl}{mv^2} \\right)^2 \\ee for the normalized variance of the velocity in $t_{dyn}$. Scaling with $N$ at fixed mass and energy (and fixed $\\ell$) requires $g/m \\sim 1/N$, and therefore  $\\frac{\\delta v^2}{v^2} \\sim N^{-1}$. It follows that the relaxation time scales linearly with $N$ in units of the $t_{dyn}$. A slightly different argument leading to the same result may be found in \\cite{tsuchiya+gouda+konishi_1996}, and a more developed analysis in \\cite{miller_1996}. In the framework of approaches based on kinetic theory, a linear scaling is obtained as collisional terms arise as the leading corrections in a formal expansion in $1/N$ which gives the collisionless (Vlasov) limit at leading order (see, e.g. \\cite{balescu,chavanis_KTb_2006, chavanis_kEqns_2010}).  This scaling linear in $N$ is to be contrasted with the case of the scaling observed for the life-time in QSS in the HMF, proportional to $N^{1.7}$. While the exponent found is not understood,  the fact that it is larger than unity is consistent with these considerations as this result applies for  spatially homogeneous QSS (which are possible in the HMF because of its periodicity) for which it has been shown that the leading collisional term vanishes (see, e.g., contributions of P.H. Chavanis, and of  F. Bouchet and J. Barr\\'e in \\cite{Assisi}).  We note that our study suggests also that the ``order parameters\" we have defined and studied may be relevant quantities for understanding relaxation in this and other long-range systems. Indeed in all cases we have  observed that, at sufficiently large $N$, these parameters start from a non-zero value in the initial QSS and evolve monotonically towards zero, i.e., the relaxation of the QSS can apparently be described as a  process of progressive ``disentanglement'' of the one particle phase space density. In this respect the very different, much less efficient, relaxation observed in the HMF might be interpreted as a result of the absence of such entanglement in spatially uniform QSS.  Further, in the case where we have enough statistics to provide a precise fit to the evolution of the parameter $\\phi_{11}$, we found that it is well fit by a simple stretched exponential form.  It would be interesting to see in further study whether this fit is more than an approximate numerical fit for the case we have studied, and, if so, whether the exponent characterizing it is the same or not. As we have remarked such a functional form has been observed in other slowly relaxing (e.g. glassy) systems and theoretical tools derived in this  context to understand relaxation may be relevant. In \\cite{almeida_etal_2001}, for example, this behaviour is linked to the existence of a fractal structure in a bounded accessible region of phase space.  \\subsection{Comparison with previous literature}  Let us now turn finally to compare our findings in greater detail with those in the previous literature.  An early numerical study by Hohl and Broaddus \\cite{hohl+broaddus_1967} which concluded a relaxation time proportional to $N^2 t_{dyn}$  was found to be incorrect by two groups, who studied the problem in greater detail (and with greater numbers of particles). However, these groups found conflicting results:  Miller et al.  found no evidence for  relaxation at all to thermal equiibrium in their simulations \\cite{wright+miller+stein_1982}, while Luwel et al. \\cite{luwel+severne+rousseeuw_1983} found relaxation on a time scales even shorter than $N t_{dyn}$. Further study (see \\cite{reidl+miller_1987, reidl+miller_1991}, which also contain a detailed synthesis of the previous literature) by Miller et al. concluded that the discrepancy was related to the very specific initial condition studied by the other group. Studying this case in detail they found that it indeed appears to thermalize very rapidly, but some further, but not completely conclusive analysis of the evolution at longer times, suggested that this thermalization was not complete.  \\begin{figure}[h!] \\begin{center} \\includegraphics[width=9cm]{./figures/V1.eps} \\caption{A ``counterstreamed'' waterbag initial condition in phase space with $R_{0}=0.3$, sampled with  $N=400$ particles.} \\label{fig_v} \\end{center} \\end{figure}  \\begin{figure}[h!] \\begin{center} \\begin{tabular}{cc} \\includegraphics[width=7.5cm]{./figures/W1.eps} & \\includegraphics[width=7.5cm]{./figures/W2.eps} \\end{tabular} \\caption{Density profile (left panel) and velocity distribution (right panel) obtained at $t=100$ from a counter-streamed initial condition with $N=100$. An average over $250$ simulations from independent realizations of the initial conditions has been performed. The solid lines correspond to the values in thermal equilibrium,  Eq.~(\\ref{rybicki-eq}). } \\label{fig_w} \\end{center} \\end{figure}  To determine whether these cases are consistent with our findings --- and see whether our analysis using the parameters $\\phi_{11}$  and $\\phi_{22}$ can throw light on these previous findings --- we have resimulated the relevant initial conditions. These are ``counterstreamed\"  waterbag initial conditions, an example of which is shown in Fig.~\\ref{fig_v}. We have simulated a range of such initial conditions, in particular the cases (one of which is that shown in the figure) considered by  \\cite{luwel+severne+rousseeuw_1983} and  \\cite{reidl+miller_1987}. Shown in Fig. \\ref{fig_w} are the density profile and velocity distribution at $t=10^2$ obtained starting from a realization of initial conditions like those shown in Fig.~\\ref{fig_v}, but with $N=100$. We see that the profiles indeed agree very well with the equilibrium ones. In Fig. \\ref{fig_x} is shown the evolution of  $\\phi_{11}$ as a function of time for the indicated values of $N$ averaged in each case over the number of realizations indicated. We observe that, although small and fluctuating, its value is clearly on average non-zero, indicating that the state, despite the good agreement of the density and velocity profiles, is not in fact an equilibrium.  Just as in the cases we studied we see clearly the relaxation towards equilibrium at longer times, and indeed that the characteristic time increases on $N$. Although we haven't done the more extensive study required to determine precisely this $N$ dependence, the results are quite consistent with Eq.~(\\ref{relaxation-scaling}) with a value of $f_{QSS}$ of order that found for the case $R_0=0$.  \\begin{figure}[h!] \\begin{center} \\includegraphics[width=9cm]{./figures/X1.eps} \\caption{Evolution as a function of time of $\\phi_{11}$ from a counterstreamed waterbag initial condition, averaged over the number of realizations of the initial conditions and particle numbers indicated. Despite the indications of the previous figure, we observe clearly that relaxation to thermal equiibrium has not taken place at $t=100$.} \\label{fig_x} \\end{center} \\end{figure}  This case illustrates  the usefulness of the parameters $\\phi_{11}$ and $\\phi_{22}$ as discriminants of relaxation: indeed we have just seen that the single measure of $\\phi_{11}$ is sufficient to discard the interpretation of Luwel et al. \\cite{luwel+severne+rousseeuw_1983} of their results. This is simply because they are physically very appropriate indicators, for the reasons we have explained in introducing them:   the property they probe --- of entanglement of the phase space distribution --- is one which must evolve significantly during relaxation, because the phase distribution must become separable. While $\\phi_{11}$ and $\\phi_{22}$  being zero does not {\\it imply} thermalization, of course,  we have not found a single QSS, despite exploring a broad range of initial conditions (considerably more extended that those reported here) in which they are {\\it both} zero (within the uncertainty of fluctuations), i.e., the only states  we have found in which  they are  both  zero are states which we have concluded, using a range of other measures, are indistinguishable  from the equilbrium state of Rybicki. It is not difficult, on the other hand, to find initial conditions which lead to a QSS in which  $\\phi_{11}\\approx 0$ {\\it or} $\\phi_{22}\\approx 0$. Indeed for the waterbag initial conditions we have studied both $\\phi_{11}$ and $\\phi_{22}$ actually change sign as  $R_0$ varies over the range we have  considered, and one can thus find by trial and error the appropriate $R_0$ which make them zero individually.  \\begin{figure}[h!] \\begin{center} \\begin{tabular}{ccc} \\includegraphics[width=5cm]{./figures/Y1.eps} & \\includegraphics[width=5cm]{./figures/Y2.eps} & \\includegraphics[width=5cm]{./figures/Y3.eps} \\\\ \\includegraphics[width=5cm]{./figures/Y4.eps} & \\includegraphics[width=5cm]{./figures/Y5.eps} & \\includegraphics[width=5cm]{./figures/Y6.eps} \\end{tabular} \\caption{Histogram $f(e)$ of individual particle energies $e$ measured at the indicated times starting from counter-streamed initial conditions sampled with $N=100$ particles. The curves are averaged over $250$ realizations. In thermal equilibrium $f(e)$ is indistinguishable from the one measured at the latest time shown. The result confirms that relaxation in fact occurs on time scales similar to those observed for the simple waterbag initial conditions.} \\label{fig_y} \\end{center} \\end{figure}  Another evident quantity to measure, which we have in fact considered systematically but have not reported in detail, is the distribution $f(e)$ of the individual particle energy $e$. This is in fact generally a better discriminant of relaxation than either $n(x)$ and $f(v)$, i.e., we have found that in quite alot of cases $n(x)$ and $f(v)$ are not easy to distinguish from the equilibrium profiles, but that $f(e)$ allows one to see more clearly that one is indeed not in the equilibrium state.  An example is the counter-streamed case just considered above. In Figs. \\ref{fig_y}  are shown, for $N=100$, the evolution of the ensemble averaged $f(e)$ at a few different times. We have not plotted the equiibrium curve, as it is indistinguishable from the measured curve at the final time shown. One can see clearly see that, despite the good agreement of $n(x)$ and $f(v)$ shown in Fig. \\ref{fig_w}, the system is not in equilibrium at the early times: $f(e)$ has a clearly visible excess of particles at high energies compared to that at the much later times at which the evolution of $\\phi_{11}$ indicated relaxation (and $f(e)$ indeed approaches very accurately its predicted equilibrium form). While such a measure of $f(e)$, averaged over a large ensemble of realizations,  can, in all the cases for which we've studied it, clearly discriminate relaxation, the use of just $\\phi_{11}$ (and possibly  $\\phi_{22}$) is an extremely efficient short-cut to ``diagnose\" relaxation.  Subsequent to  \\cite{reidl+miller_1987}, in the nineties, Tsuchiya et al. reported an analysis of larger and more importantly, longer, simulations in order to clarify the issue.  A first paper \\cite{tsuchiya+gouda+konishi_1996} they reported the evolution of a rectangular waterbag initial condition corresponding to our case $R_0=1$, and reported a detection of relaxation to thermal equilibrium. These authors made a distinction between two time scales of relaxation:  one of  ``microscopic relaxation'',   the other for ``macroscopic relaxation''. These are identified, and both found to be proportional to $N$, by considering the evolution of the mean standard deviation of the particle energies averaged over a time window $T$ from their  equipartition value. The former  is estimated from the slope at short time of this function, and the latter from the position of ``peaks\" which are observed to occur at much longer times.  While the latter is interpreted in terms of of macroscopic relaxation in the sense we have used here, the former is interpreted as a time scale on which particles sample the energy distribution but on which there is {\\it no} macroscopic evolution. The justification for these interpretations are not clear, and no direct comparison with the equilibrium distribution derived by Rybicki, Eq.~ (\\ref{rybicki-eq}),  has been reported which might show their correctness. Indeed both a subsequent article by the same authors \\cite{tsuchiya+gouda+konishi_1997} and a study by Yawn and Miller \\cite{yawn+miller_erg_1997} place in doubt the correctness of the  interpretation in terms of relaxation.  Nevertheless, in light of the results we have given here, it would be reasonable to infer that the results given by Tsuchiya et al. in  \\cite{tsuchiya+gouda+konishi_1996} are indeed correct, at least for what concerns the $N$ dependence of the relaxation. Further comparison could of course clarify the relation of the behaviour of their measured quantities and the macroscopic relaxation as we have probed it here (and should be much easier for the shorter lived, smaller $R_0$, initial conditions rather than the case $R_0=1$ studied by these authors). We do not believe, however, that there is any clear basis for either an operational or {\\it physical} distinction between ``microscopic\" and ``macroscopic\" relaxation as described by these authors: as we have discussed there is an arbitrariness in the definition of the relaxation time because of the very slow nature of this relaxation. As we have noted, we could easily, for example,  have obtained here estimates of the relaxation time differing by several of orders in magnitude in their prefactor, just like the two different time scales determined  by Tsuchiya et al.  \\cite{tsuchiya+gouda+konishi_1996}, by using slightly different definitions, or choosing to use a different order parameter.  This point can be illustrated by considering the evolution of $f(e)$ for one of the cases we have considered: shown in Figs. \\ref{fig_z} is this quantity for the case $R=0.1$ and $N=400$, averaged over $60$ realizations.  While we have associated (see Table \\ref{table-1} above) the time scale $7 \\times 10^{5}$ to the relaxation in our analysis, one can discern by inspection of these figures significant evolution (in particular of the initially clear ``core-halo\" structure)  in $f(e)$ already by $t=10^{3.5}$, i.e., there {\\it is} evolution of the energy distribution on the time scale of ``microscopic\" relaxation (of order $N t_{dyn}$) identified in \\cite{tsuchiya+gouda+konishi_1996}. While it is possible that there are different time scales associated to different physical processes as argued in \\cite{tsuchiya+gouda+konishi_1996}, it seems a more plausible interpretation to us to suppose that there is single physical relaxation process leading, albeit very slowly, to macroscopic relaxation of the system, and to characterize this relaxation by a function and the scaling of its parameters with $N$. In this respect it is interesting to note that  the specific stretched exponential form we fitted to the temporal behaviour has the known property \\cite{montroll+bendler_1984} that it is can be written as a weighted integral over simple exponentials (i.e. it can be interpreted as arising from the superposition of an infinite number of relaxation processes each with a single characteristic time).  \\begin{figure}[h!] \\begin{center} \\begin{tabular}{ccc} \\includegraphics[width=5cm]{./figures/Z1.eps} & \\includegraphics[width=5cm]{./figures/Z2.eps} & \\includegraphics[width=5cm]{./figures/Z3.eps} \\\\ \\includegraphics[width=5cm]{./figures/Z4.eps} & \\includegraphics[width=5cm]{./figures/Z5.eps} & \\includegraphics[width=5cm]{./figures/Z6.eps} \\\\ \\includegraphics[width=5cm]{./figures/Z7.eps} & \\includegraphics[width=5cm]{./figures/Z8.eps} & \\includegraphics[width=5cm]{./figures/Z9.eps} \\\\ \\includegraphics[width=5cm]{./figures/Z10.eps} \\end{tabular} \\caption{Histogram $f(e)$ of individual particle energies $e$ at the indicated time, and averaged over $60$ simulations from realizations of simple waterbag initial conditions with  $R_{0}=0.1$ and $N=400$. The curve at the latest time coincides well with that expected in thermal equilibrium. The onset of relaxation is already visible at a time of order $10^4$,  almost two orders of magnitude smaller than the time determined in Table~ \\ref{table-1}.} \\label{fig_z} \\end{center} \\end{figure}  In \\cite{tsuchiya+gouda+konishi_1997} Tsuchiya et al. have also described chaotic ``itinerant\" behaviour of these systems, starting from the same ($R_0=1$) initial conditions i.e., in which the system shows apparently stochastic {\\it macroscopic} behaviour. In our analysis  this would correspond to such behaviour for the parameters $\\phi_{11}$ or $\\phi_{22}$. While we have seen that there are indeed very significant fluctuations in these parameters, which correspond to very significant differences in the ``macroscopic\" evolution of these systems, we have studied carefully their dependence on $N$ and found them to decay monotonically. The results of \\cite{tsuchiya+gouda+konishi_1997} were obtained for $N=64$, a range in which we still see fluctuations of $\\phi_{11}$ which are order unity. Only when we reach $N$ of order several hundred do we see these fluctuations diminish significantly so that the macroscopic trajectory of the system becomes quite localized. We thus believe that as $N$ increases these effects will becomes negligible, even on the time scales on which relaxation occurs, and an effectively deterministic macroscopic evolution will occur.  It is interesting to compare our results also to those of Yawn and Miller \\cite{yawn+miller_2003}, who have analyzed in detail relaxation in a version of the sheet model in which there are sheets of different masses. In this case the relaxation towards thermal equilbrium may be clearly distinguished by testing for equipartition of the kinetic energy, and the associated spatial segregation of the mass populations. In simulations starting from waterbag type initial conditions with a virial ratio of two, for a range of different mass ratios and up to $N=128$ particles, clear evidence was found in \\cite{yawn+miller_2003} for such relaxation using appropriately defined indicators. Like the order parameters we have employed here, these show characteristic behaviours corresponding to the principal phases of the dynamical evolution (violent relaxation, QSS phase, relaxation towards thermal equilibrium). Although we cannot compare our results directly, we note that the time scales observed for relaxation of systems with $N \\sim 10^2$ particles are quite consistent with those we have observed for the equal mass system with initial virial ratio $R_0=1$.  Yawn and Miller \\cite{yawn+miller_2003} also measure temporal  correlation properties and find weak but persisting correlations characterized by a power-law decay (in time), which they interpret as evidence for the incompleteness of relaxation. In the present study we have found, in contrast, that our principal observables decay in time with a functional form which allows the identification of characteristic time scales. Further all deviations of these observables from their equilibrium values decrease clearly as $N$ increases, and thus we have interpreted  the associated ``incompleteness''  of relaxation simply in terms of finite $N$ effects. It would be interesting certainly to perform a more direct comparison of the results in the two models, and in particular to extend the study of Yawn and Miller to allow a determination of the $N$ dependence of the parameters they study. We note also that Yawn and Miller argue that the power-law decay suggests the existence of a fractal structure in phase-space, which, as we have been mentioned above, is also proposed as an explanation in \\cite{almeida_etal_2001} for the appearance of relaxation characterized by a stretched exponential behaviour.  Let us finally mention some other issues of importance concerning aspects of the dynamics of this system which have been treated elsewhere but which we have not discussed here. As we have discussed, we interpret our results in line with those of many previous studies of  this and other long-range systems: the evolution from an arbitrary out of equilibrium initial condition is characterized a first phase of relaxation to a QSS, interpreted as a finite particle sampling of a stationary solution of the Vlasov equation, on a time scale independent of $N$, followed by a slow relaxation to thermal equilibrium on an $N$-dependent time scale.  Studies of the single mass sheet model for other specific initial conditions suggest that this simple scheme may be too limiting, for this model (and possibly, for all such models). On the one hand Reidl and Miller have reported numerical results \\cite{reidl+miller_1995} for specific ``two cluster\" initial conditions which show a dependence on $N$ in the time scale for relaxation to a QSS. On the other hand,  as mentioned in the introduction, Rouet et al. \\cite{rouet+feix, mineau+rouet+feix_1990} have shown, using both particle simulations and simulations of the Vlasov equation, for yet other initial conditions that ``holes\" which rotate in phase space may be present after violent relaxation and persist on very long time scales. Although it is not  evident that there is necessarily a relation between either finding and the mechanism of relaxation to thermal equilibrium, a study incorporating such initial conditions would certainly be more complete that that reported here. Extension of the study reported here to larger $N$ still would likewise be desirable, despite the extremely rapidly growing numerical cost of such simulations with $N$.   \\subsection{Acknowledgements}  The simulations were carried out in large part at the Centre de Calcul of the Institut de Physique Nucl\\'eiare et Physique des Particules. We are particularly grateful to Laurent Le Guillou for advice and help. We thank also Duccio Fanelli for providing us with his own code which allowed us to perform checks of ours.  We thank P. Astier, J. Barr\\'e, A. Gabrielli, B. Marcos, P. Viot, F. Sicard, F. Sylos Labini for useful conversations, comments or suggestions.  \\vskip 1cm"
    },
    "1004/1004.0439_arXiv.txt": {
        "abstract": "We show an analytic method to construct a bivariate distribution function (DF) with given marginal distributions and correlation coefficient. We introduce a convenient mathematical tool, called a copula, to connect two DFs with any prescribed dependence structure. If the correlation of two variables is weak (Pearson's correlation coefficient $|\\rho| <1/3 $), the Farlie-Gumbel-Morgenstern (FGM) copula provides an intuitive and natural way for constructing such a bivariate DF. When the linear correlation is stronger, the FGM copula cannot work anymore. In this case, we propose to use a Gaussian copula, which connects two given marginals and directly related to the linear correlation coefficient between two variables. Using the copulas, we constructed the BLFs and discuss its statistical properties. Especially, we focused on the FUV--FIR BLF, since these two luminosities are related to the star formation (SF) activity. Though both the FUV and FIR are related to the SF activity, the univariate LFs have a very different functional form: former is well described by the Schechter function whilst the latter has a much more extended power-law like luminous end. We constructed the FUV-FIR BLFs by the FGM and Gaussian copulas with different strength of correlation, and examined their statistical properties. Then, we discuss some further possible applications of the BLF: the problem of a multiband flux-limited sample selection, the construction of the SF rate (SFR) function, and the construction of the stellar mass of galaxies ($M_*$)--specific SFR ($\\mbox{SFR}/M_*$) relation. The copulas turned out to be a very useful tool to investigate all these issues, especially for including the complicated selection effects. ",
        "introduction": "A luminosity function (LF) of galaxies is one of the fundamental tools to describe and explore the distribution of luminous matter in the Universe. \\citep[see, e.g.][]{binggeli88,lin96,takeuchi00a,takeuchi00b,blanton01,delapparent03,willmer06}. Up to now, studies on the LFs have been rather restricted to a univariate one, i.e.\\ LFs based on a single selection wavelength band. However, such a situation is drastically changing in the era of large and/or deep Legacy surveys. Indeed, a vast number of recent studies are multiband-oriented: they require data from various wavelengths from the ultraviolet (UV) to the infrared (IR) and radio bands. A bivariate LF (BLF) would be a very convenient tool in such studies. However, to date, it is often defined and used in a confused manner, without careful consideration of complicated selection effects in both bands. This confusion might be partially because of the intrinsically complicated nature of multiband surveys, but also because of the lack of proper recipes to describe a BLF. Then, the situation will be remedied if we have a proper analytic BLF model.  However, it is not a trivial task to determine the corresponding bivariate function from its marginal distributions, if the distribution is not multivariate Gaussian. In fact, there exist infinitely many distributions with the same marginals because the correlation structure is not specified. In general astronomical applications (not only BLFs), for instance, a bivariate distribution is often obtained by either an {\\it ad hoc} or a heuristic manner \\citep[e.g.][]{choloniewski85,chapman03,schafer07}, though these methods are quite well designed in their purposes. Further, analytic bivariate distribution models are often required to interpret the distributions obtained by nonparametric methods \\citep[e.g.][]{cross02,ball06,driver06}]. For such purposes, a general method to construct a bivariate distribution function with pre-defined marginal distributions and correlation coefficient is desired.  In econometrics and mathematical finance, such a function has been commonly used to analyze two covariate random variables. This is called ``copula''. Especially in a bivariate context, copulas are useful to define nonparametric measures of dependence for pairs of random variables \\citep[e.g.][]{trivedi05}. In astrophysics, however, it is only recently that copulas attract researchers' attention and are not very widely known yet \\citep[still only a handful of astrophysical applications:][]{benabed09,jiang09,koen09,scherrer10}. Hence the usefulness and limitations of copulas are still not well understood in the astrophysical community.  In this paper, we first introduce a relatively rigorous definition of a copula. Then, we choose two specific copulas, the Farlie-Gumbel-Morgenstern (FGM) copula and the Gaussian copula to adopt for the construction of a model BLF. Both of them have an ideal property that they are explicitly related to the linear correlation coefficient. Though, as we show in the following, the linear correlation coefficient is not a perfect measure of the dependence of two quantities, this is the most familiar and thus fundamental statistical tool for physical scientists. We focus on the far-infrared (FIR)--far-ultraviolet (FUV) BLF as a concrete example, and discuss its properties and some applications.  This paper is organized as follows: in Section~\\ref{sec:formulation} we define a copula and present its dependence measures. We also introduce two concrete functional forms, the Farlie-Gumbel-Morgenstern copula and the Gaussian copulas. In Section~\\ref{sec:blf}, we make use of these copulas to construct a BLF of galaxies. Especially we emphasize the FIR-FUV BLF. We discuss some implications and further applications in Section~\\ref{sec:discussion}. Section~\\ref{sec:conclusion} is devoted to summary and conclusions. In Appendix~\\ref{sec:jk77}, we show an iterated extension of the FGM copula. We present statistical estimators of the dependence measures of two variables in Appendix~\\ref{sec:nonparam_est} to complete the discussion.  Throughout this paper, we adopt a cosmological model $(h, \\Omega_{\\rm M0}, \\Omega_{\\rm \\Lambda0}) = (0.7, 0.3, 0.7)$ ($h \\equiv H_0/100 [\\mbox{km}\\,\\mbox{s}^{-1}]$) unless otherwise stated. ",
        "conclusions": "\\label{sec:conclusion}  In this work, we introduced an analytic method to construct a bivariate distribution function (DF) with given marginal distributions and correlation coefficient, by making use of a convenient mathematical tool, called a copula. Using this mathematical tool, we presented an application to construct a bivariate LF of galaxies (BLF). Specifically, we focused on the FUV--FIR BLF, since these two luminosities are related to the star formation (SF) activity. Though both the FUV and FIR are related to the SF activity, the marginal univariate LFs have a very different functional form: former is well described by Schechter function whilst the latter has a much more extended power-law like luminous end. We constructed the FUV-FIR BLFs by the FGM and Gaussian copulas with different strength of correlation, and examined their statistical properties. Then, we discussed some further possible applications of the BLF: the problem of a multiband flux-limited sample selection, the construction of the SF rate (SFR) function, and the construction of the stellar mass of galaxies ($M_*$)--specific SFR ($\\mbox{SFR}/M_*$) relation.  We summarize our conclusions as follows: \\begin{enumerate} \\item If the correlation of two variables is weak (Pearson's correlation coefficient $|\\rho| <1/3 $), the Farlie-Gumbel-Morgenstern (FGM) copula provides an intuitive and natural way for constructing such a bivariate DF. \\item When the linear correlation is stronger, the FGM copula becomes inadequate, in which case a Gaussian copula should be preferred. The latter connects two marginals and is directly related to the linear correlation coefficient between two variables. \\item Even if the linear correlation coefficient is the same, the structure of a BLF is different depending on the choice of a copula. Hence, a proper copula should be chosen for each case. \\item The model FIR-FUV BLF was constructed. Since the functional shape of the LF at each wavelength is very different, the obtained BLF has a clear nonlinear structure. This feature was indeed found in actual observational data \\citep[e.g.,][]{martin05}. \\item We formulated the problem of the multiwavelength selection effect by the BLF. This enables us to deal with datasets derived from surveys presenting complex selection functions. \\item We discussed the estimation of the SFR function (SFRF) of galaxies. The copula-based BLF will be a convenient tool to extract detailed information from the observationally estimated SFRF because of its bivariate nature. \\item The stellar mass--specific SFR relation was also discussed. This relation can be reduced to a BLF of luminosities at a mass-related band and a SF-related band. With an analytic BLF model constructed by a copula will provide us with a powerful tool to analyze the downsizing phenomenon with addressing the complicated selection effects. \\end{enumerate}  As the copula becomes better known to the astrophysical community and statisticians develop the copula functions, we envision many more interesting applications in the future. In a series of forthcoming papers, we will present more observationally-oriented applications of copulas."
    },
    "1004/1004.0325_arXiv.txt": {
        "abstract": "We discuss new limits on masses and radii of compact stars and we conclude that they can be interpreted as an indication of the existence of two classes of stars: ``normal'' compact stars and ``ultra-compact'' stars. We estimate the critical mass at which the first configuration collapses into the second. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.0331_arXiv.txt": {
        "abstract": "Results are presented from near-infrared spectroscopic observations of a sample of $BzK$-selected, massive star-forming galaxies (\\sbzks) at $1.5<z<2.3$ that were obtained with OHS/CISCO  at the Subaru telescope and with SINFONI at the VLT.  Among the 28 \\sbzks observed, \\hasp emission was detected in 14 objects, and for 11 of them the [\\ion{N}{2}]$\\lambda6583$ flux was also measured. Multiwavelength photometry was also used to derive stellar masses and extinction parameters, whereas \\hasp and [\\ion{N}{2}] emissions have allowed us to estimate star-formation rates (SFR), metallicities, ionization mechanisms, and dynamical masses. In order to enforce agreement between SFRs from \\hasp with those derived from rest-frame UV and mid-infrared, additional obscuration for the emission lines (that originate in \\ion{H}{2} regions) was required compared to the extinction derived from the slope of the UV continuum.  We have also derived  the stellar mass-metallicity relation, as well as the relation between stellar mass and specific SFR, and compared them to the results in other studies. At a given stellar mass, the \\sbzks appear to have been already enriched to metallicities close to those of local star-forming galaxies of similar mass. The \\sbzks presented here tend to have higher metallicities compared to those of UV-selected galaxies, indicating that near-infrared selected galaxies tend to be a chemically more evolved population. The \\sbzks show specific SFRs that are systematically higher, by up to $\\sim2$ orders of magnitude, compared to those of local galaxies of the same mass. The empirical correlations between stellar mass and metallicity, and stellar mass and specific SFR are then compared with those of evolutionary population synthesis models constructed either with the simple closed-box assumption, or within an infall scenario. Within the assumptions that are built-in such models, it appears that a short timescale for the star-formation ($\\simeq100$~Myr) and large initial gas mass appear to be required if one wants to reproduce both relations simultaneously. ",
        "introduction": "Much action takes place during the $\\sim 2$ billion years of cosmic time between $z=3$ and $z=1.4$: both the global star-formation rate (SFR) and nuclear (AGN) activity peak at $z\\sim 2$, while much of the galaxy assembly and morphological differentiation are taking place and a population of massive, passively evolving galaxies is gradually emerging. Thus a full exploration of this redshift interval is critical for our understanding of galaxy evolution. Unfortunately, this is also a relatively difficult redshift range to penetrate observationally, as strong spectral features such as \\ion{Ca}{2} H\\&K and the 4000 \\AA{} break (for passive galaxies) and [\\ion{O}{2}]$\\lambda3727$, [\\ion{O}{3}]$\\lambda5007$, \\hasp and \\hb\\ (for star forming ones) have all moved to the near-infrared (NIR).  Still, most spectroscopic surveys of galaxies at $1.4\\lesssim z\\lesssim 3$ have used optical spectrographs, relying on much weaker spectroscopic features for redshift determination and galaxy characterization, with these limitations being partly compensated by the possibility to achieve a fairly high multiplex \\citep[e.g.,][]{steidel:2004,mignoli:2005,lefevre:2005,vanzella:2005, lilly:2007,cimatti:2008,popesso:2009}. For passive galaxies at $z>1.4$ such features practically restrict to a set of absorption features at $\\lambda\\lambda \\sim 2600$--$2850$ \\AA{} due to neutral and singly ionized Mg and Fe, and for star-forming galaxies to several weak absorption lines over the rest-frame UV continuum, most of which due to the interstellar medium (ISM)  of these galaxies.  The intrinsic weakness of the absorptions and/or of the continuum made such spectroscopic observations very demanding in terms of telescope time.  Therefore, several studies of large samples of galaxies at $1.4\\lesssim z\\lesssim 2.5$ have relied on color selections and photometric redshifts.  Particularly effective has proven the $BzK$ criterion introduced by \\citet{daddi:2004bzk}, which is able to select both star-forming (called \\sbzks) as well as passively evolving galaxies (\\pbzks) over this redshift range. This has enabled estimates of SFRs, stellar masses, and clustering properties of such galaxies, with samples from $\\sim 100$'s to over $\\sim 30,000$ objects \\citep[e.g.,][]{kong:2006,daddi:2007sfr,dunne:2009,mccracken:2010}. The $BzK$ technique ensures a nearly unbiased selection of $z\\sim 2$ galaxies, including UV-selected galaxies and single color, NIR-selected galaxies \\citep[e.g.,][]{reddy:2005,mccracken:2010}.  Whereas many aspects concerning the evolution of galaxies can be investigated using only photometric redshifts, spectroscopy remains indispensable for a variety of investigations. These include full mapping of the local environment (locating clusters, groups, filaments and voids), refining SFR and mass estimates, measure stellar and ISM metallicities, and finally map the internal dynamical workings of galaxies via 3D spectroscopy. In this respect, the required telescope time is not the only drawback of the optical spectroscopy of galaxies at $1.4\\lesssim z\\lesssim 3$. Indeed, optical spectroscopy down to a limit as faint as $B\\sim 25$ does not recover but a minor fraction of the global SFR and stellar mass at $z\\sim 2$, in particular missing galaxies that are among the most massive and most intensively star-forming ones \\citep{renzini:2009b}. Moreover, high spatial resolution, 3D spectroscopy of high redshift galaxies requires the knowledge of the spectroscopic redshift, to make sure that interesting emission lines (e.g., H$\\alpha$) are free from OH and other atmospheric contaminations \\citep[see e.g.,][]{genzel:2006,forsterschreiber:2009}.  In the case of star-forming (\\sbzk) galaxies, the poor correlation of mass and SFR with $B$ magnitude is a result of high extinction. Therefore, the situation should appear more favorable in the NIR, and not only because moving to longer wavelengths should reduce the impact of extinction, but also because the most active star-forming and most massive galaxies are also among the brightest objects at these longer wavelengths, and one can access strong emission lines such as [\\ion{O}{2}] and H$\\alpha$. Yet, NIR spectroscopy of $z\\gtrsim 1.4$ galaxies is still in its infancy, especially for NIR selected samples. \\citet{erb:2006mz,erb:2006mass,erb:2006sfr} have presented results for over 100 UV-selected galaxies at $z\\sim 2$, deriving SFRs from the strength of H$\\alpha$. NIR spectroscopic observations of samples of $z\\sim 2$ galaxies selected in the NIR have been also presented by \\citet{kriek:2006a,kriek:2006b,kriek:2007,kriek:2008a,kriek:2008b}, focusing mainly (but not exclusively) on passive galaxies by detecting the 4000 \\AA{} break, and deriving spectro-photometric redshifts from it.  NIR spectroscopy of \\sbzk{} galaxies has been carried out by \\citet{hayashi:2009} for a sample of 40 \\sbzks, and detected \\hasp emission from 15 of them.  Their detections, however, are limited to $z<2$.  Finally, integral field NIR spectroscopy for  some 60 star forming galaxies at $z\\sim 2$ has been obtained by \\cite{forsterschreiber:2009},  for partly UV-selected, partly $BzK$-selected targets. Therefore, there is still just scanty spectroscopic information in the rest-frame optical wavelength for actively star-forming and heavily obscured galaxies at $z\\simeq2$, many of which would be missed by the UV-selection, and or are virtually unreachable  by current optical spectroscopy.  In the perspective of improving upon this situation, in 2004 we started NIR spectroscopic observations of $z\\sim 2$ galaxies primarily selected on \\bzk{} technique, and using a variety of NIR instruments, namely OHS, CISCO and MOIRCS, at the Subaru telescope and SINFONI at the VLT.  Our intent was to explore the effectiveness of NIR spectroscopy to improve our characterization of $z\\sim 2$ galaxies using a relatively small pilot sample of them, while assessing the feasibility of wider surveys with future instruments with higher multiplex. In this paper we present the results of observations with the OHS/CISCO instruments on Subaru and SINFONI instrument on VLT of \\sbzk galaxies, leaving the results obtained with the MOIRCS instrument for a future paper. With the observations presented here we have attempted to measure for each galaxy the dust extinction, SFR, ISM metallicity, and dynamical mass, while checking for a possible AGN contribution. Physical quantities are derived assuming the concordance cosmology, i.e., $\\Omega_{\\text{M}}=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_0 = 70$ \\kms~Mpc$^{-1}$ and photometric magnitudes are expressed in the AB system \\citep{oke:1983} if it is not explicitly noted otherwise. For the solar oxygen abundance, we use $12+\\log(\\text{O/H})_\\odot=8.69$ \\citep{allendeprieto:2001}. Emission line width are measured  assuming a Gaussian profile and the FWHM and the velocity dispersion ($\\sigma$) is given by $\\text{FWHM}=2.355\\sigma$. ",
        "conclusions": "We have conducted NIR spectroscopic observations of 28 \\sbzk galaxies at $\\simeq2$ and detected \\hasp emissions from 14 of them. By using the \\hasp and [\\ion{N}{2}] lines, we have derived \\hasp SFRs and gas phase oxygen abundances.  Stellar masses and reddening parameters have been also derived by SED fitting to the multi-wavelength photometric data. Our results are summarized as follows: \\begin{itemize} \\item Stellar masses and reddening parameters derived from SED fitting agree well with those derived by $BzK$-based recipe introduced by \\citet{daddi:2004bzk}. \\item A comparison of SFRs from different indicators (\\hasp, extinction corrected UV, and UV+MIR) indicates that additional extinction towards \\ion{H}{2} regions over that derived from the SED fitting of the stellar continuum is required for the \\hasp SFR to recover the total SFR inferred from UV+MIR. The required additional extinction is in agreement with the recipe proposed by \\citet{cidfernandes:2005}. \\item One object, D3a-11391, shows MIR-excess and a broad \\hasp component with $\\text{FWHM}\\simeq2500$\\kms, as well as a narrow line component, suggesting that it may host an AGN, which would be highly obscured at UV and X-ray wavelengths. Although with large uncertainty, we have estimated the mass of the SMBH and its intrinsic X-ray luminosity following \\citet{alexander:2008}, obtaining $M_{\\text{BH}}\\simeq (3\\text{--}9)\\times10^7\\,M_{\\odot}$, depending on the assumed geometry of the accretion disk, and $L_{2-10\\text{keV}}\\simeq10^{44}$~ergs~s$^{-1}$.  From these quantities, the mass accretion rate onto the central SMBH and Eddington factor are calculated as $\\dot{M}_{\\rm BH}\\simeq0.5\\,M_{\\odot}\\,\\text{yr}^{-1}$ and $\\eta\\simeq0.3$--$0.9$, again depending on the geometry.  These properties are in between those of normal star-forming galaxies and those of broad-line SMGs, indicating a possible evolutionary connection between normal UV/optical-selected galaxies, broad-line/MIR-excess \\sbzks, and star-forming broad-line SMGs that may result from an accelerated growth of SMBH. \\item Most of \\sbzks presented here have already reached a metallicity similar to that of local star-forming galaxies of the same mass. They tend to have higher metallicity compared to UV-selected $z\\sim 2$ galaxies, even at the top mass end. This indicates that \\sbzks are on average a more evolved population compared to that of UV-selected galaxies at the same redshift. \\item Specific SFRs of most \\sbzks are consistent with a tight \\mssfr correlation \\citep{daddi:2007sfr, pannella:2009}, with a couple of outliers reaching very high SSFRs, similar to those of  SMGs \\item Within the framework of \\pegase closed-box or infall models, the \\mz and \\mssfr relations of $z\\simeq2$ \\sbzks can be reproduced simultaneously assuming a very short star-formation or infall timescale, $\\tau\\simeq100$~Myr, and large gas mass at the beginning. However, the implied SFRs averaged over the life of the galaxies appears to be excessively large, suggesting that secularly declining SFRs may not be appropriate for star-forming galaxies at $z\\sim 2$. \\end{itemize}"
    },
    "1004/1004.0936_arXiv.txt": {
        "abstract": "Solar supergranulation plays an important role in generating and structuring the solar magnetic field and as a mechanism responsible for the 11-year solar cycle. It is clearly detected within SOHO/MDI Dopplergrams, from which a variety of properties may be derived. Techniques that extract spatial, temporal and kinematic characteristics and provide comparisons for the two most recent solar minima are described. Although supergranule lifetimes are comparable between these minima, their sizes maybe slightly smaller during the recent minimum. ",
        "introduction": "\\label{sec:Intro}  Supergranules are the largest scale component of solar convection readily seen with present observing techniques. They are typically $\\sim$30 Mm across and have been observed to live from anywhere between 1--2 days \\citep{duv80}. Their relationship with local magnetic fields have been seen via diffusion studies and spatial observations of both magnetograms and CaIIK chromospheric data.  Since its inception, SOHO/MDI \\citep{sch95} Dynamics Runs have provided yearly 60-day sets of 1-minute Dopplergrams. Analysis of the 1996 data has quantified a wide variety of supergranule characteristics \\citep[for example,][]{hat00, bec00}.  The most recent solar minimum has sparked much interest due to its lengthy minimum. Many studies are currently underway to understand the causes and consequences of this recent minimum.  This paper extends the analysis of supergranule characteristics to 2008 data and the results are compared to similar studies performed on 1996 data. ",
        "conclusions": "\\label{sec:Analysis}  While the 1/$e$ lifetimes are found to be the same for both years, the supergranule cell sizes tend to be slightly smaller during 2008. A significant discrepancy, however, is found within the velocity analysis.  Although the 1996 results are in line with those found by \\citet{hat02}, the 2008 velocity values are considerably greater. Image defocusing during 1996 \\citep{kor04} may influence the velocity values contained within the Dopplergrams by altering the resolution of the instrument. Active regions present on the disk may contribute to spurious velocity values within the calculation \\citep{liu01}. Any images displaying significant activity are currently removed from the analysis, whereas future analyses will use magnetograms to mask out these regions so the images may be included.  Comparing the past two solar minima, the decay rate and decay-times of the supergranule pattern have not changed. Missing images within a day's data currently means removing that day from the analysis, which reduces the sampling for the averaging. Future work will update the analysis process so that loss of data within a day will not mean that the whole day needs to be skipped.  The size analysis is much less problematic although it can be improved using a statistical analysis of the extracted parameters for each Dopplergram-derived spectrum. This is planned for the future along with producing similar statistics for the lifetimes and velocity analyses."
    },
    "1004/1004.2428_arXiv.txt": {
        "abstract": "We consider the Riemann problem for  relativistic flows of polytropic fluids  and find relations for the flow characteristics. Evolution of physical quantities take especially simple form for the case of   cold magnetized plasmas.  We    find exact, explicit  analytical  solutions for one dimensional  expansion of magnetized  plasma into vacuum, valid for   arbitrary magnetization.   We also consider  expansion into  cold unmagnetized  external medium both for stationary initial conditions and for initially moving plasma, as well as reflection of rarefaction wave from a wall. We also find self-similar structure of three-dimensional  magnetized  outflows into vacuum, valid close to the plasma-vacuum interface.  The key results of this work, the self-similar solutions,  were incorporated  post-initial submission   into   appendices of the published version of  Granot \\etal\\  (2010). ",
        "introduction": "Relativistic shock waves are common in different physical systems \\citep{2003LRR.....6....7M}, from heavy ion nuclear collision \\citep[\\eg][]{2009PhRvL.103c2301B} to astrophysical shocks in pulsar winds \\citep{kc84},  Active Galactic Nuclei \\citep[\\eg][]{Krolik:1999} and Gamma Ray Bursts \\citep{PiranReview}. Many modern computational algorithms are based on the  solution of Riemann problems \\citep[\\eg][]{Toro}. These algorithms are based on Godunov-type shock-capturing schemes  and do not require large artificial viscosity or smoothing operators.    Analytical solutions to the corresponding Riemann problems are then important for code testing.  Exact, explicit  non-linear solutions of relativistic fluid equations, and especially  relativistic MHD equations, are rare.  In a general form the relativistic Riemann problem was solved by \\cite{1994JFM...258..317M,Romero05}, who find the solutions for Riemann invariants and for the characteristics in quadratures. In this paper we find simple expressions for the characteristics of Riemann simple waves, and in particular for the astrophysical important case of dynamics of a cold, relativistically magnetized plasma.  Our  results can be used for  benchmark estimates of the overall dynamical behavior in   numerical simulations of relativistic flows and  strongly magnetized outflows in particular. ",
        "conclusions": "In this paper we  found exact  explicit solutions for one-dimensional relativistic   expansion of polytropic fluid into vacuum and into plasma. In particular, we discussed an astrophysical important case of strongly magnetized   outflows; in this case especially  simple analytical solutions can be obtained. We found exact solutions for one-dimensional  expansion of magnetized plasma into vacuum and into the cold medium both for stationary initial conditions and for a piston moving toward the interface. We found  exact  relations,  applicable for arbitrary magnetization, relativistic motion and  external densities. These  results can be used for  benchmark estimates of the overall dynamical behavior for the  numerical simulations of relativistic plasmas, \\eg\\ in heavy ion collisions, and in  strongly magnetized outflows in particular.    I am greatly  thankful to  Dimitros Gianios, Sergey Komisarov and Alexandre Tchekhovskoy."
    },
    "1004/1004.1154.txt": {
        "abstract": "Energetic electrons of up to tens of MeV are created during explosive phenomena in the solar corona. %Because it is now established that magnetic energies can be quickly released by magnetic reconnection, While many theoretical models consider magnetic reconnection as a possible way of generating energetic electrons, the precise roles of magnetic reconnection during acceleration and heating of electrons still remain unclear. Here we show from 2D particle-in-cell simulations that coalescence of magnetic islands that naturally form as a consequence of tearing mode instability and associated magnetic reconnection leads to efficient energization of electrons. The key process is the secondary magnetic reconnection at the merging points, or the `anti-reconnection', which is, in a sense, driven by the converging outflows from the initial magnetic reconnection regions.  By following the trajectories of the most energetic electrons, we found a variety of different acceleration mechanisms but the energization at the anti-reconnection is found to be the most important process. We discuss possible applications to the energetic electrons observed in the solar flares. We anticipate our results to be a starting point for more sophisticated models of particle acceleration during the explosive energy release phenomena. ",
        "introduction": "A solar flare is an explosive energy release phenomena on the sun %A large flare can release energies up to 10$^{33}$ergs in $\\sim$10$^3$s. Magnetic reconnection has been widely considered as a possible mechanism of the energy release process but a direct evidence of magnetic reconnection during the solar flares remains to be seen. Note that and a large fraction of the released energy appears to go to high energy, often non-thermal, particles both ions and electrons \\citep[][and references therein]{lin03}. The particle energy reaches tens of GeV for ions and tens of MeV for electrons. The mechanism of producing such energetic particles is much less understood compared to the energy release mechanism. Because magnetic reconnection is a possible mechanism of the energy release process, particle acceleration may also occur through magnetic reconnection, although difficulties remain when trying to interpret observations \\citep[][and references therein]{miller97, krucker08, benz08}. In order to explain observations, many theoretical ideas have been proposed \\citep[][and references therein]{aschwanden02}. While some theories utilize fast/slow mode shocks as well as electromagnetic waves of various scales - sometimes in a stochastic manner - many theories consider magnetic reconnection.  In this paper, we also assume {\\it a priori} that magnetic reconnection plays a role for particle acceleration and explore possible mechanism(s) of particle acceleration in association with magnetic reconnection.  In general, a test particle approach had been used to explore particle acceleration by magnetic reconnection. While some studies solved particle motion analytically \\citep[e.g.][]{litvinenko96}, a majority of studies performed test-particle simulations under model electromagnetic fields \\citep[e.g.][]{kliem94, hannah06, onofri06} or time varying fields generated by MHD simulations \\citep{sato82,scholer87, ambrosiano88}. A self-consistent, particle-in-cell (PIC) simulations are now widely used to study the detailed process of electron energization \\citep[e.g.][]{hoshino01,hoshino05, drake05, pritchett06, drake06, karlicky07, pritchett08, wan08, shinohara09}. PIC simulations have an advantage of being able to resolve the inner structure of the X-line, the so-called diffusion region. It is a scientific challenge to understand particle acceleration by magnetic reconnection that involves multi-scale.  For convenience, we categorize various theories of particle acceleration associated with magnetic reconnection into two different groups: {\\it X-type} acceleration and {\\it O-type} acceleration. The {\\it X-type} acceleration takes place at and around the X-line of magnetic reconnection or the diffusion region. Particles are unmagnetized at the X-lines and can be directly accelerated by the electric field \\cite[e.g.][]{sato82}. In the immediate downstream of the X-line are the regions with magnetic field gradient where particles further gain energy while drifting along the current sheet \\citep[e.g.][]{scholer87,kliem94,hoshino01}. %Effects of slow shocks associated with magnetic reconnection \\citep{sato82} as well as the magnetic field perpendicular to the reconnecting fields, or the guide field \\citep[e.g.][]{litvinenko96}, have also been taken into account. More recently, it was found that the in-plane, polarization electric field in the diffusion region generated by the charge separation between ions and electrons plays an important role \\citep{hoshino05}. The force by the polarization electric field can be balanced by the Lorentz force so that electrons are accelerated efficiently by the reconnection electric field while being trapped within the current sheet boundary.  Because of the trapping effect, the process was named `surfing' mechanism.  %In the mean time, a multi-scale structure of the diffusion region has been revealed. The out-of-plane current is localized only at the center of the diffusion region so that the reconnection rate remains fast while electrons form a high-velocity jet that extends large distances downstream from the diffusion region center \\citep{daughton06, fujimoto06, shay07, karimabadi07, phan07}. The region with the localized current is termed `inner diffusion region' whereas the rest of the diffusion region including the elongated electron current layer is termed `outer diffusion region'. However, it has not been made clear how this multi-scale structure affects the process of electron acceleration.  %The relative importance of the surfing mechanism also remains unclear. It is worth mentioning that magnetic reconnection with the out-of-plane component of the inital magnetic field or the guide field also produces energetic electrons \\citep{wan08}.  The {\\it O-type} acceleration takes advantage of the closed geometry of field lines in a magnetic island. In many cases, a magnetic island is bounded by two X-lines at each end. Therefore, if particles are trapped in a magnetic island, they can continue gaining energy by repetitive crossings of the gradient region, although the reconnection electric field is relatively weak inside magnetic islands \\citep{stern79,scholer87,kliem94}. In order to compensate for this weak electric field, a recent model has been developed that takes into account the dynamical, contracting motion of islands \\citep{drake06}. The time-dependent model is analogous to the energy increase of a ball reflecting between two converging walls, namely the first order `Fermi' process.  %It is important to note that the magnetic field component perpendicular to the reconnecting magnetic fields, or the guide field, may play a significant role for accelerating electrons. In fact, effects of the guide fields have been studied in both {\\it X-Type} and {\\it O-Type} acceleration scenarios. If a guide field exists, electrons tend to be more magnetized in the vicinity of X-lines and they may be energized by the reconnection electric field for relatively long time \\citep{litvinenko96, saito04, wan08}. The guide field also plays an important role in a magnetic island. It was found that density cavities generated along island separatrices support parallel electric fields that act as plasma accelerators \\citep{drake05}.  \\begin{figure} \\epsscale{1.0} \\plotone{schematic.eps} \\caption{Schematic illustration of multi-island coalescence. The thin cross marks indicate the X-lines of normal magnetic reconnection. The thick cross mark indicates the X-line of the anti-magnet-reconnection generated by the coalescence. The arrows indicate the flow directions.\\label{fig:schematic}} \\end{figure}  Despite the intensive research of magnetic reconnection, it still remains unknown which of the two different type of acceleration is important for producing energetic electrons. In this respect, multi-island coalescence has drawn considerable attentions because it potentially contains both {\\it X-Type} and {\\it O-Type} mechanisms. A coalescence is the process of merging of two magnetic islands and has been studied by both theories and simulations \\citep{finn77, pritchett79, biskamp80, tajima87, pritchett07, wan08}. Figure 1 shows a schematic illustration. In general, multiple number of X-lines and localized currents are generated from the tearing-mode instability (Figure \\ref{fig:schematic}a). If we assume an X-line at the center was relatively weak compared to the other two X-lines, the two localized currents bounding the central X-line will eventually be attracted to each other by the Lorentz force (Figure \\ref{fig:schematic}b). Each current is represented by magnetic islands and at the merging point of the two islands, a secondary magnetic reconnection or `anti-reconnection' \\citep{pritchett08} occurs (Figure \\ref{fig:schematic}c). The direction of the electric field of the anti-reconnection is reversed from the direction of the electric field of the primary reconnection. Finally, the two magnetic islands become one large magnetic island.  While the first PIC simulation of particle acceleration during multi-island coalescence was performed more than decades ago \\citep{tajima87}, a detailed study came out quite recently. %In an electron-positron system, the anti-reconnection was observed and this anti-reconnection even produced magnetic islands by the tearing instability \\citep{daughton07}. In an electron-proton system, %\\cite{saito04} focused on only two isolated magnetic islands. Each island was penetrated by strong guide fields of the same (or opposite) direction so that it would have `cohelicity' (or `counter-helicity'). In either case, they observed efficient electron acceleration at the merging point. \\cite{pritchett08} showed that electrons are energized as the number of magnetic islands is reduced by coalescence. %Dependencies on inflow plasma density as well as the guide field have also been studied. An important conclusion of the study is that the reversed electric field decelerates electrons near the anti-reconnection site. It was suggested that the main energization occurs through the {\\it O-Type} mechanism.  In this Paper, we extend the work by \\cite{pritchett08} by performing 2D PIC simulations of multi-island coalescence with no guide field. The key and rather surprising finding of our simulations is that the anti-reconnection plays an important role in accelerating electrons. %By analyzing the trajectories of energetic electrons, we found that the anti-reconnection does accelerate electrons. We have identified various other energization processes, but the most common and also efficient process was the direct acceleration by the electric field associated with the anti-reconnection. Very recently, \\cite{tanaka10} also reported intense electron energization by the anti-reconnection, but our work  provides new insights from different perspectives on this issue because we fully analyzed the trajectories of accelerating electrons and clarified the importance of the anti-reconnection with respect to the other mechanisms.  The outline of the paper is as follows. In section 2, we describe the setup of the simulation runs and the overview of results. Section 3 presents analysis of electron energy spectra. Section 4 presents the trajectory analysis of energetic electrons and describe the details of their energization processes. In section 5, we summarize the results and discuss applications to the observations of the solar flares. Finally, section 6 contains the conclusion. ",
        "conclusions": "We performed 2D PIC simulation to study electron acceleration during multi-island coalescence. By analyzing the trajectories of the most energetic electrons, we found a variety of different acceleration mechanisms such as the contracting island mechanism, ripple mechanism, island surfing mechanism, etc. However, a statistical study showed that the most important process is the energization process that takes place at the anti-reconnection region. Based on the maximum energy of electrons attained in the simulation,  we pointed out that the multi-island coalescence may play an important role in producing energetic electrons during the solar flares.      %% In this section, we use  the \\subsection command to set off %% a subsection.  \\footnote is used to insert a footnote to the text.  %% Observe the use of the LaTeX \\label %% command after the \\subsection to give a symbolic KEY to the %% subsection for cross-referencing in a \\ref command. %% You can use LaTeX's \\ref and \\label commands to keep track of %% cross-references to sections, equations, tables, and figures. %% That way, if you change the order of any elements, LaTeX will %% automatically renumber them.  %% This section also includes several of the displayed math environments %% mentioned in the Author Guide. % %"
    },
    "1004/1004.2791_arXiv.txt": {
        "abstract": "Recent observations of luminous Type IIn supernovae (SNe) provide compelling evidence that massive circumstellar shells surround their progenitors.  In this paper we investigate how the properties of such shells influence the SN lightcurve by conducting numerical simulations of the interaction between an expanding SN and a circumstellar shell ejected a few years prior to core collapse.  Our parameter study explores how the emergent luminosity depends on a range of circumstellar shell masses, velocities, geometries, and wind mass-loss rates, as well as variations in the SN mass and energy.  We find that the shell mass is the most important parameter, in the sense that higher shell masses (or higher ratios of M$_{shell}$/M$_{SN}$) lead to higher peak luminosities and higher efficiencies in converting shock energy into visual light.  Lower mass shells can also cause high peak luminosities if the shell is slow or if the SN ejecta are very fast, but only for a short time. Sustaining a high luminosity for durations of more than 100 d requires massive circumstellar shells of order 10~$M_{\\odot}$ or more.  This reaffirms previous comparisons between pre-SN shells and shells produced by giant eruptions of luminous blue variables (LBVs), although the physical mechanism responsible for these outbursts remains uncertain.  The lightcurve shape and observed shell velocity can help diagnose the approximate size and density of the circumstellar shell, and it may be possible to distinguish between spherical and bipolar shells with multiwavelength lightcurves.  These models are merely illustrative.  One can, of course, achieve even higher luminosities and longer duration light curves from interaction by increasing the explosion energy and shell mass beyond values adopted here. ",
        "introduction": "The luminosity of a supernova (SN) results from energy input by a combination of radioactive decay and shock kinetic energy \\citep[see e.g.,][]{A96}, and for a Type II SN, the shape of the light curve depends on quantities like the star's initial radius, ejecta mass, and explosion energy \\citep{A96, Y04, KW09}.  For SNe with small initial radii, like SNe of Types Ia, Ib, Ic, and peculiar SNe~II like SN~1987A that result from blue supergiants, most of the shock-deposited thermal energy imparted to the stellar envelope is converted back into kinetic energy through adiabatic expansion, so nearly all of the observed luminosity comes from the radioactive decay of $^{56}$Ni and $^{56}$Co.  In ``normal'' SNe~II-P that result from the explosions of red supergiants (RSGs), however, the large initial radius allows some modest fraction (typically 1--2\\%) of the shock-deposited thermal energy to be radiated away, powering much of the plateau of the lightcurve, although the vast majority still goes into expansion energy.  At late times, even SNe~II-P have their luminosity powered by radioactive decay (e.g., \\citealt{H03}).  Subsequently, however, as the fast SN ejecta expand, they can collide with dense circumstellar or interstellar material (CSM/ISM) that may surround the SN.  As a result, additional kinetic energy may be transformed once again back into thermal energy through shock heating, which in turn may be lost by radiative cooling if a dense radiative shock forms (e.g., \\citealt{CF08}).  This can enhance the luminosity for long after the explosion: Som SNe remain radio luminous for decades \\citep{M98, W02, V93}, and this interaction may power a visible supernova remnant (SNR) such as Cas~A for hundreds of years \\citep{C77, CO03}.  On the other hand, if the collision with dense CSM happens immediately after the explosion, it may significantly alter the spectrum and light curve of the SN itself.  This latter scenario is generally thought to be the case for the observed sub-class of Type~IIn supernovae \\citep{S90,F97}, where the ``n'' corresponds to ``narrow'' or intermediate-width H lines from the shock-heated CSM gas or decelerated SN ejecta (e.g. \\citealt{CD94,C01}).  In a normal SN, the expected results of radiative cooling and reheating of the SN ejecta due to radioactive decay yield can be estimated from analytical models of stellar structure and explosion physics \\citep{MM99}.  In SNe with strong CSM interaction such as the observed class of Type~IIn SNe, however, the effects of collisions between an expanding SN and its circumstellar gas are harder to predict with {\\it ab initio} calculations.  They depend highly on the density and morphology of the CSM, which in turn depend on the unknown mass-loss behavior of the star in the few years prior to core collapse --- potentially different for each object.  A wide variety of CSM environments are possible, leading to a wide diversity of observed lightcurves and spectral properties.  Recent observations of luminous Type~IIn supernovae such as SN~2006gy \\citep{S07,O07} and SN~2006tf \\citep{S08a} have stretched the boundaries of our understanding of SNe~IIn.  Their extreme luminosities yield strong evidence that the progenitors of these SNe were surrounded by massive shells, presumably ejected in precursor eruptions during the final years of stellar evolution \\citep{S07, S08a, S09b, SM07, WBH07}.  \\citet{S07} pointed out that the physical properties (mass, speed, H composition) of these mass ejections were analogous to those observed for giant eruptions of luminous blue variables (LBVs), and especially reminiscent of the giant 1843 eruption of $\\eta$ Carinae \\citep{S03}.  As the SN ejecta expand, they collide with the recently ejected CSM shell and this collision significantly decelerates the SN expansion, transforming kinetic energy back into thermal energy at the collision front, producing a brilliant fireworks display.  The remarkably high luminosity and long duration of the observed emission from SNe~2006gy and 2006tf imply that the circumstellar shells were very massive --- of order 10--20 $M_{\\odot}$ --- in order to sufficiently decelerate the SN blast wave and tap into its available reservoir of kinetic energy \\citep{S07,S08a,S09b,SM07,WBH07}.  \\citet{SM07} have argued based on a simplified analytical model, similar to that of \\citet{FA77}, that the high luminosity and long duration of these SNe can be explained by a SN colliding with a very massive and initially opaque CSM shell.  We explore this idea here in more detail with a variety of possible CSM environments using numerical simulations.  We suggest that the presence and shape of circumstellar shells can be a powerful tool to constrain the evolution of the progenitors of Type~IIn supernovae.  We investigate how the mass, speed, and morphology of such shells can influence the evolution of a SN lightcurve.  We undertake a parameter study of SNe with different masses interacting with a selection of possible circumstellar shells, both spherical and bipolar.  From these simulations we calculate thermal emission profiles and compare them in order to constrain how the physical properties of circumstellar nebulae can influence the SN lightcurve, and to constrain the efficiency of converting kinetic energy to light.  Our calculations are simplified in the way we treat the cooling of and radiation from the shocked gas, which we approximate as optically thin radiative cooling; by necessity; our application of these results is therefore limited in scope.  An important point to note is that our approach is to simulate a variety of hypothetical SNe to demonstrate {\\it trends} in how the lightcurve responds to changes in SN and shell properties.  We are not attempting a quantitative fit to the observed data for any individual SN.  This has been pursued for a few relatively nearby and well-observed SNe IIn, such as SN~1988Z \\citep{T93,CD94,Aetal99}, SN~1994W \\citep{Cetal04}, and SN~1998S \\citep{C01}, where the CSM properties were derived from fitting the observed light curves and spectra.  Those authors inferred massive precursor shell ejections in the few years before core collapse, although the energy demands and required shell masses for these were not as extreme as for SNe~2006tf and 2006gy.  Our work here builds upon these earlier studies.  \\smallskip  We explain our adopted initial conditions and the numerical method in \\S 2 and \\S 3, respectively, and in \\S 4 we discuss some details of the shock interaction.  In \\S 5 we discuss how the resulting light curves depend on various parameters and in \\S 6 we discuss shell velocities, and how these may help to interpret observations. Finally, in \\S 7 we interpret our results in context with the most luminous SNe IIn, and in \\S8 we provide a summary.  We include electronic datafiles containing the results of our simulations with this paper.  The L\\_....dat files contain the total luminosity [erg/s] as a function of time [s].  The V\\_....dat files contain both the volume averaged and mass averaged velocity of the shocked gas [cm/s] as a function of time [s].  A small sample of these tables is provided in Appendices \\ref{sec-lum_table} and \\ref{sec-vel_table}.  \\begin{table*} \\label{tab:sim} \\begin{centering} \\caption{Simulation input parameters} \\begin{tabular}{p{0.1\\linewidth}ccccccccl} \\hline \\noalign{\\smallskip} Name & $M_{\\rm sn}$   & $E_{\\rm sn}$  & $M_{\\rm shell}$  & $V(\\theta=0)$ & $\\dot{M}_{\\rm wind}$ & $\\Omega$ & t$_{\\rm end}$ & $dE/E_{sn}$ & $v_{final}^{1}$ [km/s] \\\\ & [$\\mso$]     &  $10^{51}$~ergs&  [$\\mso$]        & [$\\kms$]      & [$\\msoy$]            & $\\Omega$ &  [yr pre-SN]  &  \\% &  $10^3$ [km/s]\\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} O01  & 30 & 1 & N/A &200 &  $10^{-4}$ &  0.0 & N/A & 0.05 & 4.39 \\\\ O02  & 30 & 1 & N/A &200 &  $10^{-3}$ &  0.0 & N/A & 0.3  & 3.42 \\\\ O03  & 30 & 1 & N/A &200 &  $10^{-2}$ &  0.0 & N/A & 1.5  & 2.57 \\\\ O04  & 30 & 1 & N/A & 50 &  $10^{-4}$ &  0.0 & N/A & 0.108 & 3.75 \\\\ \\hline A00  & 30 & 1  &0.1 & 200 &  $10^{-4}$ &  0.0 & 2  & 0.8  & 2.85 \\\\ A01  & 30 & 1  & 1  & 200 &  $10^{-4}$ &  0.0 & 2  & 5.05 & 2.34 \\\\ A02  & 30 & 1  & 6  & 200 &  $10^{-4}$ &  0.0 & 2  & 18.7 & 1.79 \\\\ A03  & 30 & 1  & 10 & 200 &  $10^{-4}$ &  0.0 & 2  & 25.5 & 1.59 \\\\ A04  & 30 & 1  & 20 & 200 &  $10^{-4}$ &  0.0 & 2  & 36.5 & 1.27 \\\\ A05  & 30 & 1  & 10 & 200 &  $10^{-3}$ &  0.0 & 2  & 25.6 & 1.51 \\\\ A06  & 30 & 1  & 10 & 200 &  $10^{-5}$ &  0.0 & 2  & 25.3 & 1.61 \\\\ A07  & 30 & 1  & 10 &  50 &  $10^{-3}$ &  0.0 & 2  & 31.5 & 1.30 \\\\ A08  & 30 & 1  & 10 & 500 &  $10^{-3}$ &  0.0 & 2  & 16.9 & 1.66 (at500 days)\\\\ A09  & 30 & 1  & 10 &  50 &  $10^{-4}$ &  0.0 & 2  & 31.6 & 1.36 \\\\ A10  & 30 & 1  & 10 & 500 &  $10^{-4}$ &  0.0 & 2  & 16.3 & 1.80 (at500 days)\\\\ A11  & 30 & 1  & 10 &  50 &  $10^{-5}$ &  0.0 & 2  & 31.6 & 1.46 \\\\ A12  & 30 & 1  & 10 & 500 &  $10^{-5}$ &  0.0 & 2  & 16.3 & 1.80 (at 500 days)\\\\ \\hline B01  & 10 & 1  & 10 &  50 &  $10^{-4}$ &  0.0 & 2 & 54.5  & 1.52 \\\\ B02  & 10 & 1  & 10 & 200 &  $10^{-4}$ &  0.0 & 2 & 48.2  & 1.83 \\\\ B03  & 10 & 1  & 10 & 500 &  $10^{-4}$ &  0.0 & 2 & 37.0  & 2.12 (at 500 days) \\\\ B04  & 10 & 1  & 25 & 200 &  $10^{-4}$ &  0.0 & 2 ($\\Delta t=5$yr) & 65.1 &  1.08 (at 500 days) \\\\ % \\hline C01  & 60 & 1  & 10 &  50 &  $10^{-4}$ &  0.0 & 2 & 19.7 & 1.18 \\\\ C02  & 60 & 1  & 10 & 200 &  $10^{-4}$ &  0.0 & 2 & 14.5 & 1.30 \\\\ C03  & 60 & 1  & 10 & 500 &  $10^{-4}$ &  0.0 & 2 & 7.55 & 1.46 \\\\ \\hline D01  & 10 & 1  & 10 & 500 &  $10^{-4}$ &  0.9 & 2 & 42.1 &  N/A \\\\ D02  & 30 & 1  & 10 & 500 &  $10^{-4}$ &  0.9 & 2 & 20.4 &  N/A \\\\ D03  & 60 & 1  & 10 & 500 &  $10^{-4}$ &  0.9 & 2 & 10.9 &  N/A \\\\ \\hline E01  & 30 & 1  & 10 & 500 &  $10^{-4}$ &  0.0 & 4  & 14.7 & 1.79 (at 500 days) \\\\ E02  & 30 & 1  & 10 & 200 &  $10^{-4}$ &  0.0 & 10 & 24.3 & 1.62 (at 500 days) \\\\ E03  & 30 & 1  & 10 & 500 &  $10^{-4}$ &  0.0 & 10 & 13.9 & 1.79 (at 1000 days) \\\\ \\hline F01  & 30 & 0.5 & 10 & 200 &  $10^{-4}$ &  0.0 & 2 & 22.3 & 1.13 \\\\ F02  & 30 & 2   & 10 & 200 &  $10^{-4}$ &  0.0 & 2 & 27.4 & 2.19 \\\\ \\hline G01  &  6 & 1   &  6 & 200 &  $10^{-5}$ &  0.0 & 2 & 56.4 & 2.07 (at 100 days)\\\\ \\hline H01  &  1 & 1   &  1 & 200 &  $10^{-5}$ &  0.0 & 2 & 42.7 & 4.77 (at 100 days)\\\\ \\hline \\noalign{\\smallskip} \\end{tabular} \\end{centering} $(1)$ Measured at 250 days unless indicated otherwise. \\\\ \\end{table*} ",
        "conclusions": "\\label{sec-conclusion}  In this paper we describe the influence of massive circumstellar shells on core-collapse SN lightcurves, with the primary motivation of trying to understand the power source of extremely luminous Type~IIn events and their relation to the diverse population of SNe~IIn.  We show how these circumstellar shells can indeed create extreme peaks in the luminosity such as have been observed in Type~IIn supernovae like SN2006gy and SN2006tf.  The luminosity of these SNe would require extreme amounts of $^{56}$Ni if they are powered by radioactive decay, but if interactions in the CSM provide the power instead, then the shell masses and speeds that are required have reasonable precedent from observed properties of spatially resolved shells around nearby massive stars (see \\citealt{SO6} and references therein).  Our investigation is by no means exhaustive.  Pre-SN circumstellar shells may have a wide range of masses, expansion speeds, and radii, whereas we have adopted simplified shell geometries for illustrative examples.  Additionally, the underlying SNe ejecta may have wide diversity in explosion energy, mass, and ejecta speed.  In this preliminary investigation, our approach has been to vary each of these parameters individually to illustrate their influence on the light curve rather than attempting to accurately model any individual SN. We have attempted to find general ways to distinguish between different kinds of shells, using trends in the observed shapes of the lightcurves, their characteristic emission temperature, and observed shock speeds.  We find that observations of the evolution of the shock speed is necessary to help break the degeneracy in the other free parameters, while observations of the speed of the pre-shock CSM help considerably as well (see e.g., \\citealt{S02,S07,S08a,S09a,S09b,Tetal08}).  This can be used to analyze the mass-loss history of massive stars in the last years prior to the explosion, which can be a powerful tool for studying the final stages of stellar evolution.  Ultimately, we wish to know the physical origin of these SN-precursor events.  The key result is that we confirm the large masses of circumstellar shells hypothesized to account for some recent luminous SNe~IIn \\citep{SM07,S07,S08a,S09b,WBH07}, as well as the high mass and explosion energy of the underlying SNe.  One can also produce a very high peak luminosity with lower mass if the shell is slow and the SN ejecta are fast, but a lower mass shell cannot yield both a high peak luminosity and a long duration of $\\ga$100 days seen in some luminous SNe~IIn. In fact, we suspect that even larger shell masses or larger explosion energies are needed to account for the observed light curves of the most luminous SNe~IIn. Thus, more detailed attempts to model individual objects will be the focus of a second paper in this series.  \\subsection{Future developements}  Further research is required for quantitative analyses of observed SN~IIn lightcurves and to extract reliable absolute values of shell masses and SN explosion energies.  This must include adding the luminosity contribution from the underlying SN photosphere (powered by diffusion or radioactive decay) in cases where the CSM interaction luminosity is not extremely high compared to the ejecta photosphere, as well as using an improved treatment of post-shock cooling and radiative transfer at high optical depths in order to more accurately model the emergent radiation from the post-shock shells in these simulations.  As noted by \\citet{SM07} and \\citet{S08a, S09b}, it is likely that the CSM will be highly opaque, especially at the earliest phases, so the effects of radiative diffusion should be taken into account to properly model the emergent luminosity.  Finally, all our simulations have adopted a Type II-P core-collapse SN density profile, but other types of SNe with different density profiles need to be investigated in a similar manner, since any type of SN can, in principle, be a Type~IIn event if it runs into a dense H-rich environment."
    },
    "1004/1004.0088_arXiv.txt": {
        "abstract": "{} {We use the growing data sets of very-metal-poor  stars to study the impact of stellar winds of fast rotating massive stars on the chemical enrichment of the early Galaxy.}  {We use an inhomogeneous chemical evolution model for the Galactic halo to predict both the mean trend and scatter of C/O and N/O. In one set of models, we assume that massive stars enrich the interstellar medium during both the stellar wind and supernovae phases. In the second set, we consider that in the earliest phases (Z $<10^{-8}$), stars with masses above 40 $\\mathrm{M}_{\\odot}$ only enrich the interstellar medium via stellar winds, collapsing directly into black holes.} {We predict a larger scatter in the C/O and N/O ratios at low metallicities when allowing the more massive fast-rotating stars to contribute to the chemical enrichment  only via stellar winds. The latter assumption, combined with the stochasticity in the star formation process in the primordial Galactic halo can explain the wide spread observed in the N/O and C/O ratios in normal very-metal-poor stars.} {For chemical elements with stellar yields that depend strongly on initial mass (and rotation) such as C, N, and neutron capture elements, within the range of massive stars, a large scatter is expected once the stochastic enrichment of the early interstellar medium is taken into account. We also find that stellar winds of fast rotators mixed with interstellar medium gas are not enough to explain the large CNO enhancements found in most of the carbon-enhanced very-metal-poor stars. In particular, this is the case of the most metal-poor star known to date, HE~1327$-$2326, for which our models predict lower N enhancements than observed when assuming a mixture of stellar winds and interstellar medium. We suggest that these carbon-enhanced very metal-poor stars were formed from almost pure stellar wind material, without dilution with the pristine interstellar medium.} ",
        "introduction": "The very metal-poor stars of the halo play a fundamental role in chemical evolution since, at metallicities below [Fe/H] $\\approx -$2.5, only type II supernovae have had time to contribute to the interstellar medium (ISM) enrichment from which these stars formed, thus offering a way of constraining the nucleosynthesis in massive stars at low metallicities (e.g. Chiappini et al. 2005, Fran\\c cois et al. 2004).  At present, several thousand of very metal-poor stars (hereinafter VMP), i.e. stars with metallicities below [Fe/H]$=-$2. is known (e.g. Cayrel et al. 2004, Christlieb \\& Beers 2005; Masseron et al. 2009). The samples of VMP stars are expected to increase by about one order of magnitude in the next years thanks to surveys such as SDSS/SEGUE-2 and LAMOST (see Beers 2010). Larger samples of VMP stars with high-quality abundances measurements will play a fundamental role in constraining the stellar nucleosynthesis of the first generations of massive stars.  To date, around 20\\% of the VMP stars observed show unexpected large C and N enhancements with respect to solar. These stars are called carbon-enhanced metal-poor stars (CEMP). The VMP stars not showing such large C and N enhancements are called {\\it normal} VMP stars (Cayrel et al. 2004, Spite et al. 2005, 2006). In the case of the two most iron-poor stars known to date, these overabundances can reach 100 to 10000 times the ratios found in the Sun (Frebel et al. 2005, 2008, Christlieb et al. 2002, Norris et al. 2007).  The CEMP stars are classified according to the presence or absence of s- and r-process elements. The peculiar abundance of those showing over-abundances of s-process elements are usually interpreted as caused by accretion from an asymptotic giant branch (AGB) companion (e.g. Masseron et al 2009 and references therein). However, those without s-process element enhancements were most probably formed from enriched gas ejected by earlier generations of massive stars (CEMP-no).  The study of the chemical enrichment of the pristine Universe in CNO, the most abundant metals, is of particular interest. The CNO chemical enrichment could have had a important impact in shaping the early IMF as well as on the production of Li, Be, and B from spallation of C, N, and O atoms in the early Universe. Moreover, the existence of CEMP-no stars suggests the synthesis of CNO to have played an important role in the first stellar generations.  In standard nucleosynthesis (without rotation), N is essentially a secondary element in massive stars, and has both primary and secondary components in low- and intermediate-mass stars. Carbon is produced as a primary element both in massive and low-and-intermediate mass stars. Oxygen is synthesized as a primary element in massive stars. Most massive stars would end their lives as type II SNe, dispersing their chemical make-up into the interstellar medium (ISM). However, it is believed that the most massive stars would collapse directly into black holes without the ejection of a supernova, especially at very-low metallicities (Heger et al. 2003)\\footnote{Another possibility is the formation of faint supernovae, with large quantities of mixing and fall back as suggested by Nomoto and co-authors (e.g. Iwamoto et al. 2005)}. In this case, massive stars would contribute to the chemical enrichment of the ISM only via stellar winds.  The stellar yields of CNO can be affected by mass loss and rotation\\footnote{Fast-rotating massive stars can trigger stellar winds even at very-low metallicities -- see Meynet et al. (2008).}. In particular, the C and N yields of intermediate-mass stars still suffer from severe uncertainties related not only to mass loss rates, but also to the difficulties in modeling transport mechanisms (such as dredge-up episodes). In fact, there is evidence of additional mixing processes not included in standard models, triggered by rotation, gravity waves, and thermohaline mixing (e.g. Charbonnel \\& Talon 2005) in this mass range.  The measurements of C, N, and the C-isotopic ratio in normal VMP stars (Spite et al. 2005, 2006) have already set strong constraints on the nucleosynthesis of these elements by the first generations of massive stars. In fact, until 2004 no conclusive data were available for nitrogen in metal-poor halo stars.   This situation has improved with the First Stars ESO large program (Cayrel et al. 2004), which obtained CNO abundances  for the first time for a sample of giants with metallicities below [Fe/H]$=-$2.5 (Spite et al. 2005, 2006 -- see also Fabbian et al. 2009). It turned out that these VMP halo stars had N/O ratios around solar, suggestive of high levels of production of primary nitrogen in massive stars. Moreover a large true scatter (more than the uncertainties in the derived abundances) has been found for the measured N/O. A large scatter has been also found for neutron capture elements in the same stars. These findings contrast with the results for alpha elements, which instead presented striking homogeneous [$\\alpha$/Fe] ratios.  In Chiappini et al. (2005 -- hereinafter C05), the implications of the new CNO data from Spite et al. (2005) on our understanding of nitrogen enrichment in the Milky Way were investigated. By the time the latter paper was published, there was no set of stellar yields able to explain the very metal-poor data from Spite et al. (2005).  C05 concluded that the only way to account for the new data was to assume that massive stars at low metallicity rotate fast enough to produce larger amounts of nitrogen. As shown by Meynet \\& Maeder (2002), rotationally-induced mixing transports the C and O produced in the He-burning core into the H-burning shell, where they are transformed in primary $^{13}$C and $^{14}$N. The efficiency of this process increases when the initial mass and rotational velocity increase, and the metallicity decreases (but see Ekstr\\\"om et al. 2008). C05 predicted that massive stars born with metallicities below Z$ = 10^{-5}$ should produce a factor of 10 up to a few times 100 more nitrogen (depending on the stellar mass) than the values given by Meynet \\& Maeder (2002) for Z $= 10^{-5}$ and for v$_{\\mathrm{ini}}^{\\mathrm{rot}} =$ 300 km s$^{-1}$. This prediction has been confirmed by subsequent stellar evolution models computed at very low metallicities (Hirschi 2007), and higher rotational velocities, which were found to produce much more N than in the models of Meynet \\& Maeder (2002). Chemical evolution models computed with the new stellar evolution predictions turned out to not only account for the high N/O in {\\it normal}-VMP stars (Spite et al. 2005), but also for the C/O upturn and low $^{12}$C/$^{13}$C ratios (Spite et al. 2006) at very low metallicities (Chiappini et al. 2006, 2008). However, the same models cannot account for the huge CNO enhancements observed in CEMP stars.  Finally, C05 suggested that if the stellar yields of N are strongly dependent on the rotational velocities, hence on the mass of very-metal-poor massive stars, it is possible to understand the apparently contradictory finding by Spite et al. (2005) of a large scatter in N/O and the almost complete lack of scatter in [$\\alpha$/Fe] ratios found in the same very metal-poor halo stars (Cayrel et al. 2004). Although the observed scatter could be related to the distribution of stellar rotational velocities as a function of metallicity, that the neutron capture elements in the same stars also show a large scatter pointed to a strong variation in the stellar yields with the mass range of the stars responsible by the synthesis of these elements. Cescutti (2008) explains simultaneously the observed spread in the neutron capture elements and the lack of scatter in the alpha elements as being caused by the stochasticity in the formation of massive stars, combined with the fact that massive stars of different mass ranges are responsible for the synthesis of the different chemical elements, namely: only massive stars with masses between 12 and 30 M$_{\\odot}$ contribute to the neutron capture elements, whereas the whole mass range of massive stars (10 to 80 M$_{\\odot}$) contribute to the production of alpha elements.  In the present paper we study the CNO evolution in the early phases of the galactic formation by means of the inhomogeneous code developed by Cescutti (2008) with the aim of reproducing not only the mean trend in the chemical abundances of CNO as in Chiappini et al. (2006), but also the spread in these particular elements. We assume the spread to be created by the stochasticity in the formation of stars, as in Cescutti (2008). Our aim is to see whether the same process as invoked by Cescutti (2008) to explain the scatter observed in the r- and s- process elements in normal stars also applies to CNO for which a large scatter is also observed. We tested the hypothesis that in the case of CNO the scatter is created by the fact that the most massive stars in the early Universe (Z $<10^{-8}$ and M $>$ 40 M$_{\\odot}$) were fast rotators, which collapse directly into a black hole (e.g. Heger et al. 2003), but contribute to the chemical enrichment of the ISM via stellar winds (Hirschi 2007, Meynet et al. 2006). We also investigate whether the predicted scatter is enough to account for the existence of CEMP-no stars.  In Sect. 2 we introduce the observational data we have adopted in the present paper. In Sect.3 we briefly present our inhomogeneous model for the Galactic halo. The adopted stellar nucleosynthesis is described in Sect. 4. Section 5 is devoted to comparing our model predictions to the new data set of metal-poor stars in the solar vicinity. Finally, a discussion of our results and our conclusions are presented in Sect. 6. ",
        "conclusions": "In this paper we have shown the results of an inhomogeneous chemical evolution model for the mean trend and scatter of C/O and N/O in the Galactic halo, taking the contribution of fast-rotator stars into account.  The observational studies of chemical enrichment in very metal-poor stars of the halo have found a large scatter (larger than the uncertainties in the derived abundances) for the measured C/O and N/O. A large scatter was also found for neutron capture elements in the same stars. These findings contrast with the results for $\\alpha$-elements, which instead presented striking homogeneous [$\\alpha$/Fe] ratios. In Cescutti (2008), the spread in the chemical abundance ratios of neutron capture elements has been explained with a different mass range for the production of these elements (from 12 to 30 M$_{\\odot}$), compared to the whole range of massive stars for the $\\alpha$-elements.  In this work the nucleosynthesis differences for the elements C, N, and O come from the rotation of massive stars, which strongly affects the ratio of the production among these elements at low metallicity (see Hirschi et al. 2007). Moreover, we consider two possible contributions to the enrichment of the ISM by  massive  stars at very low metallicity, the usual enrichment through supernovae ejecta, and the enrichment only through their stellar winds.  We find that the assumption that the most massive fast rotators only contribute to the ISM enrichment via stellar winds leads to chemical evolution models that are able to account for both the large scatter in the N/O and C/O and the simultaneous lack of scatter in $\\alpha$-elements observed in very metal-poor normal halo stars. In fact, the stellar yields of the latter elements do not show any strong dependency on the stellar mass, contrary to what happens for N, C, and n-capture elements.  In this context, we also explored whether the scatter in N and C created by the strong yields dependence on the stellar mass, when including the fast rotators, would also account for the existence of the so-called carbon-enhanced stars (in particular the CEMP-no ones). We find that, even when considering that the more massive fast rotators enrich the early ISM mainly via stellar winds strongly enhanced in CNO, it is impossible to explain the abundances observed in CEMP-no stars. However, we point out that the latter model  marginally agrees with the abundance ratios of the only CEMP-r star for which the CNO abundances have been measured, and of the two of the three UMP stars known to date, HE~0107$-$5240 and HE~0557$-$4840. These results suggest that very metal-poor massive stars would collapse directly into black holes as predicted by Heger et al. (2003).  The discrepancy with respect to the CEMP-no stars, and in particular, HE~1327$-$2326, suggest that these objects are born from the gas expelled during the wind phase of fast rotators, without much mixing with the surrounding ISM. A way to test this hypothesis is to look for the abundances of $^{7}$Li, $^{12}$C/$^{13}$C, and helium in CEMP-no stars. In fact, the wind material is depleted in $^{7}$Li and enriched in $^{13}$C (see Meynet et al. 2010 and Chiappini et al. 2008).  This work illustrates the importance of future surveys that will enlarge the samples of very metal-poor stars with good abundance determinations. The large statistics brought by these large datasets will enable  detailed study not only of the abundance trends but also of the intrinsic scatter in those trends and its metallicity dependence. On the other hand, it is clear that high-quality abundances for key elements such as CNO, the C isotopic ratio, n-capture elements, and $^{7}$Li for a large sample of stars (including CEMP-no stars) is still needed. Spectroscopic follow-ups of metal-poor stars found in the ongoing and planned large surveys (e.g. SEGUE-2 and LAMOST) will play a crucial role in constraining the stellar nucleosynthesis of the first generation of massive stars and in unveiling the role of the fast rotators in the early Universe."
    },
    "1004/1004.2208_arXiv.txt": {
        "abstract": "{Based on the far infrared excess the Herbig class of stars is divided into a group with flaring circumstellar disks (group I) and a group with flat circumstellar disks (group II). Dust sedimentation is generally proposed as an evolution mechanism to transform flaring disks into flat disks. Theory predicts that during this process the disks preserve their gas content, \\ch{however} observations of group II Herbig Ae stars demonstrate a lack of gas.} {We map the spatial distribution of the gas and dust around the group II Herbig Ae star \\object{HD\\,95881}.} {We analyze optical photometry, Q-band imaging, infrared spectroscopy, and K and N-band interferometric spectroscopy. We use a Monte Carlo radiative transfer code to create a model for the density and temperature structure which quite accurately reproduces all the observables. } {We derive a consistent picture in which the disk consists of a thick puffed up inner rim and an outer region which has a flaring gas surface and is relatively void of \\ch{'visible'} dust grains. } {HD\\,95881 is in a transition phase from a gas rich flaring disk to a gas poor self-shadowed disk.} ",
        "introduction": "Herbig Ae (HAe) stars are known to have gas-rich, dusty disks that are the remnant of the star formation process. These disks are most likely the sites of ongoing planet formation. The processes leading to and associated with planet formation modify both the composition and the geometry of the disk. Grain growth and grain settling are expected to result in large spatial variation of the grain size distribution and the gas to dust mass ratio within the disk. The gravitational interaction of proto-planets with the disk can create gaps/holes. Also the gas and the dust chemistry is expected to vary spatially.  In order to understand planet formation, it is thus important to establish the spatial distribution of gas and dust in protoplanetary disks independently.  Observationally, the Spectral Energy Distributions (SEDs) of HAe stars have been divided into two groups, that reflect differences in the slope of the mid-IR (10-60 $\\mu$m) spectral range \\citep{2001A&A...365..476M}. Group\\,I sources have red SEDs, while group II sources have blue SEDs. These differences can be interpreted in terms of the geometry of the disk. The direct irradiation of the inner rim of a disk with an inner hole causes it to be puffed up \\citep{2001ApJ...560..957D}. This puffed up inner rim casts a shadow, and only the outer disk surface regions emerge from the shadow and receive direct stellar light. Depending on the dust opacity, some disks may never emerge from the shadow of the inner rim \\citep{2004A&A...417..159D}. This provides an elegant explanation for the observed two types of SEDs: group I sources being \\emph{flaring}, and group II sources \\emph{self-shadowed}.  This interpretation has been confirmed using spatially resolved mid-infrared (IR) imaging with the {\\em Very Large Telescope Interferometer} (VLTI; e.g. \\citealt{2004A&A...423..537L}).  A difference between group I and group II sources was also found for the strength of the mid-IR emission bands attributed to Polycyclic Aromatic Hydrocarbons (PAHs; \\citealt{2001A&A...365..476M}; \\citealt{2004A&A...426..151A}): flaring disks tend to show strong PAH emission while self-shadowed sources show weaker or no PAH emission \\ch{(\\citealt{2010ApJ...PAHs})}. A similar difference was found for the strength of the [O\\,I]~6300\\,\\r{A} line \\citep{2005A&A...436..209A}. However, there is significant scatter in these trends (see below). Both the PAHs and the [O\\,I] line strength probe the \\emph{gas} in the upper disk layers, and both require direct irradiation of the disk surface by stellar photons to be excited. PAHs mainly probe the disk on scales of several tens to 100 AU (e.g. \\citealt{2004A&A...418..177V}; \\citealt{2006Sci...314..621L}; \\citealt{2007A&A...476..279G}; \\ch{\\citealt{2010A&A...HAEBE}}), i.e. similar scales as the dust continuum emission in the 10-60 $\\mu$m wavelength range. These observations suggest that the spatial distribution of gas in group II sources is different from that of group I sources: apparently, in group II sources the gas in the surface of the outer disk does not receive direct stellar photons.  In a theoretical study, \\citet{2007A&A...473..457D} show that for disks in which the dust settles but the scale-height of the gas does not change, both group I and group II sources should show prominent PAH emission, contrary to the observed trend. However, some disks classified as group II sources (i.e. with a self-shadowed dust geometry) are observed to show prominent PAH emission and [O\\,I] line emission; examples are HD\\,98922 and HD\\,95881 \\citep{2005A&A...436..209A, 2004A&A...426..151A}. \\ch{\\cite{2010A&A...HAEBE}} show that for HD\\,95881 the PAH emission is extended at a scale similar to those of group I sources. \\citet{2008A&A...491..809F} \\ch{and \\cite{2008A&A...485..487V} studied three HAeBe stars and found that the gas and dust in their disks has a different spatial distribution.} In the case of the group II source HD\\,101412, PAH and [O\\,I] emission were detected and found to be more extended than the dust continuum at 10\\,$\\mu$m. These observations suggest that disks exist in which \\emph{the dust has settled but the scale-height of the gas is still large enough at several tens of AU distance from the star to produce substantial PAH and [O\\,I] line emission}. Such disks may provide important clues as to how gas-rich disks evolve to gas-poor debris disks.  In this paper, we study the spatial distribution of the gas and dust in the disk of HD\\,95881. This star was part of a larger study of spatially resolved mid-IR spectroscopy of HAe stars (\\ch{\\citealt{2010A&A...HAEBE}}). We use optical spectroscopy of the [O\\,I] line, the SED, infrared spectra as well as near-IR and mid-IR interferometric observations to constrain the geometry of the gas and dust in the disk. We use a hydrostatic equilibrium disk model to fit the SED and compare the predicted spatial distribution of the near-IR and mid-IR emission of the best fitting disk model to the interferometric observations. We find convincing evidence that the dust in the disk of HD\\,95881 has settled but that the gas still has a significant scale-height. We derive an estimate on the total disk mass by fitting the strength of the PAH bands.   \\parameters ",
        "conclusions": "A comprehensive study was performed to map the distribution of the gas and dust in the protoplanetary disk around HD\\,95881. In \\ch{the right panel of Fig.\\,\\ref{fig:models}} we displayed a schematic representation of the disk, which puts all results in perspective. The AMBER K-band interferometry showed that there is an extended hot inner region with emission coming from \\ch{well} within the sublimation radius. The detection of the [O\\,I]~6300\\,\\r{A} indicated that the disk has a flaring gas surface at large distances (from one to tens of AU) from the star. The finding of PAH features in the Spitzer and VISIR spectra confirmed the presence of an illuminated gas surface. The resolved VISIR spectrum traced this surface up to \\ch{radii of $\\sim$100\\,AU.}  We used the radiative transfer code MCMax (\\citealt{2009A&A...497..155M}) to create a model of the disk's density and temperature structure. Our model satisfactorily reproduced all of our observations. The main conclusions that followed from our model are that the inner disk contains most of the \\ch{visible} grains and has a puffed up inner rim, the dust grains in the outer disk have coagulated and settled towards the midplane, while the gas and PAH mixture maintain a flaring geometry. Theory predicted the existence of these type of disks (\\citealt{2007A&A...473..457D}), while observational trends showed that most of the sources with self-shadowed dust distributions have dispersed \\ch{the bulk of} their gas. In this light HD\\,95881 is a special source: it is in the transition phase from a gas rich flaring dust disk to a gas poor self-shadowed disk."
    },
    "1004/1004.0511_arXiv.txt": {
        "abstract": "{Gamma-ray binaries could be compact pulsar wind nebulae formed when a young pulsar orbits a massive star. The pulsar wind is contained by the stellar wind of the O or Be companion, creating a relativistic comet-like structure accompanying the pulsar along its orbit.} {The X-ray and the very high energy ($>$100 GeV, VHE) gamma-ray emission from the binary LS 5039 are modulated on the orbital period of the system. Maximum and minimum flux occur at the  conjunctions of the orbit,  suggesting that the explanation is linked to the orbital geometry. The VHE modulation has been proposed to be due to the combined effect of Compton scattering and pair production on stellar photons, both of which depend on orbital phase. The X-ray modulation could be  due to relativistic Doppler boosting in the comet tail where both the X-ray and VHE photons would be emitted.} {Relativistic aberrations change the seed stellar photon flux in the comoving frame so Doppler boosting affects synchrotron and inverse Compton emission differently. The dependence with orbital phase of relativistic Doppler-boosted (isotropic) synchrotron and (anisotropic) inverse Compton emission is calculated, assuming that the flow is oriented radially away from the star (\\ls) or tangentially to the orbit (\\lsi, \\psrb).} {Doppler boosting  of the synchrotron emission in \\ls\\ produces a lightcurve whose shape corresponds to the X-ray modulation. The observations imply an outflow velocity of 0.15--0.33$c$ consistent with the expected flow speed at the pulsar wind termination shock. In \\lsi, the calculated Doppler boosted emission peaks in phase with the observed VHE and X-ray maximum.} {Doppler boosting is not negligible in gamma-ray binaries, even for mildly relativistic speeds. The boosted modulation reproduces the X-ray modulation in \\ls\\ and could also provide an explanation for the puzzling phasing of the VHE peak in \\lsi. } ",
        "introduction": "Gamma-ray binaries display non-thermal emission from radio to very high energy gamma rays (VHE, $>$100 GeV). Their spectral luminosities peak at energies greater than a MeV. At present, three such systems are known: \\psrb\\ \\citep{Aharonian:2005br}, \\ls\\ \\citep{Aharonian:2005nj} and \\lsi\\ \\citep{Albert:2006wi}. A fourth system, HESS J0632+057 may also be a gamma-ray binary \\citep{Hinton:2008eg}. The systems are composed of a O or Be massive star and a compact object, identified as a young radio pulsar in PSR B1259-63. All gamma-ray binaries could harbour young pulsars \\citep{Dubus:2006lc}.  Electrons accelerated in the binary system upscatter UV photons from the companion to gamma-ray energies. The Compton scattered radiation received by the observer is anisotropic because the source of seed photons is the companion star. VHE gamma-rays will also produce $e^+e^-$ pairs as they propagate through the dense radiation field, absorbing part of the primary emission. This is also anisotropic.  Both effects combine to produce an orbital modulation of the gamma-ray flux if the electrons are in a compact enough region. This modulation depends only on the geometry. Orbital modulations of the high-energy (HE, $>$100 MeV) and VHE fluxes have indeed been observed. The modulations unambiguously identify the gamma-ray source with the binary \\citep{Aharonian:2006qw,Albert:2006wi,Acciari:2008vf}.  Synchrotron emission can dominate over inverse Compton scattering at X-ray energies, providing additional information to disentangle geometrical effects from intrinsic variations of the source.  {\\em Suzaku} and {\\em INTEGRAL} observations of \\ls\\ have revealed a stable modulation of the X-ray flux  \\citep{Takahashi:2008vu, Hoffmann:2008ys}.  Possible interpretations are discussed in \\S2. None are satisfying. The key point is that the X-ray flux is maximum and minimum at conjunctions and that this excludes any explanation unrelated to the system's geometry as seen by the observer.  In the pulsar wind scenario, the synchrotron emission is expected to arise in shocked pulsar wind material collimated by the stellar wind. This creates a cometary tail with a mildly relativistic bulk motion (Fig.~\\ref{orb_boost}). % Relativistic Doppler boosting of the emission due to this bulk motion is calculated in \\S3 with details given in Appendix A. The orbital motion leads to a modulation of the Doppler boost, as previously proposed in the context of black widow pulsars \\citep{1993ApJ...403..249A,2007A&A...463L...5H}. The calculated synchrotron modulation is similar to that seen in X-rays in \\ls. Although this is not formally confirmed due to their long orbital periods, \\lsi\\ and \\psrb\\ also appear to have modulated X-ray emission \\citep{Chernyakova:2006cu,2009MNRAS.397.2123C,2009ApJ...700.1034A,2009ApJ...706L..27A}.  The application to these gamma-ray binaries is discussed in \\S4.  \\begin{figure} \\centering \\resizebox{7cm}{!}{\\includegraphics{14023fig1}} \\caption{Geometry of Doppler boosted emission from a collimated shock pulsar wind nebula. The orbit is that of LS 5039 (to scale). The comet tail moves away from the pulsar with a speed $\\beta=v/c$ at an angle $\\theta_{\\rm flow}$. If $\\theta_{\\rm flow}=0$ then intrinsic emission in the co-moving frame is boosted in the observer frame at inferior conjunction  and deboosted at superior conjunction.} \\label{orb_boost} \\end{figure} ",
        "conclusions": "The X-ray orbital modulation of \\ls\\ peaks and falls at conjunctions, suggesting that the underlying mechanism is related to the geometry seen by the observer. Phase-dependent Doppler boosting of emission from a mildly relativistic flow provides a viable explanation. The underlying assumption is that the flow direction changes with orbital phase, so that even constant intrinsic emission becomes variable as seen by the observer. The peaks and troughs are at conjunctions for a flow directed radially away from the star, as expected if the emission arises from a shocked pulsar wind confined by the fast stellar wind of its companion  \\citep{Dubus:2006lc}. A moderate relativistic speed of $\\beta=0.15$ or 1/3 is enough to reproduce the morphology of the observed X-ray lightcurve assuming (resp.) either constant intrinsic emission or the model of \\citet{Dubus:2007oq}. Note that these values of $\\beta$ allow for quite large values of the opening angles. More detailed calculations assuming a conical geometry for the flow confirmed that the results were unchanged as long as the angular size of the flow is smaller than $1/\\Gamma$ (if larger, the modulation is dampened). Reproducing the level of X-ray emission is difficult with a one-zone model as it requires values of the magnetic field that are a factor 3 above current values, leading to cutoff in the VHE specta at energies that are too low. A more complex multi-zone model of the post-shock flow might resolve this discrepancy.  Inverse Compton scattering in the flow of external stellar photons will be modulated differently than intrinsic emission from the flow. In the case of a radial outflow, the external seed photon flux will be deboosted at all phases. However, a flow tangent to an eccentric orbit, as might arise in \\lsi\\ and \\psrb, can lead both to boosts and deboosts in the comoving frame depending on orbital phase and thus give rise to complex modulations. The calculated Doppler corrected emission in \\lsi\\  peaks in phase with the observed VHE maximum. This is noteworthy since a simple explanation had not yet been proposed for the phase of VHE (and X-ray) maximum in \\lsi. This explanation requires that the shocked pulsar wind flows along the orbit, which appears compatible with the radio VLBI images on larger scales shown in \\citet{Dhawan:2006kr}.  The present work assumed a pulsar relativistic wind in the orbital plane but microquasar models have also been proposed for both \\ls\\ and \\lsi. In this case, the emission arises from a relativistic jet.  The jet angle to the observer remains constant along the orbit and so do ${\\cal D}_{\\rm obs}$ and $F_{\\rm syn}$. Hence, no orbital modulation of intrinsic (synchrotron) X-ray emission due to Doppler boosting would be expected, apart from the possible impact of  jet precession on timescales longer than the orbital period \\citep{2002A&A...385L..10K}.  Doppler boosting in a relativistic jet cannot explain the X-ray modulation in \\ls\\ or \\lsi.  However, unless the electrons are far from the system or the system is seen pole-on, the angle of interaction between photons and electrons  $\\mathbf{e_\\star}.\\mathbf{e_{\\rm obs}}$ will change with orbital phase. A modulation in $F_{\\rm ic}$ is unavoidable. This variation in inverse Compton emission can explain the orbital modulation seen in high-energy gamma-rays from the microquasar Cygnus X-3 by  {\\em Fermi Gamma-ray Space Telescope} \\citep{2009Sci...326.1512Fa,2010arXiv1002.3888D}."
    },
    "1004/1004.5102_arXiv.txt": {
        "abstract": "A didactic introduction to current thinking on some aspects of the solar dynamo is given for geophysicists and planetary scientists. ",
        "introduction": "For a long time, solar dynamo theory was in an advantageous situation compared to planetary dynamo theory. In the Sun, dynamo generated magnetic fields can be directly measured on the boundary of the turbulent, conducting domain where they are generated; and the timescales of their variations are short enough for direct observational follow-up. Solar observations also put many constraints on the motions in the convective zone, significantly restricting the otherwise very wide range of admissible mean field dynamo models.  In the past 15 years, however, planetary dynamo theorists have turned the table. Realistic numerical simulations of the complete geodynamo have been made possible by the rapid increase in the available computing power. While the parameter range for which such simulations are feasible is still far from realistic, extrapolations based on the available results have allowed important inferences on the behaviour of planetary dynamos (\\citen{Christensen:plan.dynamo.solar}, \\citen{Petrovay:solar.planet.dynamos}).  In the light of these developments, learning from solar dynamo theory may have lost much of its former appeal to geophysicists, especially as it has become increasingly clear that the two classes of dynamos operate in fundamentally different modes, under very different conditions. Yet, keeping track with advance in the other field may not be without profit for either area. Even though the overall mechanisms may be very different, there may be many elements of each system where intriguing parallels exist.  A comprehensive review of solar dynamo theory is beyond the scope of the present article; for this, we refer the reader to papers by \\cite{Petrovay:SOLSPA}, \\cite{Charbonneau:livingreview}, \\cite{Solanki+:RPP} and \\cite{Jones+:dynamo.rev}.  Aside from reproducing the dipole dominance and other morphological traits of the geomagnetic field, two aspects of the geodynamo that are critical for judging the merits of its models are field reversals and long-term variations: whether or not a model can produce such effects in a way qualitatively, and if possible quantitatively similar to the geological record, has become a testbed for geodynamo simulations. It may thus be of special interest to review our current understanding of the analogues of these phenomena in the Sun. This is the purpose of the present paper.  In Section~2 we outline what solar observations suggest about the causes of reversals, i.e.\\ the Babcock--Leighton mechanism. Two questions naturally arising from this discussion are given further attention in Sections 3 and 4. Section~5 briefly discusses some issues related to long-term variations of solar activity, while Section~6 concludes the paper, pointing out some interesting parallel phenomena in solar and planetary dynamos. ",
        "conclusions": "At the most basic level, the same processes generate magnetic field in the Earth's core and inside the Sun: the shearing of magnetic field lines by differential rotation ($\\Omega$-effect) and their twisting by helical motions ($\\alpha$-effect) that obtain a preferred handedness from the action of Coriolis forces. Also, some magnetic phenomena may have similar causes. For example, paired flux spots at low latitudes in the geomagnetic field at the core-mantle boundary have been tentatively explained by the expulsion of toroidal flux tubes (\\citen{Bloxham86}), in analogy to the generation of sunspots. But although this interpretation is supported by some geodynamo models (\\citen{ChristensenOlson03}), other explanations for low-latitude magnetic structures have been put forward (\\citen{FinlayJackson03}). Meridional flow is thought to play the essential role for reversals of the solar magnetic field. This has also been demonstrated in a simple reversing geodynamo model (\\citen{WichtOlson04}). However, the reversal behaviour in this model is nearly cyclic as in case of the Sun, in contrast to the stochastic reversals of the geomagnetic field. In a less idealized geodynamo model with random reversals, impulsive upwellings have been identified as the cause for polarity changes (\\citen{Aubert_etal08}). These upwellings transport and amplify a multipolar magnetic field from depth to the outer boundary. While possible analogies between the solar dynamo and the geodynamo can stimulate our thinking, we must keep their limitations in mind.  Which differences in physical conditions lead to the rather distinct behaviour of the solar dynamo and the geodynamo? One important difference is that the plasma in the solar convection zone is sufficiently compressible and that the field strength is high enough so that magnetic pressure and magnetic buoyancy play an essential role for the dynamics of flux tubes. These effects are probably unimportant in Earth's core. Another difference is that the Coriolis force has a stronger influence in the geodynamo than it has in the slowly rotating Sun. This notion is supported by the observation that much more rapidly rotating stars of low mass seem to have strong large-scale magnetic fields that are frequently dominated by the axial dipole component (\\citen{Donati_etal08}). Their observed field strengths follow the same scaling law as the observed fields of Earth and Jupiter and the field intensity found in geodynamo models at sufficiently rapid rotation (\\citen{Christensen_etal09}). A third difference arises from the much slower motion and therefore lower magnetic Reynolds number in the geodynamo.  The moderate value of the magnetic Reynolds number makes the magnetic induction process in the Earth's core amenable to direct numerical simulations without the need to take recourse to turbulent magnetic diffusivities or parameterized turbulent $\\alpha$-effects. This is perhaps the most important reason for the success of geodynamo models in reproducing many observed properties of the geomagnetic field without need for ad-hoc assumptions. However, our more limited knowledge of the geomagnetic field at the top of the core, in comparison to that of the field in the solar photosphere, makes the task simpler for a geodynamo modeller. Also, helioseismology has revealed the distribution of zonal flow in the solar convection zone and a fully consistent solar dynamo model must reproduce this flow pattern as well as the magnetic field properties. Comparable information is lacking for the Earth and a geodynamo model can be declared successful once it captures the general properties of the large-scale geomagnetic field."
    },
    "1004/1004.3454_arXiv.txt": {
        "abstract": "Heating mechanisms of the solar corona will be investigated at the high-altitude solar observatory Lomnicky Peak of the Astronomical Institute of SAS (Slovakia) using its mid-size Lyot coronagraph and post-focal instrument SECIS provided by Astronomical Institute of the University of Wroc{\\l}aw (Poland). The data will be studied with respect to the energy transport and release responsible for heating the solar corona to temperatures of mega-Kelvins. In particular investigations will be focused on detection of possible high-frequency MHD waves in the solar corona. The scientific background of the project, technical  details of the SECIS system modified specially for the Lomnicky Peak coronagraph, and inspection of the test data are described in the paper. ",
        "introduction": "For more than 60 years it has been well known that the quiet solar corona is heated to a temperature of about 1--2 million Kelvins while the visible surface of the Sun is roughly 250 times cooler (Grotrian 1939; Edlen 1942; Phillips, 1995). It has been also recognized that magnetic fields or waves play a key r\\^{o}le in the heating of the solar corona so that somehow convective energy in the photosphere is converted to thermal energy in the corona via magnetic fields or wave energy. The primary energy source for this heating must lie in the convection zone below the solar photosphere (e.g. Bray et al., 1991; Golub \\& Pasachoff 1998; Aschwanden 2004) where there is 100 times as much energy available than that required to heat the corona ($\\approx$300\\,W/m$^2$: Withbroe \\& Noyes 1977, Aschwanden 2001).  Currently, the debate centres on whether the energy to heat the corona derives from dissipation of magneto-hydrodynamic (MHD) waves (e.g. Hollweg, 1981) or from numerous small-scale magnetic reconnection events giving rise to nanoflares (Parker, 1988, Aschwanden 2004). It has been found theoretically that the interaction of the magnetic field with convective flows in or below the photosphere can produce two types of magnetic disturbances in coronal structures. Firstly, the buffeting of magnetic flux concentrations in the photosphere by granulation generates MHD waves which can propagate into magnetic flux tubes and dissipate their energy in the chromosphere or corona (e.g. Ofman et al. 1998). Secondly, in coronal loops the random motions of magnetic loop foot-points can produce twisting and braiding of coronal field lines, which generates field-aligned electric currents that can be dissipated resistively (e.g. Parker 1972, 1983; Heyvaerts \\& Priest 1983; van Ballegooijen 1990). The main difference between these processes is that plasma inertia plays a key r\\^{o}le in wave propagation, but is unimportant for the dynamics of field-aligned currents along coronal loops. Thus these types of magnetic heating mechanisms can be crudely classified as either wave-heating or current-heating mechanisms.  There are theoretical arguments for both mechanisms, but the observational evidence for nano-flare heating is perhaps looking less convincing than before. Extrapolation of the number spectra of small flares down to microflares has been made to nano-flares but the total energy, while tantalisingly close, is most probably less than the required amount (Parnell et al. 2000). ",
        "conclusions": "The SECIS instrument installed at the Lomnicky Peak Observatory Lyot coronagraph will allow data to be acquired that may result in an improved knowledge of where in the corona MHD waves are generated and/or dissipated. In particular, the signatures of high-frequency MHD waves involved in coronal heating may be observed. A considerable improvement in our knowledge of a long-standing problem of solar physics could be made by such observations, with implications for the physics of active regions, flares, the solar wind, and solar activity, as well as mechanisms of solar-terrestrial relationships."
    },
    "1004/1004.1384_arXiv.txt": {
        "abstract": "{In the recent years, MRI-driven turbulent transport has been found to depend in a significant way on fluid viscosity $\\nu$ and resistivity $\\eta$ through the magnetic Prandtl number $Pm=\\nu/\\eta$. In particular, the transport decreases with decreasing $Pm$; if persistent at very large Reynolds numbers, this trend may lead to question the role of MRI-turbulence in YSO disks, whose Prandtl number is usually very small.} {In this context, the principle objective of the present investigation is to characterize in a refined way the role of dissipation. Another objective is to characterize the effect of linear (channel modes) and quasi-linear (parasitic modes) physics in the behavior of the transport.} {These objectives are addressed with the help of a number of incompressible numerical simulations. The horizontal extent of the box size has been increased in order to capture all relevant (fastest growing) linear and secondary parasitic unstable modes.} {The major results are the following:\\\\ i- The increased accuracy in the computation of transport averages shows that the dependence of transport on physical dissipation exhibits two different regimes: for $Pm \\lesssim 1$, the transport has a power-law dependence on the magnetic Reynolds number rather than on the Prandtl number; for $Pm > 1$, the data are consistent with a primary dependence on $Pm$ for large enough ($\\sim 10^3$) Reynolds numbers.\\\\ ii- The transport-dissipation correlation is not clearly or simply related to variations of the linear modes growth rates.\\\\ iii- The existence of the transport-dissipation correlation depends neither on the number of linear modes captured in the simulations, nor on the effect of the parasitic modes on the saturation of the linear modes growth.\\\\ iv- The transport is usually not dominated by axisymmetric (channel) modes. } {} ",
        "introduction": "Disks evolve on time-scales that are orders of magnitudes smaller than expected from microphysical transport processes, and various suggestions have been made over the years to explain this discrepancy. Turbulent transport, in particular, has figured among the leading candidates since the inception of the $\\alpha$-disk paradigm, and a number of hydrodynamic and MHD turbulent transport mechanisms have been proposed in the literature.  On the hydrodynamic side, subcritical turbulence (\\citealt{RZ99} and references therein), if present, is apparently too inefficient \\citep{LL05,JBSG06}. Convection was up to now found too inefficient and to transport angular momentum in the wrong direction \\citep{C96,SB96}, but a recent reinvestigation of the problem indicates that this might be an artifact of these simulations being performed too close to the stability threshold \\citep{LO10}. Two-dimensional weak turbulence driven by small-scale, incoherent gravitational instabilities (density waves) is an option \\citep{G96}. Alternatively, the baroclinic instability \\citep{KB03} may generate vorticity, and transport through the coupling with density waves, but its conditions of existence are still controversial \\citep{JG06,PSJ07}, although \\cite{LP10} have probably identified the root of this debate by pointing out the nonlinear nature of the instability; also the resulting vortices would be subject to 3D instabilities \\citep{LP09}.  \\cite{BH91a} have proposed that the magnetorotational instability (MRI) is a potentially efficient source of turbulent transport in the nonlinear regime, an expectation soon borne out in numerical simulations. This instability provides by now the most extensively studied transport mechanism, through local unstratified \\citep{HGB95}, stratified \\citep{SHGB96}, and global \\citep{H00} 3D disk simulations. These initial simulations as well as the numerous ones following them have shown that MRI turbulence is an efficient way to transport angular momentum, in the presence or absence of a mean vertical or toroidal field, with an overall transport efficiency depending on the field configuration and strength. However, the significant role played by microphysical dissipation in the resolutions accessible to date had largely been underestimated \\citep{LL07,FPLH07}.  By now, both the field strength and dissipation dependence of the simulated turbulent transport have been studied to some extent (and only in unstratified local shearing box settings for the latter one). The dependence of the Shakura-Sunyaev $\\alpha$ parameter has been characterized very early on by \\cite{HGB95} who showed that momentum transport $\\propto \\beta^{-1/2}$ both for a net vertical or toroidal field\\footnote{The definition of $\\beta$ (ratio of gas to magnetic field pressure) is based on the mean field, a conserved quantity in the settings used in these simulations.} (albeit with very different efficiencies in the two configurations), a scaling further confirmed in later simulations, as summarized in \\cite{PCP07}.  Until recently, the effect of physical viscosity ($\\nu$) and resistivity ($\\eta$) on the transport had been neglected, under the implicit assumption that these should not matter too much once inertial turbulent scales are resolved in the simulations. However, \\cite{LL07} have shown that, in the presence of a mean vertical field, the MRI-driven turbulent transport did exhibit a substantial dependence on the magnetic Prandtl number $Pm=\\nu/\\eta$, with no clear trends with respect to either viscosity of resistivity alone\\footnote{The linear stability dependence on viscosity and resistivity is totally different.}. Recently, \\cite{SH09} found similar results in shearing boxes with a mean toroidal field instead of a mean vertical one.  When the mean magnetic flux vanishes, the transport behavior is more complex. The initial investigation by \\cite{HGB96} concluded that the transport was converging to a finite value, but \\cite{GS05} found that the transport efficiency was dependent on the simulation resolution. More recently, the role of the magnetic Prandtl number $Pm$ has been identified in this setting \\citep{FPLH07}: turbulence exists only for magnetic Prandtl numbers larger than about 2, which requires the explicit inclusion of viscous and resistive terms in the fluid equations for numerical simulations to correctly capture the physics of the problem. The disappearance of turbulence at low $Pm$, as well as the need of large enough amplitudes in the initial conditions at $Pm > 2$, indicate that the zero net flux magnetized shearing box is a subcritical system rather than a linearly unstable one \\citep{LO08b,LO08a}.  Thus it appears that in all configurations explored to date, the magnetic Prandtl number plays a significant role on the existence and/or efficiency of the turbulent transport, at least at the resolutions accessible on present day computers (or equivalently, the accessible Reynolds and magnetic Reynolds numbers). This raises a number of issues.  For one, the exact role played by channel modes and parasitic modes is unclear. Although they exist only when a mean vertical field is present, they are simpler to analyze and their behavior may provide insight into the generic mechanism responsible for saturation of the linear instability. Channel modes are the axisymmetric unstable modes of the MRI \\citep{BH91a,PC08}, and are often observed both in 2D and 3D simulations with a mean vertical field; their name derives from their vertically layered characteristic channel-like radial flow. They were quickly recognized to be also nonlinear solution of the problem by \\cite{GX94}; the same authors found them to be unstable with respect to a secondary instability (parasitic modes). A few recent papers have focused on the possibility that the saturation of the channel mode by this parasitic instability might be the mechanism explaining the magnitude of the turbulent transport in MRI simulations, with diverging conclusions \\citep{PG09,LLB09}.  In relation to this, the role of the aspect ratio of the simulations has probably been underestimated in the past. Boxes with an aspect ratio $R:Z=1:1$ do not allow for the fastest parasitic modes to grow, and \\cite{BMCRF08} pointed out that narrow boxes tend to overemphasize the role of the channel modes with respect to more horizontally extended boxes. This calls for a reassessment of the Prandtl number dependence of MRI-driven transport in horizontally extended simulation boxes with a mean vertical field.  More generally, it is still unclear whether this dependence of the transport on physical dissipation is a consequence of the limited Reynolds numbers that can be achieved on present day computers. In particular, none of the published simulations has been able to capture the existence of a significant inertial range in the kinetic or magnetic energy spectrum, which makes it difficult to address this issue. The question here revolves mostly around the direction and locality of transfers and fluxes in Fourier space, and will be addressed elsewhere.  For the time being, we focus the potential role of the channel and parasitic modes in the efficiency of turbulent transport. This is explored by numerical simulations in the shearing sheet limit, with a net vertical magnetic flux, and with horizontally extended simulation boxes. The paper is organized in the following way. Our numerical method, setup, and run parameters are described in section 2. Relevant aspects of the theory of channel modes are summarized in section 3. Section 4 is the core of this paper, and discusses our numerical results; the issues bearing on the resolved linear and secondary modes are also discussed there. The implications of these results are presented in the final section along with some possible future lines of work. ",
        "conclusions": "\\label{discussion}  Perhaps the most significant new result of this work, disclosed on Figs.~\\ref{transp-pm} and \\ref{transpReRm}, is the existence of a double regime separated by a critical magnetic Prandtl number $Pm_c\\sim 1$. For $Pm < Pm_c$, at a given field strength, the transport correlates mostly with $Rm$; for $Pm > Pm_c$, the transport seems to depend mostly on $Pm$ and only weakly on either $Re$ or $Rm$ (once $Re \\gtrsim 10^3$), although a larger number of $Pm$ values need to be probed on this issue. It is tempting to assume $Pm_c\\simeq 2$, as this is the critical value for the zero mean field problem, but this identification requires further work to be substantiated. The identification of this double dissipation regime was made possible by the increased accuracy, with respect to our previous work, in the determination of the transport averages. In the small Prandtl regime, in contrast to the large one, our most resolved simulations show no sign of convergence with respect to dissipation, although values of $Rm$ up to 20000 have been reached.  The role of linear physics and parasitic modes on transport properties has also been investigated, and the major results on these questions can be summarized as follows:  \\begin{itemize} \\item The existence of the dependence of turbulent transport on dissipation is not related to the role of dissipation on linear modes growth rates. \\item The existence of the correlation of transport with dissipation is not related to the number of channel modes captured in the simulations. \\item Saturation of the channel modes growth by parasitic modes is not the responsible for correlation of transport with dissipation. This is hardly surprising as both the channel and parasitic modes are large scale modes, that are expected to be little affected by dissipation, whereas the trends of the transport with dissipation are substantial. \\item Furthermore, the transport is usually not dominated by axisymmetric modes; these modes contribute at most at the same level as non-axisymmetric modes for the strongest fields investigated, and have negligible contributions for the weakest fields. \\end{itemize}  Note that all these results were obtained while the fastest growing channel and parasitic modes were always captured, so that they do not depend on limitations on this front.  It is worth pointing out that our results do not totally disqualify a saturation of the unstable linear modes by the parasitic modes; only the relevance of this process to the relation between transport and physical dissipation has been disproved.  The role of the Prandtl number disclosed in this investigation makes more physical sense than a direct dependence of the transport on $Pm$. Indeed, it is very likely that the critical value $Pm_c$ relates to the switching of the magnetic and kinetic dissipations scales $k_\\eta$ and $k_\\nu$ in Fourier space: for $Pm < Pm_c$ (resp.\\ $Pm > Pm_c$), $k_\\eta < k_\\nu$ (resp.\\ $k_\\eta > k_\\nu$). For $Pm \\ll 1$, $k_\\eta \\ll k_\\nu$, the flow at scales $k >$ or $\\gg k_\\eta$ in Fourier space is therefore purely hydrodynamic. As information flows from small $k$ to large $k$ in this scale range, the flow should be independent of $Re$, so that in this regime $\\alpha=\\alpha(\\beta,Rm)$ is expected. The transport may even become independent of $Rm$ at large enough $Rm$ in the absence of backreaction of the small scales on the large ones, and if energy exchanges are not strongly nonlocal in Fourier space. We merely note here that in the context of homogeneous, isotropic, incompressible MHD turbulence, all energy transfers are direct (from small $k$ to large $k$); exchanges between magnetic and kinetic energy are nonlocal in Fourier space  \\citep{AMP07}, although this nonlocality seems to be of finite (albeit large) extent \\citep{AE09}. The  situation is more complex in the large Prandtl regime, as non-local transfers and small scale dynamo action might take place at scales smaller than the viscous dissipation scale, and further investigations are required to characterize the properties of the small dissipation limit, in particular concerning the locality of transfers.  The behavior of MRI-driven turbulence with respect to these various issues will be reported elsewhere, through the analysis of energy transfers in Fourier space. More generally, investigating the possible independence of transport with respect to dissipation in the $Pm \\ll Pm_c$ or $Pm \\gg Pm_c$ limits requires to achieve a double scale separation: one must first separate in Fourier space the smallest injection scales (the scales which contribute to the Reynolds and Maxwell stresses) from the largest dissipation scales, and then the two dissipation scales themselves. This is extremely demanding in terms of resolution. Our simulations in the small Prandtl regime are consistent with a separation of dissipation scales nearly achieved. As a first step at small Prandtl (the most critical regime for YSO disks), one can achieve one or the other of the two scales separation. Such investigations are underway and will be reported elsewhere.  These questions are critical for astrophysical disk dynamics, where the Reynolds numbers are always very large ($Re \\gtrsim 10^{10}$, while the Prandtl number spans very small ($\\lesssim 10^{-5}$, YSO disks) to very large values ($\\gtrsim 10^{5}$, AGN disks). Addressing these issues requires to overcome a number of limitations in the simulations. Besides the questions of scale separation and nonlocality of transfers in Fourier space just mentioned, the role of compressibility and vertical stratification on these results must be quantified. These questions can certainly be investigated in shearing boxes, but global simulations are probably still largely out of reach, due to the resolution demand. Also (and probably as much importantly), different dissipation regimes must be analyzed in the context of YSOs, namely the ambipolar and Hall regime, which are known to have substantial effect on the linear development of the MRI, while being relevant for large fractions of YSO disks."
    },
    "1004/1004.3724_arXiv.txt": {
        "abstract": "We propose to address the fine tuning problem of inflection point inflation by the addition of extra vacuum energy that is present during inflation but disappears afterwards. We show that in such a case, the required amount of fine tuning is greatly reduced. We suggest that the extra vacuum energy can be associated with an earlier  phase transition and provide a simple model, based on extending the SM gauge group to $SU(3)_C \\times SU(2)_L\\times U(1)_Y\\times U(1)_{B-L}$, where the Higgs field of $U(1)_{B-L}$ is in a false vacuum during inflation. In this case, there is virtually no fine tuning of the soft SUSY breaking parameters of the flat direction which serves as the inflaton. However, the absence of radiative corrections which would spoil the flatness of the inflaton potential requires that the $U(1)_{B-L}$ gauge coupling should be small with $g_{B-L}\\leq 10^{-4}$. ",
        "introduction": "Inflation generated at a point of inflection has the attractive feature of allowing a very low inflationary scale without compromising the amplitude of the density perturbation~\\cite{RM}. This is a direct consequence of the extreme flatness of the potential at the inflection point. A low scale seems like a necessity if we ever hope to connect cosmology with experimental particle physics.  It is well known that the scalar potential of the Minimal Supersymmetric Standard Model (MSSM) has a number of flat directions~\\cite{MSSM-REV} along which inflection points may be found. Indeed, it has been demonstrated~\\cite{AEGM,AKM,AEGJM,AJM} that inflation can occur within MSSM and its minimal extensions, with the remarkable property that the inflaton is {\\it not} an arbitrary gauge singlet. Rather, it is a $D$-flat direction in the scalar potential consisting of the supersymmetric partners of quarks and leptons\\footnote{For models of inflation where the inflaton is not a gauge singlet see~\\cite{FEW}.}. These models give rise to a wide range of scalar spectral indices~\\cite{BDL,AEGJM}, including the whole range permitted by WMAP~\\cite{WMAP7}. Since the inflaton belongs to the observable sector, its couplings to matter and decay products are known.  It is therefore possible to track the thermal history of the universe from the end of inflation. The parameter space permitting successful inflation is compatible with supersymmetric dark matter~\\cite{ADM2} (and may even lead to a unified origin of inflation and dark matter~\\cite{ADM1}).  However, MSSM inflation has one significant problem: soft SUSY breaking parameters in the Lagrangian must be tuned~\\cite{AEGJM} to a very high degree in order to have a sufficiently flat potential around the point of inflection. This tuning does not pose a problem \\textit{per se}; it is common in inflationary model building, particularly in models of low scale inflation. The fine tuning of tree-level parameters might actually reflect the theory at supergravity level and be a natural consequence of the form of the K\\\"ahler potential~\\cite{sugra}, although in that case hidden sector dynamics may also affect inflation \\cite{Lalak}. It is also possible that the proximity of the soft SUSY breaking parameters at inflationary scale can be generated dynamically by virtue of renormalization group equations~\\cite{ADS}.    By means of a simple observation, we can resolve this tuning problem. The fine tuning problem in MSSM inflation arises because the flat interval around the point of inflection is much smaller than the Vacuum Expectation Value (VEV) of the inflection point. Raising the potential during inflation will increase the ratio of the flat interval length to the inflection point VEV and ameliorate the tuning, with the exact degree of tuning dependent on the height of the potential. This also relaxes related constraints such as the $\\eta$ and initial condition problems. Additionally, obtaining acceptable density perturbations for a fixed potential height implies a smaller inflection point VEV and consequently less fine tuning. This opens up the interesting possibility that the inflection point in the potential can be determined from renormalizable couplings of the theory.  The simplest way to lift the potential is by adding vacuum energy $V_0$ which is present during inflation but disappears at the end of the inflationary era. The vacuum energy associated with the Higgs field(s) of a new symmetry will suffice (in a manner similar to hybrid inflation). Indeed, new (gauged or global) symmetries are typical in physics beyond the standard model. The simplest example is a $U(1)$ symmetry that can be implemented in a minimal extension of MSSM.  This paper is structured as follows. We begin by presenting a general analysis of inflection point inflation and its ramifications. We then underline the role of a constant term in the potential and how it can resolve the fine tuning issue. Thirdly, we discuss a possible extension of MSSM that could give rise to inflection point inflation without fine tuning, and finally we offer some concluding remarks. ",
        "conclusions": "We have proposed a solution to the problem of fine-tuning inherent in inflection point inflation, where the extreme flatness of the potential makes it unstable against radiative corrections. In  MSSM inflation models~\\cite{AEGM,AKM,AEGJM,AJM} based on the $udd$ and $LLe$ flat directions, the amount of fine-tuning required for soft SUSY breaking parameters is harsh, i.e. $A/\\sqrt{40}m\\sim \\sqrt{1-4\\beta^2}$ with $\\beta\\sim 10^{-10}$. While it might be possible to sidestep the fine tuning within the context of string landscape~\\cite{AFM}, in the present paper we offer a more mundane prescription based on the simple observation that during inflation, there can be present some vacuum energy in addition to the one given by the inflaton potential at the inflection point.  In this paper the amount of fine-tuning is quantified by the parameter $\\beta$ defined in Eq. (\\ref{defbeta}). We have shown that by adding a constant term $V_0$ to the potential, associated with some field in a false vacuum during inflation, the requisite finetuning of $\\beta$ can be much alleviated and even removed completely.  A simple realization of such a scenario is provided by extending the MSSM gauge group to either adding a singlet field as in the case of NMSSM, or $SU(3)_c \\times SU(2)_L\\times U(1)_Y\\times U(1)_{B-L}$. In either cases the inflaton can be made out of the right handed sneutrino, the Higgs and a slepton, while the extra vacuum energy during inflation is provided by the Higgs field associated with the singlet or the $U(1)_{B-L}$ and coupled to the right-handed neutrinos, which we assume to be at its false vacuum. Once the slow-roll inflation ends, this extra Higgs would settle down to its true minimum. At the same time, the right-handed majorana neutrinos become massive. In this case, there is virtually no fine-tuning of the soft SUSY breaking parameters, as we have discussed at the end of Sect. \\ref{removeft}. However, as pointed out, the gauge coupling of the $U(1)_{B-L}$ extension should be very small so that radiative corrections do not to ruin the flatness of the potential. Therefore, gauge coupling unification of $U(1)_{B-L}$ with $SU(3)_c \\times SU(2)_L\\times U(1)_Y$ appears not to be feasible, but of course as such this is no compelling argument against the inflection point inflation. Whether a $SU(3)_c \\times SU(2)_L\\times U(1)_Y\\times U(1)_{B-L}$ model with small $g^2_{B-L}$ can be naturally constructed remains an open problem."
    },
    "1004/1004.2123_arXiv.txt": {
        "abstract": "HH~110 is a rather peculiar Herbig-Haro object in Orion that originates due to the deflection of another jet (HH~270) by a dense molecular clump, instead of being directly ejected from a young stellar object. Here we present new results on the kinematics and physical conditions of HH~110 based on Integral Field Spectroscopy. The 3D spectral data cover  the whole outflow extent ($\\sim$~4.5~arcmin, $\\simeq$ 0.6 pc at a distance of 460~pc) in the spectral range 6500--7000 \\AA.  We built emission-line intensity maps of H$\\alpha$, \\nii\\ and \\sii\\ and of  their radial velocity channels. Furthermore, we analysed the spatial distribution of the excitation and electron density  from \\nii/H$\\alpha$, \\sii/H$\\alpha$, and \\sii 6716/6731 integrated line-ratio maps, as well as  their behaviour as a function of velocity, from line-ratio channel maps. Our results  fully reproduce the morphology and kinematics obtained from previous imaging and long-slit data. In addition, the IFS data revealed, for the first time, the complex spatial distribution of the physical conditions (excitation and density) in the whole jet, and their behaviour as a function of the kinematics. The results here derived give further support to the more recent model simulations that involve deflection of a pulsed jet propagating in an inhomogeneous ambient medium. The IFS data  give  richer information than that provided by  current model simulations  or laboratory jet experiments. Hence, they could provide valuable clues to constrain the space parameters in future theoretical works. ",
        "introduction": "HH~110 is a Herbig-Haro (HH) jet emerging from the southern edge of the L1617 dark cloud in the Orion B complex. Both observational and theoretical works have been carried out since its discovery by \\citet{Rei91}, because unlike most of the well-known stellar jets, HH~110 has a peculiar morphology, among other properties. The morphology of HH~110 in the H$\\alpha$ and \\sii\\ lines is very complex, starting in a collimated chain of knots.  The emission away from the collimated region has a more chaotic structure and widens within a cone of unusually large opening angle ($\\sim 10^\\circ$). HH~110 appreciable wiggles along the $\\sim$ 4~arcmin jet length. The knots are all embedded in  fainter emitting gas, which outlines the whole flow, more reminiscent of a turbulent outflow. Most of the knots detected in ground-based images are spatially resolved into several components in the higher spatial resolution HST images (see \\eg\\ \\citealp{Har09}). The transverse cross-section of HH~110 shows a significant asymmetry, the eastern border is sharp and poorly resolved, whereas the strong knotty emission mostly appears towards the western side. HH~110 is the only known HH that shows a faint filament of emission lying parallel at $\\sim$ 10 arcsec to the east of the northern A-C knots, which probably represents a weak secondary or even a fossil flow channel \\citep{Rei91}. Attempts to find the driving source of HH~110 have failed at optical, near infrared and radio continuum wavelengths.  A scenario that accounts for the singular morphology of HH~110 was first outlined by  \\citet{Rei96}. These authors proposed that HH~110 originates from the deflection (deflection angle $\\simeq 60^\\circ$) of the adjacent HH~270 jet through a grazing collision with a dense molecular clump of gas. The feasibility of this scenario has been further reinforced from the results of numerical simulations, which  model the emission arising from the collision of a jet with a dense molecular clump (see \\citealp{Rag02} and references therein). In addition, analysis of further data are also in agreement with this scenario.  First, a high-density clump of gas around the region where the HH~270 jet changes its direction to emerge as HH~110 has been detected through high-density tracer molecules (in HCO$^+$, by \\citealp{Cho01}, and in NH$_3$, by \\citealp{Sep10}). Second, results from proper motion determination are also consistent with the jet/cloud collision scenario (\\citealp{Rei96}; \\citealp{Lop05}). A recent work of \\citet{Har09} that included laboratory experiments has given further support to the jet/cloud collision scenario.  HH~110 is a good candidate to search for the observational footprints of gas entrainment and turbulence by analysing the kinematics and the excitation conditions along and across the jet flow. Some works were performed in the recent past from long-slit spectroscopy and Fabry-Perot data. The kinematics and physical conditions, both along the outflow axis and at four positions across the jet beam, were explored from long-slit spectroscopy by \\citet{Rie03a}, who found very complex structures. \\citet{Rie03b} explored the spatial distribution and the characteristic knot sizes, as well as the spatial behaviour of the velocity and line width, by performing a wavelet analysis of Fabry-Perot data, but only covering the H$\\alpha$ line. Their results indicated that most of the H$\\alpha$ kinematics can be explained by assuming an axially peaked mean flow velocity, on which are superposed low-amplitude turbulent velocities. In addition, their results are suggestive of the presence on an outer envelope that appears to be a turbulent boundary layer. Finally, \\citet{Har09} compared images from laboratory jet experiments with numerical simulations and with long-slit, high-resolution optical spectra obtained along HH~110. They found a good agreement between the shock structures observed in HH~110 and those derived from experiments of a supersonic jet deflected by a dense obstacle.  In order to further advance on our understanding of HH~110, this object was included as a target within a program of Integral Field Spectroscopy of Herbig--Haro objects using the Potsdam Multi-Aperture Spectrophotometer (PMAS) in the wide-field IFU mode PMAS fibre PAcK (PPAK). The data obtained from this IFS HH observing program give a full spatial coverage of the HH~110 emission in several lines (H$\\alpha$, \\nii\\ and \\sii), thus allowing us to perform a more complete analysis of the kinematics and physical conditions through the whole flow than all the previous works. The main results are given in this work. The paper is organized as follows. The observations and data  reduction are described in \\S\\ 2.  Results are given in \\S\\ 3: the analysis of the physical conditions in \\S\\ 3.1 and the analysis of the kinematics  in \\S 3.2. A summary with the main conclusions is given  \\S\\ 4.   \\begin{figure} \\centering \\includegraphics[width=\\hsize]{pointings.eps} \\caption{CCD image of HH~110 obtained at the Nordic Optical Telescope (NOT) through a narrow-band filter that includes the \\sii\\ 6716, 6731 \\AA\\ lines (see \\citealp{Lop05} for details). The four IFU pointing fields are superposed on the image. The knots discussed throughout this work are labeled on the image, according to the nomenclature first established by \\citealp{Rei91}). The contours of the individual fibres have been drawn on the southern IFU field. \\label{pointings}} \\end{figure}   \\section[]{Observations and Data Reduction}   Observations of HH~110 were made on 22 November 2004 with the 3.5-m  telescope of the Calar Alto Observatory (CAHA). Data were acquired with the  Integral Field Instrument Potsdam Multi-Aperture Spectrophotometer PMAS \\citep{Rot05} using the PPAK configuration that has 331 science fibres, covering an hexagonal FOV of 74$\\times$65 arcsec$^2$ with a spatial sampling of 2.7 arcsec per fibre, and 36 additional fibres to sample the sky (see Fig.\\ 5 in \\citealt{Ke06}). The I1200 grating was used, giving an effective sampling  of 0.3 \\AA\\ pix$^{-1}$  ($\\sim15$  km~s$^{-1}$ for H$\\alpha$) and covering the wavelength range $\\sim6500$--7000 \\AA, thus including characteristic HH emission lines in this wavelength range (H$\\alpha$, \\nii\\ $\\lambda\\lambda$6548, 6584 \\AA\\  and \\sii\\  $\\lambda\\lambda$6716, 6731 \\AA). The spectral resolution (\\ie\\ instrumental profile) is $\\sim2$ \\AA\\  FWHM ($\\sim90$ km~s$^{-1}$) and the accuracy in the determination of the position of the line centroid is $\\sim0.2$ \\AA\\  ($\\sim10$ km~s$^{-1}$ for the  strong observed emission lines). Four overlapped pointings were observed to obtain a mosaic of $\\sim$ 5\\arcmin $\\times$ 1$\\farcm$5 to cover  the entire emission from HH~110 (see Fig.\\ \\ref{pointings}). Table 1  lists the centre positions of each  pointing, the exposure time and  the HH~110 knots included in each pointing, following the nomenclature of \\citet{Rei96}.    Data reduction was performed using a preliminary version of the {\\small R3D} software \\citep{Sa06}, in combination with {\\small IRAF}% \\footnote{{\\small IRAF} is distributed by the National Optical Astronomy Observatories, which are operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation.} and the {\\small Euro3D} packages \\citep{Sa04}. The reduction consists of the standard steps for fibre-based integral field spectroscopy. A master bias frame was created by averaging all the bias frames observed during the night and subtracted from the science frames. The location of the spectra on the CCD was determined using a continuum-illuminated exposure taken before the science exposures. Each spectrum was extracted from the science frames by co-adding the flux  within an aperture of 5 pixels along the cross-dispersion axis for each pixel in the dispersion axis, and stored in a row-stacked-spectrum (RSS) file \\citep{Sa04}. The wavelength calibration was performed  by using the sky emission lines found in the observed wavelength range. The  accuracy achieved for the wavelength calibration was better than  $\\sim0.1$  \\AA\\ ($\\sim5$ km~s$^{-1}$). Furthermore, the large scale diffuse emission from this Orion region was subtracted from our data by using the signal acquired through the 36 additional sky fibres. Observations of a standard star were used to perform a relative flux calibration. The four IFU pointings were merged into a mosaic using our own routines, developed for this task (see \\eg\\ \\citealp{San07} and references therein). The procedure is based on the comparison and scaling of the relative intensity in a certain wavelength range for spatially coincident spectra. The pointing precision is better than 0.2 arcsec, according to the pointing accuracy of the telescope using the Guiding  System of PMAS in relative offset mode (used for these observations). The overlap between adjacent pointings is $\\sim$ 10--15\\%  the field of view (more than 30 individual spectra). The merging process was checked to have little effect on the  accuracy of the wavelength calibration of the final datacube making use of the sky lines (\\ie\\ a common reference system) present in the spectra. A final datacube  containing the 2D spatial plus the spectral information of HH~110 was then created from the 3D data by using {\\small Euro3D} tasks  to interpolate the data spatially until reaching a final grid of 2 arcsec of spatial sampling and with a spectral sampling of 0.3 \\AA\\ . Further manipulation of this datacube, devoted to obtain  integrated emission line maps, channel maps and position--velocity maps, were made using several common-users tasks of {\\small STARLINK}, {\\small IRAF} and {\\small GILDAS} astronomical packages and {\\small IDA} \\citep{Ga02}, a specific {\\small IDL} software to analyse 3D data.   \\begin{figure*} \\centering \\includegraphics[angle=-90,width=0.95\\hsize]{flux.eps} \\caption{ Integrated maps of the HH~110 emission obtained from the datacube by integrating the signal within the wavelength range including the line labeled in each panel: (H$\\alpha$ : $\\lambda$6558.3--6565.9 \\AA, \\nii\\ :\t    $\\lambda$6579.6--6585.8 \\AA, \\sii\\  $\\lambda$6716:  6712.8--6718.1 \\AA\\  and \\sii\\  $\\lambda$6731:  6727.2--6733.1 \\AA ). Fluxes have been displayed in a logarithmic scale, in  units of 10$^{-16}$ erg~s$^{-1}$~cm$^{-2}$. In all the maps, north is up and East is to the left. The HH~110 knots and the spatial scale have been labeled in the H$\\alpha$ panel (30 arcsec $\\simeq$ 0.07 pc for a distance of 460 pc; \\citealp{Rei91}). \\label{flux}} \\end{figure*}  \\begin{table} \\caption{Positions of the single pointings of the HH~110 mosaic. \\label {tjcmt}} \\begin{tabular}{@{}lccc@{}} \\hline \\multicolumn{2}{c}{Position}&Exp. time(s)&knot \\\\ \\hline 05$^{\\rm h}$51$^{\\rm m}$25$\\fs3$&+02$\\degr$55$\\arcmin$16$\\arcsec$&1800&A-E \\\\ 05$^{\\rm h}$51$^{\\rm m}$24$\\fs4$&+02$\\degr$54$\\arcmin$20$\\arcsec$&1800&I-N  \\\\ 05$^{\\rm h}$51$^{\\rm m}$23$\\fs8$&+02$\\degr$53$\\arcmin$41$\\arcsec$&2400&N-R  \\\\ 05$^{\\rm h}$51$^{\\rm m}$24$\\fs9$&+02$\\degr$52$\\arcmin$40$\\arcsec$&2700&S    \\\\ \\hline \\end{tabular} \\end{table} ",
        "conclusions": "To conclude, the IFS data confirm the peculiar nature of HH~110 as compared with most of the stellar jets already observed. The most relevant results  presented in this paper might be summarized as follows:   \\begin{itemize}  \\item These IFS data have allowed us to generate the first \\nii\\ narrow-band  image of HH~110 currently obtained, which shows a  similar jet  cross section to the \\sii\\ emission. In addition, these data allow us to confirm that  the faint filamentary emission labeled X and Y, detected to the east of knots A--C, which probably corresponds to a ''fossil'' outflow channel, mainly should arise from ''true'' H$\\alpha$ line emission. The contributions  from nearby continuum or from the \\nii\\ lines, if any, should be very weak, as no emission was detected  in \\nii, and the continuum contribution was previously removed.  \\item The line-ratio maps derived from the integrated line intensities revealed  complex spatial structures for the gas excitation and  electron density. Although the line ratios that trace gas excitation are indicative of an intermediate/high-excitation degree for most of the emission mapped, several regions of low-excitation emission are detected through the flow (\\eg\\ \\sii/H$\\alpha$~$\\geq$~1.5 around knot H). In  general, the \\sii\\ line ratios correspond to low electron density values ($n_\\rmn{e}$~$\\leq$~1000~cm$^{-3}$) and show smooth spatial variations. However, strong local enhancements of $n_\\rmn{e}$ relative to its surroundings were detected at several locations (\\eg\\ around knot E, where the $n_\\rmn{e}$ value is about three times that of the surroundings).  \\item  Radial velocities appear blueshifted, relative to the \\vlsr\\ of the cloud, the modulus increasing with the distance to knot A, although with some appreciable  local departures from this trend (in particular, knot K).  This systematic increase in radial velocity should not be interpreted as  a sign of jet acceleration, since no similar trend was derived for proper motions. In fact, the trend found for tangential velocities goes in the opposite way (decreasing velocity as a function of distance from knot A, with some appreciable local departures, \\eg\\ in knot B).  Thus, this apparent acceleration should be better attributed to a geometric effect: the interaction between the jet and the inhomogeneous ambient surroundings will lead to the deflection of the jet  and therefore, to changes in  the projection of the full spatial velocity along the line of sight (radial velocity) and along the plane of the sky (tangential velocity).  \\item In general, the line profiles are broader for H$\\alpha$ than for \\sii, and the FWHM in both  lines decreases from the northern to the southern knots, although  this trend is not fully systematic along the outflow.  \\item  We recalculated the full velocity and the  motion direction relative to the plain of sky of the HH~110 knots using  the \\sii\\ radial velocities, derived from the spectra extracted  at the position of the knots peak intensity, together with knot proper motions, obtained from ground-based, multi-epoch \\sii\\ narrow-band images.  We confirmed the general trend  of a decrease in the velocity and an increment in the projection angle with distance from knot A, as  found in previous works  that used a more limited knot sample.  Significant departures from the trend are found in some knots. These knots could be tracing the loci where a stronger interaction between the jet and the inhomogeneous medium is taking place.  \\item  The channel maps showed that the emission spreads over different velocity ranges, depending on the line considered. In spite of that, the maps also showed that the morphology of the emission in each of the   velocity channels is  similar for all the lines, the strongest emission corresponds to the velocity interval $\\simeq$--20$<$\\vlsr$<$+5~\\kms, where significant emission from all the knots  was detected. In addition,  the previously unknown kinematics of the excitation and electron density along the full spatial extent of the outflow were obtained from the channel maps.  The degree of excitation and the electron density appear to be anticorrelated with the velocity modulus, in the way that the emission from the gas moving faster (either, at red or blueshifted velocities) has higher excitation and lower density as compared to the emission from the  gas moving at slower velocities, which appear denser and less excited.  \\item From the IFS spectra extracted at appropriate spaxels, we were able to compare the physical conditions obtained from these data with  those derived from previous long-slit and Fabry-Perot observations. The general trends followed by the properties of the emission and even the derived values of the observables are in good agreement at all  the commonly sampled positions.  However, the complex structure of the HH~110 kinematics and physical conditions,  already outlined in previous works, has been now better  characterized from this more complete IFS sampling.  These data revealed significant variations of the kinematics and physical conditions  over short distances that were not properly sampled  by the long-slit data.  \\item The emission lines have asymmetric profiles. The analysis of the H$\\alpha$ line profiles performed through the line bisectors method shows that the departure of the profile from a single-Gaussian shape  varies from knot to knot. We found that the shape of the profiles seems to be related to the full velocity and the viewing angle of the knot. \\end{itemize}  The scenario first outlined by \\citet{Rei96} (\\ie\\ deflection of the HH~270 jet by a dense molecular cloud) and later modeled by \\citet{Rag02} gives a reasonable explanation on the origin of the HH~110 jet and offers a good qualitative fitting for a set of properties observed in the jet  behaviour (\\eg\\  the strong lack of symmetry, relative to the outflow axis, of the proper motions), but fail in reproducing in more detail the kinematics and some other properties through the outflow. In particular, as  was already discussed by \\citet{Lop05}, the jet/cloud collision models of \\citet{Rag02} did not predict the deceleration of the full spatial velocity along the outflow that is derived from imaging, long-slit and IFS data.  A variability in the ejection outflow velocity (\\ie\\ pulsed jet models) and/or an anomalously strong interaction between the outflow and the inhomogeneous environment were then suggested as  plausible alternatives giving rise to such kinematic behaviour. Interestingly, new modeling that incorporates these mechanisms has been recently published.   \\citet{Yir08} and \\citet{Yir09} performed more complex, two-dimensional hydrodynamical simulations of jets that propagate through a small-scale inhomogeneous environment (\\eg\\ clumps with a size smaller  than the jet beam). Their simulations predict that the collisions between slow clumps overtaken by faster ones should produce shock structures slightly smaller than the jet beam and displaced from the jet axis, like what is observed at some locations along HH~110. These simulations also predict a complex evolution of these shock structures. In particular,  properties such as excitation and velocity profiles across the jet beam, are far more complex than expected from an unperturbed clump, or from the internal working surfaces arising from  current models of a pulsed jet.  Further insights on the behaviour of deflected supersonic jets were recently obtained by \\citet{Har09} from laboratory  experiments. The experiments show how the morphology of the contact discontinuity between the jet and the obstacle develops  a complex structure of cavities that gives rise to the entrainment of clumps of material from the obstacle into the flow. The experiments also succeeded in reproducing the morphology of some filaments such as those observed in high-spatial resolution (HST)  images of HH~110, although the observed dynamics in these filamentary structures were not properly reproduced by the experiments. According to the results of these laboratory experiments, as compared to their high-resolution, long-slit data, the authors proposed that the best model for HH~110 still is that of a pulsed (from a varying driving source) jet interacting with a dense molecular clump.  The results derived from our IFS observations that can be compared  with those found from these authors (\\eg\\ the velocity behaviour through the flow and the line profiles at selected positions)  appear well consistent, thus giving further support to models that involve deflection of a pulsed jet propagating through an inhomogeneous ambient medium.   The IFS data here presented give a full spatial coverage in the H$\\alpha$, \\nii\\ and \\sii\\ emission lines of the singular outflow HH~110. They have allowed us to explore for the first time the whole spatial distribution of the physical conditions and its relationship with the kinematics of the jet emission. We would like to point out that there are very few IFS observations of stellar jets with a wide spatial coverage of the outflow in several emission lines (as \\eg\\ HH~34 by \\citealp{Be07}). Such observations are highly necessary to understand the behaviour of the physical conditions as a function of the kinematics, as well as to explore whether this behaviour varies (as in HH~34) or not (as in HH~1) with the distance from the exciting source  (see \\citealp{Gar09}). Unfortunately, the paucity of available data prevents  to establish a general picture on these topics.   Because of the rather chaotic behaviour of HH~110 outlined in previous works, as derived from more limited narrow-band  imaging and long-slit spectroscopic data, the IFS observations discussed here are particularly useful for characterizing the properties of the whole outflow. Hence, a more realistic picture  has arisen, suitable for designing   new state-of-the-art simulations to match the HH~110 scenario.  Note that IFS data give much more detailed  information  on the spatial distribution of excitation as a function of velocity  than  provided by current model simulations or laboratory jet experiments.  These data could provide valuable clues  to constrain the space parameters in future theoretical works, which are necessary to understand the origin, structure and dynamics of HH~110."
    },
    "1004/1004.5535_arXiv.txt": {
        "abstract": "We present results from a joint X-ray/Sunyaev-Zel'dovich modeling of the intra-cluster gas using XMM-Newton and APEX-SZ imaging data. The goal is to study the physical properties of the intra-cluster gas with a non-parametric de-projection method that is, aside from the assumption of spherical symmetry, free from modeling bias. We demonstrate a decrease of gas temperature in the cluster outskirts, and also measure the gas entropy profile, both of which are obtained for the first time independently of X-ray spectroscopy, using Sunyaev-Zel'dovich and X-ray imaging data. The contribution of the APEX-SZ systematic uncertainties in measuring the gas temperature at large radii is shown to be small compared to the XMM-Newton and Chandra systematic spectroscopic errors. ",
        "introduction": "Accurately determining the thermodynamic state of the intra-cluster medium (ICM) out to large radii in galaxy clusters is critical for understanding the link between the total cluster mass and X-ray observables. For over a decade, observations of the thermal Sunyaev-Zel'dovich Effect (tSZE, hereafter simply SZ; Sunyaev \\& Zel'dovich 1970) have been considered as a promising complement to X-ray observations for modeling the ICM in galaxy clusters, yet only recently has it been possible to make meaningful de-projections of gas temperature and density profiles using  SZE imaging data from multi-pixel bolometer arrays. The APEX-SZ experiment (Dobbs et al. 2006, Halverson et al. 2009) employs one of the first such powerful multi-pixel Transition-Edge Sensor (TES) bolometer cameras, enabling a joint analyses of the ICM properties using SZE and X-ray data. The first results of such analysis have been published by Nord et al. (2009) for the cluster Abell 2163, and Basu et al. (2010) for the prototypical relaxed cluster Abell 2204. ",
        "conclusions": "The potential for joint X-ray/SZ analysis of the ICM properties is shown with the help of APEX-SZ data, which allows for a non-parametric modeling of the temperature and density profiles out to the cluster virial radius. Apart from the assumption of spherical symmetry, the de-projection method is free from modeling biases. The uncertainties in the gas temperature measurements near the virial radii is dominated by the statistical uncertainties in the SZ measurements. A demonstration of the decreasing gas temperature in the cluster outskirts, and also the measurement of gas entropy profiles, are made from the SZ and X-ray imaging data without the use of X-ray spectroscopy.  \\begin{theacknowledgments}  The authors thank H. Eckmiller, V. Jaritz, T. H. Reiprich, and the members of the APEX-SZ collaboration, for their contributions to this work.  \\end{theacknowledgments}"
    },
    "1004/1004.0310_arXiv.txt": {
        "abstract": "In the future several Spallation Source facilities will be available worldwide. Spallation Sources produce large amount of neutrinos from decay-at-rest muons and thus can be well adapted to accommodate state-of-the-art neutrino experiments. In this paper low energy neutrino scattering experiments that can be performed at such facilities are reviewed. Estimation of expected event rates are given for several nuclei, electrons and protons at a detector located close to the source. A neutrino program at Spallation Sources comprises neutrino-nucleus cross section measurements relevant for neutrino and core-collapse supernova physics, electroweak tests and lepton-flavor violation searches. ",
        "introduction": "Spallation sources can be well adapted to accommodate innovative neutrino experiments. At such installation large amounts of neutrinos are produced from the decay of pions, i.e. $\\pi^+ \\rightarrow \\mu^+ + \\nu_{\\mu} $ and the subsequent decay of muons $\\mu^+ \\rightarrow \\e^+ + \\nu_e + \\bar{\\nu}_{\\mu}$; while, due to the strong anti-pion absorption in the spallation target, very few electron anti-neutrinos evolve via the charge-conjugate pion decay sequence (typically at the level of 10$^{-5}$). Spallation sources have already been used in the past to study neutrino oscillations by the LSND~\\cite{Athanassopoulos:1996jb,Athanassopoulos:1997pv,Athanassopoulos:1997er} and KARMEN \\cite{Armbruster:2002mp} collaborations at LANSCE in Los Alamos and ISIS at Rutherford Appleton Laboratory respectively. The LSND (KARMEN) experiment has used a 167 (65) tons scintillator detector located at about 17 (17.6) meters from the LANSCE (ISIS) source. In parallel the same experimental setups have been exploited to perform a variety of experiments including measurements of neutrino scattering on $^{12}$C~\\cite{Albert:1994xs,Athanassopoulos:1997rn,Athanassopoulos:1997rm,Krakauer:1991rf,Allen:1990nr,bib:nuC12,Bodmann:1994py,Drexlin:1991gx} $^{13}$C, and $^{56}$Fe~\\cite{Zeitnitz:1998qg} as well as electroweak tests (lepton flavor universality~\\cite{Bodmann:1994py}, V-A structure in muon decay~\\cite{Armbruster:1998qi}, Weinberg angle \\cite{Auerbach:2001wg}) and lepton-flavor violation searches~\\cite{Armbruster:2003pq}. New spallation sources are being constructed, are planned or under study including the Spallation Neutron Source (SNS) facility at Oak Ridge~\\cite{sns}, the European Spallation Source (ESS) facility in Lund, the Japanese Spallation Neutron Source (JSNS) at JPARC~\\cite{future} and the SPL at CERN~\\cite{Bandyopadhyay:2007kx}. They will offer the possibility to realize a low energy neutrino physics program of interest for particle physics\\footnote{Note that in \\cite{Conrad:2009mh} a new idea for the search of CP violation in the lepton sector is proposed, using spallation sources.}, neutrino astrophysics and nuclear physics. \\begin{figure}[t] \\begin{center} \\includegraphics[scale=0.7,angle=0]{dar_flux.EPS} \\caption{Neutrino fluxes from muon and pion decay-at-rest at a spallation source.} \\end{center} \\label{fig:flux} \\end{figure} The only technical alternative to produce controlled neutrino fluxes in the 100 MeV energy range is a low energy beta-beam facility~\\cite{Volpe:2003fi}, which is based on the novel method of the beta-beams~\\cite{Zucchelli:2002sa} exploiting the beta-decay of boosted radioactive ions.  While the main goal of the beta-beam  is the search of CP violation in the lepton sector, the physics program of the low energy variant covers a variety of interesting topics in nuclear physics, in the study of fundamental interactions and core-collapse supernova physics (for a review of beta-beams see~\\cite{Volpe:2006in}). An advantage of the beta-beam facility is in its capacity to produce collimated beams of both $\\nu_e$ and $\\overline{\\nu_e}$ species, as well as the possibility they provide to control the average energy of the neutrino beams. On the other hand spallation sources have significantly larger neutrino production rates.  Determination of neutrino cross sections in a broad range of energies (several hundred MeV to multi-GeV range) is currently the focus of several experiments including Miner$\\nu$a~\\cite{minerva} and Sciboone~\\cite{sciboone}. Note that recently the MiniBOONE results revealed an excess of the electron-like events for the several hundred MeV neutrino energy range~\\cite{AguilarArevalo:2008rc}. The neutrino scattering experiments that can be performed at Spallation Sources, where the electron neutrino energy distribution is described by the Michel spectrum with the endpoint at about 52.8 MeV (Figure~\\ref{fig:flux}), turns to be complementary, as they test nuclei at low energy.  In the high (low) energy range the nucleon (nuclear) degrees of freedom are involved.  The neutrino low energy range is particularly attractive for timely applications, such as the precise calibration of neutrino detectors for the future observation of neutrinos from a core-collapse supernova explosion, or of the diffuse supernova neutrino background.  While electron anti-neutrinos can be detected through proton scattering in water Cherenkov and scintillator detectors,  neutrino-nucleus scattering is necessary for the identification of the electron neutrinos. Neutrino detectors running or under study exploit nuclei. For example, a lead-based supernova observatory -- the HALO project -- is now planned at SNOLAB. It has been shown e.g. in~\\cite{Engel:2002hg} that the measurement of neutrino-scattering on lead, in coincidence with (1 or 2) neutrons, has an interesting sensitivity upon the third neutrino mixing angle. Large-size detectors, such as  GLACIER or MEMPHYS (UNO, Hyper-K) and LENA, are currently under study~\\cite{Autiero:2007zj}. In these observatories the $\\nu_e$ of the diffuse supernova neutrino background can be measured through scattering on argon~\\cite{Cocco:2004ac}, on oxygen and on carbon~\\cite{Volpe:2007qx}. Measurements of neutrino-nucleus cross sections are not only necessary to determine precisely detectors responses but for several other applications. In~\\cite{Haxton:1987bf} an original strategy for the detection of relic galactic supernova neutrinos is proposed, based on the geochemical measurement of $^{97}$Tc in $^{98}$Mo ore. As pointed out in~\\cite{Lazauskas:2009yh} a precise measurements of the $\\nu$-$^{98}$Mo and $\\nu$-$^{97}$Tc cross sections is necessary to extract unambiguously the supernova contribution. Neutrino-nucleus cross sections on stable and radioactive nuclei are one of the observables needed to understand stellar nucleosynthesis and, in particular, the r-process  (see e.g.~\\cite{Meyer:1998sn,Surman:1998eg}). Unravelling its site is one of the major open questions in nuclear astrophysics, one of the candidates being core-collapse supernova explosions (see e.g.~\\cite{Balantekin:2003ip}).  A precise knowledge of the nuclear response to neutrinos is also crucial in order to calibrate some phenomenological ingredients (e.g. knowledge of the forbidden multipoles and their possible quenching) of the microscopic approaches currently used to study neutrinoless double-beta decay~\\cite{Ejiri:2003ap,Volpe:2005iy}. The observation of neutrinoless double-beta decay would represent a major discovery, from which key information on the CP Majorana phases, on the electron neutrino effective mass and mass hierarchy can be extracted \\cite{Elliott:2004hr,Bilenky:2002aw}. However this requires reducing the current discrepancies among the half-life predictions. In this respect a step forward can be made through neutrino-nucleus measurements, combined with other measurements on the candidate emitters (beta-decay, muon capture, charge-exchange reactions and the two-neutrino double beta decay) \\cite{Zuber:2005fu}.  The theoretical description of neutrino-nucleus cross sections in the low energy range benefits from a variety of sophisticated models, including Effective Field Theories~\\cite{Kubodera:1993rk,Kubodera:2009au}, the Continuum Random-Phase-Approximation (CRPA)~\\cite{Kolbe:2003ys,Jachowicz:2002rr}, the Quasi-particle RPA (QRPA)~\\cite{Volpe:2000zn} and projected QRPA~\\cite{Samana:2008pt}, relativistic RPA~\\cite{Paar:2007fi}, the Shell Model (SM)~\\cite{Volpe:2000zn,Hayes:1999ew} and the Shell Model in the complex energy plane~\\cite{Civitarese:2007ht}. Regardless the high degree of sophistication achieved, important discrepancies remain between the predicted neutrino-nucleus cross sections. For example, in ~\\cite{Samana:2008pt} it is shown that the calculations of the energy dependent cross section on $^{56}$Fe differ as much as by factor 3-4 at a given energy, whereas the convoluted neutrino cross sections turn to be in agreement with the experimental data\\footnote{Note that the experimental uncertainties are in this case at the level of 50\\%.~\\cite{Zeitnitz:1998qg}}~\\cite{Zeitnitz:1998qg}. In ~\\cite{Paar:2007fi} the comparison of the flux-averaged neutrino-lead cross sections, associated to supernova neutrinos, show very similar discrepancies.  The realization of precise measurements of the neutrino scattering cross sections for an ensemble of nuclei therefore should help to pin down differences among the models, enabling accurate and reliable description of the isospin and spin-isospin nuclear response. This is indeed a mandatory  step for many innovative and timely applications.  \\begin{table}[tbp] \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|} \\hline \\multicolumn{1}{|c|}{Facility} & \\multicolumn{1}{|c|}{Power} & \\multicolumn{1}{|c|}{Proton energy} & \\multicolumn{1}{|c|}{Time structure} & \\multicolumn{1}{|c|}{Repetition rate} \\\\ \\hline LANSCE (USA) & 56 kW & 0.8 GeV & Continuous & N/A \\\\ ISIS (UK)  & 160 kW & 0.8 GeV & 200 ns & 50 Hz \\\\ SNS (USA) & $>$ 1 MW & 1 GeV & 380 ns & 60 Hz \\\\ JSNS (Japan)  & 1 MW & 3 GeV & 1 $\\mu$s & 25 Hz \\\\ SPL (CERN) & 4 MW & 3.5 GeV & 0.76 ms & 50 Hz \\\\ ESS (Sweden) & 5 MW & 1.3 GeV & 2 ms (1.4 $\\mu$s) & 17 Hz (50 Hz) \\\\ \\hline \\end{tabular}% \\end{center} \\par \\vskip 0.5cm \\caption{Comparison of characteristics of the past, present, and future Spallation Source Facilities in different regions of the world\\label{t:facilities}.} \\end{table}    Coherent neutrino-nucleus scattering is another important phenomenon that has never been measured. Such an experiment is useful as a test for the Standard Model but also as a technological advancement for supernova neutrino detectors and dark matter searches.  The cross section in the 50 MeV energy range can be as large as 10$^{-39}$cm$^{2}$. The CLEAR experiment is now planned at SNS to measure coherent neutrino-nucleus scattering~\\cite{Vergados:2009ei,Scholberg:2009ha}. Furthermore, neutrino capture on radioactive nuclei would open a completely new window on the Universe since this process could be used to observe the cosmological neutrino background. This exciting possibility has been pointed out in~\\cite{Cocco:2007za} and further investigated in~\\cite{Lazauskas:2007da}. Note that, since this background has a temperature of 1.96 K, the impinging neutrino energy is so low, that one can calibrate the corresponding cross sections using the inverse process (the beta-decay of the nuclei of interest).  Detailed investigations of the physics scope of spallation source facilities as far as neutrino physics, neutrino astrophysics, tests of the Standard Model are concerned, can be found in the literature (see e.g.~\\cite{Avignone:2003ep}). Ref.~\\cite{sns} is a proposal made for the SNS facility presenting also an in-depth study of the possible backgrounds and of the detectors' design. In this paper, motivated by the recently approved European Spallation Source facility, we present new predictions for neutrino-nucleus, neutrino-proton and electron scattering rates. We emphasize general aspects, relevant for the cross section measurements and for low energy tests of the Standard Model, in particular a Weinberg angle measurement and Lepton-Flavour Violation (LFV) searches. ",
        "conclusions": "New Spallation Source facilities  are under construction, planned or under study. They will produce an intense flux of neutrinos. If these facilities are designed to accommodate a neutrino detector(s) close to the source, they will offer an unique opportunity to perform low energy neutrino scattering experiments. Neutrino detectors with different active material can be used to perform neutrino nucleus cross section measurements that have a variety of innovative and timely applications going from nuclear physics to neutrino astrophysics. Neutrino-nucleus scattering data would create conditions for the breakthrough of the theoretical models, permitting to pin  down some of their phenomenological ingredients. Besides, physics beyond the standard model can be tested through the searches for non-standard contributions to neutrino electron scattering; while the identification of electron anti-neutrino scattering on protons can be used to improve established limits on rare lepton flavor violating decays, in particular $\\mu^+ \\rightarrow \\e^+ + \\bar{\\nu}_e + {\\nu}_{\\mu}$. We have presented new predictions for these processes and emphasized some general aspects, hoping one day such measurements to be realized.  \\vspace{.3cm} \\noindent {\\bf Acknowledgements} The authors thank G. Drexlin and M. Mezzetto for useful discussions. \\vspace{1.cm}"
    },
    "1004/1004.1066_arXiv.txt": {
        "abstract": "In studies of accreting black holes in binary systems, empirical relations have been proposed to quantify the coupling between accretion processes and ejection mechanisms. These processes are probed respectively by means of X-ray and radio/optical-infrared observations. The relations predict, given certain accretion conditions, the expected energy output in the form of a jet. We investigated this coupling by studying the black hole candidate Swift J1753.5-0127, via multiwavelength coordinated observations over a period of $\\sim 4$ years. We present the results of our campaign showing that, all along the outburst, the source features a jet that is fainter than expected from the empirical correlation between the radio and the X-ray luminosities in hard spectral state. Because the jet is so weak in this system the near-infrared emission is, unusually for this state and luminosity, dominated by thermal emission from the accretion disc. We briefly discuss the importance and the implications of a precise determination of both the slope and the normalisation of the correlations, listing some possible parameters that broadband jet models should take into account to explain the population of sources characterized by a dim jet. We also investigate whether our data can give any hint about the nature of the compact object in the system, since its mass has not been dynamically measured. ",
        "introduction": "\\label{par:intro} Transient black hole candidates (BHCs) are low-mass X-ray binaries that usually show relatively short (weeks to months) outbursts, separated by long periods of quiescence (see Remillard \\& McClintock 2006 for a review). When an outburst starts, BHCs go through a loop in an X-ray hardness-intensity diagram (HID), in which in many cases they draw a q-shape pattern. Four spectral states (Homan et al. 2001, Homan \\& Belloni 2005, Belloni 2009; see McClintock \\& Remillard 2006 for an alternative definition of states) can be identified in the X-ray HID: the low/hard state (LHS), the high/soft state (HSS) and two intermediate states (called hard and soft intermediate state). The LHS of BHCs is characterized by a hard power-law X-ray spectrum (with photon index $\\Gamma \\sim 1.4-2.1$; e.g. Remillard \\& McClintock 2006) which is usually interpreted as the result of comptonization of seed photons by hot electrons (the so called ``corona'', e.g. Esin et al. 1997). Observations suggest that the accretion disc is cold and truncated at large radii (e.g. McClintock et al. 2001, Tomsick et al. 2009, Done \\& Diaz Trigo 2009; but see also Miller et al. 2006b and Reis et al. 2010) at low X-ray luminosities ($L_{X} \\lesssim 1 \\% L_{Edd}$ where $L_{Edd}$ is the Eddington luminosity; Cabanac et al. 2009). The X-ray energy spectrum of BHCs in the HSS is dominated by a thermal emission below $\\sim 5$ keV, likely produced by an optically thick/geometrically thin accretion disc extending to/close to the innermost stable circular orbit (ISCO). The energy spectrum in the HSS also features a steeper power-law tail than in the LHS. The intermediate states are characterized by spectral properties in between the LHS and the HSS. In timing studies, the LHS and the intermediate states show strong quasi periodic oscillations (QPOs) and noise components while the HSS is characterized by weak/absent variability (see van der Klis 2006 and Belloni 2009 for reviews). \\begin{figure*} \\begin{tabular}{c} \\resizebox{15cm}{!}{\\includegraphics{licu_BAT_radio_SMARTS.ps}} \\end{tabular} \\caption{Lower panel: Swift/BAT light curve of Swift J1753.5-0127 for the period 2005 May 30 - 2009 October 16. Bin size is 1 day. The dashed line marks the zero of the {\\it y} axis. Upper panel: radio light curve of Swift J1753.5-0127. We plotted all our data points (VLA, MERLIN and WSRT). Bin size is a whole observation. The horizontal thick line marks the period in which we performed SMARTS OIR observations (H and I band). A complete log of all the radio and OIR observations is reported in the appendix (\\S \\ref{app_tables}).} \\label{fig:BAT_radio_SMARTS} \\end{figure*}  A characterization of the four states needs to take into account the behaviour at longer wavelengths (e.g. Fender, Belloni \\& Gallo 2004; Fender, Homan \\& Belloni 2009). In the LHS a compact, steady jet is on and substantially contributes to the total energy output of the system (see Fender 2006 for a review). The key signature of such a jet is a flat/slightly inverted spectrum with spectral index $\\alpha \\gtrsim 0$ ($S_{\\nu} \\propto \\nu^{\\alpha}$, where $S_{\\nu}$ is the radio flux density at a certain frequency $\\nu$) observed in the radio band (e.g. Hjellming \\& Wade 1971) and extending from radio to higher (sometimes to optical/infrared, OIR) frequencies (Markoff et al. 2003). Following a classical argument presented in Blandford \\& K\\\"onigl (1979), the jet spectrum can be attributed to a superposition of self-absorbed synchrotron spectra from segments of a collimated jet. In the last decade jet-dominated models have been developed (Falcke \\& Biermann 1995, Markoff, Falcke \\& Fender 2001, Giannios 2005, Maitra et al. 2009, Pe'er \\& Casella 2009; but see also Zdziarski et al. 2003, Heinz 2004, Maccarone 2005) and have been used to successfully fit broadband spectra of BHCs in the LHS and in quiescence (e.g. Gallo et al. 2007; Migliari et al. 2007). In this framework the comptonizing medium (the corona) responsible for the hard X-ray power-law tail could actually be the base of the jet (Markoff, Nowak \\& Wilms 2005; but see also Malzac \\& Belmont 2009). There is evidence that the compact jet quenches when the BHC switches to the soft states (Tananbaum et al. 1972; Fender et al. 1999). Intermittent radio emission (characterized by an optically thin spectrum with spectral index $\\alpha < 0$) is occasionally detected in the HSS (Fender et al. 2009) and it could be attributed to the interaction between ejecta detached from the jet base and the interstellar medium. However, the possibility that it originates in a weaker core jet itself can not be completely ruled out without high angular resolution VLBI (Very Long Baseline Interferometry) radio observations.  Not all the BHCs undergo outbursts drawing a q-shape path in the HID: GRS 1915+105 for example spends all its time in the intermediate states at very high X-ray luminosity (see Fender \\& Belloni 2004 for a review), while there is a number of sources that stay for the whole outburst in the LHS (Brocksopp, Bandyopadhyay \\& Fender 2004) or (possibly) in the LHS and in the hard-intermediate state (SAX J1711.6-3808, Wijnands \\& Miller 2002), without transiting to the soft states. Some of these systems underwent both ``normal'' outburst and LHS-only outbursts (e.g. XTE J1550-564, Homan et al. 2001, Belloni et al. 2002). H~1743-322 is the only known BHC that exhibited both ``normal'' outbursts (see e.g. Jonker et al. 2010) and one outburst in the LHS and in the hard-intermediate state only (Capitanio et al. 2009). The possibility of explaining both type of outbursts (with/without a transition to the soft states) is a challenge for theoretical models: both the disc-instability model (DIM, King \\& Ritter 1998, Dubus et al. 2001, Lasota 2001) and the ``diffusive'' model (Wood et al. 2001) expect the outburst to take place in the disc, but contribution from a corona and/or a jet needs to be added to explain the LHS emission and possibly the quiescent one, if the jet indeed plays a substantial role (e.g. Gallo et al. 2006).  To understand the mechanisms governing BHC outbursts it is important to quantify the contribution of the different processes (accretion/ejection) to the total energy output of the system. Hannikainen et al. (1998) and Corbel et al. (2003) found an empirical correlation between the radio flux density and the X-ray flux in the LHS of the BHC GX 339-4. Gallo, Fender \\& Pooley (2003) enlarged the sample including other sources in the LHS, covering more than three orders of magnitude in X-ray luminosity and proposed that the correlation could be universal in both slope and normalisation. Since the radio emission very likely originates in the compact jet and the X-ray emission from the accretion processes, assuming that the two are causally related, then the correlation constitutes a physical relation between the inflow and the outflow. The contribution of the compact jet to the total energy output of BHCs in the LHS sometimes extends to OIR frequencies (e.g. Markoff et al. 2003; in the optical band the emission is dominated by the companion star and the heated/irradiated accretion disc rather than the jet). Russell et al. (2006) tested whether a correlation between the X-ray flux and the OIR flux also holds, in the LHS and in quiescence. They found that a global correlation exists, extending over $\\sim$9 orders of magnitude in X-ray luminosity (2-10 keV).\\\\ Gallo et al. (2006) included in the sample data from the BHC A0620-00 in quiescence. In this way they extended the radio/X-ray correlation from typical LHS levels down to very low X-ray luminosities ($L_X \\sim 10^{-8.5} L_{Edd}$ for a distance to the source of 1.2 kpc), showing that a jet can still be detected at very low accretion rates (with a radio flux at $\\mu$Jy level at 8.5 GHz), suggesting that the physics of the inflow/outflow coupling in the LHS probably still holds in quiescence. The BHC V404 Cyg also produces a jet when in its quiescent state (Gallo et al. 2003; Miller-Jones et al. 2008), although its quiescent X-ray luminosity is rather high ($L_{X} \\sim 5.3 \\times 10^{32}$ erg/s, Corbel, K\\\"{o}rding \\& Kaaret 2008; we considered a distance to the source of $2.39\\pm0.14$ kpc, Miller-Jones et al. 2009), as expected (Lasota 2008) for a long orbital period BHC ($\\sim 155$ hours).  The same X-ray/radio scaling found for BHCs may hold for supermassive black holes in active galactic nuclei (AGN), if the mass of the black hole is taken into account. Merloni, Heinz \\& di Matteo (2003) and Falcke, K\\\"{o}rding \\& Markoff (2004) have independently shown that BHCs and AGN populate a ``Fundamental Plane'' (FP) in the log($L_R, L_X, M$) domain. This suggests that the same mechanisms govern accretion and ejection processes from black holes hold over $\\sim$ 9 orders of magnitude in mass.\\\\ K\\\"{o}rding et al. (2006b) noted that, since the accretion states of BHCs are defined according to the source position on an X-ray HID, a generalization of HID could be constructed also for AGN. In this way they showed that radio loud BHCs and AGN populate the same region of a HID, suggesting that despite the different masses involved, systems that contain a black hole display similar accretion states and jet properties.\\\\ McHardy et al. (2006) showed that BHCs in the HSS and AGN also populate a plane in the space defined by the black hole mass $M$, the accretion rate $\\dot{M}$ and a characteristic frequency $\\nu$ in the X-ray power density spectra. The plane has been extended by K\\\"{o}rding et al. (2007) by also including BHCs in the LHS (considering a constant offset for the frequencies in the two states). This also suggests that some fundamental properties of BHCs and AGN, like characteristic timescales and mass accretion rate, are related in the same way in the two classes of objects, once the mass has been taken into account.  The existence of the radio/X-ray correlation and the FP have broad implications. For example, the small scatter around them has been used as an argument by Heinz \\& Merloni (2004) to infer that jets from BHCs and AGN (once a mass-correction factor is introduced) are characterized by similar bulk velocities. However, in the last years, the supposed universality of the radio/X-ray correlation has been doubted (Xue \\& Cui 2007) and several outliers have been found (Gallo 2007). These sources seem to follow ``normal'' outbursts in the X-rays (their X-ray luminosities are similar to the other BHCs) but they are fainter in radio (at the same X-ray luminosity) than other sources. Corbel et al. (2004) and Gallo (2007) proposed that a correlation with the same slope but a lower normalisation (a factor $\\sim 20$) could describe this discrepancy, at least in a few sources (e.g. in the BHC XTE J1650-500, Corbel et al. 2004). How BHCs accreting at similar rates (and so displaying similar X-ray luminosities) can produce different ejecta has not been clarified yet. Considering the similarities in the inflow/outflow coupling between BHCs and AGN, any explanation should also be relevant for supermassive black holes (possibly as far as helping to explain the apparent radio loud:radio quiet dichotomy, e.g. Sikora, Stawarz \\& Lasota 2007, Tchekhovskoy et al. 2010).  \\subsection{Swift J1753.5-0127} \\label{par_1753} Swift J1753.5-0127 was discovered in the hard X-ray band with the Burst Alert Telescope (BAT) on board Swift on 2005 May 30 (Palmer et al. 2005). The source was also detected in the soft X-rays with the Swift/X-ray Telescope (Swift/XRT, Burrows et al. 2005) and with the Proportional Counter Array (PCA) aboard the Rossi X-ray Timing Explorer (RXTE, Morgan et al. 2005). Figure \\ref{fig:BAT_radio_SMARTS} (bottom panel) shows the Swift/BAT light curve: after the discovery, the source flux reached a peak of 200 mCrab on 2005 July 9 in the All Sky Monitor (ASM, 1.2-12 keV) on board RXTE (Cadolle Bel et al. 2007). Subsequently the flux started decreasing and then stalled at a level of $\\sim$20 mCrab (2-20 keV) for more than $\\sim$ 6 months. This is an unusual behaviour for a transient, but even more unusual is the subsequent flux rise which has been ongoing since roughly June 2006, with a steepening in June 2008 (Krimm et al. 2008), followed by a decreasing/variable trend. No transition to the soft states has been reported: the source has always been in the LHS during the whole outburst (Zhang et al. 2007; Cadolle Bel et al. 2007), although its X-ray spectral hardness has not remained constant (Zhang et al. 2007, Ramadevi \\& Seetha 2007, Negoro et al. 2009). The source was also detected in ultraviolet with the Ultraviolet/Optical telescope UVOT aboard Swift (Still et al. 2005) and in optical with the Michigan-Dartmouth-MIT 2.4 m telescope (Halpern 2005). In July 2005, near the peak of its outburst, Fender, Garrington \\& Muxlow (2005) observed Swift J1753.5-0127 at radio frequencies with the Multi-Element Radio-Linked Interferometer Network (MERLIN), tentatively detecting a point-like radio counterpart, consistent with the presence of a compact jet.\\\\ The strongest hint that the system harbours a black hole comes from the hard power-law tail in the X/$\\gamma$-ray energy spectrum up to $\\sim$600 keV, detected with INTEGRAL (Cadolle Bel et al. 2007), as no low-mass neutron star X-ray binary (NSXB) has ever been detected above $\\sim 200$ keV (di Salvo et al. 2006; Falanga et al. 2007). To date four (or possibly five) high-mass X-ray binaries have been detected at TeV energies (see e.g. Rea et al. 2010 and references therein), although the nature of the accretor is clear only in two of them (Cyg X-1 contains a black hole and PSR B1259-63 harbours a neutron star; see Paredes \\& Zabalza 2010). In all these sources, the spectrum at TeV energies is not compatible with the extrapolation of the hard X-ray spectrum (see e.g. Sidoli et al. 2006 for LS I +61 303).\\\\ The mass of the compact object in Swift J1753.5-0127 has not been dynamically measured. Bearing all these considerations in mind, we will treat Swift J1753.5-0127 as a black hole candidate.\\\\ Cadolle Bel et al. (2007) considered both the high Galactic latitude of Swift J1753.5-0127 and the low Galactic column density in its direction and concluded that its distance should not be larger than 10 kpc, most likely in the $4-8$ kpc range. Using an empirical relation that predicts the absolute magnitude of the accretion disc of a BHC in outburst (from Shahbaz \\& Kuulkers 1998), Zurita et al. (2008) derived a distance to the source $D > 7.2 kpc$. Throughout this paper we will consider a distance $D = 8$ kpc to Swift J1753.5-0127, unless it will be differently specified.  Swift J1753.5-0127 is peculiar for a number of reasons. First of all it is an outlier to the radio/X-ray luminosity correlation of Gallo et al. (2003). According to Cadolle Bel et al. (2007), Swift J1753.5-0127 should be 8-60 times more luminous in radio (depending on the distance, assuming that it must be inside our Galaxy), to fit the correlation. This BHC never left the LHS, so we expect its jet to substantially contribute to the total energy budget of the system, as observed in most of the BHCs in the LHS.\\\\ Secondly, it would be the BHC with the shortest orbital period, as claimed by Zurita et al. (2008) and Durant et al. (2009), who reported a $\\sim$3.2 hr modulation in the optical lightcurves. It is worth to note (Zurita et al. 2008) that such period, if confirmed, would be close to the orbital period of two other BHCs observed only in the LHS, XTE J1118+480 and GRO J0422+32: they respectively feature orbital periods of 4.1 hr (McClintock et al. 2001; Wagner et al. 2001) and 5.1 hr (Filippenko et al. 1995). Interestingly, all these sources are likely located in the Galactic halo (Cadolle Bel et al 2007, Zurita et al. 2008, Durant et al. 2009) and they could constitute a population of high Galactic latitude X-ray binaries, putting constraints on their formation and evolution. Hynes et al. (2009) recently performed coordinated optical and X-ray observations of Swift J1753.5-0127 at high time resolution, reporting a short delay between the optical and the X-ray emission, consistent with an orbital period as short as 3.2 hours.\\\\ From binary evolution calculations nothing prevents BHCs to have in principle orbital periods of $\\sim$2 hours and evolutionary models actually predict that short-period systems might form the majority of them (Yungelson et al. 2006). Nevertheless, such a short period would definitely be peculiar for a BHC (see e.g. Charles \\& Coe 2006). NSXBs usually feature a fainter jet than BHCs (about a factor $\\sim$30 in radio flux; Fender \\& Hendry 2000; Migliari \\& Fender 2006) at the same X-ray luminosity. This could suggest that the compact object in the system is a neutron star and not a black hole, although other evidence points towards a black-hole nature, as mentioned above.  Another interesting aspect concerning Swift J1753.5-0127 is related to the possible presence of an accretion disc extending to the ISCO even if in the LHS (Miller et al. 2006a), at variance with the ``standard'' picture in which the disc is truncated at larger radii (e.g. Esin et al. 1997) in the LHS at low luminosities (Cabanac et al. 2009; Tomsick et al. 2009). The need for a thermal component extending to the ISCO to fit the source spectra has been recently weakened by Hiemstra et al. (2009): they have shown that several spectral models, without necessarily including disc components extending to the ISCO, can fit the data equally well (but see Reis et al. 2009 and Reynolds et al. 2010).  In this paper we present the results of radio, OIR and X-ray observations of Swift J1753.5-0127 performed from the beginning of the outburst (2005 July) until 2009 June (\\S\\ref{par:data}). During this interval, we collected a large sample of (quasi) simultaneous radio/X-ray and OIR/X-ray observations to test whether the jet is systematically less luminous than predicted by the empirical correlations introduced above (\\S \\ref{par:gallo_rel} and \\S \\ref{par:russell_rel}). In 2007 July we also obtained (quasi) simultaneous multiwavelength observations from radio up to hard X-ray frequencies, that we used to produce spectral energy distributions (SEDs, \\S \\ref{par:SED}). Our results will be presented in \\S \\ref{par:results} and discussed in \\S \\ref{par:discussion}. In the last section we will summarize our conclusions (\\S \\ref{par:conclusions}). ",
        "conclusions": "\\label{par:conclusions} We observed the BHC Swift J1753.5-0127 with a campaign of coordinated multiwavelength observations in the radio, OIR and X-ray bands. Following the source for $\\sim$4 years we could clearly confirm that it features a jet that is fainter than expected from the empirical correlation between the radio and the X-ray scaled fluxes (Corbel et al. 2003; Gallo et al. 2003, 2006). We also verified that the Swift J1753.5-0127 is only slightly fainter in OIR than expected from the correlation of Russell et al. (2006): in the ($L_X,L_{OIR}$) plane the source clusters with other BHCs.\\\\ From the analysis of the SEDs in Figure \\ref{fig:SED_all} we inferred that probably the jet is not responsible for the OIR emission. Viable mechanisms are either reprocessing of the X-rays in the outer regions of the accretion disc or emission from the accretion disc itself while the possibility that the OIR emission comes from the companion star seems unlikely.\\\\ The optically thin radio spectra that we obtained suggest that the radio emission originates in blobs of plasma ejected near the peak of the outburst. However, the presence of observations characterized by flat or slightly inverted spectra is a hint that a compact, steady jet is indeed on, although the radio emission is often dominated by the discrete ejecta detached from the core.  The best-fit parameters to the Swift J1753.5-0127 data in the ($S_X, S_{radio}$) plane could only be poorly constrained, because of the lack of observations at low ($< 1$ Crab) scaled X-ray flux. One must also keep in mind that any determination of the slope of the radio/X-ray correlation from radio observations characterized by optically thin spectra optically thin spectra should be treated with care. We discussed the importance of a precise determination of both the slope $b$ and the normalisation $k$ of the radio/X-ray correlation to infer the radiative efficiency of the jet. The possibility that (at least some of) the outliers of the correlation could be fitted using a relation characterized by the same slope $b$ as the majority of the BHCs but a lower normalisation suggests that some parameters (e.g. the spin of the black hole and the jet magnetic field) might play a role in regulating the disc/jet coupling and the energy output in the form of a jet.  We also briefly discussed on the nature of Swift J1753.5-0127, since the mass of the accretor  has not been measured. From our data we can not draw any conclusion on the nature of the compact object. Although the possibility that the system is a NSXBs can not be excluded, the detection of a hard power-law tail in the X-ray/$\\gamma$-ray energy spectrum up to $\\sim$600 keV still constitutes the strongest hint that Swift J1753.5-0127 is a BHC."
    },
    "1004/1004.3313_arXiv.txt": {
        "abstract": "We report the results of abundance analyses of new samples of stars with planets and stars without detected planets. We employ these data to compare abundance-condensation temperature trends in both samples. We find that stars with planets have more negative trends. In addition, the more metal-rich stars with planets display the most negative trends. These results confirm and extend the findings of Ramirez et al. (2009) and Melendez et al. (2009), who restricted their studies to solar analogs. We also show that the differences between the solar photospheric and CI meteoritic abundances correlate with condensation temperature. ",
        "introduction": "\\citet{mel09} were the first to detect a significant correlation between elemental abundances and condensation temperature (T$_{\\rm c}$) in a sample Sun-like stars; in a follow-up study, \\citet{ram09} confirmed their findings. In particular, they found solar twins/analogs to be enhanced in refractory elements relative to the Sun. Although this was the first time such a trend had been found, it had been searched for unsuccessfully multiple times before within the context of the self-enrichment hypothesis \\citep{hu05, gg06, ecu06}. \\citet{ram09} speculate that the trends are due to planet formation processes.  In order to test the findings of \\citet{ram09} and also to expand on their analysis over a broader range in T$_{\\rm eff}$, we revisit the topic of abundance trends with T$_{\\rm c}$ among stars with planets (SWPs). We do so with a new method of analysis we introduced in \\citet{gg08} and new stellar samples, which we described in \\citet{gg10}.  Our paper is organized as follows. In Section 2 we describe the new SWP and comparison star samples. We employ them in Section 3 to search for evidence of abundance trends with T$_{\\rm c}$; we also examine the recent abundance data of \\citet{neves09} and the Solar System abundances. We discuss our results in Section 4 and present our conclusions in Section 5. ",
        "conclusions": "Using new abundance analyses of SWPs and stars without known planets, we have found that SWPs tend to have more negative [X/H]-T$_{\\rm c}$ slopes than stars without planets. Our results confirm \\citet{ram09}, who focused their study on solar analogs.  We also find that SWPs with [Fe/H] $> 0.10$ tend to have more negative [X/H]-T$_{\\rm c}$ slopes than more metal-poor SWPs, showing that the process responsible for these trends is sensitive to metallicity.  We revisited the question of the abundances in the solar photosphere relative to CI meteorites and confirmed previous findings that a significant trend with T$_{\\rm c}$ exists. The Sun is slightly ($\\sim$ 0.05 dex) enhanced in refractory elements relative to the CI meteorites, but compared to the inner Solar System meteorites, the Sun is deficient by almost a factor of two. This implies that sequestration of the refractory elements into terrestrial planets left their marks in the distribution of Solar System abundances.  It appears that both the enrichment of refractory elements in the solar photosphere via accretion and the sequestration of refractory elements into the terrestrial planets left their marks in the distribution of Solar System elemental abundances.  These results, combined with our recent findings that SWPs have lower Li abundances (confirmed by \\citet{israel09} for solar analogs) and rotate slower than comparison stars, place new stringent constraints on planet formation models \\citep{bv08}."
    },
    "1004/1004.2535_arXiv.txt": {
        "abstract": "Recent results from the Pierre Auger Observatory, showing energy-dependent chemical composition of ultrahigh-energy cosmic rays (UHECRs) with a growing fraction of heavy elements at high energies, suggest a possible non-negligible contribution of the Galactic sources.  We show that, in the case of UHECRs produced by gamma-ray bursts (GRBs) or rare types of supernova explosions that took place in the Milky Way in the past, the change in UHECR composition can result from the difference in diffusion times for different species.  The anisotropy in the direction of the Galactic Center is expected to be a few per cent on average, but the locations of the most recent/closest bursts can be associated with observed clusters of UHECRs. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.5496.txt": {
        "abstract": "{} {Previous studies have found that the coefficients and intrinsic dispersions of both the Kormendy relation and the Fundamental Plane depend on the magnitude range within which the galaxies are contained. We study whether this type of behaviour is also present for the Faber-Jackson relation.} {We take a sample of early-type galaxies from the Sloan Digital Sky Survey (SDSS-DR7, $\\sim$ 90 000 galaxies) spanning a range of approximately 7 $mag$ in both $g$ and $r$ filters and analyse the behaviour of the Faber-Jackson relation parameters as functions of the magnitude range. We calculate the parameters in two ways: i) We consider the faintest (brightest) galaxies in each sample and we progressively increase the width of the magnitude interval by inclusion of the brighter (fainter) galaxies (increasing-magnitude-intervals), and ii) we consider narrow-magnitude intervals of the same width ($\\Delta M = 1.0$ $mag$) over the whole magnitude range available (narrow-magnitude-intervals).} {Our main results are that: i) in both increasing and narrow-magnitude-intervals the Faber-Jackson relation parameters change systematically, ii) non-parametric tests show that the fluctuations in the values of the slope of the Faber-Jackson relation are not products of chance variations.} {We conclude that the values of the Faber-Jackson relation parameters depend on the width of the magnitude range and the luminosity of galaxies within the magnitude range. This dependence is caused, to a great extent by the selection effects and because the geometrical shape of the distribution of galaxies on the $M - \\log (\\sigma_{0})$ plane depends on luminosity. We therefore emphasize that if the luminosity of galaxies or the width of the magnitude range or both are not taken into consideration when comparing the structural relations of galaxy samples for different wavelengths, environments, redshifts and luminosities, any differences found may be misinterpreted.} ",
        "introduction": "\\label{sec:intro}  The structural relations of early-type galaxies (ETGs) play an important role in the study of the formation and evolution of these galaxies. Among the most important structural relations for ETGs there are: The Kormendy relation (KR; \\cite{kor77})    \\begin{equation} \\left\\langle \\mu\\right\\rangle_{e}=\\alpha+\\beta\\log(r_{e}), \\end{equation} the Faber-Jackson relation (FJR; \\cite{fab76})   \\begin{equation} \\log\\sigma_{0}=A-BM_{}, \\end{equation} and the Fundamental Plane (FP; \\cite{djo87}; \\cite{dre87})   \\begin{equation} \\log(r_{e})=a\\log(\\sigma_{0})+b\\left\\langle \\mu\\right\\rangle _{e}+c. \\end{equation} where $r_{e}$ represents the effective radius, $\\left\\langle \\mu\\right\\rangle _{e}$ is the mean effective surface brightness inside  $r_{e}$, $M$ is the total absolute magnitude, $\\sigma_0$ is the central velocity dispersion, and $\\alpha$, $\\beta$, $A$, $B$, $a$, $b$, and $c$ are scale factors.  The FP relation is a consequence of the dynamical equilibrium condition (virial theorem) and the regular behaviour of both the mass-luminosity ratio and the structure along the entire range of ETGs luminosities (homology).  Based on these considerations, we should find that (a, b)=(2, -0.4). There are, however, discrepancies between these values and those obtained from observations. These discrepancies might indicate that the FP and other structural relations are not universal, i.e., that they may depend on wavelength, environment, redshift, and/or luminosity. There have been several studies of this universality using different samples of galaxies. We now discuss briefly the most important results about the universality of the structural relations.  Several studies have demonstrated that both of the FP coefficients $a$ and $b$ remain unchanged when the relation is examined at different wavelengths (\\cite{ben92}; \\cite{BENDERETAL98}; \\cite{ber03b}; \\cite{ber03c}; \\cite{lab05}; \\cite{lab08}). However, other studies have found that the structural relations do depend on wavelength (\\cite{jor96}; \\cite{hud97}; \\cite{pah98}; \\cite{sco98}; \\cite{jun08}), and in particular, that the $a$ coefficient of the FP varies with wavelength (i.e., the coefficient is larger at longer wavelengths), while the $b$ coefficient remains stable, that is ($a$, $b$) $\\sim$ (1.2-1.6, -0.34) for wavelengths of 0.55-8.0 $\\mu m$. Some of the afore mentioned authors in this paragraph have interpreted this as an effect caused by colour gradients inside galaxies.  Several studies have also discovered that the structural relations and/or the structural parameters of galaxies are affected by the environment (\\cite{ben92}; \\cite{BENDERETAL98}; \\cite{tru01}; \\cite{tru02}; \\cite{ber03b}; \\cite{agu04}; \\cite{gut04}; \\cite{den05}; \\cite{jor05}), although other studies have reached exactly the opposite conclusion (\\cite{ros01}; \\cite{tre01}; \\cite{evs02}; \\cite{gon03}; \\cite{red04}; \\cite{red05}; \\cite{nig07}). This second set of results suggest that galaxies very quickly absorb the changes caused by gravitational interactions, not keeping memory of these changes in regard to the structural relations. Another possible explanation is that most of the perturbations affect only the external parts of galaxies which are not represented in the structural relations given that these are defined by the use of effective parameters only.  When the structural relations are constructed for galaxy samples at different redshifts, several studies report similarly (\\cite{bar98}; \\cite{zie99}; \\cite{lab03}; \\cite{bar06}) that only the zero point of these relations depends on the redshift and this dependence may be caused by the passive evolution of the stellar populations that form at high redshift ($z_{form} > 2$). Nevertheless, their results conflict with those of other authors. For example, Treu et al. (2005), J\\o rgensen et al. (2006), and Fritz et al. (2009) find that the zero point of the FP depends on redshift but that the slope of the FP is steeper for higher redshift galaxies than for galaxies in the local Universe. They interpret this as a mass dependence of the star formation history, that is, the low-mass galaxies ($10^{10.3}M_{\\odot}$) have experienced star formation as recently as $z_{form} \\sim 1.1$, while, galaxies with masses $ \\sim 10^{10.8}M_{\\odot}$ and masses $ > 10^{11.3}$ $M_{\\odot}$ had their last major star formation episode at $z_{form} > 1.25$ and $z_{form} > 1.6$ (\\cite{jor07}).  Finally, considering luminosity, Kormendy (1985) states that dwarf ellipticals do not follow the same plane as bright ellipticals. Other authors have reported similar results (\\cite{ham87}; \\cite{ben92}; \\cite{cao93}; \\cite{agu09}). Recent studies have demonstrated that some projections of the FP, such as the FJR depend on the mass, radial distance from the cluster centre (\\cite{fri05}) and on the luminosity (\\cite{des07}) in the sense that brighter galaxies have a steeper slope. These authors attribute these dependences to galaxies of different luminosities experiencing different formation histories and brighter galaxies being less affected by dissipation. Other studies indicate that it is inappropriate to draw conclusions about the physical properties of galaxies by comparing the coefficients of the structural relations for different magnitude ranges because the form of the distribution of galaxies in the space of the variables that define these relations plays a crucial role in determining their coefficients (\\cite{nig07b}; \\cite{nig08}; \\cite{nig09}). These studies find that the detected changes in the coefficients values for the structural relations obtained from fits to samples of galaxies of different luminosity are not necessarily related to differences in the intrinsic properties of these galaxies.  The results presented in Nigoche-Netro (2007) and Nigoche-Netro et al. (2008; 2009) characterise the effects produced by the magnitude range restrictions for different galaxy samples and indicate that the differences between the coefficients calculated within a magnitude range may differ by as much as 60\\% for the KR from the values calculated for another magnitude range,and up to 30\\% for the FP. Using Monte Carlo simulations, these studies also show that the changes in the values of the coefficients for the KR and the FP relations may be attributed to a `geometrical effect', in other words, the geometrical form of the galaxy distribution in the variables space defining the KR and the FP relations changes systematically as brighter galaxies are considered. Therefore, the values of the coefficients also change, because the fitting of relations to galaxy distributions with different geometrical forms produces different results. An important conclusion of these studies is that any restriction imposed on any of the variables involved in the KR and FP relations will produce similar effects as those produced by the restrictions on the magnitude. In light of this, it is reasonable to assume that part of the discrepancies in results about the universality of the structural relations may have been caused by the large majority of studies having underestimated the effects produced by magnitude restrictions, as well as all restrictions affecting every variable involved in the determination of the parameters of the structural relations.  In this paper, we present a sample of ETGs from the SDSS-DR7 that contains approximately 90 000 galaxies and covers a relatively wide magnitude range ($\\Delta M$ $\\sim 7$ $mag$ in $g$ and $r$ filters). Using these data, we analyse the behaviour of the coefficients and the intrinsic dispersion of the FJR with respect to several characteristics of the magnitude range. Throughout this paper, we use $H_{0}$ = 70 $km s^{-1}Mpc^{-1}$, $\\Omega_{m}$ = 0.3 and $\\Omega_{\\Lambda}$ = 0.7.  Our paper is organised as follows. In Sect. 2 we present the galaxy sample used to study the FJR. Section 3 presents the fitting method that we use in calculating the FJR coefficients, the results of these calculations and the analysis of the behaviour of the slope of the FJR as a function of the magnitude range. In Sect. 4 we present a discussion of the most important results of this paper. Finally, in Sect. 5 we present our conclusions. ",
        "conclusions": "\\label{sec:conclusions}  Analysing the FJR as a function of the magnitude range we obtain the following:  \\begin{itemize}  \\item The parameters of the FJR depend on the magnitude range within which the galaxies of the sample under analysis are distributed.  \\item The dependence of the FJR on the magnitude range may be explained by a `geometrical effect' (see Sects. 3.2 and 4 for full details).  \\end{itemize}  With the results given above and the data from Nigoche-Netro et al. (2008; 2009), we find that the intrinsic dispersion, magnitude segregation and, to a larger extent, the observational biases (or those biases caused by arbitrary cuts made on galaxy samples) are mainly responsible for the geometrical form of the galaxy distribution and that this form determines the values of the coefficients of the structural relations.  We find that it is risky to draw conclusions about the physical properties of galaxies by comparing the slopes of the structural relations in magnitude ranges of different widths or in magnitude ranges of the same width but of different luminosity, because, with the exception of the full magnitude interval, there is no privileged width for making comparisons. This means that, if the magnitude range is narrow the differences are negligible, whereas if it is wide, the geometrical form is dominated by the magnitude cut that could mask any differences caused by intrinsic physical properties of the galaxies. In the case of samples of galaxies in magnitude ranges of the same width and luminosity but of different wavelengths, redshifts, or environments, comparison of the slope values is also a delicate matter, since the magnitude cut could mask the intrinsic physical properties of the galaxies and the conclusions about these physical properties might be misinterpreted.  We find that, if there were differences in the total intrinsic dispersion for faint and bright galaxies and we were to compare magnitude ranges decreasing in width, the differences would not disappear (as happens with the differences for the slopes of the structural relations) and when $\\Delta M = 0$, we would find the exact values for the intrinsic dispersions for each one of the magnitudes under study. We, therefore, consider the most appropriate procedure for obtaining physical information for a sample of galaxies to be finding its intrinsic dispersion at each magnitude value and then performing comparisons of this dispersion at different luminosities, wavelengths, redshifts, or environments. A first approach to the study of galaxies of different luminosities (see Sect. 4) reveals that in the regime $M_{g}\\lesssim -20.0$ the intrinsic dispersion for bright galaxies is smaller than that for faint galaxies, confirming previous results by J\\o rgensen et al. (1996) and Hyde \\& Bernardi (2009). The study of the intrinsic dispersion as a function of luminosity requires a far more complete and detailed analysis. This analysis shall be published in a forthcoming paper.   %A preliminary analysis of galaxies of different luminosities (SEE %SECTION 4) allowed us to find that the intrinsic dispersion for %bright galaxies is smaller than that for faint galaxies, CONFIRMING %PREVIOUS RESULTS BY J\\o rgensen et al. (1996) AND Hyde \\& Bernardi %(2009).  We may conclude that part of the discrepancies found in the literature about the behaviour of the structural relations with respect to the following variables: wavelength, environment, redshift, and luminosity (see Sects. 1 and 4) must be due to the vast majority of the investigations carried out to date having underestimated the effects produced by magnitude restrictions as well as those restrictions that affect each one of the variables involved in the parameters of the structural relations. Finally, it is very important to redirect efforts towards investigating the physical processes that cause the intrinsic dispersion in the structural relations and also to bright galaxies having a different intrinsic dispersion from faint galaxies.  %THE AUTHORS OF THESE PAPERS  %DESCRIBIMOS AMLIAMENTE EN LA SECCION 1 DEL PRESENTE ARTICULO QUE %EXISTEN DISCREPANCIAS ENTRE DISTINTOS AUTORES RESPECTO AL %COMPORTAMIENTO DE LAS RELACIONES ESTRUCTURALES RESPECTO A LAS %VARIABLES wavelength, environment, REDSHIFT and luminosity. DE %MANERA QUE PODEMOS CONLCUIR THAT the majority of the DICHAS %discrepancies must be due to the fact that they have underestimated %the effects produced by magnitude restrictions as well as those %restrictions that affect each one of the variables involved on the %parameters of the structural relations. Finally, it is very %important to redirect efforts to investigate the physical processes %that cause the intrinsic dispersion of the structural relations and %also to the fact that bright galaxies have a different value for %their intrinsic dispersion from the value that faint galaxies have.  %We may conclude that the majority of the discrepancies found by %several authors about the universality of the structural relations %must be due to the fact that they have underestimated the effects %produced by magnitude restrictions as well as those restrictions %that affect each one of the variables involved on the parameters of %the structural relations. Finally, it is very important to redirect %efforts to investigate the physical processes that cause the %intrinsic dispersion of the structural relations and also to the %fact that bright galaxies have a different value for their intrinsic %dispersion from the value that faint galaxies have."
    },
    "1004/1004.2473_arXiv.txt": {
        "abstract": "{The radio source 18P87, previously thought to be a point source, has been serendipitously found to be resolved into a core-jet geometry in VLA maps. \\ion{H}{i} absorption of continuum emission (in data from the Canadian Galactic Plane Survey) appears in gas with radial velocities $>+2$ km/s but not in brightly emitting gas at lower radial velocity. Examination of further archival observations at radio, infrared and optical wavelengths suggests that the ``obvious'' interpretation as a radio galaxy requires a rather unusual object of this kind and a highly unusual local line of sight. We argue that 18P87 may be a Galactic object, a local astrophysical jet. If this is correct it could have arisen from outbursts of a microquasar. } ",
        "introduction": "Our surveys of the Cygnus X region (Wendker, Higgs, \\& Landecker \\cite{cx18}, Paper XVIII) generated several follow-up observations at higher resolution with the VLA. In one of these we serendipitously found that the source \\pen\\ was resolved into a structure that at first glance resembles a faint radio galaxy and we did not investigate it more thoroughly. However, while perusing continuum absorption by \\ion{H}{i} in data from the Canadian Galactic Plane Survey (CGPS - Taylor et al. \\cite{art}) we found that \\pen\\ shows \\ion{H}{i} absorption for only a small part of the local gas, which is quite unexpected for an extragalactic source. In this paper we present the available observations and examine the possibility that the source is Galactic. ",
        "conclusions": "The hitherto poorly studied radio source \\pen\\ has been seren\\-di\\-pitously shown to have the structure of an astrophysical jet. We suggest that it may be a Galactic object. While the distinction on the basis of physical properties is not totally conclusive, the Galactic interpretation relies on the absence of \\ion{H}{i} absorption of the continuum emission beyond the middle of the local gas. If it is indeed a Galactic structure its non-thermal radio emission classes it as a microquasar.  We suggest that it is only sporadically active and that the structures we have observed are the remnants of interaction with a dense interstellar environment."
    },
    "1004/1004.2190_arXiv.txt": {
        "abstract": "{ Small, bright stellar disks with scale lengths of a few tens of parsec are known to reside in the center of galaxies. They are believed to have formed in a dissipational process as the end result of star formation in gas either accreted during a merging (or acquisition) event or piled up by the secular evolution of a nuclear bar. Only a few of them have been studied in detail to date.} { Using archival Hubble Space Telescope (HST) imaging, we investigate the photometric parameters of the nuclear stellar disks hosted by three early-type galaxies in the Virgo cluster, NGC~4458, NGC4478, and NGC4570, to constrain the process that forms their stars.} { The central surface brightness, scale length, inclination, and position angle of the nuclear disks were derived by adopting the photometric decomposition method introduced by Scorza \\& Bender and assuming the disks to be infinitesimally thin and exponential.} { The location, orientation, and size of the nuclear disks is the same in all the images obtained with the Wide Field Planetary Camera 2 and Advanced Camera for Surveys and available in the HST Science Archive. The scale length, inclination, and position angle of each disk are constant within the errors in the observed $U$, $B$, $V$, and $I$ passbands, independently of their values and the properties of the host spheroid.} { We interpret the absence of color gradients in the stellar population of the nuclear disks as the signature that star formation homogeneously occurred along their length. An inside-out formation scenario is, instead, expected to produce color gradients and is therefore ruled out.} ",
        "introduction": "\\label{sec:intro}  The presence of small and bright stellar disks with scale lengths of a few tens of parsec was first discovered in the nuclear regions of nearby galaxies about fifteen years ago thanks to the subarcsec-resolution capabilities of the Hubble Space Telescope \\citep[HST,][]{vdbetal94, kormetal94, emsetal94, laueretal95}. Afterward, nuclear stellar disks (NSDs) and larger scale embedded disks were reported in a large number of early-type galaxies as a results of photometric \\citep{laueretal95, caroetal97, ravin01, restetal01, tranetal01, ferretal06, sethetal06, lopeetal07, kormetal09} and spectroscopic surveys \\citep{hallietal01, kuntetal06, mcdeetal06} respectively. Given that nuclear disks are easiest to detect when nearly edge-on \\citep{rixwhi90}, the observed fraction is consistent with NSDs being a common structure in early-type galaxies \\citep{ledoetal10}. At the center of spiral galaxies NSDs are also present but these are relatively rare in late types \\citep[][]{pizzetal02, falcetal06, peleetal07,balcetal07,moreetal08}.  Photometric parameters of NSDs (i.e., central surface brightness, scale-length, axial ratio, and position angle) have been derived for nine galaxies spanning a wide range of Hubble types, from ellipticals \\citep{moreetal04}, to lenticulars \\citep{vdbetal94, laueretal95, kormetal96a, vdbetal98, scovdb98}, and to early-type spirals \\citep{pizzetal02}. They are characterized by a smaller scale length ($10-30$ pc) and higher central face-on surface brightness ($15-19$ mag arcsec$^{-2}$ in the $V$ band) than the embedded stellar disks observed in ellipticals and lenticulars as well as to the large kpc-scale disks of lenticulars and spirals \\citep{vdb98, pizzetal02, moreetal04}. In particular, the embedded stellar disks represent the intermediary between the sequence of elliptical galaxies and that of disk galaxies. They have been detected in both disky elliptical and lenticular galaxies, have scale lengths of $0.1-1$ kpc, and central face-on surface brightnesses of $18-21$ mag arcsec$^{-2}$ in the $V$ band. Their angular momentum is parallel to that of their host spheroid. This demonstrates that the embedded disks are not the result of accretion or merger events but are likely to be primordial, in the same sense as the main disks of lenticular and spiral galaxies \\citep{scoben95}. The smooth variation in the scale parameters from the nuclear to the embedded and main stellar galactic disks provides observational support that galaxy properties vary with continuity along a sequence of increasing disk-to-bulge ratio \\citep[see][and references therein]{korben96}. Therefore, unveiling the formation scenario of NSDs may also improve our understanding of galaxy formation and evolution.  In the current picture, NSDs are believed to have formed in a dissipational process as the end result of star formation in gas either accreted in a merging (or acquisition) event \\citep[the most clearcut case is NGC~4698,][]{bertetal99, pizzetal02} or piled up by the secular evolution of a nuclear bar \\citep[e.g., NGC~4570,][]{vdbetal98, scovdb98, vdbems98}. Each of these scenarios is likely to be correct for some, but not all, objects. In both of them, the gas is efficiently directed toward the galaxy center \\citep[e.g.,][]{barher96, athaetal05, dottetal07, elicetal09, hopqua10}, where it first dissipates and settles onto an equilibrium plane and then forms into stars.  A striking example of an on going process of dissipational formation is provided by NGC~4486A \\citep{kormetal05}. The NSD of this low-luminosity elliptical coexists with its progenitor disk of dust and gas. The question about the origin of these dynamically-cold components is also related to their formation epoch, i.e., whether they built up in the early stages of the galaxy assembly or whether they formed later.  To address this issue, \\citet{vdbetal98}, \\citet{vdbems98}, \\citet{moreetal04}, and \\citet{krajaf04} studied the age and metallicity content of some NSDs. In some cases, the nuclear disk was found to be younger than the stellar population of their surrounding spheroidal component (e.g., NGC~4486A, \\citealt{kormetal05}), in other cases the formation of the disks occurred at the same time as the main body of the host galaxy (e.g., NGC~4342, \\citealt{vdbetal98}).  Characterizing how the star formation proceeded in nuclear disks can lead to understanding what triggered their formation. For instance, if stars formed more or less at the same rate throughout the disk from disk material that accumulated very rapidly at the center, as one expects to happen during a merging event, we should not find strong radial gradients in the properties of the stellar populations. In contrast, color gradients are expected if the disk experiences an inside-out star formation, like the disks of spiral galaxies \\citep[e.g.,][]{munoetal07} In this paper, we use HST archival images of the nuclei of three early-type galaxies in the Virgo cluster, NGC~4458, NGC~4478, and NGC~4570 to perform a multiband analysis of the photometric parameters of their NSDs. Our aim is to understand whether the star formation of the nuclear disk was homogeneous and occurred simultaneously everywhere along its length. ",
        "conclusions": "We have investigated the photometric properties of the NSDs of three early-type galaxies in the Virgo cluster, NGC~4458, NGC~4478, and NGC~4570 by analyzing the WFPC2 and ACS images of their nuclei available in the HST Science Archive. The images were unsharp masked and their visual inspection suggested that the size, orientation, and location of each NSD was independent of the observed passband. To verify this, we derived the central surface brightness, scale length, inclination, and position angle of each NSD in all the passbands by applying the photometric decomposition method of \\citet{scoben95}.  Some of the images had already been analyzed by \\citet[][ WFPC2/F555W for NGC~4570]{vdbetal98} and \\citet[][WFPC2/F814W for NGC~4458 and NGC~4478]{moreetal04}.  Since we adopted an identical procedure to that of \\citet{moreetal04}, we relied on their WFPC2/F814W decomposition of NGC~4458 and NGC~4478 and repeated the WFPC2/F555W analysis of NGC~4570 to build homogeneous set of photometric parameters and corresponding errors for all the NSDs in all the passbands.  \\begin{figure} \\includegraphics[height=9cm,angle=90]{14285fig5.ps} \\caption[]{Photometric parameters of the nuclear stellar disks of NGC~4458 ({\\em blue circles}), NGC~4478 ({\\em green squares}), and NGC~4570 ({\\em red triangles}) in the WFPC2/F336W, ACS/F475W, WFPC2/F555W, and WFPC2/F814W passbands. } \\label{fig:parameters} \\end{figure}  The structural parameters (i.e., scale length, inclination, and position angle) of each NSD are constant within the errors in all the observed passbands, independently of their values and the properties of their host spheroid. This implies an absence of color gradients in nuclear disks and can be used to constrain their star formation processes. NSDs are considered to be the end result of the formation of stars from gas of external (i.e., accreted from galaxy neighborhoods or captured during a merger with a gas rich-companion; e.g., NGC~4458, NGC~4478 \\citealt{moreetal04}) or internal origin (i.e., conveyed from outer galactic regions by a bar; e.g., NGC~4570 \\citealt{vdbems98}) piled up in the galaxy nucleus. Thus independent of the way in which the gas was originally accumulated towards the center, the dissipational formation process has to be invoked to account for the dynamically cold structure of the nuclear disks. This is hard to explain if NSDs assembled from already formed stars. We interpret the absence of color gradients in NDSs as the signature of a star formation event that occurred homogeneously over the entire disk. Although our sample is admittedly very small, an inside-out formation scenario \\citep{munoetal07} seems to be ruled out in NSDs given that this would produce radial population gradients that are not observed."
    },
    "1004/1004.5065_arXiv.txt": {
        "abstract": "{} % {NGC~2903 is a nearby barred spiral with an active starburst in the center and \\hii\\ regions distributed along its bar. We aim to analyse the star formation properties in the bar region of NGC~2903 and study the links with the typical bar morphological features.} {A combination of space and ground--based data from the far--ultraviolet to the sub--millimeter spectral ranges is used to create a panchromatic view of the NGC~2903 bar. We produce two catalogues: one for the current star formation regions, as traced by the \\ha compact emission, and a second one for the ultraviolet (UV) emitting knots, containing positions and luminosities. From them we have obtained ultraviolet colours, star formation rates, dust attenuation and \\ha EWs, and their spatial distribution have been analysed. Stellar cluster ages have been estimated using stellar population synthesis models ({\\em Starburst99}).} {NGC~2903 is a  complex galaxy, with a very different morphology on each spectral band. The CO~($J$=1-0) and the 3.6~\\micron\\ emission trace each other in a clear barred structure, while the \\ha leads both components and it has an s--shape distribution. The UV emission is patchy and does not resemble a bar. The UV emission is also characterised by a number of regions located symmetrically with respect to the galaxy center, almost perpendicular to the bar,  in a spiral shape covering the inner $\\sim2.5$kpc. These regions do not show a significant \\ha nor 24~\\micron\\ emission. We have estimated ages for these regions ranging from 150 to 320~Myr, being older than the rest  of the UV knots, which have ages lower than 10~Myr. The SFR calculated from the UV emission is  $\\sim$0.4~M$_{\\odot}$~yr$^{-1}$, compatible with the SFR as derived from \\ha calibrations ($\\sim$1~M$_{\\odot}$~yr$^{-1}$).} {} ",
        "introduction": "In the years after the launch of the \\textit{Spitzer Space Telescope} and the \\textit{Galaxy Evolution Explorer} ({\\em galex}) knowledge on the star formation (SF) in galaxies have grown considerably. The panchromatic view of nearby galaxies offered by large surveys carried out from these space telescope missions gives an extraordinary database to study star formation in galaxies. In particular, it allows us to link general galaxy properties with the the local interstellar medium (ISM) properties and the galaxy dynamics. These links are crucial for understanding SF in galaxies. In nearby galaxies these sets of multi-wavelength data give the opportunity to study in detail, with good spatial resolution, the location and properties of relatively young populations, recent massive star formation and dust attenuation (e.g. Calzetti et al. 2005; Tamura et al. 2009; Rela\\~no \\& Kennicutt 2009).\\nocite{tamura}  New dust and SF indicators have been developed and calibrated using all the available bands. The combination of 24~\\micron\\ and \\ha emission has become a reliable indicator of dust attenuation (Kennicutt et al. 2007; Rela\\~no \\& Kennicutt 2009)\\nocite{Kennicutt2007,Relano2009}. Both bands are linked to star formation processes, with the \\ha emission originating from  the recombination of hydrogen  in the surrounding medium of very recently formed (less than a few Myr)  massive stars, and the 24 \\micron\\ emission as tracing local star formation radiation obscured by dust (Calzetti et al. 2005). Extinction corrected UV emission can also be used to retrieve the star formation rate (SFR) (Kennicutt 1998)\\nocite{Kennicutt1998} and to get some insight into the star formation occurred in last  Gyr (e.g. Bianchi et al. 2005; Hibbard et al. 2005). This feature makes a combination of UV,  \\ha and  24 \\micron\\ emission ideal to reconstruct the recent star formation history in a galaxy by studying the location and properties of individual clusters and those of the gas and dust emission around them.\\nocite{bianchi2005,hibbard2005}\\\\  To this date, and to our knowledge, only a few galaxies have  been analysed in this way;  M51 (Calzetti et al. 2005), M81 (P\\'erez-Gonz\\'alez et al. 2006); M33 (Rela\\~no \\& Kennicutt 2009; Verley et al. 2009); NGC~7331 (Thilker et al. 2007).\\nocite{perezgonzalez,verley2009,thilker2007}  Barred galaxies offer a useful tool to investigate the physical conditions that favour star formation in galaxies. The motions in bars are characterised by non-circular motions that push the gas into intersecting orbits where shocks and star formation can be triggered. The position and strength of the shocks are determined by the bar potential, the global dynamics within the bar region is driven by the bar. We have now a relatively good understanding of the gas behaviour under a bar potential (e.g. P\\'erez et al. 2004)\\nocite{perez2004}  and this knowledge can be used to understand the conditions triggering star formation. Bars and their surroundings host extreme physical conditions and a variety of ISM environments. They are perfect places to study the link between the conditions favouring star formation with the galaxy dynamics. Therefore,  a panchromatic view revealing the history of star formation in bars can give a unique insight into the links of star formation and the galaxy dynamics, which could also help  to understand how bars form and evolve.  This work presents a detailed multi-wavelength study of the star formation in the bar region of NGC~2903 analysing the correlations between the location and ages of the {\\it young} stellar clusters with the morphology of the bar.  This is done by analysing the emission in \\ha, UV and 24 \\micron\\ as well as optical data from the Sloan Digital Sky Survey (SDSS), and using 8 \\micron\\ and CO~(J$=$1-0) emission as complementary data.   NGC 2903 is chosen for this research for a number of reasons: it is  close by (8.9 Mpc; Drozdovsky \\& Karachentsev 2000)  \\nocite{Drozdovsky2000}allowing us to have high spatial resolution ($\\sim$~43~pc~arcsec$^{-1}$), and  it is isolated from large companions, preventing major merger effects in the results. Recent work (Irwin et al. 2009) on the \\hi\\ content of NGC~2903 has shown that it possesses a large \\hi\\ envelope of around three times its optical size. They also found a small \\hi\\ companion 64~kpc away from the galaxy in projection,  which adds to a previously known small stellar companion. No clear sign of interaction has been found so far.\\nocite{irwin2009}  NGC~2903 is a SBd galaxy showing a symmetric strong bar considered typical for this class of galaxies (Laurikainen \\& Salo 2002).\\nocite{Laurikainen2002} Previous observations have shown large amounts  of \\ha emission along the bar and not only  at the ends of the bar and nuclear region (Sheth et al. 2002).  The CO($J=$1 - 0), \\spitzer and \\galex data available makes this galaxy an ideal object for a multi wavelength study to retrieve insight in star formation history in bars.  A previous study by Leon et al. (2008)\\nocite{Leon2008} on the NGC 2903 bar  showed that HCN(1-0) is distributed along the bar and in the center. They compared the star formation rate ratio between the bar and the center with results from numerical simulations by Martin \\& Friedli (1997)\\nocite{MartinFriedli1997}. This made them propose that the bar in NGC 2903 has an age between 200 and 600 Myr.  The plan for the article is the following: in section \\ref{ObservationsDataReduction} we present the observational data; in section~\\ref{morphology} we analyse the general morphology of NGC~2903. We  present our methodology to obtain the bar \\hii\\ regions and UV emission knots catalogues in section~\\ref{photometry}. The following section contains the main results regarding EW$_{H\\alpha}$, star formation rates, UV colours and ages of the stellar clusters.  In section~\\ref{sec:discussion} we discuss our main results and finally, we end  with a summary and conclusions. ",
        "conclusions": "We have performed a detailed multi-wavelength study from UV to sub--millimeter observations on the NGC 2903 bar and its surrounding regions. We mapped and catalogued the \\hii\\ regions of the bar and measured their \\ha equivalent widths. Furthermore we have obtained a catalogue of the UV emitting regions, and measured their peak location with respect to \\ha. The extinction has been estimated using the \\ha and 24 \\micron\\ emission. SFRs using both \\ha and UV indicators have been calculated. We have estimated the age of the regions, using the EW$_{H\\alpha}$ and the FUV-NUV colour together with stellar population synthesis models. Our main results are:   \\begin{itemize} \\item NGC 2903 is a morphologically complex galaxy. The near--infrared as well as the CO~($J$=1-0) band show a clear barred structure whereas the \\ha and UV maps show a patchy spiral like structure.  \\item There are clear spiral like UV complexes with no significant H$\\alpha$, 24~\\micron\\ and  CO~(J$=$1-0)  counterpart emission. These complexes are located Northwest and Southeast of the bar within the inner 1 arc minute radius (corresponding to $\\sim$2.5 kpc).  \\item The \\ha emission along the bar, leads the CO. The 3.6~\\micron\\ and CO~(J$=$1-0)  emission trace each other, both leading the major axis of the optical light distribution.  \\item The \\ha luminosities and EW$_{H\\alpha}$ of the bar \\hii\\ regions are within typical ranges for \\hii\\ regions in bar and unbarred spirals. The spatial distribution of the  \\ha\\ EWs does not correlate with any morphological feature of the bar.  \\item The average dust attenuation in the bar area of NGC~2903 is A$_{H\\alpha}$= 1.06, and ranges from 0 to 1.5 mag.  \\item The {\\em FUV-NUV} colour distribution is distributed in two regions, the  bluer regions range from -0.18  to 0.15 implying an age of $\\sim$3 to 10 Myr. The redder regions have colours ranging  from  0.4 to  0.85 which imply ages ranging from 150 to 400 Myr. The latter  correspond to regions with no significant \\ha nor 24~\\micron\\ emission.  \\item The SFRs of the bar region derived from \\ha and from the UV emission are  $0.9 \\pm 0.2\\, M_\\odot \\,yr^{-1}$  and {\\bf $0.4 \\pm 0.1\\, M_\\odot \\,yr^{-1}$ }  respectively. \\end{itemize}   All these results suggest that an agent triggered a SF burst a few hundred Myrs ago. Interestingly, we see some stellar clusters unrelated to the current SF locations, symmetrically located nearly perpendicular to the bar (in the inner $\\sim$2.5 kpc). The origin of these regions might be related to this SF burst.  In a following paper we will analyse the gas kinematics of the galaxy, with the aim of shedding some light on the  origin (merger vs. secular evolution) of these findings."
    },
    "1004/1004.0689_arXiv.txt": {
        "abstract": "{}{% The test-field method for computing turbulent transport coefficients from simulations of hydromagnetic flows is extended to the regime with a magnetohydrodynamic (MHD) background. }{% A generalized set of test equations is derived using both the induction equation and a modified momentum equation. By employing an additional set of auxiliary equations, we derive linear equations describing the response of the system to a set of prescribed test fields. Purely magnetic and MHD backgrounds are emulated by applying an electromotive force in the induction equation analogously to the ponderomotive force in the momentum equation. Both forces are chosen to have Roberts flow-like geometry. }{% Examples with an MHD background are studied where the previously used quasi-kinematic test-field method breaks down. In cases with homogeneous mean fields it is shown that the generalized test-field method produces the same results as the imposed-field method, where the field-aligned component of the actual electromotive force from the simulation is used. Furthermore, results for the turbulent diffusivity tensor are given, which are inaccessible to the imposed-field method. For MHD backgrounds, new mean-field effects are found that depend on the occurrence of cross-correlations between magnetic and velocity fluctuations. For strong imposed fields, $\\alpha$ is found to be quenched proportional to the fourth power of the field strength, regardless of the type of background studied. }{} ",
        "introduction": "Astrophysical bodies such as stars with outer convective envelopes, accretion discs, and galaxies tend to be magnetized. In all those cases the magnetic field varies on a broad spectrum of scales. On small scales the magnetic field might well be the result of scrambling an existing large-scale field by a small-scale flow. However, at large magnetic Reynolds numbers, i.e.\\ when advection dominates over magnetic diffusion, another source of small-scale fields is small-scale dynamo action (Kazantsev 1968). This process is now fairly well understood and confirmed by numerous simulations (Cho \\& Vishniac 2000; Schekochihin et al.\\ 2002, 2004; Haugen et al.\\ 2003, 2004); for a review see Brandenburg \\& Subramanian (2005). Especially in the context of magnetic fields of galaxies, the occurrence of small-scale dynamos may be important for providing a strong field on short time scales ($10^7\\yr$), which may then act as a seed for the large-scale dynamo (Beck et al.\\ 1994).  In contemporary galaxies the magnetic fields on small and large length scales are comparable (Beck et al.\\ 1996), but in stars this is less clear. On the solar surface the solar magnetic field shows significant amounts of small-scale fields (Solanki et al.\\ 2006). The possibility of generating such magnetic fields locally in the upper layers of the convection zone by a small-scale dynamo is sometimes referred to as surface dynamo (Cattaneo 1999; Emonet \\& Cattaneo 2001; V\\\"ogler \\& Sch\\\"ussler 2007). On the other hand, simulations of stratified convection with shear show that small-scale dynamo action is a prevalent feature of the kinematic regime, but becomes less important when the field is strong and has saturated (Brandenburg 2005a; K\\\"apyl\\\"a et al.\\ 2008).  An important question is then how the primary presence of small-scale magnetic fields affects the generation of large-scale fields if these are the result of a dynamo process that produces magnetic fields on scales large compared with those of the energy-carrying eddies of the underlying and in general turbulent flow (Parker 1979) via an instability. A commonly used tool for studying these large-scale dynamos is mean-field electrodynamics, where correlations of small-scale magnetic and velocity fields are expressed in terms of the mean magnetic field and the mean velocity using corresponding turbulent transport coefficients or their associated integral kernels (Moffatt 1978; Krause \\& R\\\"adler 1980). The determination of these coefficients (e.g., $\\alpha$ effect and turbulent diffusivity) is the central task of mean-field dynamo theory. This can be performed analytically, but usually only via approximations which are hardly justified in realistic astrophysical situations where the magnetic Reynolds numbers, $\\Rm$, are large.  Obtaining turbulent transport coefficients from direct numerical simulations (DNS) offers a more sustainable alternative as it avoids the restricting approximations and uncertainties of analytic approaches. Moreover, no assumptions concerning correlation properties of the turbulence need to be made, because a direct ``measurement\" of those properties is performed in a physically consistent situation emulated by the DNS. The simplest way to accomplish such a measurement is to include, in the DNS, an imposed large-scale (typically uniform) magnetic field whose influence on the fluctuations of magnetic field and velocity is utilized to infer some of the full set of transport coefficients. We refer to this technique as the {\\em imposed-field method}.  A more universal tool is offered by the test-field method (Schrinner et al.\\ 2005, 2007), which allows the determination of all wanted transport coefficients from a single DNS. For this purpose the fluctuating velocity is taken from the DNS and inserted into a properly tailored set of {\\em test equations}. Their solutions, the {\\em test solutions} represent fluctuating magnetic fields as responses  to the interaction of the fluctuating velocity with a set of properly chosen mean fields. These mean fields will be called {\\em test fields}. For distinction from the test equations, which are in general also solved by direct numerical simulation, we will refer to the original DNS as the {\\em main run}. This method has been successfully applied to homogeneous turbulence with helicity (Sur et al.\\ 2008, Brandenburg et al.\\ 2008a),  with shear and no helicity (Brandenburg et al.\\ 2008b), and with both (Mitra et al.\\ 2009).  A crucial requirement on any test-field method is the independence of the resulting transport coefficients on the strength and geometry of the test fields. This is immediately clear in the kinematic situation, i.e., if there is no back-reaction of the mean magnetic field on the flow. Indeed, for given magnetic boundary conditions and a given value for the magnetic diffusivity, the transport coefficients must not reflect anything else than correlation properties of the velocity field which are completely determined by the hydrodynamics alone. For this to be guaranteed the test equations have to be linear and the test solutions have to be linear and homogeneous in the test fields.  Beyond the kinematic situation the same requirement still holds, although the flow is now modified by a mean magnetic field occurring in the main run. (Whether it is maintained by external sources or generated by a dynamo process does not matter in this context.) Consequently, the transport coefficients are now functions of this mean field. It is no longer so obvious that under these circumstances a test-field method with the aforementioned linearity and homogeneity properties can be established at all. Nevertheless, it turned out that the test-field method developed for the kinematic situation gives consistent results even in the nonlinear case without any modification (Brandenburg et al.\\ 2008c). This method, which we will refer to as ``quasi-kinematic\" is, however, restricted to situations in which the magnetic fluctuations are solely a consequence of the mean magnetic field. (That is, the primary or background turbulence is purely hydrodynamic.)  The power of the quasi-kinematic method was demonstrated based on a simulation of an  $\\alpha^2$ dynamo where the main run has reached saturation with mean magnetic fields being Beltrami fields (Brandenburg et al.\\ 2008c). Magnetic and fluid Reynolds numbers up to 600 were taken into account, so in some of the high $\\Rm$ runs there was certainly small-scale dynamo action, that is, a primary magnetic turbulence $\\bbN$ should be expected. Nevertheless, the quasi-kinematic method was found to work reliably even for strongly saturated dynamo fields. This was revealed by verifying that the analytically solvable mean-field dynamo model employing the values of $\\alpha$ and turbulent diffusivity as derived from the saturated state of the main run indeed yielded a vanishing growth rate. Very likely the small-scale dynamo had saturated on a low level so the contribution to the mean electromotive force, which was not taken into account by the quasi-kinematic method, could not create a marked error.  Limitations of the quasi-kinematic test-field method were recently pointed out by Courvoisier et al.\\ (2010) and will be commented upon in more detail in the discussion section. Indeed, the purpose of our work is to propose a generalized test-field method that allows for the presence of magnetic fluctuations in the background turbulence. Moreover, its validity range should cover dynamically effective mean fields, that is, situations in which the velocity and magnetic field fluctuations are significantly affected by the mean field.  With a view to this generalization we will first recall the mathematical justification of the quasi-kinematic method and indicate the reason for its limited applicability (Sect.\\ 2). In Sect.\\ 3 the foundation of the generalized method will be laid down in the context of a relevant set of model equations. In Sect.\\ 4 results will be presented for various combinations of hydrodynamic and magnetic backgrounds having Roberts-flow geometry. The astrophysical relevance of our results and the connection with the work of Courvoisier et al.\\ (2010) will be discussed in Sect.\\ 5. ",
        "conclusions": "Having been applied to situations with a magnetohydrodynamic background where both $\\uuN$ and $\\bbN$ have Roberts geometry, the proposed method has proven its potential for determining turbulent transport coefficients. In particular, effects connected with cross-correlations between $\\uuN$ and $\\bbN$ could be identified and are in full agreement with analytical predictions as far as available. No basic restrictions with respect to the magnetic Reynolds number or the strength of the mean field in the main run, which causes the nonlinearity of the problem, are observed so far. As a next step, of course, the simplifications in the hydrodynamics we used will be dropped, thus allowing to produce more relevant results and facilitating comparisons with work already done.  Due to the fact that we have no strict mathematical proof for its correctness, there can be no full certainty about the general reliability of the method. As a hopeful indication, in many cases, all four flavors of the method produce practically identical results, but occasionally some of them show, for unknown reasons, unstable behavior in the test solutions. Clearly, further exploration of the method's degree of reliance by including three-dimensional and time-dependent backgrounds is necessary. Homogeneity should be abandoned and backgrounds which come closer to real turbulence such as forced turbulence or turbulent convection in a layer are to be taken into account.  Thus, the utilized approach of establishing a test-field procedure in a situation where the governing equations are inherently nonlinear, although by virtue of the Lorentz force only, has proven to be promising. This fact encourages us to develop test-field methods for determining turbulent transport coefficients connected with similar nonlinearities in the momentum equation. An interesting target is the turbulent kinematic viscosity tensor, and especially its off-diagonal components that can give rise to a mean-field vorticity dynamo (Elperin et al.\\ 2007; K\\\"apyl\\\"a et al.\\ 2009), as well as the so-called anisotropic kinematic $\\alpha$ effect (Frisch et al.\\ 1987; Sulem et al.\\ 1989; Brandenburg \\& von Rekowski 2001; Courvoisier et al.\\ 2010) and the $\\Lambda$ effect (R\\\"udiger 1980, 1982). Yet another example is given by the turbulent transport coefficients describing effective magnetic pressure and tension forces due to the quadratic dependence of the total Reynolds stress tensor on the mean magnetic field (e.g., Rogachevskii \\& Kleeorin 2007; Brandenburg et al.\\ 2010).  \\appendix"
    },
    "1004/1004.5253_arXiv.txt": {
        "abstract": "Various possibilities are currently under discussion to explain the observed weakness of the intrinsic magnetic field of planet Mercury. One of the possible dynamo scenarios is a dynamo with feedback from the magnetosphere. Due to its weak magnetic field Mercury exhibits a small magnetosphere whose subsolar magnetopause distance is only about 1.7 Hermean radii. We consider the magnetic field due to magnetopause currents in the dynamo region. Since the external field of magnetospheric origin is antiparallel to the dipole component of the dynamo field, a negative feedback results. For an $\\alpha\\Omega$-dynamo two stationary solutions of such a feedback dynamo emerge, one with a weak and the other with a strong magnetic field. The question, however, is how these solutions can be realized. To address this problem, we discuss various scenarios for a simple dynamo model and the conditions under which a steady weak magnetic field can be reached. We find that the feedback mechanism quenches the overall field to a low value of about 100 to 150 nT if the dynamo is not driven too strongly.\\bigskip \\begin{keywords} Mercury, magnetic field, dynamo, magnetosphere \\end{keywords}\\bigskip ",
        "introduction": "The recent flybys of the MESSENGER spacecraft at planet Mercury confirm the existence of a large scale magnetic field \\citep{messenger_2009}. The dipole surface field, however, is roughly one to two orders of magnitude too weak to be commensurable with classical dynamo theory \\citep{wicht_2007, olson_christensen_2006}.  There are several approaches to explain this disagreement \\citep{Heimpel_2005, Stanley_2005, Christensen_2006, Matsushima_2006, glassmeier_2007} with different dynamo configurations. Here, we further study the feedback dynamo scenario suggested by \\citet{glassmeier_2007} who investigated the interaction of the dynamo and the magnetospheric field. They derived two stationary solutions and ascribed the weaker solution to Mercury's magnetic field. They however do not address the question how the dynamo reaches either of these solutions. Allowing a variable magnetopause which depends on the internal field and solar wind conditions, it is so far not conceivable how a dynamo can develope into a state where it can be quenched by the external feedback field. Therefore, the present study aims at discussing conditions under which a steady and weak magnetic field can evolve when the dynamo is exposed to a magnetospheric magnetic field. ",
        "conclusions": "Using a kinematic $\\alpha\\Omega$-dynamo in a feedback configuration, we have demonstrated that the feedback of the external field on the internal dynamo mechanism can indeed result in relatively small field strengths below $150$ nT as suggested by \\citet{glassmeier_2007}. However, in our simplified kinematic dynamo model the responsible quenching would only be sufficient in a narrow regime where the dynamo number does not exceed 18\\% of its critical value. If Mercury is captured in the quenched regime our model implies that the Hermean dynamo is unique. It should be noted here that alternative explanations for the weak Hermean dynamo field \\citep[e.g.][]{Stanley_2005,Christensen_2006} also require the assumption of special conditions for Mercury. The saturation field strength strongly depends on the assumed response function describing the dependence of the external field on the internal field strength. Unfortunately, very little is known about the underlying interaction, especially for a magnetopause close to the surface which would be appropriate for Mercury which is neccessary for our suggested feedback mechanism to work.\\\\  This paper is part of a series of studies examining the model of a feedback dynamo scenario. \\citet{glassmeier_2007} made use of extensively simplified models and examined stationary dynamo solutions without addressing the question how these stationary solutions could be realized.  This problem has been addressed in this study. We further consider an analytical solution to an approximation of the kinematic dynamo problem which allows us to examine the influence of the shape of the response function on the dynamo solution. The results could be useful for the application of the idea of a feedback dynamo to other astrophysical bodies such as gas giants close to their host star. Furthermore, we address the response function (also for higher magnetic multipoles) for Mercury by using a hybrid code simulating the interaction of Mercury's magnetosphere with the solar wind. Another investigation concerns how a three-dimensional, self-consistent, numerical dynamo model in approximate magnetostrophic balance \\citep{wicht_2002} reacts to an imposed uniform and constant-in-time external field.  From the results of these simulations we will know what kind of characteristic reactions of the dynamo we can expect when examining the full time dependent, 3D model with the exact and full magnetospheric response function."
    },
    "1004/1004.1193_arXiv.txt": {
        "abstract": "In order for a \\wdf\\ to achieve the Chandrasekhar mass, $M_C$, and explode as a Type~Ia supernova (SNIa), it must interact with another star, either accreting matter from or merging with it. The failure to identify the class or classes of binaries which produce SNeIa is the long-standing ``\\pr\\ problem''. Its solution is required if we are to utilize the full potential of SNeIa to elucidate basic cosmological and physical principles.  In single-degenerate models, a \\wdf\\ accretes and burns matter at high rates. Nuclear-burning \\wdf s (NBWDs) with mass close to $M_C$ are hot and luminous, potentially detectable as supersoft x-ray sources (SSSs). In previous work we showed that $> 90-99\\%$ of the required number of \\pr s do not appear as SSSs during most of the crucial phase of mass increase.  The obvious implication might be that double-degenerate binaries form the main class of progenitors. We show in this paper, however,  that many binaries that later become double-degenerates must pass through a long-lived NBWD phase during which they are potentially detectable as SSSs. The paucity of SSSs is therefore not a strong argument in favor of \\d-d\\ models. Those NBWDs that are the progenitors of double-degenerate binaries are likely to appear as symbiotic binaries for intervals $> 10^6$~years. In fact, symbiotic pre-double-degenerates should be common, whether or not the white dwarfs eventually produce Type~Ia supernovae.  The key to solving the Type~Ia \\pr\\ problem lies in understanding the appearance of NBWDs. Most of them do not appear as SSSs most of the time. We therefore consider the evolution of NBWDs to address the question of what their appearance may be and how we can hope to detect them. ",
        "introduction": "\\t1e have been used to map the expansion history of the Universe. The results have been exciting, indicating epochs of deceleration and acceleration, and suggesting the presence of dark energy (see, e.g., Riess et al.\\, 2007; Kuznetsova et al.\\, 2008). Unfortunately, we have not yet identified the astronomical systems that produce these distinctive explosions. (See Kotak 2009 and Branch et al.\\, 1995 for reviews.) Until we do, it will be impossible to understand or quantify the systematic uncertainties and to optimize the further use of \\t1e to explore physics and cosmology. The \\t1\\ \\pr\\ problem is therefore considered to be one of the key outstanding questions in astronomy today.  We know that the explosions occur when a \\wdf\\ gains mass from a binary companion. Indications from both theory and observation are that the supernova is triggered when the \\wdf\\ reaches the Chandrasekhar mass, $M_C$ (Mazzali et al.\\, 2007). What we don't know are the characteristics of the binary. Is the donor on the main sequence, evolved, or degenerate? Whatever the nature of the donor, it must be able to contribute enough mass to the \\wdf\\ to allow it to transition from its starting mass to $M_C.$  In single-degenerate binaries, the rate of mass transfer to a \\wdf\\ from a non-degenerate donor must be high enough that matter can be burned in either a quasisteady way or else during recurrent novae, thereby eliminating opportunities for more explosive nuclear burning that can reduce the mass of the \\wdf\\ (Iben 1982; Nomoto 1982; Fujimoto 1982). That is, the \\wdf s that reach $M_C$ must process accreting material; they are nuclear-burning \\wdf s (NBWDs) for long intervals. They are therefore potentially detectable as hot, luminous  supersoft x-ray sources (SSSs) during the crucial epoch when the \\wdf 's mass is increasing\\footnote{The known SSSs typically have $30\\, {\\rm eV} < k\\, T < 100\\, {\\rm eV}$ and $10^{36} {\\rm erg~s}^{-1} < L_X < 10^{38} {\\rm erg~s}^{-1}.$ NBWDs with mass near $M_C$ have surface temperatures and luminosities at the top end of these ranges. [See Figure 1 of the companion paper (Di\\thinspace Stefano 2010).]}. Some bright SSSs may be progenitors of Type~Ia supernovae (Rappaport et al.\\, 1994; \\rd\\ \\& Rappaport 1994). Nevertheless, the companion paper (\\rd\\ 2010; see also \\rd\\ et al.\\, 2010 and \\rd\\ 2007) shows conclusively that the majority of the progenitors do not appear as bright SSSs during intervals long enough ($\\sim 10^5$~yrs) to allow quasisteady burning of the necessary amounts of accreting matter. For both spiral and elliptical galaxies, the discrepancy is at least an order of magnitude, perhaps as much as two orders of magnitude. In addition, we found that existing data already place restrictions on sub-Chandrasekhar models. These restrictions may be tightened as additional exposures with {\\it Chandra} and {\\it XMM-Newton} are taken and more data are analyzed.  The most obvious interpretation of the mismatch is that it rules out single-degenerate models. In fact, a weaker measure of the mismatch was recently derived for six early-type populations, and  was used to argue that single degenerates can produce no more than $5\\%$ of the \\t1e\\ in early-type galaxies (Gilfanov \\& Bogd{\\'a}n 2010). If single-degenerates are ruled out, then the alternative would appear to be double-degenerate models in which two carbon-oxygen (C-O) \\wdf s execute a close orbit. In order for the \\wdf s to come to interact in a Hubble time, they must have had an opportunity to spiral toward each other in a common envelope. This paper explores the epoch prior to the common envelope. In \\S 2 we find that, immediately before the common envelope phase that produces a close double-degenerate, an epoch of nuclear burning on an accreting white dwarf is expected. In \\S 3 we predict the numbers of NBWDs required if the \\d-d\\ channel is the main route to \\t1e. We then compare these numbers with the numbers of SSSs detected in external galaxies, and find a large mismatch. In \\S 4 we discuss the significance and implications of the mismatch. Section 5 focuses on the symbiotic nature of pre-double-degenerate binaries, and discusses the prospects for using the distinctive symbiotic phase to test double-degenerate models for SNIa \\pr s. Our conclusions are presented in \\S 6.  The bottom line is that, for neither young nor old populations can the absence of SSSs be interpreted as evidence for the absence of NBWDs. If the photospheres of NBWDs are large, soft x-rays may not be emitted. In fact photospheric adjustments in known SSSs seem to occur (see. e.g., Greiner \\& \\rd\\ 2002). In addition, local mass associated with the system, such as winds, can absorb radiation from the \\wdf. In fact, the very binaries most likely to produce \\t1e must eject significant winds if they are to survive (\\rd\\ et al. 1997; \\rd\\ \\& Nelson 1996; \\rd\\ 1996). ",
        "conclusions": "We have shown that many pre-double-degenerate binaries pass through an epoch during which the first-formed \\wdf\\ accretes and burns matter from a giant companion. If \\dd s with total \\wdf\\ mass greater than $M_C$ comprise the major component of \\t1e \\pr s, then the numbers of symbiotics with NBWDs must be on the order of a thousand in galaxies such as our own. If the nuclear burning episodes produce SSS-like signatures, then we should be able to identify the pre-\\d-d\\ \\pr s of \\t1e in other galaxies by identifying SSSs. In \\S 4.1 we have sketched the steps needed to determine the numbers of SSSs in external galaxies that have the luminosities and temperatures predicted for the pre-\\d-d\\ \\pr s of \\t1e. Already, however, data from M31, M101, more than $380$ additional galaxies, and from the Milky Way strongly indicate that there is a mismatch of $\\sim 2$ orders of magnitude between the predicted numbers of SSSs and the numbers we actually detect in other galaxies. The result holds for young and old stellar populations.  We have already derived an even stronger result for single-degenerate \\pr s of \\t1e (\\rd\\ 2007; \\rd\\ et al.\\, 2010; \\rd\\ 2010). We falsified the hypothesis that the \\s-d\\ channel in which \\wdf s accrete and burn enough matter to reach $M_C$ is the primary \\pr\\ channel and that the NBWDs appear as SSSs\\footnote{ A similar result was claimed for old stellar populations based on limits on the diffuse soft emission from the Bulge of M31 and several early-type galaxies Gifavov \\& Bogd{\\'a}n 2010). In these cases, however, the bright, hot NBWDs with masses near $M_C$ would have been detected directly had they been there, so the previously-existing limits apply. } Combining the results for \\d-ds \\s-ds, we find that there are not enough SSSs in our own and other galaxies to explain the observed rates of \\t1e. Since the supernovae occur, the implication is that there is a disconnect between either (1)~mass infall at high rates and nuclear burning, or (2)~nuclear-burning and SSSs.  \\noindent{\\bf (1)} If mass infall doesn't lead to nuclear burning, this would seem to imply that a change is needed in our understanding of fundamental astrophysics. An alternative is that when nuclear-burning does occur, enough energy is released to deflect winds,  providing a kind of thermostat mechanism.  \\noindent{\\bf (2)} Mass accretion and nuclear burning is not always linked to SSS-like behavior. This is already known to be the case for many systems. For example the duty cycle of SSS-like behavior is low for recurrent novae; some of these, such as RS Ophiuchi are symbiotics (Nelson et al.\\, 2009 and references therein). In addition, absorption is expected because winds from symbiotics can absorb radiation from the \\wdf . In fact the nebulae associated with symbiotics illustrate this point (see, e.g., Kenyon \\& Murdin 2000).  With regard to the question of ``hiding'' the progenitors of \\t1e, symbiotics are intriguing for three reasons. First, of course, is the likelihood that at least some symbiotics are progenitors of \\t1e. In this paper we have focused on pre-\\d-ds that may be \\t1\\ \\pr s. Even among single-degenerates, however, there are models in which the \\pr\\ passes through a phase in which a giant donates mass to a \\wdf\\ either through winds or through Roche-lobe overflow (\\rd\\ 1996).  Second, symbiotics are examples of very bright systems that have proved difficult to identify. Estimates of the numbers of Galactic symbiotics are as high as $4 \\times 10^4$ (Magrini et al.\\, 2003). In spite of these large numbers, and in spite of the fact that symbiotics are, by their very natures highly luminous, the numbers of known symbiotics had stood in the low hundreds until recently. Within the past several years, $\\sim 1000$ candidates, now being checked, have been identified (see Corradi et al. 2009). Whatever the appearance of the \\pr s of \\t1e, they too must be very bright, at least during episodes of nuclear burning. They too, appear to be underrepresented, in that too-few candidates have been identified in our Galaxy and in other galaxies. Third, whether or not specific symbiotic binaries are \\t1\\ \\pr s, many contain NBWDs. The low numbers of SSSs we find is therefore relevant for understanding the appearance of symbiotics.  \\smallskip  \\noindent {\\bf Summary:} The key issue identified by the study of galaxy populations of SSSs is that there are too few of them to serve as the \\pr s of \\t1e. This applies to early-type and late-type galaxies. It applies to \\s-d\\ and \\d-d\\ models. To understand the progenitors of \\t1e, we must be able to predict the appearance of NBWDs and to identify a larger fraction of them in our own and other galaxies.      \\bigskip \\noindent{\\bf Acknowledgements:} It is a pleasure to acknowledge helpful conversations, most recently with Scott Kenyon and Jeno Sokoloski, and also with participants (especially Ed van den Heuvel, Lev Yungelson, and Jim Liebert), of the KITP conference and workshop on {\\it Accretion and Explosion} held at UC Santa Barbara in 2007. This work was supported in part by an LTSA grant from NASA and by funding from the Smithsonian Institution.  \\bigskip"
    },
    "1004/1004.5247_arXiv.txt": {
        "abstract": "The primary goal of this paper is to provide the evidence that can either prove or falsify the hypothesis that dark matter in the Galactic halo can clump into stellar-mass compact objects. If such objects existed, they would act as lenses to external sources in the Magellanic Clouds, giving rise to an observable effect of microlensing. We present the results of our search for such events, based on the data from the second phase of the OGLE survey (1996-2000) towards the SMC. The data set we used is comprised of 2.1 million monitored sources distributed over an area of 2.4 square degrees. We found only one microlensing event candidate, however its poor quality light curve limited our discussion on the exact distance to the lensing object.  Given a single event, taking the blending (crowding of stars) into account for the detection efficiency simulations, and deriving the {\\it HST}-corrected number of monitored stars, the microlensing optical depth is $\\tau=(1.55\\pm1.55) \\times 10^{-7}$. This result is consistent with the expected SMC self-lensing signal, with no need of introducing dark matter microlenses. Rejecting the unconvincing event leads to the upper limit on the fraction of dark matter in the form of MACHOs to $f<20$ per cent for deflectors' masses around 0.4 $\\msun$ and $f<11$ per cent for masses between 0.003 and 0.2 $\\msun$ (95 per cent confidence limit). Our result indicates that the Milky Way's dark matter is unlikely to be clumpy and form compact objects in the sub-solar-mass range. ",
        "introduction": "The Magellanic Clouds are harbours to millions of stars. The light of each of these objects can be magnified if another massive object is close enough to the line-of-sight connecting the observer and a distant star. \\citet{Paczynski1986} first realised that with the advent of CCDs, forthcoming massive photometric surveys could effectively test the hypothesis that dark matter in the Galactic halo can clump and form Massive Compact Halo Objects (MACHOs). These objects, if they existed, would act as lenses to more distant LMC/SMC stars, within the reach of current observing facilities. This brilliant, yet simple idea triggered several microlensing programs to emerge. The first detections of the microlensing effect were reported by the MACHO \\citep{MACHO}, OGLE \\citep{Udalski1993}, EROS \\citep{EROS}, MOA \\citep{MOA}, Angstrom \\citep{ANGSTROM}, POINT-AGAPE \\citep{POINTAGAPE}, and WeCaPP \\citep{WECAPP} microlensing teams.  For almost two decades, microlensing as an astrophysical tool has been very successful in finding objects which do not emit any or emit little light. The OGLE group alone have discovered over 4000 ordinary microlensing events to date. A list of exotic microlensing events includes detection of black-holes (\\eg \\citealt{OGLEBH}), planets (\\eg \\citealt{UdalskiOB05071}, \\citealt{Gaudi2008}, \\citealt{DongKB07400}), binary stars \\citep{Skowron2007binaries} and also a variety of effects such as the parallax (\\eg \\citealt{Smith2003parallax}, \\citealt{Gould2009terrestialParallax}), xallarap \\citep{Assef2006MACHO97SMC1}, etc.  However, since the microlensing field has evolved into a tool nowadays primarily concentrated on finding either the most distant or the smallest known planets, Paczy{\\'n}ski's original idea has been somewhat forgotten. The primary motivation for this paper is to fully explore the existing OGLE data to search for microlensing events towards the Magellanic Clouds. As of 2010 we have collected approximately 13 seasons (4 seasons of OGLE-II and 9 seasons of OGLE-III) of data for both the LMC and SMC. In \\cite{Wyrzykowski2009} (hereafter Paper I) we presented our first estimate of the microlensing optical depth towards the LMC from the OGLE-II data. The detection of two events led to the optical depth of $\\tau_{\\rm LMC} = (0.43\\pm0.33)\\times 10^{-7}$. However, the MACHO collaboration derived the optical depth of $\\tau_{\\rm LMC} = (1.0 \\pm 0.3) \\times 10^{-7}$ based on their 10 candidates (\\citealt{AlcockMACHOLMC}, \\citealt{BennettMACHOLMC}). If this number is compared to the optical depth for the Galactic halo entirely made of MACHOs, $\\tau_{\\rm halo} \\approx 4.7 \\times 10^{-7}$ (\\citealt{BennettMACHOLMC}), it gives the fractional contribution of $f = \\tau_{\\rm LMC}$/$\\tau_{\\rm halo} \\approx$ 20 per cent. On the other hand, the EROS collaboration has derived $\\tau_{\\rm LMC} < 0.36 \\times 10^{-7}$, which translates to $f<8$ per cent only \\citep{TisserandEROSLMC}. The OGLE--II estimate of $f<10$ per cent from Paper I, favours the EROS solution but the two detected events are also consistent with the expected LMC self-lensing signal.  The SMC has received somewhat less attention in terms of microlensing studies than the LMC. So far, only the EROS data were studied systematically and the optical depth of $\\tau_{\\rm SMC} =  (1.7 \\pm 1.7) \\times10^{-7}$ was derived for one microlensing event detected in their Bright Stars Sample \\citep{TisserandEROSLMC}. Another study of 5 years of the EROS data gives $f<25\\%$ for objects with masses $10^{-7}$ to $1\\msun$ \\citep{EROSSMC2003}. On the other hand, the MACHO collaboration estimated the optical depth to be $(2-3) \\times 10^{-7}$ based on their two events \\citep{Alcock1999MACHO98SMC1}. The SMC self-lensing estimates are in a range of $(0.4$--$1.8) \\times10^{-7}$ from N-body simulations by \\cite{Graff1999} and from analytical work of \\cite{EROSSMC1998}.  In this paper we extend our work from Paper I, on search for dark matter compact objects in the Galactic halo, to an independent SMC data set collected by OGLE during its second phase in years 1996--2000. The paper has the following structure.  First, the empirical optical depth estimator is described. Then the observational data used in the analysis are presented in Section \\ref{sec:data}. Next, in Section \\ref{sec:search} the search procedure for events is described and its yield presented. The detection efficiency of events and the calculation of the optical depth is discussed in Section \\ref{sec:results}--\\ref{sec:tau}. The paper concludes with a discussion of the results. ",
        "conclusions": "OGLE--II has provided a new and independent constraint on the presence of compact dark matter objects in the Galactic halo. In the LMC there were 2 candidate events found, both most likely due to self-lensing. A single candidate event was detected in the SMC data and its presence, if of microlensing nature at all, is consistent with the self-lensing scenario, in which the source is located at the far back end of the SMC and was lensed by a lens from within the SMC. The unusual position of the source on the colour-magnitude diagram could also be explained by, \\eg binarity of the source, and is generally not very trustworthy due to numerous ambiguities of the data.  The derived optical depth estimate for the single event indicates a value of  $\\tau_{SMC}=(1.55\\pm1.55)\\times  10^{-7}$, which is very close to the previous measurements obtained with other data sets by the EROS and MACHO collaborations and is in agreement with self-lensing estimates. However, its low statistical significance prevents deriving any reasonable conclusions on its origin, which may require further and more detailed studies of the structure of the SMC and resulting self-lensing optical depth.  The detection of a single and unconvincing candidate event in the OGLE--II SMC data only strengthen our previous conclusions reached in the OGLE--II LMC data (Paper I). The hypothesis that the Galactic halo is composed of compact objects is not favoured by these results. They show that the fraction of compact objects is close to zero and the only reason why we cannot completely exclude the compact-object hypothesis is due to limitations of the survey.  This verdict, as determined from the OGLE--II data, will be further validated by analysing data from the recently-completed OGLE--III survey, which covers a much wider area and has a duration of 8 years."
    },
    "1004/1004.0462_arXiv.txt": {
        "abstract": "We present first results of a {\\em Chandra}  X-ray observation of   the  rare oxygen-type Wolf-Rayet star WR 142 (= Sand 5 = St 3) harbored in the young, heavily-obscured cluster Berkeley 87. Oxygen type WO stars are thought to be the most evolved of the WRs and progenitors of supernovae or gamma ray bursts. As part of an X-ray survey of supposedly single Wolf-Rayet stars, we observed WR 142 and the surrounding Berkeley 87 region with {\\em Chandra} ACIS-I. We detect WR 142 as a faint, yet extremely hard X-ray source. Due to weak emission, its nature as a thermal or nonthermal emitter is unclear and thus we discuss several emission mechanisms. Additionally, we report seven detections and eight non-detections by {\\em Chandra} of massive OB stars in Berkeley 87, two of which are bright yet soft X-ray sources whose spectra provide a dramatic contrast to the hard emission from WR 142. ",
        "introduction": "Wolf-Rayet (WR) stars  are massive, highly-evolved stars nearing the end of their lives as supernovae (SN) or as collapsing objects emitting gamma-ray bursts (GRB, e.g. MacFadyen \\& Woosley 1999, Postnov \\& Cherepashchuk 2001, Georgy et al. 2009). WR stars undergo  rapid mass loss through strong winds, with initial masses of $\\geq$25M$\\sun$ (Crowther 2007). The classification of WR stars is determined spectroscopically in the optical and is divided among the nitrogen-rich WNs, carbon-rich WCs, and oxygen-rich WOs. For the most part, single WR stars are thought to follow the evolutionary path: O $\\rightarrow$ (LBV/RSG) $\\rightarrow$ WN $\\rightarrow$ WC $\\rightarrow$ WO $\\rightarrow$ SN [or GRB] (Conti et al. 1983, Crowther 2007). The initial mass of the O star determines whether it passes through an intermediate luminous blue variable (LBV) or red supergiant (RSG) phase (Crowther 2007).  Sanduleak (1971) noticed a class of stars that did not have planetary nebulae (PN) but displayed a WR-like spectrum with strong O VI doublet emission similar to those found in the central stars of PN named the O VI sequence by Smith and Aller (1969). These stars were suggested to be a separate WO sequence of Wolf-Rayet stars (rather than an extension of the WC sequence) by Barlow \\& Hummer (1982). WO stars are thought to be in the late helium-burning stage, or possibly the carbon-burning stage, where the enhanced oxygen abundance compared to less evolved WR stars is revealed by mass loss stripping (Barlow \\& Hummer 1982). Of the 298 galactic WR stars in the appended VIIth catalogue of galactic Wolf-Rayet stars (van der Hucht 2001, 2006), only four are of the rare WO spectral type, including the star WR 142, also known as Sand 5 (Sanduleak 1971) and St 3 (Stephenson 1966). Crowther et al. (1998) developed a new classification scheme using primary and secondary oxygen line ratios that confirmed previous classifications of WR 142 as a member of the subclass WO2 (Barlow \\& Hummer 1982, Kingsburgh et al. 1995).  Despite new discoveries of WR stars in the galactic plane, including a WO type as the exciting star of the planetary nebula Th 2-A (Weidmann et al. 2008), little is yet understood about the mechanisms that drive high-energy processes such as X-ray emission in single WR stars, particularly so for the rare WO stars. Although WR stars have much stronger winds and are more evolved chemically, the line-driven instabilities that are thought to give rise to soft X-rays (kT $<$ 1 keV)  from  shocks in O star winds may also be present in WR winds (Gayley \\& Owocki 1995). If that is the case, then WR stars may also be capable of producing soft X-rays from radiative wind shocks (Baum et al. 1992).  Few single WR stars had been studied with high sensitivity in X-rays until recently. Thus far, several apparently single WN stars have been observed to emit X-rays (Skinner et al. 2002, Ignace et al. 2003, Oskinova 2005). An ongoing  {\\em Chandra} and {\\em XMM-Newton}  survey has recently detected the apparently single WN stars WR 2, WR 18, WR 79a, and WR 134 with X-ray luminosities log $L_{\\rm X}$ $\\approx$ 32.2 - 32.7 erg s$^{-1}$ (Skinner et al. 2010b), comparable to some WR$+$OB binaries like $\\gamma^2$ Vel (WC8$+$O7) with log L$_{\\rm X}$ $=$ 32.9 erg s$^{-1}$ (Skinner et al. 2001) or WR 147 (WN8$+$OB) with log L$_{\\rm X}$ $=$ 32.83 erg s$^{-1}$ (Skinner et al. 2007; see also Sec. 4.4). Only upper limits of log L$_{\\rm X}$ $<$ 29.82 - 30.97 erg s$^{-1}$ exist from observations of single WC stars (Oskinova et al. 2003, Skinner et al. 2006). The WC sequence is thus either faint in the X-rays or possibly X-ray quiet. WR 142 just recently became the first WO star to be detected in the X-rays using {\\em XMM-Newton} (Oskinova et al. 2009).  WR 142 resides in the open cluster Berkeley 87. Table 1 summarizes the general properties of WR 142. Berkeley 87  lies in a heavily obscured region of the Milky Way in Cygnus. Initial cluster age estimates of $\\sim$ 1 - 2 Myr (Turner \\& Forbes 1982, hereafter TF82) have now been revised upward. A study restricted to three of the highest mass cluster stars suggests a slightly older age of $\\sim$ 3 Myr (Massey et al. 2001). A more recent study (Turner et al. 2010) gives $\\sim$ 5 Myr. Berkeley 87 is an interesting cluster containing $\\approx$ 105 cluster members identified by an optical study (TF82),  B-supergiants including HD 229059, an O8.5-O9 star BD+36 4032, a possible luminous blue variable Be star V439 Cyg, and the pulsating M3-supergiant BC Cyg. Additionally, OH masers and compact HII regions trace massive star formation 3$'$ - 9$'$ north of WR 142 (Argon et al. 2000, Matthews et al. 1973). At a distance of 1230 $\\pm$ 40 pc (Turner et al. 2006), the proximity of Berkeley 87 offers ample opportunity to learn about the properties of rare objects such as WR 142 as well as the surrounding cluster.  As part of an X-ray survey aimed at determining if single Wolf-Rayet stars that are not known to be in binary systems are X-ray emitters (specifically which spectral types), we observed  WR 142 with {\\em Chandra}. We are unaware of any evidence pointing to a companion of WR 142. WR 142 was chosen over other WO stars because of its relatively nearby distance,  low A$_{\\rm V}$ compared to other WR stars, and interesting surrounding cluster. Its astonishing supersonic wind (v$_{\\infty}$ $=$ 5500 km s$^{-1}$, Kingsburgh et al. 1995) and an unusual optical detection of diffuse emission from the C IV doublet (Polcaro et al. 1991) that is a signature of shocked gas make WR 142 an even more interesting target as one of the few oxygen-type WR stars.  The primary objective of the study presented here is to determine the X-ray properties of WR 142. X-ray emission from WR 142 was previously observed by {\\em XMM-Newton} (Oskinova et al. 2009), but {\\em Chandra} has several advantages over {\\em XMM-Newton}, including improved angular resolution (by a factor of $\\sim$ 4 - 5), sharp on-axis point-spread function (PSF) to provide checks on source extent, longer uninterrupted exposure that provides a continuous light curve and more stringent test for variability, and lower instrumental background. Such benefits allow for reliable source identification, especially when compared with optical positions; negligible background subtraction for lightcurves and spectra; and searches for the presence of diffuse shock emission around WR 142. We analyze the faint X-ray spectrum of WR 142 and discuss possible X-ray emission mechanisms. We also report {\\em Chandra} detections of seven other massive OB stars in Berkeley 87 and compare the much softer spectra from two X-ray bright OB stars to the hard spectrum of WR 142. In this paper, we focus on WR 142 and massive OB stars in Berkeley 87. A first overview of the {\\em Chandra} results for the cluster as a whole was presented by Skinner et al. (2010a) and the cluster will be discussed in more detail in a future paper. ",
        "conclusions": "\\begin{enumerate}  \\item {\\em Chandra} has detected hard, heavily-absorbed X-ray emission from the rare WO-type star WR 142. No soft emission below 2 keV was detected by {\\em Chandra}.  \\item The observed X-ray flux is consistent with the flux from a previous {\\em XMM-Newton} observation, within the uncertainties, and no significant X-ray variability was detected during the {\\em Chandra} observation.  \\item Due to the faint emission from WR 142, lack of prominent emission lines, and low numbers of counts, statistics are unable to distinguish between thermal and nonthermal X-ray emission mechanisms when fitting the spectrum and both have been considered. If the emission is thermal, very high plasma temperatures are inferred.  \\item In addition to WR 142, {\\em Chandra} detected seven luminous OB stars in Berkeley 87, while eight other B stars were undetected. The hard X-rays and excess absorption of WR 142 contrast with two X-ray bright OB stars in Berkeley 87, which display predominately soft X-ray emission and absorptions consistent with optical values.  \\item Though the X-ray emission mechanism in WR 142 is unclear, the hard X-ray spectrum observed by {\\em Chandra} could be explained by a colliding wind shock onto an as yet unseen companion, though the escape of the X-rays may require a geometry with the companion in front of WR 142. A colliding wind shock interpretation would also resolve an apparent contradiction of non-detections in the X-rays of single WC stars while WR 142 (a WO star) was detected. A key point not to be overlooked is that colliding winds can produce very hot plasma (as observed) even in the absence of magnetic fields.  \\item Alternatively, mechanisms that assume stellar magnetic fields such as MCWS or inverse Compton scattering (as formulated by Chen \\& White 1991a) could play a role in the X-ray emission of WR 142. However for MCWS, very strong surface fields of tens of kG would be required to effectively confine the powerful WR 142 wind (Sec. 5.4.3), also noted  by Oskinova et al. (2009) in their MCWS interpretation. There is so far no observational evidence of such strong B-fields in WR stars, nor do the existing X-ray data show any obvious signatures of impulsive X-ray flares that often accompany magnetic reconnection. If WR 142 has even a weak surface magnetic field, then X-ray production via inverse Compton scattering provides an attractive alternative to the extreme conditions required for magnetic wind confinement. A key question is whether the X-ray spectrum of WR 142 is indeed a power-law, as expected for inverse Compton scattering. To answer this question, a higher signal-noise X-ray spectrum from a much deeper observation will be required.  \\end{enumerate}"
    },
    "1004/1004.0712_arXiv.txt": {
        "abstract": "We present a complete explicit  $N=1$, $d=4$ supergravity action in an arbitrary Jordan frame with non-minimal scalar-curvature coupling of the form $\\Phi(z, \\bar z)\\, R$.   The action is derived by suitably gauge-fixing  the superconformal action.  The theory has a modified \\K \\, geometry, and it exhibits a significant dependence on the frame function $\\Phi (z, \\bar z)$ and its derivatives over scalars, in the bosonic as well as in the fermionic part of the action. Under certain simple conditions, the scalar kinetic terms in the Jordan frame have a canonical form.  We consider an embedding of the Next-to-Minimal Supersymmetric Standard Model (NMSSM) gauge theory into supergravity, clarifying  the Higgs inflation model recently proposed by  Einhorn and Jones.  We find that the conditions for canonical kinetic terms are satisfied for the NMSSM scalars in the Jordan frame,  which leads to a simple action. However, we find that the gauge singlet field experiences a strong tachyonic instability during inflation in this model. Thus, a modification of the model is required to support the Higgs-type inflation. ",
        "introduction": "Supersymmetry imposes certain restrictions on the non-supersymmetric models of particle physics and cosmology. A well known example of such restrictions is the fact that the supersymmetric version of the Standard Model (SM) of particle physics requires at least two Higgs superfields. Meanwhile, for cosmology Einstein equations have to be solved, therefore the supersymmetry embedding of the Higgs model inflation requires local supersymmetry, \\textit{i.e.} supergravity. Thus, one can try to see how the potential discovery of supersymmetry may affect various models of inflation, derived in the past in the context of general relativity coupled to scalar fields without supersymmetry. It would be interesting to find general restrictions, as well as to study particular models.  Here we are  motivated by a particular issue in cosmology, the so-called $\\xi \\phi ^{2}R$ coupling, which attracted a lot of attention starting from the early days of inflation \\cite{Futamase:1987ua}. Recently, it became also quite important in the context of SM inflation \\cite{Sha-1}.  Until now, the $N=1$, $d=4$ supergravity action in an arbitrary Jordan frame described by the frame function $\\Phi(z,{\\bar z})$, with arbitrary K{\\\"a}hler potential $\\mathcal{K}(z, {\\bar z})$,  holomorphic superpotential $W(z)$ and holomorphic function $f_{ab}(z)$, was not known.  Here we will derive this action, which is the first goal of this paper. This will be achieved by starting with the superconformal theory developed in \\cite{Kallosh:2000ve}, and by gauge-fixing all extra symmetries in order to get a general  supergravity action in Jordan  frame.  Our results generalize the formulation of $N=1$ supergravity in Jordan frame for the particular case in which the K{\\\"a}hler potential $\\mathcal{K}$ and the frame function are related by $\\mathcal{K}(z, {\\bar z})= -3 \\log (-\\ft13 \\Phi(z, {\\bar z}))$. The corresponding action in Jordan frame was derived in components in \\cite{Cremmer:1978hn,CFGVP-1}, and in superspace in \\cite{Girardi:1984eq,Wess:1992cp}. In our treatment, we will also specify the conditions required for the frame function to make the kinetic terms of the scalar fields canonical in the Jordan frame.  The non-minimal coupling of scalar fields to curvature is allowed by all known symmetries of the SM and general relativity. If one tries to describe the early universe using the particle physics SM coupled to gravity in the Einstein frame, one finds that: 1) the coupling $\\lambda$ of the Higgs field has to be of the order $10^{-13}$; 2) the mass of the Higgs field has to be of the order $10^{13}$ GeV. These conditions may be satisfied in a general theory of a scalar field, but not in the simplest version of the standard model. However, if the $\\xi \\phi ^{2}R$ coupling is included, \\textit{i.e.} if the embedding of the particle physics SM into the Jordan frame gravity is considered, a satisfactory description of cosmology for the Higgs mass in the interval between $126$ and $194$ GeV can be found  \\cite{Sha-1}.  This is possible for very large values of the non-minimal scalar-curvature coupling $\\xi \\sim 10^{4}$. The model predicts the cosmological parameters $n_{s}\\approx 0.97$, and $r\\approx 0.003$, which are consistent with cosmological observations. Thus, this model provides very interesting predictions, which will be testable both at LHC and by a Planck satellite.  When this work was in progress,  a very interesting proposal \\cite{Einhorn:2009bh} was made how to generalize the model of Bezrukov-Shaposhnikov \\cite{Sha-1} in presence of supersymmetry. Under certain assumptions, it was found that slow regime inflation is not possible within the supergravity embedding of the Minimal Supersymmetric Standard Model (MSSM), but rather it is possible for the NMSSM (see \\textit{e.g.} \\cite{Ellwanger:2009dp} for a recent review of NMSSM).  In the present paper we will study the supergravity embedding of the NMSSM and look for a consistent cosmological models of the Higgs-type inflation.  Firstly, we will derive the complete $N=1$ action in the general Jordan frame, where it is very simple and has interesting features. This will help to clarify the meaning of the large non-minimal $\\xi \\phi ^{2}R$ coupling in the context of supergravity. In particular, the origin of the canonical kinetic terms of all scalars of the NMSSM in the Jordan frame is explained, whereas in the Einstein frame scalar kinetic terms are generally very complicated.  Secondly, we will study the theory as a function of all three chiral multiplets, namely two Higgs doublets and a singlet,  and analyze various directions in the space of scalar fields. In particular, in \\cite{Einhorn:2009bh} it was shown that a slow-roll inflationary regime is possible in NMSSM when the Higgs fields move in the  $D$-flat direction of the two Higgs doublets $H_u$ and $H_v$, assuming that the gauge singlet $S$ is small. However, it was not clear whether this last assumption is justified, {\\it i.e.} whether $S = 0$ corresponds to a minimum of the potential with respect to the field $S$ when inflation takes place in the $D$-flat direction of the two doublet Higgs fields. We will show that, unfortunately, the potential of the field $S$ has a sharp maximum near $S = 0$ in this regime. This means that the inflationary regime studied in \\cite{Einhorn:2009bh} is unstable, and a search for more general models is required to find a supersymmetric version of the Higgs-type inflation.  The paper is organized as follows. In Sec. \\ref{ss:N1Jordan} we present the complete explicit  $N=1$, $d=4$ supergravity action in an arbitrary Jordan frame with non-minimal scalar-curvature coupling of the form $\\Phi(z, \\bar z)R$. This includes the bosonic as well as fermionic action. In the special case in which the frame function $\\Phi (z,\\bar{z})$ is related to the K{\\\"a}hler potential by the relation $\\mathcal{K}(z,{\\bar z})= -3 \\log (-\\ft13 \\Phi(z,{\\bar z}))$, the action reduces to the one derived in \\cite{Cremmer:1978hn,CFGVP-1}. In the case $\\Phi=-3$, the action becomes the well known action of $N=1$ supergravity in the Einstein frame.  Sec. \\ref{ss:BosN1Frames} is devoted to a detailed discussion of the bosonic part of the supergravity action, which is especially important for cosmology. In particular, sufficient conditions for the kinetic terms of scalars to be canonical are specified.  Sec. \\ref{ss:sugraNMSSM} starts with a short description of the Higgs-type inflation with non-minimal scalar-curvature coupling. Then, we proceed with an attempt to generalize this model to the supersymmetric case. For this purpose, we study the embedding of the NMSSM into supergravity, focussing on the Einhorn-Jones cosmological model \\cite{Einhorn:2009bh}. We study this model in the Jordan as well as in the Einstein frame. The dependence of the potential on the singlet gauge field $S$, as well as at large values of the Higgs fields in a $D$-flat direction of the two Higgs doublets, is explicitly computed. We find that this potential has a maximum for small values of $S$ near the inflationary trajectory. The resulting instability disallows the inflationary regime in the model of \\cite{Einhorn:2009bh}, unless some way of stabilizing the field $S$ is found. Sec. \\ref{ss:derivN1Jordean} provides a detailed derivation of the Jordan frame supergravity action presented in Sec. \\ref{ss:N1Jordan}, by gauge-fixing the extra symmetries of the superconformal action. Finally, the Appendix contains a discussion of the cosmological behavior of the angle $\\beta$ between the two components of the Higgs field. ",
        "conclusions": "The main goal of our paper was to derive a complete formulation of $N=1$, $ d=4$ supergravity in a generic Jordan frame. We found that, in general, this formulation is very non-trivial. It involves modified {K{\\\"a}hler geometry (in the sense specified in our treatment), and it gives rise to many new complicated terms in the supergravity Lagrangian.}  However, we identified a subclass of theories where the resulting formulation is remarkably simple. This subclass includes the recently proposed model of Einhorn and Jones \\cite{Einhorn:2009bh}, which was introduced as an $N=1$ supergravity realization of the Higgs field inflation \\cite{Sha-1}. We found that the inflationary regime in this model is unstable.  Hopefully, however, the general formalism developed in our paper may allow one to find new realistic inflationary models in supergravity. As a starting approach, one can simply study in the Jordan frame several classes of inflationary models in supergravity, which were found long time ago in the Einstein frame. As shown by the example of the Higgs inflation, sometimes it is helpful to identify and study various physical features of the cosmological models by switching from one frame to another."
    },
    "1004/1004.1302_arXiv.txt": {
        "abstract": "The scientific community is presently witnessing an unprecedented growth in the quality and quantity of data sets coming from simulations and real-world experiments. To access effectively and extract the scientific content of such large-scale data sets (often sizes are measured in hundreds or even millions of Gigabytes) appropriate tools are needed. Visual data exploration and discovery is a robust approach for rapidly and intuitively inspecting large-scale data sets, e.g. for identifying new features and patterns or isolating small regions of interest within which to apply time-consuming algorithms. This paper presents a high performance parallelized implementation of Splotch, our previously developed visual data exploration and discovery algorithm for large-scale astrophysical data sets coming from particle-based simulations. Splotch has been improved in order to exploit modern massively parallel architectures, e.g. multicore CPUs and CUDA-enabled GPUs. We present performance and scalability benchmarks on a number of test cases, demonstrating the ability of our high performance parallelized Splotch to handle efficiently large-scale data sets, such as the outputs of the Millennium II simulation, the largest cosmological simulation ever performed. ",
        "introduction": "\\label{intro}  Nowadays the technological advances in instrumentation and computing capability impact profoundly on the dramatic growth in the quality and quantity of astrophysical data sets obtained from observational instruments, e.g.\\ sky surveys \\cite{sdss}, \\cite{lofar}, or large-scale numerical simulations, e.g.\\ the Millennium II simulation \\cite{2009MNRAS.398.1150B}.  The main characteristic of modern astrophysical data sets is extremely large sizes (in the order of hundreds of Gigabytes) requiring storage in extremely large-scale distributed databases. The forthcoming next-generation astrophysical data sets are expected to exhibit massively large sizes (in the order of hundreds of Terabytes), e.g.\\ \\cite{lsst}. To obtain a comprehensive insight into modern astrophysical data sets, astronomers employ sophisticated data mining algorithms, often at prohibitively high computational costs. Visual data exploration and discovery tools are then exploited in order to rapidly and intuitively inspect very large-scale data sets to identify regions of interest within which to apply time-consuming algorithms. Such tools are based on a combination of meaningful data {\\it visualizations} and user interactions with them.  This apporoach can be a very intuitive and ready way of discovering and understanding rapidly new correlations, similarities and data patterns. For on-going processes, e.g. a numerical simulation in progress, visual data exploration and discovery allow constant monitoring and - if anomalies are discovered - prompt correction of the run, thus saving valuable time and resources.  The data exploration tools traditionally employed by astronomers are limited either to processing and displaying of 2D images (see, e.g., \\cite{iraf}, \\cite{midas}, \\cite{sao}, \\cite{gaia}) or to generation of meaningful 2D and 3D plots (e.g.\\ \\cite{gnuplot}, \\cite{supermongo}, \\cite{idl}).   To overcome the shortcomings of traditional tools, a new generation of software packages is now emerging, providing astronomers with robust instruments in the context of large-scale astrophysical data sets (e.g.\\ \\cite{paraview}, \\cite{aladin}, \\cite{topcat}, \\cite{visivo1} and \\cite{visivo2}, \\cite{3dslicer}, \\cite{splash} and \\cite{visit}). The underlying principles are exploitation of high performance architectures (i.e.\\ multicore CPUs and powerful graphics boards), interoperability (different applications can operate simultaneously on shared data sets) and collaborative workflows (permitting several users to work simultaneously for exchanging information and visualization experiences).   This paper describes a high performance implementation of Splotch \\citep{2008NJPh...10l5006D}, our previously developed ray-tracing algorithm for effective visualization of large-scale astrophysical data sets coming from particle-based computer simulations. N-Body simulations constitute prime examples of particle-based simulations, typically associated with very large-scale data sets, e.g.\\ the Millennium II simulation \\citep{2009MNRAS.398.1150B}. This is a simulation of the evolution of a meaningful fraction of the universe by means of 10 billion fluid elements ({\\it particles}) interacting with each other through gravitational forces. The typical size of a snapshot of the Millennium II simulation is about 400 Gigabytes representing a particle's ID, position and velocity together with additional properties, e.g.\\ local smoothing length, density and velocity dispersion. For further details on the Millennium II simulation and other works about the visualization of its data sets, the reader is referred to \\citep{2009MNRAS.398.1150B}, \\cite{fraedrich2009} and \\cite{Szalay2008}.  The fundamentals and the traditional sequential operation of Splotch are reviewed in section 2. Section 3 discusses our strategy for parallelizing Splotch based on different approaches that are suitable for a variety of underlying architecture configurations. Our implementations are Single Instruction Multiple Data (SIMD) designs founded on the MPI library \\cite{mpi} in order to support distributed multicore CPUs and CUDA \\cite{cuda} for exploiting not only currently available but also forthcoming next-generation multiple GPUs. The advantage of adopting several parallelization solutions is that we can deploy them simultaneously on hybrid architectures, e.g.\\ mixed hardware architectures consisting of a large number of multicore CPUs and CUDA-enabled GPUs. Benchmarks for our parallelization designs and a discussion on the Millenium II visualization are presented in section 4. Finally section 5 outlines a summary of our work and includes pointers to future developments. ",
        "conclusions": "\\label{conclusions}   In this paper we have described a high performance parallel implementation of Splotch able to execute on a variety of high performance computing architectures. This is due to its hybrid nature exploiting multi-processor systems adopting an MPI based approach, multi-core shared memory processors exploiting OpenMP, and modern CUDA enabled graphics boards. This allows to achieve extremely high performance overcoming the typical memory barriers posed by small personal computing systems, commonly adopted for visualization. Finally, as parallel Splotch is implemented in ISO C++ and is completely self-contained (in other words it does not require any complementary library apart from MPI, OpenMP and CUDA), the code is highly portable and compilable over a large number of different architectures and operating systems. We discussed test results based on custom-made benchmark data sets and also the Millennium II simulation, which is the largest cosmological simulation currently available containing 10 billion particles.  Our future work will involve porting and running the parallelized Splotch on hybrid architecture computing systems containing several multiprocessor computers with CUDA enabled graphics boards, thus exploiting MPI and CUDA simultaneously. Several optimizations are also planned for our CUDA implementation, e.g.\\ unrolling short loops for improved control flow or using local/shared memory to accelerate data fetching. Mechanisms for optimal load balance between CPU and the graphics processor and between several graphics processors when these are available should also be considered. We will investigate possibilities for designing advanced scheduling mechanisms to minimize the idle CPU times contained in the current implementation. Finally, we will explore the opportunities offered by the OpenCL library, in order to exploit a wider range of computing architectures."
    },
    "1004/1004.3288_arXiv.txt": {
        "abstract": "Two populations of minor bodies in the outer Solar System remain particularly elusive: Scattered Disk objects and Sedna-like objects. These populations are important dynamical tracers, and understanding the details of their spatial- and size-distributions will enhance our understanding of the formation and on-going evolution of the Solar System. By using newly-derived limits on the maximum heliocentric distances that recent pencil-beam surveys for Trans-Neptunian Objects were sensitive to, we determine new upper limits on the total numbers of distant SDOs and Sedna-like objects. While generally consistent with populations estimated from wide-area surveys, we show that for magnitude-distribution slopes of $\\alpha \\gtrsim 0.7-1.0$, these pencil-beam surveys provide stronger upper limits than current estimates in literature. ",
        "introduction": "A number of deep, narrow-angle ``pencil-beam'' surveys for distant Solar System objects have been undertaken over the past two decades, first appearing in the literature in Tyson et al. (1992). These surveys avoid the issue of trailing losses in long exposures by taking a large number of shorter exposures, predicting the sky rate of motion of sources of interest, then compensating for this motion in software before finally stacking the images. Due to the large number of rates at which images must be combined in order to properly compensate for the range of motions real objects can have, the orbital range over which this method is applied is limited to maintain computational feasibility.  Parker \\& Kavelaars (2010 submitted; hereafter P10) re-characterize the orbital limits of several published pencil-beam surveys and show that these orbital limits are poorly characterized in literature. As the re-derived maximum heliocentric distances these surveys were sensitive to ranges from $150-400$ AU, we find that these surveys were sensitive to several dynamically interesting yet currently poorly constrained populations; namely, Scattered Disk Objects (SDOs, eg., Trujillo et al. 2000; hereafter T00) and Sedna-like objects (SLOs, eg., Schwamb et al., 2009; hereafter S09). SDOs are highly-eccentric, non-resonant objects that have perihelia that interact with Neptune, and may be the source of Jupiter-family comets (eg., Duncan \\& Levison 1997). SLOs are long-period objects that have perihelia high enough ($> 70$ AU) that they are out of reach of the giant planets but aphelia not distant enough for galactic tides to have significant effect on their orbits. Their current orbits require emplacement mechanisms that may be linked to dynamics of the Solar birth cluster (eg., Brasser et al. 2006), interactions with ``rogue'' planets (eg., Gladman \\& Chan 2006), or close stellar passages (eg., Kenyon \\& Bromley 2004). Few objects (one, in the case of Sedna) of either population have been discovered, and our understanding of these populations is severely hampered by these limited samples. Here we use the newly-derived heliocentric distance limits from P10 in conjunction with a simple survey simulator in order to determine new limits on the SDO and SLO populations. ",
        "conclusions": "We have used the re-characterization by Parker \\& Kavelaars (2010) of the orbital sensitivity of existing pencil-beam surveys in literature to derive new upper limits on distant populations of minor bodies in the Solar System. P10 considered three surveys (Fraser et al. 2008; Fraser \\& Kavelaars 2009; and Fuentes et al. 2009) and found that for each survey the outer limit was significantly more distant than claimed in the original publication. Using these newly-derived outer limits, we determined the constraints these surveys (plus Bernstein et al. 2004) put on the Scattered Disk and Sedna-like populations.  The combination of survey types allows us to limit the maximum possible size of the population for any CDF slope. Our limits constrain the Scattered Disk to be less than $5-40$ times more populous than the Main Kuiper Belt, depending on the radial distribution of material. For most plausible CDF slopes, SDOs are better constrained by relatively shallow, wide area surveys like those presented in Trujillo et al. (2000). For both a $d^{-1.5}$ and a uniform radial distribution model, we estimate the maximum population size at any CDF slope to less than $N(D>100$ km$) \\leq 3.5 - 25\\times10^5$, respectively.  A strong sensitivity to the assumed radial distribution is demonstrated by the Sedna-like population, as up to $1300$ times the $D<100$ km Main Kuiper Belt population could reside on Sedna-like orbits (as assumed by the limits presented in S09) and still remain undetected in existing surveys, whereas if the material follows a $d^{-1.5}$ radial distribution this upper limit is reduced to 40 times the current population of the Main Kuiper Belt.  At present the upper-limits on the population size determined by lack of detection in existing pencil-beam surveys do not appear to conflict with the population estimates based on wide-area, shallow surveys. This implies that the lack of distant detections in current-generation surveys are not indicative of any incongruences in survey sensitivity calibrations or other issues."
    },
    "1004/1004.5467_arXiv.txt": {
        "abstract": "We place new constraints on the primordial local non-Gaussianity parameter $f_{NL}$ using recent Cosmic Microwave Background anisotropy and galaxy clustering data. We model the galaxy power spectrum according to the halo model, accounting for a scale dependent bias correction proportional to $f_{NL}/k^2$. We first constrain $f_{NL}$ in a full $13$ parameters analysis that includes $5$ parameters of the halo model and $7$ cosmological parameters. Using the WMAP7 CMB data and the SDSS DR4 galaxy power spectrum, we find $f_{NL}=171^{+140}_{-139}$ at $68\\%$ C.L. and $-69<f_{NL}<+492$ at $95\\%$ C.L.. We discuss the degeneracies between $f_{NL}$ and other cosmological parameters. Including SN-Ia data and priors on $H_0$ from Hubble Space Telescope observations we find a stronger bound: $-35<f_{NL}<+479$ at $95\\%$. We also fit the more recent SDSS DR7 halo power spectrum data finding, for a $\\Lambda$CDM+$f_{NL}$ model, $f_{NL}=-93\\pm128$ at $68\\%$ C.L. and $-327<f_{NL}<+177$ at $95\\%$ C.L.. We finally forecast the constraints on $f_{NL}$ from future surveys as EUCLID and from CMB missions as Planck showing that their combined analysis could detect $f_{NL}\\sim 5$. ",
        "introduction": "The standard paradigm of structure formation relies on the inflation \\cite{starob, albrecht}. As shown in \\cite{M,Bardeen1983}, quantum-mechanical fluctuations in the scalar field driving inflation lead to primordial density perturbations responsible of the large scale structures we observe today. Although the simplest assumption is that these fluctuations were Gaussian distributed \\cite{bardeen}, there are several inflationary models \\cite{linde2}-\\cite{sasaki} involving the existence of a primordial non-Gaussianity. A detection or exclusion of non-Gaussianities would hence be of fundamental interest for the understanding the physics of the primordial Universe. Several cosmological observables and methods can be used to constrain non-gaussianities. \\\\Cosmic Microwave Background (CMB) anisotropies provide the most direct method for the detection of primordial non-Gaussianity (see e.g. \\cite{Bartolo:2010qu}) through, for example, measurements of the three-point correlation function \\cite{Matarrese:1997sk} (or equivalently the bispectrum) which is non-zero in presence of non-Gaussianities. The large scale structure of the Universe is also affected by non-Gaussianities that may be detected by looking at the bispectrum or the trispectrum of galaxy distribution. The abundance of galaxy clusters, that depends on the tails of the density probability distribution, is also sensitive to any deviation from gaussianity. Non-Gaussianity has a direct impact on the clustering of dark matter halos by changing their mass and correlation function (\\cite{lu_mat}-\\cite{Dalal:2007cu}).\\\\ A common way to parameterize primordial non-Gaussianities is to introduce a quadratic correction to the potential \\cite{Gangui:1993tt}\\cite{Komatsu:2000vy}: \\begin{equation}\\label{fnl} \\Phi=\\phi+f_{NL}(\\phi^2-\\langle\\phi^2\\rangle) \\end{equation} where $\\Phi$ is the primordial potential and $\\phi$ is a gaussian random field. In this case the non-Gaussianity is a local type correction whose amplitude is given by $f_{NL}$. The most recent constraint on $f_{NL}$ from CMB gives $-10<f_{NL}<+74$ at $95\\%$ C.L. from bispectrum analysis of WMAP-7 years data \\cite{Komatsu:2010fb}, improving the WMAP5 constrain ($-9<f_{NL}<+111$) \\cite{Komatsu:2008hk}. The authors of \\cite{Calabrese:2009bu} used a different estimator applied to WMAP5 data, finding $f_{NL}=-13\\pm62$ at $68\\%$ C.L., while \\cite{Rudjord:2009mh}, using the needlet bispectrum applied to the same data, found $f_{NL}=+84\\pm40$ at $68\\%$ C.L..\\\\In \\cite{Verde:1999ij}-\\cite{Reid:2010vc}-\\cite{Dalal:2007cu}-\\cite{Matarrese:2008nc}-\\cite{Slosar:2008hx} it has been shown that a quadratic correction to the potential like that of equation (\\ref{fnl}) produces a scale dependence in the bias of the galaxy clustering with respect to matter distribution. In particular, a scale dependent term $\\Delta b(k)$ arises in the halo bias on larger scales (smaller $k$) and is proportional to $f_{NL}$ ($\\Delta b(k)\\propto f_{NL}/k^2$) hence with galaxies being more (less) clustered for positive (negative) values of $f_{NL}$. The authors of \\cite{Slosar:2008hx} analyzed the galaxy power spectrum of luminous red galaxies (LRG) of the Sloan Digital Sky Survey (SDSS) to constrain this scale dependence of the bias, putting the constraint $-21<f_{NL}<+209$ at $95\\%$ C.L.. In the same work other large scale datasets have been used to constrain non-gaussianity (Quasars, integrated Sachs-Wolfe effect data and photometric LRG sample) finding $-29<f_{NL}<+70$ at $95\\%$ C.L. from the combination of all datasets. Recently, the authors of \\cite{Xia:2010yu} obtained $+25<f_{NL}<+117$ at $95\\%$ C.L. from the combination of WMAP-7 years data \\cite{Komatsu:2010fb}, Baryonic Oscillations data from SDSS and Two-degree Field Galaxy Redshift Survey (2dFGRS) \\cite{Percival:2009xn} and Supernovae distance moduli measurements \\cite{Kowalski:2008ez} with auto correlation function measurements of radio sources from NRAO VLA Sky Survey \\cite{condon98}, claiming a detection of non gaussianity at $\\sim 3\\sigma$.  In this paper we follow the methodology of \\cite{Slosar:2008hx} and constrain $f_{NL}$ by looking at the scale dependence of the bias in current galaxy surveys data. We implement the calculation of galaxy and halo power spectrum using the halo-model (see \\cite{Cooray:2002dia} and section \\ref{halomodel}) and we include in it the non-Gaussian scale-dependent correction to the bias. We then place constraints on $f_{NL}$ by comparing this model to the Sloan Digital Sky Survey (SDSS) \\cite{www.sdss.org} galaxy power spectrum data \\cite{Tegmark:2006az} and to the halo power spectrum data obtained from the luminous red galaxies (LRG) sample \\cite{Reid:2009xm}. We include in the analysis the WMAP-7 years cosmic microwave background anisotropy data \\cite{Komatsu:2010fb}. We also fit the LRG galaxy power spectrum data to the same model, including Hubble constant measurements from Hubble Space Telescope (HST, \\cite{Riess:2009pu}) and Supernovae distance moduli measurements for the Union dataset \\cite{Kowalski:2008ez}. Results are shown in section \\ref{results}. We finally forecast the power of future Galaxy surveys in constraining non gaussianity by generating mock data for galaxy power spectrum using specifications of EUCLID \\cite{Refregier:2006vt} survey combined with mock data from Planck \\cite{:2006uk} satellite and showing in section \\ref{forecasts} that the combination of data from these experiments could reach the precision required to detect even small non-Gaussianities. ",
        "conclusions": "We place new constraints on the local type non-gaussianity parameter $f_{NL}$ by looking at the scale dependence of the halo bias (at small wave vectors) in the recent galaxy and halo power spectra measurements from LRG sample of the Sloan Digital Sky Survey. We fit 2006 SDSS power spectra data with an halo model consisting of 5 parameters plus 7 cosmological paramers and $f_{NL}$. Our large parameter space and the the restriction of the dataset to relatively small scales ($k>0.01 h Mpc^{-1}$) leads to a weak constraint: $-69<f_{NL}<+492$ at $95\\%$ C.L.. We show and discuss degeneracies with halo model parameters. When including both Type Ia Supernovae and $HST$ data the $1\\sigma$ error on $f_{NL}$ is reduced of about $\\sim10\\%$. We use also 2009 halo power spectra data obtained from SDSS LRG sample finding a slightly better constraint $-327<f_{NL}<+177$ at $95\\%$ C.L., again limited by the fact that the dataset does not extend below $k\\sim0.02h Mpc^{-1}$. We also forecast the constraints obtainable from datasets of a survey like EUCLID when combined with Planck CMB data, finding that these surveys could reach the accuracy required to detect even small non-gaussianities as $f_{NL}=+5$, thus confirming the power of this method. Finally we discuss the possible systematics and theoretical uncertainties that may affect the results.  \\textit{Acknowledgements}\\\\\\\\ FDB thanks UCI Center for Cosmology for support and hospitality while this research was conducted. This work was supported by NSF CAREER AST-0645427 and NASA NNX10AD42G at UCI. This research has been partially supported by the ASI/INAF agreement I/072/09/0 for the Planck LFI Activity of Phase E2."
    },
    "1004/1004.3378.txt": {
        "abstract": "We   show  that  several   features  reminiscent  of short-hard  Gamma-ray Bursts (GRBs) arise  naturally when  Quark-Novae occur in  low-mass X-ray  binaries born with  massive neutron stars ($\\ge 1.6M_{\\odot}$) and harboring  a circumbinary disk. Near the end of the first accretion phase, conditions are  just right for  the explosive conversion of the neutron star to a quark star (Quark-Nova). In our model, the subsequent interaction of material from  the neutron star's ejected crust  with the circumbinary disk explains  the duration, variability and  near-universal nature of the prompt emission in short-hard GRBs. We also describe a statistical approach  to  ejecta  break-up  and  collision to obtain the  photon spectrum  in our  model, which turns out remarkably similar to the empirical Band function \\citep{band93}. We apply the model to the fluence and spectrum of GRB 000727, GRB 000218, and GRB980706A  obtaining excellent fits. Extended  emission (spectrum and duration)  is explained  by shock-heating  and ablation  of  the white dwarf  by  the  highly  energetic  ejecta.  Depending  on  the orbital separation   when  the  Quark-Nova occurs,  we isolate interesting regimes within  our model when  both  prompt and extended emission can occur. We find that the spectrum  can carry signatures typical of Type Ib/c SNe although these should appear less luminous than normal type Ib/c SNe. Late X-ray activity is due to accretion onto  the quark  star as well  as its spin-down luminosity. Afterglow activity arise  from the expanding shell of  material from the shock-heated  expanding circumbinary disk. We find a correlation between the duration  and spectrum of short-hard GRBs as well as modest hard-to-soft time evolution of the peak energy. ",
        "introduction": "Intensive  sampling of  gamma-ray  burst (GRB)  light  curves by  {\\it Swift}  has revealed  that at  least  in some  short GRBs  (henceforth shGRBs),  prompt emission  is  followed by  softer, extended  emission lasting tens  to hundreds of  seconds. For example, the  lightcurve of GRB 050709  \\citep{villasenor05} has  a short-hard  pulse $T_{90} \\sim  0.2$s and  a  long-soft  pulse $T_{90}  \\sim  130$s; GRB  050724 \\citep{barthelmy05} has prominent emission lasting for $\\sim 3$ s followed by  a long,  soft, less prominent  emission peaking  at $\\sim 100$  s  after  the  trigger,  while XRT  observations  reveal  strong flare-like activities  within the first  hundreds of seconds.   In GRB 060614,  the  lightcurve  has  a  short-hard episode  followed  by  an extended soft  emission component  with strong spectral  evolution. On average, the principal properties  of the intense prompt component ($< 2$  s) of short  bursts with  and without  extended emission  (EE) are indistinguishable,  suggesting  that  short bursts  are  ``universal'' \\citep{norris08}.   While  the origin  of  such  extended emission (seen roughly  in a quarter of {\\it Swift}  shGRBs) is debated, it can provide some clues to  the identity of the elusive mechanism of shGRBs.   The association of some shGRBs  with early type galaxies that have low star formation rates suggests that short bursts arise from a different progenitor mechanism than long bursts and that neutron star-black hole (NS-BH)  or double  neutron  star  (NS-NS) mergers  might  be at  play \\citep{bloom06}.  But how  does  one obtain  extended emission in  this case? Extended  accretion episodes in  NS-BH mergers have  been proposed  \\citep{barthelmy05},  while \\citet{rosswog06} suggest that  some debris may  be launched during the  merger process, which   would   fall   back    later   to   power   flares   at   late times.  Alternatively,  disk  fragmentation \\citep{perna06} or magnetic field  barrier near  the accretor \\citep{proga06} may induce intermittent  accretion that power the flares.   Other types of mergers such  as a  white dwarf-neutron star  (WD-NS) merger  \\citep{king07} have  been  advanced  to  interpret  shGRBs.  \\citet{dai06} argue that the  final product of  a NS-NS merger may  be a heavy, differentially-rotating NS, whose post-merger magnetic activity would give rise  to flares following the merger  events. The `off axis collapsar' models \\citep{lazzati10}  and the millisecond magnetar model \\citep{metzger08}  offer explanations for why the extended emission from short GRBs may resemble the prompt emission from long GRBs.  However, it is unclear how this model can explain the apparent different galactic environments in which short and long GRBs are found. In both types of bursts, the photon spectrum  is well fitted by  the Band function, hinting at a kind of universality in the underlying physical model.    Here, we will explore an alternative model based on a Quark-nova explosion occurring in a low mass X-ray binary (LMXB). Such an explosion can happen if the neutron star, in its spin-down evolution, reaches the quark deconfinement density and subsequently undergoes a phase transition to the more stable strange quark matter phase \\citep{itoh70,bodmer71,witten84}, resulting in a conversion front that propagates towards the surface in the detonative regime. This hypothesis is motivated by recent work on numerical simulations of the conversion of neutron matter to strange quark matter \\citep{niebergal10b}. Previous works that have focused either exclusively on conversion through strange matter diffusion \\citep[eg.][]{Olinto87,HB88} or use pure hydrodynamics \\citep[eg.][]{Drago07,Cho94} do not find detonation, only deflagration. However, in the most recent work on this matter \\citep{niebergal10b}, we have analyzed the issue numerically including both reaction-diffusion {\\it and} hydrodynamics consistently, albeit within a 1D approximation. We also included for the first time, the effects of neutrino cooling, finding numerical evidence for laminar speeds approaching 0.04$c$ (much higher than previous works; $c$ is the speed of light) and a possible wrinkling instability in higher dimensions. As argued recently by \\citet{Hor10}, wrinkles cause turbulence in the conversion front that serve as a platform for detonation. If this is borne out by more sophisticated simulations, we will have found potentially a new engine for GRBs that is capable of explaining several features of the GRB light curve, as we show in this work. This exciting possibility, combined with the problems faced by more conventional scenarios described in the previous paragraph, should make our model worth exploring as a possible option.  We note that \\citet{niebergal10b} is already an important step forward in resolving some of the issues with uncertainties in the dynamics of the Quark-Nova. The equation of state of quark matter, and hence the putative quark deconfinement density is also unconstrained, but recently \\citet{Ozel09} have argued that accurate measurement of mass and radius for a minimum of three neutron stars will provide strong constraints up to several times nuclear saturation density. While we must await future observations of neutron star parameters for the resolution of some of these issues on the underlying Quark-Nova dynamics to be more confident in our model assumptions, we nevertheless adopt a {\\it working hypothesis} in this paper that the detonative regime is indeed the likely result of a quark-hadron phase transition inside a neutron star. From this point on, we will follow standard arguments to demonstrate that a Quark-Nova occurring in an LMXB following  the end of  the first  accretion phase leads  to features  (prompt emission,  extended emission  and flaring) reminiscent of those  discussed above in the context  of short GRBs as observed by  {\\it Swift}. The same engine at play inside a collapsar can  naturally explain the  many similarities  between short  and long duration GRBs,  thus exhibiting universality between the two classes. In the end, we obtain many observed features of short GRBs beginning with our single hypothesis of the detonative transition, implying that shGRBs could well provide the astrophysical evidence for the Quark-nova, and that this mechanism deserves detailed study.    This paper is organized as follows: In \\S \\ref{sec:qn} we describe the LMXB  parameters  and  give  a   brief  account  of  the  Quark-Nova (henceforth QN)  and  the relativistic  ejecta  it produces.   In \\S \\ref{sec:circdisk},  we describe  characteristics of the circumbinary disk (CD). The interaction  between  the  ejecta from  the  QN and  the CD is studied in \\S  \\ref{sec:prompt} where we describe the energetics and spectral features of the prompt  emission.  Using our model, we perform fits to observed GRBs that bring out the origin of  the empirical Band function. We also derive the duration of  the prompt emission in our  model  and predict  an optical  counterpart. Extended emission  arising from ablating the white  dwarf is addressed in  \\S  \\ref{sec:extended}.   Late  X-ray  activity  ascribed  to  the spin-down  evolution of  the  quark star  and  accretion from  trapped debris is explained in \\S  \\ref{sec:xray} along with  afterglows from the shocked CD. \\S \\ref{sec:application} contains our main conclusions. ",
        "conclusions": "\\label{sec:application}   In this paper we presented the idea of a Quark-Nova occurring in an LMXB (NS-WD binary)  during or near the end of the first accretion phase.   Our basic assumption of a detonative phase transition occurring in the massive neutron star in an LMXB leads to many interesting consequences. We found that several features of the short-hard GRBs can arise naturally in our model, with the binary separation playing an important role. Prompt emission occurs due to the interaction between the Quark-Nova ejecta and the circumbinary disk. Adopting a size and speed distribution for the chunks can explain the range of energetic photons observed (keV-TeV) as well as the duration  of the prompt component. Most importantly, a statistical approach to the ejecta break-up and collision with the  circumbinary disk could in principle lead to a light curve closely resembling the empirical Band function.  Extended emission in our model comes from ablation of the WD, either thermally or via nuclear ignition. Interestingly, the spectral features depend principally on the binary separation, and are expected to display many similarities to type I SNe. Thus, our model offers a plausible unification all the Type I Supernovae. Late X-ray activity is explained in terms of Nickel decay and accretion of gravitationally trapped material onto the Quark star.  It may or may  not become a black hole, and X-ray  activity can result in either case; in  the former it will consist of  X-ray spikes while in the  latter, it will be  the quiescent spin-down  luminosity of the QS. Afterglows  are explained in  terms of synchrotron  radiation from the  rapidly  expanding  shocked   circumbinary disk  material.  In  all  cases,  the energetics provide  the correct order of magnitude.   The rarity of short GRBs calls for rare situations/events. Our model suggests Quark-Novae in LMXBs born with heavy Neutron Stars are the likely engines of the burst. The universality of our model can be confirmed if it turns out that indeed heavy neutron stars ($>1.6M_{\\odot}$) in LMXBs are  associated with or favor the formation of  circumbinary disks. The characteristics of the disk (e.g. size, density etc ...) are expected to vary slightly from one system to another  which could explain differences in burst duration, apparent energy release and spectra.  If our model is a correct representation for short-GRBs engines, most of these engines  would reside in globular clusters (GCs) where many LMXBs are seemingly found  \\citep{bogdabov06,camilo05}. It also implies that  short GRBs should be associated with early-type galaxies (no spiral arms) known to form their LMXBs in GCs.  There is some evidence that short GRBs reside in the outskirts of early-type galaxies \\citep{fox07} where many GCs are located. Finally,  the  WD ablation in our model  provides  a potential  unifying framework for the central engines of type I Supernovae. In particular,  systems  observed at high latitudes (i.e. $> H/R$) when the QN goes off, could find some interesting applications in the context of  unusual type Ia SNe (e.g. SN2005E; \\cite{perets10}) and/or subluminous type Ia SNe in general (\\cite{gonzalez10}).\u00d3     \\vskip 0.5cm"
    },
    "1004/1004.2357_arXiv.txt": {
        "abstract": "The muon and anti-muon neutrino energy spectrum is determined from 2000-2003 AMANDA telescope data using regularised unfolding. This is the first measurement of atmospheric neutrinos in the energy range 2 - 200 TeV. The result is compared to different atmospheric neutrino models and it is compatible with the atmospheric neutrinos from pion and kaon decays. No significant contribution from charm hadron decays or extraterrestrial neutrinos is detected. The capabilities to improve the measurement of the neutrino spectrum with the successor experiment IceCube are discussed. ",
        "introduction": "\\label{intro} At energies above $ 0.1$~TeV, about one cosmic ray particle per square meter per second reaches Earth. At the highest observed energies, particles reach more than $10^{20}$~eV, which is far above what can be achieved in man-made accelerators.  The origin of these charged cosmic rays is still being discussed, as their direction is scrambled by extragalactic and galactic magnetic fields. One option for identifying the origin of cosmic rays is the observation of secondary particles produced in cosmic ray interactions in the astrophysical plasmas themselves: if a proton $p$ interacts with ambient matter or photon fields $\\gamma$, pionic secondaries are produced via the processes $p+p\\rightarrow \\pi+X$ and $p+\\gamma\\rightarrow \\Delta^{+}\\rightarrow n+\\pi^{+}/p+\\pi^{0}$, respectively \\citep{ppb2008}. The charged pions subsequently decay into neutrinos, $\\pi^{\\pm}\\rightarrow \\mu^{\\pm}+\\nu_{\\mu}\\rightarrow e^{\\pm}+\\nu_{e}+\\nu_{\\mu}+\\nu_{\\mu}$, where we do not distinguish between neutrinos and anti-neutrinos. The resulting neutrino flux usually follows the spectral behaviour of the protons, which is predicted to be close to $dN/dE \\propto E^{-2}$ according to Fermi acceleration (\\cite{fermi1949,fermi1954}). The conventional atmospheric neutrino spectrum due to pion and kaon decay, on the other hand, shows a spectral behaviour of approximately $dN/dE\\propto E^{-3.7}$ \\citep[]{Honda,Volkova,gaisser2001,barr2004,honda_04}. An additional component of the atmospheric neutrino flux comes from the decays of hadrons containing charm and bottom quarks.  This flux, known as the prompt component is expected to have a spectrum close to $dN/dE\\propto E^{-2.7}$ \\citep[e.g.]{naumov_RQPM_89,Costa,naumov_RQPM_01,martin_GBW,honda_04}. The prompt atmospheric neutrino flux is lower than the conventional flux but could start to dominate the total spectrum at energies above about $100$~TeV. So far, only the conventional neutrino flux is observed \\citep{jess_diffuse}. Measurements at high energies, i.e.\\ above $10-100$~TeV, provide an opportunity to reveal an extraterrestrial or a charm component. At these energies, the Antarctic Muon And Neutrino Detector Array (AMANDA) and its successor IceCube are able to make measurements to look for deviations from the conventional atmospheric neutrino flux.  The AMANDA-II detector was designed for the detection of neutrinos above 100~GeV. It is composed of 677 Optical Modules (OMs), each containing a 8-inch, 14-dynode photomultiplier tube (PMT) and a voltage divider for the high voltage. The PMTs are optically coupled to the pressure glass sphere with a silicon gel and can be operated at a high gain of about $1 \\cdot 10^9$. The optical modules are attached to 19 vertical strings, instrumenting a cylindrical volume of $0.016$~km$^{3}$ (with a radius of 100 m and a height of 500 m), see e.g.\\ \\cite{ty_neutrino2008}. Secondary muons in the ice are produced via the process $\\nu_{\\mu}+ N\\rightarrow\\mu+X$ \\footnote{Here and throughout the paper, we use the same notation for particles and antiparticles.}. The muons produce Cherenkov radiation if they travel faster than the speed of light in ice (i.e.\\ if the muons travel faster than $v>0.8\\cdot c_0$, with $c_0$ as the speed of light in vacuum).  Additional Cherenkov radiation comes from the particles produced in  muon interactions, such as bremsstrahlung, direct pair production and photonuclear interactions, all dominating at muon energies above 1 TeV. At higher energies, the sum of the energy loss due to stochastic processes (i.e.\\ bremsstrahlung, pair production and nuclear interaction) is dominant and increases linearly with the energy. The amount of light detected with the optical modules  rises with the muon energy and therefore also with the energy of the parent neutrino. Thus the detected light amount can be used to determine the primary neutrino energy spectrum. Neutrino-induced muons can be distinguished from atmospheric muons by selecting events that traverse the Earth and arrive at the detector from below the horizon. Atmospheric muons cannot reach the detector from those directions since they are absorbed on their way through the Earth. In this energy range ($E< 200$~TeV), neutrino absorption in the Earth is not significant. Neutrinos can traverse the matter without loss and some neutrinos interact close to the detector, so that the products of these neutrino interactions can be observed.  AMANDA data from the years 2000 to 2003 are analyzed to determine the energy spectrum of neutrinos, presenting for the first time the atmospheric neutrino spectrum in the energy range $2-200$~TeV. In section \\ref{nus:sec}, predictions for atmospheric neutrinos are reviewed. In section \\ref{deconvolutionexplained}, a conceptual overview of the issues involved in deconvolving a spectrum from observed data are discussed. In section \\ref{analysis:sec}, more details of  the data reduction, simulation and  analysis method for the deconvolution of the neutrino spectrum are explained. Section \\ref{nn:sec} then describes a neural network used for the construction of an optimal energy-correlated variable, while section \\ref{unfolding:sec} shows how the atmospheric spectrum is determined by regularised unfolding, and discusses the sources of statistical and systematic uncertainties that enter the calculation. Section \\ref{results:sec} summarises the results while section \\ref{discussion:sec} discusses them in the context of other experimental results and flux",
        "conclusions": "} The unfolded muon and anti-muon neutrino energy spectrum is presented for the energy range $2$~TeV and $200$~TeV, constituting the first measurement at such high energies. The spectrum is compatible with predictions of the conventional and prompt atmospheric neutrino spectra. The AMANDA detector was switched off in May 2009 but its more than 60 times larger successor IceCube is currently being built at the same South Pole location. As of February 2010, 79 strings have been deployed and completion is planned within a year, completing an instrumented volume of $1$~km$^{3}$.   \\begin{figure}[h!] \\centering{ \\epsfig{file=fig11.eps,width=\\linewidth} \\caption{Measured muon and anti-muon neutrino spectrum and predictions of extraterrestrial neutrino fluxes. Neutrinos are expected from e.g.\\ Active Galactic Nuclei (e.g.\\ \\cite{stecker96,stecker_mod}, (1)), Gamma Ray Bursts (e.g.\\ \\cite{wb97,wb99}, (2)) as well as from the interactions of ultra high-energy cosmic rays with the cosmic microwave background (e.g.\\ \\cite{yuksel_kistler2007}, (3)). The expected sensitivity of IceCube to an $E_{\\nu}^{-2}$ neutrino spectrum is in the range of the hatched area \\citep{hoshina2008,francis_sept2008}.\\label{extragalactic:fig}}} \\end{figure}  Figure \\ref{extragalactic:fig} shows the results of this analysis together with predictions for extraterrestrial neutrino fluxes. Typical neutrino fluxes from e.g.\\ Active Galactic Nuclei or Gamma Ray Bursts are expected to follow a spectrum close to $E_{\\nu}^{-2}$, which is much harder than both the conventional ($\\sim E_{\\nu}^{-3.7}$) and the prompt ($\\sim E_{\\nu}^{-2.7}$) neutrino flux \\citep[e.g.]{halzen_hooper2002,julias_review}. This implies a flattening of the spectrum towards high energies which is much more distinct than for prompt neutrinos. IceCube has the potential to observe this flattening of the spectrum, as its main sensitivity lies in the range $10^{5}-10^{8}$~GeV \\citep{hoshina2008,francis_sept2008} and will be able to measure the high-energy neutrino spectrum with higher accuracy and towards higher energies than AMANDA within the first few years of operation. \\clearpage \\subsection*"
    },
    "1004/1004.0955_arXiv.txt": {
        "abstract": "We present 2323 High-Amplitude \\dsctsing~(HADS) candidates discovered in the Large Magellanic Cloud (LMC) by the SuperMACHO survey \\citep{Rest05}.  Frequency analyses of these candidates reveal that several are multimode pulsators, including 119 whose largest amplitude of pulsation is in the fundamental (F) mode and 19 whose largest amplitude of pulsation is in the first overtone (FO) mode.  Using Fourier decomposition of the HADS light curves, we find that the period-luminosity (PL) relation defined by the FO pulsators does not show a clear separation from the PL-relation defined by the F pulsators.  This differs from other instability strip pulsators such as type c RR~Lyrae.  We also present evidence for a larger amplitude, subluminous population of HADS similar to that observed in Fornax \\citep{Poretti08}.% ",
        "introduction": "\\dsctsing~variables populate the region of the Hertzprung-Russell diagram where the instability strip meets the main sequence.  The high-amplitude variables are generally believed to be pulsating primarily in radial modes, whereas \\dsct~with smaller amplitudes are believed to have many non-radial modes of pulsation.  \\citet{Breger} provides a thorough review of the theoretical models describing \\dsctsing~pulsation.  Recent space-based observations from the CoRoT telescope have begun to reveal the rich complexity of \\dsctsing~pulsation modes \\citep{Poretti09}.  Despite their multimode nature, HADS have been shown to define a period-luminosity relationship, allowing for their use as standard candles \\citep[and references therein]{mcnamaraLMC,Poretti08}.  Until recently, observations of large sets of HADS have been limited due to these stars' intrinsic faintness and short periods.  The majority of known HADS have been found within the Milky Way (e.g. \\citealt{machoblg}, \\citealt{MACHOblgdsct}, and \\citealt{Pigulski06}).  More recent work has revealed 90 \\dsct~(or SX~Phoenicis stars, Population II \\dsct) in Fornax \\citep[hereafter P08]{Poretti08}.  Using a subset of the MACHO project data, \\citet{machodsct} finds 101 \\dsct~in the LMC.  \\citet{mcnamaraLMC} report 24 \\dsct~in the LMC using the OGLE-II data set.  Their work also provides a summary of HADS detected by ground-based surveys.  In this paper we present analyses of a large set of \\dsct~discovered by the SuperMACHO survey of the LMC.  In Section~\\ref{sec:SM} we provide an overview of the survey and data reduction.  In Section~\\ref{sec:data} we discuss our data and HADS selection criteria, and we present our candidates.  We discuss our findings in Section~\\ref{sec:disc}.  We perform a frequency spectrum analysis of the HADS candidates to identify multimode pulsators and present evidence for a large set of FO pulsators.  We examine a subset of our candidates having larger amplitudes.  We find evidence for an excess population of faint sources, and discuss whether it is the subluminous population observed in Fornax (P08). % ",
        "conclusions": "We have discovered 2323 candidate HADS in the Large Magellanic Cloud (LMC) using the SuperMACHO set of variables.  Using the SigSpec software, we performed frequency analyses of these candidates which reveal several multimode pulsators, including 119 that are pulsating fundamental modes and 19 in first overtone modes.  We find evidence for a large set of FO pulsators within this data set. Notably, the PL-relation defined by the FO pulsators does not show a clear separation from the PL-relation defined by the F pulsators. This is not necessarily surprising, as \\dsct~occupy a region of the instability strip that spans a broad region of temperatures, and hence intrinsic color.  Though we are unable to do so with our single-epoch multi-band photometry, future surveys that obtain multi-epoch color information may be able to better define this region of the CMD.  Such data would also allow for the determination of period-luminosity-color (PLC) corrected magnitudes that would reveal whether a clear separation between F and FO pulsators does exist.  We also find that the majority of our HADS with amplitudes greater than 0.4~mag lie below the ridgeline of the PL-diagram.  By examining only these candidates, we find an excess of subluminous sources similar to that observed in Fornax (P08).  Because the separation between this population and the main PL-relation is also similar to that found in Fornax, despite its having different composition and formation history, we postulate that they form at the same time as those on the main PL-relation rather than constituting a second, older and metal-poor population.  Spectroscopic observations to measure the metallicities of these stars may help to illuminate this discussion."
    },
    "1004/1004.2161_arXiv.txt": {
        "abstract": "{} {We have analysed optical spectra of BL Lacertae, the prototype of its blazar subclass, to verify the broad H$\\alpha$ emission line detected more than a decade ago and its possible flux variation. We used the spectroscopic information to investigate the question of the BL Lacertae parent population.} {Low- and high-resolution optical spectra of BL Lacertae were acquired with the DOLORES spectrograph at the 3.58 meter Telescopio Nazionale Galileo (TNG) during four nights in 2007--2008, when the source was in a relatively faint state.  In three cases we were able to fit the complex H$\\alpha$ spectral range with multiple line components and to measure both the broad H$\\alpha$ and several narrow emission line fluxes.} {A critical comparison with previous results suggests that the broad H$\\alpha$ flux has increased by about $50\\%$ in ten years.  This might be due to an addition of gas in the broad line region (BLR), or to a strengthening of the disc luminosity, but such flux changes are not unusual in Broad Lined active nuclei. We estimated the BL Lacertae black hole mass by means of its relation with the bulge luminosity, finding 4--$6 \\times 10^8 \\, M_{\\sun}$.  The virial mass estimated from the spectroscopic data gives instead a value 20--30 times lower. An analysis of the disc and BLR properties in different AGNs suggests that this discrepancy is due to an underluminosity of the BL Lacertae BLR. Finally, we addressed the problem of the BL Lacertae parent population, comparing its isotropic quantities with those of other AGN classes. From the point of view of the narrow emission line spectrum, the source is located close to low-excitation radio galaxies. When one also considers its diffuse radio power, an association with FR~I radio galaxies is severely questioned due to the lower radio luminosity (at a given line luminosity) of BL Lacertae. The narrow line and radio luminosities of BL Lacertae instead match those of a sample of miniature radio galaxies, which however do not show a BLR. Yet, if existing, ``misaligned BL Lacertae\" objects should have entered that sample. We also rule out the possibility that they were excluded because of a QSO optical appearance.} {The observational constraints suggest that BL Lacertae is caught in a short term transient stage, which does not leave a detectable evolutionary ``trace'' in the AGN population. We present a scenario that can account for the observed properties.} ",
        "introduction": "According to the commonly accepted scenario, the central engine of active galactic nuclei (AGNs) is a supermassive black hole (SMBH) fed by infall of matter from an accretion disc. Inner fast-moving clouds produce broad emission lines, which may be obscured by absorbing material, while outer clouds are responsible for narrow emission lines.  About one fifth of AGNs is radio-loud \\citep{kellermann94}, showing plasma jets sometimes extending on Mpc scales. Among them, BL Lac objects and flat spectrum radio quasars (FSRQs) form the blazar class, characterized by variable emission from the radio to the $\\gamma$-ray band, with flux variations on time scales from hours to years, high radio and optical polarization, core-dominated radio morphology, flat radio spectra, and apparent superluminal motion of radio jet components. Their properties are explained in terms of plasma relativistic motion in a jet pointing at a small angle with the line of sight, with consequent beaming of the observed radiation \\citep{bla78}. Hence, the continuum radiation of blazars is dominated by the relativistically beamed non-thermal radiation from the jet.  In FSRQs, thermal emission from the disc may be observable in the optical--ultraviolet band when the source is not in a flaring state. Disc signatures were detected e.g.\\ in 3C 273 \\citep{smi93,von97,gra04,tur06}, 3C 279 \\citep{pia99}, PKS 1510-089 \\citep{kat08,dam09}, 3C 345 \\citep{bre86}, and 3C 454.3 \\citep{rai07b,rai08c}. Moreover, strong broad and narrow emission lines are usually present in their spectra.  As for BL Lac objects, according to the original definition, they may show at most weak emission lines, with equivalent widths not exceeding 5 \\AA\\ in the rest frame \\citep{sti91}.  This seems to be due not so much to low line fluxes, but rather to a high continuum flux \\citep{sca97}.  Indeed, strong emission lines, in particular broad ones, have occasionally been detected in the spectra of BL Lac objects in faint states.  These include BL Lacertae, the prototype of the blazar subclass named after it \\citep{ver95,cor96,cor00}, and the distant source AO 0235+164 \\citep{coh87,nil96,rai07a}.  The unified scheme for radio-loud AGNs predicts that BL Lac objects and FSRQs are the beamed counterparts of Fanaroff-Riley type I (FR I) and Fanaroff-Riley type II (FR II) radio galaxies, respectively, even if these correspondences were questioned by various observing evidences \\citep[e.g.][]{tad08}. In particular, many BL Lac objects show high, FR II-like extended radio powers and morphologies (see e.g. \\citealt{lan08} \\, and references therein). The distinction between blazars and radio galaxies, as well as that between FSRQs and BL Lac objects, has been discussed by several authors, and different criteria have been proposed, involving the value of the Ca H\\&K break \\citep{mar96}, or the strength of the oxygen-narrow emission-lines \\citep{lan04}.  A powerful way to classify AGNs is their position in diagnostic diagrams comparing selected emission line ratios \\citep{hec80,bal81}. In particular, ratios of lines close in wavelength, like [\\ion{O}{III}]/H$\\beta$, [\\ion{N}{II}]$\\lambda 6583$/H$\\alpha$, [\\ion{S}{II}]$\\lambda \\lambda 6716, 6731$/H$\\alpha$, and [\\ion{O}{I}]/H$\\alpha$ are expected to be the most reliable ones \\citep{vei87}.  The application of diagnostic diagrams to radio-loud galaxies by \\citet{lai94} confirmed former suggestions that FR II sources can be divided between high-excitation galaxies (HEG) and low-excitation galaxies (LEG).  In particular, \\citet{but10} verified this dichotomy when analysing the radio sources belonging to the well known 3CR catalogue. They found prominent broad lines in a sub-sample of HEG, but not in LEG. Moreover, they saw that HEG are associated with very powerful FR II only, while LEG are spread on a wide range of radio powers, and can be of both FR II and FR I type. Actually, the situation is even more complex, as the existence of miniature radio galaxies, characterized by extremely low radio power, relatively luminous narrow emission lines, and no BLR, demonstrates \\citep{bal09}.  An analysis of the spectroscopic properties of blazars to understand their relationship with the radio galaxies (and other AGN classes) is not an easy task, because the dramatic variability of the non-thermal continuum flux strongly affects the appearance of lines, especially in BL Lac objects. However, this analysis can help clarify whether blazars differ from radio galaxies only for their orientation with respect to the line of sight, or if they are intrinsically different sources.  We present spectroscopic observations of BL Lacertae carried out in 2007--2008 with the 3.58 m Telescopio Nazionale Galileo (TNG) on the Canary Islands, to address the problem of its parent population.  In the same period BL Lacertae was the target of a multiwavelength campaign by the Whole Earth Blazar Telescope\\footnote{{\\tt http://www.oato.inaf/it/blazars/webt/}} (WEBT), also involving three pointings by the XMM-Newton satellite. The results of the WEBT campaign have been presented by \\citet{rai09}. The source was observed in a relatively faint state at all wavelengths, and a UV excess was clearly visible in the source spectral energy distribution (SED), which was interpreted as the signature of thermal radiation from the accretion disc. ",
        "conclusions": "More than a decade ago \\citet{ver95} and subsequently \\citet{cor96,cor00} detected a broad H$\\alpha$ emission line in the spectra of BL Lacertae. This luminous line should have been detected in previous spectra, suggesting that its flux must have increased by at least a factor 5 since 1989. The luminous H$\\alpha$ line suggested a Seyfert-like nucleus in BL Lacertae, complicating the already difficult task of understanding what AGN population this object (and BL Lac objects in general) belongs to, considering that FR~I generally do not show broad lines.  To investigate this matter we acquired low- and high-resolution spectra of BL Lacertae with the TNG during four nights in 2007--2008, when the source optical brightness was $R \\sim 14$--14.5.  Our spectra confirm the presence of a luminous H$\\alpha$ broad line of $\\sim 4 \\times 10^{41} \\, \\rm erg \\, s^{-1}$ and $\\rm FWHM \\sim 4600 \\rm \\, km \\, s^{-1}$, as well as several narrow emission lines. Through a critical comparison of our data with those by \\citet{cor00}, we concluded that the BLR luminosity has increased by $\\sim 50\\%$ in about ten years. This level of variability is not unusual for Broad Lined AGN and it does not necessarily implies an evolutionary trend.  Then we examined the nuclear properties of BL Lacertae.  The relationship between the SMBH mass and bulge luminosity in AGNs allowed us to derive a mass of 4--$6 \\times 10^8 \\, M_{\\sun}$.  Using the spectroscopic information to calculate the virial mass, we instead obtained a value about 20--30 times lower. To understand the reason of this discrepancy we analysed the disc and BLR properties of other AGNs, and concluded that the BLR of BL Lacertae is underluminous by a factor 70--300.  Finally, we analysed the physical quantities that do not depend on orientation and beaming, and thus should also characterize the parent population of BL Lacertae. We defined diagnostic indices with the most reliable narrow emission lines, and found that their values provide a tentative identification of BL Lacertae as a LEG. Broad lines are instead observed only in HEG, but the diffuse radio luminosity of BL Lacertae is at least 100 times lower than in these powerful radio sources. On the other hand miniature radio galaxies are LEG, share both the narrow line and radio power properties of BL Lacertae, but they do not show a BLR. Taking into account how the miniature radio galaxy sample was selected, we expect that it should include ``misaligned BL Lacertae\" objects, unless they were excluded on the base of a QSO appearance. An analysis of the galaxy morphology, spectral features, and radio power of the QSO sources, initially discarded from the sample of miniature radio-galaxies, revealed that no object meets the requirements to represent the BL Lacertae parent population. Yet, for typical values of the Lorentz factor, we would expect 10-10$^3$ ``misaligned BL Lacertae\".  This leaves us with the only possibility that the observed properties of BL Lacertae are the result of a transient short lasting phase. We can envisage the following scenario, somewhat similar to that already suggested by \\citet{cor96}. BL Lacertae in its initial state has properties similar to the sources of the SDSS/NVSS sample. Indeed, these are massive early-type galaxies and a large number of them have narrow lines and radio luminosities similar to that of BL Lacertae. From the point of view of their optical spectra they are LEG and lack broad lines.  Subsequently (possibly $\\sim 20$ years ago), its BLR underwent an increase of luminosity due to an increased amount of cold gas in the nuclear regions and/or to a higher level of ionizing continuum. These two effects may even be related and caused by a fresh input of accreting gas. The BLR structure might not have yet reached a stable configuration, accounting for its different properties when compared to other AGN.  Also the NLR luminosity will grow with time and will also eventually change its state of ionization, but on a much larger timescale with respect to the BLR.  Based on the analysis of a single object it is clearly impossible to set a timescale for the duration of the putative bright phase. Furthermore, BL Lacertae was probably discovered since this object has been subject to repeated spectroscopic observations. However, our failure to find objects in the local Universe that might constitute its parent population suggests that the timescale associated with the period of high accretion must be orders of magnitude shorter than the lifetime of radio-loud AGN.  An alternative possibility is that the birth of the BLR marks the transition from a low-power radio galaxy to a high-power source. This would require a rapid increase in the luminosity of the large-scale radio structures to reach the level observed in e.g.\\ the HEG of the 3CR sample, within a sufficiently short time so as not to produce a substantial population of transient sources.  Instead, the available data rule out that BL Lacertae became an AGN only very recently, i.e.\\ that we are witnessing its birth, because its radio emission extends $\\sim 10$ kpc away from the core. This implies that this source is active since at least $\\sim 3 \\times 10^5$ years, assuming an expansion speed of 0.1 c.  We conclude that the parent population of BL Lacertae can be found among the large population of miniature radio-loud AGN forming the SDSS/NVSS sample, but this also requires that this object is experiencing a short transient phase. A continuation of the spectroscopic monitoring of this peculiar source caught in a crucial phase of its evolution can help us tremendously in our study of the physics and evolution of these systems."
    },
    "1004/1004.1162_arXiv.txt": {
        "abstract": "A generic expectation for gas accreted by high mass haloes is that it is shock heated to the virial temperature of the halo. In low mass haloes, or at high redshift, however, the gas cooling rate is sufficiently rapid that an accretion shock is unlikely to form. Instead, gas can accrete directly into the centre of the halo in a `cold mode' of accretion. Although semi-analytic models have always made a clear distinction between hydrostatic and rapid cooling they have not made a distinction between whether or not an accretion shock forms. Starting from the well-established {\\sc Galform} code, we investigate the effect of explicitly accounting for cold mode accretion using the shock stability model of Birnboim \\& Dekel. When we modify the code so that there is no effective feedback from galaxy formation, we find that cold mode accretion is the dominant channel for feeding gas into the galaxies at high redshifts. However, this does not translate into a significant difference in the star formation history of the universe compared to the previous code. When effective feedback is included in the model, we find that the the cold mode is much less apparent because of the presence of gas ejected from the galaxy. Thus the inclusion of the additional cold mode physics makes little difference to basic results from earlier semi-analytic models which used a simpler treatment of gas accretion. For more sophisticated predictions of its consequences, we require a better understanding of how the cold mode delivers angular momentum to galaxies and how it interacts with outflows. ",
        "introduction": "\\label{sec:Intro}  The process of galaxy formation must begin with gas, initially distributed rather smoothly, collapsing to high densities. Furthermore, to sustain ongoing star formation in galaxies requires a continued input of gas over their lifetimes. As such, the question of how galaxies get their gas has received a great deal of attention over the history of galaxy formation studies. While the initial stages of this collapse are purely gravitational (the gas being dragged along by the gravitationally dominant dark matter), after halo formation hydrodynamic forces come into play and further collapse is mitigated by the interplay of gravity, hydrodynamics and cooling processes.  An accretion shock is a generic expectation whenever the gas accretes supersonically as it will do if the halo virial temperature exceeds the temperature of the accreting gas \\pcite{binney_physics_1977}. Models of virialization shocks have been presented by several authors \\pcite{bertschinger_self-similar_1985,tozzi_evolution_2001,voit_origin_2003,book_role_2010} with the general conclusion that the shock occurs at a radius comparable to (or perhaps slightly larger than) the virial radius. Much debate has occurred over the existence of such shocks --- their existence has often been assumed in analytic models of galaxy formation\\footnote{This has always been understood to be an approximation, valid only in specific mass regimes: \\protect\\cite{white_galaxy_1991}, in discussing the fate of the gaseous component of a halo, note that ``cooling rates may be short enough for the multiphase structure to survive the shocks, and it is then unclear how the dynamics of the gas component should be modeled.''} since \\cite{rees_cooling_1977}.  Recent work has examined these issues in greater detail. Motivated by hydrodynamical simulations (\\citealt{fardal_cooling_2001}; see also \\citealt{kerevs_do_2005,ocvirk_bimodal_2008,kerevs_galaxies_2009}), which show that a significant fraction of gas in galaxies has never been shock heated, \\cite{birnboim_virial_2003} developed an analytic treatment of virialization shock stability. The virial shock relies on the presence of a stable atmosphere of post-shock gas to support itself. If cooling times in the post-shock gas are sufficiently short, this atmosphere cools and collapses and can no longer support the shock. For cosmological halos this implies that shocks can only form in halos with mass greater than $10^{11}M_\\odot$ for primordial gas (or around $10^{12}M_\\odot$ for gas of Solar metallicity). These values are found to depend only weakly on redshift and are in good agreement with the results of hydrodynamical simulations.  As a result, in low mass halos gas tends to accrete ``cold''---never being shock heated to the virial temperature and instead raining into the halo as cold clumps along filaments\\footnote{While this picture seems reasonable on theoretical grounds, it as yet has little direct observational support \\protect\\citep{steidel_structure_2010}.}. If this gas is to make it into the galaxy it must nevertheless lose its energy, either by drag processes in the halo or through a shock close to the galaxy which turns its kinetic energy into thermal energy which is immediately radiated away. Halos which do support shocks are expected to contain a quasi-hydrostatic atmosphere of hot gas. The structure of this atmosphere is determined by the entropy that the gas gains at the accretion shock and that may be later modified by radiative cooling \\pcite{voit_origin_2003,mccarthy_modelling_2007}. The transition from cold to hot mode accretion is not sharp---halos able to support a shock still experience some cold mode accretion (the relative contributions of the cold mode decreasing with halo mass).  The consequences of cold vs. hot accretion for the properties of the galaxy forming from such an accretion flow have yet to be fully worked out.  As \\cite{croton_many_2006} have stressed, the absence of an explicit cold mode may not be important since the cold gas accretion rate in small haloes is limited by the growth of the halo rather than by the system's cooling time. In contrast, \\cite{brooks_role_2009} demonstrate in hydrodynamical simulations that cold mode accretion does allow accreted gas to reach the galaxy more rapidly, by virtue of the fact that it does not have to cool but instead merely has to free-fall to the centre of the halo (starting with a velocity comparable to the virial velocity). This results in earlier star formation than if all gas were assumed to be initially shock heated to the virial temperature of the halo. It is also clear that the situation needs to be carefully reassessed in the presence of effective feedback schemes that prevent excessive star formation, particularly in the high redshift universe.  In this work we implement a treatment of cold-mode accretion into the \\gf\\ semi-analytic model of galaxy formation, following the methodology of \\cite{birnboim_virial_2003}. This will allow us to assess the importance of cold-mode accretion for cosmological populations of galaxies across a range of redshifts, and to suggest additional studies of the cold mode which might improve our understanding of its role in the process of galaxy formation. An important aspect is that we are able to explore how feedback and reheating of cold gas moderate the cold mode. We note that \\cite{cattaneo_modellinggalaxy_2006} have previously explored a simpler implementation of cold mode accretion in a semi-analytic model of galaxy formation. They parameterized cold mode accretion by defining a critical mass scale (which had some dependence on redshift) above which accretion switched from cold mode to hot mode. While motivated by the results of the \\cite{birnboim_virial_2003} calculation, this parameterization did not capture the full generality of that work as we will attempt to do herein.  The remainder of this paper is arranged as follows. In \\S\\ref{sec:Model} we describe our implementation of cold mode accretion in the \\gf\\ semi-analytic model. In \\S\\ref{sec:Results} we present our results and, finally, in \\S\\ref{sec:DiscConc} we discuss their implications. ",
        "conclusions": "\\label{sec:DiscConc}  We have described a simple implementation of cold-mode accretion in a semi-analytic model of galaxy formation, using a previously proposed analytical model with no modifications. Without any adjustment this model provides an excellent match to results from numerical simulations---this could no doubt be improved with some fine-tuning of the model, but our aim here was to demonstrate that the simulation results can be encapsulated by an easy to implement model.  We find that the inclusion of \\SNe-driven outflows has a dramatic effect on the fraction of cold mode gas present in dark matter halos---much more so than switching the cold mode accretion on or off --- suggesting that simulations must account for such outflows before they can make robust predictions for the properties of cold mode gas. The inclusion of cold mode accretion makes little difference to luminosity functions and galaxy sizes at $z=0$, implying that results from earlier semi-analytic models which did not treat cold mode accretion are still valid. Without \\SNe-driven outflows the cold mode results in a significant increase in the mass and luminosity of brighter galaxies which are able to gain some mass through the cold mode even when their hot mode has been effectively shut down by feedback from \\AGN.  \\begin{figure} \\includegraphics[width=80mm]{tdyn.pdf} \\caption{The median dynamical time of galactic disks as a function of halo virial mass at $z=6$. Dashed lines show models with no \\protect\\SNe-driven outflows while solid lines include \\protect\\SNe-driven outflows. Red lines show results for models with no cold-mode accretion, while blue lines are for models including cold-mode accretion.} \\label{fig:tdyn} \\end{figure}  Cold mode accretion is found to be effective at getting gas into the galaxy phase at high redshifts, potentially allowing for high rates of star formation during these epochs. However, we find that the star formation rate at $z\\gsim 3$ is mostly unaffected when we include cold mode accretion due to the fact that the galaxies that form are larger and lower density than they would be if cold mode accretion were neglected. We caution that the treatment of angular momentum delivery to galaxies via the cold mode is therefore of crucial importance to assessing its impact on star formation rates. To date, this has not been studied in detail in hydrodynamical simulations but should be in order to refine our understanding of how the cold mode affects galaxy formation.  In summary, cold mode accretion should be explicitly accounted for in semi-analytic models --- and will require some retuning of parameters to restore good fits to observational data --- but does not seem to qualitatively change our picture of galaxy formation, at least at the coarse-grained level studied here. More detailed results (luminosity function shapes, distribution of sizes and formation times) will require a significantly more detailed implementation of the cold mode, including how it delivers angular momentum to galaxies (e.g. \\citealt{navarro_disk_2009}). Calibration from numerical simulations should improve the accuracy of cold mode results. However, we have seen that the inclusion of cold-mode accretion in to our semi-analytic model does not produce any surprises. The changes in star formation rates and stellar mass fractions, particularly in the low redshift universe are reassuringly small, and much smaller than the differences in galaxy properties created by changes to the feedback effects of galactic winds and AGN."
    },
    "1004/1004.4747_arXiv.txt": {
        "abstract": "{In this review I concentrate on three areas related to structure of disks in spiral galaxies. First I will review the work on structure, kinematics and dynamics of stellar disks. Next I will review the progress in the area of flaring of \\HI\\ layers. These subjects are relevant for the presence of dark matter and lead to the conclusion that disk are in general not `maximal', have lower $M/L$ ratios than previously suspected and are locally stable w.r.t. Toomre's $Q$ criterion for local stability. I will end with a few words on `truncations' in stellar disks.} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.3217_arXiv.txt": {
        "abstract": "Collisionless simulations of the CDM cosmology predict a plethora of dark matter substructures in the halos of Milky Way sized galaxies, yet the number of known luminous satellites galaxies is very much smaller, a discrepancy that has become known as the `missing satellite problem'. The most massive substructures have been shown to be plausibly the hosts of the brightest satellites, but it remains unclear which processes prevent star formation in the many other, purely dark substructures. We use high-resolution hydrodynamic simulations of the formation of Milky Way sized galaxies in order to test how well such self-consistent models of structure formation match the observed properties of the Galaxy's satellite population. For the first time, we include in such calculations feedback from cosmic rays injected into the star forming gas by supernovae as well as the energy input from supermassive black holes growing at the Milky Way's centre and its progenitor systems.  We find that non-thermal particle populations quite strongly suppress the star formation efficiency of the smallest galaxies. In fact, our cosmic ray model is able to reproduce the observed faint-end of the satellite luminosity function, while models that include only the effects of cosmic reionization, or galactic winds, do significantly worse. Our simulated satellite population approximately matches available kinematic data on the satellites and their observed spatial distribution.  We conclude that a proper resolution of the missing satellite problem likely requires the inclusion of non-standard physics for regulating star formation in the smallest halos, and that cosmic reionization alone may not be sufficient. ",
        "introduction": "\\renewcommand{\\thefootnote}{\\fnsymbol{footnote}} \\footnotetext[1]{E-mail: wadepuhl@mpa-garching.mpg.de}  The leading $\\Lambda$CDM cosmology predicts that galaxies form hierarchically in a `bottom up' fashion \\citep[e.g.][]{WhiteRees1978,GalaxyFormation}, where small perturbations in the dark matter density distribution collapse earlier than larger perturbations, and low-mass halos grow by smooth accretion or mergers with other halos, successively building up ever bigger structures. But structures falling into bigger systems during this process are not always disrupted completely. As N-body simulations show, the inner cores of infalling objects often survive the various disruptive effects acting on them, like tidal truncation, tidal shocking or ram-pressure stripping. It is believed that the observed dwarf galaxies orbiting around the Milky Way (MW) are examples of such surviving remnants.  Based on the first generation of very high resolution collisionless CDM simulations, \\citet{FirstMSP} and \\citet{Moore1999} pointed out a very striking apparent discrepancy between theoretical predictions for such satellite systems and actual observations. Given the very large number of predicted dark matter substructures, there appears to be a dearth of luminous satellites in the Milky Way. In fact, the cumulative number of observed satellite galaxies and of predicted substructures above a given circular velocity value differed by a factor of $\\sim 10$.  This has become known as the `missing satellite problem'.  The initial analysis of \\citet{Moore1999} and \\citet{FirstMSP} may have overstated the magnitude of the discrepancy, both because of uncertainties in assigning correct circular velocity values to the observed satellites \\citep{Stoehr2002} and because a number of additional faint satellites have been discovered meanwhile in the MW \\citep[see for example][]{SatData9, SatData11, SatData13, SatKinematics, SatData14, SatData15, SatData16, SatData17, SatData18, SatData19}. However, there is a consensus that the many low-mass satellites predicted by the N-body simulations need to be strongly suppressed in luminosity, otherwise a significant discrepancy with the observed satellite luminosity functions results, that, if confirmed, may in principle even be used to rule out cold dark matter.   With increasing numerical resolution, the missing satellite problem has become more acute. Modern cosmological dark matter simulations of Milky Way sized halos \\citep{Diemand2008,Aquarius,Stadel2009} resolve up to $\\sim 300,000$ dark matter substructures, while the number of observed satellite galaxies still comprises just a few dozens. We note that the modern dark matter only simulations are even able to resolve substructures inside substructures, and interestingly, there is also some observational evidence for a satellite possibly orbiting around another satellite \\citep{SatData16}.   \\citet{FirstMSP} did not only raise the missing satellite problem, they were also among the first to suggest a potential solution to this issue. In particular, they proposed that star formation inside low mass halos could be suppressed because of photo evaporation of gas due to a strong intergalactic ionizing UV background. This would keep most of the orbiting satellites dark and render them visually unobservable. Indeed, the simple filtering mass model of \\citet{Gnedin2000} for the impact of a UV background on the cooling efficiency of small halos predicts a quite sizable effect, with a nearly complete suppression of cooling in all halos with circular velocity below $50\\,{\\rm km\\, s^{-1}}$. However, recent full hydrodynamical simulations have not confirmed this \\citep{Hoeft2006,Okamoto2008}. They find a considerably weaker effect, where only halos with circular velocities less than $\\sim 25\\,{\\rm km\\, s^{-1}}$ are affected. This also casts some doubt about the faint-end results of numerous semi-analytic models for the satellite population \\citep[e.g.][]{Benson2002,Kravtsov2004}, which typically employed the filtering mass formalism and hence assumed an overly strong effect of the UV background. We will reexamine this question in this work based on our cosmological hydrodynamic simulations of Milky Way formation, which include a treatment of cosmic reionization.  Another possible solution for the satellite problem was proposed by \\citet{Redef_MSP} who suggested that not the satellite mass at the present epoch determines whether a satellite would be luminous or not, but rather the maximum mass it had before accretion onto the Milky Way's halo. This is based on the idea that tidal stripping and ram pressure unbinds the gas from an infalling satellite and thus stalls any further star formation. With this assumption, the stellar mass at the time of accretion is essentially retained until the present epoch, and it becomes a question of allowing high-redshift star formation only in satellites above a sufficiently high mass threshold.  A more radical conjecture is that the properties of the dark matter particles may have to be changed. Instead of having negligible velocity dispersion at the time of decoupling, we may instead be dealing with (slightly)  warm dark matter (WDM). This can suppress the abundance of low mass structure considerably \\citep[e.g.][]{Colin2000}, but provided the particle mass is not lower than $\\sim 1\\,{\\rm keV}$ a sufficiently large number of substructures still survives to explain the observed satellite abundance \\citep{MacioFontanot2010}.  In the most recent works on the subject, a number of interesting and encouraging results have been obtained.  Observationally, it has been discovered that the satellites all have approximately the same central mass density (within 300 to 600 pc), independent of their luminosity \\citep{Gilmore2007,Strigari2008}. Explaining this central density threshold has become an important additional challenge for theoretical models. Also, a significant number of new faint satellites have been discovered with the help of the SDSS \\citep{SatData9, SatData11, SatData13, SatKinematics, SatData14, SatData15, SatData16, SatData17, SatData18, SatData19}, improving our knowledge about the full satellite population significantly, but at the same time also raising the question whether we may perhaps still be missing large numbers of satellites at ultra low surface brightnesses.  On the theoretical side, refined treatments of the effects of reionization, often coupled to the results of high-resolution collisionless simulations have been used to model the satellite population.  \\citet{Maccio10} employed a number of different semi-analytic models and low-resolution hydrodynamic simulations to study the satellite luminosity function. Despite just invoking photoheating as primary feedback process, they achieved reasonable agreement for some of their models, leading them to argue that the satellite problem may be solved. Similarly, \\citet{Li2010} invoked a strong impact of reionization in a semi-analytic model similar to those applied to the Millennium Simulation \\citep{Croton2006} to reproduce the luminosity function of galaxies around the Milky Way. \\citet{Busha2010} used simple prescriptions for the impact of inhomogeneous reionization on the satellite population, pointing out that subtle changes in the assumptions about how reionization affects star formation in small galaxies can lead to large changes in the predicted number of satellites.  Recently, \\citet{Strigari2010} examined the kinematics of five well-measured Milky Way satellite galaxies and compared them to dark matter satellites of the high-resolution simulations of the Aquarius Project \\citep{Aquarius}. They showed that these systems are fully consistent with $\\Lambda$CDM expectations and may be hosted in cored dark matter structures with maximum circular velocities in the range 10 to $30\\,{\\rm km\\,s^{-1}}$.  Interestingly, \\citet{Bullock2009} pointed out that the number of real satellite systems may in fact be much larger than commonly believed, with the majority of them being so far undetected because of their low surface brightness. In this scenario, the `common mass scale' inferred for the observed satellites may in fact just arise from a selection bias.   The first high-resolution hydrodynamic simulation able to directly resolve the satellite population has recently been presented by \\citet{Okamoto2009}. They argue that the common mass scale identified in the observations arises from early reionization at redshift around $z\\sim 12$, and that satellites that have not yet grown to a maximum circular velocity of $\\sim 12\\,{\\rm km\\, s^{-1}}$ {\\em by the time of reionization}, will not be able to make any stars later on. Even if they grow above this threshold, \\citet{Okamoto2009} predict them to remain dark.    Despite all of this progress, it is evident that there remain many open questions concerning the population of faint and ultra-faint satellite galaxies orbiting around Milky Way like galaxies. Especially the influence of different baryonic feedback processes on the luminosity function of the simulated satellites has not been investigated in sufficient detail. It is therefore far from clear whether photoheating from a UV background and ordinary supernovae feedback are indeed the correct physical solutions to the missing satellite problem. In fact, the in part contradictory results that have been obtained with analytic recipes to describe the impact of reionization suggest that more accurate methodologies are required to reliably settle the issue.  We have therefore embarked on a research program where we use high-resolution hydrodynamic simulations of the formation of Milky Way-sized halos to shed more light on these questions, in particular by investigating a variety of feedback processes known to be important in galaxy formation. Besides the impact of reionization, these include galactic winds and outflows, energy input by growing supermassive black holes, or the non-thermal support of gas by cosmic rays or magnetic fields. Ultimately we aim to reach similar numerical resolution as has been obtained for recent collisionless simulations, even though this goal may still be several years away.  In this work, we present some of our first results. We use several well resolved hydrodynamical simulations of the formation of a Milky Way sized galaxy to investigate the properties of the predicted population of satellite galaxies, for different choices of the included physics. Besides a default reference model that includes only a treatment of radiative cooling, star formation, and cosmic reionization, we consider also models that add galactic winds, supermassive black hole growth, or cosmic ray injection by supernovae shock waves. By comparing the simulation results with a comprehensive catalogue of the known Milky Way satellites, we seek to determine which of these processes is most important in shaping the satellite population.  This paper is organized as follows. In Section~\\ref{sec:methodology}, we describe the methodological details of our simulations, while the observational knowledge about the satellites is briefly summarized in Section~\\ref{sec:observations}. Sections~\\ref{sec:abundance}, \\ref{sec:history} and \\ref{sec:scaling} present the results for our simulated populations of satellite galaxies, both with respect to individual satellite histories as well as with respect to their population as a whole. Our conclusions are summarized in Section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  In this work, we studied a set of high-resolution hydrodynamical simulations of the formation of a Milky Way sized galaxy, starting from cosmological initial conditions. Such simulations are now able to reach sufficiently high resolution to directly resolve the formation of the small dwarf galaxies that orbit in the halo, thereby allowing studies of the missing satellite problem and of the properties predicted by simulations for the population of satellite galaxies. These galaxies are especially interesting both because the dark matter substructure abundance is a fundamental challenge for the $\\Lambda$CDM cosmology, and because the low star formation efficiencies of the satellites provide crucial information about the physics of feedback.  We have therefore repeated our simulations using different models for feedback physics, with the goal to test the sensitive of the results for the satellites with respect to these physics assumptions.  In the \\textsc{Ref} model, we considered only star formation and SN feedback, together with instantaneous reionization at $z=6$. The three other models included additional processes like AGN feedback (\\textsc{BH}), wind driven galactic outflows (\\textsc{Wind}) and the generation and decay of cosmic rays (\\textsc{CR}). Not unexpectedly, the \\textsc{BH} model showed no significant differences compared to the reference \\textsc{Ref} model, as most of the satellites are simply too small to grow a large supermassive black hole and are rarely affected by strong quasar feedback in neighbouring galaxies. In contrast, the \\textsc{Wind} model showed a significant reduction of the number of high mass satellites, but did not give a significantly different abundance of low mass systems.  The \\textsc{CR} model had exactly the opposite effect as it did not change the high mass satellites but suppressed star formation in low mass satellites. This made the cosmic ray model most successful in matching the faint-end of the observed satellite luminosity function. Our results further suggest that a combination of the \\textsc{Wind} and \\textsc{CR} feedback models should be able to yield a nearly perfect match of the luminosity function.  The total number of satellites observable with an SDSS-like survey covering the whole sky has been estimated to be $57$ \\citep{SatKinematics}. Interestingly, imposing the same surface brightness detection threshold on all of our simulated systems yields a prediction of $77$ observable satellites for our \\textsc{BH} model, which is only moderately higher than the observations despite the fact that this simulation overpredicts the satellite luminosity function considerably. For our \\textsc{CR} model instead, the number drops considerably, to $18$, perhaps caused in part by an overprediction of the effective stellar radii of the satellites, which could easily arise from the limited spatial resolution of our simulations. In any case, this stresses that a large number of additional satellites may actually still be hidden just below the surface brightness limit of the SDSS \\citep[see also][]{Bullock2009}.   Our simulations have also highlighted the relative importance of some of the evolutionary aspects of satellite galaxies. In particular, we do not find a very distinctive mark of the epoch of reionization on the satellites, and most satellites continue their star formation activity in our simulations to much lower redshift than $z=6$. This suggests that simplified treatments of satellite histories, where relatively high cooling thresholds due to a ionizing UV background are invoked, are not particularly realistic.  Our simulation results agree much better with the scenario outlined in \\citet{Redef_MSP}, which in fact resembles many of our simulation findings quite closely.  We find that the observed relationship between V-band luminosity and velocity dispersion is quite well reproduced by our simulations, albeit with large scatter. The small amount of reliable observational data for the velocity dispersions leaves it unclear at present whether the larger scatter we find indicates a problem of the simulations or whether is is also present in reality. What is comparatively clear though is that the observed relation between stellar mass and maximum circular velocity is really tight, a finding that is also reproduced by our simulation results. On the other hand, the correlation between present-day stellar mass and star formation rate seen in our simulations seems to be not nearly as well-defined as in the observational data. This is related to the fact that we do not find a good correlation between the present stellar and gaseous masses; many simulated satellites have comparable stellar masses but differ in their gas fractions by huge factors. Gas-rich and completely gas-depleted satellites coexist in the same total and stellar mass regime, rendering a tight correlation with the star formation rate unlikely.  But perhaps the most significant discrepancy between the simulation results and observations lies in the inferred mass-to-light ratios.  The mass-to-light ratios of the simulated galaxies are off by about a factor of $5$ when compared at face value to the observational estimates. This means that they are either too massive, or too faint for their mass. The discrepancy could also be caused by a systematic underestimate of the total satellite masses in the observations. Due to the difficulty of reliably determining the `outer edge' of the dark matter halo of an orbiting satellite, this possibility cannot be easily excluded.  In summary, we find that the current generation of cosmological hydrodynamic simulations is able to explain many properties of the observed satellite population surprisingly well. We have shown that different feedback physics affects the satellite population strongly, with respect to quantities such as luminosity function, scaling relations, or star formation histories. This emphasizes the significant potential of ``near-field cosmology'' within our Local Group to inform the general theory of galaxy formation. Our work has also shown that it is not necessarily the physics of cosmic reionization and supernova feedback alone that is responsible for resolving the missing satellite problem. In fact, the role of reionization has probably been grossly overstated in many previous works, while other important feedback, such as cosmic rays, has been ignored. It will therefore be very interesting to refine the hydrodynamical simulations further in future work, and to make them more faithful in capturing all the relevant physics."
    },
    "1004/1004.5251.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {We have performed a comprehensive multiwavelength analysis of a sample of 20 starburst galaxies that show  a substantial population of very young massive stars, most of them classified as Wolf-Rayet (WR) galaxies. } {We have analysed optical/\\NIR\\ colours, physical and chemical properties of the ionized gas, stellar, gas and dust content, star-formation rate and interaction degree (among many other galaxy properties) of our galaxy sample using multi-wavelength data. We compile 41 independent star-forming regions  --with oxygen abundances between \\mbox{\\abox= 7.58} and 8.75--, of which 31 have a direct estimate of the electron temperature of the ionized gas. } % aims heading (mandatory) {This paper, only submitted to astro-ph, compiles the most common empirical calibrations to the oxygen abundance, and presents the comparison between the chemical abundances derived in these galaxies using the direct method with those obtained through empirical calibrations, as it is published in L\\'opez-S\\'anchez \\& Esteban (2010b).} % methods heading (mandatory) % {We compared the abundances provided by the direct method with those obtained through empirical calibrations} {We find that (i) the Pilyugin method (Pilyugin 2001a,b; Pilyugin \\& Thuan 2005), %\\citep{P01a,P01b,PT05}, which considers the $R_{23}$ and the $P$ parameters, is the best suited empirical calibration for these star-forming galaxies, (ii) the relations between the oxygen abundance and the $N_2$ or the $O_3N_2$ parameters provided by Pettini \\& Pagel (2004) give acceptable results for objects with \\abox$>$8.0, and (iii) the results provided by empirical calibrations based on photoionization models (McGaugh, 1991; Kewley \\& Dopita, 2002; Kobulnicky \\& Kewley, 2004) are systematically 0.2 -- 0.3 dex higher than the values derived from the direct method. These differences are of the same order that the abundance discrepancy found between recombination and collisionally excited lines. This may suggest the existence of temperature fluctuations in the ionized gas, as exists in Galactic and other extragalactic \\HII regions. } % conclusions heading (optional), leave it empty if necessary {All these results are included in the paper \\emph{Massive Star Formation in Wolf-Rayet galaxies IV. Colours, chemical-composition analysis and metallicity-luminosity relations}, L\\'opez-S\\'anchez \\& Esteban (2010b), A\\&A, in press (Sect.~4.4 and Appendix~A). Please, if this information is used, reference that paper and NOT this document, which have been only submitted to astro-ph to emphasize these results.} % {Our detailed analysis is fundamental to understand the nature of galaxies that show strong starbursts, as well as to know their star formation %history and the relationships with the environment. This study is complementary --but usually more powerful-- to the less detailed analysis of large %galaxy samples that are very common nowadays.}  %  {We consider that it is fundamental to perform a detailed analysis %of both the photometric and the chemical properties of the star-forming galaxies %to understand the evolutionary stage of each system. Such study is absolutely %needed to globally analyse the properties of larger galaxy samples.}      \\titlerunning{Massive star formation in Wolf-Rayet galaxies IVb: Empirical calibrations }  \\authorrunning{L\\'opez-S\\'anchez \\& Esteban} ",
        "introduction": "The knowledge of the chemical composition of galaxies, in particular of dwarf galaxies, is vital for understanding their evolution, star formation history, stellar nucleosynthesis, the importance of gas inflow and outflow, and the enrichment of the intergalactic medium. Indeed, metallicity is a key ingredient for modelling galaxy properties, because it determines \\UV, optical and \\NIR\\ \\mbox{colours} at a given age (i.e., Leitherer et al. 1999), nucleosynthetic yields (e.g., Woosley \\& Weaver 1995), the dust-to-gas ratio (e.g., Hirashita et al 2001), the shape of the interstellar extinction curve (e.g., Piovan et al. 2006), or even the properties of the Wolf-Rayet stars \\citep{Crowther07}.  The most robust method to derive the metallicity in star-forming and starburst galaxies is via the estimate of metal abundances and abundance ratios, in particular through the determination of the gas-phase oxygen abundance and the nitrogen-to-oxygen ratio. The relationships between current metallicity and other galaxy parameters, such as colours, luminosity, neutral gas content, star-formation rate, extinction or total mass, constrain galaxy-evolution models and give clues about the current stage of a galaxy. For example, is still debated whether massive star formation results in the instantaneous enrichment of the interstellar medium of a dwarf galaxy, or if the bulk of the newly synthesized heavy elements must cool before becoming part of the interstellar medium (ISM) that eventually will form the next generation of stars. Accurate oxygen abundance measurements of several \\HII regions within a dwarf galaxy will increase the understanding of its chemical enrichment and mixing of enriched material.   \\onecolumn   Furthermore, today it is the metallicity (which reflects the gas reprocessed by stars and any exchange of gas between the galaxy and its environment) and not the stellar mass (which reflects the amount of gas locked up into stars) of a galaxy the main problem to get a proper metallicity-luminosity relation, so that different methods involving direct estimates of the oxygen abundance, empirical calibrations using bright emission-line ratios or theoretical methods based on photoionization models yield very different values (i.e., Yin et al. 2007; Kewley \\& Elisson, 2008).  Hence precise photometric and spectroscopic data, including a detailed analysis of each particular galaxy that allows conclusions about its nature, are crucial to address these issues. We performed such a detailed photometric and spectroscopic study in a sample of strong star-forming galaxies, many of them previously classified as dwarf galaxies. The majority of these objects are Wolf-Rayet (WR) galaxies, a very inhomogeneous class of star-forming objects which share at least an ongoing or recent star formation event that has produced stars sufficiently massive to evolve into the WR stage \\citep{SCP99}. The main aim of our study of the formation of massive stars in starburst galaxies and the role that the interactions with or between dwarf galaxies and/or low surface brightness objects have in its triggering mechanism. In Paper~I \\citep{LSE08} we described the motivation of this work, compiled the list of the 20 analysed WR galaxies (Table~1 of Paper~I), the majority of them showing several sub-regions or objects within or surrounding them, and presented the results of the optical/\\NIR\\ broad-band and \\Ha\\ photometry. In Paper~II \\citep{LSE09} we presented the results of the analysis of the intermediate resolution long-slit spectroscopy of 16 WR galaxies of our sample -- the results for the other four galaxies were published separately, see \\citet{LSER04a,LSER04b,LSEGR06,LSEGRPR07}. In many cases, two or more slit positions were used to analyse the most interesting zones, knots or morphological structures belonging to each galaxy or even surrounding objects. Paper~III \\citep{LSE10a} presented the analysis of the O and WR stellar populations within these galaxies. Paper~IV \\citep{LSE10b} globally compile and analyse the optical/\\NIR\\ photometric data and study the physical and chemical properties of the ionized gas within our galaxy sample. The results shown in this paper haven been already published in Paper~IV. The final paper of the series (Paper~V; L\\'opez-S\\'anchez 2010) compiles the properties derived with data from other wavelengths (UV, FIR, radio, and X-ray) and complete a global analysis of all available multiwavelength data of our WR galaxy sample. We have produced the most comprehensive data set of these galaxies so far, involving multiwavelength results and analysed according to the same procedures. ",
        "conclusions": "We compared the abundances provided by the direct method with those obtained using the most common empirical calibrations in our sample of star-forming regions within Wolf-Rayet galaxies --see \\citet{LSE10b}--. The main conclusions are: \\begin{itemize} \\item The Pilyugin-method of \\citet{P01a,P01b}, which considers the $R_{23}$ and the $P$ parameters and is updated by \\citet{PT05}, is nowadays the best suitable empirical calibration to derive the oxygen abundance of star-forming galaxies. The cubic fit to $R_{23}$ provided by \\citet*{Nagao06} is not valid for analysing these star-forming galaxies. \\item The relations between the oxygen abundance and the $N_2$ or the $O_3N_2$ parameters provided by \\citet{PP04} give acceptable results for objects with \\abox$>$8.0. \\item The results provided by empirical calibrations based on photoionization models \\citep{McGaugh91,KD02,KK04} are systematically 0.2 -- 0.3 dex higher than the values derived from the direct method. These differences are of the same order as the abundance discrepancy found between abundances determined from recombination and collisionally excited lines of heavy-element ions. This may suggest temperature fluctuations in the ionized gas, as they exist in Galactic and other extragalactic \\HII\\ regions. \\end{itemize}    \\twocolumn"
    },
    "1004/1004.2741_arXiv.txt": {
        "abstract": "Inflation can occur near a point of inflection in the potential of flat directions of the Minimal Supersymmetric Standard Model. In this paper we elaborate on the complementarity between the bounds from Cosmic Microwave Background measurements, dark matter and particle physics phenomenology in determining the underlying parameters of MSSM inflation by specializing to the Minimal Supergravity scenario. We show that the future measurements from the Large Hadron Collider in tandem with all these constraints will significantly restrict the allowed parameter space. We also suggest a new perspective on the fine tuning issue of MSSM inflation. With quantum corrections taken into account, the necessary condition between the soft supersymmetry breaking parameters in the inflaton potential can be satisfied at scales of interest without a fine tuning of their boundary values at a high scale. The requirement that this happens at the inflection point determines a dimensionless coupling, which is associated with a non-renormalizable interaction term in the Lagrangian and has no bearing for phenomenology, to very high accuracy. ",
        "introduction": "Inflation is the dominant paradigm of the early universe cosmology to solve the problems of the hot big-bang model and create the seeds for structure formation. Although observations strongly support a period of superluminal expansion~\\cite{WMAP}, a successful realization of inflation within particle physics has remained as a challenge. Recently it has been shown~\\cite{AEGM,AEGJM,AKM} that inflation can happen within the Minimal Supersymmetric Standard Model (MSSM) and its minimal extensions.  In these models inflation occurs near a point of inflection along a $D$-flat direction~\\cite{GKM} in the scalar potential of Supersymmetric (SUSY) partners of quarks and leptons (called squarks and sleptons respectively). The scale of inflation is very low, $H_{\\rm inf} \\sim {\\cal O}(100~{\\rm MeV})$, and the requirement to generate density perturbations of the correct size singles out two $D$-flat directions, which consist of squarks and sleptons respectively, as the inflaton candidates. Since the inflaton belongs to the observable sector~\\footnote{Low scale inflection point inflation can also happen in the hidden sector~\\cite{ADS}.}, its couplings to matter and its decay products are known, therefore it is possible to track the thermal history of the universe right from the end of inflation. Also, inflation is compatible with SUSY dark matter~\\cite{ADM1} (and even a unified origin of inflation and dark matter may emerge~\\cite{ADM2})~\\footnote{For a review on MSSM inflation, see~\\cite{Rouzbeh}.}.   MSSM inflation has remarkable features. The mere fact that the inflaton is related to squarks and sleptons implies that it can be tested outside cosmology. This is quite interesting because it gives the first example of an inflationary model with predictions for phenomenology, and hence experiments other than measurements from the Cosmic Microwave Backgrond (CMB) are needed to identify the allowed parameter space of inflation. Another feature, which is due to the fact that inflation occurs near a point of inflection, is that, unlike other models of inflation, MSSM inflation can give rise to a wide range of the scalar spectral index~\\cite{LK,AEGJM} including the whole range allowed by the WMAP data~\\cite{WMAP}. This, coming as a virtue, also raises an issue. The robustness comes at the expense of a finely tuned relationship between two dimensionful parameters (i.e. the soft SUSY breaking mass and the $A$-term of the flat direction that plays the role of the inflaton). The seriousness of the issue is that this fine tuning is not protected by a symmetry and needs to be performed to several orders in perturbations theory.  In this work we investigate these two issues in more detail. We point out that once quantum corrections are taken into account, the necessary condition between the the soft mass and $A$-term can be satisfied without a fine tuning in their input values at a high scale like the Grand Unified Theory (GUT) scale. One actually needs to tune a dimensionless coupling that controls the Vacuum Expectation Value (VEV) of the inflection point, to ensure that the relationship between the soft SUSY breaking parameters is satisfied at the right scale. This coupling represents a non-renormalizable interaction term that has no bearing for phenomenology, and its only role is to lead to successful inflation in MSSM.  We also demonstrate the complementarity of cosmological and phenomenological bounds in restricting the parameter space of MSSM inflation by performing a detailed study for the Minimal Supergravity (mSUGRA) scenario. We show that bounds from SUSY dark matter and mass measurements at the Large Hadron Collider (LHC), as well as those from the muon anomalous magnetic moment and rare decays, significantly restrict the allowed region of the parameter space. More data from different experiments can therefore pin down the model parameters in the future.  The organization of this paper is as follows. In Section II we give a brief recount of inflection point inflation in MSSM. In Section III we discuss the parameter space of MSSM inflation and constraints from the cosmological density perturbations. We discuss the fine tuning issue in light of radiative corrections in Section IV, and suggest that it can be considered as tuning of a dimensionless parameter that is relevant only for inflation. We then specialize to the mSUGRA scenario in Section V and show how various bounds (dark matter, sparticle mass spectrum, muon anomalous magnetic moment, etc) significantly restrict the allowed parameter space. We close the paper by concluding remarks in Section VI. ",
        "conclusions": "MSSM inflation represents a realistic embedding of inflation in high energy physics where the inflaton has a natural place in a well-motivated and testable model of particle physics instead of being added as an extra field. In this paper we discussed some aspects of MSSM inflation mainly focusing on two main issues.  The robustness of MSSM inflation with regard to its predictions for density perturbations, which is due to the fact that inflation occurs near a point of inflection, is a remarkable feature. However, generating acceptable perturbations requires that a very precise relation between the soft SUSY breaking parameters be satisfied up to several orders in perturbation theory.  We suggested a different perspective on this issue. The necessary relationship can be satisfied at scales, which are phenomenologically interesting, without any fine tuning between the boundary values of the soft SUSY breaking parameters at an input scale (like the GUT scale). For given boundary values, after using the relevant RGEs, we can find the scale at which the relation is satisfied (Figs.~2,~3,~4). Requiring that this scale matches a point of inflection in the potential, determines a dimensionless coupling that represents a non-renormalizable interaction term in the superpotential to very high accuracy. This coupling has no bearing for phenomenology and its only role is to give rise to a point of inflection that is suitable for inflation.  Another important feature of MSSM inflation, due to the fact that the inflaton is related to squarks and sleptons, is that it can be tested outside cosmology. Once a specific framework is supposed, we can obtain predictions of the inflationary model for phenomenology. Then the CMB measurements can be combined with various phenomenological bounds to identify the allowed parameter space of the model.  We presented a detailed study of the parameter space of MSSM inflation in the case of mSUGRA scenario. We demonstrated the complementarity between the bounds from CMB measurements and those from phenomenology (Figs.~5,~6,~7) in the $m_\\phi-\\phi_0$ plane ($\\phi_0$ and $m_\\phi$ denoting the point of inflection VEV and the inflaton mass calculated at that scale respectively). The limits from SUSY dark matter and future mass measurement of SUSY particles at the LHC, as well as muon anomalous magnetic moment and rare decays, significantly restrict the allowed range of $m_\\phi$. On the other hand, more precise determination of the scalar spectral index from CMB experiments will further narrow down the allowed range of $\\phi_0$.  More data from a whole array of experiments (PLANCK, LHC, dark matter direct detection experiments, etc) will lead to tighter constraints on the allowed region of the parameter space. Eventually, this collaboration between cosmology and particle physics can be used to determine the underlying parameters of the MSSM inflation."
    },
    "1004/1004.2788_arXiv.txt": {
        "abstract": "{} { We investigate a model for the shallow decay phases of Gamma-ray Burst (GRB) afterglows discovered by Swift/XRT in the first hours following a GRB event. In the context of the fireball scenario, we consider the possibility that long-lived energy injection from a millisecond spinning, ultramagnetic neutron star (magnetar) powers afterglow emission during this phase.} {We consider the energy evolution in a relativistic shock subject to both radiative losses and energy injection from a spinning down magnetar in spherical symmetry. We model the energy injection term through magnetic dipole losses and discuss an approximate treatment for the dynamical evolution of the blastwave. We obtain an analytic solution for the energy evolution in the shock and associated lightcurves. To fully illustrate the potential of our solution we calculate lightcurves for a few selected X-ray afterglows observed by Swift and fit them using our theoretical lightcurves.} {Our solution naturally describes in a single picture the properties of the shallow decay phase and the transition to the so-called normal decay phase. In particular, we obtain remarkably good fits to X-ray afterglows for plausible parameters of the magnetar. Even though approximate, our treatment provides a step forward with respect to previously adopted approximations and provides additional support to the idea that a millisecond spinning (1-3 ms), ultramagnetic (B$\\sim 10^{14}-10^{15}$ G) neutron star loosing spin energy through magnetic dipole radiation can explain the luminosity, durations and shapes of X-ray GRB afterglows.} {} ",
        "introduction": "\\label{introduction}  Before the launch of Swift in November 2004, X-ray afterglows of long Gamma-Ray Bursts could be pointed with X-ray telescopes not earlier than several hours after the trigger. These observations showed in most cases a smooth power-law like decay $F(t)\\propto t^{-\\alpha}$, with typical index of $\\alpha\\geq 1$. With the advent of Swift, X-ray fluxes could be monitored from a few minutes after the burst trigger. These observations have revealed a complex behavior in the first few hours after the GRB, which nonetheless displays remarkably standard properties across different events. This behavior can be described with a double broken power law, with an initial very steep decay (up to few hundreds of seconds after the trigger) with $\\alpha>2$ followed by a shallow phase, lasting $\\sim 10^3-10^4$ s, with $\\alpha<0.8$ and, later on, a steeper 'normal' decay with $\\alpha\\sim1.2-1.4$ (Nousek et al. 2006, Gehrels et al. 2009).  The X-ray spectral slope does not change between the shallow and normal decay, in marked contrast to what would be expected in case this temporal break was caused by the passage of a characteristic synchrotron frequency in the X-ray band (e.g. Sari et al. 1998). A possible interpretation requiring no spectral variations in the observed energy band invokes prolonged energy injection into the external shock that is believed to give rise to the GRB afterglow. Energy injection could come either from relativistic shells impacting the fireball at late times (e.g. Rees and Meszaros 1998, Sari and Meszaros 2000) or from a long-lived central engine (e.g. Zhang et al. 2006, Nousek et al. 2006, Panaitescu et al. 2006a).  Among different hypotheses on the nature of GRB central engines, two major classes can be identified. The first considers the formation of a black hole- debris torus system, the prompt emission being related to accretion of matter from the torus during the first $\\sim 10-100$s (Narayan, Paczynski \\& Piran 1992, Woosley 1993, Meszaros, Rees \\& Wijers 1999). In this scenario, keeping the energy production active in order to power the afterglow for at least $\\sim 10^4$ s is a difficult and far from settled matter, (Mc Fadyen et al. 2001, Ramirez-Ruiz 2004, Cannizzo \\& Gehrels 2009, Barkov \\& Komissarov 2009).  An alternative class of models invokes the formation of a strongly magnetic (B$>10^{14}-10^{15}$ G) , millisecond spinning neutron star (NS, Usov 1992, Duncan \\& Thompson 1992, Blackman \\& Yi 1998, Kluzniak \\& Ruderman 1998, Wheeler 2000). Recently, time-dependent MHD simulations have shown that long GRBs can originate from the interaction between a relativistic and strongly magnetized wind produced by a newly-born NS and the surrounding stellar envelope. NS spin periods of $\\sim 1$ ms and ultrastrong magnetic fields, \\textit{i.e.} B$\\geq 10^{15}$ G, would be required in this case  (e.g. Thompson et al. 2004, Bucciantini et al. 2006, 2008, 2009; see also Tchekhovskoy, McKinney \\& Narayan 2009).  The newly formed NS is expected to loose its initial spin energy ($>10^{52}$ erg) at a very high rate for the first few hours through magnetic dipole spin down, something that provides a long-lived central engine in a very natural way. Dai \\& Lu (1998) considered this idea in relation to possible observable effects on the afterglow emission. Zhang \\& Meszaros (2001) argued that, in this scenario, achromatic bumps in afterglow lightcurves are expected for NS spin periods shorter than a few ms and magnetic fields stronger than several times $10^{14}$ G. Interestingly, studies of the origin of NS magnetism envisage that millisecond spin period at birth is the key property that allows a proto-NS to amplify a seed magnetic field to a strength far exceeding $10^{14}$ G, through efficient conversion of its initial differential rotation energy (e.g Duncan \\& Thompson 1992, Thompson \\& Duncan 1993). Such highly magnetized, fast spinning NSs are expected to loose angular momentum at a high rate in the first decades of their life and later become slowly rotating magnetars whose major free energy reservoir is in their magnetic field (Thompson \\& Duncan 1995, 1996, 2001, cfr. Woods \\& Thompson 2006, Mereghetti 2008). We term these NSs as magnetars since their birth even though, when they spin at millisecond period, their rotational energy is still the main free energy reservoir.  After the Swift discovery of early afterglow shallow phases, the magnetar scenario has been invoked to interpret the X-ray light curve of both some short and long GRBs (e.g. 051221A by Fan and Xu 2006; 060313 by Yu and Huang 2007; GRB 050801 by De Pasquale et al. 2007; 070110 by Troja et al. 2007).  For GRB 060729 this scenario was shown to provide a good agreement with the shallow and normal decay phases in the optical and X-ray bands (Grupe et al. 2007, Xu et al. 2009).\\\\ Finally we note that, besides the interest in understanding GRB physics, the very fast spin and huge magnetic field envisaged in the magnetar formation scenario makes these objects very interesting also for gravitational wave (GW) astronomy. Different possibilities for this to occur have been investigated in the literature (Palomba 2001, Cutler 2002, Stella et al. 2005, Dall'Osso \\& Stella 2007, Dall'Osso, Shore \\& Stella 2009, Corsi \\& Meszaros 2009) showing that, in astrophysically plausible conditions, GW emission might efficiently extract spin energy from the NS, in competition with magnetic dipole losses. The study presented in this paper builds on the ansatz that millisecond spinning magnetars are formed in the events that give rise to long GRBs. We investigate the evolution of energy in a relativistic blastwave subject to radiation losses due to shock deceleration in the ISM and energy injection from a magnetically braking NS. We extend previous treatments by describing the injection term by the standard magnetic dipole formula and deriving a prediction for the evolution of energy and luminosity that can interpret the X-ray afterglows through their shallow and normal decay phases altogether. We derive an approximate solution for the blastwave luminosity which we compare with X-ray GRB afterglow lightcurves observed by Swift. We obtain a remarkably good match to these lightcurves for the range of initial spin periods and magnetic field strengths expected for magnetars at birth. These results illustrate the potential of this scenario in explaining the early afterglow observations in a simple, unified picture. ",
        "conclusions": "\\label{conclusions} In the framework of prolonged energy injection models for GRB afterglows observed by Swift, we have considered the possibility that newly born magnetars - strongly magnetized and millisecond spinning NSs - are formed in the events producing (long) GRBs. In the first hours after formation of the NSs, the high spindown luminosity caused by magnetic dipole radiation losses represents a natural mechanism for prolonged energy injection in the external shock. To assess the viability of this scenario we considered the energy balance of a blastwave subject to injection of energy by a NS spinning down through magnetic dipole radiation, along with radiative losses ($\\propto E/t$). We found an approximate expression for the (isotropic) bolometric luminosity of the blastwave as a function of time that is in substantial agreement with general properties of the shallow-decay and normal-decay phases of X-ray GRB afterglows observed by Swift.  Moreover, we have shown that individual lightcurves can be very well fitted by using our derived expression for the bolometric luminosity of the continuously-powered blastwave. In particular, our best fits provide values for the initial spin period of the NS in the range 1-3 ms, which match well the values expected in magnetar formation scenarios. Best-fit values for the magnetic dipole field,  $10^{14}-10^{15}$ G, are also in the range expected for such objects at formation and in agreement with the dipole fields estimated for Anomalous X-ray Pulsars and Soft Gamma-ray Repeaters, the candidate magnetars in our Galaxy."
    },
    "1004/1004.2094_arXiv.txt": {
        "abstract": "\\label{sect: abstract}  \\noindent The Magellanic Mopra Assessment (MAGMA) is a high angular resolution \\aco\\ mapping survey of giant molecular clouds (GMCs) in the Large and Small Magellanic Clouds using the Mopra Telescope.  Here we report on the basic physical properties of 125 GMCs in the Large Magellanic Cloud (LMC) that have been surveyed to date. The observed clouds exhibit scaling relations that are similar to those determined for Galactic GMCs, although LMC clouds have narrower linewidths and lower CO luminosities than Galactic clouds of a similar size. The average mass surface density of the LMC clouds is 50~\\mpcsq, approximately half that of GMCs in the inner Milky Way. We compare the properties of GMCs with and without signs of massive star formation, finding that non-star-forming GMCs have lower peak CO brightness than star-forming GMCs. The average CO-to-\\hh\\ conversion factor, \\xco, of non-star-forming GMCs is also $\\sim50$ per cent larger than for star-forming GMCs. We compare the properties of GMCs with estimates for local interstellar conditions: specifically, we investigate the \\hi\\ column density, radiation field, stellar mass surface density and the external pressure. Very few cloud properties demonstrate a clear dependence on the environment; the exceptions are significant positive correlations between i) the \\hi\\ column density and the GMC velocity dispersion, ii) the stellar mass surface density and the average peak CO brightness, and iii) the stellar mass surface density and the CO surface brightness. The molecular mass surface density of GMCs without signs of massive star formation shows no dependence on the local radiation field, which is inconsistent with the photoionization-regulated star formation theory proposed by \\citet{mckee89}. We find some evidence that the mass surface density of the MAGMA clouds increases with the interstellar pressure, as proposed by \\citet{elmegreen89}, but the detailed predictions of this model are not fulfilled once estimates for the local radiation field, metallicity and GMC envelope mass are taken into account. ",
        "introduction": "\\label{sect:intro}  \\noindent In the Milky Way, molecular gas is mostly located in giant molecular clouds (GMCs) with masses $M>10^{5}$~\\msol\\ \\citep[][ henceforth S87]{solomonetal87}. Understanding the physical properties of the GMCs is important because these clouds are the primary sites of star formation: the formation of GMCs and the transformation of molecular gas into stars are key processes in the life cycle of galaxies. Models of galactic evolution typically assume that GMCs are sufficiently similar across different galactic environments that a galaxy's star formation rate can be parameterised as the product of the GMC formation rate and the star formation efficiency of molecular gas \\citep[e.g.][]{ballesterosparedeshartmann07,blitzrosolowsky06}. This approach was initially justified by studies of Galactic molecular clouds, which found that the basic physical properties of GMCs in the Milky Way's disc obeyed well-defined scaling relations, often referred to as ``Larson's laws'' \\citep[e.g. S87,][]{larson81,heyeretal01}. More recently, considerable effort has been devoted to determining whether GMCs in other galaxies also follow the Larson relations \\citep[e.g.][]{rosolowskyetal03,rosolowskyblitz05,rosolowsky07}, since empirical evidence that GMC properties are uniform -- or at least exhibit well-behaved correlations with a parameter such as metallicity or pressure -- would provide valuable information for developing models of star formation and galaxy evolution through cosmic time. \\\\  \\noindent As well as furnishing galaxy evolution models with empirical inputs, studies of extragalactic GMC populations aspire to resolve long-standing questions about the physical processes that are important for the formation and evolution of molecular clouds: are GMCs quasi-equilibrium structures, for example, or transient features in the turbulent interstellar medium? Do all GMCs form stars, and if not, why not? What is the physical origin of Larson's scaling relations?  Although a number of different theories to explain molecular cloud properties and the Larson relations have been proposed \\citep[e.g. M89, E89, ][]{chieze87,fleck88}, there are few extragalactic GMC samples that are comparable to the S87 catalogue, which contains 273 clouds in the Galactic disc between longitudes 8$^{\\circ}$ and 90$^{\\circ}$, and with radial velocities between -100 and 200~\\kms. The survey of \\aco\\ emission in the LMC by NANTEN (henceforth ``the NANTEN survey'') provided the first complete inventory of GMCs in any galaxy \\citep{fukuietal08}, but did not resolve molecular cloud structures smaller than $\\sim$40~pc \\citep[we adopt 50.1~kpc for the distance to the LMC, e.g.][]{alves04}. Thorough testing of the different molecular cloud models will require deep, unbiased wide-field surveys of molecular clouds at high angular resolution across a range of interstellar conditions. Extensive surveys of this kind are only just feasible with current instrumentation, and hence the number of molecular cloud samples that can be used to falsify molecular cloud models remains frustratingly small. \\\\  \\noindent To date, studies of the CO emission in nearby galaxies have concluded that extragalactic GMCs are alike. For a sample of $\\sim70$ resolved GMCs located in five galaxies (M31, M33, IC10, and the Magellanic Clouds), \\citet{blitzetal07} found that extragalactic GMCs not only follow the Galactic Larson relations, but also that different galaxies have similar GMC mass distributions. Similar conclusions were reached by \\citet[][henceforth B08]{bolattoetal08} using a sample of $\\sim100$ resolved GMCs in twelve galaxies, although these authors noted that molecular clouds in the SMC tend to have low CO luminosities and narrow linewidths compared to GMCs of a similar size in other galaxies. By comparing tracers of star formation and neutral gas on $\\sim~1$~kpc scales for galaxies in The \\hi\\ Nearby Galaxy Survey \\citep[THINGS, which does not include the Magellanic Clouds,][]{walteretal08}, \\citet{leroyetal08} found that the star-forming efficiency of molecular gas (defined as the star formation rate surface density per unit molecular gas surface density $SFE_{\\rm H_{2}} \\equiv \\Sigma_{\\rm SFR}/\\Sigma_{\\rm H_{2}}$) is approximately constant in normal spiral galaxies, $SFE_{\\rm H_{2}} = 5.25\\pm2.5 \\times 10^{-10}$~yr$^{-1}$. As noted by the authors, this result could arise if the star-forming efficiency of an individual GMC is determined by its intrinsic properties, and if the properties of GMCs are independent of their interstellar environment \\citep[e.g.][]{krumholzmckee05}. While the existing observational evidence has so far been interpreted in favour of uniform GMC properties, a dependence of GMC properties on the local interstellar environment is by no means ruled out. A constant $SFE_{\\rm H_{2}}$ on kiloparsec scales indicates that the properties of GMC ensembles are alike on those scales; whether this conclusion can be applied to individual GMCs is far less certain. Neither B08 nor \\citet{blitzetal07} pursued the origin of the scatter in the extragalactic Larson relations that they observed, moreover, even though the mean GMC mass surface density for the galaxies in their respective samples varies by more than an order of magnitude, and the mass surface densities of the individual GMCs varies between $\\sim10$ and 1000~\\mpcsq\\ \\citep[see also][ for evidence that the mass surface density of Milky Way clouds is not constant]{heyeretal09}.  A resolved survey of a large number ($>100$) of GMCs located in a single nearby galaxy therefore remains valuable, since it eliminates the uncertainties inherent in combining heterogeneous datasets and provides a sample that is large enough to investigate both the average properties and scaling relations of an extragalactic GMC population, as well as the dispersion around overall trends and average quantities.\\\\  \\noindent In this paper, we report on some initial results from the Magellanic Mopra Assessment (MAGMA), an ongoing, high-resolution survey of the \\aco\\ emission from molecular clouds in the Magellanic Clouds using the Mopra Telescope. Here we present results from the Large Magellanic Cloud (LMC) only; a description of the molecular clouds surveyed by MAGMA in the Small Magellanic Cloud (SMC) has been presented elsewhere (Muller \\ea, accepted). While the \\aco\\ emission from molecular gas in the LMC has been the target of extensive mapping with the NANTEN telescope and Swedish-ESO Submillimetre Telescope (SEST), neither project obtained observations that were ideal for studying the Larson relations in the LMC \\citep{fukuietal08,israeletal03}. The spatial resolution of the NANTEN survey is comparable to the size of a typical Milky Way GMC \\citep[$\\sim50$~pc, e.g.][]{blitz93}; ideally, we would like to resolve structures on smaller spatial scales in order to include less massive GMCs in our analysis. Resolved observations are crucial, moreover, for accurate estimates of derived GMC quantities such as virial mass and mass surface density. The SEST Key Programme {\\it CO in the Magellanic Clouds} mapped molecular clouds in the LMC with comparable spatial resolution as MAGMA ($\\sim10$~pc), but was strongly biased towards regions associated with well-known sites of active star formation. An analysis of the striking molecular cloud complex situated south of the 30 Doradus star-forming complex using the MAGMA data has already been presented by \\citet{ottetal08} and \\citet{pinedaetal09}; in this paper, we turn our attention to the general LMC cloud population.\\\\  \\noindent This paper is structured as follows: in Section~\\ref{sect:data}, we summarise the MAGMA observing strategy and our data reduction procedure, and also describe the ancillary data that we have used in our analysis. Section~\\ref{sect:cprops} outlines the approach that we have used to identify GMCs and to measure their physical properties. In Section~\\ref{sect:younggmcs}, we compare the properties of GMCs with and without star formation. Scaling relations between the cloud properties are discussed in Section~\\ref{sect:larsonlaws}, while Section~\\ref{sect:tracers} presents a comparison between the intrinsic physical properties of the GMCs and properties of the local interstellar environment. In Section~\\ref{sect:discussion}, we discuss whether our results are consistent with i) the photoionization-regulated theory of star formation proposed by \\citet[][ henceforth M89]{mckee89} and ii) a dominant role for interstellar gas pressure in the determination of molecular cloud properties, as suggested by \\citet[][ henceforth E89]{elmegreen89}. We conclude with a summary of our key results in Section~\\ref{sect:conclusions}. ",
        "conclusions": "\\\\  \\noindent 1. The observed GMCs have radii ranging between 13 and 160~pc, velocity dispersions between 1.0 and 6.1~\\kms, peak CO brightnesses between 1.2 and 7.1~K, CO luminosities between $10^{3.5}$ and $10^{5.5}$~K~\\kms~pc$^{2}$, and virial masses between $10^{4.2}$ and $10^{6.8}$~\\msol. The clouds tend to be elongated, with a median major-to-minor axis ratio of 1.7. These values are comparable to the measured properties of Galactic GMCs. The average mass surface density of the observed clouds is $\\sim50$~\\msol~pc$^{-2}$, approximately half the value determined for GMCs in the inner Milky Way catalogue of \\citet{solomonetal87}. \\\\  \\noindent 2. The MAGMA clouds exhibit scaling relations that are similar to those previously determined for Galactic and extragalactic GMC samples \\citep[e.g.][]{solomonetal87,bolattoetal08}. However the MAGMA LMC clouds are offset towards narrower linewidths and lower CO luminosities compared to GMCs of a similar size in these samples. The scatter in the scaling relations corresponds to order of magnitude peak-to-peak variations in the CO-to-\\hh\\ conversion factor (as inferred from the ratio of the virial mass to the CO luminosity), the \\hh\\ mass surface density and the CO surface brightness of the MAGMA GMCs.\\\\  \\noindent 3. The physical properties of star-forming GMCs are very similar to the properties of GMCs without signs of massive star formation. Sightlines through non-star-forming GMCs tend to have lower peak CO brightness, suggesting that the filling fraction and/or brightness temperature of the CO-emitting substructure is lower for clouds without star formation.   \\\\  \\noindent 4. We find a significant positive correlation between the peak CO brightness and CO surface brightness of the MAGMA clouds and the stellar mass surface density. We propose that these correlations are due to an increase in the CO brightness temperature and/or an increase in the abundance of CO relative to \\hh\\ in the stellar bar region. \\\\  \\noindent 5. The velocity dispersion ($\\sigma_{\\rm v}$) of the MAGMA GMCs increases in regions with high \\hi\\ column density (\\nh). Higher volume densities and/or higher virial parameters for GMCs in regions with high \\nh\\ could produce the observed correlation, although the MAGMA data does not provide unambiguous evidence for either of these alternatives.\\\\  \\noindent 6. There is some evidence that the \\hh\\ mass surface density of the MAGMA LMC clouds increases with the interstellar kinetic pressure, $P_{ext}$. Although the molecular cloud model proposed by \\citep{elmegreen89} predicts a relation between $P_{ext}$ and the mass surface density of an atomic+molecular cloud complex, the MAGMA clouds do not fulfil the predictions of the model for reasonable values of the metallicity, radiation field and GMC envelope mass."
    },
    "1004/1004.0869_arXiv.txt": {
        "abstract": "We present the results of a comprehensive infrared, submillimetre, and millimetre continuum emission study of isolated low-mass star-forming cores in 32 Bok globules, with the aim to investigate the process of star formation in these regions. The submillimetre and millimetre dust continuum emission maps together with the spectral energy distributions are used to model and derive the physical properties of the star-forming cores, such as luminosities, sizes, masses, densities, etc. Comparisons with ground-based near-infrared and space-based mid and far-infrared images from Spitzer are used to reveal the stellar content of the Bok globules, association of embedded young stellar objects with the submm dust cores, and the evolutionary stages of the individual sources. Submm dust continuum emission was detected in 26 out of the 32 globule cores observed. For 18 globules with detected (sub)mm cores we derive evolutionary stages and physical parameters of the embedded sources. We identify nine starless cores, most of which are presumably prestellar, nine Class\\,0 protostars, and twelve Class\\,I YSOs. Specific source properties like bolometric temperature, core size, and central densities are discussed as function of evolutionary stage. We find that at least two thirds (16 out of 24) of the star-forming globules studied here show evidence of forming multiple stars on scales between 1,000 and 50,000\\,AU. However, we also find that most of these small prototstar and star groups are comprised of sources with different evolutionary stages, suggesting a picture of slow and sequential star formation in isolated globules. ",
        "introduction": "\\label{sec-intro}   Different aspects of star formation can be studied on different size scales and in different environments. The large-scale distribution of star-forming regions and the relation between molecular cloud life cycles, galactic spiral density waves, and star formation can be studied by observing nearby galaxies \\citep[e.g.][]{1998A&A...333...92B,2009A&A...494...81S}. The stellar initial mass function (IMF), which is needed to interpret these data, is usually derived from rich young stellar clusters in our own Galaxy \\citep[e.g.][]{2002Sci...295...82K,2003ApJ...586L.133C}. Dense star-forming dark cloud complexes such as the Orion and Ophiuchus molecular clouds are the places to study the relation between the molecular core mass spectrum (CMF) and the interstellar IMF \\citep[e.g.][]{1998A&A...336..150M,2007MNRAS.374.1413N,2008MNRAS.391..205S,2008A&A...477..823G}.  Nearby and more isolated star-forming cores, such as Bok globules, are the best places to study in detail the initial properties of individual star-forming cores, their chemical evolution, kinematic structure, and the physics of their collapse and fragmentation \\citep[e.g.,][]{1988ApJS...68..257C,1995MNRAS.276.1052B,1997A&A...326..329L,1997MNRAS.288L..45L, 1998A&A...338..223H,2007prpl.conf...33W,2008ApJ...687..389S}. Bok globules are small, simply-structured, relatively isolated, opaque molecular clouds that often contain only one or two star-forming core. They are often not completely isolated, but reside in the filamentary outskirts of larger dark cloud complexes \\citep{1997A&A...326..329L}, a fact that may tell something about their origin. With their size, mass, densities, etc., Bok globules resemble small clumps in larger molecular clouds \\citep[cf. ][]{2007ARA&A..45..339B}, only that they lack the surrounding cloud. Table \\ref{tbl-globprop} summarizes the average general properties of typical Bok globules and their star-forming cores.  \\placetable{tbl-globprop}  Although they are the most simple star-forming molecular clouds, many globules deviate considerably from spherical geometry. They are often cometary or irregularly shaped. The dense star-forming cores are not always located at the center of the globule, but in cometary-shaped globules are often located closer to the sharper rim at the ``head'' side \\citep[e.g.,][]{Launhardt:1996}. Similarly, pre-stellar cores, which are the earliest stage of star formation \\citep{1994MNRAS.268..276W,2007prpl.conf...33W} often appear to be fragmentary and filamentary. However, the protostellar cores and envelopes of the more evolved Class 0 \\& I YSOs \\citep{1987IAUS..115....1L,1993ApJ...406..122A} are more spherically symmetric, which can be understood as a result of the gravitational collapse of the inner dense $R\\sim5000$\\,AU region. Many of these isolated cores were found to be the sources of bipolar molecular outflows, indicating the presence of embedded protostars \\citep[e.g.,][]{1994ApJS...92..145Y,1995MNRAS.274.1219W,1996A&A...314..477B}.  In order to investigate the star-forming potential and evolutionary stages of Bok globules in the solar neighbourhood, we had surveyed a large number of globules for signs of star-forming cores, using as tracers, e.g., the mm dust continuum emission \\citep{Launhardt:1996,1997A&A...326..329L,1998A&A...338..223H,1998ApJS..119...59L}, NH$_3$\\ \\citep{1995MNRAS.276.1067B} or CS line emission \\citep{1998ApJS..119...59L}.  In this paper we present a submillimetre and millimetre -- hereafter (sub)mm -- continuum study of 32 Bok globules, which were identified from these previous surveys as promising candidates for globules with currently ongoing star formation. The (sub)mm maps are complemented by deep near-infrared (NIR) images and NIR to mm spectral energy distributions (SEDs). In Sect.\\,\\ref{sec-obs}, we describe the observations and data reduction. In Sect.\\,\\ref{sec-res}, we present the NIR and (sub)mm images, SEDs, and results on multiplicity, physical parameters, and evolutionary stages. In Sect.\\,\\ref{sec-dis1}, we discuss in particular the source properties as function of evolutionary stage and the results on multiplicity. In Sect.\\,\\ref{sec-dis}, we describe and discuss the individual globules, and Sect.\\,\\ref{sec-sum} summarizes the main results of this study. ",
        "conclusions": "\\label{sec-sum}  We have studied the dense cores of 32 Bok globules and have obtained deep NIR images and (sub)mm dust continuum maps at up to three wavelengths (0.45, 0.85, and 1.3\\,mm). With the exception of a small control sample, all sources were selected from earlier surveys which identified them as good candidates for having ongoing star formation, i.e., this is not an unbiased survey. We also compiled SEDs, taking special care of separating the flux contributions from different neighbouring sources at different wavelengths, and fitted them to derive various source quantities and evolutionary stages of the sources. The main results of this study are: \\begin{enumerate} \\item We detected (sub)mm dust continuum cores in 26 out of the 32 globules observed. The tentative, low SNR, single-dish 1.3\\,mm continuum detections in CB\\,52 and BHR\\,41, published in and earlier paper, could not be confirmed.  \\item Eight of the 26 globules with detected (sub)mm cores are not studied in further detail or evaluated in terms of their evolutionary stage because they were either at too large distances ($>1$\\,kpc) and multiple embedded sources were not resolved, or the (sub)mm maps were of too low quality, or we simply did not have enough data to draw any reliable conclusions.  \\item In 18 globules with detected (sub)mm cores, we derived evolutionary stages and physical parameters of the embedded sources. In total, we identified nine starless cores, presumably prestellar, nine Class\\,0 protostars, and eleven Class\\,I YSOs in these 18 globules.  \\item We find that the bolometric temperature is the most reliable tracer to discriminate between Class\\,0 protostars and Class\\,I YSOs and confirm the empirical boundary of 70\\,K. The spread of $L_{\\rm smm}\\,/\\,L_{\\rm bol}$\\ ratios within the Class\\,0 and Class\\,I groups is relatively large (2\\,--\\,10\\% within Class\\,0 and 0.8\\,--\\,3.5\\% within Class\\,I), with no significant correlation between $T_{\\rm bol}$\\ and $L_{\\rm smm}\\,/\\,L_{\\rm bol}$\\ within the groups. However, the three most evolved Class\\,I sources, with visible stars and compact (sub)mm emission arising presumably from circumstellar disks, also have the lowest $L_{\\rm smm}\\,/\\,L_{\\rm bol}$\\ ratios ($\\le 1.3$\\%). We take this as tentative indication that the $L_{\\rm smm}\\,/\\,L_{\\rm bol}$\\ ratio, as indicator of envelope dispersal, may better trace the evolution within the Class\\,0 and Class\\,I phases.  \\item The mean FWHM core sizes decrease from 8,000\\,AU for prestellar cores, to 3,000\\,AU for Class\\,0 protostars, to $<$2,000\\,AU for Class\\,I YSOs. The latter group also exhibits extended remnant envelopes with diameters of order 20,000\\,AU. Source-averaged volume densities, $n_{\\rm H}$, increase from $1\\times 10^7$\\,cm$^{-3}$\\ for prestellar cores, to $7\\times10^7$\\,cm$^{-3}$\\ for Class\\,0 protostars, to $>8\\times10^7$\\,cm$^{-3}$\\ for Class\\,I YSOs. The extended envelopes of the Class\\,I YSOs have a mean density of $7\\times 10^4$\\,cm$^{-3}$.  \\item At least two thirds (16 out of 24) of the star-forming globules studied here show evidence of forming multiple stars on scales between 1,000 and 50,000\\,AU, either as multiple star-forming cores, wide embedded binaries, or small star clusters. The fraction of closer binaries formed from unresolved mm cores might be higher, but remains unknown from this study.  \\item We find that the large majority of these small prototstar and star groups in globules with multiple star formation are comprised of sources with different evolutionary stages. This includes neighbouring mm sources with obviously different evolutionary stages, prestellar or protostellar cores with nearby IR sources, presumably more evolved protostars or Class\\,I YSOs, as well as NIR star clusters next to large (sub)mm cores with the potential to form more stars. In only three globules we find coeval pairs, ranging from multiple pre-stellar cores in CB\\,246, embedded Class\\,0 protostars in BHR\\,71, to an embedded Class\\,I YSOs pair in CB\\,230. This widespread non-coevality possibly suggests a picture of slow and sequential star formation in isolated globules.  \\item These findings also call for special attention when compiling SEDs and attempting to derive source properties from flux measurements with insufficient angular resolution. One may easily end up classifying the combined SED of a prestellar core and a nearby, more evolved YSO as a Class\\,0 protostar. \\end{enumerate}  While this paper presents an extensive, though not complete inventory of star-forming cores in nearby Bok globules, it remains subject of other detailed studies to investigate the physical and chemical properties, multiplicity and evolutionaly stages of individual sources, and to evaluate if and how star formation in isolated globules differs from that in larger molecular cloud complexes. We expect soon to obtain spatially resolved mid to far-infrared Herschel data for many of the globule cores presented here and plan to use these data to accurately measure the dust temparure profiles and derive more reliable density profiles. In particular for the prestellar cores, which characterize the initial conditions of the protostellar collapse, this may lead to significant corrections of previous estimates of mass and density distributions which were often hampered by the lack of temperature measurements. Furthermore, we are currently working on detailed studies of two globule cores which characterize the stage of the onset of the protostellar collapse and which may shed light on the non-coeval evolution of different sub-cores within one globule (Schmalzl et al., in prep.). Last, but not least, this paper shall provide some guidance for further follow-up studies of individual globules, e.g., with Herschel, SCUBA2, and ALMA."
    },
    "1004/1004.4810_arXiv.txt": {
        "abstract": "One of the most pressing issues in cosmology is whether general relativity (GR) plus a dark sector is the underlying physical theory or whether a modified gravity model is needed. Upcoming dark energy experiments designed to probe dark energy with multiple methods can address this question by comparing the results of the different methods in constraining dark energy parameters. Disagreement would signal the breakdown of the assumed model (GR plus dark energy). We study the power of this consistency test by projecting constraints in the $w_0-w_a$ plane from the four different techniques of the Dark Energy Survey in the event that the underlying true model is modified gravity. We find that the standard technique of looking for overlap has some shortcomings, and we propose an alternative, more powerful {\\it Multi-dimensional Consistency Test}.  We introduce the methodology for projecting whether a given experiment will be able to use this test to distinguish a modified gravity model from GR. ",
        "introduction": "General relativity (GR) is currently a bad fit to cosmological data unless a new substance, so-called {\\it dark energy}, is invoked. If GR really is an incomplete or incorrect theory and we are tasked with identifying the correct model, a major hurdle will be determining how to confront upcoming data sets in the absence of a well-understood model. What new parameters should be introduced and fit for when, e.g., data on weak gravitational lensing or galaxy clusters are analyzed?  Several authors have addressed this question~\\cite{Linder:2005in,Linder:2007hg,Zhang:2007nk,Hu:2007pj,Zhao:2009fn}, and it has recently become possible to test GR using survey data \\cite{rapetti_allen_etal_2009, daniel_linder_etal_2010, reyes_mandelbaum_etal_2010,Lombriser:2010mp}.  Here we address a slightly less ambitious question: using multiple cosmological probes, how can we determine whether cosmic acceleration is driven by dark energy or modified gravity (MG)?  One approach is to analyze the data assuming that GR is correct and see whether the constraints on dark energy parameters from different probes overlap  \\cite{Annis:2005ba, ishak_upadhye_etal_2006}. Non-overlapping constraints would be a strong signal that the underlying parameterization is wrong; i.e, that GR+dark energy cannot account for the data and that a modified theory of gravity is called for.  A similar approach is to look at parameter constraints coming from separate dynamical effects such as the cosmic expansion or perturbation growth \\cite{zhang_hui_etal_2005}.  Here we explore the former method in depth in the context of a concrete example.  Ishak et al. showed that, in principle, non-overlapping dark energy parameter constraints obtained from multiple experiments is a signature of MG \\cite{ishak_upadhye_etal_2006}.  In particular, they found that dark energy parameters obtained from a space-based supernova survey and a space-based weak lensing survey will not agree if the Universe is in fact described by the Dvali-Gabadadze-Porrati (DGP) braneworld model \\citep{dvali_gabadadze_etal_2000}.  We reexamine this general method with our own example, assuming that the universe is governed by a toy MG model and considering projections from the upcoming Dark Energy Survey (DES). We present the projected constraints from all four DES probes in the plane of dark energy parameters $w_0$ and $w_a$, where the dark energy equation of state is assumed to be $w=w_0+w_a(1-a)$ and $a$ is the scale factor of the universe. This straightforward plot is not the most powerful way to combine probes, so we introduce a more quantitative formalism that should be useful for future attempts in this direction. The formalism assigns a $\\chi^2$ for the combined probes which can be interpreted in the usual fashion so that a ``bad'' $\\chi^2$ corresponds to disagreement among the probes, and therefore a quantitative assessment of how well the model of GR+dark energy works.  Section II discusses modified gravity models in general and details the modified gravity model we adopt as our working example. Section III then presents the DES projections in the $(w_0,w_a)$ plane along with a description of the shortcomings of this approach. In Section IV, we present a more quantitative approach (see also \\cite{bernstein_huterer_2010}), which we call the {\\it Multi-dimensional Consistency Test} (MCT), illustrate how to obtain MCT projections, deal with the issue of degenerate directions, and finally conclude by applying this formalism to DES for the model under study. % ",
        "conclusions": "The consistency of different dark energy probes promises to be a powerful tool in the quest to distinguish dark energy from modified gravity. Here we have illustrated that, using the Multi-dimensional Consistency Test (MCT), future probes from the Dark Energy Survey will be able to rule out standard (GR+dark energy) if the true gravity model is only a modest modification of GR. Carrying out the MCT once the data are in reduces to computing the $\\chi^2$ of \\ec{chi2_def} while properly accounting for degeneracies as described in \\S\\ref{sec:param_degen}. Although we have not explored this in detail here, projections of the MCT might make a useful metric for future surveys when trying to understand their constraining power towards modified gravity models, complementary to the figures of merit for dark energy~\\cite{detf}."
    },
    "1004/1004.1118_arXiv.txt": {
        "abstract": "We calculate the energy spectra of cosmic rays (CR) and their secondaries produced in a supernova remnant (SNR), taking into account the time-dependence of the SNR shock. We model the trajectories of charged particles as a random walk with a prescribed diffusion coefficient, accelerating the particles at each shock crossing. Secondary production by CRs colliding with gas is included as a Monte Carlo process. We find that SNRs produce less antimatter than suggested previously: The positron/electron ratio $\\re$ and the antiproton/proton ratio $\\rp$ are a few percent  and few $\\times 10^{-5}$, respectively. Both ratios do not rise with energy. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.1604_arXiv.txt": {
        "abstract": "{Recent $Suzaku$ X-ray observations of the ejecta-dominated supernova remnant W49B have shown that in the global spectrum there is a clear indication for the presence of overionized plasma whose physical origin is still under debate.} {In order to ascertain the physical origin of such a rapidly cooling plasma, we focus on the study of its spatial localization within the X-ray emitting ejecta.} {We confirm the presence of a saw-edged excess (interpreted as a strong radiative recombination continuum) in the global spectrum above 8 keV, emerging above the ionization-equilibrium model. We produce a hardness ratio map to determine where the plasma is overionized and we perform a spectral analysis of the regions with and without strong overionization.} {We find that the overionized plasma is localized in the center of the remnant and in its western jet, while it is not detected in the bright eastern jet, where the expansion of the ejecta is hampered by their interaction with a dense interstellar cloud.} {The location of overionized plasma suggests that the inner ejecta are rapidly cooling by expansion, unlike the outer ejecta, for which expansion is hampered by interstellar clouds seen in H$_2$} ",
        "introduction": "\\label{Introduction}  X-ray observations of young supernova remnants (SNRs) allow us to probe the physical and chemical conditions of the shock-heated ejecta, and to study the physical processes involved in their interaction with the circumstellar ambient medium. %  W49B is one of the brightest ejecta-dominated SNRs observed in X-rays, where it shows a jet-like morphology. % Its global X-ray spectrum, characterized by intense emission lines from He-like and H-like ions of overabundant metals (Si, S, Ar, Ca, Fe, and also Cr and Mn) indicates a plasma at collisional ionization equilibrium (CIE, see, for example, \\citealt{fti95} and \\citealt{hph00}). These results are confirmed by spatially resolved spectroscopy with \\emph{XMM-Newton} (\\citealt{mdb06} hereafter M06) which reveals a significant Ni overabundance ($Ni/Ni_{\\odot}=10^{+2}_{-1}$) in the center. M06 found similar chemical composition and temperatures of the ejecta in the central and eastern regions, while in the western part the abundances (in particular for Fe) and the temperatures are lower. % Both the comparison of the observed ejecta abundances with the yields predicted by explosive nucleosynthesis models (performed by M06) and the multipole expansion spatial analysis (performed by \\citealt{lrb09}) concur in indicating a core-collapse origin for W49B.  The eastern border of the remnant shows bright radio emission (\\citealt{llk01}), and is spatially coincident with a shocked molecular H$_{2}$ cloud that hampers the expansion of the ejecta in the east direction (being about three order of magnitude denser than the ejecta, see \\citealt{krr07} and M06), thus distorting the eastern jet southward. Molecular hydrogen and [Fe II] and radio emission have been observed in the south-western edge of W49B, while in the center the [Fe II] emission reveals a barrel-like structure ``surrounding\" the jet (\\citealt{krr07}). \\citet{mdb08} have shown that W49B seems to be the result of an aspherical jet-like supernova explosion with explosion energy $\\la 1.5\\times 10^{51}$ erg and mass of the shocked ejecta $\\sim 6$ M$_{\\odot}$.  \\citet{kon05} analyzed the global spectrum of W49B and measured the intensity ratio of the H-like to the He-like K$\\alpha$ lines of Ar and Ca, finding that the ionization temperature $T_z$ is larger than the electron temperature, $T_e$, thus claiming the presence of overionization in W49B. % M06 found $T_z>T_e$ for Ca and no evidence of overionization for Ar. Nevertheless, they also pointed out that the estimate of the electron temperature derived from the global spectrum is not reliable, given that it is significantly lower than that obtained from the spatially resolved spectral analysis of homogeneous regions. M06 also performed the same analysis on the uniform central region of W49B. Finding no evidence for overionization in the spectrum, they concluded that the temperatures of the Ar and Ca ions are consistent with $T_e$. Recently, \\citet{oky09} (hereafter O09) analyzed the $Suzaku$ spectrum extracted from the whole remnant and revealed the presence of a saw-edged bump above 8 keV. They demonstrated that the bump is associated with a strong radiative recombination continuum of iron and produced an accurate spectral model to describe both the recombination continuum and the lines and to obtain diagnostics for the overionized ejecta (similar radiative recombination features have been also observed by \\citealt{yok09} in IC 443, but for Si and S). O09 derived that, for Fe, the ionization temperature is $kT_z\\sim2.7$ keV, while $kT_e$ is $\\sim 1.5$ keV. Moreover, O09 estimated that the overionized plasma has the same emission measure as the bremsstrahlung emitting plasma, thus suggesting a common origin for the two components. These results strongly indicate that the Fe ejecta in W49B are cooling so fast that they are overionized. Nevertheless, there are still a few important open issues.  First of all the electron temperature derived from the global spectrum by O09 is systematically lower than that observed in the uniform spectral regions discussed in M06, where $kT_e$ ranges between $\\sim1.8$ keV and $\\sim3$ keV, and in the accurately selected regions presented in \\citet{lrp09} (based on $Chandra$ data), where $kT_e=1.8-3.7$ keV. This discrepancy in $kT_e$ may be due to the fact that the global spectrum originates from physically non-uniform regions of W49B with different temperatures and abundances. Secondly, it is important to verify if all the ejecta in W49B are overionized. While M06 did not find overionization effects for Ar and Ca, O09 detected overionization for Fe. Notice that \\citet{lrp09} have shown that the iron morphology is very distinct from that of other elements, being more asymmetric and more segregated (and localized in the central and eastern parts of the remnant), while Ar and Ca appear well mixed and more isotropic. So, in principle, it is possible that the physical conditions in the Fe-rich ejecta are different from those of the other ejecta. Finally, a spatial localization of the overionized plasma may contribute to ascertain the physical origin of its rapid cooling, still not understood.  Here we present the analysis of archive \\emph{XMM-Newton} observations of W49B specifically devoted to constrain the spatial distribution of the overionized plasma. The data analysis procedure is shown in Sect. \\ref{The Data} and our results are presented and discussed in Sect. \\ref{Results} and Sect. \\ref{Discussion}, respectively. % ",
        "conclusions": "\\label{Discussion}  Our analysis confirms the presence of a radiative recombination continuum in the high energy X-ray spectrum of W49B. As explained by O09, this indicates that the ejecta (at least the Fe ions) are overionized and are rapidly cooling to reach the ionization equilibrium. Nevertheless, we have shown that the bright eastern ejecta of W49B do not present tracers of overionization. It is important to stress that the eastern jet of ejecta cannot expand freely because a large and dense molecular cloud hampers the expansion in the east direction (\\citealt{krr07}). The western jet, instead, is expanding undisturbed, since it does not interact with any dense structure.  The overionization due to a rapid cooling of the ejecta is therefore observed only where the ejecta can expand freely. Such rapid cooling is then likely associated with the adiabatic expansion of the central-western jet. A similar scenario has been invoked by \\citet{yok09} to explain the overionization effects in the supernova remnant IC 443. We expect that the overionized Fe ejecta experienced a strong heating in the early phases of the SNR evolution. This is quite reasonable considering that the explosion occurred in a very complex environment, likely shaped by winds from the progenitor stars, and with the circumstellar material showing a barrel-like morphology (\\citealt{krr07}). The interaction of the remnant with such environment can produce strong reverse and reflected shocks that can efficiently heat and ionize the ejecta. The subsequent rapid expansion of the ejecta in the center and to the West then produced their rapid cooling. We have estimated that a temperature $T\\sim5\\times10^7$ K and $\\int n_edt\\sim2.5\\times10^{12}$ s cm$^{-3}$ would be necessary, in order to have equal ionization time (to H-like Fe) and electron heating time. Nevertheless, an accurate analysis requires a modeling of the hydrodynamic evolution of the plasma including the effects of non-equilibrium of ionization (e.g. \\citealt{ro08}).  Not all the ejecta of W49B show signatures of overionization for the Fe ions and Fig. \\ref{fig:excess} shows that in regions with similar (and large) Fe abundances, like the eastern jet and the central regions (see M06 and \\citealt{lrp09}), there are different states of ionization (equilibrium to the East and overionization in the center), while the central and western regions (that have different Fe abundances, see M06 and \\citealt{lrp09}) are both overionized. We conclude that, while we derive overionization effects for the Fe ions, not all the Fe-rich ejecta are overionized. On the other hand we remind that in the center of W49B no overionization effects are visible for the Ar and Ca ions (M06). % We expect, therefore, different ionization temperatures and different ratios $kT_z/kT_e$ for different species and, at least in the case of Fe, for different regions of the remnant. In principle, it is also possible that the temperature of the electrons responsible for the Fe recombination is different from that of the electrons responsible for the bremsstrahlung emission in the $4.4-6.2$ keV band. However, from the global spectrum of W49B, O09 found that both the radiative recombination and the bremsstrahlung continua can be associated with electrons at the same temperature. Nevertheless, a spatially resolved study appears necessary in future works.  Another important issue is related to the determination of the value of the ion and electron temperatures in the ejecta. O09 have modeled both the radiative recombination continuum and line emission by deriving the ionization temperature from the global spectrum, finding $kT_z\\sim2.6$ keV that is larger than the electron temperature $kT_e\\sim1.5$ keV. In our large regions we also find similar values of $kT_e$, as shown in Table \\ref{tab:specres}. Nevertheless, these values are significantly lower than those derived from smaller regions by M06. For example, our region Y is approximately the union of regions 3, 4, 5, and 6 in M06, where they find $kT_e\\sim1.8-2.6$ keV that is significantly larger than the value $kT_e\\sim1.75$ keV reported in Table \\ref{tab:specres}. It is then important to analyze spectra extracted from relatively small regions in order to constrain accurately the ion and electron temperature distribution. Unfortunately, the available statistics does not allow us such a detailed analysis and we then stress the importance of obtaining deeper observations of W49B to address this point."
    },
    "1004/1004.3753_arXiv.txt": {
        "abstract": "Hydrodynamical models of colliding hypersonic flows are presented which explore the dependence of the resulting dynamics and the characteristics of the derived X-ray emission on numerical conduction and viscosity. For the purpose of our investigation we present models of colliding flow with plane-parallel and cylindrical divergence. Numerical conduction causes erroneous heating of gas across the contact discontinuity which has implications for the rate at which the gas cools. We find that the dynamics of the shocked gas and the resulting X-ray emission are strongly dependent on the contrast in the density and temperature either side of the contact discontinuity, these effects being strongest where the postshock gas of one flow behaves quasi-adiabatically while the postshock gas of the other flow is strongly radiative.  Introducing additional numerical viscosity into the simulations has the effect of damping the growth of instabilities, which in some cases act to increase the volume of shocked gas and can re-heat gas via sub-shocks as it flows downstream. The resulting reduction in the surface area between adjacent flows, and therefore of the amount of numerical conduction, leads to a commensurate reduction in spurious X-ray emission, though the dynamics of the collision are compromised.  The simulation resolution also affects the degree of numerical conduction. A finer resolution better resolves the interfaces of high density and temperature contrast and although numerical conduction still exists the volume of affected gas is considerably reduced. However, since it is not always practical to increase the resolution, it is imperative that the degree of numerical conduction is understood so that inaccurate interpretations can be avoided. This work has implications for the dynamics and emission from astrophysical phenomena which involve high Mach number shocks. ",
        "introduction": "\\label{sec:intro}  Colliding hypersonic flows occur in a number of astrophysical environments and over a wide range of scales, e.g. massive young stellar objects \\citep{Parkin:2009b}, astrophysical jets \\citep{Falle:1991, Shang:2006, Bonito:2007, Sutherland:2007}, colliding wind binary systems \\citep[CWBs,][]{Stevens:1992, Pittard:2009}, wind-blown bubbles around evolved stars \\citep[see][and references there-in]{Arthur:2007}, SNe \\citep[e.g.][]{Tenorio-Tagle:1991, Dwarkadas:2007}, and the cumulative outflows from young star clusters \\citep[][]{Canto:2000, Rockefeller:2005, Wunsch:2008, Rodriquez-Gonzalez:2008, Reyes-Iturbide:2009} and starburst galaxies \\citep[][]{Strickland:2000, Tenorio-Tagle:2003, Tang:2009}. Flow collisions can be subject to turbulent motions, the growth of linear and non-linear instabilities in boundary layers, and in some cases a global instability of the shocked gas \\citep[i.e. radiative overstability, ][]{Chevalier:1982}. The combination of these effects leads to complex scenarios for which numerical hydrodynamics has proved to be a useful investigatory tool.  However, in the discretization of the governing equations of hydrodynamics, additional terms are introduced which are purely numerical in origin. Depending on the order of the scheme, the appearance of these terms acts to disperse or dissipate the solution, and therefore terms such as ``numerical dispersion'', ``numerical diffusion'' or ``artificial viscosity'' are often used to describe them. The undesirable effects of numerical diffusion are minimized as one uses higher order schemes, though all schemes are only first order accurate near discontinuites such as shocks (where flow variables as well as the perpendicular velocity component, $v_{\\rm p}$, are discontinuous) and contact discontinuities (where there is a density and/or temperature jump but $v_{\\rm p}$ is unchanged). Contact discontinuities, and interfaces between different fluids, create special problems for multi-dimensional hydrodynamic codes. Unlike shocks, which contain a self-steepening mechanism, contact discontinuities spread diffusively during a calculation, and continue to broaden as the calculation progresses \\citep[see e.g. ][]{Robertson:2010}. Some schemes employ an algorithm known as a contact discontinuity steepener to limit this diffusion \\citep[e.g.][]{Fryxell:2000}. However, their use remains controversial, since the algorithm is based on empirical values with no physical or mathematical basis, and requires some care, since under certain circumstances it can produce incorrect results (i.e. ``staircasing'', Blondin, private communication).  Purely numerical effects are most prevalent when there are large density and temperature contrasts. Unfortunately, these frequently occur in practice, as when radiative cooling is effective cold dense regions of gas can form. Such regions are also inherently unstable, and compressed interface layers may be fragmented resulting in cold dense clumps/filaments residing next to hot tenous gas. When modelling such phenomena, the numerical transfer of heat from hot to cold cells can change the behaviour of the shocked gas. In particular, hot cells on one side of the contact discontinuity can reduce the net cooling rate of denser gas in adjacent cells on the other side of the contact discontinuity, and vice-versa. A further concern comes when one derives synthetic emission from the simulation output. For instance, due to the $\\rho^2$ dependence of {\\it thermal} emission, artificial heating caused by numerical conduction can cause dramatic differences in the spectral hardness and the magnitude of the integrated luminosity.  The goal of this work is to provide both a qualitative and quantitative analysis of the effects of numerical conduction and viscosity on the dynamics and observables from colliding flows as a function of the density and temperature constrast between the postshock gas. For the purposes of our investigation we have performed hydrodynamic simulations of colliding flows in plane-parallel and cylindrical geometries. In both scenarios the influence of efficient radiative cooling and powerful instabilities cause cold dense layers/clumps to reside next to hot rarefied gas. We show that the calculated X-ray emission from the postshock gas is strongly dependent on the parameters of the opposing flows.  The remainder of this paper is structured as follows: in \\S~\\ref{sec:model} we give a description of the hydrodynamics code and details of the X-ray emission calculations. In \\S~\\ref{sec:results} we present model descriptions and results, in \\S~\\ref{sec:discussion} a discussion, and we close with conclusions in \\S~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have presented hydrodynamic models of colliding hypersonic flows with the aim of examining the effects of numerical conduction on the simulation dynamics and the derived X-ray characteristics. The conduction of heat occurs across flow discontinuities due to diffusive terms introduced in the discretization of the governing flow equations. A key conclusion from this work is that the magnitude of the numerical heat conduction is strongly related to the density (and temperature) contrast between adjacent gas. X-ray calculations performed on the simulation results show that significant changes to spectra can occur by numerical conduction alone. Further tests performed with additional artificial viscosity reveal a complicated relationship between the flow dynamics, the magnitude of numerical conduction, and the resulting X-ray emission. For instance, the inherent instability of the collision regions of hypersonic flows naturally enhances the interface area between the flows, which in turn enhances the level of numerical conduction. Introducing sufficient viscosity to damp the growth of instabilities can reduce these effects, but the additional diffusion introduced into the fluid equations may increase the level of numerical heat conduction where the interface is relatively stable (e.g. near the apex of the wind-wind collision region in a colliding winds binary system). Finally, we note that while enhancing the resolution of the simulation increases the growth of small scale instabilities, and thus the area of the interface between the hot and cold phases, the overall effect of numerical conduction is reduced.  In the present work we have highlighted a fundamental problem encountered when using grid-based hydrodynamics to model fluids where high density and temperature contrasts are present - conditions which can be found in a multitude of astrophysical phenomena. Unfortunately, there is no simple fix. The brute-force approach to resolving this problem would be to employ higher simulation resolution, though this is not always a realistic option.  \\subsection*{Acknowledgements} This work was supported in part by a Henry Ellison Scholarship from the University of Leeds, and by a PRODEX XMM/Integral contract (Belspo). JMP gratefully acknowledges funding from the Royal Society and previous discussions with Robin Williams which instigated this work. 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."
    },
    "1004/1004.1756_arXiv.txt": {
        "abstract": "We present a sample of 8498 quasars with both SDSS $ugriz$ optical and UKIDSS $YJHK$ near-IR photometric data. With this sample, we obtain the median colour-z relations based on 7400 quasars with magnitude uncertainties less than 0.1mag in all bands. By analyzing the quasar colours, we propose an empirical criterion in the $Y-K$ vs. $g-z$ colour-colour diagram to separate stars and quasars with redshift $z<4$, and two other criteria for selecting high redshift quasars. Using the SDSS-UKIDSS colour-z relations, we estimate the photometric redshifts of 8498 SDSS-UKIDSS quasars, and find that 85.0\\% of them are consistent with the spectroscopic redshifts within $|\\Delta z|<0.2$, which leads to a significant increase of the photometric redshift accuracy than that based on the SDSS colour-z relations only. As two tests, we compare our colour selection criterion with a small UKIDSS/EDR quasar/star sample and a sample of 4671 variable sources in the SDSS Stripe 82 region with both SDSS and UKIDSS data. We find that they can be clearly divided into two classes (quasars and stars) by our criterion in the $Y-K$ vs. $g-z$ plot. We select 3834 quasar candidates from the variable sources with $g<20.5$ in Stripe 82, 826 of them being SDSS quasars and the rest without SDSS spectroscopy. We estimate the photometric redshifts for 3519 quasar candidates with all UKIDSS $YJHK$ data and find an accuracy of 87.5\\% within $|\\Delta z|<0.2$ with the spectroscopic redshifts of 819 SDSS-UKIDSS identified quasars among them. We demonstrate that even at the same spectroscopy limit as SDSS, with our criterion we can at least partially recover the missing quasars with $z\\sim2.7$ in SDSS. The SDSS identified quasars only take a small fraction (21.5\\%) of our quasar candidates selected from the variable sources in Stripe 82, indicating that a deeper spectroscopy is very promising in producing a much larger sample of quasars than SDSS. The implications of our current results to the future Chinese LAMOST quasar survey are also discussed. ",
        "introduction": "The number of quasars has increased substantially in the last decade primarily due to two large optical sky surveys, namely the Two Degree Field (2DF) survey (Boyle et al. 2000) and the Sloan Digital Sky Survey (SDSS)(York et al. 2000). 2DF has obtained spectroscopy of more than 20,000 quasars (Croom et al. 2004), and SDSS has identified more than 100,000 quasars (Schneider et al. 2007; Abazajian et al. 2009). Quasar candidates in these surveys were mainly selected based on their optical colours derived from the photometric data. 2DF mainly selected lower redshift ($z<2.2$) quasars with UV-excess ($u-b_J<$-0.36)(Smith et al. 2005), while SDSS adopted a multi-band optical colour selection method for quasar selection mainly by excluding the point sources within the stellar locus of the colour-colour diagrams (Richards et al. 2002). 90\\% of SDSS quasars have lower redshift ($z<2.3$), although some dedicated methods were also developed for discovering high redshift quasars (Fan et al. 2001a,b; Richards et al. 2002). The lower efficiency in identifying quasars with redshift between 2 and 3 is obvious in SDSS (Schneider et al. 2007), because these quasars usually have similar optical colours as stars and thus are mostly excluded by the SDSS quasar candidate selection algorithm (Warren et al. 2000).  In order to obtain a more complete sample of quasars, obviously we need to improve the previous quasar candidate selection methods. An important way to do this has been suggested by using the infrared K-band excess to identify the `missing' quasars with redshift around 2.7 based on the UKIRT (UK Infrared Telescope) Infrared Deep Sky Survey (UKIDSS) (Warren et al. 2000; Hewett et al. 2006; Maddox et al. 2008).\\footnote{The UKIDSS project is defined in Lawrence et al. (2007). UKIDSS uses the UKIRT Wide Field Camera (WFCAM; Casali et al. 2007) and a photometric system described in Hewett et al. 2006. The pipeline processing and science archive are described in Hambly et al. (2008).} Although the $z\\sim$2.7 quasars have similar optical colours as stars, they are more luminous in the infrared K-band. In addition, combining the optical colours in SDSS with the infrared colours in UKIDSS, it should be more efficient to separate stars from  both lower redshift quasars ($z<$3)(Chiu et al. 2007) and higher redshift ones ($z>$6) (Hewett et al. 2006) in the colour-colour diagrams. This advantage needs to be confirmed by future spectroscopic observations on the larger sample of quasar candidates selected with the SDSS and UKIDSS photometric colours.  Multi-band colours are also crucial for the photometric redshift estimations of quasars. Based on the SDSS colour-z relations, about 70\\% of the photometric redshifts of quasars are consistent with the spectroscopic ones with the difference $|\\Delta z|<0.2$ (Richards et al. 2001; Wu et al. 2004; Weinstein et al. 2004). This percentage can increase to 85\\% or higher if the extra GALEX UV photometric data (Ball et al. 2007) or Spitzer IRAC data (Richards et al. 2009b) are added. It has been demonstrated that the photo-z accuracy should be improved if combining the UKIDSS infrared photometry with the SDSS optical photometry for galaxies (Maddox et al. 2008). Such improvement can be also expected for quasar candidates with both SDSS and UKIDSS photometric data, but has not been confirmed by the photometric redshift estimations on a large SDSS-UKIDSS quasar sample so far. In addition, the reliable photometric redshift estimations based on multi-band colours are also important for preparing the quasar candidates for the future spectroscopic surveys (Richards et al. 2009a).  Large quasar survey is one of the key projects of the Chinese Large Sky Area Multi-Object Fibre Spectroscopic Telescope (LAMOST), which is a novel reflecting Schmidt telescope with 4 meter effective mirror size, 20 square degree field of view (FOV) and 4000 fibres (Su et al. 1998). As the most efficient optical spectroscopic telescope in the world, LAMOST finished its main construction in 2008 and has entered the commissioning phase. A pilot survey and the regular survey have been planned in 2010 and 2011-2015 respectively. Unlike SDSS, LAMOST does not have its own photometric survey. Therefore, the input catalogue of LAMOST quasar survey will largely rely on the photometric data from other existing surveys. Because of the large overlap of the surveyed area between SDSS and UKIDSS, a combination of SDSS and UKIDSS photometric data is expected to help us to efficiently select a large catalogue of quasar candidates and provide reliable photometric redshifts for the LAMOST quasar survey.  In this study, we first present a large quasar sample with both SDSS and UKIDSS data, and then investigate whether there is a more efficient selection criterion for quasars and whether the combination of UKIDSS data can improve the photometric redshift estimations. As two tests, the results are applied to a small UKIDSS/EDR quasar/star sample and a sample of the variable sources in the SDSS stripe 82 region. The implications of our study to the future LAMOST quasar survey are also discussed.  \\begin{table*} \\caption{Parameters of 8498 SDSS-UKIDSS quasars} \\centering \\begin{scriptsize} \\begin{tabular}{ccccccccccccc}\\\\ \\hline SDSS & SDSS & Redshift & SDSS-UKIDSS & $u$ & $g$ & $r$ & $i$ & $z$ & $Y$ & $J$ & $H$& $K$\\\\ RA & Dec & &Offset($''$)&&&&&&&&&\\\\ \\hline 0.0498394 &  0.0403587 &     0.479 &     0.1452 &       17.10 &       17.88  &   17.70     &     17.41  &   17.20   &  16.97 &  16.70 &  15.96  & 15.03 \\\\ 0.1625516&  -0.3010664 &      2.125  &    0.2016  &     18.16   &    19.04 &   18.64   &   18.44  &       18.21 &    18.07 &    17.76  &  17.43  &    16.75  \\\\ 0.2212764&  -0.6201625 &      1.321  &    0.3712   &     17.80  &    18.59  &    18.31   &    18.10  &     18.03  &    17.94 &   17.82  &   17.41 &   16.81  \\\\ 0.2371434 & -1.0693223  &     2.106  &    0.2559  &     19.36 &     20.09  &     19.36   &     18.80  &     18.24 &     18.28 &    18.02  &      17.35   &    16.52  \\\\ 0.2426759 & -0.7795238 &      1.897  &    0.6057 &      17.75 &     18.82  &     18.64   &    18.12   &      17.85 &   17.72  &    17.60  &   17.19 &  16.58 \\\\ \\hline \\end{tabular}\\\\ Note: All magnitudes are in Vega system. The full table of 8498 quasars is available in the electronic version of the paper. \\end{scriptsize} \\end{table*} ",
        "conclusions": "Quasar candidate selections largely rely on the multi-band photometric colours. We have demonstrated that combining the UKIDSS $YJHK$ data with SDSS $ugriz$ data will lead to significant improvements in both quasar selection efficiency and photometric redshift accuracy. With our quasar selection criteria involving both SDSS and UKIDSS colours, we can probably select more $z>2$ quasar candidates than SDSS even at the same spectroscopic magnitude limit. We believe that at least some SDSS missing quasars with $z\\sim2.7$ can be recovered with our selection criteria. Deeper spectroscopic observations on the quasar candidates selected with our criteria are expected to confirm our results and produce a more complete and much larger quasar sample in the future.  The quasar selection criteria we proposed in this paper are fairly simple, involving only cuts in the two-colour diagrams for optical point sources. More complicated approaches, such as the multi-dimensional search, would be more optimal. To check this, we have tried several different ways to search the criteria for separating quasars and stars, including 3-dimensional colour-colour diagrams, support vector machine (SVM) and Bayesian classification algorithm, but we did not find significant improvements. The detailed comparisons of the results by adopting these different approaches with our current work will be presented in a separated paper.  The study presented in this paper is mainly based on the SDSS-UKIDSS sample of 8498 quasars. Although the cross-identifications were done by finding the closest UKIDSS counterparts of the SDSS quasars and the offsets between the SDSS and UKIDSS positions are mostly within 0.5$''$, the mis-identifications of some sources are unavoidable. However, by checking the colour-z relations and colour-colour diagrams of these SDSS-UKIDSS quasars, we believe that the fraction of such mis-identifications should be very low. Therefore, the main results obtained in this paper will not be affected by the mis-identifications of some sources.  Because there is no quasar with $z>5.3$ in our SDSS-UKIDSS quasar sample, our quasar selection criteria can only be applied to identify $z<5.3$ quasars. Although we present a criterion in the $Y-J$ vs $i-z$ diagram for finding $z>5.6$ quasars based on the colour-z relation given by Hewett et al. (2006), this criterion obviously needs to be tested with a large sample of $z>5.6$ quasars with both SDSS and UKIDSS data. Moreover, for the high-z quasars with $z>5$, the photometric redshift estimation will largely rely on the UKIDSS colours because both $u-g$ and $g-r$ colours are not available. We also need a sample of $z>5$ quasars with both SDSS and UKIDSS photometric data to check the accuracy of our photometric redshift estimations. This is expected to be done in the future when a larger high-z quasar sample is available.  Although we made two tests using a UKIDSS/EDR sample and a sample of variable sources in SDSS Stripe 82 to check the robustness of our quasar selection criteria and the photometric redshift method, obviously we still need more tests. We have selected some quasar candidates with $i_{\\rm{AB}}<19.1$ using our criteria from the variable sources in Stripe 82, which do not have SDSS spectroscopy, and will identify them with optical spectroscopy. This will provide direct evidence for whether using our SDSS/UKIDSS selection criteria can help to find missing quasars in SDSS. On the other hand, finding quasars with certain criteria and comparing them with those in the existing spectroscopic surveys may not be enough for testing the completeness, as most known quasars were identified with the colour selection methods.  A better way to do it may be to run simulated colours of quasars, with assumptions on their continuum, emission line and reddening, for different redshift and luminosity (Fan 1999), and check how many simulated quasars can be found with the selection criteria. This will be studied in our future work.  The Chinese LAMOST quasar survey will be done in the next a few years with the most efficient spectroscopic telescope in the world, and is expected to obtain a much larger quasar sample (with the magnitude limit of $i$=20.5 or 21) than the previous surveys. The results of our current study, if confirmed with further spectroscopic observations, will be helpful for preparing the input catalog of quasar candidates and estimating their photometric redshifts for the LAMOST quasar survey, especially in the overlapped sky area between SDSS and UKIDSS. By doing necessary tests of various selection criteria with the spectroscopic observations in the pilot survey, we will determine the best criteria for quasar candidates selections and apply them to the LAMOST quasar survey. Although the UKIDSS/LAS will cover only a sky area of 4000deg$^2$, with the near-IR colours we will be able to select more complete quasar samples with redshift up to 5 than previous surveys. This large and complete quasar sample will be very important to many further studies such as those on the quasar luminosity function, quasar clustering, and large scale structure in the universe."
    },
    "1004/1004.3086_arXiv.txt": {
        "abstract": "We use current and future simulated data of the growth rate of large scale structure in combination with data from supernova, BAO, and CMB surface measurements, in order to put constraints on the growth index parameters. We use a recently proposed parameterization of the growth index that interpolates between a constant value at high redshifts and a form that accounts for redshift dependencies at small redshifts. We also suggest here another exponential parameterization with a similar behaviour. The redshift dependent parametrizations provide a sub-percent precision level to the numerical growth function, for the full redshift range. Using these redshift parameterizations or a constant growth index, we find that current available data from galaxy redshift distortions and Lyman-alpha forests is unable to put significant constraints on any of the growth parameters. For example both $\\Lambda$CDM and flat DGP are allowed by current growth data. We use an MCMC analysis to study constraints from future growth data, and simulate pessimistic and moderate scenarios for the uncertainties. In both scenarios, the redshift parameterizations discussed are able to provide significant constraints and rule out models when incorrectly assumed in the analysis. The values taken by the constant part of the parameterizations as well as the redshift slopes are all found to significantly rule out an incorrect background. We also find that, for our pessimistic scenario, an assumed constant growth index over the full redshift range is unable to rule out incorrect models in all cases.  This is due to the fact that the slope acts as a second discriminator at smaller redshifts and therefore provide a significant test to identify the underlying gravity theory. ",
        "introduction": "Since its discovery over a decade ago \\cite{acc1a}, cosmic acceleration stands as one of the most important and challenging problems in all physics, see for example the reviews \\cite{reviews} and references therein. As discussed in these reviews and others, cosmic acceleration can be caused by the presence of a dark energy component in the universe or alternatively a modification of gravity physics (namely General Relativity) at cosmological scales. A significant step toward the understanding of the cause of cosmic acceleration is to be able to distinguish between the two competing alternatives.  Indeed, one ongoing approach to understand the origin of cosmic acceleration is to constrain the equation of state of dark energy \\cite{reviews} while other approaches rely on comparisons of the cosmic expansion history to the growth rate of large scale structure, see the incomplete list  \\cite{lue,Aquaviva,gong08b,polarski,linder,Koyama,Koivisto,Daniel,knox,ishak2006,laszlo,Zhang,Hu}. Namely, it was shown in many of these references and others that two gravitational theories can have very degenerate Hubble curves but yet have distinct functions of the growth rate of large scale structure in the universe.  The growth rate can thus be used in order to constrain the underlying gravity theory and there has been much interest in providing a parameterization of the growth factor function with one or two parameters that are distinct and characteristic for a given gravity theory, again see the partial list \\cite{lue,Aquaviva,gong08b,polarski,linder,Koyama,Koivisto,Daniel,knox,ishak2006,laszlo,Zhang,Hu}.  In this paper, we compare parameterizations of the growth index as a function of the redshift to current and future observations. We also introduce a new exponential parameterization for the growth parameter which allows us to easily characterize the asymptotic value of the growth index. For current constraints, we use the growth data (mainly from galaxy redshift distortions and Lyman-alpha forests) from \\cite{porto,ness,guzzo,colless,tegmark,ross,angela,mcdonald,viel1,viel2}, the Constitution compilation of supernova data sets \\cite{Constitution}, baryon acoustic oscillation (BAO) measurement from the Sloan Digital Sky Survey (SDSS) \\cite{sdss6}, and the distance to the surface of last scattering of the CMB as measured from the Wilkinson Microwave Anisotropy Probe 5 yr data (WMAP5) \\cite{WMAP5}. We also explore, using a Monte-Carlo-Markov-Chain (MCMC) analysis, how well future growth data from galaxy redshift distortions and Lyman-alpha forests will be able to constrain the growth index parameters and to rule out incorrectly assumed underlying gravity theories. ",
        "conclusions": "We explored comparisons of redshift parameterizations of the growth factor index to current and future growth data. The first parametrization used was introduced in previous work and interpolates between a redshift dependent form at small redshifts and a constant value at high redshifts. A second parametrization based on an exponential form is introduced here and exhibits a similar redshift dependence, as it should. We found it to fit theoretical data to within $0.015\\%$ for $\\Lambda$CDM and $0.09\\%$ for DGP, over the entire redshift range up to the CMB surface. While more precise parametrizations are welcome, we consider that the more significant plus in these redshift dependent parametrizations is that they provide the slope of the growth index as a second test to the underlying gravity model. This is relevant because the slope is related to variations in $\\gamma(z)$ at small redshifts where more data can be obtained. Using redshift dependent parameterizations or constant value of the growth index, we find that current growth data from redshift distortions and Lyman alpha forests is unable to put significant constraints on the growth parameters. In order to explore how well future growth data could constrain these parameters, we  simulated growth data and ran a Monte-Carlo-Markov-Chain analysis. We find that a pessimistic or moderate scenarios for future data uncertainties will be able to rule out an incorrectly assumed theoretical model using any of the two parameterizations discussed while we find that in our pessimistic scenario a constant growth index parameter will be unable to rule out an incorrect model. This is due to the fact that the slope acts as a second discriminator at smaller redshifts."
    },
    "1004/1004.2049_arXiv.txt": {
        "abstract": "The Galactic disc is opaque to radio waves from extragalactic sources with frequencies $\\nu$ less than $\\sim 3\\ \\MHz$.  However, radio waves with kHz, Hz, and even lower frequencies may propagate through the intergalactic medium (IGM).  I argue that the presence of these waves can be inferred by using the Universe as our detector.  I discuss possible sub-MHz sources and set new non-trivial upper limits on the energy density of sub-MHz radio waves in galaxy clusters and the average cosmic background.  Limits based on five effects are considered: (1) changes in the expansion of the Universe from the radiation energy density (2) heating of the IGM by free-free absorption; (3) radiation pressure squeezing of IGM clouds by external radio waves; (4) synchrotron heating of electrons in clusters; and (5) Inverse Compton upscattering of sub-MHz radio photons.  Any sub-MHz background must have an energy density much smaller than the CMB at frequencies below 1 MHz.  The free-free absorption bounds from the Lyman-$\\alpha$ forest are potentially the strongest, but are highly dependent on the properties of sub-MHz radio scattering in the IGM.  I estimate an upper limit of $6 \\times 10^4\\ \\Lsun\\ \\Mpc^{-3}$ for the emissivity within Lyman-$\\alpha$ forest clouds in the frequency range $5 - 200$ Hz.  The sub-MHz energy density in the Coma cluster is constrained to be less than $\\sim 10^{-15}\\ \\ergcm3$.  At present, none of the limits is strong enough to rule out a maximal $T_b = 10^{12}\\ \\Kelv$ sub-MHz synchrotron background, but other sources may be constrained with a better knowledge of sub-MHz radio propagation in the IGM. ",
        "introduction": "Whenever a new wavelength window has been opened on the electromagnetic spectrum, it has led to new discoveries \\citep{Harwit81,Lawrence07}.  In the past century, almost the entire electromagnetic spectrum has been explored, with few gaps between $10~\\MHz$ and $100~\\TeV$.  The cosmic electromagnetic backgrounds at most of these frequencies\\footnote{Although the extragalactic Extreme Ultraviolet (EUV; $100 - 912$ \\AA) radiation is easily absorbed by neutral hydrogen, EUV has been detected from AGNs and galaxy clusters at wavelengths of $\\la 160\\ {\\rm \\AA}$ (see the EUV review by \\citealt*{Bowyer00}).  Closer to the Lyman limit, not even a Galactic background has been detected yet \\citep{Edelstein01}, and the Galactic neutral hydrogen would much more effectively block incoming EUV radiation.  However, the extragalactic ionizing background at a variety of redshifts is indirectly measured by the ionization state of Lyman-$\\alpha$ forest clouds (e.g., \\citealt*{Bajtlik88}; \\citealt{Shull99}).} have now been measured \\citep[as reviewed by][]{Ressell90,Fukugita04,Trimble06}.  Searches are underway for very-high energy photons, including those with PeV \\citep{Chantell97,Borione98,Schatz03} and even EeV energies \\citep[e.g.,][]{Auger08}.  These searches face several challenges.  At PeV energies, extragalactic searches are hampered by $\\gamma\\gamma$ opacity, in which PeV photons interact with the CMB to produce $e^+e^-$ pairs, and the Universe is highly opaque at these energies \\citep*{Moskalenko06}.  EeV photons are predicted to exist (e.g., \\citealt*{Wdowczyk72,Gelmini08}), and the Universe is possibly transparent to several Mpc at these energies \\citep{Protheroe96}, but very low number statistics are a problem.  Finally, above $10^{24} \\eV$, the Galaxy becomes completely opaque, as single photons pair produce $e^+e^-$ off the Galactic magnetic field \\citep{Stecker03}.  The only other frontier in terms of photon energy is at the other end of the electromagnetic spectrum, the lowest frequency radio waves.  The Universe is filled with tenuous plasma, the intergalactic medium (IGM), which prevents radio propagation below the plasma frequency $\\nu_P = \\sqrt{n_e e^2 / (\\pi m_e)}$.  Waves below this frequency evanesce within one wavelength, reflecting off the medium.  The mean baryonic density of the Universe is $\\mean{n_b} \\approx 2.5 \\times 10^{-7} (1 + z)^3~\\cm^{-3}$, although regions in the IGM may have higher or lower densities.  For IGM with electron density $\\delta_e \\mean{n_b}$, the plasma frequency is \\begin{equation} \\label{eqn:nuP} \\nu_P = 4.5 \\delta_e^{1/2} (1 + z)^{3/2}~\\Hz. \\end{equation} The density distribution in the Universe can be approximated as a lognormal distribution, with most of the volume being relatively empty \\citep{Coles91}.  Roughly $90\\%$ of the Universe's volume at $z = 0$ is predicted to have $\\delta_e \\ga 0.002$ \\citep{Bi97}, which corresponds to a plasma frequency of $\\nu_{P} = 0.2~\\Hz$.  This is essentially the low end of the cosmic electromagnetic spectrum, below which no electromagnetic wave can ever travel.\\footnote{There is a loophole, since electromagnetic waves below the plasma frequency will evanesce with a scale of about one wavelength: if the electromagnetic wave has a cosmologically long wavelength, it can stretch over the entire Universe.  Hawking radiation from the accelerating expansion of the Universe actually does have cosmological wavelengths \\citep[e.g.,][]{Gibbons77}, but the energy density of Hawking radiation is negligible.}  Although there might be electromagnetic radiation with $\\nu \\la \\Hz$ in the Universe, direct observations at these frequencies are impossible.  Observations down to $10\\ \\MHz$ will soon be routine with LOFAR \\citep{Rottgering03}.  However, the Earth's ionosphere has a typical plasma frequency of $10\\ \\MHz$, depending on the time of day and other conditions.  Similarly, the Moon may have an ionosphere with a plasma frequency of a few hundred kHz, which sets a lower limit to the frequency of lunar-based observatories \\citep{Jester09}.  Satellites have measured the Galactic radio emission down to $\\sim 100\\ \\kHz$ \\citep{Brown73,Novaco78}.  Space-based observatories near the Earth, such as the proposed ALFA \\citep{Jones00a,Jones00b}, can potentially observe down to $30\\ \\kHz$, which is the plasma frequency of the Solar wind at Earth's orbit.  Further out from the Sun, the plasma frequency continues to drop, and the \\emph{Voyager} probes took advantage of this to detect kHz emission \\citep{Kurth84}.  However, the interstellar medium itself has a plasma frequency of 2 kHz, which serves as a hard limit for direct observations of Galactic and extragalactic sources.  The situation is even worse for extragalactic and distant galactic sources.  The warm ionized medium (WIM) forms a disc with scale height $h \\approx 1~\\kpc$ and a typical density of $n_{\\rm WIM} \\approx 0.01~\\cm$.  The WIM is opaque to low frequency radio emission because of free-free absorption.  Although this absorption has uses for tomography of the Galactic ISM \\citep{Peterson02}, it prevents all direct extragalactic observations at $\\nu \\la 3\\ \\MHz$.  There may be a few low density `chimneys' that allow in some lower frequency radio emission \\citep{Jester09}, but for practical purposes, the extragalactic sky at the lowest radio frequencies will be shrouded from our direct view for the foreseeable future.  But there are ways around this limit to locate or place bounds on \\emph{extragalactic} sub-MHz radio sources.  Free-free absorption and other opacity sources become more effective at low frequency; this is why they prevent sub-MHz radiation from reaching Earth.  Yet, this also means that the lowest frequency radio waves are tightly coupled with the matter in the Universe.  Observations of intergalactic matter therefore constrain these radio waves.  Synchrotron absorption and Inverse Compton scattering also place bounds on sub-MHz radio in regions with cosmic rays and magnetic fields, like galaxy clusters.  I will argue that sub-MHz radio emission does not need to reach us, because \\emph{the Universe is our detector}.  I will first discuss our expectations for the sub-MHz sky (\\S~\\ref{sec:SubMHzSky}), including postulated sub-MHz radio sources (\\S~\\ref{sec:SubMHzSources}), the IGM phases that can interact with sub-MHz radio waves (\\S~\\ref{sec:IGMPhases}), and sub-MHz radio propagation through the IGM (\\S~\\ref{sec:IGMPropagation}).  I then set new limits on the extragalactic background at sub-MHz frequencies.  The first limit I consider is the weak bound from the expansion history of the Universe (\\S~\\ref{sec:OmegaR}).  The second bound is from the heating of the IGM by free-free absorption of sub-MHz radio waves (\\S~\\ref{sec:FFAbsorption}).  A third bound on incident radiation at the lowest frequencies comes from the radiation force exerted on IGM clouds (\\S~\\ref{sec:CloudCrush}).  Two more bounds can be set for clusters with cosmic rays: bounds on synchrotron heating by low frequency radio waves (\\S~\\ref{sec:SynchHeat}) and bounds on Inverse Compton upscattered radio waves (\\S~\\ref{sec:ICBound}).  Finally, if extragalactic photons with energy $\\ga 10^{20} \\eV$ are ever detected, they will set extremely strong limits on the kHz to MHz radio background (\\S~\\ref{sec:UHELimits}).  My goal throughout this paper is to derive upper bounds on the sub-MHz emission with order of magnitude accuracy, where possible with current knowledge.  In some cases even this is not possible -- the radio scattering properties of the IGM over large distances are not well known, and this is a huge source of uncertainty in arguments that rely on radiative transfer.  Future theoretical work may reduce these uncertainties.  Throughout this work, I consider the time-averaged sub-MHz background at $z \\approx 0$, though these methods may be applied to other redshifts. ",
        "conclusions": "\\label{sec:Conclusion} \\begin{figure*} \\centerline{\\includegraphics[width=18cm]{f4.eps}} \\caption{The $z = 0$ cosmic backgrounds for the electromagnetic spectrum.  The cosmic plasma frequency (assuming $\\delta_e \\approx 0.002$) makes propagation impossible at the lowest frequencies (\\emph{dark grey}); we expect there to be \\emph{no} radio background below this cutoff.  The $\\Omega_R$ bound \\citep{Zentner02} is in grey at top (\\S~\\ref{sec:OmegaR}).  Free-free absorption bounds (\\S~\\ref{sec:FFAbsorption}) from Ly$\\alpha$ clouds of $\\delta = 0.1$ and $1.0$ and from a weak Mg II/C IV absorber are shown in pink/red, assuming $J_{\\nu}$ is a bump at $\\nu$.  The solid red bounds are for the case when the radiation fills the cloud evenly.  The shaded regions are bounds on an incident external radiation field; lighter shading (\\emph{dashed boundaries}) for no scattering in the cloud, while darker shading (\\emph{dotted boundaries}) when eq.~\\ref{eqn:mfpScattering} describe the scattering. Radiation pressure can crush a cloud (\\S~\\ref{sec:CloudCrush}); the bounds on cloud crushing are shown in brown.  The solid line and shading assumes there is no scattering, while the dotted line assumes that eq.~\\ref{eqn:mfpScattering} describe the scattering.  The radiation pressure bounds are bolometric below $\\nu_{\\rm int}$. I also plot naive upper expectations on the radio background expected from several sources (\\S~\\ref{sec:SubMHzSources}), not accounting for IGM absorption: the maximum synchrotron brightness temperature ($T_b \\approx 10^{12}\\ \\Kelv$), pulsars, and gravitational wave conversion. See Table~\\ref{table:BigFigKey} for a full legend with references. \\label{fig:EMBackgrounds}} \\end{figure*}  The extragalactic sub-MHz background is invisible to direct observation from Earth, but we can still detect its effects on intergalactic matter.  I have placed new limits on the magnitude of the sub-MHz radio background, using various IGM phases and clusters as radio detectors.  Figure~\\ref{fig:EMBackgrounds} (full legend in Table~\\ref{table:BigFigKey}) summarises the bounds on the radio background from the IGM thermal state (\\S~\\ref{sec:FFUBound}; \\emph{red}) and the radiation pressure exerted on the IGM  (\\S~\\ref{sec:CloudCrush}; \\emph{brown}).  A sub-MHz background with an energy density as large as the CMB at any frequency is easily ruled out, and energy densities comparable to the cosmic starlight backgrounds are also not allowed at almost all frequencies.  \\begin{table*} \\begin{minipage}{140mm} \\caption{Legend for Figure~\\ref{fig:EMBackgrounds}.} \\begin{tabular}{llll} \\hline Wavelength band & Method/Instrument & Reference & Symbol\\\\ \\hline All & $\\Omega_R$ & \\citet{Zentner02} & \\emph{light grey shading}\\\\ Sub-MHz radio & IGM pressure & This work (\\S~\\ref{sec:CloudCrush}) & \\emph{brown shading}\\\\ & IGM thermal state & This work (\\S~\\ref{sec:FFUBound}) & \\emph{pink shading}\\\\ Radio & Theoretical prediction & \\citet{Protheroe96} & \\emph{dotted black lines}\\\\ & ARCADE2 & \\citet{Fixsen09} & \\emph{open triangles}\\\\ IR to UV & Theoretical prediction & \\citet*{Franceshini08} & \\emph{solid black line} \\\\ Infrared & FIRAS & \\citet{Fixsen98} & \\emph{orange dotted line}\\\\ & BLAST & \\citet{Marsden09} & \\emph{orange circles}\\\\ & DIRBE & \\citet{Wright04} & \\emph{orange pentagons}\\\\ & \\emph{Spitzer} & \\citet{Dole06} & \\emph{green lower limits}\\\\ & \\emph{Spitzer} & \\citet{Papovich04} & \\emph{orange triangle}\\\\ & \\emph{Spitzer} & \\citet{Savage05} & \\emph{green squares}\\\\ Optical & Galaxy counts with \\emph{Hubble} & \\citet{Madau00} & \\emph{filled triangles} \\\\ UV & GALEX & \\citet{Xu05} & \\emph{violet triangles}\\\\ EUV & Ly$\\alpha$ forest ionization & \\citet{Shull99} & \\emph{open circle}\\\\ X-rays & \\emph{XMM-Newton} (Lockman hole) & \\citet{Worsley05} & \\emph{grey 6-stars}\\\\ & \\emph{Chandra} & \\citet{Hickox06} & \\emph{5-stars}\\\\ & \\emph{Swift} & \\citet{Moretti09} & \\emph{dotted blue line}\\\\ & \\emph{Swift} & \\citet{Ajello08} & \\emph{solid blue line}\\\\ & RXTE & \\citet{Revnivtsev03} & \\emph{solid green line}\\\\ & HEAO1 & \\citet{Kinzer97} & \\emph{grey solid line}\\\\ MeV $\\gamma$-rays & SMM & \\citet{Watanabe00} & \\emph{solid cyan line}\\\\ & COMPTEL & \\citet{Weidenspointner00} & \\emph{crosses}\\\\ GeV $\\gamma$-rays & EGRET & \\citet*{Strong04} & \\emph{open squares}\\\\ $\\ge \\TeV$ $\\gamma$-rays & GeV background & \\citet{Coppi97} & \\emph{violet shading}\\\\ TeV $\\gamma$-rays & HESS & \\citet{HESS09} & \\emph{solid black line}\\\\ & HESS & \\citet{Aharonian08} & \\emph{solid violet line}\\\\ & GRAPES-3 & \\citet{Hayashi03} & \\emph{blue arrows}\\\\ PeV $\\gamma$-rays & CASA-MIA & \\citet{Chantell97} & \\emph{grey arrows}\\\\ EeV $\\gamma$-rays & Auger & \\citet{Auger09} & \\emph{black arrows}\\\\ & Auger & \\citet{Abraham08} & \\emph{black arrows}\\\\ \\hline \\end{tabular} \\label{table:BigFigKey} \\\\The PeV $\\gamma$-ray limits do not correct for pair-production absorption from the CMB.\\\\ \\end{minipage} \\end{table*}  Low frequency radio waves can heat the IGM through free-free absorption.  Observations of the IGM thermal state constrain the amount of heating from the extragalactic sub-MHz radio background.  Free-free absorption bounds (\\S~\\ref{sec:FFUBound}) are potentially the strongest of all of the limits, but are highly model dependent at low frequencies.  If we are considering a radio bath that pervades the entire cloud evenly (\\emph{dashed red lines}), then the entire cloud is heated up and the bounds are very strong at low frequency in terms of energy density.  In fact, the energy density constraints within the Lyman-$\\alpha$ forest just above its plasma frequency would be the strongest of any photon energy, as seen in Figure~\\ref{fig:EMBackgrounds}.  If a background was simply incident on the IGM clouds, then when the cloud becomes optically thick, the outside of the cloud will be heated but the interior will not.  In this case, if there is no scattering of radio waves within the cloud (\\emph{solid red lines, light pink shading}), the clouds usually remain optically thin down to sub-kHz frequency, and the energy density bounds remain strong.  However, scattering will increase the effective absorption optical depth.  Naively applying the scattering mean free path in equation~\\ref{eqn:mfpScattering} considerably weakens the energy density bounds (\\emph{dotted red lines, darker pink shading}).  However, the approximations in equation~\\ref{eqn:mfpScattering} may break down at high scattering optical depth (\\ref{sec:IGMPropagation}; \\citealt{Cohen74}).  I have also not considered the opacity of any remaining shell of `evaporated' material around an optically thick cloud.  In order to set firm limits from free-free absorption, we need to understand the radiative transfer of sub-MHz radio waves through the IGM better.  I have also set an upper limit on the luminosity density \\emph{within} each IGM phase from free-free absorption (\\S~\\ref{sec:FFEpsilonBound}), assuming the radio waves fill the IGM cloud evenly.  From the existence and temperature of the Lyman-$\\alpha$ forest, I infer that at frequencies of $5 - 100$ Hz, these clouds have a maximum emissivity of $6 \\times 10^4\\ \\Lsun\\ \\Mpc^{-3}$.  Again, a better understanding of the radiative transfer and the scattering in particular of sub-MHz radio waves is needed to set more firm limits on the luminosity density of the Universe at low frequencies.  At the lowest frequencies, there is a window in the free-free absorption constraints.  This is because voids are extremely underdense, with a very low plasma frequency.  Sub-Hz radio could be generated in the voids and would simply reflect off the more condensed structures that would otherwise be heated by them.  Free-free absorption constraints are weak because scattering or absorption of the radio waves shields the interior of each IGM cloud.  At the very lowest frequencies, below the plasma frequency of the cloud, the radio waves simply reflect off it.  There are none the less constraints even at these lowest frequencies, because the reflection, scattering, or absorption of these waves squeezes IGM clouds (\\S~\\ref{sec:CloudCrush}).  These radiation pressure bounds (Figure~\\ref{fig:EMBackgrounds}, \\emph{brown}) are somewhat weak but are still strong enough to rule out an average sub-Hz background as large as the starlight backgrounds.  Unlike the free-free absorption bounds, the maximum $u_{\\nu}$ are not model dependent.  However, the frequency range over which they are applicable also depends on the radio scattering properties of the IGM.  At the very least, the radiation pressure bounds apply until the cloud is optically thin to free-free absorption.  Galaxy clusters have hot and dense gas, which makes them poor free-free absorption detectors for any sub-MHz radio waves within them.  However, they also contain magnetic fields and cosmic rays, which can interact with sub-MHz radio waves in additional ways and provide additional limits (Figure~\\ref{fig:ClusterLimits}).  Low frequency radio waves above the Razin frequency ($\\sim 20\\ \\kHz$) can actually heat CR electrons/positrons, and would create a peak into the observed MHz to GHz synchrotron radio spectra of clusters (\\S~\\ref{sec:SynchHeat}).  The lack of such a peak rules out sub-MHz radio backgrounds as small as $\\sim 10^{-15} \\ergcm3$ in the Coma cluster, comparable to the energy density in starlight.  CR electrons/positrons can also Inverse Compton scatter low frequency radio waves to observable MHz to GHz frequencies (\\S~\\ref{sec:ICBound}).  The observed radio spectrum again constrains the sub-MHz radio background in Coma to be as small as $\\sim 10^{-15} \\ergcm3$ for $\\nu \\la \\kHz$.  Some relatively weak statements can be made about whether the backgrounds described in \\S~\\ref{sec:SubMHzSources} exist.  The most exotic sources of sub-MHz radio waves are constrained.  The radiation pressure bound from weak Mg II/C IV absorbers are sufficiently strong to exclude a radio background as large as the LIGO upper limits on a 100 Hz stochastic gravitational wave background.  A background of pulsar waves, if they somehow escaped into the IGM and did not suffer absorption, would be weaker still; the free-free absorption bounds without scattering are strong enough to rule out such backgrounds from all young pulsars and all MSPs in Galactic discs.  These maximum estimates of the sub-MHz radio background are probably unrealistic anyway (see the discussion in \\S~\\ref{sec:ExoticSources}).  Unfortunately, a more realistic synchrotron background also seems to be out of reach by the free-free absorption bounds; a $T_b \\le 10^{12} \\Kelv$ background is ruled out only for $\\nu \\ga \\MHz$, which can already be directly observed.  However, if the scattering properties of the IGM are similar to those in the ISM, such that eq.~\\ref{eqn:mfpScattering} holds, then the radiation pressure bounds from the Lyman-$\\alpha$ forest will rule out maximal synchrotron backgrounds at $\\sim 100\\ \\kHz$.  There are several ways to reduce the uncertainties in these bounds.  Knowledge of the low frequency scattering properties of the IGM is essential for the free-free absorption bounds, which would otherwise be strong (\\emph{solid red lines} in Figure~\\ref{fig:EMBackgrounds}).  This knowledge can also help us determine the frequency range the radiation pressure bounds apply over.  Strong scattering will \\emph{weaken} the free-free absorption bounds, because radiation cannot diffuse deep into the cloud and heat it; but strong scattering \\emph{strengthens} the radiation pressure bounds, because radiation can then efficiently couple with the cloud exterior and squeeze it.  Low frequency radio observations of galaxy clusters and other environments with CRs are especially useful for the Inverse Compton and synchrotron heating bounds.  Observations of extragalactic ultra high energy photons, if they exist, would set extremely strong constraints on the extragalactic sub-MHz background (\\S~\\ref{sec:UHELimits}), especially at high frequencies where the other bounds are weakest.  It would be a simple matter to apply similar constraints to sub-MHz emission within the Galaxy itself, which will be done in a future paper.  Not only are the thermal properties of the interstellar medium relatively well characterised, but the Galaxy has a well known CR electron spectrum and magnetic field.  Therefore, we could apply synchrotron heating and IC upscattering arguments, which are not dependent on the scattering of low frequency radio waves.  This would allow us to probe distant regions of the Galaxy that are not visible at low radio frequencies because of free-free absorption, such as the Galactic Centre, which is obscured by free-free absorption at frequencies as high as 330 MHz \\citep[e.g.,][]{Pedlar89}.  Although the bounds in this paper may not strongly constrain expected sources like a synchrotron background, the extragalactic sub-MHz sky is not \\emph{completely} unknowable.  Instead of disregarding the low frequency emission of radio sources, it is possible to consider the effects of the emission on their surroundings.  These effects may prove to be important to our understanding of the regions around sub-MHz sources.  Even if the extragalactic sub-MHz sky is forever invisible to us directly, its presence can still be seen."
    },
    "1004/1004.2755_arXiv.txt": {
        "abstract": "We present the Sloan Low-mass Wide Pairs of Kinematically Equivalent Stars ({\\slowpokes}), a catalog of 1342 very-wide (projected separation $\\gtrsim 500$~AU), low-mass (at least one mid-K -- mid-M dwarf component) common proper motion pairs identified from astrometry, photometry, and proper motions in the  Sloan Digital Sky Survey. A Monte Carlo based Galactic model is constructed to assess the probability of chance alignment for each pair; only pairs with a probability of chance alignment $\\leq 0.05$ are included in the catalog. The overall fidelity of the catalog is expected to be 98.35\\%. The selection algorithm is purposely exclusive to ensure that the resulting catalog is efficient for follow-up studies of low-mass pairs. The {\\slowpokes} catalog is the largest sample of wide, low-mass pairs to date and is intended as an ongoing community resource for detailed study of {\\em bona fide} systems. Here we summarize the general characteristics of the {\\slowpokes} sample and present preliminary results describing the properties of wide, low-mass pairs.  While the majority of the identified pairs are disk dwarfs, there are 70 halo subdwarf pairs and 21 white dwarf--disk dwarf pairs, as well as four triples. Most {\\slowpokes} pairs violate the previously defined empirical limits for maximum angular separation or binding energies. However, they are well within the theoretical limits and should prove very useful in putting firm constraints on the maximum size of binary systems and on different formation scenarios. We find a lower limit to the wide binary frequency for the mid-K -- mid-M spectral types that constitute our sample to be 1.1\\%. This frequency decreases as a function of Galactic height, indicating a time evolution of the wide binary frequency. In addition, the semi-major axes of the {\\slowpokes} systems exhibit a distinctly bimodal distribution, with a break at separations around 0.1~pc that is also manifested in the system binding energy. Comparing with theoretical predictions for the disruption of binary systems with time, we conclude that the {\\slowpokes} sample comprises two populations of wide binaries: an ``old\" population of tightly bound systems, and a ``young\" population of weakly bound systems that will not survive more than a few Gyr.  The {\\slowpokes} catalog and future ancillary data are publicly available on the world wide web for utilization by the astronomy community. ",
        "introduction": "\\label{Sec: intro}  The formation and evolution of binary stars remains one of the key unanswered questions in stellar astronomy. As most stars are thought to form in multiple systems, and with the possibility that binaries may host exoplanet systems, these questions are of even more importance. While accurate measurements of the fundamental properties of binary systems provide constraints on evolutionary models \\citep[e.g.][]{Stassun2007}, knowing the binary frequency, as well as the distribution of the periods, separations, mass ratios, and eccentricities of a large ensemble of binary systems are critical to understanding binary formation \\citep[][and references therein]{Goodwin2007}. To date, multiplicity has been most extensively studied for the relatively bright high- and solar-mass local field populations \\citep[e.g.][hereafter \\citetalias{Duquennoy1991}]{Duquennoy1991}. Similar studies of low-mass M and L dwarfs have been limited by the lack of statistically significant samples due to their intrinsic faintness. However, M dwarfs constitute $\\sim$70\\% of Milky Way's stellar population \\citep[][hereafter \\citetalias{Bochanski2010}]{Miller1979, Henry1999, Reid2002, Bochanski2010} and significantly influence its properties.  Since the pioneering study of \\citet{Heintz1969}, binarity has been observed to decrease as a function of mass: the fraction of primaries with companions drops from 75\\% for OB stars in clusters \\citep{Gies1987, Mason1998, Mason2009} to $\\sim$60\\% for solar-mass stars (\\citealt{Abt1976}, \\citetalias{Duquennoy1991}, \\citealt{Halbwachs2003}) to $\\sim$30--40\\% for M dwarfs (\\citealt[][hereafter \\citetalias{Fischer1992}]{Fischer1992}; \\citealt{Henry1993, Reid1997, Delfosse2004}) to $\\sim$15\\% for brown dwarfs \\citep[BDs;][]{Bouy2003, Close2003, Gizis2003, Martin2003}. This decrease in binarity with mass is probably a result of preferential destruction of lower binding energy systems over time by dynamical interactions with other stars and molecular clouds, rather than a true representation of the multiplicity at birth \\citep{Goodwin2005}. In addition to having a smaller total mass, lower-mass stars have longer main-sequence (MS) lifetimes \\citep*{Laughlin1997} and, as an ensemble, have lived longer and been more affected by dynamical interactions. Hence, they are more susceptible to disruption over their lifetime. Studies of young stellar populations (e.g.\\ in Taurus, Ophiucus, Chameleon) appear to support this argument, as their multiplicity is twice as high as that in the field \\citep{Leinert1993, Ghez1997, Kohler1998}. However, in denser star-forming regions in the Orion Nebula Cluster and IC 348, where more dynamical interactions are expected, the multiplicity is comparable to the field (\\citealt*{Simon1999}; \\citealt{Petr1998}; \\citealt*{Duchene1999b}). Hence, preferential destruction is likely to play an important role in the evolution of binary systems.  \\citetalias{Duquennoy1991} found that the physical separation of the binaries could be described by a log-normal distribution, with the peak at $a \\sim$30~AU and $\\sigma_{\\log a} \\sim$1.5 for F and G dwarfs in the local neighborhood. The M dwarfs in the local 20-pc sample of \\citetalias{Fischer1992} seem to follow a similar distribution with a peak at $a \\sim$3--30~AU, a result severely limited by the small number of binaries in the sample.  Importantly, both of these results suggest the existence of very wide systems, separated in some cases by more than a parsec. Among the nearby ($d<$ 100~pc) solar-type stars in the {\\em Hipparcos} catalog, \\citet*{Lepine2007a} found that 9.5\\% have companions with projected orbital separations $s>$ 1000~AU. However, we do not have a firm handle on the widest binary that can be formed or on how they are affected by localized Galactic potentials as they traverse the Galaxy. Hence, a sample of wide binaries, especially one that spans a large range of heliocentric distances, would help in (i) putting empirical constraints on the widest binary systems in the field \\citep[e.g.][]{Reid2001b, Burgasser2003, Burgasser2007b, Close2003, Close2007}, (ii) understanding the evolution of wide binaries over time \\citep*[e.g.,][]{Weinberg1987, Jiang2009}, and (iii) tracing the inhomogeneities in the Galactic potential (e.g.~\\citealt*{Bahcall1985} \\citealt{Weinberg1987}; \\citealt*{Yoo2004}; \\citealt{Quinn2009}).  Recent large scale surveys, such as the Sloan Digital Sky Survey \\citep[SDSS;][]{York2000}, the Two Micron All Sky Survey \\citep[2MASS;][]{Cutri2003}, and the UKIRT Infrared Deep Sky Survey \\citep[UKIDSS;][]{Lawrence2007}, have yielded samples of unprecedented numbers of low-mass stars. SDSS alone has a photometric catalog of more than 30 million low-mass dwarfs \\citepalias{Bochanski2010}, defined as mid-K -- late-M dwarfs for the rest of the paper and a spectroscopic catalog of more than 44000 M dwarfs \\citep{West2008}. The large astrometric and photometric catalogs of low-mass stars afford us the opportunity to explore anew the binary properties of the most numerous constituents of  Milky Way, particularly at the very widest binary separations.  The orbital periods of very wide binaries (orbital separation $a>$ 100~ AU) are much longer than the human timescale ($P=$ 1000~yr for \\Mtot $=$ 1\\Msun\\ and $a=$ 100~AU). Thus, these systems can only be identified astrometrically, accompanied by proper motion or radial velocity matching. These also remain some of the most under-explored low-mass systems. Without the benefit of retracing the binary orbit, two methods have been historically used to identify very wide pairs: \\begin{enumerate} \\item \\citet{Bahcall1981} used the two-point correlation method to argue that the excess of pairs found at small separations is a signature of physically associated pairs; binarity of some of these systems was later confirmed by radial velocity observations \\citep{Latham1984}. See \\citet{Garnavich1988} and \\citet{Wasserman1991} for other studies that use this method. \\item To reduce the number of false positives inherent in the above, one can use additional information such as proper motions. Orbital motions for wide systems are small; hence, the space velocities of a gravitationally bound pair should be the same, within some uncertainty. In the absence of radial velocities, which are very hard to obtain for a very large number of field stars, proper motion alone can be used to identify binary systems; the resulting pairs are known as common proper-motion (CPM) doubles. \\citet{Luyten1979a, Luyten1988} pioneered this technique in his surveys of Schmidt telescope plates using a blink microscope and detected more than 6000 wide CPM doubles with $\\mu>$100\\masyr\\ over almost fifty years. This method has since been used to find CPM doubles in the AGK~3 stars by \\citet{Halbwachs1986}, in the revised New Luyten Two-Tenths \\citep[rNLTT;][]{Salim2003} catalog by \\citet{Chaname2004}, and among the {\\em Hipparcos} stars in the Lepine-Shara Proper Motion-North \\citep[LSPM-N;][]{Lepine2005} catalog by \\citet{Lepine2007a}. All of these studies use magnitude-limited high proper-motion catalogs and, thus, select mostly nearby stars. \\end{enumerate}  More recently, \\citet*[][hereafter \\citetalias{Sesar2008}]{Sesar2008} searched the SDSS Data Release Six \\citep[DR6;][]{Adelman-McCarthy2008} for CPM binaries with angular separations up to 30$\\arcsec$ using a novel statistical technique that minimizes the difference between the distance moduli obtained from photometric parallax relations  for candidate pairs. They matched proper motion components to within 5\\masyr\\ and identified $\\sim$22000 total candidates with excellent completeness, but with a one-third of them expected to be false positives. They searched the SDSS DR6 catalog for pairs at all mass ranges and find pairs separated by 2000--47000 AU, at distances up to 4 kpc.  Similarly, \\citet{Longhitano2010} used the angular two-point correlation function to do a purely statistical study of wide binaries in the $\\sim$675 square degrees centered at the North Galactic Pole using the DR6 stellar catalog and predicted that there are more than 800 binaries with physical separations larger than 0.1~pc but smaller than 0.8~pc. As evidenced by the large false positive rate in \\citetalias{Sesar2008}, such large-scale searches for wide binaries generally involve a trade-off between completeness on the one hand and fidelity on the other, as they depend on statistical arguments for identification.  Complementing this type of ensemble approach, a high-fidelity approach may suffer from incompleteness and/or biases; however, there are a number of advantages to a ``pure\" sample of {\\em bona fide} wide binaries such as that presented in this work. For example, \\citet{Faherty2010} searched for CPM companions around the brown dwarfs in the BDKP catalog \\citep{Faherty2009} and found nine nearby pairs; all of their pairs were followed up spectroscopically and, hence, have a much higher probability of being real. As mass, age, and metallicity can all cause variations in the observed physical properties, e.g.\\ in radius or in magnetic activity, their effects can be very hard to disentangle in a study of single stars. Components of multiple systems are expected to have been formed of the same material at the same time, within a few hundred thousand years of each other (e.g.~\\citealt{White2001}; \\citealt*{Goodwin2004a}; \\citealt{Stassun2008}). Hence, binaries are perfect tools for separating the effects of mass, age, and metallicity from each other as well as for constraining theoretical models of stellar evolution. Some examples include benchmarking stellar evolutionary tracks (e.g.~\\citealt{White1999}; \\citealt*{Stassun2007}; \\citealt{Stassun2008}), investigating the age-activity relations of M dwarfs \\citep*[e.g.][]{Silvestri2005}, defining the dwarf-subdwarf boundary for spectral classification \\citep[e.g.][]{Lepine2007b}, and calibrating the metallicity indices \\citep{Woolf2005, Bonfils2005}. Moreover, equal-mass multiples can be selected to provide identical twins with the same initial conditions (same mass, age, and metallicity) to explore the intrinsic variations of stellar properties. In addition, wide binaries (a $>$ 100~AU) are expected to evolve independently of each other; even their disks are unaffected by the distant companion \\citep{Clarke1992}. Components of such systems are effectively two single stars that share their formation and evolutionary history. In essence they can be looked at as {\\em coeval laboratories} that can be used to effectively test and calibrate relations measured for field stars. Finally, as interest has grown in detecting exoplanets and in characterizing the variety of stellar environments in which they form and evolve, a large sample of {\\em bona fide} wide binaries could provide a rich exoplanet hunting ground for future missions such as SIM.  In this paper, we present a new catalog of CPM doubles from SDSS, each with at least one low-mass component, identified by matching proper motions and photometric distances. In \\S~\\ref{Sec: observation} we describe the origin of the input sample of low-mass stars; \\S~\\ref{Sec: method} details the binary selection algorithm and the construction of a Galactic model built to assess the fidelity of each binary in our sample. The resulting catalog and its characteristics are discussed in \\S~\\ref{Sec: catalog}. We compare the result of our CPM double search with previous studies in \\S~\\ref{Sec: discussion} and summarize our conclusions in \\S~\\ref{Sec: conclusions}. ",
        "conclusions": "\\label{Sec: conclusions} We have created the {\\slowpokes} catalog, comprising 1342 CPM binary pairs, identified through statistical matching of angular separation, photometric distances, and proper motion components. We have sifted the sample of chance alignments using a Galactic model based on empirical observations of stellar spatial and kinematic distributions. With the objective that each pair can be confidently used to investigate various science questions regarding low-mass stars, we have adopted a very restrictive set of selection criteria. This approach clearly underestimates the number of binary systems. Moreover, the sample includes several biases, the most important of which are a lack of systems with physical separations smaller than $\\sim 1000$~AU (arising from a strict bias against angular separations smaller than 7$\\arcsec$), and the exclusion of certain types of higher-order multiples (e.g.\\ triples) due to the strict photometric distance matching. However, as a consequence the catalog should contain very few false positives, making follow-up studies efficient.  We built a Monte Carlo-based six-dimensional Galactic model that is able to replicate the positional and kinematic properties of the stars in the Milky Way. In its current incarnation, we used it to calculate the the number of stars within a certain spatial volume in the Galaxy and the likelihood that those stars have common kinematics (proper motions) by chance. One of the things this model underscores is how difficult it is to find two physically unassociated stars close together in space: along a typical SDSS LOS, there are expected to be only 0.52 chance alignments within 15$\\arcsec$ and a minuscule 0.03 chance alignments if the volume is considered. The additional matching of proper motions gives each of the accepted {\\slowpokes} binaries a very low probability of being a false positive.  Due to their intrinsic faintness and the resulting small numbers, binarity studies of low-mass stars have been limited in scope. However, with the advent of large-scale deep surveys, detailed and statistically significant studies of M (and L) dwarfs are being done. {\\slowpokes} is now the largest sample to date of very wide, low-mass binaries. In particular, {\\slowpokes} provides a large sample of systems with physical separations up to $a\\sim$ 1~pc that will be useful for putting firm constraints on the maximum size of physically associated systems. How the widest of these systems form, and how long they survive, is in particular an interesting question that {\\slowpokes} is well suited to address. While numerical calculations suggest that approximately half of {\\slowpokes} systems can remain bound for at least 10~Gyr \\citep{Weinberg1987}, previously proposed empirical limits are violated by many {\\slowpokes} systems.  Indeed, the distribution of {\\slowpokes} binary separations is distinctly bimodal, suggesting the presence of (i) a population of tightly bound systems formed with sufficient binding energy to remain intact for the age of the Galaxy and (ii) a population of weakly bound systems that recently formed and that are unlikely to survive past 1--2~Gyr. Recent N-body simulations \\citep[e.g.][]{Kouwenhoven2010, Jiang2009} in fact predict a bimodal distribution of binary separations on the scales probed by the widest {\\slowpokes} systems.  We observed a wide binary frequency of $\\sim$1.1\\%, which is likely to be a minimum given the nature of our sample. While this is consistent with the results from \\citetalias{Sesar2008} who found 0.9\\% of stars at $Z=$ 500~pc had wide companions., it is significantly lower than $\\sim$9.5\\% of nearby solar-type {\\em Hipparcos} stars having wide companions \\citep{Lepine2007a}. While the incompleteness involved in the searching for companions at large distances probably causes some of this, both this study and \\citetalias{Sesar2008} saw a decrease in binary fraction as a function of Galactic height, evidence of dynamical destruction of older systems. Hence, the wide binary fraction around our initial sample of low-mass stars might actually be significantly lower than the {\\em Hipparcos} stars.  Besides the importance of {\\slowpokes} for constraining models of formation and evolution of binary stars, {\\slowpokes} systems are coeval laboratories---sharing an identical formation and evolutionary history without affecting each other---making them ideal for measuring and calibrating empirical relationships between rotation, activity, metallicity, age, etc. We have started programs to test and calibrate the age--activity relationship measured by \\citet{West2008} and to explore whether gyrochronology \\citep{Barnes2003a, Barnes2007} can be applied in the fully-convective regime. As coeval laboratories allow for the removal of one or more of the three fundamental parameters (mass, age, and metallicity), much more science can be done with a large sample of such systems.  Future astrometric missions, such as the Space Interferometry Mission (SIM), should provide exquisite astrometry, perhaps enabling us to trace the orbits of some of the {\\slowpokes} systems. While tracing orbits with periods $\\gtrsim10^{4-6}$~years sounds ambitious, with SIM's microarcsecond level (or better) astrometry \\citep{Unwin2008} combined with SDSS, DSS, and/or other epochs, it is not unrealistic. Similarly, the ``identical'' twins in {\\slowpokes} would be ideal sites to probe for the presence and differences in the formation mechanism of planets. As each identical twin in a CPM double provides similar environment for the formation and evolution of planets, these systems can be ideal sites to study planetary statistics. Due to their large separations the stars are not expected to influence each others evolution but have similar mass, age, and metallicity, as we noted earlier in \\S~\\ref{Sec: intro}.  The {\\slowpokes} catalog, as the name suggests, only contains systems for which kinematic information is available. We can, however, use the results from the Galactic model to identify pairs at the small separations ($\\dtheta<7\\arcsec$), albeit with a larger uncertainty, for which no kinematic information is available. Similarly companions which are fainter than $r=$ 20 can also be identified as the SDSS photometry is complete to $r=$ 22.5. The latter systems are likely to be skewed towards late-type dMs and unequal-mass pairs. A follow-up paper will study such systems and will add a large proportion of wide systems."
    },
    "1004/1004.0420_arXiv.txt": {
        "abstract": "We present photometric and spectral observation for four novae: V2362 Cyg, V2467 Cyg, V458 Vul, V2491 Cyg. All objects belongs to the ``fast novae'' class. For these stars we observed different departures from a typical behavior in the light curve and spectrum. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.5033_arXiv.txt": {
        "abstract": "We derive accurate proper motions of the \\M\\ 12~GHz masers towards the \\T\\ UC~H{\\sc ii} region, employing seven epochs of VLBA observations spanning a time interval of about 10~yr. The achieved velocity accuracy is of the order of 0.1~\\kms, adequate to precisely measure the relative velocities of most of the 12~GHz masers in \\T, with amplitude varying in the range \\ 0.3--3~\\kms. Towards \\T, the most intense 12~GHz masers concentrate in a small area towards the north (the northern clump) of the UC~H{\\sc ii} region. We have compared the proper motions of the \\M\\ 12~GHz masers with those (derived from literature data) of the OH~6035~MHz masers, emitting from the same region of the methanol masers. In the northern clump, the two maser emissions emerge from nearby (but likely distinct) cloudlets of masing gas with, in general, a rather smooth variation of line-of-sight and sky-projected velocities, which suggests some connection of the environments and kinematics traced by both maser types. The conical outflow model, previously proposed to account for the 12~GHz maser kinematics in the northern clump, does not reproduce the new, accurate measurements of 12~GHz maser proper motions and has to be rejected. We focus on the subset of 12~GHz masers of the northern clump belonging to the ``linear structure at P.A. = 130\\degr--140\\degr'', whose regular variation of LSR velocities with position presents evidence for some ordered motion. We show that the 3-dimensional velocities of this ``linear distribution'' of 12~GHz masers can be well fitted considering a flat, rotating disk, seen almost edge-on. ",
        "introduction": "\\label{intro}  Studying the evolution of  H{\\sc ii} regions and their complex interaction with the surrounding environment, is important in the context of star formation. H{\\sc ii} regions emit intense far-ultraviolet (FUV) radiation which photoionizes the circumstellar gas and evaporates more volatile molecules frozen on dust grain mantles. The physical and chemical properties of the gas in the natal molecular core (with typical size of about 0.1~pc) are greatly altered, with a major impact on the subsequent star-formation conditions. The radiation pressure exerted by the H{\\sc ii} regions onto dust grains can eventually significantly contribute to the disruption of the molecular core, too.  Hyper-compact (HC) and ultra-compact (UC)  H{\\sc ii} regions are clear indicators of forming (or recently born) massive stars. In regions of massive star-formation are often observed intense maser emissions of OH (at frequency of 1.6~and~6.0~GHz) , \\W\\ (22~GHz) and \\M\\ (6.7~and~12~GHz).  Using multi-epoch Very Long Baseline Interferometry (VLBI) observations, one can accurately determine absolute positions and velocities of maser spots (the single maser emission centers), providing unique information on the kinematics of the molecular gas around massive Young Stellar Objects (YSO).   Learning from the few sources studied in details so far, it appears that \\W\\ 22~GHz and OH 1.6~GHz masers could be tracing respectively  the fast (50--100~\\kms) and slow (5~\\kms) expansion of  compact H{\\sc ii} regions in different evolutionary stages \\citep{Tor03, Mos07, Fis07b}. Accurate measurements of \\emph{internal} motions of \\M\\ 6.7~and~12~GHz masers are still lacking in the literature. In several sources (single-epoch) VLBI observations have shown linear or arc-like distributions of maser spots, which have been interpreted in terms of edge-on rotating toroids or disks \\citep{Nor98,Pes05,Pes09}. However, in a few cases such elongated maser distributions lie in projection on the sky parallel to typical outflow tracers (such as the H$_2$ 2.12$\\mathrm{\\mu}$m line), rather suggesting association with outflowing gas \\citep{Deb03}. Measuring \\emph{relative} proper motions of \\M\\ maser spots appears the most direct way to discriminate between rotation and outflow.    \\T\\ is one of the best studied UC~H{\\sc ii} region. It is associated with a far-infrared source with a total luminosity of 10$^5$~L$_{\\sun}$ \\citep{Cam89}, which corresponds to a main-sequence O7 star. It harbours a plethora of maser transitions, among which the most intense and best studied by means of VLBI are those of OH at 1.6 \\citep{Blo92,Wri04,Fis06}, 6.0 \\citep{Des98, Fis07a} and 13~GHz \\citep{Bau98}, and those of \\M\\ at 6.7 \\citep{Men92} and 12~GHz \\citep[][hereafter MMWR1, MMWR2 and XRZM, respectively]{Mos99, Mos02,Xu06}. OH 1.6~and~6.0~GHz and \\M\\ 6.7~GHz masers are distributed across an area of about 2\\arcsec in size, covering all the western half of the continuum emission, where the H{\\sc ii} region is confined by relatively denser molecular gas. In contrast, most of \\M\\ 12~GHz maser emission comes from a small (diameter of about 200~mas) cluster (the ``northern clump'') to the north of the UC H{\\sc ii} region \\citep[][Fig.~3]{Mos99}, in correspondence with the northern ionized clump observed in high frequency (15--23~GHz) continuum images of \\T. The highly excited OH 13~GHz masers also emerge from a compact region (extended about 100~mas) of the northern ionized clump. This area, including highly excited maser transitions, the brightest continuum emission, and  showing also intense magnetic fields \\citep{Wri04}, is the most active of the whole UC H{\\sc ii} region and could host the main source of excitation in \\T.  By comparing two Very Long Baseline Array (VLBA)\\footnote{The VLBA is operated by the National Radio Astronomy Observatory (NRAO).} epochs separated by about 5~yr, MMWR2 first measured relative proper motions of \\M\\ 12~GHz masers in \\T\\ and found typical velocity amplitudes of a few \\kms. Positions and velocities of the 12~GHz masers in the northern clump have been fitted with a narrow conical outflow model oriented at close angle with the line of sight. However, most of the measured proper motions have large ($\\ga$50\\%) uncertainties and the outflow model was effectively constrained by the (well-known) maser line-of-sight velocities only. Recently, by using 5 VLBA epochs,  XRZM have measured the parallax of the \\M\\ 12~GHz masers in W3(OH), deriving a very accurate source distance of $1.95\\pm0.04$~kpc. To better constrain the kinematics of the 12~GHz masers, in this work we combine the MMWR2 and XRZM 12~GHz maser observations (forming a dataset of seven VLBA epochs spanning about 10~yr) and derive accurate, relative proper motions for the persistent maser spots.  Section~\\ref{sum} of this paper briefly describes the VLBA observations of 12~GHz methanol masers which we employ for proper motion derivation. Section~\\ref{var} presents a study of the 12~GHz maser variability over the 10~yr period spanned by the VLBA observations. In Sect.~\\ref{met_pm} we derive the 12~GHz maser proper motions and compare the new values with the previous measurements by MMWR2. Basing on two published VLBI observations of OH 6.0~GHz masers in \\T, Sect.~\\ref{oh_pm} derives proper motions of this maser transition and compares the velocity distributions of OH 6.0~GHz and \\M\\ 12~GHz masers in the  northern clump.  Finally,  in Sect.~\\ref{mod_kin}, the conical outflow model by MMWR2 is tested against the more accurate 12~GHz maser proper motions derived in Sect.~\\ref{met_pm}. Conclusions are drawn in Sect.~\\ref{conclu}. ",
        "conclusions": "\\label{conclu}  Employing seven VLBA epochs spanning a lapse of time of about 10~yr, this work determines accurate proper motions of the \\M\\ 12~GHz masers in the northern clump of \\T. The main features of the 12~GHz maser emission are persisting and most of the strongest maser spots ($\\ge$5~Jy) are detected at all the epochs. The best measured sky-projected velocities have amplitudes in the range \\ 0.3--2~\\kms\\ and the corresponding uncertainties are as small as 0.1~\\kms.  The conical flow model proposed by MMRW2 to interpret the kinematics of the \\M\\ 12~GHz masers in the northern clump of \\T, is tested with the new, more accurate measurements of spot proper motions.  The result shows that the conical flow model is not able to reproduce adequately the 12~GHz maser sky-projected velocities and has to be discarded. We consider a subset of 12~GHz masers in the northern clump belonging to the ``linear structure at P.A. = 130\\degr--140\\degr'', whose regular variation of LSR velocities with position presents evidence for some ordered motion. Positions and 3-dimensional velocities of this group of 12~GHz maser spots are well fitted with a flat disk geometry and a pure rotation field.  This ``linear distribution'' of 12~GHz masers could trace a self-gravitating, low-mass circumstellar disk, in the phase of being photo-evaporated by the strong UV-radiation field excaping from the \\T\\ UC H{\\sc ii} region.  Using literature data, we have derived proper motions of the OH 6.7~GHz masers in the northern clump of \\T. Comparing the overall distribution of positions and velocities of the OH and \\M\\ maser emissions in the northern clump, it seems that OH~6035~MHz and \\M\\ 12~GHz masers complement each other, emerging from nearby (but likely distinct) cloudlets of masing gas with, in general, a rather smooth variation of line-of-sight and sky-projected velocities."
    },
    "1004/1004.3035_arXiv.txt": {
        "abstract": "A distinguishable and observable physical property of Naked Singular Regions of the spacetime formed during a gravitational collapse has important implications for both experimental and theoretical relativity. We examine here whether energy can escape physically from naked singular regions to reach either a local or a distant observer within the framework of general relativity. We find that in case of imploding null dust collapse scenarios field outgoing singular null geodesics including the cauchy horizon can be immersed between two Vaidya spacetimes as null boundary layers with non vanishing positive energy density. Thus energy can transported from the naked singularity to either a local or a distant observer. And example illustrating that similar considerations can be applied to dust models is given. ",
        "introduction": "A star with sufficient remnant mass ($ \\gtrsim 3 M_{\\odot}$), on completion of its nuclear fuel cycle, must enter the phase of a continuous gravitational collapse. Once the nuclear fuel is exhausted gravitational forces become all powerful and hence star's internal pressure can not sustain the equilibrium resulting in a continued collapse \\cite{dutt,oppen}. In the late stages of collapse the gravitational forces become dominant and the physics of collapse is determined mainly by the theory of general relativity. Under quite general and physical situations general relativity predicts that such a collapse must end in a singularity, i.e., a region of spacetime with extreme curvatures \\cite{pen65,HW66,HW}. Physically one could describe singularity as a region of space with vanishing volume and unbounded gravitational forces. General relativity, however, does not say anything about the nature or physical properties of such a singularity. This is partially due to the fact that mathematical structure breaks down preventing analysis at and beyond the singularity. One could perhaps argue that as collapse progresses and matter is condensed in a region comparable to Planck length the quantum physical properties of spacetime would become dominant, thus preventing the formation of singularity. But this picture may not hold also since gravity as a force is very different in its nature in comparison to other forces and has a geometrical interpretation as curvature of spacetime. Moreover, despite numerous efforts, a viable quantum theory of gravity is not in sight. Hence for such regions of spacetime, whether relativity theory or quantum physics would determine the physics is still an open question.  To fill in the gap in our understanding of spacetime singularities in a mathematical consistent manner, a cosmic censorship conjecture, that all gravitational collapse must end in a black hole was proposed \\cite{pen65,Pen}. The physical consequence of such a hypothesis is that even before the formation of a singularity a trapped surface develops covering the singularity from the outside world. Hence from a physical point of view singularity is hidden from the outside world. Initial studies in censorship were directed towards formulating the conjecture in a mathematically precise manner which could then possibly be proven \\cite{wald_rev}. This also led to formulation of other conjectures like, hoop conjecture by Kip Thorne and Siefert's conjecture \\cite{kip,seifert}. However, extensive studies in collapse with various forms of matter fields have shown that under fairly generic reasonable physical conditions both naked singularity and black holes would form as an end state of collapse, depending on various initial and boundary conditions \\cite{JJS}. It is still not very clear how to classify either matter or the initial and boundary conditions in a satisfactory way which would end in either state of singularity (naked or covered). Thus from the studies this far almost all physically reasonable matter fields lead to both naked and covered singularities during collapse (see, \\cite{joshi_book}, and references therein).  Considerable work has since been done on naked singularities from the point of view of giving counterexamples to cosmic censorship but also on the study of their nature and structure. Having established their existence it is important to study the phenomena of formation of naked singularity from a more astrophysical perspective. One could look for a possible observable signature of naked singularity distinguishing them from other compact strong gravity objects, like black holes. In the studies carried out this far the stress has been towards showing that for a naked singularity to be ``observable'' a family of lightlike geodesics must terminate at the singularity \\cite{Christo_all,jd}. Optical appearance and redshift for such possible radiation has also been studied \\cite{red-shift}. However, from the point of view of general relativity the first null ray coming out of singularity forms a Cauchy horizon (CH), and the spacetime model cannot remain valid after its formation. Therefore, without any consistent extension of spacetime beyond CH the validity and usefulness of all such geodesic analysis becomes doubtful. The basic question of the existence of the spacetime structure after the CH is unaddressed (it is difficult to provide extensions of spacetimes, for example, even for shell-crossing singularities which are gravitationally weak \\cite{clarke_extension}) which is of utmost importance if we want to talk about families of geodesics ending at singularity in the past, making it a possible astrophysical source.  In this paper we wish to study the structure of the spacetime from this perspective.  Is it possible to connect the two spacetimes before and after with Cauchy horizon as the boundary? Whether the resulting spacetime after the CH has formed can still have the same symmetry? Does relativity theory allows such continuation of spacetime through CH and whether boundary conditions pose any restrictions? Furthermore, can these boundary layers carry energy from naked singularity to a distant observer? Earlier Hiscock et al. has considered a model spacetime in which cauchy horizon ultimately becomes the event horizon of the schwarzschild black hole with non vanishing surface energy density and where it could be visible to observers falling into the blackhole \\cite{Hiscock}.  If indeed the formation of a naked singularity is a physical phenomenon then the CH would represent a null surface layer emanating from the naked singularity, and reaching the distant observer separating the two spacetimes. It has been suggested in various studies that naked singularities may be responsible for various high energy phenomena in our universe (for example gamma ray bursts etc. \\cite{gammaburst}). It has also been suggested that in the late stages of collapse, when spacetime shrinks to size of the order of planck length quantum effects would play a dominant role resulting in either a burst of particle creation or preventing the formation of singularity all together \\cite{quantumeffects}. Our aim in this paper is to examine two examples of naked singularities within the frame work of general relativity and whether this allows such a scenario as emission of a impulsive null wave carrying energy from the naked singularity. The result of such a study would have manifold implications. First does there exist a spacetime after the formation of a naked singularity which can be joined satisfactorily together with the original model separated by the null shockwave (CH)? If such a spacetime exists then whether it allows the existence of outgoing families of geodesics terminating at the singularity in the past. Second, and equally important, question is the structure of the CH itself. Whether this null surface `boundary layer' is allowed to carry huge amounts of energy along the null ray to distant observer? And, if the answer is in affirmative, what is its structure and whether this scenario can be called a valid solution to the Einstein equations? ",
        "conclusions": ""
    },
    "1004/1004.4086_arXiv.txt": {
        "abstract": "{We report results of the INTEGRAL Target of Opportunity observations of the transient X-ray burster XTE J1810-189.  The observations were performed on April 3--6, 2008, soon after the discovery of the source and near the peak of its outburst. That time the source had a flux of about 50 mCrab and exhibited a hard Comptonized X-ray spectrum extending well above 100 keV. Being approximated by a power law with an exponetial cut-off in the broad 3--100 keV energy band it gave the average photon index $\\Gamma\\simeq 1.6$ and $kT_{cutoff}\\simeq 67$ keV. We found only slight indications for changes in the index during the observation ($\\Gamma$ first steady decreased from $\\sim 2.0$ to $\\sim1.3$ and then increased back to $\\sim 2.0$). However the $N_{\\rmn H}$ value measured by absorption in the low energy part of the spectrum changed drastically and very irregularly (from $\\sim 4\\times 10^{22}$ till $\\sim 100\\times 10^{22}$ cm$^{-2}$). There were 10 type I X-ray bursts detected from the source during these TOO observations. Assuming that the Eddington luminosity was reached during the burst with the highest peak flux we get an upper estimate for a distance to the source $D=6.4\\pm0.6$ kpc. From the X-ray burst parameters we conclude that this LMXB harboures an evolved star.}  \\FullConference{The Extreme sky: Sampling the Universe above 10 keV - extremesky2009,\\\\ October 13-17, 2009\\\\ Otranto (Lecce) Italy}   \\begin{document} ",
        "introduction": "The X-ray transient XTE~J1810-189 was discovered in the spring of 2008 during RXTE/PCA monitoring scans of the Galactic ridge region \\cite{markwardt08}. A pointed observation started on March 10, at 21:05 UTC revealed a variable source (30\\% r.m.s. fluctuations). The emission spectrum was consistent with an absorbed power law ($N_{\\rmn H}\\simeq1\\times10^{22}~\\mbox{cm}^{-2}$, photon index $\\Gamma\\simeq1.9$). The PCA flux history suggested its gradual rise since March 5. The source was observed on March 12-15 with INTEGRAL that measured a slightly steeper spectrum $\\Gamma=2.26\\pm0.12$ in the hard $>20$ keV IBIS/ISGRI energy band \\cite{neronov08}. The observation of XTE~J1810-189 on March 17, 2008 with the Swift/XRT telescope revealed the similar spectrum but the higher absorption $N_{\\rmn H}=(4.2\\pm0.7)\\times10^{22}~\\mbox{cm}^{-2}$ indicating that there might be an internal source of absorption in the system \\cite{krimm08}. In a pointed observation of the source on March 26, at 12:47 UTC the RXTE/PCA detected a type I X-ray burst identifying a compact object in the system as a neutron star. Assuming the Eddington peak luminosity, the upper limit  for a distance to the source was obtained $D\\la 11.5$~kpc \\cite{markwardt08b}.  In this paper we report the results of TOO (Target of Opportunity) observations of XTE J1810-189 performed with INTEGRAL on April 3-6, 2008. ",
        "conclusions": ""
    },
    "1004/1004.2496_arXiv.txt": {
        "abstract": "The first light from a supernova (SN) emerges once the SN shock breaks out of the stellar surface. The first light, typically a UV or X-ray flash, is followed by a broken power-law decay of the luminosity generated by radiation that leaks out of the expanding gas sphere. Motivated by recent detection of emission from very early stages of several SNe, we revisit the theory of shock breakout and the following emission, paying special attention to the photon-gas coupling and deviations from thermal equilibrium. We derive simple analytic light curves of SNe from various progenitors at early times. We find that for more compact progenitors, white dwarfs, Wolf-Rayet stars (WRs) and possibly more energetic blue-supergiant explosions, the observed radiation is out of thermal equilibrium at the breakout, during the planar phase (i.e., before the expanding gas doubles its radius), and during the early spherical phase. Therefore, during these phases we predict significantly higher temperatures  than previous analysis that assumed equilibrium. When thermal equilibrium prevails, we find the location of the thermalization depth and its temporal evolution. Our results are useful for interpretation of early SN light curves. Some examples are: (i) Red supergiant SNe have an early bright peak in optical and UV flux, less than an hour after breakout. It is followed by a minimum at the end of the planar phase (about 10 hr), before it peaks again once the temperature drops to the observed frequency range. In contrast WRs show only the latter peak in optical and UV. (ii) Bright X-ray flares are expected from all core-collapse SNe types. (iii) The light curve and spectrum of the initial breakout pulse holds information on the explosion geometry and progenitor wind opacity. Its spectrum in more compact progenitors shows a (non-thermal) power-law and its light curve may reveal both the breakout diffusion time and the progenitor radius. ",
        "introduction": "A breakout of a shock through the stellar surface is predicted to be the first electro-magnetic signal heralding the birth of a supernova (SN) \\citep{Colgate74,Falk78,Klein78,Imshennik81,Ensman92,MatznerMcKee99}. Before breakout the shock is propagating through the opaque stellar envelope. The shock is radiation dominated (i.e., the energy density behind the shock is dominated by radiation) and it accelerates while propagating through the decreasing density profile of the envelope, leaving behind the shock an expanding radiation dominated gas. Following the shock breakout, photons continue to diffuse out of the expanding stellar envelope producing a long lasting emission that slowly decays with time \\citep[e.g.,][]{Grassberg71,Chevalier76,Chevalier92,Chevalier08,Piro09}. The typical frequency of the breakout emission ranges from far ultra-violet to soft $\\gamma$-rays in core collapse SNe, and as we show here, is in $\\gamma$-rays in type Ia SNe. The typical frequency of the following emission decreases to the visible-near UV bands after a day. The energy released during the breakout increases with the progenitor radius and can reach $\\sim 0.1\\%$ of the SN explosion energy in a red supergiant. The luminosity of core collapse SNe after a day is $\\sim 10^{41}-10^{42} {\\rm~ erg/s}$. Thus, the shock breakout and the emission through the first day can be detected out to the nearby Universe, but without any preceding knowledge of where to look, their detection is challenging. Nevertheless, the search worth the effort as this emission bears direct information on the properties of the progenitor and the explosion, which are difficult to obtain in any other way.   The development during the recent decade of sensitive UV, X-ray and soft gamma-ray detectors, with relatively large fields of view, lead to the discovery of several shock breakout candidates \\citep{Campana06,Soderberg08,Gezari08,Schawinski08,Modjaz09}. Motivated by these, and by the  rising potential for future detection of shock breakouts from various progenitors, we revisit this topic. We develop an analytic model that provides light curves (luminosity and temperature) starting from the breakout, through the quasi-planar expansion phase to the spherical expansion phase, until recombination and/or radioactive decay start playing a significant role. These phases were explored in previous works, where the most updated analytic study of the spherical phase was carried-out by \\cite{Chevalier92,Chevalier08}, and \\cite{Waxman07,Rabinak10}. The study of the planar phase was carried out only very recently by \\cite{Piro09} in the context of Type Ia shock breakout. The advantage of our model is that we follow the photon-gas coupling within the expanding gas. At each stage of the evolution we find the location at which the observed temperature is determined. We find whether the radiation at this place is in thermal equilibrium or not and calculate the observed temperature. It turns out that the temperature evolution during the planar phase and the early spherical phase depends strongly on the thermal equilibrium of the radiation just behind the shock in the breakout layer. In radiation dominated shocks, radiation is out of thermal equilibrium if the shock velocity is high enough \\citep{Weaver76}, which is the case in shock breakout from more compact progenitors and more energetic explosions \\citep{Katz09}. We provide the first calculation (analytic or numerical) of the light curve in case that the observed radiation is out of thermal equilibrium at the source. We also carry out the first analytic calculation of the evolution of the location of the thermalization depth, through the different phases, when the radiation  is in thermal equilibrium. We use our model to examine the light curve and spectrum of the initial pulse that is strongly affected by light travel time effects and opacity of the progenitor stellar wind. Finally, we use our model to explore the properties of early SNe light curves resulting from various progenitor types including red supergiants (RSG), blue supergiants (BSG), Wolf-Rayet stars (WR) and white dwarfs (WD). We provide  simple formula of early SN light curves for these different progenitors.  We present our model and its general results in section \\ref{SEC theory}. The bolometric luminosity and spectrum of the initial pulse are discussed in section \\ref{SEC initial pulse}. Early SNe light curves resulting from various progenitors are presented in section \\ref{SEC lightcurve}. A reader that is interested only at the final light curves should go directly to this section. In section \\ref{SEC comparison} we compare our calculations to previous analytic and numerical studies. We summarize our main results in section \\ref{SEC Summary}. We ignore in this paper any cosmological redshift effects. ",
        "conclusions": "\\label{SEC Summary} We derive analytic SNe light curves at early times, as long as recombination and radioactive decay do not play an important role. These light curves are valid while the observed temperature is above about $1$ eV and before injection by radioactive decay becomes important. These conditions hold during the first day after the explosion of a typical SN. The main advantage of our analysis over previous ones is the account for the radiation-gas coupling, which leads to determination of the observed temperature when the radiation at the color shell is out of thermal equilibrium. It also corrects previous estimates of the observed temperature, and the color shell location, when the radiation at the color shell is in thermal equilibrium. We define a thermal coupling coefficient, $\\eta$, and find that the temperature evolution can follow two very different tracks, depending on $\\eta_0$, i.e., the value of $\\eta$ in the breakout shell at the breakout time. When the breakout shell is out of thermal equilibrium ($\\eta_0>1$) the observed temperature starts high above the value obtained when thermal equilibrium is assumed and it drops faster than the case that the breakout shell is in thermal equilibrium. Thermal equilibrium is typically gained (when $\\eta_0>1$) only at early stages of the spherical phase.  We discuss the luminosity and spetral evolution during the initial pulse and derive early SN light curves for various SN progenitors as a function of the explosion energy and the progenitor mass and radius. These are useful for interpretation of SNe light curves during the first day, which can teach us about properties of the progenitor star and potentially to lead to its identification. Additionally, it can be used to evaluate the effect of the early emission (e.g., ionization of the circum burst medium), in case that its detection is missed, on the environment at the SN vicinity. Finally it is useful for planning targeted searches of shock breakouts of various SNe types. The theory we discuss here can be also applied in some cases to shock breakout from non-SN stellar explosions. For example, the explosion of solar-like star by tidal forces in the vicinity of a super-massive black hole discussed recently by \\cite{Guillochon09}.  The main conclusions based on our analysis are: \\begin{itemize} \\item It was shown that shock breakout radiation from WDs, WRs and some BSGs is out of thermal equilibrium \\citep{Katz09}. We show that it typically remains out of thermal equilibrium throughout the planar phase and until the early spherical phase. In SN from these compact progenitors the observed temperature at this time is significantly higher than the one obtained when thermal equilibrium is assumed. The observed temperature falls as $t^{-\\alpha}$, where $1/3<\\alpha<2/3$, during the planar phase, and once the evolution becomes spherical it plunges down (roughly as $t^{-2}$) until we observe radiation that is in thermal equilibrium at the source. \\item Breakouts from RSGs and some BSGs are in thermal equilibrium. The flux at frequencies below $T_{obs}$ (e.g., optical/UV) starts with a bright initial pulse and then it decays during the planar phase reaching a minimum at $t_s$. The flux is rising during the planar phase reaching a second maximum when $T_{obs}$ falls into the observed frequency. \\item In cases where the radiation is in thermal equilibrium at the source, the location of the thermalization depth, $r_{cl}$ is not trivial \\citep[e.g.,][]{Ensman92}. The assumptions used in some previous analytic calculations, such as $r_{cl}=r(\\tau=1)$ \\citep[e.g., ][]{Chevalier08} or $r_{cl}=\\rh$ \\citep[e.g., ][]{Piro09} are incorrect. Instead, in this case $\\rh<r_{cl}<r(\\tau=1)$ and it satisfies $r_{cl} =r({\\eta=1})$, where $T_{obs} = (L\\tau_{cl}/4\\pi \\sigma r_{cl}^{~2})^{1/4}$. \\item The initial pulse, which is smeared by light travel-time and asphericity, contains emission from gas at different temperatures. The resulting combined spectrum is not a black-body. Instead, below the peak of $\\nu F_\\nu$ the spectrum is a power-law $F_\\nu \\propto \\nu^{-\\beta}$ where $0<\\beta<1/2$. If asphericity is negligible then this effect is negligible in RSGs, minor in BSGs and significant in WRs, where a power-law $F_\\nu \\propto \\nu^{-\\beta}$, ranges over one or two orders of magnitude in $\\nu$. \\item  The bolometric luminosity is dictated by the balance between the radiation diffusion time scale and the dynamical time scale \\citep[e.g.,][]{Chevalier92}. Namely, the source of the observed energy at any time is the shell where these two timescales are comparable. Since the photons dominate the heat capacity, the bolometric luminosity does not depend on the photon-gas coupling,  and it evolves independently of the thermal coupling. \\item The bolometric light curve and spectral evolution of the initial pulse depends on the explosion asphericity and wind transparency. A ``bare\" spherical explosion has a unique and therefore identifiable light curve. It shows multiple time scales as well as spectral evolution during the initial pulse. It rises quickly, over $t_0$, and lasts for $R_*/c$. A mildly opaque wind ($1<\\tau_w<c/\\vt$) results in a pulse of characteristic time $\\tau_w R_*/c$ during which the spectrum does not evolve significantly. A ``bare\" aspherical explosion produces a time evolving spectrum, whose light curve and spectral evolution depends on the asphericity details. Much can be learned from observation of the initial pulse! \\item The initial pulse of a RSG may release $\\sim 3 \\cdot 10^{46}$ erg over a duration of $\\sim 1000$ sec in soft X-rays $\\approx 0.2$ keV. A BSG shock breakout releases a comparable amount of energy in harder X-rays over $\\sim 100$ sec. The breakout from a WR releases $\\sim 10^{45}$ erg within $\\sim 10$ sec in the form of hard x-rays or soft $\\gamma$-rays. Thus, not only WR stars, but also breakouts from BSGs and potentially RSGs are predicted to produce strong X-ray flares that can be detected by X-ray telescopes to large distances. WR breakouts may also be hard enough to be detectable by current soft gamma-ray detectors to a distance of $\\sim 10$ Mpc, and mimic a single pulse long gamma-ray burst. \\item The early emission from a type Ia SNe is in $\\gamma$-rays. It is detectable within our galaxy by current detectors and may mimic a very short ($10-20$ ms) gamma-ray burst. \\end{itemize}  We thank Carlos Badenes, Orly Gnat, Chris Hirata, Amir Levinson, Tsvi Piran, Sterl Phinney  and Amiel Sternberg for helpful discussions. E.N. was partially supported by the Israel Science Foundation (grant No. 174/08) and by an IRG grant. R.S. was partially supported by ERC and IRG grants, and a Packard Fellowships. \\renewcommand{\\theequation}{A -\\arabic{equation}} \\setcounter{equation}{0}  %"
    },
    "1004/1004.4825_arXiv.txt": {
        "abstract": "The direct detection of exoplanets has been the subject of intensive research in the recent years. Data obtained with future high-contrast imaging instruments optimized for giant planets direct detection are strongly limited by speckle noise. Specific observing strategies and data analysis methods, such as angular and spectral differential imaging, are required to attenuate the noise level and possibly detect the faint planet flux. Even though these methods are very efficient at suppressing the speckles, the photometry of the faint planets is dominated by the speckle residuals. The determination of the effective temperature and surface gravity of the detected planets from photometric measurements in different bands is then limited by the photometric error on the planet flux. In this work we investigate this photometric error and the consequences on the determination of the physical parameters of the detected planets. We perform detailed end-to-end simulation with the CAOS-based Software Package for spectro-polarimetric high-contrast exoplanet research (SPHERE) to obtain realistic data representing typical observing sequences in Y, J, H and Ks bands with a high contrast imager. The simulated data are used to measure the photometric accuracy as a function of contrast for planets detected with angular and spectral+angular differential methods. We apply this empirical accuracy to study the characterization capabilities of a high-contrast differential imager. We show that the expected photometric performances will allow the detection and characterization of exoplanets down to the Jupiter mass at angular separations of 1.0\\as and 0.2\\as respectively around high mass and low mass stars with 2 observations in different filter pairs. We also show that the determination of the planets physical parameters from photometric measurements in different filter pairs is essentialy limited by the error on the determination of the surface gravity. ",
        "introduction": "\\label{sec:introduction}  Since the detection of the first exoplanet orbiting a main sequence star, 51~Peg, a large population of these objects has been discovered covering a wide range of masses and orbital periods mostly with indirect methods such as radial velocities measurements and transits (see e.g. \\citealt{santos2008} for a review). Although mainly sensitive to planets with period of less than 10~years, radial velocities surveys have found stars which start to show long term trends indicating possible low mass companions orbiting at large orbital separations \\citep{wittenmyer2007}. The wide use of adaptive optics (AO) systems and coronagraphy in large telescopes instrumentation for high-contrast imaging has allowed to start probing the vicinity of nearby stars for low mass companions at large orbital separations. Over the last decade, a handful of objects close to the planetary mass regime have been imaged with existing instruments, such as 2M~1207~b \\citep{chauvin2005a}, DH~Tau~B \\citep{itoh2005}, GQ~Lup~b \\citep{neuhauser2005}, AB~Pic~b \\citep{chauvin2005b}, CHXR~73~B \\citep{luhman2006}, and more recently Fomalhaut~b \\citep{kalas2008}, 1RXS~J1609~b \\citep{lafreniere2008}, $\\beta$~Pic~b \\citep{lagrange2008} and the triple system around HR~8799 \\citep{marois2008b}. However, the large uncertainty on the mass of these objects may place some of them in the sub-stellar rather than the planetary mass regime.  The next generation of planet finding instruments currently being built will combine: (i) extreme AO systems with a large number of actuators \\citep{angel1994,stahl1995,langlois2001} to reach very high corrections in the near-infrared \\citep{fusco2006,allercarpentier2008}; and (ii) high-efficiency coronagraphs such as the apodized pupil Lyot coronagraph (\\citealt{soummer2005} and references therein) or the achromatic 4-quadrant phase mask \\citep{rouan2000,mawet2006} to obtain optimal star extinction. GPI (Gemini Planet Imager) for Gemini South \\citep{macintosh2006} and SPHERE (Spectro-Polarimetric High-contrast Exoplanet REsearch) for the ESO-VLT \\citep{beuzit2006} are the two leading instruments of that category. These will both start operation in 2011, along with HiCIAO (High-contrast Coronagraphic Imager for Adaptive Optics) for Subaru \\citep{hodapp2008}. They will aim at detecting exoplanets down to the Jupiter mass (\\MJup) around nearby young stars by reaching contrast values of 15 to 17.5~mag ($10^{-6}$ to $10^{-7}$) at angular separations of $\\sim$0.1\\as. Both GPI and SPHERE will incorporate diffraction limited integral field spectrographs (IFS) in the near infrared, allowing to obtain images simultaneously at several wavelengths. SPHERE will also incorporate a differential imager named IRDIS (InfraRed Dual Imaging Spectrograph, \\citealt{dohlen2008a}) that will provide simultaneous images at two close wavelengths in either one of its 5 different filter pairs over the Y to Ks bands.  These instruments will allow to use different observing strategies such as Spectral Differential Imaging (SDI, \\citealt{racine1999}) or Angular Differential Imaging (ADI, \\citealt{marois2006a}) to obtain data which will be analyzed using advanced methods such as Spectral Deconvolution \\citep{sparks2002} or Localy Optimized Combination of Images (LOCI, \\citealt{lafreniere2007}) for IRDIS data. Specific methods of signal extraction have also been developed within the SPHERE consortium \\citep{mugnier2008,smith2009} to be used in the data reduction pipeline of IRDIS. Data analysis methods are of extreme importance when it comes to the detection of faint objects in coronagraphic images dominated by speckles. In particular, the precise estimation of the object flux after applying these methods is critical for calibration of planetary mass objects model atmospheres (\\citealt{allard2001,allard2003,allard2007}, and in preparation; \\citealt{burrows2006,ackerman2001,tsuji2005}) based on effective temperature (\\teff) and surface gravity (\\logg), and the corresponding evolutionary models \\citep{burrows1997,chabrier2000,baraffe2003,fortney2005}. The only currently available method for mass estimation, when no dynamical mass estimations are known for the object, consists in incorporating photometric measurements and their error assessment into these physical models. Although spectroscopy is the method of choice for characterization, it may not be possible to obtain high-quality spectra of extremely faint sources \\citep{vigan2008,janson2010}.  In this paper we investigate the photometric limitations in high-contrast data obtained at different wavelengths with a dual-band imager like IRDIS, and we study how it translates in terms of characterization of planetary mass objects. After reminding in Sect.~\\ref{sec:limitations_high_contrast_imaging} the origin of the speckle noise which is the fundamental limitation in high-contrast imaging and the methods to overcome it, we present in Sect.~\\ref{sec:end_to_end_simulations} the end-to-end simulations of IRDIS which were performed to obtain a realistic observing sequence. Section~\\ref{sec:photometric_accuracy} briefly describes the detection limits obtained with ADI and SDI+ADI data analysis methods before studying the performances in terms of photometric accuracy. Finally, in Sect.~\\ref{sec:photometric_characterization} we advocate a filter pair procedure for IRDIS and present an analysis of the characterizations which will be possible from aperture photometry. ",
        "conclusions": "\\label{sec:conclusions}  Next generation instruments for ground based exoplanet direct imaging such as SPHERE and GPI will provide data intrinsically limited by the speckle noise. This noise needs to be attenuated using \\emph{a posteriori} data analysis methods, such as spectral and angular differential imaging. In this paper, we have quantified the exoplanet characterization capabilities of IRDIS, the differential imager of SPHERE, using photometric and differential photometric information.  The photometric performances have been evaluated with aperture photometry on the detectable planets as a function of contrast and wavelength for a standard test case. In particular we have shown that the photometric performance strongly depends on wavelength because of the PSF chromaticity and on the position with respect to the AO control radius. With ADI, a photometric accuracy of 0.2~mag is reached inside the AO control radius for contrast values of 10 to 11~mag between the star and planet, while at larger radius the precision can be reached for contrasts up to 15~mag. With SDI+ADI, the overall photometric performance is improved, increasing by 1.5 to 2.5~mag the contrast at which a 0.2~mag precision is reached.  We have defined empirical photometric accuracies for IRDIS in its different filter pairs, which were used to test the characterization capabilities of IRDIS with all possible combinations of filter pairs. Priorities for characterization have been set on the different filter pairs by finding the pairs that maximize the number of possible characterizations in various conditions. We showed that when there is no \\emph{a priori} knowledge on the planet, the filter pairs Y2Y3 and J2J3 allow a larger number of characterizations than the pairs H3H4 and K1K2. Then we showed that using filter pair H2H3, it will be possible to characterize planets with \\teff~$\\simeq$~900~K around high mass stars at small angular separation and \\teff~$\\simeq$~700~K around lower mass stars. Adding Y2Y3 and J2J3 filter pairs allows to decrease the characterizable \\teff by 200~K at all separations and for all stellar magnitudes, while considerably decreasing the number of non-characterizations for warmer planets.  Finally, we showed that non-unique characterizations, i.e. planets for which the \\teff and \\logg values could not be determined exactly (within the limits of the grid of atmosphere models) are mostly dominated by errors on the determination of \\logg. In particular, we showed that around low mass stars, where the contrast is more favorable, the determination of \\teff is largely improved, while the errors on \\logg remain identical around a high mass star. Considering evolutionary models, and including typical ages from the future SPHERE target list, we showed that such errors on the determination of a low mass planet (2~\\MJup) results in a large uncertainty around a high mass bright star, but is very close to the limits fixed by the models grid around a fainter low mass star.  With this work we showed that IRDIS, the dual-band imager of SPHERE, should be able to fulfill the goal set for a high-contrast imager, that is the ability to detect and characterize planetary companions down to the Jupiter mass around nearby young stars. Similar developments could also be performed for an Integral Field Spectrograph in the future, allowing to quantify precisely the performances of SPHERE in the near-infrared, and to work on the aspect of characterization strategy for the detected objects."
    },
    "1004/1004.3772_arXiv.txt": {
        "abstract": "{ Dictated by the string theory and various higher dimensional scenarios, black holes in $D>4$-dimensional space-times must have higher curvature corrections. The first and dominant term is quadratic in curvature, and  called the Gauss-Bonnet (GB) term. We shall show that although the Gauss-Bonnet correction changes black hole's  geometry only softly, the emission of gravitons is suppressed by many orders even at quite small values of the GB coupling. The huge suppression of the graviton emission is due to the multiplication of the two effects: the quick cooling of the black hole when one turns on the GB coupling and the exponential decreasing of the grey-body factor of the tensor type of gravitons at small and moderate energies. At higher $D$ the tensor gravitons emission is dominant, so that the overall lifetime of black holes with Gauss-Bonnet corrections is many orders larger than was expected. This effect should be relevant for the future experiments at the Large Hadron Collider (LHC).} }  \\begin{document}  \\PRDonly{ \\title{Long life of Gauss-Bonnet corrected black holes} \\author{R. A. Konoplya}\\email{konoplya_roma@yahoo.com} \\affiliation{Department of Physics, Kyoto University, Kyoto 606-8501, Japan\\\\ \\& \\\\ \\mbox{Theoretical Astrophysics, Eberhard-Karls University of T\\\"{u}bingen,}\\\\ T\\\"{u}bingen 72076, Germany} \\author{A. Zhidenko}\\email{zhidenko@fma.if.usp.br} \\affiliation{Instituto de F\\'{\\i}sica, Universidade de S\\~{a}o Paulo,\\\\ C.P. 66318, 05315-970, S\\~{a}o Paulo-SP, Brazil}  \\begin{abstract} Dictated by the string theory and various higher dimensional scenarios, black holes in $D>4$-dimensional space-times must have higher curvature corrections.  The first and dominant term is quadratic in curvature, and  called the Gauss-Bonnet (GB) term. We shall show that although the Gauss-Bonnet correction changes black hole's  geometry only softly, the emission of gravitons is suppressed by many orders even at quite small values of the GB coupling. The huge suppression of the graviton emission is due to the multiplication of the two effects: the quick cooling of the black hole when one turns on the GB coupling and the exponential decreasing of the grey-body factor of the tensor type of gravitons at small and moderate energies. At higher $D$ the tensor gravitons emission is dominant, so that the overall lifetime of black holes with Gauss-Bonnet corrections is many orders larger than was expected. This effect should be relevant for the future experiments at the Large Hadron Collider (LHC). ",
        "introduction": "During the past decade high energy physics received a great impact from theories implying existence of extra dimensions in the world. These are the string theory \\cite{strings} and higher dimensional brane-world scenarios \\cite{BWS}. The low energy limit of the string theory can be described by the slope expansion in powers of the inverse string tension (or of the inverse square of the fundamental string scale $\\ell_{s}^{-2}$) that produces higher curvature corrections to the Einstein action. The quadratic term in curvature (given by the so-called Gauss-Bonnet invariant) is the leading correction that can affect the graviton excitation spectrum near the flat space.   The extra dimensional scenarios also suggest that the fundamental gravity scale $M_{*}$ might be around the weak scale $\\sim TeV$. Thus, at particle collisions with the cross section $\\sim \\pi r_{s}^{2}$, where $r_s$ is the Schwarzschild radius, and energies larger than $M_{*}$, the production of mini-black holes should start. These black holes are intrinsically higher dimensional and usually modeled by the Tangherlini metric, which is the solution of the D-dimensional Einstein equations. However, in order to have a mathematically noncontradictory gravity in higher dimensions, one has to take account of higher curvature corrections of the same form as those appearing in the slope expansion of the string theory. The spherically symmetric solution describing neutral static black holes in the D-dimensional Einstein gravity with the GB corrections was obtained in \\cite{Deser}. This solution contains small corrections to the D-dimensional Schwarzschild-Tangherlini geometry and consequently properties of such Gauss-Bonnet corrected black holes were expected to differ only slightly from the Schwarzschild's ones. This happens, for instance, for the spectrum of proper oscillations of these black holes \\cite{GBQNMs}.  Unlike astrophysical black holes, whose Hawking evaporation is negligibly small, mini-black holes are intensively evaporating what leads to the very short lifetime of these black holes, once they are created. The latter is due-to strong production of various particles from the vacuum around a black hole and emission of them through the mechanism of Hawking radiation \\cite{Hawking:1974sw}. At large number of space-time dimensions $D$, the specific ``tensorial'' type of gravitons  (respectively the $D-2$ rotation group) dominates in the emission process \\cite{Cardoso:2005vb}. Up to now, an impressively extensive literature is devoted to the calculations of Hawking evaporation of the Schwarzschild-Tangherlini and Myers-Perry black holes \\cite{Jung-rot,Chen,charybdis2,IOP2,Nomura}, while evaporation of their higher curvature corrected generalizations was touched upon only in a couple of works \\cite{Rizzo:2006uz,Grain:2005my}. In particular, T.~Rizzo estimated the energy-emission rate for the higher curvature corrected black holes, assuming that the grey-body factor equals unity \\cite{Rizzo:2006uz}. This was expected to give the correct answer about the order of the intensity of the Hawking emission. Though, as we shall show in this paper, the contribution due to the grey-body factors can also considerably change the results. In \\cite{Grain:2005my}, the scattering of Standard Model particles around Gauss-Bonnet black holes was considered, though the calculations were terminated at the grey-body factors and the numbers of particles per frequency. Thus, none of the above works calculated the energy-emission rate for Gauss-Bonnet black holes that is necessary for the estimation of the total emission of energy and thus of the black hole lifetime. Here we shall fill this gap and calculate the energy-emission rates for fields of various spin, including gravitons, and thus shall estimate the lifetime of Gauss-Bonnet black holes.  In this work, we shall show that due to a number of reasons, the emission of the tensor type of gravitons is greatly (in fact exponentially) suppressed when one turns on the GB coupling $\\alpha'$. Thus, even at small values of the GB coupling constant $\\alpha'$ \\emph{the graviton emission is suppressed by many orders}. This means that small GB corrections lead to a much longer life of higher dimensional black holes than was expected \\cite{Cardoso:2005vb,Kanti:2008eq,Kanti:2009sn}. At first sight this enormous suppression would not seem trustworthy: why do slight corrections of geometry produce a very strong effect on the evaporation process? The reasons for this are ``multiplication'' of the two factors. First, the black hole gets much colder when one turns on the GB coupling and the emission rate is quadratic in temperature. Second, the emission is proportional to the grey-body factor which is exponentially suppressed for tensorial gravitons. This explanation certainly did not make us trust the result immediately. Therefore we reproduced our accurate numerical calculations by the semianalytical WKB estimations.  The paper is organized as follows: Sec \\ref{sec:wave-like} briefly discuss the deduction of wave equations for perturbations of fields of various spin. Sec \\ref{sec:methods} is devoted to numerical calculation of the coefficients of transmission, while Sec. \\ref{sec:WKB} gives WKB values of the coefficients. In Sec. \\ref{sec:results} the obtained scattering data are used for the calculations of the energy-emission rates. Using the WKB arguments Sec. \\ref{sec:discussions} explains why the found enormous suppression of Hawking evaporation occurs. In Sec. \\ref{sec:conclusions} we estimate the lifetime of Gauss-Bonnet corrected black holes and outline the future perspective for this direction.  We shall consider the canonical ensemble, which leads to the same results as the microcanonical one if the black hole mass $M$ is at least a few times larger than $M_*$ \\cite{Rizzo:2006uz}.  \\begin{figure*} \\includegraphics[width=.5\\textwidth,clip]{WKBcheck.eps}\\includegraphics[width=.5\\textwidth,clip]{WKBdiff.eps} \\caption{Square root of the reflection coefficient $|Z_o/Z_i|$ for $D=6$, $\\alpha=2$, $l=2$ tensor-type gravitational perturbations. The left panel shows the coefficient calculated by fitting the numerically solved equation (blue) and using the 6-th order WKB formula (red). In the right panel we plot the difference between the coefficients calculated using these two methods.}\\label{fig-WKBcheck} \\end{figure*}  \\begin{figure*} \\includegraphics[width=.5\\textwidth,clip]{GBscalarpot.eps}\\includegraphics[width=.5\\textwidth,clip]{GBscalarfac.eps} \\caption{The effective potential $V(r)$ (left panel) and the square root of the reflection coefficient $|Z_o/Z_i|$ (right panel) for $D=6$, $\\alpha=2$, $l=2$ scalar-type gravitational perturbations.}\\label{fig-scalarpot} \\end{figure*}   \\begin{widetext} ",
        "conclusions": "\\label{sec:conclusions} We have shown that the widely accepted approximation of higher dimensional black holes by their classical Schwarzschild-Tangherlini model is not good, when one considers the Hawking radiation around a black hole. Intensive Hawking emission of gravitons, as well as of other particles, is suppressed by many orders, when one takes into consideration small quantum Gauss-Bonnet corrections. Consequently, the lifetime of quantum corrected black holes is many orders larger than it is expected according to the current literature \\cite{Hawking-raznoe}. This makes further investigations of Hawking radiation of higher curvature corrected black holes appealing."
    },
    "1004/1004.2774_arXiv.txt": {
        "abstract": "In this paper, we combine the the latest observational data, including the WMAP five-year data (WMAP5), the baryon acoustic oscillations (BAO) and type Ia supernovae (SN) ``union\" compilation, and use the Markov Chain Monte Carlo method to determine the dark energy parameters. We pay particular attention to the Integrated Sache-Wolfe (ISW) data from the cross-correlations of cosmic microwave background (CMB) and large scale structure (LSS). In the $\\Lambda$CDM model, we find that the ISW data, as a complement to the WMAP data, could significantly improve the constraint of curvature $\\Omega_k$. We also check the improvement of constraints from the new prior on the Hubble constant and find this new prior could improve the constraint of $\\Omega_k$ by a factor of 2. Finally, we study the dynamical evolving EoS of dark energy from the current observational data. Based on the dynamical dark energy model, parameterizing as $w(a)=w_0+w_a(1-a)$, we find that the $\\Lambda$CDM model remains a good fit to the current data. When taking into account the ISW data, the error bars of $w_0$ and $w_a$ could be shrunk slightly. Current constraints on the dynamical dark energy model are not conclusive. The future precision measurements are needed. ",
        "introduction": "\\label{Int} Unveiling the origin of the current accelerating expansion of our Universe is a big challenge for the modern cosmology either theoretically or observationally. The origin source which drives the expansion could be attributed to a mysterious budget, dark energy. Thus, the nature of dark energy is one of the biggest unsolved problems in modern physics and has been extensively investigated in recent years.  The measurements of CMB \\cite{WMAP5GF1,WMAP5GF2,WMAP5Other}, LSS surveys \\cite{SDSS,2df} and SN \\cite{Union,cfa} have provided a lot of high quality data at present. These data have been widely used to constrain various cosmological models. However, one should keep in mind that the degeneracies of cosmological parameters generally exist in almost all cosmological observations, i.e., they are not sensitive to single parameters but to some specific combinations of them. These degeneracies could weaken constraints on the cosmological parameters. It is therefore highly necessary to combine different probes to break parameter degeneracies so as to achieve tight constraints. Furthermore, different observations are affected by different systematic errors, and it is thus helpful to reduce potential biases by combining different probes.  One of the useful complementary probe is the late-time ISW effect \\cite{sachswolfe67}. This ISW effect is produced by the CMB photons passing through the time-evolving gravitational potential well, when dark energy or curvature becomes important at later times. Therefore, the ISW effect provides a promising probe for studying the acceleration mechanism of our universe, especially for the dark energy and the curvature of Universe. Cross correlating CMB with tracers of LSS surveys for detecting the ISW effect \\cite{Crittenden:1995ak} has been widely investigated in the literature \\cite{Boughn:1997vs,Zaldarriaga:1998te,Hu:2001kj,Song:2002sg, Hu:2001fb,Hu:2001tn,Afshordi:2003xu,Gaztanaga:2004sk, Vielva:2004zg,Pietrobon:2006gh,McEwen:2006my,Giannantonio:2006du, Rassat:2006kq,Ho:2008bz,Granett:2008ju,Xia:2009dr,Pogosian:2005ez}.  In this paper, we will present the constraints on various cosmological models from the current observations, including the WMAP5 data, SN ``union\" compilation, and recently released BAO data from SDSS DR7, as well as the ISW data. The structured of the paper is as following: in Section II we describe the method and the data sets; in Section III we present our numerical results and discussion; finally we give a summary and outlook in Section VI.   \\begin{figure*}[htbp] \\begin{center} \\includegraphics[scale=0.7]{lcdm_omk_omx_1D.eps} \\caption{One dimensional distributions of $\\Omega_k$ and $\\Omega_{\\Lambda}$ from different data combinations: WMAP5 (balck solid lines), WMAP5+ISW (red dashed lines), WMAP5+HST (blue dash-dotted lines), and WMAP5+HST+ISW (purple dotted lines). \\label{fig1}} \\end{center} \\end{figure*} ",
        "conclusions": "\\label{Sum}  In this paper, we study the constraints on cosmological parameters from the recently released CMB, BAO and SNIa data. Here, we pay particular attention to the current ISW data which is the cross correlations of CMB with LSS surveys. In the $\\Lambda$CDM model, the ISW data and the new HST prior are very helpful to break the degeneracy between $\\Omega_k$ and $\\Omega_\\Lambda$. The constraints on $\\Omega_k$ and $\\Omega_\\Lambda$ significantly improve, when compared to the constraints from WMAP5 alone.  More importantly, we consider the constraints on the dark energy parameters in the CPL dark energy model. Here, we fully include the perturbations of dark energy. This result implies that the dynamical dark energy models are not excluded, the $\\Lambda$CDM model, however, is still a good fit right now.  We find that the ISW data could give slight improvement of the constraints on CPL dark energy model. But this does not mean that the ISW data can not constrain the dynamical dark energy models efficiently. Actually, it should be useful for constraining the dynamical dark energy models whose EoS $w(z)$ deviate from the cosmological constant boundary obviously, for example, the early dark energy models \\cite{Xia:2009dr}, or the dark energy model with its EoS $w(z)$ transits sharply during its evolution. Besides the dark energy models, the ISW data could be also helpful for testing the modified gravity theories \\cite{Jain:2007yk,Afshordi:2004kz,Zhang:2005vt,Schmidt:2007vj}, massive neutrinos \\cite{Lesgourgues:2007ix}, the primordial non-gaussianity \\cite{slosar,xia10}, and so on.  Furthermore, we compare the constraints between from the new HST prior and from the old one. One can see that the new HST prior gives the tighter constraints on the cosmological parameters. Finally, we check the capability of current observational data to constrain the dark energy sound speed $c_s^2$. We find that the sound speed is weakly constrained by current observations, and thus futuristic precision measurements of the CMB on a very large angular scale (low multipoles) are necessary."
    },
    "1004/1004.2318_arXiv.txt": {
        "abstract": "We summarize and critically evaluate the available data on nuclear fusion cross sections important to energy generation in the Sun and other hydrogen-burning stars and to solar neutrino production.   Recommended values and uncertainties are provided for key cross sections, and a recommended spectrum is given for $^8$B solar neutrinos. We also discuss opportunities for further increasing the precision of key rates, including new facilities, new experimental techniques, and improvements in theory.  This review, which summarizes the conclusions of a workshop held at the Institute for Nuclear Theory, Seattle, in January 2009, is intended as a 10-year update and supplement to Reviews of Modern Physics {\\bf 70} (1998) 1265. ",
        "introduction": "\\label{sec:intro} In 1998 the Reviews of Modern Physics published a summary and critical analysis of the nuclear reaction cross sections important to solar burning.  That effort, \\citet{Adel98} and denoted here as Solar Fusion I, began with a meeting hosted by the Institute for Nuclear Theory, University of Washington, 17-20 February 1997. A group of international experts in the nuclear physics and astrophysics of hydrogen-burning stars met to begin critical discussions of the existing data on relevant nuclear reactions, with the aim of determining ``best values\" and uncertainties for the contributing low-energy S-factors.  The group also considered opportunities for further improvements in both measurements and theory.  Such data and related nuclear theory have been crucial to the standard solar model (SSM) and the neutrino fluxes it predicts.  Indeed, measurements of nuclear reactions gave the field its start.  In 1958 \\citet{HJ58,HJ59} showed that the rate for $^3$He+$^4$He $\\rightarrow$ $^7$Be +$\\gamma$ was $\\sim$ 1000 times larger than expected, and thus that the pp chain for $^4$He synthesis would have additional terminations beyond $^3$He+$^3$He $\\rightarrow$ $^4$He + 2p.  This result led Davis to recognize that his chlorine detector might be able to see the higher energy neutrinos from these other terminations, and spurred Bahcall and others to develop a quantitative model of the Sun capable of predicting those fluxes \\cite{bd}.  At the time of the 1997 meeting, three decades of effort in solar neutrino physics had produced four measurements that were at variance with the SSM and the standard model of electroweak interactions.  The measurements came from the pioneering work of Ray Davis, Jr. \\cite{Davis68, Davis94}; the observation of $^8$B neutrinos in the Kamiokande water Cerenkov detector \\cite{Kamiokande96}; and the GALLEX \\cite{Kirsten03} and SAGE \\cite{Gavrin03} radiochemical detectors sensitive primarily  to pp and $^7$Be neutrinos.  The resulting pattern of fluxes that emerged from these experiments was difficult to reconcile with any plausible variation in the SSM, requiring a much sharper reduction in the $^7$Be neutrino flux than in the $^8$B flux, despite the greater sensitivity of the latter to changes in the solar core temperature.  For this reason it was argued in Solar Fusion I that the measurements provided evidence for new physics beyond the standard model.  New solar neutrino experiments that promised much more precise data -- the 50-kiloton successor to Kamiokande, Super-Kamiokande, and the heavy-water-based Sudbury Neutrino Observatory (SNO), with sensitivity to both electron and heavy-flavor neutrinos -- were then underway.   The authors of Solar Fusion I, recognizing that the impact of these new experiments would depend in part on the quality of the nuclear microphysics input to the SSM, thus undertook an extended study of the key reaction rates for the pp chain and CNO bi-cycle. The effort appears to have been of some value to the community, as Solar Fusion I has become one of the most heavily cited papers in nuclear astrophysics.  \\subsection{Solar Fusion II: the 2009/10 effort} Ten years after publication of Solar Fusion I a proposal was made to the INT to revisit this process, in order to produce a new evaluation that would reflect the considerable progress made in the past decade, as well as new motivations for further constraining the SSM.  Examples of advances in the nuclear physics include the LUNA II program at Gran Sasso \\cite{LUNAII}, which has provided remarkable low-energy measurements of key reactions such as $^3$He($\\alpha$,$\\gamma$)$^7$Be and $^{14}$N(p,$\\gamma$)$^{15}$O; several high-precision measurements addressing the key pp-chain uncertainty identified in Solar Fusion I, $^7$Be(p,$\\gamma$)$^8$B; the application of new theoretical techniques to the p+p and hep neutrino reactions; and the resolution of several unresolved questions about screening corrections in plasmas.  The context for these measurements has also changed.  In 1997 the field's central concern was, in some sense, a qualitative one, the origin of the solar neutrino problem.  This question was answered in spectacular fashion by the dual discoveries of Super-Kamiokande \\cite{SK} and SNO \\cite{SNO} -- two distinct neutrino oscillations  responsible for the missing atmospheric and solar neutrinos, largely determining the pattern of the light neutrino masses.   But issues remain, and most of these require precision.  There is intense interest in extending direct measurements to the low-energy portion of the solar neutrino spectrum ($\\lsim$ 2 MeV), where experiments with good energy resolution can determine the separate contributions of pep, CNO, $^7$Be, and pp neutrinos.  There is the potential to further constrain the solar neutrino mixing angle $\\theta_{12}$:  the solar luminosity determines the pp flux to high accuracy, and the low-energy spectrum lies in the vacuum region of the MSW triangle, in contrast to the high-energy $^8$B neutrinos, where matter effects are significant.  Thus precise low-energy measurements have considerable ``leverage\" to test $\\theta_{12}$ and the consistency of the conclusions we have drawn from SNO, Super-Kamiokande, and the KamLAND reactor neutrino experiment.  Borexino, now entering its calibration phase, is the first effort in this program of high-precision spectroscopy of low-energy solar neutrinos.  But the resolution of the solar neutrino problem has also returned the field to its roots:  Davis built the chlorine detector to probe the interior of the Sun and thereby test directly the theory of stellar evolution and nuclear energy generation \\cite{bd}. Davis was diverted from that goal by the missing solar neutrinos.  But as the weak interaction effects responsible for that anomaly are now reasonably well understood, solar neutrinos again have become a quantitative tool for astronomy. Indeed, the program carried out by SNO and Super-Kamiokande has already yielded one remarkable constraint on the Sun, a direct determination of the core temperature to high precision, through measurement of the $^8$B neutrino flux ($\\phi(^8$B) $\\propto T_c^{22}$).  The 8.6\\%  precision of the SNO NCD-phase results \\cite{NCD}, $\\phi(^8$B) = $(5.54 {}^{+0.33}_{-0.31} {}^{+0.36}_{-0.34}) \\times 10^6$/cm$^2$/s, implies a sensitivity to core temperature of $\\sim$ 0.5\\%.  New questions have arisen about the Sun that neutrinos could potentially address, provided the associated laboratory astrophysics has been done. One important success of the SSM in the 1990s was in predicting the local sound speed $c(r)$.  Comparisons between $c(r)$ deduced from helioseismology and the predictions of the SSM yielded agreement at $\\sim$ 0.2\\% throughout much of the Sun.  Bahcall and others argued \\cite{Bahcall2001} that helioseismology is a more severe and detailed test of the SSM than neutrino production, so that SSM success in reproducing $c(r)$ made a particle-physics resolution of the solar neutrino problem more likely.  The sound speed is a function of the Sun's interior pressure and density profiles, which in turn reflect thermal transport properties that depend on the Sun's metal content, through the opacity. Thus the comparison between helioseismology and the SSM tests a key assumption of the SSM, that the metals are distributed uniformly throughout the Sun, apart from small corrections due to diffusion.  This assumption allows one to equate SSM interior metal abundances to convective-zone abundances deduced from analyses of photospheric absorption lines.  Such analyses had been based on 1D models of the photosphere.  Recently {\\it ab initio} 3D analyses have been developed, yielding significant improvements in predicted line shapes and in the consistency of metal abundance determinations from various atomic and molecular lines.  However, this work also reduced metallicity estimates from Z $\\sim$  0.0169 to $\\sim$ 0.0122 \\cite{Asplund}, destroying the once excellent agreement between helioseismology and the SSM.  It has been suggested that this difficulty may reflect, contrary to the SSM, differences in solar core and convective-zone metallicities that could have arisen from the late-stage evolution of the solar disk: as a great deal of metal was scoured out of the disk by the formation of the giant planets, the last few percent of gas deposited onto the Sun could have been depleted of metals \\cite{hax08}.  Indeed, recent studies of ``solar twins\" show abundance trends that correlate with the existence of planets \\cite{asplanet,liplanet}.  \\citet{hax08} argued that a direct measurement of solar core metallicity could be made by observing CNO solar neutrinos.  In both of the above examples -- using neutrinos to determine the solar core temperature and metallicity -- nuclear physics uncertainties remain one of the limiting factors in the analyses.  The proposal to revisit in 2009 the deliberations of 1997 thus had several motivations: \\begin{itemize} \\item providing a set of standard S-factors and uncertainties that reflect the progress made in laboratory and theoretical nuclear astrophysics over the last decade; \\item enabling more precise analyses of solar neutrino experiments designed to constrain neutrino oscillations and other new physics, e.g., future pp and pep neutrino experiments that exploit these well understood fluxes;  and \\item enabling analyses in which solar neutrinos are used as a probe of the solar core. \\end{itemize} The 2009 INT workshop\\footnote[1]{The workshop was proposed in a letter to the Institute for Nuclear Theory's National Advisory Committee (NAC) and approved by the NAC and INT Director at the time of the NAC's August 2008 annual meeting.  Wick Haxton (lead), Eric Adelberger, Heide Costantini, Peter Parker, R. G. Hamish Robertson, Kurt Snover, Frank Strieder, and Michael Wiescher formed the organizing committee and served as co-editors of this paper. Additional community members joined this group to act as working group heads: Jiunn-Wei Chen, Barry Davids, Stuart Freedman, Alejandro Garcia, Uwe Greife, Michael Hass, Gianluca Imbriani, Kuniharu Kubodera, Daniela Leitner, Laura Marcucci, Filomena Nunes, Tae-Sun Park, Paolo Prati, Hanns-Peter Trautvetter, and Stefan Typel.  The working group heads were responsible for organizing discussions, creating section drafts, and responding to subsequent criticisms of the drafts.  Organizing committee members, in their capacity as co-editors, were responsible for creating from the drafts a coherent document, and for addressing any issues unresolved by the working groups. Workshop presentations are archived on the INT's web site, http://www.int.washington.edu/PROGRAMS/solar\\_fusion.html.} was modeled after that of 1997, with invitations extended to and accepted by representatives from most of the experimental groups active in the nuclear physics of hydrogen burning stars.  There was also active involvement of theorists, reflecting the progress that has been made in {\\it ab initio} calculations. The workshop participants are the authors of this manuscript.  As in 1997, early organizing included the selection of working group leaders who identified key papers, which were then entered in a database for review, prior to the start of the workshop.   These materials were then summarized and discussed during the workshop, as the various working groups considered the state of the data and outlined any additional work that would be needed for this review. The process of critically analyzing both new and older data and working toward a consensus on best-value cross sections and uncertainties continued throughout 2009. A few new topics not considered in 1997 but now recognized to be quite important, such as the shape of the $^8$B neutrino spectrum, were addressed. (The $^8$B neutrino spectrum is one of the inputs to SNO and Super-Kamiokande analyses.) The workshop included working groups on indirect techniques for constraining cross sections, to summarize the progress that has been made in validating such approaches, and on new facilities and instrumentation, in view of the facility investments that are being considered in laboratory nuclear astrophysics (above and below ground).  \\subsection{Contents of this review} The review begins in Section II with a description of hydrogen burning by the pp chain and CNO bi-cycle, and the neutrino byproducts of these reaction chains.  The role of S-factors and the associated questions of screening and of extrapolating data to the solar Gamow peak are discussed.  We provide a fairly complete overview of progress in theory, which in some cases provides our only estimate of S-factors, and in other cases determines the forms of the functions that are needed for data extrapolations.  Discussions of individual reactions are organized by chapter: Secs. III-IX discuss the pp chain reactions p+p $\\rightarrow$ d+e$^+$+$\\nu_e$; d+p $\\rightarrow$ $^3$He+$\\gamma$; $^3$He+$^3$He $\\rightarrow$ $^4$He+p+p; $^3$He+$^4$He $\\rightarrow$ $^7$Be+$\\gamma$; $^3$He+p $\\rightarrow$ $^4$He+e$^+$+$\\nu_e$; $^7$Be, pp, and CNO nuclei electron capture; and $^7$Be+p $\\rightarrow$ $^8$B+$\\gamma$. Sec. X discusses the spectrum of $^8$B neutrinos produced in the $\\beta$ decay to a broad resonance in $^8$Be. Sec. XI discusses $^{14}$N+p $\\rightarrow$ $^{15}$O+$\\gamma$ and other reactions contributing to the CNO cycles.  Sec. XII describes the progress that has been made in developing and validating indirect methods, while Sec. XIII describes future facilities and instrumentation that could further advance the field.  The conclusions of this review, in some cases, required the working groups to make some judgments.  There are discrepant data sets, and there are cases where data extrapolations have some dependence on models.   We have tried to treat such questions as consistently as possible, aware that excessively optimistic treatments of uncertainties could be misleading, while excessively conservative treatments would degrade the value of the best experiments done in the field. In most cases our working groups were able to reach consensus.  In cases where significant differences remained among the experts, we have tried to identify the source of the disagreement, so that ``consumers\" will be aware that full consensus may have to await future measurements.  Table \\ref{tab:summary} summarizes the conclusions of this review.  \\begin{table*} \\caption{The Solar Fusion II recommended values for S(0), its derivatives, and related quantities, and for the resulting uncertainties on S($E$) in the region of the solar Gamow peak -- the most probable reaction energy -- defined for a temperature of 1.55 $\\times$ 10$^7$K characteristic of the Sun's center.  See the text for detailed discussions of the range of validity for each S($E$). Also see Sec. \\ref{sec:ec} for recommended values of CNO electron capture rates, Sec. \\ref{sec:N114other} for other CNO S-factors, and Sec. \\ref{sec:spectrum} for the $^8$B neutrino spectral shape. Quoted uncertainties are 1$\\sigma$.} \\label{tab:summary} \\begin{tabular}{lccccc} \\hline\\hline Reaction~~~~~~ & ~~~~~Section~~~~~&~~~~~S(0)~~~~~ & ~~~~~S$^\\prime$(0)~~~~~ & ~~~~~S$^{\\prime \\prime}$(0)~~~~~& ~~~~~~Gamow peak~~~~~~ \\\\ &  & (keV-b) & (b) & (b/keV) & uncertainty (\\%) \\\\ \\hline p(p,e$^+\\nu_e$)d & \\ref{sec:s11} & (4.01 $\\pm$ 0.04)$\\times$10$^{-22}$ & (4.49 $\\pm$ 0.05)$\\times$10$^{-24}$ & $-$ & $\\pm$ 0.7 \\\\ d(p,$\\gamma$)$^3$He & \\ref{sec:dp} & (2.14$^{+0.17}_{-0.16}$)$\\times$10$^{-4}$& $(5.56^{+0.18}_{-0.20})\\times$10$^{-6}$& $(9.3^{+3.9}_{-3.4})\\times$10$^{-9}$ & ~~$\\pm$ 7.1~\\footnote{Error from phenomenological quadratic fit.  See text.} \\\\ ${}^3$He(${}^3$He,2p)${}^4$He & \\ref{sec:s33}& (5.21 $\\pm$ 0.27) $\\times$ 10$^3$ & $-$4.9 $\\pm$ 3.2 & (2.2 $\\pm$ 1.7) $\\times$ 10$^{-2}$ & ~~$\\pm$ 4.3~$^a$ \\\\ ${}^3$He(${}^4$He,$\\gamma$)${}^7$Be &\\ref{sec:s34} &0.56 $\\pm$ 0.03 & ($-$3.6 $\\pm$ 0.2)$\\times$10$^{-4}$~\\footnote{S$^\\prime$(0)/S(0) taken from theory; error is that due to S(0).  See text.} & (0.151 $\\pm$ 0.008)$\\times$10$^{-6}$~\\footnote{S$^{\\prime \\prime}$(0)/S(0) taken from theory; error is that due to S(0).  See text.} &  $\\pm$ 5.1\\\\ ${}^3$He(p,e$^+\\nu_e$)${}^4$He &\\ref{sec:hep} & (8.6 $\\pm$ 2.6)$\\times$10$^{-20}$ &$-$ &$-$ & $\\pm$ 30~ \\\\ ${}^7$Be(e$^-,\\nu_e$)${}^7$Li &\\ref{sec:ec} & See Eq.~(\\ref{eq:bec})& $-$& $-$ & $\\pm$ 2.0 \\\\ p(pe$^-$,$\\nu_e$)d & \\ref{sec:ec} & See  Eq.~(\\ref{eq:pep3}) &$-$  &$-$ & ~~$\\pm$ 1.0~\\footnote{Estimated error in the pep/pp rate ratio. See Eq.~(\\ref{eq:pep3})} \\\\ ${}^7$Be(p,$\\gamma$)${}^8$B & \\ref{sec:s17} & (2.08 $\\pm$ 0.16)$\\times$10$^{-2}$~\\footnote{Error dominated by theory.} &($-$3.1 $\\pm$ 0.3)$\\times$10$^{-5}$ & (2.3 $\\pm$ 0.8)$\\times$10$^{-7}$ & $\\pm$ 7.5 \\\\ ${}^{14}$N(p,$\\gamma$)${}^{15}$O & \\ref{sec:N114} & 1.66 $\\pm$ 0.12 & ($-$3.3 $\\pm$ 0.2)$\\times$10$^{-3}$~$^b$ & (4.4 $\\pm$ 0.3)$\\times$10$^{-5}$~$^c$ & $\\pm$ 7.2\\\\ \\hline\\hline \\end{tabular} \\end{table*} ",
        "conclusions": "\\begin{tabular}{lccccc} \\hline\\hline Reaction~~~~~~ & ~~~~~Section~~~~~&~~~~~S(0)~~~~~ & ~~~~~S$^\\prime$(0)~~~~~ & ~~~~~S$^{\\prime \\prime}$(0)~~~~~& ~~~~~~Gamow peak~~~~~~ \\\\ &  & (keV-b) & (b) & (b/keV) & uncertainty (\\%) \\\\ \\hline p(p,e$^+\\nu_e$)d & \\ref{sec:s11} & (4.01 $\\pm$ 0.04)$\\times$10$^{-22}$ & (4.49 $\\pm$ 0.05)$\\times$10$^{-24}$ & $-$ & $\\pm$ 0.7 \\\\ d(p,$\\gamma$)$^3$He & \\ref{sec:dp} & (2.14$^{+0.17}_{-0.16}$)$\\times$10$^{-4}$& $(5.56^{+0.18}_{-0.20})\\times$10$^{-6}$& $(9.3^{+3.9}_{-3.4})\\times$10$^{-9}$ & ~~$\\pm$ 7.1~\\footnote{Error from phenomenological quadratic fit.  See text.} \\\\ ${}^3$He(${}^3$He,2p)${}^4$He & \\ref{sec:s33}& (5.21 $\\pm$ 0.27) $\\times$ 10$^3$ & $-$4.9 $\\pm$ 3.2 & (2.2 $\\pm$ 1.7) $\\times$ 10$^{-2}$ & ~~$\\pm$ 4.3~$^a$ \\\\ ${}^3$He(${}^4$He,$\\gamma$)${}^7$Be &\\ref{sec:s34} &0.56 $\\pm$ 0.03 & ($-$3.6 $\\pm$ 0.2)$\\times$10$^{-4}$~\\footnote{S$^\\prime$(0)/S(0) taken from theory; error is that due to S(0).  See text.} & (0.151 $\\pm$ 0.008)$\\times$10$^{-6}$~\\footnote{S$^{\\prime \\prime}$(0)/S(0) taken from theory; error is that due to S(0).  See text.} &  $\\pm$ 5.1\\\\ ${}^3$He(p,e$^+\\nu_e$)${}^4$He &\\ref{sec:hep} & (8.6 $\\pm$ 2.6)$\\times$10$^{-20}$ &$-$ &$-$ & $\\pm$ 30~ \\\\ ${}^7$Be(e$^-,\\nu_e$)${}^7$Li &\\ref{sec:ec} & See Eq.~(\\ref{eq:bec})& $-$& $-$ & $\\pm$ 2.0 \\\\ p(pe$^-$,$\\nu_e$)d & \\ref{sec:ec} & See  Eq.~(\\ref{eq:pep3}) &$-$  &$-$ & ~~$\\pm$ 1.0~\\footnote{Estimated error in the pep/pp rate ratio. See Eq.~(\\ref{eq:pep3})} \\\\ ${}^7$Be(p,$\\gamma$)${}^8$B & \\ref{sec:s17} & (2.08 $\\pm$ 0.16)$\\times$10$^{-2}$~\\footnote{Error dominated by theory.} &($-$3.1 $\\pm$ 0.3)$\\times$10$^{-5}$ & (2.3 $\\pm$ 0.8)$\\times$10$^{-7}$ & $\\pm$ 7.5 \\\\ ${}^{14}$N(p,$\\gamma$)${}^{15}$O & \\ref{sec:N114} & 1.66 $\\pm$ 0.12 & ($-$3.3 $\\pm$ 0.2)$\\times$10$^{-3}$~$^b$ & (4.4 $\\pm$ 0.3)$\\times$10$^{-5}$~$^c$ & $\\pm$ 7.2\\\\ \\hline\\hline \\end{tabular} \\end{table*}"
    },
    "1004/1004.3364_arXiv.txt": {
        "abstract": "In this Letter, a modified Chaplygin gas (MCG) model of unifying dark energy and dark matter with the exotic equation of state $p_{MCG}=B\\rho_{MCG} -\\frac A{\\rho_{MCG}^\\alpha }$ is constrained from recently observed data: the 182 Gold SNe Ia, the 3-year WMAP and the  SDSS baryon acoustic peak. It is shown that the best fit value of  the three parameters ($B$,$B_{s}$,$\\alpha$) in MCG model are (-0.085,0.822,1.724). Furthermore, we find the best fit $w(z)$ crosses -1 in the past and the present best fit value $w(0)=-1.114<-1$, and the $1\\sigma$ confidence level of $w(0)$ is $-0.946\\leq w(0)\\leq-1.282$. Finally, we find that the MCG model has the smallest $\\chi^{2}_{min}$ value in all eight given models. According to the Alaike Information Criterion (AIC) of model selection, we conclude that recent observational data support the MCG model as well as other popular models. ",
        "introduction": "{\\small {~~~}}~The type Ia supernova (SNe Ia) explorations \\cite{[1]}, the cosmic microwave background(CMB) results from WMAP \\cite{[2]} observations, and surveys of galaxies \\cite{[3]} all suggest that the universe is speeding up rather than slowing down. The accelerated expansion of the present universe is usually attributed to the fact that dark energy is an exotic component with negative pressure. Many kinds of dark energy models have already been constructed such as $\\Lambda$CDM \\cite{[4]}, quintessence \\cite{[5]}, phantom \\cite{[6]}, generalized Chaplygin gas (GCG) \\cite{[7]}, quintom \\cite{[8]}, holographic dark energy \\cite{[9]}, and so forth.  On the other hand, to remove the dependence of special properties of extra energy components, a parameterized equation of state (EOS) is assumed for dark energy. This is also commonly called the model-independent method. The parameterized EOS of dark energy which is popularly used in parameter best fit estimations, describes the possible evolution of dark energy. For example, $w=w_{0}$=const \\cite{[10]}, $w(z)=w_{0}+w_{1}z$ \\cite{[11]}, $w(z)=w_{0}+\\frac{w_{1}z}{1+z}$ \\cite{[12]}, $w(z)=w_{0}+\\frac{w_{1}z}{(1+z)^{2}}$ \\cite{[13]},  $ w(z)=\\frac{1+z}{3}\\frac{A_{1}+2A_{2}(1+z)}{X}-1$ (here $X\\equiv A_{1}(1+z)+A_{2}(1+z)^{2}+(1-\\Omega_{0m}-A_{1}-A_{2})$) \\cite{[14]}. The parameters $w_{0}$, $w_{1}$, or $A_{1}$, $A_{2}$ are obtained by the best fit estimations from cosmic observational datasets.   It is well known that the GCG model has been widely used to interpret the accelerating universe. In the GCG approach, dark energy and dark matter can be unified by using an exotic equation of state. Also, a Modified Chaplygin gas (MCG) as a extension of the generalized Chaplygin gas  model has already been applied to describe the current accelerating expansion of the universe \\cite{[15]} \\cite{[16]} \\cite{[17]} \\cite{[18]}. The constraint on parameter B in MCG model, i.e., the added parameter relative to GCG model, is discussed briefly by using the location of the peak of the CMB radiation spectrum in Ref. \\cite{[19]}. In this Letter, we study the constraints on the best fit parameters ($B,B_{s}$,$\\alpha$) and EOS in the MCG model from recently observed data: the latest observations of the 182 Gold type Ia Supernovae (SNe) \\cite{[20]}, the 3-year WMAP CMB shift parameter \\cite{[21]} and the baryon acoustic oscillation (BAO) peak from Sloan Digital Sky Surver (SDSS) \\cite{[22]}. The result of this study indicates that the best fit value of parameters ($B,B_{s}$,$\\alpha$) in MCG model are (-0.085,0.822,1.724). Furthermore, we find the best fit $w(z)$ crosses -1 in the past and the present best fit value $w(0)=-1.114<-1$, and the $1\\sigma$ confidence level of $w(0)$ is $-0.946\\leq w(0)\\leq-1.282$. At last, because the emphasis of the ongoing and forthcoming research is shifting from estimating specific parameters of the cosmological model to model selection  \\cite{[23]}, it is interesting to estimate which model for an accelerating universe is distinguish by statistical analysis of observational datasets out of a large number of cosmological models. Therefore, by applying the recent observational data to the Alaike Information Criterion (AIC) of model selection, we compare the MCG model with other seven general cosmological models to see which model is better. It is found that the MCG model has almost the same support from the data as other popular models. In the Letter, we perform an estimation of model parameters using a standard minimization procedure based on the maximum likelihood method.  The Letter is organized as follows. In section 2, the MCG model is introduced  briefly. In section 3, the best fit value of parameters ($B,B_{s}$,$\\alpha$) in the MCG model are given from the recent observations of SNe Ia, CMB and BAO, and we present the evolution of the best fit of $w(z)$ with $1\\sigma$ confidence level with respect to redshift $z$. The preferred cosmological model is discussed in section 4 according to  the AIC. Section 5 is the conclusion.\\\\ ",
        "conclusions": "~~~~In summary, the constraints on the MCG model, proposed as a candidate of the unified dark matter-dark energy scenario, has been studied in this Letter. We obtained the best fit value of the three parameters ($B$,$B_{s}$,$\\alpha$) in the MCG model (-0.085,0.822,1.724). Meanwhile, it is easy to see that the best fit $w(z)$ can cross -1 as it evolves with the redshift $z$, and the present best fit value $w(0)=-1.114<-1$. Furthermore, it is shown that the $1\\sigma$ confidence level of $w(0)$ is $-0.946\\leq w(0)\\leq-1.282$, and the possibility of $w(0)>-1$ cann$^{,}$t be excluded in $1\\sigma$ level. We can see that the cosmological constant model (i.e.,$w(z)=-1$) is not in $1\\sigma$ confidence contour of the best fit dynamical $w(z)$.  Finally, in order to find the status of MCG scenario in a large number of cosmological models, we compared the MCG model with other seven popular ones offering explanation of current acceleration of the universe in terms of the values of $\\chi^{2}_{min}$ and AIC quantity. We find that, as the quantity $\\chi^{2}_{min}$ measures the quality of model fit, the MCG model is preferred  by recent observational data because of its a small minimum $\\chi^{2}$ value. On the other hand, it is shown that the MCG model has a slightly high value of AIC due to its many parameters. However, according to the rules of judgment of the AIC model selection, we conclude that recently observed data supports the MCG model as well as other popular models, because the value of $\\bigtriangleup_{i}$ for it is in the range 0-2 relative to the best model. In addition, the result of study shows that the recent observational data equivalently supports all of the models in Table 2 except for the case of $w(z)=w_{0}+w_{1}z$. We expect the new probers such as SNAP and Planck surveyor can provide more accurate data and further explore the nature of dark energy.       \\textbf{\\ Acknowledgments } The research work is supported by NSF (10573003), NSF (10747113), NSF (10573004), NSF (10703001), NSF (10647110), NBRP (2003CB716300) and DUT(893326) of P.R. China."
    },
    "1004/1004.0960_arXiv.txt": {
        "abstract": "We use the Om statistic and the Genetic Algorithms (GA) in order to derive a null test on the spatially flat cosmological constant model $\\Lambda$CDM. This is done in two steps: first, we apply the GA to the Constitution SNIa data in order to acquire a model independent reconstruction of the expansion history of the Universe $H(z)$ and second, we use the reconstructed $H(z)$ in conjunction with the Om statistic, which is constant only for the $\\Lambda$CDM model, to derive our constraints. We find that while $\\Lambda$CDM is consistent with the data at the $2\\sigma$ level, some deviations from $\\Lambda$CDM model at low redshifts can be accommodated. ",
        "introduction": "In the previous decade it was discovered that the Universe is undergoing an accelerated expansion (Riess et al. 2004; Spergel et al. 2007; Readhead et al. 2004). This acceleration is usually attributed either to a cosmic fluid with negative pressure dubbed Dark Energy or to an IR modification of gravity. In order to identify the properties of Dark Energy or the structure of the IR modification of gravity it is necessary to know to a high precision the rate of the expansion of the Universe, parameterized as $H\\equiv\\frac{\\dot{a}}{a}$ where $a=\\frac{1}{1+z}$ is the scale factor and $z$ is the redshift of the cosmological probe, as it measured by the observations.  The behavior of the expansion of the Universe can be identified by studying two functions, the Equation of State (EoS) $w(z)\\equiv \\frac{P}{\\rho}$ which can be rewritten as \\be w(z)\\,=\\,-1\\,+\\frac{1}{3}(1+z)\\frac{d\\ln (\\delta H(z)^2)}{d z}\\,,{\\label{wzh1}} \\ee where $\\delta H(z)^2=H(z)^2/H_0^2-\\Omega_{\\rm 0m} (1+z)^3$ accounts for all terms in the Friedmann equation not related to matter and the deceleration parameter $q(z)\\equiv -\\frac{\\ddot{a}}{a \\dot{a}^2}$ which can be rewritten as \\be q(z)\\,=\\,-1\\,+(1+z) \\frac{d\\ln (H(z))}{d z}\\,,{\\label{qzh1}} \\ee Obviously, the cosmological constant ($w(z)=-1$) corresponds to a constant dark energy density, while in general $w(z)$ can be time dependent. Also, an important parameter is the value of the deceleration parameter today, ie $q(z=0)\\equiv q_0$, which for the cosmological constant model in GR it is $q_0=-1+3 \\Omega_{\\rm 0m}/2$.  However, despite all the recent progress the origin of the accelerated expansion of the universe still remains unknown with many possibilities still remaining open, see for example Perivolaropoulos (2006). The simplest choice that agrees well with the data is a positive cosmological constant which has to be small enough to have started dominating the universe at late times. As it was demonstrated by the Seven-Year WMAP data (Komatsu et al. 2010), the cosmological constant remains the best candidate and has the advantage of having only one free parameter related to the properties of the Dark Energy. Nonetheless, this model fails to explain why the cosmological constant is so small that it can only dominate the universe at late times, a problem known as the {\\it coincidence problem} and there are a few cosmological observations which differ from its predictions (Perivolaropoulos 2008; Perivolaropoulos \\& Shafieloo 2008).  A very important complication in the investigation of the behavior of dark energy occurs due to the bias introduced by the parameterizations used. At the moment, there is a multitude of available phenomenological ans\\\"atze for the dark energy equation of state parameter $w$ or dark energy density, each with its own merits and limitations (see Sahni \\& Starobinsky (2006) and references therein). The interpretation of the SNIa data has been shown to depend greatly on the type of parametrization used to perform a data fit (Sahni \\& Starobinsky 2006; Shafieloo, Sahni \\& Starobinsky 2009). Choosing a priori a model for dark energy can thus adversely affect the validity of the fitting method and lead to compromised or misleading results.  The need to counteract this problem paved the way for the consideration of a complementary set of non-parametric reconstruction techniques (Daly \\& Djorgovski 2003; Wang \\& Mukherjee 2004; Saini 2003; Wang 2009; Clarkson \\& Zunckel) and model independent approaches (Sahni et al. 2002; Alam et al. 2003; Shafieloo et al. 2006, Shafieloo 2007; Wang \\& Tegmark 2005; Shafieloo \\& Clarkson 2010). These try to minimize the ambiguity due to a possibly biased assumption for $w$ by fitting the original datasets without using any parameters related to some specific model. The result of these methods can then be interpreted in the context of a dark energy model of choice. Non-parametric reconstructions can thus corroborate parametric methods and provide more credibility. However, they too suffer from a different set of problems, mainly the need to resort to differentiation of noisy data, which can itself introduce great errors.  In this paper, we present a method that can be used as a model independent approach in testing the standard cosmological model. This is done by using the Genetic Algorithms (GA) technique, first used in the analysis of SNIa data in Bogdanos \\& Nesseris (2009). The GAs represent a method for non-parametric reconstruction of the dark energy equation of state parameter $w$, based on the notions of genetic algorithms and grammatical evolution. GAs are more useful and efficient than usual techniques when \\begin{itemize} \\item The parameter space is very large, too complex or not enough understood, as is the case with dark energy. \\item Domain knowledge is scarce or expert knowledge is difficult to encode to narrow the search space. \\item Traditional search methods give poor results or completely fail. \\end{itemize} Naturally, therefore, they have been used with success in many fields where one of the above situations is encountered, like the computational science, engineering and economics. Recently, they have also been applied to study high energy physics (Becks, Hahn \\& Hemker 1994; Allanach, Grellscheid \\& Quevedo 2004; Rojo \\& Latorre 2004) gravitational wave detection (Crowder, Cornish \\& Reddinger 2006) and gravitational lensing (Brewer \\& Lewis 2005). Since the nature of Dark Energy still remains a mystery, this makes it for us an ideal candidate to use the GAs as a means to analyze the SNIa data and extract model independent constraints on the behavior of the Dark Energy. In Section 2 we briefly describe the Om statistic while in Section 3 we provide an overview of the general methodology of the GA paradigm and finally, we present our results in Section 4. ",
        "conclusions": "We used the Om statistic and the GAs in order to derive a null test on the cosmological constant model $\\Lambda$CDM. Our interest in the GAs stems from the fact that they represent a method for non-parametric reconstruction of the expansion history of the universe, based on the notions of genetic algorithms and grammatical evolution. These kinds of algorithms are more useful and efficient than usual techniques especially when the problem under examination is not well understood, as is the case with dark energy. Since the nature of dark energy still remains a mystery, this makes it for us an ideal candidate to use the GAs as a means to analyze the SNIa data and extract model independent constraints on the behavior of the Dark Energy. On the other hand, the $Om$ diagnostic (Sahni et al. 2008) enables us to distinguish $\\Lambda$CDM from other dark energy models without directly involving the cosmic EoS.  Our methodology is completely model independent and it can be summarized in two steps: first, we applied the GA to the Constitution SNIa data in order to acquire a model independent reconstruction of the expansion history of the Universe $H(z)$. After we reconstructed $H(z)$, the second step was to use it in conjunction with the Om statistic and derive our null test.  Our main results was that $\\Lambda$CDM remains consistent with the data at the $2\\sigma$ level, see Fig.~\\ref{fig6}, but at the same time, some deviations from the standard $\\Lambda$CDM model can be accommodated and this is also in accordance with the data especially at low redshifts (see the behavior of the best fit in Fig.~\\ref{fig6}).  However, we should mention that the slowing down at low redshifts for the Constitution set mentioned earlier can also be seen with the standard analysis with the CPL ansatz (see Sanchez, Nesseris \\& Perivolaropoulos (2009)) but only for the Constitution set. For the other SNIa datasets the reverse trend (speeding up the acceleration at low z) was observed. For this reason and due to the fact that the current data still have quite large errors we believe that the full potential of our method, while it is very promising, will only be realized in the near future when more high quality SNIa data will become available. When this happens, the error region at low $z$ in Fig.~\\ref{fig6} will be small enough for the Om statistic to completely discriminate between $\\Lambda$CDM and the various Dark Energy models."
    },
    "1004/1004.0221_arXiv.txt": {
        "abstract": "% We study the three-dimensional distribution of matter at $z\\sim2$ using high resolution spectra of QSO pairs and simulated spectra drawn from cosmological hydro-dynamical simulations. We present a sample of 15 QSOs, corresponding to 21 baselines of angular separations evenly distributed between $\\sim 1$ and 14 arcmin, observed with the Ultraviolet and Visual Echelle Spectrograph (UVES) at the European Southern Observatory-Very Large Telescope (ESO-VLT). The observed correlation functions of the transmitted flux in the \\HI\\ \\Lya\\ forest transverse to and along the \\los\\ are in agreement, implying that the distortions in redshift space due to peculiar velocities are relatively small and - within the relatively large error bars - not significant. The clustering signal is significant up to velocity separations of $\\sim 300$ \\kms, corresponding to about 5 $h^{-1}$ comoving Mpc. Compatibility at the $2\\ \\sigma$ level has been found both for the Auto- and Cross-correlation functions and for the set of the Cross correlation coefficients. The analysis focuses in particular on two QSO groups of the sample, the Sextet and the Triplet. Searching for alignments in the redshift space between \\Lya\\ absorption lines belonging to different \\loss, it has been possible to discover the presence of a wide \\HI\\ structures extending over about ten Mpc in comoving space, and give constraints on the sizes of two cosmic under-dense regions in the intergalactic medium, which have been detected with a 91\\% and 86\\% significance level, respectively in the Sextet and in the Triplet. ",
        "introduction": "%  The understanding of the \\Lya\\ forest has dramatically improved in the recent decade, both on the theoretical and the observational side. Semi-analytical and hydro-dynamical simulations have outlined a new picture where the \\Lya\\ forest is due to the fluctuations of the intermediate and low-density intergalactic medium (IGM), arising naturally in the hierarchical process of structure formation. Relatively simple physical processes impact on the thermal state of the gas, which, on scales larger than the Jeans length, effectively traces the underlying distribution of dark matter. Support for this scenario is given by the satisfactory reproduction by semi-analytical and hydro-simulations of many properties of the \\Lya\\ forest (from the column density and the Doppler parameter distribution to the number density and effective opacity evolution) derived from the analysis of high resolution, high signal-to-noise ratio (SNR) QSO spectra obtained at 8-10m class telescopes \\citep[e.g.,][]{dave99,brianmachacek00,kim01,bianchi03,janknecht06}. A step forward in the study of the IGM with QSO absorption spectra is represented by the use of multiple \\loss\\ at close angular separations,  which allows information about the transverse direction to be obtained. Common absorption features observed in the spectra of multiply lensed quasars \\citep[e.g.,][]{smette92} and of close quasar pairs \\citep[e.g.,][]{vale98,aracil02} have provided evidence that the \\Lya\\ absorbers have dimensions of a few hundred kpc, in agreement with the predictions of simulations. Recently, lensed and more widely separated QSO pairs have been used to recover the kinematics of the gaseous cosmic web \\citep{rauch05}, confirming that the Hubble expansion and gravitational instability are the main processes influencing the \\Lya\\ forest gas.  A critical test of the nature of the \\Lya\\ absorbers, as proposed by simulations, comes from the determination of their spatial distribution properties. This has been done by analysing a great number of uncorrelated QSO lines of sight and computing the flux correlation function by averaging over many spectra \\citep[e.g.,][]{tripp98,savaglio99,croft02}. In this observational approach, however, the three-dimensional information is convolved with distortions in redshift space, due to peculiar motions and thermal broadening. Multiple \\loss\\ at small angular separations offer an invaluable alternative to address the spatial distribution of the absorbers, enabling a more direct interpretation of the observed correlations.  The final goal of this work is to investigate the distribution properties of matter in the IGM applying the modern interpretation of the \\Lya\\ forest to a sample of close QSO groups. The computation and comparison of the flux correlation function along and across the lines of sight was carried out in a previous paper \\citep[][Paper I]{vale06}. In this paper, we update the previous results using higher SNR spectra and, in particular, we investigate the coincidences of absorbers among three or more close lines of sight in order to detect cosmological structures extending to large scales. Finally, we look for extended under-dense regions in the IGM using multiple lines of sight.  Note that tomographic studies based on Lyman-alpha lines similar to those presented here will be of great importance in the near future due to the large numbers of quasar spectra that will be collected using low and medium resolution spectrographs (BOSS survey and the X-Shooter instrument for example) that aim at constraining spatial matter correlations in the transverse and longitudinal directions to ultimately measure baryonic acoustic oscillations and perform the Alcock-Paczyinski test at high redshift.  The paper structure is the following: in Section 2 we describe the observed data sample, the reduction procedure and the simulated spectra. Section 3 is devoted to the computation of the auto- and cross-correlation functions and the cross correlation coefficients of both the observed and simulated spectra and to their comparison. In Section 4, we analyze the coincidences in redshift space of the \\HI\\ and \\CIV\\ absorbers in order to detect structures extending over several comoving Mpc, while the under-dense regions common to multiple lines of sight are studied in Section 5. The conclusions are then drawn in Section 6. Throughout this paper we adopt $\\Omega_{0\\rm{m}} = 0.26,\\ \\Omega_{0\\Lambda}=0.74$, and  $h=H_0/(72$ km s$^{-1}$ Mpc$^{-1})$.   \\begin{table} \\centering \\caption{Characteristics of the observed QSO spectra.} \\label{obs:qso} \\begin{tabular}{@{}rcccc@{}} \\hline Object & $z$ & M$_{\\rm B}$ & \\Lya\\ & SNR  \\\\ & & & range & per pixel \\\\ \\hline Pair A\\ \\ \\ \\ \\ \\ \\ \\ \\  PA1 & 2.645 & 19.11 & 2.094--2.585 & 8--12 \\\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ PA2 & 2.610 & 19.84 & 2.094--2.550 & 3.5--6.5 \\\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\   &   &   &   &   \\\\ Triplet\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ T1 &  2.041 & 18.20 & 1.633--1.991 & 3--15 \\\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ T2 &  2.05  & 18.30 & 1.592--1.999 & 4--15 \\\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ T3 &  2.053 & 18.10 & 1.665-2.002 & 2.5--7 \\\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\   &   &   &   &   \\\\ Sextet\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ S1 &  1.907 & 19.66 & 1.665--1.859 & 2--7 \\\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ S2 &  2.387 & 19.53 & 1.858--2.331 & 3--8\\\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ S3 &  2.102 & 19.31 & 1.633--2.051 & 4--12 \\\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ S4 &  1.849 & 19.59 & 1.575--1.802 & 3--9 \\\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ S5 &  2.121 & 18.85 & 1.633--2.069 & 4--12\\\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ S6 &  2.068 & 20.19 & 1.592--2.017 & 3--10 \\\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\   &   &   &   &   \\\\ Pair U\\ \\ \\ \\ \\ UM680  & 2.144 & 18.60 & 1.653--2.092 & 6.5--17 \\\\ \\ \\ \\ \\ \\ UM681  & 2.122 & 19.10 & 1.634--2.070 & 7--17 \\\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\ \\   &   &   &   &   \\\\ Pair Q Q2343+12 & 2.549 &17.00 & 1.994--2.490 & 13--23 \\\\ Q2344+12 & 2.773 & 17.50 & 2.183--2.711 & 12--18 \\\\ \\hline \\\\ \\end{tabular} \\end{table} ",
        "conclusions": "In this paper, we exploited the capabilities of high resolution UVES spectra of QSO pairs to study the 3-dimensional distribution properties of baryonic matter in the IGM as traced by the transmitted flux in the QSO \\Lya\\ forests. Our sample is formed by 21 QSO pairs evenly distributed between angular separations of $\\sim 1$ and 14 arcmin, with \\Lya\\ forests at a median redshift $z \\simeq 1.8$. By calculating the correlation functions we compared the observed sample with a set of mock spectra drawn from a cosmological hydro-simulation run in a box of $120\\ h^{-1}$ comoving Mpc, adopting the cosmological parameters of the concordance model. The simulated sample reproduces 50 different realizations of the observed sample (see Section 2 for details). Furthermore, particular emphasis has been given to the search for alignments and other particular features of the \\Lya\\ forests in the two QSO groups present in our sample.  In the following, we summarize our main results:   \\noindent{\\em Two-point statistics} \\begin{enumerate} \\item{The computed correlation functions are in substantial agreement with those obtained in our previous paper \\citep[][Paper I]{vale06}. There is consistency between the clustering properties of matter in the IGM calculated in the direction parallel and transverse to the line of sight using the parameters of the concordance cosmology to map the angular distance into velocity separation. This is also due to the relatively large error bars of the computed quantities. As an implication, peculiar velocities in the absorbing gas are likely to be smaller than $\\sim 100$ \\kms.  Matter in the IGM is clustered on scales smaller than $\\sim 300$ \\kms\\ or about $4\\ h^{-1}$ comoving Mpc. The simulated correlation functions are consistent with the observed analogous quantities at the $1-2\\,\\sigma$ level for this particular sample. } \\item{Thanks to the increased SNR for some of the spectra in our sample the enhanced clustering signal measured in Paper I with the cross correlation coefficient at a transverse velocity separation $\\Delta v_{\\pe} \\sim 500$ \\kms\\ is no longer significant.} \\smallskip \\par\\noindent {\\em Three or more point statistics} \\item{Significant coincidences of \\Lya\\ absorptions have been detected among the \\loss\\ forming the Sextet implying the presence of coherent gas structures extending $\\sim14 h^{-1}$ comoving Mpc. In particular an excess of triplets and quadruplets of lines within $\\Delta v = 100$ \\kms\\ has been measured at a significance of 16 and 9 $\\sigma$. Besides, one group of five coincident lines in the S2-S3-S4-S5-S6 QSOs is observed, an occurrence that has a probability P=0.013 to arise from a random distribution of lines.} \\item{A method for the detection of under-dense regions in relatively low SNR spectra has been developed. One cosmic common under-dense region has been detected in each QSO group; in the Sextet the under-dense region has a dimension of 10.7 $h^{-1}$ comoving Mpc and a significance level of 91\\%, while in the Triplet it has a dimension of 8.8 $h^{-1}$ comoving Mpc and a significance level of 86\\%. These values are significantly smaller than those typical of under-dense regions detected along single lines of sight. The under-dense region common to the lines of sight of the Sextet can be  parametrized by a sphere of radius $6.75\\ h^{-1}$ comoving Mpc. } \\end{enumerate}  This study of the cosmic web environment at $z\\sim 2$ will soon be extended with many more QSO spectra either at medium resolution with the X-Shooter spectrograph \\citep{kaper09} or/and with the $R\\sim 10^5$ QSO spectra provided by SDSS-III \\citep{schlege07}. The ultimate goal is the characterization of the topology of the IGM in real space in a variety of environments by using \\Lya\\ and metal lines. On this basis, both  hydrodynamical simulations of structure formation and high-resolution spectroscopic samples, like the one presented here, will provide the link between observational quantities and the underlying density field and shed light on the impact of systematic effects."
    },
    "1004/1004.5318_arXiv.txt": {
        "abstract": "Where does solar flare energy come from? More specifically, assuming that the ultimate source of flare energy is mechanical energy in the convection zone, how is this translated into energy dissipated or stored in the corona? This question appears to have been given relatively little thought, as attention has been focussed predominantly on mechanisms for the rapid dissipation of coronal magnetic energy by way of MHD instabilities and plasma micro instabilities. We consider three types of flare theory: the steady state ``photospheric dynamo'' model in which flare power represents coronal dissipation of currents generated simultaneously by sub-photospheric flows; the ``magnetic energy storage'' model where sub-photospheric flows again induce coronal currents but which in this case are built up over a longer period before being released suddenly; and ``emerging flux'' models, in which new magnetic flux rising to the photosphere already contains free energy, and does not require subsequent stressing by photospheric motions. We conclude that photospheric dynamos can power only very minor flares; that coronal energy storage can in principle meet the requirements of a major flare, although perhaps not the very largest flares, but that difficulties in coupling efficiently to the energy source may limit this mechanism to moderate sized flares; and that emerging magnetic flux tubes, generated in the solar interior, can carry sufficient free energy to power even the largest flares ever observed. ",
        "introduction": "\\label{section:intro} It is generally accepted that solar flares occur only in regions of highly stressed magnetic field, and it is widely supposed that the flare energy is stored in situ in the coronal magnetic field, or at least that the magnetic field serves as a conduit for the energy flux supplying the flare. Although many mechanisms for releasing the energy stored in stressed coronal fields have been investigated, and models have been developed for the ways in which magnetic fields can be stressed by motions of their photospheric footpoints, the fundamental processes responsible for moving the footpoints have received relatively little attention. For instance the steady state ``photospheric dynamo'' models assume a given photospheric flow field, neglecting the reaction of the flow to generated $\\vecJ \\times \\vecB$ forces. Under this assumption, an infinite energy is available! Here we address the question of the origin of flare power, a question which appears to have been considered only crudely in the past: \\cite{Stenflo1969} concluded that kinetic energy in the convection zone beneath a sunspot was sufficient to supply a major flare, while \\cite{Spicer1982} commented that there was probably insufficient kinetic energy in the convection zone to directly power a flare.  In this paper, our goal is to compute upper limits to the available energy flux in as general a way as possible, avoiding many of the considerations which would have to be resolved in the construction of a realistic model. We examine the coupling of mechanical energy in sub-photospheric flows to the coronal magnetic field. The flows are assumed to differ at the two footpoints of a magnetic arch or arcade of loops, so that the differential motion of the footpoints exerts a twisting or shearing force on the coronal magnetic field. We evaluate the mechanical energy of turbulent motions in the upper convection zone and of large-scale differential rotation as possible sources of flare energy. We also estimate the free energy carried by a newly emerging flux tube, generated in the solar cycle dynamo region, which is independent of near-surface flows.  Observations suggest that magnetic flux at the photospheric level is concentrated into tight bundles or ``flux tubes.'' The amount of mechanical energy which can be intercepted depends on the geometry of the sub-surface magnetic flux tube, the height and time scales over which coherent motions extend, and the time scale for communication of forces to the solar surface. It is believed that isolated flux tubes extend at least part way down through the convection zone and that they are in hydrostatic equilibrium with magnetic pressure plus internal gas pressure equal to external pressure \\citep{Fisher1989}. We will adopt this flux tube model and assume that flux tubes are of roughly circular cross section. The assumptions of pressure confinement and a compact cross section impose the main limiting factor on the power supplied to flux tubes by convective motions.  We calculate the work done on the magnetic field in the following two limits. If currents induced by mechanical stresses are simultaneously dissipated resistively, the plasma slips across the magnetic field, as in an MHD generator, and the maximum power which can be extracted from the flow is a fraction of the kinetic energy flux intercepted by the flux tube. This is the ``photospheric dynamo'' model. We assume optimum impedance matching, so that the maximum energy flux can be extracted as flare power.  On the other hand, if resistive dissipation is negligible, the flow distorts the magnetic field and the work done by the gas is stored inductively. We presume that in this ``energy storage'' model the flow couples to the magnetic flux tube through aerodynamic drag. The available power is computed in the same way as for the ``dynamo'' model. To make use of this power, however, the coronal magnetic field must match the ``impedance'' (ratio of force to velocity) of the source. We find that optimum power transfer occurs when the opposing force due to bending of the coronal magnetic field lines allows the flux tube to move at 1/3 of the convective driving velocity, and that 4/27 of the kinetic energy flux can be transferred to the corona.  \\begin{figure} \\includegraphics[width=5.5in]{anm_yosemite_figure.eps} \\caption{ The kinetic energy flux of convective motions which can be intercepted by a magnetic flux tube above depth $z$ beneath the photosphere for an unrealistic flux tube of uniform magnetic field strength (1000 G) and for more realistic pressure-confined model flux tubes of photospheric field strengths 100, 500 and 2000 G. Open circles mark the maximum depths from which energy can propagate on a flare time scale of 1 hour, and filled circles mark the depths from which stresses can propagate on a coronal energy accumulation time scale of 1 day. (All the pressure-confined flux tubes can propagate stresses from the base of the convection zone in 1 day.) } \\label{figure:intro} \\end{figure} ",
        "conclusions": "\\label{section:conclusions}  We have considered three modes by which mechanical energy of the solar convection zone could be converted to transient energy release in solar flares. The mechanical energy due to differential rotation was found to be much too small to be significant, as was the energy in near-photospheric convective flows. However, the kinetic energy of motions in the upper convection zone can power moderate flares, while motions deep in the convection zone appear to be able to generate the energy required by large flares.  First, we examined the ``photospheric dynamo,'' which has been proposed by a number of authors in analogy with magnetospheric substorms and MHD generators. Here, flare energy is dissipated in the corona simultaneously with its generation by sub-photospheric flows across the magnetic field lines at the footpoints of coronal magnetic fields. We find that the kinetic energy flux in convective motions is far too small to power flares in this way, except perhaps for minor subflares.  Next we considered ``coronal energy storage,'' in which the kinetic energy of convective motions is intercepted in a similar way to the photospheric dynamo. In this case, however, the energy is stored in distortions of the coronal magnetic field, rather than being dissipated instantaneously. If energy is allowed to accumulate for periods of order 1 day, this mechanism has the potential to fuel even major flares. However, it appears to fall short by an order of magnitude of the energy thought to be released in the very largest observed flares. Moreover, it seems to be difficult to efficiently couple the energy source to the coronal magnetic field in order to extract the energy required for major flares on time scales of order 1 day, suggesting that convective motions may be restricted to powering only moderate sized flares.  Last, we computed the free energy available in an ``emerging flux tube'' and concluded that a large flux tube should carry enough free energy to account easily for even the largest flares. This theory is less satisfactory than the previous two in that the source of the free energy is the mechanism which generates magnetic flux in the first place, about which little is understood. But it is the only mechanism which appears able to fully account for the energy release in a large flare.  In conclusion, we note that solar flares exhibit an amazing variety of phenomena, and it is quite possible that many different forms of energization can result in a phenomenon which we recognize as a ``flare.'' But the extreme energy requirements of a major flare narrow the range of candidates considerably."
    },
    "1004/1004.0017_arXiv.txt": {
        "abstract": "We propose a method for setting upper limits to the extragalactic background light (EBL). Our method uses simultaneous {\\em Fermi}-LAT  and ground-based TeV observations of blazars and is based on the assumption that the intrinsic spectral energy distribution (SED) of  TeV blazars lies below the extrapolation of the {\\em Fermi}-LAT SED  from GeV to TeV energies. By extrapolating the {\\em Fermi}-LAT spectrum, which for TeV blazars is practically unattenuated by photon-photon pair production with EBL photons,  a firm upper limit on the intrinsic SED at TeV energies is provided. The ratio of the extrapolated spectrum to the observed TeV spectrum provides  upper limits to the optical depth for the propagation of the TeV photons due to pair production on the EBL, which in turn sets firm upper limits  to EBL models. We demonstrate our method using simultaneous observations from {\\em Fermi}-LAT and  ground-based TeV telescopes of the blazars \\object{PKS~2155-304} and \\object{1ES~1218+304}, and show that  high EBL density models are disfavored. We also discuss  how our method can be optimized and how {\\em Fermi } and X-ray monitoring observations of TeV blazars can guide future TeV campaigns, leading to potentially much stronger constraints on  EBL models. ",
        "introduction": "}   The EBL reflects the cosmologically important time-integrated history of light production and re-processing in the Universe. For this reason, measuring its intensity is highly desirable. The two components of the EBL are dust emission peaking at $\\lambda\\sim 100 \\mu m$ and starlight  peaking at $\\sim 1 \\mu m$. The actual level of the EBL is very difficult to measure, due to the dominance of foreground emission, mostly from interplanetary dust in our solar system \\citep[for reviews see][]{hauser01,kashlinsky05}, and the EBL level remains unknown within a factor of $\\sim$ few. Model-independent lower limits to the EBL, based on galaxy counts (e.g. Dole et al. 2006, B\\'ethermin et al. 2010), are the only strict lower limits on the EBL to date.  Modeling the EBL can lead to definite prediction, however uncertainties in the star formation rate, initial mass function, dust extinction, and how they evolve with redshift, has led to significant discrepancy among models \\citep{salamon98, stecker06, primack05, gilmore09, kneiske02, kneiske04, razzaque09, finke10_model, franceschini08}. Recently, \\cite{georganopoulos08} proposed a new method based on detecting as GeV emission EBL radiation that has been inverse Compton-scattered by relativistic electrons in the lobes of nearby radio galaxies such as Fornax~A.    The EBL in the 1-10 $\\mu$m range can in principle be obtained by using the TeV blazars as background light sources and modeling its attenuation due to pair production with EBL photons: by assuming that the intrinsic TeV spectrum is known from modeling of the broadband blazar SED, one can derive the mid-IR EBL by comparing the observed to the presumed intrinsic spectrum  \\citep[e.g.,][]{stecker92,stanev98,renault01}.  Because it is not possible to determine with confidence the intrinsic TeV spectrum, a variation of this method has been proposed that  sets limits on the EBL by assuming that the intrinsic blazar TeV spectrum cannot be arbitrarily hard:  from simple shock acceleration arguments one would not expect an intrinsic TeV photon index harder than $\\Gamma_{TeV}=1.5$ \\citep[e.g.,][]{aharonian06}. Detailed shock acceleration simulations, however,  indicate that harder VHE spectral indices may be possible \\citep{stecker07}. A large lower electron Lorentz factor \\citep{katarzynski06}, and Compton scattering of the cosmic microwave background in an extended jet may also lead to hard TeV spectra  \\citep{boett08}. These  considerations significantly relax the EBL limits derived  by  assuming $\\Gamma_{TeV}\\geq 1.5$  \\citep{mazin07,finke09}.    { Methods that constrain the EBL through purely spectral arguments are free of the uncertainties of adopting a particular physical model. Such methods, based solely on TeV data, have been proposed by \\cite{dwek05}, who considered unnatural TeV SEDs that exhibit an exponential increase at their high energy end    and by \\cite{schroedter05} who assumed  that all TeV blazars in flaring states have TeV spectra with the same   maximum intrinsic photon index $\\Gamma_{TeV}=1.8$. Here we present a spectral method for obtaining upper limits to the EBL energy density that makes use of simultaneous LAT and TeV observations. In \\S \\ref{method} we describe our method  and apply it to recent simultaneous LAT and TeV observations observations of PKS~2155-304,   in \\S \\ref{disc} we discuss how we can produce stronger constraints on the EBL and demonstrate this using  simultaneous GeV/TeV observations of 1ES 1218+304, and in \\S \\ref{conc} we conclude. ",
        "conclusions": "}      We presented a simple, model-independent method for setting upper limits to the EBL. Our method is based on the assumption that the level of the intrinsic TeV emission of blazars is below the extrapolation of the LAT SED to TeV energies. We applied our method to PKS 2155-304 and 1ES 1218+304, the only two TeV blazars of known redshift for which  published  simultaneous LAT-TEV observations currently exist. Even with these first  applications of our method, the highest level EBL models are disfavored. Future LAT-TeV simultaneous  observations hold the promise of pushing the EBL upper limits much further. We argued that it is important to devote TeV time not only to nearby sources, but also to more distant sources, hoping in both cases to observe high TeV states that for a given source will exhibit  smaller GeV to TeV spectral break $\\Delta \\alpha$, and will, therefore, produce stronger constraints.  Because the most difficult observations to obtain are in the TeV band, requests to monitor particular sources with TeV facilities can be triggered from {\\sl Fermi} observations of high states of TeV sources.  A complementary approach is X-ray monitoring of the TeV blazars.  In this case, because the X-ray emission of most TeV blazars is the high energy tail of the synchrotron component, high X-ray states are, in general,  a good proxy for  high TeV states (e.g. \\cite{aharonian09_2155}).  Such X-ray monitoring holds the promise of catching states in which, while the GeV emission does not increase substantially, the TeV emission does.  We note here that our method assumes that the entire spectral break from the LAT to the TeV bands is due to EBL  pair production absorption. This is an extreme assumption and it is highly probable that a substantial fraction of the break is intrinsic to the source. This in turn means that the actual level of the EBL may be significantly lower than the upper limits produced by our method. It would be very exciting and possibly hinting to new physics \\citep[e.g.,][]{amelino01} if the lowest collective values of $\\tau_{max}$ that our method will produce, challenge the lower level on the EBL inferred by galaxy counts \\citep[e.g.,][]{madau00, fazio04,bethermin10}. We anticipate that current and upcoming TeV-GeV blazar monitoring campaigns will provide plenty of opportunity for applying our method."
    },
    "1004/1004.5410_arXiv.txt": {
        "abstract": "We have discovered strong gravitational lensing features in the core of the nearby cluster Abell~3827 by analyzing Gemini South GMOS images. The most prominent strong lensing feature is a highly-magnified, ring-shaped configuration of four images around the central cD  galaxy. GMOS spectroscopic analysis puts this source at $z \\sim 0.2$. Located $\\sim$ 20\\arcsec~away from the central galaxy is a secondary tangential arc feature which has been identified as a background galaxy with $z \\sim 0.4$. We have modeled the gravitational potential of the cluster core, taking into account the mass from the cluster, the brightest cluster galaxy (BCG) and other galaxies. We derive a total mass of $(2.7\\,\\pm\\,0.4) \\times 10^{13}$ \\Msol\\ within 37 h$^{-1}$ kpc. This mass is an order of magnitude larger than that derived from X-ray observations. The total mass derived from lensing data suggests that the BCG in this cluster is perhaps the most massive galaxy in the nearby universe. ",
        "introduction": "\\label{sec:intro}  As the densest galaxy environments known, the cores of massive clusters are expected to host the strongest dynamical evolution. Such cluster cores are found to be dominated by early-type, D or cD galaxies that are also the brightest cluster galaxies (BCGs). These have extended luminous haloes \\citep{schombert1987}, which are not necessarily smooth \\citep{johnstone1991} and contain multiple or complex nuclei \\citep{rood1979}.  They are located close to the peak of the X-ray emission \\citep{jones1984} and near the kinematical centers of their clusters \\citep{quintana1982,quintana2000}.  There has been a long-running debate over the extent to which BCGs can be assembled through continuous merging of galaxies in the cluster potential (``galactic cannibalism'') versus early merging during cluster collapse \\citep[e.g.][]{west1994,dubinski1998}. However, recent work has shown that BCGs are likely  to form via ``dry'' or dissipationless major mergers of smaller early-type galaxies \\citep[e.g.][]{vandokkum2005,bell2006,delucia2007,whitaker2008}. These dry mergers are thought to be the primary mechanism through which massive galaxies grow from $z\\sim1$ to the present day and continue populating the upper end of the mass function without changing the overall mass density of elliptical systems. It has been pointed out that the merger rate increases both with the stellar mass of galaxies and with age \\citep{liu2009}. Therefore, there is a large fraction of major dry mergers in the nearby universe. Indeed, as many as 3.5\\% of BCGs show ongoing evidence of mergers \\citep{liu2009}, and about 10\\%~- 20\\%~of massive galaxies have undergone a dry merger in the last gigayear \\citep{kochfar2009}.  Thus, dry mergers ultimately lead to the formation of very massive and dense BCGs with large mass-to-light ratios.  We report here the discovery of a multi-component BCG that is, to our knowledge, the most massive galaxy ever seen in the local universe. This is located in the core of Abell 3827, a massive galaxy clusters in Abell's cluster catalog (richness class 2 and Bautz-Morgan type I), with an X-ray emission of $L_{X}$(0.1--2.4 keV) $=$ 2.1 $\\times$ 10$^{44}$ erg s$^{-1}$. This BCG is perhaps the most extreme example of ongoing galaxy cannibalism known: a super-giant elliptical that appears to be in the throes of devouring at least four other galaxies.  Evidence for a recent major merger is also supported by the appearance of an extended asymmetric halo at the center of the cluster. The super-giant galaxy also shows features arising from strong gravitational lensing, the most prominent being a surrounding, magnified ring-shape configuration of four similarly-shaped images. The existence of such features provides a unique opportunity to study the mass distribution and the evolution of BCGs with unprecedented spatial detail.  This Letter is organized as follows. In section 2, we summarize our observations and data reduction. In section 3, we describe our results, including the strong lensing modeling. In section 4 we present further discussion. Throughout this Letter we use a standard cosmology of $H_{0}=70$ $h$ km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{m}=0.3$ and $\\Omega_{\\Lambda}=0.7$. At the redshift of Abell 3827, 1\\arcsec\\ corresponds to 1.83~$\\!h^{-1}$ kpc. ",
        "conclusions": "\\label{sec:discussion}  Using the strong-lensing model presented above, we have analyzed the mass distribution to unprecedented spatial resolution. The cluster core is spatially very concentrated. Even if we assume a conservative fraction for the mass of the central cD galaxy (see above), this galaxy is very massive. The cD galaxy could be an extreme example of the effects of dry mergers on the mass of BCGs; dry mergers can produce an increase (by a factor of up to 3) in the dark-matter-to-stellar mass ratio for the most massive systems at present \\citep{ruszkowski2009}.  \\begin{figure}[!htb] \\centering \\includegraphics[width=0.95 \\columnwidth]{f3.eps} \\caption[]{Mass profile derived from the best fit model for the X-ray gas distribution (continuous line). The dash-dotted line shows the profile of the X-ray distribution for a more typical cluster. The short-dashed line shows the mass profile derived from strong lensing. The vertical solid line shows the location of the tangential arc (B.1) at 37 h$^{-1}$ kpc. \\label{mass}} \\end{figure}  We have compared the total mass derived from our strong-lensing analysis with that derived from the X-ray gas. We have calculated the total  mass as a function of radius for Abell 3827, based on a recent X-ray analysis by L. Valkonen et al. (2010, private communication) . We derive the total cluster mass by assuming hydrostatic equilibrium and using their best fit for the three-dimensional distribution of X-ray emitting gas. L. Valkonen et al. (2010, private communication) report $M_{200} = 9.7 \\times 10^{14}$ M$_{\\sun}$ inside a radius $r_{200}=2.84$ h$^{-1}$ Mpc, derived from the $M-T$ relation by \\citet{salhen2009} for a cluster temperature of $kT = (7.15 \\pm 0.18)$ keV.  Fig.~\\ref{mass} shows the radial distribution of the total mass derived from the X-ray gas and from the strong-lensing analysis.  For the X-ray gas, we plot mass profiles: (1) for a typical cluster with $r_{core} = 250$ h$^{-1}$ kpc and $\\beta = 0.7$ (dash-dot line), and (2) for the derived total mass of Abell 3827 within 37 h$^{-1}$ kpc (solid vertical line). Finally, we plot the mass profile for Abell 3827 derived from strong lensing as a dashed line. At all radii, the mass derived from strong lensing is at least a factor of 10 larger than that derived from X-ray data. Discrepancies between strong-lensing cluster masses and X-ray cluster masses for the same clusters have been reported previously in the literature \\citep[e.g.][]{gitti2007,halkola2008}. Moreover, significant biases of up to $\\sim 50$\\% may be introduced into strong-lensing mass estimates when models are extrapolated outside the Einstein ring \\citep{meneghetti2009}.  Analysis using both strong-lensing and X-ray methods at small radii show that differences can grow by as much as an order of magnitude \\citep{verdugo2007}.  How can we explain this discrepancy? For the lower limit on mass, within $3\\sigma$ error of the strong-lensing fit, we obtain $6.3 \\times 10^{12}$ M$_{\\sun}$ inside the ring structure of radius 10\\arcsec~($\\sim~19$~h$^{-1}$ kpc). Hence, the most likely uncertainties, related to extrapolation of the strong-lensing model to larger radii, are unable to explain our factor-of-ten discrepancy.  Assuming a spherical mass distribution, the mass inside system A is $(9.6 \\pm 0.9) \\times 10^{12}$ M$_{\\sun}$. The large difference between the strong-lensing mass and X-ray mass could be related to the dynamical state of the cluster, to the mass inferred from the X-ray gas, or to both effects. Total X-ray cluster masses derived from flat-core density profiles may be underestimated by at least a factor of 2 within the central $\\sim$ 30 h$^{-1}$ kpc \\citep{voigt2006}. Simulations \\citep{meneghetti2009} of strong-lensing models have shown that physical substructures along the line of sight to the cluster can overestimate the derived total mass by a factor of 2; in fact, our GMOS spectra reveal a bi-modality in the velocity distribution of cluster galaxies, suggesting that Abell 3827 is presently merging.  In all, these effects can produce a factor of about 4, still insufficient to explain the large discrepancy between strong-lensing and X-ray derived masses. Moreover, if we take out the factor of 4 due to the two effects described above, we are left with a cD galaxy with a mass of  $\\sim 2 \\times 10^{12}$ M$_{\\sun}$,  which appears to be one of the most massive galaxies known in the local universe. Detailed studies of the central objects at different wavelengths will reveal insights into the origins and subsequent evolution of this peculiarly massive galaxy."
    },
    "1004/1004.0832_arXiv.txt": {
        "abstract": "The  initial  mass  function  determines  the  fraction  of  stars  of different intial mass born per  stellar generation.  In this paper, we test  the effects  of the  integrated galactic  initial  mass function (IGIMF)  on the chemical  evolution of  the solar  neighbourhood.  The IGIMF (Weidner \\& Kroupa 2005) is computed from the combination of the stellar intial mass  function (IMF), i.e. the mass  function of single star clusters,  and the embedded  cluster mass function, i.e.  a power law with index $\\beta$.  By taking into account also the fact that the maximum achievable stellar mass is a function of the total mass of the cluster, the  IGIMF becomes  a time-varying  IMF which depends on  the star formation rate.   We applied this  formalism to a chemical evolution model for  the solar neighbourhood and compared the results obtained  by assuming three  possible values for  $\\beta$ with the  results obtained by  means of  a standard,  well-tested, constant IMF.  In general, a lower  absolute value of $\\beta$ implies a flatter IGIMF,  hence  a larger  number  of  massive  stars and  larger  metal ejection rates.  This translates into  higher type Ia and II supernova rates,  higher mass  ejection rates  from massive  stars and  a larger amount  of  gas  available  for  star formation,  coupled  with  lower present-day   stellar  mass  densities.    Lower  values   of  $\\beta$ correspond  also  to  higher  metallicities and  higher  [$\\alpha$/Fe] values at  a given  metallicity. We consider  a large set  of chemical evolution  observables and test  which value  of $\\beta$  provides the best  match  to  all  of  these  constraints.   We  also  discuss  the importance  of  the  present  day  stellar  mass  function  (PDMF)  in providing a way to  disentangle among various assumptions for $\\beta$. Our results indicate  that the model adopting the  IGIMF computed with $\\beta  \\simeq 2 $  should be  considered  the best  since it  allows us  to reproduce the  observed PDMF and to  account for most  of the chemical evolution constraints considered in this work. ",
        "introduction": "The initial stellar mass function is one of the major ingredients of chemical evolution models. Moreover, observed chemical abundances allow one to put robust constraints on both the normalization and the slope of the initial mass function (IMF; Chiappini et al. 2000; Romano et al. 2005). The environment providing most observational constraints for chemical evolution studies is the solar neighbourhood (S. N. hereinafter), for which a large set of observables are available. These observables include diagrams of abundance ratios versus metallicity, particularly useful when they involve two elements synthesised by stars on different timescales. An example is the [$\\alpha$/Fe] vs [Fe/H] diagram, since $\\alpha$ elements are produced by massive stars on short ($<0.03$ Gyr) timescales, while type Ia supernovae (SNe) produce mostly Fe on timescales spanning from 0.03 Gyr up to one Hubble time (Matteucci 2001). This diagnostic is a strong function of the IMF, but depends also on the assumed star formation history (Matteucci 2001; Calura et al. 2009). Another fundamental constraint is the metallicity distribution of living stars, which provides us with fundamental information on the IMF and on the infall history of the studied system. Another  diagnostic, depending both on the IMF and the past star formation history, is the present-day mass function, which represents the mass function of living stars observed now in the Solar Vicinity (Elmegreen \\& Scalo 2006). Other important observables useful for chemical evolution studies include the type Ia and type II SN rates, as well as the surface density of stars and gas, depending on the IMF and on the rate at which the gas has been processed into stars and remnants in the past, i.e. on the SFR.  In a previous paper, Recchi et al. (2009) considered a star-formation dependent IMF, called the integrated galactic initial mass function (IGIMF), which originates from the combination of the stellar IMF within each star cluster and the embedded cluster mass function. \\footnote{Embedded clusters are stellar clusters that are partially or fully encased in interstellar gas and dust within molecular clouds, therefore often visible only in the infrared.  It is supposed that all (or the large majority of) the stars form originally in embedded clusters (Lada \\& Lada 1991), but then they can loose their cocoon of gas because of the feedback of O stars (see e.g. Boily \\& Kroupa 2003a,b).  We will name therefore hereafter ``embedded clusters'' also the clusters which have lost their envelope, but still retain all of their stars, to be consistent with the terminology used in the original papers describing the IGIMF theory (e.g. Kroupa \\& Weidner 2003; Weidner \\& Kroupa 2005).}  Within each star cluster, the IMF can be well approximated by a two-part power-law form, $\\xi(m) \\propto m^{-\\alpha}$ (e.g. Pflamm-Altenburg, Weidner \\& Kroupa 2007). Massey \\& Hunter (1998) have shown that for stellar masses $m>$ a few M$_\\odot$, a slope similar to the Salpeter (1955) index (i.e. $\\alpha=2.35$) can approximate well the IMF in clusters and OB associations for a wide range of metallicities. Other studies have shown that the IMF flattens out below $m$ $\\sim$ 0.5 M$_\\odot$ (Kroupa, Tout \\& Gilmore 1993; Chabrier 2003). On the other hand, the embedded cluster mass function is well approximated by a single slope power law. This implies that small embedded clusters are more numerous in galaxies and they lock up most of of the stellar mass. However, the most massive stars tend to form preferentially in massive clusters (Weidner \\& Kroupa 2006). The integrated IMF in galaxies, the IGIMF, is a function of the galactic star formation rate (SFR) and, as a consequence of the embedded cluster mass function, it is steeper than the stellar IMF within each single star cluster (Kroupa \\& Weidner 2003; Weidner \\& Kroupa 2005).  Recchi et al. (2009) studied the effects of the IGIMF on the evolution of the SN rates in galaxies and on the chemical evolution of elliptical galaxies, showing how the IGIMF naturally accounts for the relation between the [$\\alpha$/Fe] and the stellar velocity dispersion observed in local elliptical galaxies. In this paper, we consider the effects of the IGIMF on the chemical evolution of the solar neighbourhood. As already stressed, the advantage of this approach is the availability of a large set of observational constraints, useful to test the IGIMF and, most importantly, to constrain its main parameter, i.e.  the index $\\beta$ of the power law expressing the embedded clusters mass function.  We will compare the results computed by means of a standard IMF, similar to the one by Scalo (1986), successful in reproducing most of the chemical evolution properties of the Solar Neighbourhood, with the results computed by means of the IGIMF. This Paper is organized as follows. In Section 2 we describe the IMF used in standard chemical evolution models and the formalism behind the IGIMF. In Section 3 we present a brief description of the chemical evolution model of the Solar Neighbourhood. In Sect. 4 we present our results and in Sect. 5 we draw our conclusions. ",
        "conclusions": "The initial mass function regulates the number of stars of different intial mass born per stellar generation, hence it plays a fundamental role in galactic chemical evolution studies. The aim of this paper was to test the effects of adopting the integrated galactic initial mass function (IGIMF) on the chemical evolution of the solar neighbourhood. The IGIMF is computed from the combination of the stellar intial mass function, i.e. the mass function in single star clusters, and the embedded cluster mass function, i.e. a power law with index $\\beta$, and taking into account that within each single cluster, the maximum stellar mass is a function of the total mass of the cluster. The result is a time-varying IMF which is a function of the galactic star formation rate. We applied the formalism developed by Weidner \\& Kroupa (2005) to a chemical evolution model for the solar neighbourhood, based on the two-infall model by Chiappini et al. (1997). For the embedded cluster mass function, we tested three different values of $\\beta$ and studied various physical quantities  and abundances for various chemical elements, comparing our results to the ones obtained by means of a standard IMF, constant in time and similar to the Scalo (1986) IMF. A statistical test to determine which is the best model in reproducing the set of observational data considered in this work has been developed. Our results can be summarized as follows:\\\\ 1) The value of $\\beta$ has important effects on the predicted star formation history. Also the effects of the SF threshold may be weaker or stronger, depending on the value chosen for  $\\beta$. In general, lower absolute values of $\\beta$ imply a flatter IGIMF, hence a larger number of massive stars and larger mass ejection rates. This translates into a larger amount of  gas available for star formation and gas density values never below the threshold and a smooth star fomation history. \\\\ 2) The value of $\\beta$ has an obvious deep impact on the predicted SN rates. Beside this, also other quantities can be influenced by this parameter. In general,  a lower  $\\beta$ implies higher mass ejection rates from massive stars, hence at any time a  larger gas mass density. \\\\ 3) The interstellar abundances are strongly influenced by the parameter $\\beta$. In general, lower absolute values of $\\beta$  imply higher metallicities and higher [O/Fe] values at a given metallicity. \\\\ 4) We have considered several chemical evolution constraints, including the observed local SN rates, local gas and stellar surface densities, the abundance ratios in local stars and the stellar metallicity distribution and, by varying the SF efficiency, we tested which assumption for the embedded cluster mass function exponent $\\beta$ provides the best simultaneous match to all of these observables. Our fitness test indicates that the best result has been obtained by assuming $\\beta =2 $, which produces an IMF resembling that of our standard model and that allows us to account for most of the  observables.\\\\ 5) Useful hints on the initial mass function come from the study of the present day mass function, i.e. the mass function of living stars observed today in the solar neighbourhood. The PDMF is a very good test of the chosen star formation history (Elmegreen \\& Scalo 2006) once an IMF has been assumed and a valuable test for the IMF once the SFR is assumed. In the case of the standard IMF, the star fomation history assumed here predicts a present-time SFR in good agreement with the observations, and is based on the results by Kennicutt (1989). Models with the IGIMF and different assumptions for $\\beta$ provide different results. Lower values for $\\beta$ produce a PDMF richer in massive stars, whereas higher values of  $\\beta$ imply steeper IGIMFs and a lower relative number of massive stars. The model with $\\beta =2 $ should be considered the best since it allows to reproduce at the same time the observed PDMF and to account for most of the chemical evolution constraints considered in this work.  In the future, we plan to extend our study of the effects of the IGIMF in galaxies of different morphological types. As forthcoming step, we will investigate how the IGIMF affects the chemical evolution of  dwarf galaxies."
    },
    "1004/1004.2006_arXiv.txt": {
        "abstract": "A procedure for unfolding the true distribution from experimental data is presented. Machine learning methods are applied for simultaneous identification of an apparatus function and solving of an inverse problem. A priori information about the true distribution from theory or previous experiments is used for Monte-Carlo simulation of the training sample. The training sample can be used to calculate a transformation from the true distribution to the measured one. This transformation provides a robust solution for an unfolding problem with minimal biases and statistical errors for the set of distributions used to create the training sample. The dimensionality of the solved problem can be arbitrary. A numerical example is presented to illustrate and validate the procedure. ",
        "introduction": "An experimentally measured distribution differs from the true physical distribution because of the limited efficiency of event registration and the finite resolution of a particular set-up. To identify a physical distribution, an unfolding procedure is typically applied \\cite{zhigunov2, blobel, correcting, gagunashvili,schmelling,hocker,zech,agost, blobel2, gagunashvili_phystat,albert}. Unfolding is an underspecified problem. Any approach to solving the problem requires a priori information about the solution. Methods for unfolding differ, directly or indirectly, in the use of this a priori information.  Unfolding when the apparatus function or transformation model for a true distribution from the measured one is unknown has been considered previously \\cite {gagunashvili, gagunashvili_phystat}. In this paper these ideas are further developed and the problem of simultaneously identifying a transformation model and inverse problem is solved. To obtain a robust solution for an unfolding problem, information about the shape of the distribution to be measured is used to create a training sample in Monte-Carlo simulations of an experiment. An approximation of the apparatus function is calculated for the set of distributions for the training sample. Use of this type approximation can minimize the statistical errors and biases of the unfolded distribution for  distributions used to create the training sample.  There is no restriction on the size and shape of bins, linearization of the problem is simple (if the set-up has non-linear distortions), and multidimensional data can be unfolded. A machine learning approach provides a method for validating the unfolding procedure and for improving the results.  The remainder of the paper is organized as follows. In Section 2 the main equation for solving an unfolding problem is proposed. A formal method for solving the unfolding problem and estimating the statistical errors for the unfolded distribution is discussed. Section 3 presents the algorithm for calculating the transformation matrix. In Section 4 the overall unfolding procedure is described. This consists of bin choice, system identification, solution of the basic equation and validation  of the unfolding procedure. Section 5 presents a numerical example. For comparison, an example reported elsewhere is used \\cite{blobel, schmelling, hocker}. To investigate biases  in the unfolding distribution, a numerical experiment with 1000 runs is performed. The results show that biases for the unfolded distribution is small. To demonstrate the robustness of the unfolding method for  distributions used to create the training sample, the same investigation is performed for eight distributions randomly chosen from training sample. The results reveal that there are small biases and low statistical errors for all the unfolding distributions, which confirms that the procedure is robust. Statistical errors are as small as possible in all cases because of application of the least mean square method and the method for system identification. ",
        "conclusions": "The main difficulties of the unfolding problem, which is a particular case of the inverse problem, are widely known. Information is lost in measuring owing to the inefficiency of registration in the frequency domain because of the low-pass filter defined by the resolution function and to the inefficiency of events registration defined by the acceptance function of the set-up. Thus, there are an infinite number of true distributions that give the same measured distribution and therefore a priori information about the solution must be used to solve an unfolding problem (inverse problem).  One way to solve an unfolding problem is to replace the original problem by a problem for a smoothed original true distribution and to use a sliding window (bin) for a smoothing. This is equivalent to solving the unfolding problem for the true distribution in some binning. Smoothing is low-pass filtering and the loss of information for a smoothed distribution due to the resolution function effect, which is another low-pass filter, is lower than for the original true distribution. Solution of the unfolding problem is easier, but no information is obtained about the structure of the original true distribution inside the bin.  In practical applications of the unfolding procedure, the transformation matrix $P$ must be calculated. Simulation of the measurement process is used for this, especially in nuclear and particle physics. This process is very time-consuming and the sample size for simulated events is often of the same order as for measured events. The calculated matrix will have many noisy matrix elements in this case, which is another source of instability in solving the inverse problem.  Main points related with difficulties of the unfolding problem have formulated above on physical level of rigor permit us summarize results of given paper and define place of proposed unfolding method among other known methods.  The method presented here is a completely new approach to unfolding problems using machine learning concepts, including a training sample, a validation procedure and boosting. All a priori information about the solution is contained in the training sample, which is a set of physically motivated true distributions known from theory and other experiments. Methods for selecting distributions for the training sample were presented in Section 3 and are supported by previous research \\cite{goodness,gagunashviliph}.  In the proposed method, an unfolded distribution can be calculated for a grid of points or for bins. There are no restrictions imposed by the dimensionality of the problem or the configuration of the bins or the grid. The method for identification provides a linear approximation of a transformation from the true distribution to the measured distribution if this transformation is non-linear.  The numerical example presented demonstrates the robustness of the new unfolding procedure and the possibility of unfolding a whole set of distributions with a single calculated matrix for transformation $P$. The set is defined as distributions used to create the training sample. Biases and statistical errors for components of the unfolded distribution were calculated using a Monte-Carlo method with 1000 runs. The examples demonstrate that the bias is small for components of the unfolded distribution and for estimates of the statistical errors. It should be noted that such biases were investigated for unfolding for the first time. The unfolding procedure is validated using a machine learning approach and has a good statistical interpretation. The proposed method has wide potential for applications in nuclear and particle physics, where  models for training samples can be proposed and Monte-Carlo simulations can be used to calculate transformation matrices."
    },
    "1004/1004.0003_arXiv.txt": {
        "abstract": "We present a detailed description of a phenomenological $\\H2$ formation model and local star formation prescription based on the density of molecular (rather than total) gas. Such approach allows us to avoid the arbitrary density and temperature thresholds typically used in star formation recipes.  We present results of the model based on realistic cosmological simulations of high-$z$ galaxy formation for a grid of numerical models with varied dust-to-gas ratios and interstellar far UV (FUV) fluxes. Our results show that both the atomic-to-molecular transition on small, $\\sim 10\\dim{pc}$ scales and the Kennicutt-Schmidt (KS) relation on $\\sim\\dim{kpc}$ scales are sensititive to the dust-to-gas ratio and the FUV flux. The atomic-to-molecular transition as a function of gas density or column density has a large scatter but is rather sharp and shifts to higher densities with decreasing dust-to-gas ratio and/or increasing FUV flux. Consequently, star formation is concentrated to higher gas surface density regions, resulting in steeper slope and lower amplitude of the KS relation at a given $\\Sgas$, in less dusty and/or higher FUV flux environments. These trends should have a particularly strong effect on the evolution of low-mass, low surface brightness galaxies which typically have low dust content and anemic star formation, but are also likely to be important for evolution of the Milky Way-sized systems.  We parameterize the dependencies observed in our simulations in convenient fitting formulae, which can be used to model the dependence of the KS relation on the dust-to-gas ratio and FUV flux in semi-analytic models and in cosmological simulations that do not include radiative transfer and $\\H2$ formation. ",
        "introduction": "\\label{sec:intro}  Conversion of gas into stars is one of the major sources of uncertainty in modeling formation of galaxies.  This uncertainty reflects our incomplete understanding of the process of star formation both locally and on global scales.  Traditionally, star formation is included in cosmological simulations and simulations of isolated galaxies by using simple phenomenological prescriptions that relate local rate of star formation to the local density of gas, with some additional criteria such as temperature and density thresholds for the gas to be eligible for star formation.  The parameters of these prescriptions are chosen so that the empirical power law relation between the {\\it surface density} of star formation, $\\Ssfr$, and surface density of (hydrogen) gas averaged on kpc scales, $\\Sgas$, $\\Ssfr\\propto \\Sgas^n$ with $n\\approx 1 - 1.4$, \\citep{schmidt59,sfr:k98a,sfr:blwb08} observed in $z\\approx 0$ galaxies is reproduced \\citep[see, e.g.,][for a recent overview]{schaye_dallavecchia08}.  However, both theoretical considerations and observational evidence indicate that such approach may miss some important environmental trends. For example, relation between the local star formation recipe and the large-scale Kennicutt-Schmidt (KS) relation is not trivial and depends on the density and thermal structure of the interstellar medium \\citep[ISM,][]{sims:k03,tassis07,wada_norman07,robertson_kravtsov08,schaye_dallavecchia08,saitoh_etal08}. This is because for a given large-scale gas surface density the fraction of dense, star forming gas is determined by the gas density distribution function, which, in turn, depends on the thermal state of the ISM \\citep{wada_norman01,robertson_kravtsov08}. For the same reason, the global rate of star formation may be controlled by the rate with which dense gas is formed by the ISM, rather then by the assumed local efficiency of the gas \\citep{saitoh_etal08}. This implies that star formation parameters tuned to reproduce the empirical KS relation in one situation \\citep[e.g., in controlled simulations of isolated disks][]{springel_hernquist03,schaye_dallavecchia08} may not reproduce this relation in galaxies with significantly different ISM density distributions.  In addition, there is a growing observational evidence that the KS relation is more complex than previously thought \\citep{heyer_etal04,boissier_etal03,sfr:blwb08}. For example, instead of a well-defined surface density threshold at low $\\Sgas$ below which $\\Ssfr$ drops to zero \\citep{martin_kennicutt01}, observations indicate continuous relation between star formation rate and gas surface densities \\citep{boissier_etal07} down to small $\\Sgas$, albeit with a steeper slope \\citep[e.g.,][]{sfr:blwb08}. Likewise, studies of individual dwarf galaxies, which typically have low gas surface densities ($\\Sgas\\lesssim 10-20\\Msun\\dim{pc}^{-2}$) throughout their disks, show that the KS relation in such galaxies is generally characterized by a considerably steeper slope, $n\\approx 2-4$, than the canonical value of 1.4 \\citep{heyer_etal04,sfr:blwb08,verley_etal10}. Moreover, recent detailed study of the global star formation relation by \\citet{sfr:blwb08} shows that a single power law is in general a poor description of the KS relation over the entire range of surface densities. Instead, the slope of the $\\Ssfr-\\Sgas$ relation may vary from the steep values of $n\\approx 2-4$ at $\\Sgas\\la 10 \\Msun\\dim{pc}^{-2}$ to linear $n\\approx 1$ at $\\Sgas\\sim 10-100 \\Msun\\dim{pc}^{-2}$ and then possibly steepening again to $n\\approx 1.5-2$ at $\\Sgas\\gtrsim 100\\Msun\\dim{pc}^{-2}$.  Finally, the growing evidence indicates that in high-redshift galaxies ($z\\gtrsim 3$) the KS relation is significantly steeper and has an order of magnitude lower amplitude at $\\Ssfr\\lesssim 100\\Msun\\dim{pc}^{-2}$ \\citep[][see also Fig. 3 in \\citeauthor{ng:gk10a} \\citeyear{ng:gk10a}]{sfr:wc06,rafelski_etal10}.  This complex behavior of the star formation rate density with the density of the neutral gas ($\\HI$+$\\H2$) can be understood if star formation occurs only in the molecular gas \\citep{robertson_kravtsov08,ng:gtk09,krumholz_etal09,pelupessy_popadopoulos09,ng:gk10a}. Indeed, detailed observations of nearby galaxies show that star formation correlates most strongly with the molecular gas \\citep[e.g.,][]{wong_blitz02,sfr:blwb08}, especially with the densest gas traced by HCN emission \\citep{gao_solomon04,wu_etal05}, while it only correlates weakly, if at all, with the density of atomic gas \\citep{wong_blitz02,kennicutt_etal07,sfr:blwb08}. We can thus expect that the relationship between the star formation rate density and gas density $\\Sgas=\\Smol+\\Shi$ (the KS relation) varies depending on the molecular fraction of the gas $f_\\H2=\\Smol/\\Sgas$.  Several factors may control the molecular fraction in the gas on different spatial scales. On small scales of individual molecular complexes it is primarily the cosmic dust abundance and the interstellar FUV radiation that control the atomic-to-molecular transition \\citep[e.g.,][see \\citeauthor{stahler_palla05} \\citeyear{stahler_palla05} for pedagogical review]{elmegreen93,sfr:kmt08}. On larger ($\\sim\\dim{kpc}$) scales the fraction of dense, molecular gas in a patch of gas of a given $\\Sgas$ is expected to depend on the density distribution of gas in that patch \\citep[e.g.,][]{elmegreen02}. The density distribution itself depends on thermodynamics of gas \\citep[see, e.g.,][]{robertson_kravtsov08} and metallicity, as more metal rich gas may be more efficient in building regions of higher densities via radiative shocks arising in the highly turbulent medium of gaseous disks. The density PDF should also reflect the global dynamics of gas in galactic disks in general. For example, spiral density wave will compress the gas facilitating its cooling and conversion of atomic gas into molecular form. Likewise, large-scale instabilities seed the turbulence in the disk that can shape the global density PDF \\citep{wada_norman01,elmegreen02,sims:k03,sfr:km05}.  Although observational studies of environmental dependence of the KS relation on gas metallicity, interstellar FUV radiation, and other properties of galaxies are in their early stages \\citep[e.g.,][]{sfr:blwb08,sfr:kept09,rafelski_etal10}, it is clear that such strong dependences can have important implications for our understanding of galaxy evolution \\cite[see discussion in][]{ng:gk10a}. For example, given that observations indicate that star formation in low-metallicity, high-UV flux environments of high-redshift galaxies is concentrated to significantly higher gas surface densities \\citep{sfr:wc06,rafelski_etal10}, stars in these galaxies should be confined to the high surface density regions and should therefore be more resistant against dynamical heating in mergers. At the same time, the longer gas consumption time scales in lower density regions of high-$z$ gaseous disks along with high accretion rate would keep them gas rich and more resilient to mergers as well \\citep[e.g..][]{robertson_etal04,robertson_etal06,springel_hernquist05}. This can help to resolve one of the major puzzles of hierarchical galaxy formation: prevalence of thin disks at low redshifts in the face of high merger rates at high redshifts.  It is thus important to explore potential effects and implications of the enviromental dependence of the KS relation for the evolution of galaxies. However, to capture the key physics responsible for this dependence in cosmological simulations of galaxy formation is challenging, because this requires high spatial resolution to model dynamics of interstellar medium in the hierarchically forming galaxies, 3D radiative transfer to model local UV radiation flux, and formation of molecular hydrogen. The latter is mediated by dust grains which catalyze H$_2$ formation and provide the initial key shielding from interstellar FUV radiation. This shielding allows build-up of molecular fraction sufficient for H$_2$ self-shielding, which in turn shapes the sharp transition of atomic to molecular gas.  Although fully self-consistent modeling of dust chemistry and H$_2$ formation is still far beyond reach, phenomenological model capturing the essential metallicity and UV flux dependence of molecular fraction can be used to model H$_2$ in self-consistent, high-resolution cosmological simulations \\citep[][]{ng:gtk09,ng:gk10a}. In this study we present a detailed description of such H$_2$ formation model and local star formation prescription based on the density of molecular (rather than total) gas. We present results for a grid of numerical models with varied dust-to-gas ratios and interstellar FUV radiation fluxes and explore the dependence of atomic-to-molecular transition on small, molecular cloud scales, on these variables and the effect this dependence has on the Kennicutt-Schmidt relation on large $\\sim\\dim{kpc}$ scales. We parameterize the dependencies observed in our simulations in convenient fitting formulae, which can be used to model the metallicity and UV flux dependence of the KS relation in semi-analytic models and in cosmological simulations that do not include radiative transfer and $\\H2$ formation. ",
        "conclusions": "\\label{sec:discussion}  We have presented results of a phenomenological model for formation of molecular hydrogen and have illustrated the dependence of molecular fraction on the gas density, dust-to-gas ratio, and far UV radiation flux. We have also presented the large-scale Kennicutt-Schmidt relation arising in our simulated galaxies when the local star formation is based on the density of molecular (rather than total) gas. Such approach allows us to avoid arbitrary density and temperature thresholds typically used in star formation recipes. Our results show that both the molecular fraction and the KS relation are sensititive to the dust-to-gas ratio and the FUV flux, although the sensitivity of the KS relation to the dust-to-gas ratio is stronger than to the FUV flux.  We parameterize the dependencies observed in our simulations by fitting formulae (\\S~\\ref{sec:fh2} and \\ref{sec:sfl}), which can be used to approximately account for $\\H2$ formation and $\\H2$-based star formation in simulations, which do not include a full $\\H2$ formation model and radiative transfer (see \\S~\\ref{sec:SIMrecipe}). We demonstrate that our fitting formulae, when applied to realistic simulations, produce results that are close to those obtained in simulations with the full $\\H2$ formation model and radiative transfer (Figure~\\ref{fig:sflpars2}).  We also provide fitting formulae for the dust-to-gas and the FUV radiation flux dependence of the KS relation that can be used in the semi-analytic models of galaxy formation (\\S~\\ref{sec:SAMrecipe}). One recent example of a model where such dependendcies can be relevant is the study of \\citet{dutton_etal10}.  The results of that study indicate that the redshift evolution of SFR-$M_{\\ast}$ relation of galaxies depends on the evolution of the relation between stellar and molecular masses. \\citet{dutton_etal10} find that, in their model, the effective surface density of atomic hydrogen is $\\Shi\\approx 10\\rm\\ M_{\\odot}\\,yr^{-1}$ and does not evolve with redshift. Our results, however, indicate that $\\Shi$ should increase with increasing redshift, as metallicities (and, hence, the dust abundance) of galaxies decrease and their FUV fluxes increase. Conversely, the $M_{\\ast}-M_{\\H2}$and SFR-$M_{\\ast}$ relations should evolve differently if their expected dependence on the dust-to-gas ratio and the FUV flux is taken into account. Given that at lower metallicities (and, hence, the dust abundance) we expect smaller star formation rate for the same amount and spatial distribution of neutral gas, the trends described in this paper may potentially explain why the model of \\citet{dutton_etal10} overpredicts the specific star formation rate ($\\dim{SSFR}\\equiv \\dim{SFR}/M_{\\ast}$) of small-mass galaxies at $z\\gtrsim 3$.  One of the most interesting results of our simulations is that significant amounts of ionized gas can be present around high redshift gaseous disks. This ionized gas is akin to the diffuse ionized gas observed in local galaxies \\citep[e.g.,][]{hoopes_etal03,haffner_etal09} and the Milky Way \\citep{reynolds89,reynolds91,gaensler_etal08}.  Our results indicate that the ionized gas may dominate the gas mass at low surface densities ($\\Sigma\\lesssim 10\\Msun\\dim{yr}^{-1}$). Furthermore, our simulations show that ionized gas can remain a significant mass component at higher gas surface densities in environments with low dust content and/or high FUV fluxes (e.g., compare gas surface densities for a given $\\Ssfr$ in Figures~\\ref{fig:sfltot} and \\ref{fig:sflntr}). One has to keep in mind the possible presence of significant amounts of ionized gas in theoretical interpretations of the KS relation and observational estimates of the total gas mass. The significantly different KS relation in the low dust-to-gas ratio, high FUV flux environments of high-redshift galaxies may also strongly bias gas mass estimates that use $z=0$ calibration of that relation \\citep[e.g.,][]{erb_etal06,manucci_etal09}.  As we discussed in \\citet{ng:gk10a}, the dust-to-gas ratio and the FUV flux dependence of the KS relation that we observe in our simulations has a number of important implications for galaxy evolution, such as a lower efficiency of star formation in DLA systems, star formation confined to the highest gas surface densities of high-$z$ disks, and generally longer gas consumption time scales in gaseous disks of high-redshift galaxies. The latter can be, at least partly, responsible for the prevalence of disk-dominated galaxies at low redshifts. This is because low efficiency of star formation can maintain disks gas rich until major mergers become rare. The outer, mostly \\emph{gaseous} regions of high-redshift disks should be more resistant against dynamical heating in mergers \\citep[e.g.,][]{robertson_etal04,robertson_etal06,springel_hernquist05} and would help maintain forming stellar disks dynamically cold during minor mergers \\citep{moster_etal09b} at later epochs. Moreover, minor mergers of forming disks should be largely gaseous, and gas brought in by such mergers should be deposited at large radii as it is ram pressure stripped by interaction with the gaseous disk and/or halo around it. This should prevent formation of large bulges, which was plaguing galaxy formation models, and instead lead to formation of more extended, higher-angular momentum disks. This scenario is borne out in recent galaxy formation simulations of \\citet*{agertz_etal10}, who show that low efficiency of star formation at high redshifts leads to more realistic disks and smaller bulge-to-disk ratios.  Another interesting consequence of the complex dependence of the KS relation on the dust-to-gas ratio and the FUV flux may be relevant to our own backyard. Recently, \\citep{dsh:ogws08} noted that star formation histories of Milky Way satellites can only be explained by a KS relation (Equation (\\ref{eq:sfitk})) with the sharp threshold if the threshold varies semi-randomly within a modest dispersion of about 0.1 dex. This variation is consistent with the variation given by Equation (\\ref{eq:sstar}) for the values of $\\D$ and $\\U$ typical for dwarf galaxies ($\\D\\gtrsim0.1$, $\\U\\gtrsim1$). Since star formation histories of galactic satellites are known to be highly variable \\citep{dsh:m98,dsh:dwsh05}, the FUV flux is expected to vary accordingly; such variations may be responsible for the needed variation of the threshold in the KS relation, or, more precisely, the characteristic surface density $\\Sigma_{\\ast}$ from Equation (\\ref{eq:sstar}).  The high mass-to-light ratios (and hence low star formation efficiencies) of the Local Group dwarf spheroidal galaxies may also be partially explained by the environmental dependence of $\\H2$ abundance and, hence, star formation. Star formation in such low metallicity, low dust content dwarf galaxies should be confined only to the highest gas surface densities (i.e., the central regions) while leaving the bulk of the gas at lower gas surface densities inert to star formation. This is consistent with observations of local dwarf low surface brightness galaxies which exhibit very low molecular gas fractions and anemic star formation rates \\citep{matthews_etal05,das_etal06,boissier_etal08,wyder_etal09,roychowdhury_etal09}.  The examples described above illustrate the importance of further investigation of the effects of environmental dependencies of the KS relation discussed in this paper. The results and fitting formulae that we present should aid in implementing such dependencies in both cosmological simulations and semi-analytic models and should thus help to explore a wide range of possible effects."
    },
    "1004/1004.0529_arXiv.txt": {
        "abstract": "We present UV broadband photometry and optical emission-line measurements for a sample of 32 Brightest Cluster Galaxies (BCGs) in clusters of the Representative XMM-Newton Cluster Structure Survey (REXCESS) with $z=0.06-0.18$. The REXCESS clusters, chosen to study scaling relations in clusters of galaxies, have X-ray measurements of high quality. The trends of star formation and BCG colors with BCG and host properties can be investigated with this sample. The UV photometry comes from the XMM Optical Monitor, supplemented by existing archival GALEX photometry. We detected  H$\\alpha$ and forbidden line emission in 7 (22\\%) of these BCGs, in optical spectra obtained using the SOAR Goodman Spectrograph. All of these emission-line BCGs occupy clusters classified as cool cores based on the central cooling time in the cluster core, for an emission-line incidence rate of 70\\% for BCGs in REXCESS cool core clusters. Significant correlations between the H$\\alpha$ equivalent widths, excess UV production in the BCG, and the presence of dense, X-ray bright intracluster gas with a short cooling time are seen, including the fact that all of the H$\\alpha$ emitters inhabit systems with short central cooling times and high central ICM densities. Estimates of the star formation rates based on H$\\alpha$ and UV excesses are consistent with each other in these 7 systems, ranging from $0.1-8$ solar masses per year. The incidence of emission-line BCGs in the REXCESS sample is intermediate, somewhat lower than in other X-ray selected samples ($\\sim35\\%$), and somewhat higher than but statistically consistent with optically selected, slightly lower redshift BCG samples ($\\sim10-15\\%$). The UV-optical colors (UVW1-R $\\sim 4.7 \\pm 0.3$)  of REXCESS BCGs without strong optical emission lines are consistent with those predicted  from templates and observations of ellipticals dominated by old stellar populations. We see no trend in UV-optical colors with optical luminosity, $R-K$ color, X-ray temperature, redshift, or offset between X-ray centroid and X-ray peak ($\\langle w \\rangle$). The lack of such trends in these massive galaxies, particularly the ones lacking emission lines, suggests that the proportion of UV-emitting (200-300 nm) stars is insensitive to galaxy mass, cluster mass, cluster relaxation, and recent evolution, over the range of this sample. ",
        "introduction": "In current models of hierarchical galaxy formation, Brightest Cluster Galaxies (BCGs) are the trash heaps of the universe. As galaxies fall through and past the center of a cluster of galaxies, tidal forces strip them of their stars, and ram pressure stripping by the dense intracluster medium (ICM) removes some of their gas. The BCG settles to the cluster's center of mass and accretes stars and gas, developing a huge extended stellar halo characteristic of central dominant (cD) galaxies. These galaxies are therefore not simply overgrown massive ellipticals, but represent a category of galaxies that appear to have very special growth histories.  Modern simulations have great difficulty reproducing the observed properties of these singular objects, which form deep in the potential well of a cluster. Apparently, the star formation in these systems must be quenched in order to explain their optical luminosities and colors \\citep[e.g., ][]{1998MNRAS.294..705K}. In the most massive galaxies, the primary agent of feedback at late times is thought to be the central supermassive black hole.  The vast majority of the stars in the most massive elliptical galaxies were already present billions of years ago, based on studies of their color-magnitude relation and spectral energy distributions \\citep[e.g., ][]{1992MNRAS.254..601B,1998A&A...334...99K, 2008MNRAS.386.1045A}, and the mass of a galaxy's central black hole tends to be about 0.2\\% the mass of its spheroidal component \\citep[e.g., ][]{1998AJ....115.2285M,2000ApJ...539L..13G,2000ApJ...539L...9F}. Somehow, formation and growth of the black hole is coupled to the formation and growth of the galaxy.  Meanwhile, the downsizing phenomenon \\citep{1999AJ....118..603C}, wherein massive galaxies stop forming stars earlier than low-mass galaxies, suggests that star formation somehow decouples from the growth of the galaxy by hierarchical accretion \\citep[e.g., ][]{2002MNRAS.333..156B}.  Galaxy formation models without feedback from an active galactic nucleus (AGN) predict high-mass galaxies that are far too blue and luminous \\citep{1998MNRAS.294..705K}, but the situation may be resolved when (sufficient) feedback from an AGN is included \\citep[e.g., ][]{1997ApJ...487L.105C, 1998A&A...331L...1S, 2004MNRAS.347.1093B}.  The star formation in today's BCGs, if present at all, is a mere shadow of what it must have been billions of years ago when most of the stars were formed.  Nevertheless, some BCGs apparently do persist in forming stars, with the trend that the clusters with the shortest cooling times appear to be much more likely to host emission line systems and blue cores \\citep[e.g., ][]{1983ApJ...272...29C, 1989ApJ...338...48H, 1992ApJ...393..579M,2006ApJ...652..216R,2008ApJ...683L.107C, 2008ApJ...681L...5V}. Such clusters comprise about half of X-ray luminous ($L_X \\gtrsim 10^{44}$ erg s$^{-1}$) clusters at low redshift \\citep{1992MNRAS.258..177E,1992ApJ...385...49D,1999MNRAS.306..857C}. \\citet[][]{2005ApJ...635L...9H}, in an XMM Optical Monitor study of 33 galaxies, of which 9 were BCGs, showed a connection between excess ultraviolet (UV) emission and recent X-ray estimates of mass cooling rates. There appears to be a clear empirical connection between the state of the hot gas and the activity in the BCG.  Observationally, multi-wavelength studies of clusters of galaxies and their BCGs provide some insight into how this feedback process may proceed. Cavities in the X-ray intracluster gas that surrounds radio lobes provided smoking-gun evidence for AGN feedback \\citep[e.g., ][]{1993MNRAS.264L..25B,2005Natur.433...45M}. Estimates of the $PdV$ work required to inflate such cavities have shown that the kinetic energy outputs from these AGNs are surprisingly high compared to their radio luminosities and estimated lifetimes \\citep[e.g., ][]{2007ARA&A..45..117M}. Central radio sources are quite common at the centers of galaxy clusters with cool, dense ICM cores \\citep{1990AJ.....99...14B}. Sporadic AGN outbursts with kinetic outputs of $\\sim 10^{45} \\, {\\rm erg \\, s^{-1}}$ can plausibly stabilize cooling and star formation in the BCGs of those clusters, explaining why those galaxies are less red and less luminous than BCGs modeled without AGN feedback \\citep[e.g., ][]{2005ApJ...634..955V}. While this is not necessarily the feedback mode that quenches star formation in elliptical galaxies at high redshift, this is the only form of AGN feedback that we can currently study in such detail.  Multi-wavelength studies of a well-chosen X-ray sample without any particular morphological criteria, with uniform and accurate X-ray measurements, are the best way to study the relationship of star formation, intracluster gas, and AGN feedback at low redshift. Such studies shed light on what regulates star formation in large galaxies at earlier times in the universe. Ultraviolet observations are sensitive to the continua of OB stars, and H$\\alpha$ is produced by photoionization of the interstellar medium (ISM) by O-stars.  Elliptical galaxies, and therefore older stellar populations, exhibit an increasing UV emission towards shorter UV wavelengths (``UV upturn'') from extreme horizontal branch stars.  However, recent discussions suggest that relatively tiny amounts of young stars can contribute to the observed scatter in this component for UV starlight \\citep[e.g., ][]{2007MNRAS.380.1098H}.  In this work we investigate the UV properties of BCGs including both UV from recent star formation and from older stellar populations. We present new UV observations from the {\\em XMM-Newton} Optical Monitor and archival results from the Galaxy Evolution Explorer (GALEX), together with  a census of BCG emission-line activity from ground-based observations from the SOAR Goodman Spectrograph. These observations are of BCGs in the Representative {\\em XMM--Newton} Cluster Structure Survey (REXCESS) \\citep{2007A&A...469..363B}. We provide a small number of supplemental observations to complete the spectroscopic coverage of a comparison sample of BCGs. We assume cosmological parameters $H_0=70$ km s$^{-1}$ Mpc$^{-1}$ ($h=0.7$) and a flat geometry with $\\Omega_M=0.3$ throughout. ",
        "conclusions": "We report the results of UV broad-band photometry from the XMM Optical Monitor and from the GALEX mission and of long-slit optical emission-line spectroscopy for 32 BCGs in REXCESS. Seven of these clusters exhibit classic signatures of star formation activity associated with a cool-core cluster as a host. Indeed these seven BCGs inhabit the 10 clusters in the REXCESS sample with the shortest cooling times in the hot gas, as inferred from XMM observations, and the two most luminous of these in H$\\alpha$ are the most prominent cool core clusters in REXCESS. The incidence rate of emission-line BCGs is intermediate between that found in the B55 or EMSS cluster samples and that found in optically-selected SDSS clusters, possibly a consequence of the lack of morphological bias in the selection of REXCESS clusters. The BCGs with the largest H$\\alpha$ equivalent widths are also the BCGs with the bluest UV-optical colors. We report a correlation between the BCG H$\\alpha$ equivalent width and -- to a lesser degree -- the UV-optical color with conditions in the intracluster gas in the cluster core: the scaled core electron density and the gas cooling time. The incidence rates and the correlations suggest a physical connection between activity (emission-line excitation, recent star formation) in the BCG and the cooling time of the intracluster gas. We see no correlation between H$\\alpha$ equivalent width and X-ray temperature or cluster mass; we also see no correlation between the UV-optical color and cluster mass, temperature, or degree of relaxation, suggesting that whatever ignites BCG activity at low redshift is relatively insensitive to halo properties. This insensitivity is puzzling, given the difference in the incidence of emission-line BCGs in optically selected vs. X-ray selected samples of clusters of galaxies. A decreasing incidence (or strength) of emission-line BCGs with decreasing halo mass or BCG mass could have explained the discrepancy. However, we do not have evidence for a correlation of that nature, at least over the range probed by the REXCESS cluster sample. However, the strong correlation between the presence of low entropy gas and the appearance of star formation signatures suggests there is a connection between the nature of the cluster and conditions in its BCG. These star-forming BCGs also tend to inhabit clusters with low $\\langle w \\rangle$, characteristic of clusters with relatively relaxed dynamical states. Optically selected cluster samples may include a wider range of morphologies and dynamical states than that sampled by X-ray selected cluster samples.   Almost all ($29/31$) of the BCGs were detected in the OM UVW1 band, regardless of their emission-line activity. The UVW1-optical colors of these galaxies are consistent with the UV light of an old population, with excess UV coming from recent star formation. The UVW1 photometry is not strongly affected by the UV upturn stars, and therefore turns out to be a decent choice to study UV associated with star formation. We analyze archival GALEX and XMM OM UVM2 data at shorter wavelengths.  The inactive BCGs in our sample, classified as such by the lack of H$\\alpha$ emission, define a relatively homogeneous sample.  The observed UV-optical colors of these BCGs (UVW1-R$ \\sim 4.8$) are independent of redshift, BCG absolute rest frame R magnitude, and rest frame R-K color, suggesting that the UVW1-optical colors of inactive BCGs are also unrelated to BCG mass and metallicity. Their colors are independent of cluster temperature (a measure of the depth of the cluster's gravitational potential) and degree of relaxation, as measured by the offset of the X-ray centroid from the X-ray emission peak. This lack of correlation means that the UVW1-R colors are not affected by the gravitating mass of the cluster or its dynamical state."
    },
    "1004/1004.0974_arXiv.txt": {
        "abstract": "A detection or nondetection of primordial non-Gaussianity in the CMB data is essential not only to test alternative models of the physics of the early universe but also to discriminate among classes of inflationary models. Given this far reaching consequences of such a non-Gaussianity detection for our understanding of the physics of the early universe, it is important to employ alternative indicators in order to have further information about the Gaussianity features of CMB that may be helpful for identifying their origins. In this way, a considerable effort has recently gone into the design of non-Gaussianity indicators, and in their application in the search for deviation from Gaussianity in the CMB data. Recently we have proposed two new large-angle non-Gaussianity indicators which provide measures of the departure from Gaussianity on large angular scales. We have used these indicators to carry out analyses of Gaussianity of the single frequency bands and of the available foreground-reduced \\emph{five-year} maps with and without the \\emph{KQ75} mask. Here we extend and complement these studies by performing a new analysis of deviation from Gaussianity of the \\emph{three-year} harmonic ILC (HILC) foreground-reduced  full-sky and \\emph{KQ75} masked maps obtained from WMAP data. We show that this full-sky foreground-reduced maps presents a significant deviation from Gaussianity, which is brought down to a level of consistency with Gaussianity when the \\emph{KQ75} mask is employed. ",
        "introduction": "A detection or nondetection of primordial non-Gaussianity in the CMB data is crucial not only to discriminate inflationary models but also to test some alternative scenarios for the physics of the early universe. However, the extraction of primordial non-Gaussianity is a difficult enterprise since several effects of non-primordial nature can produce non-Gaussianity in the CMB data. Clearly the study of detectable non-Gaussianities in the WMAP data must take into account that they may have non-cosmological origins as, for example, unsubtracted foreground contamination, unconsidered point sources emission and systematic errors.\\cite{Chiang-et-al2003,Naselsky-et-al2005,Cabella-et-al2009} Deviation from Gaussianity may also have a cosmic topology origin (see, e.g., the review articles Refs.~\\refcite{CosmTopReviews} and  related Refs.~\\refcite{TopDetec}). If, on the one hand, different statistical tools can in principle provide information about distinct forms of non-Gaussianity, on the other hand one does not expect that a single statistical estimator can be sensitive to all possible forms of non-Gaussianity in CMB data. It is therefore important to test CMB data for Gaussianity by using different statistical indicators to shed some light on its possible causes. In view of this, a great deal of effort has recently gone into verifying the existence of non-Gaussianity by employing several statistical estimators.\\cite{Some_non-Gauss-refs}  Recently have we proposed\\cite{Bernui-Reboucas2009a} two new large-angle non-Gaussianity indicators, based on skewness and kurtosis of large-angle patches of CMB maps, which provide measures of the departure from Gaussianity on large angular scales. We used these indicators to search for the large-angle deviation from Gaussianity in the three and five-year single frequency K, Ka, Q, V, and W maps with and without a \\emph{KQ75} mask.  We have found strong deviation from Gaussianity in the unmasked maps, whereas a \\emph{KQ75} mask lowers significantly the level of non-Gaussianity (for details see Ref.~\\refcite{Bernui-Reboucas2009a}).  Motivated by the fact that most of Gaussianity analyses with Wilkinson Microwave Anisotropy Probe (WMAP) data have been carried out by using CMB frequency bands masked maps, and that sky cut can in principle induce bias in Gaussianity analyses, in a more recent paper\\cite{Bernui-Reboucas2009b} we have carried out an analysis of Gaussianity of the available full-sky foreground-reduced \\emph{five-year} CMB maps~\\cite{ILC-5yr-Hishaw,HILC-Kim,NILC-Delabrouille} by using the statistical indicators of Ref.~\\refcite{Bernui-Reboucas2009a}.  We have shown that the available full-sky five-year foreground-reduced maps present a significant deviation from Gaussianity, which varies with the foreground-cleaning procedures. We have also shown that there is a substantial reduction in the level of deviation from Gaussianity in these full sky maps when a \\emph{KQ75} mask is used. Our main aim here is to extend and complement our previous work\\cite{Bernui-Reboucas2009b} by performing a similar analysis of Gaussianity of the \\emph{three-year} harmonic ILC (HILC) maps\\cite{HILC-Kim} foreground-reduced  full-sky and \\emph{KQ75} masked maps. To this end, in the next section we give an account of the large-angle non-Gaussianity indicators of Ref.~\\refcite{Bernui-Reboucas2009a}, while in the last section we apply our indicators to perform a Gaussianity analysis of the HILC three-year full-sky and \\emph{KQ75} cut-sky maps, and present our main results. Our principal conclusion is that the HILC full-sky foreground-reduced maps presents a significant deviation from Gaussianity, which is reduced to a level of consistency with Gaussianity when the \\emph{KQ75} mask is employed. ",
        "conclusions": "In this section we shall report the results of our Gaussianity analysis performed with   $S = S(\\theta,\\phi)$ and $K = K(\\theta,\\phi)$ indicators calculated from the foreground reduced HILC  full-sky and $KQ75$ masked maps computed from three-year WMAP data.% \\footnote{We note that in the analysis of Gaussianity with the \\emph{KQ75} masked maps the implementation of the mask is made by removing the pixels inside the masked regions from the set of pixels of the each cap whose intersection with the mask is not empty. Thus, the values $S_j$ and $K_j$ for a  $j^{\\,\\rm{th}}$ cap (with pixels in the mask region) are calculated with small number $N_{\\rm p}$ of pixels.}   \\begin{figure*}[htb!] \\begin{center} \\includegraphics[width=6cm,height=3.8cm]{reboucasFig1a.ps}  % \\hspace{0.3cm} \\includegraphics[width=6cm,height=3.8cm]{reboucasFig1b.ps}   % \\caption{\\label{Fig1} Low $\\ell $ \\emph{differential} power spectra of skewness $|S_{\\ell} - \\overline{S}_{\\ell}|$ (left) and kurtosis (right) $|K_{\\ell} - \\overline{K}_{\\ell}|$ calculated from the  foreground-reduced HILC full-sky and \\emph{KQ75} masked maps. The $95\\%$ confidence level (obtained from Monte-Carlo Gaussian maps) is indicated by the dashed line. \\vspace{-0.6cm} } \\end{center} \\end{figure*}  To minimize the statistical noise, in the calculations of $S-$map and $K-$map from the HILC  foreground-reduced three-year map, we have scanned the celestial sphere with $12\\,288$  spherical caps of aperture  $\\gamma = 90^{\\circ}$, centered at points homogeneously generated on the two-sphere by using  HEALPix\\cite{Gorski-et-al-2005}.  Figure~\\ref{Fig1} shows the differential power spectrum of the skewness $S_{\\ell}$ (left panel) and kurtosis $K_{\\ell}$ (right panel) indicators for $\\,\\ell=1,\\,\\,\\cdots,10\\,$, calculated from \\emph{full-sky} and \\emph{KQ75} \\emph{cut-sky} three-year foreground-reduced HILC maps. The $95\\%$ confidence level, obtained from $S$ and $K$ maps calculated from Monte-Carlo (MC) statistically Gaussian CMB maps, is indicated in this figure.% \\footnote{For details on the calculation of these (data and MC) maps and the associated power spectra we refer the readers to Ref.~\\refcite{Bernui-Reboucas2009a} and Ref.~\\refcite{Bernui-Reboucas2009b}.} To the extent that the deviations $|S_{\\ell} - \\overline{S}_{\\ell}|$ and $|K_{\\ell} - \\overline{K}_{\\ell}|$ for these maps are not within $95\\%$ of the mean MC value, Fig.~\\ref{Fig1} shows an important deviation from Gaussianity in full-sky foreground-reduced HILC three-year map. This figure also shows a significant reduction in the level of large-angle deviation from Gaussianity when the \\emph{KQ75} mask is used.  To have an overall assessment power spectra $S_\\ell$ and $K_\\ell$ calculated from the HILC foreground-reduced three-year full and cut map, we have performed a $\\chi^2$ test to find out the goodness of fit for $S_{\\ell}$ and $K_{\\ell}$ multipole values as compared to the expected multipole values obtained from $S$ and $K$ maps calculated from Monte-Carlo (MC) statistically Gaussian CMB maps. This gives a number for each case that quantifies collectively the deviation from Gaussianity. For the power spectra $S_\\ell$ and  $K_\\ell$ we found the values given in Table~\\ref{table1} for the ratio $\\chi^2/\\text{dof}\\,$ (dof stands for degrees of freedom) for the power spectra calculated from three-year HILC foreground-reduced full-sky and cut-sky maps.  \\begin{table}[th] \\tbl{$\\chi^2$ test goodness of fit for $S_{\\ell}$ and $K_{\\ell}$ calculated from the HILC  full-sky and cut-sky three-year maps as compared with the expected values $\\overline{S}_{\\ell}$ and  $\\overline{K}_{\\ell}$ obtained from MC maps  .} {\\begin{tabular}{@{}lcc@{}} \\toprule % Map  & $\\chi^2$ for $S_{\\ell}$  & $\\chi^2$ for $K_{\\ell}$  \\\\ \\colrule HILC full-sky           & $1.6 \\times 10^3$  & $   8.8 \\times 10^5 $   \\\\ HILC $KQ75$ cut-sky     & 0.8                & 1.5   \\\\ \\botrule % \\end{tabular} \\label{table1}} \\end{table}  Clearly, the greater is the values for $\\chi^2/\\text{dof}\\,$  the smaller are the $\\chi^2$ probabilities, that is the probability that the power spectra $S_{\\ell}$ and $K_{\\ell}$ and the expected  MC power spectra agree. Thus, from  Table~\\ref{table1} is one concludes that the HILC presents the substantial level of deviation from Gaussianity as detected by the indicators, which is reduced to a level that can be considered consistent with Gaussianity when the \\emph{KQ75} mask is employed.  Finally we note that the relative deviation of the full-sky power spectrum from the cut-sky spectrum can be calculated with no reference to the Gaussian MC spectra. To this end, we have performed a $\\chi^2$ test to find out the goodness of fit for $S_{\\ell}$ and $K_{\\ell}$ multipole values for the full-sky maps as compared to the corresponding cut-sky values. For this relative assessment of power spectra $S_\\ell$ and $K_\\ell$ we have found that $\\chi^2/\\text{dof}\\,$ are $1.4 \\times 10^3$ and $7.2 \\times 10^5$. These values make apparent the significant effect of the mask in the reduction of the deviation from Gaussianity in the full-sky HILC three-year map, and give information % on reliability of the HILC full-sky foreground-reduced three-year map as Gaussian reconstruction of the whole CMB sky."
    },
    "1004/1004.0690_arXiv.txt": {
        "abstract": "Previous studies of the active galactic nuclei (AGN) contribution to the cosmic X-ray background (CXB) consider only observable parameters such as luminosity and absorbing column.  Here, for the first time, we extend the study of the CXB to physical parameters including the Eddington ratio of the sources and the black hole mass.  In order to calculate the contribution to the CXB of AGN accreting at various Eddington ratios, an evolving Eddington ratio space density model is calculated.  In particular, Compton thick (CT) AGN are modeled as accreting at specific, physically motivated Eddington ratios instead of as a simple extension of the Compton thin type 2 AGN population. Comparing against the observed CT AGN space densities and $\\log N$--$\\log S$ relation indicates that CT AGN are likely a composite population of AGN made up of sources accreting either at $>$ 90$\\%$ or $<$ 1$\\%$ of their Eddington rate. ",
        "introduction": "\\label{sect:intro}  It is believed that all forms of active galactic nuclei (AGN) are variations of the same phenomenon, namely accreting supermassive black holes \\citep{R84}. The most basic observational description of AGN is their bolometric luminosity, $L_{bol}$, and the obscuring column density, $N_{\\mathrm{H}}$, through which they are observed. AGN are known to radiate at luminosities $L_{bol}\\approx$ 10$^{40}$--10$^{48}$ erg s$^{-1}$, but this is more physically described by the Eddington ratio, defined as $L_{bol}/L_{Edd}$, where $L_{Edd}\\equiv 4\\pi GM_{\\bullet}m_pc/\\sigma_T$ is the Eddington luminosity of a black hole with mass $M_{\\bullet}$.  Observed Eddington ratios range roughly ten orders of magnitude with a maximum of $\\sim$1 \\citep{CX07}.  There is a similarly wide range of observed column densities ($20\\leq\\log N_{\\mathrm{H}}\\la25$; e.g., Tueller \\etal 2008).  The most widely accepted theory to explain the observed range of obscuration is known as the unified model and claims that the obscuration of an AGN is dependent solely on the orientation of the AGN with respect to the line of sight to the observer (e.g., Antonucci 1993).  However, there is evidence that suggests that Compton thick (CT) obscuration, when $\\log N_{\\mathrm{H}}\\gtrsim 24$, is linked to an evolutionary phase during which the black hole and the host stellar spheroid accrete rapidly (e.g., Fabian 1999; Page \\etal 2004; Rigopoulou \\etal 2009).  As the black hole grows, the outflow from the black hole strengthens and ejects the obscuring material, revealing an unobscured quasar.  This outflow also halts the growth of the host spheroid causing the observed proportionality between black hole mass and host spheroid mass (e.g., Crenshaw \\etal 2003; Page \\etal 2004; Alexander \\etal 2010).  Galaxy merger simulations show that mergers between gas rich galaxies will cause gas to flow into the nuclear region, igniting a nuclear starburst and highly obscured, high Eddington ratio quasar activity (e.g., Hopkins et al. 2006).  \\citet{Fab09} also show that if an AGN is accreting at close to Eddington, radiation pressure on dust will blow out any column density with $\\log N_{\\mathrm{H}}\\la 24$. Therefore, if an AGN is obscured and has a high Eddington ratio, it must be CT.  Observational evidence also supports this evolutionary scenario.  \\citet{P04} found that star formation in the hosts of unobscured quasars has already peaked, whereas the hosts of obscured quasars are still undergoing massive amounts of star formation indicating that their galactic spheroids are still forming.  \\citet{T09a} found that the space density of luminous CT AGN candidates evolves strongly at $z=$ 1.5--2.5, so that the CT quasar population seems to peak at a slightly higher redshift than the population of unobscured quasars.  This evolutionary theory is difficult to observationally test since the high levels of obscuration found in CT AGN make them nearly invisible below 10 keV in their rest-frame.  If being CT is an early evolutionary phase of powerful quasars, CT AGN should be very rare at $z\\sim0$; however, {\\em Swift}/BAT and {\\em INTEGRAL} have detected several CT AGN at $z<0.025$ (Tueller \\etal 2008; Malizia \\etal 2010).  According to black hole masses and X-ray luminosities reported by \\citet{Gult09}, the two brightest CT AGN, the Circinus Galaxy and NGC 4945, have $\\lambda=-2.0$ and $\\lambda=-5.6$, respectively, where $\\lambda\\equiv\\log L_{bol}/L_{Edd}$. Therefore the scenario where CT AGN are high Eddington ratio quasars does not completely describe the entire CT AGN population.  As the majority of the cosmic X-ray background (CXB) $\\la$ 10 keV has been resolved into AGN by deep observations by {\\em ROSAT}, {\\em Chandra}, and XMM-{\\em Newton} (see Brandt \\& Hasinger 2005), the CXB provides a complete census of AGN.  Due to the observational constraints of highly obscured sources, CT AGN are difficult to study even in the local universe; thus, historically CT AGN have been invoked to augment the observed AGN population in population synthesis models to make these models properly fit the observed, but not resolved, peak of the CXB around 30 keV.  As there is a dearth of observational information about CT AGN, previous population synthesis models have treated CT AGN as a simple extension of type 2 AGN (e.g., Ueda \\etal 2003; Treister \\& Urry 2005; Ballantyne \\etal 2006; Gilli \\etal 2007); however, if CT AGN are part of an evolutionary sequence, then they would not evolve in the same manner as the less extremely obscured type 2 AGN.  Indeed, fitting the CXB in that manner seems to overpredict the local observed space density of CT AGN (see Treister et al. 2009a).  In this letter, we compute the contributions of AGN with various Eddington ratios to the CXB and model CT AGN using a physically motivated Eddington ratio distribution. We then compare our model predictions with observed cumulative and decadal space densities and the {\\em Swift}/BAT CT $\\log N$--$\\log S$ relation.  We find that CT AGN are a composite population of AGN accreting at both very high Eddington ratios and very low Eddington ratios.  When necessary, a $\\Lambda$CDM cosmology is assumed with $h_0=0.7$, $\\Omega_M=0.3$, and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "\\label{sect:sum}  We find that current data suggests that the majority of AGN accreting at close to the Eddington limit are CT, in agreement with galaxy merger simulations (e.g., Hopkins \\etal 2006) that show that mergers funnel gas into the central regions of galaxies igniting both starbursts and high Eddington ratio accretion onto the central black hole.  These AGN are able to accrete rapidly because of a large abundance of gas and dust in the central region of the host bulge. This abundance of gas and dust naturally leads to a very high obscuring column density \\citep{F99}.  Furthermore, radiation pressure from a rapidly accreting AGN will blow out any dusty gas with $\\log N_{\\mathrm{H}}\\la 24$ \\citep{Fab09}. Thus, in accordance with our findings, the vast majority of rapidly accreting, obscured AGN must be CT.  A new result is that the CT AGN population, in addition to including sources with $\\lambda>-0.05$, also includes sources with $\\lambda\\la-2.0$. While \\citet{H08} explains that most observed low Eddington ratio sources tend to be unobscured, \\citet{TW03} find evidence for a population of highly obscured low Eddington ratio AGN.  Indeed \\citet{M10} find an obscuration-luminosity relation in their sample of {\\em INTEGRAL} sources which claims that lower luminosity sources tend to be more obscured.  As these observed lower luminosity sources are local, these AGN are likely to have low Eddington ratios.  Since low Eddington ratio AGN are only weakly accreting, they will have very little effect on their environment.  Therefore molecular clouds would be able to come deep into the core of the bulge without being affected.  Indeed \\citet{TW03} find that the highly obscured low Eddington ratio sources in their sample might exhibit variable obscuration, as would be expected if the obscuration were due to molecular clouds near the black hole. Several studies have been conducted on so called 'changing-look' AGN (see Bianchi \\etal 2005 and references therein), which are AGN whose spectrum changes from Compton thin to reflection dominated on the time scale of years.  It is believed that the changes in the spectrum of these local, low luminosity AGN is due to either the nucleus being obstructed by a cloud with a CT column density or the nucleus switching off, so that the reflection dominated state is actually the echo of a previous accretion episode.  \\citet{B05} find that it is most likely that the four AGN in their sample are in CT states.  In a study of 82 LINERs, \\citet{G09} find about half of the LINERs in their sample appear to be CT and that there is a higher percentage of CT LINERs than CT Seyferts.  The average Eddington ratio of the CT LINERs in the sample used by \\citet{G09} is $\\lambda=$ -4.7.  This finding is consistent with the composite model for CT AGN discussed here which finds that $\\sim$60$\\%$ of AGN with $\\lambda < -2.0$ are CT.  It is clear that previous attempts to model the elusive CT AGN population as an extension of the Compton thin type 2 population fail to explain the observed CT AGN space density and $\\log N$--$\\log S$ relation.  \\citet{T09b} suggest that CT AGN only contribute $\\sim$10$\\%$ of the CXB and that the normalization of the CXB is overestimated by population synthesis models which predict a large population of CT AGN. However, changing the normalization of the CXB does not account for the strong evolution of CT AGN at $z=$ 1.5--2.5 found by \\citet{T09a}.  Here we assume that CT AGN contribute $\\sim$20$\\%$ of the CXB based on the findings of \\citet{DB09}, but the composite model presented here does account for the strong evolution of CT AGN at high redshift while not overestimating the local CT AGN space density.  Additionally, models which include CT sources with $-2.0<\\lambda<-0.5$ overpredict locally observed CT AGN space density.  Models which only include CT sources with $\\lambda<-2.0$ underpredict the density of high luminosity CT sources.  If CT sources are assumed to only have $\\lambda>-0.5$, the model overpredicts the observed $\\log N$--$\\log S$ relation.  A composite model, where the CT AGN population includes both sources accreting at close to Eddington and sources accreting at $<$1$\\%$ of Eddington, best explains observations.  As CT AGN are rare and difficult to observe due to high obscuration, it is difficult to use observations to directly decipher the physical phenomenon which gives rise to the extreme levels of obscuration found in CT AGN (e.g., Triester et al. 2009b; Rigby et al. 2009). Thus physical models of CT AGN must use indirect observational constraints, like observed space densities and number counts, to uncover the nature of the extreme obscuration of these sources. This letter presents a model which constrains the physical parameter of the Eddington ratios of CT AGN to $\\gtrsim$90$\\%$ and $\\la$1$\\%$.  In the future the understanding of CT AGN will be expanded through IR and sub-mm studies of the reprocessed radiation from CT AGN and very hard X-ray imaging with missions like {\\em NuSTAR}."
    },
    "1004/1004.0373_arXiv.txt": {
        "abstract": "The scintillation light yield of liquid argon from nuclear recoils relative to electronic recoils has been measured as a function of recoil energy from 10 keVr up to 250 keVr at zero electric field.  The scintillation efficiency, defined as the ratio of the nuclear recoil scintillation response to the electronic recoil response, is \\LeffErr~above 20 keVr. ",
        "introduction": "\\label{sec:intro}  A number of existing and proposed experiments use liquefied noble gases as detection media for Weakly Interacting Massive Particles (WIMPs)~\\cite{Aprile:2005a, Cline:2003, Brunetti:2005, Benetti:2007, Rubbia:2006}, a well motivated dark matter candidate~\\cite{Jungman:1996}. Liquefied noble gases have a high scintillation yield, are relatively simple to purify of both radioactive contaminants and light absorbers, and should be easily scalable to the large masses required for very sensitive detectors. Although the best limit for the spin-independent WIMP-nucleon cross section is currently set by the germanium-based CDMS experiment~\\cite{Ahmed:2009} at $3.8 \\times 10^{-44}$ cm$^2$ for a 70-GeV WIMP mass, the XENON-10 experiment has set a comparable limit of $8.8 \\times 10^{-44}$ cm$^{2}$ for a 100-GeV WIMP mass~\\cite{Angle:2007}, showing that liquefied noble gases are viable dark matter targets.  Events in a noble liquid dark-matter detector may arise from scattering off of the nucleus or atomic electrons; dark matter will only scatter off the nucleus to an appreciable extent.  The ratio of the scintillation light yield for nuclear recoil events relative to electronic recoil events is defined as the scintillation efficiency or \\SE.  A WIMP dark matter search requires an energy threshold on the order of tens of keV, and it is necessary to measure the scintillation efficiency down to this energy threshold so as to quantify the WIMP detection sensitivity. In order to make this measurement of \\SE, a D-D neutron generator was used to produce neutrons that scattered from a liquid argon detector into an organic liquid scintillator detector used as a coincidence trigger.  The organic scintillator was placed at a series of known angles, and the energies of the selected nuclear recoils in the liquid argon were kinematically determined. The scintillation efficiency was determined from the ratio of the measured electron-equivalent recoil energy at a given scattering angle to the expected nuclear recoil energy (keVr) at that angle. Details of this measurement in a 4-kg liquid argon detector are presented in this paper, along with scintillation efficiency results for nuclear recoil energies between 10 and 250 keVr at zero electric field. ",
        "conclusions": ""
    },
    "1004/1004.4766_arXiv.txt": {
        "abstract": "Polarimetry is widely considered a powerful observational technique in X-ray astronomy, useful to enhance our understanding of the emission mechanisms, geometry and magnetic field arrangement of many compact objects. However, the lack of suitable sensitive instrumentation in the X-ray energy band has been the limiting factor for its development in the last three decades. Up to now, polarization measurements have been made exclusively with Bragg diffraction at 45$^\\circ$ or Compton scattering at 90$^\\circ$ and the only unambiguous detection of X-ray polarization has been obtained for one of the brightest object in the X-ray sky, the Crab Nebula. Only recently, with the development of a new class of high sensitivity imaging detectors, the possibility to exploit the photoemission process to measure the photon polarization has become a reality. We will report on the performance of an imaging X-ray polarimeter based on photoelectric effect. The device derives the polarization information from the track of the photoelectrons imaged by a finely subdivided Gas Pixel Detector. It has a great sensitivity even with telescopes of modest area and can perform simultaneously good imaging, moderate spectroscopy and high rate timing. Being truly 2D it is non-dispersive and does not require any rotation. This device is included in the scientific payload of many proposals of satellite mission which have the potential to unveil polarimetry also in X-rays in a few years. ",
        "introduction": "X-ray polarimetry was born together with X-ray astronomy. First pioneering experiments were carried out in the seventies with polarimeters based on Bragg diffraction at 45$^\\circ$ or Compton scattering at 90$^\\circ$ on-board sounding rockets and first results were quite encouraging: already Novick et al. \\citep{Novick1972} reported a marginal yet significant detection of polarization in the emission of the Crab Nebula, confirmed a few years later with high significance by the Bragg polarimeter on-board OSO-8 \\citep{Weisskopf1978}. This observation was favoured by the intense flux of the source and the high degree of polarization, signature of synchrotron emission ($\\mathcal{P}\\approx$20\\%), and, as a matter of fact, it has remained unique. Only upper limits were derived for other astrophysical objects because of the combined effect of a lower flux and an inferior polarization degree \\citep{Long1979,Hughes1984}.  Unfortunately no other tool dedicated to X-ray polarimetry has been launched after OSO-8. The Stellar X-ray Polarimeter on-board the Spectrum-X-Gamma mission, although the flight model was ready and calibrated, was never put in orbit because of the collapse of Soviet System. The proposals to include polarimeters on-board observatories like XMM or AXAF, which were the only opportunities to have a sufficient collecting area, has never been carried out: instruments exploiting Bragg diffraction or Compton scattering didn't look attractive because, while imaging and spectroscopic devices promised an enormous increase of sensitivity, polarimetry would be limited to a few bright sources even in the focus of these large telescopes. Moreover ``classical'' polarimeters were cumbersome because they need to be rotated around the direction of incident photons.  Conversely, the lack of experimental feedback has not prevented the development of a rich literature on the basis of which we expect that almost all sources in the X-ray sky should emit partially polarized radiation \\citep[for reviews see][]{Rees1975,Meszaros1988,Weisskopf2009}. The study of the state of polarization would unveil the magnetic field and the geometry of the sources and it would pinpoint the emission processes at work, discriminating among competitive models otherwise equivalent from the spectral or the timing point of view. This is the case of emission geometry in pulsars \\citep{Dyks2004} or X-ray pulsars in binaries \\citep{Meszaros1988}, but peculiar signatures are also expected for isolated neutron stars because of the different opacity of the two normal modes in a magnetized plasma and because of vacuum polarization \\citep{Canuto1971,Pavlov2000,Lai2002,Heyl2003}. Moreover polarimetry is a powerful probe to investigate fundamental theories. General Relativity in the strong field regime can be tested by means of the rotation of the plane of polarization with energy expected for stellar-mass black-holes, and the amplitude of the effect would provide a measurement of the spin \\citep{Stark1977,Dovciak2008,Li2009}. Instead a rotation of the polarization angle increasing with distance could tightly constrain the vacuum birefringence expected in some theories of Quantum Gravity \\citep{Gambini1999,Mitrofanov2003,Kaaret2004}.  As a proof for the impelling interest in X-ray polarimetry, many authors have attempted to extract the polarization information as a byproduct of existing imaging devices. By selecting those events which are scattered and detected between two adjacent pixels, any imaging instrument is in principle a Compton polarimeter because the line connecting the hit pixels approximates the scattering direction. While some results have been achieved for the Crab Nebula with INTEGRAL \\citep{Dean2008,Forot2008}, measurements of this kind may be affected by strong systematic effects, are limited to strong sources because of the low Compton scattering probability in the detector and, in many case, remains questionable \\citep{Coburn2003,Rutledge2004}.  Today gas detectors able to image the tracks of photoelectrons provide a valuable alternative to classical techniques. Basically, they promise a jump in sensitivity thanks to the much higher capability to collect photons, obtained with a larger energy band with respect to Bragg polarimeters and a lower energy threshold than Compton instruments. In the following we present the Gas Pixel Detector (GPD hereafter), one of the first devices able to resolve photoelectron paths in a gas mixture at atmospheric pressure even at low energy and specifically designed for astrophysical application. ",
        "conclusions": "X-ray polarimetry is a fundamental tool for increasing our knowledge of emission processes acting in compact sources and for investigating theories of fundamental physics. Today gas detectors which resolve photoelectron tracks provide a valuable alternative to classical techniques of Bragg diffraction and Thomson scattering and promise to perform polarimetry of tens of galactic sources even within a small mission to be launched in a few years. The Gas Pixel Detector is in an advanced phase of development: the sensitivity expected on the basis of Monte Carlo simulations is confirmed by measurements and the detector survived irradiation, vibrations and thermo-vacuum tests. Once in-orbit with an X-ray telescope, the GPD will contemporaneously measure the polarization, the energy, the time and the direction of arrival of photons. Measurements at the level of 1\\% and below are within the possibility of the GPD because systematic effects are well under control thanks to the intrinsic azimuthal symmetry of the instrument. On-board of small missions, the GPD could reach between 2 and 10~keV a minimum detectable polarization of about 3\\% for 10~mCrab sources in 100~ks of observation and NHXM offers the possibility to extend polarimetry up to 35~keV with a similar sensitivity. The GPD will also be part of the scientific payload of the large mission IXO, whose large collecting area will allow for the measurements of polarization even of faint extragalactic sources (1\\% for 1~mCrab source in 100~ks) with a fine angular resolution."
    },
    "1004/1004.3693_arXiv.txt": {
        "abstract": "We explain the motivation and main idea of our work in Ref.~\\refcite{RenauxPetel:2009sj}. We present a simple model of multifield Dirac-Born-Infeld inflation whose bispectrum exhibits a linear combination of the equilateral and local shapes, which are usually considered as separate possibilities. We also point out the presence of a particularly interesting component of the primordial trispectrum. ",
        "introduction": "The study of the non-Gaussian properties of the primordial density fluctuations is a very promising tool with which to further discriminate between competing scenarios of the very early universe \\cite{Komatsu:2009kd}. The departure from Gaussianity is usually measured by means of connected $n$-point correlation functions of the primordial curvature perturbation $\\zeta$. Chief amongst these is the three point-function, or in Fourier space the bispectrum. The sum of the corresponding three momenta vanishes by virtue of translational invariance, thus forcing the momenta to form a triangle. The shape of the triangles for which the bispectrum signal is largest turns out to be a key discriminant between models \\cite{Babich:2004gb}. In particular, in the two well known types of scenarios that can generate large non-Gaussianities -- multiple field models and single-field models with non-standard kinetic terms --  the bispectrum peaks respectively for squeezed (the so-called local shape) and equilateral triangles, their amplitude being respectively characterized by the parameters $\\Floc$ and $\\Feq$. The best motivated model that produces equilateral non-Gaussianities is the so-called Dirac-Born-Infeld (DBI) scenario \\cite{Alishahiha:2004eh}. This is a string inspired model in which an extended object of three spatial dimensions (a D3-brane) evolves in the six extra dimensions of string theory. Its motion is governed by a non-standard Dirac-Born-Infeld action -- hence its name -- resulting in potentially large non-Gaussianities. Actually, the simplest models, in which the brane moves along a single radial direction, are already under strain because the predicted $\\Feq$ is in excess of the observational bound \\cite{Baumann:2006cd}. However, the brane can a priori move and fluctuate in each of the six extra dimensions, in which case a multifield description is required. It was shown in a model-independent way that the amplitude of equilateral non-Gaussianities in this more general framework is reduced compared to the single-field case \\cite{Langlois:2008wt,Langlois:2008qf,Langlois:2009ej}, which is therefore of crucial importance for model-building. Here we report on an example of a multifield DBI inflationary scenario where it is possible to assess the importance of this supression. We also point out that, because of multiple fields effects, the large scale evolution can be highly nonlinear, resulting in possibly large local non-Gaussianities, thus giving the first example in which both equilateral and local non-Gaussianities -- usually considered as separate possibilities -- are present at an observable level in the same model. Their combined presence also manifests itself non-trivially in the connected 4-point function of $\\zeta$, the trispectrum. ",
        "conclusions": "The common presence of non-standard kinetic terms and multiple-field effects implies that multifield DBI inflation is able to produce both large equilateral and local non-Gaussianities in the bispectrum. Moreover, due to the presence of light scalar fields with non standard kinetic terms, other than the inflaton, it can be shown that the trispectrum acquires a component with a particular momentum-dependence whose amplitude is given by the product $\\Floc \\, \\Feq$ \\cite{RenauxPetel:2009sj}. This consistency relation is valid for every DBI model, not restricting to a radial trajectory nor to a specific scenario for the entropy to curvature transfer. It thus constitutes an interesting observational signature of multifield DBI inflation which one can hope to test with forthcoming experiments."
    },
    "1004/1004.1143_arXiv.txt": {
        "abstract": "It has been suggested that moons around transiting exoplanets may cause observable signal in transit photometry or in the Rossiter-McLaughlin (RM) effect. In this paper a detailed analysis of parameter reconstruction from the RM effect is presented for various planet-moon configurations, described with 20 parameters. We also demonstrate the benefits of combining photometry with the RM effect. We simulated 2.7$\\times 10^9$ configurations of a generic transiting system to map the confidence region of the parameters of the moon, find the correlated parameters and determine the validity of reconstructions. The main conclusion is that the strictest constraints  from the RM effect are expected for the radius of the moon. In some cases there is also meaningful information on its orbital period. When the transit time of the moon is exactly known, for example, from transit photometry, the angle parameters of the moon's orbit will also be constrained from the RM effect. From transit light curves the mass can be determined, and combining this result with the radius from the RM effect, the experimental determination of the density of the moon is also possible. ",
        "introduction": "The number of known transiting exoplanets is rapidly increasing, which has recently inspired significant interest as to whether they can host a detectable moon (e.g. Szab\\'o et al. 2006, Simon et al. 2007, Kipping 2008, 2009, Simon et al. 2009, Kipping et al. 2009). Historically, our Moon has constantly inspired scientific research and it has played a key role in supporting life on Earth (e.g. Wagner 1936, Asimov 1979). It may be that the presence of a large exomoon is a {\\it sine qua non} requirement for the development of an intelligent civilization on an exoplanet.  Although there has been no such example where the presence of a satellite was proven, several methods have already been investigated for such a detection in the future (barycentric Transit Timing Variation, TTV Sartoretti \\&{} Schneider 1999, Kipping, 2008; photocentric Transit Timing Variation, TTV$_{\\rm p}$ Szab\\'o et al. 2006, Simon et al. 2007; Transit Duration Variation, TDV, Kipping, 2009; Time-of-Arrival analysis of pulsars, Lewis et al. 2008; microlensing, Liebig \\& Wambsganss 2009). Deviations from perfect periodic timing of transits might suggest the presence of a moon (D\\'\\i{}az et al. 2008), perturbing planets (Agol et al. 2005) or indicate periastron precession (P\\'al \\&{} Kocsis 2008).  All these methods (excluding microlensing) rely on transit photometry. In the era of ultraprecise space photometry (CoRoT, Kepler), one can expect accurate light curves to 0.1 mmag  that promises the discovery of Moon-like satellites of Earth-like planets (Szab\\'o et al. 2006, Kipping et. al 2009). Additionally, the $\\sim$1 cm/s velocimetric accuracy is promised with laser frequency combs (Li et al. 2008). Radial velocity measurements during a transit already play an important role in understanding the planet via its Rossiter-McLaughlin (RM) effect (Gaudi \\&{} Winn 2007), and the prospects of exomoon detection in this way are quite encouraging.   Here we continue our previous investigations by invoking the RM effect as a possible tool in characterising exomoons. Earlier we described a photometric method, the Photocentric Transit Timing Variation, TTV$_{\\rm p}$ (Szab\\'o{} et al. 2006) that is very sensitive to the presence of transiting moons. In Simon et al. (2007) we examined which moons can be detected with that method in space observatory measurements, with respect to different values of M masses, R radii, and P orbital periods. In Simon et al. (2009) we demonstrated that another method, based on the Rossiter-McLaughlin effect of the moon, is also capable of attaining observational signature of a possible satellite. Now we give a full description of the parameter reconstruction from the RM effect. An error analysis is also presented: from simulated transits we examine which parameters of the satellite can be recovered at certain S/N of the measurements. ",
        "conclusions": "\\begin{figure} \\includegraphics[width=8cm]{hist_r.eps} \\includegraphics[width=8cm]{hist_p.eps}% \\caption{Posterior probability distributions of $R$ and $P$, marginalized from the fit likelihoods in the $P$--$R$ subspace with uniform priors. The three different distributions refer to S/N levels of 1, 2 and 5 with the same colour coding as in Figs. 3--4--5. The arrow shows the input value of the moon parameters.} \\end{figure}   The joint confidence intervals of the parameter pairs are shown in 20 panels in Figs. 4 and 5. Fig. 4 shows sections where the plotted parameters promise a reliable reconstruction. The best results are given for the size of the moon, that is closely linked to the amplitude of the residuals in the RM curve, after subtracting the best-fit planet. The radius is well reproduced in all sections and did not suffer serious degenerations. The topmost row suggests that the size is a little biased and the inclination of the moon is essentially unknown. This is because the residuals of the planet fitting is forced to be close to zero, consequently the size of the moon is slightly underestimated. For this smaller moon, a lower inclination parameter is preferred.  The ascending node of the orbit is somewhat correlated with the radius. Interestingly, once other parameters are well constrained, $\\Omega$ can be determined very accurately as Fig. 4 shows (see the right panel in the second row). The 3rd and 5th rows of Fig. 4 demonstrate that there is no information on the mass of the satellite from the analysis of radial velocity, all probed masses are equally probable. The most important conclusion is that the mass of the satellite must be omitted from the analysis, i.e. a fixed value for the mass or a fixed assumption on the density of the satellite will lead to equally reliable solutions in the other parameters with a much faster process. The abscissa of the 4th and 5th panels in Fig. 4 show the orbital period of the moon. Surprisingly, there is some information in the light curve for this period, as values between 2--10 days are preferred (the model moon has 5 days period). Because of the little position change of the moon during the transit, the residual RM effect due to the moon will last a somewhat shorter or longer time than that of the planet, which may be detected.  Fig. 5 shows further sections where degenerations are prominent. The joint confidence intervals evidently show that the angle parameters, $\\iota$, $\\phi$ and $\\Omega$ are seriously interrelated. Acceptable solutions can be characterized with very different values of $\\Omega$ or $\\phi$, because the orbital period of the moon is unknown (see the 2nd and 3rd rows). The 4th row shows that $\\iota$ data do not constrain $\\phi$ well. We suggest that a reliable value for $\\iota$ will have to be assumed in practical applications, and the other two angle parameters should be taken out of fitting (they can be either marginalized or evaluated along some prior with Bayesian analysis).  However, these limitations of parameter determination do not lessen dramatically the power of radial velocity analysis in the exploration of exomoons. This method nicely completes the evaluation of photometry, since the mass can be determined from (barycentric) TTV. Simon et al. (2007) showed that from photocentric TTV$_{\\rm p}$ we get some information on the radius of the satellite, but this is correlated with the density of the moon. The analysis of RM effects give prior information on the size of the moon, and the combination of all methods, theoretically, may lead to the direct experimental determination of the density of the exomoons.  In Fig. 6, we show the posterior probability distributions of $P$ and $R$, marginalized from the likelihood data with assuming uniform priors to all variables. The reconstruction of $R$ is satisfactory, although the radius is somewhat biased toward smaller sizes. There is some information for the period, too, which is somewhat surprising as the orbital period of the moon is $\\approx 1/12$ transit duration. The increasing noise level does not bias the mode of the distributions, but gives the wings slightly more weight.  \\subsection{The general case of main sequence dwarf stars}  We have shown the reconstruction of parameters for a 0.8 M$_\\odot$ main-sequence star, which has a spectral type of K0 in case of solar metallicity. It is very important to note that our results are general and indicative for other types of stars. Although the signal will be smaller for bigger stars, the shape of the curves do not change significantly, hence error propagation will follow the same scenario as presented in Figs. 4-5. Consequently, the general behaviour of the parameters concerning their stability and degenerations will be the same. For a general case, we propose that the moon's radius is the best parameter for reconstruction; in some cases, the orbital period might be also constrained.  Earlier type stars have larger radius, hence the area eclipsed by the moon is a smaller fraction of the projected stellar disk. The other parameter that determines the $A_{RM}$ half-amplitude of the RM effect is the $v_{rot} \\sin i$ rotation velocity of the star, such that \\begin{equation} A_{RM} \\ \\propto \\ {\\left( R \\over R_*\\right)^2 } \\ v_{rot}\\sin i, \\end{equation} where $R$ denotes the radius of the moon, $R_*$ is the stellar radius and we can assume that $\\sin i \\approx 1$ for those systems that display transits and RM effect. Equality is true in Eq. 4 if we neglect limb darkening. Limb darkening can reduce the amplitude of the RM effect by 20--40\\%, depending on the exact intensity profile of the stellar disk. We can combine the RM effect of a moon and a planet together, \\begin{equation} A_{RM,m+p} \\propto {R_p^2 + R^2 \\over R_*^2} \\ v_{rot}. \\end{equation} Now an upper estimate can be given for the size of the moon if it is not detected in the residuals of the RM curve. In this case the RM effect of the moon is hidden in the scatter of the radial velocity data, i.e. $3 \\sigma_{v_{rad}}$ is larger than the satellite's effect: \\begin{equation} 3 \\sigma_{v_{rad}}>A_{RM}={R^2\\over R_*^2}v_{rot}. \\end{equation} Rearranging this for the size of the moon, the upper limit is given as $R<\\sqrt{3 \\sigma_{v_{rad}} / v_{rot}} R_*$, or simply substituting the amplitude of the measured RM effect (and assuming $A_{RM, m+p} \\approx A_{RM,planet} = R_p^2/R_*^2 v_{rot}$), \\begin{equation} R<\\sqrt{3 \\sigma_{v_{rad}} \\over A_{RM, planet}} R_p. \\end{equation}  The confidence of this estimate is 99.9 \\%, i.e. $3 \\sigma$ confidence.  Which spectral types represent the best candidates for a successful detection of exomoons with the RM effect? To answer this question, we performed simple calculations for stars in the Pleiades open cluster, assuming that every star has a planet with a Ganymede-sized moon in central transit. Our intention was to predict the amplitude of the satellite's RM effect (assumed to be independent of that of the planet), as a function of stellar mass.  The $B-V$ colours and $v\\sin i$ data for Pleiades stars were taken from Queloz et al. (1998). Stellar masses and radii were estimated from the $B-V$ colour, using the latest Padova isochrones (Bertelli et al. 2008), adopting 70 Myr for the age and Z=0.017 for the metallicity (Boesgaard \\& Friel 1990). We calculated the amplitude of the RM effect according to Eq. 4, inserting the stellar radius and rotation velocity for each star, and substituting the size of Ganymede. The results are plotted in Fig. 7 with open circles.  Although there are hints of a tendency, the scatter is large, which can be explained by the different spin axis orientations and different rotation evolution for each star. The shape of the distribution in Fig. 7 can be better seen using a statistical relationship for the rotation period of main-sequence stars determined by Barnes (2007): $P_{rot} \\propto (B-V-0.4)^{0.601} t^{0.52}\\ $ days, where $t$ is the age of the star. This formula is singular at $B-V=0.4$, thus it is valid for stars later than F5, approximately. Similarly to the individual stars, the $B-V$ colours of the isochrone points have been converted to masses and radii. The rotation velocity has then been calculated as $v_{rot}=2\\pi R_*/P$ (we kept assuming $\\sin i\\approx 1$). To estimate RM amplitudes, we again used Eq. 4 and the size of Ganymede. In Fig. 7, the solid line shows the resulting average RM amplitude. We conclude that stars below 0.6--0.8 M$_\\odot$ offer the best opportunity to detect the RM effect of the exomoons. Stars with masses greater than 1.2 M$_\\odot$ do also show larger effect, however, stellar variability quickly becomes dominant with the increasing mass.  \\begin{figure} \\includegraphics[angle=270,bb= 153 100 534 580,width=8cm]{rmeffect.ps} \\caption{Half amplitude of the Rossiter-McLaughlin effect due to Ganymede-sized moons of planets orbiting G, K and M dwarfs. Solid line: empirical model based on Barnes (2007) rotation model and isochrones. Open circles: individual stars of the Pleiades. The dashed curve shows the amplitude of solar-like oscillations.} \\end{figure}  This variability has two dominant components: the jitter due to convective motions (including the convectively excited solar-like oscillations) and the rotational modulation due to stellar activity. Among the brighter dwarf stars, old, inactive G and K dwarfs offer the best performance: a few stars are known to have $<$1 m/s jitter, while they typically have jitter levels in the 1--5 m/s regime (Saar et al. 2003, Wright 2005, O'Toole et al. 2008). On the contrary, some F-type stars display large jitter that even challenges asteroseismology (see the example of Procyon in Arentoft et al. 2008) and the detection of long-period planets (Lagrange et al. 2009). The jitter from solar-like oscillations can be estimated via the scaling relation of the velocity amplitude, which depends on the light-to-mass ratio of the star: \\begin{equation} v_{osc}={L/L_\\odot \\over M/M_\\odot }\\left( 23.4\\pm 1.4 \\right) \\ \\ {\\rm cm/s} \\end{equation} (Kjeldsen \\& Bedding 1995). The predicted oscillation velocity amplitudes are also plotted in Fig. 7 with the dashed line (blue in colour). Since M-dwarf stars have very small $L/M$ ratio, and consequently the amplitude of solar-like oscillations is tiny, they are promising candidates to be quested for exomoons. These stars are faint in the visual, but recent work of  Bean et al. (2009), for example, has opened the door to the sub-m/s velocimetric accuracy in the infrared with CRIRES on VLT.  It is worth noting that late M-dwarfs (beyond M4) exhibit larger rotation velocities (Jenkins et al. 2009), which makes them difficult targets for high-precision RV measurements because rapid rotation washes out the spectral features (Bouchy et al. 2001). Hence, the best targets for exomoon exploration are the K and early M dwarf stars, for which both rotation and activity reach a minimum level (Jenkins et al. 2009, Wright 2005).  In case of higher stellar activity, the situation is not entirely hopeless, as illustrated by a recent study by Queloz et al. (2009), who filtered out activity with a Fourier polynomial using the first rotation harmonics. This way they pushed the residuals from $\\pm 20$ m/s to $\\pm$ 5 m/s. Similar residual levels in fitting the RM effect were reached by Triaud et al. (2009).  It is known that the frequency of giant planets increases linearly with the parent-star mass for stars between 0.4 and 3 $M_{\\odot }$ (Ida \\&{} Lin 2005, Kennedy \\&{} Kenyon 2008), with e.g., 6$\\%$ frequency of giant planets around 1 $M_{\\odot }$ and 10$\\%$ frequency around 1.5 $M_{\\odot }$. However, we know planets around red dwarfs and there is observational indication for a few multiple planetary systems among them (e.g. Rivera et al. 2005). The possible detection of exomoons around planets of larger stars, if these satellites ever exist, is a significant challenge for signal processing to minimise the ambiguity caused by the higher level of velocity jitter.   A more elaborated distinction between signals of exoplanets and signals from stellar physics requires a deep analysis that is beyond the scope of this paper. There is reason for some optimism because the time-scales of the exoplanet-exomoon systems and that of the stellar signals are usually very different. Moreover, rapid development in instrumentation may reach levels of precision that were unimaginable even a few years ago."
    },
    "1004/1004.4620_arXiv.txt": {
        "abstract": "\\vskip .3cm Effects of quantized free scalar fields in cosmological spacetimes with Big Rip singularities are investigated.  The energy densities for these fields are computed at late times when the expansion is very rapid.  For the massless minimally coupled field it is shown that an attractor state exists in the sense that, for a large class of states, the energy density of the field asymptotically approaches the energy density it would have if it was in the attractor state. Results of numerical computations of the energy density for the massless minimally coupled field and for massive fields with minimal and conformal couplings to the scalar curvature are presented. For the massive fields the energy density is seen to always asymptotically approach that of the corresponding massless field.  The question of whether the energy densities of quantized fields can be large enough for backreaction effects to remove the Big Rip singularity is addressed. ",
        "introduction": "\\label{introduction}  Surveys of type Ia supernovae and detailed mappings of the cosmic microwave background provide strong evidence that the universe is accelerating \\cite{Supernova}.  To explain this acceleration within the framework of Einstein's theory of General Relativity requires the existence of some form of ``dark energy'' which has positive energy density and negative pressure \\cite{DE-Review}.  One common model for dark energy \\cite{DE-Review} is to treat it as a pervasive, homogenous perfect fluid with equation of state $p = w \\rho$.  Cosmic acceleration demands that $w < -1/3$, and observations from the Wilkinson Microwave Anisotropy Probe in conjunction with supernova surveys and baryon acoustic oscillations measurements place the current value at $w=-0.999^{+0.057}_{-0.056}$~\\cite{WMAP}.  Although this is consistent with the effective equation of state for a cosmological constant, $w=-1$, we cannot rule out the possibility that our Universe contains ``phantom energy'', for which $w < -1$.  If $w$ is a constant and $w < -1$ then general relativity predicts that as the universe expands the phantom energy density increases with the result that in a finite amount of proper time the phantom energy density will become infinite and the universe will expand by an infinite amount.  All bound objects, from clusters of galaxies to atomic nuclei, will become unbound as the Universe approaches this future singularity, aptly called the ``Big Rip''~\\cite{Caldwell}.  A simple model of a spacetime with a Big Rip singularity can be obtained by considering a spatially flat Robertson-Walker spacetime, with metric \\eq ds^2 &=& -dt^2 + a^2(t)(dx^2 + dy^2 + dz^2). \\label{FRW} \\qed At late times the phantom energy density is much larger than the energy density of classical matter and radiation and the solution to the classical Einstein equations is \\eq a(t) \\approx a_1 (t_r - t)^{-\\sigma},  \\label{abigrip} \\qed with \\eq \\sigma = -2/\\left(3+3w\\right) > 0 \\;. \\qed Here $a_1$ is a constant and $t_r$ is the time that the Big Rip singularity occurs.  The phantom energy density is \\eq \\rho_{ph} &=& \\frac{3}{8\\pi}\\sigma^2 (t_r-t)^{-2}.  \\label{rhoph} \\qed Note that both the scale factor and the phantom energy density become infinite at $t=t_r$.  In addition to the singularity in the model described above, other classes of models containing somewhat milder phantomlike future singularities have been identified. Barrow~\\cite{Barrow} constructed a class of models called ``sudden singularities'' in which $w$ is allowed to vary as $w \\propto (t_r - t)^{\\alpha-1}$ for $0 < \\alpha < 1$. In these models, the pressure and scalar curvature diverge at time $t_r$ but the energy density and scale factor remain finite. Other types of singular behavior were found in Refs.~\\cite{Barrow2, NO2, stefancic}. Nojiri, Odintsov and Tsujikawa~\\cite{NOT} came up with a general classification scheme for future singularities.  The strongest singularities are classified as type I and Big Rip singularities fall into this class.  The sudden singularities are examples of type II singularities.  Two other classes, type III and type IV were also identified.  For type III singularities both the energy density $\\rho$ and the pressure $p$ diverge at the time $t_r$ but the scale factor $a$ remains finite. For type IV singularities, the scale factor remains finite, the energy density $\\rho$ and the pressure $p$ go to zero at the time $t_r$, but divergences in higher derivatives of $H = \\dot{a}/a$ occur.  Other classification schemes for cosmological singularities have been given in Refs.~\\cite{cattoen-visser, dabrowski, fernandez-lazkoz}.  At times close to $t_r$ in each of the above cases it is possible that quantum effects could become large and that the backreaction of such effects could moderate or remove the final singularity. One way to investigate whether this would occur is to compute the energy density for the quantum fields in the background geometry of a spacetime with a final singularity. Then a comparison can be made between the phantom energy density and the energy density of the quantized fields. A second way is to solve the semiclassical backreaction equations to directly see what effects the quantum fields have.  Nojiri and Odintsov~\\cite{NO1,NO2} studied the backreaction of conformally invariant scalar fields in the cases of sudden singularities and Big Rip singularities and found that quantum effects could delay, weaken or possibly remove the singularity at late times.  In Ref.~\\cite{NOT}  Nojiri, Odintsov and Tsujikawa used a model for the dark energy with an adjustable equation of state to find examples of spacetimes with type I, II, and III singularities.  They then solved the semiclassical backreaction equations and found that the singularities were usually either moderated or removed by quantum effects.  Calder\\'{o}n and Hiscock~\\cite{Hiscock-Calderon} investigated the effects of conformally invariant scalar, spinor, and vector fields on Big Rip singularities by computing the stress-energy of the quantized fields in spacetimes with constant values of $w$.  Their results depend on the value of $w$ and on the values of the renormalization parameters for the fields.  For values of $w$ that are realistic for our universe they found that quantum effects serve to strengthen the singularity. Calder\\'{o}n~\\cite{Calderon} made a similar computation in spacetimes with sudden singularities and found that whether the singularity is strengthened or weakened depends on the sign of one of the renormalization parameters.  Barrow, Batista, Fabris and Houndjo~\\cite{Barrow-Batista} considered models with sudden singularities when a massless, minimally coupled scalar field is present.  They found that in the limit $t \\rightarrow t_r$, the energy density of this field remains small in comparison with the phantom energy density.  Thus quantum effects are never important in this case. Batista, Fabris and Houndjo~\\cite{Batista} investigated the effects of particle production when a massless minimally coupled scalar field is present in spacetimes where $w$ is a constant.  To do so they used a state for which  Bunch and Davies~\\cite{bunch-davies} had previously computed the stress-energy tensor.  They found that the energy density of the created particles never dominates over the phantom energy density.  Pavlov~\\cite{Pavlov} computed both the number density of created particles and the stress-energy tensor for a conformally coupled massive scalar field for the case in which $w = -5/3$.  It was found that backreaction effects are not important for masses much smaller than the Planck mass and times which are early enough that the time until the Big Rip occurs is greater than the Planck time.  In this paper we compute the energy densities of both massless and massive scalar fields with conformal and minimal couplings to the scalar curvature in spacetimes with Big Rip singularities in which the parameter $w$ is a constant. While our calculations are for scalar fields, it is worth noting that both massive and massless conformally coupled scalar fields can be used to model spin $1/2$ and spin $1$ fields, and in homogeneous and isotropic spacetimes the massless minimally coupled scalar field can serve as a model for gravitons~\\cite{Grischuck,Ford-Parker}.  For conformally invariant fields the natural choice of vacuum state in homogeneous and isotropic spacetimes is the conformal vacuum~\\cite{b-d-book}.  For all other fields there is usually no natural choice.  However, it is possible to define a class of states called adiabatic vacuum states which, when the universe is expanding slowly, can serve as reasonable vacuum states~\\cite{b-d-book}.  They can be obtained using a WKB approximation for the mode functions, and they are specified by the order of the WKB approximation. It has been shown that the renormalized stress-energy tensor for a quantum field is always finite if a fourth order or higher adiabatic vacuum state is chosen.\\footnote{In this paper we generalize the definition of an n'th order adiabatic vacuum state to include all states whose high momentum modes are specified by an n'th order WKB approximation but whose other modes can be specified in any way.}  Here we compute the renormalized energy densities of the quantum fields in fourth order or higher adiabatic states and investigate their behavior as the universe expands.  One focus is on the differences that occur for the same field in different states.  We find in all cases considered that the asymptotic behavior of the energy density is always the same for fields with the same coupling to the scalar curvature, regardless of whether they are massless or massive and regardless of what states the fields are in.  Fields with minimal coupling to the scalar curvature have a different asymptotic behavior than those with conformal coupling.  We also address the question of whether and under what conditions the energy density of the quantized fields becomes comparable to the phantom energy density.  We find that for fields in realistic states for which the energy density of the quantized fields is small compared to that of the phantom energy density at early times, and for spacetimes with realistic values of $w$, there is no evidence that quantum effects become large enough to significantly affect the expansion of the spacetime until the spacetime curvature is of the order of the Planck scale or larger, at which point the semiclassical approximation breaks down.  In Sec. II the quantization of a scalar field in a spatially flat Robertson-Walker spacetime is reviewed along with a method of constructing adiabatic states. In Sec. III, the energy density for massless scalar fields with conformal and minimal coupling to the scalar curvature is discussed and a comparison is made with the phantom energy density. For the massless minimally coupled scalar field a proof is given that one particular state serves as an attractor state in the sense that for a large class of states, the energy density of the field asymptotically approaches the energy density it would have if it was in the attractor state. In Sec. IV numerical calculations of the energy density for massive scalar fields with conformal and minimal coupling to the scalar curvature are discussed. A comparison is made with both the phantom energy density and the energy density of the corresponding massless scalar field. Our main results are summarized and discussed in Sec. V. Throughout units are used such that $\\hbar = c =  G = 1$ and our sign conventions are those of Misner, Thorne, and Wheeler~\\cite{MTW}. ",
        "conclusions": "We have computed the energy densities of both massless and massive quantized scalar fields with conformal and minimal coupling to the scalar curvature in spacetimes with big rip singularities. We have restricted attention to states which result in a stress-energy tensor which is homogeneous and isotropic and free of ultraviolet and infrared divergences.  For the numerical computations we have further restricted attention to states for which, near the initial time of the calculation, the energy density of the quantum field is much less than that of the phantom field.  For the massless minimally coupled scalar field we have shown that the energy density for the field in any fourth order or higher adiabatic state for which the stress-energy tensor is homogeneous, isotropic, and free of infrared divergences, always asymptotically approaches the energy density which this field has in the Bunch-Davies state.  In this sense the Bunch-Davies state is an attractor state.  For massive minimally coupled scalar fields numerical computations have been made of the energy density for different fourth order adiabatic states and in every case considered the energy density approaches that of the Bunch-Davies state for the massless minimally coupled scalar field at late times. For conformally coupled massive scalar fields numerical computations have also been made of the energy density for different fourth order adiabatic states.  In each case considered the energy density asymptotically approaches that of the massless conformally coupled scalar field in the conformal vacuum state.  Thus it appears that the asymptotic behavior of the energy density of a quantized scalar field in a spacetime with a big rip singularity depends only upon the coupling of the field to the scalar curvature and not upon the mass of the field or which state it is in, at least within the class of states we are considering.  Analytic expressions for the energy densities of both the massless conformally coupled scalar field in the conformal vacuum state and the massless minimally coupled scalar field in the Bunch-Davies state, in spacetimes with big rip singularities have been previously obtained~\\cite{b-d-book, bunch-davies} and are shown in Eqs.~\\eqref{ccm0} and \\eqref{BD}. To investigate the question of whether backreaction effects are important in these cases, the energy density of the scalar field can be compared to the phantom energy density to see if there is any time at which they are equal or at least comparable.  Then one can determine whether the semiclassical approximation is likely to be valid at this time by evaluating the scalar curvature and seeing whether or not it is well below the Planck scale.  We have done this and find that for the conformally coupled field the two energy densities are equal at the point when the scalar curvature is at the Planck scale if $w \\sim -1500.$  For the minimally coupled scalar field this combination occurs if $w \\sim -140$. Thus for $w \\ll -1500$ for the conformally coupled field and $w \\ll -140$ for the minimally coupled field, one expects that backreaction effects may be important at times when the scalar curvature is well below the Planck scale. However, the values of $w$ which satisfy these constraints are completely ruled out by cosmological observations.  Another way in which the energy density of a quantum field can be comparable to the phantom energy density at scales well below the Planck scale is to construct a state for which this is true.  There is no doubt that such states exist. The analytic and numerical evidence we have is that over long periods of time the energy density of a conformally or minimally coupled scalar field in such a state would decrease and at late enough times become comparable to that of a massless field in the conformal or Bunch-Davies vacuum state respectively.  More importantly one can ask whether any such states exist which are realistic for the universe that we live in.  This seems unlikely because today is certainly an early time compared to $t_r$ if our universe does have a big rip singularity in its future; however the energy densities of quantized fields today are much less than that of the dark energy.  Thus it would be necessary for the energy density of a quantum field to be small compared to the phantom energy density today and then to grow fast enough to become comparable to it well before the Planck scale is reached. This type of behavior seems highly artificial, particularly since it is not what happens for a massless scalar field in the conformal or Bunch-Davies vacuum states.  Thus we find no evidence which would lead us to believe that backreaction effects due to quantum fields would remove a big rip singularity in our universe if, indeed, such a singularity lies in our future."
    },
    "1004/1004.1233_arXiv.txt": {
        "abstract": "We show that oblique propagation of electrons in crystals of Ge and Si, where the electron velocity does not follow the electric field even on average, can be explained using standard anisotropic theory for indirect gap semiconductors.  These effects are pronounced at temperatures below $\\sim$1~K and for electric fields below $\\sim$5~V/cm because inter-valley transitions are energetically suppressed forcing electrons to remain in the same band valley throughout their motion and the valleys to separate in position space. To model, we start with an isotropic approximation which incorporates the average properties of the crystals with one phonon mode, and include the ellipsoidal electron valleys by transforming into a momentum space where constant energy surfaces are spheres.  We include comparisons of simulated versus measured drift velocities for holes and electrons, and explain the large discrepancy between electrons and holes for shared events in adjacent electrodes. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.3831_arXiv.txt": {
        "abstract": "{One of the most intriguing features revealed by the \\emph{Swift} satellite is the existence of flares superimposed to the Gamma-Ray Burst (GRB) X-ray light curves. The vast majority of flares occurs before $1000$ s, but some of them can be found up to $10^6$ s after the main event. } {In this paper we shed light on \\emph{late} time (i.e. with peak time $t_{pk} \\gtrsim 1000$ s) flaring activity.  We address the morphology and energetic of flares in the window $\\sim 10^3-10^6$ s to put constraints on the temporal evolution of the flare properties and to identify possible differences in the mechanism producing the early and late time flaring emission, if any. This requires the complete understanding of the observational biases affecting the detection of X-ray flares superimposed on a fading continuum at $t > 1000$ s. } {We consider all the \\emph{Swift} GRBs that exhibit late time flares. Our sample consists of $36$ flares, $14$ with redshift measurements. We inherit the strategy of data analysis from  Chincarini et al. (2010) in order to make a direct comparison with the early time flare properties.} {The morphology of the flare light curve is the same for both early time and late time flares, while they differ energetically. The width of late time flares increases with time similarly to the early time flares. Simulations confirmed that the increase of the width with time is not due to the decaying statistics, at least up to $10^4$ s. The energy output of late time flares is one order of magnitude lower than the early time flare one, being $\\sim 1\\% E_{prompt}$.  The evolution of the peak luminosity as well as the distribution of the peak flux-to-continuum ratio for late time flares indicate that the flaring emission is decoupled from the underlying continuum, differently from early time flares/steep decay. A sizable fraction of late time flares are compatible with afterglow variability.} {The internal shock origin seems the most promising explanation for flares. However, some differences that emerge between late and early time flares suggest that there could be no unique explanation about the nature of late time flares.} ",
        "introduction": "One of the most intriguing and unexpected features revealed by the X-Ray Telescope \\citep[XRT,][]{2005SSRv..120..165B} on board the \\emph{Swift} satellite \\citep{2004ApJ...611.1005G} is the existence of flares superimposed to the Gamma-Ray Burst (GRB) X-ray light curves \\citep{2005Sci...309.1833B,2006ApJ...641.1010F,2007ApJ...671.1903C,2007ApJ...671.1921F}. The vast majority of flares occurs before $1000$ s \\citep{2007ApJ...671.1903C}, but some of them can be found up to $10^6$ s after the main event.  Recent analyses of the flare temporal and spectral properties (\\citealp{chinca10}, hereafter C10; \\citealp{giantflares10}) of a large sample of \\emph{early} time (i.e. with peak time $t_{pk} \\lesssim 1000$ s) flares and of a subsample of bright flares revealed close similarities between them and the prompt emission pulses, pointing to an internal origin of their emission \\citep{2007ApJ...671.1903C,chinca10}. Therefore, the central engine itself should remain active and variable for long time. Alternatively, flaring emission can be powered by delayed magnetic dissipation during the deceleration of the ejecta \\citep{2006A&A...455L...5G}. Despite many efforts from both the observational and the theoretical point of view, a clear explanation of the X-ray flares is still missing (C10).  In this paper we shed light on \\emph{late} time (i.e. $t_{pk} \\gtrsim 1000$ s) flaring activity observed by XRT in the $0.3-10$ keV energy band. Its existence poses constraints on the theoretical models since late time flares require huge releases of energy ($\\sim 10^{50}$ erg) up to $1$ month after the main event. A previous work on late time X-ray flares has been presented by \\citet{2008A&A...487..533C}. These authors analysed a sample of $7$ GRBs that exhibit flares after $10^4$ s by fitting them with a gaussian profile plus power-law underlying continuum. The purpose was to compare the temporal and spectral properties of the underlying continuum and of the temporal and flux amplitude variability of these flares with the results found in \\citet{2007ApJ...671.1903C} and with the prescription of the internal and the external shock models. They concluded that late time flares are not different from the early time ones. However, due to the small number of flares in their sample, this statement needed further investigation.  The present work addresses the morphology of flares in the temporal window $\\sim 10^3-10^6$ s of a larger sample than in \\citet{2008A&A...487..533C}. The good fraction of flares of the sample with redshift measurement allows us also to characterise the energetic properties at very late times. We inherit the C10 strategy of data analysis to make a direct comparison with the results obtained for early time flares and put stringent constraints on the temporal evolution of the flare properties. The aim is to identify possible differences in the mechanism producing the observed early and late time flare emission, if any.  This requires the complete understanding of the observational biases affecting the detection of X-ray flares superimposed on a fading continuum at $t > 1000$ s.  In Sect. \\ref{sample} we present the late time flare sample and we describe the fitting procedure. In Sect. \\ref{analysis} we show the main results of the present analysis of the late time flares temporal profiles and we compare them with those obtained from the C10 sample. In Sect. \\ref{discussion} we discuss our findings. In Sect. \\ref{conclusions} we summarize the results obtained and draw our main conclusions. We adopt standard values of the cosmological parameters: $H_\\circ=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.27$ and $\\Omega_{\\Lambda}=0.73$.  Errors are given at $1\\, \\sigma$ confidence level unless otherwise stated. ",
        "conclusions": "We analysed $23$ GRB X-ray light curves observed by XRT that exhibit late time (i.e. $t_{pk}  \\gtrsim 1000$ s) flares. The present work inherits the C10 strategy of data analysis but extends the temporal window of investigation of three orders of magnitude, from $10^3$ s to $10^6$ s, with a larger sample than in previous works on late time flares \\citep{2008A&A...487..533C}. We found that late time flares are similar to early time flares for the following reasons:  \\begin{itemize} \\item the width of late time flares increases with time similarly to the early time flares. This is a real property of flares, at least up to $10^4$ s; \\item the \\emph{self-similar} behaviour of the flare profile, being the decay proportional to the rise time, is also a property shared by both early time and late time flares. This is a more stringent requirement than being simply \"asymmetrical pulses\". \\end{itemize} However, they differ for the following characteristics: \\begin{itemize} \\item the median energy output of late time flares is one order of magnitude lower than the early time flare one, being $\\sim 1\\% E_{prompt}$; \\item the bias introduced by the shallower underlying continuum after $1000$ s allows the detection of only the brightest flares. However it seems unlikely that the flares observed at $t\\gtrsim 10^5$ s are compatible with the extrapolation of the behaviour $t^{-2.7}$ found for early time flares; \\item the decoupling of the evolution of the peak luminosity of late time flares from the underlying continuum can account for the fainter distribution of the peak flux to continuum ratio $\\Delta F/F$ of late time flares; \\item a sizable fraction of late time flares are compatible with variability in the GRB afterglow. \\end{itemize}  The global pictures that emerges is that the morphology of the flare light curve is the same for both early and late time flares, while they differ energetically. The similarities of late time flares with the early time ones and with the prompt emission pulses as well as their energetic can be fairly well explained in the framework of the accretion models. The existence of flares in this context is ascribed to instabilities either in the disk  \\citep{2006ApJ...636L..29P} or in the fall-back material \\citep{2008MNRAS.388.1729K,2007MNRAS.376L..48R}. This fragmentation can in principle account for the observed energetic and the longer duration of the accretion episodes at later times \\citep{2006ApJ...636L..29P}. Intermittence can be achieved also in presence of magnetic fields \\citep{2003ApJ...592..767P,2006MNRAS.370L..61P}. Different explanations are provided by the magnetic reconnection model \\citep{2006MNRAS.369L...5L,2006A&A...455L...5G,2009ApJ...695L..10L} that do not require a long lasting activity of the central engine. However, the presence of $\\sim 86\\%$ of the late time flares of our sample that can be ascribed to variability in the GRB afterglow emission leads us to conclude that there could be no unique explanation about the nature of late time flares."
    },
    "1004/1004.2410_arXiv.txt": {
        "abstract": "Cosmic Microwave Background satellite missions as the on-going Planck experiment are expected to provide the strongest constraints on a wide set of cosmological parameters. Those constraints, however, could be weakened when the assumption of a cosmological constant as the dark energy component is removed. Here we show that it will indeed be the case when there exists a coupling among the dark energy and the dark matter fluids. In particular, the expected errors on key parameters as the cold dark matter density and the angular diameter distance at decoupling are significantly larger when a dark coupling is introduced. We show that it will be the case also for future satellite missions as EPIC, unless CMB lensing extraction is performed. ",
        "introduction": "Current cosmological measurements point to a \\emph{flat} universe whose mass-energy includes $5\\%$ ordinary matter and $22\\%$ non-baryonic dark matter, but is dominated by the \\emph{dark energy} component, identified as the engine for the accelerated expansion~\\cite{wmap5,Wood-VaseySN07,TegmarkLRGDR4,PercivalLRG,reid,percival,wmap7}. The most economical description of the cosmological measurements attributes the dark energy to a Cosmological Constant (CC) in Einstein's equations, representing an invariable vacuum energy density. The equation of state of the dark energy component $w$ in the CC case is constant and $w=-1$. However, from the quantum field approach, the bare prediction for the current vacuum energy density is $\\sim 120$ orders of magnitude larger than the observed value. This situation is the so-called CC problem. In addition, there is no proposal which explains naturally why the matter and the vacuum energy densities give similar contributions to the universe's energy budget at this moment in the cosmic history. This is the so-called \\emph{why now?} problem, and a possible way to alleviate it is to assume a time varying, dynamical fluid. The quintessence option consists on a cosmic scalar field $\\phi$ (called \\emph{quintessence} itself) which changes with time and varies across space, and it is slowly approaching its ground state. Also, the quintessence equation of state is generally not constant through cosmic time~\\cite{RatraPeebles88a,RatraPeebles88b,Wetterich95,Caldwell98, Quint99,Wang00}. In principle, the quintessence field may couple to the other fields. In practice, observations strongly constrain the couplings to ordinary matter~\\cite{carroll}. However, interactions within the dark sectors, i.e. between dark matter and dark energy, are still allowed. This could change significantly the universe and the density perturbation evolution, the latter being seeds for structure formation. For models equivalent to the one studied here, see {\\it e.g.} Refs.~\\cite{amendola,Valiviita:2008iv,He:2008si,Jackson:2009mz,Gavela:2009cy,CalderaCabral:2009ja,Valiviita:2009nu,Majerotto:2009np}. \\\\  In this paper we investigate how allowing for a feasible interacting dark matter and dark energy model will affect the cosmological constraints expected from future CMB experiments. The Planck satellite mission, for example, is expected to provide high quality constraints on several key parameters (see e.g. \\cite{Perotto:2006rj}). However, those forecasts are usually performed under  the assumption that the dark energy component is either a cosmological constant or a fluid with constant, redshift independent equation of state $w=P/\\rho$. It is therefore timely to investigate if the assumption of a more elaborate dark energy component with a coupling with the dark matter could have an impact on these constraints. Here we indeed focus on the future CMB data constraints on interacting dark matter-dark energy models, exploiting, in particular, the gravitational CMB lensing signal. The structure of the paper is as follows. Section \\ref{sec:seci} presents the background and the linear perturbations of the interacting dark matter-dark energy model explored here. Sections \\ref{sec:secii} and \\ref{sec:seciii} describe the CMB lensing extraction method and the future CMB data simulation used in our numerical analysis, respectively. We present our results in Sec.~\\ref{sec:seciiii} and draw our conclusions in Sec.~\\ref{sec:seciiiii}. ",
        "conclusions": "\\label{sec:seciiiii} The current accelerated expansion of the universe is driven by the so-called dark energy. This negative pressure component could be interpreted as the vacuum energy density, or as a cosmic, dynamical scalar field.  If a cosmic quintessence field is present in nature, it may couple to the other fields in nature. While the couplings of the quintessence field to ordinary matter are severely constrained, an energy exchange among the dark matter and dark energy sectors is allowed by current observations.  The major goals of the on-going Planck and the future EPIC experiments are to determine the nature of the dark energy component and to measure the remaining cosmological parameters with unprecedented precision. Several studies in the literature have been devoted to explore the performance of Planck and EPIC experiments in the dark energy scenario, see for instance Ref.~\\cite{Perotto:2006rj}. In this paper we have explored the performance of Planck and EPIC experiments in alternative dark energy cosmologies, more concretely, in a universe with a coupling $\\xi$ among the dark energy and dark mater components~\\cite{Gavela:2009cy}.  We have generated mock data for the Planck and EPIC experiments. CMB gravitational lensing extraction has also been included in the analysis. The lensing noise has been computed by means of the minimum variance estimator method of Ref.~\\cite{lensextr}. The mock data have then been analyzed using MCMC techniques to compute the errors on the several cosmological parameters considered here. We find that relevant degeneracies are present among the coupling $\\xi$ and some other cosmological parameters, as the cold dark matter density $\\Omega_{c}h^2$. Therefore, in the presence of a coupling, the expected Planck or EPIC errors on quantities as the cold dark matter energy density or the angular diameter distance at decoupling $\\theta_s$ are one order of magnitude larger than in standard cosmologies with $\\xi=0$.  When gravitational CMB lensing extraction is included in the analysis, Planck results remain unchanged, due to the high level of lensing noise for this experiment. However, the EPIC mission, which will benefit from a much lower lensing noise level, can (a) provide tighter constraints on the cosmological parameters, even in the presence of a coupling, and (b) distinguish among coupled and uncoupled models that would look identical if they were fitted to Planck (lensed or unlensed) data."
    },
    "1004/1004.0309_arXiv.txt": {
        "abstract": "It is known that strong electric fields produce electron and positron pairs from the vacuum, and due to the back-reaction these pairs oscillate back and forth coherently with the alternating electric fields in time. We study this phenomenon in spatially inhomogeneous and bound electric fields by integrating the equations of energy-momentum and particle-number conservations and Maxwell equations. The space and time evolutions of the pair-induced electric field, electric charge- and current-densities are calculated. The results show that non-vanishing electric charge-density and the propagation of pair-induced electric fields, differently from the case of homogeneous and unbound electric fields. The space and time variations of pair-induced electric charges and currents emit an electromagnetic radiation. We obtain the narrow spectrum and intensity of this radiation, whose peak $\\omega_{\\rm peak}$ locates in the region around $4$ KeV for electric field strength $\\sim E_c$. We discuss their relevances to both the laboratory experiments for electron and positron pair-productions and the astrophysical observations of compact stars with an electromagnetic structure. \\comment{ Based on the Maxwell equations, energy-momentum and particle-number conservations, we study the back-reaction of electron-positron pairs created in spatially inhomogeneous electric fields.  Numerically integrating a set of partial differential equations, we find, in inhomogeneous case, there would be net charges at local slices instead of everywhere is neutral in spatially uniform electric field. Then the involved physical variables (electric field, number density, energy density and momentum density of positrons and electrons etc.)can evolve to more complicated sub structure in space. The most interesting thing is that the electric field and current not only synchronously oscillate in time, but also there is a wave propagating along the inverse direction of external electric field gradient in space. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.5376_arXiv.txt": {
        "abstract": "We use {\\it Chandra} and {\\it XMM-Newton} to study the hot gas content in a sample of field early-type galaxies.  We find that the L$_{\\rm X}$--L$_{\\rm K}$ relationship is steeper for field galaxies than for comparable galaxies in groups and clusters.  The low hot gas content of field galaxies with L$_{\\rm K} \\lessapprox L_{\\star}$ suggests that internal processes such as supernovae driven winds or AGN feedback expel hot gas from low mass galaxies. Such mechanisms may be less effective in groups and clusters where the presence of an intragroup or intracluster medium can confine outflowing material.  In addition, galaxies in groups and clusters may be able to accrete gas from the ambient medium.  While there is a population of L$_{\\rm K} \\lessapprox L_{\\star}$ galaxies in groups and clusters that retain hot gas halos, some galaxies in these rich environments, including brighter galaxies, are largely devoid of hot gas. In these cases, the hot gas halos have likely been removed via ram pressure stripping.  This suggests a very complex interplay between the intragroup/intracluster medium and hot gas halos of galaxies in rich environments with the ambient medium helping to confine or even enhance the halos in some cases and acting to remove gas in others.  In contrast, the hot gas content of more isolated galaxies is largely a function of the mass of the galaxy, with more massive galaxies able to maintain their halos, while in lower mass systems the hot gas escapes in outflowing winds. ",
        "introduction": "It has been known since the mid-1980's that hot gas halos are common in early-type galaxies \\citep{forman85}.  From the earliest studies with {\\it Einstein}, it was apparent that the X-ray luminosity of early-type galaxies is correlated with the stellar luminosity \\citep{forman85,tf85,c87}.  The correlation between these two quantities suggests that the origin of the hot gas must be related to the stellar content of the galaxy. It is generally believed that the hot gas originates from stellar mass lost from evolved stars and planetary nebula \\citep{m90,mb03,BP09}. However, the scatter in the L$_{\\rm X}$-L$_{\\rm B}$ relationship is very large. Since the B band can be strongly affected by both recent star formation and dust, it may not be a good measure of the true stellar content of a galaxy in some cases.  The large scatter in the relationship remains, however, when L$_{\\rm K}$ is used as a proxy for stellar light \\citep{ellis06}, suggesting the scatter is dominated by variations in the X-ray properties of galaxies. The large range in X-ray luminosity for a given stellar luminosity could be due to either intrinsic differences in galaxy properties or environmental effects.  There has been considerable effort by the astronomical community to understand the scatter in the L$_{\\rm X}$-L$_{\\rm B}$ and L$_{\\rm X}$-L$_{\\rm K}$ relationships with different authors often reaching opposing conclusions. For example, some authors have found evidence for the X-ray luminosities of early-type galaxies to vary with environment \\citep{ws91,brown00}, while others find no such trend \\citep{o01,h01,ellis06}.  One problem that plagued these earlier studies was the inability to cleanly separate out the thermal emission in galaxies from other contributions to the X-ray emission. In particular, the contribution of X-ray binaries and an active galactic nucleus (AGN) can be substantial in some galaxies. The broader bandpasses and superior spatial resolution of {\\it Chandra} and {\\it XMM-Newton} allow a much cleaner measurement of the thermal gas component than was possible with earlier telescopes \\citep{kf03,diehl07}.   The superb spatial resolution of {\\it Chandra} is also important because it allows one to separate out an individual galaxy's hot gas halo from the more extended intragroup or cluster medium. Recent studies of early-type galaxies in groups and clusters indicate that a large fraction of such galaxies retain their hot gas halos even in these dense environments \\citep{sun07,jbm08,sun09}.  The presence of hot gas halos in group and cluster galaxies has important consequences for galaxy evolution.  In nearly all models of galaxy formation, the condensation of hot halo gas is a primary driver for the build up of massive galaxies \\citep{wf91,cole00,bower06}.  Since the pioneering work of \\citet{wf91}, it has generally been assumed that a galaxy's hot gas halo is stripped completely when a galaxy enters a group or cluster. With the hot gas halo removed, there is no new source for gas and the star formation rate quickly declines and the galaxy becomes red.  The {\\it Chandra} observations of hot gas halos in groups and clusters demonstrate that the assumption of complete stripping is overly simple. Some authors have started to incorporate more sophisticated stripping prescriptions in to the semi-analytic models and a more realistic treatment of the hot gas appears to alleviate some problems that were present in the earlier versions of the models \\citep{kang08,mccarthy08,font08}.  To quantify the importance of stripping in groups and clusters, the properties of ellipticals in these rich environments must be compared to the properties of galaxies in environments where stripping is not important, i.e. the field.  \\citet{jbm08} attempted to make this comparison and found that group and cluster galaxies appear to be X-ray-faint compared to field galaxies. However, this result was based on {\\it ROSAT} observations of field galaxies, so the contribution of X-ray binaries and AGN had to be estimated.  In this {\\it Letter}, we use {\\it Chandra} and {\\it XMM-Newton} observations of field early-type galaxies to cleanly measure their hot gas content for the first time.  We then compare the X-ray properties of these field galaxies to similar galaxies in groups and clusters to study how environment impacts hot halos in galaxies.  We adopt H$_{0}$=70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega$$_{\\rm M}$ = 0.3 and $\\Omega$$_{\\Lambda}$= 0.7 throughout this {\\it Letter}. ",
        "conclusions": "In the previous section we show that the L$_{\\rm X}-L_{\\rm K}$ relation is steeper for field galaxies than for galaxies in groups and clusters. This result has important implications for the role of environment on the hot gas halos of galaxies. In particular, our study suggests there is a population of early-type galaxies with L$_{\\rm K}$ $\\lessapprox$ L$_{\\star}$ in groups and clusters that retain substantial hot gas reservoirs, while their counterparts in the field are mostly devoid of gas.  The lack of a substantial hot gas component in field galaxies with L$_{\\rm K}$ $\\lessapprox$ L$_{\\star}$ could reflect a fundamental difference in the global properties of field and group/cluster early-type galaxies.  For example, previous studies of early-type galaxies have suggested a possible trend between the X-ray luminosity and the age of a galaxy estimated from dynamical or spectroscopic indicators \\citep{o01}, with luminous X-ray emission apparently restricted to galaxies with ages greater than a few Gyrs.  These observations are consistent with a scenario where hot gas is initially removed during major mergers and the hot gas halos take several gigayears to build up \\citep{cox06}.  The difference between the field galaxies and those in richer environments could therefore reflect a difference in age.  In fact, there is some indication from simulations that isolated ellipticals should be on average younger than their counterparts in groups and clusters \\citep{sami10}.  However, \\citet{reda04} have studied the stellar populations in several of the field galaxies in our sample and found that the bulk of the stars in these galaxies are very old. In fact, for the galaxies they studied, \\citet{reda04} found that the formation epoch of field and cluster ellipticals appears similar.  This suggests that age differences are unlikely to be the explanation for the observed differences in the hot gas content.  Another possibility is that field galaxies lack substantial dark matter halos and are therefore unable to keep X-ray halos \\citep{o04}. Minimal dark matter halos have been implied for some elliptical galaxies from kinematic studies of planetary nebula at large radii \\citep{r03,d07}.  However, the low velocity dispersions derived from the planetary nebula could be due to halo stars on radial orbits and not small dark matter halos \\citep{dekel05}.  To explain the observed X-ray/optical relationships, dark matter halos of low L$_{\\rm K}$ ellipticals in groups and clusters would need to be more substantial than their counterparts in the field.  However, one might naively expect the dark matter halos of ellipticals in groups and clusters to be reduced relative to those in the field since tidal stripping is much more likely in these denser environments.  Thus, the differences in the scaling relations of field and group/cluster galaxies most likely do not reflect differences in the dark matter halos.  In addition to potential intrinsic differences between field and group/cluster early-types, environmental processes could be important. However, environmental processes that remove gas from galaxies (such as ram pressure stripping) are likely only important in richer environments, where there is a substantial intragroup or intracluster medium \\citep{km08,bekki09}. This suggests that internal processes must be responsible for removing hot gas from low mass field galaxies. Most likely gas has been expelled from these galaxies by stellar winds or AGN feedback. Detailed studies of low X-ray luminosity ellipticals have concluded that winds sustained by Type 1a supernovae are likely the dominant mechanism by which galaxies lose their hot gas, although AGN outbursts may also be important in some cases \\citep{david06,pel07,trinch08}.  In our field sample, there is little evidence for significant AGN activity. Approximately half of our sample galaxies are detected in radio continuum in the NVSS \\citep{condon98}, but in nearly all cases the emission is very weak.  However, the weak radio emission in our field sample does not necessarily mean that AGN feedback is not important in these systems, since there is little correlation between 1.4 GHz radio luminosity and disturbed X-ray morphologies (i.e. cavities) in many nearby ellipticals \\citep{dong10}.  While supernovae or AGN driven outflows can explain the low hot gas content of field galaxies, such mechanisms may be less effective in groups and clusters where the presence of an ambient medium may stifle such winds \\citep{babul92,babul99,brown98,brown00}.  In addition, early-type galaxies in groups and clusters may be able to accrete gas from the intragroup or intracluster medium \\citep{bm98,bm99,brown00}. The combination of these two effects likely accounts for the population of low L$_{\\rm K}$ group and cluster galaxies that still contain significant amounts of hot gas.  There are several observations that could help test the relative importance of gas accretion from the ambient medium versus the suffocation of outflowing winds. If accretion of gas is the dominant mechanism by which low L$_{\\rm K}$ cluster galaxies maintain halos, we might expect the metallicity of the gas to be lower than if the gas is produced internally in the galaxies.  Deeper X-ray observations of low L$_{\\rm K}$ group and cluster galaxies with halos should allow this test to be performed.  We note that the L$_{\\rm X}-L_{\\rm K}$ relationships shown in Figure 1 were derived for galaxies with a detected hot gas halo.  Given that it is more difficult to detect individual hot halos in groups and clusters (because of the higher \\lq\\lq background\\rq\\rq \\ from the intragroup/intracluster medium), we likely could not detect the very low L$_{\\rm X}$ halos in these richer environments that we detect in the field.  In fact, there are many galaxies in groups and clusters where a thermal component has not been detected \\citep[see][]{sun07,jbm08} and in some cases the limits on the hot gas luminosity would place these galaxies well-below the group/cluster L$_{\\rm X}-L_{\\rm K}$ relationship shown in Figure 1.  This suggests that although there is a population of low L$_{\\rm K}$ galaxies in groups and clusters that retain hot gas halos, there are other galaxies, including brighter galaxies, in these rich environments that have likely lost their hot gas halos to ram pressure stripping.  Our study therefore suggests a very complex interplay between the intragroup/intracluster medium and the hot gas halos of galaxies in rich environments: the presence of an ambient medium can act to maintain or even enhance a hot halo in some galaxies and remove halo gas in other cases. In contrast, the hot gas content of more isolated galaxies is largely a function of the mass of the galaxy, with more massive galaxies able to maintain their halos, while the hot gas is expelled in lower mass systems. To better understand the importance of the various environmental processes at play in groups and clusters, studies of how the properties of hot gas halos vary spatially in groups and clusters would be valuable."
    },
    "1004/1004.3285_arXiv.txt": {
        "abstract": "We explore the reconstruction of the gravitational lensing field of the cosmic microwave background in real space showing that very little statistical information is lost when estimators of short range on the celestial sphere are used in place of the customary estimators in harmonic space, which are nonlocal and in principle require a simultaneous analysis of the entire sky without any cuts or excisions. Because virtually all the information relevant to lensing reconstruction lies on angular scales close to the resolution scale of the sky map, the gravitational lensing dilatation and shear fields (which unlike the deflection field or lensing potential are directly related to the observations in a local manner) may be reconstructed by means of quadratic combinations involving only very closely separated pixels. Even though harmonic space provides a more natural context for understanding lensing reconstruction theoretically, the real space methods developed here have the virtue of being faster to implement and are likely to prove useful for analyzing realistic maps containing a galactic cut and possibly numerous small excisions to exclude point sources that cannot be reliably subtracted. ",
        "introduction": "Recently there has been much discussion of how the intervening mass from clustered matter between the last scattering surface (at $z\\approx 1100$) and an observer at the present time distorts the CMB anisotropy by means of gravitational lensing \\cite{blanchard,cole,tomita,bernardeau,van-waerbeke}. On the level of the two-point correlation function, this effect distorts the TT (temperature-temperature) correlation power spectrum \\cite{seljak} and mixes the EE and BB polarization power spectra as well as distorting them \\cite{zald,benabed-bis,seljak-ter}. Lensing also introduces non-Gaussianities that manifest themselves in the higher-point correlation functions \\cite{zaldarriaga,cooray-ter}. At the level of the three-point correlation function, to leading order there is no nonzero expectation value if we regard the lensing potential as a random field \\cite{cooray,kesden}. However if we consider the CMB lensing potential as fixed, we find that expectation values of the form \\ba \\left< T (\\boldsymbol{\\theta}) T(\\boldsymbol{\\theta}') \\right> _{\\Phi (\\boldsymbol{\\theta}^{\\prime \\prime })} \\ea do not vanish, and this property may be exploited to recover or ``reconstruct'' the lensing field using estimators quadratic in $T$ (or in $E$ and $B$) \\cite{hu2001a}.  Much effort has been devoted to developing an optimal reconstruction of the lensing potential in harmonic space, which implicitly assumes full sky coverage with no galactic cut, no bad pixels due to point sources that must be excised, and no nonuniform weighting to account for uneven sky coverage \\cite{challinor,lesgourgues,bode}. For a reconstruction based on the temperature anisotropy alone, it has been shown how to construct the optimal quadratic estimator in this idealized context \\cite{hu2001b}, and the improvement that can be gained from using an even more optimal maximum likelihood estimator is marginal \\cite{seljak-bis}, because the distortion due to lensing is small compared to the intrinsic cosmic variance and noise of the experiment (although this assumption is less valid at very large $\\ell $ for very clean maps where the lensing signal is dominant). For the exploitation of polarized anisotropies, the situation is somewhat more complicated. When the experimental noise is large compared to the B polarization mode generated by lensing, the situation is essentially the same as for a reconstruction using the temperature data \\cite{hu-okamoto,cooray-bis}. However at a higher sensitivity where the B signal is essentially entirely due to lensing, the quadratic estimator underperforms because the actual multipole moments rather than their averages should be used for the optimal weighting. In this case, the higher order corrections to the quadratic estimator present in the maximum likelihood estimator are no longer negligible \\cite{hirata-seljak}.  In this paper we investigate a real space approach to lensing reconstruction. Under the ideal conditions often assumed and described above, this approach would naturally yield the same result as the conventional approach in harmonic space. Our interest however lies in considering slightly non-optimal estimators that have been modified to have a finite range so that cuts, excisions of pixels, and non-uniform coverage may be included in a simple and flexible way. We believe that such non-ideal but more robust local estimators defined in real space may prove superior for confronting the complications inherent in analyzing real data \\cite{smith,miller,perotto}. Another advantage of the approach here is that the dilatation and pure shear provide separate and essentially independent lensing reconstructions which may be confronted with each other. This feature may prove useful as a way of diagnosing spurious signals, which are unlikely to affect the two reconstructions in the same way. Moreover the presence of two shear components enables one to estimate the noise of the reconstruction through the implied transverse displacement field, which is forbidden in weak lensing.  Before proceeding to the details of this program, it is useful to consider the relation between the various descriptions of the lensing and the deformation of the anisotropies in real space. It is also useful to consider which angular scales contribute the most statistical weight to the lensing reconstruction.  The lensing distortion of the CMB anisotropy on the surface of last scatter may be described in three ways: by a lensing potential $\\Phi ,$ by a deflection field $\\boldsymbol{\\xi }=\\nabla \\Phi ,$ or by the three components of the shear tensor \\ba \\kappa = \\begin{pmatrix} \\kappa _0+\\kappa _{+}& \\kappa _{\\times }\\cr \\kappa _{\\times }& \\kappa _0-\\kappa _{+}\\cr \\end{pmatrix} =\\nabla _a\\nabla _b\\Phi . \\ea Even if we have simultaneous access to the entire sky, the descriptions $\\Phi $ and $\\boldsymbol{\\xi }$ suffer from an ambiguity. $\\Phi $ cannot be distinguished from $\\Phi +\\textrm{(constant)}$ and the vector field $\\boldsymbol{\\xi }$ can be measured only up to a constant translation (or more properly a rotation in the presence of sky curvature). This is because if we know only the CMB power spectrum, a patch of sky and its translation necessarily have the same likelihood on account of isotropy of the underlying stochastic process. Consequently, the absolute translation due to lensing cannot be observed. By contrast, locally the shear and dilatation (which are gradients of the translation vector field) are completely well defined. This can easily be seen by considering the effect of a constant deformation described by a deformation matrix $S$ relating the angular coordinates $\\boldsymbol{\\theta },$ the actual coordinates on the celestial sphere of a point on the last scattering surface, to the coordinates $\\boldsymbol{\\theta }'$, the coordinates that the same point would have in the absence of lensing. \\footnote{In the sequel we shall, unless otherwise indicated, employ the flat sky approximation where the vector ${\\thetab }$ represents a point on the flattened celestial sphere and ${\\ellb }$ represents a wavevector. At times summations over $(l,m)$ shall also be used.} We have $\\boldsymbol{\\theta '}= S\\boldsymbol{\\theta }$ where $S=\\exp [-\\boldsymbol{\\kappa }].$ (Note that we employ the flat sky approximation and assume that the deformation is small so that a linear treatment is adequate.) To linear order, the power spectrum is modified in the following way by this linear deformation, which preserves the homogeneity but not the isotropy of the underlying statistical process: \\ba C(\\vert {\\ellb }\\vert ) \\to C(\\ell ) \\biggl[ 1 + \\kappa _0 \\left( \\frac{d(\\ln[C(\\ell )])}{d(\\ln [\\ell ])} + 2\\right) +\\left( \\frac{\\kappa _+(\\ell _1^2-\\ell _2^2)+\\kappa _\\times  (2\\ell _1\\ell _2)}{\\ell ^2} \\right) \\frac{d(\\ln [C(\\ell )])}{d(\\ln [\\ell ])} \\biggr] . \\ea For the case of perfect scale invariance (i.e., a power law of the form $C(\\ell )\\propto \\ell ^{-2}$) there is no change in the correlations due to the dilatation component of $S,$ and similarly for a perfect white noise spectrum (i.e., a power law of the form $C(\\ell )\\propto \\ell ^0$) there is no sensitivity to the pure shear components $\\kappa _+$ and $\\kappa _\\times $ in the anisotropic $(m=\\pm 2)$ correlations.  To the extent that the shear-dilatation components are slowly varying, we may construct estimators of $\\kappa _0,$ $\\kappa _+,$ and $\\kappa _\\times $ as follows \\ba \\hat \\kappa _0 &=& {N_0}^{-1} \\int _A d^2\\theta ~ \\int _A d^2\\theta '~ \\left[ T(\\thetab )~T(\\thetab ')- \\left< T({\\thetab })~T({\\thetab '})\\right> _{\\kappa =0} \\right] \\cr &&\\times \\int \\frac{d^2\\ell }{(2\\pi )^2} \\exp [i{\\ellb }\\cdot ({\\thetab }-{\\thetab '})]~ \\frac{C(\\ell )}{[C(\\ell )+N(\\ell )]^2} \\left( \\frac{d(\\ln [C(\\ell )])}{d(\\ln [\\ell ])} + 2\\right)\\cr &=& {N_0}^{-1}{\\cal A}^{-1} \\int _A d^2\\theta ~ \\int _A d^2\\theta '~ \\left[ T({\\thetab })~T({\\thetab '})- \\left< T({\\thetab })~T({\\thetab '})\\right> _{\\kappa =0} \\right] K_0({\\thetab }-{\\thetab '})\\cr &=& \\frac{1}{\\cal A} \\int _A d^2\\theta ~ \\left[ T({\\thetab })~({\\cal F}_{\\kappa _0}T)({\\thetab }) -\\textrm{(constant)} \\right] \\ea and similarly \\ba \\begin{pmatrix} \\hat \\kappa _{+}\\cr \\hat \\kappa _{\\times }\\cr \\end{pmatrix} &=& {N_{+,\\times }}^{-1} \\int _A d^2\\theta ~ \\int _A d^2\\theta '~T(\\thetab )~T({\\thetab '})~\\cr &&\\times \\int \\frac{d^2\\ell }{(2\\pi )^2} \\exp [i{\\ellb }\\cdot ({\\thetab }-{\\thetab '})]~ \\frac{C(\\ell )}{[C(\\ell )+N(\\ell )]^2} \\frac{d(\\ln [C(\\ell )])}{d(\\ln [\\ell ])} \\begin{pmatrix} \\cos (2\\vartheta )\\cr \\sin (2\\vartheta )\\cr \\end{pmatrix}\\cr &=& {N_{+,\\times }}^{-1}{\\cal A}^{-1} \\int _A d^2\\theta ~ \\int _A d^2\\theta '~T({\\thetab })~T({\\thetab '})~ \\begin{pmatrix} K_+({\\thetab }-{\\thetab '})\\cr K_\\times ({\\thetab }-{\\thetab '})\\cr \\end{pmatrix}, \\ea where the normalization factors are given by \\ba N_0&=&{\\cal A} \\int \\frac{d^2\\ell }{(2\\pi )^2} \\frac{[C(\\ell )]^2}{[C(\\ell )+N(\\ell )]^2} \\left( \\frac{d(\\ln [C(\\ell )])}{d(\\ln [\\ell ])} + 2\\right)^2,\\cr N_+&=&N_\\times =\\frac{\\cal A}{2} \\int \\frac{d^2\\ell }{(2\\pi )^2} \\frac{[C(\\ell )]^2}{[C(\\ell )+N(\\ell )]^2} \\left( \\frac{d(\\ln [C(\\ell )])}{d(\\ln [\\ell ])}\\right)^2 \\ea and ${\\cal A}$ is the area of the domain. Here $N(\\ell )$ is the noise of the experiment being considered and serves as a cut-off at large $\\ell ,$ above which there is very little exploitable information because of the low signal-to-noise of the CMB maps.  The above expressions assume a spatially flat, two-dimensional domain formally of large but finite area subject to a spatially uniform linear deformation. We may consider a toroidal domain in the limit that the periods of the torus become arbitrarily large (though the toroidal domain is not essential for the application of the real space estimator presented in this paper). Before proceeding it is useful to consider the relation of this simplified problem to the real problem of reconstructing the lensing field starting from a CMB map of finite resolution. Because of the smallness of the lensing distortion of the CMB anisotropy, the idealized situation considered above is less different from the actual situation, where one is dealing with a curved sky and a varying dilatation and shear fields, than one might at first sight suppose.  \\begin{figure} \\begin{center} \\includegraphics[width=17cm]{lens_potl_derivs_v2.pdf} \\end{center} \\caption{\\baselineskip=0.5cm{ {\\bf Lensing power spectrum.} The lensing field power spectrum is shown represented in several manners. The three panels (from left to right) illustrate the lensing field expressed as potential, a deflection field, and dilatation/shear field, respectively. Plotted are $[\\ell (\\ell +1)C_{XX}/(2\\pi )]^{1/2}$ where $XX=\\Phi \\Phi ,$ ${\\boldsymbol{\\xi \\xi }},$ $\\kappa \\kappa .$ Here $C_{\\ell }^{\\boldsymbol{\\xi \\xi }}=\\ell (\\ell +1) ~C_{\\ell }^{\\Phi \\Phi },$ and $C_{\\ell }^{\\kappa \\kappa }=\\ell ^2(\\ell +1)^2 ~C_{\\ell }^{\\Phi \\Phi }/4.$ Of the three, the last power spectrum is more directly related to the observed distortion. }} \\label{Fig:LensingFields} \\end{figure}  The rightmost panel of Fig.~\\ref{Fig:LensingFields} shows the shear-dilatation power spectrum as a function of multipole number $\\ell ,$ and we observe that for $\\ell <100,$ the distortion is always less than about $1.5\\% .$ This implies that in order to attain an $(S/N)$ of approximately unity it is necessary to consider a region containing at least $O(10^3)$ resolution elements, where a resolution element is a pixel of the map of sufficient size so that the noise and angular resolution of the survey give $S/N\\approx 1.$ Consequently, there is little point to trying to reconstruct the lensing field over a region not having at least 30 resolution elements on a side. If the distortion from lensing were greater, the situation would be different.  For the ideal linear estimator \\ba \\frac{1}{\\sigma ^2_{\\hat \\kappa _{0,ideal}}}= \\left( \\frac{S}{N} \\right) ^2 &=& \\sum _{\\ell , m} \\frac{C_\\ell ^2} {2(C_\\ell +N_\\ell )^2} \\left[ \\frac{d(\\ln [C])}{d(\\ln [\\ell ])}+2 \\right] ^2\\cr &\\approx & \\frac{{\\cal A}}{(4\\pi )} \\int _0^\\infty d\\ell ~\\ell ~ \\frac{C(\\ell )^2} {(C(\\ell )+N(\\ell ))^2} \\left[ \\frac{d(\\ln [C(\\ell )])}{d(\\ln [\\ell ])}+2 \\right] ^2 \\label{SoverN} \\ea where ${\\cal A}$ is the area of the sky patch considered.  We now consider other unbiased estimators of $\\kappa _0$ and $\\kappa_+$ where a slight increase in the variance is compensated for by other desirable properties. We are presently interested in estimators for which the real space filtering kernel is short range. To this end it is useful to define the inner product on the space of weight vectors ${w}=\\{ w_\\ell \\}$: \\ba \\left< w^A, w^B \\right> &=& \\sum _{\\ell , m} \\frac{1}{2(C_\\ell +N_\\ell )^2} w^A_{\\ell ,m} w^B_{\\ell ,m}\\cr &=& {\\cal A}\\int \\frac{d^2\\ell }{(2\\pi )^2} \\frac{1}{2[C(\\ell )+N(\\ell )]^2} w^A(\\ellb ) w^B(\\ellb) \\ea where we give both the spherical and flat sky continuum forms. If we set $w_{ideal}(\\ell )=\\delta C(\\ell )_{th, \\kappa _0=1}= C(\\ell)[\\frac{d \\ln C_\\ell}{d \\ln \\ell}+2],$ then in terms of the above inner product \\ba \\hat \\kappa _{0,ideal}= \\frac{ \\left< w_{ideal}, \\delta C_{obs}\\right> }{ \\left< w_{ideal},  w_{ideal} \\right> } \\ea and \\ba \\left( \\frac{S}{N}\\right) ^2_{\\kappa _0, ideal}= \\left<\\delta C_{th, \\kappa _0=1}, \\delta C_{th, \\kappa _0=1}\\right> . \\ea Given an arbitrary weight vector $w,$ using the above inner product we may define the following unbiased estimator of $\\kappa _0$ \\ba \\hat \\kappa _0(w)= \\frac{ \\left< w, \\delta C_{obs}\\right> }{ \\left< w,  w_{ideal} \\right> } \\ea provided that the denominator does not vanish, and its variance is given by $\\left< w,  w\\right> /\\left< w_{ideal},  w\\right> ^2,$ so that the increase in variance with respect to the ideal estimator is given by following geometric expression for the secant squared \\ba \\frac{ \\textrm{Var}(\\hat \\kappa _0(w)) }{ \\textrm{Var}(\\hat \\kappa _0(w_{ideal})) } = \\frac{ \\left< w,  w\\right> \\left< w_{ideal},  w_{ideal} \\right> }{ \\left< w,  w_{ideal} \\right> ^2 }~\\rm{sec}^2(\\chi). \\label{VVar} \\ea For $\\kappa _+$ and $\\kappa _\\times $ analogous formulae may be derived straightforwardly.  \\begin{figure} \\begin{center} \\includegraphics[width=7.5cm]{lensed_CMB.pdf} \\includegraphics[width=7.5cm]{lensing_dimensionless_spectral_index.pdf}\\\\ \\includegraphics[width=7.5cm]{lensing_cumulative_information.pdf} \\includegraphics[width=7.5cm]{kappa_plus_cumulative_information_B.pdf} \\end{center} \\caption{\\baselineskip=0.5cm{ {\\bf Character of the signal.} The top row shows as a function of multipole number $\\ell $ the temperature power spectrum and local spectral index, defined as $d\\ln C_\\ell /d\\ln \\ell,$ for the standard cosmology (WMAP best-fit model). On the bottom row, the left panel shows the normalized cumulative $\\chi ^2$ as a function of $\\ell $ integrated both from the left and from the right using the sensitivity and resolution parameters for the PLANCK experiment (where the 100, 143 and 217 GHz channels have been combined in quadrature) \\cite{bluebook}. We observe that the central 80\\% of the information is concentrated in the range $\\ell =800$--$1600.$ Smaller $\\ell $ contribute almost no information because there are comparatively very few independent multipoles, and moreover the angular spectrum in the sky is very nearly scale invariant. At much larger $\\ell $ instrument noise and beam smearing wash out the usable signal. In the intermediate range a structure of plateaus connected by steep rises can be observed. This structure is a direct result of the Doppler oscillations. Around the crests and troughs the spectrum is almost scale invariant and hence does not contain any information for determining the dilatation. The right panel shows the corresponding plot for the shear, where the plateaus are less pronounced. }} \\label{Fig:Info} \\end{figure}  We consider how the information (or $S^2/N^2$) contained in the ideal estimator is distributed over the various multipoles. In Fig.~\\ref{Fig:Info} [panels (c) and (d)] we plot the quantity \\ba F_<(\\bar \\ell )= \\frac{ \\int _0^{\\bar \\ell }\\ell ~d\\ell ~ \\frac{ {C(\\ell) }^2}{({C(\\ell)}+{N(\\ell )})^2} \\left[ \\frac{d(\\ln [C(\\ell )])}{d(\\ln [\\ell ])}+2 \\right] ^2 }{ \\int _0^{\\infty }\\ell ~d\\ell ~ \\frac{ {C(\\ell )}^2}{({C(\\ell )}+{N(\\ell )})^2} \\left[ \\frac{d(\\ln [C(\\ell )])}{d(\\ln [\\ell ])}+2 \\right] ^2 } \\ea and $F_>(\\bar \\ell )=1-F_<(\\bar \\ell )$ where \\ba N_\\ell =N_0 ~\\exp \\left[ +\\ell ^2\\theta ^2_{beam}\\right] =N_0 ~\\exp \\left[ +\\ell ^2/\\ell ^2_{beam}\\right] \\ea and $\\ell _{beam}=(810)(10'/\\theta _{beam}^{fwhm}).$ The corresponding quantity is also shown for the shear.  In their present state, the ideal minimum variance kernels for the estimators $\\hat \\kappa _0,$ $\\hat \\kappa _+,$ and $\\hat \\kappa _\\times $ have their support sharply peaked at small separations, but nevertheless there is still some small support for large separation. We now investigate how much information is lost if the support at large separation is completely cut away. The quantitative way to characterize this loss is to ask by what factor the variance of the estimator is increased relative to the optimal estimator after our pruned estimator has been renormalized to render it unbiased.  \\def\\half{\\frac{1}{2}}  In real space the minimum variance full-sky estimator kernel for $\\kappa _0$ is given by \\ba K_{ideal}(\\theta ) &=& \\int _0^\\infty \\ell ~d\\ell ~ K_{ideal}(\\ell )\\cr &=& \\int _0^\\infty \\ell ~d\\ell ~ {J_0(\\ell \\theta )} \\frac{1}{N_{ideal}} \\frac{C(\\ell )} {[C(\\ell) +N(\\ell )]^2}~ \\left[ \\frac{d(\\ln [C(\\ell )])}{d[\\ln (\\ell )]} +2 \\right] . \\ea This kernel may be inverted using the following inverse Bessel transform: \\ba K_{ideal}(\\ell )=\\frac{1}{2\\pi }\\int _0^\\infty \\theta ~d\\theta ~ J_0(\\ell \\theta )~ K_{ideal}(\\theta ). \\ea We limit the support of the kernel by requiring that $K(\\theta )$ be nonzero only for $\\theta \\le \\theta _{max}$ where $\\theta _{max}$ is varied. This is accomplished numerically by expressing $J_0(\\ell \\theta )$ as a linear combination of cubic spline basis functions spanning the interval $\\theta \\in [0, \\theta _{max}]$ and optimizing for the shape that minimizes the variance calculated according to eqn. (\\ref{VVar}). Analogous expressions may be obtained for the shear by replacing $J_0$ with $J_2.$ Fig. \\ref{Fig:TruncEst} shows the variance ratio as a function of $\\theta _{max}$ for the dilatation and shear estimators. We observe that the increase in variance at small separations is more modest for the shear estimator.  \\begin{figure} \\begin{center} \\includegraphics[width=5.3cm]{idl_lambda_vs_theta_max.pdf} \\includegraphics[width=5.3cm]{idl_filter_theta.pdf} \\includegraphics[width=5.3cm]{idl_filter_ell.pdf}\\\\ \\vspace{0.5cm} \\includegraphics[width=5.3cm]{idl_lambda_vs_theta_max_shear.pdf} \\includegraphics[width=5.3cm]{idl_filter_theta_shear.pdf} \\includegraphics[width=5.3cm]{idl_filter_ell_shear.pdf} \\end{center} \\caption{\\baselineskip=0.5cm{ {\\bf Performance of estimator with truncated angular support.} We indicate how limiting the angular support of the filter in our estimator increases the noise. The panels on the top row refer to the dilatation filter, while the panels on the bottom row refer to the shear filter. Panel (a) indicates how the estimator variance (with the estimator normalized to be unbiased) increases as the angular support (disk radius in degrees) is reduced. Panels (b) and (c) indicate the profiles of the optimal truncated estimators in both angular space and harmonic space. }} \\label{Fig:TruncEst} \\end{figure} ",
        "conclusions": "We have demonstrated how to reconstruct in real space using a filter of compact support the weak gravitational lensing field, here represented as three fields, a dilatation field $\\kappa _0(\\thetab )$ and the two components of the pure shear distortion field $\\kappa _+(\\thetab )$ and $\\kappa _\\times (\\thetab ).$ The three fields are related by a set of nonlocal consistency conditions, which may subsequently be exploited to reduce the noise of the reconstruction. Except for an integration constant and two translational and one rotational zero modes, the weak lensing may alternatively and equivalently be described by either (1) a gravitational lensing potential $\\Phi (\\thetab ),$ (2) a displacement field ${\\boldsymbol {\\xi}}(\\thetab ),$ $\\Phi (\\thetab ),$ (2) a displacement field ${\\boldsymbol{\\xi}(\\thetab )},$ or (3) the dilatation field $\\kappa _0(\\thetab )$ and the two components of the pure shear distortion field $\\kappa _+(\\thetab )$ and $\\kappa _\\times (\\thetab ).$ In this paper we argue that for the purpose of reconstruction the representation (3) is advantageous because this is the representation for which the lensing field bears a local relation to the real space CMB maps. This locality comes at a price because the three components are not independent and subject to nonlocal consistency conditions, which may be exploited to improve the reconstruction. Locality allows different regions of the sky to be analyzed independently in a natural way, quite unlike the quadratic reconstruction in harmonic space, where the entire sky must be analyzed simultaneously. This approach and variations thereof hold promise for dealing with partial sky coverage and excised point sources. For the filters developed in this paper there is very little loss of information for the lensing field at low wavenumbers. However at larger wavenumbers the lensing signal is attenuated according to a wavenumber dependent form factor, which can be deconvolved by applying a correction filter.  \\vspace{0.6cm} \\noindent \\textbf{Acknowledgements:} MB and KM acknowledge support from a joint CNRS/NRF travel grant. The work of MB, CSC and MR was supported in part by the Projet Blanc VIMS-PLANCK of the Agence Nationale de la Recherche. KM and CSC are supported by the South African National Research Foundation."
    },
    "1004/1004.3766_arXiv.txt": {
        "abstract": "We carry out global three-dimensional radiation hydrodynamical simulations of self-gravitating accretion discs to determine if, and under what conditions, a disc may fragment to form giant planets.  We explore the parameter space (in terms of the disc opacity, temperature and size) and include the effect of stellar irradiation.  We find that the disc opacity plays a vital role in determining whether a disc fragments.  Specifically, opacities that are smaller than interstellar Rosseland mean values promote fragmentation (even at small radii, $R < 25$AU) since low opacities allow a disc to cool quickly.  This may occur if a disc has a low metallicity or if grain growth has occurred.  With specific reference to the HR~8799 planetary system, given its star is metal-poor, our results suggest that the formation of its imaged planetary system could potentially have occurred by gravitational instability.  We also find that the presence of stellar irradiation generally acts to inhibit fragmentation (since the discs can only cool to the temperature defined by stellar irradiation).  However, fragmentation may occur if the irradiation is sufficiently weak that it allows the disc to attain a low Toomre stability parameter. ",
        "introduction": "\\label{sec:intro}  There are two ways in which giant planets have been hypothesised to form: core accretion \\citep{Safronov_CA,Goldreich_Ward_CA,Pollack_etal_CA} and gravitational instability (\\citealt{GI_Cameron,Boss_GI}; review by \\citealt{GI_review_Durisen_PPV}).  The former has been favoured but historically has had difficulties for two reasons: the first is a temporal issue where the timescales required to form planets may be too large such that the gas in the disc is depleted before the gas giant planet is fully formed.  Secondly, simulations of this method model the growth of planets typically starting with kilometre-sized planetesimals, but while the growth of particles from small grains to metre-sized objects appears to be straight-forward, how to get from metre-sized objects to kilometre-sized planetesimals is unknown.  Gravitational instability, on the other hand, eliminates the timescale problem, forming gas giant planets in $\\lesssim$~$\\rm O(10^4)$ years.  Such planets may not have solid cores, which may well be the case for Jupiter \\citep{Saumon_Guillot2004}, though the capture of solid material to form a core after the fragmenting stage in a gravitationally unstable disc has also been proposed \\citep*{Solid_capture_postGI}.  However, since gravitational instability is not thought to operate close to the central star, it has not been thought to be the dominant mechanism by which giant planets form as it was unable to describe the observations of close-in giant planets.  Recent observations of planets at large distances \\citep{Fomalhaut,HR8799} encourage us to revisit this theory.  \\cite{Nero_Bjorkman_GI_analysis} have argued analytically that Fomalhaut b and at least the outer planet of the HR~8799 system could have formed by gravitational instability as the cooling timescales are likely to be small enough such that fragmentation is possible.  Moreover, discs in their early stages are thought to be massive \\citep{Eisner_Carpenter_massive_discs} suggesting that gravitational instability must play a role in the evolution of a disc in the late Class I and early Class II stages.  It has also been proposed that core accretion may be a method by which planets may form at small radii ($\\sim \\rm{O}(10) \\rm{AU}$) whilst gravitational instability may be the mechanism by which planets may form at larger radii ($\\gtrsim \\rm{O}(100) \\rm{AU}$) \\citep[e.g.][]{Boley_CA_and_GI} though a hybrid scenario of forming gas giants in the same system by both core accretion and gravitational instability has yet to be modelled.  There are two quantities that have historically been used to determine whether a disc is likely to fragment.  The first is the stability parameter \\citep{Toomre_stability1964},  \\begin{equation} \\label{eq:Toomre} Q=\\frac{c_{\\rm s}\\kappa_{\\rm ep}}{\\pi\\Sigma G}, \\end{equation} where $c_{\\rm s}$ is the sound speed in the disc, $\\kappa_{\\rm ep}$ is the epicyclic frequency, which for Keplerian discs is approximately equal to the angular frequency, $\\Omega$, $\\Sigma$ is the surface mass density and $G$ is the gravitational constant.  Therefore, once the surface mass density and the rotation of the disc have been established, the stability is purely dependent on the disc temperature.  \\cite{Toomre_stability1964} showed that for an infinitesimally thin disc to fragment, the stability parameter must be less than a critical value, $Q_{\\rm crit} \\approx 1$.  \\cite{Gammie_betacool} showed that in addition to the stability criterion above, the disc must cool at a fast enough rate.  Using shearing sheet simulations, he showed that if the cooling timescale can be parametrized as  \\begin{equation} \\label{eq:beta} \\beta = t_{\\rm cool}\\Omega, \\end{equation} where  \\begin{equation} \\label{eq:tcool} t_{\\rm cool} = u \\big(\\frac{du_{\\rm cool}}{dt}\\big)^{-1}, \\end{equation} $u$ is the internal energy and $du_{\\rm cool}/dt$ is the total cooling rate, then for fragmentation we require $\\beta \\lesssim 3$, for a ratio of specific heats $\\gamma = 2$ (in two dimensions).  \\cite*{Rice_beta_condition} carried out three-dimensional simulations using a smoothed particle hydrodynamics (SPH) code and showed that this cooling parameter is dependent on the equation of state so that fragmentation can occur for discs with $\\gamma = 5/3$ and $7/5$ if $\\beta \\lesssim 7$ and $13$, respectively.  \\cite*{Libby_MSci} showed that the critical value of $\\beta$ (below which fragmentation will occur if the stability criterion is met) may depend on the disc's thermal history: if the timescale on which the disc's cooling timescale is decreased is slower than the cooling timescale itself (i.e. a gradual decrease in $\\beta$) then the critical value may decrease by up to a factor of 2.  More recently, \\cite*{Cossins_opacity_beta} showed that the critical value varies with the temperature dependence of the cooling law.  In general though, the critical value is thought to be of the order of the dynamical timescale.  The above fragmentation criteria are based on the assumption that the dominant form of dissipation in the disc is due to internal heating processes.  Previous simulations without external irradiation have considered isolated discs with simple cooling prescriptions \\citep[e.g.][]{Lodato_Rice_original} and with radiative transfer \\citep[e.g.][]{Boss_RT,Cai_etal_RT,Mayer_etal_RT}.  \\cite{Johnson_Gammie} suggested that discs with external irradiation are likely to be effectively isothermal and can therefore be treated as such.  \\cite{Matzner_Levin2005} analytically considered externally irradiated discs and concluded that stellar irradiation quenches fragmentation.  \\cite{Cai_envelope_irradiation} carried out simulations with external irradiation and found that their discs are more resistant to fragmentation and proposed that these results may be extended to discs with stellar irradiation.   \\cite{Stamatellos_no_frag_inside_40AU} also carried out simulations taking into account the effects of stellar irradiation and found this to be a stabilising factor.  \\cite{Rafikov_SI} analytically explored fragmentation in gravitoturbulent discs including the effects of stellar irradiation and suggested that fragmentation can only occur beyond $\\approx 120\\rm{AU}$.  \\cite{Dodson-Robinson_HR8799} carried out a linear stability analysis on irradiated discs to show that gravitational instability takes place for systems with a large disc to star mass ratio.  However, whilst fragmentation in gravitationally unstable discs may be less likely than previously thought, we still do not know in \\emph{what situations} discs may fragment when modelling them realistically with radiative transfer and by considering the effects of stellar irradiation.  It is therefore important to deduce when fragmentation may occur when simulating discs with more detailed energetic conditions, and just how realistic or unrealistic fragmentation is in real discs.   \\cite{Boss_metallicity} carried out simulations of gravitationally unstable discs and varied the opacity from $0.1\\times$ to $10\\times$ the Rosseland mean opacities and found that the fragmentation results were insensitive to the dust grain opacity.  However, given that a reduced opacity is more likely to allow energy to stream out of a disc more easily causing it to cool and promote fragmentation, whilst in a high opacity disc the converse is true, it is interesting to consider what opacity values allow and do not allow fragmentation.  Given that a disc's opacity gives somewhat an indication of how metal-rich it is or how large or small the grain sizes are, we may then make preliminary conclusions on the disc conditions that are likely to promote fragmentation, which is a key focus of this paper.   In this paper, we model the evolution of massive self-gravitating discs using a global three-dimensional SPH code including radiative transfer and the effects of stellar irradiation.  In particular, we explore the parameter space in terms of dust opacity, disc temperature and size in order to scope out if, and under what conditions, a self-gravitating disc may fragment.  In Section~\\ref{sec:numerical_setup}, we describe the code used to carry out our simulations.  In Section~\\ref{sec:sim}, we outline our simulations, including the disc setup and discuss the parameter space.  In Section~\\ref{sec:results}, we present our results, while we compare with previous studies and make conclusions in Sections~\\ref{sec:disc} and~\\ref{sec:conc}, respectively. ",
        "conclusions": "\\label{sec:conc} We have carried out radiation hydrodynamical simulations to investigate the evolution of massive self-gravitating discs ($M_{\\rm disc}/M_{\\star} = 0.1$).  We consider discs with opacities ranging from $0.01\\times$ to $10\\times$ the interstellar Rosseland mean values.  We also consider the effects of changing the initial and boundary temperatures of the discs as well as simulating different disc sizes (with outer disc radii, $R_{\\rm out} = 25$ and 300 AU).  We find that the disc opacity is very important in determining whether a disc is likely to fragment.  In particular, we find that fragmentation is promoted in discs with opacity values lower than interstellar Rosseland mean values since this allows radiation to leave the disc quickly.  Low opacities may exist in low metallicity discs or discs with larger grain sizes.  This is a particularly important and timely result given the recent discoveries of wide orbit planets \\citep{Fomalhaut,HR8799} and the future emphasis for surveys of planets on such wide orbits.  We show that it is possible for fragmentation to occur in gravitationally unstable discs even at radii where the innermost planet of the HR~8799 system is located ($R \\gtrsim 24$~AU).  Furthermore, HR~8799 is known to be a metal-poor, $\\lambda$ Bootis star with metallicity $[M/H] = -0.47$ \\citep{HR8799_metallicity} so it is reasonable to assume that its disc was similarly metal-poor.  We have shown that such a scenario favours fragmentation and therefore, our results indicate that all three planets of the HR~8799 system may well have formed via gravitational instability.  Though a hybrid core accretion and gravitational instability scenario for planet formation may also be a possibility for this system, our calculations show that such a hybrid scenario may not be necessary.  We find that the presence of a \\emph{thermal blanket} as a result of the stellar irradiation inhibits fragmentation since the discs are only able to cool to the boundary temperature.  However, we also show that under certain circumstances, fragmentation may occur.  Our results demonstrate that for fragmentation, weak irradiation is required such that the boundary temperature and hence Toomre stability parameter is low, in addition to low enough opacities (even in large, cool discs) since this allows more efficient cooling so that the disc's temperature does not increase significantly (due to internal dissipation) above the boundary temperature."
    },
    "1004/1004.3882_arXiv.txt": {
        "abstract": "The transition density $n_t$ and pressure $P_t$ at the inner edge between the liquid core and the solid crust of a neutron star are analyzed using the thermodynamical method and the framework of relativistic nuclear energy density functionals. Starting from a functional that has been carefully adjusted to experimental binding energies of finite nuclei, and varying the density dependence of the corresponding symmetry energy within the limits determined by isovector properties of finite nuclei, we estimate the constraints on the core-crust transition density and pressure of neutron stars: $0.086 \\ {\\rm fm}^{-3} \\leq n_t < 0.090 \\ {\\rm fm}^{-3}$ and $0.3\\ {\\rm MeV \\ fm}^{-3} < P_t \\leq 0.76 \\ {\\rm MeV \\ fm}^{-3}$.    \\vspace{0.3cm}  PACS number(s): 21.30.Fe, 21.60.Jz, 26.60.Gj, 26.60.Kp  \\\\  Keywords: Nuclear density functional, Equation of state, Neutron star crust. ",
        "introduction": "Neutron stars are extraordinary astronomical laboratories for the physics of dense neutron-rich nuclear matter \\cite{Shapiro-83,Haensel-07}. They consists of several distinct layers: the atmosphere, the surface, the crust and the core. The latter, divided into the outer core and inner core, has a radius of approximately 10 km and contains most of the star's mass. The crust, of $\\approx 1$ km thickness and containing only a few percent of the total mass, can also be divided into the outer crust and inner crust.  Although less exotic and smaller in size than the core, the crust is nevertheless crucial for the understanding of the physics of neutron stars. It represents the interface between the observable surface phenomena and the invisible core. The structure of the crust can be related to some peculiar phenomena, such as pulsar glitches, thermal relaxation after matter accretion, quasi periodic oscillations and anisotropic surface cooling \\cite{Lattimer-00,Lattimer-01,Link-99}. A very important ingredient in the study of the structure and various properties of neutron stars is the equation of state (EOS) of neutron-rich nuclear matter \\cite{Lattimer-07}.   One of the most important prediction of a given EOS is the location of the inner edge of a neutron star crust. The inner crust comprises the region from the density at which neutrons drip-out of nuclei, to the inner edge separating the solid crust from the homogeneous liquid core. At the inner edge, in fact, a phase transition occurs from the high-density homogeneous matter to the inhomogeneous matter at lower densities. In the transitional region nuclear matter exhibits instability against clusterization into a two-phase system: neutron-rich  nuclei immersed in dripped neutrons (and sometimes protons). As nuclei are arranged in a lattice, they form solid state crust covering the star's core, which is considered to be a homogeneous liquid \\cite{Kubis-07}. The uniform matter is nearly pure neutron matter, with  a proton fraction of a few percent, determined by the condition of $\\beta$-equilibrium. The transition density takes its critical value $n_c$ when the uniform neutron-proton-electron matter (npe) becomes unstable with respect to the separation into two coexisting phases (one corresponding to nuclei, the other to a nucleonic sea) \\cite{Lattimer-07}.  While the density at which neutrons drip-out of nuclei is rather well determined, the transition density $n_t$ at the inner edge is much less certain due to our insufficient knowledge of the EOS of neutron-rich nuclear matter. The value of $n_t$ determines the structure of the inner part of the crust. If sufficiently high, it is possible for non-spherical phases, with rod- or plate-like nuclei, to occur before the nuclei dissolve. If $n_t$ relatively low, then matter makes a direct transition from spherical nuclei to uniform nucleonic fluid. The extent to which non-spherical phases occur will have important consequences for other properties that are determined by the solid crust \\cite{Pethick-95b}.  In general, the determination of the transition density $n_t$ itself is a very complicated problem because the inner crust may have a very complicated structure. A well established approach is to find the density at which the uniform liquid first becomes unstable against small-amplitude density fluctuations, indicating the formation of nuclear clusters. This approach includes the dynamical method \\cite{Pethick-95b,Baym-71a,Baym-71b,Pethick-95a,Douchin-00, Oyamatsu-07,Ducoin-07, Xu-09-1,Xu-09-2}, the thermodynamical method \\cite{Lattimer-07,Kubis-07,Kubis-04,Worley-08}, and the random phase approximation (RPA) \\cite{Horowitz-01,Carriere-03}.  All theoretical studies have shown that the core-crust transition density and pressure are very sensitive to the density dependence of the nuclear matter symmetry energy. The EOS of neutron-rich nuclear matter has been constrained by using results from heavy-ion reaction studies \\cite{LCK08}. In particular, it has been shown that the $E_{sym}(\\rho)$ constrained in the same sub-saturation density range as the neutron star crust by the isospin diffusion data in heavy-ion collisions at intermediate energies \\cite{Tsa04,Che05a,LiBA05c}, limits the transition density and pressure to $0.040$ fm$^{-3}$ $\\leq \\rho _{t}\\leq 0.065$ fm$^{-3}$ and $0.01$ MeV/fm$^{3}$ $\\leq P_{t}\\leq 0.26$ MeV/fm$^{3}$, respectively . These constrained values appear to be significantly lower than their fiducial values currently used in the literature. In a very recent study \\cite{Xu-09-2}, the core-crust transition density and pressure have been systematically analyzed using the dynamical and thermodynamical methods with a modified Gogny (MDI) and a set of $51$ different Skyrme interactions. Most of these interactions predict values for the transition density and pressure that are considerably higher than the intervals cited above.  In a recent work \\cite{Niksic-08} we have explored a particular class of empirical relativistic nuclear energy density functionals, with parameters adjusted to experimental binding energies of a large set of axially deformed nuclei. Starting from microscopic nucleon self-energies in nuclear matter, and empirical global properties of the nuclear matter equation of state, the coupling parameters of the functional have been determined in a careful comparison of the predicted binding energies with data, for a set of 64 axially deformed nuclei in the mass regions $A\\approx 150-180$ and $A\\approx 230-250$. The isovector channel, in particular, has been carefully adjusted to reproduce available data in medium-heavy and heavy nuclei, including neutron-skin thickness and excitation energies of isovector dipole resonances. It will be interesting, therefore, to apply this class of relativistic density functionals in a systematic investigation of the transition density $n _{t}$ and pressure $P_{t}$ at the inner edge separating the liquid core from the solid crust of neutron stars. In the present study the thermodynamical method will be used.  In recent years, there has been an increased interest in studies of the relationship between the size of the neutron-skin in heavy nuclei, and the symmetry energy at subsaturation densities \\cite{Horowitz-01,Brown-00,Furnstahl-02,Dieperink-03,Steiner-05, Todd-05,Chen-05,Sammarruca-09a, Vidana-09,Centelles-09,Yoshida-04,Klimkiewicz-07}. Studies have also been reported on the correlation between the size of the neutron-skin and properties of a neutron star crust. This was pioneered by Horowitz et. al. \\cite{Horowitz-01}, who used the Random Phase Approximation  based on the Relativistic Mean-Field (RMF) framework for nuclear matter and finite nuclei. An almost linear correlation was established between the predicted core-crust transition density $n_t$ and the size of the neutron-skin. More recently such studies have been carried out by by Xu et. al. \\cite{Xu-09-1,Xu-09-2}, confirming this linear correlation.  The article is organized as follows. In Sec. II we review the thermodynamical method used for locating the inner edge of a neutron star crust. Sec. III contains a brief description of relativistic density functionals that will be used to analyze the constraints on the core-crust transition density and pressure of neutron stars. The results are presented and discussed in Sec. IV, and Sec. V summarizes the present study. ",
        "conclusions": "To adjust the functional DD-PC1, in  Ref.~\\cite{Niksic-08} sets of effective interactions with different values of the volume $a_v$, surface $a_s$, and symmetry energy $a_4$ in nuclear matter were generated, and the corresponding binding energies of deformed nuclei with $A\\approx 150-180$ and $A\\approx 230-250$ were analyzed. The nuclear matter saturation density, the Dirac mass, and the compression modulus, were kept fixed: $n_0 =0.152~\\textnormal{fm}^{-3}$ in accordance with values predicted by most modern relativistic mean-field models, $m^*_D = 0.58 m$ in the narrow interval of values allowed by the empirical energy spacings between spin-orbit partner states in finite nuclei, and $K_{nm} = 230$ MeV to reproduce experimental excitation energies of isoscalar giant monopole resonances in relativistic (Q)RPA calculations.  Nuclear structure data do not constrain the nuclear matter EOS at high nucleon densities. Therefore, two additional points on the $E(\\rho)$ curve in symmetric matter were fixed to the microscopic EoS of Akmal, Pandharipande and Ravenhall \\cite{Akmal-98}, based on the Argonne V$_{18}$ NN potential and the UIX three-nucleon interaction. This EOS has extensively been used in studies of high-density nucleon matter and neutron stars. At almost four times nuclear matter saturation density, the point $n =0.56~\\textnormal{fm}^{-3}$ with $E/A = 34.39$ MeV was chosen and, to have an overall consistency, one point at low density: $n =0.04~\\textnormal{fm}^{-3}$ with $E/A = -6.48$ MeV (cf. Table VI of Ref.~\\cite{Akmal-98}).  The calculated binding energies of finite nuclei are very sensitive to the choice of the nuclear matter volume energy coefficient $a_v$. In fact, one of the important results of analysis of deviations between calculated and experimental masses (mass residuals) of Ref.~\\cite{Niksic-08}, is the pronounced isospin and mass dependence of the residuals on the nuclear matter volume energy at saturation. To reduce the absolute mass residuals to less than 1 MeV, and to contain their mass and isotopic dependence, $a_v$ had to be constrained to a narrow interval of values: $-16.04$ MeV $\\leq a_v \\leq  -16.08$ MeV. Experimental masses do not place very strict constraints on the parameters of the expansion of $E_{sym}(n)$ (cf. Eq.~(\\ref{Esym-expa})), but self-consistent mean-field calculations show that binding energies can restrict the values of $E_{sym}$ at nucleon densities somewhat below saturation density, i.e. at $ n  \\approx 0.1~\\textnormal{fm}^{-3}$. Additional information on the symmetry energy can be obtained from data on neutron-skin thickness and excitation energies of giant dipole resonances. Recent studies have shown that relativistic effective interactions with volume asymmetry $a_4$ in the range $31\\;\\textnormal{MeV} \\le a_4 \\le 35\\; \\textnormal{MeV}$ predict values for neutron-skin thickness that are consistent with data, and reproduce experimental excitation energies of isovector giant dipole resonance~\\cite{VNR.03}. Therefore, in the construction of the functional DD-PC1 in Ref.~\\cite{Niksic-08}, the volume asymmetry was held fixed at $a_4=33$ MeV, and the symmetry energy at a density that corresponds to an average nucleon density in finite nuclei: $\\langle n \\rangle =0.12\\;\\textnormal{fm}^{-3}$ was varied. The quantity $E_{sym}(n=0.12\\;\\textnormal{fm}^{-3})$ will be denoted $\\langle S_2 \\rangle$.  Starting from the relativistic energy density functional DD-PC1, in this work we examine the sensitivity of the core-crust transition density $n_t$ and pressure $P_t$ of neutron stars, on the density dependence of the corresponding symmetry energy of nucleonic matter. In Ref.~\\cite{Niksic-08} the value of $\\langle S_2 \\rangle$ was varied in a rather narrow interval of values 27.6 MeV $\\le \\langle S_2 \\rangle \\le $ 28.6 MeV, constrained by the empirical values of binding energies and ground-state isovector properties of finite nuclei. Fig.~\\ref{Fig_A} displays the corresponding symmetry energy curves $E_{sym}$ as a function of the baryon density $n$. For $a_v=-16.06$ meV (DD-PC1) the minimum $\\chi^2$ deviation of the theoretical binding energies from data is obtained when $\\langle S_2 \\rangle = 27.8$ MeV.  Table \\ref{t:1} and Fig.~\\ref{Fig_B} display the values of the transition density $n_t$ (in fm$^{-3}$) and transition pressure $P_t$ (in Mev$\\cdot$fm$^{-3}$), calculated in the thermodynamical model, as functions of $<S_2>$ for three values of the nuclear matter volume energy coefficient $a_v$. For a given value of the parameter $a_v$, the values of $n_t$ rise with increasing $<S_2>$, whereas the opposite is found for the values of $P_t$. For the considered interval of $<S_2>$, however, the changes are small. An increase of $3.5 \\%$ in $<S_2>$ leads to an increase of $1.5 \\%$ in the value of $n_t$. The transition pressure exhibits a somewhat more pronounced  dependence (the corresponding decrease is around $16-20 \\%$). Both $n_t$ and $P_t$ display a negligible dependence on $a_v$, even though $a_v = -16.02$ MeV and $a_v = -16.14$ MeV lie outside the interval of values for which the absolute deviations between calculated and experimental masses are smaller than 1 MeV.  In Fig.~\\ref{Fig_C} we plot the transition pressure $P_t$ as a function of the transition density $n_t$ for the three sets of nuclear matter EOS and symmetry energy described above, in comparison with results of recent calculations performed using an isospin and momentum-dependent modified Gogny effective interaction (MDI) \\cite{Xu-09-2,Krastev-010}. The different values of the parameter  $x$ in the MDI model correspond to various choices of the density dependence of the nuclear symmetry energy. In Refs.~\\cite{LiBA05c,LiBA05d} it has been shown that only $-1 \\leq x \\leq 0$ leads to a density dependence of the symmetry energy in the sub-saturation density region that is consistent with isospin diffusion data and the empirical value of the neutron-skin thickness in $^{208}$Pb. In addition to the MDI EOS, in Fig.~\\ref{Fig_C} we also show the result obtained by Akmal et al. \\cite{Akmal-98} with the $A18+\\delta v+UIX^*$ interaction (ARP), and the value obtained in the recent Dirac-Brueckner-Hartree-Fock (DBHF) calculation \\cite{Sammarruca-09a} with the Bonn B One-Boson-Exchange (OBE) potential (DBHF+Bonn B) \\cite{Machleidt-89}.  A distinctive feature of the present analysis is the narrow interval of allowed values $(n_t,P_t)$ that results from the rather stringent constraints on the parameter $<S_2>$. The effect of varying the volume energy at saturation $a_v$ is almost negligible. The present results for $n_t$ and $P_t$ lie in  the region constrained by the measure of the current uncertainty in the density dependence of the symmetry energy \\cite{Lattimer-07}, and are found very close to the result of Akmal et al. \\cite{Akmal-98}. We note that all the results shown in Fig.~\\ref{Fig_C} are obtained using the parabolic approximation for the EOS of isospin-asymmetric nuclear matter. The transition density and pressure have also been estimated using the full equation of state and employing both the dynamical and thermodynamical methods \\cite{Xu-09-1,Xu-09-2}.  As explained above, the rather narrow interval of  $<S_2>$, for this type of nuclear energy density functionals, has been constrained by the empirical values of binding energies and ground-state isovector properties of finite nuclei. The symmetry energy at saturation density, $a_4=33$ MeV, was fixed in Ref.~\\cite{Niksic-08} to obtain the best results for the neutron-skin thickness in Sn isotopes and $^{208}$Pb, and for the excitation energies of isovector dipole resonances. However, because of large experimental uncertainties, especially for the neutron-skin thickness, good agreement with data can also be obtained for other values of $a_4$. This is shown in Fig.~\\ref{Fig_D}, where we plot the predictions for the differences between neutron and proton $rms$ radii of Sn and Pb isotopes, in comparison with available data \\cite{Kra.99,SH.94,Kra.94}, for different choices of the symmetry energy at saturation density. The self-consistent mean-field calculations have been performed using the relativistic Hartree-Bogoliubov (RHB) model \\cite{VALR.05}, with pairing correlations described by the pairing part of the finite-range Gogny interaction. The isoscalar channel of the particle-hole interaction corresponds to the DD-PC1 functional, and in the isovector channel $<S_2>$ is kept fixed at 27.8 MeV (DD-PC1), whereas  $a_4$ is varied in the interval between 30 MeV and 35 MeV. The corresponding symmetry energy as a function of the nucleon density is shown in Fig.~\\ref{Fig_E}.  We notice, therefore, that by keeping $<S_2>$ constant and varying $a_4$ in the interval between 30 MeV and 35 MeV,  the density dependence of the symmetry energy can be modified in a controlled way, i.e. the corresponding energy density functionals still reproduce ground-state properties of finite nuclei in fair agreement with data. In Table \\ref{t:2} and Fig.~\\ref{Fig_F} we display the corresponding values of the transition density $n_t$ (in fm$^{-3}$) and transition pressure $P_t$ (in Mev$\\cdot$fm$^{-3}$) as functions of $a_4$ for three values of the nuclear matter volume energy coefficient $a_v$. The transition pressure $P_t$ as a function of the transition density $n_t$ for the three sets of nuclear matter EOS and symmetry energy is plotted in Fig.~\\ref{Fig_G}. Not surprising, considering the symmetry energy curves of Fig.~\\ref{Fig_E}, the constraints on $n_t$  and $P_t$ have been relaxed in this case, and the allowed values span a much larger interval of values compared to the restricted variation of $<S_2>$ shown in Figs.~\\ref{Fig_A} and \\ref{Fig_B}.   To be able to compare the present results for the transition density and transition pressure with recent studies \\cite{Xu-09-2}, in Fig.~\\ref{Fig_H} we plot the calculated values of $n_t$ and $P_t$ as functions of the slope parameter of the symmetry energy (cf. Eq.~(\\ref{L-1})), for the two sets of effective interactions described above. $n_t$ is a monotonously decreasing, and $P_t$ monotonously increasing function of $L$. In the small interval of $L$ values determined by the variation of $<S_2>$ between 27.6 MeV and 28.6 MeV, both $n_t$ and $P_t$ display a linear dependence on $L$. In the much larger interval determined by the variation of $a_4$ from 30 MeV to 35 MeV, a weak parabolic dependence of $n_t$ and $P_t$ is found. We note that transport model studies of the isospin diffusion data in heavy-ion reactions have constrained the slope parameter L to the values $88 \\pm 25$ MeV \\cite{Xu-09-2}. Considering that we can also, most probably, exclude the value $a_4 = 30$ MeV for the asymmetry at saturation density (cf. Figs.~\\ref{Fig_D} and \\ref{Fig_F}), because it implies an unrealistically small value of $< 0.1$ fm for the neutron-skin thickness of  $^{208}$Pb, the present analysis places the following constraints on the core-crust transition density and pressure of neutron stars: $0.086 \\ {\\rm fm}^{-3} \\leq n_t < 0.090 \\ {\\rm fm}^{-3}$ and $0.3\\ {\\rm MeV \\ fm}^{-3} < P_t \\leq 0.76 \\ {\\rm MeV \\ fm}^{-3}$.   Finally, in Fig.~\\ref{Fig_I} we compare the present prediction for the range of values of the transition density $n_t$ with the results of Horowitz and Piekarewicz who, in Ref.~\\cite{Horowitz-01}, also used the framework of relativistic mean-field effective interactions to study the relationship between the neutron-skin thickness of a heavy nucleus and the properties of neutron star crusts. Starting from the NL3 meson-exchange effective interaction \\cite{NL3}, the density dependence of the symmetry energy was varied by adding nonlinear couplings between the isoscalar and isovector mesons to the original interaction. The variation was carried out in such a way to enhance the changes in the neutron density and neutron-skin thickness, while keeping small the corresponding changes in the binding energy and proton density distribution. For the solid crust of a neutron star, the effective RMF interactions were used in a simple RPA calculation of the transition density below which uniform neutron-rich matter becomes unstable against small amplitude density fluctuations. The resulting transition densities are plotted in  Fig.~\\ref{Fig_I} as a function of the predicted difference between neutron and proton $rms$ radii in $^{208}$Pb. This inverse correlation was parameterized \\cite{Horowitz-01} \\begin{equation} n_t\\approx 0.16-0.39(R_n-R_p), \\label{Hor-1} \\end{equation} with the skin thickness expressed in fm. In the present analysis, using a different type of relativistic effective interactions and varying the density dependence of the symmetry energy by explicitly modifying $<S_2>$ or $a_4$, we find a much weaker dependence $n_t$ on the neutron-skin thickness of $^{208}$Pb."
    },
    "1004/1004.1908_arXiv.txt": {
        "abstract": "In order to determine (as best as possible) the nature of the dark matter particle (mass and decoupling temperature) we compare the theoretical evolution of density fluctuations computed from first principles since the end of inflation till today to the observed properties of galaxies as the effective core density, core radius and surface density. We match the theoretically computed surface density to its observed value in order to obtain: (i) the decreasing of the phase-space density since equilibration till today (ii) the mass of the dark matter particle and the decoupling temperature $ T_d $ (iii) the type of density profile (core or cusp). The dark matter particle mass turns to be between 1 and 2 keV and the decoupling temperature $ T_d $ turns to be above 100 GeV. keV dark matter particles necessarily produce cored density profiles while wimps ($ m \\sim 100 $ GeV, $ T_d \\sim 5 $ GeV) inevitably produce cusped profiles at scales about 0.003 pc. We compute in addition the halo radius $ r_0  $, the halo central density $ \\rho_{0} $ and the halo velocity $ \\sqrt{{\\overline {v^2}}_{halo}} $: they all reproduce the observed values within one order of magnitude. These results are independent of the particle model and vary very little with the statistics of the dark matter particle. Moreover, our results for the surface density using typical CDM wimps agree with the values obtained from CDM simulations (which are five orders of magnitude larger than the surface density observed values). DM particles with mass in the keV scale well reproduce the observed  properties of galaxies as the effective core density, cored profiles and the value of the surface density. The linear framework presented here turns to apply both to keV and to GeV mass scale DM particles. keV scale particles reproduce all observed galaxy magnitudes within one order of magnitude while GeV mass particles disagree with observations in up to eleven orders of magnitude. ",
        "introduction": "Since several years and more recently \\citep{disvdb,disvdb2,disvdb3,disvdb4} it has been  stressed that basic galaxy parameters as mass, size, baryon-fraction, central density, are not independent from each other but in fact all of them do depend on one parameter that works as a galaxy identifier. In fact there exist functional relations that constrain the different galaxy parameters in such a way that the galaxy structure depends essentially on one parameter (\\citep{sal07} and references therein).  These functional relations may play for galaxies the r\\^ole that the equations of state play in thermodynamical systems.  First, let us remind that the density of DM halos around galaxies is usually well reproduced by means of dark halos with a cored distribution \\citet{deb,sfm}, where $ r_0 $ is the core radius, $ \\rho_0 $  is the central density $ {\\displaystyle \\lim_{r \\to 0}} \\; \\rho(r)  = \\rho_0 $ and $ \\rho(r) $ for $ r<r_0 $ is approximately constant. Recent findings highlight the quantity $ \\mu_0 \\equiv  r_0  \\; \\rho_0  $ proportional to the halo central surface density defined as $$ 2 \\int_0^{\\infty}  \\rho(0,0,x_3) \\; dx_3 \\quad {\\rm where} \\quad {\\vec r}=(x_1,x_2,x_3) \\; $$ where $ x_3 $ goes along the line of sight. The quantity $ \\mu_0 $ is found nearly {\\bf constant} and independent  of luminosity in  different galactic systems (spirals, dwarf irregular and spheroidals, elliptics) spanning over $14$ magnitudes in luminosity and  over different Hubble types. More precisely, all galaxies seem to have the same value for $ \\mu_0 $, namely $ \\mu_0 \\simeq 120 \\; M_\\odot /{\\rm pc}^2 $ \\citep{kor,dona,span}.   It is remarkable that at the same time other important structural quantities as $ r_0 , \\; \\rho_0 $, the baryon-fraction and the galaxy mass vary orders of magnitude from one galaxy to another.  The constancy of $ \\mu_0 $ is unlikely to be a coincidence and probably has a deep physical meaning in the dynamics of galaxy formation. It must be stressed that $ \\mu_0 $ is the only dimensionful quantity which is constant among galaxies.  By analogy with the theory of phase transitions in statistical physics we find useful to call 'universal' those quantities which take the same value for a large set of galaxies while non-universal quantities vary orders of magnitude from one galaxy to another. In this context the quantities called universal take the same value up to $ \\pm 20 \\% $ for different galaxies. That is, they are only approximately 'universal'. Other known universal quantity in the above sense is the shape of the density profile when expressed as a function of $ r/r_0 $ and normalized to unit at $ r = 0 $.  \\medskip  We use here the linear approximation to the Boltzmann-Vlasov equation \\citep{dod,kt} to compute the DM density profile and from it the surface density $ \\mu_0 $.    In this paper, from first principles and using the primordial power spectrum we follow the evolution of the gravitational collapse of a perturbation of mass $ M \\sim 3 \\times 10^{12} M_{\\odot} $ and derive  the resulting linear halo density profile. This  reproduces  the phase of fast accretion found in $N$-body  simulations.  As a result we obtain  robust predictions for the properties of DM halos from first principles. In the case of  $\\Lambda$CDM our results agree with the  $N$-body $\\Lambda$CDM simulations and in the case of $\\Lambda$WDM our results agree with the observations.     We obtain a very good fit of the linear profile to the Burkert profile which determines the relation between $ r_0 $ and the free-streaming length. This shows that our linear profile can appropriately fit halo density data. The linear approximation to the Boltzmann-Vlasov equation improves for large halo radius $ r_0 $ (see fig. \\ref{rlin}). Since universal quantities take the same value for large or small galaxies one can expect good results from the linear approximation which improves for large and less compact galaxies.  \\medskip  We also compute non-universal galaxy quantities as the halo radius, galaxy mass, halo central density and squared halo velocity. We find that the linear approximation provides halo central densities smaller than or in the range of the observations, and halo velocities larger than the observed ones by a factor between 1 and 10. Notice that the determination of these  non-universal galaxy quantities is unrelated to the computation of the DM particle mass. Hence, our linear calculations amply suffice for our specific goals to determine properties of the DM particles.   \\medskip  Clearly, for non-universal quantities one has to include the nonlinear effects (including for instance mergers) to exactly reproduce the observations. Notice that general arguments based on the Boltzmann-Vlasov equation show that cores cannot become cusps through mixing and mergers and that cusps do not become steeper neither shallower through  mixing and mergers \\citep{wd}.    Moreover, recent $N$-body $\\Lambda$CDM simulations (Acquarius) have found that the DM halos form in a sort of \"monolithical\" way. Their inner regions,  that contain the visible galaxies, are found to be stable since early times and  contrary to previous believes, major mergers (i.e. those with progenitor mass ratios greater than 1:10)  are found to contribute little to their total mass growth \\cite{wang}. This indicates that nonlinearities (i.e. mergers) have a reduced importance. Minor mergers, secondary infall, rare major mergers are certainly important for details, but the gross features of DM halos are determined during  the fast-accretion phase of their gravitational collapse, as  the history of the quasar-galaxies coevolution also seems to indicate \\cite{dane}.  The halo formation essentially consists of two main phases: A first fast accretion phase (that can be treated by the linear approximation), and a second subsequent slow accretion phase with mergers and infalls, that have a random character and that can only be described by numerical simulations. This second phase does not have an essential influence in the shape of the halo profile. Thus, in order to explain the observed halo profiles one just needs to describe the first phase of halo formation, as we do here in this paper.    \\medskip  Therefore, the cored and cusped character we find for the linear profiles depending on the DM particle mass considered (keV and GeV mass scale, respectively) should remain valid after mixing and mergers are taken into account.    Clearly, to build the theory of galaxy formation simulations are needed using the correct linear primordial power spectrum as input. Our work is only indicative in that respect.    \\medskip  We combine in this paper observed properties of galaxies as the effective core density and the core radius with the theoretical linear evolution of density fluctuations computed from first principles since the end of inflation till today. We compute the galaxy surface density and match it with the observed values. Within the same scheme, we derive analytically the halo radius $ r_0 $ and the factor $ Z $ characterizing the reduction of the phase-space density since equilibration till today. In this way, we find the value of the mass of the dark matter particle which turns to be between 1 and 2 keV, and the number of ultrarelativistic degrees of freedom of the dark matter coupling at decoupling $ g_d $, or similarly, the decoupling temperature $ T_d $ which turns to be above 100 GeV. Evidence based on the phase space density pointing towards a DM particle mass in the keV scale was presented in refs. \\cite{dm1,dm2}.  We consider in this paper the whole range of galaxy virial masses going from 5 to 300 $ \\times 10^{11} \\; M_\\odot $.  Non-universal quantities are reproduced in the linear approximation within a factor of ten.   The analytic expressions we derived for the density profile, and the mass of the dark matter particle also imply that keV dark matter particles always produce cored density profiles while heavy dark matter particles as wimps ($m = 100$ GeV, $ T_d $ = 5 GeV) inevitably produce cusped profiles at scales of 0.03 pc. These results are independent of the particle model, vary very little with the statistics of the dark matter particle and are all obtained in the linear framework.  \\medskip   Analytic methods have been used since some time to derive galaxy properties using the primordial power of the density fluctuations [see for example \\cite{hs,pjep}] and using the spherical model  \\citet{EBa,EBb}.  \\medskip  The theoretical treatment presented here captures many essential features of dark matter, allowing to determine its nature. Our treatment also applies to CDM: if we use the CDM surface density value obtained from CDM simulations \\cite{yh}, we determine [sec. \\ref{conclu}] a dark matter particle mass in the wimps mass scale (GeV), fully consistent with CDM simulations.  \\medskip  This paper is organized as follows: Sec. 2 presents galaxy data and empirical formulas relating basic galaxy parameters; sec. 3 deals with the phase-space density; sec. 4 contains our theoretical results for the density profile from the linearized Boltzmann-Vlasov equation. In sec. 5 we derive the DM particle mass and the decoupling temperature from the theoretical and observed galaxy surface density, in sec. 6 we compute non-universal galaxy properties and in sec. 7 we derive the profiles for keV scale DM particles and for wimps (cored vs. cusped profiles). In sec. 8 we present our conclusions. ",
        "conclusions": "Dark matter is characterized by two basic quantities: the DM particle mass $ m $ and the number of ultrarelativistic degrees of freedom at decoupling $ g_d $ (or, alternatively  the decoupling temperature $ T_d $). By computing the evolution of the density fluctuations since the end of inflation till today we obtain the density profiles and theoretical relations between $ m $ and $ g_d $ involving the observable densities $ \\rho_{DM} $ and  $ \\mu_0 $ eqs.(\\ref{Qprim}), (\\ref{gdu}) and (\\ref{muteo}). Inserting the observed values of $ \\rho_{DM} $ and  $ \\mu_0 $ in these theoretical relations yields $ m, \\;  g_d $ and $ Q_p $ eqs. (\\ref{dm})-(\\ref{mcero}) and (\\ref{gdu2}), respectively.  We derive the values of  the halo radius $ r_0 $ and the values of the factor $ Z $ characterizing the reduction of the phase-space density since the equilibration time till today. For these results we use the observed values of the halo phase-space density $ Q_{halo} $.  \\medskip  From the observed values of the surface density we present here clear evidence that the mass of the DM particle is about one or two keV. Evidence based on the phase space density pointing towards a DM particle mass in the keV scale was presented in refs. \\cite{dm1,dm2}.  In addition, one can wonder what would be the results for heavy wimps. For example, for wimps at $ m_{wimp} = 100 $ GeV the characteristic scale $ r_{lin} $ eq.(\\ref{rlin1}) takes the value given by eq.(\\ref{rlinW}). For such {\\bf small} $ r_{lin} $ the linear profile $ \\rho_{lin}(r)_{wimp} $ appears as a {\\bf cusped} profile when observed at scales from $ 0.003$ pc $ 1 $ pc as shown in fig. \\ref{perT}. Cusped profiles are thus clearly associated to heavy DM particles with a huge mass $ m_{wimp} $ well above the physical keV scale while cored profiles are associated to DM particles with mass in the keV scale.  \\medskip  Notice that the linear density profile turns out to be cored or cuspy depending on the DM particle mass $ m $. For $ m \\sim $ keV the resulting linear density profile is cored as depicted in fig. \\ref{plgg} while for  $ m \\gtrsim $ GeV the linear density profile turns to be cusped as shown in fig. \\ref{perT}. Figs. \\ref{plgg}-\\ref{perT} show that the linear density profiles for a 1-2 keV DM particle are similar to Burkert (within a factor 2-3, irrelevant for the aims of this paper) while for a wimp DM particle, the linear density profile is similar to a NFW profile.  \\medskip   Interestingly enough, it is possible to derive the value of the surface density $ \\mu_0 $ from CDM simulations. Values of the product $ r_s \\; \\rho_s $ from NFW fits to CDM simulations for galaxies were reported in \\cite{yh}. From these values of $ r_s \\; \\rho_s $ we can derive the surface density $ \\mu_0 $, since $ \\mu_0 = \\rho_0 \\; r_0 \\simeq 25 \\; r_s \\; \\rho_s $ with the result \\be\\label{mucdm} \\mu_0^{CDM} \\simeq 10^7 \\;  M_\\odot /{\\rm pc}^2 \\; . \\ee [Notice that $ \\rho_s $ in \\cite{yh} differs by a factor four from eq.(\\ref{nfw})].  We see that the surface density from CDM simulations is {\\bf five orders} of magnitude {\\bf larger} than the observed surface density $ \\mu_{0 \\, obs} \\simeq 120 \\; M_\\odot /{\\rm pc}^2 $ \\citep{kor,dona,span}.  It is illuminating to insert in eq.(\\ref{eqq}) the above value of the CDM surface density $ \\mu_0^{CDM} $ instead of the observed value $ \\mu_{0 \\, obs} $. This gives for the mass of the CDM particle $ m^{CDM} \\sim 60 $ GeV which is a typical wimp mass. Therefore, the linear approximation also provides a consistent value for the mass of the CDM particles from simulations.  \\medskip  These results show that our theoretical treatment captures many essential features of dark matter, allowing to determine its nature. When contrasted to the CDM surface density value obtained from CDM simulations (instead of the surface density value obtained from observations), our approach  gives for the dark matter particle mass the typical CDM wimps mass scale (GeV), fully consistent with CDM simulations.   \\medskip  It must be stressed that the linear framework presented here applies to any kind of DM particles: particles with mass in the keV scale reproduce all observed galaxy magnitudes within one order of magnitude, while wimps ($ m  \\sim 100 $ GeV) present discrepancies with observations of up to eleven orders of magnitude. This is a robust indication that the DM particle mass is in the keV scale.    \\begin{figure} \\begin{turn}{-90} \\psfrag{\"wim100.dat\"}{$ \\log_{10} \\rho_{lin}(r)_{wimp} $ vs. $ r $ in kpc} \\psfrag{\"nfwG.dat\"}{NFW profile} \\includegraphics[height=9.cm,width=9.cm]{perLnfw2} \\end{turn} \\caption{The linear density profile for 100 GeV wimps (broken green line) and the NFW profile (solid red line) for the same galaxy mass as the Burkert profile in fig. \\ref{plgg}. In all cases the densities are in g/${\\rm cm}^3 $ and $ r $ in kpc. The wimps linear density profile follows eq.(\\ref{asifin}). The wimp linear profile exhibits a cusp like the NFW profile.} \\label{perT} \\end{figure}"
    },
    "1004/1004.4000_arXiv.txt": {
        "abstract": "We present high resolution (R=55,000) optical spectra obtained with MIKE on the 6.5 m Magellan Clay Telescope as well as \\textit{Spitzer} MIPS photometry and IRS low resolution (R$\\sim$ 60) spectroscopy of the close (14 AU separation) binary, HD 101088, a member of the $\\sim$12 Myr old southern region of the Lower Centaurus Crux (LCC) subgroup of the Scorpius-Centaurus OB association. We find that the primary and/or secondary is accreting from a tenuous circumprimary and/or circumsecondary disk despite the apparent lack of a massive circumbinary disk. We estimate a lower limit to the accretion rate of \\.M $> 1\\times10^{-9} ~M_{\\odot} ~yr^{-1}$, which our multiple observation epochs show varies over a timescale of months. The upper limit on the 70 $\\mu$m flux allows us to place an upper limit on the mass of dust grains smaller than several microns present in a circumbinary disk of 0.16 M$_{moon}$. We conclude that the classification of disks into either protoplanetary or debris disks based on fractional infrared luminosity alone may be misleading. \\end {abstract} ",
        "introduction": "Most young stars are initially surrounded by optically thick accretion disks \\citep{beckwith90}. Though the frequency of binarity is a function of mass and formation environment, one-third of all main sequence stars in the Galactic disk are in binaries \\citep{lada06}. In nearby star forming regions, more than half \\citep{ghez93,simon95} of young stars have been observed to be members of binary systems. Since the formation of a binary is a common outcome of the star formation process, studying the structure and evolution of disks in binary systems is an important part of developing a complete understanding of the planet formation process. In the case of young binary systems, disks may surround the primary, the secondary, and/or both components \\citep{artymowicz94}. The separation between stars in a binary system is an important parameter that affects the geometry and structure of any disk material surrounding the stars. Optically thick disks around each member of young binaries have been observed for stars separated by as little as 14 AU \\citep{hartigan03}. The outer radii of disks in close binary systems are truncated through gravitational interactions between the binary components, limiting the amount of material available for planet formation and accretion onto the stars.  Optically thick accretion disks around single stars appear to dissipate within a few million years \\citep{haisch01}. Since the disks around each star in close binary systems are truncated to smaller outer radii than disks around single stars, they may disperse faster. \\citet{bouwman06} conducted a survey with the \\textit{Spitzer Space Telescope} of the $\\sim$8 Myr old $\\eta$ Chamaeleontis cluster and found that circumstellar disks were detected around 80\\% of single stars yet absent around 80\\% of the close binary stars. This is suggestive of a shorter timescale for disk removal in close binaries although the sample size was small and the binaries were not spatially resolved. Additional observations of disks in close binary systems are needed to confirm these results.  Here we present Magellan Inamori Kyocera Echelle (MIKE) R=55,000 optical spectroscopy along with \\textit{Spitzer} Infrared Spectrometer (IRS; Houck et al. 2004) and Multiband Imaging Photometer for \\textit{Spitzer} (MIPS; Rieke et al. 2004) observations of the close binary star, HD 101088, a member of the $\\sim$12 Myr old southern region of the Lower Centaurus-Crux (LCC) subgroup of the Scorpius-Centaurus OB association \\citep{dezeeuw99}. HD 101088 consists of a F5 primary star \\citep{houk75} and a secondary of unknown spectral type according to \\textit{Hipparcos} astrometry. The two components are separated by 0.15\\arcsec or 14 AU at a distance of 94 pc \\citep{vanleeuwen07}. Our high resolution optical spectra reveal broad, spatially unresolved H$\\alpha$ emission from this source which is indicative of ongoing stellar accretion. The \\textit{Spitzer} IRS spectrum shows there is very little, if any, hot circumstellar dust. The lack of strong emission in the mid- and far-infrared indicates the absence of a cold outer disk. ",
        "conclusions": "We have obtained high resolution optical spectroscopy with MIKE on the Magellan Clay telescope as well as \\textit{Spitzer} IRS spectroscopy and MIPS 24,70 $\\mu$m photometry of the close binary, HD 101088.  1. Broad H$\\alpha$ emission is present in our spectra and reveals ongoing stellar accretion and the presence of circumstellar gas.  We derive a lower limit on the accretion rate of \\.M $> 1\\times10^{-9} ~M_{\\odot} ~yr^{-1}$.  2. The truncated disks in such a close binary have viscous accretion lifetimes much shorter than the age of the system, suggesting that some source of replenishment must be present to maintain the ongoing accretion.  3. The upper limit at 70 $\\mu$m leads to a constraint on the dust mass in a circumbinary disk if it is present of $<0.16$ M$_{moon}$.  4. HD 101088 would be classified as having a debris disk based on its small fractional infrared luminosity, 7.0$\\times$10$^{-4}$.  The presence of ongoing accretion shows that the classification of disks into either protoplanetary or debris disks based on fractional infrared luminosity alone may be misleading.  5. We find that the primary and/or secondary is accreting from a tenuous circumprimary and/or circumsecondary disk despite the apparent lack of a massive circumbinary disk. This unique situation merits further study."
    },
    "1004/1004.1454_arXiv.txt": {
        "abstract": "Using multi-wavelength observations of SoHO/MDI, SOT-Hinode/blue-continuum (4504 \\AA), G-band (4305 \\AA), Ca II H (3968 \\AA) and TRACE 171 \\AA, we present the observational signature of highly twisted magnetic loop in AR 10960 during the period 04:43 UT-04:52 UT at 4 June, 2007. SOT-Hinode/blue-continuum (4504 \\AA) observations show that penumbral filaments of positive polarity sunspot have counter-clock wise twist, which may be caused by the clock-wise rotation of the spot umbrae. The coronal loop, whose one footpoint is anchored in this sunspot, shows strong right-handed twist in chromospheric SOT-Hinode/Ca II H (3968 \\AA) and coronal TRACE 171 \\AA\\, images. The length and the radius of the loop are $L\\sim$80 Mm and $a\\sim$4.0 Mm respectively. The distance between neighboring turns of magnetic field lines (i.e. pitch) is estimated as $\\approx$ 10 Mm. The total twist angle, $\\Phi\\sim$12$\\pi$ (estimated for the homogeneous distribution of the twist along the loop), is much larger than the Kruskal -Shafranov instability criterion. We detected clear double structure of the loop top during 04:47-04:51 UT on TRACE 171 \\AA \\ images, which is consistent with simulated kink instability in curved coronal loops (T{\\\"o}r{\\\"o}k et al. 2004). We suggest, that the kink instability of this twisted magnetic loop triggered B5.0 class solar flare, which occurred between 04:40 UT and 04:51 UT in this active region. ",
        "introduction": "Solar coronal magnetic field has complex topology which is caused due to photospheric motions and emergence of new magnetic flux. The complex configurations often lead to various instability processes, which eventually trigger solar flares and CMEs (Coronal Mass Ejections). Kink instability is one of those processes and it is connected to the azimuthal twist of magnetic tubes. The exact amount of twist required to trigger the kink instability depends on various factors including loop geometry and overlying magnetic fields (e.g., Hood and Priest, 1979; Lionello et al., 1998; Baty et al., 1998; Baty, 2001; T{\\\"o}r{\\\"o}k et al., 2004; Fan and Gibson, 2003, 2004, Leka et al., 2005 and references therein). Recent observations of kink instability accompanied with full filament eruption (Williams et al., 2005), partial cavity eruption (Liu et al., 2007), partial filament eruption (Liu et al. 2008), and failed filament eruption (Alexander et al. 2006) indicate to its importance in filament interaction with magnetic environment. Kink instability is also found to be an efficient mechanism for solar eruptive phenomena, e.g, triggering solar flares and CMEs ( Sakurai 1976; Hood, 1992; T{\\\"o}r{\\\"o}k and Kliem 2005; Kliem and T{\\\"o}r{\\\"o}k, 2006 and references therein).  Although we have few observational evidences of kink instability in various magnetic structures (e.g., coronal loops, filaments), the theory is much more established in terms of modeling and numerical simulations. Previous theoretical models, based on straight tube assumption, have studied intensively various aspects of kink instability in the solar corona including the formation of current sheets (Baty and Heyvaerts 1996; Gerrard et al., 2001; Gerrard and Hood, 2003; Hynes and Arber 2007) and magnetic topology (Baty, 2000; Lionello et al., 1998). Later on, more sophisticated models based on curved flux tube geometry, addressed the formation of current sheets (e.g. T{\\\"o}r{\\\"o}k et al., 2004), loop response to injected twist through its footpoints (e.g., Klimchuk, 2000; Tokman and Bellan, 2002; Aulanier et al., 2005 and references cited there), and eruption of kink unstable loops through overlaying arcades (e.g. T{\\\"o}r{\\\"o}k and Kliem, 2005; Fan, 2005).  In this paper, we present the observational evidence of highly twisted coronal loop in the AR NOAA 10960 as observed on 04 June, 2007 between 04:43 UT and 04:52 UT, which probably caused B5.0 class flare during this time.  We use observations from several different instruments in order to cover almost whole solar atmosphere. SOHO/MDI and SOT-Hinode/blue-continuum (4504 \\AA) have been used to observe the photospheric part of the active region. We used data from SOT-Hinode Ca II H (3968 \\AA) to study the chromospheric level. Finally, TRACE 171 \\AA \\ observations are used to study the coronal structure. In Section 2, we describe multi-wavelength observations of twist and kink instability in AR 10960 on 04 June, 2007. In Section 3, we present some theoretical interpretations. In the last section, we present some discussions and conclusions. ",
        "conclusions": "Using multiwavelength observations of SoHO/MDI, SOT-Hinode/blue-continuum (4504 \\AA), G-band (4305 \\AA), Ca II H (3968 \\AA) and TRACE 171 \\AA , we find the observational signature of highly twisted coronal loop in AR 10960 during the period of 04:43 UT-04:52 UT at 04 June, 2007. This twisted loop was probably unstable to the kink instability, which triggered small B-class flare during the same time. SOT-Hinode/blue-continuum (4504 \\AA) images show that the penumbral filaments of small positive polarity sunspot have counter-clock wise twist around its center (Figure 3). The twist is probably caused by the clockwise rotation of the sunspot umbrae as expected in the southern hemisphere.  The observed coronal loop whose one footpoint is anchored in this spot, shows strong right-handed twist in chromospheric (Ca II H 3968 \\AA \\ , Figure 4) and coronal (TRACE 171 \\AA \\ , Figure 5) images. We estimated the loop length and radius as $\\sim$ 80 and $\\sim$ 4 Mm respectively. From TRACE 171 \\AA\\ time sequence we estimate the pitch of twisted loop as $\\sim$ 10 Mm. The twist is assumed to be homogeneous along the whole loop as any asymmetry can be quickly smoothed out over the Alfv\\'enic time, which is estimated as $\\sim$80 s for our loop. The total twist angle for the whole loop is $\\Phi$$\\sim$12$\\pi$ in the case of homogeneous distribution. This is much larger than the Kruskal-Shafranov kink instability criterion, and approximately equal to the more conservative thin-tube estimate for the kink instability. The loop top shows clear double structure during 04:47-04:51 UT, which was accompanied by sudden enhancement of soft X-ray flux around 04:49 UT observed by XRT/Hinode.  The strong twist of the coronal loop probably leads to the kink instability, although no clear displacement of the loop axis or sigmoid structure have been detected during the flaring time. However, the observed double structure of the loop top probably indicates to the current sheet induced by the kink instability as it was suggested by numerical simulations (T{\\\"o}r{\\\"o}k et al. 2004). This current sheet probably triggered the B-class flare via reconnection.  Instability of internal kink mode (Arber et al. 1999; Haynes and Arber, 2007), where the kink structure is not apparent from the global field shape of the active region, may be considered as a triggering mechanism for the B-class flare. Haynes and Arber (2007) considered the example of short coronal loop ($L$= 10 Mm) with zero net current. This means that the magnetic twist changes the sign along the loop radius, which causes the confinement of unstable kink mode in space. This process has also been shown by Arber et al. (1999) to release sufficient energy to cause a transient brightening of confined loops. However, there are some questions related to the simulation of Haynes and Arber (2007). First, our observed loop is longer compared with the simulated one by a factor of 8. Therefore, it is unclear if the same mechanism may work for longer loops. Second, it is unclear why the twist should have opposite signs along the tube radius. Haynes and Arber (2007) considered photospheric motion as a generator of this configuration, however no clear process has been suggested.  Twisted coronal loops can be also subject to the sausage instability. Linear sausage pinch instability occurs when $B^2_{\\phi}> 2 B^2_z$ (Aschwanden 2004). It gives the critical twist angle of $\\sim 9 \\pi$ in our loop parameters. Nonlinear analysis of Zaqarashvili et al. (1997) shows even smaller critical twist angle. The observed twist in our loop is much larger. However, the sausage pinch does not lead to observed double structure at the loop top, which rules out the occurrence of sausage instability in the loop.  It is quiet possible that the observed twist of coronal loop corresponds to the helical structure formed after the kink instability. Then, the radius of the loop will be bit smaller than $\\sim 4$ Mm. However, this means that the wave length of unstable mode is $\\sim 15$ Mm (i.e. twice the distance between neighboring bright areas along the loop, see Figure 5), which is much less comparing to the loop length. It means that the higher order harmonic rather than fundamental one was unstable to the kink instability. This unexpected fact needs special justification, which may require detailed numerical simulations.  It should be mentioned, that the small B5.0 flare was precursor for stronger M-class flare, which occurred just 15 min later at 05:06 UT. It is possible that the magnetic field reconfiguration due to the small B5.0 flare caused global kink instability with consequent reconnection and global energy release. It is interesting to check whether the big flares are usually preceded with smaller energy release. Liu et al. (2003) and Gary and Moore (2004) have observed recurrent flare activity of the active region AR 10030 in the northern hemisphere on 15 July 2002, which was accompanied by CME. The large flare/CME was preceded by a small flare and the authors supposed the rise of helical flux rope as the source for the activity. Their observations show strongly twisted loop in TRACE images (see Fig. 1 in Gary and Moore, 2004), which arose upwards and disappeared during short time interval ($\\sim$ 50 s). One may suggest that the strong twist of our loop may be also formed due to the similar rise of magnetic flux rope. However, there are significant differences between the observations. First, no new magnetic flux emergence is seen at the photospheric level in our case. And second, our twisted loop stays unchanged during longer time interval (at least 5 min) and no sign of upward motion is detected. Therefore, we think that the rise of helical flux rope is not relevant to our observations, which probably suggests at least two different triggering mechanisms for solar flares. The similarity between the two events is that the small flares seem to be the pre-cursors of large flares. On the other hand, the flare in the active region AR 10030 on 15 July 2002 was accompanied by CME (Liu et al. 2003, Gary and Moore 2004), while the flare in AR 10960 on 04 June 2007 was not (Kumar et al., 2010). Therefore, we suggest that the large flares accompanied with energetic CMEs can be triggered by flux-rope eruption with significant changes in the photospheric fields, while the moderate flares without CME are triggered by some instabilities (e.g., kink instability). More statistical study is required to make a firm conclusion.        In conclusion, we observe the strong twist of coronal loop in AR 10960 during small B5.0 flare between 04:43 UT--04:52 UT at 04 June, 2007. The loop top shows clear double structure, which is consistent with simulated kink instability of curved coronal loop (T{\\\"o}r{\\\"o}k et al. 2004). We suggest that the current sheet formed at the loop top due to the kink instability was the reason for the B5.0 flare."
    },
    "1004/1004.3360.txt": {
        "abstract": "{In this work, we study the large scale structure formation in the modified gravity in the framework of Palatini formalism and compare the results with the equivalent smooth dark energy models as a tool to distinguish between these models. Through the inverse method, we reconstruct the dynamics of universe, modified gravity action and the structure formation indicators like the screened mass function and gravitational slip parameter. Consequently, we extract the matter density power spectrum for these two models in the linear regime and show that the modified gravity and dark energy models predictions are slightly different from each other at large scales. It is also shown that the growth index in the modified gravity unlike to the dark energy models is a scale dependent parameter. We also compare the results with those from the modified gravity in the metric formalism. The modification on the structure formation can also change the CMB spectrum at large scales however due to the cosmic variance it is hard to detect this signature. We show that a large number of SNIa data in the order of 2000 will enable us to reconstruct the modified gravity action with a suitable confidence level and test the cosmic acceleration models by the structure formation.}   % %PACS numbers: 04.50.+h, 95.36.+x, 98.80.-k   %Department of Physics, Sharif University of Technology, P.O.Box %11365--9161, Tehran, Iran      %\\baselineskip=12pt %\\def\\be{\\begin{equation}} %\\def\\ee{\\end{equation}} %\\def\\bea{\\begin{eqnarray}} %\\def\\eea{\\end{eqnarray}} %\\def\\orc{\\Omega_{r_c}} %\\def\\om{\\Omega_{\\text{m}}} %\\def\\E{{\\rm e}} %\\def\\bearst{\\begin{eqnarray*}} %\\def\\eearst{\\end{eqnarray*}} %\\def\\peleven{\\parbox{11cm}} %\\def\\peffec{\\peight{\\bearst\\eearst}\\hfill\\peleven} %\\def\\pspace{\\peight{\\bearst\\eearst}\\hfill} %\\def\\ptwelve{\\parbox{12cm}} %\\def\\peight{\\parbox{8mm}} %\\twocolumn[\\hsize\\textwidth\\columnwidth\\hsize\\csname@twocolumnfalse\\endcsname  %\\title{Distinguishing modified gravity from smooth dark energy by the structure %formation probes in Palatini formalism}    %\\newpage % ] \\begin{document} ",
        "introduction": "For more than a decade, the positive acceleration of the universe is one of the challenging questions in  physics \\cite{R04,Dun09}. The physics and the mechanism behind this phenomenon is unknown. The Cosmological Constant (CC) is the most straightforward suggestion to describe an accelerating universe \\cite{Adam98,Davis07}. However the fundamental questions like the fine tuning and coincidence problems opened new horizons to introduce the alternative models \\cite{Wienberg89} like Dark Energy (DE) \\cite{Peeb03} and models of the Modified Gravity (MG) \\cite{Carroll04}. The other motivation for the modified gravity models is unifying the dark matter and the dark energy problems in a unified formalism \\cite{rahsaff}.    The MG theories need to be tested in two domains of the observations: (a) cosmological scales \\cite{Jain08} and (b) local gravity \\cite{Faraoni07}. The simplest modification of gravity is the extension of Einstein-Hilbert action where instead of Ricci scalar, we use a generic function of $f(R)$. These modified gravity theories are divided into the Metric and Palatini formalism \\cite{Tsujikawa2010}. There are a lot of debates in the literature on the viability of this type of modified gravity model concerning the local gravity tests \\cite{Faraoni04} where a vast range of them had been falsified \\cite{Barausse2008,Kobayashi2008}. Also many models are proposed to evade the falsification criteria \\cite{Upadhye2009,HuSawicki2007,Starobinsky}. The Palatini formalism is a second order differential equation which has more simpler structure than the fourth order metric formalism. Since all the conventional  field equations  are described by the second order differential equation, it seems that gravity should also follow this rule. So this is the main advantage of working with the Palatini formalism. From cosmological point of view, we will show that for a given dynamics of the Universe there is a one-to-one map to the corresponding action in the Palatini formalism. Also in this formalism we don't have the instability problem as in the metric formalism. While there are advantages in this formalism, there is challenging questions in the Palatini formalism as the microscopic behavior of the matter fields where in the Einstein frame, these models disagrees with the standard scalar-matter coupling theories \\cite{Flanagan2004}. Means that we can not apply the perturbation theory in these scales as the gravity field is directly relates to the energy-momentum tensor, in contrast to the Einstein gravity where the amount of the perturbation of the metric at a given point is averaged over all the space. Although there are debates on how we should do averaging procedure over the microscopic scales \\cite{Li}, it has been shown that the problematic microscopic behavior of this theory can be solved in the case of $f(R)$ very near to the $\\Lambda$CDM model \\cite{Tsujikawa2010}.  %In a series of works  it have been shown that the masses, geodesics %of the particles in the palatini formalism is not distinguishable GR %plus cosmological constant.      Keeping in mind that there are many fundamental questions in the Palatini formalism , here in this work we mainly focus on the behaviors of the modified gravity models at the cosmological scales and their predictions on the structure formation. The main question we address in this work is the investigation of cosmological probes in the structure formation to distinguish between CC, DE and MG models \\cite{Bertschi2008}. We first reconstruct the dynamics of the background and the corresponding MG action via the inverse method from the SNIa data, described in \\cite{Baghram09}. The dynamics of the background (i.e Hubble parameter $H=H(z)$) can only distinguish between CC and the alternative models \\cite{Baghram09} however in the level of probing the expansion history of the Universe, MG and DE models are not distinguishable. Consequently, we use the structure formation probes as the promising tool to distinguish between alternative models \\cite{Pogosian2008,Lue2004}. It is worth to mention that the alternative models can be extended to a complicated forms such as interacting dark energy-dark matter models \\cite{Wei2008} or clustering DE theories,  which can not be distinguished from the MG models even in the level of the structure formation probes \\cite{Kunz2007}.    In this work we focus on the structure formation issue in the modified gravity--Palatini formalism and compare the observational effects with that in the smooth dark energy models (sDE). We derive the screened mass function defined as the fraction of effective gravitational constant to the standard gravitational theory as $Q\\equiv {G_{eff}}/{G_{N}}$. The effective gravitational constant appears in the Poisson equation depends on the scale of the structure. The other relevant parameter appear in our calculation is the gravitational slip parameter $\\gamma=-{\\Phi}/{\\Psi}$, which is the fraction of spatial perturbation of the metric to the perturbation of the time-time element. The screened mass function as well as the gravitational slip parameter in DE models are equal to one in GR while in MG models they are time and scale dependent parameters  . This scale dependance of structure formation in modified gravity theories is a well known effect discussed in literature \\cite{HuSawicki2007,Starobinsky,Zhang08}. A combination of screened mass and gravitational slip parameters recently have been used for distinguishing the MG in metric formalism with the alternative models via analyzing the structure formation and weak lensing \\cite{rey2010}. We show that the derived reconstructed matter density power spectrum and growth index in this model are affected by the screened mass function. Also we study the effect of gravitational slip parameter on the Integrated Sachs-Wolf (ISW) and compare it with recent observations of WMAP data. As a comparison with the metric formalism, we compare our results with that of structure formation in the metric formalism.  The structure of this article is as follows: In section \\ref{section2} we review the standard equations governing the relativistic structure formation. In section \\ref{section3} we re-derive the structure formation equations for DE and Palatini MG models. Also in this section the metric formalism structure formation equations are derived for comparison with that in the Palatini formalism. %and compare their predictions. In section \\ref{section4} we extract the screened mass function and gravitational slip parameters, reconstructed from a modified gravity equivalent to a dark energy model and the effect of screened mass function on the power spectrum of the structures as well as the growth index is studied. Also the screened mass in the metric formalism is obtained in order to show the scale dependence of structure's growth. Details of the calculation in metric formalism is given in Appendix A. In section \\ref{section5} we investigate the effect of gravitational slip parameter on CMB power spectrum due to the ISW effect and compare the results with recent WMAP data. The conclusion is presented in section \\ref{conc}. %In Appendix A, we study the structure formation %in Metric formalism and obtained the power-spectrum and growth index %in order to compare with Palatini formalism results. ",
        "conclusions": "\\label{conc}   One of the important questions in the problem of the acceleration of the universe is either this acceleration is produced by a dark energy fluid or that is a manifestation of the gravity modification. It has been shown that the expansion history of the universe can not dynamically distinguish between these two models and at this stage they are equivalent.  In this work we compared various  aspects of the structure formation in the Modified Gravity (MG) and smooth Dark Energy (sDE) models. We used an ansatz for the equation of state of a dark energy model and obtained the equivalent modified gravity . Using a synthetic SNIa data set from the SNAP observation, we did the inverse problem approach and obtained the effective action of the gravity (i.e. $f(R)$). It has been shown that the structure formation unlike to the background dynamics can distinguish between these two models.  In the modified gravity, the modification of the Einstein-Boltzman equation for the structure formation has extra terms compare to the sDE model. This modification can be given in terms of screened mass function and gravitational slip parameter. We obtained these two parameters in the MG model and solved the differential equation for the evolution of the structure formation. Using the Harrison-Zeldovich initial condition for the structures, finally we obtained the power spectrum and the growth index of the structures in the two scenarios in the linear regime of the structure formation. We have shown that for the structures with the larger size, the effective spectral index at the present time is slightly more than one. From the observational point of view, sampling of large scale structure more than $100$Mpc is needed to test this effect. On the other hand, we showed that the growth index parameter in the MG unlike to the sDE models is a scale dependent parameter. Next generation of peculiar velocity survey may measure this effect. Finally we obtained the effect of the MG structure formation on the Integrated Sachs-Wolfe effect on CMB map. We showed that the deviation of the CMB power spectrum due to the MG is larger in the small $l$s, however it is not too large to distinguish between the two models due to the uncertainty from the cosmic variance.  \\appendix"
    },
    "1004/1004.1662_arXiv.txt": {
        "abstract": "We carry out a linear analysis of the vertical normal modes of axisymmetric perturbations in stratified, compressible, self-gravitating gaseous discs in the shearing box approximation. An unperturbed disc has a polytropic vertical structure that allows us to study specific dynamics for subadiabatic, adiabatic and superadiabatic vertical stratifications, by simply varying the polytropic index. In the absence of self-gravity, four well-known principal modes can be identified in a stratified disc: acoustic p-modes, surface gravity f-modes, buoyancy g-modes and inertial r-modes. After classifying and characterizing modes in the non-self-gravitating case, we include self-gravity in the perturbation equations and in the equilibrium and investigate how it modifies the properties of these four modes. We find that self-gravity, to a certain degree, reduces their frequencies and changes the structure of the dispersion curves and eigenfunctions at radial wavelengths comparable to the disc height. Its influence on the basic branch of the r-mode, in the case of subadiabatic and adiabatic stratifications, and on the basic branch of the g-mode, in the case of superadiabatic stratification (which in addition exhibits convective instability), does appear to be strongest. Reducing the three-dimensional Toomre's parameter $Q_{3D}$ results in the latter modes becoming unstable due to self-gravity, so that they determine the onset criterion and nature of the gravitational instability of a vertically stratified disc. By contrast, the p-, f- and convectively stable g-modes, although their corresponding $\\omega^2$ are reduced by self-gravity, never become unstable however small the value of $Q_{3D}$. This is a consequence of the three-dimensionality of the disc. The eigenfunctions corresponding to the gravitationally unstable modes are intrinsically three-dimensional. We also contrast the more exact instability criterion based on our three-dimensional model with that of density waves in two-dimensional (razor-thin) discs. Based on these findings, we comment on the origin of surface distortions seen in numerical simulations of self-gravitating discs. ",
        "introduction": "Self-gravity plays an important role in a variety of astrophysical systems. It is a main agent determining the dynamical evolution of star clusters, galaxies, various types of accretion discs, etc. Particularly in protoplanetary discs, that are the central subject of our study, self-gravity provides one of the main source of outward angular momentum transport through the excitation of density waves \\citep{PS91,LB94,LR04,LR05} and is able to cause fragmentation of a disc into bound clumps, or planets, under certain conditions \\citep{Gam01,Riceetal03,RLA05,B04,Mayeretal07,Raf07}. Starting with the seminal paper by \\citet{T64}, there have been a large number of studies of the stability of self-gravitating gaseous discs, both in the linear \\citep[e.g.,][]{GLB65a,GLB65b,GT78,ARS89,Bertetal89,PL89,LKA97} and non-linear regimes, including other relevant physical factors (e.g., heating, cooling, radiation transport) with up-to-date numerical techniques \\citep[e.g.,][]{PS91,LB94,B98,Petal00,Petal03,Gam01,B03,JG03,Riceetal03,RLA05,B04,Metal05,Betal06,Mayeretal07, SW09}.  Linear stability analysis in a vast majority of cases is restricted, for simplicity, to razor-thin, or two-dimensional (2D) discs that are obtained by vertically averaging all quantities. In other words, perturbations are assumed to have large horizontal scales compared with the disc thickness. In this limit, the well-known Toomre's parameter $$ Q_{2D}=\\frac{c_s\\Omega}{\\pi G \\Sigma} $$ controls the stability of self-gravitating discs \\citep{T64}. In this case, a density wave, which is the only mode in a 2D disc, is influenced by self-gravity and thus can become unstable, as the local dispersion relation for the latter clearly demonstrates \\citep{GT78,BT87,Bertetal89}.  Stability analysis in a more realistic case of self-gravitating three-dimensional (3D) discs is more complicated. The disc is vertically stratified due to both its own self-gravity and the vertical component of the gravity of a central object. Depending on the nature of the stratification, there exists a whole new set of various vertical modes in the system (see below), some of which can become unstable due to self-gravity on horizontal length scales comparable to the disc thickness. In this situation, the vertical variation of perturbations is important and for a correct characterization of the gravitational instability it is necessary to introduce another parameter not involving height-dependent variables, such as the sound speed in Toomre's parameter. Furthermore, not all types of stratification permit two-dimensional modes, that is, modes with no vertical motions commonly occurring in the 2D treatment. For example, in non-self-gravitating discs with polytropic vertical structure, there are no 2D modes \\citep{LPS90,LP93a,KP95} implying that the dynamics does not always reduce to that of the 2D case. Therefore, a more accurate stability analysis of self-gravitating discs should necessarily be three-dimensional.  Obviously, before studying the gravitational instability of stratified discs, one must first classify and characterize vertical normal modes of perturbations in the simplified case of no self-gravity. Analysis of the modal structure of stratified, polytropic, compressible, non-self-gravitating discs has been done in several papers: \\citet[][hereafter RPL]{RPL88}, \\citet[][hereafter KP]{KP95}, \\citet{Og98}. In convectively stable discs, i.e., with subadiabatic vertical stratification, four principal types of vertical modes can be distinguished. These modes are: acoustic p-modes, surface gravity f-modes, buoyancy g-modes and inertial r-modes. The modes are named after their corresponding restoring forces, which can be well identified for each mode at horizontal wavelengths smaller than the disc height and are provided by one of the following: compressibility/pressure, displacements of free surfaces of a disc, buoyancy due to vertical stratification and inertial forces due to disc rotation, respectively, for the p-, f-, g- and r-modes. In the case of superadiabatic stratification, the r- and g-modes merge and appear as a single mode, which becomes convectively unstable for horizontal wavenumbers larger than a certain value (RPL); the p- and f-modes remain qualitatively unchanged. For neutral/adiabatic stratification, buoyancy is absent and the g-mode disappears. \\emph{The main purpose of this paper is to investigate how self-gravity modifies the frequencies and the structure of the eigenfunctions of these modes, which mode acquires the largest positive growth rate due to self-gravity and, therefore, determines the onset criterion and nature of the gravitational instability of a stratified disc.} So, the mode dynamics in the 3D case can appear more complex than that in the 2D one, where only the density wave mode can be subject to gravitational instability. Previously, \\citet[][hereafter GLB]{GLB65a} considered gravitational instability in a uniformly rotating gaseous slab with an adiabatic vertical stratification, thereby leaving out all modes associated with buoyancy. Other studies also considered the gravitational instability of 3D galactic discs, however, the analysis was essentially 2D, finite-thickness effects were only taken into account by means of various reduction factors in 2D dispersion relations \\citep{Shu68,Rom92,Rom94}. In all these studies, as in GLB, the main focus was on finding the criterion for the onset of gravitational instability, so that a full analysis of various types of vertical normal modes existing in stratified self-gravitating discs was not carried out. Actually, we generalize the study of GLB to subadiabatic and superadiabatic stratifications having different modal structure.  Another motivation for our study is that the f-mode is thought to play an important dynamical role in self-gravitating discs. The non-linear behaviour of 3D perturbations involving large surface distortions, as seen in numerical simulations, has been attributed to the surface gravity f-mode \\citep{Petal00}. However, this was done without analysing the behaviour of other vertical modes under self-gravity. It was shown that the f-mode leads to a large energy dissipation in the vicinity of the disc surface, which may facilitate disc cooling, because the energy is deposited at smaller optical depth where it can be radiated away more quickly \\citep[see e.g.,][]{JG03,Betal06}. Later it was realized that in fact the non-linear vertical motions in self-gravitating discs can be much more complex than just the f-mode and can have a shock character \\citep[shock bores,][]{BD06}. Thus, in the 3D case, the dynamics of self-gravitating discs is much richer and diverse than that of idealized 2D ones and requires further study. To fully understand the origin of such three-dimensional effects and what type of instability they are associated with, one should start with a rigorous linear study of the characteristic properties of all the types of vertical normal modes mentioned above, not only the f-mode, in the presence of self-gravity. The present work is just a first step in this direction.  Numerical simulations of self-gravitating discs are often in the context of global discs \\citep[e.g.,][]{Petal00,Petal03,Riceetal03,RLA05,LR04,LR05,Betal06,Bol09} and, therefore, are not always able to well resolve vertical motions, which, as shown in the present study, inevitably arise during the development of the gravitational instability associated with intrinsically three-dimensional modes. So, these simulations may not quite accurately capture all the subtleties of the gravitational instability in 3D discs. In this connection, we should mention the work by \\citet{N06} that extensively discusses the issue of vertical resolution and its importance in the outcome of the gravitational instability in numerical simulations of self-gravitation discs. Resolving and analysing vertical motions are also crucial for properly understanding cooling processes in discs and, particularly, whether convection is able to provide sufficiently effective cooling for disc fragmentation to occur, which is still a matter for debate in the literature \\citep{B04,Mayeretal07,Betal06,Betal07,Raf07}. In addition, these studies, for simplicity, use the criterion for gravitational instability based on the two-dimensional Toomre's parameter, which, as we will demonstrate, cannot always be uniquely mapped into an analogous three-dimensional stability parameter and give a precise criterion for the onset of gravitational instability.  In this paper, following other works in a similar vein: \\citet[][hereafter LP]{LP93a}, KP, \\citet[][hereafter LO98]{LO98}, we adopt the shearing box approximation and consider the linear dynamics of vertical normal modes of perturbations in a compressible, stratified, self-gravitating gaseous disc with Keplerian rotation. In the unperturbed disc, pressure and density are related by a polytropic law, which is a reasonably good approximation for optically thick discs (see e.g., LO98). This allows us to consider the specific features of the dynamics for subadiabatic, adiabatic and superadiabatic vertical stratifications by simply varying the polytropic index. As a first step towards understanding the effects of self-gravity on the perturbation modes, we restrict ourselves to axisymmetric perturbations only. The linear results obtained here will form the basis for studying the non-linear development of gravitational instability in the local shearing box approximation that allows much higher numerical resolution than global disc models can afford.  The paper is organized as follows. The physical model and basic equations are introduced in Section 2. The classification of vertical modes in the absence of self-gravity is performed in Section 3. Effects of self-gravity on all normal modes are analysed in Section 4. In Section 5, we focus on the properties of gravitational instability in 3D. Comparison between the criteria of gravitational instability in 3D and 2D is made in Section 6. Summary and discussions are given in Section 7. ",
        "conclusions": "In this paper, we have analysed the axisymmetric normal modes of perturbations in stratified, compressible, self-gravitating gaseous discs with subadiabatic, adiabatic and superadiabatic vertical stratifications. First, we performed a classification of perturbation modes in stratified discs in the absence of self-gravity to compare with previous calculations. Four well-known main types of modes can be distinguished: acoustic p-modes, surface gravity f-modes, buoyancy g-modes and inertial r-modes. The restoring forces for these modes for large radial wavenumbers are mainly provided by one the following: pressure/compressibility, displacements of the disc surface, buoyancy and inertial forces due to disc rotation for the p-,f-,g- and r-modes, respectively. For smaller wavenumbers, the restoring force for each mode is a combination of these forces. In the case of adiabatic stratification, buoyancy is zero and, therefore, the g-mode disappears, while other modes remain qualitatively unchanged. For superadiabatic stratification, the g-mode becomes convectively unstable and merges with the r-mode, so that only a single convectively unstable mode appears in the dispersion diagram at $\\omega^2 \\leq \\kappa^2$, which we still call the g-mode. Due to the reflection symmetry of the equilibrium vertical structure with respect to the midplane, each mode comes in even and odd pairs. By our terminology, for even (odd) modes, pressure and gravitational potential perturbations are even (odd), while the perturbations of vertical velocity and derivative of potential are odd (even) functions of the vertical coordinate. After classifying and characterizing modes in the absence of self-gravity, we introduced self-gravity in the perturbation equations and investigated how it modifies the properties of these modes. We found that self-gravity, to some extent, reduces the frequencies of all normal modes at radial wavelengths comparable to the disc height, but its influence on the basic even r-mode, in the case of subadiabatic and adiabatic stratifications, and on the basic even g-mode, in the case of superadiabatic stratification, appears to be strongest. With decreasing $Q_{3D}$, these modes become unstable due to self-gravity and thus determine the gravitational instability of a vertically stratified disc. The basic even g-mode also exhibits convective instability due to a negative entropy gradient but, unless the disc is strongly self-gravitating, these two instabilities grow concurrently in the linear regime, because their corresponding radial scales are separated. We also obtained the corresponding criterion for the onset of gravitational instability in 3D, which is more exact than the standard instability criterion in terms of the 2D Toomre's parameter, $Q_{2D}<1$, for axisymmetric density waves in razor-thin discs. By contrast, the p-, f- and convectively stable g-modes have their $\\omega^2$ reduced by self-gravity, but never become unstable for any value of $Q_{3D}$. This is a consequence of the three-dimensionality of the disc. The eigenfunctions associated with the gravitationally unstable modes are intrinsically three-dimensional, that is, have non-zero vertical velocity and all perturbed quantities vary over the whole vertical extent of the disc. In this regard, we would like to mention that resolving the gravitationally most unstable mode in numerical simulations thus reduces to properly resolving the disc height (together with resolving the corresponding radial scale, which, as shown here, appears somewhat larger than the disc height). So, the criterion of \\citet{N06} that at least $\\sim 4$ particle smoothing lengths should fit into per scale height may apply in three-dimensional SPH simulations. This implies a substantially larger number of SPH particles per vertical column because the disc itself may extend over many scale heights. He also shows that a similar criterion applies to grid-based simulations.  Here for simplicity, and also to gain the first insight into the effects of self-gravity on the vertical modes in stratified discs, we considered only axisymmetric perturbations. Non-axisymmetric perturbations are dynamically richer, though more complicated, because the phenomena induced by Keplerian shear/differential rotation -- strong transient amplification of perturbations and \\emph{linear} coupling of modes (not to be confused with non-linear mode-mode interactions) -- come into play for these type of perturbations. Transient (swing) amplification of perturbations (density waves) has been studied previously in razor-thin 2D approximation \\citep{GLB65b,GT78,T81,MC07}. From the analysis presented here, we may expect that in the linear regime, the non-axisymmetric basic even r-mode can undergo larger transient amplification due to self-gravity than other modes in the disc. This transient amplification of perturbations may be important for explaining the large burst phases seen in numerical simulations at the initial stages of the development of gravitational instability in discs \\citep{Riceetal03,RLA05,LR04,LR05,Betal06}. So, in this respect one should analyse and quantify the transient amplification of non-axisymmetric perturbations in stratified 3D self-gravitating discs starting with linear theory. As for the linear coupling of modes, it was demonstrated that in non-self-gravitating stratified discs, Keplerian shear causes rotational (vortex) mode perturbations to couple with and generate g-mode perturbations \\citep{TCZ08}. In the context of 2D discs, it was shown that vortex mode perturbations can also excite density waves due to shear \\citep{Bodoetal05,MC07,MR09,HP09}. In the 3D case, there are a larger number of modes in a disc and it is quite possible that some of them may appear linearly coupled due to shear and, therefore, be able to generate each other during evolution, especially the f-mode and the basic branch of the r-mode (because they vary on comparable vertical scales in the presence of self-gravity, Figs. 9, 10). Another related problem also of interest is the interaction between self-gravity and the MRI in magnetized discs \\cite[see e.g.,][]{F05}. In particular, how self-gravity can modify the growth rates of magnetic normal modes responsible for the MRI. Actually, this will be the generalization of the extensive analysis of normal modes in magnetized discs by \\cite{Og98}.  \\cite{LO98} showed that in non-self-gravitating discs with polytropic vertical stratification, an external forcing preferentially excites the f-mode, because it has the largest responsiveness to an external driving compared to other modes. This mode, propagating through a disc, results in energy dissipation near the disc surface. Based on these results and partly on the properties of f-modes in stellar dynamics, \\citet{Petal00} identified the behaviour of 3D perturbations in self-gravitating discs involving large surface distortions and the resulting energy dissipation in the upper layers, with f-mode dynamics. The self-stimulated potential was thought to play the role of an external/tidal force. However, it is not obvious that the effect of a self-stimulated potential is the same as that of the external potential. In fact, our analysis has revealed that in self-gravitating discs, in addition to the f-mode, the r-mode can also be dynamically important, because this mode appears to be subject to gravitational instability, while the f-mode is not. The eigenfunctions of the gravitationally unstable basic even r-mode differ from those of the r-mode in non-self-gravitating discs in that they are no longer concentrated near the midplane and behave somewhat similarly to the eigenfunctions of the f-mode: they vary over the whole vertical extent of the disc and also involve noticeable perturbations of the disc surface. Consequently, like the f-mode, the gravitationally unstable r-mode can, in principle, also induce gas motion causing large surface distortions and resultant energy dissipation in the upper layers of the disc, which is thought to play a role in enhancing disc cooling \\cite[because the energy is deposited in the upper layers with small optical depth, it can be radiated away more quickly and effectively cool the disc, but this is a subject of further study, see e.g.,][]{JG03,Betal06}. Furthermore, in the case of non-axisymmetric perturbations, as mentioned above, because of shear, the gravitationally unstable r-mode can couple with and generate the strong f-mode. So, the surface distortions may be caused by a combination of the f- and r-modes. In order to explore where dissipation can predominantly occur in a self-gravitating disc, one needs to generalize the analysis of \\cite{LPS90}, LP, LO98, \\cite{Bateetal02} on the propagation of waves in stratified non-self-gravitating discs and consider the propagation of non-axisymmetric modes in stratified self-gravitating discs. The dispersion properties of modes in the presence of self-gravity as found here are one of the necessary things for studying mode propagation.  Another point we want to raise concerns the spatial distribution of temperature. In order to realistically model the cooling of protoplanetary discs, \\citet{Betal07} employed the radiative transfer technique. In the vertical $z-$direction, the radiative transfer equation was solved exactly assuming a plane-parallel atmosphere approximation, but in the radial direction only the radiation diffusion approximation was employed. However, as our linear results (Fig. 11) and other non-linear simulations \\citep[e.g.,][]{Metal05,Betal06} demonstrate, temperature and, therefore, opacity may vary on comparable scales in both the radial and vertical directions and have very complex structure in the non-linear regime. This implies that a more general radiative transfer treatment based on solving the ray equation in all directions, rather than using the diffusion approximation in either direction, would be more appropriate for better understanding cooling processes."
    },
    "1004/1004.4101_arXiv.txt": {
        "abstract": "{The International Gamma-Ray Astrophysics Laboratory (INTEGRAL) is discovering a large number of new hard X-ray sources, many of them being HMXBs. The identification and spectral characterization of their optical/infrared counterparts is a necessary step in undertaking a detailed study of these systems.} {In a previous paper, we presented spectral analyses and classifications of six newly discovered \\emph{INTEGRAL} sources. In this paper, we extend the analysis to IGR~J16493--4348.} {We used the ESO/VLT ISAAC spectrograph to observe the proposed IR counterpart to the source, obtaining a $K_\\mathrm{s}$ medium-resolution spectrum ($R = 500$) with a signal-to-noise ratio (S/N) $\\gtrsim$ 150. We classified the source by comparing with published atlases.} {We spectrally classified the source as a {B0.5-1 supergiant} and estimated its interstellar extinction. We compared the extinction  derived from X-ray data with effective interstellar extinction obtained from our data, discussing the absorption component associated with the circumstellar environment. } {} ",
        "introduction": "High mass X-ray binaries (HMXBs) are composed of an early-type massive star and an accreting compact object, a neutron star, or a black hole. The majority of known systems are Be/X-ray binaries (BeXRBs), consisting of a neutron star accreting matter from the circumstellar equatorial disk of a Be star. Most of them are transient, exhibiting short and bright outbursts in the X-ray band. In the second major class of HMXBs, the supergiant X-ray binaries (SGXRBs), the compact object accretes matter from an early supergiant star through its radially outflowing stellar wind. As a consequence, most of SGXRBs are persistent systems, with $L_{X}\\sim 10^{36}$ erg s$^{-1}$. \\\\  The \\emph{INTEGRAL} survey of the Galactic plane and central regions is helping to substantially improve our knowledge of Galactic X-ray binaries \\citep{bird07,bir10}. A large fraction of the newly discovered sources are heavily obscured supergiant massive X-ray binaries  \\citep[first suggested by][]{rev03}, exhibiting much larger column densities ($N_\\mathrm{H}\\gtrsim 10^{23}$ cm$^{-2}$) than  expected along the line of sight \\citep[see][]{kuulkers05}. These sources were missed by previous high-energy missions, whose onboard instruments were sensitive to a softer energy range. Optical counterparts to these obscured sources are also barely detectable because of the high interstellar extinction, $A_{V}$ being in excess of up to $\\sim20$ mag.\\\\  In this context, infrared spectroscopy is an important tool for characterizing these systems \\citep[see also][hereafter referred to as Paper I]{nes08}. With high-energy data, it helps us to identify the HMXB subclass the sources belong to and the mass-transfer process of the system, providing information about the intrinsic physics of the X-ray binary. \\\\   In Paper I, we presented IR spectroscopy of six \\emph{INTEGRAL} sources, classifying their counterparts and estimating both their distance and interstellar extinction. In this paper, we extend the analysis including one more \\emph{INTEGRAL} source, IGR~J16493-4348, located in the direction of the Norma-arm tangent region.\\\\  The source was discovered by \\citet{gre05}. Subsequent RXTE observations by \\citet{mark05} found that the mean spectrum is consistent with a heavily absorbed power law with $N_\\mathrm{H} \\sim 10^{23}$ cm$^{-2}$ and a photon index of 1.4. The measured flux was 1.0, 1.3, and 2.1 $\\times 10^{-11}$ erg cm$^{-2}$ s$^{-1}$ in the 2--10, 10--20, and 20--40 keV energy bands, respectively.  \\citet{kui05} performed Chandra imaging of the field of IGR~J16493--4348 for 4.1 ks. They detected a single point source within the $2'$ error circle of the \\emph{INTEGRAL} source at R.A. =$16^h 49^m 26.92^s$, \\mbox{Dec =$-43^\\circ 49' 8.96''$}, with a 0.6$''$ error in each coordinate. No spectrum could be extracted from the data and \\citet{kui05} noted that previous measurements of IGR~J16493--4348 by RXTE may be contaminated by another X-ray source, 1RXS~J164913.6--435527, located $\\sim$6.7$'$ away.  They also pointed out that no significant source can be found at the position of the Chandra source in the optical DSS maps, indicating strong absorption in the direction of the source. From \\emph{INTEGRAL} data, \\citet{hil08} discarded the previously proposed association with the free radio pulsar PSR~J1649--4349; the best-fit model of the X-ray spectrum obtained by the same authors included an absorbed cut-off powerlaw with $N_\\mathrm{H} = 5.4 \\times 10^{22}$ cm$^{-2}$ and $\\Gamma = 0.6$. From Suzaku data, \\citet{mor09} obtained $N_\\mathrm{H} = 2.6 ^{+0.9}_{-0.8} \\times 10^{23}$  cm$^{-2}$ and $\\Gamma = 2.4$. Although having different parameters, both results are consistent with an accreting neutron star.\\\\  The infrared counterpart to the source was proposed by \\citet{kui05}, who reported a single 2MASS source, 2MASS~J16492695--4349090, compatible with both the Chandra and Swift/XRT positions. \\citet{kui05} observed the source in the $K_\\mathrm{s}$ band and found a magnitude of 12, consistent with the 2MASS magnitude. No optical/IR spectra of the counterpart are available, and the {nature of the system}, although its position and X-ray behavior {suggest} that it is a HMXB system, remains unproven.\\\\  {In the next section, we describe the observations and data reduction. In Sect.~\\ref{results}, we present the obtained spectrum, analyze its features to propose a classification, and calculate the interstellar hydrogen column density. In sect.~\\ref{discussion}, we discuss our results, before concluding.} Preliminary results of our data analysis were published in \\citet{nes08_b}. ",
        "conclusions": "\\label{discussion}   Using $K$-band spectroscopy, we classified the counterpart to IGR J16493--4348  by comparison with published atlases. We found that the system  has a supergiant companion, estimating its type to be {B0.5--1 Ia-Ib}. Combined with information from X-ray data, our results permit us to classify the system as a SGXRB.\\\\  This work allowed us to calculate the extinction from IR data. Previous works have pointed out \\citep[see][]{kuulkers05,cha07} that \\emph{INTEGRAL} {is detecting} a new class of highly obscured supergiant HMXBs. The origin and position (around the compact object only, or enveloping the entire system) of the absorbing material remain unclear, and only multiwavelength studies can help us to address the problem, by distinguishing between the absorption in X-ray and in the IR/optical bands. \\\\  We calculated the effective interstellar extinction $A_{V}$ and converted it into a hydrogen column density of $N_\\mathrm{H} = 2.92\\pm1.96 \\times10^{22}$. This value is compatible with estimates of the weighted average neutral hydrogen density in a cone of radius $1^{\\circ}$ along the line of sight of IGR J16493--4348 \\citep[$\\sim$ 1.4--1.8$\\times10^{22}$ cm$^{-2}$,][]{kal05,dic90}. Our results were compared with the values obtained from X-ray data. If the two derived values were compatible within their corresponding errors, we would have to consider two possible scenarios: either the source of absorption is just the interstellar medium, or there is a contribution from an extensive envelope around the whole binary system, if the extinction is systematically higher by some orders of magnitude than the estimated interstellar value. In contrast, if the reddening (measured from the IR colors) were low compared to the measured $N_\\mathrm{H}$ from X-rays, this would imply that there is an additional source of extinction, which only affects the compact object in which the X-ray emission originates and can be assumed to be the absorbing material around it.\\\\  In our case, the extinction measured at high energy was found to be one order of magnitude higher than that obtained from IR data, which indicates that the material absorbing in the X-ray is concentrated around the neutron star. The same results were found in Paper I for the following systems: IGR J16465--4507, IGR J16479--4514, AX J1841.0--0536, and IGR J19140+0951.\\\\  The so-called highly absorbed IGR sources are usually identified to be those for which the measured $N_\\mathrm{H}$ is $\\gtrsim$$10^{23}$ cm$^{-2}$, \\emph{i.e.}, one to two orders of magnitude higher than the assumed Galactic value of \\mbox{$\\sim10^{22}$ cm$^{-2}$} \\citep{kuulkers05}. As in Paper I, the comparison of X-ray data, which are sensitive to the absorption caused by the environment of the compact object, with infrared data, which indicate in general only the radiation absorbed by the interstellar medium, was a powerful, alternative criterion for identifying this class of highly absorbed sources.  \\\\"
    },
    "1004/1004.2382_arXiv.txt": {
        "abstract": "{} {This letter investigates the transport properties of MHD turbulence induced by the magnetorotational instability at large Reynolds numbers $Re$ when the magnetic Prandtl number $Pm$ is larger than unity.} {Three MHD simulations of the magnetorotational instability (MRI) in the unstratified shearing box with zero net flux are presented. These simulations are performed with the code Zeus and consider the evolution of the rate of angular momentum transport as $Re$ is gradually increased from $3125$ to $12500$ while simultaneously keeping $Pm=4$. To ensure that the small scale features of the flow are well resolved, the resolution varies from $128$ cells per disk scaleheight to $512$ cells per scaleheight. The latter constitutes the highest resolution of an MRI turbulence simulation to date.} {The rate of angular momentum transport, measured using the $\\alpha$ parameter, depends only very weakly on the Reynolds number: $\\alpha$ is found to be about $7 \\times 10^{-3}$ with variations around this mean value bounded by $15\\%$ in all simulations. There is no systematic evolution with $Re$. For the best resolved model, the kinetic energy power spectrum tentatively displays a power-law range with an exponent $-3/2$, while the magnetic energy is found to shift to smaller and smaller scales as the magnetic Reynolds number increases. A couple of different diagnostics both suggest a well--defined injection length of a fraction of a scaleheight.} {The results presented in this letter are consistent with the MRI being able to transport angular momentum efficiently at large Reynolds numbers when $Pm=4$ in unstratified zero net flux shearing boxes.} ",
        "introduction": "\\label{intro}  Angular momentum transport in accretion disks has been an outstanding issue in theoretical astrophysics for decades. To date the most likely mechanism appears to be MHD turbulence driven by the magnetorotational instability \\citep[MRI,][]{balbus&hawley91,balbus&hawley98}. Several numerical simulations have been performed to study its properties. The most popular approach is to work in the local approximation, using the shearing box model, as pioneered by \\citet{hawleyetal95}, \\citet{hawleyetal96} or \\citet{brandenburgetal95}. These early simulations have shown that MRI--powered MHD turbulence is a robust mechanism that transports angular momentum outward. The rate of transport, measured by the famous $\\alpha$ parameter \\citep{shakura&sunyaev73} depends on the field geometry but is always positive, indicating outward flux of angular momentum. The results obtained in the 1990's however obviously suffered from the limited computational resources available at that time. With no mean vertical magnetic field threading the shearing box (a field geometry referred to as {\\it the zero net flux} case), \\citet{fromang&pap07} recently demonstrated with the code Zeus \\citep{hawley&stone95} that it is indeed a problem: $\\alpha$ decreases by a factor of two each time the resolution is doubled. This behavior has since been shown to be very robust as it has been confirmed by simulations performed with codes using different algorithms \\citep{simonetal09,guanetal09}. This result, although it raised the concern that MRI--induced transport could vanish at infinite resolution, was interpreted as an indication that the small scale behavior of the flow is an important ingredient to determine the rate of MRI--induced angular momentum transport: small scale explicit dissipation coefficients, namely viscosity and resistivity, need to be included in the simulations. With such calculations \\citet{lesur&longaretti07} showed that, for a nonzero vertical mean magnetic field, $\\alpha$ rises with the magnetic Prandtl number $Pm$, the ratio of viscosity over resistivity. This result is actually very general: it is independent of the field geometry and was also found for a mean toroidal magnetic field \\citep{simon&hawley09} and in the zero net flux case of interest here \\citep{fromangetal07}. Recently \\citet{simonetal09} measured the {\\it numerical} dissipation properties of the code Athena \\citep{gardiner&stone08,stoneetal08}. They found that an increase in resolution amounts to an increase of the {\\it numerical} Reynolds numbers, while keeping the {\\it effective} magnetic Prandtl number (i.e. the ratio between the numerical viscosity and the numerical resistivity) roughly constant and equal to about two. In light of these results a possible interpretation of the findings of \\citet{fromang&pap07} is that $\\alpha$ is decreasing when the {\\it physical} Reynolds number increases at fixed $Pm$. If unchecked, this decreasing $\\alpha$ would mean that MRI--induced MHD turbulence is ineffective at transporting angular momentum without a mean flux, even in systems that have $Pm$ values higher than unity. Here, high resolution numerical simulations in which $Re$ and $Rm$ are simultaneously increased while keeping their ratio $Pm$ constant are used to examine if this is indeed the case. ",
        "conclusions": "\\label{conclusion_section}  Here zero net flux high resolution numerical simulations of MRI--driven MHD turbulence are used to demonstrate this result: when $Pm$$=$$4$, the dependence of $\\alpha$ on the Reynolds number is very weak. In all models, $\\alpha \\sim 7 \\times 10^{-3}$ to within about $15\\%$. This result unambiguously shows that the decrease of $\\alpha$ with resolution reported by \\citet{fromang&pap07} is a numerical artifact that contains no physical information about the nature of the MHD turbulence in accretion disks. Quite differently, the present simulations are consistent with a nonzero value of $\\alpha$ at infinite Reynolds numbers for a magnetic Prandtl number higher than unity. Note that this weak dependence of $\\alpha$ with $Re$ for $Pm>1$ is also suggested by the data recently reported by \\citet{simon&hawley09} and \\citet{longaretti&lesur10} respectively for a mean azimuthal and vertical magnetic field.  In addition, a number of statistical properties of the turbulence are reported. The kinetic energy power spectrum of the turbulence and the two--points--correlation function of the velocity both suggest a well--defined injection length $l_{inj}$ of a few tens of a scaleheight. For the range of the Reynolds numbers $Re$ that can be probed with current resources, $l_{inj}$ seems to be independent of $Re$. At the highest resolution achieved here, the kinetic energy power spectrum displays a power--law scaling over almost a decade in wavenumber. However, given the limited extent of the power--law range, the precise exponent of this power--law cannot be accurately determined: an exponent of $-3/2$ appears to be consistent with the data, while a $-5/3$ exponent seems too steep. Nevertheless, as suggested in Sect.~\\ref{correlation_sec}, the separation between the forcing and the dissipative scales might still be marginal. This is why a detailed comparison of these exponents with existing MHD turbulence theories \\citep{iroshnikov63,kraichnan65,goldreich&sridhar95} is probably premature at this stage. Higher resolution simulations are definitively needed. Finally, the shape of the magnetic energy power spectrum shows that magnetic energy is mostly located at small scales and shifts to smaller and smaller scales as $Rm$ increases, as expected from small scale dynamo theory \\citep{schekochihinetal02a}. This is consistent with the scenario postulated by \\citet{rinconetal08} of a large scale MRI forcing that generates and coexists with a small scale dynamo."
    },
    "1004/1004.0949_arXiv.txt": {
        "abstract": "We present a catalog of 5324 massive stars in the Small Magellanic Cloud (SMC), with accurate spectral types compiled from the literature, and a photometric catalog for a subset of 3654 of these stars, with the goal of exploring their infrared properties. The photometric catalog consists of stars with infrared counterparts in the {\\it Spitzer}\\, SAGE-SMC survey database, for which we present uniform photometry from $0.3-24$ $\\mu$m in the $UBVIJHK_{s}$+IRAC+MIPS24 bands. We compare the color magnitude diagrams and color-color diagrams to those of stars in the Large Magellanic Cloud (LMC), finding that the brightest infrared sources in the SMC are also the red supergiants, supergiant B[e] (sgB[e]) stars, luminous blue variables, and Wolf-Rayet stars, with the latter exhibiting less infrared excess, the red supergiants being less dusty and the sgB[e] stars being on average less luminous. Among the objects detected at 24~$\\mu$m in the SMC are a few very luminous hypergiants, 4 B-type stars with peculiar, flat spectral energy distributions, and all 3 known luminous blue variables. We detect a distinct Be star sequence, displaced to the red, and suggest a novel method of confirming Be star candidates photometrically. We find a higher fraction of Oe and Be stars among O and early-B stars in our SMC catalog, respectively, when compared to the LMC catalog, and that the SMC Be stars occur at higher luminosities. We estimate mass-loss rates for the red supergiants, confirming the correlation with luminosity even at the metallicity of the SMC.  Finally, we confirm the new class of stars displaying composite A \\& F type spectra, the sgB[e] nature of 2dFS1804 and find the F0 supergiant 2dFS3528 to be a candidate luminous blue variable with cold dust. ",
        "introduction": "\\label{section:intro}  The {\\em Spitzer Space Telescope} Legacy Surveys SAGE \\citep[``Surveying the Agents of a Galaxy's Evolution'',][]{Meixner06} and SAGE-SMC \\citep{Gordon10} have for the first time made possible a comparative study of the infrared properties of massive stars at a range of metallicities, by imaging both the Large and Small Magellanic Clouds (LMC and SMC). In \\citet[][hereafter Paper I]{Bonanos09a}, we presented infrared properties of massive stars in the LMC (at 0.5\\,\\zsun). The motivation for that study was twofold: (a) to use the infrared excesses of massive stars to probe their winds, circumstellar gas and dust, and (b) to provide a template for studies of other, more distant, galaxies. Paper I was the first major compilation of accurate spectral types and multi-band photometry from 0.3$-$24 $\\mu$m for massive stars in any galaxy, increasing by an order of magnitude the number of massive stars for which mid-infrared photometry was available. The recently completed SAGE-SMC survey offers the opportunity to extend the study of infrared properties of massive stars to a metallicity of approximately 0.2\\,\\zsun~\\citep[see e.g.][]{Hunter07}, and enables the investigation of their dependence on metallicity, at least over the range 0.2--0.5\\,\\zsun.  Infrared excess in hot massive stars is primarily due to free free emission from their ionized, line-driven, stellar winds. \\citet{Panagia75} and \\citet{Wright75} first computed the free-free emission from ionized envelopes of hot massive stars, as a function of the mass-loss rate ($\\dot{M}$) and the terminal velocity of the wind ($v_{\\infty}$). The properties of massive stars, and in particular their stellar winds (which affect their evolution) are expected to depend on metallicity ($Z$). For example, \\citet{Mokiem07a} found empirically that mass-loss rates scale as $\\dot{M} \\sim Z^{0.83\\pm 0.16}$, in good agreement with theoretical predictions \\citep{Vink01}. The expectation, therefore, is that the infrared excesses of OB stars in the SMC should be lower than in the LMC, given that $\\dot{M}$ is lower in the SMC\\footnote{The dependence of $v_{\\infty}$ on $Z$ is negligible for this range of $Z$ \\citep[see][]{Leitherer92, Evans04b}.}. Furthermore, there is strong evidence that the fraction of classical Be stars among B-type stars is higher at lower metallicity \\citep[possibly due to faster rotation, as measured by][]{Martayan07b}.  \\citet{Grebel92} were the first to find evidence for this, by showing that the cluster NGC\\,330 in the SMC has the largest fraction of Be stars of any known cluster in the Galaxy, LMC or SMC. More recent spectroscopic surveys \\citep{Martayan10} have reinforced this result. In Paper I, we showed that the Be stars in the LMC are easily discriminated by their mid-infrared colors; we are therefore interested in a comparison with the SMC to quantify the global dependence of the Be star fraction on metallicity.  The incidence of Be/X-ray binaries is also much higher in the SMC than in the LMC \\citep{Liu05}, while the incidence of Wolf-Rayet (WR) stars is much lower; therefore, a comparison of infrared excesses for these objects is also of interest. Finally, there is interest in the metallicity dependence of the mid-infrared colors of red supergiants, which probe circumstellar dust.  Following the same strategy as in Paper I, we compiled a catalog of massive stars, which we cross-matched in the SAGE-SMC database to study their infrared properties. This paper is organized as follows: Section~\\ref{sec:catalog} describes our spectroscopic and photometric catalogs of massive stars in the SMC, Section~\\ref{sec:cmd} presents the resulting color--magnitude and two--color diagrams, and Section \\ref{sec:hotstars} presents the infrared excesses detected in various types of massive stars. Section~\\ref{sec:rsg_mdot} investigates the mass-loss rates in red supergiants, and Section \\ref{sec:summary} summarizes our results. ",
        "conclusions": "\\label{sec:summary}  This paper presents the first catalogs of accurate spectral types and multi-wavelength photometry of massive stars in the SMC, which are used to study their infrared properties. The spectroscopic catalog contains 5324 massive stars, with accurate positions and spectral types compiled from the literature, and constitutes the largest such catalog currently available for any galaxy. The photometric catalog comprises uniform 0.3--24~$\\mu$m photometry in the $UBVIJHK_s$+IRAC+MIPS24 bands for a subset of 3654 stars that were matched in the SAGE-SMC database. The low foreground reddening toward the SMC, and the identical distance of the stars minimize systematic errors due to reddening and enable the investigation of infrared excesses. As in Paper I, we construct CMDs and TCDs, and discuss the position of O, early and late-B stars, WR, LBV, sgB[e], classical Be stars, RSGs, AFG supergiants, and Be/X-ray binaries on them. These diagrams are useful for interpreting infrared photometry of resolved massive stars in nearby galaxies at low metallicity.  A comparison of the infrared colors of the massive stars in the SMC to those of their counterparts in the LMC (presented in Paper I) reveals differences that are due to the different evolution at SMC metallicity. The main results of our study concern the emission line Oe and Be stars, and the RSGs. We detected a clear bimodal distribution of early-B stars, with the redder sequence corresponding to classical Be stars and therefore propose that Be stars (and similarly Oe stars) can be easily discriminated photometrically, by their infrared colors.  We find the fraction of emission line stars in the SMC ($10\\pm2\\%$ for Oe and $27\\pm2\\%$ for Be) to be double that of the LMC ($5\\pm1\\%$ for Oe and $16\\pm2\\%$ for Be), when including the ``photometric'' Oe and Be stars. This is the first time the frequency of Oe stars is determined beyond the Galaxy and at subsolar metallicities. We also find Be stars and Be/X-ray binaries to occur at higher luminosity, sgB[e] stars to be on average less luminous than their counterparts in the LMC (at [3.6]), and WR stars to have smaller excess, all due to the different evolution at the lower metallicity of the SMC.  The infrared colors of the RSGs in the SMC are found in most cases to be consistent with little dust, with only the most luminous sources showing excess emission presumably from circumstellar dust. This is in contrast to the generally dusty RSGs in the LMC, and agrees with the expectation that the dust content in metal-poor RSGs is lower. We find that the mass-loss rates in SMC RSGs correlate positively with luminosity, as in the LMC. Finally, we confirm the astrophysical peculiarity of the composite A \\& F type spectra discovered by \\citet{Evans04}, the sgB[e] nature of 2dFS1804, and find the F0 supergiant 2dFS3528 to be an LBV candidate. This paper thus demonstrates the wealth of information contained in the SAGE-SMC survey, and enables studies of the infrared properties of massive stars as a function of metallicity, in combination with the SAGE survey of the LMC. Studies of particular classes of massive stars and whole massive star populations in nearby galaxies, as outlined in Paper I, are obvious directions for follow-up."
    },
    "1004/1004.3876_arXiv.txt": {
        "abstract": "{} { We present the first results of the ongoing Canada-France Brown Dwarfs Survey-InfraRed, hereafter CFBDSIR, a near infrared extension to the optical wide-field survey CFBDS. Our final objectives are to constrain ultracool atmosphere physics by finding a statistically significant sample of objects cooler than 650K and  to explore the ultracool brown dwarf mass function building on a well-defined sample of such objects.} {We identify candidates in CFHT/WIRCam $J$ and CFHT/MegaCam $z'$ images using optimised psf-fitting, and follow them up with pointed, near-infrared imaging with SOFI at the NTT. We finally obtain low-resolution spectroscopy of the coolest candidates to characterise their atmospheric physics.} { We have so far analysed and followed up all candidates on the first 66 square degrees of the 335 square degree survey. We identified 55 T-dwarfs candidates with $z'-J>3.5$ and have confirmed six of them as T-dwarfs, including 3 that are strong later-than-T8 candidates, based on their far-red  and NIR colours. We also present here the NIR spectra of one of these ultracool dwarfs, CFBDSIR1458+1013, which confirms it as one of the coolest brown dwarf known, possibly in the 550-600K temperature range. } {From the completed survey we expect to discover 10 to 15 dwarfs later than T8, more than doubling the known number of such objects. This will enable detailed studies of their extreme atmospheric properties and provide a stronger statistical basis for studies of their luminosity function. }  \\date{} ",
        "introduction": "The significant improvement in detector technology, data storage, and analysis facilities in the past decade has made it possible to carry out wide-field surveys covering a large fraction of the sky instead of targeting specific sources. The wealth of data from these surveys necessitate a complex dedicated computer analysis to single out relevant scientific information. These surveys, such as DENIS \\citep{Epchtein.1997}, SDSS \\citep{York.2000}, 2MASS \\citep{Skrutskie.2006}, UKIDSS \\citep{Lawrence.2007}, and CFBDS \\citep{Delorme.2008b} contain hundreds of millions of astrophysical sources and led to many advances in various fields, notably to identify extremely rare objects and build robust statistical studies.  The survey we are presenting here, the Canada-France Brown Dwarfs Survey-InfraRed aims at finding ultracool brown dwarfs (T$_{eff}<$ 650K) of which only 6 are currently published by \\citet{Warren.2007,Delorme.2008a,Burningham.2008,Burningham.2009,Burningham.2010sub}. \\citet{Lucas.2010sub} have very recently identified a probably even cooler object. These rare objects are in many ways the intermediate ``missing link\" between the cold atmospheres of the Solar System's giant planets and cool stellar atmospheres. The physics and chemistry of their atmospheres, dominated by broad molecular absorption bands, are very planetary-like \\citep[see][for instance]{Kulkarni.1997} and the cool brown dwarfs spectra are the key to constraining planetary and stellar atmosphere models.  Nowadays, the Teff$<$700K atmosphere temperature regime is troublesome for modellers. A few ultracool late T brown dwarfs have now been discovered  with effective temperatures below 650K. These discoveries step into unexplored territory and a new generation of models is emerging. It is facing several difficulties. 1) Out-of-equilibrium chemistry plays an important role, resulting for instance in NH$_3$ being less prevalent than expected \\citep{Cushing.2006}. 2) Fine details of convection control both the L/T transition and the dredge up of hot chemical species in late T atmospheres. 3) Water cloud formation and dust nucleation play important roles. 4) Line opacities of several molecules, in particular NH$_3$ and, to a lesser extent, CH$_4$ are unknown and cause important spectroscopic feature mismatches.  As a good example of the need for refined models, \\citet{Burningham.2009} have recently determined that a T8.5 dwarf companion to an M star was actually $\\sim$15\\% cooler than model fitting would have predicted and \\citet{Dupuy.2009} and  \\citet{Liu.2008} reached similar conclusions when studying brown dwarf binaries with dynamical masses.  Under those circumstances, observations are key to the development of the models. Only five brown dwarfs with temperatures below 650K (T8.5) are currently known from recent discoveries by UKIDSS and CFBDS, and this small number prevents discerning general trends from individual peculiarities. \\citet{Kirkpatrick.1999} and \\citet{Burgasser.2002} could rely on samples of 20-25 objects to respectively define the L and T spectral types. In this article we present the CFBDSIR, a near infrared (hereafter NIR) extension to the CFBDS that will provide a WIRCAM \\citep{Puget.2004} $J$-band coverage down to $J_{vega}=20.0$ for 10$\\sigma$ detections atop 335 square degrees of CFBDS MegaCam \\citep{Boulade.2003proc} $z'$-band imaging with a 5$\\sigma$ detection limit of $z'_{AB}$=23.25-24.05. All optical magnitudes presented in this article are AB magnitudes, while all NIR magnitudes are Vega magnitudes. When the CFBDSIR is complete, we hope to achieve a threefold increase in sample size to 15-20 characterised ultracool brown dwarfs and possibly find a few substellar objects significantly cooler than 500K of which none is known yet outside the solar system. This will define general trends and dispersions around them, permitting the study of ultracool dwarfs not only as individual interesting objects, but as a population. This will help define the T/Y spectral transition that is expected to occur in this temperature range. In section 1 we present the rationale of this new wide-field survey and the observations at its core. In section 2 we describe the data reduction and the data analysis methods we used to identify ultracool brown dwarfs candidates. Finally, we present the spectra and the photometry of the first ultracool brown dwarf identified with CFBDSIR in section 3. ",
        "conclusions": "We have described CFBDSIR, a new NIR survey dedicated at finding ultracool brown dwarfs and using WIRCam camera on the the CFHT. Complementing existing deep far-red data by new $J$-band observations, we select brown dwarfs candidates on their very red $z'-J$ colour. A robust PSF analysis allows us to derive reliable colours and to distinguish point-source-like brown dwarfs from most contaminants. The candidates are then confirmed by follow-up pointed NIR observations in $J$-band and confirmed ultracool brown dwarfs are imaged in $H$ and $K_{\\rm s}$ bands. We used these photometric measurements to identify several new brown dwarfs, including 3 objects likely as cool as and possibly even cooler than any published brown dwarfs.  We  presented CFBSIR1458, the first CFBDSIR ultracool brown dwarf confirmed by spectroscopy. The analysis of its $H$-band spectra, though at relatively low-signal-to-noise, robustly confirms it as later than T8 spectral type and hints at a temperature in the 550-600K range, so among the coolest brown dwarfs discovered. When the 335 square degree survey is completed, we expect to discover a sample of 10 to 15  ultracool brown dwarfs, more than doubling the currently known population of later than T8 objects and enabling study of them as a population rather than extreme individual objects. This will put strong constraints on cool stellar and planetary atmosphere, and with additional mid-infrared follow-up, will help to define the selection criteria for the upcoming WISE survey."
    },
    "1004/1004.1568_arXiv.txt": {
        "abstract": "The temporal and spectral analysis of 9 bright X-ray flares out of a sample of 113 flares observed by \\emph{Swift} reveals that the flare phenomenology is strictly analogous to the prompt $\\gamma$-ray emission: high energy flare profiles rise faster, decay faster and peak before the low energy emission. However, flares and prompt pulses differ in one crucial aspect: flares evolve with time. As time proceeds flares become wider, with larger peak lag, lower luminosities and softer emission. The flare spectral peak energy $E_{\\rm{p,i}}$ evolves to lower values following an exponential decay which tracks the decay of the flare flux. The two flares with best statistics show higher than expected isotropic energy $E_{\\rm{iso}}$ and peak luminosity $L_{\\rm{p,iso}}$ when compared to the $E_{\\rm{p,i}}-E_{\\rm{iso}}$ and $E_{\\rm{p,i}}-L_{\\rm{iso}}$ prompt correlations. $E_{\\rm{p,i}}$ is found to correlate with $L_{\\rm{iso}}$ within single flares, giving rise to a time resolved $E_{\\rm{p,i}}(t)-L_{\\rm{iso}}(t)$. Like prompt pulses, flares define a lag-luminosity relation: $L_{\\rm{p,iso}}^{0.3-10 \\,\\rm{keV}}\\propto t_{\\rm{lag}}^{-0.95\\pm0.23}$. The lag-luminosity is proven to be a fundamental law extending $\\sim$5 decades in time and $\\sim$5 in energy. Moreover, this is direct evidence that GRB X-ray flares and prompt gamma-ray pulses are produced by the same mechanism. Finally we establish a flare- afterglow morphology connection: flares are preferentially detected superimposed to one-break or canonical X-ray afterglows. ",
        "introduction": "\\label{Sec:Introduction} The high temporal variability was one of the first properties to be attributed to the Gamma-ray burst (GRB) prompt emission in the $\\gamma$-ray energy band (\\citealt{Klebesadel73}). The advent of \\emph{Swift} (\\citealt{Gehrels04}) revealed that a highly variable emission characterises also the early time X-ray afterglows in the form of erratic flares. This established the temporal variability as one of the key features in interpreting the GRB phenomena.  GRB\\,050502B and the X-ray flash 050406 (\\citealt{Falcone06}; \\citealt{Romano06b}; \\citealt{Burrows05b}) provided the first examples of dramatic flaring activity superimposed to a smooth decay: in particular, GRB\\,050502B demonstrated that flares can be considerably energetic, with a 0.3-10 keV energy release comparable to the observed prompt fluence in the 15-150 keV band. Thanks to the rapid re-pointing \\emph{Swift} capability, it was later shown that flares are a common feature of the early X-ray afterglows, being present in the $\\sim 33\\%$ of X-ray light-curves (\\citealt{Chincarini07}, hereafter C07; \\citealt{Falcone07}, hereafter F07). On the contrary, a convincing optical flare, counterpart to a detected X-ray flare is still lacking, suggesting that the detected optical afterglow contemporaneous to the high-energy flares is dominated by a different emission component (see e.g. GRB\\,060904B,  \\citealt{Klotz08} but see also \\citealt{Greiner09} where an optical flare was probably detected but, unfortunately, contemporaneous X-ray coverage is lacking).  Based on the  temporal and spectral study of a statistical sample of X-ray flares within GRBs, C07 and F07 showed that the flares share common properties and that the flare phenomenology can be described using averaged properties (see C07 and F07 and references therein): \\begin{itemize} \\item The same GRB can show multiple flares (see e.g. GRB\\,051117A which contains a minimum of 11 structures in the first 1 ks of observation); \\item The underlying continuum is consistent with having the same slope before and after the flare, suggesting that  flares constitute a separate component in addition to the observed continuum; \\item Each flare determines a flux enhancement evaluated at the peak time $\\Delta F/ F$ between $\\sim1$ and $\\sim1000$, with a fluence that competes in some cases (e.g. GRB\\,050502B) with the prompt $\\gamma$-ray fluence. The average flare fluence is $\\sim 10$\\% the 15-150 keV prompt fluence; \\item Flares are sharp structures, with $\\Delta t/t \\sim 0.1$, a fast rise and a slower decay; \\item Each flare determines a hardening during the rise time and a softening during the decay time (F07), reminiscent of the prompt emission (e.g. \\citealt{Ford95}): the result is a hardness ratio curve that mimics the flare profile (see e.g. GRB\\,051117A, \\citealt{Goad07}, their figure 9). In this sense flares are spectrally harder than the underlying continuum; \\item The spectrum of a consistent fraction of flares is better fitted by a Band (\\citealt{Band93}) model, similarly to prompt emission pulses (see e.g. \\citealt{Kaneko06}). The flare spectral peak energy is likely to be in the soft X-ray range (a few keV). The spectrum evolves with time as testified by the hardness ratio curve and by accurate spectral modelling. During the decay time a clear softening is detected (e.g. \\citealt{Krimm07}; \\citealt{Godet07}); \\item There is no correlation between the number of flares and the number of prompt emission pulses; \\item The vast majority of flares are concentrated in the first 1 ks after trigger.  However, late-time flares ($t_{\\rm{peak}}\\sim 10^5-10^6$ s) are present as well: flares are not confined to the steep decay phase, but can happen during the plateau and  the normal decay phases. Their temporal properties are consistent with those of early flares (\\citealt{Curran08}), even if their lower brightness prevents a detailed comparison with the entire set of early time flare properties (this is especially true from the spectral point of view); \\item Flares happen both in low-z and high-z environments: the record holder GRB\\,090423 at z$\\sim8.2$ (\\citealt{Salvaterra09}; \\citealt{Tanvir09}) shows a prominent flare with standard properties when compared to the sample of X-ray flares of \\cite{Chincarini10} (C10, hereafter); \\item Flares  have been detected both in hard and soft events such as X-Ray Flashes  (e.g. XRF\\,050406); \\item Variability has also been  detected in the X-ray afterglows of \\emph{short} GRBs (GRB with a prompt emission duration $T_{90}<2$ s, \\citealt{Kouveliotou93}). However, given the lower brightness associated to these events it is still unclear if  what is currently identified as a short GRB flare emission, quantitatively shares the very same properties as the population of flares detected in \\emph{long} GRBs. GRB\\,050724  (\\citealt{Barthelmy05b}) constitutes a good example of short GRB with late-time variability. \\end{itemize}  From the systematic study of 113 flares in the XRT 0.3-10 keV energy band, as well as in 4 sub-energy bands, C10 demonstrated that: \\begin{itemize} \\item Flares are asymmetric with an average asymmetry parameter similar to the prompt emission value; no flare is found rising slower than decaying; \\item The flare width evolves linearly with time $w\\propto 0.2\\, t_{\\rm{peak}}$. This is a key point which clearly distinguishes the flares from the prompt emission, where no evolution of the pulse-width has ever been found (see e.g. \\citealt{Ramirez00}); \\item The width evolution is the result of the linear evolution of both the rise and the decay times: $t_{\\rm{r}}\\propto0.06\\, t_{\\rm{peak}}$; $t_{\\rm{d}}\\propto 0.14\\,t_{\\rm{peak}}$. \\item The asymmetry does not evolve with time. Instead the rise over decay time ratio is constant with time, implying that both time scales are stretched of the same factor. Furthermore $t_{\\rm{d}}\\sim2\\,t_{\\rm{r}}$. Flares are \\emph{self-similar} in time. \\item At high energy the flares are sharper with shorter duration:  $w\\propto E^{-0.5}$. Prompt pulses share the same property, with a similar dependence on the energy band (\\citealt{Fenimore95}; \\citealt{Norris96}); \\item The flare peak luminosity decreases with time. Accounting for the sample variance the best fit relation reads: $L_{\\rm{peak}}\\propto t_{\\rm{peak}} ^{-2.7\\pm0.5}$. The average flare luminosity declines as a power-law in time $<L>\\propto t^{-1.5}$ (\\citealt{Lazzati08}); \\item The isotropic 0.3-10 keV flare energy distribution is a log normal peaked at $\\sim 10^{51}$ erg (this can be viewed as the typical flare isotropic energy). The distribution further shows a hint of bimodality. \\item In multiple-flare GRBs, the flares follow a softening trend which causes later time flares to be softer and softer: the emitting mechanism keeps track of the previous episodes of emission. \\end{itemize}  Starting from these pieces of evidence, C07, F07 and C10 concluded that in spite of the softness and width evolution, the X-ray flares and the prompt pulses are likely to share a common origin. However  important questions are still to be addressed: do flares follow the entire set of temporal and spectral relations found from the analysis of prompt emission pulses? In particular: is it possible to define a flare peak-lag? Do flares follow a lag-luminosity relation? Is the lag-time linked to other temporal properties (e.g. the flare duration, asymmetry, etc.)? What can be said about the pulse start conjecture for flares (\\citealt{Hakkila09})? Is it possible to quantify the evolution of the flare temporal properties as a function of the energy band? Do the rise and decay times evolve differently with energy band? Is it possible to track and quantify the evolution of the spectral peak energy $E_{\\rm p}$ during the flare emission? Is there any connection between the temporal and spectral properties of the flares? What is the position and the track of the flares in the % $E_{\\rm{p}}-L_{\\rm{iso}}$ (spectral peak energy-isotropic luminosity) plane? Is there any link between the flares and the underlying X-ray afterglow morphology?  This set of still open questions constitutes the major motivation for undertaking the present investigation. The primary goal of this paper is to model the X-ray flare profiles and constrain their evolution with energy to obtain parameters that uniquely qualify the shape and spectrum of the flares and compare those values to the well known signatures of the prompt emission pulses. Since this is often difficult because of low statistics, overlap or non trivial estimate of the underlying continuum, the present work concentrates on  very bright and isolated flares for which the underlying continuum does not play a major role neither from the temporal point of view nor from the spectral point of view. This work is observationally driven: a critical review of theoretical models in the light of the present results is in preparation.  This work is organised as follows: the sample selection and data reduction are described in Sect. \\ref{Sec:datared}. In Sect. \\ref{Sec:tempan} we perform the flare temporal analysis, while to the spectral properties of the sample is dedicated Sect. \\ref{Sec:spec}. The discussion follows (Sect. \\ref{Sec:disc}). Conclusions are drawn in Sect. \\ref{Sec:conc}.   The phenomenology of the  different GRBs is presented in the observer frame unless otherwise stated.  The 0.3-10 keV energy band is adopted unless specified. The zero time is assumed to be the trigger time. The convention $F(\\nu,t)\\propto \\nu^{-\\beta}t^{-\\alpha}$ is followed, where $\\beta$ is the spectral index, related to the photon index $\\Gamma$ by $\\Gamma=\\beta+1$. All the quoted uncertainties are given at 68\\% confidence level (c.l.): a warning is added if it is not the case. Standard cosmological quantities have been adopted: $H_{0}=70\\,\\rm{Km\\,s^{-1}\\,Mpc^{-1}}$, $\\Omega_{\\Lambda }=0.7$, $\\Omega_{\\rm{M}}=0.3$.\\\\ ",
        "conclusions": "\\label{Sec:conc} \\begin{figure} \\vskip -0.0 true cm \\centering \\includegraphics[scale=0.43]{flare_paradigma.eps} \\caption{Flare paradigm. Dashed curves: high energy profile. Solid line: low energy profile. High energy profiles rise faster, decay faster and peak before the  low energy emission. As time proceeds, flares becomes wider, with higher peak lag,  lower peak luminosities and softer emission. However, at the 0th level of approximation the self-similarity is preserved both in energy and in time (with $t_r/t_d\\sim0.5$). We stress that this is an \\emph{average} behaviour. The thick solid arrow underlines the dependence of the flare properties on the \\emph{time}, in strong opposition to prompt pulses. } \\label{Fig:flareparadigma} \\end{figure}  The GRB X-ray flare phenomenology shows a list of properties strictly analogous to the gamma-ray prompt emission (see \\citealt{Hakkila08} for a recent study on individual prompt pulses). However, it differs in one crucial aspect: flares \\emph{evolve} with \\emph{time}\\footnote{We parenthetically note that this difference could be due to the fact that the relevant time for a peak (prompt pulse or flare) is the time since the shell ejection in a standard internal shock scenario, and not the time since trigger. While these two time scales are completely unrelated for the prompt pulses, they are more closely related for flares (since the time is large the difference between the two becomes negligible). See \\cite{Willingale10} and Willingale et al., in prep. for further details.}. Flares evolve with time to lower peak intensities, larger widths, larger lags and softer emission. Fig. \\ref{Fig:flareparadigma} illustrates the flare paradigm. As the prompt emission pulses, flares have correlated properties: short-lag flares have shorter duration, are more luminous and harder than long-lag flares. In particular the following properties add  to the lists of Sec. \\ref{Sec:Introduction}: \\begin{itemize} \\item Flares define a lag-luminosity relation: \\\\$L_{\\rm{p,iso}}^{\\rm{0.3-10\\,keV}}\\propto t_{\\rm{lag}}^{-0.95\\pm0.23}$. The best fit slope is remarkably consistent with the prompt findings  (see \\citealt{Ukwatta10} and references therein for a recent study); \\item The lag is found to correlate with the flares width while it is inversely correlated to the flare spectral hardness; \\item The flares temporal profiles in different band-passes are only a nearly exact time stretched version of one another: the rise and decay times evolve following slightly different power-laws in energy: $t_r\\propto E^{-\\alpha_{tr}}$, $t_d\\propto E^{-\\alpha_{td}}$ with $\\alpha_{tr}>\\alpha_{td}$. The result is that flares are on average more asymmetric at high energy; \\item The rate of evolution of a flare profile in different band-passes anti-correlates with its spectral hardness: the harder the flare the lower is the rate of evolution from one energy band to the other; \\item The flare spectral peak energy $E_{p}(t)$ evolves to lower values following an exponential decay which tracks the decay of the flare flux. The detected evolution is faster and inconsistent with the $\\sim t^{-1}$ behaviour even when the zero time is re-set to the beginning of the flare emission; \\item The two flares with best statistics show higher than expected $E_{\\rm{iso}}$ and $L_{\\rm{p,iso}}$ values when compared to  the $E_{\\rm{p,i}}-E_{\\rm{iso}}$ and $E_{\\rm{p,i}}-L_{\\rm{iso}}$ prompt correlations; \\item The rest frame spectral peak energy $E_{\\rm{p,i}}$ correlates with the isotropic luminosity $L_{\\rm{iso}}$ within single flares, giving rise to a time resolved  $E_{\\rm{p,i}}-L_{\\rm{iso}}$ correlation; \\item The flare emission preferentially builds up at lower energies: flares do not seem to be consistent with the Pulse Start Conjecture; \\item Among the different types of X-ray afterglow light-curves, the simple power-law afterglows are under-represented in the flare sample. Flares are preferentially detected superimposed to one-break or canonical light-curves. \\end{itemize} The strict analogy between the prompt pulses and flares phenomenology strongly suggests a common origin of the two phenomena."
    },
    "1004/1004.4451_arXiv.txt": {
        "abstract": "An analysis of the fluorine abundance in Galactic AGB carbon stars {\\bf (24 N-type, 5 SC-type and 5 J-type)} is presented. This study uses the state-of-the-art carbon rich atmosphere models and improved atomic and molecular line lists in the $2.3~\\mu$m region. F abundances significantly lower are obtained in comparison to previous study in the literature. The main reason of this difference is due to molecular blends. In the case of carbon stars of SC-type, differences in the model atmospheres are also relevant. The new F enhancements are now in agreement with the most recent theoretical nucleosynthesis models in low-mass AGB stars, solving the long standing problem of F in Galactic AGB stars. Nevertheless, some SC-type carbon stars still show larger F abundances than predicted by stellar models. The possibility that these stars are of larger mass is briefly discussed. ",
        "introduction": "The first  observational evidence of $^{19}$F stellar nucleosynthesis was reported  by Jorissen, Smith  \\& Lambert  (1992, hereafter JSL). These authors derived F  enhancements up to a factor 50 solar in a sample of Galactic AGB stars, and found a  correlation  between  this  enhancement and the  C/O  ratio. Since the C/O  is expected to  increase as a consequence  of third dredge  up  (TDU)  episodes  during  the  AGB  phase  (e.g.  Busso  et al. 1999), this was interpreted as a clear evidence of F production in these stars. Further observational  evidence of such a production exists from studies of post-AGB stars (Werner, Rauch \\& Kruk 2005)   and  planetary  nebulae   (Otsuka  et al. 2008).  Other  sites for F  production have also been proposed: Wolf-Rayet stars (Meynet \\& Arnould 2000) and neutrino spallation  in core-collapse supernovae (Woosley \\& Haxton 1988). However, the role of these sources in the F budget is  still uncertain  (Cunha et al.  2003; Palacios,  Arnould \\& Meynet 2005).  Nevertheless,  from Galactic chemical evolution models, Renda et al. (2004) concluded that all three of these three sources are required to explain the observed Galactic evolution of F, as deduced from  abundance determinations in  field stars  (Cunha \\&  Smith 2005; Cunha, Smith \\& Gibson 2008), although,  only in AGB stars there exist an observational confirmation that F production is an ongoing process.  Fluorine can be produced in AGB stars from the nuclear chain $^{14}N(\\alpha,\\gamma)^{18}F(\\beta^+)^{18}O(p,\\alpha)^{15}N\\\\(\\alpha,\\gamma)^{19}F$; where  protons are mainly  provided by $^{14}N(n,p)^{14}C$  and neutrons by  $^{13}C(\\alpha,n)^{16}O$. However, the theoretical attempts made to explain the JSL results are  unsatisfactory. The  problem  is that current AGB models fail  to explain the highest F enhancements  found in the C-rich objects of the JSL sample (Goriely \\& Mowlavi 2000; Lugaro et al.  2004).  This  discrepancy has  led  to  a  deep revision  of  the uncertainties  in   the  nuclear   reaction  rates  involved   in  the synthesis of F in AGB stars (Lugaro et al. 2004; Stancliffe et  al.  2005),  to argue  for  alternative nuclear  chains and/or  to propose  the  existence   of  non-standard  mixing/burning  processes, similar to  those commonly used to  explain some of  the isotopic anomalies found in dust grains formed in the envelopes of AGB stars  (e.g. Nollett  et al.  2003; Busso  et al.  2007).  However, no solution has been found up to  date, leaving the subject of the origin of this element open.  Very  recently,  Abia et  al.  (2009,  hereafter  Paper I)  derived  F abundances from  VLT spectra in  three Galactic AGB C-stars  in common with  the JSL  sample,  namely: AQ  Sgr,  TX Psc  (N-type)  and R  Scl (J-type), by using mainly the  HF R9 line at $\\lambda 2.3358~\\mu$m. In that work, we showed that this HF  line is almost free of blends in AGB C-stars  and thus is probably  the  best tool  for F  abundance determinations  in  these  stars   (see also Uttenthaler et al. 2008). For these three stars, we derived  F abundances  $\\sim 0.7$  dex lower  in average  than those  obtained by JSL. We ascribed this difference to molecular blends (mainly of C$_2$ and CN). These new F  abundances  are in   better  agreement  with the  most  recent theoretical nucleosynthesis  predictions in low-mass  ($<3$ M$_\\odot$) TP-AGB models  (Cristallo et al.  2009). Motivated by this result, we have reanalysed  the  F abundances  in the  whole C-stars sample studied by JSL using the same  tools as in Paper I.  In this Letter, we  confirm the results found in Paper I and show  that the new F abundances nicely agree with the theoretical  predictions of  low-mass TP-AGB  models.  This apparently solves  the long  standing  problem  of F  enhancements  in AGB  stars without requiring  the existence  of any extra  mixing/burning process. ",
        "conclusions": "The main consequence of the new  F abundances is that the large [F/Fe] (or [F/O]) ratios (up to 1.8 dex) found by JSL in Galactic  AGB C-stars are  systematically reduced.  The  largest F enhancements  are  now close to  $\\sim  1$  dex.  These  enhancements  can  be accounted  for by  current low-mass  TP-AGB nucleosynthesis  models of solar metallicity, as  we will show below. As noted  in \\S 1, during the  ascension along the  AGB,  fresh carbon  is mixed within the envelope due to TDU episodes. Eventually,  an O-rich AGB star becomes  a  C-star  when  the   C/O  ratio  in  the  envelope  exceeds 1. Similarly, F is also expected  to increase in the envelope during the  AGB phase, thus  a fluorine vs. carbon correlation should  exist. Figure  2 shows  the observed relationship derived  in this study.  We have also included in  this figure the  intrinsic\\footnote{Intrinsic  AGB   stars  have  their  envelopes polluted by nucleosynthesis products made {\\it in situ}. Extrinsic  stars, on the contrary, own  their  chemical peculiarities to a  mass  transfer  episode  in a  binary system.  All  the  O-rich  stars  shown in  Fig.  2  show  $^{99}$Tc ($\\tau_{1/2}\\sim  2\\times 10^5$  yr)  in their  envelopes (Smith  \\& Lambert  1990), revealing their intrinsic nature.}  O-rich  AGB stars studied  by  JSL (not  analysed  here, open circles). Excluding  the J-type\\footnote{The origin  of J-type  C-stars is still  unknown. They are  slightly less luminous  than N-type  C-stars and  show chemical peculiarities  rather  different  with respect to these stars  (e.g.  Abia  \\&  Isern 2000). It has  been suggested that they might  be the descendants of the early R-type stars, however  this has been questioned recently (Zamora 2009).}  stars (triangles), a clear increase of the  F abundance with the C  abundance can be seen.  This behaviour is  well reproduced by theoretical AGB models. Lines in Fig. 2 show the predicted F and C content in the envelope (Cristallo et al. 2010) for a  1.5 M$_\\odot$ with different metallicities (continuous  lines) and  2  M$_\\odot$ model (dashed  line) with solar metallicity. These metallicities match those of the stellar sample ($-0.5\\leq$[Fe/H]$\\leq  0.1$). The theoretical curves start  with an envelope composition as determined by the first dredge-up and end at the last TDU episode. On the other hand, Fig. 2 seems to indicate that O-rich stars and SC stars (squares) have F abundances systematically larger than N-type stars (dots) for a given C abundance. Similarly to the intrinsic C-stars, also in the case  of the  O-rich stars this may be simply due to a systematic error  affecting the  analysis:  in fact, for stars of spectral types  K and  M (with [Fe/H]$\\sim  0.0$), JSL derived an average F abundance $\\sim0.13$ dex higher (4.69) than the solar F abundance adopted here (4.56). These K and M stars are sub-giants or RGB stars, thus they are not expected to present F enhancements. Excluding that the Sun might be anomalous with respect to nearby stars of similar metallicity (F determinations in unevolved dwarf stars are needed in this sense), blends may also play a role in the analysis of F lines in O-rich AGB stars. Decreasing the  F abundances by 0.13 dex  in them, the agreement with theoretical predictions would be definitely better. However, it does not apply to the case of the star BD$+48^o1187$ (at log $\\epsilon$(F)$\\sim  5.5$ in Fig. 2); the large F enhancement in this star can be only obtained in metal-poor AGB models, but this  star has [Fe/H]$=0.07$.  Something similar  seems to  occur with SC-type stars. SC  stars are AGB stars with C/O$\\approx 1$. According to the accepted chemical (spectral) evolution along the AGB (M$\\rightarrow$MS$\\rightarrow$S$\\rightarrow$SC$\\rightarrow$N), these stars should  present  equal-or-slightly lower F enhancements than  N-type stars. However, some of the studied SC stars clearly deviate from this sequence, showing larger F abundances with respect to N stars. This is reinforced by WZ Cas\\footnote{This star  has also been classified  as J-type.} ([Fe/H]$\\sim  0.0$, at the  right upper corner  of  Fig.  2), which shows a  huge  F enhancement. We remind, however, that this  star is indeed peculiar: it  is one of the few  super   Li-rich  C-stars  known,  also presenting   very  low $^{12}$C/$^{13}$C  (4.5) and  $^{16}$O/$^{17}$O (400)  ratios.  The fact that most  of the SC-stars studied here show low or no $s-$element enhancements  (see below), makes the evolutionary status of these stars very uncertain.  The connection between the F and the $s$-process is less straightforward. There are two different contributions to the F production in low mass AGB stars. The first comes from the $^{15}$N production in the radiative $^{13}$C pocket, the site where the s-process main component is built up: when, during the interpulse period,  the $^{13}C(\\alpha,n)^{16}O$ is activated within the pocket, some protons are released by the main poisoning reaction,  $^{14}N(n,p)^{14}C$, thus producing $^{15}$N via $^{18}O(p,\\alpha)^{15}N$. The second contribution involves the same reaction chain, but it is activated when the $^{13}$C left by the advancing H-burning shell is engulfed into the convective zone generated by a thermal pulse. In the latter case, the correlation with the $s-$process is less stringent, because the resulting neutron flux is not enough to give rise to a sizable production of elements beyond iron. At nearly solar metallicity, the latter contribution accounts for $\\sim$70\\% of the total F production {\\bf(e.g. Cristallo 2006)}. Nevertheless, a correlation between the F and the $s-$element overabundance is in any case expected if a large enough $^{13}$C pocket forms after each dredge up episode. Figure  3 shows  the  new [F/Fe] ratios  in  C-stars vs.  the observed average $s-$element enhancement {\\bf(from Abia et al. 2002; Zamora 2009 or derived in this work, see \\S 1 )}.  J-type stars are not shown in this figure as they do not  have $s-$element enhancements (see  Table 1;  Abia \\&  Isern  2000). Similarly  to Fig.  2, we  have included  the  intrinsic O-rich  stars  analysed  by  JSL taken  their $s-$element content from Smith \\& Lambert (1990). It is clear that the F and  $s-$nuclei overabundances  correlate and  that this correlation  is  nicely reproduced by theoretical TP-AGB stellar  models  (Cristallo et  al. 2010).  The theoretical predictions in Fig. 3 correspond to 1.5, 2 and 3 M$_\\odot$ TP-AGB  models   with  Z$=0.008$  ([Fe/H]$\\sim   -0.25$,  the  average metallicity of the stars analysed  here) starting from the 1$^{st}$ TP pulse. The number  of TPs calculated until the  occurrence of the last TDU is indicated in the figure for each  model. Models with a metallicity within $\\sim \\pm 0.2$ dex  of this value do not significantly differ from those in  Figure 3. In addition, with the combination of models with different stellar mass and metallicity, it is possible not only to reproduce the new observed [F/Fe]  vs. $s-$element relationship, but also the  detailed  F  and  $s-$element  abundance pattern  observed  in  a particular star (see e.g. Fig. 3 in  Paper I). Note that this is not possible for most of the stars when adopting the F determinations of JSL.  As already noted, SC-type stars (except GP  Ori) show larger F abundances than predicted by low-mass AGB models with the same $s-$element  enhancement. A possible explanation could be that the SC stars are more massive (in average) than the bulk of the N-type C-stars. From AGB models with mass of about 3-5 M$_\\odot$, we expect smaller $^{13}$C pockets and weaker $s-$element surface enrichments, even if the number of TPs and TDU episodes is larger. It occurs because the larger the core mass is, the smaller the He-rich intershell is with a steeper pressure gradient. Notwithstanding, the $^{13}$C left in the ashes of the H-shell still provide a substantial contribution to the fluorine production. Recent studies of the luminosity function of Galactic C-stars (Guandalini 2008) support such a hypothesis. They indicate that the SC-type stars are among the most luminous AGB C-stars. Once again, the case of WZ Cas is particularly extreme. This  star shows a huge F abundance ([F/Fe]$=+1.15$) with almost no $s-$element enhancement. Its large Li abundance might be interpreted as a consequence of the hot bottom burning (HBB) mechanism which operates only for M$\\geq 5$ M$_\\odot$. {\\bf However, the operation of the HBB would be at odds with the presence of F in the envelope since $^{19}F(p,\\alpha)^{16}$O might destroy $^{19}$F}. Detailed abundance studies (including  F) in a larger sample of SC-stars is needed to confirm this hypothesis."
    },
    "1004/1004.1797_arXiv.txt": {
        "abstract": "We construct Super-Eddington Slim Disks models around both stellar and super-massive black holes by allowing the formation of a porous layer with a reduced effective opacity. We show that at high accretion rates, the inner part of the disks become radiation pressure dominated. However, unlike the standard scenario in which the disks become thick, super-Eddington disks remain slim. In addition, they accelerate a significant wind with a ``thick disk\" geometry. We show that above about 1.5 times the standard critical mass accretion rate (needed to release the Eddington luminosity), the net luminosity released is above Eddington. At above about 5 times the standard critical rate, the central BH accretes more than the Eddington accretion rate. Above about $20 \\dot{m}_{crit}$, the disk remains slim but the wind becomes spherical, and the present model breaks down. ",
        "introduction": "The accretion of matter onto compact objects can often be described using the run-of-the-mill thin accretion disk model of \\citet[][S\\&S]{Shakura1973}. Because the accretion disk is optically thick, matter can radiate the potential energy it dissipates and remain cold, thus forming a geometriclaly ``thin\" disk. The turbulent viscosity responsible for the dissipation is often described through the standard $\\alpha$-model, and it is also responsible for the transport of angular momentum outwards.  The S\\&S thin disk model applies to a wide range of conditions found in nature, however, when one of the underlying assumptions break down, so does the model.  At sufficiently low accretion rates, the disk becomes optically thin and it cannot radiate the energy dissipated. This energy is therefore advected with the flow, forming the so called Advection Dominated Accretion Flow \\citep{Ichi:1977,NarayanYi,Abra:95}. Because of the high temperatures and pressures, ADAFs inflate to become geometrically ``thick\" and sub-Keplerian. Another interesting aspect of disks in this accretion regime is the possibility of generating significant outflows \\citep{adios,gonewiththewind}.  The inability to radiate enough energy also arises for very high accretion rates, giving rise to advection dominated flows once the disk becomes radiation pressure dominated \\citep{Paczynsky1980,Jaro:1980}.  The energy release rate in a Keplerian disk, down to a radius $r$, is given by $GM_\\mr{BH}\\dot{m}/2r$, where $M_\\mr{BH}$ is the black hole mass, and $\\dot{m}$ is the mass accretion rate. This energy should be compared to the Eddington luminosity of the BH, defined as: \\begin{equation} L_\\mr{Edd}\\equiv \\frac{4\\pi cGM_\\mr{BH}}{\\kappa}, \\label{eq:ledd} \\end{equation} from which a critical accretion rate can be define to be \\begin{equation} \\dot{m}_\\mr{crit}\\equiv\\frac{L_\\mr{Edd}}{\\eta_0 c^2}, \\label{eq:mdotcrit} \\end{equation} with $\\eta_0=1/16$ being the standard efficiency for accretion around a Schwarzschild BH. If $\\dot{m}$ is large enough, that is, $\\dot{m} \\gtrsim \\dot{m}_\\mr{crit}$ then the energy released will approach the Eddington luminosity before reaching the inner radius of the disk. From that point inwards, the radiation pressure will be dominant. The disk will become geometrically thick (having a scale height $H(r) \\sim r$), and the local flux will approach the Eddington flux defined as \\begin{equation} F_\\mr{Edd}\\equiv\\frac{cGM_\\mr{BH}z}{\\kappa R^3}, \\label{eq:Fedd} \\end{equation} with $R \\equiv \\sqrt{r^2 + z^2}$ (where $r$ and $z$ are the cylindrical coordinates). However, because the Eddington flux cannot be surpassed,  the dissipated energy cannot be entirely radiated away and part of it should therefore be advected inwards. The high pressure also implies that the disk rotates at sub-Keplerian velocities.  Other interesting aspects of these disks is their possible ability to accelerate very luminous jets \\citep{AP1980}, and the fact that although the local flux in the disk does not surpass the Eddington flux, the overall luminosity can surpass the Eddington luminosity, simply due to the disk geometry \\citep{Jaro:1980, Paczynsky1980}.  Roughly, the Eddington luminosity can be surpassed by a factor of $\\sim \\ln r_{out}/r_{in}$, where the radii denote the inner and outer extents of the radiation pressure disk.  For nearly critical accretion, an intermediate type of solutions exists, that of ``slim-disks\". However, because it exists only near the last stable orbit (thus allowing for advection of heat into the Roche Lobe overflow), its solution requires relativistic corrections \\citep{Abramowicz1988}.  The above models assume, however, that the Eddington flux cannot be surpassed locally. Nevertheless, it was shown that super-Eddinigton states do naturally arise in nature \\citep{ShavivNovae}, allowing for high luminosities, while generating optically thick winds. Our goal in the present work is to consider the recent advances in the understanding of how super-Eddington atmospheres arise, and what they look like, and to incorporate these ideas into models for very high accretion rate accretion disks.  In \\S\\ref{sec:background}, we begin by reviewing our present understanding of how super-Eddington states arise. In \\S\\ref{sec:model} we describe our model for super-Eddington accretion, and in \\S\\ref{sec:numerical} we describe the numerical solution. In \\S\\ref{sec:results} we describe the numerical results, and end with a discussion in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}   In the present analysis, we  searched numerically for possible solutions describing super-critical accretion disks around black holes, while allowing for the formation of ``porous\" layers with a reduced opacity. We found solutions with significantly super-critical accretion rates, in which the vertical disk height is  smaller than the radius, that is, {\\em slim} disks. Because the super-Eddington state excites a strong wind, the actual mass accretion onto the BH can be notably smaller. Solutions were found to exist with accretion rates ranging between about 0.5 ${\\dot m}_{crit}$ to about 20${\\dot m}_{crit}$. In all cases, there is a photon tired continuum driven wind.  At the low range, the disks are overall sub-Eddington, but locally the flux can surpass the critical value, and therefore it can excite continuum driven winds.  Namely, the critical accretion rate given by eq.\\ \\ref{eq:mdotcrit} is not the lower limit for wind generation. At the low range, almost all the energy dissipated is either radiated away or transferred to the wind. Moreover, the wind is then optically thin and the photosphere coincide with the sonic point of the wind.  As the outer mass accretion rate increases, the winds become more massive, thereby reducing the fraction of mass accreted onto the black hole. Also, some of the dissipated energy is then advected with the flow into the BH.  For very high accretion rates surpassing about 20$\\dot{m}_{crit}$, the photosphere which resides in the wind is found to be located at $z \\gtrsim r$. This implies that the 1D+1D type of solution described here breaks down. Instead, one has to look for a solution in which the semi-hydrostatic inflow has a disk-like solution, while the super-sonic outgoing wind has a spherical-like solution.  The description of such extremely high accretion rate disks is the subject of a future publication.  One of the uncertainties in the model is the opacity law behaviour of the porous atmosphere. The reasons it is not known well is because it depends on the nonlinear radiative hydrodynamic configuration the atmosphere will reach, and unfortunately, there are still no numerical simulations or empirical data which can constrain the effective opacity law. It is for this reason that we parameterized the effective opacity (see eq.\\ \\ref{eq:kappa_eff})  As can be seen in figs.\\ \\ref{fig:mparameter}-\\ref{fig:Lparameter} almost all disks characteristics have either a small or a modest sensitivity to the changes in the opacity law. For example $\\dot{\\mr{m}}_{\\mr{real}}/\\dot{\\mr{m}}_{\\mr{out}}$ varies between 0.43 to 0.44, or the total luminosity varies between $\\sim 0.570 \\ledd$ to $0.576 \\ledd$, while changing the opacity parameter $B$ by 50\\%. This implies that the uncertainties do not undermine the model predictions. On the other hand, it would be impossible to use super-Eddington accretion disks to constrain the relevant parameters."
    },
    "1004/1004.4517_arXiv.txt": {
        "abstract": "Numerical values of charged-particle thermonuclear reaction rates for nuclei in the A=14 to 40 region are tabulated. The results are obtained using a method, based on Monte Carlo techniques, that has been described in the preceding paper of this series (Paper I). We present a low rate, median rate and high rate which correspond to the 0.16, 0.50 and 0.84 quantiles, respectively, of the cumulative reaction rate distribution. The meaning of these quantities is in general different from the commonly reported, but statistically meaningless expressions, ``lower limit\", ``nominal value\" and ``upper limit\" of the total reaction rate. In addition, we approximate the Monte Carlo probability density function of the total reaction rate by a lognormal distribution and tabulate the lognormal parameters $\\mu$ and $\\sigma$ at each temperature. We also provide a quantitative measure (Anderson-Darling test statistic) for the reliability of the lognormal approximation. The user can implement the approximate lognormal reaction rate probability density functions directly in a stellar model code for studies of stellar energy generation and nucleosynthesis. For each reaction, the Monte Carlo reaction rate probability density functions, together with their lognormal approximations, are displayed graphically for selected temperatures in order to provide a visual impression. Our new reaction rates are appropriate for {\\it bare nuclei in the laboratory}. The nuclear physics input used to derive our reaction rates is presented in the subsequent paper of this series (Paper III). In the fourth paper of this series (Paper IV) we compare our new reaction rates to previous results. ",
        "introduction": "In the preceding work, referred to as Paper I, we presented a method of evaluating thermonuclear reaction rates that is based on the Monte Carlo technique. The method allows for calculating statistically meaningful values: the recommended reaction rate is derived from the median of the total reaction rate probability density function, while the 0.16 and 0.84 quantiles of the cumulative distribution provide values for the {\\it low rate} and the {\\it high rate}, respectively. We refer to such rates as ``Monte Carlo reaction rates\" in order to distinguish them from results obtained using previous techniques (which we call ``classical reaction rates\"). As explained in Paper I, we will strictly avoid using the statistically meaningless expressions ``lower limit\" (or ``minimum\") and  ``upper limit\" (or ``maximum\") of the total reaction rate. For detailed information on our method, see Paper I.  In the present work, referred to in the following as Paper II, we present our numerical results of charged-particle thermonuclear reaction rates for A=14 to 40 nuclei on a grid of temperatures ranging from T=0.01 GK to 10 GK. These reaction rates are assumed to involve {\\it bare nuclei in the laboratory}. The rates of reactions induced on lighter target nuclei, A$<$14, are not easily analyzed in terms of the present techniques and require different procedures (see, for example, Descouvemont et al. \\cite{Des04} for an evaluation of Big Bang nuclear reaction rates using R-matrix theory). The higher target mass cutoff, A=40, was entirely dictated by limitations in resources and time. For use in stellar model calculations, the results presented here must be corrected, if appropriate, for (i) electron screening at elevated densities, and (ii) thermal excitations of the target nucleus at elevated temperatures. Details will be provided below. We emphasize that the present reaction rates are overwhelmingly based on {\\it experimental nuclear physics information}. Only in exceptional situations, for example, when a certain nuclear property had not been measured yet, did we resort to nuclear theory. In the subsequent work (Paper III) we will provide the complete nuclear physics data input used to derive our new Monte Carlo reaction rates. In the fourth paper of this series (Paper IV) we compare our new reaction rates to previous results.  Paper II is organized as follows. In Sec. 2 we summarize briefly our Monte Carlo technique. An overview of the literature sources for the nuclear physics input data used to derive our results is provided in Sec. 3. Detailed examples for how to interpret the new Monte Carlo reaction rates are given in Sec. 4. The extrapolation of the laboratory reaction rate to high temperatures is described in Sec. 5, while modifications of the reaction rate that are necessary for use in stellar model calculations are discussed in Sec. 6. The calculation of reverse rates is described in Sec. 7. A summary is given in Sec. 8. Appendix A contains information regarding statistical hypothesis tests. Our Monte Carlo reaction rates are presented in tabular and graphical format in App. B. ",
        "conclusions": "We tabulate charged particle thermonuclear reaction rates and probability density functions for target nuclei in the A=14 to 40 mass range, which are overwhelmingly based on {\\it experimental nuclear physics information}. The results are obtained using a Monte Carlo method, which is based on physically motivated probability density functions of all nuclear physics input quantities. For the first time, an evaluation is presented that provides statistically meaningful low, recommended and high reaction rates. They are obtained at each stellar temperature from the 0.16, 0.50 and 0.84 quantiles of the cumulative reaction rate distribution (corresponding to a coverage probability of 68\\%). Graphs of reaction rate probability density functions at selected temperatures and of reaction rate uncertainties are provided for visual inspection.  We approximate the Monte Carlo probability density function of the total reaction rate by a lognormal distribution. At each temperature, the lognormal parameters $\\mu$ and $\\sigma$ are tabulated. A quantitative measure (Anderson-Darling test statistic) is provided for assessing the reliability of the lognormal approximation. The reasons why the actual Monte Carlo reaction rate probability density function {\\it frequently follows} a lognormal distribution are discussed in detail (Sec. \\ref{sec:si28pg} and Sec. 5.4 of Paper I). The user can implement the approximate lognormal reaction rate probability density functions directly in a stellar model code for studies of stellar energy generation and nucleosynthesis.  We provide clear evidence that the reaction rate probability density function does {\\it not always follow} a lognormal distribution, as is sometimes erroneously claimed in the literature. This is especially the case when the uncertainty in a resonance energy is large and the rate is calculated from a particle partial width, or if upper limits of partial widths influence the total rate. If a particular reaction rate probability density deviates significantly from lognormality and it turns out that stellar model calculations are sensitive to the {\\it shape} of this distribution, then either new measurements need to be performed in order to improve uncertainties of important input quantities, or another analytical function must be used  that is more complicated than a lognormal distribution, but approximates the actual reaction rate probability density function more closely.  The results tabulated here are appropriate for {\\it bare nuclei in the laboratory}. For use in stellar model calculations, they must be corrected, if appropriate, at elevated temperatures for thermal target excitations and at elevated densities for electron screening effects. These corrections introduce another source of uncertainty, which is not easily quantified at present. Thus our {\\it laboratory} reaction rate uncertainties and probability density functions may not provide, in general, an impression of the {\\it stellar} reaction rate uncertainties and probability density functions. More work is called for to assess the reliability of Hauser-Feshbach reaction rates in the region of the light nuclei (A$\\leq$40). The probability density function of a reverse rate is easily calculated from the tabulated lognormal parameters of the corresponding forward rate. However, in an actual reaction network study the forward and reverse rates should not be sampled independently since they are correlated. The nuclear physics input used to derive our results is presented in the subsequent paper of this series (Paper III). In the fourth paper of this series (Paper IV) we compare our reaction rates to previous results."
    },
    "1004/1004.3337_arXiv.txt": {
        "abstract": "The mass function of galaxy clusters is a powerful tool to constrain cosmological parameters, e.g., the mass fluctuation on the scale of 8~$h^{-1}$~Mpc, $\\sigma_8$, and the abundance of total matter, $\\Omega_m$. We first determine the scaling relations between cluster mass and cluster richness, summed $r$-band luminosity and the global galaxy number within a cluster radius. These relations are then used to two complete volume-limited rich cluster samples which we obtained from the Sloan Digital Sky Survey (SDSS). We estimate the masses of these clusters and determine the cluster mass function. Fitting the data with a theoretical expression, we get the cosmological parameter constraints in the form of $\\sigma_8(\\Omega_m/0.3)^{\\alpha}=\\beta$ and find out the parameters of $\\alpha=$0.40--0.50 and $\\beta=$0.8--0.9, so that $\\sigma_8=$0.8--0.9 if $\\Omega_m=0.3$. Our $\\sigma_8$ value is slightly higher than recent estimates from the mass function of X-ray clusters and the Wilkinson Microwave Anisotropy Probe (WMAP) data, but consistent with the weak lensing statistics. ",
        "introduction": "Precise determination of cosmological parameters is an important goal in astrophysics. In the linear theory, the present root-mean-square (rms) mass fluctuation on the scale of 8~$h^{-1}$~Mpc, $\\sigma_8$, is one of fundamental parameters \\citep[see][]{svp+03} to describe the power spectrum of mass fluctuations in the universe. It is one of key parameters in the large scale structure simulations \\citep[e.g.,][]{jfp+98}. The $\\sigma_8$ can be determined by galaxy-galaxy correlations \\citep[e.g.,][]{tbs+04,cpp+05}, fluctuations in the cosmic microwave background \\citep{svp+03,sbd+07,kdn+09}, gravitational lensing statistics \\citep[e.g.,][]{hmv+06,kht+07,bhs+07}, cluster mass function \\citep[e.g.,][]{wef93,bf98,rb02}, Ly$\\alpha$ forest \\citep{jnt+05,msc+05} and galaxy peculiar velocities \\citep{fjf+03}.  The cluster mass function can be determined by the estimated masses for a sample of clusters \\citep[e.g.,][]{dah06}, or by the X-ray luminosity and temperature function with a prior scaling relation \\citep{vl96,asf+03}. Fitting the cluster mass function with a theoretical expression can provide constraint on $\\sigma_8$. Generally, $\\sigma_8$ is coupled with $\\Omega_m$, the abundance of present total matter, in the form of $\\sigma_8(\\Omega_m/0.3)^{\\alpha}=\\beta$. Previous studies have found $\\alpha$ in the range 0.3--0.6 and $\\beta$ in the range 0.6--1.2 (see Table~\\ref{comp} in Section~\\ref{discu}). The determined $\\sigma_8$ in recent years (2002--2009) has a mean value of 0.73$\\pm$0.05 assuming $\\Omega_m=0.3$, which is in agreement with the WMAP data \\citep{kdn+09}, but lower than those by weak lensing statistics \\citep{hss+07}, galaxy-galaxy correlations \\citep{tbs+04,cpp+05} and Ly$\\alpha$ forest \\citep{jnt+05,msc+05}.  The amplitude of cluster mass function has large uncertainties, mainly caused by the uncertain normalization of the mass scaling relation \\citep[e.g., ][]{hen04}. Other uncertainties come from the scatter of mass scaling relation and the incompleteness of the X-ray flux-limited cluster samples \\citep{rb02}. The cluster mass function may be underestimated if only X-ray clusters are used. \\citet{evm+00} and \\citet{dpl+03} have noticed the existence of a class of X-ray-underluminous massive clusters. \\citet{pbb+07a} found that 40\\% of Abell clusters have a low level or no detection in X-rays. A large complete volume-limited sample of clusters is crucial for the purpose. Using the photometric redshifts of galaxies, we found 39,668 clusters in the redshift range $0.05<z<0.6$ \\citep{whl09}. Clusters are approximate volume-limited complete in the redshift range $0.05<z<0.42$. The richnesses and the summed luminosities of clusters are estimated from their luminous members, and they are tightly related to cluster mass. In Section 2, we carefully determine the scaling relation for cluster mass. In Section 3, we get the cluster mass function for a local sample of clusters and a sample at mediate redshifts, and then fit the cluster mass function with a theoretical expression for constraints on cosmological parameters, $\\Omega_m$ and $\\sigma_8$. Discussions and conclusions are given in Section 4.  Throughout this paper, we assume a flat $\\Lambda$CDM cosmology, taking $H_0=$100 $h$ ${\\rm km~s}^{-1}$ ${\\rm Mpc}^{-1}$, with $h=0.72$, $\\Omega_m=1-\\Omega_{\\Lambda}$.  \\begin{table*} \\begin{minipage}{160mm} \\caption[]{Cluster masses from literature and the mass tracer values for 53 clusters (richness $R\\ge8$ and $0.03\\le z\\le0.3$, sorted with $z$) in the field of the SDSS DR6.} \\begin{center} \\begin{tabular}{lrrccccrlr} \\hline \\mc{1}{l}{Name}&\\mc{1}{c}{R.A.}&\\mc{1}{c}{Decl.}&\\mc{1}{c}{$z$}&\\mc{1}{c}{$R$}& \\mc{1}{c}{$L_r$} &\\mc{1}{c}{$GGN/r_{\\rm GGN}$}& \\mc{1}{c}{M$_{\\rm vir}$} & \\mc{1}{c}{Method} & \\mc{1}{c}{Ref.} \\\\ &\\mc{1}{c}{(deg)}& \\mc{1}{c}{(deg)} & & &\\mc{1}{c}{($10^{10}~h^{-2}~L_{\\odot}$)} & \\mc{1}{c}{(Mpc$^{-1}$)} & \\mc{1}{c}($10^{14}~h^{-1}~M_{\\odot}$) & &\\\\ \\hline Abell 2199     &   247.15930 &   39.55121 &  0.030 & 17.76 & 28.49 & 18.10 &   $2.39^{+0.38}_{-0.38}$ &X-ray &  1 \\\\%xu01x a &             &            &        &       &       &       &   $3.42^{+0.26}_{-0.26}$ &X-ray &  2 \\\\%rei02xb Abell 2052     &   229.18536 &    7.02162 &  0.035 &  9.50 & 21.97 & 11.62 &   $1.47^{+0.28}_{-0.28}$ &X-ray &  1 \\\\%xu01x a &             &            &        &       &       &       &   $1.58^{+0.05}_{-0.06}$ &X-ray &  2 \\\\%rei02xb Abell 2063     &   230.77209 &    8.60922 &  0.035 & 15.30 & 22.96 & 19.93 &   $2.31^{+0.18}_{-0.16}$ &X-ray &  2 \\\\%rei02xb &             &            &        &       &       &       &   $2.78^{+0.42}_{-0.42}$ &X-ray &  1 \\\\%xu01x a Abell 2147     &   240.57094 &   15.97465 &  0.035 & 12.22 & 22.50 & 14.67 &   $2.46^{+0.83}_{-0.52}$ &X-ray &  2 \\\\%rei02xb Abell 2151w    &   241.14914 &   17.72156 &  0.037 & 11.08 & 22.16 & 13.29 &   $1.23^{+0.09}_{-0.09}$ &X-ray &  2 \\\\%rei02xb MKW9           &   233.13339 &    4.68100 &  0.040 &  8.88 & 15.77 & 11.44 &   $1.01^{+0.25}_{-0.53}$ &X-ray &  3 \\\\%poi05xc &             &            &        &       &       &       &   $1.06^{+0.27}_{-0.27}$ &X-ray &  4 \\\\%arn05xd Abell 1983     &   223.23048 &   16.70286 &  0.044 &  9.27 & 11.76 & 11.45 &   $1.39^{+0.53}_{-0.53}$ &X-ray &  3 \\\\%poi05xc &             &            &        &       &       &       &   $1.41^{+0.55}_{-0.55}$ &X-ray &  4 \\\\%arn05xd Abell 160      &    18.24822 &   15.49129 &  0.045 & 16.61 & 34.96 & 16.54 &   $0.96^{+0.13}_{-0.14}$ &X-ray &  5 \\\\%xun09xe Abell 85       &    10.46029 & $-$9.30312 &  0.056 & 23.61 & 54.65 & 21.85 &   $4.23^{+1.15}_{-1.15}$ &X-ray &  1 \\\\%xu01x a &             &            &        &       &       &       &   $5.55^{+0.58}_{-0.53}$ &X-ray &  2 \\\\%rei02xb Abell 1991     &   223.63122 &   18.64232 &  0.059 & 19.22 & 39.13 & 17.48 &   $1.35^{+0.19}_{-0.53}$ &X-ray &  3 \\\\%poi05xc &             &            &        &       &       &       &   $1.44^{+0.17}_{-0.17}$ &X-ray &  4 \\\\%arn05xd &             &            &        &       &       &       &   $1.52^{+0.21}_{-0.21}$ &X-ray &  6 \\\\%vik06xf &             &            &        &       &       &       &   $1.66^{+0.32}_{-0.24}$ &X-ray &  5 \\\\%xun09xe Abell 1795     &   207.21877 &   26.59293 &  0.063 & 19.56 & 39.70 & 19.56 &   $5.42^{+0.65}_{-0.65}$ &X-ray &  1 \\\\%xu01x a &             &            &        &       &       &       &   $7.58^{+1.92}_{-1.70}$ &X-ray &  7 \\\\%sch07xg &             &            &        &       &       &       &   $7.86^{+0.70}_{-0.70}$ &X-ray &  6 \\\\%vik06xf &             &            &        &       &       &       &   $7.94^{+1.64}_{-1.51}$ &X-ray &  2 \\\\%rei02xb Abell 2092     &   233.31403 &   31.14515 &  0.067 &  9.27 & 15.93 & 11.45 &   $1.09^{+0.23}_{-0.20}$ &X-ray &  5 \\\\%xun09xe Abell 2065     &   230.60008 &   27.71436 &  0.072 & 40.86 & 76.31 & 40.86 &  $12.11^{+15.48}_{-4.88}$&X-ray &  2 \\\\%rei02xb ZwCl 1215      &   184.42134 &    3.65584 &  0.075 & 20.69 & 43.81 & 18.71 &   $7.54^{+2.55}_{-1.91}$ &X-ray &  2 \\\\%rei02xb Abell 1800     &   207.34822 &   28.10732 &  0.075 & 19.77 & 44.73 & 18.83 &   $4.14^{+4.31}_{-1.65}$ &X-ray &  2 \\\\%rei02xb Abell 1775     &   205.45477 &   26.37347 &  0.076 & 14.52 & 32.62 & 17.76 &   $3.07^{+0.43}_{-0.29}$ &X-ray &  2 \\\\%rei02xb Abell 2029     &   227.73376 &    5.74478 &  0.077 & 24.87 & 60.63 & 27.47 &   $7.06^{+0.69}_{-0.54}$ &X-ray &  7 \\\\%sch07xg &             &            &        &       &       &       &   $9.30^{+1.93}_{-1.93}$ &X-ray &  1 \\\\%xu01x a &             &            &        &       &       &       &   $9.76^{+1.77}_{-1.65}$ &X-ray &  2 \\\\%rei02xb &             &            &        &       &       &       &  $10.55^{+1.01}_{-1.01}$ &X-ray &  6 \\\\%vik06xf Abell 2255     &   258.11996 &   64.06072 &  0.080 & 40.23 & 83.40 & 33.51 &   $9.70^{+1.05}_{-0.87}$ &X-ray &  2 \\\\%rei02xb Abell 1650     &   194.67288 & $-$1.76146 &  0.084 & 17.08 & 42.92 & 20.68 &   $8.13^{+4.21}_{-2.53}$ &X-ray &  2 \\\\%rei02xb Abell 1692     &   198.05661 & $-$0.97448 &  0.085 & 14.40 & 26.42 & 14.73 &   $1.19^{+0.38}_{-0.24}$ &X-ray &  5 \\\\%xun09xe Abell 1750     &   202.71080 & $-$1.86197 &  0.088 & 24.35 & 58.05 & 27.58 &   $4.78^{+2.64}_{-2.64}$ &WL    &  8 \\\\%oka08 h Abell 2142     &   239.58334 &   27.23341 &  0.090 & 39.33 & 80.90 & 39.88 &  $11.00^{+2.85}_{-1.93}$ &X-ray &  2 \\\\%rei02xb &             &            &        &       &       &       &  $11.95^{+5.24}_{-5.24}$ &WL    &  8 \\\\%oka08 h Abell 2244     &   255.67705 &   34.05999 &  0.097 & 31.22 & 62.35 & 28.67 &   $7.23^{+9.45}_{-3.25}$ &X-ray &  2 \\\\%rei02xb Abell 2034     &   227.54883 &   33.48646 &  0.113 & 32.00 & 68.97 & 31.08 &   $7.17^{+4.30}_{-4.30}$ &WL    &  8 \\\\%oka08 h Abell 1068     &   160.18541 &   39.95313 &  0.138 & 16.38 & 38.87 & 16.38 &   $4.87^{+0.42}_{-0.53}$ &X-ray &  3 \\\\%poi05xc &             &            &        &       &       &       &   $5.12^{+0.45}_{-0.45}$ &X-ray &  4 \\\\%arn05xd Abell 1413     &   178.82501 &   23.40491 &  0.143 & 37.17 & 93.39 & 41.82 &   $5.70^{+0.57}_{-0.53}$ &X-ray &  3 \\\\%poi05xc &             &            &        &       &       &       &   $5.87^{+0.59}_{-0.59}$ &X-ray &  4 \\\\%arn05xd &             &            &        &       &       &       &   $8.72^{+0.80}_{-0.70}$ &X-ray &  2 \\\\%rei02xb &             &            &        &       &       &       &   $9.76^{+2.54}_{-1.89}$ &X-ray &  7 \\\\%sch07xg &             &            &        &       &       &       &   $9.95^{+1.04}_{-1.04}$ &X-ray &  6 \\\\%vik06xf RXJ1720.1$+$2637&  260.04184 &   26.62557 &  0.164 & 26.87 & 69.11 & 27.30 &   $5.88^{+4.98}_{-4.60}$ &WL    &  9 \\\\%ped07 i &             &            &        &       &       &       &   $5.84^{+4.30}_{-4.30}$ &WL    & 10 \\\\%dah06 j Abell 1914     & 216.48611   & 37.81645  & 0.171  & 34.16 &  84.15 &  38.07  & $ 4.66^{+3.78}_{ -3.49}$ &WL   &  9 \\\\%ped07 i &             &           &        &\t &        &         & $ 6.14^{+3.19}_{ -3.19}$ &WL   &  8 \\\\%oka08 h &             &           &        &\t &        &         & $ 9.08^{+5.66}_{ -5.66}$ &WL   & 10 \\\\%dah06 j &             &           &        &\t &        &         & $18.40^{+2.20}_{ -1.86}$ &X-ray&  2 \\\\%rei02xb MS 0906.5$+$1110& 137.30312  & 10.97475  & 0.176  & 36.00 & 119.08 &  38.61  & $ 8.30^{+2.30}_{ -2.30}$ &WL   & 11 \\\\%hoe07 k \\end{tabular} \\end{center} \\label{mass} \\end{minipage} \\end{table*}  \\addtocounter{table}{-1} \\begin{table*} \\begin{minipage}{160mm} \\caption[]{-- {\\it continued} } \\begin{center} \\begin{tabular}{lrrccccrlr} \\hline \\mc{1}{l}{Name}&\\mc{1}{c}{R.A.}&\\mc{1}{c}{Decl.}&\\mc{1}{c}{$z$}&\\mc{1}{c}{$R$}& \\mc{1}{c}{$L_r$} &\\mc{1}{c}{$GGN/r_{\\rm GGN}$}& \\mc{1}{c}{M$_{\\rm vir}$} & \\mc{1}{c}{Method} & \\mc{1}{c}{Ref.} \\\\ &\\mc{1}{c}{(deg)}& \\mc{1}{c}{(deg)} & & &\\mc{1}{c}{($10^{10}~h^{-2}~L_{\\odot}$)} & \\mc{1}{c}{(Mpc$^{-1}$)} & \\mc{1}{c}($10^{14}~h^{-1}~M_{\\odot}$) & &\\\\ \\hline Abell 1689     & 197.87291  &$-$1.34108 & 0.184  & 62.74 & 136.85 &  50.32  & $10.50^{+1.91}_{ -1.91}$ &WL   & 12 \\\\%cor09 l &            &           &        &\t &        &         & $11.69^{+1.80}_{ -1.80}$ &SL+WL& 13 \\\\%lim07 m &            &           &        &\t &        &         & $13.24^{+3.99}_{ -3.99}$ &X-ray& 14 \\\\%zha07xn &            &           &        &\t &        &         & $13.43^{+0.93}_{ -0.96}$ &X-ray&  2 \\\\%rei02xb &            &           &        &\t &        &         & $13.51^{+1.40}_{ -1.40}$ &WL   & 15 \\\\%bro05 o &            &           &        &\t &        &         & $14.70^{+1.40}_{ -1.40}$ &WL   & 16 \\\\%ume08 p &            &           &        &\t &        &         & $17.55^{+3.04}_{ -3.04}$ &WL   & 17 \\\\%bar07 q &            &           &        &\t &        &         & $20.53^{+1.74}_{ -1.74}$ &WL   & 18 \\\\%hal06 r &            &           &        &\t &        &         & $25.50^{+5.30}_{ -4.50}$ &WL   & 11 \\\\%hoe07 k Abell 963      & 154.26515  & 39.04705  & 0.206  & 43.00 &  93.83 &  38.52  & $ 3.45^{+0.80}_{ -0.80}$ &WL   & 17 \\\\%bar07 q &            &           &        &\t &        &         & $ 6.50^{+2.00}_{ -1.90}$ &WL   & 11 \\\\%hoe07 k &            &           &        &\t &        &         & $ 6.53^{+2.00}_{ -2.00}$ &X-ray& 14 \\\\%zha07xn &            &           &        &\t &        &         & $ 7.00^{+2.45}_{ -1.60}$ &X-ray&  7 \\\\%sch07xg &            &           &        &\t &        &         & $ 8.01^{+8.14}_{ -6.36}$ &WL   &  9 \\\\%ped07 i &            &           &        &\t &        &         & $ 8.38^{+7.46}_{ -7.46}$ &WL   & 10 \\\\%dah06 j Abell 1423     & 179.32219  & 33.61092  & 0.214  & 20.78 &  64.82 &  21.16  & $14.98^{+7.67}_{ -7.67}$ &WL   & 10 \\\\%dah06 j RX J1504.1$-$0248&226.03130 &$-$2.80460 & 0.215  & 20.08 &  73.32 &  26.82  & $15.10^{+9.40}_{ -4.30}$ &X-ray&  7 \\\\%sch07xg Abell 773      & 139.47261  & 51.72704  & 0.217  & 66.83 & 138.50 &  55.08  & $10.63^{+3.27}_{ -3.27}$ &X-ray& 14 \\\\%zha07xn &            &           &        &\t &        &         & $15.33^{+6.95}_{ -6.95}$ &WL   & 10 \\\\%dah06 j &            &           &        &\t &        &         & $24.70^{+9.44}_{-11.91}$ &WL   &  9 \\\\%ped07 i Abell 1682     & 196.70833  & 46.55927  & 0.226  & 40.76 & 112.14 &  39.49  & $ 3.96^{+3.52}_{ -2.40}$ &WL   &  9 \\\\%ped07 i &            &           &        &\t &        &         & $ 5.49^{+3.72}_{ -3.72}$ &WL   & 10 \\\\%dah06 j Abell 1763     & 203.83372  & 41.00115  & 0.228  & 34.10 & 109.20 &  32.77  & $ 6.23^{+1.92}_{ -1.92}$ &X-ray& 14 \\\\%zha07xn &            &           &        &       &        &         & $ 8.91^{+4.60}_{ -5.70}$ &WL   &  9 \\\\%ped07 i &            &           &        &\t &        &         & $10.54^{+5.16}_{ -5.16}$ &WL   & 10 \\\\%dah06 j &            &           &        &\t &        &         & $12.28^{+2.37}_{ -2.37}$ &WL   & 17 \\\\%bar07 q &            &           &        &       &        & \t    & $13.50^{+3.70}_{ -3.30}$ &WL   & 11 \\\\%hoe07 k Abell 2219     & 250.08253  & 46.71148  & 0.228  & 45.62 &  94.82 &  45.62  & $ 8.19^{+5.69}_{ -5.69}$ &WL   & 10 \\\\%dah06 j &            &           &        &\t &        &         & $ 8.57^{+4.46}_{ -5.45}$ &WL   &  9 \\\\%ped07 i &            &           &        &\t &        &         & $11.30^{+3.20}_{ -2.70}$ &WL   & 11 \\\\%hoe07 k &            &           &        &\t &        &         & $18.66^{+3.94}_{ -3.94}$ &WL   & 17 \\\\%bar07 q Abell 2111     & 234.91872  & 34.42426  & 0.229  & 47.58 & 112.27 &  43.38  & $ 6.12^{+3.64}_{ -3.64}$ &WL   & 10 \\\\%dah06 j &            &           &        &       &        &         & $ 6.92^{+3.39}_{ -3.95}$ &WL   &  9 \\\\%ped07 i Abell 267      &  28.17483  &  1.00711  & 0.230  & 31.53 &  86.85 &  32.62  & $ 2.36^{+1.29}_{ -1.29}$ &WL   & 17 \\\\%bar07 q &            &           &        &       &        &         & $ 5.36^{+1.70}_{ -1.70}$ &X-ray& 14 \\\\%zha07xn &            &           &        &\t &        &         & $ 7.50^{+2.30}_{ -2.20}$ &WL   & 11 \\\\%hoe07 k &            &           &        &\t &        &         & $14.98^{+5.63}_{ -5.63}$ &WL   & 10 \\\\%dah06 j &            &           &        &\t &        &         & $16.34^{+5.57}_{ -6.75}$ &WL   &  9 \\\\%ped07 i Zw 1231.4$+$1007 & 188.57277  &  9.76623  & 0.231  & 32.00 &  98.28 &  36.99  & $ 3.15^{+2.65}_{ -2.65}$ &WL   & 10 \\\\%dah06 j MS 1231.3$+$1542 & 188.48055  & 15.43305  & 0.234  & 13.82 &  35.53 &  17.40  & $ 1.50^{+0.90}_{ -0.90}$ &WL   & 11 \\\\%hoe07 k RX J2129.6$+$0005& 322.41649  &  0.08920  & 0.234  & 24.00 &  80.65 &  20.45  & $ 5.96^{+8.54}_{ -2.70}$ &X-ray&  7 \\\\%sch07xg &            &           &        &\t &        &         & $ 7.51^{+5.24}_{ -5.24}$ &WL   & 10 \\\\%dah06 j &            &           &        &       &        &         & $15.49^{+8.97}_{ -9.24}$ &WL   &  9 \\\\%ped07 i Abell 1835     & 210.25863  &  2.87846  & 0.252  & 51.57 & 129.59 &  47.96  & $ 8.19^{+3.07}_{ -3.07}$ &WL   & 12 \\\\%cor09 l &            &           &        &\t &        &         & $ 8.41^{+2.57}_{ -2.57}$ &X-ray& 14 \\\\%zha07xn &            &           &        &\t &        &         & $10.61^{+5.69}_{ -5.69}$ &WL   & 10 \\\\%dah06 j &            &           &        &\t &        &         & $15.62^{+8.56}_{ -6.35}$ &WL   &  9 \\\\%ped07 i &            &           &        &\t &        &         & $17.00^{+3.10}_{ -3.40}$ &X-ray&  7 \\\\%sch07xg &            &           &        &\t &        &         & $24.21^{+3.76}_{ -3.76}$ &WL   & 17 \\\\%bar07 q MS 1455.0$+$2232 & 224.31295  & 22.34288  & 0.258  & 26.00 &  72.03 &  26.00  & $ 5.75^{+4.11}_{ -3.26}$ &WL   &  9 \\\\%ped07 i &            &           &        &\t &        &         & $ 7.30^{+1.90}_{ -1.80}$ &WL   & 11 \\\\%hoe07 k &            &           &        &\t &        &         & $12.88^{+6.07}_{ -6.07}$ &WL   & 10 \\\\%dah06 j Abell 2631     & 354.41554  &  0.27138  & 0.277  & 46.22 & 113.32 &  44.94  & $ 6.12^{+4.25}_{ -4.25}$ &WL   & 10 \\\\%dah06 j Abell 1758N    & 203.16007  & 50.55992  & 0.279  & 42.24 & 107.00 &  41.36  & $ 5.26^{+5.70}_{ -5.70}$ &WL   &  8 \\\\%oke08 h &            &           &        &       &        &         & $26.91^{+9.57}_{ -9.57}$ &WL   & 10 \\\\%dah06 j &            &           &        &       &        &         & $39.09^{+12.77}_{-13.16}$&WL   &  9 \\\\%ped07 i \\end{tabular} \\end{center} \\end{minipage} \\end{table*}  \\addtocounter{table}{-1} \\begin{table*} \\begin{minipage}{160mm} \\caption[]{-- {\\it continued} } \\begin{center} \\begin{tabular}{lrrccccrlr} \\hline \\mc{1}{l}{Name}&\\mc{1}{c}{R.A.}&\\mc{1}{c}{Decl.}&\\mc{1}{c}{$z$}&\\mc{1}{c}{$R$}& \\mc{1}{c}{$L_r$} &\\mc{1}{c}{$GGN/r_{\\rm GGN}$}& \\mc{1}{c}{M$_{\\rm vir}$} & \\mc{1}{c}{Method} & \\mc{1}{c}{Ref.} \\\\ &\\mc{1}{c}{(deg)}& \\mc{1}{c}{(deg)} & & &\\mc{1}{c}{($10^{10}~h^{-2}~L_{\\odot}$)} & \\mc{1}{c}{(Mpc$^{-1}$)} & \\mc{1}{c}($10^{14}~h^{-1}~M_{\\odot}$) & &\\\\ \\hline Abell 697      & 130.73982  & 36.36646  & 0.282  & 27.72 &  77.66 &  37.83  & $22.92^{+9.78}_{ -9.46}$ &WL   &  9 \\\\%ped07 i &            &           &        &       &        &         & $26.10^{+9.54}_{ -9.54}$ &WL   & 10 \\\\%dah06 j Abell 959      & 154.39984  & 59.56710  & 0.285  & 52.77 & 139.82 &  48.82  & $17.38^{+6.81}_{ -5.85}$ &WL   &  9 \\\\%ped07 i Abell 611      & 120.23674  & 36.05655  & 0.288  & 27.86 &  91.23 &  30.92  & $ 6.18^{+3.82}_{ -1.81}$ &X-ray&  7 \\\\%sch07xg &            &           &        &       &        &         & $ 6.54^{+4.43}_{ -4.43}$ &WL   & 10 \\\\%dah06 j &            &           &        &       &        &         & $ 6.90^{+5.65}_{ -5.11}$ &WL   &  9 \\\\%ped07 i Abell 781      & 140.20117  & 30.47176  & 0.288  & 27.17 &  77.19 &  33.22  & $12.67^{+5.86}_{ -5.86}$ &WL   & 10 \\\\%dah06 j Zw3146         & 155.91515  &  4.18629  & 0.291  & 31.76 &  70.81 &  30.33  & $11.31^{+5.44}_{ -5.44}$ &WL   & 10 \\\\%dah06 j &            &           &        &       &        &         & $13.99^{+7.99}_{ -6.17}$ &WL   &  9 \\\\%ped07 i Zw1459.4+4240  & 225.34604  & 42.34448  & 0.292  & 27.67 &  81.26 &  29.89  & $10.59^{+6.00}_{ -6.00}$ &WL   & 10 \\\\%dah06 j Abell 1576     & 189.24684  & 63.18658  & 0.300  & 42.00 & 156.45 &  38.40  & $16.01^{+6.60}_{ -4.86}$ &WL   &  9 \\\\%ped07 i &            &           &        &       &        &         & $18.45^{+6.03}_{ -6.03}$ &WL   & 10 \\\\%dah06 j \\hline \\end{tabular} \\end{center} { Note for Method: X-ray stands for the mass determined by X-ray measurements; WL stands for weak lensing, WL+SL stands for weak lensing combined with strong lensing; References for mass estimates: [1] \\citet{xjw01}, [2] \\citet{rb02}, [3] \\citet{pap05}, [4] \\citet{app05}, [5] \\citet{svd+09}, [6] \\citet{vkf+06}, [7] \\citet{sa07}, [8] \\citet{ou08}, [9] \\citet{pd07}, [10] \\citet{dah06}, [11] \\citet{hoe07}, [12] \\citet{ckc09}, [13] \\citet{lrj+07}, [14] \\citet{zfb+07}, [15] \\citet{btu+05}, [16] \\citet{ub08}, [17] \\citet{bsk+07}, [18] \\citet{hsp06}. } \\end{minipage} \\end{table*}  \\begin{figure*} \\centering \\includegraphics[width = 17.cm]{fig1.eps} \\caption{Correlations between cluster mass $M_{\\rm vir}$ and richness $R$, summed luminosity $L_r$ and $GGN/r_{GGN}$ for 24 nearby clusters ($R\\ge8$ and $z\\lesssim0.1$). The solid line is the best fit as given in Equation~(\\ref{mr})--(\\ref{mggn}).} \\label{richmass} \\end{figure*}  \\begin{figure*} \\centering \\includegraphics[width = 17.cm]{fig2.eps} \\caption{Correlations between cluster mass $M_{\\rm vir}$ and richness $R$, summed luminosity $L_r$ and $GGN/r_{GGN}$ for 17 clusters in the redshift range $0.17<z<0.26$. The black dots are the clusters with more than three estimates of their masses. The solid lines are the same shown in Figure~\\ref{richmass}. The dashed line is the new scaling relation with the same slope but different offsets determined from the data.} \\label{richmass2} \\end{figure*} ",
        "conclusions": "\\label{discu}  \\begin{figure} \\centering \\includegraphics[width = 7.cm]{fig7.eps} \\caption{The value of $\\sigma_8$ (with $\\Omega_m=0.3$ fixed) from cluster masses based on the $M_{\\rm vir}$--$L_r$ relation varies with a possible systematic bias on mass conversion of $\\gamma=M_{\\rm vir}/M_{\\rm vir, true}$.} \\label{gamma} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width = 7.cm]{fig8.eps} \\caption{The value of $\\sigma_8$ (with $\\Omega_m=0.3$ fixed) from cluster masses based on the $M_{\\rm vir}$--$L_r$ relation varies with a possible systematic bias on the slope of the scaling relation by $\\Delta \\nu$.} \\label{slope} \\end{figure}  \\begin{table*} \\begin{minipage}{160mm} \\caption[]{Comparison of results on $\\sigma_8$--$\\Omega_m$ derived from cluster mass function (upper part) and cosmic microwave background (CMB) measurement (middle part). See Table 5 of \\citet{hss+07} for the results derived from weak lensing statistics.} \\begin{tabular}{lllll} \\hline \\mc{1}{c}{Reference}&\\mc{1}{c}{Sample}&\\mc{1}{c}{No. of clusters}&\\mc{1}{c}{$\\sigma_8$--$\\Omega_m$ relation}& \\mc{1}{c}{$\\sigma_8$ ($\\Omega_m=0.30$)}\\\\ \\mc{1}{c}{}    &\\mc{1}{c}{or method}&\\mc{1}{c}{or observation} & \\mc{1}{c}{} & \\mc{1}{c}{} \\\\ \\hline  \\citet{vl96}       & X-ray   &25  & $\\sigma_8=0.60\\Omega_m^{-0.59+0.16\\Omega_m-0.06\\Omega_m^2}$          &1.16 \\\\ \\citet{ecf96}      & X-ray   &25  & $\\sigma_8$=(0.52$\\pm$0.04)$\\Omega_m^{-0.52+0.13\\Omega_m}$          &0.93$\\pm$0.07 \\\\ \\citet{mar98}      & X-ray   &30  & $\\sigma_8=0.78\\pm0.04$ with $\\Omega_m=0.30$ fixed                &0.78$\\pm0.04$ \\\\ \\citet{pen98}      & X-ray   &25  & $\\sigma_8$=0.53$\\Omega_m^{-0.53}$                                &1.00 \\\\ \\citet{brt+99}     & X-ray   &70  & $\\sigma_8$=0.58$\\pm$0.06$\\Omega_m^{-0.47\\pm0.16\\Omega_m}$          &0.96 \\\\ \\citet{vl99}       & X-ray   &10  & $\\sigma_8$=0.56$\\Omega_m^{-0.47}$                                &0.99 \\\\ \\citet{bsb+00}     & X-ray   &25  & $\\sigma_8=0.96$ with $\\Omega_m=0.30$ fixed                       &0.96 \\\\ \\citet{oa01}       & X-ray   &69  & $\\sigma_8$=0.59$\\Omega_m^{-0.57+1.45\\Omega_m-3.48\\Omega_m^2+3.77\\Omega_m^3-1.49\\Omega_m^4}$&0.91\\\\ \\citet{wu01}       & X-ray   &25  & $\\sigma_8$=0.477$\\Omega_m^{-0.3-0.17\\Omega_m^{0.34}-0.13\\Omega_{\\Lambda}}$&0.87 \\\\ \\citet{brt+01}     & X-ray   &103 & $\\sigma_8$=0.66$^{+0.06}_{-0.05}$, $\\Omega_m=0.35^{+0.13}_{-0.10}$ &-- \\\\ \\citet{psw01}      & X-ray   &30  & $\\sigma_8$=$(0.495^{+0.034}_{-0.037})\\Omega_m^{-0.60}$             &$1.02^{+0.07}_{-0.08}$ \\\\ \\citet{vnl02}      & X-ray   &452 & $\\sigma_8$=0.38$\\Omega_m^{-0.48+0.27\\Omega_m}$                     &0.61 \\\\ \\citet{rb02}       & X-ray   &106 & $\\sigma_8$=0.43$\\Omega_m^{-0.38}$                                &0.68\\\\ \\citet{sel02}      & X-ray   &30  & $\\sigma_8(\\Omega_m/0.3)^{0.44}=0.77\\pm0.07$                      &0.77$\\pm$0.07 \\\\ \\citet{vkl+03}     & X-ray   &40  & $\\sigma_8=0.78^{+0.30}_{-0.06}$ with $\\Omega_m=0.35$ fixed        &--\\\\ \\citet{sbc+03}     & X-ray   &452 & $\\sigma_8$=0.711$^{+0.039}_{-0.031}$, $\\Omega_m=0.341^{+0.031}_{-0.029}$&--\\\\ \\citet{pbs+03}     & X-ray   &63  & $\\sigma_8$=0.77$^{+0.05}_{-0.04}$ with $\\Omega_m=0.30$ fixed      &0.77$^{+0.05}_{-0.04}$\\\\ \\citet{asf+03}     & X-ray   &111 & $\\sigma_8$=(0.508$\\pm$0.019)$\\Omega_m^{-0.253\\pm0.024}$           &0.69$\\pm$0.05\\\\ \\citet{hen04}      & X-ray   &51  & $\\sigma_8=0.66\\pm0.16$ with $\\Omega_m=0.30$ fixed               &0.66$\\pm$0.16 \\\\ \\citet{dah06}      & X-ray   &35  & $\\sigma_8(\\Omega_m/0.3)^{0.37}=0.67^{+0.04}_{-0.05}$              &$0.67^{+0.04}_{-0.05}$ \\\\ \\citet{rdn07}      & X-ray   &66  & $\\sigma_8=0.92^{+0.24}_{-0.19}$, $\\Omega_m=0.24^{+0.14}_{-0.09}$   &$0.84\\pm0.03$\\\\ \\citet{heh+09}     & X-ray   &48  & $\\sigma_8(\\Omega_m/0.32)^{0.30}=0.86\\pm0.04$ for $\\Omega_m\\le0.32$& 0.88$\\pm$0.04\\\\ &         &    & $\\sigma_8(\\Omega_m/0.32)^{0.41}=0.86\\pm0.04$ for $\\Omega_m\\ge0.32$& \\\\ \\citet{vkb+09}     & X-ray   &49  & $\\sigma_8(\\Omega_m/0.25)^{0.47}=0.813\\pm0.027$                   &0.75$\\pm$0.02\\\\ \\citet{wef93}      & Optical &    & $\\sigma_8=0.57\\Omega_m^{-0.56}$                                  &1.12 \\\\ \\citet{bf98}       & Optical & 3  & $\\sigma_8\\Omega_m^{0.29}=0.8\\pm0.1$                              &1.13$\\pm$0.14 \\\\ \\citet{gbg+98}     & Optical &152 & $\\sigma_8$=(0.60$\\pm$0.04)$\\Omega_m^{-0.46+0.09\\Omega_m}$          &1.01$\\pm$0.07 \\\\ \\citet{bdb+03}     & Optical &300 & $\\sigma_8\\Omega_m^{0.60}=0.33\\pm0.03$                            &0.68$\\pm$0.06\\\\ \\citet{ebc+06}     & Optical &    & $\\sigma_8$=0.25$\\Omega_m^{-0.92-4.5(\\Omega_m-0.22)^2}$             &0.78 \\\\ \\citet{rwr+10}     & Optical &13832 & $\\sigma_8(\\Omega_m/0.25)^{0.41}=0.83\\pm0.03$                     &0.76$\\pm$0.03\\\\ \\hline \\citet{kdn+09}     & CMB    & WMAP5   & $\\sigma_8=0.81\\pm0.03$, $\\Omega_m=0.27\\pm0.01$    &  \\\\ \\citet{ll09}       & CMB    & WMAP5   & $\\sigma_8=0.92\\pm0.04$, $\\Omega_m=0.32\\pm0.03$    &  \\\\ \\citet{ldh+10}     & CMB    & WMAP7   & $\\sigma_8=0.80\\pm0.03$, $\\Omega_m=0.26\\pm0.01$    &  \\\\ \\citet{rmc+04}     & CMB    & CBI     & $\\sigma_8=0.96^{+0.06}_{-0.07}$  ~~(68\\% confidence level)           & \\\\ \\citet{dhc+06}     & CMB    & BIMA    & $\\sigma_8=1.03^{+0.20}_{-0.29}$  ~~(68\\% confidence level)           & \\\\ \\citet{rab+09}     & CMB    & ACBAR   & $\\sigma_8=0.93^{+0.04}_{-0.05}$            & \\\\ \\citet{smw+09}     & CMB    & CBI     & $\\sigma_8=0.92^{+0.05}_{-0.05}$            & \\\\ \\citet{sgr+09}     & CMB    & Bolocam & $\\sigma_8<1.57$ ~~~~~~~~~~~(90\\% confidence level)           & \\\\ \\citet{vac+09}     & CMB    & BOOMERANG & $\\sigma_8<0.92$ ~~~~~~~~~~~(95\\% confidence level)           & \\\\ \\hline This work ($0.05<z<0.1$)&$R$          &56 & $\\sigma_8(\\Omega_m/0.3)^{0.42}=0.82\\pm0.04$                  &0.82$\\pm$0.04\\\\ &$L_r$        &   & $\\sigma_8(\\Omega_m/0.3)^{0.46}=0.90\\pm0.04$                  &0.90$\\pm$0.04\\\\ &$GGN/r_{GGN}$&   & $\\sigma_8(\\Omega_m/0.3)^{0.40}=0.83\\pm0.04$                  &0.83$\\pm$0.04\\\\ This work ($0.2<z<0.25$)&$R$          &810& $\\sigma_8(\\Omega_m/0.3)^{0.42}=0.85\\pm0.02$                  &0.85$\\pm$0.02\\\\ &$L_r$        &   & $\\sigma_8(\\Omega_m/0.3)^{0.46}=0.94\\pm0.02$                  &0.94$\\pm$0.02\\\\ &$GGN/r_{GGN}$&   & $\\sigma_8(\\Omega_m/0.3)^{0.39}=0.82\\pm0.02$                  &0.82$\\pm$0.02\\\\ \\hline \\end{tabular} \\label{comp} \\end{minipage} \\end{table*}  Cluster mass function can be accurately determined from a complete volume-limited sample. The scaling relations of cluster mass have been determined for three optical observations, cluster richness, summed luminosity and $GGN/r_{GGN}$. The scaling relations are then used to estimate cluster mass for two samples of rich clusters. We get cluster mass functions and fit them with a theoretical expression. Cosmological parameters are constrained in the form of $\\sigma_8(\\Omega_m/0.3)^{\\alpha}=\\beta$, with $\\alpha=$0.40--0.50 and $\\beta=$0.8--0.9. For $\\Omega_m=0.3$, we get $\\sigma_8=$0.8--0.9 using different mass tracers or using the rich cluster samples at different redshift ranges.  The $\\sigma_8$ values from the mass tracers of richness $R$ and $GGN/r_{GGN}$ obtained using both cluster samples are consistent, while $\\sigma_8$ values derived from $L_r$ are higher. This discrepancy may come from some potential systematic bias on the mass scaling relations. If the $M_{\\rm vir}$--$L_r$ relations for both samples are really unbiased, then the cluster masses tracer by richness $R$ and $GGN/r_{GGN}$ are systematically underestimated. However, it is hard to assess which one is a better mass tracer. Given the scarce of mass estimates from different methods for the same clusters in Table~\\ref{mass} for the scaling relations, it is also hard to estimate the systematic bias on these mass estimates due to different methods (X-ray or weak lensing).  In our work, one potential systematic bias may come from the conversion of cluster mass from measured radii to the virial radius. Here, we use $\\gamma=M_{\\rm vir}/M_{\\rm vir, true}$ to stand for the systematic bias of masses in Table~\\ref{mass}, where $M_{\\rm vir, true}$ stands for the true virial mass of a cluster. Assuming a $\\gamma$, we get $M_{\\rm vir, true}$ and then fit the mass function of clusters to obtain $\\sigma_8$. Figure~\\ref{gamma} shows the variation of $\\sigma_8$ (with $\\Omega_m=0.3$ fixed) as a function $\\gamma$ based on the $M_{\\rm vir}$--$L_r$ relation. We are only concerned about the cases $\\gamma\\ge 1$.  For example $\\gamma=1.3$, i.e., masses systematically overestimated by 30\\%, the values of $\\sigma_8$ are lower by about 10\\%. In fact, the deviation of $\\gamma$ from 1.0 is related to the uncertainty of intercept in the logarithm scaling relations in Equation~(\\ref{mr})--(\\ref{mggn2}). The other possible systematic bias on $\\sigma_8$ may come from the slope uncertainties of the scaling relations. Here, we illustrate the dependence of $\\sigma_8$ on the slope uncertainty, $\\Delta \\nu$. We only apply to the $M_{\\rm vir}$--$L_r$ relation, for example. Given a $\\Delta \\nu$, i.e., $M_{\\rm vir}=A\\,L_r^{1+\\nu+\\Delta \\nu}$, here $\\nu=1.49$ according to Equation~(\\ref{ml}) and (\\ref{ml2}), we fit the power law with the data in Figure~\\ref{richmass} and \\ref{richmass2} to get $A$, and then get the cluster mass function and fit for $\\sigma_8$. Figure~\\ref{slope} shows the $\\sigma_8$ value varies with $\\Delta \\nu$. We find that the $\\sigma_8$ from the cluster sample of $0.05<z<0.1$ does not change significantly with $\\Delta \\nu$, while the $\\sigma_8$ decreases from 1.05 to 0.81 for the cluster sample of $0.2<z<0.25$ when the slope varies by $\\Delta \\nu$ from -0.4 to 0.4.  We can compare our results of $\\sigma_8$ with previous determinations from cluster mass function, as listed in Table~\\ref{comp}. Most of previous results are based on X-ray flux-limited cluster samples. Our results are systematically larger than those from the mass function of X-ray clusters.  \\citet{rwr+10} used the largest number of clusters from SDSS maxBCG catalog to determine the amplitude of cluster mass function. They did not estimate the mass for each cluster, but gave a statistical mass for clusters within a richness bins by weak lensing. They got $\\sigma_8=0.76\\pm0.03$ assuming $\\Omega_m=0.30$. The maxBCG clusters were selected based on the red brightest cluster galaxies (BCGs). However, the maxBCG method may miss about 25\\% clusters in which the BCGs have emission line and blue colors \\citep{kma+07}. We notice that about 15\\% rich clusters ($R\\ge16.7$) are missing by the maxBCG method compared to our sample in the redshift range of $0.2< z<0.25$. However, the systematic incompleteness only induces an underestimate of 3\\% for $\\sigma_8$.  Therefore, the discrepancy probably comes from the uncertainty of mass scaling relations.  If we take $\\Omega_m=0.26$ derived from WMAP7, then our values of $\\sigma_8$ should become larger by a factor of $(0.26/0.3)^{\\sim0.42}=1.06$, roughly equal to adding 0.05 to the our $\\sigma_8$ value in Table~\\ref{comp}.  Therefore, the $\\sigma_8$ values we derived from galaxy clusters are slightly larger than the those from the WMAP data \\citep{kdn+09,ldh+10}. While some reanalysis of the WMAP5 data independently \\citep{lls+09} gives $\\sigma_8=0.921\\pm0.036$ for $\\Omega_m=0.32\\pm0.03$ \\citep[see][]{ll09}. Some studies of cosmic microwave background at small scales also give higher values of $\\sigma_8$ than that from WMAP \\citep{rmc+04, dhc+06,rab+09,smw+09}.  Our result of $\\sigma_8$ are consistent with many recent studies using other methods. For example, the $\\sigma_8$ by weak lensing method has a mean value of 0.85$\\pm$0.03 \\citep[see previous results in Table 5 of ][]{hss+07}, which is higher than previous results from X-ray clusters. \\citet{tbs+04} studied the power spectrum of galaxies from the SDSS to constrain cosmological parameters. They obtained $\\sigma_8=0.89\\pm0.02$ and $\\Omega_m=0.30\\pm0.03$. \\citet{lee09} studied the normalization of the power spectrum via the ellipticity function of giant galaxy voids from SDSS DR5 and obtained $\\sigma_8=0.90\\pm0.04$. \\citet{jnt+05} used the Ly$\\alpha$ data and found $\\sigma_8=0.9$ and $\\Omega_m=0.27$. \\citet{fjf+03} used the galaxy peculiar velocities to probe the growth rate of the structure and found that $\\sigma_8=1.13^{+0.22}_{-0.23}$ and $\\Omega_m=0.30^{+0.17}_{-0.07}$.  In this work, we get six values of $\\sigma_8$ by cluster mass function. Basically, the results are consistent. However, the precise value of $\\sigma_8$ is still to be determined since our constraint is not only coupled with $\\Omega_m$, but also has large uncertainties on the scaling relations."
    },
    "1004/1004.2331_arXiv.txt": {
        "abstract": "Abell 3667 is the archetype of a merging cluster with radio relics. The NW radio relic is the brightest cluster relic or halo known, and is believed to be due to a strong merger shock.  We have observed the NW relic for $\\sim 40$ ksec of net XMM time.  We observe a global decline of temperature across the relic from 6 to 1 keV, similar to the Suzaku results.  Our new observations reveal a sharp change of both temperature and surface brightness near the position of the relic.  The increased X-ray emission on the relic can be equivalently well described by either a thermal or nonthermal spectral model. The parameters of the thermal model are consistent with a Mach number ${\\cal{M}}\\sim2$ shock and a shock speed of $\\sim$1200 km s$^{-1}$.  The energy content of the relativistic particles in the radio relic can be explained if they are (re)-accelerated by the shock with an efficiency of $\\sim$0.2\\%.. Comparing the limit on the inverse Compton X-ray emission with the measured radio synchrotron emission, we set a lower limit to the magnetic field in the relic of 3 $\\mu G$.  If the emission from the relic is non-thermal, this lower limit is in fact the required magnetic field. ",
        "introduction": "\\label{sec:intro}  Many clusters of galaxies appear to be forming at the present time through massive cluster mergers. Major cluster mergers are the most energetic events which have occurred since the Big Bang, involving total energies of $\\sim$$10^{64}$ ergs. Merger shocks driven into the intracluster gas are the primary heating mechanism of the gas in massive clusters. Chandra and XMM-Newton X-ray observations have provided beautiful images and spectra of merger hydrodynamical effects, including cold fronts and merger bow shocks \\citep[e.g.,][]{Mar+00,MGD+02}. However, the number of clusters with well-observed merger shocks is limited.  \\begin{figure*} \\includegraphics[width=17.cm]{a3667_520_raw_v2.ps} \\caption{XMM-Newton mosaic image of A3667 in the 0.5-2 keV energy band. The gain in spatial coverage of the cluster due to the addition of our new northwest observation is evident. The image has been background subtracted and exposure corrected. \\label{fig:mosaic}} \\end{figure*}   Enigmatic extended cluster radio sources with very steep spectra and no clear optical counterparts have been known for over 30 years (\\citealt{Wil70}, for a review see \\citealt{Fer08}). These sources occur in a small fraction of rich galaxy clusters.  They have such steep radio spectra that they can be detected only at lower frequencies ($\\la$5 GHz).  Fairly symmetric sources which are projected on the cluster center are often referred to as ``radio halos'' \\citep[e.g., the Coma cluster; ][]{DRL+97}, while similar elongated sources usually located on the cluster periphery are called ``relics.''  The extremely steep radio spectra flatten somewhat at low frequencies, which suggest that the relativistic electrons have undergone significant losses due to synchrotron and inverse Compton (IC) emission.  In every case, such diffuse cluster radio sources have been found in irregular clusters which are apparently undergoing mergers.  This suggests that the radio emitting electrons are accelerated or re-accelerated by shocks or turbulence associated with these cluster mergers.  One possible theoretical picture is that the radio halos are accelerated by turbulence following the passage of merger shocks, while the relics are the direct result of merger shock acceleration \\citep[e.g.,][]{Fer08}.  The same relativistic electrons which produce the radio synchrotron radiation will produce hard X-ray (HXR) emission by inverse Compton scattering of Cosmic Microwave Background (CMB) photons.  However, the detection of this emission has been difficult and remains controversial \\citep[e.g., compare][]{FOB+04,RM04}.  In many ways, Abell 3667 is the ideal site to study mergers and radio relics.  It is a very bright X-ray cluster at a low redshift ($z = 0.0552$). The ROSAT and ASCA observations showed that it is a spectacular merger with shock heated gas \\citep{MSV99}.  The galaxy spatial and redshift distributions are also strongly bimodal \\citep{OCN09}.  A3667 was one of the first clusters in which a merger cold front was found with Chandra \\citep{VMM01a,VMM01b}, and the term was coined based on the A3667 observations.  The very sharp density and temperature gradients in this cold front have provided very strong limits on the role of transport processes (like heat conduction) in clusters.  Subsequent Chandra and XMM/Newton observations \\citep{MFV02,BFH04} have provided spectacularly detailed images and information on the dynamics of the merger.  However, most of the Chandra and XMM observations study the interior regions of the cluster, not the outer regions near the radio relics.  Abell 3667 contains a pair of curved cluster radio relics \\citep{RWH+97}. Their location on either side of the cluster center at 1.8Mpc distance ($29^\\prime$) and sharp, inwardly curved outer edges are exactly what is expected for merger shocks and shock particle acceleration.  Models which reproduce the optical and inner X-ray properties of A3667 predict shocks at or near these locations \\citep{RBS99,RS01}.  The radio spectra steepen with distance from the outer edge \\citep{RWH+97} as expected if the electrons are accelerated there, and the higher energy electrons lose energy due to synchrotron and IC emission as they are advected away from the shock.  The northwest radio relic in Abell 3667 is the brightest (highest flux) cluster radio relic or halo source which is known, with a flux at 20 cm of 3.7 Jy \\citep{Joh-Hol04}.  It is also one of the largest relics, with a total extent of 33\\arcmin\\ or 2.1 Mpc. Since the electrons which produce the IC HXR emission are basically the same ones which emit radio synchrotron, for a given magnetic field the HXR flux should be nearly proportional to the radio flux.  Thus, Abell 3667 might be expected to be the brightest nonthermal (NT) cluster hard X-ray source.  Since this relic is at a large projected radius from the cluster center, if anything the magnetic field should be lower than in other objects, implying an even larger HXR flux.  \\begin{figure*} \\includegraphics[width=17.cm]{a3667_wv_kmm_cb_2relics_300dpi_cross.ps} \\caption{Wavelet reconstruction of the 0.5-2 keV X-ray mosaic of A3667 from XMM. There is a sharp discontinuity in the X-ray surface brightness to the northwest;  The blue ellipses show the location, name and size of three main galaxy components (KMM5, KMM2 and KMM4) of A3667 from the analysis of Owers et al. (2009). The black cross marks the center of KMM2 group. The white contours show the wavelet reconstruction of the SUMMS 843 MHz radio image in the region of radio relics.  The contours are drawn at the 1, 3, 9 and 20 mJy/beam levels. \\label{fig:mosaic-wv}} \\end{figure*}  We have a long Suzaku observation of the NW radio relic region of Abell 3667 with Suzaku \\citep{Nak+09}, as well as two observations of more central regions of the cluster.  While hard X-ray emission is seen coming from the cluster center with HXD intrument \\citep{Nak+09}, it could be thermal emission, while there is no clear Suzaku detection of inverse Compton hard X-ray emission from the NW radio relic, given the systematic uncertainties in the non-X-ray background and thermal emission from the cluster.  Our Suzaku observation also provides a CCD X-ray image of a portion of the radio relic region with the X-ray Imaging Spectrometer (XIS).  However, this covers only a part of the region of the radio relic and possible merger shock.   Here, we present the results of a new XMM-Newton observation of Abell~3667 which is centered on the NW radio relic.  We assume a concordance cosmology with $H_0 = 71$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.27$, and $\\Omega_\\lambda=0.73$.  At the redshift of Abell 3667 of $z=0.05525$, the angular diameter distance is $d_A = 212.6$ Mpc, and 1\\arcmin\\ corresponds to 61.8 kpc.  For a mean cluster temperature of 7.2 keV, the virial radius should be $r_{180} \\approx 2.33$ Mpc $= 38\\arcmin$ \\citep{MFS+98}. Unless otherwise specified, we provide confidence intervals at the 68\\% level. ",
        "conclusions": "We have made a new X-ray observation of the northwest radio relic region in the Abell 3667 cluster in order to gain insight on both thermal and non-thermal phenomena that should be associated with it.  We detect a jump in both the surface brightness and the hardness of the X-ray emission at the outer edge of the relic.  The X-ray spectra of these regions is consistent with the extra emission on the relic being either thermal emission from a merger shock or inverse Compton emission from relativistic electrons in the radio relic.  There are some arguments that favor the shock model, which implies a shock with a Mach number of ${\\cal M} \\approx 2$ and a shock speed of $\\sim$1200 km s$^{-1}$.  The energy content of the relativistic particles in the radio relic can be explained if they are (re)-accelerated by the shock with an efficiency of $\\sim$0.2\\%.  The merger shock model and shock acceleration of the relativistic electrons can explain, at least roughly, the radio surface brightness and spectral distributions in the relic.  Alternatively, the jump in the brightness and hardness of the X-rays at the outer edge of the relic might be due to inverse Compton emission from the relic. Assigning all the emission from the relic to the IC component, we derive a lower limit on the magnetic field in the relic of $\\ga$3 $\\mu$G.  These observations have yielded the first robust characterization of the X-ray properties near the northwest radio relic and possible merger shock in Abell 3667. Unfortunately, the data are not sufficiently deep to determine if the relative roles the thermal shock emission and nonthermal IC emission. Recently, we were approved for a very long observation of this regions as an XMM-Newton Large Program.  These data should allow us to determine the efficiency of particle acceleration and the relative roles of thermal and nonthermal processes in this merger shock."
    },
    "1004/1004.2936_arXiv.txt": {
        "abstract": "We present observations of \\pwn\\ with the \\fermi-LAT. The source is detected with high confidence as an emitter of high-energy gamma-rays. The spectrum lacks any evidence for the characteristic cutoff associated with emission from pulsars, indicating that the emission arises primarily from the pulsar wind nebula. Broadband modeling implies an evolved nebula with a low magnetic field resulting in a high $\\gamma$-ray to X-ray flux ratio. The \\fermi\\ emission exceeds predictions of the broadband model, and has a steeper spectrum, possibly resulting from a distinct excess of low energy electrons similar to what is inferred for both the Vela X and Crab pulsar wind nebulae. ",
        "introduction": "The basic structure of a pulsar wind nebula (PWN) is determined by the spin-down energy injected by the central pulsar and the interaction of the nebula with the interior regions of the supernova remnant (SNR) in which it evolves. Losses from synchrotron radiation in the nebular magnetic field, whose strength depends both on the nature of the injected wind and on the evolving size of the PWN, inverse-Compton scattering of ambient photons by the energetic electron population within the nebula, and adiabatic expansion as the nebula sweeps up the surrounding supernova ejecta, all combine to determine emission structure and long-term evolution of the nebula. (See Gaensler \\& Slane 2006 for a review.) Multiwavelength observations of PWNe provide crucial information on the underlying particle spectrum which, in turn, strongly constrains both the magnetic field strength and the stage of evolution. Of particular importance is the spectrum of particles injected into the PWN. While this is typically assumed to be a power law, it has long been known that the spectrum of the Crab Nebula is not fully consistent with such a structure; changes in the spectral index in the optical band seem to indicate inherent features in the injection spectrum, and the large population of radio-emitting particles may have a completely distinct origin from that of the higher energy particles.  Complex structure in the observed spectrum of a PWN can originate in a number of ways that are associated with the long-term evolution, including synchrotron losses and interaction with the SNR reverse shock (Reynolds \\& Chevalier 1984; Gelfand, Slane, \\& Zhang 2009). In addition, recent studies of the spectrum immediately downstream of the wind termination shock show that, at least in some cases, the injection spectrum itself deviates significantly from a simple power law (Slane et al.  2008), and particle-in-cell simulations of the acceleration process produce spectra that are well-described by a Maxwellian population with a power law tail (Spitkovsky 2008).  Any such structure in the spectrum of the injected particles imprints itself on the broadband emission spectrum of the entire nebula. The resulting breaks or regions of curvature in the emission spectrum, and the frequencies at which they are observed, depend upon the energy at which features appear in the electron spectrum as well as the means by which the photons are produced (e.g. synchrotron radiation or inverse-Compton emission).  To fully understand the nature of the particle injection, as well as the long-term evolution of PWNe, it is thus crucial to study the emission structure over the entire electromagnetic spectrum.  Gamma-ray observations have provided an important window into the late-phase structure of PWNe. While the X-ray emission from older PWNe is relatively weak due to the long-term decline in magnetic field strength, inverse-Compton scattering of the cosmic microwave background (CMB), as well as other ambient photons, by the energetic particles in the nebula produce very energetic photons, providing unique discovery space for these systems. Even in younger nebulae, the $\\gamma$-ray emission provides a crucial probe of the particle spectrum in energy regions that are not accessible with other observations. The inverse-Compton emission from photons with a characteristic temperature $T$ peaks at an energy $\\epsilon_{ic} \\approx 5 \\gamma^2 kT$ for scattering from electrons of energy $\\gamma m_e c^2$.  CMB photons scattered into the bandpass of the \\fermi\\ Large Area Telescope (LAT), for example, originate from interactions with electrons in the approximate energy range 0.25 -- 4 TeV. Such particles produce synchrotron radiation with photon energies in the range $\\epsilon_{s} \\approx 0.03 - 7 B_{10}$~eV, where $B_{10}$ is the magnetic field strength in units of 10$\\mu$G. Depending on the magnetic field, which can range from $>100 \\mu$G or higher for young PWNe, down to $\\sim 5 \\mu$G for highly evolved systems, the synchrotron emission from particles in this energy range can be difficult to detect due to instrumentation limitations (at long radio wavelengths), absorption (in the optical/UV range) or confusion with the bright sky Galactic background (in the infrared).  \\fermi\\ measurements can thus provide a unique probe of the emission from a significant population of the particle spectrum in PWNe.  \\pwn\\ (see Figure 1) is an extended source of very high energy $\\gamma$-ray emission discovered with the High Energy Stereoscopic System (H.E.S.S.) during a survey of the Galactic plane (Aharonian et al.  2006).  Centered within the radio SNR G338.3$-$0.0 (Shaver \\& Goss 1970), the deconvolved TeV image of the source has an RMS width of $2.7 \\pm 0.5$~arcmin (Funk et al. 2007). HI measurements show absorption against G338.3$-$0.0 out to velocities corresponding to the tangent point, indicating a distance of at least 8 kpc (Lemiere et al. 2009), and thus implying a rather large size for the PWN ($R_{PWN} > 6.4 d_{10} $~pc, where $d_{10}$ is the distance in units of 10~kpc).  X-ray observations with \\xmm\\ (Funk et al. 2007) and Chandra (Lemiere et al. 2009) establish the presence of an accompanying X-ray nebula and an X-ray point source that appears to be the associated neutron star. In addition, Aharonian et al. (2006) noted the presence of the unidentified {\\em EGRET} source 3EG~J1639$-$4702 located at a nominal position 34$^\\prime$ from \\pwn, but the very large error circle on its position made any association with \\pwn\\ highly uncertain.  Here we report on observations of \\pwn\\ with the \\fermi-LAT.  The observations and data analysis are summarized in Section 2, and a discussion of the observed $\\gamma$-ray emission is discussed in the context of the evolutionary state of \\pwn\\ in Section 3. Our conclusions are summarized in Section 4.  \\begin{figure}[t] \\epsscale{1.15} \\plotone{f1.ps} \\epsscale{1.0} \\caption{ \\fermi\\ LAT image (2 - 200 GeV) of \\pwn. The cyan circle indicates the uncertainty in the centroid of the \\fermi\\ LAT source, the magenta circle indicates the 95\\% encircled flux distribution of the HESS image, and the white circle indicates the 95\\% probability contour for the position of 3EG J1639$-$4702. The white contours outline radio emission from G338.3$-$0.0 while the black contours at the center outline extended X-ray emission observed with \\xmm. A compact X-ray source detected with \\chandra\\ resides within the X-ray contours. } \\end{figure} ",
        "conclusions": "Broadband studies of \\pwn\\ have identified this source as a likely PWN, with X-ray observations providing images of an extended nebula as well as the putative pulsar powering the system. Modeling of the PWN evolution based on inferred parameters of the pulsar imply detectable emission in the GeV $\\gamma$-ray band, and our investigations of the \\fermi-LAT observations of this region reveal clear evidence of such emission, consistent with the source 1FGL~J1640.8$-$4634. The flux and spectrum we derive are consistent with that of the previously-identified source 3EG~J1639$-$4702, and the much-improved position makes it likely that the emission arises from \\pwn. The lack of a spectral cutoff rules out an association with emission directly from the pulsar that powers the nebula, and the flux is inconsistent with $\\gamma$-rays from \\snr, in which \\pwn\\ resides.  We have investigated the radio, X-ray, and $\\gamma$-ray emission from \\pwn\\ in the context of a simple one-zone model in which power is injected into a nebula at a time-dependent rate consistent with the current observed X-ray emission from the system. We find that models constrained by the observed size of the associated SNR, as well as limits on the size of the PWN, require an approximate age of 10~kyr and a current magnetic field strength of only $\\sim 4 \\mu$G, consistent with expectations for the late-phase evolution of a PWN. The conditions in such an evolved PWN yield a considerably higher $\\gamma$-ray flux, relative to the X-ray flux, than in younger systems where the higher magnetic field results in significant synchrotron radiation.  The observed \\fermi-LAT emission from \\pwn\\ significantly exceeds that predicted by our broadband models. We propose that the excess emission is a signature of a distinct population of low-energy electrons similar to that inferred from studies of the Crab Nebula and Vela~X, although the nature of this electron component is not well constrained. Deeper radio observations are needed to place stronger constraints on this population. Sensitive searches for pulsations from the central pulsar are of particular importance to constrain the spin-down properties of the system, which can only be very roughly constrained at present. There has been considerable success in identifying such pulsars with the \\fermi-LAT, but the lack of an obvious pulsar-like spectrum in \\pwn\\ may argue for more likely success with deep radio timing searches."
    },
    "1004/1004.2041_arXiv.txt": {
        "abstract": "In recent years argument has been made that a high fraction of early-type galaxies in the local universe experience low levels ($\\lesssim 1 M_{\\odot} {\\rm yr}^{-1}$) of star formation (SF) that causes strong excess in UV flux, yet leaves the optical colors red. Many of these studies were based on \\galex\\ imaging of SDSS galaxies ($z \\sim 0.1$), and were thus limited by its $5''$ FWHM. Poor UV resolution left other possibilities for UV excess open, such as the old populations or an AGN. Here we study high-resolution far-ultraviolet \\hst/ACS images of optically quiescent early-type galaxies with strong UV excess. The new images show that three-quarters of these moderately massive ($\\sim 5\\times 10^{10}M_{\\odot}$) early-type galaxies shows clear evidence of extended SF, usually in form of wide or concentric UV rings, and in some cases, striking spiral arms. SDSS spectra probably miss these features due to small fiber size. UV-excess early-type galaxies have on average less dust and larger UV sizes ($D>40$ kpc) than other green-valley galaxies, which argues for an external origin for the gas that is driving the SF. Thus, most of these galaxies appear `rejuvenated' (e.g., through minor gas-rich mergers or IGM accretion). For a smaller subset of the sample, the declining SF (from the original internal gas) cannot be ruled out. SF is rare in very massive early-types ($M_*>10^{11} M_{\\odot}$), a possible consequence of AGN feedback. In addition to extended UV emission, many galaxies show a compact central source, which may be a weak, optically inconspicuous AGN. ",
        "introduction": "In this paper we present a discovery of significantly extended regions of star formation in some early-type galaxies (ETGs)--galaxies usually thought to lie on the passive side of galaxy bimodality. Bimodality in terms of morphology and color has been known since the earliest studies of galaxies, but it was not until the massive datasets of the Sloan Digital Sky Survey (SDSS) that fuller implications in terms of galaxy evolution became evident \\citep{strateva,kauff03}. Optical colors reflect the mean age of stellar populations and are therefore sensitive only to high levels of continuous SF \\citep{kauff07}. If the SFR per unit stellar mass (specific SFR) drops below some threshold, optical colors become uniformly red and SDSS photometry cannot distinguish a truly passive galaxy from one that also contains a young population. These limitations are alleviated when ultraviolet (UV) photometry, dominated by young stars, is available. Early results from \\galex\\ showed that a surprisingly high fraction (15\\%) of optically red SDSS ETGs exhibit strong UV excess \\citep{yi}. \\citet{rich} found strong far-UV (FUV) excess even when selecting ETGs with no \\ha\\ emission in SDSS spectra. Is this UV excess due to star formation, as assumed by \\citet{yi,kaviraj07,kaviraj09}? While SF and molecular gas have been studied in nearby early-type galaxies for some time (e.g., \\citealt{bregman}), their significance as a possible {\\it phase} in galaxy evolution or a {\\it mode} of galaxy growth requires the large samples we have today. Before considering such far-reaching implications one must ask whether other explanations for the UV flux exist? After all, nearby ellipticals are known to exhibit a moderate UV excess (the ``UV upturn'', \\citealt{code}), that comes from old stars (presumably hot horizontal branch), and not massive young stars \\citep{greggio}. Also, a continuum from a weak AGN could in principle produce an UV excess \\citep{agueros}.  With $5''$ FWHM, \\galex\\ imaging marginally resolves SDSS galaxies at $z \\sim 0.1$ (angular diameter $<20''$), which is why we turned to {\\it high-resolution} FUV imaging with the Solar Blind Channel (SBC) of the ACS. \\hst\\ images of our sample of massive ETGs with strong UV excess and no obvious optical signs of SF reveal a surprise: they are dominated by {\\it extended} star formation on scales of 10--50 kpc, and with rates of up to 2 $M_{\\odot} {\\rm yr^{-1}}$. ",
        "conclusions": "Early-type galaxies with strong far-UV excess are dominated by SF, except when they are red in the ${\\rm NUV}-r$, in which case the FUV comes from a compact source, possibly an optically weak AGN. Extended SF around ETGs has recently been found in some {\\it individual} nearby galaxies, but these have stellar masses one to two orders of magnitude smaller \\citep{thilker10,donovan}. We extended SF in galaxies up to $M_*\\sim 10^{11} M_{\\odot}$. However, strong UV excess is rare in yet more massive ETGs--possible consequence of AGN feedback and a reflection of an `E-E' dichotomy. In the majority of cases the SF gas is probably external in origin (minor mergers or IGM accretion), making these ETGs `rejuvenated'. Most common UV morphology are rings, which likely result from resonances (from yet undetected bars or other disturbances). A smaller fraction of UV-excess ETGs may be more typical green-valley galaxies, where we see the declining original SF activity. Similar conclusions regarding the heterogeneity of evolutionary paths of transiting galaxies (but not necessarily ETGs) were recently reached by \\citet{cortese} based on HI content of nearby ($<25$ Mpc) galaxies."
    },
    "1004/1004.2088_arXiv.txt": {
        "abstract": "We investigated the dependence of the solar magnetic parity between the hemispheres on two important parameters, the turbulent diffusivity and the meridional flow, by means of axisymmetric kinematic dynamo simulations based on the flux-transport dynamo model. It is known that the coupling of the magnetic field between hemispheres due to turbulent diffusivity is an important factor for the solar parity issue, but the detailed criterion for the generation of the dipole field has not been investigated. Our conclusions are as follows. (1) The stronger diffusivity near the surface is more likely to cause the magnetic field to be a dipole. (2) The thinner layer of the strong diffusivity near the surface is also more apt to generate a dipolar magnetic field. (3) The faster meridional flow is more prone to cause the magnetic field to be a quadrupole, i.e., symmetric about the equator. These results show that turbulent diffusivity and meridional flow are crucial for the configuration of the solar global magnetic field. ",
        "introduction": "The solar magnetic $11$-year cycle is thought to be sustained by the dynamo motion of the internal ionized plasma \\citep{1955ApJ...122..293P}. Based on the internal structure of the velocity field, i.e., the meridional flow and the differential rotation revealed by helioseismology \\citep[see review by][]{2003ARA&A..41..599T}, the flux-transport dynamo was suggested \\citep{1995A&A...303L..29C,1999ApJ...518..508D,2001A&A...374..301K,2010ApJ...709.1009H}, as a model to successfully explain some features of solar activity such as the equatorward migration  of sunspots and the poleward migration of the surface field.\\par The solar global field has a distinct parity: the poloidal field is a dipole, i.e., antisymmetric about the equator. The polar fields almost always have the different sign between hemispheres, even though they show the occasional weak north-south asymmetry in phase and amplitude. In addition, Hale's polarity law states that the sunspots between hemispheres are nearly always antisymmetric about the equator \\citep{1908ApJ....28..315H}. It can then be interpreted that the toroidal fields ($B_\\phi$) below the surface are antisymmetric about the equator. This interesting feature is, however, not axiomatically explained by the flux transport dynamo model since this model significantly depends on three free parameters, i.e., the $\\alpha$-effect, the meridional flow, and the turbulent diffusivity. \\par It has been suggested that the $\\alpha$-effect around the base of the convection zone leads to the generation of the global dipolar magnetic field. \\citep{2001ApJ...559..428D,2002A&A...390..673B,2004A&A...427.1019C}. The existence of the poloidal fields around the tachocline and the coupling of these fields between hemispheres are significant factors for the generation of the dipole field. A detailed explanation of this process is given in the next paragraph. \\cite{2004A&A...427.1019C} also suggested, however, that the dipole field can be obtained with the strong diffusivity in the convection zone without the presence of the the $\\alpha$-effect around the tachocline. Hence the exact necessity of the $\\alpha$-effect in generating the dipole field is still inconclusive. \\par The dependence of the parity on these parameters can be explained when we understand the role of the turbulent diffusivity in the solar magnetic parity issue. If the global magnetic field is antisymmetric, i.e. is a dipole like our sun, the $\\phi$ component of the magnetic vector potential in each hemisphere has the same sign (Fig. \\ref{explain}a). When the cyclic phase in one hemisphere slightly differs from the other, the coupling effect by the turbulent diffusivity of the poloidal field distinguishes the phase difference in the vector potential and causes the magnetic field to be a dipole. On the other hand, when the magnetic field is symmetric, i.e., is a quadrupole, this effect does not occur. Therefore the substantial coupling of the poloidal field generates the antisymmetric (a dipole) magnetic field. The sign of the toroidal field in one hemisphere differs from that in the other hemisphere when the global magnetic field is a solar-like dipole. With the same logic posited above, it is obvious that the substantial diffusive coupling of the toroidal field between the hemispheres helps the magnetic field to become symmetric (a quadrupole; Fig \\ref{explain}b). In summary, the parity of the stellar global magnetic field depends on which field, the toroidal or the poloidal, is more coupled by the turbulent diffusivity between the hemispheres. Detailed systematic parameter studies are necessary to understand for the parity issue.\\par In this study, we investigate the dependence of the solar magnetic parity on the distribution of the turbulent diffusivity and the amplitude of the meridional flow. The obtained constraint on the turbulent diffusivity is important since it is one of the key components of the solar dynamo model, although it is difficult to measure by direct observations. ",
        "conclusions": "We investigated the dependence of the global magnetic parity on the distribution of the diffusivity (the amplitude and the surface depth) and the amplitude of the meridional flow. Three results were obtained. First, the model shows that the stronger diffusivity near the surface acts to make the magnetic field a dipole. The diffusivity near the surface enhances mainly the coupling of the poloidal field near the surface between the hemispheres, leading to the generation of dipolar magnetic field. The second result is that the thinner layer of the strong surface diffusivity also works to cause the magnetic field to become dipolar. The thinner surface depth suppresses the coupling of the toroidal field between the hemispheres since most of the toroidal field exists around the tachocline. The third result is that the fast meridional flow causes the magnetic field to become a quadrupole. The fast meridional flow prevents the poloidal field from coupling  near the surface of the equator because the flow transports the poloidal field poleward. In addition, the flow transports the toroidal field around the tachocline equatorward, thus causing the coupling of the toroidal field. These three results quantitatively constrain the distribution and the amplitude of turbulent diffusivity, which cannot be determined by observation and is a important factor for the dynamo problem. \\par In this study, we did not investigate the dependence of the parity on the $\\alpha$-effect in the convection zone, which may be a strong factor in causing the magnetic field to become a dipole. The poloidal field generated by this effect around the tachocline is transported equatorward by the meridional flow, and this process enhances the coupling of the poloidal field between the hemispheres \\citep{2001ApJ...559..428D,2002A&A...390..673B,2004A&A...427.1019C}. It is possible that the criterion for a dipole field we obtain in this study may be modified with this type of $\\alpha$-effect. We will discuss the possibility of the existence and the influence of the $\\alpha$-effect in a forthcoming paper. Another interesting issue to be addressed is the possibility that the variation of the velocity field in the solar cycle affects the parity. In the calculations for the earth dynamo there is the significant difference between the kinematic and the MHD cases in the parity issue \\citep{nishikawa2008simulation}. Thus, in the future we will investigate the parity issue with the Lorentz feedback \\citep{2006ApJ...647..662R}."
    },
    "1004/1004.2263_arXiv.txt": {
        "abstract": "We report the first results of imaging the Carina Nebula (NGC~3372) with the Infrared Array Camera (IRAC) onboard the {\\it Spitzer Space Telescope}, providing a photometry catalog of over 44,000 point sources as well as a catalog of over 900 candidate young stellar objects (YSOs) based on fits to their spectral energy distributions (SEDs).  We discuss several aspects of the extended emission, including the structure of dozens of dust pillars that result when a clumpy molecular cloud is shredded by feedback from massive stars. There are surprisingly few of the ``extended green objects'' (EGOs) that are normally taken as signposts of outflow activity in {\\it Spitzer} data, and not one of the dozens of Herbig-Haro jets detected optically are seen as EGOs.  EGOs are apparently poor tracers of outflow activity in strongly irradiated environments, due to the effects of massive star feedback.  A population of ``extended red objects'' tends to be found around late O-type and early B-type stars, some with clear bow-shock morphology.  These are dusty shocks where stellar winds collide with photoevaporative flows off nearby clouds.  Finally, the relative distributions of O-type stars, small star clusters, and sub-clusters of YSOs as compared to the dust pillars shows that while some YSOs are located within dust pillars, many more stars and YSOs reside just outside pillar heads. We suggest that pillars are transient phenomena, part of a continuous outwardly propagating wave of star formation driven by feedback from massive stars.  As the pillars are destroyed, they leave newly formed stars in their wake, and these are then subsumed into the young OB association.  The YSOs are found predominantly in the cavity between pillars and massive stars, arguing that their formation was in fact triggered.  Altogether, the current generation of YSOs shows no strong deviation from a normal initial mass function (IMF).  The number of YSOs is consistent with a roughly constant star-formation rate over the past $\\sim$3 Myr, implying that propagating star formation in pillars constitutes an important mechanism to construct unbound OB associations.  These accelerated pillars may give birth to massive O-type stars that, after several Myr, could appear to have formed in isolation. ",
        "introduction": "Most stars form in OB associations, where the stellar winds, UV radiation, and eventual supernova (SN) explosions from the massive members significantly impact the environment.  The feedback from newly born massive stars will ultimately destroy their natal molecular cloud, clearing away the dust and gas and thereby shutting off star formation and determining the star formation efficiency.  In the process, the same feedback may simultaneously trigger the birth of new generations of stars, allowing star formation to propagate continuously from one point to the next (Elmegreen \\& Lada 1977). This distributes the star formation in space and time, giving rise to significant age spreads in the resulting stars.  This mode of star formation may give rise to large OB associations, rather than compact single star clusters.  The southern Milky Way provides a striking example of this contrast: the Carina Nebula and the cluster NGC~3603 have roughly the same number of O-type stars and roughly the same ionizing photon output (Smith 2006; Crowther \\& Dessart 1998), but NGC~3603 is dominated by a single dense cluster within a radius of $\\sim$1 pc, whereas the O-type stars in Carina are distributed in several sub-clusters over 10--20 pc.  Evidence suggests that the dust pillars commonly seen at the borders of H~{\\sc ii} regions may be prime sites for this mode of propagating star formation (Bally \\& Reipurth 2003; McCaughrean \\& Andersen 2002; Jiang et al.\\ 2002; Stanke et al.\\ 2002; Smith et al.\\ 2005; Rathborne et al.\\ 2004).  Whether or not the second generation of stars forming in the pillars was triggered is still unclear: it is difficult to verify whether the stars began to form spontaneously due to initial clumps in the cloud and were simply uncovered by the advancing ionization front, whether their parent clumps were agglomerated in the ``collect and collapse'' scenario (Elmegreen 1992) where these clumps subsequently formed stars under their own self-gravity, or if they were triggered directly via radiation-driven implosion (Oort \\& Spitzer 1955; Kahn 1969; Dyson 1973; Elmegreen 1976; Bertoldi 1989; Bertoldi \\& McKee 1990; Williams et al.\\ 2001; Gorti \\& Hollenbach 2002).  Star formation in these environments plays a critical role in the evolution of H~{\\sc ii} regions and the interstellar medium (ISM) (e.g., Elmegreen \\& Scalo 2004).  The Carina Nebula is a special case among massive star-forming regions in the Milky Way: like the Orion Nebula, it is near enough (2.3 kpc; Smith 2006) to facilitiate studies of the details of faint nebular emission and the population of low to intermediate mass stars forming alongside massive stars, but unlike Orion, Carina has over 70 O-type stars and samples the top end of the stellar mass function.  This provides a large region over which to investigate the impact of massive-star feedback.  The global properties of the Carina Nebula have been discussed recently by Smith \\& Brooks (2007).  See Smith (2006) for a census of the massive stars that power Carina, and see Walborn (1995, 2009) for excellent reviews of the remarkable collection of extreme O-type stars in the region.  The stellar content of the massive central clusters Trumpler (Tr) 14 and 16 includes $\\eta$~Carinae, three WNH stars, and a number of the most extreme early O-type stars known (e.g., Walborn et al.\\ 2002a).  In fact, Carina is the first region where the earliest spectral type O stars were recognized (Walborn 1973).  All other regions in the Milky Way with a comparable stellar population are more distant and more obscured by dust.  Thus, Carina provides a laboratory for low-mass star formation and protoplanetary disks in regions analogous to much more extreme starbursts that are far too distant for such studies.    \\begin{figure}\\begin{center} \\includegraphics[width=3.1in]{fig01.eps} \\end{center} \\caption{A large-scale 8.6~$\\mu$m image of the Carina Nebula obtained with {\\it MSX}, from Smith \\& Brooks (2007).  The regions mapped with IRAC are shown.  The clusters Tr~14 and Tr~16 are located between these two fields amid the brightest diffuse emission.}\\label{fig:map} \\end{figure}  \\begin{figure*}\\begin{center} \\includegraphics[width=6.7in]{fig02.eps} \\end{center} \\caption{Color image of the South Pillars made from IRAC data with 3.6 $\\mu$m (blue), 4.5 $\\mu$m (green), 5.8 $\\mu$m (orange), and 8.0 $\\mu$m (red).  Several previously studied features are labeled.  The size scale and field of this image are shown in Figure~\\ref{fig:map}.  Tr~14 and Tr~16 are located off the top of this page.  The center of this field is at roughly 10$^{\\rm h}$47$^{\\rm m}$, --60$\\arcdeg$05$\\arcmin$ (J2000). }\\label{fig:colorSP} \\end{figure*}   Although a great number of massive stars have formed in the past $\\sim$3 Myr in Carina, the current hot-bed of star formation activity has migrated to the southern part of the nebula, in a region called the South Pillars (Smith et al.\\ 2000).  The widespread ongoing star-formation activity in the South Pillars, as well as the structure of the pillars themselves, was first recognized based on wide-field IR images (Smith et al.\\ 2000) obtained with the {\\it Midcourse Space Experiment} ({\\it MSX}).  (There appear to be several other regions of ongoing star formation to the north and west as well, which are less active than the South Pillars.)  An 8.6 $\\mu$m {\\it MSX} image of the Carina Nebula is shown in Figure~\\ref{fig:map}.  The presence of active, ongoing star formation in this region has since been confirmed by obervations of embedded young clusters like the ``Treasure Chest'' and others (Smith et al.\\ 2005; Rathborne et al.\\ 2002, 2004; H\\\"agele et al.\\ 2004) and a large number of outflows in the form of irradiated Herbig-Haro (HH) jets throughout the region (Smith et al.\\ 2004a, 2010).  There are many dense cores detected in C$^{18}$O that are spread across the region, providing sites of ongoing and potential future star formation, with some evidence for molecular outflow activity (Yonekura et al.\\ 2005).  The South Pillars are intriguing because they provide a powerful laboratory in which to study the details of the feedback mechanism by which massive stars destroy their natal molecular clouds.  In the process, the massive stars influence subsequent generations of stars and planets forming in the surrounding region.  Perhaps the birth of some of those newly formed stars was even directly triggered by the external feedback.  The massive stars also clear away and illuminate the surrounding gas as we witness the assemblage of an OB association. Study of regions like the South Pillars, where young stars are forming in the immediate vicinity of several dozen stars that will explode as supernovae in the next 1--2 Myrs, may have implications for our own Solar System (see Smith \\& Brooks 2007).  In this paper we present the first results from observations of the Carina Nebula obtained with the {\\it Spitzer Space Telescope} (Werner et al.\\ 2004; Gehrz et al.\\ 2007).  This is the first systematic study of the IR properties of the stars, dust, and gas in the South Pillars of Carina, so our aim here is to give a broad overview of the region as seen at 3-8 $\\mu$m; more detailed multi-wavelength studies will follow.  We present our observations in \\S 2, we discusss the photometry of point sources and the distribution of observed properties among the point source spectral energy distributions (SEDs) in \\S 3, and in \\S 4 we discuss various properties of the extended gas and dust emission and how it relates to nearby stars.  This includes extended green objects (\\S 4.2), extended red objects (\\S 4.3), relationships between dust pillars and massive stars (\\S 4.4 and 4.5), and the discussion of a newly recognized star cluster near Tr~16 that was obscured behind a dark dust lane at visual wavelengths (\\S 4.6). In \\S 5 we synthesize the main results, and in \\S 6 we recap with a list of our main conclusions.  There are several additional and more detailed results and implications from these data, which cannot be described in full here, but will be included in future papers.   \\begin{figure}\\begin{center} \\includegraphics[width=3.0in]{fig03.eps} \\end{center} \\caption{Same as Figure~\\ref{fig:colorSP}, but showing the western field with an adjusted intensity scale.  The greenish/red region at the top was at the edge of the field and was observed in Bands 2 (4.5 $\\mu$m) and 4 (8.0 $\\mu$m), but not 1 (3.6 $\\mu$m) and 2 (5.8 $\\mu$m).  We include it here because it contains interesting structure near Tr~14.  The orientation and field of view are shown in Figure~\\ref{fig:map}.}\\label{fig:colorWest} \\end{figure} ",
        "conclusions": "We have analyzed multiwavelength images of two regions in the Carina Nebula obtained with the IRAC camera onboard {\\it Spitzer}, including both point sources and the diffuse emission from dust and PAHs.  We have conducted photometry and produced a point source catalog merged with data from the 2MASS point source catalog, and we have used this to produce a list of highly reliable YSOs selected on the basis of fits to their SEDs.  Here we provide a brief list of several conclusions from this work.  1.  We provide a merged IRAC+2MASS catalog of over 44,000 point sources detected in at least 4 of 7 filters.  Most of these are foreground and background sources.  We identify 909 YSOs, selected based on fits to their SEDs.  We note that this is a severe underestimate of the true number of YSOs for several reasons, including our rigorous selection criteria for point sources.  2.  Our YSO sample is also an underestimate of the true number of YSOs because it misses faint sources due to sensitivity at stellar masses below 2--3 $M_{\\odot}$.  Correcting for these missing sources using an Orion-like IMF (Muench et al.\\ 2002), we find that there should be more than 5,000 YSOs in the South Pillar region of Carina. Considering that even this is an underestimate, we find it likely that star formation has continued at a relatively constant level over the 3 Myr lifetime of the region.  The current star formation occurring in dust pillars may therefore represent an important mode for the continual buildup of large OB associations.  3.  Comparing the $K$-band and 3.6 $\\mu$m luminosity functions (LFs) of our YSO sample to the LFs from Orion, we find no compelling evidence that the IMF is different among the collective generation of stars currently forming in Carina.  There are slight deficits in Carina at the bright end of the LFs, which are subject to low number statistics from the Orion sample, but we speculate that the apparent deficit of more luminous YSO sources in Carina compared to Orion may result if more massive stars shed their own disks more quickly than lower mass stars.  The 10--20 O-type stars that are spatially associated with the YSOs in the South Pillars probably formed along with them, and none of these have dusty disks.  4.  A corollary to the previous point is that while the aggregate population shows no convincing deviations from a standard IMF, there do appear to be possible fluctuations in the IMF from one sub-cluster to the next.  This is indicated by the apparent ratio of the numbers of O-type stars to associated YSOs, although this may also be an age effect.  5.  We detect surprisingly few of the so-called extended green objects (EGOs; thought to be molecular outflows), given the large number of YSOs and the known outflow activity traced by optical HH jets (Smith et al.\\ 2010).  We attribute this lack of EGO sources to photodissociation of the molecular outflows by the UV radiation field in Carina and to the added difficulty of detecting excess Band 2 flux amid the bright background emission in Carina.  6.  A population of extended red objects (EROs) exhibit diffuse emission that is not consistent with the colors expected for PAH emission.  The EROs are found to be associated with late O-type or early B-type stars, and several show a bow-shock morphology with the apex of the shock pointing inward to the first generation massive stars.  We suggest that these EROs represent thermal dust emission in shocks at the interface between stellar winds and dense photoevaporative flows in the South Pillars, akin to the similar structures seen in M~17 (Povich et al.\\ 2008).  7.  Judging by qualitative aspects of the observed structures of pillars seen in PAH emission, their filamentary structure is consistent with shocked or photoablated clouds in advanced stages of destruction.  The smaller pillars have probably been accelerated outward from their initial positions.  8.  We have analyzed the directions of the pillar axes for the ensemble of pillars in Carina.  While large pillars generally point inward to the first generation O-type stars in Tr~14 and Tr~16, many of the smaller pillars and cometary clouds point in other directions. We suggest that as a cloud is sculpted into the shape of a pillar by the first generation stars, it may also succumb to the influence of local O-type stars that were born in the second generation or that moved into its vicinity, effectively bending the pillar to a different orientation.  Local O-type stars may also be more influential when the most massive stars like $\\eta$ Car reach the ends of their lives and their ionizing flux drops.  9.  The relative spatial distribution of YSOs and dust pillars reveals that while several Stage~I YSOs are indeed located within the heads of dust pillars, many more Stage~I, II and III YSOs are scattered outside of pillars.  The YSOs show some small-scale clumping, but in general they form a large association of YSOs occupying a cavity that is bounded by pillars.  In several cases, smaller sub-clusters of YSOs are found just interior to the current locations of the heads of dust pillars.  10. We draw attention to one subcluster of stars, which was shown previously to have several X-ray sources (Sanchawala et al.\\ 2007). We argue that this cluster (Tr 16 SE) is not an embedded cluster, but is instead a young cluster that is obscured behind (not within) the dark dust lane that bisects the Carina nebula, with a likely age of $\\sim$1--2 Myr.  11. We propose a scenario where pillars are transient features in a continually outwardly propagating wave of star formation.  As the dust pillars are accelerated outward and destroyed by feedback from first-generation massive stars, they also form new stars in the process and leave behind a wake of YSOs.  The YSOs formed over time by an ensemble of pillars are subsumed into a young OB association formed over a span of $\\sim$10$^6$ yr, but the youngest ones still show some hierarchical sub-clustering and spatial association near the heads of pillars.  12.  The current star formation rate compared to the existing population of massive stars in the Carina Nebula is roughly consistent with a relatively {\\it constant} rate of star formation averaged over $\\sim$3 Myr.  Roughly 10--20 of the region's 70 O stars are closely associated with gas and dust pillars in the South Pillar region, and appear to have formed there recently (in the last $\\sim$1 Myr, rather than 3 Myr ago like Tr16), based on the Stage II and III YSOs that they are associated with).  This is consistent with our finding above that there is no compelling evidence for a substantially altered IMF. Thus, the current star formation in the Pillars of Carina is likely to represent an important mode of star formation for the gradual buildup of an OB association.  13.  We find that the fledgling OB association is likely to be unbound.  When all star formation in the region has ceased and the gas is removed, we will likely be left with a dispersing OB association where the massive stars have an age spread of 3--4 Myr.  There will be a ``halo'' of stars spread across $\\sim$100 pc, surrounding a pair of central clusters (Tr~14 and 16) that will be marginally bound; altogether it may appear much like h and $\\chi$ Persei (Currie et al.\\ 2009).  The imminent onset of the supernova driven phase may impact this picture, however, as discussed by Smith \\& Brooks (2007). We suggest that the nature as an unbound association is likely related to its formation mechanism in an outwardly propagating wave distrubuted over space rather than in a single compact cluster.  14.  Feedback-driven star formation may result in a substantial fraction of the massive-star content being spatially dispersed, spread across $\\sim$100 pc after several Myr.  Migrating from their birth site, these massive stars may falsely appear to have formed in isolation, and their SNe will not necessarily be coincident with an obvious H~{\\sc ii} region or star cluster. More detailed consideration of these points will be discussed in a future paper.    \\smallskip\\smallskip\\smallskip\\smallskip \\noindent {\\bf ACKNOWLEDGMENTS} \\smallskip \\scriptsize  This work was based on observations made with the Spitzer Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA. Support for this work was provided by NASA through awards issued by JPL/Caltech as part of GO programs 3420 and 20452.% M.S.P.\\ is supported by an NSF Astronomy \\& Astrophysics Fellowship under award AST-0901646.  B.A.W.\\ was supported by NASA through the Spitzer Space Telescope Theoretical Research Programs, through a contract issued by the JPL/Caltech under a contract with NASA. R.D.G.\\ was supported by NASA through contracts No.\\ 1256406 and 1215746 issued by JPL/Caltech to the University of Minnesota."
    },
    "1004/1004.2922_arXiv.txt": {
        "abstract": "We have recently completed a 64-night spectroscopic monitoring campaign at the Lick Observatory 3-m Shane telescope with the aim of measuring the masses of the black holes in 12 nearby ($z < 0.05$) Seyfert~1 galaxies with expected masses in the range $\\sim 10^6$--$10^7$~M$_{\\odot}$ and also the well-studied nearby active galactic nucleus (AGN) NGC\\,5548.  Nine of the objects in the sample (including NGC\\,5548) showed optical variability of sufficient strength during the monitoring campaign to allow for a time lag to be measured between the continuum fluctuations and the response to these fluctuations in the broad H$\\beta$ emission, which we have previously reported.  We present here the light curves for the H$\\alpha$, H$\\gamma$, \\ion{He}{2} $\\lambda 4686$, and \\ion{He}{1} $\\lambda 5876$ emission lines and the time lags for the emission-line responses relative to changes in the continuum flux.  Combining each emission-line time lag with the measured width of the line in the variable part of the spectrum, we determine a virial mass of the central supermassive black hole from several independent emission lines.  We find that the masses are generally consistent within the uncertainties.  The time-lag response as a function of velocity across the Balmer line profiles is examined for six of the AGNs.  We find similar responses across all three Balmer lines for Arp\\,151, which shows a strongly asymmetric profile, and for SBS\\,1116+583A and NGC\\,6814, which show a symmetric response about zero velocity.  For the other three AGNs, the data quality is somewhat lower and the velocity-resolved time-lag response is less clear.  Finally we compare several trends seen in the dataset against the predictions from photoionization calculations as presented by \\citeauthor{korista04}. We confirm several of their predictions, including an increase in responsivity and a decrease in the mean time lag as the excitation and ionization level for the species increases.  Specifically, we find the time lags of the optical recombination lines to have weighted mean ratios of $\\tau{\\rm (H\\alpha)} : \\tau{\\rm (H\\beta)} : \\tau{\\rm (H\\gamma)} : \\tau$(\\ion{He}{1}) $ : \\tau$(\\ion{He}{2}) $ = 1.54 : 1.00 : 0.61 : 0.36 : 0.25$.  Further confirmation of photoionization predictions for broad-line gas behavior will require additional monitoring programs for these AGNs while they are in different luminosity states. ",
        "introduction": "Active galactic nuclei (AGNs) are some of the most energetic objects in the Universe, radiating at luminosities above $10^{42}$\\,erg\\,s$^{-1}$, and yet their continuum emission is known to vary on timescales as short as days.  The size constraints set by such rapid variability mean that the extreme energy output of AGNs, often comparable to or more than the energy output of all the stars in a typical galaxy, must originate within a region whose size is $\\sim 0.01$\\,pc (approximately the size of our Solar System).  This large amount of energy arising from such a small region is theorized to be the result of gravitational accretion onto a supermassive black hole (e.g., \\citealt{rees84}).  For even the nearest AGNs, the region in which the continuum emission arises is only microarcseconds in angular size and is therefore unresolvable with current imaging detectors.  Dedicated monitoring programs have instead taken advantage of the fast, and often dramatic, variability of AGNs to completely revise our understanding of the physical conditions present in the gas on these small scales.  Early monitoring programs with monthly sampling found that variations in the broad emission lines promptly followed variations in the continuum flux, putting an upper limit on the size of the broad-line region (BLR) of only a light-month for typical Seyfert galaxies (e.g., NGC\\,4151: \\citealt{antonucci83}; Ark\\,120: \\citealt{peterson85}). Especially in the case of Ark\\,120, this upper limit was surprising, as the size of the BLR was expected to be an order of magnitude larger, based on photoionization models (e.g., \\citealt{kwan81,ferland82}).  Higher temporal sampling has since confirmed the size of the BLR for typical nearby Seyferts to be only a few light-days.  In addition, densely sampled monitoring programs have discovered that higher ionization lines respond more promptly (and more strongly) to continuum variations than lower ionization lines (e.g., \\citealt{clavel91}), indicating radial ionization stratification throughout the BLR, contrary to the previous single-cloud models where all emission lines were thought to arise from the same location.  More recent models such as the ``locally optimally emitting cloud'' (LOC) model \\citep{baldwin95} predict ionization stratification as a natural outcome.  In the LOC model, a range of cloud parameters is present in the BLR and the emission that we happen to see as observers arises from selection effects working within the BLR such that the majority of the emission from a specific line will come from a location where the parameters are most conducive to the production of that line.  A further discovery of monitoring programs is that the BLR appears to be virialized; the distance to a specific region in the BLR is inversely proportional to the square of the gas velocity in that region.  This was first conclusively shown for the most well-studied AGN, NGC\\,5548 (\\citealt{peterson99}; see also \\citealt{krolik91}), where $\\tau \\propto v^{-2}$, with $\\tau$ the broad emission line time lag relative to changes in the continuum flux (i.e., the BLR light-crossing time), and $v$ the velocity width of the broad line. Subsequent studies have also shown this to be true for several additional AGNs (e.g., \\citealt{onken02,kollatschny03}).  This behavior is consistent with the fact that the BLR gas is under the gravitational dominance of the central supermassive black hole, and so the response of the BLR gas can be used to learn about the mass of the black hole.  To date, black hole masses have been determined for some 44 AGNs (\\citealt{peterson04,peterson05,bentz09c}).  The most recent additions come from the Lick AGN Monitoring Program (LAMP), a dedicated 64-night spectroscopic monitoring campaign using the Lick Observatory 3-m Shane telescope and supplemented by four small-aperture telescopes employed in photometric monitoring.  First results from LAMP were presented by \\citet{bentz08} (hereafter Paper~I), followed by a full presentation of the photometric light curves (\\citealt{walsh09}; hereafter Paper~II) and the H$\\beta$ light curves and analysis (\\citealt{bentz09c}; hereafter Paper~III), and a re-examination of the \\msigma\\ relationship for AGNs (\\citealt{woo10}; hereafter Paper~IV). In this work, we present the light curves and analysis for the additional broad optical emission lines in the LAMP sample, namely H$\\alpha$, H$\\gamma$, \\ion{He}{2} $\\lambda 4686$, and \\ion{He}{1} $\\lambda 5876$.  We compare the results for these optical emission lines with results from previous monitoring campaigns, as well as with recent theoretical predictions of BLR behavior based on photoionization models. ",
        "conclusions": "The LAMP sample of AGNs was originally chosen for spectroscopic monitoring in order to extend to lower masses the range of black hole scaling relationships in AGNs.  With the high-quality spectroscopic dataset obtained at Lick Observatory, we are able to go beyond the original goals of LAMP and begin to examine the details of the BLR physics in these AGNs in the following ways:  \\begin{itemize}  \\item We have presented time-delay measurements and line widths for all of the optical H and He recombination lines in the spectra of the LAMP sample of AGNs: H$\\alpha$, H$\\beta$, H$\\gamma$, \\ion{He}{2} $\\lambda 4686$, and \\ion{He}{1} $\\lambda 5876$.  \\item Comparisons of the black hole masses determined from multiple emission lines are consistent within individual sources, even when assuming a single scaling factor $f$.  For at least the optical recombination lines, it appears that the scaling factor is not heavily dependent on the specific emission line when determining black hole masses from reverberation mapping.  \\item The time lag versus the line-width measurements for multiple emission lines in an individual source are generally consistent with a virial relationship ($\\tau \\propto v^{-2}$).  Virial relationships have been seen in other AGNs with high-quality spectroscopic monitoring data, upholding the use of reverberation-mapping results as a probe of the gravitational influence of the supermassive black hole on the BLR gas.  \\item For six of the LAMP AGNs, we have examined the velocity-resolved time-lag response across the broad H$\\alpha$, H$\\beta$, and H$\\gamma$ lines.  In three of the AGNs, we find a significant trend in the delay versus the velocity across the line profiles of all three Balmer lines.  In the other three AGNs, there is no significant trend in delay across the line profile, which may, in fact, argue for evidence of circular motions in a Keplerian potential.  We are currently investigating whether more detailed decompositions of the velocity-resolved time-lag response in these objects may be accomplished using the maximum entropy method \\citep{horne91,horne94}.  \\item We are able to confirm several trends in the behavior of the broad optical recombination lines that are expected from recent photoionization calculations and have also typically been seen in other AGN monitoring campaigns.  Specifically, we confirm an increase in responsivity and a decrease in the mean time lag as the excitation and ionization level for an emission line increases. This is manifest as $\\tau{\\rm (H\\alpha)} > \\tau{\\rm (H\\beta)} > \\tau{\\rm (H\\gamma)} > \\tau$(\\ion{He}{1}) $ > \\tau$(\\ion{He}{2}) and $\\eta{\\rm (H\\alpha)} < \\eta\\rm{ (H\\beta)} < \\eta{\\rm (H\\gamma)} < \\eta$(\\ion{He}{1})$ < \\eta$(\\ion{He}{2}).  Agreement with these photoionization calculations argues for optical-depth effects that appear to ``fine tune'' the responses of the optical recombination lines, as expected under the LOC model for AGN BLRs.  \\end{itemize}  Many of the additional predictions of \\citet{korista04} for optical recombination lines in AGN BLRs require multiple monitoring campaigns of multiple emission lines from a single AGN in different flux states. The investment of time to examine these predictions is both warranted and necessary.  The optical recombination line emissivities and responsivities depend on the local continuum flux (i.e., radius) for a fixed continuum luminosity, and thus the optical recombination lines are important to include in quasar tomography for mapping out the physical parameters of BLR \\citep{horne03}.  The recovery of a velocity-delay map for a single emission line, such as H$\\beta$, is a key goal of reverberation mapping and would allowing insight into the geometry and kinematics of the BLR.  The simultaneous recovery of velocity-delay maps for several emission lines could set much stronger constraints on, and perhaps break degeneracies between, the physical parameters of the line-emitting gas in the BLR and may usher in yet another new era of understanding for this spatially unresolved region in AGNs."
    },
    "1004/1004.2325_arXiv.txt": {
        "abstract": "We extend existing semi-analytic models of galaxy formation to track atomic and molecular gas in disk galaxies. Simple recipes for processes such as cooling, star formation, supernova feedback, and chemical enrichment of the stars and gas are grafted on to dark matter halo merger trees derived from the Millennium Simulation. Each galactic disk is represented by a series of concentric rings. We assume that surface density profile of infalling gas in a dark matter halo is exponential, with scale radius $r_{\\rm d}$ that is proportional to the virial radius of the halo times its spin parameter $\\lambda$. As the dark matter haloes grow through mergers and accretion, disk galaxies assemble from the inside out. We include two simple prescriptions for molecular gas formation processes in our models: one is based on the analytic calculations by Krumholz, McKee \\& Tumlinson (2008), and the other is a prescription where the $\\rm{H_2}$ fraction is determined by the pressure of the interstellar medium (ISM). Motivated by the observational results of Leroy et al. (2008), we adopt a star formation law in which $\\Sigma_{\\rm{SFR}}\\propto\\Sigma_{\\rm{H_2}}$ in the regime where the molecular gas dominates the total gas surface density, and $\\Sigma_{\\rm{SFR}}\\propto \\Sigma_{\\rm gas}^2$ where atomic hydrogen dominates. We then fit these models to the radial surface density profiles of stars, HI and $\\rm{H_2}$ drawn from recent high resolution surveys of stars and gas in nearby galaxies. We explore how the ratios of atomic gas, molecular gas and stellar mass vary as a function of global galaxy scale parameters, including stellar mass, stellar surface density, and gas surface density. We elucidate how the trends can be understood in terms of three variables that determine the partition of baryons in disks: the mass of the dark matter halo, the spin parameter of the halo, and the amount of gas recently accreted from the external environment. ",
        "introduction": "Before we can reliably compute how galaxies form stars and evolve as a function of cosmic time, we must understand the physical processes that regulate the balance between neutral and molecular gas in their interstellar media. Only if $\\h2$ forms, will gravitationally unstable clouds cool and collapse to high enough densities to trigger star formation in the first galaxies. It is also generally believed that star formation occurs exclusively in molecular clouds in all galaxies at all epochs.  Many galaxy formation models adopt the so-called ``Kennicutt-Schmidt'' law (hereafter K-S law, Schmidt 1959, Kennicutt 1998) to prescribe the rate at which a disk galaxy of given cold gas mass and scale radius will form its stars. This has the form \\begin{equation}\\label{eq:kslaw} \\Sigma_{\\rm{SFR}}\\propto\\sgas^{n} \\end{equation} where $\\Sigma_{\\rm{SFR}}$ represents the star formation rate surface density, $\\sgas$ is the total surface density of the cold gas in the disk, and the exponent $n=1.4$ is often adopted. Some semi-analytic models also account for a critical density below which disks become gravitationally stable and star formation no longer occurs (e.g Kauffmann 1996, De Lucia \\& Blaizot 2007). In this case, \\begin{equation}\\label{eq:kslaw1} \\Sigma_{\\rm{SFR}}\\propto[\\sgas-\\scrit] \\end{equation} where the critical density $\\scrit$ is evaluated using the disk stability criterion given in Toomre (1964). In both Equations (\\ref{eq:kslaw}) and (\\ref{eq:kslaw1}), the star formation rate surface density is proportional to the total surface density of cold gas (i.e. both HI and H$_2$ components) in the galaxy. This prescription was motivated by the analysis of 97 nearby galaxies by Kennicutt (1998), which showed that star formation is more tightly correlated with $\\sgas$ than with $\\Sigma_{\\h2}$. There have been studies in apparent disagreement with these conclusions; for example, Wong \\& Blitz (2002) found that the relation between $\\Sigma_{\\rm{SFR}}$ and $\\Sigma_{\\h2}$ is stronger than that between $\\Sigma_{\\rm{SFR}}$ and $\\Sigma_{\\rm{gas}}$ in galaxies with high molecular gas fractions. In recent years, high quality, spatially-resolved maps of the cold gas have become available for samples of a few dozen nearby galaxies. Examples of such data include HI maps from The HI Nearby Galaxy Survey (THINGS) and CO maps from the Berkeley-Illinois-Maryland Association Survey of Nearby Galaxies (BIMA SONG) and HERA CO-Line Extragalactic Survey (HERACLES). Measurements of the rate at which stars are forming at different radii in the galaxy are provided by Spitzer and GALEX observations. The combination of these different data sets has led to important new constraints on the relationship between star formation and gas in galactic disks. Bigiel et al. (2008) studied 18 disk galaxies and showed that $\\rm{H_2}$ forms stars at a roughly constant efficiency in spirals at radii where it can be detected. Their results suggest a star formation law of the form \\begin{equation}\\label{eq:kslawh2} \\Sigma_{\\rm{SFR}}\\propto\\Sigma_{\\rm{H_2}}^{1.0\\pm0.2} \\end{equation}  Motivated by these findings, galaxy formation modelers are now progressing beyond a simple single-component view of the cold phase of the interstellar medium, and are attempting to model the formation of molecular hydrogen in galaxies. Gnedin, Tassis \\& Kravtsov (2009) included a phenomenological model for $\\h2$ formation in hydrodynamic simulations of disk galaxy formation. Their model includes nonequilibrium formation of $\\h2$ on dust and approximate treatment of both its self-shielding, and shielding by dust from the dissociating UV radiation field. Dutton (2009) and Dutton \\& van den Bosch (2009) utilized the empirically-motivated hypothesis of Blitz and Rosolowsky (2004,2006) that hydrostatic pressure alone determines the ratio of atomic to molecular gas averaged over a particular radius in the disk in their analytic models of disk formation in a $\\Lambda$CDM cosmology. They analyzed the radial distribution of stars and star formation in their disks, but did not focus very much on gas properties in their model. There have also been some attempts to predict the balance between atomic and molecular gas in galaxies at different redshifts by post-processing the publically available outputs of semi-analytic galaxy formation models (e.g. Obreschkow et al. 2009). This work also used the same Blitz and Rosolowsky (2004, 2006) prescription to predict the fraction of molecular gas in disks. However, the Obreschkow et al. approach is not self-consistent, because the simulations have been run assuming a ``standard'' Kennicutt-Schmidt law for star formation and the presence or absence of molecular gas has no influence on the actual evolution of the galaxies in the model.  In this paper, we develop new semi-analytic models that follow gas cooling, supernova feedback, the assembly of galactic disks, the conversion of atomic gas into molecular gas as a function of radius within the disk, and the conversion of the gas into stars. In the 1990's, semi-analytic models of galaxy formation were developed into a useful technique for interpreting observational data on galaxy populations (e.g. Kauffmann, White \\& Guiderdoni 1993; Cole et al. 1994; Somerville \\& Primack 1999). In the first decade of the new Millennium, considerable effort went into grafting these models on to large N-body simulations of the dark matter component of the Universe. These efforts began with relatively low resolution simulations (Kauffmann et al. 1999), but have rapidly progressed to simulations with high enough resolution to follow the detailed assembly histories of millions of galaxies with luminosities well below $L_*$ (Croton et al. 2006; Bower et al. 2006; De Lucia \\& Bliazot 2007; Guo et al. 2010).  Our new models are an extension of the techniques described in Croton et al. (2006) and De Lucia \\& Blaizot (2007) and are implemented using the merger trees from the Millennium Simulation (Springel et al. 2005). We explore two different ``recipes'' for partitioning the cold gas into atomic and molecular form: a) a prescription based on the analytic models of $\\h2$ formation, dissociation and shielding developed by Krumholz, McKee \\& Tumlinson (2009), in which the molecular fraction is a local function of the surface density and the metallicity of the cold gas, b) the same pressure-based formulation explored by Obreschkow et al. (2009).  We first use our models to calculate the HI, $\\h2$, stellar mass and SFR surface density profiles of disk galaxies that form in dark matter haloes with circular velocities $v_{\\rm{cir}} \\sim$ 200 km/s (i.e. galaxies comparable to the Milky Way) and we compare our results to the THINGS/HERACLES observations presented in Bigiel et al. (2008).  We then turn to the issue of the predicted {\\em scaling relations} between atomic gas, molecular gas and stars for an ensemble of disk galaxies forming in dark matter haloes spanning a range of different circular velocities. We currently enjoy a rich and diverse array of scaling laws that describe the stellar components of galaxies. For example, the Tully-Fisher relation and the size-mass relation for local spiral galaxies play a crucial role in constraining current theories of disk galaxy formation. Likewise, the scaling laws of bulge-dominated galaxies (the Fundamental Plane) provide important constraints on how these systems may have assembled through merging. In contrast, few well-established scaling laws exist describing how the cold gas is correlated with other global physical properties of galaxies. Surveys of atomic and molecular gas in well-defined samples of a few hundred to a thousand galaxies are currently underway, and this paper will explore what can be learned about disk galaxy formation from the results.  Our paper is organized as follows. In section 2, we briefly describe the simulation used in our study as well as the semi-analytic model used to track the formation of galaxies in the simulation. In section 3, we describe the new aspects of the models presented in this paper, including our spatially resolved treatment of disk formation in radial bins, the recipes that prescribe how atomic gas is converted into molecular gas, and our new prescriptions for star formation and feedback. In section 4, we compare the radial profiles in our models to observations from the THINGS/HERACLES surveys, and present the global gas properties of the galaxies in our model, such as atomic and molecular gas mass functions. In section 5, we introduce a set of scaling relations for the atomic and molecular gas fractions of galaxies and we clarify which aspects of the input physics are responsible for setting the slope and the scatter of these relations. Finally, in section 6 we summarize our work and discuss our findings. ",
        "conclusions": "  (i) A simple star formation law in which $\\Sigma_{\\rm{SFR}} \\propto \\Sigma_{\\h2}$ leads to gas consumption time-scales in the inner disk that are too short. In this paper, we simply patch over this problem by decreasing the efficiency of supernova feedback in the inner disk.  (ii) The {\\em mean} stellar, HI and $\\h2$ surface density profiles of the disk galaxies in our model are only weakly sensitive to the adopted $\\h2$ fraction prescription. The reason for this is that for typical $L_*$ disk galaxies, the local gas surface density is the main controlling parameter for both recipes. At low gas surface densities, the $\\h2$ fraction depends sensitively on metallicity for the Krumholz et al. prescription, but considerably less sensitively on stellar surface density $\\mu_*$ for the pressure-based prescription. As a result, the correlation between molecular-to-atomic fraction and $\\mu_{\\rm gas}$ for local disk galaxies exhibits more scatter if the Krumholz et al. model is correct.  (iii) Our results indicate that galaxies that have recently accreted a significant amount of gas from the external environment are characterized by higher-than-average {\\em total} cold gas content. If the galaxy has high gas surface density, then this excess gas is an unambiguous signature of a recent accretion event, because the time-scale over which gas is consumed into stars is short in such systems. On the other hand, if the galaxy has low surface density, a higher-than-average total cold gas content could indicate a recent accretion event, but it may also mean that the galaxy has a higher-than-average spin parameter. Higher spin parameters result in disk galaxies with more extended distributions of cold gas, lower-than-average molecular-to-atomic ratios, and low star formation efficiencies. For these ambiguous systems, one must seek additional evidence that the outer disks were assembled {\\em recently}.   Although these conclusions are somewhat open-ended, they do suggest avenues for further research. We believe that a more realistic way forward to solving the gas consumption timescale problem would be to model radial inflow of the gas. Attempts have been made to construct phenomenological models that do include radial mixing of the stars and gas in disks as well as the effect of this mixing on the chemical evolution of the stars formed in the solar neighbourhood (e.g Sch{\\\"o}nrich \\& Binney 2009). Results from hydrodynamical simulations also indicate that the gas tends to flow inwards, while the stars migrate outwards (e.g Ro\\v{s}kar et al. 2008). The main way to distinguish between different scenarios may be the predicted metallicity gradients. We intend to explore these issues in more detail in future work.  In our model results, although the surface density profiles from the two $\\h2$ fraction prescriptions are very similar, the models indicate that one should, in principle, be able to confirm the metallicity-dependence of the molecular gas fraction predicted by the Krumholz prescription, if one measures the average gas-phase metallicities of nearby disk galaxies using emission lines. Alternately, one can break the degeneracy by observing systems where the metallicity is low but the pressure is high (Fumagalli, Krumholz \\& Hunt 2010).  Another interesting issue is whether a galaxy's location in the gas scaling relation diagrams can serve as a diagnostic as to whether it has accreted gas from the external environment. Although the theory of gas accretion in galaxies has received considerable attention of late (e.g. Kere{\\v s} et al. 2005; Dekel \\& Birnboim 2006; Dekel et al. 2009), there is little {\\em direct} observational evidence that this occurs in practice. This is true both for galaxies in the local Universe and at high redshifts, where gas accretion rates are expected to be much higher. Although average gas accretion rates are expected to be low at the present day, precise quantification of the expected scaling relations for equilibrium disk galaxies may allow us to identify a subset of systems which deviate significantly from the mean in terms of their gas content. Following the conclusion (iii), one may try to gain a better understanding of the observationally detectable signatures of a recent gas accretion episode. Possible ways forward would be to look for signatures of recent accretion in the observed age gradients of the stars or in the metallicity gradients of the gas in the disk. One could also look for accretion signatures in the kinematics of the stars and the gas in the outer disks. Alternatively, one could search for evidence of complex structure (e.g. tidal streams or shells) in the stellar haloes of gas-rich galaxies (Cooper at al 2010). We intend to explore these possibilities in more detail in future work.  Ongoing and future surveys, such as the Galex Arecibo Sloan Survey (GASS) (Catinella et al. 2010) and the COLD GASS survey carried out at the IRAM 30m telescope (Saintonge et al. in preparation) will enable us to quantify the scaling relations discussed in this paper in considerable detail. These surveys will provide interesting targets for follow-up programs, which may help us understand that extent to which galaxies still accrete gas at the present day. In the next few years, it will become possible to observe gas in galaxies at higher redshifts using facilities such as ALMA and Square Kilometer Array pathfinder experiments such as ASKAP or MEERKAT. We are certain that our simplified treatment of disk formation in concentric rings that undergo no radial mixing will not be a good way to describe the assembly of the clumpy, highly turbulent disks that are now known to exist at $z \\sim 2$ (e.g. Genzel et al. 2008). Nevertheless, we believe that our models may still be useful in elucidating the gaseous and chemical evolution of disks over a somewhat smaller range in lookback time."
    },
    "1004/1004.5219_arXiv.txt": {
        "abstract": "{ eROSITA (extended ROentgen Survey with an Imaging Telescope Array) is the core instrument on the Russian Spektrum-Roentgen-Gamma (SRG) mission which is scheduled for launch in late 2012. eROSITA is fully approved and funded by the German Space Agency DLR and the Max-Planck-Society.  The design driving science is the detection of 50 - 100 thousands Clusters of Galaxies up to redshift $z\\sim$1.3 in order to study the large scale structure in the Universe and test cosmological models, especially  Dark Energy. This will be accomplished by an all-sky survey lasting for four years plus a phase of pointed observations. eROSITA consists of seven Wolter-I telescope modules, each equipped with 54 Wolter-I shells having an outer diameter of 360 mm. This would provide and effective area at 1.5 keV of $\\sim$ 1500 cm$^{2}$ and an on axis PSF HEW of 15$^{\\prime\\prime}$ which would provide an effective angular resolution of 25$^{\\prime\\prime}$-30$^{\\prime\\prime}$. In the focus of each mirror module, a fast frame-store pn-CCD will provide a field of view of 1$^{\\circ}$ in diameter for an active FOV of $\\sim$0.83 deg$^{2}$. At the time of writing the instrument development is currently in phase C/D. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.1921_arXiv.txt": {
        "abstract": "One of the sources of gravitational waves for the proposed space-based gravitational wave detector, the Laser Interferometer Space Antenna (LISA), are the inspirals of compact objects into supermassive black holes in the centres of galaxies --- extreme-mass-ratio inspirals (EMRIs). Using LISA observations, we will be able to measure the parameters of each EMRI system detected to very high precision. However, the statistics of the set of EMRI events observed by LISA will be more important in constraining astrophysical models than extremely precise measurements for individual systems. The black holes to which LISA is most sensitive are in a mass range that is difficult to probe using other techniques, so LISA provides an almost unique window onto these objects. In this paper we explore, using Bayesian techniques, the constraints that LISA EMRI observations can place on the mass function of black holes at low redshift. We describe a general framework for approaching inference of this type --- using multiple observations in combination to constrain a parameterised source population. Assuming that the scaling of EMRI rate with black hole mass is known and taking a black hole distribution given by a simple power law, ${\\rm d}n/{\\rm d}\\ln M = A_0 (M/M_*)^{\\alpha_0}$, we find that LISA could measure the parameters to a precision of $\\Delta(\\ln A_0)\\sim 0.08$, and $\\Delta(\\alpha_0)\\sim 0.03$ for a reference model that predicts $\\sim1000$ events. Even with as few as $10$ events, LISA should constrain the slope to a precision $\\sim0.3$, which is the current level of observational uncertainty in the low-mass slope of the black hole mass function. We also consider a model in which $A_0$ and $\\alpha_0$ evolve with redshift, but find that EMRI observations alone do not have much power to probe such an evolution. ",
        "introduction": "The proposed Laser Interferometer Space Antenna (LISA)~\\cite{lisaCQG} will be sensitive to gravitational waves from systems containing supermassive black holes in the $10^4M_{\\odot}$--$10^7M_{\\odot}$ range. This mass range is very hard to probe electromagnetically, and only five such systems have been robustly identified from dynamical measurements \\citep{gultekin2009}, including the black hole in the centre of our own galaxy. Placing constraints on the mass function of low-mass black holes has, however, key astrophysical implications.  It has been suggested \\citep{volonteri2008} that the shape of the mass function of massive black holes at the low-mass end is a key diagnostic to derive constraints on the mechanism that formed black hole seeds. Seed black holes are predicted to form in the mass range $\\sim 100-10^5 M_\\odot$, depending on the specific physical process involved \\citep{volonteri2008}.  As black holes grow from low-mass seeds, it is natural to expect that at least some black holes have not grown efficiently and still trace the initial conditions.  Clearly, black holes at the high-mass end of the mass function have increased their mass by accretion, or they have experienced mergers and dynamical interactions. Any dependence of the black hole mass on the initial seed mass is erased. However, the distribution of low-mass black holes still retains some ``memory'' of the original seed mass distribution. The expectation is that ungrown seeds produce a peak in the mass function that corresponds to the peak of the seed mass function.  The higher the efficiency of seed formation, the more pronounced is the peak.  Additionally, the mass function of low-mass black holes can provide insights on the co-evolution of black holes and their hosts. Observations show that the masses of black holes correlate with the mass, luminosity and the stellar velocity dispersion of the host \\citep[and references therein]{gultekin2009}.  These correlations imply that black holes evolve along with their hosts throughout cosmic time. Lauer et al.~\\cite{Lauer2007} suggest that at least some of these correlations break down at the largest galaxy and black hole masses \\citep[but see][]{Bernardi2007}. One unanswered question is whether this symbiosis extends down to the lowest galaxy and black hole masses \\citep{Greene2008}, due to changes in the accretion properties \\citep{Mathur2005}, dynamical effects \\citep{Volonteri2007}, or cosmic bias \\cite{Volonteri2009}.  Since current measurements of black hole masses extend barely down to $M_{\\rm bh}\\sim 10^6 M_\\odot$, these features cannot be observationally tested with present data. LISA observations will significantly improve our understanding of the astrophysics of black holes in this mass range. LISA will detect mergers between supermassive black holes with these masses out to very high redshift, which will probe the early assembly of these systems and their host galaxies. LISA will also detect gravitational waves generated when compact objects (white dwarfs, neutron stars or black holes) are captured by and inspiral into supermassive black holes in the centres of galaxies~\\cite{astrogr}, which are referred to as extreme-mass-ratio inspirals (EMRIs). The EMRI events will mostly be at low redshift, $z \\lesssim 1$, and can therefore be used to probe the quiescent population of $\\sim10^4M_{\\odot}$--$10^7M_{\\odot}$ black holes that remain today.  In this paper, we focus on this second type of source and explore what constraints LISA might be able to place on the low redshift population of black holes in this mass range.  A typical EMRI event will have frequency $\\sim 1$mHz and will be observed by LISA for $\\sim 1$ year, and so we will detect $\\sim 10^5$ cycles of the waveform. This allows LISA to make very precise measurements of the parameters of the host system~\\cite{AK,HG09}, and it is hoped that EMRI observations can be used to carry out high precision tests of the spacetime geometry of the central object~\\citep[and references therein]{astrogr}. For a typical EMRI event with signal-to-noise ratio (SNR) of $\\sim 30$, we would expect to recover the redshifted mass of the central black hole to a precision $\\Delta\\ln((1+z) M) \\sim 10^{-4}$ and the distance to the source to a precision of $\\sim 3\\%$. The spin of the central black hole, the redshifted mass of the inspiralling object and the orbital parameters (initial radius and eccentricity) should also be measured to very high accuracy, $\\Delta\\ln X\\sim10^{-4}$--$10^{-3}$. While such precise measurements for individual systems are interesting and very important if the data is to be used for high precision tests of relativity, astrophysically it is the statistics of the set of EMRI events that LISA observes that will be of greatest value in constraining models. It is this application of LISA EMRI observations that we focus on in the current paper.  The distribution of events that LISA detects will depend on three factors --- the number density of possible source systems; the rate at which EMRIs occur in systems with particular parameters; and the sensitivity, in terms of distance reach, of the LISA detector to particular types of EMRI event. The last effect can be estimated theoretically in advance of the LISA mission, so we concentrate on what LISA can tell us about the first two effects. A particular model for the black hole distribution and the rate of inspirals per black hole does not precisely predict the number of events that LISA will observe, since inspirals start stochastically in any given galaxy. However, a particular model does predict the rate at which observable EMRIs of particular type will occur, and hence the probability distribution of the observed events. The LISA observation is a sample from this distribution, and we wish to infer the parameters of the underlying model from this sample. This can be done with a simple application of Bayes Theorem. Bayesian techniques have been employed widely in the context of gravitational wave data analysis~\\cite[see, for instance,][]{cornishMCMC,christensenMCMC}, but this has been primarily to make statements about the parameters of individual sources. We can also use Bayesian methods to make statements about the underlying population from which the sources we detect are drawn and it is that which we will do here. Using LISA supermassive black hole merger events to constrain astrophysical models was considered in~\\cite{plowmanSMBH}. The focus of that work was to derive an ``error kernel'' which would map the intrinsic distribution of source parameters onto the observed distribution of source parameters, and for model selection they used a variant of the Kolmogorov-Smirnov test. Our work differs from that approach not only in the fact that we consider EMRIs, but in the use of a Bayesian framework for the analysis and a parameterised model for the underlying distribution that we wish to constrain.  LISA will be able to measure the product of the EMRI rate per black hole with the number density of black holes, but it is not clear if it will be able to decouple the two effects. The scaling of the EMRI rate per black hole with the black hole parameters can, in principle at least, be constrained in advance through numerical simulations. One such analysis was carried out in~\\cite{hopman09}. Various assumptions were made in that analysis which may not be valid for black holes in the LISA mass range, and we will discuss these issues further in Section~\\ref{sec:probdist}. However, for the purpose of this work we will assume that the black hole rate scaling is known, and is given by the results in~\\cite{hopman09}. This allows us to interpret our results in terms of the underlying distribution of black hole masses. An alternative interpretation that does not rely on this assumption is that we are constraining the convolution of the black hole number density with the EMRI rate per black hole.  The parameters of individual sources will not be measured perfectly by LISA, due to noise in the detector and confusion from other sources in the data stream. It is possible to include parameter errors consistently when making statements about the population, and we will describe how this is done in practice in Section~\\ref{parerr}. However, this requires having properly sampled posterior distributions for all of the sources in the data set. An alternative approach is to bin the sources according to their maximum likelihood parameters. Although sources may end up in the wrong bins due to the parameter errors, the majority of sources will be correctly classified. We use the binning approach in this paper, since it allows us to assess LISA's ability to constrain the black hole population without requiring the computationally expensive simulation of LISA noise and recovery of posteriors for hundreds of EMRI sources. We will demonstrate that our conclusions are not affected by the inclusion of reasonable parameter estimation errors, which indicates that our results are an accurate reflection of what will be achievable in practice.  In this paper we take a single power law model for the black hole mass function and consider both a redshift-independent case of the form ${\\rm d}n/{\\rm d}\\ln M = A_0 (M/M_*)^{\\alpha_0}$, and a redshift-dependent case of the form ${\\rm d}n/{\\rm d}\\ln M = A_0 (1+z)^{A_1} (M/M_*)^{\\alpha_0-\\alpha_1 z}$. In the former case, we find that LISA will be able to measure the parameters to precisions $\\Delta(\\ln A_0) \\sim 0.08$ and $\\Delta(\\alpha_0) \\sim 0.03$, and in the latter case to precisions $\\Delta(\\ln A_0) \\sim 0.2$, $\\Delta(\\alpha_0) \\sim 0.06$, $\\Delta(A_1) \\sim 0.7$ and $\\Delta(\\alpha_1) \\sim 0.2$. These precisions scale with the number of observed events like $N_{\\rm obs}^{-1/2}$, and have been normalised to a reference case that predicts $\\sim1000$ events. We find that changing our assumptions about the performance of the LISA mission affects these conclusions only through the change in the number of events predicted. We also find that the precisions are somewhat improved if the black holes in the EMRI systems tend to have large spins.  The paper is organised as follows. In Section~\\ref{theory} we describe the theoretical framework that we employ in this analysis. This includes a discussion of Bayes Theorem, a description of the probability distribution for EMRI events that LISA will detect, the proper treatment of parameter measurement errors and a summary of existing constraints on the shape of the black hole mass function. In Section~\\ref{sec:res}, we present our results for both the redshift-independent and redshift-dependent models. This section includes a discussion of the effect of measurement errors and variation in the bin size used in the analysis. We also demonstrate how the results change as we vary the true black hole population in the Universe and as we change our assumptions about the performance of the LISA detector and the spin of the black holes. Finally, in Section~\\ref{sec:discuss} we summarise our results and discuss directions for future research. ",
        "conclusions": "\\label{sec:discuss} We have explored the ability of the proposed space-based gravitational wave detector, LISA, to probe the properties of black holes through observations of extreme-mass-ratio inspiral events. We have presented a general framework for addressing such questions, and considered a particular special case in which we imagine that the events observed by LISA are divided into bins in mass and redshift. Assuming that the only model uncertainty is in the unknown number density of black holes in the LISA range, $10^4 M_{\\odot} < M < 10^7 M_{\\odot}$, and taking a simple, non-evolving, power-law model for the black hole mass function, ${\\rm d}n/{\\rm d}\\ln M = A_0 (M/M_*)^{\\alpha_0}$, we conclude that LISA will be able to measure the amplitude of the mass function to a precision $\\Delta(\\ln A_0) \\sim 0.08$ and the slope to a precision $\\Delta(\\alpha_0) \\sim 0.03$. These precisions scale with the number of observed events like $N_{\\rm obs}^{-1/2}$ and have been normalised to a reference model with $N_{\\rm obs} \\approx1000$. The present uncertainty in the slope of the mass function in the LISA range is of the order of $\\pm 0.3$, so LISA will beat this with as few as $10$ detected events.  If we allow for the mass function to evolve with redshift, using the ansatz ${\\rm d}n/{\\rm d}\\ln M = A_0 (1+z)^{A_1} (M/M_*)^{\\alpha_0-\\alpha_1 z}$, but assume the Universe has a non-evolving mass function with $A_1=\\alpha_1=0$, we find LISA will be able to measure the parameters of the distribution to a precision $\\Delta(\\ln A_0) \\sim 0.2$, $\\Delta(\\alpha_0) \\sim 0.06$, $\\Delta(A_1) \\sim 0.7$ and $\\Delta(\\alpha_1) \\sim 0.2$. These again show a scaling like $N_{\\rm obs}^{-1/2}$ and are normalised to a black hole distribution that predicts $1000$ events. The errors in the evolution parameters, $A_1$ and $\\alpha_1$, are sufficiently large that we must conclude that EMRI observations alone will be unable to detect evolution in the black hole mass function.  The work presented here has made use of various simplifications, which we now discuss.  \\subsection{EMRI parameter space} We have considered measurements of the central black hole mass and the source redshift only, and have taken all EMRIs to be circular, equatorial inspirals into black holes with the same spin. We also used signal-to-noise ratios that were averaged over sky position and orientation when assessing the detectability of different sources. This allowed us to use the observable lifetime functions given in~\\cite{gairEMRIastro}. It will be important to extend this work to generic EMRIs, which will require the extension of the observable lifetime calculation to generic orbits. In general, the addition of extra parameters to a model tends to decrease the precision to which the model parameters can be measured, as we saw when we allowed the parameters of the mass function to vary with redshift. However, in this case, we will not only be adding additional parameters, but additional measurements, since LISA EMRI observations will also measure the black hole spin, orbital eccentricity etc. to high precision~\\cite{AK}. We would therefore not expect our conclusions about the precision to which we can measure the black hole mass function to change significantly. We have already seen that we can measure the slope of the mass function to higher precision if black holes tend to be more rapidly spinning, as the presence of spin increases the range in mass and redshift within which EMRIs can be detected. Eccentricity may have a similar effect, as the additional waveform harmonics that arise due to eccentricity will enter the LISA frequency band earlier, thus tending to enhance the signal-to-noise ratio. These effects must be quantified in the future.  The most important consequence of including new parameters will be the ability to ask additional astrophysics questions. EMRIs will have considerable power to probe the spin distribution of black holes in the appropriate mass range at low redshift. Measurements of the masses of the inspiraling objects and the orbital eccentricities will provide constraints on the processes occurring in stellar clusters in the centres of galaxies, such as mass segregation and the efficiency of the various channels that lead to EMRIs~\\cite{astrogr}. It is important to understand what LISA will be able to tell us about these various processes and how correlations could affect our ability to address these questions independently.  \\subsection{Data analysis technique} There are various assumptions in the analysis technique described here that could be relaxed. We have ignored parameter errors, other than to verify they did not significantly affect our results. This was possible because we were using binning of the observations before analysing the data. We have described how errors can be included properly and how the data can be analysed in a continuum limit, but this requires having properly sampled posterior pdfs for the EMRI sources. While we do not anticipate that our results will change significantly under such an analysis, it might be possible to explore this using pdfs for EMRI sources constructed in the context of the Mock LISA Data Challenges~\\cite{mldc4}. Additionally, the construction of the observable lifetime used here assumes that the LISA observation is $100\\%$ complete for sources with SNR $>$ 30, and $0\\%$ complete for sources with SNR$<30$. The exact LISA completeness function will depend on the algorithms employed to carry out data analysis, which are rather uncertain at present~\\cite{BGPemri,Cemri}. Our present model is a reasonable approximation if an SNR cut is used for source selection for the follow-up. It is unlikely that LISA data analysis preparation will reach a point where we will have a better model for the completeness anytime soon, but it will be straightforward to recompute the effective observable lifetime once one is available. Finally, our analysis has used the same model, based on a Poisson probability distribution, to generate the data sets we have searched and for the posterior construction.This is likely to be a good model for the EMRIs occurring in a given galaxy, but there are various complicating factors, since it takes some time for galactic centres to become relaxed after galaxy mergers. Although the galaxies in which LISA can detect EMRIs are at low redshift and therefore are unlikely to have undergone recent mergers, the Poisson model could be checked by using numerical simulations, of the type described in~\\cite{hopman09}. We are also assuming that all systems of a given type have the same EMRI rate. While there will certainly be an average rate for systems of particular type, there may be rate ``noise'' which we have ignored here. This could be included in the generation of the sets of events to search, but we would require a reasonable model for the amount of noise to include.  \\subsection{EMRI rate uncertainties} One significant uncertainty in these results is in the correct interpretation of what we are measuring. As discussed earlier, the rate of EMRIs of a particular type depends on both the number density of black holes, and on the intrinsic rate of inspirals per black hole. We assumed that the second of these was known, using results in~\\cite{hopman09}, and therefore that the EMRI observations were telling us about the number density of black holes. In practice, we will be measuring the convolution of these two effects to the precisions quoted here. As discussed in Section~\\ref{sec:probdist}, observations of low-mass galaxies and further theoretical work should better inform our understanding of the EMRI rate per galaxy and the level of our uncertainty in it before LISA flies. However, there will be some residual uncertainty in the model assumptions, and there may be correlations between the rate and the black hole mass, or even other parameters such as spin, which might not have been included in the models. LISA is unlikely to be able to tease apart these two effects directly, which must be borne in mind when the results are interpreted. The use of other observations, either of gravitational waves or in electromagnetic wavebands, might help with the final interpretation.  \\subsection{Future applications} The work described in this paper has illustrated the potential power of LISA observations for studying the astrophysics of black holes in the range $10^4M_{\\odot} < M < 10^7 M_{\\odot}$. We have focussed on EMRI observations as an illustration, but the same approach can be used for the interpretation of LISA observations of mergers between supermassive black holes (SMBHs). SMBH mergers and EMRIs probe different subsets of the same population of black holes --- SMBH events will probe the mergers between black holes up to the highest redshift and earliest cosmic times, while EMRIs will probe the whole population of black holes (not necessarily active or merging) at $z<1$ that are the end products of such SMBH mergers. Models that predict the rates of LISA SMBH mergers will therefore also make predictions for the number density of these low redshift black holes that play host to EMRIs. Thus, it should be possible to derive more powerful constraints on the models from combined observations. In particular, it might be possible to use SMBH mergers at lower redshift in conjunction with EMRI observations to place constraints on the evolution of the black hole mass function with redshift, which we have seen is not possible using EMRI observations alone. This should be explored in more detail in the future."
    },
    "1004/1004.3579_arXiv.txt": {
        "abstract": "Convective driving, the mechanism originally proposed by \\citet{Brickhill91a,Brickhill83} for pulsating white dwarf stars, has gained general acceptance as the generic linear instability mechanism in DAV and DBV white dwarfs.  This physical mechanism naturally leads to a nonlinear formulation, reproducing the observed light curves of many pulsating white dwarfs.  This numerical model can also provide information on the average depth of a star's convection zone and the inclination angle of its pulsation axis. In this paper, we give two sets of results of nonlinear light curve fits to data on the DBV GD~358. Our first fit is based on data gathered in 2006 by the Whole Earth Telescope (WET); this data set was multiperiodic, containing at least 12 individual modes. Our second fit utilizes data obtained in 1996, when GD~358 underwent a dramatic change in excited frequencies accompanied by a rapid increase in fractional amplitude; during this event it was essentially monoperiodic. We argue that GD~358's convection zone was much thinner in 1996 than in 2006, and we interpret this as a result of a short-lived increase in its surface temperature.  In addition, we find strong evidence of oblique pulsation using two sets of evenly split triplets in the 2006 data. This marks the first time that oblique pulsation has been identified in a variable white dwarf star. ",
        "introduction": "White dwarf stars offer several advantages for astrophysical study. First, they are the evolutionary endpoint of about 97\\% of all stars and are therefore representative of a large fraction of the stellar population. Second, the source of their pressure support is electron degeneracy \\citep{Chandrasekhar39} so their bulk mechanical structure is well understood. Third, nuclear reactions, if any, contribute a negligible amount to their energy, so their evolution is dominated by simple cooling \\citep{Mestel52}. Finally, they are observed to pulsate in specific temperature ranges. The pulsators are believed to be typical in every other way, so what we learn \\emph{asteroseismologically} about them should apply to all white dwarf stars \\citep[for recent reviews, see][]{Winget08,Fontaine08}.  In addition to learning about the stars themselves, the relative simplicity of white dwarfs makes them ideal laboratories for testing and constraining poorly-understood physical processes. One such process, convection, is an important energy transfer process in most stars, yet it remains one of the largest sources of uncertainty in stellar modeling. For instance, main-sequence stars at least 20\\% more massive than the Sun have convective cores, and the amount of convective overshoot and mixing is the primary factor that determines their main sequence lifetimes \\citep[see, e.g.,][]{DiMauro03}.  In addition, red giants and AGB stars have large convective envelopes, and the details of convection play a role in the evolution of their surface abundances and in their overall evolution \\citep{Bertelli09}.  We have developed a method which uses the pulsations of white dwarf stars to measure fundamental parameters of their convection zones. The physical idea is that the pulsations cause local surface temperature (``\\teff'') variations that lead to local variations in the depth of the convection zone. As the convection zone waxes and wanes it both absorbs and releases energy, modulating the local energy flux \\citep{Brickhill91a,Goldreich99a}. Due to the extreme temperature sensitivity of convection, finite amplitude pulsations can lead to highly nonlinear light curves \\citep{Brickhill92a,Wu01,Ising01}.  In \\citet{Montgomery05a} we showed how a simple numerical model could be used to obtain not only good light curve fits but also information on the average depth of a star's convection zone and the inclination angle of its pulsation axis. ",
        "conclusions": "In this paper we have extended our nonlinear light curve fitting technique to the multiperiodic pulsator GD~358. Our fit to the 2006 WET data provides a good match to the light curves and we find that the thermal response time of its convection zone is $\\tau_0 = 572.9 \\pm 6.1$~s. This is considerably larger than that of the star PG1351+489, for which $\\tau_0 \\sim 100$~s \\citep{Montgomery05a}.  This difference in $\\tau_0$ is consistent with the effective temperatures of these stars: the pure He solution for PG1351+489 yields a \\teff\\ which is $\\sim 2,000$~K hotter than that for GD~358.  We also obtained a fit to the light curve of GD~358 during the \\emph{sforzando} event in 1996. These fits showed that GD~358 had $\\tau_0 \\sim 42 \\pm 2$~s, a value much less than that determined from the 2006 data. This suggests that its effective temperature was approximately 2,000~K hotter in 1996 than in 2006, and this is consistent with the estimate of \\citet{Weidner03} that the light curve shape suggests $\\teff \\sim 27,000$~K. Independent evidence of GD~358's brightness relative to comparison stars is also consistent with such a temperature increase at the time of the \\emph{sforzando} \\citep{Provencal09}. The physical origin of this temperature increase will be the subject of future work.  As expected, these data indicate an increase in the depth/mass of the convection zone with decreasing \\teff. A similar trend is given by ML2/$\\alpha = 1.1$ convection \\citep{Bohm71}, although the slope of the theoretical relation appears less steep than that of the data. In addition, lower values of $\\alpha \\la 0.6$ are excluded.  In general, $\\tau_0$ is also a function of $\\log g$, albeit a somewhat weaker one.  Our ultimate goal is to map $\\tau_0$ as a function of both \\teff\\ and $\\log g$ for both the DBV and DAV instability strips. These data will provide insight into the physics of convection, still one of the largest uncertainties in stellar modeling. They will also serve as important constraints for new hydrodynamic simulations of convection which are starting to come online.  Multiple lines of evidence point to some of GD~358's modes undergoing oblique pulsation, in particular, the peaks in the $k=12$ region. First, these peaks can be fit with two sets of exactly evenly spaced triplets. Second, the relative phases of each of the components within the triplet indicate that each originates as a single $m=-1$ or +1 mode aligned with the magnetic axis; as the star rotates, the magnetic axis precesses around the rotation axis, generating a triplet for each intrinsic mode. Finally, the oblique pulsator model qualitatively and quantitatively fits the amplitudes of the peaks seen in the Fourier transform. Taken together, this marks the first time that oblique pulsation has been seen in a white dwarf variable.  Having now identified the characteristics of oblique pulsation in GD~358 we now know what to look for in other white dwarf variables; we have found preliminary indications of it in other stars and in other data sets of GD~358. As discussed in the previous sections, oblique pulsation may prove to be a diagnostic of both the magnetic field and its changes as well as a diagnostic of differential rotation. This opens an exciting chapter in the seismology of these objects."
    },
    "1004/1004.2443_arXiv.txt": {
        "abstract": "Intermediate between the prestellar and Class\\,0 protostellar phases, the first core is a quasi-equilibrium hydrostatic object with a short lifetime and an extremely low luminosity. Recent MHD simulations suggest that the first core can even drive a molecular outflow before the formation of the second core (i.e., protostar). Using the Submillimeter Array and the {\\it Spitzer Space Telescope}, we present high angular resolution observations towards the embedded dense core IRS2E in L1448. We find that source L1448 IRS2E is not visible in the sensitive $Spitzer$ infrared images (at wavelengths from 3.6 to 70\\,$\\mu$m), and has weak (sub-)\\,millimeter dust continuum emission. Consequently, this source has an extremely low bolometric luminosity ($<$\\,0.1\\,$L_\\odot$). Infrared and (sub-)\\,millimeter observations clearly show an outflow emanating from this source; L1448 IRS2E represents thus far the lowest luminosity source known to be driving a molecular outflow. Comparisons with prestellar cores and Class\\,0 protostars suggest that L1448 IRS2E is more evolved than prestellar cores but less evolved than Class\\,0 protostars, i.e., at a stage intermediate between prestellar cores and Class\\,0 protostars. All these results are consistent with the theoretical predictions of the radiative/magneto hydrodynamical simulations, making L1448 IRS2E the most promising candidate of the first hydrostatic core revealed so far. ",
        "introduction": "Stars form by the gravitational collapse of dense cores in molecular clouds. A comprehensive understanding of the formation and evolution of dense cores is thus a necessary prerequisite to the understanding of the origin of stellar masses, multiple systems, and outflows. Over the past decade, observational studies of (low-mass) dense cores have made significant progress (see, e.g., Reipurth et al. 2007 for recent reviews). Representing the earliest phase of star formation, both prestellar and protostellar cores have been extensively observed and studied using large (sub)\\,millimeter telescopes (e.g., JCMT and IRAM-30m) and infrared telescopes (e.g., {\\it Spitzer Space Telescope}). In practice, however, it is still difficult to distinguish the two types of cores because of the lack of readily observable differences between them. This is illustrated by the fact that several ``prestellar\" cores, like L1014, were found to harbor very low-luminosity protostars in sensitive $Spitzer$ observations (see Young et al. 2004). Consequently, despite all of the observational advances in the past decade, we still do not have a good understanding of the evolutionary process that turns a prestellar core into a protostar.  On the theoretical side, the collapse and evolution from prestellar cores to Class\\,0 protostars have been long studied since the pioneering work of Larson (1969). Theoretical calculations and simulations in fact predict two successive collapse phases, before and after the dissociation of molecular hydrogen, resulting in two different hydrostatic objects (see, e.g., Larson 1969; Masunaga et al. 1998, 2000; Andr\\'{e} et al. 2008). The collapsing prestellar core is initially optically thin to the thermal emission from dust grains, and the compressional heating rate by the collapse is much smaller than the cooling rate by the thermal radiation. The collapse is therefore isothermal at the very beginning. This condition is broken when the compressional heating rate surpasses the radiative cooling rate, and the central temperature increases gradually above 10\\,K. The collapse is then decelerated and forms a shock at the surface of a quasi-adiabatic hydrostatic object, the so-called ``first hydrostatic core\" or ``first core\", which consists mainly of hydrogen molecules. The inward motion at this phase is called the ``first collapse\". When the central temperature reaches about 2000\\,K, hydrogen molecules begin to dissociate into atoms, which acts as an efficient coolant of the gas. When released gravitational energy is consumed by the dissociation, the gas pressure cannot increase rapidly enough to support the first core against its self-gravity, the ``second collapse\" begins. After the dissociation is completed, the ``second core\", a truly hydrostatic protostellar object, forms in the center. Most, if not all, Class\\,0 protostars observed so far actually belong to the population of the ``second cores\" (Ph. Andr\\'{e}, private communication).  Intermediate between the prestellar and Class\\,0 protostellar phases, the first core is a transient object accreting from the surrounding dense envelope; the lifetime of the first core is calculated to be only 10$^3$ to 10$^4$ years (Boss \\& Yorke 1995; Masunaga et al. 1998; Machida et al. 2008). Based on radiative hydrodynamical (RHD) simulations, Boss \\& Yorke (1995) and Masunaga et al. (1998) modelled the spectral energy distribution (SED) of the first core, and found that it should have an extremely low bolometric luminosity ($<$\\,0.1\\,$L_\\odot$), and have no detectable infrared emission at wavelengths shorter than $\\sim$\\,30\\,$\\mu$m with current telescopes. Furthermore, recent magneto-hydrodynamical (MHD) simulations have found that the first core can even drive a molecular outflow before the formation of the second core (i.e., protostar) (Tomisaka 2002; Banerjee \\& Pudritz 2006; Machida et al. 2008). Therefore, the observational detection of the first core would not only confirm the predictions of RHD models but also set strong constraints on MHD models of protostellar outflows. Unfortunately, due to its short lifetime and extremely low luminosity, no first core has been observationally found as yet.  In this paper, we present Submillimeter Array\\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.} (SMA; Ho et al. 2004) and {\\it Spitzer Space Telescope} ($Spitzer$) observations towards an embedded dense core in the L1448 region ($d$\\,=\\,240\\,$\\pm$\\,20\\,pc; Hirota et al. 2008). As a bridge between the isolated star-forming cores and the large-scale clusters, L1448 is an excellent region for studying star formation on the intermediate scale and has been observed extensively in the past two decades (see, e.g., Bally et al. 2008 and reference therein). L1448~IRS2, in the western part of the L1448 filament, was classified as a Class\\,0 protostar by O'Linger et al. (1999). Located $\\sim$\\,50$''$ to the east of IRS2, another dense core was revealed in the SCUBA submm images in O'Linger et al. (1999), and was formally cataloged as SCUBA core No.\\,31 in Hatchell et al. (2005) and SMM\\,J032543+30450 in Kirk et al. (2006). This core was found to have a mean kinetic gas temperature of $T_{\\rm kin}\\approx 11$\\,K, and the observed width of NH$_{3}$\\,(1,\\,1) is $\\sim$\\,0.16\\,km\\,s$^{-1}$ (Rosolowsky et al. 2008). We refer to this dense core as L1448 IRS2E in this work. ",
        "conclusions": "\\subsection{Spectral Energy Distribution}  Table~1 lists the (sub-)\\,mm fluxes of L1448 IRS2E, estimated from the SCUBA, Bolocam, and SMA images. Since there is no local emission peak at the position of the SMA compact source in the SCUBA/Bolocam images (see Fig.\\,2), the estimated fluxes per beam around IRS2E in these images represent conservative upper limits to the fluxes from the embedded source. The 3\\,$\\sigma$ upper limits in the $Spitzer$ images are also listed in Table~1. Based on these data points, we constructed the spectral energy distribution (SED) of IRS2E (plot not shown here). To estimate the luminosity of IRS2E, we first interpolated and then integrated the SED (all the upper limits were used), always assuming spherical symmetry. Interpolation between the flux densities was done by a $\\chi$$^2$ single-temperature grey-body fit to all points at $\\lambda$\\,$\\geq$\\,70\\,$\\mu$m, using the same method as described in Chen et al. (2008). A simple logarithmic interpolation was performed between all points at $\\lambda$\\,$\\leq$\\,70\\,$\\mu$m. The estimated bolometric luminosity of L1448 IRS2E is less than 0.1\\,$L_\\odot$.  Although only an upper limit to the bolometric luminosity could be derived, we can still use it to further constrain the evolutionary stage of L1448 IRS2E. If we assume a steady mass-accretion rate given by $\\dot{M}$\\,=\\,0.975$c_{\\rm s}^3$/$G$ (Shu 1977), where $c_{\\rm s}$ is the effective sound speed, for a gas temperature of 10\\,K the accretion rate is $\\sim$\\,2\\,$\\times$\\,10$^{-6}$\\,$M_\\odot$\\,yr$^{-1}$. The accretion luminosity is calculated as $L_{\\rm acc}$\\,=\\,$GM_*\\dot{M}$/$R_*$, where $M_*$ is the stellar mass and $R_*$ is the stellar radius. The bolometric luminosity being $<$\\,0.1\\,$L_\\odot$ implies a protostellar mass of $<$\\,0.01\\,$M_\\odot$, assuming a radius of 2\\,$R_\\odot$. The age of a $<$\\,0.01\\,$M_\\odot$ `protostar' under the assumption of a constant mass-accretion rate of 2\\,$\\times$\\,10$^{-6}$\\,$M_\\odot$\\,yr$^{-1}$ is then calculated to be $<$\\,5000\\,yr, which is consistent with the outflow age estimated above ($\\geq$\\,1800\\,yr).  The estimated low luminosity and age suggest that L1448 IRS2E is a very young object, in which star formation has just started. Nevertheless, it must be noted that uncertainties remain in our estimates due to the limited observations available. More information, such as {\\it Herschel Space Observatory} imaging at 75$-$300\\,$\\mu$m, is needed to constrain the SED of L1448 IRS2E in order to address more precisely its evolutionary status.  \\subsection{Comparisons to Prestellar, Class\\,0, and VeLLO Objects}  \\noindent{\\bf Comparison to Prestellar Cores:} Prestellar cores are dense ($n$$_{\\rm H}$\\,$\\sim$\\,10$^{4}$--10$^{6}$\\,cm$^{-3}$) cores which are self-gravitating and evolve toward higher degrees of central condensation, but no central hydrostatic protostellar object exists yet within the core (Andr\\'{e} et al. 2000; 2008). Although the properties of L1448 IRS2E are still poorly known, its observed narrow width of the NH$_{3}$ line ($\\sim$\\,0.16\\,km\\,s$^{-1}$; Rosolowsky et al. 2008), as well as the fact that no point-like source is detected in the $Spitzer$ images, resemble the properties of prestellar cores (see Andr\\'{e} et al. 2008). However, as suggested by the SMA CO\\,(2--1) observations, L1448 IRS2E appears to drive a molecular outflow, which implies ongoing accretion onto a central condensation and has never seen before in prestellar cores. Furthermore, the estimated ratio of $I$[C$^{18}$O(1$-$0)] (Hatchell et al. 2005) to $I$[N$_2$H$^+$(1$-$0)] (Kirk et al. 2007) in the IRS2E core is $\\sim$\\,0.26, similar to that of `evolved' prestellar cores, like L1544 (see Tafalla 2005), which suggests that the IRS2E core is chemically evolved and probably already passed the last stage of the prestellar phase.  \\noindent{\\bf Comparison to Class\\,0 Objects:} Class\\,0 objects are the youngest accreting protostars with an age of a few~$\\times$\\,10$^{4}$\\,yr. These objects are in an early evolutionary stage, right after point mass formation, when most of the mass of the system is still in the surrounding dense core/envelope (Andr\\'{e} et al. 2000). They represent the truly hydrostatic protostellar objects (i.e., the second core) formed in dense cores. So far at least 50 Class\\,0 protostars have been identified (Andr\\'{e} et al. 2000; Froebrich 2005). Most of them are detectable in the $Spitzer$ images (at least in the MIPS bands), and are associated with strong submm and mm dust continuum emission (in both single-dish and interferometric maps). Although the collimated outflow from IRS2E possesses the typical properties of an outflow from a Class\\,0 protostar (see Arce et al. 2007), an obvious difference between L1448 IRS2E and known Class\\,0 protostars (e.g., L1448C, IRS3, and IRS2) is that IRS2E is not visible in the sensitive $Spitzer$ images, has weak dust continuum emission, and consequently has an extremely low bolometric luminosity ($<$\\,0.1\\,$L_\\odot$). The estimated age of L1448 IRS2E (a~few\\,$\\times$\\,10$^{3}$\\,yr) is also much less than those of the Class\\,0 protostars, suggesting that IRS2E is younger (less-evolved) than Class\\,0 protostars.  Furthermore, we compare L1448 IRS2E to another source in the Perseus molecular cloud: SVS\\,13B (see Chen et al. 2009 and references therein). Like L1448 IRS2E, SVS\\,13B has no point-like infrared emission at wavelengths from 3.6 to 70\\,$\\mu$m in the $Spitzer$ images (also c2d data). However, it must be noted that SVS\\,13B is located $\\sim$\\,15$''$ to the south of the bright Class\\,I object SVS\\,13A, and thus the detection limits in the $Spitzer$ images around SVS\\,13B are about three times worse than those in the L1448 images (because the imaging backgrounds around SVS\\,13B were raised by the bright source SVS\\,13A). Interestingly, SVS\\,13B is also driving a collimated outflow seen in the high angular resolution SiO and CO images (Bachiller et al. 1998; 2000). In contrast to L1448 IRS2E, SVS\\,13B has much stronger dust continuum emission at submm and mm wavelengths, and correspondingly has much higher gas mass ($>$\\,1\\,$M_\\odot$)  and bolometric luminosity ($>$\\,1\\,$L_\\odot$). In addition, the kinematic properties of SVS\\,13B, e.g., fast rotation and subsonic turbulence (see Chen et al. 2009), are similar to those of Class\\,0 protostars (e.g., Chen et al. 2007). Therefore, SVS\\,13B is very likely more evolved than L1448 IRS2E and has already formed an extremely young Class\\,0 protostar.  \\noindent{\\bf Comparison to Known VeLLOs:} The extremely low luminosity of L1448 IRS2E is similar to what is seen in the so-called very low luminosity objects (VeLLOs), an interesting subset of embedded, low-luminosity protostars (see Dunham et al. 2008 and references therein). However, non-detection at both 24 and 70\\,$\\mu$m bands distinguishes L1448 IRS2E from all VeLLOs revealed thus far (Dunham et al. 2008). Direct observations, together with radiative transfer modelling, have shown that young (sub-)\\,stellar objects have already formed in these VeLLOs. In contrast, there is yet no clear evidence for the presence of a protostar in L1448 IRS2E, even though the sensitivities of the $Spitzer$ images of L1448 IRS2E are comparable to those used to detect the known VeLLOs (see Dunham et al. 2008 and references therein).  The evolutionary status and eventual final state of VeLLOs are still unclear. Some of them, e.g., IRAM\\,04191+1522 (see Dunham et al. 2006), represent typical Class\\,0 low-mass protostars, while others, e.g. L1014-IRS (see Bourke et al. 2005), could represent precursors of sub-stellar objects (i.e., proto-brown dwarfs). In the case of L1448 IRS2E, it is more likely that we are catching the very first moments of low-mass star formation because L1448 IRS2E already has about 0.04\\,$M_\\odot$ of gas estimated from the SMA dust continuum observations, and more gas in the outer envelope/core can continue accreting onto it. If we assume a steady accretion rate and a core-to-star efficiency of 15--30\\% (Evans et al. 2009), then it is very probable that a low-mass star ($\\geq$\\,0.1\\,$M_\\odot$) will eventually form in the L1448 IRS2E core.  \\subsection{A Candidate First Hydrostatic Core}  The observational detection of the first hydrostatic core is of prime importance for understanding the early evolution of star-forming dense cores and the origin of outflows. Encouraged by these facts, searches for the first core have been undertaken over the past decade. Based on the HCO$^+$/H$^{13}$CO$^+$ observations, Onishi et al. (1999) suggested that L1521F could be a first core candidate, but $Spitzer$ observations soon found that L1521F harbors a low luminosity protostar (Bourke et al. 2006). More recent studies suggest that the evolutionary stage of L1521F is similar to or younger than the Class\\,0 phase, and may be consistent with the early second collapse phase (Shinnaga et al. 2009; Terebey et al. 2009). Another promising object was Cha-MMS1, suggested by Belloche et al. (2006) from the measurement of the deuterium fractionation. However, a mid-infrared source was detected by $Spitzer$ MIPS observations, indicating a compact hydrostatic object had already formed in Cha-MMS (see Belloche et al. 2006 for more details).  Based on the SMA and $Spitzer$ observations, we find that source L1448 IRS2E has the following characteristics: (1) it is not visible in the sensitive $Spitzer$ infrared images (from 3.6 to 70\\,$\\mu$m); (2) has very weak (sub-)\\,mm dust continuum emission, and consequently has an extremely low bolometric luminosity ($<$\\,0.1\\,$L_\\odot$); and (3) appears to drive a molecular outflow. Comparisons with prestellar cores and Class\\,0 protostars suggest that L1448 IRS2E is more evolved than prestellar cores but less evolved than Class\\,0 protostars, i.e., at a stage intermediate between prestellar cores and Class\\,0 protostars. These results are consistent with the theoretical predictions in the RHD/MHD models for the first hydrostatic core (see Section\\,I)\\footnote{In the MHD model of Machida et al. (2008), the outflow driven by the first core has a slow speed of $\\sim$\\,3\\,km\\,s$^{-1}$. However, in their models the first core only has a mass of 0.01\\,$M_\\odot$ at the end of the calculations. Since the first core will grow in mass by at least 1--2 orders of magnitude in the subsequent gas accretion phase, the relatively high outflow velocity of L1448 IRS2E ($\\sim$\\,25\\,km\\,s$^{-1}$; about 8 times larger than that in the model) could be explained by the relatively larger (gas) mass of the L1448 IRS2E core ($\\sim$\\,0.04\\,$M_\\odot$), i.e., a deeper gravitational potential and a faster escape speed.}, making L1448 IRS2E the most promising first hydrostatic core candidate thus far.  However, it must be noted that the nature of source L1448 IRS2E is not definitive. More observations are needed to constrain its SED and to refine its outflow maps. Detections of other objects like L1448 IRS2E will be important for understanding the process of dynamical collapse and the origin of outflows. Sensitive surveys at wavelengths from far-infrared (e.g., $Herschel$) to (sub-)\\,mm continuum (e.g., SCUBA) are needed to search for more first core candidates in nearby clouds. We also speculate that some of the objects in the current sample of prestellar cores may already harbor first cores, which drive molecular outflows hidden within the extended cloud emission and are therefore not revealed in low resolution single-dish observations. A systematic high-resolution interferometric CO survey toward these cores is needed to search for potential outflow activity."
    },
    "1004/1004.2650.txt": {
        "abstract": "Luminous infrared galaxies ($L_{\\rm{IR}}>10^{11} L_{\\odot}$) are often associated with interacting galactic systems and are thought to be powered by merger--induced starbursts and/or dust--enshrouded AGN. In such systems, the evolution of the dense, star forming molecular gas as a function of merger separation is of particular interest. Here, we present observations of the CO(3-2) emission from a sample of luminous infrared galaxy mergers that span a range of galaxy-galaxy separations. The excitation of the molecular gas is studied by examining the CO(3-2)/CO(1-0) line ratio, $r_{31}$, as a function of merger extent. We find these line ratios, $r_{31}$, to be consistent with kinetic temperatures of $T_k$=(30--50)~K and gas densities of $n_{\\rm{H}_2}=10^3 \\, \\rm{cm}^{-3}$. We also find weak correlations between $r_{31}$ and both merger progression and star formation efficiency ($L_{\\rm{fIR}} / L_{\\rm{CO(1-0)}}$). These correlations show a tendency for gas excitation to increase as the merger progresses and the star formation efficiency rises. To conclude, we calculate the contributions of the CO(3-2) line to the 850~$\\mu$m fluxes measured with SCUBA, which are seen to be significant ($\\sim$24\\%). ",
        "introduction": "Luminous Infrared Galaxies (LIGs) ($L_{\\rm{IR}}>10^{11} L_{\\odot}$) and Ultra Luminous Infrared Galaxies (ULIGs) ($L_{\\rm{IR}}>10^{12} L_{\\odot}$) are the dominant population for galaxies with bolometric luminosities greater than $10^{11} L_{\\odot}$ in the local universe $(z \\lesssim 0.3)$ \\citep{1986ApJ...303L..41S}. Observations have shown that most are associated with merging/interacting galaxies which are rich in molecular gas (see \\citet{1996ARA&A..34..749S} for an extensive review). Analysis of the original Infrared Astronomical Satellite (IRAS) data showed that LIGs are quite rare in the local universe but increase significantly with increasing redshift \\citep{1987ApJ...316L..15H,2005ApJ...622..772C}. Studies of \\emph{local} luminous infrared galaxies have, therefore, become essential for a wider understanding of the role of mergers and starbursts both in the local universe and at higher redshifts, where such objects are more commonly found.  Optical studies of local LIGs have shown that many are associated with interacting or merging galactic systems \\citep{1987AJ.....94..831A,1998ApJ...492L.107S,2000ApJ...545..228A,2002ApJS..143..315V}. It is now thought that as a galactic merger progresses, gas and dust from the parent galaxies dissipate energy through shocks, giving rise to an in-fall of gas and dust towards the centre of gravity of the interacting system \\citep{1996ApJ...471..115B,1996ApJ...464..641M,2006MNRAS.373.1013C}. The concentration of cold, dense material then acts to fuel starburst activity, or possibly an active galactic nucleus (AGN), or both \\citep{1998ApJ...508..627K}. The energy produced by the starburst or AGN in the optical and ultraviolet is absorbed and then re-radiated by dust, leading to an intense infrared luminosity.   The effect of a galactic merger upon star formation, as well as its effect upon the total mass and properties of the molecular and atomic gas within the interacting system, has been the subject of several observational studies. \\citet{1999ApJ...512L..99G} used the CO(1-0) spectral line strength as an indicator of total molecular gas mass, $M_{\\rm{H}_2}$ and plotted this quantity against the projected distance of the merging galactic nuclei at optical wavelengths, the \\emph{projected component separation}, in a merging sample of LIGs. The authors found a positive correlation between these two quantities. The authors also found an anti-correlation between the ratio of infrared luminosity to molecular gas mass, $L_{\\rm{IR}}/M_{\\rm{H}_2}$, a measure of star formation efficiency, and the component galaxy separation for their sample. Both observations are consistent with a depletion of gas as a huge increase in star formation (to around $\\sim 100 \\, M_\\odot \\, \\rm{yr}^{-1}$) occurs. In a survey of literature data, \\citet{2000MNRAS.318..124G} found strong evidence for gas depletion and increased star formation efficiency in a merging galaxy sample which included pre--merger and post--merger candidates. The authors also found an increase in the central molecular hydrogen surface density, as traced by CO(1-0), and a decrease in the fraction of cold gas mass as the galactic mergers progressed. This result is consistent with large molecular gas inflows, as predicted by numerical simulations, the conversion of neutral hydrogen to molecular hydrogen, stars and hot gas and molecular hydrogen depletion due to ongoing star formation.  Merging and starburst activity also affects the atomic component of the ISM. \\citet{1989ApJ...340L..53M} have found that the ratio of molecular hydrogen to atomic hydrogen gas mass, $M_{\\rm{H}_2}/ M_{\\rm{HI}}$, increases with $L_{\\rm{IR}}$. It is currently unclear whether this increase is chiefly due to a depletion of HI via ejection from the interacting system, depletion of HI due to photo-ionisation caused by the formation of young stars or enhanced formation of molecular clouds in merger induced shocks. Further observations of $M_{\\rm{HI}}$ have been attempted to investigate whether this depletion of atomic gas is also seen across a merging sequence of LIGs \\citep{2001A&A...368...64V}, but no conclusive answer has yet been derived.  The evolution of molecular hydrogen gas, which is associated with star formation, within merging luminous infrared galaxies is of particular interest. Observations of the CO(1-0) rotational transition are often used to trace the total molecular hydrogen gas mass e.g. \\citet{young:scoville}. After $\\rm{H}_2$, CO is the most abundant molecule in the ISM ($[\\rm{CO}/\\rm{H}_2] \\sim 10^{-4}$), and the CO(1-0) transition is typically thermalised in molecular clouds, making it a very convenient transition to observe. The low excitation energy, $E_{10}/k_B = 5.5$ K and the low critical density $n_{\\rm{H}_2,\\rm{crit}} \\sim 410\\, \\rm{cm}^{-3}$ of the CO(1-0) transition make it ideal for tracing the bulk metal--enriched $\\rm{H}_2$ in galaxies. These same properties, however, make CO(1-0) insensitive to the physical conditions of the molecular gas. If we require constraints on the density and kinetic temperature, we can observe the higher-J CO lines, (i.e. $\\rm{CO}_{J+1 - J}$, where $J\\geq2$) \\citep{2003ApJ...588..243Z}.  These transitions will only become thermalised and luminous for the denser and warmer gas components and it is these dense components which are a direct indication of the star formation potential of the interacting systems.  The higher order CO(3-2) transition, in particular, is commonly used to trace the warmer, denser components of the ISM associated with star formation \\citep{1999A&A...341..256M,2006A&A...460..467B,2009ApJ...693.1736W} and constrain the kinetic temperature and density of the molecular gas. It is likely that there will be multiple density components of molecular gas within the ISM of luminous infrared galaxies, and the critical density of CO(3-2) \\citep[$n_{\\rm{H}_2,\\rm{crit}}> 8.4 \\times 10^3 \\, \\rm{cm}^{-3}$,][]{1995PhDT.......310J} is well matched to the density of the \\emph{star forming} component. The relative ratios of the higher-order CO line intensities depend on both the density and the kinetic temperature of the molecular gas \\citep{1974ApJ...187L..67S,1974ApJ...189..441G}.  By measuring a range of CO line ratios one can therefore begin to constrain the likely conditions of the molecular interstellar medium. Observing the CO(3-2) transition is convenient as it traces mostly the starforming molecular gas and has relatively large luminosities compared to other commonly used tracers, such as HCN ($n_{\\rm{H}_2,\\rm{crit}}> 10^5 \\, \\rm{cm}^{-3}$), which have higher dipole moments, but lower abundances in the ISM. Measuring the CO(3-2) emission is also important since this emission can contaminate 850~$\\mu$m continuum observations of the dust component on the ISM for these galaxies \\citep{2000ApJ...537..631P}. Correcting this 850~$\\mu$m flux for CO emission is important when using this flux to determine the dust mass and kinetic temperature \\citep{2004MNRAS.349.1428S}.  In this paper, we present observations of the CO(3-2) rotational transition to investigate the evolution of the \\emph{dense} ($n_{\\rm{H}_2,\\rm{crit}}> 8.4 \\times 10^3 \\, \\rm{cm}^{-3}$) molecular gas component as a function of merger extent in a sample of LIG mergers. Our aim is to study the evolution of merging luminous infrared galaxies by observing a LIG sample at differing stages of interaction, and thus obtain a detailed understanding of how the starburst evolves as the interaction proceeds. Such observations are key to obtaining a \\emph{temporal} understanding of the evolution of the molecular gas reservoir and how it relates to star formation activity as the merging component galaxies become increasingly tidally disrupted. Our study builds on work reported by \\citet{1999ApJ...512L..99G}, in which a sample of luminous galaxies was chosen to represent a merging sequence with a range of nuclear separations of the merging components. We observe the CO(3-2) transition in a subset of this merging sample of 49 LIGs to trace the evolution of the warmer, denser molecular gas component as a function of merger separation. In particular, we investigate how the excitation of the molecular gas changes by examining the evolution of the {CO(3-2)/CO(1-0)} line ratio, $r_{31}$, as merging progresses. $r_{31}$ will be larger for warmer, denser gas and one might expect its value to increase with decreasing nuclear separation, as an increasing fraction of the total molecular gas becomes concentrated in star forming regions. We then examine how the line ratio varies with merging extent, far-infrared luminosity ($L_{\\rm{fIR}}$, defined in Section \\ref{observations}) and star formation efficiency ($L_{\\rm{fIR}} / L_{\\rm{CO(1-0)}}$), and compare our sample with two samples with smaller infrared luminosities.  Finally we determine contamination arising from the CO(3-2) line flux upon the 850~$\\mu$m continuum observations of the dust component of the ISM for these galaxies. Section \\ref{observations} outlines the observations that were made, including the sample selection criterion, the choice of pointings and calibration of the data. The data reduction and results of the observations are presented in Section \\ref{results} and their interpretation and correlations with other data for the sample are discussed in Section \\ref{interpretation}.  Section \\ref{conclusions} outlines the conclusions which may be drawn from this study. ",
        "conclusions": "\\label{conclusions}   We have observed the CO(3-2) emission from a sample of IR-selected merging galaxies using the single-pixel 350 GHz receiver at the JCMT. We have found that the CO(3-2) to CO(1-0) line ratio, $r_{31}$, is less than unity for all but one object in our merging sample with a sample average of $r_{31}=0.47$. LVG modelling indicates that this is consistent with $T_k$=(30--50) K, $n=10^3 \\rm{cm}^{-3}$ and $\\Lambda \\equiv [\\rm{CO/H_2}]/(dV/dR) =10^{-5}$ $(\\rm{km \\, s^{-1}}\\, \\rm{pc}^{-1})^{-1} $, although warmer ($T_k \\sim$ 80--110~K) and more diffuse $(n \\sim 3 \\times 10^2 \\rm{cm}^{-3})$ conditions for the gas phase are not ruled out.  We have calculated the CO(3-2) contribution to the 850~$\\mu$m continuum flux for the galaxies in our sample. We have found that this contribution can be a large fraction of the 850~$\\mu$m flux underlining the importance of making CO(3-2)-based flux corrections when deriving SED fits or dust masses from 850~$\\mu$m fluxes. Our correction will be presented in a forthcoming paper presenting the 850~$\\mu$m SCUBA data for this galaxy set.  The spread of line ratios observed in our merging sample was broadly similar to that measured in both an IR-selected subsample of the SCUBA Local Universe Galaxy Survey (SLUGS) \\citep{2003ApJ...588..771Y}, and a sample of normal galaxies \\citep{1999A&A...341..256M}. We observe a weak anti-correlation between $\\log(r_{31})$ and $\\log(\\rm{sep; kpc})$ in our sample, a result consistent with gas excitation increasing as the merger progresses. A plot of line ratio versus star formation efficiency ($\\log(r_{31})$ vs.  $\\log(L_{\\rm{fIR}} / L_{\\rm{CO(1-0)}} )$) also shows a weak correlation, as gas excitation increases with increased star formation efficiencies. These correlations are suggestive of an increase in the excitation of the molecular gas, as measured by the CO(3-2)/CO(1-0) line ratio, with the progression of galactic mergers. However, studies of line ratios of these kind are often unavoidably affected by uncertainties in line luminosities arising both from random errors and systematic errors in the absolute flux calibration, which can be as high as 20\\%. Uncertainties also arise from assumptions that often must be made about the angular size and shape of the CO(3-2) emission regions relative to the telescope's beam. These latter uncertainties will hopefully be reduced as more small sources ($D_A \\sim$ a few arcseconds) are mapped interferometrically, and as more extended sources ($D_A \\gtrsim 10''$) are mapped with the next generation of multi-element focal-plane heterodyne arrays, such as HARP-B on the JCMT.  Our future work includes an LVG analysis in which we will combine our CO(3-2) data with HCN line measurements for the same galaxy set to systematically probe the gas properties across the merging sequence. We also intend improve our total luminosity estimates for our more extended sources by mapping them with the newly commissioned HARP-B focal-plane array.     \\appendix"
    },
    "1004/1004.4584.txt": {
        "abstract": " ",
        "introduction": "Anomalies in the aphelia distribution and orbital elements of Outer Oort cloud comets led to the suggestion that $\\approx$ 20\\% of these comets were made discernable due to a weak impulse from a bound Jovian mass body  (\\cite{mww99}). Since that time the data base of comets has doubled. Further motivation for an updated analysis comes from the recent launch of the Wide-field Infrared Survey Explorer (WISE; \\cite{wise}), which could easily detect the  putative companion orbiting in the outer Oort cloud. Such an object would be incapable of creating comet ``storms\". To help mitigate popular confusion with the Nemesis model (\\cite{wj84}, \\cite{dhm84}) we use the name recently suggested by Kirkpatrick and Wright (2010), Tyche, (the good sister of Nemesis) for the putative companion.   The outer Oort cloud (OOC) is formally defined as the ensemble of comets having original semimajor axes $A \\geq 10^4$ AU (\\cite{oort50}). It has been shown that the majority of these comets that are made discernable are first-time entrants into the inner planetary region (\\cite{fern81}) and these comets are therefore commonly referred to as {\\em new}. The dominance of the galactic tide in making OOC comets discernable at the present epoch has been predicted on theoretical grounds (\\cite{ht86}). Observational evidence of this dominance has been claimed to be compelling (\\cite{del87,mw92,wt99,ml04}).  Matese and Lissauer (2004) adopted an  {\\em in situ} energy distribution similar to the initial distribution of Rickman {\\em et al.} (2008) and took the remaining phase space external to a ``loss cylinder\"  to be uniformly populated at the present epoch.  The distribution of cometary orbital elements made discernable from the tide alone was then obtained and compared with observations.  Similar modeling (\\cite{ml02}) had been performed including single stellar impulses which mapped the comet flux over a time interval of 5 Myr, in 0.1 Myr intervals. Peak impulsive enhancements $\\ge 20\\%$ were found to have a half-maximum duration of $\\approx 2$ Myr and occurred with a mean time interval of $\\approx 15 $ Myr. Various time-varying distributions of elements were compared with the modeled tide-alone results  and inferences about the signatures of a weak stellar impulse were drawn.    In Section 2 we review  a discussion (\\cite{ml04}) of a subtle characteristic of galactic tidal dominance which is difficult to mimic with observational selection effects or bad data. Along with the more well known feature of the deficiency of major axis orientations in the direction of the galactic poles and equator, we compare with observations these predictions based on the tidal interaction alone and show that the data are of sufficiently high quality to unambiguously demonstrate the dominance of the galactic tide in making comets discernable at the present epoch. A critique of objections to this assertion (\\cite{rffv08}) is also presented. More recent detailed modeling (\\cite{kaib09}) provide important insights into the evolving populations of the {\\em in situ} and discernable populations of the Oort cloud. We comment further on these works in this section.  In Section 3 we describe the theoretical analysis combining a secular approximation for the galactic tide and for a point mass perturber, describing how a {\\it weak} perturbation of OOC  comets would manifest itself observationally. Evidence suggesting that there is such an aligned impulsive component of the observed OOC comet flux has been previously reported (\\cite{mww99}). It has been found that none of the known observational biases can explain the alignment found there (\\cite{he02}). The size of the available data  has since doubled which leads us to review the arguments here.  In Section 4 we present the supportive evidence that an impulsive enhancement in the new comet flux of $\\approx 20\\%$ persists in the updated data. We also discuss dynamical and observational limits on parameters describing the putative companion. Section 5 summarizes our results and presents our conclusions. ",
        "conclusions": "We have described how the dynamics of a dominant galactic tidal interaction, weakly aided by an impulsive perturbation, predicts specific properties for observed distributions of the galactic orbital elements of outer Oort cloud comets. These subtle predictions have been found to be manifest in high-quality observational data at statistically significant levels, suggesting that the observed OOC comet population  contains an $\\approx 20\\%$ impulsively produced excess. The extent of the enhanced arc is inconsistent with a weak stellar impulse, but is consistent with a Jovian mass solar companion orbiting in the OOC. A putative companion with these properties may also be capable of producing detached Kuiper Belt objects such as Sedna and has been given the name Tyche. Tyche could have significantly depleted the inner Oort cloud over the solar system lifetime requiring a corresponding increase in the inferred primordial Oort cloud population. A substantive difficulty with the Tyche conjecture is the absence of a corresponding excess in the presumed IOC daughter population.   \\begin{center} {\\bf ACKNOWLEDGMENTS} \\end{center} The authors thank Jack J. Lissauer for his continuing interest and for his contributions to this research."
    },
    "1004/1004.0247_arXiv.txt": {
        "abstract": "We study, using both theory and molecular dynamics simulations, the relaxation dynamics of a microcanonical two dimensional self-gravitating system. After a sufficiently large time, a gravitational cluster of $N$ particles relaxes to the Maxwell-Boltzmann distribution. The time to reach the thermodynamic equilibrium, however, scales with the number of particles. In the  thermodynamic limit, $N\\rightarrow\\infty$ at fixed total mass, equilibrium state is never reached and the system becomes trapped in a non-ergodic stationary state.   An analytical theory is  presented which allows us to quantitatively described this final stationary state, without any adjustable parameters. ",
        "introduction": "Systems interacting through long-range forces behave very differently from those in which particles interact through short-range potentials. For systems with short-range forces, for arbitrary initial condition, the final stationary state corresponds to the thermodynamic equilibrium and can be  described equivalently by either microcanonical, canonical, or grand-canonical ensembles. On the other hand, for systems with unscreened long-range interactions, equivalence between ensembles breaks down~\\cite{Gibbs,Barre2001}. Often these systems are  characterized by a negative specific heat~\\cite{Ly67,Thir70,Ly77} in the microcanonical ensemble and a broken ergodicity~\\cite{Muka2005,Rami2008}.  In the infinite particle limit, $N \\rightarrow \\infty$, these systems never reach the thermodynamic equilibrium and become trapped in a stationary out of equilibrium state (SS)~\\cite{Kl54,Bouch2005}.  Unlike normal thermodynamic equilibrium, the SS does not have Maxwell-Boltzmann velocity distribution. For  finite $N$, relaxation to  equilibrium proceeds in two steps.  First, the system relaxes to a quasi-stationary state (qSS), in which it stays for time $\\tau_\\times(N)$, after which it crosses over to the normal thermodynamic equilibrium with the Maxwell-Boltzmann (MB) velocity distribution~\\cite{Kav2007}. In the limit $N \\rightarrow \\infty$, the life time of qSS diverges, $\\tau_\\times \\rightarrow \\infty$, and the thermodynamic equilibrium is never reached.   Unlike the equilibrium state, which only depends on the global invariants such as the total energy and momentum and is independent of the specifics of the initial particle  distribution, the SS explicitly depends  on the initial condition.  This is the case for self-gravitating systems~\\cite{Pa90}, confined one component plasmas~\\cite{Levinprl,Rizz09},  geophysical systems~\\cite{Cha05}, vortex dynamics~\\cite{Miller90,Cha96,Venaille09}, etc~\\cite{Campa09}, for which the SS state often has a  peculiar core-halo structure~\\cite{Levinprl}. In the thermodynamic limit, none of these systems can be described by the usual equilibrium statistical mechanics, and new methods must be developed.  In this paper we will restrict our attention to self-gravitating systems. Unfortunately, it is very hard to study  these systems in 3d~\\cite{Levingrav,Joyce09}. The reason for this is that the 3d Newton potential is not confining.  Some particles can gain enough energy to completely escape from the gravitational cluster, going all the way to infinity. In the thermodynamic limit, one must then consider three distinct populations: particles which will relax to form the central core, particles which will form the halo, and  particles which will completely evaporate.  Existence of three distinct classes of particles, makes the study of 3d systems particularly difficult. On the other hand, the interaction potential in 2d is logarithmic, so that all the particles remain gravitationally bound. Similar to  magnetically confined plasmas the stationary state of a 2d gravitational system should, therefore, have a core-halo structure~\\cite{Levinprl}. We thus expect that the  insights gained from the study of confined plasmas might  prove to be useful to understand the 2d gravitational systems. ",
        "conclusions": "We have studied the thermodynamics of 2d self-gravitating system in the microcanonical ensemble. It was shown that the gravitational clusters containing  finite number of particles relax to the equilibrium state characterized by the MB distribution.  Prior to achieving the thermodynamic equilibrium, however, these systems become trapped in a quasi-stationary state,  where they stay for time $\\tau_\\times$, which diverges as $N^{1.35}$ for large $N$.  Thus, in the limit $N \\rightarrow \\infty$ at fixed total mass $M$, thermodynamic equilibrium can not be reached in a finite time.  A new approach, based of the conservation properties of the Vlasov dynamics and on the theory of parametric resonances, is formulated and allows us to quantitatively predict the one particle distribution function in the non-equilibrium stationary state. Finally, it is curious to consider what will happen to a self-gravitating system in a contact with a thermal bath --- the canonical ensemble. In Appendix \\ref{Virial}, it is shown that for a 2d self-gravitating system a stationary state is possible, if and only if, $\\langle v^2 \\rangle$=1/2, i.e. when the kinetic temperature is $T=1/4$.  If such system is put into contact with a thermal bath which has $T>1/4$ there will be a constant heat flux from the reservoir into the system. This heat will be converted into the gravitational potential energy --- since the kinetic energy is fixed by the virial condition ---  making the cluster expand without a limit. Conversely if the bath temperature is  $T<1/4$, the heat flux will be from the system into the bath. Again, since the system can only exist in a stationary state if $T=1/4$, the energy for the heat flux can come only from the gravitational potential.  In this case the gravitational cluster will contract without a limit, concentrating all of its mass at the origin.  Thus,  in the canonical ensemble no thermodynamic equilibrium is possible, unless the reservoir is at exactly $T=1/4$. We hope that the present work will also help shed new light on the collisionless relaxation in 3d self-gravitating systems.  Unfortunately the 3d problem  is significantly more difficult, since besides the core-halo formation, one must also account for the particles evaporating to infinity.  This work was partially supported by the CNPq, INCT-FCx, and by the US-AFOSR under the grant FA9550-09-1-0283.  \\clearpage \\appendix"
    },
    "1004/1004.0137_arXiv.txt": {
        "abstract": "Elliptical, lenticular, and early-type spiral galaxies show a remarkably tight power-law correlation between the mass $M_\\bullet$ of their central supermassive black hole (SMBH) and the number $N_{GC}$ of globular clusters: $M_\\bullet = m_{\\bullet/\\star} \\times N_{GC}^{1.08 \\pm 0.04}$ with $m_{\\bullet/\\star} = 1.7\\times 10^5M_{\\odot}$. Thus, to a good approximation the SMBH mass is the same as the total mass of the globular clusters.  Based on a limited sample of 13 galaxies, this relation appears to be a better predictor of SMBH mass (rms scatter 0.2 dex) than the $M_\\bullet$--$\\sigma$ relation between SMBH mass and velocity dispersion $\\sigma$. The small scatter reflects the fact that galaxies with high globular cluster specific frequency $S_N$ tend to harbor SMBHs that are more massive than expected from the $M_\\bullet$--$\\sigma$ relation. ",
        "introduction": "Supermassive black holes (SMBHs) have been detected in the centers of many nearby galaxies \\citep{kr95,mag98,gul09}. The SMBH masses are correlated with several properties of their host galaxies \\citep{nov06}, in particular the velocity dispersion (the $M_{\\bullet}$--$\\sigma$ relation, e.g., \\citealt{fm00,geb00,tre02,gul09}) and the mass and luminosity of the spheroidal component---the entire galaxy in the case of ellipticals or the bulge in the case of lenticular and spiral galaxies \\citep{kor93,kr95,mh03,hr04}. \\cite{sf09} find a tight correlation between SMBH and dark matter halo masses. As the dark halo properties are inferred from the number of globular clusters in a galaxy this also indicates a connection between globular clusters and SMBHs. These correlations suggest a strong link between SMBH formation and galaxy formation, although the nature of this link is poorly understood.  Numerous authors have investigated the possibility that the growth of galaxies and their SMBHs is regulated by their interactions \\citep[e.g.,][]{hk00,bk01,som08,cat09}. SMBHs grow by several mechanisms, including accretion of gas, swallowing stars whole, or merging with other SMBHs acquired through a merger of their host galaxies. The So\\l tan argument \\citep{sol82,yt02} suggests that gas accretion is the dominant contributor to the SMBH mass budget, and both observations \\citep{san88} and simulations \\citep{hop06} suggest that much or most of this accretion occurs during mergers. SMBH growth through gas accretion can release substantial amounts of energy that can heat the interstellar gas, quench star formation, or even drive a wind that sweeps the galaxy free of gas, thereby halting star formation completely. Simulations of gas-rich galaxy mergers, including seed black holes, can reproduce the observed $M_{\\bullet}$--$\\sigma$ relation remarkably well given the simplicity of the empirical prescriptions used to model the accretion and feedback processes \\citep{smh05,cox06,cro06,hop08,joh09a,joh09b}.  The origin of the seeds of SMBHs is controversial.  According to one hypothesis, the seeds were remnants of the first generation of metal-free, massive stars that formed at high redshifts. However, the existence of bright quasars at redshifts of $z\\gtrsim 6$ demonstrates that SMBHs with masses exceeding $10^9M_\\odot$ were already in place less than a billion years after the Big Bang.  It is difficult for black-hole remnants from first-generation stars to grow fast enough to explain these observations \\citep{may09}. A second hypothesis is that the seeds are much larger (``intermediate-mass'') black holes of $10^2$--$10^5M_\\odot$, perhaps formed by the direct collapse of gas at the centers of protogalaxies.  Globular clusters (GCs) are among the oldest stellar systems in the universe and may have formed at the same time as the first stars. Their high stellar densities, sometimes exceeding $10^5 M_\\odot\\hbox{ pc}^{-3}$, lead to a variety of complex dynamical phenomena \\citep{spi87,hh03,bt08}. Among these is mass segregation, through which heavy, compact, stellar remnants---neutron stars and black holes---spiral into the center by dynamical friction. Once these arrive at the center it is possible, though far from certain, that they merge to form an intermediate-mass black hole \\citep{lee87,qs87,pz04,ku05}.  There is significant observational evidence for intermediate-mass BHs with masses $4\\times10^3$--$4\\times10^4M_\\odot$ in the centers of several GCs \\citep{ger02,grh05,noy08,vdm10}, but this evidence is still controversial\\footnote{An additional uncertainty is whether some of these systems might be tidally stripped dwarf galaxies masquerading as GCs.}.  The number of globular clusters in a galaxy, $N_{GC}$, is roughly proportional to the total luminosity of the galaxy's spheroidal component. This relation was quantified by \\cite{hvdb81}, who introduced the specific globular cluster frequency $S_N$, defined as the number of GCs per unit absolute visual magnitude $M_V=-15$, \\begin{equation} S_N \\equiv N_{GC} \\times 10^{0.4(M_V+15)} \\label{eq:sn} \\end{equation} where $M_V$ is the magnitude of the spheroidal component.  \\cite{bs06} have summarized the progress that has been made in the quarter-century since the work by Harris \\& van den Bergh. It has become clear that star cluster populations are powerful tracers of galaxy evolution and that the observed correlations between globular cluster and galaxy properties provide valuable information about their joint formation.  One of the most comprehensive studies of early-type galaxies is by \\cite{peng08}, who measured specific frequencies for the globular cluster systems of 100 elliptical and lenticular galaxies in the Virgo cluster. They find that early-type galaxies with intermediate luminosities ($-22 < M_V < -18$) typically have $S_N \\sim 1.5$, while luminous galaxies have $S_N\\sim 2$--5. The dominant galaxy M87 has an even larger specific frequency \\citep{rac68}, estimated by Peng et al.\\ to be $S_N\\simeq 13$.  The formation of GCs is not well-understood (see, e.g., \\citealt{bs06} for a review). An important clue is that gas-rich merging galaxies contain large numbers of young massive star clusters that presumably formed in the merger \\citep{sch87,ws95}. As this population of clusters ages it is likely to evolve into a population of ``normal'' GCs \\citep{fz01}. Another scenario is the combined formation of SMBH seeds and globular clusters in super star-forming clumps of gas-rich galactic disks at $z\\sim2$ \\citep{sgf10,mp96}.  In summary, (i) both the SMBH mass $M_\\bullet$ and the total number of GCs $N_{GC}$ are roughly proportional to the total luminosity of the spheroidal component in early-type galaxies; (ii) GCs may provide the black-hole seeds from which SMBHs grow; (iii) both the growth of SMBHs and the formation of GCs appear to be associated with major mergers or global gravitational instabilities in gas-rich protogalaxies. Given these observations, it is natural to ask how the properties of the GC population in early-type galaxies are correlated with the properties of their associated SMBHs.  In this paper we show that there is a tight, power-law relation between the mass of SMBHs and the total number of globular clusters in elliptical, lenticular and early-type spiral galaxies. Remarkably, this relation appears to have even less scatter than the classic relation between SMBH mass and the velocity dispersion of the host galaxy. The relation can be approximately characterized by the statement that the SMBH mass equals the total mass in GCs. ",
        "conclusions": "We have found that there is a strong correlation between the number of globular clusters and the mass of the central SMBH in early-type galaxies.  This correlation appears to be at least as tight as the well known correlation between velocity dispersion and SMBH mass, although this conclusion is based on only 13 galaxies. To a reasonably good approximation, the BH-GC correlation simply says that the mass of the central SMBH in an early-type galaxy is equal to the mass of its GCs (eq.\\ \\ref{eq:five}). We suspect that the proportionality of the SMBH mass to the total globular-cluster mass offers insight into their formation processes, but the near-equality of the masses is a coincidence.  Most galaxies have GC populations with a bimodal color distribution: there are red (metal-rich) and blue (metal-poor) peaks, presumably reflecting two sub-populations of GCs (e.g., \\citealt{bs06}). It is interesting to investigate whether the SMBH mass is correlated with one or the other of these sub-populations. Table 1 shows the red cluster fraction $f_{\\rm red}$ for 11 galaxies, taken from \\cite{peng08} and \\cite{rz04}. Note that $f_{\\rm red}$ is rather constant, with mean and standard deviation $0.3\\pm0.1$. The lower right panel of Figure \\ref{fig:two} shows the $M_{\\bullet}$--$N_{GC}$ correlation separately for the blue (triangles) and red (circles) clusters. As expected from the small rms variation in the red cluster fraction both correlations are of similar quality.  The origin of the $M_{\\bullet}$--$N_{GC}$ relation is obscure. One possibility is that both the growth of SMBHs and the formation of GCs are associated with major mergers, so that galaxies that experienced a recent major merger will have anomalously large SMBH masses and GC populations. Another possibility is the correlated formation of SMBH seeds and globular clusters in gas-rich young galaxies.  An important next step is to expand the sample of galaxies having both reliable SMBH masses and reliable globular-cluster populations.  \\noindent {\\bf Acknowledgments:} We thank Jeremiah Ostriker, Karl Gebhardt, Simon White and Leslie Sage for interesting discussions and Karl Gebhardt for the use of unpublished data. We also thank the referee, John Kormendy, for comments that substantially improved the paper. The research of A.B. is supported by a Max Planck Fellowship and by the DFG Cluster of Excellence ``Origin and Structure of the Universe''. S.T.'s research is supported by NSF grant AST-0807432 and NASA grant NNX08AH24G."
    },
    "1004/1004.4714_arXiv.txt": {
        "abstract": "Recent observations of the supernova remnant W44 by the \\emph{Fermi }spacecraft observatory strongly support the idea that the bulk of galactic cosmic rays is accelerated in such remnants by a Fermi mechanism, also known as diffusive shock acceleration. However, the W44 expands into weakly ionized dense gas, and so a significant revision of the mechanism is required. In this paper we provide the necessary modifications and demonstrate that strong ion-neutral collisions in the remnant surrounding lead to the steepening of the energy spectrum of accelerated particles by \\emph{exactly one power}. The spectral break is caused by Alfven wave evanescence leading to the fractional particle losses. The gamma-ray spectrum generated in collisions of the accelerated protons with the ambient gas is also calculated and successfully fitted to the Fermi Observatory data. The parent proton spectrum is best represented by a classical test particle power law $\\propto E^{-2}$, steepening to $E^{-3}$ at $E_{br}\\approx7GeV$ due to deteriorated particle confinement. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.3302_arXiv.txt": {
        "abstract": "Single field inflationary models predict nearly Gaussian initial conditions and hence a detection of non-Gaussianity would be a signature of the more complex inflationary scenarios. In this paper we study the effect on the cosmic microwave background and on large scale structure from primordial non-Gaussianity in a two-field inflationary model in which both the inflaton and curvaton contribute to the density perturbations.  We show that in addition to the previously described enhancement of the galaxy bias on large scales, this setup results in large-scale stochasticity.  We provide joint constraints on the local non-Gaussianity parameter $\\tilde f_{\\rm NL}$ and the ratio $\\xi$ of the amplitude of primordial perturbations due to the inflaton and curvaton using WMAP and SDSS data. ",
        "introduction": "One of the most important questions that cosmology faces today is the origin of structure in the universe. The generally accepted paradigm is that of inflation~\\cite{Infl1,Infl2,Infl3,Infl4} which produces small adiabatic perturbations that evolve into the observed structure. The inflationary paradigm is extremely powerful as it remedies most of the problems of the original Big Bang scenario and also has a set of predictions that are well confirmed by current observations. On the other hand, although the generic predictions of inflation are quite clear, the nature of specific physical processes that govern inflation are still poorly understood.  The major obstacle in understanding inflation is that it can not be directly observed either in the laboratory or with telescopes. This problem is at the same time a virtue of inflation as it allows to indirectly probe physics at energies and time-scales that are far beyond the reach of current facilities. By comparing astrophysical observations with predictions of various inflationary models one can expect to distinguish between different extensions of the Standard Model of particle physics~\\cite{LythReview}. Understanding of the reheating phase of inflation can provide a link between scalar fields driving inflation and the observable Universe that consists of dark and baryonic matter.  One of the many possible ways to deeper understand inflation is by studying the primordial density fluctuations. The usual inflationary model of a slowly-rolling inflaton field requires that the perturbations are highly Gaussian~\\cite{Gaus1,Gaus2,Gaus3,Gaus4} and hence the detection of non-Gaussianity in either the cosmic microwave background (CMB) spectrum or the large scale structure (LSS) distribution would be a clear evidence that the physics driving inflation is more complicated than the standard inflaton scenario.  Non-Gaussianity naturally arises in inflationary models with more than one field~\\cite{MField1,MField2,MField3,MField4}. One of the most studied models is the curvaton model~\\cite{MField4,Curv1,Curv2,Curv3,Curv4,Curv5}, in which initial perturbations are generated by the curvaton field after inflation is over. In this model significant non-Gaussianity can be generated since the predicted curvature perturbation is proportional to the square of the curvaton field (as distinct from single-field inflation, where the required smoothness of the inflaton potential renders the curvature perturbation very nearly linear in the field fluctuations).  Most attempts to constrain non-Gaussianity have used the so-called ``local-type'' or $f_{\\rm NL}$ parameterization~\\cite{Komatsu00} in which one includes a quadratic term into the primordial potential, $\\Phi = \\phi + f_{\\rm NL}\\phi^2$.  In this parametrization both linear and quadratic terms in the potential originate from the same Gaussian field, e.g. a curvaton field, and the contributions from perturbations in other fields (e.g. the inflaton field responsible for inflation itself) are negligible.  The signature of local-type non-Gaussianity in the CMB has been described at length \\cite{CMBNG1}. It has also been established that $f_{\\rm NL}$ has an effect on the galaxy bispectrum \\cite{Verde00,Sefusatti07,Jeong09}. The effect on the large-scale galaxy power spectrum has been considered only recently \\cite{Dalal08, Slosar08, Carbone08, Afshordi08, McDonald08}, but it rapidly became clear that the method was competitive, stimulating work on $N$-body simulations of halo formation in non-Gaussian cosmologies \\cite{Desjacques09, Grossi, P10, Reid}. Recent constraints have been derived from the CMB bispectrum as measured by WMAP \\cite{Komatsu03, Spergel07, Yadav08, Komatsu08, Komatsu10, Smith09,i3} and from large scale structure in the Sloan Digital Sky Survey (SDSS) \\cite{Slosar08}.  Recently, $\\sim3\\sigma$ evidence for excess clustering consistent with non-Gaussianity has been identified in the NRAO VLA Sky Survey (NVSS) \\cite{XiaNVSS}.  In this paper we extend the formalism to include the case where both the inflaton and curvaton contribute significantly to the curvature perturbation.  Perturbations generated by the inflaton field are purely Gaussian, while curvaton fluctuations can result in non-Gaussianity if the conversion from curvaton fluctuation $\\delta\\sigma$ to primordial potential $\\Phi$ contains quadratic terms.  The ratio of inflaton to curvaton contributions $\\xi$ is arbitrary: the framework of the curvaton model allows it to take on any positive value.  Usually one takes $\\xi\\gg 1$ since in the opposite limit ($\\xi\\ll 1$) the curvaton has no observable effect on the primordial perturbations.  Here we investigate the consequences of general $\\xi$ -- including values of order unity -- for the CMB and LSS.  The type of non-Gaussianity generated could be called ``local-stochastic,'' in that it results from local nonlinear evolution of the inflaton and curvaton fields (and thus the primordial bispectrum will have the local-type configuration dependence), but that the full nonlinear potential $\\Phi$ is not a deterministic function of the linear potential.    Studying non-Gaussianity is particularly important in the face of the current generation of CMB projects~\\cite{CMBTF} such as the {\\slshape Planck} satellite as well as ongoing and future LSS projects.  To fully exploit the potential of the future probes it is imperative to investigate theoretically the range of types of non-Gaussianities that can be produced in unconventional inflation (e.g. multi-field models), and understand what effect they have on the CMB and LSS.  The rest of the paper is organized as follows. In Sec.~\\ref{sec:theory} we discuss the generation of non-Gaussian primordial perturbations in the inflationary model with both inflaton and curvaton fields contributing to the curvature perturbation. In Sec.~\\ref{sec:cmb} we describe the effect of two-field models on the CMB bispectrum. In Sec.~\\ref{sec:halo} we derive the halo power spectrum using the peak-background split formalism \\cite{Cole}, and in Sec.~\\ref{sec:lss} we consider the angular power spectrum of galaxies. Section~\\ref{sec:constraints} provides the constraints on the two-field model from existing data, and we conclude in Sec.~\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}  This paper has extended the analysis of non-Gaussianity constraints into a two field inflationary models. In most previous studies of non-Gaussianity it was assumed that primordial density perturbations were generated either by inflaton field, in which case they are perfectly Gaussian, or only by the second field (for example curvaton) which contains quadratic part and generates non-Gaussian initial conditions.  It is important, however, to realize the possibility of an intermediate case where part of the curvature perturbation is derived from quantum fluctuations of the inflaton field, while an additional part is associated with a second field and converted to an adiabatic perturbation upon its decay.  This results in a peculiar type of non-Gaussian initial condition (which we may call ``local-stochastic'' since the field $\\phi_2$ entering in the nonlinear term is correlated with but not identical to the linear potential) that is both observable and distinguishable from the curvaton-only ``local-deterministic'' or $f_{\\rm NL}$ form.  This type of non-Gaussianity has two parameters: a nonlinear coupling coefficient $\\fnl$, and the ratio $\\xi$ of inflaton to curvaton contributions to the primordial density perturbation spectrum. We connect these parameters with parameters characterizing inflationary fields in Eqs.~(\\ref{eq:fnl}) and (\\ref{eq:xi}).  \\begin{figure} \\includegraphics[width=3.3in]{fnl4.eps} \\caption{\\label{fig:2dplot}Constraints in the $(x_1,x_2)$ plane, including both the CMB bispectrum and the galaxy power spectrum.} \\end{figure}  Using the power spectrum and bispectrum constraints from SDSS and WMAP we are able to constrain these parameters.  Adding two sets of constraints together allows us to break the degeneracy in the $(\\fnl, \\xi)$ parameters that exists with the CMB bispectrum alone.  If non-Gaussianity in the CMB is ever detected, and the bispectrum has the local configuration dependence, this will enable us to measure the relative contributions of the inflaton and curvaton.  We have found that in contrast to the local-deterministic non-Gaussianity, whose main effect on large scale structure is a scale-dependent increase in the bias, the local-stochastic non-Gaussianity can introduce stochasticity between the matter and halo distributions.  It can also lead to relative stochasticity between haloes of different masses, since Eq.~(\\ref{eq:chi-k}) depends on the Gaussian bias $b_{\\rm g}$ of the haloes (e.g. $\\chi\\rightarrow 1$ if $b_{\\rm g}\\rightarrow 0$).  The potential use of these effects to directly test the hypothesis of multiple fields contributing to the primordial perturbations is left to future work."
    },
    "1004/1004.3558_arXiv.txt": {
        "abstract": "It has recently been suggested that the presence of a plenitude of light axions, an Axiverse, is evidence for the extra dimensions of string theory. We discuss the observational consequences of these axions on astrophysical black holes through the Penrose superradiance process. When an axion Compton wavelength is comparable to the size of a black hole, the axion binds to the black hole ``nucleus\" forming a gravitational atom in the sky. The occupation number of superradiant atomic levels, fed by the energy and angular momentum of the black hole, grows exponentially. The black hole spins down and an axion Bose-Einstein condensate cloud forms around it. When the attractive axion self-interactions become stronger than the gravitational binding energy, the axion cloud collapses, a phenomenon known in condensed matter physics as ``Bosenova\". The existence of axions is first diagnosed by gaps in the mass vs spin plot of astrophysical black holes. For young black holes the allowed values of spin are quantized, giving rise to \"Regge trajectories\" inside the gap region. The axion cloud can also be observed directly either through precision mapping of the near horizon geometry or through gravitational waves coming from the Bosenova explosion, as well as axion transitions and annihilations in the gravitational atom. Our estimates suggest that these signals are detectable in upcoming experiments, such as Advanced LIGO, AGIS, and LISA. Current black hole spin measurements imply an upper bound on the QCD axion decay constant of $2\\cdot 10^{17}$ GeV, while Advanced LIGO can detect signals from a QCD axion cloud with a decay constant as low as the GUT scale. We finally discuss the possibility of observing the $\\gamma$-rays associated with the Bosenova explosion and, perhaps, the radio waves from axion-to-photon conversion for the QCD axion. ",
        "introduction": "Black holes are among the most fascinating systems in astrophysics, and the most mysterious objects in quantum gravity and string theory, for a long time serving as  principal ``theoretical laboratories\"  for exploring non-perturbative gravitational dynamics. The purpose of  this paper is to initiate a detailed study of the exciting possibility~\\cite{Arvanitaki:2009fg} that astrophysical black holes may serve as  {\\it actual}   laboratories for the discovery of new elementary particles.  There are several reasons why we believe this possibility is realistic. On a purely phenomenological side, black hole observations are routine practice in nowaday astronomy (see, e.g.,  \\cite{Narayan:2005ie} for a review). About 40 stellar mass black holes in X-ray binaries in the Milky Way and neighboring galaxies have been identified with masses in the range $\\sim 5\\div 20 M_\\odot$. Supermassive black holes with masses $\\sim 10^5\\div 10^{10}~M_{\\odot}$ have been found in centers of many galaxies including the Milky Way   and believed to be hosted by nearly all of the galaxies. Also, the first intermediate mass ($\\sim 100\\div 10^{5}~M_{\\odot}$) candidates have been identified.  Following the evolution of binary systems or measuring the velocity dispersion of stars rotating around  galactic centers allows to determine black hole masses. Most crucially for what follows, recent advances in X-ray astronomy and in numerical magnetohydrodynamical simulations of the accreting gas in the Kerr metric open the possibility for  a detailed exploration of  the near-horizon region and, as a consequence, for high precision neasurements of black hole spins~\\cite{McClintock:2009as,Brenneman:2009hs}. First estimates for the angular momentum of several black holes have already been delivered~\\cite{McClintock:2009dn}, often suggesting high values for the spin, although at the moment different techniques sometimes give rise to conflicting results \\cite{Blum:2009ez}.  In the future, apart from improvements of traditional astronomical techniques for observing the near horizon environment and its better theoretical modeling, a unique probe of the black hole geometry will be provided by low frequency gravitational waves observatories, such as LISA \\cite{LISA1} or  AGIS, a gravitational wave detector based on atom interferometric techniques \\cite{Dimopoulos:2008sv,Dimopoulos:2007cj}. For the purpose of testing the near horizon geometry the most promising candidates are the so-called extreme mass ratio inspirals---stellar mass compact objects captured by supermassive black hole in the galactic center (see, e.g., \\cite{Hughes:2006pm}). LISA and AGIS are expected to detect about a hundred of such events per year.  Each such measurement allows not only to determine the black hole spin and mass with an exquisite accuracy, $10^{-3}\\div 10^{-5}$, depending on the details of a particular event, but also to check whether higher order  metric moments, up to $6\\div 7$, agree with their values for the Kerr geometry.  This ongoing observational progress indicates that we are witnessing the dawn of  precision black hole physics. Undoubtedly, black hole observations will be of great value for astrophysics, however it is natural to inquire whether these data may be useful for beyond the Standard Model physics as well, given that  it will provide a rare test of non-linear gravity. However, possibly contrary to naive expectation, it turns out quite challenging to find modified gravity theories which would predict deviations from general relativity near astrophysical black holes and would not contradict current gravity tests. One candidate class of modified theories of gravity affecting black hole dynamics are models of Higgs phases of gravity, where black hole no-hair theorems can be violated~\\cite{Dubovsky:2007zi}.  In this paper we explore a less exotic possibility  to test fundamental physics with precision black hole observations. It is related to the famous Penrose process, a mechanism to extract energy and angular momentum from rotating black holes \\cite{Penrose:1969pc,Christodoulou:1970wf}. As reviewed in detail below, this process, known as superradiance, when applied to waves rather than particles  \\cite{Zeld,Misner:1972kx,Starobinskii}, gives rise to a spin-down instability of a rotating black hole \\cite{Damour:1976kh,Ternov:1978gq,Zouros:1979iw,Detweiler:1980uk}, if a massive boson with a Compton wavelength of order the black hole gravitational radius is present in nature. As we will see, this instability turns rotating astrophysical black holes into sensitive detectors of bosons with masses in the range $\\mu\\sim 10^{-9}\\div 10^{-21}$~eV. Before focusing on the observational consequences of the superradiant instability, let us review why it is natural to expect ultra-light bosons in the theory that transform astrophysical black holes in probes of fundamental physics.   A natural situation giving rise to  a particle of a small, but non-vanishing mass is when this particle is a (pseudo)Goldstone boson of a spontaneously broken global symmetry, which is also explicitly broken by non-perturbative effects. Probably the best motivated candidate for such a particle is the QCD axion $\\phi_a$---a pseudoscalar particle coupled to the QCD instanton number density via \\be \\label{axionaction} S_\\theta={1\\over 32\\pi^2f_a}\\int d^4x\\;\\phi_a\\epsilon^{\\mu\\nu\\lambda\\rho}\\Tr \\,G_{\\mu\\nu}G_{\\lambda\\rho} \\, . \\ee Note that at the classical level $S_\\theta$ is invariant under the Peccei--Quinn (PQ) symmetry $\\phi_a\\to \\phi_a+const$, so that the QCD axion is indeed a (pseudo)Goldstone boson with $f_a$ being the scale of spontaneous symmetry breaking. This symmetry is explicitly broken by the QCD instanton effects that generate the axion potential giving rise to a solution for the strong CP problem---the primary motivation for the QCD axion. As a result the QCD axion acquires a mass equal to \\be \\label{QCDmass} \\mu_a\\approx 6\\cdot 10^{-10} \\mbox{eV} \\l{10^{16}\\mbox{GeV}\\over f_a}\\r  \\;. \\ee The Compton wavelength of the QCD axion with a high symmetry breaking scale $f_a\\gtrsim 10^{16}$~GeV matches the size of stellar mass black holes and, consequently, can affect their dynamics, suggesting that this part of the parameter space for the QCD axion can be explored through black hole observations.  There are several reasons why this conclusion is very important. First, non-gravitational interactions of the QCD axion with the rest of the Standard Model particles are very suppressed at these high values of $f_a$. As a result this part of the parameter space can not be easily probed by any other means, either laboratory or astrophysical. Second, in many ``generic\" string constructions, i.e., in compactifications where the extra-dimensional manifold is neither highly anisotropic, nor strongly warped, the values of $f_a$ are naturally around the grand unification scale $M_{GUT}\\simeq2\\times 10^{16}$~GeV \\cite{Svrcek:2006yi}. Finally, as elaborated in more detail in section~\\ref{anthropic}, finding the QCD axion with $f_a\\sim M_{GUT}$ would indicate that the baryon-to-dark matter ratio varies on length scales longer than the observed part of the Universe and its local value is determined by anthropic reasoning. Discovery of the QCD axion in this regime would be further evidence for enviromental selection already suggested by the cosmological constant problem, and  by the string landscape.   There is an even stronger and more direct link between the QCD axion and the landscape of string vacua, a link that gives rise to the expectation of a {\\it plenitude} of light axion-like particles, an axiverse~\\cite{Arvanitaki:2009fg} -- this same link also suggests the existence of many massless vectors, whose massive superpartners may be discovered at the LHC~\\cite{Arvanitaki:2009hb}. In string constructions, an axion usually arises as a Kaluza--Klein (KK) zero mode of a higher-dimensional  antisymmetric form field. Such zero modes have a purely topological origin: they are labeled by non-contractable cycles in the extra-dimensional manifold. Non-contractable cycles allow for non-trivial gauge field configurations with a vanishing field strength, the so called Wilson lines. These configurations do not carry energy and correspond to zero KK modes at the perturbative level. They only acquire a mass due to non-perturbative effects.  Interestingly, the very same ingredients that give rise to the string axiverse, higher-dimensional form fields and non-trivial cycles in the compactification manifold allowing also to turn on gauge fluxes, also give rise to the string landscape of $10^{500}$ or so vacua. In order to allow for the tuning of the cosmological constant at the $\\sim 10^{-120}$ level, as required by observations, the compactification manifold should contain of order few hundred cycles, given that the total amount of flux quanta for a cycle is typically limited by a number around ten in order to stay in a  perturbative regime. Consequently, one may expect hundreds of axion-like particles in a given string compactification. However, a plenitude of cycles does not yet guarantee the presence of a plenitude of axions.  There is a number of effects in string theory that could produce a large axion mass, such as branes wrapping the cycles, and fluxes. One can roughly estimate the number of light axions as being determined by the number of cycles without fluxes---presumably, around one tenth of the total number of cycles. Still this leaves us with the expectation of several tens of axion-like particles.  The discovery of a plenitude of particles in {\\it our vacuum} with similar properties but different masses supports the idea of a plenitude of {\\it vacua}, as both the axiverse and the multiverse are dynamical consequences of the same fundamental ingredients.  The masses of string axions are exponentially sensitive to the sizes of the corresponding cycles, so one expects them to be homogeneously distributed on the logarithmic scale. However, given that the QCD $\\theta$-parameter is constrained to be less than $10^{-10}$, non-perturbative string corrections to the QCD axion potential should be at least ten orders of magnitude suppressed as compared to the QCD generated potential. It is then natural to expect many of the axions to be much lighter than the QCD axion; these are the axions whose mass is dominated only by these small non-perturbative string effects.  The implicit, and very plausible assumption behind this line of reasoning is that there is no anthropic reason for the existence and properties of the QCD axion. Consequently, these properties should follow from the dynamics of the compactification manifold, rather than being a result of fine-tuning, and the QCD axion should be a typical representative among other axion-like fields. A priori we expect tens (or even hundreds) of light axions, it would be really surprising if the QCD axion turned out to be the single one. \\begin{figure}[t!] % \\begin{center} \\includegraphics[width=5in]{cloud.pdf} \\caption{{\\bf Axionic Black Hole Atom:} The spinning black hole ``feeds\" superradiant states forming an axion Bose-Einstein condensate. The resulting bosonic atom will emit gravitons through axion transitions between levels and annihilations and will emit axions  as a consequence of self-interactions in the axion field.} \\label{carnotcycle} \\end{center} \\end{figure}  These arguments motivate us to look not only for the QCD axion, but for axions in the entire mass range  $\\mu\\sim 10^{-9}\\div 10^{-21}$~eV, where they can affect stellar or galactic astrophysical black holes through superradiance. Let us summarize now the major features of superradiance and its principal observational consequences.  Superradiance \\cite{Zeld,Misner:1972kx,Starobinskii} is the phenomenon of wave amplification during scattering off a rotating black hole which takes place whenever the wave frequency $\\omega$ and the magnetic quantum number $m$ satisfy the superradiance condition \\be \\label{Omegacond} 0<\\omega<mw_+\\;, \\ee where $w_+$ is the angular velocity of the black hole horizon defined as \\be w_+ = \\frac{1}{2r_g} \\frac{{a}/{r_g}}{1+\\sqrt{1 - (a/r_g)^2}} \\label{spinvelo} \\ee with $r_g$ being the gravitational radius of the black hole and $0<a<r_g$ is the spin-to-mass ratio. Superradiant amplification may lead to an instability under certain circumstances. One example, admittedly not a very practical one, is the ``black hole bomb\"  by Press and Teukolsky \\cite{Press:1973zz,Press:Nature}: a rotating black hole surrounded by a spherical mirror. A single photon introduced in the system, or created by quantum fluctuations, with quantum numbers satisfying the superradiance condition (\\ref{Omegacond}) gives rise to an exponentially growing number of photons inside the mirror through a chain of consequent amplifications at the black hole horizon and reflections from the mirror.  Remarkably, nature provides such a mirror in the presence of a massive boson  \\cite{Damour:1976kh,Ternov:1978gq,Zouros:1979iw,Detweiler:1980uk}. Massive bosons, in our case axions, have bound Keplerian levels in the gravitational field of a black hole. This allows for a black hole to release its spin by populating levels satisfying the condition (\\ref{Omegacond}), see Fig.~\\ref{carnotcycle}. This creates an axionic Bose-Einstein condensate (BEC) cloud rotating around the black hole. The process is only efficient if the Compton wavelength of an axion is comparable to the black hole size. As a result, the energy spectrum of superradiant levels is quantized and very close to the spectrum of a hydrogen atom. The superradiant axion cloud loses its energy and momentum by gravitational wave emission associated with axion transitions between different  ``atomic\" levels, and with axion annihilations to gravitons. Another important loss mechanism is related to non-linearities in the axion potential and results in the emission of axions. Finally, the whole system may be fueled by energy and spin inflow from the matter accreting onto the black hole.  \\begin{figure}[t] % \\begin{center} \\includegraphics[width=6in,trim=70 50 50 100]{reach.pdf} \\caption{The part of the black hole and axion parameter space potentially affected by superradiance. For axion and black hole masses in the colored region the time required  to create a substantial axion cloud is shorter than the age of the Universe. For masses in a light colored region the superradiance rate is faster than the Eddington accretion rate. } \\label{aeq1plot} \\end{center} \\end{figure} In Fig. \\ref{aeq1plot} we show the region of black hole mass and axion mass parameter space that is affected by superradiance. The dark area outlines the region where a superradiant cloud has enough time to be built up during the lifetime of the Universe for a maximally spinning black hole. The lighter region shows the part of the parameter space where the superradiant spindown rate is faster than the spinup rate due to Eddington accretion, so that Eddington accreting black holes in this mass range lose their spin as the cloud develops.  There are three major classes of observational signatures associated with the axion cloud that we are going to discuss in the current paper: \\begin{itemize} \\item Gaps in the black hole ``Regge plot\" (spin vs mass plane)---the absence of rapidly rotating black holes when their gravitation radius matches the Compton wavelength of an axion. \\item Direct gravitational wave signals from the rotating axion cloud. For the QCD axion this signal falls into the sensitivity band of the Advanced LIGO interferometer. \\item Modification of the near-horizon metric due to the presence of the axion cloud. \\end{itemize} The main goal of this paper is to evaluate the observability of these classes of signatures in near future experiments. Detailed quantitative predictions for gravitational wave emission rates and waveforms, as well as the change of the templates for the near horizon metric is likely to require numerical work. The reason is that the axion BEC cloud is a rich and complicated dynamical system with many processes that need to be taken into account in order to have an accurate description of its behavior at cosmological time scales. One important process that poses a challenge for an analytical treatment is the effect of axion self-interactions on  superradiance. Here we limit ourself to a qualitative, semi-analytical discussion of the major physical processes involved. Our analysis indicates that depending on the relation between the axion mass and the black hole size all three classes of signatures can be observable for different systems. This  provides strong motivation for further numerical analysis of axionic superradiance.  \\begin{figure}[t!] % \\begin{center} \\includegraphics[width=5in,trim= 0 0 0 40]{spinplotqcd1.pdf} \\end{center} \\begin{center} \\includegraphics[width=5in,trim= 0 0 0 70]{spinplotgalactic.pdf} \\caption{The regions in the black hole Regge plane affected by  superradiance for the QCD axion with $m=3\\cdot 10^{-11}$~eV (the upper panel) and for a lighter axion with $m=10^{-17}$~eV. The data points correspond to spin measurements obtained by fitting the thermal continuum X-ray spectra \\cite{McClintock:2009dn}. Old black holes are expected to be found in the shaded region, where they are not affected by superradiance. Young black holes may be found also on the dashed colored lines inside the gap. Different colors correspond to superradiant levels with different orbital quantum numbers $l$. } \\label{fig_summary} \\end{center} \\end{figure}  Let us briefly summarize our main results. In Fig. \\ref{fig_summary} we show the black hole Regge plot for two different axion masses. The upper plot corresponds to an axion mass of $\\mu_a=3\\cdot 10^{-11}$ eV, which is the mass of the QCD axion for $f_a=2\\cdot 10^{17}$ GeV, and the lower illustrates the effect of a lighter axion with $\\mu_a=10^{-17}$ eV. Black holes in the shadowed region are not affected by superradiance during the age of the Universe. Sufficiently old black holes are expected to be found in this region. In the plot corresponding to the QCD axion we also present existing spin measurements; this data suggest an upper bound on the QCD axion decay constant of $2 \\cdot 10^{17}$ GeV.  Note that the uncolored gap regions exhibit internal structure (dashed lines), reflecting the quantized behavior of the gravitational black hole atom. Black holes may stay on these lines, the ``Regge trajectories\", for cosmological time scales. In order to understand this behavior, we should note that there is a number of different superradiant bound states labeled by different angular quantum number $l$. The instability time-scale rapidly drops down as the orbital quantum number $l$ increases. Therefore, the spin-down process of the black hole is initially driven by the level with the minimum value of $l$ for which the superradiance condition (\\ref{Omegacond}) is satisfied, and stops when enough spin has been extracted, so that the superradiance instability rate goes to zero --- the superradiance condition is now saturated. At this point the black hole finds itself on one of the dashed lines of the Regge plot.  Next, we should expect that the black hole spin-down proceeds with the growth of the level with the next-to-largest superradiance rate, i.e., the $l+1$ level. However, further spin-down is damped when the black hole reaches one of the Regge trajectories as a consequence of axion self-interactions. The axions bound to the level that is no longer superradiating serve as an axially asymmetric perturbation to the system that mixes superradiant with non-superradiant levels and shuts off the black hole spin-down process. This is similar to introducing a non-spherical defect on the mirror in the Press-Teukolsky black hole bomb: photons reflected from the mirror exit the superradiant region (\\ref{Omegacond}) and are now absorbed by the black hole, turning the instability off. As a result of these axion non-linearities, black holes stay on the Regge trajectories until various axion loss processes, or a violent accretion event, dissipate the cloud down to a small enough size, such that non-linearities no longer inhibit the instability. Then the black hole rapidly jumps to the next Regge trajectory by populating the level with a larger orbital quantum number.  Axion non-linearities trigger yet another dramatic effect during transitions between Regge trajectories. Namely, attractive axion self-interactions result in a catastrophic instability of the axion cloud and its subsequent collapse. The analogous effect has been observed in laboratory BEC systems with attractive self-interactions and is known under the name ``Bosenova\" \\cite{Bosenova}. We find that, depending on the parameters, transitions between Regge trajectories proceed through a series of tens to hundreds of Bosenova events. These events produce gravitational waves, and, in the case of the QCD axion, gamma- and X-rays that may be detectable on the Earth.  The black hole may also emit gravitational waves at the observable level due to axion transitions between different levels and annihilations to gravitons. Accurate prediction of the strength and duration of the signal requires a more detailed analysis of the dynamics; in the current paper we limit ourselfs to qualitative estimates which imply that this signal is detectable. It is especially exciting that the gravitational wave signal from the QCD axion cloud around stellar mass black holes falls into the Advanced LIGO frequency band and turns it  into a particle discovery machine.  In Fig.~\\ref{QCDplot}  we present the estimated signal strength as a function of the axion and black hole masses for a source at a 20~Mpc distance from the Earth. We choose a coherent integration time for the signal of $10^4$ second , when drawing the sensitivity curves. \\begin{figure}[t] % \\begin{center} \\includegraphics[width=6in]{gwtransitions2qcd.pdf} \\caption{The contour plot of constant gravitational wave signal from axion transitions between the $6g$ and the $5g$ levels for a black hole located at 20~Mpc away from the Earth. The projected sensitivity curves of Advanced LIGO \\cite{advLIGO} and Einstein Telescope \\cite{ET} assume $10^4$ seconds of a coherent integration time.} \\label{QCDplot} \\end{center} \\end{figure}  The paper is organized as follows. We start with explaining the basics of superradiance in Section~\\ref{spectroscopy}. We argue that the non-relativistic approximation is accurate for describing superradiant levels for most of the parameter space, and adequate at least qualitatively at all values of the parameters.   We present semi-analytical results for the superradiant rates, that are accurate over a large part of the parameter space and enough for our purposes.  In Section~\\ref{dynamics} we discuss the processes that determine the dynamics of the superradiant cloud---gravitational wave emission and axion non-linearities. We provide approximate expressions for the relevant transition rates, and for characteristic timescales and  masses of the superradiant cloud when different processes, such as Bosenova collapse, happen.  In Section~\\ref{signatures} we combine all of the above ingredients and discuss how the superradiant instability develops and what are the associated observational signatures. We discuss in detail the black hole spin-down and of the Regge trajectories. We then estimate the gravitational wave signals during transitions between Regge trajectories and briefly comment on prospects for directly detecting the axionic cloud around supermassive black holes during extreme mass ratio inspirals.  We also focus on the potential reach for the QCD axion, and some possible QCD axion specific signatures related to its direct coupling to Standard Model fields. A particularly intriguing effect of the QCD axion is that the effective value of the QCD $\\theta$-parameter may become of order one inside the cloud. This possibly destabilizes nuclei in the accretion disk and results in $\\gamma$-ray signals and exotic X-ray lines from the nuclear decay products in the black hole vicinity. We finally entertain the possibility of radio waves from axion-to-photon conversion in the near-horizon magnetic field.   Section~\\ref{anthropic} elaborates on an issue which is  aside from the main line of the paper, but still a very important part of the theoretical motivations for string axions in the mass range probed by the black hole superradiance. Namely, all these axions are ``anthropic\"---their initial misalignment angle needs to be tuned to an atypically small value in the observed part of the Universe. It has been known for a long time~\\cite{Linde:1987bx}, that this is not a problem for a single QCD axion, and here we discuss what changes if several anthropic axions are present. We conclude in Section~\\ref{fuuh}. ",
        "conclusions": "\\label{fuuh}  We hope to have convinced the reader that black hole superradiance for axions is an extremely rich phenomenon that has good chances to be observed in near future measurements of black hole properties. Ongoing black hole spin measurements may trace gaps or Regge trajectories in the spectrum of rapidly spinning black holes. Advanced LIGO may observe gravitational wave signals from the QCD axion cloud around stellar mass black holes as far as the Virgo cluster for masses down to $10^{-10}$ eV which correspond to an axion decay constant close to the grand unification scale. In a more distant future, gravitational waves may be observed for supermassive black holes at lower frequencies by experiments such as LISA and AGIS. The low frequency gravitational wave detectors may also see the effect of the cloud on the waveforms during extreme mass ratio inspirals. Finally, for the QCD axion the superradiant cloud might also give rise to direct photon signals.    In this paper our main goal has been to develop a general intuition about superradiance development and its consequences without going into an extensive numerical work. Given the richness of the system it seems inevitable that detailed numerical simulations will be required in the future to obtain accurate quantitative predictions. This is especially important for predicting the strength, duration and precise waveform of gravitational wave signals from superradiant clouds. Simulations are needed both for the accurate prediction of superradiance rates, which to large extent has already been accomplished in \\cite{Dolan:2007mj}, and most importantly to get an accurate description of the superradiant cloud including axion self-interactions. It's worth stressing, however, that some of our result are very robust. For instance, the black hole Regge trajectories of Fig.~\\ref{fig_summary} are mostly determined by the basic superradiance condition (\\ref{Omegacond}) and do not depend on the above uncertainties (apart from the left most declining segments of these curves; however, in that region even larger uncertainties are likely to come from variations in the accretion rate for different black holes).  Of course,  we expect also other qualitative results obtained here to reproduce well the gross features of the system, although given its richness more surprises are possible. Most importantly, our results appear encouraging for prospects to observe superradiance with future astrophysical data, and this justifies further theoretical efforts for better understanding of this fascinating process."
    },
    "1004/1004.1288_arXiv.txt": {
        "abstract": "It is becoming ever clearer that the neutrino signal from the next supernova in our Galaxy can reveal missing information about the neutrino as well as allowing us to probe the explosion of the star by decoding the temporal and spectral evolution of the flavor composition of the signal. But this information may be lost if turbulence in the supernova `depolarizes' the neutrinos so that the observed flux for each flavor is an equal mixture of the initial - unencoded - spectra. Determining if depolarization occurs is one of the most pressing issues of this field. The most difficult aspect of studying the effect of turbulence upon the neutrinos is the lack of any theoretical models that allow us to understand the results of numerical studies. This paper makes the suggestion that Random Matrix Theory (RMT) may shine some light in this direction and presents support for this the possibility by comparing the distribution of crossing and survival probabilities obtained numerically for some `test case' calculations with the distributions one expects from RMT in the calculable limit of depolarization of N neutrino flavors. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.0698_arXiv.txt": {
        "abstract": "{ The solar chromosphere is a vigorously dynamic region of the sun, where waves and magnetic fields play an important role. To improve chromospheric diagnostics, we present new observations in \\ion{Ca}{II}~8542 carried out with the SST/CRISP on La Palma, working in full-Stokes mode. We measured Stokes line profiles in active regions. The line profiles observed close to the solar limb show signals in all four Stokes parameters, while profiles observed close to disk center only show signals above the noise level in Stokes I and V. We used the NLTE inversion code 'NICOLE' to derive atmospheric parameters in umbral flashes present in a small round sunspot without penumbra. ",
        "introduction": "The magnetic structure of the solar chromosphere is currently the subject of intense computational  and observational research. Numerical MHD  simulations have  improved in the last few years, becoming more and more realistic, i.e, \\cite{2004A&A...414.1121W}, \\cite{2007A&A...473..625L} and \\cite{2009ApJ...701.1569M}. Unfortunately, progress on the observational side has arguably been somewhat slower. The few optical spectral lines probing the chromosphere are typically characterized by: \\begin{enumerate} \\item A vast formation range, from the photosphere to the chromosphere. \\item NLTE. \\item A weak polarization signal (especially in Q and U). \\item A large spectral width. \\end{enumerate}  Inversion codes have been commonly used to derive atmospheric parameters, usually assuming  LTE or considering a Milne-Eddington atmosphere. However, this approximation cannot be used  for modeling chromospheric lines, such as the Ca II infrared triplet. A method for NLTE inversions of Zeeman-induced Stokes profiles was described by \\citet{2000ApJ...530..977S} and used in a study by \\citet{2007ApJ...670..885P} of slit-based spectroscopic observations. The development of new instrumentation at large aperture telescopes (like SST/CRISP, DST/IBIS, VTT/TESOS) has made it possible to observe the chromosphere with very high spatial and temporal resolution. Here, we present inversions of chromospheric full-Stokes CRISP data. ",
        "conclusions": "We have reduced polarimetric CRISP observations of the \\ion{Ca}{II}~854.2~nm line and analyzed a small, round pore located at $\\mu=0.81$ using an adapted version of the NLTE inversion code NICOLE. The code is not able to reproduce discontinuities  as the model is smoothed with a spline fit during the inversion. Nevertheless, the shock leaves a clear fingerprint in the derived velocity, density and magnetic field strength. Since the observed Stokes V profile of this line is known to be properly described by Zeeman splitting in such a strong magnetic field, the spread in the magnetic field strength is probably mostly caused by noise. The amount of cross-talk between other thermodynamic parameters  is still unknown and will be the subject of future research."
    },
    "1004/1004.2979_arXiv.txt": {
        "abstract": "We study Kelvin-Helmholtz (KH) instability at the interface of a disc and corona system by doing a linear perturbation analysis. The disc is assumed to be thin, however, the corona is considered to be nearly quasispherical because of its high temperature.  Under these circumstances, the interface is subject to the KH instability for a given set of the input parameters. Growth rates of the KH unstable modes are calculated for a wide range of the input parameters. We show that for a certain range of the perturbations, the unstable KH perturbations are growing with time scales comparable to the inverse of the angular velocity of the accretion disc (dynamical time scale). Thus, KH instability at the interface of a disc-corona may have enough time to affect the dynamical structure of its underlying accretion disc by possible exchange of the mass, angular momentum or even energy.  Our linear analysis shows that KH instability may provide a mechanism for such exchanges between a disc and its corona. ",
        "introduction": "The standard theory of the accretion discs \\citep*{Shakura1973} has been used widely over recent three decades. Assuming the disc is radiating like a black body object, we can explain the thermal emission of some of the accretion discs using the standard theory. However,  non-thermal emission is also detected from some of the accretion discs, namely discs around compact objects or even active galactic nuclei (AGN). One can hardly explain the non-thermal emission based on the standard theory of the accretion discs and for this reason non-standard models are proposed \\citep*[e.g.,][]{Kaw2001}. In fact, X-ray of many accreting binaries and AGN imply a hot region around the system and these considerations have motivated some authors to propose a two-component model for these accreting systems, in which a thin  disc is surrounded by a hot and tenuous corona \\citep*[e.g.,][]{Haardt91,Nak1993,Kaw2001,Merloni2003}. On the other hand, it is now commonly accepted that magneto-rotational instability (MRI) is the main driving mechanism of the turbulence in the accretion discs \\citep*{Balbus1991}. Numerical simulations of MRI show the magnetic field lines may come out of the disc  because of the buoyancy force and generate the magnetic loops outside the disc and dissipate the energy, mainly through the  reconnection mechanism \\citep*[e.g.,][]{Miller2000,Liu2002}. Then, the region around the disc is heated significantly because of the magnetic energy dissipation of the loops \\citep*[e.g.,][]{Heyv1989,Field1993,Black2000,Kun2004}.  However, there is also another scenario for the formation of a corona around an accretion disc. Some authors believe that the coronae are created by  the radiative pressure driven winds, which produce a hot plasma above an accretion disc \\citep*[e.g.,][]{ohsuga2005}.    Since the first phenomenological model of \\citet{Haardt91}, the model has been extended in various directions. However, it is still uncertain if the coronae are formed because of the magnetic field or radiation driven winds or even a combination of both mechanisms.  Whatever is the formation mechanism of a corona, the existence of a hot gaseous system above the disc will modify the structure of the underlying accretion disc because of the mass, energy and even angular momentum exchange between the disc and corona. Although the true nature of these exchanges is yet to be understood, one can consider physical arguments to present models for disc-corona systems including possible exchanges. The basic work about the mass exchange between the disc and corona was given by \\citet{meyer94}. \\citet{czer2000} studied a hot advective corona in pressure equilibrium with the cold disc. In their model, the corona exchanges energy with the disc through the radiative interaction and conduction. In another study, \\citet{krolik} presented a model  for the structure of a hot corona by considering the possible effects of thermal conduction. Most of the semi-analytical models for disc-corona systems are following a similar approach \\citep[e.g.,][]{czr1999,roz1999,Liu2002,hof2006,Mayer2007}.  Recently, \\citet{Mayer2007} presented time-dependent simulations of a two-phase accretion flow around a black hole. They also considered mass and energy exchanges between disc and corona.   In most of the previous studies, for simplicity, it is assumed the corona is rotating with a Keplerian rotational velocity like the disc. Also, the geometrical shape of the corona is not investigated in detail. However, the temperature of the corona is much larger than the disc because the observed high energy cut-off in X-ray spectra usually is about 100 keV. Moreover,  the corona is not thin comparing to its cold disc and the gradient of the pressure in the corona can {\\it not} be neglected. Therefore, in our analysis, a corona tends to have  a quasi-spherical shape with a non-Keplerian rotational velocity because of the gradient of pressure. However, in the underlying cold disc the gradient of pressure can be neglected, and it  rotates with Keplerian velocity. In other words, not only there is a significant density contrast between the disc and corona, a velocity drift should exist between the disc and corona.  When two fluids are in relative motion on either side of a common boundary, the Kelvin-Helmholtz (KH) instability can occur. Thus, we lead to a simple question: Is the interface of a disc and corona system unstable because of KH instability? Our goal is to address this question in this paper. KH instability has already been studied in different astrophysical systems from  the interactions between the solar wind and planetary magnetospheres which are studied based on the KH instability \\citep[e.g.,][]{nagano79} to jets and outflows \\citep[e.g.,][]{hardee,rossi} and dusty layer in protoplanetary discs  \\citep[e.g.,][]{michikoshi}. Possible instabilities at the interface between an accretion disc and a magnetosphere surrounding the accreting mass have been studied by many authors \\citep[e.g.,][]{narayan2004}. It  was shown the magnetospheric interface is generally Rayleigh-Taylor unstable and may also be KH unstable \\citep{narayan2004}. But the KH instability at the interface between a disc and its corona has not been studied as far as we know. In the next section, general equations of our model are presented. Analysis of the linear perturbation analysis is done in section 3. A summary of the results  are discussed in section 4. ",
        "conclusions": "We have shown that the interface between a disc and its corona is unstable subject to the KH instability and the growth rate of the unstable modes was calculated. Different time scales are defined for an accretion disc like dynamical and viscous time scales. In most of the analytical models of the accretion discs, dynamical time scale is less than viscous time scale.  For some of the input parameters and a range of the wavelength of the perturbations, the time scale of growing KH unstable modes is comparable to the inverse of the angular velocity of the disc (i.e., dynamical time scale). It means that KH instability has enough time to affect the structure of the disc.   Growth rate of the perturbations versus wavenumbers shows a very complicated profile depending on the input parameters. However, when the wavelength of the perturbations is around $\\lambda = 2\\pi H_{d} (\\rho_{c} / \\rho_{d} )$ the time scale of growth is of the order of the dynamical time scale.  It is discussed that the corona may form because of a large scale magnetic field \\citep*[e.g.,][]{black2009}. The exact amount of the energy transfer from the disc to the corona by the magnetic flux tubes is yet to be understood theoretically, though there are some estimates based on the physical considerations \\citep*[][]{Merloni2003,Kun2004,black2009}. This possible exchange of the energy not only modifies the structure of the corona, but may have significant effect on the structure of the disc as well \\citep*{khaj}.  Our analysis proposes another mechanism to be considered for the exchange of the energy or mass between the disc and corona.  Evidently, our analysis does not allow to estimate possible momentum and energy exchanges between the disc and corona because of KH instability. Further numerical simulations are needed to address this issue."
    },
    "1004/1004.2238_arXiv.txt": {
        "abstract": "An update of \\citet{owe08} shows that the relationship between the coronal mass ejection (CME) rate and the heliospheric magnetic field strength predicts a field floor of less than 4 nT at 1 AU. This implies that the record low values measured during this solar minimum do not necessarily contradict the idea that open flux is conserved. The results are consistent with the hypothesis that CMEs add flux to the heliosphere and interchange reconnection between open flux and closed CME loops subtracts flux. An existing model embracing this hypothesis, however, overestimates flux during the current minimum, even though the CME rate has been low. The discrepancy calls for reasonable changes in model assumptions. ",
        "introduction": "Using an analytical, empirically-based model, \\citet{owe06,owe07} demonstrated the feasibility of the idea that the rise and fall of magnetic field strength in the heliosphere over the solar cycle could be caused solely by the addition of closed magnetic loops from coronal mass ejections (CMEs) coupled with a loss of flux through reconnection at the Sun. Either of two patterns of reconnection can satisfy the constraints posed by observations of suprathermal electrons, which can sense magnetic topology: 1) CME loops gradually open by interchange reconnection with open field lines, or 2) open field lines of opposite polarity reconnect, thus disconnecting flux from the Sun.  The two options for flux loss through reconnection have different advantages. Option 1 is attractive because it restricts the rise and fall of field magnitude to a single phenomenon, CMEs, and it conserves open flux, which serves as a backdrop. On the other hand, because option 2 calls upon a phenomenon separate from CMEs for the loss of flux, that is, disconnection elsewhere, it has the flexibility of reducing the heliospheric flux to zero at solar minimum rather than to some backdrop value. In view of the weakness of the heliospheric magnetic field strength during the present solar minimum, unprecedented in the $\\sim$50-year history of in situ measurements, we ask whether option 1 is still viable.  A preliminary answer of ``yes'' to the question of option 1 viability was provided by \\citet{owe08} using data through 5 July 2008. Since the heliospheric field strength continued to drop after that date, this report updates the results of \\citet{owe08} using data through 30 May 2009. ",
        "conclusions": "The unprecedented low magnetic field strength in the heliosphere during the current solar minimum does not eliminate the possibility that open flux is conserved. Nor does it render invalid the hypothesis that the heliospheric field strength depends solely upon the rate of injection of loops from CMEs and the rate those loops open by interchange reconnection. According to this view, the field floor reflecting the background open flux is never reached during solar minima owing to a residual influx of CME loops. As modeled by \\citet{owe06}, however, the extended lower CME rate during the current minimum cannot fully account for the depressed field strength. Assumptions of constant flux per CME and/or rate of interchange reconnection must be relaxed, as seems reasonable owing to apparent changes in CME properties during the current minimum."
    },
    "1004/1004.4905_arXiv.txt": {
        "abstract": "We propose a scale dependent analytic approximation to the exact linear growth of density perturbations in Scalar-Tensor (ST) cosmologies. In particular, we show that on large subhorizon scales, in the Newtonian gauge, the usual scale independent  subhorizon growth equation does not describe the growth of perturbations accurately, as a result of scale-dependent relativistic corrections to the Poisson equation. A comparison with exact linear numerical analysis indicates that our approximation is a significant improvement over the standard subhorizon scale independent result on large subhorizon scales. A comparison with the corresponding results in the Synchronous gauge demonstrates the validity and consistency of our analysis. ",
        "introduction": "Cosmological data from a wide range of sources including type Ia supernovae \\cite{union08, perivol, hicken}, the cosmic microwave background \\cite{Komatsu}, baryon acoustic oscillations \\cite{bao,percival}, cluster gas fractions \\cite{Samushia2007,Ettori} and gamma ray bursts \\cite{Wang,Samushia2009} seem to indicate that at least 70\\% of the energy density in the universe is in the form of an exotic, negative-pressure component, called dark energy. While the standard $\\Lambda$CDM framework is the minimal model that successfully accounts for observations \\cite{Komatsu}, there remain numerous viable alternatives that also pass  current experimental tests.  These alternative models can be broadly categorized as quintessence \\cite{RatraPeebles,Caldwell:1997ii,SteinhardtWangZlatev,WangSteinhardt:1998,dutta} or modified gravity models \\cite{Parker:1999td,Perivolaropoulos:2005,CaldwellCoorayMelchiorri,JainZhang,NesserisPerivolaropoulos:2006,WangHuiMayHaiman,Tsujikawa,HeavensKitchingVerde,EspositoFarese:2000ij,Dvali:2000hr,Freese:2002sq,Carroll:2004de,OCallaghan:2009bu, HL, FT}, and both categories can, with sufficient tuning, replicate the expansion history of the universe in consistence with observations. (See \\cite{Copeland, HutererTurner} for recent reviews). In order to distinguish between these two categories of models, it is therefore important to look beyond the expansion rate.  The growth of structure offers a hope of breaking this degeneracy since different growth histories can arise from models which have similar expansion histories.  To examine the growth of structure one examines the evolution of the linear matter density contrast $\\delta \\equiv \\delta\\rho/\\rho$ which is given in terms of the background density $\\rho$ and the perturbation $\\delta\\rho$. For scales much smaller than the horizon, $\\delta$ satisfies a simple equation called the growth equation, which is scale-independent. Note that in what follows we use the usual definition of 'scale' as either the distance $\\lambda_p$ in physical FRW coordinates or the corresponding wavenumber $k=\\frac{2\\pi}{\\lambda_p}$. There are many investigations which attempt to characterize the evolution of $\\delta$ through the use of a growth parameterization which assumes a different value depending on the cosmological model used, thereby allowing for models to be distinguished (e.g. \\cite{WangSteinhardt:1998,LinderCahn:2007,Gong:2009sp}).  The standard definition of the growth parameter, $\\gamma$, in terms of the growth function $f$, the matter density $\\Omega_m$ and the scale factor $a$ is given as \\be \\label{stdpar} f(a)\\equiv \\frac{d \\textrm{ln} \\delta}{d\\textrm{ln}a} \\equiv \\Omega_{m}(a)^{\\gamma}\\,. \\ee Once this parameter is determined (for earlier theoretical developments on the parameterization of the growth parameter and experimental constraints on $\\gamma$, see \\cite{dgp,LinderCahn:2007,Uzan:2006mf,Polarski:2007rr,Nesseris:2006er,Hawkins:2002sg,Viel:2004bf,Viel:2005ha,Kaiser:1996tp,Seo:2003pu,Mantz:2007qh,Pogosian:2005ez,Nesseris:2007pa}) one may then be in a position to determine whether the standard general relativistic (GR) framework of $\\Lambda$CDM is responsible for the acceleration of the universe, or some other, more exotic process is at work.  However, in \\cite{DentDutta:2008,DentDuttaPerivolaropoulos:2009}, it was demonstrated that in the Newtonian gauge (for another interesting look at gauge issues see \\cite{Wenbinlin:2001}) the above parameterization can become inaccurate for large subhorizon scales $\\gtrsim 100h^{-1}$Mpc. In other words, if the physical growth of structure is correctly described by the Newtonian gauge, then it would show up as inconsistent with scale-independent parameterization in \\eqn{stdpar} and (mistakenly) appear to be caused by exotic physics. The reason for the discrepancy was shown to be the scale dependence of the growth of $\\delta$, which becomes important for large subhorizon scales ($\\gtrsim 100h^{-1}$Mpc). An improved version of the growth equation was derived in \\cite{DentDutta:2008} which incorporates the scale-dependence.  In \\cite{DentDuttaPerivolaropoulos:2009}, a new scale-dependent parametrization of the growth function $f$ was proposed,  which was shown to account for the evolution of $f(a)$ on these large scales with considerably greater accuracy than \\eqn{stdpar}.  In the present work, we focus on the growth of perturbations in scalar-tensor (ST) theories of gravity. It is well known that in the sub-Hubble approximation the growth of perturbations in these theories is also described by an equation similar to the growth equation in GR, up to a redefinition of the gravitational constant. Working in the Newtonian gauge, we show that the usual growth equation approximation for $\\delta$ becomes unreliable on large scales, for the same reason as in the GR case, i.e. the effects of scale dependence. We derive an improved version of the growth equation relevant for these models and propose a more accurate scale-dependent parameterization for growth.  The layout of our paper is as follows. In Section II we discuss the growth of perturbations in ST theories, demonstrating the failure of the usual growth equation approximation and introducing an improved growth equation and a new parameterization for the growth function in these models. In Section III we compare these approximations to exact solutions to demonstrate their accuracy. Our conclusions can be found in Section IV. ",
        "conclusions": "\\label{conclusions} We have derived and tested numerically a simplified ordinary differential equation whose solution describes fairly accurately the growth of linear cosmological perturbations in ST cosmologies up to Hubble scales (beyond $10^3 h^{-1}$Mpc). A corresponding analytic form of the scale dependent growth rate $f(a,k)$ was also presented and tested numerically in comparison with the exact linear result. This is a significant improvement over the previously known sub-Hubble which breaks down on scales larger than about $10^2 h^{-1}$Mpc.  We have thus demonstrated that in the context of the Newtonian gauge, the comparison of cosmological large scale structure data with corresponding theoretical predictions in ST based cosmologies on scales larger than a few hundred Mpc should not be based on the sub-Hubble scale independent approximation. Instead it requires either use of the full linear numerical solution of the cosmological perturbation equations in each theory or the imposed scale dependent approximation presented in the present study. Thus, previous work using the scale independent sub-Hubble approximation to make predictions for the growth of perturbations in ST theories is reliable on scales up about $200h^{-1}Mpc$. However, such results should be used with care on larger scales. For example, on scales of a few hundred $Mpc$ the error induced in the growth rate by not using the scale dependent effects is about $10\\%$ on redshifts $z>3$ while on scales larger than $1Gpc$ the corresponding error exceeds $50\\%$ even at low redshifts $z\\simeq 1$.  We have also shown that our results are consistent with a corresponding calculation in the synchronous gauge and verified that the quantity $\\frac{\\delta \\rho}{\\rho}$ is gauge dependent and there can be a significant difference between its forms in different gauges on large scales."
    },
    "1004/1004.3839_arXiv.txt": {
        "abstract": "We present our recently developed {\\em galcon} approach to hydrodynamical cosmological simulations of galaxy clusters - a subgrid model added to the {\\em Enzo} adaptive mesh refinement code - which is capable of tracking galaxies within the cluster potential and following the feedback of their main baryonic processes. Galcons are physically extended galactic constructs within which baryonic processes are modeled analytically. By identifying galaxy halos and initializing galcons at high redshift ($z \\sim 3$, well before most clusters virialize), we are able to follow the evolution of star formation, galactic winds, and ram-pressure stripping of interstellar media, along with their associated mass, metals and energy feedback into intracluster (IC) gas, which are deposited through a well-resolved spherical interface layer. Our approach is fully described and all results from initial simulations with the enhanced {\\em Enzo-Galcon} code are presented. With a galactic star formation rate derived from the observed cosmic star formation density, our galcon simulation better reproduces the observed properties of IC gas, including the density, temperature, metallicity, and entropy profiles. By following the impact of a large number of galaxies on IC gas we explicitly demonstrate the advantages of this approach in producing a lower stellar fraction, a larger gas core radius, an isothermal temperature profile in the central cluster region, and a flatter metallicity gradient than in a standard simulation. ",
        "introduction": "\\label{sec:intro}  Hydrodynamical simulations of galaxy clusters, incorporating semi-analytic models for star formation and galactic feedback processes, show an appreciable level of inconsistency with observational results. This is particularly apparent in the simulated properties of intracluster (IC) gas - temperature, metallicity and entropy profiles - and the stellar mass density at high redshift (see the review by Borgani et al. 2008, and references therein). Statistical properties of clusters, such as X-ray luminosity-temperature, entropy-temperature, and mass-temperature relations, seem also to be discrepant when compared with high-precision optical and X-ray observations (e.g., Kay et al. 2007, Nagai et al. 2007, Tornatore et al. 2007, Kapferer et al. 2007; for a recent review, see Borgani et al. 2008). The mismatch between simulated and observed cluster properties (e.g., Evrard \\& Henry 1991, Cavaliere, Menci \\& Tozzi 1998, Tozzi \\& Norman 2001) is largely due to insufficient accounting for essential physical processes, unrealistic simplifications of the evolution of star formation and feedback processes, and inadequate level of spatial resolution.  Some of the relevant physical processes that affect cluster properties and the dynamical and thermal state of IC gas are mergers of subclusters, galactic winds, ram-pressure stripping and gravitational drag. These have been partly implemented (e.g., Kapferer et al. 2005, Domainko et al. 2005, Bruggen \\& Ruszkowski 2005, Sijacki \\& Springel 2006, Kapferer et al. 2007) with some success in predicting IC gas properties. Different combinations of these processes and the various ways they are included in simulation codes generally result in quite different gas properties. An example is star formation (SF), whose self-consistent modeling (in cluster simulations) requires a prohibitively high level of spatial resolution which cannot be achieved with the current computing resources. Because of this limitation, most current simulations use a SF prescription that follows the formation of collisionless star `particles' in a running simulation (e.g., Cen \\& Ostriker 1992; Nagai \\& Kravtsov 2005), an approach which leads to an overestimation of the evolution of the SF rate (SFR) (Nagamine et al. 2004), and a higher than expected stellar to gas mass ratio. In addition, as we will show, in this particular implementation of SF the impact of the process remains spatially localized, resulting much lower mass (including metals) and energy ejection out of cluster galaxies, and consequently insufficient suppression of cooling and gas overdensity in cluster cores. This difficulty reflects the complexity of structure and SF processes, underlying the fact that a full implementation of these processes in hydrodynamical simulations is indeed a challenging task that nonetheless motivates attempts to develop a new approach in cluster simulations.  Considerations of galaxy clustering and star formation episodes at high redshift and the inclusion of heating to suppress gas cooling and condensation, lead us to identify Lyman Break Galaxies (LBGs) at $z \\geq 3$ as early (`pre-heating') sources of IC gas. Implementation of longer episodes of SF in these galaxies induces stronger winds. As a result, the amount of energy and metal-rich gas ejected to IC space is higher, as required for consistency with observations. Gas dispersal is further enhanced by ram-pressure stripping. Incorporating these baryonic processes motivated us to develop a new approach in the description of the evolution of IC gas, one that is based on the powerful adaptive mesh refinement (AMR) cosmological hydrodynamical simulation code - {\\em Enzo} (Bryan \\& Norman 1997, O'Shea et al. 2004). We have modified and improved {\\em Enzo} such that it is capable of following more realistically the hierarchical formation of structure through the inclusion of the most essential physical phenomena. This is accomplished by modeling the baryonic contents of galactic halos at high redshift by an extended `galaxy construct', which we refer to as {\\it galcon}. The new {\\em Enzo-Galcon} code does not require additional computational resources compared to the original {\\em Enzo} code. Because SF and feedback are modeled analytically, the level of resolution required to achieve improved results (compared to the standard simulation) is not extreme. This is the first time that most of the known processes are included in a hydrodynamical non-adiabatic simulation which is also capable of achieving high resolution ($\\leq$ 10 kpc). Initial results from the first implementation of our galcon approach were briefly described by Arieli, Rephaeli \\& Norman (2008, hereafter ARN).  In this paper a more complete description is given of our galcon approach, and an expanded analysis of the first simulations with the new code, including a wider range of IC gas properties than presented in ARN. In Section 2 we quantitatively describe the main baryonic processes included in our code. The galcon approach is introduced in Section 3, and results from the first {\\em Enzo-Galcon} simulations are presented in Section 4, with a detailed comparison with the corresponding results from a (`standard') {\\em Enzo} simulation using popular SF and feedback prescription. We end with a summary in Section 5. ",
        "conclusions": ""
    },
    "1004/1004.4072_arXiv.txt": {
        "abstract": "{Linearly polarized Galactic synchrotron emission provides valuable information about the properties of the Galactic magnetic field and the interstellar magneto-ionic medium, when Faraday rotation along the line of sight is properly taken into account. } {We aim to survey the Galactic plane at $\\lambda$6\\ cm including linear polarization. At such a short wavelength Faraday rotation effects are in general small and the Galactic magnetic field properties can be probed to larger distances than at long wavelengths. } {The Urumqi 25-m telescope is used for a sensitive $\\lambda$6\\ cm survey in total and polarized intensities. WMAP K-band (22.8~GHz) polarization data are used to restore the absolute zero-level of the Urumqi $U$ and  $Q$ maps by extrapolation. } {Total intensity and polarization maps are presented for a Galactic plane region of $129\\degr \\leq \\ell \\leq 230\\degr$ and $|b| \\leq 5\\degr$ in the anti-centre with an angular resolution of $9\\farcm5$ and an average sensitivity of 0.6~mK and 0.4~mK $\\rm T_{B}$ in total and polarized intensity, respectively. We briefly discuss the properties of some extended Faraday Screens detected in the $\\lambda$6\\ cm polarization maps. } {The Sino-German $\\lambda$6\\ cm polarization survey provides new information about the properties of the magnetic ISM. The survey also adds valuable information for discrete Galactic objects and is in particular suited to detect extended Faraday Screens with large rotation measures hosting strong regular magnetic fields.}  \\keywords {Polarization -- Surveys -- Galaxy: disk -- ISM: magnetic fields -- Radio continuum: general -- Methods: observational}  \\titlerunning{A Sino-German $\\lambda$6\\ cm polarization survey of the Galactic plane II.} \\authorrunning{X. Y. Gao et al.} ",
        "introduction": "Surveys of the Galactic plane at several frequencies are required to disentangle the individual star formation complexes, or thermal \\ion{H}{II} regions, non-thermal supernova remnants (SNRs) and extragalactic sources.  The diffuse emission associated with the Galactic disk is produced by relativistic electrons spiraling in magnetic fields and by thin ionized thermal gas. Both the diffuse non-thermal emission and the SNRs have significant linear polarization. Mapping of the Galactic plane at several radio frequencies including linear polarization offers a method to separate these non-thermal components as well as allowing a delineation of the Galactic magnetic field.  The Galactic plane has been surveyed from 22~MHz up to 10~GHz, albeit usually without polarization measurements. Sensitive Galactic polarization plane surveys began in the 1980s. A 2.7~GHz survey using the Effelsberg 100-m telescope by \\citet{Junkes87} showed a section of the Galactic plane with $4\\farcm3$ angular resolution. Further Northern sky Galactic plane surveys at 2.7~GHz \\citep{Reich9011,Fuerst90,Duncan99} were complemented by 2.4~GHz Southern Galactic plane surveys using the Parkes 64-m telescope \\citep{Duncan95,Duncan97}. To achieve angular resolution of arc minutes at  lower frequencies synthesis radio telescopes had to be used for surveys: e.g. the Westerbork Synthesis Radio Telescope at 350~MHz \\citep{Wieringa93,Haverkorn031,Haverkorn032}, the Dominion Radio Astrophysical Observatory synthesis telescope at 408~MHz and 1.4~GHz (Canadian Galactic Plane Survey, CGPS) \\citep{Taylor03}, and the Australian Telescope Compact Array at 1.4~GHz (Southern Galactic Plane Survey, SGPS) \\citep{Gaensler01,Haverkorn06}. Most of the mentioned surveys only cover a narrow strip along the Galactic plane. To overcome this deficiency the Galactic plane was mapped at 1.4~GHz with the Effelsberg 100-m telescope for $|b| \\leq 20\\degr$. First maps from this survey were shown by \\citet{Uyaniker99} and by \\citet{Reich04}. To study the nature of sources and the properties of the magnetic field, polarization surveys at higher radio frequencies are needed. Valuable information about diffuse polarized Galactic emission was provided by WMAP at 22.8~GHz and higher frequencies \\citep{Hinshaw09}, although the angular resolution of $50\\arcmin$ at 22.8~GHz is in general too coarse to resolve the complex Galactic structures in the Galactic plane.  The Sino-German $\\lambda$6\\ cm survey, covering a 10$\\degr$ wide strip of the Galactic plane, has been carried out since 2004 using the 25-m radio telescope of the Urumqi Observatory, National Astronomical Observatories, CAS. This survey fills the existing gap in frequency coverage by providing maps of the Galactic plane from $10\\degr \\leq \\ell \\leq 230\\degr$ and $|b| \\leq 5\\degr$ with an angular resolution of $9\\farcm5$. The survey maps and a list of compact sources will be released after completion of the $\\lambda$6\\ cm survey project expected for the end of 2010. The first results have already been presented by \\citet{Sun07} (hereafter called Paper~I), including details of the survey concept, the observing and calibration methods and the reduction process. In Paper~I, covering the longitude range from $122\\degr$ to $129\\degr$, we illustrated the scientific potential provided by the $\\lambda$6\\ cm survey by delineating new faint \\ion{H}{II} regions, studied spectra of SNRs, discovered Faraday Screens as well as traced the magnetic fields in this section of the Galactic plane. Most remarkable discoveries are two extended Faraday Screens located at the Perseus arm. One of them is caused by a previously unknown faint \\ion{H}{II} region. Both Faraday Screens host strong regular magnetic fields with rotation measures ($RM$) of the order of 200~rad\\ m$^{-2}$. They are not visible at low frequencies because such high $RM$s cause a polarization angle rotation by more than $180\\degr$, or they are beyond the polarization horizon. This proves the value of a sensitive $\\lambda$6\\ cm polarization survey to detect them in the magnetized interstellar medium. The commonly adopted picture of the Galactic magnetic field in the thin disk to consist of a regular component following basically the spiral arms of the Galaxy together with a turbulent magnetic field component of about similar strength might be modified in case numerous extended Faraday Screens with a uniform regular magnetic field exist. The origin of such magnetic bubbles acting as Faraday Screens is not clear so far.  Here we present the second section of the $\\lambda$6\\ cm survey for the outer Galaxy covering the region $129\\degr \\leq \\ell \\leq 230\\degr$. In Sect.~2 observation and data processing details for this survey area are discussed. In Sect.~3 we present the total power and polarization maps (Sect.~3.1), followed by a brief discussion on the survey's potential to study and detect SNRs (Sect.~3.2) and  \\ion{H}{II} regions (Sect.~3.3), while in Sect.~3.4 we focus  on newly detected and prominent Faraday Screens in the interstellar medium. Results are summarized in Sect.~4. ",
        "conclusions": "In Paper~II, we present the second section covering the outer Galaxy for the area $129\\degr \\leq \\ell \\leq 230\\degr$, $|b| \\leq 5\\degr$ of the Sino-German $\\lambda$6\\ cm polarization survey of the Galactic plane at an angular resolution of $9\\farcm5$. It is the ground-based polarization survey at the highest frequency for the Galactic anti-centre region. The observed polarization data have been restored to an absolute level by adding extrapolated large-scale components from the WMAP K-band polarization maps \\citep{Hinshaw09}.  Numerous newly detected Faraday Screens indicate the presence of large magnetic bubbles in the ISM hosting regular magnetic fields of a few $\\mu$G. A simple model fit to selected Faraday Screens, which also includes \\ion{H}{II} regions, was used to estimate their physical parameters. Our main results are:  \\begin{enumerate}  \\item We note that the remarkable polarized ``lens'' Faraday Screen in front of W5 detected by \\citet{Gray98} at $\\lambda$21\\ cm becomes invisible at $\\lambda$6 \\ cm, while previously unknown polarized structures were detected at $\\lambda$6\\ cm at the boundaries of W5.  \\item The Faraday Screen model fits for LBN~676, LBN~677 and LBN~679 indicate that besides the established Perseus arm objects LBN~676 and LBN~679 also LBN~677 is located in the Perseus arm rather than at 0.8~kpc distance, because its polarized foreground level is as that of the other two objects. The parameters of the three LBNe are listed in Table~3. For LBN~676 the model fit indicates a magnetic field direction opposite to those of the other two LBNe. A similar case was noted by \\citet{Mitra03} for a few Perseus arm \\ion{H}{II} regions based on pulsar $RM$s shining through them.  \\item The newly discovered extended Faraday Screen G146.4-3.0 is likely quite local. $B_{\\parallel}$ is estimated to be about $-8.6~\\mu$G if located at 690~pc, which is most likely an upper limit. The field strength within G146.4-3.0 will increase in case its distance is smaller.  \\item The two huge polarized bubbles located at $\\ell = 165\\degr$ as revealed by \\citet{Kothes04} at $\\lambda$21\\ cm become very faint at $\\lambda$6\\ cm.  \\item An extended blob showing excessive polarized emission is detected in the lower area of the ``bow-tie'' shaped \\ion{H}{II} region complex around $\\ell = 173\\degr$. Absorption at lower radio frequencies coincides with the $PI$ excess. We find evidence that this is a local Faraday Screen with a likely distance smaller than 300~pc.  \\end{enumerate}  For most of the selected Faraday Screens the polarized emission becomes weaker compared to their surroundings. This is expected when the $PA$ of the background polarization is rotated away from the foreground direction. The polarized emission exceeds that of the surroundings only in case the difference of the $PA$s is reduced.  The selected Faraday Screens from the $\\lambda$6\\ cm polarization survey demonstrate the existence of numerous high-$RM$ features in the interstellar medium. These structures cover a significant fraction of the surveyed area. $RM$ studies based on pulsars or extragalactic sources aiming to derive the parameters of the large-scale Galactic magnetic field need to take the $RM$-contribution from Faraday Screens into account. The formation of strong regular magnetic fields in thermal low-density regions exceeding the interstellar value needs to be investigated. We note that most of the Faraday Screens visible at $\\lambda$6\\ cm are not seen at longer wavelengths, where their $RM$ causes polarization angle rotations exceeding $180\\degr$. Missing depolarization implies that small-scale fluctuations across the beam may not be significant."
    },
    "1004/1004.4591_arXiv.txt": {
        "abstract": "We present the results of simulations of forced turbulence in a slab where the mean kinetic helicity has a maximum near the mid-plane, generating gradients of magnetic helicity of both large and small-scale fields. We also study systems that have poorly conducting buffer zones away from the midplane in order to assess the effects of boundaries. The dynamical $\\alpha$ quenching phenomenology requires that the magnetic helicity in the small-scale fields approaches a nearly static, gauge independent state. To stress-test this steady state condition we choose a system with a uniform sign of kinetic helicity, so that the total magnetic helicity can reach a steady state value only through fluxes through the boundary, which are themselves suppressed by the velocity boundary conditions. Even with such a set up, the small-scale magnetic helicity is found to reach a steady state. In agreement with earlier work, the magnetic helicity fluxes of small-scale fields are found to be turbulently diffusive. By comparing results with and without halos, we show that artificial constraints on magnetic helicity at the boundary do not have a significant impact on the evolution of the magnetic helicity, except that ``softer\" (halo) boundary conditions give a lower energy of the saturated mean magnetic field. ",
        "introduction": "Stars with outer convection zones tend to possess magnetic fields that display spatio-temporal order with variations that are often cyclic and, in the case of the Sun, antisymmetric with respect to the equatorial plane. Simulations now begin to reproduce much of this behavior \\citep[see, e.g.,][]{Brown10,Kapyla10,GCS10}. A useful tool for understanding the outcomes of such models is mean-field dynamo theory. A central ingredient of this theory is the $\\alpha$ effect. This effect quantifies a component of the mean electromotive force that is proportional to the mean magnetic field \\citep{Mof78,KR80}.  Mean-field theory gives meaningful predictions when to expect cyclic or steady solutions, and what the symmetry properties with respect to the equator are \\citep{B98}. Even in the nonlinear regime, the simple concept of $\\alpha$ quenching, which reduces $\\alpha$ locally via an algebraic function of the mean magnetic field, tends to give plausible results. However, under some circumstances, it becomes quite clear that this simple-minded approach must be wrong. Such a special case is that of a triply-periodic domain. Astrophysically speaking, such a model is quite unrealistic, but it is often employed in numerical simulations. It was also employed as the primary tool to compute $\\alpha$ quenching from simulations \\citep{CH96}. These simulations suggest that $\\alpha$ quenching would set in once the mean field becomes comparable to a small fraction ($\\Rm^{-1/2}$, where $\\Rm$ is the magnetic Reynolds number) times the equipartition value. If this were true also for astrophysical bodies such as the Sun, the $\\alpha$ effect could not be invoked for understanding the dynamics of the Sun's magnetic field.  Later it became clear that there are counter examples to the simple idea that $\\alpha$ is quenched just depending on the local field strength. Surprisingly, simulations later suggested that even in a triply-periodic domain a large-scale magnetic field can be generated that can even exceed the equipartition value \\citep{B01}. However, it would take a resistive time-scale to reach these field strengths, so there was still a problem. Around the same time, the idea emerged that open boundaries might help \\citep{BF00a,BF00b,KMRS00,KMRS02}. This is connected with the fact that an $\\alpha$ effect dynamo produces magnetic helicity of opposite sign at large and small scales \\citep{See96,Ji99}. The magnetic helicity at small scales is an unwanted by-product that can feed back adversely on the dynamo. The resistively slow saturation phase in periodic-box simulations can then be understood in terms of the time it takes to dissipate this small-scale magnetic helicity. It is indeed a particular property of triply-periodic domains that magnetic helicity is strictly conserved at large magnetic Reynolds numbers. A possible remedy might then be to consider open domains that allow magnetic helicity fluxes.  The first simulations with open domains were not encouraging. While it was possible to reach saturation more quickly, the field was found to level off at a value that becomes progressively smaller at larger magnetic Reynolds numbers \\citep{BD01,BS05}. A possible problem with these simulations might be the absence of magnetic helicity fluxes within the domain. Indeed, \\cite{BD01} considered a kinetic helicity distribution that was approximately uniform across the domain, so there were no gradients except in the immediate proximity of boundaries, where boundary conditions on the velocity prevent turbulent diffusion. The situation improved dramatically when simulations with shear were considered \\citep{B05,Kapyla08,HP09}. Shear provides not only an additional induction effect for the dynamo, but it might also lead to an additional source of magnetic helicity flux within the domain \\citep{VC01,SB04,SB06}. More recently it turned out that, even without shear, diffusion down the gradient of small-scale magnetic helicity could, at least in principle, help avoid vanishingly small saturation levels of the mean magnetic field when the magnetic Reynolds number becomes large \\citep{BCC09,Mitra10}.  An important goal of the present paper it to revisit this issue using direct simulations of turbulent dynamos without shear, and even with the same sign of magnetic helicity everywhere, but with a spatial modulation of the helicity within the domain. In other words, the level of turbulence is maintained at a high level throughout the domain, but the amount of swirl diminishes toward the boundaries. In most of the simulations we include a turbulent halo outside the dynamo domain where the Ohmic resistivity is enhanced. This might be important as several simple boundary conditions such as pseudo-vacuum (or vertical field) conditions fix the value of the magnetic helicity artificially, and if fluid motions through the boundary are prohibited, turbulent transport there is impossible.  Our simulations also allow us to make contact with nonlinear mean-field phenomenology where the evolution of the small-scale magnetic helicity is taken into account. This leads then to an evolution equation for an additional contribution to the $\\alpha$ effect, $\\alphaM$. This approach is referred to as dynamical $\\alpha$ quenching. In the present paper we will also attempt to assess the validity of some of the corner stones of dynamical $\\alpha$ quenching. Firstly, there is the magnetic $\\alpha$ of \\cite{PFL}, where the fluctuating magnetic field generates an $\\alphaM$ that is proportional to the current helicity of the fluctuating field. This $\\alphaM$ counteracts the kinetic $\\alpha$, and so saturates the dynamo. Secondly there is magnetic helicity conservation which notes that the total magnetic helicity is nearly conserved under common conditions, and so the helicity in the fluctuating field can be related to the helicity in the large-scale field. Finally, there is the assumption that the mean current helicity of the fluctuating field is proportional to the mean magnetic helicity in the fluctuating field.  As noted above, a problematic prediction of dynamical $\\alpha$ quenching is that rapid (exponential) growth of mean magnetic fields will be halted below equipartition with the turbulent energy.  The export of small-scale helicity could provide a release from this constraint but will likely occur side-by-side with export of the mean field. The interplay between these effects can smother the dynamo even in the presence of small-scale helicity transport.  Treatment of large-scale helicity transport proves significantly more complicated than that of the small-scale helicity, but we will draw some preliminary conclusions.  In \\Sec{dyn} we discuss the dynamical $\\alpha$ quenching phenomenology. In \\Sec{numerics} we describe the numerical setup of the simulations whose results are analyzed in \\Sec{analysis}.  Mean-field models of the systems are discussed in \\Sec{meanfield} and we conclude in \\Sec{conclusions}. ",
        "conclusions": "\\label{conclusions}  Confirming earlier work of \\cite{Mitra10}, we have demonstrated the existence of a diffusive flux $\\meanFFFF_{\\rm f}$ of mean magnetic helicity of the small-scale field. In the present case, however, the Weyl-gauged magnetic helicity of the large-scale field never reaches a steady state. Nevertheless, the magnetic helicity density of the small-scale magnetic field is found to be statistically steady, so the corresponding magnetic helicity flux must be gauge-independent \\citep{Mitra10}. This supports the validity of using the small-scale magnetic helicity as a meaningful proxy for the small-scale current helicity, and hence the magnetic correction to the $\\alpha$ effect.   Understanding the transport of magnetic helicity of the large-scale field, $\\meanFFFF_{\\rm m}$, would be useful for creating analytic post-kinematic models. Although we have not converged on a formula for this flux, it is certainly finite and apparently $\\Rm$ dependent. It is not yet clear whether this flux will converge to a diffusive one for large $\\Rm$. Our mean-field simulations reproduce the final field strength well, reinforcing the conclusion that post-kinematic dynamical $\\alpha$ quenching can be used as part of a mean-field simulation.  The preliminary evidence on the use of small-scale helicity fluxes to escape the small predicted post-kinematic mean fields is negative: the observed flux of large-scale helicity, while poorly modeled, is larger than the flux of the small-scale helicity. If this holds for larger $\\Rm$, it would have the unfortunate result of closing escape holes from $\\alpha$ quenching opened by $\\meanFFFF_{\\rm f}$, but would also imply that dynamo systems with more realistic profiles than simple homogeneity will reach $\\Rm$ independent behavior for high but currently nearly numerically achievable $\\Rm$. It is likely that conclusive evidence for or against $\\Rm$-dependent quenching requires values of $\\Rm$ in the range between $10^3$ and $10^4$ \\citep{BCC09,Mitra10}."
    },
    "1004/1004.1643_arXiv.txt": {
        "abstract": "We report results from a deep high-frequency search for pulsars within the central parsec of Sgr~A* using the Green Bank Telescope.  The observing frequency of $15$\\,GHz was chosen to maximize the likelihood of detecting normal pulsars (i.e. with periods of $\\sim 500$\\,ms and spectral indices of $\\sim -1.7$) close to Sgr~A*, that might be used as probes of gravity in the strong-field regime; this is the highest frequency used for such pulsar searches of the Galactic Center to date. No convincing candidate was detected in the survey, with a $10\\sigma$ detection threshold of $\\sim 10 \\mu$Jy achieved in two separate observing sessions.  This survey represents a significant improvement over previous searches for pulsars at the Galactic Center and would have detected a significant fraction ($\\gtrsim 5$\\%) of the pulsars around Sgr~A*, if they had properties similar to those of the known population. Using our best current knowledge of the properties of the Galactic pulsar population and the scattering material toward Sgr~A*, we estimate an upper limit of 90 normal pulsars in orbit within the central parsec of Sgr~A*. ",
        "introduction": "The detection of radio-emitting neutron stars within the central parsec of our Galaxy would provide us with an unprecedented opportunity to study the super-massive black hole \\sgra\\ and its environs. For example, a single orbiting pulsar would yield a direct probe of the magneto-ionized accretion environment around a black hole, through measurements of temporal changes in the dispersion and rotation measures \\citep{cl97}. Pulsars orbiting within the curved space-time around \\sgra\\ (with orbital periods of $\\lesssim 100$~years) could serve as probes of gravity in the strong-field regime, at field strengths far larger than those accessible with neutron star binaries. The long-term timing of such pulsars, supplemented by accurate astrometry, would allow precise determination of their three-dimensional orbital motion around \\sgra.  Depending on the properties of the pulsars and their orbits, it should be possible to measure subtle general relativistic deviations from Keplerian orbits (e.g.  time dilation, gravitational redshifts, frame dragging, Shapiro delays, etc; e.g. \\citealp{ckl+04,pl04}), and it may even be possible to determine the spin of the black hole (e.g. \\citealp{wk99,kbc+04}).  While theoretical estimates indicate that $100 - 1000$ radio pulsars with periods $\\lesssim 100$~years should be orbiting \\sgra\\ \\citep{pl04}, the observational evidence for neutron stars at the Galactic Center (GC) is mostly indirect. For example, recent studies have found a number of dense clusters of young, massive stars within $\\sim 1$~pc of \\sgra\\ \\citep{sog+03,gsh+05,pgm+06}, while \\citet{wlg06} report X-ray observations of a pulsar wind nebula near the massive stellar complex IRS~13, with properties consistent with it being powered by a young neutron star.  Long term monitoring by \\chandra\\ has revealed an excess of transient sources within a parsec of \\sgra, interpreted by \\citet{mpb+05} as a population of X-ray binaries. The flaring radio and X-ray source detected by \\citet{bry+05}, $\\sim 0.1$~pc from \\sgra, is also likely to be an X-ray binary.  Despite the above evidence for massive stars around \\sgra, there is a remarkable dearth of radio pulsar detections there, despite several deep searches (e.g.  \\citealp{jwv+95,jkl+06,dcl09}). The closest known radio pulsars are 11\\arcmin\\ from \\sgra, and less than one percent of the known pulsar population lies within a degree of the Galactic Center, despite indications of a large population in its environs \\citep{dcl09}. The reason for this deficit is well understood: hyper-strong scattering of radio waves by the turbulent, ionized gas within the central 100~pc of \\sgra, which results in temporal smearing of pulsed signals. This pulse broadening has a strong frequency dependence, $\\propto \\nu^{-4}$, making it near-impossible to detect pulsars at the typical observing frequencies of $\\lesssim 1.4$~GHz (e.g. \\citealp{lc98}).  To overcome the effects of temporal smearing, searches for pulsars at the GC have been carried out at progressively higher observing frequencies over the last few years (e.g. \\citealp{jkl+06,dcl09}), albeit as yet without a detection in the central 25~pc. In this work, we report results from a deep {Green Bank Telescope} (GBT) search for pulsars toward \\sgra\\ at $\\sim 15$~GHz, the highest observing frequency used till date. The choice of this frequency is motivated in Section~\\S\\ref{sec:strat}, and the observations and results described in Section~\\S\\ref{sec:obs}. Finally, Section~\\S\\ref{sec:dis} discusses the constraints placed by our observations on the GC pulsar population, and the prospects for pulsar detections in future surveys. ",
        "conclusions": "\\label{sec:dis}  \\subsection{The 607\\,ms candidate of 2006}\\label{sec:candidate}  The high statistical significance of the pulsed signal detected in the datasets of 2006 implies that it arises either from a genuine pulsar towards the GC or as an artifact of a terrestrial signal (e.g. RFI). While the non-detection in the 2008 dataset might indicate the latter possibility, it should be emphasized that the GC environment is very different from the environments of typical pulsars. Specifically, a pulsar on a short-period ($< 100$~year) orbit around the GC could easily have its emission beam precess away from our line of sight over a time-scale of 2~years. This implies that caution must be used while dismissing possible pulsar candidates towards the GC, although retaining skepticism about their reality. We will hence summarize the characteristics of the pulsed signals seen in the 2006 datasets, and discuss the possibilities that they might arise from a real pulsar or RFI.  Figs.\\,\\ref{fig:Res1} and \\ref{fig:Res2} show that the detection $\\chi^2$ in the two long datasets of 2006 increases steadily over the course of each observing session, indicating that, if the signals are spurious or local RFI, they are at least persistent both over the course of each observation, and over multiple observing epochs. It is curious, however, that the period of the pulsed signal is different in the sessions on 28-29~June and 10-11~August; for RFI, this would require either that we have detected two distinct but alternately intermittent RFI signals, or that the period of the RFI itself is changing. Interpreted in terms of Doppler shifts, the period change corresponds to a velocity change of 1500~km/s, much larger than that associated with the Earth's motion around the Sun or the motions of terrestrial objects, but not implausible for a pulsar orbiting around the GC.  Next, the S/N ratios of our candidates peak at dispersion measures of $\\approx 3000-4000$\\,pc~cm$^{-3}$, comparable to values expected for pulsars at the GC.  Unfortunately, our small fractional bandwidth means that the dispersion in the signal across the band is very small ($\\sim 3.6$\\,ms over the 800~MHz bandwidth). This means that, unlike the situation in low-frequency pulsar surveys, dispersion cannot be used to test whether the signal is of extra-terrestrial origin. Finally, the pulse profiles of the candidates are extremely broad, with a duty cycle of $\\sim 50$\\%, unlike the narrow profiles expected for high-frequency pulsar emission. However, the pulse properties too could be affected by the unique GC environment. For example, the thin-screen approximation might not be applicable for the scattering, or the screen could be much closer to the pulsar than typical estimates of $\\sim 100$~pc; both of these would increase the scattering time and broaden the pulse profile, even at such a high frequency. Specifically, the scattering timescale at a frequency of 14.6~GHz for an object at the Galactic Center is $2.5/D_{\\rm scat}$~seconds, where $D_{\\rm scat}$ is the distance, in pc, of the scattering medium from \\sgra. While the best estimate of $D_{\\rm scat}$ is $\\sim 100$~pc from angular broadening measurements of \\sgra\\ and nearby masers \\citep{cl97}, the effect of scattering material close to \\sgra\\ is much stronger on temporal smearing than on the angular broadening of background sources. As such, the angular broadening estimates of $D_{\\rm scat}$ do not rule out a substantial contribution to the pulse broadening from material closer to \\sgra. One may hence have a sizeable contribution to the pulse broadening from material at $D_{\\rm scat} \\lesssim 10$~pc (e.g. \\citealt{mb06}). The expected temporal smearing timescale would then be $\\gtrsim 250$~ms, comparable to that needed to explain the pulse shape of the 607~ms candidate. The large observed duty cycle of the candidate thus does not rule out the possibility that the signal arises from a genuine pulsar.  It thus appears very difficult to rule out the reality of the candidate on the basis of the 2006 data alone, and, as noted above, the non-detection in 2008 could arise due to precession of the pulsar beam away from our sightline. Thus, while we remain skeptical about the reality of these signals, we conclude that further observations are needed to test the possibility that they arise from a genuine pulsar at the GC.  \\subsection{Constraints on the GC pulsar population}\\label{sec:constraints}  There is compelling but indirect evidence for a substantial population of neutron stars at the Galactic Center. However, strong interstellar scattering along the line of sight has limited past searches for radio pulsars. To overcome these effects, we have used the superb sensitivity of the GBT to carry out a deep search for pulsars in the central parsec of the GC at 15~GHz --- the highest observing frequency at which a search has been carried out to date.  Despite this, we find no convincing pulsar candidates. Was our survey sufficiently sensitive to detect a population of pulsars around \\sgra?  The total number of pulsars detectable at the GC depends on the total number of pulsars accumulated in the region, and the fraction of these objects that would be detectable given our survey sensitivity, and the S/N considerations of Section~\\ref{sec:strat}. The detectable fraction depends particularly on the number of pulsars with flat spectral indices, since these objects influence the pulsar luminosity function most strongly at frequencies $> 10$\\,GHz where they are most easily detectable towards \\sgra.  A simple estimate of the number of detectable pulsars can be obtained by positing that the \\sgra\\ pulsar population has similar properties to those of the {\\it known} population of pulsars and to estimate the fraction of the known population that would be detectable at the GC with our survey. This is done in Fig.\\,\\ref{fig:sens}, where we have plotted pulsars with measured 1.4\\,GHz luminosities [from the \\citet{mhth05} catalog] on a period-luminosity diagram. The solid red line shows the pulsar sensitivity curve of our 14.6\\,GHz survey, obtained using equations\\,(1-3) with a 10$\\sigma$ detection threshold of 10 $\\mu$Jy, and assuming a 10\\% pulsar duty cycle, a scattering screen distance $D_{\\rm scat}=133$\\,pc (\\S\\ref{sec:strat}) and a GC dispersion measure of $1700$\\,pc~cm$^{-3}$.  The sensitivity curve has been scaled to 1.4\\,GHz using a mean pulsar spectral index of $-1.7$. The cutoff in period where most of the sensitivity is lost is taken to be at P$_{\\rm spin}$=$2\\times\\tau_{\\rm scat}$. This is less severe than the scatter-based sensitivity cutoff in Fig.~\\ref{fig:ScatEffectsPlot} but does reflect the fact that some partially recycled or young pulsars (i.e. P$<50$ msec) would be detectable if they were much brighter than our noise threshold.  For comparison purposes, this figure also shows the 8.4\\,GHz sensitivity curve for the Parkes GC survey \\citep{jkl+06}, and the $5\\sigma$ noise threshold for a deep imaging survey of the GC at 22.5\\,GHz, using the Very Large Array \\citep{zmga09}, again scaling both of these to a frequency of 1.4\\,GHz using a mean spectral index of $-1.7$.  We also highlight the seven known pulsars within one degree of \\sgra, including four new ones from \\citet{dcl09} and \\citet{cng+09}.  \\begin{figure*} \\includegraphics[angle=0,width=0.9\\textwidth]{l1400.eps} \\caption{The 1.4 GHz luminosities of the known sample of pulsars versus pulse periods (blue dots). Larger circles (red dots) indicate those pulsars within one degree radius of \\sgra. The $10\\sigma$ pulsar sensitivity of our 14.6\\,GHz search is shown by the solid red line. This was obtained by using the flux density limit (10 $\\mu$Jy at 10$\\sigma$) of our survey to calculate the luminosity limit at the distance of the Galactic Center and then scaling the result to a frequency of 1.4\\,GHz, using an average spectral index of $\\langle\\alpha\\rangle=-1.7$. We also show sensitivity curves derived in the same manner for a Parkes 8.4\\,GHz pulsar survey of the GC (the dashed red line, with a $10\\sigma$ detection threshold of 200\\,$\\mu$Jy; \\citealp{jkl+06}), and a deep 22.5\\,GHz VLA image of the GC (the dotted green line, with a $10\\sigma$ detection threshold of 200\\,$\\mu$Jy; \\citealp{zmga09}).} \\label{fig:sens} \\end{figure*}  A more rigorous estimate can be obtained by computing the fraction of pulsars detectable above some flux density cutoff by considering the pulsar luminosity function at $\\nu_0=1.4$\\,GHz, $f_0 (L)$, combined with the spectral index distribution, $p(\\alpha)$. This is the approach followed by \\citet{pl04} and \\citet{cl97}, but updated with the most recent results on pulsar luminosity functions \\citep{fk06,lfl+06}. We model the 1.4\\,GHz luminosity function as a power law between lower and upper cutoffs $L_{\\rm min}$ and $L_{\\rm max}$ respectively: \\begin{eqnarray} f_0(L) = A L^{-\\beta},  \\qquad A=(1-\\beta) \\left[ L_{\\rm max}^{1-\\beta} - L_{\\rm min}^{1-\\beta} \\right]^{-1}  \\;\\;, \\end{eqnarray} where the normalization is chosen so that the integral over all luminosities is unity, such that $f dL$ is interpreted as the fraction of all pulsars with luminosities between $L$ and $L+dL$. Recent studies suggest that $L_{\\rm min}=0.01\\,$mJy\\,kpc$^2$, $L_{\\rm max}=32\\,$Jy\\,kpc$^2$ and $\\beta = 1.2-2$ \\citep{fk06,lfl+06}. Following \\citet{sks+09}, the spectral index distribution is modeled as a Gaussian \\begin{eqnarray} p(\\alpha)= \\frac{1}{\\sqrt{2 \\pi \\sigma_\\alpha^2}} \\exp \\left[ - \\frac{(\\alpha-\\alpha_m)^2}{2 \\sigma_\\alpha^2}\\right] \\;\\;, \\end{eqnarray} with mean spectral index $\\alpha = -1.7$ and standard deviation $\\sigma_\\alpha =0.35$.  The pulsar luminosity function at some arbitrary frequency is then \\begin{eqnarray} f(\\nu,L ) = \\int_{-\\infty}^{\\infty} d \\alpha \\, p(\\alpha) f \\left( L \\left( \\frac{\\nu}{\\nu_0} \\right)^{-\\alpha} \\right), \\label{lumfrac} \\end{eqnarray} where, for a given spectral index $\\alpha$ chosen from the distribution, the luminosity function $f$ has lower and upper cutoffs \\begin{eqnarray} L_{\\rm min}'  = L_{\\rm min} \\left( \\frac{\\nu}{\\nu_0} \\right)^{\\alpha}, \\qquad L_{\\rm max}' =  L_{\\rm max} \\left( \\frac{\\nu}{\\nu_0} \\right)^{\\alpha}, \\end{eqnarray} and the normalization constant $A$ is modified to \\begin{eqnarray} A = (1 - \\beta) \\left( \\frac{\\nu}{\\nu_0} \\right)^{-\\alpha \\beta} \\left( L_{\\rm max}'^{1 - \\beta} - L_{\\rm min}'^{1-\\beta} \\right)^{-1} \\end{eqnarray} in order to ensure that $f dL$ may be interpreted as the fraction of all pulsars in the luminosity range $L$ to $L+dL$. We integrate equation\\,(\\ref{lumfrac}) to obtain the total fraction of pulsars above a given flux density threshold $L_{\\rm cut} = d^2 S_{\\rm cut}$.   Of course, this threshold depends on both the pulsar period and observing frequency due to propagation effects and changes in the system temperature. Figure \\ref{FracDetect2} hence plots the fraction of the pulsar population detectable above $10\\sigma$ significance at 15\\,GHz versus pulse period.  \\begin{figure*}[t] \\begin{center} \\includegraphics[angle=0,width=0.9\\textwidth]{FracDetect2.eps} \\caption{The fraction of Galactic Center pulsars that would be detectable at $10\\sigma$ significance for a 10\\,h integration on the GBT at 15~GHz as a function of spin period. We assume an intrinsic pulse width of 10\\% of the spin period, and take $D_{\\rm scat}=100\\,$pc.  The curves, from top to bottom, correspond to luminosity function indices of $\\beta=1.2, 1.5$ and $1.8$, with a mean spectral index of $\\alpha = -1.6$ and dispersion $\\sigma_\\alpha=0.35$, as discussed in the text.} \\label{FracDetect2} \\end{center} \\end{figure*}  Note that the present survey is sensitive only to the slower (P$>40$\\,ms) and more-luminous (L$>$40\\,mJy\\,kpc$^2$) pulsars. The millisecond pulsars that are presumably powering the low-mass X-ray binaries near \\sgra\\ \\citep{mpb+05} and the low-luminosity tail of young pulsars \\citep{cng+09} are out of the reach of our survey. However, it is also clear from the figures that this is the first survey capable of peering past the ``fog'' of scattering material and detecting a significant number of pulsars within a parsec of the GC with properties similar to the known pulsar population. Past high-frequency pulsar searches \\citep{jkl+06} or imaging searches \\citep{zmga09} have not had the requisite temporal or flux density sensitivity to detect a significant fraction of the known population. The GBT search thus represents a significant improvement over past pulse searches and imaging efforts.  We note, in passing, that this implicitly assumes that pulsar spectral indices do not typically steepen at high frequencies, $\\gtrsim 5$~GHz.   It is clear from Figs.\\,\\ref{fig:sens}-\\ref{FracDetect2} that the 15\\,GHz GBT survey could have detected a significant fraction ($\\sim$ 1\\%-15\\%) of the pulsars around \\sgra, if they had properties similar to those of the known population. The estimate obtained from Fig.\\,\\ref{fig:sens} is at the high end (15\\%) and it is likely biased by luminosity-dependent completeness limits in pulsar surveys. The lowest estimate ($\\sim$1\\%) comes from the curve in Figs.\\,\\ref{FracDetect2} with the steepest luminosity slope ($\\beta=1.8$). We adopt a nominal value of 5\\% from the intermediate curve ($\\beta=1.5$) which is based on our best current knowledge of the properties of the pulsar population and the scattering material toward \\sgra. Given this detection fraction and our null detection we can use straightforward binomial statistics to estimate the size of the putative pulsar population at the GC. If the probability of detecting a normal pulsar is 5\\%, the non-detection of any pulsars in our survey implies the upper limit (at 99\\% confidence level) of 90~normal pulsars within the 1~pc region around the GC encompassed by the GBT beam. Taking the full allowed range of the detection fraction (1\\% to 15\\%), the upper limit on the number of normal pulsars ranges from 460 to 30, respectively, again at 99\\% confidence level.  Although our estimate is both approximate and subject to much uncertainty, we note that it is significantly lower than the $\\sim 100-1000$ pulsars derived by \\citep{pl04} for the normal pulsar population with orbits of $\\leq$100 yrs (i.e. a radius 50 times smaller than the size of the GBT search area).  Finally, we have shown that the frequency range $10-16$\\,GHz is optimal for searches for ``normal'' pulsars at the GC. The GBT remains the most powerful high-frequency instrument capable of detecting the GC pulsar population for at least the next decade, until the advent of next-generation telescopes like the Square Kilometer Array. The primary limitation of the GBT (and of the present survey) is the relatively small instantaneous bandwidth (800~MHz) available for such searches, resulting in a small fractional bandwidth. An increased bandwidth at the GBT would imply not only an improvement in sensitivity, but also a better rejection of terrestrial signals, using the dispersive sweep of genuine signals across the band. Fig.\\,\\ref{fig:sens} shows that an improvement in sensitivity by merely a factor of $2-3$ would push the GBT into the bulk of the pulsar population. Future GBT experiments should hence aim to utilize the full frequency coverage available with the high-frequency receivers, for both better discrimination against systematic effects and improved sensitivity."
    },
    "1004/1004.1475_arXiv.txt": {
        "abstract": "We compile a sample of 38 galaxy clusters which have both X-ray and strong lensing observations, and study for each cluster the projected offset between the dominant component of baryonic matter center (measured by X-rays) and the gravitational center (measured by strong lensing). Among the total sample, $45\\%$ clusters have offsets $>10''$. The $>10''$ separations are significant, considering the arcsecond precision in the measurement of the lensing/X-ray centers. This suggests that it might be a common phenomenon in unrelaxed galaxy clusters that gravitational field is separated spatially from the dominant component of baryonic matter. It also has consequences for lensing models of unrelaxed clusters since the gas mass distribution may differ from the dark matter distribution and give perturbations to the modeling. Such offsets can be used as a statistical tool for comparison with the results of $\\Lambda$CDM simulations and to test the modified dynamics. ",
        "introduction": "Seventy years after Zwicky's first piece of evidence for dark matter (DM) in galaxy clusters, the physics model for DM still remains ambiguous, ranging from Weakly-Interacting Massive Particles (WIMPs) to gravity-modifying Tensor-Vector-Scalar fields (TeVeS, Bekenstein 2004).  In-between these seemingly conflicting theories, we also have models where the DM changes properties in different environments due to gravitational polarization or interactions with a dark energy field (Blanchet \\& Le Tiec 2009, Li \\& Zhao 2009). These in-between DM models explain why DM in galaxies seems to satisfy the MOND formulae of Milgrom (1983) with a common empirical scale $a_0 \\sim \\sqrt{\\Lambda} \\sim 10^{-10}{\\rm m}\\,{\\rm s}^{-2}$ found by fitting galaxy rotation curves.  They are also consistent with the cosmic microwave background, and the late time dark energy effect $\\Lambda$ or order $a_0^2$.  MOND has gained enormous momentum, partly for its success in making reasonable stable galaxies and explaining galactic phenomenology (e.g., Wang et al. 2008; Wu et al. 2009; Gentile et al. 2009). MOND can also give high velocity encounters of galaxy clusters (Llinares et al. 2009). However, it does not fully account for the discrepancy between the X-ray and dynamical mass in rich clusters of galaxies (Gerbal et al. 1992; The \\& White 1988; Aguirre, Schaye \\& Quataert 2001; Sanders 2003; Tian, Hoekstra, \\& Zhao 2009), which Sanders (2003) explained by introducing a 2~eV neutrino component. We also should note that the environmental-dependent DM of Li \\& Zhao (2009) and the DM dipoles of Blanchet \\& Le Tiec (2009), both mimic MOND for galaxies, might resolve the apparent contradiction of MOND for clusters.  A big challenge to modified gravity is the observations of the bullet cluster 1E0657-56 (Bradac et al. 2006; Clowe et al. 2006; Markevitch et al. 2006). Weak lensing observations of the bullet cluster, combined with earlier X-ray measurements, clearly indicated that the gravitational field of the cluster has an obvious offset from its ordinary matter distribution.  One immediate question one may ask is: Is the bullet cluster the only system that uniquely shows the spatial separation between dark and ordinary matter? Is the phenomenon of DM-baryon separation so rare in the universe, or could it be more common?  In galaxy clusters, most baryons (or ordinary matter) exist in the form of diffuse X-ray emitting gas. The stellar component is larger at the cluster center where bright galaxies are concentrated but DM is still the dominant component. This has been demonstrated by Gavazzi et al. (2003) and Gavazzi (2005): for the inner $\\lesssim 100$~kpc regions of a lensing cluster, the stellar component only occupies a few percent of the total mass. Therefore, the X-ray images could be used as a reasonable approximation of the ordinary matter distribution in a cluster. X-ray observations measure directly the ordinary matter distributions, while the total projected mass distributions (mainly DM) can be measured by gravitational lensing. Thus, a comparison between X-ray and lensing observations of galaxy clusters potentials may unveil possible differences between dark and ordinary matter distributions, just as the observations of the bullet cluster have revealed.  Strong lensing has the potential to determine the cluster mass center with arcsecond precision. The $0.5''$ high spatial resolution of the Chandra X-ray satellite means that we may determine an accurate position of a cluster's baryonic center, though the X-ray data processing may eventually give rise to a larger uncertainty of no more than a few arcseconds (see e.g. Smith et al. 2005). Thus, we can compare the gravitational and baryonic centers of galaxy clusters by investigating a fairly large sample. If the gravitational center of a lensing cluster does not match the ordinary matter center, we could say that a separate DM component might exist. Indeed we expect this result for cluster on-going major merger.   Offsets between lensing and X-ray centers were incidently noticed a decade ago by Allen (1998) when he was studying a sample of 13 clusters. However, no attempt has been made to use such offsets as a dynamical signature, which might be used as quantitative measure of the quality of the DM model by comparing with similar statistics coming from $\\Lambda$CDM simulations.  In this paper, we compile a sample of 38 galaxy clusters that have both strong lensing (SL) and X-ray observations. We carefully check the lensing hypothesis and location of the main potential of the lens if several deflectors are considered. Combining the lensing data with X-ray data, we obtain for each cluster an offset between the lensing center and X-ray center on the projected plane. We use this offset as an ``indicator'' to describe the dark matter-ordinary matter separation. Our data strongly support the idea that the gravitational potential in clusters is mainly due to a non-baryonic fluid, and any exotic field in gravitational theory must resemble that of CDM fields very closely. Moreover, we find that unrelaxed clusters have larger offsets than relaxed clusters.  Interestingly, simulations of CDM+baryon for galaxy clusters in the standard $\\Lambda$CDM cosmology has recently found offsets between the baryonic and DM centers in clusters (Forero-Romero et al. 2010). The offset can be as large as 100~kpc. This roughly supports our findings, though the details of the cluster distribution functions between the simulations and our observational data show some difference. ",
        "conclusions": "\\label{sect:discussion}  The mass in stars of a cluster is small compared to its X-ray gas. Therefore the X-ray images can be used as a reasonable approximation of ordinary matter distribution in a cluster. Similarly the stellar mass component is not dominant in SL modeling even if it often seems the SL and stellar mass centers peak at the same place, the BCG center. Consequently we have used the X-ray images (which is the result of the diffuse intracluster X-ray gas) to find the center of ordinary matter distribution and the lensing mass to find the center of dark matter distribution. Indeed this is an approximation and we have explained that the error in the separation angle shall be much less than $10''$. Therefore the separation between DM and ordinary matter is highly significant for the unrelaxed clusters.  It is noticeable that all the clusters in our sample are clusters with $z>0.1$ (most of them have $z>0.2$). This selection effect is caused by strong lensing clusters because higher redshift clusters have higher lensing probabilities. In order to have high precision determinations of the gravitational center of clusters, we have to focus on strong lensing clusters---which means that our cluster sample has to be a high redshift sample instead of containing many local clusters. In principle, we should have a more unbiased sample with sufficient low redshift clusters which are measured by, e.g., high quality weak lensing.  However, current weak lensing determination of the lensing center is much less robust compared with strong lensing. Indeed, we have investigated clusters with both X-ray and weak lensing observations (but not strong lensing) and found the lensing/X-ray offsets. For example, cluster MS1054 (Jee et al. 2005) has an offset of $19.5''$. MS1008 (Ettori \\& Lombardi 2003) has an offset of $5.43''$. But unfortunately, the errors from weak lensing are much larger. Nevertheless, high quality weak lensing observations of clusters, especially low redshift clusters, will be of particular interest.  It should be noted that the offset here is only 2-D, i.e., the separation on the projected plane. The true separation (3-D) could be much larger. One extreme example is cluster CL0024+17. The redshift distribution of the cluster's member galaxies revealed that the configuration of this cluster is along the line of sight (Czoske et al. 2001). Moreover, recent studies suggested that this cluster may have undergone a head-on collision along the line of sight (Jee et al. 2007; Qin et al. 2008). So its 3-D separation is probably much larger than the 2-D offset in Table~1. The other extreme case is the bullet cluster, where the configuration (as well as the head-on collision) is roughly on the projected plane, indicating that its 3-D separation is close to the 2-D offset.  The bullet cluster has provided us a system that clearly shows the existence of a DM component. RX J1347.5-1145 (Bradac et al. 2008) is similar but with another line of sight projection: the east massive clump has lost all its X-ray gas. For the bullet cluster, Angus et al. (2007) have pointed out that MOND could be rescued if DM is made of 2 eV neutrinos, following the $\\mu$HDM model introduced by Sanders (2003). Indeed such an observation seems to match structure formation in both $\\mu$HDM and $\\Lambda$CDM cosmologies. Meanwhile, Knebe et al. (2009) found that MOND can in principle produce offsets of effective DM, but the offsets are often small, about 1~kpc.  There are also implications in the CDM framework. The significant offset found here is a signal of the merging process and is a measure of the departure from equilibrium, and consequently the X-ray determined dynamical mass based on equilibrium would under-predict compared to lensing-determined mass, which does not use assumptions of equilibrium (Allen 1998; Smith et al. 2005; Zu Hone et al. 2009).  Concerning the SL modeling, it is necessary to identify all the mass distribution that might induce an external shear on the arc system. But it might be also important to include the shear-like perturbation of an offset gas component, a possible systematic effect neglected in previous models. All SL modeling (except the bullet cluster) has been continuously done with a smooth halo component which includes DM component plus the gas with the implicit hypothesis that the gas follows the DM distribution, which is not true except for a fully relaxed cluster. Most sophisticated models also considered the mass perturbation of the stellar component and subhalo associated to early type galaxy members (sometimes with a separate stellar component of the BCG) to improve the modeling of the arc configuration (Kneib et al. 1996; Meneghetti et al. 2003; Keeton 2003; Limousin et al. 2007; Natarajan et al. 2007). Despite that Allen (1998) has explicitly noted the offset of the ordinary gas matter, it is remarkable that nobody has considered the SL modeling with an offset of a large proportion of the ordinary mass, on the same footing as the member galaxies effect. In this paper, we strongly argue that the offset of the ordinary gas matter, which represents $10\\%-20\\%$ of the total cluster mass, should be figured out explicitly for unrelaxed clusters.  In summary, our finding of a high percentage ($45\\%$) of clusters with offsets $>10''$ in the whole sample suggests that it might be a common phenomenon that in unrelaxed galaxy clusters the gravitational field is more or less separated from the ordinary matter distribution. Such separation is best explained if a non-baryonic matter component (DM) does exist. The separations are probably due to dynamical relaxations, as all the clusters with offsets $>10''$ in the sample are unrelaxed clusters.  Indeed simulations of the cluster gas in the CDM framework do show the existence of the large offsets (Ferero-Romero et al. 2010). The simulation (Gottloeber \\& Yepes 2007), called The Marenostrum Universe, was run using the code GADGET2 and followed the adiabatic evolution of gas and dark matter from $z=40$ to $z=0$ in a comoving cube of $500 \\rm h^{-1}$~Mpc. The simulation predicts a median offset of about 18~kpc, which is in general agreement with our whole sample. Nevertheless, the profiles of these offset distributions (i.e., cluster distribution functions with respect to offset) are nonidentical, as shown in Figure~1. A K-S test shows that the significance of differences are $99.2\\%$ and $99.6\\%$ for the whole sample and the subsample of relaxed clusters, respectively.  From Figure~1, the biggest difference between the simulation results and our total cluster sample is that our sample has more clusters with large offsets. The main reasons could be as follows:  (1) Our sample of lensing clusters is biased towards higher redshift while the simulations are not. Obviously, dynamical relaxation will reduce the offset in clusters, leading to fewer clusters with large offsets as compared with the high redshift sample. In other words, most clusters in our sample have $z>0.2$, and it is possible that these high redshift clusters are not fully evolved dynamically as compared with low redshift clusters. Also our sample of 38 clusters is still not big enough, and a larger sample is needed to draw a more robust conclusion.  (2) When a cluster of a given total mass is not fully relaxed and has still large merging clumps about to merge at the center, the length of the caustic lines is increased as compared to a fully relaxed cluster. So it could result in a higher probability to produce lensing arcs, as has been pointed out by Meneghetti et al. (2007). Therefore the fact that we observe more distant clusters which are less relaxed than at $z=0$ has a double bias effect.  (3) The gas physics in galaxy clusters is complicated and less well-understood. Obviously, different treatment of the gas could result in different cluster mass profiles and baryonic distributions, which could give different offset values.  There are other possible explanations, e.g., the merging history model may need to be revised, or the simulations may not have the complete recipes for the physics of DM. Another example is that DM could be possibly coupled to a dark energy scalar field, which distorts the non-linear dynamics at the centers of halos without ruining the large scale success of the standard cold DM. In short, more comparisons of new simulations and larger samples of clusters would be very rewarding in the future, as the offsets observed here provide astrophysicists a new quantitative tool to evaluate different cosmologies.   \\chapter{\\flushright{\\bf"
    },
    "1004/1004.2706_arXiv.txt": {
        "abstract": "The observed angular correlation function of the cosmic microwave background has previously been reported to be anomalous, particularly when measured in regions of the sky uncontaminated by Galactic emission. Recent work by Efstathiou {\\it et al.} presents a Bayesian comparison of isotropic theories, casting doubt on the significance of the purported anomaly. We extend this analysis to all anisotropic Gaussian theories with vanishing mean ($\\langle \\delta T \\rangle = 0$), using the much wider class of models to confirm that the anomaly is not likely to point to new physics.  On the other hand if there is any new physics to be gleaned, it results from low-$\\ell$ alignments which will be better quantified by a full-sky statistic.  We also consider quadratic maximum likelihood power spectrum estimators that are constructed assuming isotropy. The underlying assumptions are therefore false if the ensemble is anisotropic. Nonetheless we demonstrate that, for theories compatible with the observed sky, these estimators (while no longer optimal) remain statistically superior to pseudo-$C_{\\ell}$ power spectrum estimators. ",
        "introduction": "Observations of the cosmic microwave background (CMB) by the {\\it Wilkinson Microwave Anisotropy Probe} (WMAP; {\\it e.g.} \\cite{2003ApJS..148....1B,2010arXiv1001.4744J}) are widely interpreted as confirming the standard model of cosmology in which inflation generates a homogeneous and isotropic background and seeds isotropic, nearly scale-free perturbations.  Yet a variety of tests suggest that, on large scales, something may be amiss \\cite{Copi:2003kt,Eriksen:2003db,deOliveiraCosta:2003pu,2005PhRvL..95g1301L,2006PhRvD..74b3005D,2007PhRvD..75b3507C,Hansen:2008ym,2008arXiv0808.3767C}. (For a wide-ranging assessment of such anomalies in the 7-yr WMAP data see Ref. \\cite{Bennett:2010jb}.) The interpretation of these results is complicated by the {\\it a posteriori} nature of anomaly hunting: any large dataset will contain statistical flukes which, in isolation, can be made to look unacceptable. This is a particularly pernicious problem in the context of large-scale cosmology: with only one sky to observe, frequentist statistics are almost impossible to interpret.  Frequentist results can be made into more concrete Bayesian statements by considering specific alternative CMB theories or classes of theories (see {\\it e.g.} Refs. \\cite{2009arXiv0911.0150G,2009ApJ...690.1807G,Zheng:2010ty}). But a single, fixed dataset can still contribute overwhelming evidence in favour of or against the very same theory, depending on the alternatives against which we are judging (for an elucidation of this point, see Ref. \\cite{2003prth.book.....J}, Sec. 5.5).  In other words there is no unique way to ascribe significance to departures from the standard theory.  This does not imply we should abandon critical evaluations of WMAP and other data: if we simply accept we have an `unlikely' realization of our favoured theory, we might miss the opportunity to discover new physics (or instrumental systematics). Thus frequentist results cannot be dismissed out-of-hand; but we would advocate their interpretation as pointers to interesting areas of work, rather than quantifiable death-knells of existing models or theories.  In the present work, we will consider a long-standing debate about the nature of the angular correlation function $\\mathcal{C}(\\theta)$ of the CMB. The argument is usually phrased in terms of the statistic $\\shalfcut$, which traces the extent to which temperature fluctuations (outside a Galactic mask) are correlated between points separated by $60^{\\circ}$ or more. For a quantitative definition, see Section \\ref{sec:background-notation}.  A number of recent works have attempted to assess the significance of the purportedly anomalous value of $\\shalfcut$, reaching essentially contradictory conclusions. In particular, the frequentist $P$-value \\cite{2008arXiv0808.3767C} suggests the observed sky is highly anomalous, while a Bayesian analysis of the optimally reconstructed sky by Efstathiou {\\it et al.}  suggests the opposite \\cite{2009arXiv0911.5399E}; see also Ref. \\cite{Bennett:2010jb}. However, any Bayesian result pivots crucially on the alternative models considered; the assumptions in Ref. \\cite{2009arXiv0911.5399E} mean that only isotropic models are considered. This is a significant omission, since it leaves open the possibility that suboptimal estimates of $\\shalf$ formed from cut sky data can be reframed as useful measures of anisotropy.  The present work rectifies that omission. The anomaly is analysed from within harmonic space, and then anisotropic theories which make our CMB realization more probable are considered.  The $\\shalfcut$ anomaly is found to be uninformative in the following two senses: \\begin{enumerate} \\item The trivial maximum likelihood anisotropic Gaussian theory for our observed sky\\footnote{Namely, that with covariance matrix $\\mathbf{C}=\\vec{a}\\vec{a}^{\\dagger}$ where $\\vec{a}$ is the observed sky data vector.} does not lead to substantially better likelihoods for the single statistic $\\shalfcut$; \\item Theories constructed specifically to maximize the likelihood of $\\shalfcut$ (ignoring the rest of the information on the sky) also yield little gain. \\end{enumerate} These failures arise from the large variance inherent in using a statistic, such as $\\shalfcut$, which is quartic in the data. Overall, then, the present work reinforces the view that the frequentist `unlikeliness' of the observed sky must be regarded as a statistical fluke.  Some broader results arise from our study. First, we consider the effect of an anisotropic theory on quadratic maximum likelihood (QML) estimates of the power spectrum.  The QML estimators are derived under the (in this context false) assumption of isotropy; despite this, they typically remain superior to pseudo-$C_{\\ell}$ approaches to power spectrum estimation (Section \\ref{sec:background-notation}, with detail in Appendix \\ref{sec:relat-betw-qml}). Second, we present an extremely fast method for finding the maximum angular momentum direction of a CMB map (Appendix \\ref{sec:rapid-calc-l2_m}). Third, we demonstrate that cut-sky correlation functions can be exactly reproduced from the pseudo-$C_{\\ell}$ power spectrum (Appendix \\ref{sec:estim-ctheta-s_12}). This final result, applicable also for weighted data, has been reported previously \\cite{2004PhRvD..69h3524A} but ignored by recent work; to our knowledge no proof appears in the existing literature.  The paper is structured as follows. Section \\ref{sec:background-notation} introduces the necessary background and notation. In Section \\ref{sec:why-mathc-small} we consider, from a harmonic-space perspective, the origin of the low observed $S_{1/2}^{\\mathrm{cut}}$. Anisotropic, Gaussian theories which reproduce this result are considered in Section \\ref{sec:theories}, and show that even the best conceivable fit to the observed CMB makes no substantial improvement to the $S_{1/2}^{\\mathrm{cut}}$ likelihood. Finally, the work is summarized in Section \\ref{sec:conclusions}. Further details and discussion are contained in appendices. ",
        "conclusions": "\\label{sec:conclusions} \\vspace{-0.2cm} There is a classic difficulty in understanding large and complex datasets such as those produced by WMAP and, in the future, {\\it Planck}: they contain so much information that statistical anomalies can be found without any difficulty. We have taken as an example the purported anomalous aspects of the angular correlation function.  Some previous work claims that, after considering these anomalies, the entire cosmological paradigm is to be doubted \\cite{2008arXiv0808.3767C}; other authors claim that apparent anomalies can be dismissed as the product of {\\it a posteriori} analysis \\cite{2009arXiv0911.5399E}. Yet {\\it a posteriori} reasoning must be allowed in science, since otherwise we would rarely, if ever, recognize failings of our existing knowledge.  The contrary statistical claims relating to $S_{1/2}^{\\mathrm{cut}}$ are reconciled by appreciating that, without an alternative theory to test against, there is no unambiguous significance to any anomaly. We have therefore presented an alternative approach to this puzzle: we examined the origin of the low $S_{1/2}^{\\mathrm{cut}}$ in harmonic space, and then attempted to find theories that reproduce the required patterns.  In the process we noted that the cut-sky correlation function contains identical information to the PCL power spectrum estimates. We therefore used the PCL estimates for the majority of our results, but also demonstrated that the standard QML techniques provide more reliable reconstructions of the full sky, even when anisotropy is suspected.  We informed our intuition about the behaviour of the estimators by considering simple anisotropic modifications to the concordance models (contamination, Bianchi and quadrupolar modulation theories). This showed explicitly that the QML estimator biases introduced by anisotropic theories were smaller than or comparable to the PCL case.  Then, by attempting to construct anisotropic Gaussian theories which improve the likelihood of the low $\\shalfcut$, we demonstrated that no significant gains in likelihood for this single statistic are available.  Since there is no suggestion in the observed sky that the underlying ensemble is significantly non-Gaussian \\cite{2010arXiv1001.4538K}, it is implausible that post-Gaussian corrections would substantially change our results. We therefore conclude that the $\\shalfcut$ anomaly is not likely to point to new physics.  If it does have any meaning, the $\\shalfcut$ anomaly (and the underlying shortfall of power seen by PCL estimators) does not indicate a vanishing large-scale correlation function, but rather is related to alignments of low-$\\ell$ power on the full sky (Section \\ref{sec:why-mathc-small}).  It is likely that full-sky statistics can be constructed which capture these unexpected correlations better than $\\shalfcut$ -- and these could evade our likelihood limits.  However, we argued that more trivial modifications (such as taking the ratio $\\shalfcut/\\shalf$) which sidestep our constraint by attaining an infinite likelihood under the `picture' theory ($\\mathbf{C}=\\vec{a}\\vec{a}^{\\dagger}$) are not helpful; see Section \\ref{sec:summary-discussion}. In other words it is highly desirable to choose statistics, such as $\\shalfcut$, that do allow for a finite limit to be placed on the Bayesian statistical gain available under a wide class of alternative straw-man models.  Considering the magnitude of that limit is then, in our view, a plausible way to probe the significance of {\\it a posteriori} anomalies.  \\vspace{-0.4cm} \\subsection*"
    },
    "1004/1004.5583_arXiv.txt": {
        "abstract": "{} {We report on evidence of the inhomogeneity (multiplicity) of the stellar population in the Galactic globular cluster (GC) NGC 3201, which is irregularly reddened across its face.} {We carried out a more detailed and careful analysis of our recently published new multi-color photometry in a wide field of the cluster with particular emphasis on the $U$ band.} {Using the photometric data corrected for differential reddening, we found for the first time two key signs of the inhomogeneity in the cluster's stellar population and of its radial variation in the GC. These are (1) an obvious trend in the color-position diagram, based on the $(U-B)$ color-index, of red giant branch (RGB) stars, which shows that the farther from the cluster's center, the bluer on average the $(U-B)$ color of the stars is; and (2) the dependence of the radial distribution of sub-giant branch (SGB) stars in the cluster on their $U$ magnitude, where brighter stars are less centrally concentrated than their fainter counterparts at a confidence level varying between 99.2\\% and 99.9\\% depending on the color-index used to select the stars. The same effects were recently found by us in the GC NGC 1261. However, contrary to NGC 1261, we are not able to unambiguously suggest which of the sub-populations of SGB/RGB stars can be the progenitor of blue and red horizontal branch stars of the cluster. Apart from M4, NGC 3201 is another GC very probably with an inhomogeneous stellar population, which has essentially lower mass than the most massive Galactic GCs where multiple stellar populations were unambiguously detected for the first time.} {}      \\keywords {globular clusters: general -- globular clusters: individual: NGC 3201} ",
        "introduction": "\\label{introduc}  The southern Galactic globular cluster (GC) NGC 3201, known not only by its peculiar kinematic characteristics but also by irregular differential reddening across its face, was the subject of our recent study (Kravtsov et al. \\cite{kravtsovetal09}) based on a new multi-color photometry in a 14$\\arcmin$x14$\\arcmin$ field of the GC. In that study, where we primarily dealt with some aspects of the properties and characteristics of the cluster stellar population, we also allowed a possible spread in the population, but did not examine it. In a later more detailed analysis of the same data, we were able to find not only apparent manifestations, but also stronger and more objective evidence of the inhomogeneity in the stellar population and of its radial variation in the cluster. The present letter reports on these findings in detail. The obtained results contribute more to our past (Alca\\'ino et al. \\cite{alcainoetal99}) and recent (Kravtsov et al. \\cite{kravtsovetal10}) studies of the inhomogeneity (multiplicity) of the stellar populations in the populous Large Magellanic Cloud (LMC) cluster NGC 1978 and Galactic GC NGC 1261, respectively and to the rapidly growing body of photometric and spectroscopic evidence about multiple stellar populations in both Magellanic Clouds star clusters (e.g., Mackey et al. \\cite{mackeyetal08}; Milone et al. \\cite{milonetal09a}; and references therein) and Galactic GCs (some relevant publications are referred to elsewhere in the paper). ",
        "conclusions": "Based on a more detailed analysis of a new multi-color photometry in an extended field of the differentially reddened GC NGC 3201, we found the following signs of the inhomogeneity (multiplicity) of the cluster's stellar population. First, there is an obvious dependence of the radial distribution of SGB stars in the cluster on their $U$ magnitude: brighter stars are less centrally concentrated than their fainter counterparts at a confidence level fluctuating above 99.2\\% in relation to the color-index of CMD relied on to isolate SGB stars. Second, RGB stars exhibit a systematically different location in both the $U$-$(U-B)$ CMD and the $R$-$\\delta(U-B)$ color-position diagram at different radial distances from the cluster center: the proportion of stars bluer in the $(U-B)$ increases towards the cluster outskirts. We note (1) the same kind of photometric inhomogeneity of RGB and SGB stars in NGC 3201 and in another Galactic GC, NGC 1261 (Kravtsov et al. \\cite{kravtsovetal10}), and also (2) a very similar radial trend in both GCs.  Finally, it is worth mentioning that NGC 3201 after M4, is the second non-massive Galactic GC that very probably has an inhomogeneous stellar population."
    },
    "1004/1004.4646_arXiv.txt": {
        "abstract": " ",
        "introduction": " ",
        "conclusions": "While current data are consistent with a cosmological constant as a source for dark energy, a cornucopia of other physical origins are in agreement as well.  There are many ways to leave $\\Lambda$ as an explanation for cosmic acceleration, some without the fine tuning and other issues.  We briefly outlined approaches based on the microphysical properties of the dark energy, on a high energy physics origin, and on extending gravity beyond general relativity. All are valid possibilities.  The exciting goal of future observations is to explore this wonderland of physics.  We have seen that for the dynamical aspects, next generation measurements of the equation of state and its time variation, $w$ and $w'$, in the calibrated form of $w_0$ and $w_a$ describe the experimental reach to better than observational accuracy. Comparison of tests of growth and expansion could give key clues to the underlying physics, as can contrasting the density, velocity, and gravitational potential fields of large scale structure.  These should be enabled by future wide field imaging and spectroscopic surveys.  To give a more speculative view, the rich variety of information within the CMB, to be revealed by Planck and ground based polarization experiments, can explore signatures of early dark energy.  If the early dark energy density at CMB last scattering is at much higher levels than the part in a billion in the cosmological constant model, then this would be a major clue to the physical origin (note percent level contributions can be accommodated within the barotropic model of Sec.~\\ref{sec:micro}). Lensing of the CMB, and weak lensing of galaxies, can probe aspects of dark energy clustering and interaction.  Eventually we can hope to have as wide an array of aspects of dark energy to probe as have been developed for inflation.  We are very much at the beginning of our explorations of the physics behind cosmic acceleration."
    },
    "1004/1004.1549_arXiv.txt": {
        "abstract": "We present a VLT/FORS1 imaging and spectroscopic survey of the Wolf-Rayet (WR) population in the Sculptor group spiral galaxy NGC 7793. We identify 74 emission line candidates from archival narrow-band imaging, from which 39 were observed with the Multi Object Spectroscopy (MOS) mode of FORS1. 85\\% of these sources displayed WR features. Additional slits were used to observe H\\,{\\sc ii} regions, enabling an estimate of the metallicity gradient of NGC 7793 using strong line calibrations, from which a central oxygen content of $\\log$ (O/H) + 12 = 8.6 was obtained, falling to 8.25 at R$_{\\rm 25}$. We have estimated WR populations using a calibration of line luminosities of Large Magellanic Cloud stars, revealing $\\sim$27 WN and $\\sim$25 WC stars from 29 sources spectroscopically observed. Photometric properties of the remaining candidates suggest an additional $\\sim$27 WN and $\\sim$8 WC stars. A comparison with the WR census of the LMC suggests that our imaging survey has identified $\\sim$80\\% of WN stars and $\\sim$90\\% for the WC subclass. Allowing for incompleteness, NGC 7793 hosts $\\sim$105 WR stars for which N(WC)/N(WN)$\\sim$0.5. From our spectroscopy of H\\,{\\sc ii} regions in NGC~7793, we revise the global H$\\alpha$ star formation rate of Kennicutt et al. upward by 50\\% to 0.45 M$_{\\odot}$ yr$^{-1}$. This allows us to obtain N(WR)/N(O) $\\sim$0.018, which is somewhat lower than that resulting from the WR census by Schild et al. of another Sculptor group spiral NGC 300, whose global physical properties are similar to NGC 7793. Finally, we also report the fortuitous detection of a bright ($m_{\\rm V}$ = 20.8 mag) background quasar Q2358-32 at $z \\sim 2.02$ resulting from C\\,{\\sc iv} $\\lambda$1548-51 redshifted to the $\\lambda$4684 passband. ",
        "introduction": "Classical Wolf-Rayet (WR) stars are helium burning stars descended from massive O stars. Their strong stellar winds produce a unique broad emission-line spectrum, making WR stars easily identifiable in both Local Group \\citep{Massey&Johnson1998} and more distant star-forming galaxies \\citep{Conti&Vacca1990}. WR stars contribute significantly to the chemical evolution of the interstellar medium (ISM) via stellar winds and core-collapse supernova (ccSNe, \\citealt{Dray&Tout2003}). Indeed WR stars are believed to be the progenitors of Type Ib/c supernova and some long Gamma-Ray Bursts (GRBs), however a direct observational link is yet to be established \\citep{Woosley&Bloom2006}.  Wolf-Rayet stars can be divided into subtypes that are nitrogen-rich (WN) or carbon-rich (WC). Metal-rich environments are observed to favour WC stars due to stronger, metal-driven winds during both the WR phase \\citep{Crowther2002} and the progenitor O star phase \\citep{Mokiem2007}, while we expect to find a higher fraction of WN stars in metal-poor environments \\citep{Massey&Johnson1998}. We can investigate the distribution of WR subtypes with respect to metallicity by studying galaxies spanning a range of metallicities. Indeed, many spiral galaxies possess a super-solar nuclei and sub-solar outer regions (eg. \\citealt{Pagel&Edmunds1981, Magrini2007}).  It is thought that WN and WC stars are the progenitors of Type Ib and Ic SNe, respectively. The advent of 8-m class telescopes has allowed searches for WR populations to move beyond the Local Group \\citep{Schild2003}. The identification of a Type Ib/c supernova progenitor is the long-term aim of our survey. The survey consists of 10 nearby star-forming galaxies, and one dwarf irregular galaxy, which were largely chosen based on criteria such as distance, star-formation rate and orientation. To date five galaxies in our sample have been completed (\\citealt{Schild2003, Hadfield2005, Hadfield2007, CrowtherBibby2009} \\& this work), whilst three are underway, and three are in the preliminary stages.  By surveying $\\sim$10 galaxies within 10\\,Mpc, our overall aim is to produce a complete catalogue of $\\geq$10$^{4}$ WR stars which can be referred to when a Type Ib/c supernova occurs. O stars have lifetimes of 3--10\\,Myr, of which $\\sim$0.5\\,Myr is spent in the WR phase \\citep{Crowther2007}. Given this short lifetime, statistically we would expect at least one of the stars in our sample to undergo core-collapse producing a Type Ib (H-poor) or Type Ic (H, He-poor) SNe within the next few decades. \\citet{Kelly2008} investigate the location of different classes of supernovae relative to the light distribution of the host galaxy which supports different progenitors for Type Ib and Ic SNe. \\citet{LeLoudas2010} extend this investigation to the distribution of WR subtypes with respect to the light distribution of two galaxies (M83 and NGC 1313) in our sample. They find WC stars to favour the brighter regions, consistent with the prediction that WC stars are progenitors of Type Ic SNe. Moreover, early-type WN (WNE) stars are found to be more consistent with the distribution of Type Ib SNe, and are ruled out as Type Ic SNe progenitors.  NGC 7793 is a SA(s)d galaxy \\citep{deVaucouleurs1991} that is part of the Sculptor group of galaxies at a distance of 3.91\\,Mpc \\citep{Karachentsev2003}. Despite its relatively low star-formation rate (0.3\\,M$_{\\odot}$yr$^{-1}$, \\citealt{Kennicutt&Lee2008}) its low distance and favourable orientation make it an appropriate addition to our galaxy survey. Previous spectroscopic observations (using the Anglo-Australian 4m telescope) of 4 H\\,{\\sc ii} regions within NGC~7793 have detected weak, broad He\\,{\\sc ii}\\,$\\lambda$\\,4686 emission \\citep{Chun1983}. However no comprehensive WR survey has been undertaken to date. Previous, albeit few, observations of H\\,{\\sc ii} regions within NGC~7793 suggest that it has a shallow metallicity gradient \\citep{Webster1983}.  In this paper we use Very Large Telescope (VLT) optical imaging and spectroscopy, combined with archival VLT and Hubble Space Telescope (HST) images to determine the massive stellar content of NGC~7793. Details of observations of NGC~7793 are presented in Section 2, including details of WR candidate selection. In Section 3 we discuss the properties of the nebular, whilst stellar properties and WR subtypes are determined in Section 4. Section 5 provides a comparison between ground and space-based observations and addresses the completeness of our survey in relation to WR stars in the Large Magellanic Cloud (LMC). A discussion of Giant H\\,{\\sc ii} regions follows in Section 6, whilst Section 7 reports the serendipitous detection of a background quasar Q2358-32. Section 8 discusses the global WR population of NGC~7793 and is compared with the WR content of NGC 300, another Sculptor group spiral, and other nearby galaxies. The paper concludes with a brief summary in Section 9. ",
        "conclusions": "We present the results of a VLT/FORS1 imaging and spectroscopic survey of the WR population of the Sculptor group spiral galaxy NGC~7793. \\begin{enumerate} \\item From archival narrow-band imaging, we identify 74 candidate emission line regions, of which 39 have been spectroscopically observed with the Multi Object Spectroscopy (MOS) mode of FORS1. Of these, 85\\% of these sources exhibited WR features above a 3 $\\sigma$ level.  Additional slits were used to observe H\\,{\\sc ii} regions, enabling an estimate of the metallicity gradient of NGC~7793 using strong line calibrations, from which $\\log$ (O/H) + 12 = 8.61 $\\pm$ 0.05 - (0.36 $\\pm$ 0.10) r/R$_{\\rm 25}$ was obtained. We have estimated WR populations using a calibration of line luminosities of Large Magellanic Cloud stars, revealing $\\sim$27 WN and $\\sim$25 WC stars for sources spectroscopically observed. \\item Photometric properties of the remaining candidates suggest an additional $\\sim$27 WN and $\\sim$8 WC stars. In addition, a comparison with LMC WR stars degraded to the spatial resolution achieved for NGC~7793 suggests that our imaging survey has identified $\\sim$80\\% of WN stars and $\\sim$90\\% for the WC subclass, from which a total of 68 WN and 37 WC stars are inferred within NGC~7793, i.e. N(WC)/N(WN)$\\sim$0.5 \\item Our H\\,{\\sc ii} region spectroscopy permits an updated star formation rate of 0.45 $M_{\\odot}$ yr$^{-1}$ with respect to \\citet{Kennicutt&Lee2008}, from which N(WR)/N(O)$\\sim$0.017 is obtained, assuming N(O)/N(O7V)$\\sim$1.5. A comparison between the WR census of NGC~7793 and survey of the central region of NGC 300 by \\citet{Schild2003} is carried out. This reveals somewhat higher N(WR)/N(O) and N(WC)/N(WN) ratios in NGC~300, in part anticipated owing to its slightly higher mean metallicity. \\item NGC~7793 represents the fourth of ten star-forming spiral galaxies within 2--8 Mpc whose WR populations that we are surveying. Once completed, these will provide a database from which the nature of a future Type Ib/c supernova can be investigated. \\item Therefore, we have considered biases arising from differences between the intrinsic line strengths of WN and WC stars plus ground-based spatial resolution limitations. Therefore, we consider (a) the LMC WR population degraded to that if it was located at a distance of 4 Mpc; (b) differences to the apparent magnitude to the subset of NGC~7793 sources resulting from higher spatial resolution HST/ACS imaging. Upcoming narrow-band WFC3 Hubble Space Telescope imaging of the grand-design spiral galaxy M101 (P.I. M.~Shara) will therefore provide the ideal dataset with which to assess the nature of its future core-collapse SN. This complements both our ground-based studies and the WFPC2 survey of the nearby late-type spiral NGC~2403 by \\citet{Drissen99}. \\item Finally, we also report the fortuitous detection of a bright ($m_{\\rm V}$ = 20.8 mag) background quasar Q2358-32 at $z \\sim 2.02$ resulting from C\\,{\\sc iv} $\\lambda$1548-51 redshifted to the $\\lambda$4684 passband. \\end{enumerate}"
    },
    "1004/1004.4700_arXiv.txt": {
        "abstract": "{Greaves (2006) proposed that three red, high proper motion stars within 10$^{\\circ}$ of 51 Peg (NLTT 54007, 54064, \\& 55547) are co-moving companions to this famous exoplanet host star. While the stars clearly have proper motions similar to 51 Peg, the inferred kinematic parallaxes for these stars produce extremely inconsistent color-magnitude positions 2 to 4 magnitudes below the main sequence. All three stars are likely to be background stars unrelated to 51 Peg.} ",
        "introduction": "In October 1995, Mayor \\& Queloz (1995) reported the discovery of a Jovian-mass companion orbiting the solar-type star 51 Peg in a 4.2 day orbit. The companion 51 Peg b is the prototype of the 'hot Jupiter' class, and the 51 Peg system has been the source of intense study over the past decade and a half. Thus far, surveys have failed to identify any reliable stellar companions, either within a few arcseconds (Luhman \\& Jayawardhana 2002), and out to $\\sim$5' (Raghavan et al. 2006). The existence of low-mass stellar companions to exoplanet host stars are of dynamical interest in the quest to understand the diversity of planetary systems (e.g. Desidera \\& Barbieri 2007).  Greaves (2006) reported that three faint stars might be co-moving with 51 Peg: NLTT 55547 (1$^{\\circ}$.1 away from 51 Peg), NLTT 54064 (7$^{\\circ}$.1) and NLTT 54007 (8$^{\\circ}$.4). It is extremely unlikely a priori that any of these stars would be {\\it bound} companions to 51 Peg as the maximum observed separation for $\\sim$1 M$_{\\sun}$ stars is in the range $\\sim$2500-10000 AU ($\\sim$0.012-0.048 pc; Abt 1988, Close et al. 2003), and the projected separation of the nearest of these (NLTT 55547) is 0.28 pc ($\\sim$58000 AU) if codistant with 51 Peg. The association of these stars with 51 Peg predicates not only on the proper motions of these stars, but also their color- magnitude data (and of course parallax and radial velocity). In this contribution I conclude that for all three stars the color-magnitude data combined with the kinematic distances (predicted by combining their proper motions and the space velocity of 51 Peg) are sufficient to rule out companionship to 51 Peg. ",
        "conclusions": "There is one last possibility to consider for salvaging the hypothesis that these stars could be co-moving with 51 Peg: the possibility that these stars are white dwarfs. However none of the synthetic atmosphere models for degenerate stars by Bergeron et al. (1995) produce objects as red as these three stars (V-K$_s$ $\\simeq$ 4.2-4.6), and indeed the observed V-K$_s$ colors of both DA and non-DA white dwarfs are generally less than V-K$_s$ $<$ 2.2 (Bergeron et al. 1997). Hence there is no reason to believe a priori that the three stars could be cool white dwarf companions to 51 Peg either.  I have demonstrated that while the proper motions of the three M dwarfs are similar to that of 51 Peg, their inferred kinematic and photometric distance estimates are very discordant, and hence none of the stars are likely to be comoving stellar 'siblings' with 51 Peg. Despite the lack of trigonometric parallax measurements for these faint Luyten proper motion stars (NLTT 55547, 54064, and 54007), it appears that the available color-magnitude and astrometric data are probably sufficient to convincingly rule out Greave's (2006) claim of physical association between these stars and 51 Peg. I estimate photometric distances of $\\sim$80, $\\sim$34, and $\\sim$87 pc for NLTT 55547, 54064, and 54007, respectively. The exercise demonstrates the dangers of relying too heavily on proper motions alone on interpreting the nature of widely separated stars of similar projected motion."
    },
    "1004/1004.3911_arXiv.txt": {
        "abstract": "We present global VLBI observations of the first-excited state OH masers in the massive star-forming region Onsala~1 (ON~1).  The 29 masers detected are nearly all from the 6035~MHz transition, and nearly all are identifiable as Zeeman pair components.  The 6030 and 6035~MHz masers are coincident with previously published positions of ground-state masers to within a few milliarcseconds, and the magnetic fields deduced from Zeeman splitting are comparable.  The 6.0~GHz masers in ON~1 are always found in close spatial association with 1665~MHz OH masers, in contrast to the situation in the massive star-forming region W3(OH), suggesting that extreme high density OH maser sites (excited-state masers with no accompanying ground-state maser, as seen in W3(OH)) are absent from ON~1.  The large magnetic field strength among the northern, blueshifted masers is confirmed. The northern masers may trace an outflow or be associated with an exciting source separate from the other masers, or the relative velocities of the northern and southern masers may be indicative of expansion and rotation.  High angular resolution observations of nonmasing material will be required to understand the complex maser distribution in ON~1. ",
        "introduction": "\\object{Onsala~1} (ON~1) is a kinematically complex site of massive star formation.  In the 1.6 and 6.0~GHz transitions of hydroxyl (OH), masers are seen in two disjoint velocity ranges: $-2$ to $+6$~km\\,s$^{-1}$ in the north of the source and $+11$ to $+16$~km\\,s$^{-1}$ projected atop and to the south of the \\ion{H}{2} region \\citep{argon00,fish05b,nammahachak06,fish07,fishreid07}.  A 4765~MHz excited-state maser at $+24.1$~km\\,s$^{-1}$ was also reported by \\citet{gardner83} but never redetected.  The highly excited 13\\,441~MHz transition shows similar velocity structure as the 1.6 and 6.0~GHz masers \\citep{baudry02,fish05}, although an interferometric map of the 13~GHz masers has not yet been published.  The systemic velocity of the \\ion{H}{2} region as measured by \\citet{zheng85} in H76$\\alpha$ emission is approximately $+5$~km\\,s$^{-1}$, which led the authors to conclude that the redshifted masers seen atop the \\ion{H}{2} region are undergoing infall.  However, recent proper motion measurements strongly suggest that the masers are tracing slow ($\\sim 5$~km\\,s$^{-1}$) expansion in ON~1 \\citep{fishreid07}, similar to what is seen in W3(OH) \\citep{bloemhof92,wright04,fishsjouwerman07}.  The ground-state OH masers in ON~1 have been observed with connected-element interferometry \\citep{ho83,argon00,nammahachak06,green07} and with very long baseline interferometry (VLBI) \\citep{fish05b,fishreid07} on multiple occasions. However, the only VLBI observations of the excited-state 6.0~GHz masers were taken by \\citet{desmurs98} with a three-station array. Due to the poor image fidelity inherent with such a sparse array, they were only able to detect 7 masers at 6035~MHz and 2 at 6030~MHz, ranging in flux density from 0.5 to 7.1~Jy.  The large increase in the number of telescopes with 6.0~GHz capability has since made high-fidelity imaging of northern 6.0~GHz maser sources possible.  A prime motivation for observing ON~1 with a high sensitivity, high angular resolution, high spectral resolution VLBI array derives from the experience observing W3(OH), another nearby massive star forming region.  In W3(OH), it was found that the 6030 and 6035~MHz OH masers traced some areas where no ground-state masers have been observed and highlighted portions of other areas where ground-state masers are abundant, allowing a greater understanding of the large-scale (Galactic cloud) structure delineated by molecules in the source \\citep{fishsjouwerman07}.  Of the massive star forming regions visible from the north, ON~1 is one of the brightest 6.0~GHz OH maser sources, with numerous masers detected in the \\citet{baudry97} survey.  Thus, in addition to being an interesting source per se, ON~1 is also an interesting test case to determine to what degree OH masers in the first excited state can provide information not available from the ground-state masers alone. ",
        "conclusions": "We have imaged the 6030 and 6035~MHz OH masers in ON~1 at VLBI resolution.  The distribution of these excited-state masers is similar to that of the ground-state masers.  Unlike in the similar massive star-forming region W3(OH), 6.0~GHz masers are not found in the absence of 1665~MHz masers, perhaps suggesting that ON~1 does not have analogues of the highest-density knots found in W3(OH).  The 6.0~GHz masers are spatially coincident with 1665~MHz masers to within a few milliarcseconds.  Magnetic fields strengths obtained from Zeeman splitting at 6.0 and 1.6~GHz are usually consistent to better than 1~mG, suggesting that multitransition Zeeman associations are acceptable for obtaining estimates of the local magnetic field.  Our observations confirm the existence of a strong ($-12$~mG) magnetic field among the northern, blueshifted masers, as suspected from early EVLA observations of ON~1 \\citep{fish07}.  The large magnetic field here, as well as the distribution of methanol maser spots and presumed distribution of the highly-excited 13\\,441~MHz OH masers, may indicate that the northern and southern masers are excited by two different sources within the \\ion{H}{2} region or that the northern masers trace and outflow.  However, other scenarios, such as large-scale expansion and rotation of all groups of masers in ON~1, cannot be ruled out from the data at hand.  The overall structure of ON~1 remains uncertain in the absence of sensitive, high-resolution observations of the nonmasing material."
    },
    "1004/1004.4536_arXiv.txt": {
        "abstract": "{I highlight the synergies of the Wide Field X-ray Telescope (WFXT) with the next generation radio surveys, including those to be obtained with the Australian Square Kilometre Array Pathfinder and the Square Kilometre Array, and discuss the overlap between the X-ray and radio source populations. WFXT will benefit greatly from the availability of deep radio catalogues with very high astrometric precision, while on the other hand WFXT data will be vital for the identification of faint radio sources down to $\\approx 50~\\mu$Jy. ",
        "introduction": "The Wide Field X-Ray Telescope (WFXT)\\footnote{http://www.eso.org/$\\sim$prosati/WFXT/Over\\-view.html} is a medium-class mission designed to be about two orders of magnitude more sensitive than any previous or planned X-ray mission for large area surveys and to match in sensitivity the next generation of wide-area optical, infrared, and radio surveys \\citep[see][for details]{gia09,mur09}  In five years of operation, WFXT will carry out three extragalactic surveys:  \\begin{itemize}  \\item a wide survey, which will cover most of the extragalactic sky ($\\sim 20,000$ deg$^2$) at $\\sim 500$ times the sensitivity, and $20$ times better angular resolution ($\\sim 5$\") of the ROSAT All Sky Survey;  \\item a medium survey, which will map $\\sim 3000$ deg$^2$ to deep Chandra or XMM - COSMOS sensitivities;  \\item a deep survey, which will probe $\\sim 100$ deg$^2$, or $\\sim 1000$ times the area of the Chandra Deep Fields, to the deepest Chandra sensitivity.  \\end{itemize}  I explore here the possible WFXT synergies with future radio surveys. Sect. 2 describes the current status of radio surveys, while a selection of up-coming and future radio projects is described in Sect. 3. Sect. 4 deals with the source population in deep radio and X-ray surveys, while the X-ray/radio synergy is discussed in Sect. 5. My conclusions are summarised in Sect. 6 As this is {\\it not} a review of future radio projects, only basic information on them will be provided. Readers wanting to know more should consult the relevant references and World Wide Web pages. ",
        "conclusions": "Radio astronomy is at the verge of revolutionary advances, which over the next ten years or so will allow the detection of radio sources as much as $\\ga 100$ times fainter than currently available.  Although at present X-ray and radio surveys detect somewhat different sources, with AGN making up most of the deep X-ray sky while sharing this role with star-forming galaxies in the radio band, synergy between the two bands is already required since, for example, X-ray information is vital to establish the nature of faint radio sources.  The availability of deep radio catalogues with very accurate source positions will be a huge asset to WFXT. Similarly, WFXT data will provide vital help with the identification of faint radio sources down to $\\approx 50~\\mu$Jy. At lower flux densities the X-ray counterparts of most radio sources are expected to be fainter than the WFXT deepest limit.  In summary, the combination of future deep radio surveys with WFXT will shed light on the nature of very faint X-ray and radio sources."
    },
    "1004/1004.1914_arXiv.txt": {
        "abstract": "{The interstellar medium is enriched primarily by matter ejected from evolved low and intermediate mass stars. The outflow from these stars creates a circumstellar envelope in which a rich gas-phase chemistry takes place. Complex shock-induced non-equilibrium chemistry takes place in the inner wind envelope, dust-gas reactions and ion-molecule reactions alter the abundances in the intermediate wind zone, and the penetration of cosmic rays and ultraviolet photons dissociates the molecules in the outer wind region.} {Little observational information exists on the circumstellar molecular abundance stratifications of many molecules. Furthermore, our knowledge of oxygen-rich envelopes is not as profound as for the carbon-rich counterparts. The aim of this paper is therefore to study the circumstellar chemical abundance pattern of 11 molecules and isotopologs ($^{12}$CO, $^{13}$CO, SiS, $^{28}$SiO, $^{29}$SiO, $^{30}$SiO, HCN, CN, CS, SO, SO$_2$) in the oxygen-rich evolved star IK~Tau.} {We have performed an in-depth analysis of a large number of molecular emission lines excited in the circumstellar envelope around IK~Tau. The analysis is done based on a non-local thermodynamic equilibrium (non-LTE) radiative transfer analysis, which calculates the temperature and velocity structure in a self-consistent way. The chemical abundance pattern is coupled to theoretical outer wind model predictions  including photodestruction and cosmic ray ionization. Not only the integrated line intensities, but also the line shapes, are used as diagnostic tool to study the envelope structure.} {The deduced wind acceleration is much slower than predicted from classical theories. SiO and SiS are depleted in the envelope, possibly due to the adsorption onto dust grains. For HCN and CS a clear difference with respect to inner wind non-equilibrium predictions is found, either indicating uncertainties in the inner wind theoretical modeling or the possibility that HCN and CS (or the radical CN) participate in the dust formation. The low signal-to-noise profiles of SO and CN prohibit an accurate abundance determination; the modeling of high-excitation SO$_2$ lines is cumbersome, possibly related to line misidentifications or problems with the collisional rates. The SiO isotopic ratios ($^{29}$SiO/$^{28}$SiO and $^{30}$SiO/$^{28}$SiO) point toward an enhancement in $^{28}$SiO compared to results of classical stellar evolution codes. Predictions for H$_2$O emission lines in the spectral range of the Herschel/HIFI mission are performed.} {} ",
        "introduction": "Asymptotic Giant Branch (AGB) stars are well known to release significant amounts of gas and dust in the interstellar medium via (copious) mass loss. This mass loss dominates the evolution of the star and ultimately, when the stellar envelope is exhausted, causes the star to evolve off the AGB into the post-AGB phase.  The outflow from these evolved stars creates an envelope, which fosters gas-phase chemistry. The chemical complexity in circumstellar envelopes (CSEs) is thought to be dominated by the elemental carbon to oxygen ratio: oxygen-rich M-stars have a C/O ratio less than unity, carbon-rich C-stars have C/O $>1$, and for S-stars C/O is $\\sim$1.  Many have focused on the CSEs of carbon-rich stars in which a rich chemistry takes place. This is reflected by the detection of over 60 different chemical compounds, including unusual carbon chain radicals, in the CSE of \\object{IRC~+10216}, the prototype of carbon stars \\citep[e.g.][]{Cernicharo2000A&AS..142..181C}. In contrast, only 10--12 compounds have been identified in the chemically most interesting oxygen-rich evolved stars, such as \\object{IK~Tau} and \\object{VY CMa} \\citep[e.g.][]{Ziurys2007Natur.447.1094Z}. The first observations of carbon-bearing molecules (other than CO) in oxygen-rich AGBs were somewhat unexpected \\citep[e.g.][]{Deguchi1985Natur.317..336D, Jewell1986Natur.323..311J}. Nowadays, the formation of carbon molecules is thought to be the result of shock-induced non-equilibrium chemistry in the inner circumstellar envelope \\citep[e.g.][]{Duari1999AandA...341L..47D} and/or a complex chemistry in the outer envelope triggered by the penetration of cosmic rays and ultra-violet radiation \\citep[e.g.][]{Willacy1997AandA...324..237W}. Recently, a new interstellar molecule, PO ($X\\,^2\\Pi_r$), has been detected toward the envelope of the oxygen-rich supergiant \\object{VY~CMa} \\citep{Tenenbaum2007ApJ...666L..29T}. Phosphorus monoxide is the first interstellar molecule detected that contains a P--O bond, a moiety essential in biochemical compounds. It is also the first new species identified in an oxygen-rich, as opposed to a carbon-rich, circumstellar envelope. These results suggest that oxygen-rich shells may be as chemically diverse as their carbon counterparts.  Circumstellar molecules have been extensively observed, both in the form of surveys of a single molecular species and in the form of searches for various molecular species in a limited number of carefully selected sources. The aim of these studies was to derive \\emph{(i.)} the mass-loss rate (from CO rotational lines) or \\emph{(ii.)} molecular abundances. For this latter purpose, several methods exist, each with varying degrees of complexity. \\emph{(1.)}~For example, \\citet{Bujrrabal1994AandA...285..247B} and \\citet{Olofsson1998A&A...329.1059O} showed that simple molecular line intensity ratios, if properly chosen, may be used to study the chemical behaviour in CSEs. The use of line intensity ratios has the advantage of requiring no assumptions about a circumstellar model, but it also limits the type of conclusions that can be drawn. \\emph{(2.)} Several authors have derived new constraints on chemical and circumstellar models based on the simplifying assumption of unresolved optically thin emission thermalized at one excitation temperature \\citep[e.g.][]{Lindqvist1988AandA...205L..15L, Omont1993AandA...267..490O, Bujrrabal1994AandA...285..247B, Kim2009}. \\emph{(3.)} Later on, observations were (re)-analyzed based on a non-LTE (non local thermodynamic equilibrium) radiative transfer model \\citep[e.g.][]{Bieging2000ApJ...543..897B, Teyssier2006AandA...450..167T, Schoier2007AandA...473..871S}. In this study, we will go one step further and abandon or improve few of the assumptions still made in many non-LTE analyses. \\begin{enumerate} \\item Quite often, the temperature structure --- being the most important factor determining the molecular line excitation --- is approximated using a power-law  \\citep[e.g.][]{Bieging2000ApJ...543..897B, Teyssier2006AandA...450..167T}. Effects of different heating and cooling mechanisms are hence not properly taken into account. For instance, in the outermost parts of the envelope the temperature profile deviates from a power law distribution once the influence of photoelectric heating by the external interstellar radiation field becomes important \\citep[e.g.][]{Crosas1997ApJ...483..913C, Justtanont1994ApJ...435..852J, Decin2006A&A...456..549D}. \\item The shell is often assumed to expand at a constant velocity \\citep[e.g.][]{Bieging2000ApJ...543..897B, Schoier2007AandA...473..871S}. However, for molecular lines primarily formed in the wind acceleration zone, the effect of a non-constant velocity structure on the derived molecular abundance may be significant. \\item The fractional abundances are estimated to follow an exponential or Gaussian distribution, assuming that the molecules are formed in the inner envelope, and photodissociated or absorbed onto dust grains further out \\citep[e.g.][]{Bieging2000ApJ...543..897B, GonzalesDelgado2003AandA...411..123G, Schoier2007AandA...473..871S}. The effect of extra formation and/or depletion processes in the envelope can hence not be taken into account. \\item Often, a maximum of two molecules (CO and one other) is analyzed at once \\citep[e.g.][]{GonzalesDelgado2003AandA...411..123G, Schoier2007AandA...473..871S}. \\item Integrated line intensities are often used as a criterion to analyse the circumstellar chemical structure. However, line shapes provide us with strong diagnostic constraints to pinpoint the wind acceleration, which in turn has an influence on the deduced fractional abundances. \\end{enumerate}   In this paper, we will study the circumstellar chemical abundance fractions of eleven different molecules and isotopologs in the oxygen-rich AGB star \\object{IK~Tau} based on the non-LTE radiative transfer code GASTRoNOoM \\citep{Decin2006A&A...456..549D, Decin2007A&A...475..233D}, which computes the temperature and velocity structure in the envelope in a self-consistent way. Chemical abundance stratifications are coupled to theoretical non-equilibrium (non-TE) predictions in the outer envelope by \\citet{Willacy1997AandA...324..237W} and compared to the shock-induced non-TE inner wind predictions by \\citet{Duari1999AandA...341L..47D} and \\citet{Cherchneff2006AandA...456.1001C}. \\object{IK~Tau} has been chosen for study because of the wealth of observations which are available for this target and the fact that its envelope is thought to be (roughly) spherically symmetric \\citep{Lane1987ApJ...323..756L, Marvel2005AJ....130..261M, Hale1997ApJ...490..407H, Kim2009}.  \\object{IK~Tau}, also known as NML~Tau, was discovered in 1965 by \\citet{Neugebauer1965ApJ...142..399N}. It is an extremely red Mira-type variable with spectral type ranging from M8.1 to M11.2 and a period around 470 days \\citep{Wing1973ApJ...184..873W}. From dust shell motions detected at 11\\,$\\mu$m using the ISI interferometer, \\citet{Hale1997ApJ...490..407H} deduced a distance of 265\\,pc. This is in good agreement with the results of \\citet{Olofsson1998A&A...329.1059O} who computed a distance of 250\\,pc from integrated visual, near-infrared and IRAS data using a period-luminosity relation. Estimated mass-loss rates range from $3.8 \\times 10^{-6}$ \\citep{Neri1998AandAS..130....1N}  to $3 \\times 10^{-5}$\\,\\Msun/yr \\citep{GonzalesDelgado2003AandA...411..123G}. IK Tau's proximity and relatively high mass-loss rate (for a Mira) facilitates the observation of molecular emission lines.   In Sect.~\\ref{data}, we present the molecular line observational data used in this paper. Sect.~\\ref{analysis} describes the background of the excitation analysis: the radiative transfer model used, the molecular line data and the theoretical ideas on molecular abundance stratification in the envelope. Sect.~\\ref{results} describes the results: we first focus on the velocity structure in the envelope with special attention to the acceleration zone, after which the derived stellar parameters are discussed. Thereafter, the abundance structure for each molecule is derived and compared to the theoretical inner and outer wind predictions and observational results found in the literature. The time variability and SiO isotopic ratios are discussed in Sect.~\\ref{discussion} and water line predictions are performed in Sect.~\\ref{H2O}. We end with some conclusions in Sect.~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  In this paper, we have for the first time performed a self-consistent, non-LTE radiative transfer analysis on 11 different molecules and isotopologs ($^{12}$CO, $^{13}$CO, SiS, $^{28}$SiO, $^{29}$SiO, $^{30}$SiO, HCN, CN, CS, SO, SO$_2$) excited in the circumstellar envelope around the oxygen-rich AGB star \\object{IK~Tau}. In contrast to previous studies, the temperature and velocity structure in the envelope are computed self-consistently, the circumstellar fractional abundances are linked to theoretical outer wind non-chemical equilibrium studies and the full line profiles are used as criteria to deduce the abundance structure. The Gaussian line profiles of HCN and SiO clearly point toward  formation partially in the region where the wind has not yet reached its full velocity. Using the HCN line profiles as criterion, we can deduce that the wind acceleration is slower than deduced from classical theories \\citep[e.g.][]{Goldreich1976ApJ...205..144G}. For a few molecules, a significantly different result is obtained compared to previous, more simplified, studies. SiO and SiS seem to be depleted in the intermediate wind region due to adsorption onto dust grains. The HCN and CS intermediate wind abundance around 50-300\\,\\Rstar\\ is clearly below the inner wind theoretical predictions by \\citet{Duari1999AandA...341L..47D} and \\citet{Cherchneff2006AandA...456.1001C}, which may either indicate a problem in the theoretical shock-induced modeling or possibly that, contrary to what is thought, HCN and CS do participate in the dust formation, maybe via the radical CN through which both molecules are formed. The lack of high signal-to-noise data for CN and SO prevent us from accurately determining the circumstellar abundance stratification. It turned out to be impossible to model all the SO$_2$ line profiles, particularly a few of the high-excitation lines. This may be due to a misidentification of the lines or to problems with the collisional rates. The SiO isotopic fractions point toward high $^{28}$SiO/$^{29}$SiO and $^{28}$SiO/$^{30}$SiO ratios, which are currently not understood in the framework of nucleosynthesis altering the AGB isotopic fractions, but seem to reflect the chemical composition of the interstellar cloud out of which the star is born. Finally, in Sect.~\\ref{H2O}, we present H$_2$O line profile predictions for a few lines which will be observed with the Herschel/HIFI instrument (launched on May, 14 2009).  \\begin{appendix}"
    },
    "1004/1004.0123_arXiv.txt": {
        "abstract": "{Herschel and Planck are surveying the sky at unprecedented angular scales and sensitivities over large areas. But both experiments are limited by source confusion in the submillimeter. The high confusion noise in particular restricts the study of the clustering properties of the sources that dominate the cosmic infrared background. At these wavelengths, it is more appropriate to consider the statistics of the unresolved component. In particular, high clustering will contribute in excess of Poisson noise in the power spectra of CIB anisotropies.}{These power spectra contain contributions from sources at all redshift. We show how the stacking technique can be used to separate the different redshift contributions to the power spectra.}{We use simulations of CIB representative of realistic Spitzer, Herschel, Planck, and SCUBA-2 observations. We stack the 24~$\\mu$m sources in longer wavelengths maps to measure mean colors per redshift and flux bins. The information retrieved on the mean spectral energy distribution obtained with the stacking technique is then used to clean the maps, in particular to remove the contribution of low-redshift undetected sources to the anisotropies.}{Using the stacking, we measure the mean flux of populations 4 to 6 times fainter than the total noise at 350~$\\mu$m at redshifts $z=1$ and $z=2$, respectively, and as faint as 6 to 10 times fainter than the total noise at 850~$\\mu$m at the same redshifts. In the deep Spitzer fields, the detected 24~$\\mu$m sources up to z$\\sim$2 contribute significantly to the submillimeter anisotropies. We show that the method provides excellent (using COSMOS 24~$\\mu$m data) to good (using SWIRE 24~$\\mu$m data) removal of the $z<2$ (COSMOS) and $z<1$ (SWIRE) anisotropies.}{Using this cleaning method, we then hope to have a set of large maps dominated by high redshift galaxies for galaxy evolution study (e.g., clustering, luminosity density). } ",
        "introduction": "The first observational evidence of the cosmic infrared background (CIB) was reported by \\citet{1996A&A...308L...5P} and confirmed by \\citet{1998ApJ...508..123F} and \\citet{1998ApJ...508...25H}. The CIB is composed of the relic emission at infrared wavelengths of the formation and evolution of galaxies and consists of contributions from infrared starburst galaxies and to a lesser degree from active galactic nuclei. Deep cosmological surveys of this background have been carried out with ISO  \\citep [see] [for reviews] {2000ARA&A..38..761G,2005SSRv..119...93E} mainly at 15~$\\mu$m with ISOCAM \\citep [e.g.,] [] {2002A&A...384..848E}; at 90 and 170~$\\mu$m with ISOPHOT \\citep [e.g.,] [] {2001A&A...372..364D}; with Spitzer at 24, 70, and 160~$\\mu$m  \\citep [e.g.,] [] {2004ApJS..154...70P,2004ApJS..154...87D} and with ground-based instruments SCUBA \\citep [e.g.,] [] {2002PhR...369..111B}, LABOCA \\citep [e.g.,][]{2008A&A...485..645B} ,and MAMBO \\citep [e.g.,] []{2000astro.ph.10553B} at 850, 870, and 1300~$\\mu$m respectively. The balloon-borne experiment BLAST performed the first deep extragalactic surveys at wavelengths 250-500$\\mu$m capable of measuring large numbers of star-forming galaxies, and their contributions to the CIB \\cite[][]{2009Natur.458..737D}. These surveys allowed us to obtain a far clearer understanding of the CIB and its sources \\citep [see][for a general review] {2005ARA&A..43..727L}  but many questions remain unanswered such as the evolution of their spatial distribution with redshift. \\\\  The spatial distribution of infrared galaxies as a function of redshift is a key component of the scenario of galaxy formation and evolution. However, its study has been hampered by high confusion and instrumental noise and/or by the small size of the fields of observation. Tentative studies, with a small number of sources at 850~$\\mu$m \\citep{2004ApJ...611..725B}, found evidence of a relationship between submillimeter galaxies and the formation of massive galaxies in dense environments. Works by \\citet{2006ApJ...641L..17F} and \\citet{2008MNRAS.383.1131M} measured a strong clustering of ultra luminous infrared galaxies (ULIRG) detected with Spitzer at high redshifts. Alternatively, the infrared background anisotropies could also provide information about the correlation between the sources of the CIB and dark matter \\citep{2000ApJ...530..124H,2001ApJ...550....7K,2007ApJ...670..903A}, and its redshift evolution. \\citet{2007ApJ...665L..89L} and \\citet{2009arXiv0904.1200V} reported the detection of a correlated component in the background anisotropies using Spitzer/MIPS (160~$\\mu$m) and BLAST (250, 350, and 500~$\\mu$m) data. These authors found that star formation is highly biased at z$>$0.8. The strong evolution of the bias parameter with redshift, caused by the shifting of star formation to more massive halos with increasing redshift, infers that environmental effects influence the vigorous star formation.\\\\  To improve our understanding of the formation and evolution of galaxies using CIB anisotropies, we need more information about the redshift of the sources contributing to the CIB. We also need a method that allows to go deeper than the confusion noise level. In this context, an invaluable tool is the stacking technique, which allows a statistical study of groups of sources that cannot be detected individually at a given wavelength. Its requires the knowledge of the positions of the sources being ``stacked'' as inferred from their individual detection at another wavelength. This knowledge is then used to stack the signal of the sources at the wavelength at which they cannot be detected individually. Since the signal of the sources increases with the number of sources N and the noise (if Gaussian) increases with $\\sqrt{N}$, the signal-to-noise ratio will increase with $\\sqrt{N}$. For an additional description of the basics of stacking techniques we refer to for example \\citet[][]{2006A&A...451..417D} and \\citet[][]{2009arXiv0904.1205M}.\\\\  Stacking was used to measure the contribution of 24~$\\mu$m galaxies to the background at 70 and 160~$\\mu$m using MIPS data \\citep{2006A&A...451..417D}. Contribution from galaxies down to 60 $\\mu Jy$ at 24~$\\mu$m is at least 79\\% of the 24~$\\mu$m, and 80\\% of the 70 and 160~$\\mu$m backgrounds, respectively. At longer wavelengths studies used this technique to determine the contribution of populations selected in the near- and mid-infrared to the FIRB (far-infrared background) background: 3.6~$\\mu m$ selected sources to the 850~$\\mu$m background \\citep{2006ApJ...647...74W} and 8~$\\mu$m and 24~$\\mu$m selected sources to the 850~$\\mu$m and 450~$\\mu$m backgrounds \\citep{2006ApJ...644..769D, 2008MNRAS.386.1907S}. Finally, \\citet[][]{2009arXiv0904.1205M} measured total submillimeter intensities associated with all 24~$\\mu$m sources that are consistent with 24 micron-selected galaxies generating the full intensity of the FIRB. Similar studies with Planck and Herschel will provide even more evidence about the nature of the FIRB sources.\\\\  Theoretically, a stacking technique also could be used to study the mean SED (spectral energy distribution) of the stacked sources \\citep[e.g.,][]{2007ApJ...670..301Z}. The main potential limitations would be caused by the errors in the redshifts of the sources and an insufficiently large number of sources to stack per redshift bin. The observation of sufficiently large fields to which the technique can be applied is now assured by the to Spitzer legacy surveys FIDEL, COSMOS,  and SWIRE\\footnote{http://ssc.spitzer.caltech.edu/legacy/} and Planck and Herschel surveys. Advances in the measurement of the redshift have also been accomplished, although for very small fields for sources up to $z\\sim2$ \\citep [e.g.][] {2006ApJ...637..727C}, and for the larger COSMOS fields up to $z\\sim1.3$ with very high accuracy \\citep{2009apj...690.1236I}. Future surveys are planned to measure the redshifts in larger fields such as the dark energy survey (DES\\footnote{http://www.darkenergysurvey.org/}) or the GAMA spectroscopic survey \\citep [e.g.][] {2008IAUS..245...83B}.\\\\  The difficulties in separating the contribution to the signal coming from different redshifts have handicapped the study of CIB anisotropies. However, once the mean SEDs of infrared galaxies per redshift bin are obtained we can use this information to analyze CIB anisotropies. The SEDs obtained with the stacking technique can be used to ``clean'' the low-redshift anisotropies (or at least a significant part of them) from the CIB maps. This can be performed by subtracting the undetected low-redshift ($z<1-2$) populations from the maps using their mean colors and thus build maps dominated by sources at higher redshifts. This also facilitates the study of the evolution of large-scale structures at high redshift by removing the noise coming from low redshifts. \\\\  In this paper, we use the simulations and catalogs presented in \\citet{2008A&A...481..885F}\\footnote{The simulations are publicly available at http://www.ias.u-psud.fr/irgalaxies} to study the limitations of stacking techniques in CIB anisotropy analysis. We stack 24~$\\mu$m sources detected with MIPS in Planck, Herschel, and SCUBA-2 simulated observations. The catalogs and maps were created for different levels of bias between the fluctuations of infrared galaxy emissivities and the dark matter density field. We use a bias $b=1.5$, which is very close to that measured by \\citet{2007ApJ...665L..89L}.\\\\  The paper is organized as follows. In Sect. \\ref{sec:Methods}, we explain the method used to study the capabilities of the stacking once the redshift of the sources is known. Section \\ref{sec:Limits} details the elements that limit the accuracy of the stacking technique. In Sect. \\ref{sec:Test}, we test the technique for studying the mean SEDs of galaxies. In Sect. \\ref{sec:Cleaning}, the feasibility of using information about the SEDs to clean the observations of low-redshift anisotropies is studied. The results are summarized in Sect. \\ref{sec:Summary}. Throughout this paper, the cosmological parameters are assumed to be $h=0.71,\\Omega_{\\Lambda}=0.73,\\Omega_{m}=0.27$. For the dark-matter linear clustering, we set the normalization to be $\\sigma_{8}=0.8$. ",
        "conclusions": "}  We have described a stacking algorithm and illustrated its capabilities using Spitzer observations. We have studied the accuracy of the stacking method as a means of determining the average fluxes of classes of undetectable sources at long wavelengths. The results show that the technique will be capable of measuring accurate fluxes at both far-infrared and submillimeter wavelgnths for sources as faint as 80~$\\mu$Jy at 24~$\\mu$m using average colors.\\\\  With the successful commissioning of the Planck and Herschel missions, large maps (even all-sky for Planck) from 250~$\\mu$m to the millimeter wavelength range are now available. SCUBA-2 and other submillimeter cameras (e.g., LABOCA) will provide data of higher angular resolution in the submillimeter. We have applied the stacking method to the Herschel, Planck, and SCUBA-2 simulated data and measured the full average SED of populations of sources detected at 24~$\\mu$m. The strong variation in the $S_{24}/S_{\\lambda}$ color with redshift requires us to define the populations to which the method will be applied not only in ranges of $S_{24}$ but also in terms of (photometric) redshift. We show we are able to measure the mean flux of populations 4 to 6 times fainter  than the total noise at 350~$\\mu$m at redshifts $z=1$ and $z=2$, respectively, and 6 to 10 times fainter than the total noise at 850~$\\mu$m, at the same redshifts. We have been able to reproduce the SED at wavelengths 70, 160, 250, 350, 500, and 850~$\\mu$m of a population of sources with mean flux $S_{24}=0.11$~mJy and $S_{24}=0.135$~mJy at redshifts $z=1$ and $z=2$, respectively. \\\\  In the deep Spitzer fields, the detected 24~$\\mu$m sources constitute a large fraction of the anisotropies. We have shown that the method presented in this paper enables an excellent (350-850 COSMOS) to good (350-850 SWIRE) removal of both the Poissonian and correlated low-z anisotropies. The relative contribution of sources to the background anisotropies up to $z=2$ decreases with wavelength in the model. This property is expected to remain valid independently of the details of the model from 250~$\\mu$m to the millimeter range. Although the accuracy of the subtracted map is lower at  850~$\\mu$m, the cleaning of the power spectrum is quite effective (because the contribution of the low-redshift sources is small at these submillimeter wavelengths). \\\\  The same technique could also be used to remove from the observations all the contributions from sources for which we have estimated a flux, to decrease the confusion noise caused by infrared galaxies. This would be interesting for the detection of other types of sources (for example, SZ sources in Planck data).\\\\  The method allows us to build $z\\gtrsim1-2$ CIB maps from the submillimeter to the millimeter. We have found that the method can also be successfully applied at the other Herschel and Planck wavelengths than those tested in this paper. The longer wavelengths at which this can be achieve will depend on the success of the component separation and not on the removal of the $z<2$ sources. We can then hope to have a set of large CIB maps dominated by high-redshift galaxies. This set of CIB maps at different wavelengths dominated by $z>2$ sources will be a powerful tool for studying the evolution of the large-scale structure of infrared galaxies. The effect of the K-correction ensures that each of these maps (at different wavelengths) are dominated by particular high-redshift ranges. Methods of independent component separation based on the correlation matrix between these maps \\citep [e.g.,][] {2003MNRAS.346.1089D} should allow us to extract maps and power spectra for a number of redshift ranges equal to the number of maps. This last step will fulfill the main objective of this work. It will allow the study of the evolution of the IR galaxy clustering at high redshifts by means of the power spectrum analysis of CIB anisotropies. These maps may also be used to help us understand the contribution of high-z IR galaxies both to the CIB and the star-formation history."
    },
    "1004/1004.0912_arXiv.txt": {
        "abstract": "Observing inversion lines of ammonia (NH$_3$), complemented by rotational lines of NH$_3$ and other molecular species, provides stringent constraints on potential variations of the proton-to-electron mass ratio, $\\mu$. While a limit of $|\\Delta$$\\mu|$/$\\mu$ $\\sim$ 10$^{-6}$ is derived for a lookback time of 7$\\times$10$^9$\\,yr, nearby dark clouds might show a significant variation of order (2--3)$\\times$10$^{-8}$, possibly being related to chameleon fields. The detection of radio-loud quasars with strong molecular absorption lines at redshifts $z$ $>$ 1 as well as the identification of a larger sample of nearby molecular clouds with exceptionally narrow lines ($\\Delta V$ $<$ 0.2\\,km\\,s$^{-1}$) would be essential to improve present limits and to put the acquired results onto a firmer statistical basis. ",
        "introduction": "\\label{sec1} Comparing redshifts of various spectral lines, observed toward the same distant highly redshifted object, has the potential to yield important contraints on temporal variations of fundamental constants of the standard model over timescales of billions of years, even surpassing the age of the solar system. A pre-condition is that the lines in question have different dependencies on the respective constant. Among the most commonly studied fundamental parameters are the dimensionless fine structure constant, $\\alpha$, and the proton-to-electron mass ratio, $\\mu$ (e.g., Uzan 2003; Garc{\\'i}a-Berro et al. 2007; Dent 2008). It may be possible that variations in $\\mu$ are more pronounced than those in $\\alpha$ (e.g., Flambaum 2008). Therefore, studies contraining $\\mu$ may be particularly rewarding. ",
        "conclusions": "\\label{sec5}  Measurements of ro-vibrational H$_2$ quasar absorption spectra yield $|\\Delta \\mu|$/$\\mu$ $<$ 10$^{-5}$ over the last 80\\% of the age of the Universe. Radio data including NH$_3$ inversion lines result in $|\\Delta \\mu|$/$\\mu$ $\\la$ 10$^{-6}$ over the last 50\\% of the age of the Universe. Similar data obtained from local dark clouds suggest a potential variation of order (2--3) $\\times$ 10$^{-8}$."
    },
    "1004/1004.0586_arXiv.txt": {
        "abstract": "We argue that observations of old neutron stars can impose constraints on dark matter candidates even with very small elastic or inelastic cross section, and self-annihilation cross section. We find that old neutron stars close to the galactic center or in globular clusters can maintain a surface temperature that could in principle be detected. Due to their compactness, neutron stars can accrete WIMPs efficiently even if the WIMP-to-nucleon cross section obeys the current limits from direct dark matter searches, and therefore they could constrain a wide range of dark matter candidates. ",
        "introduction": "Since the initial discovery of the ``missing mass'' problem by Zwicky in the 30's, a lot of theoretical, experimental, and observational effort has been put in unveiling the mystery of dark matter. A number of possibilities have been proposed, including modifications of the gravitational theory, hidden sector(s), primordial black holes and other massive objects, and new dark matter particles.  An attractive solution of the dark matter problem within the context of particle physics can be provided by a class of models with Weakly Interacting Massive Particles (WIMP). The Standard Model does not have a WIMP with the required characteristics, which means that WIMPs are probably related to physics beyond the Standard Model. There are several dark matter propositions according to what extension of the Standard Model one selects: supersymmetry \\cite{Jungman:1995df,Bertone:2004pz}, hidden sectors~\\cite{Pospelov:2007mp,Hambye:2008bq}, technicolor ~\\cite{Gudnason:2006yj,Kouvaris:2007iq,Ryttov:2008xe,Sannino:2010ia}, etc.  All currently existing evidence in favor of dark matter (as, for example, WMAP~\\cite{Dunkley:2008ie}) is of gravitational origin.  In order to distinguish between the dark matter models, a direct (non-gravitational) detection of dark matter particles is required. The most important parameters that determine the perspectives of the direct detection are the cross section $\\sigma_N$ of the dark matter-to-nucleon interaction, and the dark matter self-annihilation cross section $\\sigma_A$, or the decay rate in models with decaying dark matter. Underground direct search experiments such as CDMS~\\cite{Ahmed:2009zw} and Xenon~\\cite{Angle:2008we} have put tight constraints on the spin-independent and spin-dependent cross sections of WIMPs scattering off nuclei targets at the level of $\\sigma_N\\lesssim 10^{-43}{\\rm cm}^2$. Interestingly, the DAMA collaboration~\\cite{Bernabei:2010mq} claims the observation of an annual modulated signal with high statistical significance.  A possible reconciliation of all the underground search experiments points to the existence of dark matter with excited states, in which case WIMPs can interact also inelastically~\\cite{TuckerSmith:2001hy,TuckerSmith:2004jv}, or to less mainstream scenarios as in~\\cite{Khlopov:2007ic,Khlopov:2008ty}.  In the last twenty years, there have been several attempts to constrain the properties of WIMPs by looking at signatures related to the accretion and/or annihilation of WIMPs inside stars. This includes the capture of WIMPs in the Earth and the Sun~\\cite{Press:1985ug,Gould:1987ju,Gould:1987ww}, the self-annihilation of WIMPs that can lead to an observable neutrino spectrum~\\cite{Jungman:1994jr,Nussinov:2009ft}, the effect of dark matter in the evolution of low-mass stars~\\cite{Casanellas:2009dp,Casanellas:2010sj}, and the study of the WIMP accretion and/or annihilation inside compact stars such as neutron stars ~\\cite{Goldman:1989nd,Kouvaris:2007ay,Sandin:2008db} and white dwarfs~\\cite{Bertone:2007ae,McCullough:2010ai}.  Compact objects, and in particular, neutron stars constitute a potentially promising way of constraining dark matter models. Firstly, the high baryonic density in compact stars increases the probability of WIMP scattering within the star and eventually the gravitational trapping. This is crucial in view of the tiny value of $\\sigma_N$. It should be noted that in the models with the inelastic dark matter interactions, the elastic and inelastic cross sections of the WIMP scattering inside the star are of the same order, because the WIMP velocity is much higher that the asymptotic value of $220 \\text{km/s}$, and its kinetic energy is therefore much larger than the splitting between the WIMP excited and ground states. Secondly, at the late stages of their evolution, neutron stars can be rather cold objects due to lack of possible burning or heating mechanisms, and therefore heating by annihilation of the dark matter could produce an observable effect.  Close cousins of the neutron stars are the white dwarfs, the second most compact objects. They are easier to observe due to their larger surface area. However, they are lighter and less dense than neutron stars. For an efficient capture, a dark matter particle has to collide at least once per star crossing. For a neutron star, this requires the cross section to satisfy $\\sigma_N \\gtrsim 10^{-45}\\text{cm}^2$, while for a solar mass white dwarf of radius $5000$~km one should have $\\sigma_N \\gtrsim 10^{-39}\\text{cm}^2$. As a result, neutron stars can probe much smaller values of the WIMP-to-nucleon cross section.  In this paper we consider constraints on the dark matter parameters that may arise from the neutron star cooling. This question has been addressed previously~\\cite{Kouvaris:2007ay}. Here we concentrate specifically on the effect of the dark-matter-rich environments such as the Galactic center or cores of the globular clusters, and on the role of the neutron star progenitor.  In section II we review the accretion and annihilation rates of dark matter WIMPs relevant for neutron stars. In section III we study how the accretion of dark matter on the progenitor of a neutron star can affect the accretion and annihilation rates of WIMPs in a neutron star emerging after the collapse of its progenitor. In section IV, we present lower bounds for the surface temperature of a neutron star as a function of its position in the galaxy, and in section V we derive similar limits for neutron stars in globular clusters. We conclude in section VI. ",
        "conclusions": "In this paper we examined the effect of WIMP annihilation on the temperature of a neutron star. We estimated the surface temperatures of old neutron stars according to their location in the galaxy or in a globular cluster. We also investigated the effect of a neutron star progenitor on the accretion of WIMPs onto the neutron star.  We found that, although a considerable number of WIMPs is accumulated by the progenitor during the evolution preceeding the formation of a neutron star, the effect of this accumulation is observable only in cases where the annihilation cross section is extremely small.  We argued that observations of neutron stars with low (of order $10^5$~K or lower) surface temperature will put constraints on a large set of dark matter candidates.  Due to their high density, the neutron stars accrete the dark matter at a significant rate even when the WIMP-to-nucleon cross section (elastic or inelastic) is as low as $10^{-45}\\text{cm}^2$, which is two orders of magnitude lower than the current experimental limit. Even for lower values of the cross section, the effects of WIMP accretion and annihilation may be observable in neutron stars which are situated in dark-matter-rich environments such as the galactic center and cores of globular clusters. Thus, the neutron stars can probe much smaller WIMP cross sections than less dense objects such as, e.g., white dwarfs.  The WIMP constraints that we presented are valid even if the WIMPs have a very small annihilation cross section as low as $10^{-57}\\text{cm}^2$ (or even lower for large local dark matter densities). This means that our constraints hold also for a variety of WIMP candidates that are produced non-thermally, for which the annihilation cross section is, in  general,  a free parameter.  Perfect candidates to test the WIMP-burning heating mechanism are isolated neutron stars (i.e., not showing accretion of ordinary matter from other objects), which are old but appear warmer than predicted by the conventional cooling models. There is a couple of examples of such candidates. One of them is J0437-4715, a few billion years old neutron star with a roughly $10^5$~K temperature~\\cite{Kargaltsev:2003eb}. Although this temperature can be sustained by WIMP burning, it would require a substantial local dark matter density, which is unlikely as J0437-4715 is only 140 pc from the Earth. Unless there is a peak in the dark matter density at the position of J0437-4715, WIMP burning cannot explain this temperature. A similar candidate is J0108-1431 at 130~pc from Earth, with a temperature $\\sim 9 \\times 10^4$ K~\\cite{Mignani:2008jr}. Like in the case of J0437-4715, this temperature is still higher than what the dark matter burning can provide, assuming the dark matter density at the location of J0108-1431 is the same as around the Earth. Candidates like the above, with smaller temperatures or in rich dark matter regions such as globular clusters, might make it possible to constrain a large class of dark matter WIMP scenarios."
    },
    "1004/1004.5085_arXiv.txt": {
        "abstract": "We have used the Extreme Ultraviolet Imaging Spectrometer (EIS) on the {\\it Hinode} spacecraft to observe large areas of outflow near an active region. These outflows are seen to persist for at least 6 days.  The emission line profiles suggest that the outflow region is composed of multiple outflowing components, Doppler-shifted with respect to each other.  We have modeled this scenario by imposing a double-Gaussian fit to the line profiles. These fits represent the profile markedly better than a single Gaussian fit for Fe~{\\sc xii} and {\\sc xiii} emission lines. For the fastest outflowing components, we find velocities as high as 200~km~s$^{-1}$. However, there remains a correlation between the fitted line velocities and widths, suggesting that the outflows are not fully resolved by the double-Gaussian fit and that the outflow may be comprised of further components. ",
        "introduction": "\\label{sec:intro}  One of the most significant discoveries of the Extreme Ultraviolet Imaging Spectrometer (EIS) on the {\\it Hinode} spacecraft is the detection of large areas of outflowing plasma at the boundaries of active regions \\citep{Dosc07a, Saka07a, Harr08a, Dosc08a}. These outflowing regions were found to occur in areas of low line emission intensity, often adjacent to coronal loops. \\citet{Dosc07a} also found the outflowing regions to exhibit larger spectral line widths than found in the much brighter active region closed loops.  These line widths are in excess of pure thermal Doppler broadening. \\citet{Dosc08a} subsequently found a strong positive correlation between the outflow velocity (the Doppler shift of the line emission) and the non-thermal velocity (the width of the emission line).  \\citet{Harr07a} have related the outflowing regions to coronal mass ejections and \\citet{Harr08a} and \\citet{Dosc08a} have postulated that the outflowing regions could be contributors to the solar wind.  This conclusion is consistent with results based on completely independent studies of the heliospheric magnetic field, for example by \\citet{Schr03a}.  Given such implications, the further characterization  of outflowing regions could be important for the understanding of fundamental physical processes involved in production of the solar wind and mass flow into the corona.  Emission line widths in excess of their thermal Doppler widths have been observed since the rocket flight analysis of \\citet{Bola75a}.  However, the origin of the spectral broadening remains unclear; possible explanations include   turbulence in the atmosphere due to magnetic reconnection \\citep{Park88a} and the presence of coronal waves \\citep{Mari92a}.  Earlier studies of non-thermal line broadening did not link the broadening to bulk mass flows because there were few high resolution EUV-UV solar spectra coupled to images such that either the broadening or bulk mass flows could be related to particular coronal structures.  The outflows under present discussion would have been difficult to find in earlier studies because of  their appearance in areas of low intensity in coronal spectral lines, making their measurement difficult. It was not until the EIS instrument that the required spatial and spectral resolution became available to accurately determine the spectral properties of the outflowing region and marry these with magnetic structures of the solar atmosphere. Of particular interest among the EIS findings is the correlation of line shift with line width found by \\citet{Dosc08a} and \\citet{Hara08a}.   The fact that those emission lines that show the largest non-thermal Doppler velocities also display the largest widths raises the possibility that the outflows may result from multiple flow sites, all Doppler shifted relative to one another.   Spectral lines at the EIS spatial and spectral resolution might be convolutions of line emission from multiple unresolved flow sites. This possibility was suggested by both \\citet{Dosc08a}  and \\citet{Hara08a} but was not explored further; their spectral analyses assumed simple Gaussian line profiles.  An object of the present paper is to determine whether the apparent  excess line widths can be attributed to the line emission being poorly represented by a single Gaussian. We expand upon the previous works dealing with outflows in EIS spectra by attempting to model the outflowing plasma as a blend of outflow sites with different flow speeds.  This method of analysis depends on the different flow velocities being sufficiently shifted in relation to one another as to be spectrally resolved. We thus focus on observations that have previously been identified as displaying highly asymmetrical line profiles \\citep{Dosc08a}. We also limit our present model to an outflow region that can be well-represented by a sum of two flow velocities, with the emission line profiles modeled as a sum of two Gaussian components. The remainder of this paper is organized as follows:  in Section~\\ref{sec:obs} we outline the observations and explain the  EIS data reduction procedures. In Section~\\ref{sec:results} we present the results of the double Gaussian line fitting technique and discuss possible interpretations in Section~\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  As shown in Figure~\\ref{fig:spectra}, a double Gaussian representation of the emission line profiles is a more accurate fit to the data than a single Gaussian representation for certain emission lines in the outflowing region. From the results presented in Section~\\ref{sec:results}, we see that the double Gaussian representation is only applicable to the outflowing regions. Within these regions we find that the outflow is not only due to the fast moving secondary Gaussian component, but the primary component is also blue-shifted (see Figure~\\ref{fig:velocity}).  By its nature the double Gaussian fit results in a reduction in the Doppler velocity of the primary component compared to that of a single Gaussian fit.  We find velocities for this component of the order of 10~km~s$^{-1}$, compared to $\\sim20$~km~s$^{-1}$ found by \\citet{Dosc08a} for the same outflow. However, the secondary component exhibits significantly greater outflow velocities, often as large as 200~km~s$^{-1}$.  Outflows are found to persist for the 6 day duration of the observations. In Figure~\\ref{fig:variation}, we show the variations in the derived flow velocities over this period.  Here, we have compared the velocities of the primary and secondary flow components in both the eastern and western outflowing regions, and also indicated the standard deviation of velocities throughout the respective regions. Despite the observations spanning a significant extent of the solar rotation, there is relatively little change in the line-of-sight velocities over this time.  The western outflow region (black lines in Figure~\\ref{fig:variation}) show a slight decreasing trend with time (for both the primary and secondary components). The eastern outflow region (red lines) shows a more pronounced, increasing velocity, trend over time.   From the geometry of the AR, these apparent flow speeds are what we would  expect, although we do note that there is little statistical significance to the trends. It seems evident that the outflow is not confined to a narrow cross-sectional area but, rather, emanates over a wide `cone' of emission. The large standard deviation in the velocities shown in Figure~\\ref{fig:variation} further supports this interpretation. As discussed in Section~\\ref{sec:results}, the outflowing regions are also obscured by the overlying AR loop system, which further complicates the analysis.  While the double Gaussian fit appears to represent the \\fexii\\  spectra well, the results displayed in Figure~\\ref{fig:velocity vs width main} suggest that the double Gaussian does not describe the outflow completely. In all but the first two observations there is a correlation between the primary line velocity and width in at least one of the two outflowing regions.   If the assumption that the excess line width is due to multiple flow components then this remaining correlation suggests that a double Gaussian fit does not resolve the flow components.  The EIS rasters for observations 1 and 2 do not cover the entire outflowing region, Figures~\\ref{fig:intensity} and \\ref{fig:intensity secondary} show that the outflow region is to the north-west of the raster and very likely extends beyond the raster boundary.  Given that we then do not have data for the entire outflowing region, it may not be surprising that the same correlation between velocity and width is not found for these two observations.  This correlation that is seen in the primary Gaussian component is not seen for the secondary component (see Figure~\\ref{fig:velocity vs width minor}). One should be careful on the conclusions drawn from the width in a quantitative sense since the width of the primary and secondary components were set as equal in the fitting algorithm.  However, the fact that there is no correlation between the velocity and width of the secondary Gaussian component suggests that if the outflow is indeed comprised of more than two components then the extra, unresolved, flow components are to be found within the primary component. The only exception to this is observation 5 where we see a complex relation between velocity and width, suggesting that we are not fully resolving the spectral data.  Observations 5, 6, 7, and 8 all show some evidence of regions with high velocities but low FWHM for the secondary component. The preponderance of these pixels are found  in areas where the primary outflow speed is relatively low, away from the loop footpoints. This effect could be due to the direction of the outflowing material, with the primary and secondary components diverging on moving away from the footpoints. However, this remains speculation without knowing the precise topology of the outflowing plasma.  Using a different analysis technique on a different AR, \\citet{Pont09a} have also measured upflows. In contrast to our results, they found upflows throughout the AR.  While our analysis technique will most readily identify a secondary component that is not excessively weaker than the primary, we do not believe we have missed the detection of a secondary component in the AR core. The primary intensity is a factor $\\sim10$ larger in the brightest areas of the AR core than in the region where we observe the outflow.  If the secondary component had the same intensity in the AR as found in the outflowing region,  our fitting technique would find a secondary component of this magnitude should it exist. However, we see no indication of a secondary component. Further, the primary component in the AR does not show the same blue-shift ($\\sim10$~km~s$^{-1}$) as the primary component in the outflowing region.  Should a  secondary component be present in the AR with similar velocity to that found in the outflowing region it would be easier to distinguish it from the primary  given the larger difference in velocities of the two components. In addition, a single Gaussian fit to the emission in the outflowing region results in excess line widths in comparison to areas of the observation that show no outflow.  The AR does not exhibit such excess line widths.  Finally, we note that a double Gaussian fit is a better representation than a single Gaussian only for the emission from the Fe~{\\sc xii} and {\\sc xiii} lines.  Emission from other ions is either too closely affected by nearby (in wavelength) emission lines or too weak in intensity to determine whether a double Gaussian fit is accurate.  The presence of high velocity outflowing plasma of a duration of several days leads to the question of where this material is deposited.  It does not appear to be confined to the loop arcades of the immediate AR; rather it is seen to move along either open or highly extended field lines.  This indicates significant mass flow into the corona and, in the case of open field lines, a possible contributor to the solar wind mass flow.  The velocity that we find for the slow Gaussian component compares well with models of the slow solar wind by \\citet{Wang94a} and \\citet{Wang09a}. These authors determined the outflow velocity from the slow solar wind originating from small coronal holes in the vicinity of active regions---closely matching the conditions of the observations analyzed in the present paper.  They predicted source region outflow velocities of 11.3~km~s$^{-1}$ \\citep{Wang94a} and 6.0~km~s$^{-1}$ \\citep{Wang09a}.  This is in good agreement with our results of $\\sim10$~km~s$^{-1}$ for the slow moving component of the outflow. Comparison of {\\it in situ} measurements of the solar wind are difficult because of large coronal holes both proceeding and preceding the active region in question. It is more likely that any detected solar wind originates in these coronal holes rather than the outflowing regions analyzed here.  At the 1~arcsec spatial scale of the EIS observations discussed in this paper, we are unable to spatially resolve the separate components of the outflowing material.  Thus, observations of a resolution significantly better than the 1\\arcsec\\ level will be required for future instrumentation  in order to  resolve these structures. Also, given the temperature dependence of these fast outflows, spectroscopic information on such a spatial scale is needed to fully determine the nature of such phenomena."
    },
    "1004/1004.2036_arXiv.txt": {
        "abstract": "{{\\small We analyze the recently published Fermi-LAT diffuse gamma-ray measurements in the context of leptonically annihilating or decaying dark matter (DM) with the aim to explain simultaneously the isotropic diffuse gamma-ray and  the PAMELA, Fermi and HESS (PFH) anomalous $e^\\pm$ data. Five different DM annihilation/decay channels $2e$, $2\\mu$, $2\\tau$, $4e$, or $4\\mu$ (the latter two via an intermediate light particle $\\phi$) are generated with PYTHIA. We calculate both the Galactic and extragalactic prompt and inverse Compton (IC) contributions to the resulting gamma-ray spectra. To find the Galactic IC spectra we use the interstellar radiation field model from the latest release of GALPROP. For the extragalactic signal we show that the amplitude of the prompt gamma-emission is very sensitive to the assumed model for the extragalactic background light. For our Galaxy we use the Einasto, NFW and cored isothermal DM density profiles and include the effects of DM substructure assuming a simple subhalo model. Our calculations show that for the annihilating DM the extragalactic gamma-ray signal can dominate only if rather extreme power-law concentration-mass relation $C(M)$ is used, while more realistic $C(M)$ relations make the extragalactic component comparable or subdominant to the Galactic signal. For the decaying DM the Galactic signal always exceeds the extragalactic one. In the case of  annihilating DM the PFH favored parameters can be ruled out by gamma-ray constraints only if power-law $C(M)$ relation is assumed.  For DM decaying into $2\\mu$ or $4\\mu$ the PFH favored DM parameters are not in conflict with the gamma-ray data.  We find that, due to the (almost) featureless Galactic IC spectrum and the DM halo substructure, annihilating DM may give a good simultaneous fit to the isotropic diffuse gamma-ray and to the PFH $e^\\pm$ data without being in clear conflict with the other Fermi-LAT gamma-ray measurements. }} ",
        "introduction": "During the last few years several experiments have shown an anomalous excesses in the cosmic electron and positron spectra. The PAMELA satellite has observed a steep rise of positron fraction $e^+/(e^-+e^+)$ at energies above 10 GeV with no significant excess in the cosmic antiproton flux \\cite{Adriani:2008zr,Adriani:2008zq}. The Fermi satellite and the HESS atmospheric Cherenkov telescope have measured an excess of high-energy $(e^-+e^+)$ flux with a cut-off of around 800 GeV \\cite{Abdo:2009zk,Aharonian:2009ah}. The ATIC and PPB-BETS balloon measurements indicate a similar excess \\cite{Chang:2008zzr,Torii:2008xu}. Most excitingly, the excess might originate from the annihilation or decay of the dark matter (DM) particles. The nature of those signatures requires the properties of DM to deviate strongly from the standard freeze-out predictions. The thermally averaged DM annihilation cross-section $\\cs$ has to be boosted some orders of magnitude over the standard freeze-out value  $\\cs_{\\rm std} \\simeq 3 \\times 10^{-26}$ cm$^3$s$^{-1}$, which might be achieved, e.g., through the Sommerfeld effect~\\cite{Hisano:2003ec} (see, e.g., \\cite{Cirelli:2008pk,ArkaniHamed:2008qn,Slatyer:2009vg} for the related phenomenological studies) or through the Breit-Wigner resonant enhancement \\cite{Feldman:2008xs,Ibe:2008ye,Guo:2009aj,Bi:2009uj}.  On the other hand, the decaying DM \\cite{Buchmuller:2007ui,Ibarra:2008jk,Nardi:2008ix,Arvanitaki:2008hq,Ibarra:2009dr} can explain the excess independent of the freeze-out constraints. In both cases, the annihilation or decay of DM should favorably occur only through the leptonic channels~\\cite{Cholis:2008hb,Cirelli:2008pk,Donato:2008jk,Cholis:2008qq,Cholis:2008wq}, as no excess in the hadronic channels has been observed. Alternatively, the excess of $e^+$ can potentially be explained by modifying or adding astrophysical sources, e.g. pulsars \\cite{Hooper:2008kg,Yuksel:2008rf,Profumo:2008ms,Malyshev:2009tw}.  The high energy leptons of the DM annihilation/decay are inevitably accompanied by the gamma-rays due to the final state radiation of charged leptons and decays of subproducts (``prompt gamma-rays'') and due to the upscattered background photons from the inverse Compton (IC) scattering (``IC gamma-rays''). Thus, the observed gamma-ray fluxes strongly constrain the above mentioned cosmic ray anomalies from DM annihilation/decay. The strongest gamma-ray constraints should arise from the observations of the Galactic center (GC) \\cite{Bergstrom:1997fj,Dodelson:2007gd}, as the density of DM is very high and it is relatively close to us. Those analyses take into account both the prompt and IC gamma-ray contributions \\cite{Bertone:2008xr,Bergstrom:2008ag,Bell:2008vx,Cirelli:2009vg,Meade:2009iu,Cholis:2009gv}. On the other hand, the GC is densely populated by different astrophysical objects, which contaminate the gamma-ray signal and introduce significant uncertainty in the derived constraints.  The Galactic gamma-ray signal of DM annihilation/decay at higher latitudes is considerably weaker than the signal from the GC. Despite being considerably weaker, the suppressed contamination by the Galactic astrophysical sources partially compensates the weakness of the signal. In addition, adding the Galactic DM substructure into the picture \\cite{Diemand:2007qr,Springel:2008by,Kuhlen:2009is,Kamionkowski:2010mi,Cline:2010ag} may significantly change both the magnitude and the morphology of the induced gamma-ray signal as the diffuse signal from DM subhalos can be essentially isotropic. The Galactic gamma-ray signal of DM annihilation/decay at higher latitudes can also include a considerable extragalactic contribution. It can originate from different sources: DM annihilation/decay in cosmological distances, active galactic nuclei (AGN), structure formation shocks, starburst galaxies, etc.  Fermi-LAT collaboration has also derived constraints on the annihilating DM properties from the Galactic dwarf spheroidal galaxies \\cite{Abdo:2010ex} (see \\cite{Essig:2009jx,PalomaresRuiz:2010pn} and references therein for the other recent studies). Non-detection of gamma-ray signal towards several nearby galaxy clusters has a potential to severely constrain the annihilating DM models \\cite{Pinzke:2009cp,Yuan:2010gn}. However, the present Fermi results \\cite{FermiDM:2010aa} still allow for the DM parameters giving the best fit to the PAMELA, Fermi, and HESS (PFH) $e^\\pm$ anomalies even if the smallest DM substructures are present. In addition to gamma-rays, the energetic leptons from DM decay/annihilation produce other accompanying signatures: synchrotron radiation \\cite{Bertone:2008xr,Bringmann:2009ca,Crocker:2010gy}, and neutrinos from $\\mu$ or $\\tau$ decay \\cite{Desai:2004pq,Covi:2008jy,Hisano:2008ah,Liu:2008ci,Mandal:2009yk}. The gamma-rays ionize and heat the intergalactic gas. The additional electrons released in the ionization process change the scattering optical depth of the CMB photons \\cite{Padmanabhan:2005es,Mapelli:2006ej,Zhang:2006fr,Belikov:2009qx,Galli:2009zc,Slatyer:2009yq,Huetsi:2009ex,Cirelli:2009bb,Kanzaki:2009hf,Yuan:2009xq,Iocco:2009ch,Natarajan:2010dc}. Also, the energetic $e^\\pm$ released in the annihilation/decay process induce the nonthermal component to the Sunyaev-Zel'dovich effect in galaxy clusters \\cite{Colafrancesco:2005ji,2009JCAP...10..013Y,2010JCAP...02..005L}. The energy injection from the DM sector modifies the accurately calculable\\footnote{See, e.g., \\cite{RubinoMartin:2009ry,Chluba:2010fy} for the latest developments.} standard cosmic recombination process (e.g. \\cite{Chen:2003gz,Padmanabhan:2005es,Chluba:2010aa}) and the Big Bang Nucleosynthesis (e.g. \\cite{Hisano:2008ti,Hisano:2009rc,Jedamzik:2009uy}).  As in, e.g. Regis \\& Ullio \\cite{Regis:2009md}, in this study we focus on the gamma-ray signal of DM annihilation/decay at higher Galactic latitudes. Fermi-LAT collaboration published recently their measurement of the diffuse gamma-ray emission at several Galactic regions and their estimation of the ``extragalactic diffuse gamma-ray background'' \\cite{collaboration:2010nz}. The published spectra are more constraining and have considerably smaller error bars than the older analogous spectra from the EGRET experiment \\cite{Sreekumar:1997un,Strong:2004ry}. The main aim of this paper is to find  out whether the new Fermi-LAT diffuse gamma-ray measurements can rule out or, instead, to support the PAMELA, Fermi, and HESS  favored models of DM annihilation/decay as given in \\cite{Meade:2009iu}\\footnote{For an earlier analysis, similar to \\cite{Meade:2009iu}, see \\cite{Bergstrom:2009fa}.}.  Our study is ``model independent'': we assume that DM particles annihilate/decay into Standard Model charged leptons, $\\ell^{-},\\ell^{+}$;  $\\ell=e,\\,\\mu,\\,\\tau,$ extending the analyses presented in series of works, e.g. \\cite{Cirelli:2008pk,ArkaniHamed:2008qn,Cirelli:2009vg,Meade:2009iu,Borriello:2009fa,Papucci:2009gd,Cirelli:2009dv,Strumia:2010zz}. We analyze five channels presented in Table~\\ref{tab:channels}. Our general formalism for the extragalactic gamma-rays is presented in our previous work \\cite{Huetsi:2009ex}. The new ingredient  in this paper is to include the effect of the absorption of extragalactic gamma-rays due to the pair production on extragalactic background light\\footnote{See \\cite{Ibarra:2009nw}, where the authors have included the background light while discussing the prospects for detecting the decaying DM with the Fermi-LAT.}.  For that we test different models for the extragalactic background light and show that the signal of prompt gamma-rays from the DM annihilation/decay can be significantly reduced in some cases.  For calculation of the Galactic IC spectra we use the interstellar radiation field (ISRF) model from the latest release of GALPROP \\cite{Porter:2005qx}. We show that inclusion of the realistic ISRF including CMB, infrared radiation from dust, and stellar light backgrounds are complicated enough to smear out any spectral feature of the individual components and the resulting spectrum is essentially power-law-like. To model the Galactic DM halo we use three density profiles: Einasto, NFW and cored isothermal. We first perform our analyses assuming that our Galaxy consist of one structureless DM halo with the given density profile. This is probably unrealistic but often used approximation.  After that we repeat our study using a more realistic model including DM subhalos. The technical details of our analyses are collected in Appendices A, B, C.   After Fermi collaboration has made their data publicly available, several papers have appeared that estimate the constraints on annihilating/decaying DM properties \\cite{Papucci:2009gd,Cirelli:2009dv,Chen:2009uq,Zhang:2009ut,Abazajian:2010sq,Abdo:2010dk}. The first two papers \\cite{Papucci:2009gd,Cirelli:2009dv} use the preliminary Fermi-LAT data. In both of these studies the authors have neglected the potential extragalactic contribution while deriving the constraints for the annihilating DM. In \\cite{Cirelli:2009dv} the approximate ISRF of our Galaxy is used. However, we show in this paper that the use of more realistic ISRF can significantly change the (line of sight) integrated IC spectra and in some cases the extragalactic diffuse gamma-ray signal may dominate over the Galactic one.  Abazajian et al. \\cite{Abazajian:2010sq} use the results of \\cite{collaboration:2010nz} but consider only the prompt gamma-ray spectra to constrain their DM models. This is clearly unsatisfactory because, as we show in agreement with \\cite{Abdo:2010dk}, in general the bounds are dictated by the IC contribution to the diffuse gamma-ray spectrum while the prompt contribution plays just a subdominant role. The two-peak structure of the diffuse gamma-ray spectrum in DM annihilations into $2\\mu$ channel presented in Fig. 4 of  Abdo et al. \\cite{Abdo:2010dk} also shows that those authors do not use realistic ISRF to calculate the IC spectrum but approximate it with the CMB component. As we show in this paper, this approximation does not qualitatively change the bounds derived in that paper but may qualitatively affect the whole interpretation of the Fermi-LAT experimental results. In addition,  both Abazajian et al. \\cite{Abazajian:2010sq} and Abdo et al. \\cite{Abdo:2010dk} study only  the annihilating DM. In this work, similarly to  Refs.~\\cite{PalomaresRuiz:2010pn,Boehm:2010qt}, we study the difference between annihilating and decaying DM gamma-ray signals.   The most important new result of this work is  that the inclusion of realistic IRSF and Galactic DM subhalo structure into the analyses opens up a new interpretation of the Fermi-LAT isotropic diffuse gamma-ray measurement.  We find that $(i)$ the leptonically annihilating DM models can give  good fits to the Fermi-LAT isotropic diffuse gamma-ray data if the extragalactic contribution is subdominant compared to the Galactic one; $(ii)$ the best-fit regions of the preferred DM mass and annihilation cross section has an overlap with the best-fit regions of the PFH anomaly  \\cite{Meade:2009iu} if the boost factor from the Galactic DM substructure is $B_{{\\rm sub}}\\sim {\\cal O}(10)$; $(iii)$ in the case of decaying DM the fits to the isotropic diffuse data are worse and there is no overlap with the best-fit regions of the PFH anomaly. This result implies that a significant fraction of the isotropic diffuse gamma-ray signal observed by Fermi-LAT may actually be of the Galactic origin. A generic issue that for a good fit the central regions of our Galaxy should give less prominent contribution to the DM annihilation signal than generally expected is (partially) solved with the DM halo substructure. The simultaneous  fitting of PFH electron/positron and Fermi-LAT diffuse gamma ray data gives a preference to the DM models with Sommerfeld enhancement due to a new light intermediate particle $\\phi$ \\cite{ArkaniHamed:2008qn} (see Table~\\ref{tab:channels}). If the intermediate particle $\\phi$ is long-lived \\cite{Rothstein:2009pm}, a scenario not considered in this work, their long diffusion length further smears any DM annihilation signal and helps to resolve potential conflicts with the diffuse gamma-ray data measurements from the central Galactic regions.   Our paper is organized as follows. In Section 2 we describe the energy and particle input from the DM annihilation/decay. In Section 3 we solve the radiative transfer equation for the Galaxy and the intergalactic medium. Section 4 presents the constraints as inferred from the Fermi-LAT diffuse gamma-ray data. Section 5 discusses the possibility of fitting the models to the Fermi-LAT isotropic diffuse data. Our summary and final discussion is given in Section 6. ",
        "conclusions": " \\begin{itemize}  \\item For the annihilating DM models the extragalactic signal can dominate only if power-law concentration-mass relation $C(M)$ is assumed. Once arguably more realistic Maccio et al. $C(M)$ relation is used the Galactic signal always exceeds the extragalactic one. Due to those reasons our calculated bounds on annihilating DM have significant dependence on the assumed extragalactic UV background model (which strongly effects the level of prompt emission) only in the case of the power-law $C(M)$ relation. In the case of decaying DM, where the signal scales proportionally to the DM density, and thus is not very sensitive to the details of the structure formation model, the Galactic signal is always stronger than the extragalactic one. As a result the calculated bounds in this case are rather insensitive to the choice of the extragalactic UV model.  \\item The derived constraints are very conservative, as the only thing we have required is that none of the models should exceed the measured diffuse signals in any of the observed energy bins. It is clear that in reality those diffuse signals can have various Galactic and extragalactic contributions other than the potential contribution from the annihilating/decaying DM. Our constraints for the annihilating and decaying DM models are given in Figs.~\\ref{fig4} and \\ref{fig6}, respectively. The red elliptical regions in those figures show the PFH $e^\\pm$ favored regions of Meade et al. \\cite{Meade:2009iu}. It turns out that for the power-law $C(M)$ relation the PFH favored regions are quite convincingly ruled out. On the other hand, if Maccio et al. $C(M)$ is used, we find that almost always the PFH favored regions easily survive except for the tau final states that produce too many prompt gamma-rays.  For the decaying DM only the $2\\tau$ channel is significantly constrained, while for the other two cases, i.e. $2\\mu$ and $4\\mu$, the exclusion zones are still below the PFH $e^\\pm$ ellipses.  \\item Once realistic models for the ISRF are used with all the relevant components: (i) stellar light, (ii) infrared radiation from dust, (iii) CMB, it turns out that the resulting gamma-ray spectra from the Galactic annihilating/decaying DM halo can be remarkably close to a seemingly featureless power law over the relevant energy range, as seen  in Figs.~\\ref{fig1} and \\ref{fig2}. Thus, there might be a possibility that those models provide an acceptable fit to the isotropic diffuse gamma-ray data if extragalactic contribution can be kept low, which is surely true for the Maccio et al. $C(M)$ relation. Indeed, this turns out to be the case: in Figs.~\\ref{fig4} and \\ref{fig6} the blue ellipses show the 2- and 1-sigma best-fit regions for the scenario neglecting Galactic DM halo substructure. Only the $\\tau$ lepton final states fail to give reasonable fits to the data due to hard prompt photons. However,  there is no overlap between the gamma-ray and PFH best-fit regions. Several best fitting spectra corresponding to this case are plotted with solid lines on the upper left-hand panels of Figs.~\\ref{fig7} and \\ref{fig8} for the annihilating and decaying DM, respectively. It is clear from Fig.~\\ref{fig7} that the best fitting annihilating models are in conflict with other measurements from other regions, especially with the ones from the $10<b<20$ Galactic region. However, in the case of the decaying DM there is no immediate conflict (i.e. the model spectra do not exceed the data points) with the other measurements, as seen from Fig.~\\ref{fig8}.  \\item All the above results did not include the possible contribution from the Galactic DM substructure. We have devised a simple model for the halo substructure (see Appendix~\\ref{appc}), where all the specific model details can be absorbed into a single effective substructure density parameter $\\rho_{{\\rm sub}}^{{\\rm eff}}$. The inclusion of substructure has a twofold effect: (i) it leads to the increase of the total amplitude of gamma-ray emissivity, (ii) it makes the spatial emission profiles shallower compared to the smooth halo case. Due to those reasons the inclusion of substructure lowers the blue best-fit regions of Fig.~\\ref{fig4}, and at the same time helps to reduce the conflict with the measurements of other Galactic regions seen on the lower panels of Fig.~\\ref{fig7}. The green best-fit regions of Fig.~\\ref{fig4} and the dashed lines of Fig.~\\ref{fig7} correspond to the case with $\\rho_{{\\rm sub}}^{{\\rm eff}}=\\rho_{\\odot}=0.4$ GeV cm$^{-3}$, which in loose terms represents the case with substructure boost factor $\\simeq 15$. Those levels of substructure boosts are indeed possible, as discussed, e.g., in Martinez et al. \\cite{Martinez:2009jh}. If we give up our simplifying assumption that substructure follows the density profile of the main halo, and allow an increased subhalo destruction rate in the center, it should be possible to make the spatial emission profiles even shallower, and thus potentially remove the remaining small conflict on the lower right-hand panel of Fig.~\\ref{fig7}. The emission profiles can also be made shallower if one allows the intermediate particle $\\phi$ (see Section~\\ref{sec1}) to have a sufficiently long lifetime, so that it can move significantly away from the initial place of birth before decaying, which would effectively reduce the steepness of the resulting gamma-ray profile.  It is clear that diffuse gamma-ray background has other contributions. In this work we consider the standard Galactic background model $44\\_500180$ from Strong et al. \\cite{Strong:2004de}. In Figs. ~\\ref{fig9} and \\ref{fig10} we have added this background to the DM-induced signal. It is quite remarkable that the models, which provide the best fits to the isotropic diffuse data, automatically get quite close to the observed data points in other Galactic regions, once the standard background model of Strong et al. \\cite{Strong:2004de} is included. It is evident that by a relatively modest tuning of the Galactic background model satisfactory fits in all of the regions should be possible.  We conclude that it seems quite realistic to have annihilating DM models, which simultaneously: (i) are compatible with the PFH $e^\\pm$ data, and (ii) provide acceptable fits to the Fermi-LAT isotropic diffuse background, without being in clear conflict with the other Fermi-LAT gamma-ray measurements. The preferred DM annihilation final states consistent with this analysis and the PFH fits are $2\\mu,$ $4\\mu$ and $4e$.  \\item In our Galactic IC calculations we have neglected the effect of electron diffusion, i.e. in our treatment the IC spectrum is produced ``on spot''. For more precise treatment one should solve the diffusion-loss equation for electrons to determine the electron energy distribution, which is used as an input in IC calculations. Some estimates of Papucci \\& Strumia \\cite{Papucci:2009gd} have shown that by including this diffusion effect the IC signal gets reduced by a factor of around $2-3$ depending on the Galactic region. However, the full calculation of electron diffusion is quite complicated with rather large uncertainties due to, e.g., the knowledge of the magnetic fields. Even in our simple treatment there is a large uncertainty due to the assumed ISRF model. We have used the latest ISRF model from the GALPROP code, which is based on results from Porter \\& Strong \\cite{Porter:2005qx}. However, in their study the authors do not provide any errors for the estimated ISRF. It is clear that their results are highly model dependent, as from the direct measurements we can only get the ISRF in our location, while the values at other Galactic locations can only be obtained by solving the radiative transfer problem, which clearly involves several model-dependent assumptions. According to Strong et al. \\cite{Strong:2004de} a factor of two uncertainty for the ISRF is quite possible. Due to those reasons we feel that the accurate treatment for the $e^\\pm$ diffusion problem can probably wait before there are better models with realistic uncertainties available for the ISRF. Unfortunately, except for rare examples of \\cite{Porter:2005qx} and \\cite{Strong:1998fr}, the problem of determining the ISRF as a function of Galactic position has not gained much attention so far.  \\item In Introduction we referred to the other constraints on annihilating or decaying DM. Compared to the CMB constraints (see e.g. \\cite{Slatyer:2009yq,Huetsi:2009ex,Cirelli:2009bb}) the bounds derived in this paper can be stronger or weaker, depending on what type of concentration-mass relation $C(M)$ one assumes. If Maccio et al. \\cite{2008MNRAS.391.1940M} $C(M)$ is used the bounds are generally weaker, while the power-law $C(M)$ leads to stronger constraints than obtainable from the CMB measurements. The CMB constraints, being mostly determined by the free electron fraction at redshifts $100 \\lesssim z \\lesssim 1000$ \\cite{Huetsi:2009ex}, are quite insensitive to the uncertainties of the nonlinear structure formation and other complicated low-redshift astrophysical phenomena. In that respect they can be considered significantly more robust than the diffuse gamma-ray bounds derived in this paper.  The constraints from the Galactic center radio and gamma-ray measurements are generally stronger than ours, however those are hampered by the obscurity of the DM density profile, ISRF and magnetic field in the very central region (see, e.g. \\cite{Bergstrom:1997fj,Dodelson:2007gd,Bertone:2008xr,Bringmann:2009ca,Crocker:2010gy}). For example, there is no clear understanding on how much various physical processes involving baryonic component can impact on the the central cusp of the DM profile (see e.g. \\cite{deBlok:2009sp} for a recent review).  The constraints derived from the Galactic dwarf spheroidal galaxies and the nearby galaxy clusters are comparable or slightly weaker than our restrictions \\cite{Essig:2009jx,Abdo:2010ex,FermiDM:2010aa}.   \\end{itemize}  \\subsection*{Note added} When the research presented in this work was completed, Ref.~\\cite{Lin:2010fb} appeared that also performs fits of the cosmic ray and gamma-ray data. Although their approach and aim of the paper is different form ours, the main conclusions that only DM annihilations can potentially provide a satisfactory fit agree with each other."
    },
    "1004/1004.1203_arXiv.txt": {
        "abstract": "We present a calculation of the atomic and low-ionisation emission line spectra of photoevaporating protoplanetary discs. Line luminosities and profiles are obtained from detailed photoionisation calculations of the disc and wind structures surrounding young active solar-type stars. The disc and wind density and velocity fields were obtained from the recently developed radiation-hydrodynamic models of Owen et al., that include stellar X-ray and EUV irradiation of protoplanetary discs at various stages of clearing, from primordial sources to inner hole sources of various hole sizes.  Our models compare favourably with currently available observations, lending support to an X-ray driven photoevaporation model for disc dispersal. In particular, we find that X-rays drive a warm, predominantly neutral flow where the OI~6300{\\AA} line can be produced by neutral hydrogen collisional excitation. Our models can, for the first time, provide a very good match to both luminosities and profiles of the low-velocity component of the OI~6300{\\AA} line and other forbidden lines observed by Hartigan et al., which covered a large sample of T-Tauri stars.  We find that the OI~6300{\\AA} and the NeII~12.8$\\mu$m lines are predominantly produced in the X-ray-driven wind and thus appear blue-shifted by a few km/s for some of the systems when observed at non-edge-on inclinations. We note however that blue-shifts are only produced under certain conditions: X-ray luminosity, spectral shape and inner hole size all affect the location of the emitting region and the physical conditions in the wind. We caution therefore that while a blueshifted line is a tell-tale sign of an outflow, the lack of a blueshift should not be necessarily interpreted as a lack of outflow. Comparison of spectrally resolved observations of multiple emission lines with detailed model sets, like the one presented here, should provide useful diagnostics of the clearing of gaseous discs. ",
        "introduction": "The circumstellar environments of newly-born low to intermediate mass stars have recently been the object of much interest for their potential of hosting planetary systems. These (protoplanetary) discs of dust and gas hold the reservoir of material from which planets may form and provide the medium through which newly formed planets and planetesimals migrate towards or away from their parent star. For this reason no theory of planet formation can be complete without an understanding of how protoplanetary discs evolve and finally disperse.  The dispersal of protoplanetary discs has been studied observationally, through statistics of disc fractions in young clusters (e.g. Haisch et al 2001; Hernandez et al 2008; Mamajek 2009) and on an object-by-object basis through observations of individual so-called transition discs (e.g. Calvet et al 2002), i.e. discs that appear to be optically thin in the dust in their inner regions, which is generally interpreted as evidence of partial (inside-out) clearing. Whether these discs with inner-holes really represent a stage in the evolution/dispersal of protoplanetary discs is currently still uncertain. Currently, the consensus in the literature is that the dispersal of discs around T-Tauri stars is characterised by a rather fast transition timescale, evidenced by the fact that very few objects (typically 10\\%) are caught in this transition phase (Strom et al. 1989; Strutskie et al. 1990; Ercolano, Clarke \\& Robitaille 2009; Kim et al 2009; Luhman et al 2010; Muzerolle et al 2010)\\footnote{There have been claims of longer transition timescales for M-stars or older clusters (Sicilia-Aguilar et al. 2008; Currie et al 2009), these results, however, are still object of debate in the literature (Ercolano et al 2009, Muzerolle et al. 2009, Luhman et al 2010)}.  On the basis of these observational results, theoretical efforts have thus focused on identifying a mechanism to provide a two-timescale dispersal phenomenon, i.e. a mechanism that would allow protoplanetary discs to be detected as optically thick (in the dust) for a few million years and then cause them to disperse very rapidly in about a tenth of their total lifetime, i.e. a timescale much shorter than what would be expected by viscous evolution alone (Hartmann et al. 1998, Clarke et al. 2001).  The proposed mechanisms include planet formation itself (e.g. Armitage \\& Hansen 1999), grain growth (Dullemond \\& Dominik 2005), photophoresis (Krauss et al 2007), MRI-driven winds (Suzuki \\& Inutsuka 2009), and photoevaporation (Clarke et al 2001; Alexander et al 2006, Richling \\& Yorke 2000, Ercolano et al 2008b, 2009a, Gorti \\& Hollenbach 2009; Owen et al. 2010a). Understanding which (or what combination) of the proposed mechanisms dominates dispersal is key to understanding disc evolution and thus provides important constraints on planet formation theories by setting the timescale over which planets may form and migrate.  Photoevaporation by soft X-ray radiation (100~eV $<$ E $<$ 1keV) has been shown to drive powerful winds (Ercolano et al. 2008b, 2009a, Gorti \\& Hollenbach 2009), able to disperse the discs % rapidly once the accretion rates have fallen to values lower than the photoevaporation rates ($\\sim$10$^{-8}$M$_{\\odot}$/yr for X-ray luminosities of 2$\\times$10$^{30}$erg/s; Owen et al. 2010a; from now on O10). The % mass-loss rates are roughly proportional to the impinging X-ray  % luminosity (Ercolano 2008b), meaning that the expected mass-loss rates (and hence accretion % rates) for solar type stars span over the two-three decade range of X-ray% luminosities(e.g. Preibisch et al. 2005; G\\\"udel et al. 2007).% The recent observation of shorter disc lifetimes in low-metallicity environments (Yasui et al 2009) has been interpreted as further evidence for X-ray photoevaporation acting as the main disc dispersal mechanism in those clusters (Ercolano \\& Clarke 2009).  While encouraging, the evidence in favour of photoevaporation produced so far consists of statistics derived from observations % of disc dispersal based on the evolution of their dust component. However, dust represents only a percent or so of the total disc mass and the evolution of the dust and gas component of a disc do not necessarily go hand in hand (e.g. Alexander \\& Armitage 2007; Takeuchi, Clarke \\& Lin 2004). In the last few years, observations of % emission lines from warm atomic and low ionisation gas in the inner region of discs are starting to provide some helpful clues of the physical conditions of the gaseous component in these planet-forming regions (Hartigan et al 1995;Herczeg et al 2007; Pascucci et al 2007; Razka et al 2007; Espaillat et al 1995; Najita et al 2009; Pascucci \\& Sterzik 2009; Flaccomio et al 2009; G\\\"udel et al 2010, submitted). Line emission from low ionisation and atomic species is of particular interest, as it is likely to trace material that is located in the warm upper layers of the disc atmosphere or in the flow itself, allowing the response of the gas to irradiation from the central star to be probed and understood. In particular Pascucci \\& Sterzik (2009) reported evidence of a blue-shifted NeII fine structure line at 12.8$\\mu$m, the blue-shift is consistent with the line being formed in a slow moving (a few km/s) disc wind (Alexander 2008). Earlier spectrally resolved observations of optical forbidden lines in T-Tauri stars (Hartigan et al., 1995, from now on HEG95), including OI 6300~{\\AA}, also revealed blue-shifted low velocity components, interpreted as evidence for a disc wind. The interpretation of these atomic and low-ionisation emission line profiles is key to understand the kinematic structure of the warm upper disc layers and its wind. In this paper we provide the tools to exploit currently available and future observations using the framework of the X-ray driven photoevaporation model of O10 (further discussed in Section~\\ref{s:m}).  Furthermore, with the {\\it Herschel Space Telescope} and, later, {\\it James Webb Space Telescope}, committed to studying planet formation through the characterisation of discs, it is both timely and important to provide a solid theoretical framework for the interpretation of these new% observations. In this context this paper has two main aims: (i) To % provide an atlas of atomic and low ionisation emission line intensities and profiles predicted for the photoevaporating protoplanetary discs of O10 at various stages in their evolution; (ii) To compare the % predictions of O10's X-ray driven photoevaporation model with currently available observations and search for the `smoking gun' of photoevaporation. We therefore list line intensities for the brightest lines and calculate line profiles for some of the more promising wind/disc tracers. We show results for discs at different stages of clearing,  i.e. photoevaporating primordial discs where an inner hole is yet to form and systems in a more advanced clearing stage with inner holes (in dust and gas) of increasing radius. We note that, while several authors have presented theoretical calculations of emission line luminosities and profiles from T-Tauri discs (e.g. Glassgold et al 2007, Meijerink et al. 2008, Ercolano et al 2008b, Gorti \\& Hollenbach 2008, Hollenbach \\& Gorti 2009, Schisano, Ercolano \\& G\\\"udel 2010), this is the first time that a fully self-consistent radiation-hydrodynamical and photoionisation calculation is attempted, which provides us with the three-dimensional distribution of gas velocities and emissivities in the heated layers of the disc and its wind.  A comparison of the calculations presented in this paper with available observations suggests that X-ray heated discs with hydrodynamically escaping winds produce emission line intensities and profiles that are consistent with those that are currently being observed. It is hoped that the results presented here may be used for the interpretation of future observational data which will be crucial to further test the validity of the X-ray photoevaporation scenario as a dominant disc clearing mechanism.  The paper is organised as followed. Section~2 briefly describes the radiation-hydrodynamic disc models and lists the input parameters. Sections~3 presents the results, including the emission line intensities and the predicted profiles for selected transitions. Sections~4 and~5 deal with a comparison of our results with currently available data. The main points of the paper are briefly summarised in Section~6. ",
        "conclusions": "X-rays have recently been shown to drive powerful outflows (Ercolano et al. 2008b, 2009a, Owen et al 2010a). The recent radiation-hydrodynamic calculation of Owen et al. (2010a) have demonstrated that that the two time-scale dispersal of protoplanetary discs around solar-type stars is successfully reproduced by this model. Further evidence is provided by the observation of shorter disc lifetimes in low-metallicity environments (Yasui et al. 2009), a behaviour that is expected if X-ray driven photoevaporation dominates the dispersal of discs and argues against planet formation as the main dispersal mechanism (Ercolano \\&  Clarke, 2009). Furthermore the X-ray spectra of T-Tauri stars are fairly well characterised (e.g. G\\\"udel et al 2007) and the observed luminosities imply that significant X-ray photoevaporation must occur, unless the inner source is completely obscured due to accretion funnels thicker than 10$^{22}$cm$^{-2}$ (Ercolano et al. 2009a) and covering all lines of sight from the source to the disc.  Encouraged by these results, in this paper we have presented an atlas of emission lines from atomic and low-ionisation species emitted by photoevaporating protoplanetary discs around solar type stars. We have used disc and wind density distributions obtained by the radiation-hydrodynamic calculations of Owen et al (2010a) in the X-ray photoevaporation paradigm and have considered discs irradiated by a range of X-ray luminosities and at various stages of clearing, from primordial (i.e. extending all the way into the dust destruction radius) to inner hole sources with various hole sizes. Line profiles for the NeII~12.8$\\mu$m, OI~6300{\\AA} and HI recombination lines have also been calculated and should provide guidance for the interpretation of high-resolution spectroscopy of YSOs.  We compare the predictions of our models to available data in the literature and find agreement for the line intensities and profiles. We draw particular attention to the good agreement with the observations of Hartigan et al (1995) of the OI~6300{\\AA} line intensity and profile in the spectra of T-Tauri stars. This had been a problem for previous, EUV-driven photoevaporation models (e.g. Font et al. 2004) which underestimated its luminosity by a couple of orders of magnitude. The problem with these previous models lay in the fact that EUV winds are, by construction, ionised, while X-ray photoevaporation provides a mechanism to drive warm neutral winds, a condition necessary for the production of blue-shifted collisionally excited OI~6300{\\AA} and OI~5577{\\AA}.  As noted by Font et al (2004), magnetic wind models systematically underpredict the luminosities of forbidden lines and cannot fit the observed profiles (e.g. Garcia et al 2001). The main problem with magnetic wind models is that they are too cold -- an X-ray driven photoevaporative wind naturally produces temperature, ionisation degrees and consistent kinematic patterns able to produce forbidden line intensities and profiles in agreement with the observations. These results suggest that, while magnetic disc winds may work in conjunction with photoevaporation, they are unlikely to be operating alone.   We analyse our results for the NeII~12.8$\\mu$m line profiles in the context of recent high resolution observations of Pascucci et al (2009) and Najita et al. (2009). We discuss how the detection of a blue-shifted line is a clear sign of an outflow, but caution against interpreting non-blue-shifted profiles or non-detections as absence of a disc wind. We argue that the non-detections for the primordial discs in the Pascucci et al (2009) sample could be explained by insufficient sensitivity of the observations that are only able to probe the upper end of the NeII luminosity distribution. The non-blue-shifted line centroids in the Najita et al. (2009) sources are probably due to unfavourable (high inclination) viewing angles or to the effect of seeing blue and redshifted emitting material through a large inner hole, as may be the case of GM~Aur, although we caution that this explanation is only valid if the gas, known to be present in the inner (dust) hole of this object, is optically thin to X-ray radiation and to the line in question.  Finally, our models cannot reproduce the observed emission line luminosities of IR hydrogen recombination lines, which for some objects are reported to be as bright as the NeII~12.8$\\mu$m (Pascucci et al 2007; Ratzka et al 2007). As suggested by Hollenbach \\& Gorti (2009), these lines are probably produced in a dense, ionised region that is very close to the central star. High resolution spectroscopy of these lines is needed to further constrain their origin."
    },
    "1004/1004.3801_arXiv.txt": {
        "abstract": " ",
        "introduction": "A central challenge of inflation is its sensitivity to ultraviolet (UV) physics. Successful inflation requires that the inflaton mass is much smaller than the Hubble scale, {\\it i.e.}~that the parameter \\beq \\eta = M_{\\rm pl}^2 \\frac{V''}{V} \\simeq \\frac{m_\\phi^2}{3H^2} \\eeq is much less than unity. The eta problem refers to the fact that UV corrections tend to drive $\\eta$ to large values. Generically, even Planck-suppressed corrections to the inflaton potential can be important and have to be understood. This is the main motivation for pursuing a realization of inflation in a UV-complete framework like string theory. In that case, the relevant contributions can either be computed directly, or their absence can be given a natural explanation (see \\cite{McAllister:2007bg, Burgess:2007pz, Baumann:2008aq, Baumann:2009ni} for recent reviews of inflation in string theory). In this paper we instead propose a mechanism that solves the eta problem purely in effective field theory. By coupling the inflaton field to a conformal sector with suitable properties, we show that renormalization group (RG) flow naturally suppresses all dangerous UV corrections to the inflaton self-interactions.\\footnote{The effects of RG flow on inflation and its observable signatures has been previously studied (see {\\it e.g.}~\\cite{Stewart:1996ey,Stewart:1997wg, Lyth:1998xn}).  However, to our knowledge, the role of RG effects in solving the eta problem was not studied systematically.} We consider this a dynamical, low-energy solution to the eta problem.  \\vskip 4pt If the inflaton field respects an approximate shift symmetry in the infrared (IR)\\footnote{By the IR we here mean a scale close to the inflationary energy scale (see Fig.~1).} then the eta problem is significantly ameliorated~\\cite{Freese:1990rb}. However, even in that case, one has to worry that higher-dimension operators in the UV break the shift symmetry and lead to large corrections.  Without knowledge of the UV-completion, one should not assume that irrelevant operators respect the shift symmetry.  Moreover, if the shift symmetry has an associated conserved charge, then black hole evaporation suggests that the UV-completion should indeed break the symmetry \\cite{Kallosh:1995hi,Kamionkowski:1992mf, Holman:1992us, Kamionkowski:1992ax}.\\footnote{A famous consequence of the black hole no-hair theorem is the fact that the global charge of a black hole is not defined. Consider therefore a scattering process which involves an initial state with definite global charge and an intermediate state of virtual black holes. The black holes evaporate into a a final state with different global charge indicating that the global symmetry is broken. At sub-Planckian energies the symmetry breaking is described by higher-dimensiuon operators in the effective theory. Standard Wilsonian thinking then suggests that the higher-dimension operators are suppressed by appropriate powers of the Planck mass, and that all dimensionless couplings are of order one.}  It has been suggested that these arguments can be avoided if the global symmetry is gauged in the UV.  This possibility has been realized in string theory \\cite{Eva, Liam, Berg:2009tg} and in extra-dimensional constructions \\cite{ArkaniHamed:2003mz, Kaplan:2003aj}, but its validity ultimately still requires knowledge of the UV-completion. Here we will take a different approach: we will show how to control shift symmetry breaking UV operators. Our models will have an approximate symmetry in the IR that is badly broken by non-renormalizable corrections in the UV, but RG flow naturally suppresses the couplings of all dangerous operators. In this way, a rather generic effective field theory in the UV flows to a theory with small eta parameter in the IR.  Similar ideas have been put forward in particle physics to explain hierarchies that are much more severe than the hierarchy between the inflaton mass and the Hubble scale during inflation (typically, particle physicists use RG flow to explain hierarchies as large as $10^{-32}$~\\cite{Susskind:1978ms, Weinberg:1975gm, Witten:1981nf} and no smaller than $10^{-3}$~\\cite{Nelson:2000sn}). For instance, in models of supersymmetry breaking, Planck-suppressed operators can give rise to flavor-violating soft masses $m^2_{\\rm FV}$. Unless generic UV contributions are suppressed, the typical size of $m_{\\rm FV}^2$ is inconsistent with experiment.  In `conformal sequestering'~\\cite{Luty:2001jh,Luty:2001zv} flavor-blind supersymmetry breaking is achieved via a mechanism that is completely analogous to the mechanism that we propose for inflation in this paper. By coupling the supersymmetry breaking fields to a conformal sector, operators transmitting supersymmetry breaking can receive large anomalous dimensions and the flavor-violating masses ($m^2_{\\rm FV}$) are suppressed relative to the flavor-diagonal soft masses ($m^2_{\\rm FD}$). However, in that case, the suppression has to be by a very large factor, $\\frac{m^2_{\\rm FV}}{m^2_{\\rm FD}} < 10^{-6}$, while for inflation the required effect is rather mild, $\\eta = \\frac{m^2_\\phi}{3H^2} \\lesssim 10^{-2}$. This lets us hope that our mechanism can operate at weak coupling using only perturbative physics, while particle phenomenology is driven to the strong coupling regime.  Our basic idea is very simple and we believe it to apply to a wide range of inflationary mechanisms. For concreteness, however, we will study a specific supersymmetric model in detail. In that case there will be dangerous dimension-five and dimension-six operators in the K\\\"ahler potential which we control be a combination of a discrete $\\mathbb{Z}_2$ symmetry and conformal sequestering. Our main motivation for writing down the specific model in \\S\\ref{sec:susy} is to give an `existence proof' of the sequestering mechanism, while allowing for exact computations. However, since the basic idea of our paper is very intuitive, we believe that simpler (but maybe less exact or controlled) examples should exist.   \\vskip 4pt  The plan of the paper is as follow: In \\S\\ref{sec:outline} we review the eta problem and explain the basic mechanism by which we aim to suppress dangerous operators. Next, in \\S\\ref{sec:susy}, we introduce a concrete supergravity model in which i) the problem is apparent and ii) the solution can be discussed in detail. We compute the anomalous dimension of the inflaton field induced by coupling the theory to $SU(N_c)$ Yang-Mills with $3 N_c >  N_f > 3 N_c /2$ flavors. We show that our idea is safely in the perturbative regime. In \\S\\ref{sec:linear} we illustrate that the idea isn't guaranteed to work by explaining its failure for a theory with linear superpotential. We summarize the necessary conditions that are required for a successful decoupling of dangerous operators. In \\S\\ref{sec:nonSUSY} we extend our considerations to non-supersymmetric models, where we discuss the renormalization of curvature couplings in detail. In \\S\\ref{sec:gauged} we contrast our models with models involving gauged symmetries in the UV. We make some concluding remarks in \\S\\ref{sec:conclusion}.   \\newpage ",
        "conclusions": "\\label{sec:conclusion}  It is rare that low-energy physics depends sensitively on Planck-suppressed contributions. Inflation is one of the few examples where understanding these corrections to the action is absolutely essential. It is important to realize that the eta problem is independent of the energy scale of inflation and is equally severe for high-scale and low-scale models. In this paper we have presented a new solution to this problem.  While most solutions to the eta problem assume low-energy symmetries and the absence of symmetry breaking operators in the UV, we have shown that appropriate couplings of the inflaton to a conformal sector allow control over these corrections in effective field theory. This has allowed us to relax some of the commonly made assumptions about the UV structure of the theory. We have presented explicit examples to illustrate how conformal sequestering can suppress the effects of shift symmetry violating terms in the UV. The low-energy theory then remains approximately shift symmetric and has a small eta parameter even though the shift symmetry is badly broken in the UV.  We summarize what we have learned as a guide for future applications of our idea. % The theory has to contain the following elements to allow a successful decoupling of higher-dimension corrections to the inflationary action:  \\begin{enumerate} \\item {\\it Symmetries of the renormalizable action}  In the basis where the inflaton has canonical kinetic term, the potential may be split into a renormalizable part, $V_0(\\phi)$, and non-renormalizable corrections, $\\delta V(\\phi)$, \\beq V(\\phi) = V_0(\\phi) + \\delta V(\\phi)\\, . \\nonumber \\eeq Technical naturalness and radiative stability require that the renormalizable action has certain symmetries. Different models will achieve these desirable features in different ways. In this paper we used a combination of supersymmetry and an approximate shift symmetry to protect the renormalizable part of the potential.   \\item {\\it Symmetries of the coupling to the CFT}  To prevent generating dangerous operators via the coupling to the CFT itself, we require that the couplings to the conformal sector respect the same symmetries as the renormalizable action. This ensures that the couplings of the dangerous operators flow to zero and not to some finite value at the fixed point.   \\item {\\it UV corrections}  In the non-renormalizable part of the potential, $\\delta V(\\phi)$, we allow arbitrary breaking of the symmetries of the renormalizable part of the potential. RG flow will suppress the couplings of these higher-dimension operators, so that the full action in the IR has the (accidental) symmetries of the renormalizable potential.  \\end{enumerate}  We have shown in a variety of examples that these requirements can be fullfilled in a technically natural way. In our most explicit example, in \\S\\ref{sec:susy}, we computed the anomalous dimension of the inflaton---exactly via $a$-maximization and at one loop---and showed that our mechanism involves only weakly-coupled physics. Going to stronger coupling is likely to increase the efficiency of sequestering, but reduces the control over the field theory computations. It might be interesting to study this regime in the gravity dual \\cite{AdSCFT}.  \\vskip 4pt Finally, we would like to be clear that our work is not meant to be read as claiming that understanding the UV-completion of inflationary models is not important ({\\it cf.}~\\S\\ref{sec:comments}).  Our goal has been to explore the possibility of solving the eta problem while being agnostic about the effects of a UV-completion. In small-field models, we believe that this is possible through RG flow in the low-energy effective theory, provided approximate continuous  symmetries in the IR and a discrete symmetry in the UV.  However, even in this context, it would be very useful to understand the origin of approximate symmetries within a UV-complete framework.  It may be the case that engineering these approximate symmetries requires special features that ultimately suppress dangerous Planck-suppressed contributions to the potential.  Nevertheless, the K\\\"ahler corrections in our models are controlled in field theory and one only must only explain the origin of the superpotential.  For this reason, one could hope to build a model in string theory using only topological information (see, for example, \\cite{Blumenhagen:2005mu, Douglas:2006es, Blumenhagen:2006ci, Denef:2008wq}). Furthermore, we have exhibited classes of inflationary models for which UV corrections cannot be decoupled by RG flow. For these models understanding the UV-completion is essential.    \\subsubsection*{Acknowledgements}  We are grateful to Nima Arkani-Hamed, Nathaniel Craig, Anatoly Dymarsky, Jonathan Heckman, Liam McAllister, Michele Papucci, Soo-Jong Rey, Leonardo Senatore, Eva Silverstein, Matt Sudano, Tomer Volansky, and Brian Wecht for discussions. We thank Liam McAllister for extremely helpful comments on a draft. D.B.~wishes to express special thanks to Anatoly Dymarsky, Shamit Kachru, Igor Klebanov and Liam McAllister for collaboration on related questions. D.B.~thanks the Mitchell Institute for Fundamental Physics and Astronomy at Texas A\\&M for hospitality and the opportunity to present this work. D.G.~thanks the Kavli Institute for Theoretical Physics for hospitality while this work was completed. The research of D.B.~is supported by the National Science Foundation under PHY-0855425, AST-0506556 and AST-0907969. The research of D.G.~is supported by the Department of Energy under grant number DE-FG02-90ER40542.   \\newpage"
    },
    "1004/1004.3340_arXiv.txt": {
        "abstract": "In this paper, the holographic dark energy model with new infrared (IR) cut-off for both the flat case and the non-flat case are confronted with the combined constraints of current cosmological observations: type Ia Supernovae, Baryon Acoustic Oscillations, current Cosmic Microwave Background, and the observational hubble data. By utilizing the Markov Chain Monte Carlo (MCMC) method, we obtain the best fit values of the parameters with $1\\sigma, 2\\sigma$ errors in the flat model: $\\Omega_{b}h^2=0.0233^{+0.0009 +0.0013}_{-0.0009 -0.0014}$, $\\alpha=0.8502^{+0.0984 +0.1299}_{-0.0875 -0.1064}$, $\\beta=0.4817^{+0.0842 +0.1176}_{-0.0773 -0.0955}$, $\\Omega_{de0}=0.7287^{+0.0296 +0.0432}_{-0.0294 -0.0429}$, $\\Omega_{m0}=0.2713^{+0.0294 +0.0429}_{-0.0296 -0.0432}$, $H_0=66.35^{+2.38 +3.35}_{-2.14 -3.07}$. In the non-flat model, the constraint results are found in $1\\sigma, 2\\sigma$ regions: $\\Omega_{b}h^2=0.0228^{+0.0010 +0.0014}_{-0.0010 -0.0014}$, $\\Omega_k=0.0305^{+0.0092 +0.0140}_{-0.0134 -0.0176}$, $\\alpha=0.8824^{+0.2180 +0.2213}_{-0.1163 -0.1378}$, $\\beta=0.5016^{+0.0973 +0.1247}_{-0.0871 -0.1102}$, $\\Omega_{de0}=0.6934^{+0.0364 +0.0495}_{-0.0304 -0.0413}$, $\\Omega_{m0}=0.2762^{+0.0278 +0.0402}_{-0.0320 -0.0412}$, $H_0=70.20^{+3.03 +3.58}_{-3.17 -4.00}$. In the best fit holographic dark energy models, the equation of state of dark energy and the deceleration parameter at present are characterized by $w_{de0}=-1.1414\\pm0.0608, q_0=-0.7476\\pm0.0466$ (flat case) and $w_{de0}=-1.0653\\pm0.0661, q_0=-0.6231\\pm0.0569$ (non-flat case). Compared to the $\\Lambda \\textmd{CDM}$ model, it is found the current combined datasets do not favor the holographic dark energy model over the $\\Lambda \\textmd{CDM}$ model. ",
        "introduction": "Since 1998, the type Ia supernova (SNe Ia) observations \\cite{ref:Riess98,ref:Perlmuter99} have shown that our universe has entered into a phase of accelerating expansion. During these years from that time, many additional observational results, including current Cosmic Microwave Background (CMB) anisotropy measurement from Wilkinson Microwave Anisotropy Probe (WMAP)\\cite{ref:Spergel03,ref:Spergel06}, and the data of the Large Scale Structure (LSS) from Sloan Digital Sky Survey (SDSS) \\cite{ref:Tegmark1,ref:Tegmark2}, also strongly support this suggestion. These observational results have greatly inspirited theorists to understand the mechanism of the accelerating expansion of the universe, which is usually attributed to an exotic energy component with negative pressure, dubbed dark energy (DE). The simplest but most natural candidate of DE is the cosmological constant $\\Lambda$, with the constant equation of state (EOS) $w=-1$. As we know, the cosmic concordance model confronts with two difficulties: the fine-tuning problem and the cosmic coincidence problem. Both of these problems are related to the DE density. In order to solve or alleviate cosmological constant puzzles, many dynamical DE models are proposed, where the DE density and its EOS are time-varying. However, the predictions of the cosmological constant model still fit to the current observations \\cite{ref:LCDM1,ref:LCDM2,ref:LCDM3}. Therefore the dynamical DE models being proposed should not be far away from the cosmological constant model, such as quintessence \\cite{ref:quintessence01,ref:quintessence02,ref:quintessence1,ref:quintessence2,ref:quintessence3,ref:quintessence4}, phantom \\cite{ref:phantom}, quintom \\cite{ref:quintom}, K-essence \\cite{ref:kessence}, tachyon \\cite{ref:tachyon}, ghost condensate \\cite{ref:ggc}, holographic DE \\cite{ref:holo1,ref:holo2} and agegraphic DE \\cite{ref:age1,ref:age2} etc. Although many DE models have been presented, the nature of DE is still a conundrum. This puzzle can not be understood before a complete theory of quantum gravity is established. But the two additional aspects from the current cosmological observations and some basic quantum gravitational principles may shed light on probing the nature of DE.  On the one hand, provided that we know little on the theoretical nature of DE at present, the combined cosmic observations can play an important role in understanding the nature of DE. The cosmological parameters space in the DE model can be determined by the constraints of the data combinations. Recently, the 397 SN Ia data was compiled in Ref. \\cite{ref:Condata} by adding CfA3 sample from the CfA SN Group to the Union set by Ref. \\cite{ref:Kowalski}, which include 250 SN Ia at high redshift but only 57 at low redshift, to form the Constitution set. Aside from the SN Ia data, the combined analysis is required in order to break the degeneracy between the cosmological parameters, which includes cosmic observations from baryon acoustic oscillations (BAO), CMB and the observational Hubble data (OHD). The BAO are detected in the clustering of the combined 2dFGRS and SDSS main galaxy samples or the SDSS luminous red galaxies and measure the distance-redshift relation. From these samples, the values of $[r_s(z_d)/D_V(0.2), r_s(z_d)/D_V(0.35)]$ and their inverse covariance matrix in the measurement of BAO can be obtained \\cite{ref:Percival2}. For the measurement of CMB, we utilize the shift parameter $R$ at the photon decoupling epoch $z_\\ast$, the acoustic scale $l_A(z_\\ast)$, and together with the physical baryon density parameter multiplied by 100, thus it is $100\\Omega_bh^2$ \\cite{ref:Komatsu2008, ref:Bueno Sanchez}. Here, it is worth noting that the WMAP distance information $R(z_\\ast)$ and $l_A(z_\\ast)$ can not be measured by WMAP directly, but are derived from making a global fitting constraint with MCMC method by using the full WMAP data on the assumption that a certain cosmological model has been given in advance \\cite{ref:lAR}. Although in theory the inverse covariance matrix on $R(z_\\ast)$ and $l_A(z_\\ast)$ is model dependent, it is feasible to use the derived results about $R(z_\\ast)$ and $l_A(z_\\ast)$ to constrain the parameters in another DE model since $R(z_\\ast)$ and $l_A(z_\\ast)$ do not depend strongly on the DE model which is not far away from the cosmological constant model \\cite{ref:lAR}. What is more, the paper \\cite{ref:ywang} has been demonstrated that $[R(z_\\ast), l_A(z_\\ast), 100\\Omega_bh^2]$ effectively provide a good summary of CMB data when the DE model parameters are constrained. In addition, we employ the OHD at twelve different redshifts determined by using the differential ages of passively evolving galaxies in Ref. \\cite{ref:0907}, where the value of the Hubble constant is replaced by $H_0=74.2\\pm3.6$ in Ref. \\cite{ref:0905}, and add the three more observational data $H(z=0.24)=79.69\\pm2.32, H(z=0.34)=83.8\\pm2.96,$ and $H(z=0.43)=86.45\\pm3.27$ in \\cite{ref:0807}. Since the constraint results of a given model are dependent on the combined data \\cite{ref:Gdata, ref:XUdata}, in this paper we use a fully combined observations from the 397 SN Ia standard candle data, the value of $[r_s(z_d)/D_V(0.2), r_s(z_d)/D_V(0.35)]$ and their inverse covariance matrix in the measurement of BAO, the values of $[R(z_\\ast), l_A(z_\\ast), 100\\Omega_bh^2]$ and their inverse covariance matrix in the measurement of CMB, and the fifteen OHD.  On the other hand, the models which are constructed in light of some fundamental principle are more charming, since this kind of DE model may exhibit some underlying features of DE, for instance the holographic DE model \\cite{ref:holo1,ref:holo2} and the agegraphic DE model \\cite{ref:age1,ref:age2}. The holographic DE model is built on the basis of holographic principle and some features of quantum gravity theory. The agegraphic DE model is derived from taking the combination between the uncertainty relation in quantum mechanics and general relativity into account. In this paper, we focus on the holographic DE model, which is considered as a dynamic vacuum energy. According to the holographic principle, the number of degrees of freedom in a bounded system should be finite and is related to the area of its boundary. By applying the principle to cosmology, one can obtain the upper bound of the entropy contained in the universe. For a system with size $L$ and UV cut-off $\\Lambda$ without decaying into a black hole, it is required that the total energy in a region of size $L$ should not exceed the mass of a black hole of the same size, thus $L^3\\rho_\\Lambda\\leq LM_{pl}^2$. The largest $L$ allowed is the one saturating this inequality, thus we obtain the holographic DE density \\begin{eqnarray} &&\\rho_\\Lambda=\\frac{3c^2M_{pl}^2}{L^2}, \\end{eqnarray} where c is a numerical constant and $M_{pl}$ is the reduced Planck Mass $M_{pl}\\equiv1/\\sqrt{8\\pi G}$. It just means a duality between UV cut-off and IR cut-off. The UV cut-off is related to the vacuum energy, and IR cut-off is related to the large scale of the universe, for example Hubble horizon, particle horizon, event horizon, Ricci scalar or the generalized functions of dimensionless variables as discussed by \\cite{ref:holo1,ref:holo2,ref:holo3,ref:EPJCXU}. Next, we give a brief review on the main results when Hubble horizon, particle horizon, event horizon or Ricci scalar are taken as the IR cut-off, respectively.  $\\bullet$ $L^{-2}=H^2$. As pointed in \\cite{ref:holo2}, it is found that the holographic DE density is in proportion to $H^2$, the same as dark matter density, i.e. $\\rho_{de}/c^2=\\rho_m/(1-c^2)\\propto H^2$. It appears that it is natural to solve the coincidence problem. However, Hsu \\cite{ref:holo1} pointed out that the dark energy EOS $w_{de}=0$ was obtained in this instance. It is obvious that this result is not consistent with the current observations. This bad situation can be changed by considering the holographic DE with Hubble horizon as the time variable cosmological constant. More detailed analysis is presented in Ref. \\cite{ref:XU071709}.  $\\bullet$ $L^{-2}=R_{ph}(a)=a\\int_0^t\\frac{dt'}{a(t')}=a\\int_0^a\\frac{da'}{Ha'^2}$. As shown in paper \\cite{ref:holo2}, Li pointed out that this yields the dark energy EOS is not less than $-1/3$. Thus the current accelerated expansion of our universe can not be well explained. However, this result in \\cite{ref:holo2} is obtained on the assumption that DE dominates. The holographic DE model with particle horizon has been discussed in detail by \\cite{ref:XU054772}.  $\\bullet$ $L^{-2}=R_{eh}(a)=a\\int_t^\\infty\\frac{dt'}{a(t')}=a\\int_a^\\infty\\frac{da'}{Ha'^2}$. The holographic DE model with event horizon can reveal the dynamic nature of the vacuum energy and provide a desired EOS of the holographic DE with the model parameter $c$. Furthermore, the holographic DE behaves like quintessence, cosmological constant and phantom respectively for the different values of the model parameter: $c\\geq1$, $c=1$ and $c\\leq1$ \\cite{ref:holo2}. Therefore, the value of model parameter $c$ plays a crucial role in determining the property of holographic DE in this case. However, this model is confronted with the causality problem: why should the present density of DE be determined by the future event horizon of the universe.  $\\bullet$ $L^{-2}=R=-6(\\dot{H}+2H^2+\\frac{k}{a^2})$. In \\cite{ref:holo3}, it has shown that this model can avoid the causality problem and naturally solve the coincidence problem of dark energy after Ricci scalar is taken as the IR cut-off and the parameters have been well constrained by the combined astronomical observations \\cite{ref:MPLAXU,ref:LiRicci}.  Subsequently, In \\cite{ref:holo4}, Granda and Oliveros generalized the form of the IR cut-off on the basis of the Ricci scalar: \\begin{eqnarray} &&L^{-2}=\\alpha H^2+\\beta \\dot{H}, \\end{eqnarray} where there are two independent model parameters $\\alpha$ and $\\beta$, which can be determined by using the combined constraints of the thorough observational datasets. In this paper, we consider the holographic DE model with new IR cut-off in both flat and non-flat case. The performance of a global fitting will be made by using the Markov Chain Monte Carlo (MCMC) method. In this way, we can work in the framework of multi-parameter freedoms, including the basic cosmological parameters ($\\Omega_bh^2, \\Omega_ch^2, \\Omega_k$) and the new-added model parameters ($\\alpha, \\beta$).  The paper is organized as follows. In next section, we briefly review the holographic DE model with new IR cut-off. In section III, we perform the cosmic observation constraint on the holographic DE model. The last section is the conclusion. ",
        "conclusions": "In summary, in this paper we have performed a global fitting on the parameters in the holographic DE model with new IR cut-off for the flat case and the non-flat case, using a combined cosmic observations from type Ia supernovae, baryon acoustic oscillations, Cosmic Microwave Background and the observational Hubble data. The same constraints are performed on the flat and non-flat concordance models by using the same combined datasets. According to the Markov Chain Monte Carlo (MCMC) analysis, it is shown that the best fitting values of the model parameters ($\\alpha,\\beta$) in the flat holographic DE model with new IR cut-off tend to be smaller than those in the non-flat case. In the holographic DE models, the non-flat case with a smaller value of $\\chi^2/dof$ is much supported by the observations. In the non-flat cases, we have obtained the constraint values of the curvature terms $\\Omega_k=0.0305^{+0.0092 +0.0140}_{-0.0134 -0.0176}$ for the holographic DE model with new IR cut-off and $\\Omega_k=-0.0013^{+0.0070 +0.0103}_{-0.0076 -0.0108}$ for the concordance model. These results indicate the two kinds of the non-flat background geometries in the two models. Then by using the best fit parameters, we plot the evolutions of the dark energy EOS and deceleration parameter with errors. From Fig. \\ref{fig:wqflat} and Fig. \\ref{fig:wqnonflat}, it is found that the EOS of the holographic DE with new IR cut-off can cross the phantom divide $-1$, respectively with the current best values $w_{de0}=-1.1414$ (flat case) and $w_{de0}=-1.0653$ (non-flat case). Comparing the flat and non-flat holographic DE models with the corresponding cases in the $\\Lambda \\textmd{CDM}$ model, we can find that the current combined observations do not favor the holographic DE model with new IR cut-off over the $\\Lambda \\textmd{CDM}$ model."
    },
    "1004/1004.4756_arXiv.txt": {
        "abstract": "We investigate the performance of the parametric Maximum Likelihood component separation method in the context of the CMB B-mode signal detection and its characterization by small-scale CMB suborbital experiments. We consider high-resolution (FWHM$=8'$) balloon-borne and ground-based observatories mapping low dust-contrast sky areas of 400 and 1000 square degrees, in three frequency channels, 150, 250, 410 GHz, and 90, 150, 220 GHz, with sensitivity of order 1 to 10 $\\mu$K per beam-size pixel.  These are chosen to be representative of some of the proposed, next-generation, bolometric experiments.  We study the residual foreground contributions left in the recovered CMB maps in the pixel and harmonic domain and discuss their impact on a determination of the tensor-to-scalar ratio, $r$. In particular, we find that the residuals derived from the simulated data of the considered balloon-borne observatories are sufficiently low not to be relevant for the B-mode science. However, the ground-based observatories are in need of some external information to permit satisfactory cleaning. We find that if such information is indeed available in the latter case, both the ground-based and balloon-borne experiments can detect the values of $r$ as low as $\\sim 0.04$ at $95$\\% confidence level. The contribution of the foreground residuals to these limits is found to be then subdominant and these are driven by the statistical uncertainty due to CMB, including E-to-B leakage, and noise. We emphasize that reaching such levels will require a sufficient control of the  level of systematic effects present in the data. ",
        "introduction": "Astrophysical foregrounds are commonly recognized as one of the major obstacles on the way to first detecting and later exploiting the scientific potential of the Cosmic Microwave Background (CMB) polarization signal. This is particularly the case with so called B-mode polarization \\citep{1997PhRvD..55.1830Z} due to its minute amplitude as compared to the foreground contributions as well as CMB total intensity and E-mode polarization signals. In fact current foreground models \\citep{2007ApJS..170..335P} generally indicate that the foreground B-mode signal may be comparable or exceed the CMB signal by a factor of a few in a broad range of angular scales even in the cleanest available sky areas.  Some kind of foreground cleaning or separation procedure will therefore be necessary and its impact on the final `cleaned' map of the presumed CMB sky needs to be understood and properly taken into account in its subsequent studies. Developing such an understanding is also already of importance for the designing and optimization of the future CMB experiments.  This has been recognized for some time and a number of studies have been performed and published, and which have treated the problem on different levels of generality and detail. The major challenge here is two-fold. Firstly, there is no general recipe for propagating errors incurred during the component separation step, i.e. for including both the statistical uncertainty and foreground residual uncertainty. Secondly, there is no easily calculable metric measuring the impact of the component separation on the B-mode measurement, as both the power spectrum or tensor-to-scalar ratio, $r$, require a proper evaluation of the E-to-B leakage \\citep{2003PhRvD..67b3501B}.  \\citet{2005MNRAS.360..935T} have performed a Fisher analysis of the problem treating the foreground residuals as an additional source of noise, and then estimated the expected limits on $r$. In their approach the starting point was a single foreground contaminated science channel and a noisy foreground template channel from which the level of foreground residual was estimated. Although this allows to avoid specifying in great detail a foreground cleaning technique, no direct connection exists between the noise values they assume and properties of any specific experiment.  They have also neglected the impact of the E-to-B leakage. A similar approach has been followed by \\citet{2006JCAP...01..019V}, who have attempted to link their Fisher matrix considerations to specific, fiducial, multi-frequency data sets. The simplified error propagation they have adopted implicitly bypasses any realistic component separation approach, and so fails to include properly its effects on their final results. They also neglect the presence of the E-to-B leakage. \\citet{2005PhRvD..72l3006A} performed a Fisher analysis as well, but use specific parameters anchored in those of the multi-frequency data set assumed.  This last work together with \\citet{2007PhRvD..75h3508A} and \\citet{2009A&A...503..691B} come the closest in the spirit to what we discuss in this paper, although neither of the latter two works includes an actual power spectrum estimator accounting for the leakage, what is justified at least in part by their focus on full-sky observations. \\cite{2006MNRAS.372..615S} studied an application of an Independent Component Analysis based approach to the component separation of polarized partial-sky maps, resorting to the cleaned map `pseudo-spectra' as a basis for a {\\em qualitative} assessment of its performance and relevance for the B-mode work.  \\citet{2009AIPC.1141..222D} presented a review of most of those earlier approaches, including those incorporating a parametric approach similar to the one considered in this work, and presented their applications in the context of a potential future CMB B-mode satellite mission.  The approach we propose here is more specific. We focus on a particular component separation method and power spectrum estimation approach, which we then use to investigate the impact of the foreground separation on the CMB B-mode detection and characterization. The component separation method is a maximum likelihood (ML) parametric approach~\\citep{2006ApJ...641..665E} in a two-step implementation of \\citet{2009MNRAS.392..216S}.  The power spectrum estimator is a `pure' pseudo-spectrum approach introduced by \\citet{2006PhRvD..74h3002S} \\citep[see also][]{2007PhRvD..76d3001S} and elaborated on by \\citet{2009PhRvD..79l3515G}. Strictly speaking, our results will therefore be specific to these two choices. However, given that these two methods are working, promising algorithms to be implemented in the data analysis pipelines of current and future CMB experiments, the results should be of practical relevance for many efforts currently going on in the field.  We note also that whenever the frequency scaling laws can be assumed to be nearly perfectly known, as in one of the cases we study, and in particular in the small-sky, and thus potentially statistics-starved limit, the parametric maximum likelihood (ML) method would likely become a method of the choice, potentially supplemented by some priors, e.g, spatial templates for all or some of the components \\citep{2009MNRAS.397.1355E}.  The results derived here can therefore be regarded as representative and realistic expectations for the performance of classes of the future experiments we consider. Moreover, part of the analysis presented here can be straightforwardly applied to any component separation method in which foreground spectral and amplitude parameters are estimated in separate steps.  Our focus in this work is on suborbital experiments. Those have a potential advantage of selecting the cleanest sky areas, but suffer due to the cut-sky effects.  They also usually have a limited number of frequency channels with which to observe the sky.  We consider two kinds of experiments: those with an access to the high frequencies ($\\simgtalt 250$ GHz) referred to as balloon-borne, and those with access limited to frequencies lower than $250$ GHz, referred to as ground-based. We will also consider some combination and extensions of these two cases.  We then apply our proposed analysis chain to simulated data for different foreground case studies, allowing for different levels of mismatch between the assumptions made on the analysis and simulation stages, in order to evaluate the impact of the component separation residuals first on the recovered B-mode power spectrum and later on the value of a $r$ which can be derived from such data.  The paper is organized as follows. In Sections~\\ref{sec:miramare} and \\ref{sec:xpure}, we first provide brief descriptions of the specific data analysis techniques and their implementation, used throughout this paper. In Section~\\ref{sec:mockData} we describe our simulated sky model, and in Section~\\ref{sec:mock} we define the experimental characteristics and a set of foreground case studies. Our results are presented in Section~\\ref{sec:Results}, and their analysis, concerning the residuals and their impact on the cosmological B-mode detection, is given in Sections~\\ref{sect:anaRes} and~\\ref{sec:tensor2scalar}, respectively. ",
        "conclusions": "In this paper we study the performance of the maximum likelihood parametric component separation method from the point of view of its application to the CMB B-mode polarization analysis. We investigate the residuals left over from the separation in both the pixel and harmonic domains. We propose an efficient framework for evaluating the pixel domain residuals in the simulation, and show how it can be used to gain important insights into the separation process. We then compute the power spectra of the recovered CMB maps, as well as maps of the residuals, using the pure pseudo spectra technique, and estimate their variances using Monte Carlo simulations. Finally, we propose a Fisher-like approach to evaluate the effects of the foreground residuals on the $r$ parameter and use the latter to derive some estimates of typical values of $r$, which are potentially detectable by the considered experiments at the $95$\\% confidence level. The latter estimates thus include the uncertainties due to sampling variance, noise scatter, E-to-B leakage, and foreground residuals, all of which are consistently propagated through the proposed pipeline.  We focus here on the small-scale, bolometric experiments, broadly dividing them into two classes, referred to as balloon-borne and ground-based setups, both observing in three different frequency bands. We find that the balloon-borne case, with frequency bands at $150$, $250$, and $410$ GHz, provides a robust experimental setup for the detection of the B-mode polarization. The foreground residuals in the recovered CMB maps derived in this case are found to be usually subdominant. This is true whenever the assumed data model is indeed correct, but it also holds when some small systematic effects are permitted. Selected effects of this kind considered in this work include relative calibration errors, unmodelled spatial variation of the spectral parameters, and spectral mismatch between assumed and true spectral scaling laws.  We emphasize that all these systematics, though manageable if sufficiently small, may lead to spurious effects in general, and therefore need to be controlled in actual experiments with a sufficient precision, which need to be determined specifically for any experiment.  The success of the considered balloon-borne cases is related to the wide frequency range available to such experiments, which permits selecting frequency bands at the sufficiently high frequencies to avoid the unwanted residual synchrotron. The latter is found to be a dominant source of the bias for the ground-based experiments. In the balloon case we can also afford a long leverage arm between the lowest (CMB-dominated) and the highest (dust-dominated) frequency bands, which plays a pivotal role in setting tight constraints on the spectral parameters of the dust. From our Fisher-like analysis, we show that one could detect $r$ values as low as $0.04$ at the $95\\%$ confidence level with such experiments, if both our models and measurements are sufficiently well characterized.   For the ground-based case the atmospheric loading restricting the available frequency window proves to be a significant limitation. We find that even in an absence of any systematic effects with three frequency bands set at $90$, $150$, and $220$ GHz, it is generally not possible to produce sufficiently clean CMB maps. This is due to the unmodelled, and thus not separated, synchrotron contribution, which is significant enough (if the polarized emission is at the level suggested by WMAP) at these frequencies to bias the estimation of the dust spectral parameters. We point out that this contribution has been neglected in some earlier works, which consequently has arrived at a different conclusion. This  therefore emphasizes the importance of accurate sky modelling for this kind of the analysis. Nevertheless, we find a satisfactory cleaning can be achieved in such a case if some external information is available. In particular, we discuss the extended ground experiment analysis allowing for the presence of the extra lower frequency channels, rough synchrotron templates, and priors on the dust scalings. We find that in such cases, and under realistic assumptions, the ground-based experiments should reach a sensitivity roughly matching those found in the balloon-borne case, in terms of a detectable $r$ parameter. We also conclude that, for both these types of the experiments, the foreground contamination anticipated in low-contrast dust regions,  should not be an obstacle preventing them from exploring the parameter space of $r$ down to the values $\\sim 0.04$. Indeed, for the considered experimental setups this limit is determined by the uncertainty due to the CMB itself and the instrumental noise, with the effects of the residual foregrounds found to be sub-dominant. In the realm of the actual observations, whether these limits are reached will be crucially dependent on the control of systematic effects. We note here however that the limits derived on $r$ are strongly dependent on the level of the noise assumed in the input single-channel maps. These can be therefore improved upon, if a deeper integration of the same field is performed. However, if no additional external information is used, those limits will remain appropriately higher than the $r_{res}$ values obtained earlier, and below which foreground bias would become significant.  In this context we point out that the framework described in this paper provides a blueprint for similar studies focused this time on systematic effects. It can be also extended to perform a realistic experiment optimization procedure from the viewpoint of detection of the B-mode signal of cosmological origin."
    },
    "1004/1004.1459_arXiv.txt": {
        "abstract": "We have conducted $N$-body simulations of the growth of Milky Way-sized halos in cold and warm dark matter cosmologies. The number of dark matter satellites in our simulated Milky Ways decreases with decreasing mass of the dark matter particle. Assuming that the number of dark matter satellites exceeds or equals the number of observed satellites of the Milky Way we derive lower limits on the dark matter particle mass. We find with $95\\%$ confidence $m_s > 13.3$~keV for a sterile neutrino produced by the Dodelson and Widrow mechanism, $m_s > 8.9$~keV for the Shi and Fuller mechanism, $m_s > 3.0$~keV for the Higgs decay mechanism, and $m_{WDM} > 2.3$~keV for a thermal dark matter particle. The recent discovery of many new dark matter dominated satellites of the Milky Way in the Sloan Digital Sky Survey allows us to set lower limits comparable to constraints from the complementary methods of Lyman-$\\alpha$ forest modeling and X-ray observations of the unresolved cosmic X-ray background and of dark matter halos from dwarf galaxy to cluster scales. Future surveys like LSST, DES, PanSTARRS, and SkyMapper have the potential to discover many more satellites and further improve constraints on the dark matter particle mass. ",
        "introduction": "\\label{sec:1}  Cold dark matter (CDM) is extremely successful at describing the large scale features of matter distribution in the Universe but has problems on small scales. Below the Mpc scale CDM predicts numbers of satellite galaxies for Milky Way-sized halos about an order of magnitude in excess of the number observed. This is the `missing satellites' problem \\cite{kly1999, moo1999}. One proposed solution is that, due to feedback mechanisms, some dark matter satellites do not form stars and are nonluminous dark halos \\cite{efs1992,tho1996,bul2001,ric2004,ric2005A}. Another solution is the power spectrum of density fluctuations may be truncated which may arise if the dark matter is `warm' (particle mass $\\sim 1$~keV) instead of `cold' (particle mass $\\sim 1$~GeV). Warm dark matter (WDM) particles decouple from the other particle species in the early Universe with relativistic velocities and only become nonrelativistic when about a Galactic mass ($\\sim 10^{12} M_{\\odot}$) is within the horizon. Streaming motions while the particles are still relativistic can erase density fluctuations on sub-Galactic scales and reduce the number of satellites. WDM models have been studied by a number of authors \\cite{col2000, avi2001, bod2001, kne2002, kne2003, zen2003, mac2009} in relation to the missing satellites problem and other issues with CDM such as the apparent density cores in spiral and dwarf galaxies \\cite{van2001,swa2003,wel2003,don2004,gen2005,sim2005,gen2007,sal2007,kuz2010}.    $N$-body simulations of WDM cosmologies are complicated by numerical artifacts produced by the discrete sampling of the gravitational potential with a finite number of particles (see \\cite{mel2007} for a review). Matter perturbations collapse and form filaments with nonphysical halos separated by a distance equal to the mean particle spacing (see Fig.~\\ref{figI1}) \\cite{wan2007,mel2007}. These halos are numerical artifacts. The ability of these halos to survive disruption as they accrete from filaments onto Milky Way-sized halos has not been studied and they may contaminate the satellite abundances and distributions in WDM simulations.  \\begin{figure}[!th] \\includegraphics*[width=3in]{fig1.eps} \\caption{Nonphysical halos formed along a filament and accreting onto a larger halo at $z=1$ in a WDM simulation ($m_{WDM}=1$~keV). These halos are numerical artifacts.\\label{figI1}} \\end{figure}  In the past few years, 16 new dwarf spheroidal galaxies have been discovered in the Sloan Digital Sky Survey (SDSS) \\cite{cas1998} (see Table 3 and references therein). After correcting for completeness the estimated number of Milky Way (MW) satellites is $>60$ (see Sec.\\ \\ref{sec:4}). These new dwarfs have low luminosities, low surface brightnesses, and most appear to be dark matter dominated. Since the number of dark matter halos must be greater than or equal to the number of observed satellites, the new data from the SDSS may provide improved limits on the mass of the dark matter particle independent of complementary techniques.  Motivated by the recent increase in the number of observed Milky Way satellites, we have performed new simulations of the growth of Milky Way-like galaxies in CDM and WDM cosmologies for a variety of WDM particle masses. Our goal is to constrain the dark matter particle mass by comparing the number of satellite halos in the simulated Milky Ways to the observed number of luminous satellites for the actual Milky Way. Macci{\\`o} and Fontanot \\cite{mac2009} combined $N$-body simulations with semianalytic models of galaxy formation to compare the simulated and observed Milky Way satellite luminosity functions for CDM and WDM cosmologies. In this work, we do not make any assumptions on how we populate dark matter halos with luminous galaxies. We simply impose that the number of observed satellites is less than or equal to the number of dark matter halos for a range of Galactocentric radii. This guarantees a robust lower limit on the dark matter particle mass. ",
        "conclusions": "We found that a model with $m_{wdm}=4$~keV produces the best fit to observations at $< 50$~kpc, i.e.\\ this model has a number of dark matter satellites equal to the number of observed luminous satellites. However, due to the large uncertainties in the number of observed satellites due to partial sky coverage and on the number of simulated satellites due to Poisson and intrinsic scatter, that partially reflects observational uncertainties on the mass and $v_{max}$ of the Milky Way, we find much weaker lower limits on $m_{wdm}$ than $4$~keV. In the future however, the lower limit on $m_{wdm}$ will improve as observations of MW satellites become more complete. The scatter of the simulation can also be reduced using constrained simulations of the Local Group (also including the effect of baryons) in combination with more accurate determination of the mass, rotation curve, and concentration of the Milky Way.  Considering the various uncertainties in the number of observed and simulated satellites, we found a conservative lower limit of $m_{WDM} > 2.3$~keV ($2\\sigma$) on the dark matter particle mass. We also found the $1$~keV WDM simulations have too few satellites to match the Milky Way observations. This agrees with the semianalytic modeling and Milky Way satellite luminosity functions in WDM cosmologies work of Macci{\\`o} and Fontanot \\cite{mac2009}; however, we only apply a cut to the simulated halos to avoid numerical effects and do not make assumptions on how the dark matter halos are populated by luminous galaxies.  Our result can also be compared to limits on the particle mass from the Lyman-$\\alpha$ forest in high redshift quasars. Lyman-$\\alpha$ absorption by neutral hydrogen along the line of sight to distant quasars over redshifts 2--6 probes the matter power spectrum in the mildly nonlinear regime on scales 1--80 Mpc/$h$. Viel et al.~\\cite{vie2005, vie2006, vie2008} have numerically modeled the Lyman-$\\alpha$ forest flux power spectra for varied cosmological parameters and compared to observed quasar forests to obtain lower limits on the dark matter particle mass. Their 2006 work \\cite{vie2006} used low resolution spectra for 3035 quasars ($2.2 < z < 4.2$) from the SDSS \\cite{mcd2006} and found a 2$\\sigma$ lower limit of $2$~keV for a thermal WDM particle. This limit agrees with our results that a $2.3$~keV particle is the lower limit that can reproduce the observed number of Milky Way satellites and agrees with the Lyman-$\\alpha$ work of Seljak et al.~\\cite{sel2006} who find a 2$\\sigma$ limit $> 2.5$~keV for a thermal particle. The latest work of Viel et al.\\ \\cite{vie2008} uses high resolution spectra for 55 quasars ($2.0 < z < 6.4$) from the Keck HIRES spectrograph in addition to the SDSS quasars. With the new data they report a lower limit of $4$~keV (2$\\sigma$). A caveat arises in Viel et al.~\\cite{vie2009}, who show the flux power spectrum from the SDSS data prefer larger values of the intergalactic medium (IGM) temperature at mean density than expected from photoionization. The flux power spectrum temperature is also higher than that derived from an analysis of the flux probability distribution function of 18 high resolution spectra from the Very Large Telescope and also higher than constraints from the widths of thermally broadened absorption lines \\cite{ric2000, sch2000}. This could be explained by an unaccounted for systematic error in the SDSS flux power spectrum data which may also affect the derived dark matter particle mass limits.  Using the scaling relation for sterile neutrinos we find a lower limit $m_s > 13.3$~keV with $95\\%$ confidence for a DW produced sterile neutrino particle. Scaling to the other production mechanisms we get $m_s > 8.9$~keV for the SF mechanism and $m_s > 3.0$~keV for Higgs decay sterile neutrinos; however, we note this is not based on transfer function calculations for the SF and Higgs mechanisms but assumes a simple scaling for the average momentum for the different production mechanisms \\cite{kus2009}. The Lyman-$\\alpha$ forest observations discussed above in the context of a thermal particle also set limits on the sterile neutrino mass. The 2006 work of Viel et al.~\\cite{vie2006} sets $m_s > 11$~keV and is similar to the Seljak et al.~\\cite{sel2006} limit $m_s > 14$~keV. The 2008 work of Viel et al.~\\cite{vie2008} sets the highest limit of $m_s > 28$~keV but is subject to the caveats mentioned above.  Sterile neutrinos are expected to radiatively decay to a lighter mass neutrino and a X-ray photon with energy $E_{\\gamma}=m_s/2$. X-ray observations of the diffuse X-ray background \\cite{boy2006A} and dark matter halos in clusters \\cite{aba2006A,boy2006,rie2007,boy2008}, M31 \\cite{wat2006}, dwarf spheroidal galaxies \\cite{boy2007,rie2009,boy2009,loe2009}, and the halo of the Milky Way \\cite{rie2006,aba2007,boy2007} have all been used to set constraints on the sterile neutrino mass. Observations of the diffuse X-ray background have set $m_s < 9.3$~keV~\\cite{boy2006A}, while the Virgo and Coma clusters have been used to set $m_s < 6.3$~keV~\\cite{aba2006A} which also agrees with limits from the bullet cluster, 1E~0657-56, $m_s < 6.3$~keV~\\cite{boy2008} and is close to results from the Milky Way halo $m_s < 5.7$~keV~\\cite{aba2007}. Tighter constraints have been determined from M31 observations $m_s < 3.5$~keV~\\cite{wat2006} and from the dwarf spheroidal Ursa Minor $m_s < 2.5$~keV~\\cite{loe2009}.  These upper limits are well below the lower limits derived in this work and from Lyman-$\\alpha$ observations and seem to rule out the DW and SF production mechanisms. However, all of these mass limits, including the constraints set in this work, are model dependent and make certain assumptions. In general X-ray constraints depend on the sterile neutrino mass, the mixing angle with active neutrinos $\\theta$, and the cosmic matter density of sterile neutrinos $\\Omega_s$. There are also assumptions about the initial conditions, that there were no sterile neutrinos in the early Universe at temperatures $> 1$~GeV, there was no entropy dilution after creation, and no coupling to other particles. There are also uncertainties with the calculation of production rates because these occur at temperatures where the plasma is neither well described by hadronic nor quark models \\cite{asa2006,boy2006A}. Depending on the assumptions made and the adopted production model the relationship between $m_s$, $\\theta$, and $\\Omega_s$ changes so that robust constraints cannot be placed on any one model parameter.  There has also been a report of a detection of a dark matter X-ray emission line from Willman~1 consistent with $m_s=5.0 \\pm 0.2$~keV~\\cite{loe2010,kus2010}. This detection is provisional \\cite{boy2010} but if confirmed the limits derived in this work imply the sterile neutrinos are not produced by the DW mechanism or do not constitute the entirety of the dark matter."
    },
    "1004/1004.1173_arXiv.txt": {
        "abstract": "We use Shen et al.'s (2009) measurements of luminosity-dependent clustering in the SDSS Data Release 5 Quasar Catalog, at redshifts $0.4 \\leq z \\leq 2.5$, to constrain the relation between quasar luminosity and host halo mass and to infer the duty cycle \\fopt, the fraction of black holes that shine as optically luminous quasars at a given time.  We assume a monotonic mean relation between quasar luminosity and host halo mass, with log-normal scatter $\\Sigma$. For specified \\fopt\\ and $\\Sigma$, matching the observed quasar space density determines the normalization of the luminosity-halo mass relation, from which we predict the clustering bias. The data show no change of bias between the faint and bright halves of the quasar sample but a modest increase in bias for the brightest 10\\%. At the mean redshift $z=1.45$ of the sample, the data can be well described either by models with small intrinsic scatter ($\\Sigma=0.1$ dex) and a duty cycle \\fopt$=6\\times 10^{-4}$ or by models with much larger duty cycles and larger values of the scatter. ``Continuity equation'' models of the black hole mass population imply $\\mfopt \\geq 2\\times 10^{-3}$ in this range of masses and redshifts, and the combination of this constraint with the clustering measurements implies scatter $\\Sigma \\geq 0.4$ dex. These findings contrast with those inferred from the much stronger clustering of high-luminosity quasars at $z\\approx 4$, which require minimal scatter between luminosity and halo mass and duty cycles close to one. ",
        "introduction": "\\label{sec|intro}  The strong correlations between the masses of central black holes (BHs) and the luminosities, dynamical masses, and velocity dispersions $\\sigma$ of their host galaxies imply that the growth processes of BHs and their hosts are intimately linked (e.g., Magorrian et al. 1998; Ferrarese \\& Merritt 2000; Gebhardt et al.\\ 2000; Ferrarese 2002; Ferrarese \\& Ford 2005; Graham 2007; Tundo et al. 2007; Shankar et al. 2009b). However, constraining the cosmological evolution of BHs remains a challenge.  Although a variety of theoretical models may roughly match observations, the underlying physical assumptions on BH growth can vary drastically from one model to another (e.g., So{\\l}tan 1982; Silk \\& Rees 1998; Salucci et al. 1999; Cavaliere \\& Vittorini 2000; Kauffmann \\& Haehnelt 2000; Yu \\& Tremaine 2002; Steed \\& Weinberg 2003; Wyithe \\& Loeb 2003; Granato et al. 2004, 2006; Marconi et al. 2004; Merloni et al. 2004; Yu \\& Lu 2004; Miralda-Escud\\`{e} \\& Kollmeier 2005; Murray et al. 2005; Cattaneo et al. 2006; Croton et al. 2006; Hopkins et al. 2006; Lapi et al. 2006; Shankar et al. 2004, 2006, 2009a; Malbon et al. 2007; Monaco et al. 2007; Croton 2009; Cook et al. 2009). Quasar clustering provides additional, independent constraints on the BH population, helping to discriminate among otherwise viable models. As outlined by Martini \\& Weinberg (2001) and Haiman \\& Hui (2001; see also Wyithe \\& Loeb 2005; Lidz et al. 2006; Hopkins et al. 2007a; White et al. 2008; Shen et al. 2009a,b; Shankar et al. 2009c; Wyithe \\& Loeb 2009; Bonoli et al. 2009), the clustering is an indirect measure of the masses, and therefore number densities, of the halos hosting the quasars. In turn, the ratio between the quasar luminosity function and the halo mass function provides information on the duty cycle, i.e., the fraction of halos that host active quasars at a given time. In general terms, stronger clustering implies that quasars reside in rarer, more massive hosts, and matching the observed quasar space density then requires a higher duty cycle.  In this paper, we model Shen et al.'s (2009a; S09 hereafter) recent measurements of luminosity-dependent quasar clustering derived from the quasar redshift survey (Schneider et al.\\ 2007) of the Sloan Digital Sky Survey (SDSS; York et al. 2000) Data Release 5 (DR5; Adelman-McCarthy et al.\\ 2007). Ross et al.\\ (2009) also analyze the clustering of this quasar survey, concentrating on redshift evolution, but here we focus on the S09 results because they isolate the luminosity dependence of clustering. Our aim is to answer basic questions about the evolution of the AGN and supermassive BH population at $z \\leq 2.5$. Does the duty cycle depend on quasar luminosity and/or redshift? What is the underlying relation between quasar luminosity and halo mass? Does it have scatter? More generally, what combinations of duty cycle and scatter are allowed by the measurements?  Throughout the paper we adopt $\\Omega_m=0.26$, $\\Omega_\\Lambda=0.74$, $h\\equiv H_0/100\\, {\\rm km\\, s^{-1}\\, Mpc^{-1}}=0.7$, $\\Omega_b=0.0435$, $n_s=0.95$, $\\sigma_8=0.78$, and the transfer function of Eisenstein \\& Hu (1999; with zero neutrino contribution), which matches the cosmology used by S09.  \\begin{figure*}% \\includegraphics[width=17.5truecm]{bLYueFig.eps} \\caption{Bias as a function of bolometric luminosity. In all panels, the \\emph{solid circles} are the mean bias measured by S09 from the quasar auto-correlation function for the faint, bright, and brightest subsamples in their analysis. The \\emph{open circle} with \\emph{dashed} error bars is the bias measured for the brightest subsample including in the fits the bins with negative correlation function. The \\emph{open square} is the bias computed from the cross-correlation of the most luminous sources with the rest of the sample. The data are compared with predictions of several models for the mean bias at $z=1.45$, the average redshift of the S09 sample. \\emph{Left panel}: Comparison among models with the same value for the duty cycle \\fopt$=6\\times 10^{-4}$ and different values of the intrinsic Gaussian scatter $\\Sigma$ (in dex) in the quasar luminosity-host halo relation, as labeled. \\emph{Central panel}: Comparison among models with the same scatter $\\Sigma=0.1$ dex but different values of the duty cycle \\fopt, as labeled. \\emph{Right panel}: Comparison among three different models: one with constant \\fopt$=6\\times 10^{-4}$ at all luminosities and scatter $\\Sigma=0.1$ (\\emph{solid} line); another with equal scatter but with a decreasing duty cycle \\fopt$=6\\times 10^{-4}$ at $\\log (L/{\\rm erg\\, s^{-1}})=46.31$ and \\fopt$/2$ and \\fopt$/4$ at $\\log (L/{\\rm erg\\, s^{-1}})=46.56$ and 46.84, respectively (\\emph{dotted} line); and finally a model with \\fopt$=2\\times 10^{-3}$ and $\\Sigma=0.5$ (\\emph{long-dashed} line). \\label{fig|bLz}} \\end{figure*} ",
        "conclusions": "\\label{sec|discu}  Studies of quasar clustering have generally failed to find any significant dependence of clustering strength on quasar luminosity, at least at $z \\leq 2.5$.  The S09 study is one of the first to separate luminosity dependence from redshift evolution, and it mostly confirms this basic finding, except for the $\\sim 2\\sigma$ increase in bias for the brightest 10\\% of the quasars at $z \\leq 2.5$.  At first glance, the absence of luminosity dependence appears to contradict models like those of Martini \\& Weinberg (2001) or Haiman \\& Hui (2001), which assume a monotonic relation between quasar luminosity and host halo mass and therefore predict a stronger bias for more luminous quasars. However, Figures~\\ref{fig|bLz} and~\\ref{fig|WpYue} show that the S09 results can be reproduced by a model with constant duty cycle for optical quasar activity, $\\mfopt \\approx 6\\times 10^{-4}$, and minimal scatter between luminosity and halo mass. The weakness of the predicted luminosity dependence arises because, even with the large size of the SDSS quasar survey, the dynamic range of luminosity at fixed redshift is not very large ($\\approx 0.5$ dex), and the host halos at these luminosities and redshifts are not on the extreme, steeply rising tail of the $b(M)$ relation.  Croton (2009) reaches a similar conclusion (comparing to other data sets), with a model that is different in technical implementation from ours but similar in practice.  However, the S09 bias measurements can also be fit by models with a higher duty cycle and substantial scatter between luminosity and halo mass.  The increase in bias for S09's highest luminosity bin implies an upper limit on scatter, but this increase is only marginally detected depending on which method is used to estimate the bias.  Figure~\\ref{fig|bLz} shows an explicit example of an acceptable model with $\\mfopt=2\\times 10^{-3}$ and log-normal scatter $\\Sigma=0.5$ dex, and Figure~\\ref{fig|chi2} shows the regions of the $\\mfopt-\\Sigma$ parameter space that yield acceptable agreement with the S09 bias measurements. As discussed in \\S\\ref{subsec|conteq}, models of the quasar population that infer the duty cycle by evolving the BH mass function and comparing to the quasar luminosity function imply $\\mfopt \\ga 2\\times 10^{-3}$. Taken together, the clustering constraints and the continuity equation models imply substantial scatter in the luminosity-halo mass relation, with $\\Sigma \\geq 0.4$ dex.  Applying linewidth estimators of BH mass in the AGN and Galaxy Evolution Survey (AGES), Kollmeier et al.\\ (2006) infer a scatter in quasar Eddington ratios of $\\sigma_\\lambda \\leq 0.3$ dex, though Netzer et al.\\ (2007) and Shen et al.\\ (2008) argue for somewhat larger scatter based on other data sets.  The total scatter between luminosity and halo mass is a combination (in quadrature) of the scatter in Eddington ratios and the scatter between halo mass and BH mass. Physically, many models of quasar activity predict broad Eddington ratio distributions as a consequence of ``post-peak'' accretion onto a central BH, after a rapid growth phase in which the BH mass grows exponentially at a near-Eddington accretion rate (e.g., Yu \\& Lu 2008; Hopkins \\& Hernquist 2009; Shen 2009). Various authors have argued that such prolonged post-peak activity is the key to reconciling the faint end of the AGN luminosity function with measurements of quasar bias at low redshift (e.g., the above papers and Lidz et al.\\ 2006; Marulli et al.\\ 2008; Bonoli et al.\\ 2009; Shankar et al., in prep.). We conclude that scatter of $0.4-0.6$ dex in the luminosity-halo mass relation at these redshifts is plausible on both theoretical and observational grounds.  Several groups have recently tried to measure, or limit, redshift evolution of the scaling between BH mass and host galaxy properties. As several recent papers have pointed out (e.g., Lauer et al. 2007; Merloni et al. 2010; Shankar et al. 2009b; Shen \\& Kelly 2009), a large scatter between quasar luminosity and the galaxy scaling property (such as stellar mass or velocity dispersion $\\sigma$) can bias such measurements.  These biases arise from a combination of flux-limit effects, rapidly falling stellar mass (or velocity dispersion) functions of galaxies, and intrinsic scatter in the scaling relations themselves, which conspire to cause an apparent rise in the mean BH mass at fixed galaxy properties with increasing redshift. Merloni et al. (2010) note that an increasing scatter with increasing $z$ could be enough to explain the trend of evolving black hole mass over galaxy mass ratio measured in their data. Decarli et al. (2010; see also Bennert et al. 2010 for similar conclusions at lower redshifts) argue that strong evolution in the black hole mass-galaxy mass relation is still present even after carefully accounting for flux-limit effects, although they did not allow the possibility of redshift-dependent scatter in the relations (see also discussion in Shen \\& Kelly 2009). The substantial scatter inferred from our analysis shows that biases associated with this scatter must be carefully assessed in studies of the evolution of scaling relations.  A large dispersion between quasar luminosity and host halo mass cannot be the general rule at all redshifts and luminosities. In particular, explaining the high clustering amplitude measured for SDSS quasars at $z \\approx 4$ by Shen et al.\\ (2007) requires both minimal scatter and duty cycles close to one (White et al.\\ 2008; Shankar et al.\\ 2009c; Bonoli et al.\\ 2009). The quasars in this $z\\approx 4$ sample are considerably more luminous than the lower redshift quasars whose clustering is modeled here, so in principle the difference in scatter could reflect either redshift dependence or luminosity dependence. Fine et al.\\ (2008) claim direct empirical evidence for a decrease of $\\Sigma$ with increasing quasar luminosity, based on linewidth estimates of BH mass, and a decrease of this sort is also found in numerical simulations of merger-driven quasar activity (e.g., Hopkins \\& Hernquist 2009, and references therein). Assuming that $\\lambda \\approx 1$ sets an upper limit on BH luminosity, decreasing scatter at high luminosity is plausible because the BH mass function declines rapidly at high masses, so that the most luminous quasars will almost always be powered by BHs radiating near the maximum allowed Eddington ratio. (These arguments address only the scatter in $\\lambda$, not the scatter in BH mass at fixed halo mass.) We have checked that we can fit the $b(L)$ data in Figure~\\ref{fig|bLz} using models with $\\mfopt \\approx 10^{-3}$ and decreasing scatter at high luminosity, e.g., $\\Sigma(L) = 0.6$, 0.3, 0.1 dex for the three bins of increasing luminosity, or even $\\Sigma(L)=0.4$, 0.2, 0.1 dex. However, the bias in the highest luminosity bin, which is rather uncertain at present, can significantly constrain such models.  The duty cycles inferred from our analysis at $z\\approx 1.45$ are substantially lower than the values $f \\approx 0.2$ and $f\\approx 1$ inferred from the Shen et al.\\ (2007) measurements of the clustering of quasars at $z\\approx 3$ and $z\\approx 4$ (see Shen et al.\\ 2007; White et al.\\ 2008; Shankar et al.\\ 2009c). This decline in duty cycle at low redshifts is expected from continuity equation models: the BH mass function grows in time, but the observed quasar luminosity function declines at $z < 2$, so a lower duty cycle is required to reconcile them (see, e.g., figure 7 of SWM). Our current analysis does not constrain duty cycle evolution at $z<2$, but strong evolution over this interval is predicted by the SWM model and is consistent with the S09 correlation function data (see Figure~\\ref{fig|WpYue}).  The measurements in S09 provide significant constraints on the relation between quasar luminosity and halo mass, though leaving substantial degeneracy between the duty cycle and the scatter in this relation.  Reducing statistical errors and remaining systematic uncertainties, especially for the brightest luminosities, would tighten these constraints; in particular, an unambiguous and precise measurement of luminosity-dependent bias would place much tighter restrictions on scatter. The quasar catalog from SDSS DR7 (Adelman-McCarthy et al.\\ 2008) should yield noticeable improvements, with roughly 50\\% smaller error bars and fewer issues with internal boundaries in the survey region.  Since the SDSS quasar sample has high completeness and (with DR7) covers most of the high-latitude northern sky, it will be difficult to go much further with autocorrelation measurements in the S09 luminosity and redshift range. Cross-correlation against denser samples of objects --- fainter AGN or bright galaxies --- could yield higher precision clustering measurements, perhaps with photometric samples from surveys such as Pan-STARRS and LSST, but perhaps requiring spectroscopic samples like those envisioned for ambitious baryon acoustic oscillation experiments. The constraints on host halo populations can also be improved by extending clustering measurements to smaller scales, where quasar pairs from the same halo contribute, and to fainter luminosities, such as those probed by the 2dF Quasar Redshift Survey, the SDSS photometric quasar catalog, and X-ray surveys (e.g., Hennawi et al. 2006; Myers et al. 2007; Plionis et al. 2008, Hennawi et al. 2009); for example, Shen et al.\\ (2009b) use small scale measurements to put constraints on the duty cycle of BHs in satellite galaxies. Quasar clustering as a cosmological tool has moved from a prospect (Osmer 1981) to reality, and the growing precision and dynamic range of these measurements --- in luminosity, redshift, and lengthscale --- will teach us about the growth of supermassive black holes and the mechanisms that transform them from dormant monsters to brilliant beacons, and back."
    },
    "1004/1004.0049_arXiv.txt": {
        "abstract": "{} {The aim of this study is to understand the chemical conditions of ices around embedded young stellar objects (YSOs) in the metal-poor Large Magellanic Cloud (LMC).} {We performed near-infrared (2.5--5$\\mu$m) spectroscopic observations toward 12 massive embedded YSOs and their candidates in the LMC using the Infrared Camera (IRC) onboard  {\\it AKARI}. We estimated the column densities of the H$_2$O, CO$_2$, and CO ices based on their 3.05, 4.27, and 4.67$\\mu$m absorption features, and we investigated the correlation between ice abundances and physical properties of YSOs. } {The ice absorption features of H$_2$O, CO$_2$, $^{13}$CO$_2$, CO, CH$_3$OH, and possibly XCN are detected in the spectra. In addition, hydrogen recombination lines and PAH emission bands are detected toward the majority of the targets. The derived typical CO$_2$/H$_2$O ice ratio of our samples ($\\sim$0.36 $\\pm$ 0.09) is greater than that of Galactic massive YSOs ($\\sim$0.17 $\\pm$ 0.03), while the CO/H$_2$O ice ratio is comparable. It is shown that the  CO$_2$ ice abundance does not correlate with the observed characteristics of YSOs: the strength of hydrogen recombination line and the total luminosity. Likewise, clear correlation is not seen between the CO ice abundance and YSO characteristics, but it is suggested that the CO ice abundance of luminous samples is significantly lower than in other samples. } {The systematic difference in the CO2 ice abundance around the LMC's massive YSOs, which was suggested by previous studies, is confirmed with the new near-IR data. We suggest that the strong ultraviolet radiation field and/or the high dust temperature in the LMC are responsible for the observed high abundance of the CO$_2$ ice. It is suggested that the internal stellar radiation does not play an important role in the evolution of the CO$_2$ ice around a massive YSO, while more volatile molecules like CO are susceptible to the effect of the stellar radiation. } ",
        "introduction": "% The infrared spectra of embedded young stellar objects (YSOs) show absorption features originating in various ice species \\citep[solid state molecules, e.g.,][]{Ger99, Gib04, Boo08}. It is believed that a large amount of heavy elements and complex molecules of the interstellar medium (ISM) are preserved as ices in the dense and cold regions (n$_H$ $\\geq$ 10$^4$ cm$^{-3}$, T $\\sim$ 10 -- 20K), such as an envelope of a deeply embedded YSO \\citep{vDB98, BE04}. Since the star formation proceeds in the region of dense molecular clouds, these ices dominate the chemical evolution of YSOs. Absorption profiles of ices are known to be sensitive to a chemical composition and a temperature of dust grains, and ices are important tracers of the thermal history of circumstellar environments of YSOs \\citep[e.g.,][]{Pon08,Zas09}. Thus, investigating the compositions of ices as a function of physical environments is crucial for understanding the chemical evolution of YSOs and is one of the key topics in the current astrochemistry. Ices are also detected toward solar system objects such as comets, icy satellites, and Mars. Interstellar ices are thought to be taken into planets and comets as a result of subsequent planetary formation activities \\citep{Ehr00}. Therefore the chemical evolution of ices around YSOs is also of interest in terms of chemical conditions of the planetary formation.   Observations of ices around extragalactic YSOs are one of the challenges in current ice studies \\citep{ST,Oli09,Sea09,vanL05,vanL10}. So far, infrared spectroscopic observations of extragalactic embedded YSOs are still limited by observational difficulties, and their circumstellar chemistry is poorly understood. However, it is  probable that different galactic environments (e.g., metallicity or radiation field) could affect the properties of circumstellar material.  Thus detailed investigations of extragalactic YSOs provide us important information for understanding the diversity of chemical conditions of YSOs in the universe. Thanks to the progress of infrared space telescopes, such as {\\it AKARI} \\citep{Mur07} and {\\it Spitzer} \\citep{Wer04}, we are now able to extend the study of ices around YSOs to extragalactic objects.   The Large Magellanic Cloud (LMC), the nearest irregular galaxy to our Galaxy, offers an ideal environment for studying extragalactic star formations since individual YSOs can be identified with reasonable spatial resolution (1\"$\\sim$0.25pc) thanks to its proximity \\citep[$\\sim$50kpc;][]{Alv04}. A unique metal-poor environment of the LMC \\citep[Z $\\sim$ 0.3Z$_{\\odot}$,][]{Luc98} results in a generally high UV radiation field over the galaxy, and these factors can affect the chemical conditions of circumstellar materials. The nearly face-on viewing angle of this galaxy \\citep[$\\sim$35 degrees,][]{MC01} allows correlation studies of individual YSOs and the ISM. Because of these advantages, various types of surveys have been performed toward the LMC, which provide information of both ISM and stars over a wide wavelength range \\citep[e.g.,][]{Zar04,Mei06,Kat07}.   \\citet{ST} spectroscopically confirmed seven massive YSOs in the LMC based on the {\\it AKARI} LMC spectroscopic survey \\citep{Ita08} and reported the detection of the H$_2$O and CO$_2$ ices toward these objects. \\citet{Oli09} analyzed the 15.2$\\mu$m feature of the CO$_2$ ice for 15 embedded YSOs in the LMC and discuss the chemical properties of the ice mantles. The H$_2$O and CO$_2$ ices are ubiquitous and are the major components of interstellar ices. Based on the low-resolution (R $\\sim$ 20) spectra obtained in \\citet{ST}, we derived a typical CO$_2$/H$_2$O ice ratio and showed that the ratio around the LMC's massive YSOs (0.45$\\pm$0.17) is higher than previously reported toward Galactic massive YSOs \\citep[0.17$\\pm$0.03][]{Ger99}. The result suggests the chemically different nature of YSOs in the metal-poor galaxy.   An accurate determination of column densities of major ice species is crucial when discussing the variation of ice abundances between the objects. However, the uncertainties in the derived column densities from our previous low-resolution spectra are large, and the number of samples for which the column densities of major ice species are determined is still small. In addition, the features of relatively minor ice species, such as CH$_3$OH and CO are difficult to identify by the low-resolution spectra. Investigation of these minor ices should help in constraining the processing of ices in a YSO envelope \\citep[e.g.,][]{Dar99,Thi06}. Thus follow-up observations with higher spectral resolution are required to further understand the chemical conditions of extragalactic YSOs.    In the present study, we used the higher spectral resolution (R $\\sim$ 80) mode of the Infrared Camera \\citep[IRC,][]{TON07} onboard {\\it AKARI} to observe massive embedded YSOs and their candidates in the LMC. The correlation between the chemical properties of ices and the YSO characteristics, which could not be discussed in \\citet{ST} due to the large uncertainties of the observed spectra, is discussed in this paper. In addition, the improved calculation method of the ice column density by using the curve-of-growth is presented in an Appendix. ",
        "conclusions": "%  \\subsection{Large CO$_2$/H$_2$O ice ratio in the LMC}%  The derived column densities of the H$_2$O and CO$_2$ ices are plotted in Figure \\ref{NG_H2O_CO2}. The average CO$_2$/H$_2$O ice column density ratio of our samples is calculated to be 0.36 $\\pm$ 0.09. The present samples are massive YSOs and should be compared to Galactic massive samples. Thus the column densities of Galactic massive YSOs \\citep[CO$_2$/H$_2$O $\\sim$ 0.17 $\\pm$ 0.03,][]{Ger99,Gib04} are also plotted in the figure for comparison. It is reported that Galactic quiescent molecular clouds show a similar CO$_2$ ice abundance \\citep[$\\sim$ 0.18 $\\pm$ 0.04,][]{Whit07}. As can be seen in the figure, the typical abundance\\footnote{The ``abundance'' of a given ice is generally defined as the ratio of the column density relative to the H$_2$O ice since it is the most abundant ice species.} of the CO$_2$ ice toward the LMC's massive YSOs is larger than for Galactic massive YSOs. Also, the scatter of the CO$_2$/H$_2$O ratio is larger than for Galactic samples.  \\begin{figure}[!htb] \\includegraphics[width=\\hsize, angle=0]{13815fg5.eps} \\caption{ CO$_2$ ice vs. H$_2$O ice column density in units of 10$^{17}$ cm$^{-2}$. Open squares with error bars represent the results of this study. Filled squares represent those of Galactic massive YSOs \\citep{Ger99, Gib04}. An average CO$_2$/H$_2$O ice ratio of our samples (0.36) is plotted as a solid line. Dotted and dashed lines represent CO$_2$/H$_2$O ratio $\\sim$ 0.32 and 0.17, which correspond to the typical CO$_2$/H$_2$O ratio of Galactic low- and intermediate-mass YSOs, and massive YSOs, respectively. The dotted line also corresponds to the CO$_2$/H$_2$O ratio of the LMC's massive YSOs derived by \\citet{Oli09}. } \\label{NG_H2O_CO2} \\end{figure}  \\citet{ST} reported that the typical CO$_2$/H$_2$O ice ratio in the LMC is 0.45 $\\pm$ 0.17 based on the low-resolution spectra of the LMC's YSO (targets included in the present samples). The present result is consistent with the CO$_2$/H$_2$O ratio of \\citet{ST} within the uncertainties, but it is slightly lower than our previous result. It is mostly attributed to the effect of the PAH 3.3\\,$\\mu$m emission, which was not resolved and thus not correctly counted in the previous study. \\citet{Oli09} have estimated the column densities of the CO$_2$ ice for some common objects with the present targets based on the 15.2$\\mu$m feature and also revised some of the column density estimates of \\citet{ST}. The CO$_2$ ice column densities estimated in \\citet{Oli09} are slightly lower than our estimate in Table \\ref{tbl_NG_ice}, but generally consistent within the uncertainties. They also suggest that the CO$_2$ ice abundance of the LMC samples is higher than in massive YSOs in our Galaxy (CO$_2$/H$_2$O $\\sim$ 0.32, Fig.~\\ref{NG_H2O_CO2}).   The high CO$_2$/H$_2$O ratio in the LMC indicated by the previous studies is now confirmed with the increased number of samples. In addition, the uncertainties of the derived column densities decreased from those of \\citet{ST} thanks to the higher S/N, the higher spectral resolution, and the improved treatment of the curve-of-growth method. It is also worth mentioning that a relatively high CO$_2$/H$_2$O ratio of 0.32 $\\pm$ 0.02 is observed toward Galactic low-- and intermediate-- mass YSOs \\citep{Pon08}. Their samples also show a large scatter in the CO$_2$ ice abundance as seen in the present samples. A similar trend in the CO$_2$ ice abundance between the LMC's massive YSOs and Galactic low-- and intermediate--mass YSOs may indicate some similarities of star formation conditions between these different samples, but a qualitative explanation is an issue for future investigation.   The derived CO/H$_2$O ice ratio of the LMC's YSO ranges from $\\sim$0.01 to 0.2 (Table \\ref{tbl_NG_ice}), and this is similar to what is observed toward Galactic massive YSOs \\citep[$\\sim$0.02 -- 0.2,][]{Chi98,Gib04}. Also, the distribution range of the H$_2$O ice column densities of our samples is nearly comparable to that of the Galactic YSOs. Therefore, we can assume that there should be a different physical or chemical environment that selectively enhances the production of CO$_2$ ice in the LMC.   It is widely accepted that the CO$_2$ ice in the cold envelope of a YSO is formed by grain surface reactions since gas-phase reactions are not able to produce the observed abundance of the CO$_2$ ice, while molecules like CO are mainly formed by the gas-phase reaction \\citep[e.g., ][]{Mil91}. However, the formation mechanism of the CO$_2$ ice still remains unclear, and it is one of the key topics in astrochemistry. Laboratory experiments indicate that the CO$_2$ ice is efficiently produced by the UV photon irradiation to H$_2$O-CO ice mixtures \\citep[e.g., ][]{Ger96,Wat07}. It is reported by several authors that the LMC has an order-of-magnitude stronger UV radiation field than our Galaxy owing to its active star formation and its metal-poor environment \\citep{Isr86, Lee07}. Thus the strong UV radiation field is able to help the efficient formation of the CO$_2$ ice in the LMC.   On the other hand, an alternative mechanism of the CO$_2$ ice formation has been proposed to explain the detection of the abundant CO$_2$ ice toward quiescent molecular clouds, where energy sources of UV photon are not expected \\citep{Whit98}. A theoretical study suggests that the CO$_2$ ice can also be formed through the diffusive surface chemistry without UV irradiation \\citep{Ruf01}. This model suggests that a sufficient abundance of the CO$_2$ ice can be produced at relatively high dust temperatures, but this process is quite sensitive to the dust temperature and other assumed conditions. Several studies report that the interstellar dust temperature in the LMC is generally higher than in our Galaxy based on far-infrared to submillimeter observations of diffuse emission \\citep[e.g., ][]{Agu03, Sak06}. In addition, \\citet{vanL10,vanL10_b} derived the dust temperatures of YSOs in the LMC and Small Magellanic Cloud (SMC) and suggest that the dust temperatures become higher in more metal-poor environments of the SMC. Thus the dust temperatures of the present samples may be systematically higher than their Galactic counterparts, which are in more metal-rich environments. Therefore the high dust temperature in the LMC can also be a possible cause for the enhanced production of the CO$_2$ ice. From the discussion, we conclude that the general properties of the LMC's environments, the strong UV radiation field, and the high dust temperature are responsible for the observed high CO$_2$ ice abundance.   One of the interstellar molecular features that likely develop from UV photolysis is the 4.62$\\mu$m ``XCN'' feature, which is probably due to the OCN$^{-}$ \\citep[e.g., ][]{Pen99,Ber00}. This feature provides a diagnostic of the local radiation field of objects \\citep{Spo03}. In our galaxy, the XCN feature is generally very weak or absent toward YSOs and dark clouds, but is very strong toward a few high-mass YSOs and Galactic center objects \\citep{Gib04}. Thus this feature could suggest that the LMC's strong UV radiation field is a dominant factor for the high CO$_2$ ice abundance in the LMC. Some of our samples show a hint of this XCN feature on the shorter side of the CO ice feature (ST6, 7, 10). There may also be an additional component seen in the spectrum of ST4 and ST8, but the peak wavelength seems to be slightly shorter than the XCN feature. No clear trends are seen in the CO$_2$/H$_2$O ratio of the objects with possible detection of the XCN. However, we cannot separate this feature from unresolved CO gas lines because the spectral resolution is not high enough, and the detection of XCN is still tentative, so that no definite conclusion can be drawn from the present spectra. Future observations should provide the needed resolution.     \\subsection{Correlation of ice abundance with YSO properties}  % As described in $\\S$2, the present samples basically possess similar characteristics (high-mass Class I objects). Despite this, column densities of the H$_2$O, CO$_2$ and CO ices and their ratio show large variations. To investigate which physical condition is dominant in the formation and evolution of ices, we  compare the chemical conditions of ices and the properties of individual YSOs. It is probable that the radiation from the central star dominates the circumstellar environment of a massive YSO. If the formation of ices is controlled by the radiation from the star, there should be some parameters of YSOs that correlate well with the ice abundances.  \\begin{figure*}[!htb] \\begin{center} \\includegraphics[width=15cm, angle=0]{13815fg6.eps} \\caption{The equivalent width of the hydrogen recombination line (Br$\\alpha$) vs. CO$_2$/H$_2$O ratio (upper left) and CO/H$_2$O ratio (upper right). And a total luminosity of sample YSOs vs. CO$_2$/H$_2$O ratio (bottom left) and CO/H$_2$O ratio (bottom right). A YSO number given in Table \\ref{tbl_NG} is plotted for the objects discussed in the text. A value of ST6 is not plotted due to its large uncertainties. ). } \\end{center} \\label{corl} \\end{figure*}  First, we compared the strength of the hydrogen emission lines with the CO$_2$ and CO ice abundance. Although it is not seen in the low-resolution spectra of \\citet{ST}, the present spectra clearly detect hydrogen recombination lines toward the target YSOs. It is known that massive YSOs in their more evolved stage ionize their circumstellar gas by stellar UV radiation and form compact HII regions around the central star \\citep{vanA00}. A variation in their hydrogen line strength thus indicates that the present targets are in their various evolutionary stages and possess different radiation environments (especially as for UV) in their circumstellar region, as discussed in \\citet{Sea09}. We measured the equivalent width of the Br$\\alpha$ line at 4.05$\\mu$m, which is the strongest hydrogen recombination line in the wavelength range of our observations.   Figure \\ref{corl} shows the comparison of the equivalent width of the Br$\\alpha$ line with the CO$_2$/H$_2$O and CO/H$_2$O ice ratio. It seems that the abundance of the CO$_2$ ice does not correlate with the strength of the Br$\\alpha$ line. The objects with similar CO$_2$ abundance show various Br$\\alpha$ line strengths. This result indicates that the UV radiation from the central star is a less dominant factor for the evolution of the CO$_2$ ice. However, in the previous section, we mentioned that the UV irradiation can help form CO$_2$ ice. Several authors have demonstrated by simple calculation that the interstellar radiation field can penetrate the site where the CO$_2$ ice forms \\citep[e.g.,][]{Whit98,Wat07}. We therefore suggest that the external (interstellar) UV radiation field may play an important role in CO$_2$ ice formation. On the other hand, it seems that objects with strong Br$\\alpha$ line strengths (ST2, ST11) show a small CO ice abundance (upper limit), and one with the weakest Br$\\alpha$ line (ST10) shows the highest CO ice abundance. Although the CO ice abundance of relatively weak Br$\\alpha$ objects is rather scattered, the result indicates that the CO ice around YSOs is more likely to be affected by the stellar radiation than the CO$_2$ ice.    A similar comparison was also made for the total luminosity of the YSOs, which is also a good estimate of the radiation environment of a YSO. The total luminosity was derived by the SED fitting described in $\\S$2. As seen in Fig.~\\ref{corl}, no clear correlation is seen between the CO$_2$ ice abundance and total luminosity of YSOs. The trend seen in the plot of the CO ice against the total luminosity is somewhat similar to what is seen in the comparison with the Br$\\alpha$ strength; i.e., the overall correlation is weak, but the most luminous objects in the present sample (ST2 and ST11) show a small abundance of CO.   The low CO ice abundance toward luminous and probably evolved (indicated by strong Br$\\alpha$ emission) YSOs can be explained by the low sublimation temperature of the CO ice. The sublimation temperature of the pure CO ice is lower than that of the pure H$_2$O and CO$_2$ ices \\citep[16K, 50K, and 90K, for CO, CO$_2$, and H$_2$O,][]{Tie05}. Although the sublimation temperature changes depending on the chemical compositions of ices, their general trend does not change. The CO ice is therefore more easily affected by the stellar radiation owing to its low sublimation temperature. Given a similar geometry and distance between the central star and the circumstellar dust, the dust around more luminous YSOs is expected to be warmer since the radiation from the central star is the dominant heating source of their circumstellar dust. Likewise, dust around more evolved objects is expected to be warmer thanks to the dissipation of the circumstellar dust. The relatively high luminosity of ST2 and ST11 suggests that these objects might be clusters. However, we regard them here as YSOs of a similar kind as in other samples, because their infrared SEDs are well-fitted by the single YSO model of \\citet{Rob07}, as well as in other samples. We thus suggest that the CO ice around the present luminous samples (ST2 and ST11) may be sublimated thanks to the intense radiation from the central star. The CO abundance of less luminous objects also has a scatter, which suggests that other factors than the stellar radiation may be important for the CO ice chemistry as for these objects."
    },
    "1004/1004.5200_arXiv.txt": {
        "abstract": "CCD $BVRI$ photometry is presented for type Ia supernova 2008gy. The light curves match the template curves for fast-declining SN Ia, but the colors appear redder than average, and the SN may also be slightly subluminous. SN 2008gy is found to be located far outside the boundaries of three nearest galaxies, each of them has nearly equal probability to be the host galaxy. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.3610_arXiv.txt": {
        "abstract": "Dawn is the first NASA mission to operate in the vicinity of the two most massive asteroids in the main belt, Ceres and Vesta. This double-rendezvous mission is enabled by the use of low-thrust solar electric propulsion. Dawn will arrive at Vesta in 2011 and will operate in its vicinity for approximately one year. Vesta's mass and non-spherical shape, coupled with its rotational period, presents very interesting challenges to a spacecraft that depends principally upon low-thrust propulsion for trajectory-changing maneuvers. The details of Vesta's high-order gravitational terms will not be determined until after Dawn's arrival at Vesta, but it is clear that their effect on Dawn operations creates the most complex operational environment for a NASA mission to date. Gravitational perturbations give rise to oscillations in Dawn's orbital radius, and it is found that trapping of the spacecraft is possible near the {\\tt 1:1} resonance between Dawn's orbital period and Vesta's rotational period, located approximately between 520 and 580 km orbital radius. This resonant trapping can be escaped by thrusting at the appropriate orbital phase. Having passed through the {\\tt 1:1} resonance, gravitational perturbations ultimately limit the minimum radius for low-altitude operations to about 400 km, in order to safely prevent surface impact. The lowest practical orbit is desirable in order to maximize signal-to-noise and spatial resolution of the Gamma-Ray and Neutron Detector and to provide the highest spatial resolution observations by Dawn's Framing Camera and Visible InfraRed mapping spectrometer. Dawn dynamical behavior is modeled in the context of a wide range of Vesta gravity models. Many of these models are distinguishable during Dawn's High Altitude Mapping Orbit and the remainder are resolved during Dawn's Low Altitude Mapping Orbit, providing insight into Vesta's interior structure. Ultimately, the dynamics of Dawn at Vesta identifies issues to be explored in the planning of future EP missions operating in close proximity to larger asteroids. ",
        "introduction": "The Dawn Discovery mission was successfully launched on September 27, 2007, and is the first NASA science mission making use of solar electric propulsion (EP), enabled by the earlier Deep Space 1 technology demonstration mission \\citep{1999BAAS...31.1127L}. As a consequence of the efficiency of this low-thrust, low-acceleration system, Dawn is able to rendezvous with the two most massive objects in the asteroid belt, Vesta then Ceres. These targets were selected in order to study the earliest stages of planetary evolution for an object that formed dry (Vesta) and another that formed with substantial amounts of water (Ceres) \\citep{2004P&SS...52..465R}.  Dawn's first target, Vesta, has a semi-major axis of 2.36 AU and is located in the inner main asteroid belt. It is unique in having a basaltic crust that has survived over the age of the solar system, providing important constraints on models of the collisional evolution of the asteroid belt \\citep[e.g.,][]{1985Icar...62...30D}. It has been spectroscopically linked to HED meteorites on the Earth \\citep{1970Sci...168.1445M}, which represent approximately 6\\% of all meteorite falls today \\citep{1999nssy.book..351M}. It has been inferred from those meteorites that Vesta is a differentiated object with an iron-rich core \\citep{1984LPI....15..603N,1996LPI....27..407G,1996Icar..124..513R}. Hubble observations of Vesta revealed an object with an equatorial radius around 289 km and polar radius of 229 km, but with a substantial impact crater covering much of its southern hemisphere and distorting its shape \\citep{1997Icar..128...83Z,1997Icar..128...88T,1997Sci...277.1492T}. This giant impact likely gave rise to the Vesta collisional family, which spans the inner main belt from the $\\nu_6$ secular resonance with Saturn on its inner edge to the {\\tt 3:1} mean motion resonance with Jupiter, separating it from the outer main belt. Some Vesta material entering these resonances would have their orbits pumped into Mars- and eventually Earth-crossing orbits, resulting in the HED meteorites recovered on the Earth \\citep[e.g.,][]{1997M&PS...32..903M}. Dawn may provide connections between specific areas of Vesta's surface and HED (and possibly other) meteorites.  Dawn executed a Mars gravity assist on February 17, 2009, to align its orbital inclination with that of Vesta. Dawn arrives at Vesta in July 2011 when it will enter an initial orbit having a radius of 2700 km (Survey orbit), from which it will obtain a preliminary shape model using the Framing Camera (FC) and spectrally map the entire illuminated surface using the Visible InfraRed mapping spectrometer (VIR) \\citep{2006AdSpR..38.2043R,2007EM&P..101...65R}. Assuming the rotational pole of \\cite{1997Icar..128...88T}, Vesta's obliquity is $27.2^\\circ$. Dawn arrives at Vesta at the time of maximum illumination of the southern hemisphere and its large crater. After completing the Survey orbit phase, Dawn uses its EP thrusters to descend to a High Altitude Mapping Orbit (HAMO) at approximately 950 km radius from which it will use the FC to map Vesta's surface and determine its global shape and local topography using stereophotoclinometric techniques \\citep[e.g.,][]{2008M&PS...43.1049G}. Dawn will then descend to its Low Altitude Mapping Orbit (LAMO) of around 460 km radius from which it will map Vesta's elemental composition using the Gamma-Ray Neutron Detector (GRaND). The Survey, HAMO, and LAMO phases are nominally 7, 27, and 90 days in duration \\citep{2007AdSpR..40..193R}, but this is subject to further planning for the 8 month stay at Vesta, which may be extended to a year. Because Dawn uses its EP thrusters for orbit transfers, transitions between these different phases are expected to take around a month \\citep{2007EM&P..101...65R}.  To maximize the science return from the mission, we are interested in determining the lowest orbital radius from which Dawn can safely execute LAMO. The closer to the surface we can make observations, the better the spatial resolution of FC and VIR observations. However, a lower LAMO most benefits observations by GRaND. The spatial resolution of the GRaND instrument is approximately 1.5 times the altitude \\citep{2003ITNS...50.1190P}. At nominal 460 km orbital radius for LAMO, this altitude will vary between 175 km near the equator and 231 km near the pole. By decreasing LAMO to 400 km, the number of GRaND resolved elements on Vesta increases by more than 50\\%, improving our ability to identify geochemical units and relate them to HED meteorites \\citep{2010LPI....41.2299P}. Going lower also improves GRaND signal-to-noise and may enable an accurate determination of Mg and Si, which are important discriminators among the various rock types expected on Vesta (Prettyman, priv.~comm., and \\citealt{2010LPI....41.2299P}). However, lower orbits and corresponding decreased orbital periods increases the need for desaturation of the spacecraft angular momentum, increasing the operational burden, which is not explored here.  A Vesta gravity model is also greatly improved with reduced altitude, allowing for better detection of mascons and determination of Vesta's higher order gravitational terms. At HAMO, the gravity field is determined to at least degree 4 and at LAMO this is expected to improve to at least degree 10 \\citep{2007EM&P..101...65R}. Depending on the accuracy of Doppler and Doppler-rate data, simulations later in this manuscript show that these numbers can be significantly improved, reaching degree 10 at HAMO and degree 20 at LAMO.  While there are science benefits from a minimum radius LAMO, Vesta's large mass and deviation from sphericity raises the question of how its gravity field will constrain the lowest orbital radius at which Dawn can safely operate. In addition, the long transfer times between Survey, HAMO and LAMO mean that Dawn will be slowly transiting commensurabilities between its orbital period and Vesta's rotational period, where perturbations on Dawn's orbit may be significant.  With this work we explore the dynamics of the Dawn spacecraft in a polar orbit within 1000~km from Vesta. The gravitational potential of Vesta is determined assuming diverse and extreme scenarios for its interior structure, to ensure the dynamical environment is sufficiently explored and that the results are representative of what the mission is likely to experience once there. The orbital maneuver from HAMO to LAMO using EP to slowly spiral in is also simulated, to assess the effect of mean motion resonances. ",
        "conclusions": "Vesta presents novel and interesting operational challenges to the Dawn mission. Because of the low-thrust electric propulsion system, Dawn will pass through Vesta spin/Dawn orbit commensurabilities very slowly as it moves from its High Altitude Mapping Orbit to Low Altitude Mapping Orbit, maximizing the effects of the perturbation. In addition to the rapid oscillations in Dawn's orbital radius as a consequence of Vesta's complex gravity field, there is the potential that Dawn could be trapped near the {\\tt 1:1} resonance, as it slowly decreases its radius through 550 km. Dawn can escape trapping by increasing thrust at the appropriate orbit libration phase. This is not an issue that is expected to recur at Ceres, which is observed to have a simple oblate spheroid shape. Once through this resonance, Dawn can continue to decrease its orbital radius to around 400 km (60 km less than the current LAMO) before the effect of perturbations begin to progressively increase orbital radial oscillations until impact with the surface becomes a hazard as Dawn reaches an average orbital radius of around 370 km."
    },
    "1004/1004.3426_arXiv.txt": {
        "abstract": "We consider the horizontal branch (HB) of the Globular Cluster Terzan~5, recently shown to be split into two parts, the fainter one ($\\delta M_K \\sim 0.3$mag) having a lower metallicity than the more luminous. Both features show that it contains at least two stellar populations. The separation in magnitude has been ascribed to an age difference of $\\sim$6~Gyr and interpreted as the result of an atypical evolutionary history for this cluster. We show that the observed HB morphology is also consistent with a model in which the bright HB  is composed of second generation stars that are metal enriched and with a helium mass fraction larger (by $\\delta Y \\sim$0.07) than that of first generation stars populating the fainter part of the HB. Terzan 5 would therefore be anomalous, compared to most ``normal\" clusters hosting multiple populations, only because its second generation is strongly contaminated by supernova ejecta; the previously proposed prolonged period of star formation, however, is not required. The iron enrichment of the bright HB can be ascribed either to contamination from Type Ia supernova ejecta of the low--iron, helium rich, ejecta of the massive asympotic giant branch stars of the cluster, or to its mixing with gas, accreting on the cluster from the environment, that has been subject to fast metal enrichment due to its proximity with the galactic bulge. The model here proposed requires only a small age difference, of $\\sim$100~Myr. ",
        "introduction": "\\label{sec:intro} Recent years have witnessed exciting developments both in the observations and theoretical modelling of the abundance star to star variations within most of the well studied Globular Clusters (GCs). In most GCs star to star abundance variations are limited to the light elements that are susceptible to abundance changes from proton-capture reactions, such as the pp, CN, ON, NeNa, and MgAl cycles \\citep[for the most recent spectroscopic survey, see, e.g.][]{carretta2009a,carretta2009b}, but a few clusters, such as M22 \\citep{marino2009}, or perhaps NGC 1851 \\citep{han2009} are now known to exhibit variations in heavier elements \\citep[see also][]{carretta2009ferro}, and, more than the others, in $\\omega$Cen the heavy elements spreads \\citep[e.g., among others][]{norrisdacosta1995}, and the HR diagram morphologies clearly show that we are dealing with several stellar generations, enriched by the supernova (SN) ejecta \\citep[e.g.][]{sollima2005, villanova2007}. In addition, the cluster M~54, immersed in the nucleus of the Sagittarius dwarf galaxy, presently disrupting in our Galaxy \\citep{ibata1994,bellazzini2008} has been recently found to show a metallicity spread similar to \\ocen\\ \\citep{carrettam542010}  Concerning the spread in light elements, their observation at the turnoff and among the subgiant stars \\citep[e.g.,][]{gratton2001} showed that these anomalies must be attributed to self--enrichment occurring at the first stages of the life of the cluster. The most peculiar finding of the latest years is the presence of a very helium rich population in the most massive clusters: this is revealed by the presence of multiple main sequences in \\ocen\\ and NGC~2808, indicating a helium content Y=0.38--0.40 \\citep{norris2004,dantona2005,piotto2007}, and by the extreme morphology of the horizontal branch (HB) of the two massive clusters  NGC~6388 and NGC~6441. In these latter clusters, a red clump is expected as HB, due to their large metallicity \\citep[{[Fe/H]$\\sim$--0.4}, see][]{carretta2009ferro}. On the contrary, their HB is extended towards the blue, and the RR Lyr variables have so long periods that they must be highly overluminous. \\cite{caloi2007} and \\cite{dc2008} show that the HB morphology and RR Lyr's periods of these two clusters may be explained if a large fraction of the HB stars have Y$\\simgt$0.35. The presence of much more moderate helium spreads is probable in most of the other smaller clusters \\cite[e.g.][]{dc2008}. The quasi--constancy of heavy metals in most GCs leads to hypothesize that the abundance variations must be due to very peculiar chemical evolution, not or scarcely affected by SN ejecta, but involving formation of a ``second generation\" (SG) of stars from matter processed into the ``first generation\" (FG) stars. On the other hand, the numerical consistency of the SG is so high \\citep[$>$50\\%][]{dc2008,carretta2009a} that any formation model must include the hypothesis that the mass contained in FG stars ---that contribute to this second phase of star formation--- is {\\sl initially} much larger than the mass present today in the cluster. Models for the formation of these multiple generations are still in their infancy. We can divide them into two main categories: the models in which clusters are born inside dwarf galaxies, so that the polluting matter on the forming GC comes from a much larger environment \\citep[e.g.][]{bekki2007}, and the models in which there is an initial cluster 10--20 times more massive than todays'. In the latter case, the SG forms from the ejecta of the FG stars mainly in the central cluster parts, and the first dynamical phases of evolution lead to a preferential loss of the FG stars \\citep{dercole2008}. In these models, it is very difficult to accomodate large age differences between the first and second generation stars.  Consequently, the recent observations of the color magnitude diagram features and chemistry of Terzan 5 may constitute a benchmark in our understanding of GC formation. In fact, \\cite{ferraro2009} show that the cluster HB stars are divided into two clumps separated by $\\delta$M$_K \\sim 0.3$mag, and that the more luminous stars have a much larger iron content ([Fe/H]$\\sim$+0.3$\\pm$0.1) with respect to the lower HB ([Fe/H]$\\sim$--0.2$\\pm$0.1). This result shows that the evolution of this cluster is atypical, and that matter forming the SG stars (populating the upper HB clump), has been affected by SN contamination, as it occurred in \\ocen. On the other hand, \\cite{ferraro2009}, comparing the HB data to stellar isochrones of the correct metallicity, conclude that the SG must be $\\sim$6~Gyr younger than the FG. This huge age difference is very difficult to be understood in any formation framework, and this would be the first evidence for such a young age among bulge stars and clusters \\citep[e.g.][]{feltzing2000, origlia2008}.  We re--examine the problem and show that the HB morphology can also be explained by two coeval populations having an helium difference of $\\delta$Y$\\sim 0.07$, thus reaching a value not as extreme as in the cases quoted above, so its formation does not present particular problems (Sect. \\ref{sec:disc}). At the super--solar metallicity of the SG of Terzan~5 the possible helium enrichment in the SG does not produce a blue extension of the HB. In Sect. \\ref{sec:disc} we discuss some possibilities for the chemical evolution of the cluster. We finally remark that the different space distribution of the two populations might imply a mass difference between them, and that further observations and dynamical modelling may allow to choose between models based on age-- or helium--differences. ",
        "conclusions": "We have shown that the split HB of Terzan~5 can be interpreted as due to two populations differing in helium content and metallicity, and not much different in age ($\\delta$(age)$\\simlt$100Myr). Massive AGB stellar models for the chemistry of the FG are compatible with the required helium enhancement, but we need that 1) either the AGB matter itself is strongly polluted by SNIa ejecta, before the second stage of star formation begins, or 2) the AGB matter is diluted with accreted gas, fastly processed to very high metallicity in the bulge stellar environment. This suggestion may help to understand the ``true\" birth of the double population of this cluster, maybe as a mix of age and helium difference in the subsequent star formation events.  We conclude by pointing out that the two alternative scenarios (age or helium difference) predict different values for the bright HB and faint HB masses. Specifically while in the merging scenario the younger age of the bright HB implies that this population would be $\\sim$0.25\\msun\\ more massive than the faint HB (in Fig.\\ref{f1}, the mass difference between points C and D), in the scenario proposed in this paper the two populations would be almost coeval and their red giant progenitors would have only a small mass difference (in the example of Fig. \\ref{f2}, M=0.996\\msun\\ for the bright HB and  M=0.979\\msun for the faint HB). Further dynamical modelling will help to shed further light on the plausibility of the two scenarios and on the possible dynamical histories leading to the observed differences in the spatial distribution of the two populations."
    },
    "1004/1004.4495_arXiv.txt": {
        "abstract": "{High deuterium fractionation is observed in various types of environment such as prestellar cores, hot cores and hot corinos. It has proven to be an efficient probe to study the physical and chemical conditions of these environments. The study of the deuteration of different molecules helps us to understand their formation. This is especially interesting for complex molecules such as methanol and bigger molecules for which it may allow to differentiate between gas-phase and solid-state formation pathways.} {Methanol exhibits a high deuterium fractionation in hot corinos. Since CH$_3$OH is thought to be a precursor of methyl formate  we expect that deuterated methyl formate is produced in such environments. We have searched for the singly-deuterated isotopologue of methyl formate, DCOOCH$_3$, in IRAS 16293-2422, a hot corino well-known for its high degree of methanol deuteration. } {We have used the IRAM/JCMT unbiased spectral survey of IRAS 16293-2422 which allows us to search for the DCOOCH$_3$ rotational transitions within the survey spectral range (80-280 GHz, 328-366 GHz). The expected emission of deuterated methyl formate is modelled at LTE and compared with the observations.} {We have tentatively detected DCOOCH$_3$ in the protostar IRAS 16293-2422. We assign eight lines detected in the IRAM survey to DCOOCH3. Three of these lines are affected by blending problems and one line is affected by calibration uncertainties, nevertheless the LTE emission model is compatible with the observations.  A simple LTE modelling of the two cores in IRAS 16293-2422, based on a previous interferometric study of HCOOCH$_3$,  allows us to estimate the amount of DCOOCH$_3$ in IRAS 16293-2422.  Adopting an excitation temperature of 100 K and a source size of 2\\arcsec and 1\\farcs5  for  the A and B cores, respectively,  we find that N$_\\mathrm{A, DCOOCH3}$ = N$_\\mathrm{B, DCOOCH3}$ $\\sim$6 $\\times$ 10$^{14}$ cm$^{-2}$. The derived deuterium fractionation is $\\sim$ 15 \\%, consistent with values for other deuterated species in this source and much greater than that expected from the deuterium cosmic abundance. } { DCOOCH$_3$, if its tentative detection is confirmed, should now be considered in theoretical models that study complex molecule formation and their deuteration mechanisms. Experimental work is also needed to investigate the different chemical routes leading to the formation of deuterated methyl formate.} ",
        "introduction": "IRAS 16293-2422 (hereafter IRAS 16293) is a complex source hosting two hot corinos (called \"A\" and \"B\") in which many complex organic molecules have been observed: acetonitrile (CH$_3$CN), methyl formate (HCOOCH$_3$), ketene (H$_2$CCO) formic acid (HCOOH), ethanol (C$_2$H$_5$OH), ethyl cyanide (C$_2$H$_5$CN), etc. \\citep{cazaux2003,bottinelli2004,bisschop2008}. This source is also characterised by a large degree of deuteration with singly, doubly or triply deuterated molecules. HDCO, CH$_3$OD, CH$_2$DOH,  are detected in IRAS 16293 with a deuterium fractionation of $\\sim$ 15 \\%, 1.8 \\% and 37 \\%, respectively, doubly-deuterated formaldehyde (D$_2$CO) and methanol (CHD$_2$OH) are observed with a deuteration of 5 and 7.4 \\% \\citep{parise2006}. Triply deuterated methanol was also observed in IRAS 16293 with a deuterium fractionation ratio of 1.4 \\% \\citep{parise2004}.  These values are much greater than expected from the cosmic D/H ratio ($\\sim$ 10$^{-5}$). The process of deuteration is a long-standing and challenging issue. High deuterium fractionation is observed in cold regions such as pre-stellar cores \\citep{pillai2007,bacmann2003}, in warmer regions such as hot cores (eg. Orion, \\citealt{jacq1993}) and hot corinos  and in photo dissociation regions (PDRs) \\citep{leurini2006,pety2007}. The detection of deuterated molecules allows one to study the chemical pathways leading to their formation and to trace the chemical and physical conditions of the observed environments. Whereas in cold environments the deuterium molecular enrichment is explained by gas phase chemistry \\citep{roberts2003}, in hot regions it is certainly driven by chemistry on the dust icy surface \\citep{tielens1983,ceccarelli2001}. This is well illustrated by the detection of multiply deuterated methanol since methanol cannot be efficiently formed in the gas-phase \\citep{turner1998}. Since methanol is thought to be involved in methyl formate formation on grains \\citep{bennett2007}, the high deuterium fractionation of CH$_3$OH suggests that deuterated methyl formate should be produced together with methyl formate.  In this study we present the tentative detection of deuterated methyl formate, DCOOCH$_3$, in IRAS 16293 based on observations from the millimeter/submillimeter spectral survey performed toward the low-mass Class 0 protostar IRAS16293 \\citep{caux2010}. The observations are described in Sect. \\ref{obs}. The tentative detection of DCOOCH$_3$ is presented in Sect. \\ref{res} and its formation is discussed in Sect. \\ref{discussion}. ",
        "conclusions": "We have tentatively detected the singly-deuterated isotopologue of methyl formate, DCOOCH$_3$, in the protostar IRAS16293. We have assigned eight observed lines to DCOOCH$_3$ transitions. Four lines are not blended with other species (229.448, 239.654, 249.860, 260.064 GHz). Three lines are blended (209.031, 219.242, 270.267 GHz) with other species. Among the latter, the first two lines are blended with HCOOCH$_3$ and CH$_3$COCH$_3$, respectively. They are well-reproduced by an emission model at LTE. We were not able to identify the species responsible for the blending of the 270.267 GHz line however the LTE model is compatible with  the observations and does not contradict the tentative identification of DCOOCH$_3$ in IRAS 16293. The LTE emission model is compatible with the last observed line at 198.824 GHz but we cannot conclude on its detection because of calibration uncertainty. From a basic modelling in LTE we estimate the abundance of deuterated methyl formate to be N$_\\mathrm{A+B, DCOOCH3}$ $\\sim$1.2 $\\times$ 10$^{15}$ cm$^{-2}$. This leads to a DCOOCH$_3$/HCOOCH$_3$ ratio of  $\\sim$ 15 \\%, consistent with the deuteration fractionation of other singly deuterated species in this source.  Additional observations with better spectral resolution and higher sensitivity are needed for several reasons. First, considering the low spectral resolution and the relatively small signal-to-noise ratio of the lines assigned to DCOOCH$_3$, such new observations would strengthen this tentative detection. Second, it would allow to detect additional lines among the hundreds weaker lines from DCOOCH$_3$ that are present in the spectral survey. Lastly, it would be the opportunity to search for other deuterated isotopologues of methyl formate. High angular resolution observations would also be helpful to investigate the existence of two emission sources for DCOOCH$_3$ as it is observed for HCOOCH$_3$ and to conclude about the reality of the physical differences between the two cores A and B. Such observations are also necessary to properly estimate the methyl formate deuteration and, together with new experimental and theoretical work, to understand the deuteration mechanisms of complex organic molecules."
    },
    "1004/1004.2479.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) {To study the high-mass star-forming process, we have started a large project to unveil the gas kinematics close to young stellar objects (YSOs) through the Very Long Baseline Interferometry (VLBI) of maser associations. By comparing the high spatial resolution maser data, tracing the inner kinematics of the (proto)stellar cocoon, with interferometric thermal data, tracing the large-scale environment of the hot molecular core (HMC) harbouring the (proto)stars, we can investigate the nature and identify the sources of large-scale motions. The present paper focuses on the high-mass star-forming region \\G16.} % methods heading (mandatory) {Using the VLBA and the EVN arrays, we conducted phase-referenced observations of the three most powerful maser species in \\G16: H$_2$O at 22.2~GHz (4 epochs), CH$_3$OH at 6.7~GHz (3 epochs), and OH at 1.665~GHz (1 epoch). In addition, we performed high-resolution ($\\geq0\\farcs1$), high-sensitivity ($<0.1$~mJy) VLA observations of the radio continuum emission from the star-forming region at 1.3 and 3.6~cm.} % results heading (mandatory) {This is the first work to report accurate measurements of the \\emph{relative} proper motions of the 6.7~GHz CH$_3$OH masers. The different spatial and 3-D velocity distribution clearly indicate that the 22~GHz water and 6.7~GHz methanol masers are tracing different kinematic environments.  The bipolar distribution of 6.7~GHz maser l.o.s. velocities and the regular pattern of observed proper motions suggest that these masers are tracing rotation around a central mass of about 35~M$_{\\odot}$. The flattened spatial distribution of the 6.7~GHz masers, oriented  NW--SE, suggests that they can originate in a disk/toroid rotating around the massive YSO which drives the $^{12}$CO~(2--1) outflow, oriented NE--SW, observed on arcsec scale. The extended, radio continuum source observed close to the 6.7~GHz masers could be excited by a wide-angle wind emitted from the YSO associated with the methanol masers, and such a wind is proven to be sufficiently energetic to drive the NE--SW $^{12}$CO~(2--1) outflow. The H$_2$O masers distribute across a region offset about 0\\farcs5 to the NW of the CH$_3$OH masers, in the same area where emission of high-density molecular tracers, typical of HMCs, was detected. We postulate that a distinct YSO, possibly in an earlier evolutionary phase than that exciting the methanol masers, is responsible for the excitation of the water masers and the HMC molecular lines.} % conclusions heading (optional), leave it empty if necessary {} ",
        "introduction": "The process of massive star formation is based on a fine tuning between the gravitational force, that sets in the collapse of a Jeans-critical cloud, against the different types of pressures, thermal, magnetic, turbulence, and radiation, that regulate the time scale of the accretion (for recent reviews see, e.g., \\citealt{Zinnecker2007,Beuther2007}). To characterize the relationship between accretion-ejection phenomena it is essential to resolve the dynamical structures close to the central engine. Maser emission from several molecular lines is an useful signpost of the hot, dusty environment (i.e. hot molecular core, HMC) where massive star formation takes place, which, because of the extremely high dust extinction, could not be investigated with optical and near-infrared (NIR) diagnostics. Maser emission is concentrated in narrow (usually $\\lesssim 1$~\\kms\\ broad), strong (up to 10$^6$~Jy) lines and arises from compact emission centers (``maser spots'') which have typical sizes of a few AU (e.g., \\citealt{Minier2002,Moscadelli2003}). The strong brightness of the maser emission allows us to use the Very Long Baseline Interferometry (VLBI) technique to determine the position of the maser spots with milli-arcsec accuracy, and comparing positions at different VLBI epochs gives accurate measurement of the spot proper motions. Combining proper motions with the maser line-of-sight (l.o.s) velocities (derived via Doppler shift of the observed maser frequency),  one can reconstruct the full 3-D kinematics of the masing gas. With this in mind, we have started a large project whose final aim is drawing a comprehensive picture of the dynamical processes in high-mass star-forming regions (HMSFRs), by relating the pc-scale phenomena, traced by thermal continuum and molecular lines, with the AU-scale kinematics of the gas around young stellar objects (YSOs), traced by maser emissions.  In particular, observing targets where different molecular maser species appear to be spatially associated, we can hope to better sample the (proto)stellar environment (e.g., \\citealt{Goddi2007, Moscadelli2007}). As a first step, we focus on the three most powerful maser transitions, that of water (H$_2$O) at 22.2~GHz, methanol (CH$_3$OH) at 6.7~GHz, and hydroxyl (OH) at 1.665~GHz. H$_2$O masers typically trace fast shocks (of several tens of \\kms) excited by either collimated jets found at the base of molecular outflows, or wide-angle winds emitted from the YSOs (e.g., \\citealt{Gwinn1992, Torrelles2003, Moscadelli2007}). OH masers are likely associated with the slow expanding (typically a few \\kms) ionization front of H~\\textsc{ii} regions (e.g., \\citealt{Fish2007}). The origin of CH$_3$OH masers, instead, is currently a matter of debate. To date, single-epoch VLBI observations have provided accurate spatial and l.o.s. velocity distribution of the 6.7~GHz masers towards a few tens of sources. To account for the ordered gradient of l.o.s. velocities observed in some objects, three hypotheses have been proposed: circumstellar (Keplerian) discs seen edge-on (e.g., \\citealt{Norris1998, Minier2000}); outflows (e.g., \\citealt{DeBuizer2003});  propagating shock fronts through rotating dense cores and ring-like structures (e.g., \\citealt{Bartkiewicz2009}). The only way to discriminate among these different interpretations is deriving maser proper motions via multi-epoch VLBI observations.    The present paper focuses on the HMSFR \\object{G16.59$-$0.05}. In Sect.~2, we provide an up-to-date review of the single-dish and interferometric observations towards this region. Section~3 describes our VLBI observations of the 22.2~GHz H$_2$O, 6.7~GHz CH$_3$OH, and 1.665~GHz OH masers, together with the new Very Large Array (VLA) A-configuration observations of the radio continuum emission at 1.3 and 3.6~cm, which complemented VLA-C archival data at 0.7, 1.3, and 3.6~cm. Details of our data analysis are given in Sect.~4. In Sect.~5, we illustrate the spatial morphology, kinematics, and time-variability of individual maser species, and present results from our and archival VLA observations, constraining the properties of the radio continuum observed in correspondence of the masers. %\\emph{Noticeably, here, we report the first multi-epoch VLBI analysis of the proper motions of 6.7~GHz CH$_3$OH masers.} Section~6 discusses the spatial association of the maser species and their overall kinematics, and draws a comprehensive picture of the phenomena observed in the HMSFR \\G16 on angular scales from a few mas to tens of arcsec.  The main conclusions are summarized in Sect.~7. ",
        "conclusions": "Using the VLBA and the EVN interferometers, we observed the high-mass star-forming region \\G16 in the three most powerful maser transitions: 22.2~GHz H$_2$O, 6.7~GHz CH$_3$CH, and 1.665~GHz OH. The radio continuum emission toward \\G16 was also observed with the most extended VLA configuration to compare its brightness structure with the VLA--C archival data available. From our observations, we draw the following main conclusions:  \\begin{enumerate}   \\item In the present work, we have collected evidence that methanol masers, alike water emission, are suitable tracers of the kinematics around massive young stellar objects. With three VLBI epochs over a time baseline of 2~yr, we have measured accurate ($< 30\\%$) proper motions of the 6.7~GHz masers. The l.o.s and sky-projected velocity distribution of methanol masers indicates that they are rotating with an average velocity of about 7~\\kms\\ at distances of about 600~AU from a YSO. The inferred dynamical mass (35~M$_{\\odot}$) is consistent with 6.7~GHz masers tracing a high-mass YSO. The elongated NW--SE distribution of 6.7~GHz masers suggests they can originate in a flattened structure (disk/toroid) and that the YSO at the center of their distribution can be the one driving the motion of the NE--SW $^{12}$CO~(2--1) outflow observed at larger angular scale.  \\item Close to the 6.7~GHz masers, a compact (or slightly resolved) continuum source is observed with the VLA--C at 3.6~and~1.3~cm, and at 7~mm. We interpret this radio continuum in terms of free-free emission of gas ionized by a wide-angle wind emitted by the YSO associated with the 6.7~GHz masers and interacting with the surrounding dense gas. The momentum rate of this  wind is consistent with that measured for the  NE--SW $^{12}$CO~(2--1) outflow.  \\item Water masers are found offset by more than 0\\farcs5 from the center of the 6.7~GHz maser distribution, in the same region where a weak 7~mm VLA source and emission in high-density molecular lines are observed. The derived absolute velocities of the water masers, although affected by large errors, seem to indicate fast motions with an average amplitude of about 60~\\kms. We postulate the presence in this region of a distinct YSO, responsible for driving the motion of water masers and exciting  the continuum and molecular line emissions.   \\end{enumerate}  Previous SMA observations in several molecular lines have revealed multiple molecular outflows from the \\G16 star-forming region, suggesting the presence of several massive YSOs across a region of $\\approx0.1$~pc. This work shows that maser VLBI data are a powerful tool to identify the massive YSOs driving the large scale outflows, to study the gas kinematics close to the forming stars, and to detect disks/toroids of size of hundreds AU at typical distances of several kiloparsecs. In the next future, when VLBI maser observations will be complemented with data of new generation (sub)millimeter interferometers (i.e. ALMA) at comparable angular resolution, it will be possible to better constrain the physical and kinematical properties of both outflow(s) and core(s) hosting maser activity."
    },
    "1004/1004.3574_arXiv.txt": {
        "abstract": "{In general relativity coupled to Maxwell's electromagnetism and charged matter, when the gravitational potential $W^2$ and the electric potential field $\\phi$ obey a relation of the form $W^{2}= a\\left(-\\epsilon\\, \\phi+ b\\right)^2 +c$, where $a$, $b$ and $c$ are arbitrary constants, and $\\epsilon=\\pm1$ (the speed of light $c$ and Newton's constant $G$ are put to one), a class of very interesting electrically charged systems with pressure arises. We call the relation above between $W$ and $\\phi$, the Weyl-Guilfoyle relation, and it generalizes the usual Weyl relation, for which $a=1$. For both, Weyl and Weyl-Guilfoyle relations, the electrically charged fluid, if present, may have nonzero pressure. Fluids obeying the Weyl-Guilfoyle relation are called Weyl-Guilfoyle fluids.  These fluids, under the assumption of spherical symmetry, exhibit solutions which can be matched to the electrovacuum Reissner-Nordstr\\\"om spacetime to yield global asymptotically flat cold charged stars. We show that a particular spherically symmetric class of stars found by Guilfoyle has a well-behaved limit which corresponds to an extremal Reissner-Nordstr\\\"om quasiblack hole with pressure, i.e., in which the fluid inside the quasihorizon has electric charge and pressure, and the geometry outside the quasihorizon is given by the extremal Reissner-Nordstr\\\"om metric. The main physical properties of such charged stars and quasiblack holes with pressure are analyzed.  An important development provided by these stars and quasiblack holes is that without pressure the solutions, Majumdar-Papapetrou solutions, are unstable to kinetic perturbations. Solutions with pressure may avoid this instability. If stable, these cold quasiblack holes with pressure, i.e., these compact relativistic charged spheres, are really frozen stars.}  \\pacs{04.40.Nr,04.20.Jb, 04.70.Bw} ",
        "introduction": "\\label{sec-basicequations}  The cold charged fluids considered in the present work are described by Einstein-Maxwell equations, which can be written as \\begin{eqnarray} & &G_\\munu=  8\\pi \\, \\left( T_\\munu+ E_\\munu\\right)\\, , \\label{einst}\\\\ & & \\nabla_\\nu F^\\munu = 4\\pi\\,J^\\mu\\,, \\label{maxeqs} \\end{eqnarray} where Greek indices $\\mu, \\nu$, etc., run from $0$ to $3$. $G_\\munu=R_\\munu-\\frac{1}{2}g_\\munu R$ is the Einstein tensor, with $R_\\munu$ being the Ricci tensor, $g_\\munu$ the metric tensor, and $R$ the Ricci scalar. We have put both the speed of light $c$ and the gravitational constant $G$ equal to unity throughout. $E_\\munu$ is the electromagnetic energy-momentum tensor, which can be written in the form \\begin{equation} 4\\pi\\, E_\\munu= {F_\\mu}^\\rho F_\\nu{_\\rho} -\\frac{1}{4}g_\\munu F_{\\rho\\sigma} F^{\\rho\\sigma}\\, ,\\label{maxemt} \\end{equation} where the Maxwell tensor is \\begin{equation} F_\\munu = \\nabla_\\mu {\\cal A}_\\nu -\\nabla_\\nu {\\cal A}_\\mu\\, , \\label{ddemfield} \\end{equation} $\\nabla_\\mu$ representing the covariant derivative, and ${\\cal A}_\\mu$ the electromagnetic gauge field. In addition, \\begin{equation} J_\\mu = \\rho_{\\rm e}\\, U_\\mu\\, ,\\label{current} \\end{equation} is the current density, $\\rho_{\\rm e}$ is the electric charge density, and $U_\\mu$ is the fluid velocity. $T_\\munu$ is the material energy-momentum tensor given by \\begin{equation} T_\\munu = \\left(\\rho_{\\rm m}+p\\right)U_\\mu U_\\nu +p  g_\\munu \\, ,\\label{fluidemt} \\end{equation} where $\\rho_{\\rm m}$ is the fluid matter-energy density, and $p$ is the fluid pressure.   We assume the spacetime is static and the metric \\begin{equation} ds^2 = g_{\\mu\\nu}\\,dx^\\mu\\,dx^\\nu\\, \\label{metrico} \\end{equation} can be written in the form \\begin{equation} ds^2 = - W^2 dt^2 + h_{ij}\\,dx^idx^j\\, , \\quad\\quad i,\\,j=1,\\, 2,\\, 3\\,  . \\label{metricg} \\end{equation} The gauge field ${\\cal A}_\\mu$ and four-velocity $U_\\mu$ are then given by \\beqa & &{\\cal A}_\\mu = -\\phi\\,\\delta_\\mu^0\\, ,\\label{gauge1}\\\\ & &U_\\mu =  -W\\, \\delta_\\mu ^0\\, .\\label{veloc1} \\eeqa The spatial metric tensor $ h_{ij}$, the metric potential $B$ and the electrostatic potential $\\phi$ are functions of the spatial coordinates $x^i$ alone.  The particular relativistic cold charged stars we are going to study here belong to a special class of systems in which the metric potential $W$ and the electric potential $\\phi$ are functionally related through a Weyl-Guilfoyle relation (see Guilfoyle \\cite{guilfoyle}, see also Lemos and Zanchin \\cite{lemoszanchin2009}) \\begin{equation} W^2= a \\left(-\\epsilon\\,\\phi+b\\right)^2+c \\, , \\label{weylgrel1} \\end{equation} with $a$, $b$ and $c$ being arbitrary constants and $\\epsilon=\\pm 1$, the parameter $a$ being called the Guilfoyle parameter. Then, the charged pressure fluid quantities $\\rho_{\\rm m}$, $p$ and $\\rho_{\\rm e}$ satisfy the constraint \\cite{lemoszanchin2009} \\beq a\\,b\\, \\rho_{\\rm e}=\\epsilon \\left[\\,\\rho_{\\rm m}+ 3p+\\left(1-a\\right)\\, \\rho_{\\rm em}\\right]W +\\epsilon a\\, \\rho_{\\rm e}\\phi \\,, \\label{wgcondition1} \\eeq or \\beq a\\, \\rho_{\\rm e}\\left(-\\epsilon\\,\\phi +b\\right)  = \\epsilon\\, \\left[\\rho_{\\rm m}+ 3 p+ \\epsilon\\left(1-a\\right) \\rho_{\\rm em}\\,\\right] W \\,, \\label{qgcondition2} \\eeq with $\\rho_{\\rm em}$ standing for the electromagnetic energy density defined by \\beq \\rho_{\\rm em } = \\frac{1}{8\\pi}\\frac{ \\left(\\nabla_i\\phi\\right)^2}{W^2}\\,. \\label{emdensity} \\eeq Such matter systems we call Weyl-Guilfoyle fluids. ",
        "conclusions": "\\label{sec-conclusion}  We have analyzed the class Ia of solutions provided by Guilfoyle \\cite{guilfoyle}.  Such spherically symmetric relativistic charged fluid distributions are bounded by a surface of radius $r_0$.  The interior region is filled with a fluid characterized by its mass and charge densities and by a nonzero pressure.  The spacetime in the exterior region is represented by the Reissner-Nordstr\\\"om metric. These global solutions represent relativistic stars, i.e., relativistic cold charged spheres with pressure.  Besides the mass $m$ (or charge $q$, which are related to each other) and the radius of the star $r_0$, this class of solutions is characterized by another free parameter, the Guilfoyle parameter $a$. This parameter is related to the pressure: for $a<1$ the stars are supported by tension; for $a=1$ the stars have no pressure (they are Bonnor stars), and for $a>1$ the stars are supported by pressure.  The interval of the free parameter $a$ can be fixed in such a way that the fluid satisfies the energy conditions, and other physical requirements for a relativistic cold star.  We have then studied relativistic stars within the interval $1<a\\leq 4$.  We have found that these cold stars show a mixed behavior, in one instance they behave as main sequence stars in another instance as white dwarfs.  Indeed, the mass to radius relation of these cold stars is analogous to the behavior of the mass to radius relation in main sequence stars and contrary to the mass to radius relation in white dwarfs, whereas the central pressure of these cold stars has an analogous behavior to the central pressure of white dwarfs, and a contrary behavior to the central pressure of main sequence stars.  We have also shown that, in the interval $1< a\\leq 4$, the most compact configuration is a quasiblack hole with pressure. Thus, quasiblack holes without pressure ($a=1$) studied previously, as well as quasiblack hole with pressure ($1< a\\leq 4$) studied here, can be found within the context of general relativity. As the most compact configuration is a cold star it can be called a frozen star."
    },
    "1004/1004.2869_arXiv.txt": {
        "abstract": "{We present an analysis of the variability of the solar oscillation spectrum during solar cycle 23 and its extended minimum. We use simultaneous observations of the low-degree solar p modes collected by the space-based, Sun-as-a-star GOLF (radial velocity) and VIRGO (intensity) instruments, and by the ground-based, multi-site network GONG. We investigate in particular the response of the p-mode eigenfrequencies to the observed peculiar deep solar minimum of surface activity of 2007--2009 as compared with the previous solar cycle 23. We study the different temporal variations of the p-mode frequencies with individual angular degrees. } ",
        "introduction": "The {\\it new millennium solar activity minimum} of 2007--2009 has shown the quietest Sun in almost a century with a delayed onset of solar cycle~24. Helioseismic observations ha\\-ve been used to study the response of the solar oscillations to this unusually extended minimum of solar surface activity. While the temporal variations of the p-mode frequencies are closely correlated with solar surface activity proxies during the past cycles at low-  \\citep[e.g.,][]{gelly02,salabert04} and higher-angular degrees \\citep[e.g.,][]{chano01,salabert06a}, significant variations of the p-mode frequencies during the minimum of cycle~23 in contrast to the surface activity observations over the same period have been reported \\citep{salabert09, broom09}. Furthermore, \\citet{howe09} showed that the lack of sunspots and the low-activity levels during the cycle~23 minimum can be explained by a slower than usual jet stream associated to the production of the sunspots. These streams originating from the poles every 11 years migrate slowly below the surface towards the equator. Other p-mode parameters, such as the mode power and lifetime for instance, were also proven to be sensitive to the solar activity cycle in both Sun-as-a-star \\citep[e.g.,][]{salabert03, chano04} and spatially-resolved \\citep[e.g.,][]{komm00,salabert06b} observations. Thus, their response to the unusually long minimum of cycle~23 is also of great interest and is currently under investigation.  \\begin{figure*}[t] \\centering \\includegraphics[scale=0.28]{golf_l0_all_proc}\\includegraphics[scale=0.29]{golf_l1_all_proc}\\includegraphics[scale=0.30]{golf_l2_all_proc} \\caption{Frequency shifts of the $l = 0, 1,$ and 2 solar p modes (left to right panels) extracted from the analysis of the 365-day (black dots) and the 91.25-day (blue plus signs) GOLF spectra. The associated error bars of the 365-day frequency shifts are represented. The corresponding 10.7-cm radio flux averaged over the same 91.25-day timespan is shown as a proxy of the solar surface activity (red solid line). } \\label{fig:fig1} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[scale=0.75]{virgo_proc} \\caption{Frequency shifts of the $l = 0, 1,$ and 2 solar p modes (left to right panels) extracted from the analysis of the 365-day VIRGO spectra. The frequency shifts measured from the blue, green, and red VIRGO channels are represented from top to bottom respectively. The corresponding GOLF frequency shifts are shown for comparison (black dots). The associated error bars are also represented.} \\label{fig:fig2} \\end{figure*}   \\begin{figure*} \\centering \\includegraphics[scale=0.295]{gong_l0_proc}\\includegraphics[scale=0.295]{gong_l1_proc}\\includegraphics[scale=0.295]{gong_l2_proc}\\caption{Frequency shifts of the $l = 0, 1,$ and 2 solar p modes (left to right panels) extracted from the analysis of the 365-day integrated GONG spectra (red diamonds). The corresponding GOLF frequency shifts are shown for comparison (black dots). The associated error bars are also represented. } \\label{fig:fig3} \\end{figure*} ",
        "conclusions": "We analyzed simultaneous observations of the low-degree solar p modes collected by the space-based GOLF (radial velocity) and VIRGO (intensity) instruments on board the {\\it SoHO} spacecraft, and by the ground-based, multi-site network GONG in order to investigate the response of the p-mode frequencies to the unusual deep and long minimum of solar surface activity of cycle 23 during 2007--2009. We observed different temporal variations of the p-mode frequencies between individual angular degrees. These variations are identical in the GOLF, VIRGO, and GONG instruments. The differences between individual angular degrees might be interpreted as different geometrical responses to the spatial distribution of the solar magnetic field beneath the surface of the Sun. A more detailed intercomparison between several helioseismic instruments is underway.  After CoRoT revealed a first stellar activity cycle in a Sun-like star using asteroseismology \\citep{garcia10}, the NASA {\\it Kepler} mission \\citep{koch10} will reinforce the study of stellar magnetic cycles \\citep{karoff09} thanks to long-duration observations (over 3 years) with a very high signal-to-noise ratio \\citep[e.g.,][]{chaplin10,bedding10}. Thus, seismology will provide very detailed inferences of the stellar structures, in particular the depth of the convective zones. Such information will be of key importance to properly understand the physical mechanisms driving the magnetic cycles and thus, will help to better understand and predict the dynamo processes taking place in the Sun."
    },
    "1004/1004.0577_arXiv.txt": {
        "abstract": "We present X-ray emission characteristics of the massive O-type stars DH Cep and HD 97434 using archival XMM-Newton observations. There is no convincing evidence for short term variability in the X-ray intensity during the observations. However, the analysis of their spectra reveals X-ray structure being consistent with two-temperature plasma model. The hydrogen column densities derived from X-ray spectra of DH Cep and HD 97434 are in agreement with the reddening measurements for their corresponding host clusters NGC 7380 and Trumpler 18, indicating that the absorption by stellar wind is negligible. The X-ray emission from these hot stars is interpreted in terms of the standard instability-driven wind shock model. ",
        "introduction": "\\label{sec:intro}  Massive, hot and luminous O-type stars have powerful stellar winds, with mass loss rates of the order of 10$^{-6}$ to 10$^{-4 }$ M$_\\odot$ yr$^{-1}$ and velocities of several thousand km s$^{-1}$. These winds play a crucial role both in the stellar evolution as well as the galactic evolution. The existence of strong winds is supposed to be a primary reason for X-ray emission in the early type stars. Study of the X-ray emission from O-type stars is important for understanding the high energy processes that generate hot gas in the wind or in the  atmosphere of hot stars.  X-ray emission from O-type stars  was first discovered serendipitously with the EINSTEIN observatory (Harnden et al. 1979).  Most O-type stars are reasonably bright ($\\rm{10^{31} < L_{X} < 10^{34}~erg~s^{-1}}$ ) and have a soft thermal spectrum with kT $<$ 1 keV. The X-ray luminosity of these sources is found to scale with their bolometric luminosity as $\\rm{L_X /L_{bol}} \\sim 10^{-6}$ to 10$^{-8}$ (Long \\& White 1980, Pallavicini et al. 1981, Singh \\& Naranan 1982, Bergh$\\rm{\\ddot{o}}$fer et al. 1997; Sana et al. 2006). It is generally thought to be produced as a result of shocks, with velocity jumps up to a few hundred km s$^{-1}$ , generated throughout the stellar wind because of dynamic instabilities (Lucy \\& White 1980; Owocki \\& Cohen 1999; Kudritzki \\& Puls 2000).  This scenario is now called a standard model for X-ray emission from early-type stars (Rauw 2008).  Although, most of the X-ray emission is dominated by soft X-rays as expected from wind-shock model (e.g. Feldmeier, Puls \\& Pauldrach 1997), a significant and an unexpected harder component is also seen in the X-ray spectra (e.g. Leutenegger \\& Khan 2003; Bross et al. 2007; Stelzer et al. 2005).  High quality observational data from XMM-Newton and Chandra observatories are indeed crucial to provide a strong test of the latest models. In order to understand the X-ray emission processes from O-type stars, we have analyzed XMM-Newton data for two O-type stars DH Cep and HD 97434. The results of our analysis are presented here.  DH Cep is an O6 V + O7 V spectroscopic binary system and a member of the cluster NGC 7380 (distance = 3.7 kpc, age = 10 Myr; Underhill 1969; Massey, Johnson \\& De Gioio-Eastwood 1995; Hilditch, Harries, \\& Bell 1996).  It has an orbital period of 2.11 d (Penny, Gies \\& Bagnuolo 1997). A low signal, Einstein X-ray observation is reported by Chlebowski, Harden \\& Sciortino (1989) with $\\rm{L_X\\sim5.37\\times10^{33}~erg~s^{-1}}$ in the energy band 0.2--3.5 keV.  HD 97434 is an O9 V star (Skiff 2007), with a membership probability of 78\\% (Dias et al. 2006) in the cluster Trumpler 18 (distance = 1.5 kpc; age = 15 Myr; V$\\rm{\\acute{a}}$zquez \\& Feinstein 1990). The Einstein X-ray data for the source are cataloged by Chlebowski, Harden \\& Sciortino (1989) with $\\rm{L_X<3.4\\times10^{32}~erg~s^{-1}}$ (corrected for the distance of 1.5 kpc) in the energy band 0.2--3.5 keV. ",
        "conclusions": "\\label{sec:discuss}  \\subsection{X-ray Variability} The wind-shock model predicts a significant short-term variability in the X-ray flux as the shocks fade and grow on short time scales. However, it has not been observed so far.  We have also explored the possibility of short-term X-ray variability of DH Cep and HD 97434 over the observation time span ($<$40 ks).  The timing analysis reveals a lack of short term variability in the X-ray intensity for both the massive stars. The lack of short term variability could be due to the wind fragmentation, so that individual X-ray fluctuations are smoothed out over the whole emitting volume, leading to a rather constant output ( Feldmeier, Puls \\& Pauldrach 1997).  Several other massive  stars, e.g., HD 159176 (De Becker et al. 2004) and HD 47129 (Linder et al. 2006) have also shown lack of variability in their light curves.  The longer-term variability can not be probed using present data.  The long-term variability (few days) from O-type stars was reported by Snow, Cash \\& Grady (1981) for the first time, but their results were  based on the low sensitivity detectors available at that time. No significant variability was detected by the ROSAT All-Sky Survey in the X-ray emission of 57 OB stars over a time-scale of $\\approx$2 d (Bergh$\\rm{\\ddot{o}}$fer \\& Schmitt 1995).  \\subsection{X-ray temperatures of plasma}  Our analysis of XMM-Newton data reveals reveal two-component X-ray temperature structure for massive stars.  The temperatures corresponding to the cool and the hot components of HD 97434, and the cool component of DH Cep are found to be less than 1 keV. However, the temperature corresponding to the hot component of DH Cep was found to be more than 1.89 keV.  The derived values of temperatures are similar to  the O-type stars in the NGC~6231 cluster (Sana et al. 2006) and in the Carina OB1 association (Antokhin et al. 2008). Generally for O-type stars, the best fits to good quality data can usually be achieved by the sum of two thermal components at about 0.3 and 0.7--1\\,keV (G${\\rm{\\ddot{u}}}$del \\& Naz${\\rm{\\grave{e}}}$ 2009). The observed soft X-ray emission could originate from radiation-driven instabilities in stellar winds (Lucy 1982).  The wind-shock model predicts the intrinsic instability of the line driving mechanism.  Indeed, the velocity in an unstable wind is not the same everywhere and the fast-moving parcels of material will overcome the slow moving material, generating shocks between them.  This model  estimates the shock velocities (v$\\rm{_{shock}}$) by the relation  (Lucy 1982, Luo, McCray \\& Maclow 1990)   \\begin{equation} \\rm{v_{shock} =\\sqrt{\\frac{kT_{sh}}{1.95\\mu}}} \\end{equation}  Adopting the mean particle weight $\\rm{\\mu\\approx}$0.62 for O-type stars (Cassinelli et al. 2008) we derive the \"average\" value of shock velocities correspond to the cool components (kT$_1$) $\\rm{ 716^{+12}_{-12}  {~km~s^{-1} }}$ and $\\rm{ 514^{+54}_{-162}  {~km~s^{-1} }}$ for the stars DH Cep and HD 97434, respectively. However, the shock velocities correspond to the hot components (kT$_2$) are found to be $\\rm{ >1563 {~km~s^{-1} }}$ and $\\rm{ 782^{+66}_{-66}  {~km~s^{-1} }}$ for DH Cep and HD 97434, respectively. These values are about a factor of 2 larger than the predicted by radiative shock  model of Lucy (1982). However, the evolved version of the standard model by Owocki, Castor \\& Rybicki (1988) predicts X-ray emission upto 1 keV from wind shocks model. The temperature corresponding to the hot component of DH Cep is found to be more than 1.89 keV. It may arise from wind collision zone of the binary system (see Bhatt et al. 2010 and references therein), but we are not able to draw any firm conclusion  on the basis of it alone. Therefore, it appears that the cool as well as hot temperature components from DH Cep and HD 97434  could be generated by instabilities in radiation-driven wind shocks.   \\subsection{X-ray Luminosity} For massive stars DH Cep and HD 97434, the ratios of X-ray to bolometric luminosities, $\\rm{log(L_X /L_{bol}}$ ), are found to be -6.7 and -7.3 (see Table~\\ref{tab:spec_massive_bestfit} ) in the energy band 0.3-7.5 keV, respectively, which are broadly consistent with the relation derived for similar kind of O-type stars (Sana et al. 2006).  X-ray luminosities derived from XMM-Newton data are found to be an order of magnitude lower than the $\\rm{L_X}$ derived from Einstein observations for both the massive stars (see Table~\\ref{tab:spec_massive_bestfit} and \\S\\ref{sec:intro}). These discrepancies can not be explained by difference in sensitivity ranges of the various instruments. A similar results have been found by De Becker et al. (2004) for the star HD 159176 after the comparison of XMM-Newton, Einstein and ROSAT data. We found the X-ray luminosity  $\\rm{4.82\\times10^{32}}$ and $\\rm{1.95\\times10^{31}}$ in energy range 0.3--2.0 keV for the DH Cep and HD 97343, respectively, using XMM-Newton data. The choice of the energy range is nearly a similar to the energy range in Einstein observations, i.e., 0.2--3.5 keV. We further converted the Einstein IPC count rates into flux using the WebPIMMS\\footnote{http://heasarc.gsfc.nasa.gov/Tools/w3pimms.html}, where we assumed a thermal plasma model with temperature of 0.5 keV, and the hydrogen column density of $3.9\\times10^{21}$ for DH Cep and $1.8\\times10^{21}$ for HD 97434 (see Table 2). The Einstein IPC count rates are 0.0166 and $<0.0103$ counts s$^{-1}$ for the stars DH Cep and HD 97343, respectively (Chlebowski et al. 1989).  The X-ray luminosity thus calulated was $\\rm{2.3\\times10^{32}}$ and $<\\rm{1.2\\times10^{31}}$ for the stars DH Cep and HD 97434, respectively. These value are consistent with that obtained from XMM-Newton obsertions. Therefore, it appears that the model used by Chlebowski et al. (1989)  to convert the  Einstein count rates into luminosities may be responsible for these discrepancies."
    },
    "1004/1004.3481_arXiv.txt": {
        "abstract": "{Measuring the surface abundances of AGB stars is an important tool for studying the effects of nucleosynthesis and mixing in the interior of low- to intermediate mass stars during their final evolutionary phases. The atmospheres of AGB stars can be strongly affected by stellar pulsation and the development of a stellar wind, though, and the abundance determination of these objects should therefore be based on dynamic model atmospheres.} {We investigate the effects of stellar pulsation and mass loss on the appearance of selected spectral features (line profiles, line intensities) and on the derived elemental abundances by performing a systematic comparison of hydrostatic and dynamic model atmospheres.} {High-resolution synthetic spectra in the near infrared range were calculated based on two dynamic model atmospheres (at various phases during the pulsation cycle) as well as a grid of hydrostatic COMARCS models with effective temperatures $T_{\\rm eff}$ and surface gravities log\\,$g$ over an adequate range. Equivalent widths of a selection of atomic and molecular lines (Fe, OH, CO) were derived in both cases and compared with each other.} {In the case of the dynamic models, the equivalent widths of all investigated features vary over the pulsation cycle. A consistent reproduction of the derived variations with a set of hydrostatic models is not possible, but several individual phases and spectral features can be reproduced well with the help of specific hydrostatic atmospheric models. In addition, we show that the variations in equivalent width that we found on the basis of the adopted state-of-the-art dynamic model atmospheres agree qualitatively with observational results for the Mira R\\,Cas over its light cycle.} {The findings of our modelling form a starting point to deal with the problem of abundance determination in strongly dynamic AGB stars (i.e., long-period variables). Our results illustrate that some quantities such as the C/O ratio can probably still be determined to a reasonable accuracy, but the measurement of other quantities will be hampered by the dynamics. The qualitative agreement with observations of R\\,Cas opens promising possibilities for a forthcoming quantitative comparison of our synthetic spectra with observed ones of AGB variables in the globular cluster 47\\,Tuc.} ",
        "introduction": "Measurements of abundances in atmospheres of highly evolved red giants contribute to the understanding of two major astrophysical aspects, namely the processes inside red giant stars such as nucleosynthesis and mixing, and, secondly, the enrichment of the interstellar medium by freshly produced elements due to the mass loss of these stars. The determination of accurate element abundances is complicated by the large number of spectral lines from neutral atoms and molecules that give rise to heavy blending, especially in the optical range of the spectrum. The problem is less severe in the near infrared range, which has become increasingly important for this kind of analysis.  When a low or intermediate mass star evolves to the Asymptotic Giant Branch (AGB) phase, a second major problem occurs concerning the spectral analysis. Stars in this part of the HRD exhibit large amplitude radial pulsations. Being members of the variability class of so-called long period variables (LPVs), these objects change their visual brightnesses by up to several magnitudes over a few hundred days. The pulsating stellar interiors lead to complex velocity fields in the stellar atmospheres (e.g., Nowotny et al. \\cite{NowAH09}) and subsequently to phase-dependent line profile variations (asymmetric and split lines or emission features) as seen in observed spectra of these stars (e.g.\\,Hinkle et al.\\,\\cite{HHR82}). Thus, a hydrostatic description of these objects is inappropriate. However, it is this evolutionary phase where major nucleosynthesis and mixing processes take place, and therefore great interest in measuring stellar abundances exists.  In the past, a few attempts have been made to investigate the complicated profiles of individual lines observed in spectra of variable AGB stars on the basis of model calculations. Using the dynamic models for pulsating red giants without mass loss of Bessell \\& Scholz (\\cite{BessS89}), Scholz (\\cite{Schol92}) showed how velocity fields in Mira photospheres distort line profiles of different species and affect measurements of equivalent widths and curves of growth for abundance analyses from these spectra. This fundamental paper, also referencing earlier works in this field, investigated the possibility to \"predict accurately the equivalent width of lines which are little affected by outflow/infall velocities by means of a conventional hydrostatic model in place of a dynamical model\" (Scholz \\cite{Schol92}). He found that lines formed in deep layers below the arising shock fronts may be accessible to a classical curve of growth study, but that in some cases this approach seriously fails, so that abundance determination for large amplitude red variables remains a difficult problem.  In this context, a major challenge for dynamic model atmospheres was to reproduce the line profile variations observable in spectra of Mira type variables. First results concerning this were presented by Bessell et al. (\\cite{BesSW96}), who calculated synthetic line profiles (CO, Fe\\,I) based on the models of Bessell \\& Scholz (\\cite{BessS89}). They could qualitatively reproduce some of the observed line profile characteristics (asymmetries, shifted, or doubled lines).  McSaveney et al. (\\cite{MWSLH07}) attempted to derive abundances from high resolution spectra of pulsating AGB stars in the Magellanic Clouds by using the same type of model atmospheres. They succeeded in fitting the spectra of some stars at selected phases during the pulsation cycle. However, they had to restrict themselves to cases of relatively smooth atmospheric velocity fields. They were unable to derive abundances for model phases with complex atmospheric structures (shock waves), which cause distortions in the line profiles (emission components, line splitting).  While the Australia-Heidelberg model atmospheres (e.g., Bessell \\& Scholz \\cite{BessS89}, Bessell et al. \\cite{BesSW96}, Hofmann et al. \\cite{HofSW98}) represent pulsating red giants without mass loss, the Berlin models (e.g. Fleischer et al. \\cite{FleGS92}, Winters et al. \\cite{WiFLS97}+\\cite{Winters00}) describe the optically thick dusty outflows of evolved C-rich AGB stars with high mass-loss rates. Based on the latter, Winters et al. (\\cite{Winters00}) presented synthetic fundamental and first overtone CO line profiles. By comparing reasonably with observed ones for the obscured \\mbox{C-rich} Mira IRC+10216, these spectral lines can be used to investigate the radial structure and the temporal variation in the dusty envelope of these objects.  The non-grey dynamic model atmospheres of H\\\"ofner et al. (\\cite{HoGAJ03}) represent a combination of both approaches and describe the outer layers of an AGB star from the pulsating deep photosphere (piston driven), to the dust forming region where the stellar wind is triggered and beyond to the region of the steady outflow (Nowotny et al. \\cite{Nowo05a}). Keeping in mind the work of Scholz (\\cite{Schol92}), we conducted a similar exercise to investigate the problem of deriving abundances for LPVs by utilising the advanced dynamic models of H\\\"ofner et al. (\\cite{HoGAJ03}).  The major aim of the work presented here was to study the effect of atmospheric dynamics (pulsation, stellar wind) on the intensity of spectral lines for typical AGB stars. For this purpose, we compared selected spectral features in synthetic spectra computed on the basis of the aforementioned dynamic atmospheric models (Sect.\\,\\ref{s:dynvsstat}) with the corresponding results based on hydrostatic model atmospheres (Sect.\\,\\ref{s:hydro}). In a second step, we investigate observed high-resolution spectra of a typical Mira to see if the effects we found by the modelling are also recognisable in the observational data (Sect.\\,\\ref{s:obsvsmod}). ",
        "conclusions": "The aim of our study has been to explore the possibilities and caveats of deriving stellar abundances from observed spectra of AGB stars that are strongly affected by pulsation and mass loss. We can summarise our findings as follows:  \\begin{itemize}  \\item  There is not a single phase, neither for model~P nor for model~PM, where all investigated features would be in agreement with a single hydrostatic model at the same time. This suggests that there is a fundamental problem in fitting the complete spectrum of a dynamical star with one hydrostatic model.  \\item For several features, the dynamic models, especially model~PM, touch areas in the ($J$--$K$) versus equivalent width plane that seem to be unreachable by hydrostatic models even when modifying the chemical composition. Further tests are necessary to test the effects of e.g. metallicity on this finding.  \\item Spectral peculiarities due to stellar dynamics seem to be less significant in the $H$-band than in the $K$-band.  \\item For dynamic atmospheres with mass loss (represented by model~PM), the rising, post-minimum branch seems to be more promising for deriving element abundances with the help of hydrostatic models. This is, however, untrue for some of the features studied (namely the iron lines, OH 15540 and $^{12}$CO 23137).  \\item For dynamical atmospheres without mass loss (model~P), the loops are less extended and the detected trends of equivalent width with colour seem to resemble the hydrostatic case in many respects. For several features, we find an offset between the relation for model~P and the hydrostatic reference case with log\\,$g=-$0.25. This offset, however, is not the same for all features.  \\item Including the emission lines in the measurement of the strength of a spectral feature leads to an extension of the loop completed in the colour versus equivalent width plane. For high resolution spectroscopy, it is therefore better to exclude any clearly separated emission component when deriving the stellar abundance. If low resolution spectra were used, in which the emission component cannot be detected separately, we expect to see extensive loops such as those found in this study.  \\item Our results are qualitatively  similar to the findings of Scholz (\\cite{Schol92}) and others.  \\end{itemize}  We performed some tests to see if a combination of various spectral features (e.g. a ratio of the equivalent widths of two lines) could reduce the complexity of the variation with colour. But as the size and slope of the loops are different for each of the features observed, none of the combinations were found to be useful.  Deriving stellar abundances for pulsating stars without an outflow (simulated in this study by model~P with a bolometric light amplitude of 0.86$^{\\rm mag}$) may well be possible if we take into account a set of offsets that still need to be determined. Abundance determinations for these kind of stars performed so far are probably affected by systematic error, depending on the feature used.  If we wish to derive stellar abundances from dynamic atmospheres, we have to search for features that are far more sensitive to changes in the chemical composition than to dynamical effects. An example would be the $^{12}$CO 3-0 band head. As described above, the loop (Fig.\\,\\ref{f:COband_dyn}) covers a range in the colour vs. equivalent width diagram similar to a change in log\\,$g$ by 0.5. This is clearly less than the change from C/O$=$0.48 to C/O$=$0.25 (Fig.\\,\\ref{f:CO_band}). Therefore we have a good chance of determining whether, e.g., the surface abundances of an AGB star have already been modified from the first dredge-up composition by the third dredge-up (e.g. Lattanzio \\& Wood \\cite{LW04}). Having measurements of different phases would naturally help us to improve the accuracy of the result.  We emphasize that our approach includes some simplifications that need to be explored in more detail. We did not test the effect of a variation in metallicity. Furthermore, a change in C/O may have a more complex effect on dynamic atmospheres than known from the hydrostatic case. Finally, a comparison with observed time series of a mira reveals that some features are in common with the model (e.g. the occurrence of loops), while other aspects are different (e.g. the dependency on phase). At present, it is not clear if this is due to a general problem in the models or a dependence of these characteristics on the parameters of the model (note that R\\,Cas is very different from the models we used in our analysis). However, we point out that the values found for the equivalent widths in R\\,Cas are in good agreement with what we would expect from an extension of the detected trends towards redder colours.  In a forthcoming paper, we will continue our comparison of high resolution synthetic spectra from dynamical and hydrostatic model atmospheres to identify the most well-suited features for measuring stellar parameters and abundances for cool large amplitude variables. We also plan to extend the study to dynamic models with different composition. Finally, we intend to directly compare the observed spectra of AGB variables in the globular cluster 47 Tuc with the results of hydrostatic and dynamic model atmospheres."
    },
    "1004/1004.1167_arXiv.txt": {
        "abstract": "We present the fifth edition of the Sloan Digital Sky Survey (SDSS) Quasar Catalog, which is based upon the SDSS Seventh Data Release. The catalog, which contains 105,783 spectroscopically confirmed quasars, represents the conclusion of the SDSS-I and SDSS-II quasar survey. The catalog consists of the SDSS objects that have luminosities larger than \\hbox{$M_{i} = -22.0$} (in a cosmology with \\hbox{$H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$,} \\hbox{$\\Omega_M$ = 0.3,} and \\hbox{$\\Omega_{\\Lambda}$ = 0.7),} have at least one emission line with FWHM larger than 1000~km~s$^{-1}$ or have interesting/complex absorption features, are fainter than \\hbox{$i \\approx 15.0$,} and have highly reliable redshifts. The catalog covers an area of~$\\approx$~9380~deg$^2$. The quasar redshifts range from~0.065 to~5.46, with a median value of~1.49; the catalog includes 1248 quasars at redshifts greater than four, of which 56 are at redshifts greater than five. The catalog contains 9210 quasars with \\hbox{$i < 18$;} slightly over half of the entries have \\hbox{$i< 19$.} For each object the catalog presents positions accurate to better than~0.1$''$~rms per coordinate, five-band ($ugriz$) CCD-based photometry with typical accuracy of~0.03~mag, and information on the morphology and selection method. The catalog also contains radio, near-infrared, and X-ray emission properties of the quasars, when available, from other large-area surveys. The calibrated digital spectra cover the wavelength region 3800--9200 \\AA\\ at a spectral resolution \\hbox{of $\\simeq$ 2000;} the spectra can be retrieved from the SDSS public database using the information provided in the catalog. Over~96\\% of the objects in the catalog were discovered by the SDSS. We also include a supplemental list of an additional 207 quasars with SDSS spectra whose archive photometric information is incomplete. ",
        "introduction": "This paper describes the Fifth Edition of the Sloan Digital Sky Survey (SDSS; York et al.~2000) Quasar Catalog.  Previous versions of the catalog (Schneider et al.~2002, 2003, 2005, 2007; hereafter Papers~I, II, III, and~IV) were published with the SDSS Early Data Release (EDR; Stoughton et al.~2002), the SDSS First Data Release (DR1; Abazajian et al.~2003), the SDSS Third Data Release (DR3; Abazajian et al.~2005), and the SDSS Fifth Data Release (DR5; Adelman-McCarthy et al. 2007), and contained 3,814, 16,713, 46,420, and 77,429 quasars, respectively. The current catalog is the entire set of quasars from the SDSS-I and SDSS-II Quasar Survey; the SDSS-II was completed on 15 July~2008 and the Seventh Data Release (DR7; Abazajian et al.~2009) was made public on 31~October~2008.  The catalog contains 105,783 quasars.  The bulk of the quasars were identified as part of the SDSS Legacy Survey, which consisted of a systematic program to obtain spectra of one million galaxies and 100,000 quasars in two regions: an approximately 7,600~deg$^2$ area centered on the North Galactic Pole (hereafter this region will be referred to as the North Galactic Cap) and approximately 800~deg$^2$ in three narrow regions in the proximity of the Celestial Equator centered near 0$^{\\circ}$ in right ascension (see York et al.~2000 and Abazajian et al.~2009).  With this catalog, the spectroscopy in the North Galactic Cap region is now contiguous.  Outside of the Legacy area quasars were found, frequently serendipitously, on a series of ``Special Plates\" (see Adelman-McCarthy et al.~2006) that were observed when Legacy observations could not be performed. A large fraction of the Special Plates were obtained as part of two programs: Extensions of the Legacy target selection algorithms along the Fall equatorial region (Adelman-McCarthy et al.~2006), and the Sloan Extension for Galactic Understanding and Exploration (SEGUE; Yanny et al.~2009), a project to study stellar populations in the Galactic halo. The distribution of the DR7 quasars on the celestial sphere is displayed in Figure~1.  The catalog in the present paper consists of quasars that have a luminosity larger than \\hbox{$M_{i} = -22.0$}  (calculated assuming an \\hbox{$H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$,} \\hbox{$\\Omega_M$ = 0.3,} \\hbox{$\\Omega_{\\Lambda}$ = 0.7} cosmology, which will be used throughout this paper). The objects are denoted in the catalog by their DR7 J2000 coordinates; the format for the object name is \\hbox{SDSS Jhhmmss.ss+ddmmss.s}.  Since the image data used for the astrometric information can change between data releases (e.g., a region with poor seeing that is included in an early release is superseded by a newer observation in good seeing), the coordinates for an object can be modified at the~$\\approx$~0.1$''$ level (although on rare occasions (see \\S 6) the change in position can exceed~1$''$); hence the designation of a given source can change between data releases. When merging SDSS Quasar Catalogs with previous databases one should always use the coordinates, not object names, to identify unique entries.  Recent additional compilations of quasars based upon SDSS data include the DR5 Broad Absorption Line (BAL) Quasar Catalog (Gibson et al. 2009), the Two Degree Field-SDSS Luminous Red Galaxy and Quasar Survey (2SLAQ) spectroscopic catalog (Croom et al.~2009), and a sample of approximately one million ``photometric\" quasars (i.e., without spectroscopic confirmation) selected based upon their colors (Richards et al. 2009) drawn from the SDSS Data Release~6 (DR6; Adelman-McCarthy et al. 2008).  The DR7 catalog does not include classes of Active Galactic Nuclei (AGN) such as Type~2 quasars, Seyfert galaxies (i.e., low-luminosity AGN), and BL~Lacertae objects. Recent studies of these sources in the SDSS can be found in Reyes et al.~(2008) (Type~2), Hao et al.~(2005) (Seyferts), and Plotkin et al.~(2010) (BL~Lacs). No emission line equivalent width limit is imposed for inclusion in the present catalog; there are some quasars with low equivalent width emission lines in the catalog, and they tend to be among the brighter objects (the higher the quality of the spectra, the easier the identification of small equivalent width features), or at relatively high redshift, where redshifts can be measured using absorption from the Lyman~$\\alpha$ forest (e.g., Fan et al.~1999). The highest redshift SDSS quasars \\hbox{($z > 5.7$;} e.g., Fan et al.~2003, 2006; Jiang et al. 2009) were identified as candidates in the SDSS imaging data, but the spectra were not obtained with the SDSS spectrographs, so they are not included in the catalog.  The basic format and analysis of this work closely follow that of the previous four editions of the SDSS Quasar Catalog. The observations used to produce the catalog are presented in Section 2; the construction of the catalog and the catalog format are discussed in Sections 3 and~4, respectively.  Section~5 presents an overview of the catalog, Section~6 contains a discussion of some issues relevant to statistical samples, and a summary is given in Section~7. The catalog is presented in an electronic table in this paper and can also be found at an SDSS public web site.\\footnote{\\tt http://www.sdss.org/dr7/products/value$\\_$added/qsocat$\\_$dr7.html} ",
        "conclusions": "The DR5 database is entirely contained in that of DR7, but there are 181 quasars from Paper~IV (0.17\\%) that do not have a counterpart within~1.0$''$ of a DR7 quasar.  (This fraction is similar to the 0.15\\% dropped between Papers~III and~IV.)  Of the 181, there are ten that are included in the DR7 list that match at larger offsets; eight have matches within~2$''$ and the other two have separations of 2.0$''$ and~5.3$''$.  These ten separations are much larger than one would expect to find based on the 0.1$''$ rms scatter in the positions; the large coordinate offsets presumably arise from different versions of the imaging pipeline software (in particular the deblender).  The majority of the dropped objects are low-redshift, low-luminosity sources ($\\approx$~80\\% have redshifts below one). Most of the changes were introduced because 1)~the spectra did not appear to satisfy our quasar criteria (primarily the minimum line width limit) based upon the PCA analysis or 2)~small changes in the photometric measurements dropped the luminosity below our absolute magnitude criterion.  Although the quasars in the catalog have been selected using a variety of techniques (optical colors, radio properties, X-ray emission), the information contained in the catalog allows one to easily construct well-defined subsamples.  For example, to construct an optically-selected set of quasars (e.g., Strateva et al. 2005), one can either require that the low-redshift or high-redshift target selection flag is set, or, alternatively, require that the object was not chosen solely by radio/X-ray criteria (depending on whether or not you wish to have AGN found by the serendipity, galaxy, or star algorithms), using the quantities in \\hbox{columns 55--61} in the catalog.  The primary purpose of this catalog is to compile all of the quasars discovered/recovered by the SDSS survey in a complete and robust manner, allowing the general user to avoid the pitfalls that can arise from data mining such a large repository.  All of the spectra have been inspected by eye, frequently by more than one of the authors.  It may be that other users desire a more inclusive list of AGNs; for example, by relaxing our cuts on luminosity or line width.  In creating an expanded AGN list, it is important to consider two important issues in terms of the compilation and statistical analysis of such samples.  First, it is tempting to eliminate the labor-intensive visual examination stage and rely on the {\\tt zconf} flag as a means of restricting the AGN sample to the most robust objects.  However, {\\tt zconf} is not a good measure of the reliability of quasar redshifts: it depends strongly on redshift, as different emission lines enter and leave the SDSS spectral coverage.  For example, {\\tt zconf} drops dramatically in the mean from $z \\sim 0.7$ to $z \\sim 0.9$ as the H$\\beta$ feature leaves the SDSS spectral bandpass. The left panel of Figure~7 shows {\\tt zconf} as a function of redshift for bona-fide quasars whose spectra have been confirmed by eye.  The red histogram in the right panel in Figure~7 demonstrates the result of applying an arbitrary \\hbox{{\\tt zconf} $>0.95$} cut, independent of redshift, to the DR7 quasar sample.  The redshift dependence of {\\tt zconf} introduces an artificial apparent periodicity in the redshift distribution.  The second issue has to do with the effects of emission lines on quasar photometry. The SDSS quasar selection will include intrinsically fainter objects whenever a strong emission line enters the $i$ bandpass, making the quasar appear brighter than the same quasar at a redshift where the observed $i$ filter covers only continuum emission (Richards et al. 2006). When we restrict the sample to $i < 19.1$ and correct for the emission line k-correction (green histogram in Figure~7), the redshift distribution of the DR7 quasars is quite smooth."
    },
    "1004/1004.1398_arXiv.txt": {
        "abstract": "Broad-band images in the Ca~II H line, from the BFI instrument on the Hinode spacecraft, show emission from spicules emerging from and visible right down to the observed limb.  Surprisingly, little absorption of spicule light is seen along their lengths.  We present formal solutions to the transfer equation for given (ad-hoc) source functions, including a stratified chromosphere from which spicules emanate.  The model parameters are broadly compatible with earlier studies of spicules. The visibility of Ca~II spicules down to the limb in Hinode data seems to require that spicule emission  be Doppler shifted relative to  the stratified atmosphere, either by supersonic turbulent or organized spicular motion.  The  non-spicule component of the chromosphere is almost invisible in the broad band BFI data, but we predict that it will be clearly visible in high spectral resolution data.  Broad band Ca II H limb images give the false impression that the chromosphere is dominated by spicules. Our analysis serves as a reminder that the absence of a signature can be as significant as its presence. ",
        "introduction": "The Hinode spacecraft is a stable platform from which unique high resolution, seeing-free images of the Sun can be acquired \\citep{Kosugi+others2007}. The BFI instrument, fed by the Solar Optical Telescope (SOT) on Hinode \\citep{Tsuneta+others2008}, can observe a 3 \\AA{} wide spectral bandpass centered at the H line of Ca~II. Over this bandpass, the line forms in both the photosphere in the wings, and chromosphere in the core. Movies of such Ca II images have revealed a remarkably dynamic, spicule-dominated limb. The observed spicules have smaller diameters, higher apparent velocities and smaller lifetimes \\citep{DePontieu+others2007} than was previously thought \\citep{Beckers1968,Beckers1972}.  Figure~\\pref{fig:snapshot} shows a typical snapshot from a series of Ca~II BFI images acquired on 7 November 2007, in the northern polar coronal hole.    We selected a coronal hole because spicules there are longer than elsewhere, thereby providing a broad background of spicule emission against which a stratified atmosphere might easily be identified. A smooth radial gradient has been divided out of the data to enhance the emission high above the limb. The zero point of the height scale $z=0$ (along $x=0$) corresponds to the standard formation height of the 5000 \\AA{} continuum, when observed vertically, as derived by \\citet{Bjolseth2008}. Henceforth we will refer to heights on this standard scale. Bj\\o{}lseth found that the Ca~II limb lies $0.45\\pm0.034$  Mm above the blue limb.    Since the continuum at the limb forms about 0.375 Mm higher than at disk center, (e.g., \\citealp{Athay1976}), the Ca II limb forms near heights of 0.825 Mm.  Curiously, such limb images show little or no signature of a bulk, stratified chromosphere\\footnote{By ``chromosphere'' we refer not only to the traditional definition of H$\\alpha$ emitting plasma seen during eclipse flashes, but all the material lying between the quiet photosphere, with densities above $\\sim 5\\times10^{-9}$ g~cm$^{-3}$, and the corona with densities below $\\sim 10^{-13}$ g~cm$^{-3}$.  } which, as we discuss below, should have a thickness between 1 and 2 Mm.  The BFI instrument's resolution ($\\sim 0\\farcs1$) is ample to resolve structure on scales of 1-2 Mm ($0\\farcs1 \\equiv 0.0725$ Mm).  Yet a striking feature of these images is the continuous emission seen along each spicule all the way down to, and sometimes across, the limb.  Where then is the stratified chromosphere, and why are spicules so obviously dominant that one might conclude that the chromosphere itself consists of little more than a collection of spicules? In this paper we explore what these observations imply in terms of the structure of the chromosphere. ",
        "conclusions": "Hinode BFI Ca~II images obtained at the solar limb are consistent with the presence of the stratified chromosphere when spicular emission is Doppler shifted relative to the stratified material.  This can be achieved most naturally using broad and/or Doppler shifted spicule line profiles of magnitudes compatible with observed motions. The picture presented here can be tested directly using very stable spectra at the solar limb, to see for example if the behavior modeled in Figure~\\pref{fig:standard} is qualitatively correct.  This is a very challenging observation to make from the ground, but should be possible under conditions of outstandingly good seeing and with modern adaptive optics systems.  The calculations reinforce a commonly known problem regarding broad band spectral imagers: one must be very careful taking care of physical effects such as Doppler motions which are not spectrally resolved by the instrument. BFI Ca II limb observations are largely blind to the bulk of the chromosphere itself. This fact is a sobering reminder that the absence of a signature can be as significant as its presence."
    },
    "1004/1004.1217_arXiv.txt": {
        "abstract": "We discuss the recent progress in calculating the properties of `{\\it hybrid stars}' (stellar objects similar to neutron stars, classified by the incorporation of non-nucleonic degrees of freedom, including but not limited to hyperons and/or a quark-matter core) using the octet-baryon Quark-Meson Coupling (QMC) model. The version of QMC used is a recent improvement which includes the in-medium modification of the quark-quark hyperfine interaction. ",
        "introduction": "The study of the QCD phase diagram is of great interest to the scientific community. Lattice QCD produces simulations which provide an insight into the \\mbox{$\\{T>0, \\mu\\simeq 0\\}$} region of the phase diagram, whilst properties of QCD at finite chemical potential can be calculated using various models for dense hadronic matter (at $T\\ge 0, \\mu \\gg 0$). In order to qualify the success of these models we require a method to test the predictions via experiment and observation. For finite chemical potentials, this requires tests at various relevant density scales. For low-density predictions we can compare calculations to experiments involving finite nuclei, but no experiment (to date) can probe the entire range of extreme densities believed to exist within neutron stars\\footnote{In Ref.~\\cite{QMC2007} the authors note that heavy-ion collisions may be able to provide insight into the properties of matter up the density range of 2--3$\\rho_0$.}. We therefore turn our attention to proxy measurements which may support or refute our predictions, such as the masses and radii of pulsars.\\par  We focus our attention on a particular model\\emdash the Quark-Meson Coupling model (QMC), which we shall describe in Section~\\ref{sec:QMC} \\emdash that has seen much success at each of these density scales, and focus in particular on the high-density Equation of State (EoS) for infinite matter (see Section~\\ref{sec:infinite}). In order to examine the effects of quark degrees of freedom at extremely high densities we model a phase transition between this hadronic EoS and a quark-matter EoS as described by several models (see Section~\\ref{sec:phasetrans}).  We examine the predictions for neutron star properties that arise from calculations involving the QMC EoS and the effect that the phase transition has on these properties in Section~\\ref{sec:stars}. ",
        "conclusions": "The efficiency with which QMC allows one to model hyperon contributions highlights the importance of baryon structure. The lack of a phase transition between hyperonic QMC and NJL quark-matter indicates the importance of both hyperon degrees of freedom and dynamic chiral symmetry breaking, both of which must be treated with care. A transition is possible when one or both of these factors is neglected, but the physics must be {\\em removed} by hand.\\par  In order to achieve a Gibbs phase transition between a baryon phase and a quark-matter phase in which dynamical chiral symmetry breaking gives rise to constituent-quark masses, a moderately stiff baryon EoS is required, otherwise no such transition is possible. Such is the case for hyperonic QMC and NJL quark-matter.\\par  While the QMC model at Hartree level successfully reproduces many properties of finite nuclei and stellar objects, the softness of the hyperonic/hybrid EoS (as evidenced by the maximum stellar mass) is unable to fully account for the most massive pulsars observed to date. We note that further improvements such as calculations at Hartree--Fock level (of particular interest are the Fock terms corresponding to the $\\pi$ meson) are currently in development.   \\begin{theacknowledgments} This research was supported by the Australian Research Council. The author would like to thank A.~W.~Thomas for his guidance and support, as well as D.~B.~Leinweber and A.~G.~Williams for their helpful discussions. \\end{theacknowledgments}"
    },
    "1004/1004.3738_arXiv.txt": {
        "abstract": "{ } { We measured the radial velocity of 139 stars in the region of \\object{NGC 6253}, discussing cluster's membership and binarity in this sample, complementing our analysis with photometric, proper motion, and radial velocity data available from previous studies of this cluster, and analyzing three planetary transiting candidates we found in the field of \\object{NGC~6253}. } { Spectra were obtained with the UVES and GIRAFFE spectrographs at the VLT, during three epochs in August 2008. } { The mean radial velocity of the cluster is ($\\overline{RV_{cl}}\\pm\\overline{\\sigma_{cl}})=(-29.11\\pm0.85)$ km/s. Using both radial velocities and proper motions we found 35 cluster's members, among which 12 are likely cluster's close binary systems. One star may have a sub-stellar companion, requiring a more intensive follow-up. Our results are in good agreement with past radial velocity and photometric measurements. Furthermore, using our photometry, astrometry and spectroscopy we identified a new sub-giant branch eclipsing binary system, member of the cluster. The cluster's close binary frequency at (29$\\pm$9)$\\%$ (34$\\%$$\\pm$10$\\%$ once including long period binaries), appears higher than the field binary frequency equal to (22$\\pm$5)$\\%$, though these estimates are still consistent within the uncertainties. Among the three transiting planetary candidates the brightest one ($V=15.26$) is worth to be more intensively investigated with higher percision spectroscopy. } { We discussed the possibility to detect sub-stellar companions (brown dwarfs and planets) with the radial velocity technique (both with UVES/GIRAFFE and HARPS) around turn-off stars of old open clusters. We isolated 5 stars that are optimal targets to search for planetary mass companions with HARPS. Our optimized strategy minimizes the observing time requested to isolate and follow-up best planetary candidates in clusters with high precision spectrographs, an important aspect given the faintness of the target stars. } ",
        "introduction": "\\label{s:introduction}    The advent of multi-object spectrographs feeding large aperture telescopes offers the possibility to obtain simultaneous, repeated, and accurate spectroscopic measurements of a large sample of stars. The application of this technique to the study of open clusters is of particular interest. While open clusters photometric surveys are now routinely performed spanning typically several consecutive nights, there are relatively few multi-epoch radial velocity surveys targeting these objects (Mermilliod et al. 2009; Hole et al. 2009). Besides allowing further culling of cluster's members, multiple radial velocity measurements allow a more complete census of binary and multiple systems. Spectroscopic follow-up of eclipsing binary systems that belong to open clusters gives the opportunity to derive accurate stellar masses, constraining models of stellar evolution (e. g. Southworth \\& Clausen 2006). Moreover, the increasing stability and accuracy of modern spectrographs is driving toward detection of binary systems with low mass companions (such as brown dwarfs and planets) also in open clusters, where target stars are typically fainter than those observed in common radial velocity planet searches in the solar surrounding.  In this work, we focused our attention on the old and metal-rich open cluster \\object{NGC 6253} ($\\alpha_{2000}=16^{h}\\,59^{m}\\,05^{s}, \\delta_{2000}=-52^{\\circ}\\,42\\arcmin\\,30\\arcsec, l=335\\fdg5,b=-6\\fdg3$). This cluster was selected as part of our project aimed at searching for transiting hot-jupiter planets in metal-rich open clusters (Montalto et al. 2007). In 2004, we performed a photometric transit search toward \\object{NGC 6253} using the WFI at the 2.2m Telescope for ten consecutive nights (Montalto et al. 2009), identifying three transiting planet candidates.  In the solar neighborhood, FGK metal-rich dwarf stars have higher probability to host jupiter-like planets ( Gonzalez 1998, Santos et al.~2001, Fischer \\& Valenti 2005). Accordingly to the most accredited explanation that the planet-metallicity correlation is of primordial origin, we expect a higher planet discovery rate targeting preferentially stars borned in metal-rich environments. While at solar metallicity the frequency of planets (with period P $<$ 4 years, and radial velocity semi-amplitude $K>30$ m/s) around FGK dwarf stars in the solar neighbourhood is about 3$\\%$, at the metallicity of \\object{NGC~6253}, the expected frequency is around $18\\%$, as can be deduced from Fischer \\& Valenti (2005).  Since the large expected frequency of planets in metal rich clusters, these objects are top targets for planet searches and offer an excellent natural laboratory to test ideas of planet formation and evolution, while probing at the same time the effects of the environment.   Searches for planets around dwarf stars in clusters have not yet provided even bona-fide planetary candidates. This result may still be due to statistical problems, since in particular open clusters are tipically loosely populated and the most widely used technique to detect planets in such environments is the transit method. However, the increasingly larger number of surveys and the lack of detections could start to reveal some fundamental differences between planet formation processes around field and cluster's dwarf stars. It is then of primary importance to continue monitoring clusters and estabilish on a firmly observational basis if there is indeed a difference between the planet frequency in these environments and in the field.  In this work we present some preliminary results regarding our survey toward \\object{NGC~6253}. In August 2008, we followed-up with VLT three candidate transiting planets we found in the region of this cluster. We will dedicate a forthcoming contribution to the study of planet frequency in the field and in the cluster. Our observational strategy was tied also to the study of cluster's members and surrounding field stars, and here we present the results of the complete multi-epoch radial velocity campaign toward \\object{NGC 6253}. We used FLAMES in MEDUSA mode, targeting a total of 204 stars in the region of the cluster.  In this work we then explore the possibility to search for planetary companions around dwarf stars of old open clusters by means of the radial velocity technique. This detection method would greatly increase our chances to detect planetary mass companions, allowing us to overcome the problem of the small sample of stars available in clusters. However such technique has been applyied systematically only to the Hyades cluster so far (Cochran et al. 2002, Paulson et al. 2002, Paulson et al. 2003) given the closeness and consequently brightness of its members. In young clusters planet detection is significantly hampered and complicated by stellar activity (Paulson et al. 2004a, Paulson et al. 2004b). Old open clusters should be better targets, although the faintness of their members seems to significantly limit the application of the radial velocity technique, requiring in general a very large investment of observing time. Here we discuss an optimized observing strategy aimed at isolating only the most promising objects to follow-up with high precision spectroscopy, involving the use of photometry, astrometry and multi-object spectroscopy. Beside resulting in a great improvement in the knowledge of cluster's properties this method minimizes the time needed to identify and follow-up best candidates. The application of this technique to the particular case of \\object{NGC~6253} is of special interest given the characteristics of this cluster.   \\object{NGC 6253} has been studied by several authors in the past, both photometrically (Bragaglia et al. 1997, Piatti et al. 1998, Sagar, Munari, \\& de Boer 2001, Twarog, Anthony-Twarog \\& De Lee 2003, Anthony-Twarog, Twarog, \\& Mayer 2007, Montalto et al. 2009, De Marchi et al. 2009), and spectroscopically (Carretta et al. 2000, Carretta, Bragaglia \\& Gratton 2007, Sestito et al. 2007). In addition, in Montalto et al.~(2009) we calculated proper motion membership probabilities.  These studies have demonstrated that \\object{NGC 6253} is an old ($\\sim$3.5 Gyr, e. g. Montalto et al.~2009), and metal-rich cluster ([Fe/H]=+0.39$\\pm$0.07 Sestito et al.~2007; [Fe/H]=+0.46 Carretta, Bragaglia \\& Gratton~2007), being in fact one of the most metal-rich open clusters of the Galaxy. It is also one of the few old open clusters located inward the solar ring (Carraro et al. 2005a, 2005b), at a Galactocentric distance of around 6 kpc, where more prohibitive environmental conditions in general prevent clusters' survival (Wielen 1971). \\object{NGC 6253} it is also important in the more general context of stellar population studies, offering an homogeneous sample of coeval metal-rich stars against which stellar models at this extreme metallicity can be tested and compared. Its peculiar location in the Galactic disk gives the opportunity to extend toward the inner regions of the Galaxy the baseline for the study of the Galactic disk radial abundance, which is a basic ingredient of Galactic chemical evolution models (Tosi 1996).   This paper is organized as follows: in Sect.~\\ref{s:observations}, we describe our observations; in Sect.~\\ref{s:reduction}, we give a detailed description of reductions and calibrations; in Sect.~\\ref{s:results}, we discuss cluster's membership and binarity in our sample; in Sect.~\\ref{s:transits}, we analyze our three transiting planetary candidates; in Sect.~\\ref{s:EA}, we focus our attention on four detached eclipsing binary systems; in Sect.~\\ref{s:EB}, we discuss a new eclipsing binary system located at the sub-giant branch of \\object{NGC~6253}; in Sect.~\\ref{s:RV_search}, we present an optimized strategy to search for sub-stellar objects with the radial velocity technique around turn-off stars of old open clusters; in Sect.~\\ref{s:simulation}, we discuss a method to constrain the minimum mass and period of cluster's close binary systems, detected by RV surveys; finally in Sect.~\\ref{s:conclusions}, we summarize and conclude. ",
        "conclusions": "\\label{s:conclusions}  We have presented the results of the first multi-epoch radial velocity survey toward the old metal-rich open cluster \\object{NGC 6253}. The mean radial velocity of the cluster is $\\overline{RV_{cl}}\\pm\\overline{\\sigma_{cl}}=(-29.11\\pm0.85)$ km/s. Using our photometry, proper motions, and radial velocities we identified 35 likely cluster's members, populating the turn-off, sub-giant, red-giant, red clump, and blue straggler regions of the cluster. Among this sample, 12 objects are likely cluster's close binary systems.  We detected one object that may have a companion with a minimum mass in the sub-stellar regime, and that needs to be further investigated: star 39810.  We isolated five turn-off stars that are optimal targets to search for planetary mass companions with the high precision spectroscopy: stars 40519, 45512, 44104, 45523, 45267.  Among the three planetary candidtes we found in the region of \\object{NGC~6253}, star 171895 is the most interesting objects, requiring a more intensive photometric and spectroscopic follow-up in particular with high precision spectroscopy.  We found a new sub-giant branch eclipsing binary system: star 45368.  We compared our results with previous literature radial velocity measurements of a few cluster's members, finding in general a good agreement with our study.  The close binary frequency among turn-off and evolved cluster's stars is f$_{cl}=$(29$\\pm$9)$\\%$ (f$_{bin}=$[34$\\pm$10]$\\%$,  once including also longer period binaries), appears larger than the field binary frequency f$_{cl,field}=$(22$\\pm$5)$\\%$, although still consistent with that estimate considering the uncertainties.  We showed that searching for sub-stellar objects (brown dwarfs and jupiter planets) with the radial velocity technique around old open clusters turn-off stars, appears feasible also with present day instrumentation, but it is necessary: (i) an efficient pre-selection of candidate cluster's members; (ii) further identification of likely massive sub-stellar objects by means of multi-object spectroscopy; (iii) subsequent follow-up of the best target stars by means of high precision spectroscopy. We derived that three observing epochs uniformly distributed over a period of four years give a good compromise between expected number of detectable planets, and observing time.  In Table 12 we report the cross-correlation among the star's ID introduced in Montalto et al.~(2009) and other authours and catalogs.  Very recently Anthony-Twarog et al.~(2010) have presented new radial velocity data on \\object{NGC~6253} members. These measurements have not been included in our analysis. Merging the results presented in that work with our own results, will allow further identification of the best targets to follow-up more intensively with high resolution spectroscopy to search for sub-stellar companions around cluster's turn-off stars."
    },
    "1004/1004.5602_arXiv.txt": {
        "abstract": "We review the recently found large-scale anomalies in the maps of temperature anisotropies in the cosmic microwave background. These include alignments of the largest modes of CMB anisotropy with each other and with geometry and direction of motion of the Solar System, and the unusually low power at these largest scales. We discuss these findings in relation to expectation from standard inflationary cosmology, their statistical significance, the tools to study them, and the various attempts to explain them. ",
        "introduction": "The Copernican principle states that the Earth does not occupy a special place in the universe and that observations made from Earth can be taken to be broadly characteristic of what would be seen from any other point in the universe at the same epoch.  The microwave sky is isotropic, apart from a Doppler dipole and a microwave foreground from the Milky Way. Together with the Copernican principle and some technical assumptions, an oft-inferred consequence is the so-called cosmological principle.  It states that the distributions of matter and light in the Universe are homogeneous and isotropic at any epoch and thus also defines what we mean by cosmic time.  This set of assumptions is a crucial, implicit ingredient in obtaining most important results in quantitative cosmology.  For example, it allows us to treat cosmic microwave background (CMB) temperature fluctuations in different directions on the sky as multiple probes of a single statistical ensemble, leading to the precision determinations of cosmological parameters that we have today.  Although we have some observational evidence that homogeneity and isotropy are reasonably good approximations to reality, neither of these are actual logical consequences of the Copernican principle. For example the geometry of space could be homogeneous but anisotropic --- like the surface of a sharp mountain ridge, with a gentle path ahead but the ground dropping steeply away to the sides.  Indeed, three-dimensional space admits not just the three well known homogeneous isotropic geometries (Euclidean, spherical and hyperbolic -- $E^3$, $S^3$ and $H^3$),  but five others which are homogeneous but anisotropic.  The two simplest are $S^2\\times E^1$ and $H^2\\times E^1$. These spaces support the cosmological principle but have preferred directions.  Similarly, although the Earth might not occupy a privileged place in the universe, it is not necessarily true that all points of observation are equivalent.  For example, the topology of space may not be simply-connected --- we could live in a three-dimensional generalization of a torus so that if you travel far enough in certain directions you come back to where you started.  While such three-spaces generically admit locally homogeneous and isotropic geometries, certain directions or points might be singled out when non-local measurements are considered.  For example the length of the shortest closed non-trivial geodesic through a point depends on the location of that point within the fundamental domain.  Similarly, the inhomogeneity and anisotropy of eigenmodes of differential operators on such spaces are likely to translate into statistically inhomogeneous and anisotropic large scale structure, in the manner of Chladni figures on vibrating plates.  The existence of non-trivial cosmic topology and of anisotropic geometry are questions that can only be answered observationally.  In this regard, it is worth noting that our record at predicting the gross properties of the universe on large scales from first principles has been rather poor. According to the standard concordance model of cosmology, over $95$\\% of the energy content of the universe is extraordinary --- dark matter or dark energy whose existence has been inferred from the failure of the Standard Model of particle physics plus General Relativity to describe the behavior of astrophysical systems larger than a stellar cluster --- while the very homogeneity and isotropy (and inhomogeneity) of the universe owe to the influence of an inflaton field whose particle-physics-identity is completely mysterious even after three decades of theorizing.  The stakes are set even higher with the recent discovery of dark energy that makes the universe undergo accelerated expansion. It is known that dark energy can affect the largest scales of the universe --- for example, the clustering scale of dark energy may be about the horizon size today. Similarly, inflationary models can induce observable effects on the largest scales via either explicit or spontaneous violations of statistical isotropy. It is reasonable to suggest that statistical isotropy and homogeneity should be substantiated observationally, not just assumed.  More generally, testing the cosmological principle should be one of the key goals of modern observational cosmology.  With the advent of high signal-to-noise maps of the cosmic microwave background anisotropies and with the conduct of nearly-full-sky deep galaxy surveys, statistical isotropy \\textit{has} begun to be precisely tested. Extraordinary full-sky temperature maps produced by the Wilkinson Microwave Anisotropy Probe (WMAP), in particular, are revolutionizing our ability to probe the universe on its largest scales \\cite{Bennett2003, Spergel2003,Hinshaw2003,Spergel2006,wmap5,wmap7}.  In the near future, these will be joined by higher resolution temperature maps and high-resolution polarization maps and, eventually, by deep all-sky surveys, and perhaps by tomographic 21-cm line observations that will extend our detailed knowledge of the universe's background geometry and fluctuations into the interior of the sphere of last scattering.  \\begin{figure} \\includegraphics[width=\\linewidth]{ComScale.pdf} \\caption{The comoving length within the context of the concordance model of an arc seen at an fixed angle and the comoving Hubble length as functions of redshift. Linear perturbation theory is expected to work well outside the shaded region, in which the large scale structure (LSS) forms.} \\label{fig:comscale} \\end{figure}  In this brief review, we describe the large-scale anomalies in the CMB data, some of which were first reported on by the Cosmic Background Explorer (COBE) Differential Microwave Radiometer (DMR) collaboration in the mid 1990's. In particular, we report on alignments of the largest modes of CMB anisotropy with each other, and with geometry and direction of motion of the Solar System, as well as on unusually low angular correlations at the largest angular scales. We discuss these findings and, as this is not meant to be a comprehensive review and we emphasize results based on our own work in the area, we refer the reader to literature for all developments in the field. This review extends an earlier review on the subject by \\citet{Huterer_NewAst_review}, and complements another review on statistical isotropy in this Special Issue \\cite{Abramo_AdvAstro}.  The paper is organized as follows. In Sec.~\\ref{sec:expect} we describe the statistical quantities that describe the CMB, and the expectations for their values in the currently favored $\\Lambda$CDM cosmological model. In Sec.~\\ref{sec:align}, we describe the alignments at the largest scales, as well as multipole vectors, which is a tool to study them. In Sec.~\\ref{sec:2pt}, we describe findings of low power at largest scales in the CMB. Section \\ref{sec:explain} categorizes and covers the variety of possible explanations for these anomalies. We conclude in Sec.~\\ref{sec:discussion}. ",
        "conclusions": "The CMB is widely regarded as offering strong substantiating evidence for the concordance model of cosmology.  Indeed the agreement between theory and data is remarkable --- the patterns in the two point correlation functions (TT, TE and EE) of Doppler peaks and troughs are reproduced in detail by fitting with only six (or so) cosmological parameters.  This agreement should not be taken lightly; it shows our precise understanding of the causal physics on the last scattering surface. Even so, the cosmological model we arrive at is baroque, requiring the introduction at different scales and epochs of three sources of energy density that are only detected gravitationally --- dark matter, dark energy and the inflaton.  This alone should encourage us to continuously challenge the model and probe the observations particularly on scales larger than the horizon at the time of last scattering.  At the very least, probes of the large-angle (low-$\\ell$) properties of the CMB reveal that we do not live in a typical realization of the concordance model of inflationary $\\Lambda$CDM.  We have reviewed a number of the ways in which that is true: the peculiar geometry of the $\\ell=2$ and $3$ multipoles --- their planarity, their mutual alignment, their alignment perpendicular to the ecliptic and to the dipole; the north-south asymmetry; and the near absence of two-point correlations for points separated by more than $60^o$.  If indeed the observed $\\ell=2$ and $3$ CMB fluctuations are not cosmological, one must reconsider all CMB results that rely on the low $\\ell$, e.g.~the measurement of the optical depth from CMB polarization at low $\\ell$ or the spectral index of scalar perturbations and its running.  Moreover, the CMB-galaxy cross-correlation, which has been used to provide evidence for the Integrated Sachs-Wolfe effect and hence the existence of dark energy, also gets contributions from the lowest multipoles (though the main contribution comes from slightly smaller scales, $\\ell\\sim 10$).  Indeed, it is quite possible that the underlying physical mechanism does not cut off abruptly at the octopole, but rather affects the higher multipoles.  Indeed, several pieces of evidence have been presented for anomalies at $l>3$ (e.g.\\ \\cite{Land2005a,Land:2006bn}). One of these is the parity of the microwave sky.  While the observational fact that the octopole is larger than the quadrupole ($C_3 > C_2$) is not remarkable on its own, including higher multipoles (up to $\\ell \\sim 20$) the microwave sky appears to be parity odd at a statistically significant level (since WMAP 5yr) \\cite{Land:2005jq,Kim:2010gf,Kim:2010gd}. It is hard to imagine a cosmological explanation for a parity odd universe, but the same holds true for unidentified systematics or unaccounted astrophysical foregrounds, especially as this recently noticed puzzle shows up in the very well studied angular power spectrum.   While the further WMAP data is not expected to change any of the observed results, our understanding and analysis techniques are likely to improve. Much work remains to study the large-scale correlations using improved foreground treatment, accounting even for the subtle systematics, and in particular studying the time-ordered data from the spacecraft. The Planck experiment will be of great importance, as it will provide maps of the largest scales obtained using a very different experimental approach than WMAP --- measuring the absolute temperature rather than temperature differences. Polarization maps, when available at high enough signal-to-noise at large scales (which may not be soon), will be a fantastic independent test of the alignments, as each explanation for the alignments, in principle, also predicts the statistics of the polarization pattern on the sky."
    },
    "1004/1004.2352_arXiv.txt": {
        "abstract": "We report on a detailed investigation of the high-energy $\\gamma$-ray emission from NGC\\,1275, a well-known radio galaxy hosted by a giant elliptical located at the center of the nearby Perseus cluster. With the increased photon statistics, the center of the $\\gamma$-ray emitting region is now measured to be separated by only $0.46$\\,arcmin from the nucleus of NGC\\,1275, well within the $95\\%$ confidence error circle with radius $\\simeq 1.5$\\,arcmin. Early {\\it Fermi}-LAT observations revealed a significant decade-timescale brightening of NGC\\,1275 at GeV photon energies, with a flux about seven times higher than the one implied by the upper limit from previous EGRET observations. With the accumulation of one-year of {\\it Fermi}-LAT all-sky-survey exposure, we now detect flux and spectral variations of this source on month timescales, as reported in this paper. The average $>$100 MeV $\\gamma$-ray spectrum of NGC\\,1275 shows a possible deviation from a simple power-law shape, indicating a spectral cut-off around an observed photon energy of $\\varepsilon_{\\gamma} = 42.2 \\pm 19.6$ GeV, with an average flux of $F_{\\gamma} = (2.31 \\pm 0.13) \\times 10^{-7}$\\,ph\\,cm$^{-2}$\\,s$^{-1}$ and a power-law photon index, $\\Gamma_{\\gamma} = 2.13 \\pm 0.02$. The largest $\\gamma$-ray flaring event was observed in April--May 2009 and was accompanied by significant spectral variability above $\\varepsilon_{\\gamma} \\gtrsim 1-2$ GeV. The $\\gamma$-ray activity of NGC\\,1275 during this flare can be described by a hysteresis behavior in the flux versus photon index plane. The highest energy photon associated with the $\\gamma$-ray source was detected at the very end of the observation, with the observed energy of $\\varepsilon_{\\gamma} = 67.4$\\,GeV and an angular separation of about $2.4$\\,arcmin from the nucleus. In this paper we present the details of the {\\it Fermi}-LAT data analysis, and briefly discuss the implications of the observed $\\gamma$-ray spectral evolution of NGC\\,1275 in the context of $\\gamma$-ray blazar sources in general. ",
        "introduction": "\\label{sec:intro} With the successful launch of the {\\it Fermi} Gamma-ray Space Telescope, we have a new opportunity to study the $\\gamma$-ray emission from different types of extragalactic sources --- not only blazars, but also radio galaxies and possibly other classes of active galactic nuclei (AGN) --- with much improved sensitivity than previously available (Abdo et al. 2010a). During the initial all-sky survey performed during the first four months after its launch, the {\\it Fermi} Large Area Telescope (LAT) detected only two radio galaxies at high significance ($\\geq 10\\sigma$), namely NGC\\,1275 (Abdo et al. 2009a; hereafter Paper-I) and Cen\\,A (Abdo et al. 2009b, 2009c). More recently, the detection of MeV/GeV emission from yet another famous radio galaxy M\\,87, an established TeV source, was reported based on ten-months of all-sky survey {\\it Fermi}-LAT data (Abdo et al. 2009d). Yet the detection of NGC\\,1275 was particularly noteworthy  because this source, unlike Cen\\,A or M\\,87, was previously undetected in $\\gamma$-rays, neither by {\\it CGRO}/EGRET during its $\\sim$10 years of operation, nor by ground-based Cherenkov Telescopes. The $\\gamma$-ray flux of NGC\\,1275 detected by the {\\it Fermi}-LAT was about seven times higher than the one implied by the $2\\sigma$ upper limit reported by EGRET, namely $F_{\\varepsilon_{\\gamma} > 100\\,{\\rm MeV}} < 3.72 \\times 10^{-8}$\\,ph\\,cm$^{-2}$\\,s$^{-1}$ (Reimer et al. 2003). We note that {\\it COS B} data taken between 1975 and 1979 (Strong et al. 1982; Mayer-Hasselwander et al. 1982) showed a $\\gamma$-ray excess coincident with the position of this galaxy, although the evidence for the claimed high-energy source to be \\emph{uniquely} related to NGC\\,1275 is ambiguous.  NGC\\,1275 is a giant elliptical galaxy located at the center of the Perseus cluster\\footnote{The Perseus cluster (Abell 426): redshift $z=0.0179$, luminosity distance $d_{\\rm L} = 75.3$\\, Mpc, scale $21.5$\\,kpc arcmin$^{-1}$ (for flat cosmology with $H_{\\rm 0} = 71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$ and $\\Omega_{\\rm M} = 0.27$).}. This cluster is the brightest cluster of galaxies in the X-ray band (e.g., B\\\"{o}hringer et al. 1993; Fabian et al. 2003, 2006), and as such it has been the focus of several extensive research programs over many years and across the entire available electromagnetic spectrum. When observed at radio wavelengths, NGC\\,1275 hosts the exceptionally bright radio FR\\,I radio galaxy Perseus\\,A = 3C\\,84 (e.g., Vermeulen et al. 1994, Taylor et al. 1996, Walker et al. 2000, Asada et al. 2006). Although high-energy $\\gamma$-rays may in general be produced within the intergalactic/interstellar medium of the Perseus cluster, in Paper-I, we argued that the inner radio jet of 3C\\,84 was the most likely source of the observed $\\gamma$-ray photons because of the variability of the MeV/GeV flux on year/decade timescales implied by the EGRET and early {\\it Fermi}-LAT observations. Specifically, these measurements implied the $\\gamma$-ray emission region size in NGC\\,1275 has a radius, $R \\lesssim ct_{\\rm var} \\simeq 1$\\,pc. Note, however, that no significant variability was detected within the four-month long {\\it Fermi}-LAT dataset analyzed in Paper-I. Also, the LAT error circle determined from these initial data was too large to exclude possible contributions from other galaxies to the observed $\\gamma$-ray flux (specifically, NGC~1273, 1274, 1277, 1278 and 1279, were still within the previously determined $95\\%$ error circle with $R=5.2$\\,arcmin).  Meanwhile, follow-up observations by the VERITAS Cherenkov telescope (Weekes et al. 2002) put strong constraints on the very high-energy (VHE) $\\gamma$-ray emission from NGC\\,1275 above 100 GeV. In particular, no VHE $\\gamma$-ray emission from NGC\\,1275 was detected by VERITAS, with a $99\\%$ confidence level upper limit of $2.1\\%$ of the Crab Nebula flux, corresponding to $19\\%$ of the power-law extrapolation of the MeV/GeV flux observed during the first four months of the {\\it Fermi}-LAT observations (assuming the photon index, $ \\Gamma_{\\gamma} \\simeq 2.2$; Acciari et al. 2009b). This naturally indicates a deviation from the pure power-law spectrum in the VHE regime, possibly compatible with the presence of an exponential cutoff around or below photon energies $\\varepsilon_{\\gamma} \\simeq 100$\\,GeV (Acciari et al. 2009b). The MAGIC Cherenkov telescope also recently measured upper limits for the VHE $\\gamma$-ray emission of NGC\\,1275, namely $F_{\\varepsilon_{\\gamma} > 100\\,{\\rm GeV}} < (4.6-7.5) \\times 10^{-12}$\\,ph\\,cm$^{-2}$\\,s$^{-1}$ for the photon indices ranging from $1.5$ up to $2.5$ (Aleksi\\'c et al. 2010). Thus, the implied deviation in the $\\gamma$-ray spectrum of 3C\\,84 from a simple power-law form, as well as a possibility for a short-timescale ($<$ month) variability of the Perseus\\,A $\\gamma$-ray flux, may now be finally addressed and re-examined by {\\it Fermi}-LAT, due to the much improved photon statistic (especially above $10$\\,GeV) after the one-year-all-sky survey. Obviously, such deep studies of NGC\\,1275 in the MeV/GeV photon energy range are of the major importance for understanding the whole class of $\\gamma$-ray emitting radio galaxies in general.  Firmly motivated, we performed a detailed investigation of NGC\\,1275 in $\\gamma$-rays based on the accumulation of one-year of {\\it Fermi}-LAT all-sky survey data. In particular, we aimed to address the following problems: (i) presence of short-timescale $\\gamma$-ray flux variability, (ii) positional coincidence of the $\\gamma$-ray emitting center with the active nucleus of NGC\\,1275, and (iii) spectral curvature and spectral evolution of NGC\\,1275 in the MeV/GeV photon energy range. In $\\S$\\,2, we describe the {\\it Fermi}-LAT $\\gamma$-ray observations and data reduction procedure. The results of the analysis are given in $\\S$\\,3, and the discussion and conclusions are presented in $\\S$\\,4. ",
        "conclusions": "\\label{sec:discussion}  In the previous sections, we reported on the analysis of the $\\gamma$-ray emission from NGC\\,1275 observed with {\\it Fermi}-LAT during its one-year-all-sky survey. We showed that with the increased photon statistics, the positional center of the $\\gamma$-ray emission is now much closer to the NGC\\,1275 nucleus as compared to that reported in Paper-I. In addition, we have shown that the average $\\gamma$-ray spectrum of NGC\\,1275 reveals a significant deviation from a simple power-law above photon energies $\\varepsilon_{\\gamma} \\sim 1-2$\\,GeV. That is, the observed {\\it Fermi}-LAT spectrum is best fitted by a power-law function ($\\Gamma_{\\gamma} \\simeq$ 2.1) with an exponential cutoff at the break photon energy $\\varepsilon_{\\rm c} = 42.2 \\pm 19.6$\\,GeV. Finally, we argued that significant flux and spectral changes of NGC\\,1275 are detected with {\\it Fermi}-LAT on a timescale of a few months, although the possibility for even shorter variability remains uncertain.  We also reported the detection of an interesting spectral evolution, consisting of a \\emph{persistent} (over more than a few months) spectral hardening (from $\\Gamma_{\\gamma} \\simeq 2.2$ to $\\Gamma_{\\gamma} \\simeq 2$) after the largest flaring event observed in April--May 2009. During this flat-spectrum/low-flux-level epoch the highest energy photon ($\\varepsilon_{\\gamma} \\simeq 67.4$\\,GeV) was detected from the direction of NGC\\,1275. All these new findings basically support the idea put forward in Paper-I that the observed $\\gamma$-ray emission from the Perseus system originates in the (sub) pc-scale radio jet of NGC\\,1275, and is therefore most likely analogous to high-energy emission observed in blazars. In fact, as shown in Paper-I, the overall $\\nu F_\\nu$ spectral energy distribution (SED) of NGC\\,1275 constructed with multi-frequency radio to $\\gamma$-ray data shows a close similarity to the ``two-bump'' SEDs of so-called low-frequency peaked BL Lac objects (hereafter LBLs). Till now, only a few LBLs have been detected at TeV photon energies -- BL Lac (Albert et al. 2007),  3C\\,66A (Acciari et al. 2009a), S5 0716+714 (Anderhub et al. 2009), and also W~Comae from an IBL class (Acciari et al. 2008; 2009d)  -- and more similar discoveries are expected. Hence, NGC\\,1275 itself has been suggested to be a potential TeV source as well, thus motivating deep VERITAS (Acciari et al. 2009b) and MAGIC (Aleksi\\'c et al. 2010) observations. These observations so far resulted only in upper limits to the VHE $\\gamma$-ray emission of the studied region. One should note, however, that both the low- and high-energy peaks of NGC\\,1275 (in the $\\nu F_\\nu$ representations) are located at substantially lower frequencies than those of typical LBLs. Indeed, the low-energy (synchrotron) emission components of 3C\\,66A and BL Lac peak around $10^{13-15}$\\,Hz, while it is around $10^{12}$\\,Hz in the case of NGC\\,1275 (Paper-I). Correspondingly, the high-energy emission component of 3C\\,66A (but not necessarily of BL Lac) peaks at higher photon energies than that observed in NGC\\,1275. Within such an interpretation, this is consistent with the softer MeV/GeV spectrum in NGC\\,1275 ($\\Gamma_{\\gamma} = 2.13 \\pm 0.02$) compared to 3C\\,66A ($\\Gamma_{\\gamma} = 1.97 \\pm 0.04$; Abdo et al. 2009b) as measured by the {\\it Fermi}-LAT. Note also that the X-ray spectrum of NGC\\,1275 is clearly ``rising'' in the $\\nu F_\\nu$ representation ($\\Gamma_{\\rm X}$ = 1.6$-$1.7; Balmaverde et al. 2006, Ajello et al. 2009), indicating that it is dominated by the low-energy portion of the IC emission, whilst the X-ray spectra of LBLs are typically very flat, suggesting a transition between the synchrotron and IC components ($\\Gamma_{\\rm X}$ $\\simeq$ 2; e.g., Ghisellini et al. 1998). Such observational differences may indicate that, unlike in the case of LBLs, the high-energy spectrum of NGC\\,1275 does not extend up to TeV photon energies.  An interesting comparison can be made to another nearby radio galaxy detected by {\\it Fermi}-LAT, namely M\\,87 which resembles NGC\\,1275 in many respects, and is an established TeV source (Abdo et al. 2009d). While M\\,87 is closer to us than NGC\\,1275 ($d_{\\rm L} = 16$\\,Mpc versus $d_{\\rm L} = 75.3$\\,Mpc), the GeV flux of M\\,87 is much lower than that of NGC\\,1275 ($F_{\\varepsilon_{\\gamma}>100\\,{\\rm MeV}} = (2.45 \\pm 0.63) \\times 10^{-8}$\\,ph\\,cm$^{-2}$\\,s$^{-1}$ versus $(2.31 \\pm 0.13) \\times 10^{-7}$\\,ph\\,cm$^{-2}$\\,s$^{-1}$), with similar LAT measured spectra ($\\Gamma_{\\gamma} = 2.26 \\pm 0.13$ versus $2.13 \\pm 0.02$). Note also that both radio galaxies are located at the centers of rich clusters, that the synchrotron emission components in both sources peak in the far infrared ($\\sim 10^{12}-10^{13}$\\,Hz), and the estimated jet powers are similar, $L_{\\rm j}\\sim 10^{44}$ erg s$^{-1}$ (Owen et al. 2002, Dunn \\& Fabian 2004). Thus, it may be surprising that only one of these has so far been detected at TeV photon energies. The detailed analysis of the spectral evolution of NGC\\,1275 within the {\\it Fermi}-LAT range reported in this paper may provide a viable explanation for such a behavior. In particular, as already emphasized above, we found that the epochs characterized by the flattest GeV continuum of this source, as well as the arrival times of the highest energy photons from the direction of NGC\\,1275, do not coincide with the epochs of the highest photon flux above $100$\\,MeV. This is clearly illustrated in Figure\\,5, which shows a correlation between $F_{\\varepsilon_{\\gamma}>100\\,{\\rm MeV}}$ and the photon indices emerging from the power-law fits. Because no significant flux or spectral changes were observed during the pre- and post-flare epochs A ($0-252$\\,day) and C ($294-374$\\,day) (see Figure\\,4), the average fluxes and photon indices for these time periods are reported. Note that the average fluxes of these pre- and post-flare epochs are comparable, whilst the corresponding photon indices differ significantly by $\\Delta \\Gamma_{\\gamma} \\simeq 0.2$. Moreover, the observed spectral evolution during one year of the {\\it Fermi}-LAT exposure reveals a hysteresis-like character, more clearly seen for the flaring period (epoch B, with a time bin of 2 weeks), followed by a gradual flattening in the subsequent decay phase. Further insights into the spectral evolution of NGC\\,1275 within the {\\it Fermi}-LAT photon energy range are provided by Figure\\,6, which shows the two SEDs for pre- and post-flare epochs A and C. Here the dotted line corresponds to the best ``power-law with an exponential cutoff'' fit function determined from an average $\\gamma$-ray spectrum, as given in Figure\\,3. Interestingly, the difference between the two SEDs consists of an excess at photon energies $\\varepsilon_{\\gamma} \\simeq 1-2$\\,GeV, with the low-energy $\\gamma$-ray flux remaining essentially unchanged between the two epochs. This implies that (1) the $\\gamma$-ray variability in NGC\\,1275 (and possibly other radio galaxies) may be restricted to $\\geq$\\,GeV photon energies, and that (2) the position of the peak in the high-energy spectral component (in the $\\nu F_{\\nu}$ representation) may change substantially even within the same object with no accompanying significant flux changes. Note in this context that both the VERITAS and MAGIC non-detections were obtained during the pre-flare epoch A ($164-206$\\,day; see Acciari et al. 2009b; Aleksi\\'c et al. 2010). Hence, the emerging conclusion is that \\emph{it is not the total flux above $100$\\,MeV which should play a major role in triggering TeV observations of steep-spectrum {\\it Fermi}-LAT sources, but instead it is the flux and photon index determined at higher photon energies ($\\geq$\\,GeV)}.  On the theoretical side, this conclusion could be possibly justified by noting that after a new episode of injection of freshly accelerated electrons into the emission zone (e.g., a downstream region of a shock), higher energy electrons may lag behind the lower-energy electrons. In such a case, as discussed previously in the context of blazar modeling (e.g., Kirk et al. 1998; Sato et al. 2008), an ``anti-clockwise'' hysteresis in the flux versus photon index plane may arise, similar to what we observed in NGC\\,1275 at $\\gamma$-ray photon energies.  Here, we comment on the question if NGC\\,1275 -- being a representative example of a low-power radio galaxy -- may be considered as a misaligned blazar, most likely of the LBL type. Assuming a homogeneous jet model, one should expect its jet Doppler factor $\\delta = \\Gamma_{\\rm j}^{-1} \\, (1-\\beta_{\\rm j} \\, \\cos \\theta)^{-1} = \\Gamma_{\\rm j}^{-1} \\, (1-\\sqrt{1-\\Gamma_{\\rm j}^{-2}} \\, \\cos \\theta)^{-1} \\simeq 1-2$, for the typically expected jet viewing angle of $\\theta \\simeq 20^{\\circ}-30^{\\circ}$ and jet bulk Lorentz factors $\\Gamma_{\\rm j} \\simeq 10$. Indeed, modeling of the broad-band emission of LBLs (beamed counterparts of radio galaxies such as NGC\\,1275 by assumption) requires $\\delta \\sim \\Gamma_{\\rm j} \\sim 10$. Thus, if the only difference between radio galaxies and blazars is due to the viewing angle,  this should manifest in (i) different observed positions of the spectral peaks in the $\\nu F_\\nu$ representation ($\\nu \\propto \\delta$), (ii) different observed variability patterns ($t_{var} \\propto \\delta^{-1}$, assuming the emission region is a moving source), and finally in (iii) different observed luminosities ($L_{obs} \\propto \\delta^4$ for a moving blob case, or $\\propto \\delta^3/\\Gamma_{\\rm j}$ for the steady jet (see Sikora et al. 1997).  Indeed, both the low- and high-energy $\\nu F_\\nu$ peaks of NGC\\,1275 are located at substantially lower frequencies than those of typical LBLs. Taking the difference of beaming factors into account (i.e., $\\delta_{\\theta = 0}/\\delta_{\\theta = 20^{\\circ}} \\sim 10$), the broad-band SED of the NGC\\,1275 would be similar to the SEDs of blazars such as BL Lacertae or 3C\\,66A. Note however that the position of the high-energy spectral peak may change substantially in a single object even for comparable flux levels, at least in NGC\\,1275, so the diagnostics related to the location of the spectral peaks may not be very conclusive.  Variability as short as day timescales is often observed in LBLs. For example, during the historical flare of BL Lacertae in 1997, correlated $\\gamma$-ray and optical flares were observed, with the $\\gamma$-ray flux increasing by a factor of $2.5$ within a day (Bloom et al. 1997). Similarly, daily variability has been discovered in both {\\it Fermi}-LAT and VERITAS observations of 3C\\,66A (Reyes et al. 2009). Hence, in the simple unification scheme outlined above, we should expect NGC\\,1275 to vary in $\\gamma$-rays on timescales of several days. This, however, is difficult to test even with the excellent sensitivity of the {\\it Fermi}-LAT instrument, due to the limited photon statistics for weekly time bins. The analysis presented in this paper gives instead only a robust upper limit $t_{\\rm var} \\le$ a few months. However, note that day-timescale variability has been detected at TeV photon energy range for the M\\,87 radio galaxy (Acciari et al. 2009c). Thus, more frequent monitoring of NGC\\,1275 by ground-based Cherenkov Telescopes would be valuable in this respect.  The observed photon flux of $F_{\\rm \\varepsilon_{\\gamma}>100 MeV} = (2.19 \\pm 0.13) \\times 10^{-7}$\\,ph\\,cm$^{-2}$\\,s$^{-1}$ implies an \\emph{observed} (isotropic) $\\gamma$-ray luminosity of NGC\\,1275, $L_{\\gamma} \\simeq 4 \\pi d_{\\rm L}^2 \\, (\\Gamma_{\\gamma}-1) \\, F_{\\varepsilon_{\\gamma}>\\varepsilon_0} \\, \\int_{\\varepsilon_0}^{\\varepsilon_{\\rm c}} (\\varepsilon/\\varepsilon_0)^{1-\\Gamma_{\\gamma}} \\, d \\varepsilon \\sim 10^{44}$\\,erg\\,s$^{-1}$, for $\\Gamma_{\\gamma} \\simeq 2.0-2.2$ and $\\varepsilon_{\\rm c} \\simeq 42$\\,GeV. This  is already comparable to the typical observed $\\gamma$-ray luminosities of LBLs, which range between $L_{\\gamma} \\sim 10^{44} - 10^{46}$\\,erg\\,s$^{-1}$. Note in particular that in the framework of the simple unification scheme discussed above, the beamed analog of NGC\\,1275 would then be characterized by the observed $\\gamma$-ray luminosity larger by a factor of $[\\delta_{\\theta = 0}/\\delta_{\\theta = 20^{\\circ}}]^4 \\sim 3 \\times 10^4$ (or at least $[\\delta_{\\theta = 0}/\\delta_{\\theta = 20^{\\circ}}]^3 \\sim 3 \\times 10^3$) than this, i.e. $L_{\\gamma} > 10^{47}$\\,erg\\,s$^{-1}$. Such luminosities are not expected for LBL-type blazars (see Abdo et al. 2010a). On the other hand, the observed $\\gamma$-ray luminosity of NGC\\,1275 is not energetically problematic, since the total \\emph{emitted} $\\gamma$-ray power in this source seems rather moderate \\emph{as long as the emitting plasma moves with highly relativistic bulk velocities}, namely $L_{\\rm \\gamma,\\,em} \\simeq (\\Omega_{\\rm j}/4 \\pi) \\, L_{\\gamma} \\simeq L_{\\gamma} / 4 \\Gamma_{\\rm j}^2 < 10^{42}$\\,erg\\,s$^{-1}$ for $\\Gamma_{\\rm j} \\sim 10$, where $\\Omega_{\\rm j} \\simeq \\pi \\theta_{\\rm j}^2$ is the solid angle defined by the jet opening angle $\\theta_{\\rm j}$, for which we assumed $\\theta_{\\rm j} \\sim 1/\\Gamma_{\\rm j}$. Such a relatively small emitted power would constitute less than $1\\%$ of the total kinetic power of the NGC\\,1275 jet, estimated by Dunn \\& Fabian (2004) to be roughly $L_{\\rm j} \\sim (0.3-1.3) \\times 10^{44}$\\,erg\\,s$^{-1}$. Yet the problem of an unexpectedly large observed $\\gamma$-ray luminosity of the beamed analog of NGC\\,1275 remains, and poses a serious challenge to the simplest version of the AGN unification scheme. In this context, a viable explanation for this problem would be to postulate that the high-energy emission observed from ``misaligned'' blazars such as NGC\\,1275 (or M\\,87) is dominated not by a jet ``spine'' characterized by large bulk Lorentz factors ($\\Gamma_{\\rm j} \\sim 10$ as is the case in bona-fide blazars), but by the slower jet boundary layers ($\\Gamma_{\\rm j} \\sim$ few) as discussed by several authors (e.g., Celotti et al. 2001, Stawarz \\& Ostrowski 2002, Ghisellini et al. 2005).  Alternatively, one may propose that the $\\gamma$-ray emission observed from radio galaxies is not produced within the ``proper'' blazar  emission zone, but at larger distances from the active center characterized by slower bulk velocities (say, $\\Gamma_{\\rm j} \\simeq$ few) of the emitting plasma. Yet another possibility may be that the inner jets in NGC\\,1275 (and also in similar objects) are in general intrinsically less relativistic than the ones in bona-fide blazars; this would be consistent with the conclusions of Lister \\& Marscher (1997), who argue that radio-loud AGN with nuclear jets characterized by $\\Gamma_{\\rm j} \\geq 10$ must be rather rare among the general population. Whichever scenario is correct, the {\\it Fermi} results seem to indicate that low-power radio galaxies are most likely not simple off-axis analogs of BL Lac objects in terms of their $\\gamma$-ray properties."
    },
    "1004/1004.2487_arXiv.txt": {
        "abstract": "A method for implementing cylindrical coordinates in the \\ath magnetohydrodynamics (MHD) code is described.  The extension follows the approach of \\athnosp's original developers and has been designed to alter the existing Cartesian-coordinates code \\citep{sto08} as minimally and transparently as possible.  The numerical equations in cylindrical coordinates are formulated to maintain consistency with constrained transport, a central feature of the \\ath algorithm, while making use of previously implemented code modules such as the Riemann solvers.  Angular-momentum transport, which is critical in astrophysical disk systems dominated by rotation, is treated carefully.  We describe modifications for cylindrical coordinates of the higher-order spatial reconstruction and characteristic evolution steps as well as the finite-volume and constrained transport updates.  Finally, we present a test suite of standard and novel problems in one-, two-, and three-dimensions designed to validate our algorithms and implementation and to be of use to other code developers.  The code is suitable for use in a wide variety of astrophysical applications and is freely available for download on the web. ",
        "introduction": "The \\ath code (\\citealt{gar05}, hereafter \\citetalias{gar05}; \\citealt{gar08}, hereafter \\citetalias{gar08}; \\citealt{sto08}) is a new, second-order Godunov code for solving the equations of ideal magnetohydrodynamics (MHD).  Among its salient features are that it preserves the divergence-free constraint, $\\nabla \\cdot \\mathbf{B} = 0$, to within machine round-off error via unsplit evolution of the magnetic field, and that it employs fully conservative updates of the MHD equations.  This last feature distinguishes \\ath from its predecessor, {\\it Zeus} \\citep{sto92a,sto92b}, which also preserves the divergence-free constraint, but employs operator-split finite-difference methods.  \\ath has been extensively tested via both comparison to analytic solutions, and comparison to the results of other numerical MHD codes.  The code package is freely available to the community, and is highly portable and easily configurable, as it is self-contained and does not rely on outside libraries other than MPI for computation on multi-processor distributed memory platforms.  The equations of ideal MHD consist of eight coupled partial differential equations, which are not analytically solvable in general, and fully three-dimensional numerical solutions can be quite costly.  For many astrophysical systems of interest, however, the computational cost for certain problems can be reduced by exploiting geometric symmetry.  For example, the high angular velocity of the plasma in accreting systems implies that most of the mass is confined within a disk.  If the properties are statistically independent of azimuthal angle, $\\phi$, these disks can be studied using radial-vertical ($R$-$z$) models, and if vertical variations are of lesser importance, these disks can be studied using radial-azimuthal ($R$-$\\phi$) models.  The dynamical properties of winds and jets from astrophysical systems can also be analyzed using axisymmetric models.  Exploiting symmetry in this way to reduce the effective dimension of the problem can greatly simplify the calculations involved and allow finer resolution when and where needed.  In addition, for either reduced-dimensional or fully three-dimensional problems, using a curvilinear coordinate system for rotating, grid-aligned flow is superior for preservation of total angular momentum, and renders imposition of boundary conditions much simpler compared to the Cartesian-grid case.  There are several other publicly available high-resolution shock-capturing codes for astrophysical MHD in wide use, including {\\it VAC} \\citep{tot96}, {\\it BATS-R-US} \\citep{pow99}, {\\it FLASH} \\citep{fry00}, {\\it RAMSES} \\citep{tey02}, {\\it NIRVANA} \\citep{zie04}, and {\\it PLUTO} \\citep{mig07}, to name a few.  Although these and other codes enjoy increasing popularity within the community, as of this writing only {\\it VAC} and {\\it PLUTO} have the capability for MHD in curvilinear coordinates.  In this paper, we describe our adaptation of \\ath to support cylindrical geometry, and present a suite of tests designed to validate our algorithms and implementation.  These tests include standard as well as novel problems, and may be of use to other code developers.  A guiding principal of our approach is to alter the existing \\ath code as minimally and as transparently as possible.  This will involve a careful formulation of the MHD equations so that the finite-volume algorithm remains consistent with constrained transport, and so that the built-in Riemann solvers (as well as computation of wavespeeds and eigenfunctions) need not be changed.  Finally, we pay particular attention to angular-momentum transport, which is critical in systems dominated by rotation.  The plan of this paper is as follows:  In \\S\\ref{eqns}, we describe the conservative system of mathematical equations that we shall solve, and in \\S\\ref{algover}, we briefly outline the main steps used in \\ath to evolve the system numerically.  In \\S\\ref{lin}, we describe the projected primitive variable system used in the reconstruction step.  In \\S\\S\\ref{recon} and~\\ref{charevo}, we describe the modifications needed for cylindrical coordinates in the higher-order spatial reconstruction and characteristic evolution steps, respectively.  In \\S\\S\\ref{fvm} and~\\ref{ct}, we describe the implementation in cylindrical coordinates of the finite volume and constrained transport updates, respectively, and then in \\S\\ref{alg}, we summarize the steps of the whole algorithm in detail.  In \\S\\ref{tests}, we present code verification tests and results, and we conclude in \\S\\ref{conclusion}.  Our version of the code, including the suite of test problems we have developed, is freely available for download on the Web.   { ",
        "conclusions": "\\label{conclusion}} We have described an adaptation of the \\ath astrophysical MHD code for cylindrical coordinates.  The original Cartesian code uses a combination of higher-order Godunov methods (based on the CTU algorithm of \\citetalias{col90}) to evolve the mass-density, momenta, and total energy, and constrained transport \\citep{eva88} to evolve the magnetic fields.  We have described modifications to the second- and third-order reconstruction schemes, the finite-volume and finite-area formulations of the MHD equations, and the inclusion of geometric source terms.  Our approach is advantageous in that it does not require modification to the majority of the existing code, in particular to the Riemann solvers and eigensystems.  Furthermore, our approach to implementing cylindrical coordinates could be applied in a straightforward manner to enable other curvilinear coordinate systems, such as spherical coordinates, in the \\ath code as well as other higher-order Godunov codes.  Finally, our code and test suite are publicly available for download on the Web.  The code is currently being used for a variety of applications, including studies of global accretion disks, and we hope it will be of use to many others studying problems in astrophysical fluid dynamics."
    },
    "1004/1004.4696_arXiv.txt": {
        "abstract": "Globular clusters have proven to be essential to our understanding of many important astrophysical phenomena. Here we analyse spectroscopic observations of ten Halo globular clusters to determine their dark matter content, their tidal heating by the Galactic disc and halo, describe their metallicities and the likelihood that Newtonian dynamics explain their kinematics.  We analyse a large number of members in all clusters, allowing us to address all these issues together, and we have included NGC 288 and M30 to overlap with previous studies.  We find that any flattening of the velocity dispersion profiles in the outer regions of our clusters can be explained by tidal heating.  We also find that all our GCs have M/L$\\rm _V\\lesssim5$, therefore, we infer the observed dynamics do not require dark matter, or a modification of gravity.  We suggest that the lack of tidal heating signatures in distant clusters indicates the Halo is not triaxial.  The isothermal rotations of each cluster are measured, with M4 and NGC 288 exhibiting rotation at a level of $0.9\\pm0.1$\\,km\\,s$^{-1}$ and $0.25\\pm0.15$\\,km\\,s$^{-1}$, respectively. We also indirectly measure the tidal radius of NGC 6752, determining a more realistic figure for this cluster than current literature values. Lastly, an unresolved and intriguing puzzle is uncovered with regard to the cooling of the outer regions of all ten clusters. ",
        "introduction": "Globular clusters (GCs) are often used as tracers of the gravitational potentials of galaxies and galaxy clusters \\cite[e.g.][]{Kissler-Patig99,Cote03,Wu06,Quercellini08,Gebhardt09}. Although this has been applied to theoretical Galactic potentials \\cite[e.g.][]{Allen06}, the actual Milky Way (MW) potential has not yet been analysed in this way. The tidal forces of spiral galaxies are thought to be strongest near the disc because the concentrated mass in that region (gas and stars) has a larger density gradient than the more slowly varying density of the dark matter (DM) halo. Interestingly, many distant MW objects such as GCs and dwarf galaxies are known to be tidally stripped, despite being far enough from the Disc that they should not directly interact with it.  For example, NGC 7492 is $\\sim3$\\,kpc further from the Galactic centre than the most distant detection of the Monoceros Ring, an object on the very outskirts of the Disc \\cite[$\\sim22$\\,kpc;][]{Conn07}, and exhibits clear evidence of tidal interaction with the Galaxy \\cite[][]{Lee04}. In this paper we consider 10 GCs at varying Galactocentric and Planar distances, allowing the inference of properties of the potential of the MW for the first time, by looking for signatures of tidal heating in these clusters.  Many globular clusters exhibit internal accelerations below $a_0\\approx1.2\\times10^{-10}$\\,m\\,s$^{-2}$, the level at which either modified gravity \\cite[e.g. MOND;][]{Milgrom83} or dark matter is required to reconcile the observed kinematics of elliptical galaxies with theory. Near the tidal radius ($r_t$) it is likely that most stars in GCs feel accelerations below this level, making them an ideal testing ground for low-acceleration gravity \\cite[][and references therein]{Sollima10}.  Furthermore, if all GCs exhibit similar behaviour, Galactic influences can not be the primary cause. In this final paper in the series [see also \\citealt{Lane09} (hereafter Paper I) and \\citealt{Lane10a} (hereafter Paper II)], we present the velocity dispersions and mass-to-light profiles of four GCs, namely M4, M12, NGC 288 (chosen for comparison with earlier studies) and NGC 6752, bringing the total for this project to 10. This sample allows statistically significant conclusions to be made on the dark matter content of Halo GCs, and on whether a modification of gravity is required to reconcile their internal kinematics with Newtonian gravitational theory.  Our sample of GCs contains three close to the Galaxy (M55, M12 and M22), four at intermediate distances (NGC~6752, M4, M30 and 47~Tuc) and three that are distant (M68, NGC~288 and M53).  We define `close' to be $R<5$\\,kpc, `intermediate' as $5<R<10$\\,kpc, and `distant' to be $R>10$\\,kpc, following \\cite{Harris96}.  NGC 288 was chosen, in part, because it is one of the GCs analysed by \\cite{Scarpa07b} who found it to have a flat velocity dispersion profile, similar to that of Low Surface Brightness galaxies which are thought to be DM dominated through to their cores.  Our targets were then analysed in separate studies (Papers I and II and the current paper) ensuring a mix of nearby, intermediate and distant GCs to ensure any Galactic influences, if any, would be clearly observed. See Table \\ref{metaltable} for the estimated acceleration, due to the cluster, of the most distant cluster member for all ten clusters analysed in this project. Note that the three distant clusters all experience accelerations due to the Galaxy of $\\sim a_0$.  \\section[]{Data Acquisition and Reduction}\\label{data}  We used AAOmega, a double-beam, multi-object spectrograph on the 3.9m Anglo-Australian Telescope (AAT) at Siding Spring Observatory in New South Wales, Australia, to obtain the data for this survey.  AAOmega covers a two-degree field of view, and is capable of obtaining spectra for 392 individual objects over this field. We used 30 sky fibres used for optimal sky subtraction, and 5--8 fibres for guiding.  The positional information for our targets was taken from the 2MASS Point Source Catalogue \\citep{Skrutskie06} which has an accuracy of $\\sim0.1''$.  Observations of M4 were performed on February 15--17, 2008, with $1.5''-2.5''$ seeing.  The data for M12 were taken over two observing runs: 7 nights on August 12--18 2006, and a further 8 nights on August 30 -- September 6 2007, both with mean seeing of $\\sim1.5''$. NGC 288 was observed during the 2006 run and NGC 6752 during the 2007 run. For all observations we used the 2500V grating in the blue arm, resulting in spectra between $4800{\\rm \\AA}$ and $5150{\\rm \\AA}$ with $\\lambda/\\Delta\\lambda = 8000$.  In the red arm we used the 1700D grating, which is optimized for the CaII IR triplet region.  The red spectra cover $8350-8790{\\rm \\AA}$, with $\\lambda/\\Delta\\lambda = 10000$.  This setup returns the highest spectral resolution available with AAOmega, and is suitable for measuring stellar radial velocities. We selected targets for this campaign by matching the $J-K$ colour and $K$ magnitude range of the red giant branch (RGB) of each cluster.  To minimize scattered-light cross-talk between fibres, each configuration was limited to 3 magnitudes in range.  We obtained 718, 2826, 1223 and 3664 spectra in the M4, M12, NGC 288 and NGC 6752 fields, respectively.  Flat-field and arc-lamp exposures were used to ensure accurate data reduction and wavelength calibration.  Data reduction was performed with the {\\tt2dfdr} pipeline\\footnote{http://www2.aao.gov.au/twiki/bin/view/Main/CookBook2dfdr}, which was specifically developed for AAOmega data. We checked the efficacy of the pipeline with a comparison of individual stellar spectra.  Radial velocities and atmospheric parameters were obtained through an iterative process, taking the best fits to synthetic spectra from the library by \\cite{Munari05}, degraded to the resolution of AAOmega, and cross-correlating this model with the observed spectra to calculate the radial velocity [a process very similar to that used by the Radial Velocity Experiment \\cite[RAVE;][]{Steinmetz06,Zwitter08} project]. We used the same spectral library as the RAVE studies; this process is outlined in detail by \\cite{Kiss07}.  \\subsection{Cluster Membership}\\label{Cluster Membership}  We determined cluster membership using four parameters: the equivalent width of the calcium triplet lines, surface gravity, radial velocity and metallicity ([m/H]).  Stars matching all criteria were judged to be members. Only stars having $\\log g<4.0$ and $\\log g<4.6$ were selected for NGC 6752 and M12, respectively, ensuring the majority of Galactic contaminants were removed before further selection criteria were applied.  This probably removed some genuine cluster members but was necessary to ensure our sample was as free from Galactic field stars as possible.  For several clusters studied in \\citetalias{Lane09}, a cutoff of $T_{\\rm eff}\\gtrsim9000$\\,K was necessary to remove hot horizontal branch (HB) stars.  These have radial velocities with large uncertainties due to the calcium triplet in very hot stars being replaced by hydrogen Paschen lines.  No cuts were made on $T_{\\rm eff}$ for any of the current clusters because no stars with $T_{\\rm eff}\\gtrsim7425$\\,K (for M4), $T_{\\rm eff}\\gtrsim5600$\\,K (for M12), $T_{\\rm eff}\\gtrsim7000$\\,K (for NGC 288) or $T_{\\rm eff}\\gtrsim5500$\\,K (for NGC 6752) remained after our selection process.  In total, 200, 242, 133 and 437 stars were found to be members of M4, M12, NGC 288 and NGC 6752, respectively.  Figure \\ref{members} shows the relative locations of the observed stars and highlights those found to be members. Note that 19 stars in the M4 field were found to be members of the globular cluster NGC 6144. For M4, we find that 88.0\\% of the selected members fall within within $2\\sigma$ of $all$ selection parameters and 100\\% within $3\\sigma$. For M12 these values are: 94.2\\% and 100\\%, for NGC 288: 89.5\\% and 100\\% and for NGC 6752: 97.5\\% and 100\\%. Based on this we see no statistical reason to think there is significant Galactic contamination in our final samples.  \\begin{figure*} \\begin{centering} \\includegraphics[angle=-90,width=0.48\\textwidth]{figures/M4_members.ps} \\includegraphics[angle=-90,width=0.48\\textwidth]{figures/M12_members.ps} \\includegraphics[angle=-90,width=0.48\\textwidth]{figures/NGC288_members.ps} \\includegraphics[angle=-90,width=0.48\\textwidth]{figures/NGC6752_members.ps} \\caption{Distribution on the sky of the stars observed in the four fields, with axes in both degrees and parsecs from the cluster centre. Circled points indicate stars that we determined to be cluster members (see text).  The large circle is the tidal radius of the cluster from \\citet{Harris96}, with the smaller thinner circle in the lower right panel being our derived value for $r_t$ for NGC 6752 (see text).  The large points in the upper right of the M4 field are the stars we determined to be members of NGC 6144.  In each panel, North is up and East is to the left.} \\label{members} \\end{centering} \\end{figure*} ",
        "conclusions": "In the current paper we have analysed four GCs (M4, M12, NGC 288 and NGC 6752) to determine their velocity dispersion and M/L$_{\\rm V}$ profiles, bringing the total to 10 for this project. We have included GCs that have external accelerations extending from above $a_0$ down to $a_0$ and we find no deviation from our Plummer models at lower accelerations. Therefore, we see no indication that DM, or a modified version of gravitational theory, is required to reconcile GC dynamics with Newtonian gravity. This adds to the growing body of evidence that GCs are DM-poor, and that our understanding of weak-field gravitation is not incomplete. Within the stated uncertainties, the dynamics of all these clusters are well described by purely analytic \\citet{Plummer11} models, which indicates that Newtonian gravity adequately describes their velocity dispersions, and we observe no breakdown of Newtonian gravity at $a_0\\approx1.2\\times10^{-10}$\\,m\\,s$^{-2}$, as has been claimed in previous studies.  Despite this, we see the intriguing possibility of an unknown cooling process in the outskirts of GCs; the external regions of our GCs seem to cool much faster following tidal Disc shocks than the relaxation time of the clusters.  Because it is highly unlikely that a MONDian process, or a significant DM component, is the cause of this cooling (based on our velocity dispersion and M/L$_{\\rm V}$ profiles), further work is required to solve this puzzle. Furthermore, the lack of tidal heating events in the distant clusters provides some indication that the dark Halo is not triaxial.  The Plummer model was also used to determine the total mass, scale radius, and M/L$_{\\rm V}$ profile for each cluster.  We find that none of our clusters have M/L$_{\\rm V}\\gg1$, further evidence that DM does not dominate.  We have produced M/L$_{\\rm V}$ profiles, rather than quoting a single value based on the central velocity dispersion and central surface brightness.  This method is used because it describes the M/L$_{\\rm V}$ of the entire cluster, rather than only its core. This is particularly important for post-core-collapsed GCs, where crowding and confusion effects introduce significant uncertainty into luminosity and kinematic measurements at small radii.  Within the uncertainties, our estimated cluster masses all match those in the literature except for M4, which we calculate to have a total mass about twice that of the literature values.  The reason for this discrepancy is that the tidally heated cluster has an increased velocity dispersion in its outer regions, flattening the Plummer fit, increasing the value of $r_s$, and therefore, increasing the mass estimate.  Another important result from this study is the measured rotations of our clusters.  Of the four clusters studied here, M4 and NGC 288 show clear rotation, M12 may have some rotation, and NGC 6752 displays no rotation signature.  Throughout this project we have found similar results for the dark matter content, and Newtonian kinematics, of our 10 GCs, all at varying distances from the Galactic centre and Disc, including three that experience external accelerations due to the Galaxy of $\\sim a_0$. All data were acquired using the same instrument (AAOmega on the Anglo-Australian Telescope), reduced using the same pipeline ({\\tt 2dfdr}), and analysed in the same way.  This homogeneous approach is vital to a large project such as this, to ensure all systematics are accounted for in a similar fashion.  Because of all these factors, our results from the three papers are $strongly$ indicative that the current picture of globular clusters being dark-matter poor, and with dynamics explained by standard Newtonian theory, is correct."
    },
    "1004/1004.4249_arXiv.txt": {
        "abstract": "Fermi Gamma ray Space Telescope measurements of spectra,  variability time scale, and maximum photon energy give lower limits to the apparent jet powers and, through $\\gamma\\gamma$ opacity arguments, the bulk Lorentz factors of relativistic jets.  The maximum cosmic-ray particle energy is limited by these two quantities in Fermi acceleration scenarios. Recent data are used to constrain the maximum energies of cosmic-ray protons and Fe nuclei accelerated in colliding shells of GRBs and blazars. The Fermi results indicate that Fe rather than protons are more likely to be accelerated to ultra-high energies in AGNs, whereas powerful GRBs can accelerate both protons and Fe to $\\gtrsim 10^{20}$ eV. Emissivity of nonthermal radiation from radio galaxies and blazars is estimated from the First Fermi AGN Catalog, and shown to favor BL Lac objects and FR1 radio galaxies over flat spectrum radio quasars, FR2 radio galaxies, and long-duration GRBs as the sources of UHECRs. ",
        "introduction": "\\citet{hil84} pointed out an essential requirement for acceleration of ultra-high-energy cosmic rays (UHECRs), namely that the particle Larmor radius $r_{\\rm L}\\cong E/QB$ must be smaller than the size scale of the acceleration region. Here $E$ is the particle energy, $Q=Ze$ is its charge, and $B$ is the magnetic field in the acceleration zone. This limitation applies to Fermi acceleration scenarios where a particle gains energy while diffusing through a magnetized region. Additional limitations due, for example, to radiative losses or available time, further restrict the maximum energies and therefore the allowed sites of UHECR acceleration.  Two plausible classes of astrophysical accelerators of extragalactic UHECRs are active galactic nuclei (AGN) \\citep{mb89,bgg02} and gamma-ray bursts (GRBs) \\citep{wax95,vie95,mu95} \\citep[see also][]{rm98,hh02,dm09}, though other types of sources, including young, highly magnetized neutron stars \\citep{ghi08} and structure formation shocks \\citep{ino08} remain viable. The announcement by the \\citet{Auger07} of anisotropy in the arrival directions of cosmic rays with  energies $E \\gtrsim 6\\times 10^{19}$ eV, even given the reduced correlation in the latest data from the Pierre Auger Observatory \\citep{abr09}, is compatible with the production of UHECRs in many source classes, including GRBs and blazars. Because of the GZK effect involving photohadronic interactions of protons or ions  with CMB radiation \\citep{gre66,zk66,ste68}, higher-energy cosmic rays with $E\\gtrsim 10^{20}\\,\\rm eV$ must be produced by sources located within distances $d\\lesssim 100$ Mpc in order to reach us without losing significant energy \\citep[e.g.,][]{nw00,hmr06}.  The most powerful AGNs and long-duration GRBs are found far outside the GZK radius, at redshifts $z\\gtrsim 0.1$. It is therefore of interest to re-examine Fermi acceleration requirements to determine if there are AGN and GRB sources with appropriate properties within the GZK radius. Here we make a detailed examination to justify a simple derivation of maximum particle energy relating apparent source power and bulk Lorentz factor $\\Gamma$ in the framework of Fermi acceleration in colliding shells. (Note that these arguments do not apply to non-Fermi type mechanisms, for example, electrodynamic acceleration in pulsar magnetospheres.) The derived limits are compared with values implied by Fermi data, yielding constraints on UHECR acceleration in these sources. We then use the First Fermi Large Area Telescope (LAT) AGN Catalog (1LAC) \\citep{abd10a} to estimate the nonthermal emissivity of AGNs.  We find that the lower luminosity BL Lac objects and FR1 radio galaxies are more likely to be the sources of UHECRs than the rare, powerful flat spectrum radio quasars (FSRQs) and FR2 radio galaxies, and are more likely to accelerate Fe than protons to ultra-high energies. GRBs, on the other hand, can accelerate both protons and Fe nuclei to ultra-high energies, but are rare within the GZK volume. ",
        "conclusions": "Fermi observations shown in Fig.\\ 3 indicate that FR1 radio galaxies and misaligned BL Lac objects located within the GZK radius have sufficient emissivity to power the UHECRs. With typical Lorentz factors $\\approx 2$ -- 10, and apparent jet powers $\\approx 10^{44}$ -- $10^{45}$ erg s$^{-1}$ (which could exceed $10^{46}$ erg s$^{-1}$ and large Lorentz factors during flaring episodes), Fig.\\ 2 shows that acceleration of Fe nuclei in FR1 radio galaxies is possible in colliding shells made in the jets of these galaxies. Given the favorable circumstances needed for colliding shells to accelerate UHECRs, including large Lorentz factor contrast and short times between shell ejections, the acceleration of protons is less likely. The $L-\\Gamma$ constraint is unfavorable for UHECR acceleration at sites with low apparent luminosity, such as starburst galaxies or the lobes of radio galaxies.  Long-duration GRBs have sufficient power to accelerate cosmic rays to ultra-high energies, but their local photon luminosity density in photons, $\\sim 10^{43}$ -- $10^{44}$ erg Mpc$^{-3}$ yr$^{-1}$, implies comparable or large baryon loading in most models for UHECR acceleration. When compared with the clearly nonthermal Fermi LAT flux, the required baryon-loading becomes significant, as shown by \\citet{eic10}. The local nonthermal luminosity density of FR1 radio galaxies and BL Lac objects by far dominates that of GRBs, especially when compared only with the LAT fluxes from GRBs and blazars. This circumstance favors UHECR acceleration by the supermassive black-hole engines in radio galaxies and blazars, provided that UHECRs are predominantly Fe ions."
    },
    "1004/1004.3375_arXiv.txt": {
        "abstract": "Multivariate methods have been recently introduced and successfully applied for the discrimination of signal from background in the selection of genuine very-high energy gamma-ray events with the H.E.S.S. Imaging Atmospheric Cerenkov Telescope. The complementary performance of three independent reconstruction methods developed for the H.E.S.S. data analysis, namely Hillas, model and 3D-model suggests the optimization of their combination through the application of a resulting efficient multivariate estimator. In this work the boosted decision tree method is proposed leading to a significant increase in the signal over background ratio compared to the standard approaches. The improved sensitivity is also demonstrated through a comparative analysis of a set of benchmark astrophysical sources. ",
        "introduction": "In the past decade, a new astronomical window has been opened thanks to the last generation of ground-based Imaging Atmospheric Cerenkov Telescopes (IACTs). Before the construction of IACT arrays such as H.E.S.S., M.A.G.I.C. and V.E.R.I.T.A.S., only a few very-high energy (VHE) \\gr\\ sources ($>$100~GeV) were known. This new generation of experiments has resulted in the discovery of many tens of galactic and extra-galactic \\gr\\ sources.  The H.E.S.S. system is currently the most efficient instrument to look at the inner part of the Galactic plane. The system is composed of four IACTs and provides a sensitivity to a 1\\% of Crab Nebula flux in around 25 h of observations~\\cite{Vincent08}. A systematic survey of about a third of the Galactic plane has been conducted since the beginning of the observations in full operation mode in 2004, leading to  the discovery of more than 50 sources within our Galaxy~\\cite{Aharonian05}\\cite{Aharonian06c}.  IACTs detect the Cerenkov light emitted by the secondary particle showers generated by the interaction of the incoming \\gr\\ into the atmosphere. They face a dominant background due to the hadron induced showers in the research of \\gr\\ signal. Three alternative reconstruction and discrimination methods have been developed and applied so far to the H.E.S.S. data analysis, namely Hillas, model and 3D-model. They have been individually improved and updated in the last years. The Xeff multivariate analysis method has been recently introduced~\\cite{Dubois09} in the \\hess\\ data analysis, increasing the discrimination power of genuine VHE gamma-ray event signals from the cosmic-ray background and improving the reconstruction performance (e.g. energy and direction reconstructions) through the combination of the three methods together. In this work the optimization of the multivariate analysis is presented through the application of a boosted decision tree (BDT) method leading to a single alternative discriminating estimator. After describing the methodology, some examples of application of the proposed multivariate method are presented in order to demonstrate the achieved gain in terms of sensitivity and precision. ",
        "conclusions": "The discrimination between \\gr\\ events and hadron induced background events is a key issue for ground based Cerenkov telescopes such as \\hess . A multi-variate analysis based on boosted decision trees has been studied. Three analysis methods are currently at work for the analysis of \\hess\\ data. The main discriminating variables from these original methods have been combined. The discrimination has been increased including the difference between the reconstructed direction of the various methods. The boosted decision trees have been trained in several bands in zenith angle and reconstructed energy in order to improve the rejection all over the energy range of the experiment and in all the observation conditions. This leads to a sizable improvement of the sensitivity. The signal over background ratio is dramatically increased compared to the original methods. The method has been tested on several reference sources which represent a wide range of sources in term of extension, nature, and observations conditions. The improvement in term of signal over background ratio and significance of the sources is illustrated. The application of this methods results also in a broader energy range for the spectral fit of faint sources compared to the previous methods. The robustness of the analysis in term of spectral reconstruction has been also demonstrated. The improved discrimination brings also a substantial gain in the angular resolution of the analysis."
    },
    "1004/1004.4625_arXiv.txt": {
        "abstract": "{Chemical element abundances for distant Galactic globular clusters (GCs) hold important clues to the origin of the Milky Way halo and its substructures.} {We study the chemical composition of red giant stars in Pal~4 --- one of the most remote GCs in the Milky Way --- and compare our abundance measurements to those for both low surface brightness dwarf galaxies, and GCs in the inner and the outer halo.} {By co-adding high-resolution, low-S/N Keck/HIRES spectra of 19 stars along the red giant branch, we estimate chemical abundance ratios of 20 $\\alpha$-, iron peak-, and neutron-capture elements. Our method gives total uncertainties on most element-to-iron ratios of typically 0.2 dex.} {We measure ${\\rm [Fe/H]} = -1.41\\pm0.04~{\\rm (statistical)} \\pm0.17~{\\rm (systematic)}$ and an $\\alpha$-enhancement of [$\\alpha$/Fe] = +0.38$\\pm0.11$~dex, which is consistent with the canonical value of $\\sim$ +0.4 dex found for Galactic halo field stars and most halo GCs at this metallicity. Although Pal~4 has higher enhancements in the heavier elements with respect to the halo, the majority of the element ratios are, within the measurement errors, consistent with those for  local halo field stars. We find, however, evidence for a lower [Mg/Ca] ratio than in other halo clusters.} {Based on the available evidence, we conclude that the material from which Pal~4 and the Galactic halo formed experienced similar enrichment processes, despite the apparently younger age of this cluster. Within the limitations of our methodology we find no significant indication of an iron spread, as is typical of genuine GCs of the Milky Way. However, abundance ratios for individual stars in Pal~4 and other distant satellites are urgently needed to understand the relationship, if any, between remote GCs and other halo substructures (i.e., luminous and ultra-faint dwarf spheroidal galaxies).} ",
        "introduction": "As the oldest readily identifiable stellar systems in the universe, globular clusters (GCs) are important tracers of the formation and early evolution of galaxies, the Milky Way (MW) included. Noting the apparent lack of a metallicity gradient among remote Galactic GCs, Searle \\& Zinn (1978) proposed an accretion origin for the Galactic halo extending over a period of several Gyr.  Evidence for this picture of hierarchical halo growth has come from the existence of a second-parameter problem among outer halo GCs (e.g., Catelan 2000; Dotter et al. 2010), which points to a significant age range within this population.  The remote GC Pal 4 is such an example of a second parameter cluster. At a Galactocentric distance of $R_ G = 109$ kpc (Stetson et al. 1999), it is one of only a few halo GCs at distances of  $\\sim$100 kpc or beyond. With a half-light radius of $r_h \\approx 23$ pc, it is also one of the most extended Galactic GCs currently known, being significantly larger than ``typical\" GCs in the Milky Way or external galaxies (which have $\\langle{r_h}\\rangle \\approx 3$~pc; see, e.g., Jord\\'an et~al. 2005). In fact, with a total luminosity of just $L_V \\sim 2.1\\times10^4~L_{V,{\\odot}}$, it is similar in several respects to some of the more compact ``ultra-faint'' dwarf spheroidal (dSph) galaxies (Simon \\& Geha 2007) that are being discovered  in the outer halo with increasing regularity (e.g., Belokurov et~al. 2007). Since almost nothing is known about their proper motions and internal dynamics, the relationship of faint, extended GCs like Pal~4 to such low-luminosity galaxies is currently an open question.  While there is a general consensus that Pal~4 is likely to be $\\approx$ 1--2 Gyr younger than inner halo GCs of the same metallicity, such as M5, age estimates in the literature do not fully agree (e.g., Stetson et al. 1999 vs. Vandenberg 2000). In particular, Stetson et al. (1999) note that an age difference with respect to the inner halo GCs could be explained if ``either [Fe/H] or [$\\alpha$/Fe] for the outer halo clusters is significantly lower than ... assumed''. Conversely, Cohen \\& Mel\\'endez (2005a) found that the outer halo GC NGC~7492 (R$_{\\rm GC}$ = 25 kpc) shows chemical abundance patterns that are very similar to inner halo GCs like M3 or M13. This similarity in the chemical enrichment now appears to extend into the outermost halo for at least some GCs: it has recently been shown that the abundance ratios in the remote (R$_{\\rm GC}$ = 92 kpc) cluster Pal~3 (Koch et al. 2009; hereafter Paper~I) bear a close resemblance to those of inner halo GCs. The chemical abundance patterns of remote halo GCs like Pal~3 and Pal~4 are important clues to the formation of the Milky Way, as they allow for direct comparisons to those of the  dSph galaxies, which are widely believed to have been accreted into the halo (e.g., Klypin et~al. 1999; Bullock et~al. 2001; Font et~al. 2006; Li et~al. 2009). In this spirit, Mackey \\& Gilmore (2004) conclude that all young halo clusters (i.e., 30) did not originate in the MW but were donated by at least seven mergers with ``cluster-bearing'' dSph-type galaxies.  There are, however, no high-dispersion abundance data yet published for this remote cluster. Previous low-resolution spectroscopic and photometric studies have established Pal~4 as a mildly metal-poor system, with [Fe/H] estimates ranging from $-1.28$ to $-1.7$ dex (Armandroff et al. 1992; Stetson et al. 1999; Kraft \\& Ivans 2003). In this paper, one of a series, we aim to extend the chemical element information for GCs in the Galactic halo out to the largest possible distances, and to carry out a first analysis of Pal~4's chemical abundance patterns. As we have shown in Paper~I, which presented a similar analysis for Pal~3, it is possible to derive reliable abundance measurements for remote Galactic GCs by performing an integrated analysis of stacked, low signal-to-noise (S/N) --- but high-resolution --- spectra (see also McWilliam \\& Bernstein 2008). Note, however, that this method presupposes that there is no significant abundance scatter present along the RGB and that all stars have the same mean abundances for all chemical elements. We have therefore no means of distinguishing Pal~4 as a genuine GC with no internal abundance spread from a dSph that may have a very broad abundance range (e.g., Shetrone et~al. 2001, 2003; Koch 2009), nor of discerning any intrinsic abundance variations (e.g., Lee et al. 2009).  We will return to this question in Section~5.2. Neverthess, such studies can provide an important first step towards an overall characterization of the chemical element distribution, and enrichment history, of the outer halo. ",
        "conclusions": "Motivated by the good agreement between the abundance ratios measured from high-S/N spectra of individual stars in Pal~3 and those found using co-added, low-S/N spectra (Paper~I), we have used the same technique to measure chemical abundance ratios in the remote halo GC Pal~4. Although systematic uncertainties and the low S/N ratios complicate such studies, an accuracy of 0.2 dex is possible for most abundance ratios, sufficient to place such faint and remote systems into a context with both the inner and outer halo GCs, as well as dSph and UF-dSph galaxies. In the future, this technique may enable the global abundance patterns to be characterized in additional remote systems, allowing a first reconnaissance of the chemical enrichment histories of remote Galactic satellites.  Perhaps the most striking finding in Pal~4 is the subsolar [Mg/Ca] ratio, which is not observed in the sample of reference GCs that span a broad range of Galactocentric distances. Despite an overlap of our observed ratio with the halo field population, its low value may rather resemble the low-[Mg/Ca] tail of the distribution for dSph stars. In contrast, we see tentative evidence for a solar [Ba/Y] ratio, which militates against a slow chemical evolution and accompanying AGB enrichment as suggested by enhanced [Ba/Y] values in about two thirds of the dSph stars studied to date. Overall, most of the element ratios determined in this study overlap with the corresponding measurements for halo field stars, although a few ratios seem to fall above the halo star trends (see \\S5). This favors a scenario in which the material from which both Pal~4 and the Galactic halo formed underwent rather similar enrichment processes.  In their analysis of the CMD of Pal~4, Stetson et al. (1999) state that the cluster is younger than the inner halo GC M5 by about 1.5 Gyr (at [Fe/H]=$-$1.33 dex; Ivans et al. 2003; Koch \\& McWilliam 2010) {\\em if} they ``all have the same composition -- and [...] this means both [Fe/H] and [$\\alpha$/Fe]''. Our work has shown that Pal~4 is enhanced by +0.38$\\pm$0.11~dex in the $\\alpha$-elements, which is  consistent with the value of 0.3 dex   assumed in the above CMD modeling. On the other hand, the CMD analysis suggested an [Fe/H] of $-$1.28 dex, which is slightly more metal rich than what we found in the present spectroscopic study: $\\langle$[Fe/H]$\\rangle$ = $-$1.41 dex. As noted in Vandenberg (2000), ``an increase in [Fe/H] or [$\\alpha$/Fe] would result in slightly younger [...] ages'' for Pal~4 (as determined via the magnitude offset between the horizontal branch and main-sequence turnoff). This would imply that Pal~4 is slightly older than found in Stetson et~al (1999) and hence more similar in age to the older halo population. This, however, is in contradiction to the younger age suggested by its peculiar (i.e., red) horizontal branch morphology, unless further parameters, such as red giant mass loss, are invoked  (Catelan 2000).  Based on the evidence at hand, Pal~4 seems to have an abundance pattern that is typical of other remote GCs in the outer halo. An open question, given the nature of our analysis which relies on co-adding individual RGB star spectra, is whether Pal~4 is monometallic or, like dSph and UF-dSph galaxies, shows an internal spread in metallicity. We argued in Sect.~5.2. judging from our limited quality spectra, however, that it is unlikely that this object exhibits any significant intrinsic iron scatter. It is clear that high-quality abundance ratio measurements for individual stars in Pal~4 and other remote substructures are urgently needed to understand the relationship, if any, between remote GCs and other substructures in the outer halo."
    },
    "1004/1004.4139_arXiv.txt": {
        "abstract": "{The distribution of chemical abundances and their variation in time are important tools to understand the chemical evolution of galaxies: in particular, the study of chemical evolution models can improve our understanding of the basic assumptions made for modelling our Galaxy and other spirals.} {To test a standard chemical evolution model for spiral disks in the Local Universe and study the influence of a threshold gas density and different efficiencies in the star formation rate (SFR) law on radial gradients (abundance, gas and SFR). The model will be then applied to specific galaxies.} {We adopt a one-infall chemical evolution model where the Galactic disk forms inside-out by means of infall of gas, and we test different thresholds and efficiencies in the SFR. The model is scaled to the disk properties of three Local Group galaxies (the Milky Way, M31 and M33) by varying its dependence on the star formation efficiency and the time scale for the infalling gas into the disk. } {Using this simple model we are able to reproduce most of the observed constraints available in the literature for the studied galaxies. The radial oxygen abundance gradients and their time evolution are studied in detail. The present day abundance gradients are more sensitive to the threshold  than to other parameters, while their  temporal evolutions are more dependent on the chosen SFR efficiency. A variable efficiency along the galaxy radius can reproduce the present day gas distribution in the disk of spirals with prominent arms. The steepness in the distribution of stellar surface density differs from massive to lower mass disks, owing to the different star formation histories. } {The most massive disks seem to have evolved faster  (i.e. with more efficient star formation) than the less massive ones, thus suggesting a downsizing in star formation for spirals. The threshold and the efficiency of star formation play a very important role in the chemical evolution of spiral disks and an efficiency varying with radius can be used to regulate the star formation. The oxygen abundance gradient can steepen or flatten in time depending on the choice of this parameter.} ",
        "introduction": "The study of the chemical evolution of nearby spiral galaxies is very important to improve our knowledge about the main ingredients used in chemical evolution models and to test the basic assumptions made for modelling our Galaxy. M31 and M33 are other spiral members of the Local Group of galaxies and during recent years many observational studies have been made to investigate the chemical and dynamical properties of these neighbouring systems. New surveys (Braun et al. 2009, Magrini et al. 2007,2008) contributed to the analysis of different  stellar populations and provided more accurate data to constrain  the chemical evolution models.  The disks of M31 and M33 have many similarities with the Milky Way disk but some observational constraints like the present day gas distribution can only be explained by assuming different star formation histories for these galaxies. The SFR is one of the most important parameters regulating the chemical evolution of galaxies (Kennicutt 1998, Matteucci 2001, Boissier et al. 2003) together with the initial mass function (IMF).    Another important mechanism is the \"inside-out\" disk formation that is very important to reproduce the radial abundance gradients (see Colavitti et al. 2008 for the most recent paper on the subject). A faster formation of the inner disk relative to the outer disk was originally proposed by Matteucci \\& Fran\\c cois (1989)and supported in the following years by Boissier \\& Prantzos (1999) and Chiappini et al. (2001).    The chemical evolution of M31 in comparison with that of the Milky Way has been already discussed by Renda et al. (2005) and Yin et al. (2009). Renda et al. (2005) concluded that while the evolution of the MW and M31 share similar properties, differences in the formation history of these two galaxies are required to explain the observations in detail. In particular, they found that the observed higher metallicity in the M31 halo can be explained by either (i) a higher halo star formation efficiency, or (ii) a larger reservoir of infalling halo gas with a longer halo formation phase. These two different pictures would lead to (a) a higher [O/Fe] at low metallicities, or (b) younger stellar populations in the M31 halo, respectively. Both pictures result in a more massive stellar halo in M31, which suggests  a possible correlation between the halo metallicity and its stellar mass. Yin et al. (2009) concluded that M31 must have been more active in the past than the Milky Way although its current SFR is lower than in the Milky Way, and that our Galaxy must be a rather quiescent galaxy, atypical of its class (see also Hammer et al. 2007). They also concluded that the star formation efficiency in M31 must have been higher by a factor of two than in the Galaxy. However, by adopting the same SFR as in the Milky Way they failed in reproducing the observed radial profile of the star formation and of the gas, and suggested that possible dynamical interactions could explain these distributions.  Magrini  et al. (2007) computed the chemical evolution of the disk of M33: they claimed to reproduce the observational features by assuming a continuous almost constant infall of gas.  In this work we present a one-infall chemical evolution model for the Galactic disk based on an updated version of the Chiappini et al. (2001) model.  This model can predict the evolution of the abundances of 37 chemical elements from the light to the heavy ones. We use this model to reproduce the chemical evolution of the Milky Way disk and that of the two nearby spiral galaxies (M31 and M33). To do that,  we assume that the disk of each galaxy  formed by gas accretion and vary the star formation efficiency as well as the gas accretion timescale. The similarities and the differences between the chemical evolution of these objects and the Milky Way are discussed to provide a basis for the understanding of the chemical evolution of disks.  The paper is organized as follows: in section 2 we describe our chemical evolution model and the assumptions made for each galaxy. In section 3 we present the results for the models and these results are discussed in detail in section 4. Finally in section 5 we summarize our conclusions. ",
        "conclusions": "\\subsection{The MW}  We found that the oxygen gradient along the disk of the Milky Way is well reproduced if an inside-out disk formation is assumed together with a threshold in the star formation of $7M_{\\odot}pc^{-2}$ or $4M_{\\odot}pc^{-2}$ , in agreement with previous works (CMR2001, Colavitti et al. 2008). The present time radial oxygen gradient is very dependent on the threshold in the star formation while it seems not to be so sensitive to the efficiency ($\\nu$) of the SFR.  The oxygen gradient can either flatten or steepen in time according to the assumption made on the star formation efficiency as a function of galactocentric distance. Models with a constant $\\nu$ tend to predict a steepening of the gradients in time, whereas those with a $\\nu$ decreasing with the radius  tend to flatten (in agreement with some recent observations of Maciel et al. 2003) Clearly the gradient evolution with time is strongly related to the assumed history of star formation in the disk.  The present-day gas profile in the MW is better reproduced by the model with a threshold of $4M_{\\odot}pc^{-2}$ and $\\nu(R)$. All models predict a lower SFR for the inner disk of the Galaxy but are in very good agreement with the observed data in the solar neighbourhood and in the outer parts of the disk. The higher SFR in the inner parts of the MW disk can be due to the presence of a bar, as suggested in Portinari \\& Chiosi (2000), and therefore cannot be reproduced by simple chemical evolution models.  The stellar surface density of the MW is in agreement with the observed values but models with a smaller threshold (MW-A2 and MW-B) overestimate the stellar content in the outer region of the disk. Unlike other observational constraints, the variable efficiency in the SFR does not play an important role for the results relative to the stellar sufarce density, indicating that the threshold is a stronger mechanism to regulate it.  In summary, the model that fits best the observational constraints for the Milky Way is the model with variable efficiency for the the star formation (MW-B).   \\subsection {M31}  The evolution of the disk of M31 is well reproduced by assuming a faster evolution (faster means a more intense SFR which is due both to the higher efficiency of SF and to the shorter infall timescale) than in the disk of the MW and a higher star formation threshold. Since the disk of M31 is more massive than the MW one, this implies that more massive disks should form faster and therefore that they are older than less massive ones (see also Boissier et al. 2003). The O abundance gradient from HII regions is well reproduced by all models for M31. This result confirms the predictions for the MW, showing that the present day abundance gradient is not so sensitive to the changes in the star formation efficiency. On the other hand, the time evolution of the O gradient is very dependent on this efficiency, steepening or flattening in time according to the chosen $\\nu$.  This fact can be confirmed with the results obtained for M31-Bk1.25 where we used the same efficiency but a different exponent for the SFR.   Models with constant efficiency in the star formation (M31-A1 and M31-A2) provide an exponential distribution of the present day gas surface density,  while models M31-B and M31-Bk1.25 with variable efficiency predict a more realistic scenario with a peak in the gas distribution around 12 kpc which can be related to M31 spiral arms.  The stellar density profile is flatter than the one predicted for the MW and M33 and all models show a similar distribution.  The predicted SFR for M31 is very similar in all models, specially M31-A2 and M31-B that present a smaller star formation after the peak in 12 kpc  In summary, the best model for the M31 disk is also the one with a star formation efficiency varing through the disk with a lower exponent in the SF law (M31-Bk1.25).  \\begin{figure}[ht!] \\includegraphics[width=0.45\\textwidth]{stars_v10.eps} \\caption{Comparative plot of stellar surface density for all galaxies in function of Radius in kpc. For the MW solid line for MW-A1, dotted line for MW-A2 and dashed line for MW-B, for M31 solid line represents M31-A1, dotted line M31-A2, dashed line M31-B and  dot-dashed line for M31-Bk1.25, finally for M33 solid line is for M33-A05, dotted line for M33-01 and dashed line for M33-B. See table 2 for more details. The shaded area corresponds to a scaled exponential distribution with $R_{Dstars} = 2.5$kpc (Freudenreich 1998) using the local values for the stellar density ($35 M_{\\odot}pc^{-2}$, Gilmore et al. 1989).} \\label{stars} \\end{figure}  \\begin{figure}[ht!] \\includegraphics[width=0.45\\textwidth]{deut_v11.eps} \\caption{Deuterium gradient for all galaxies in function of radius in kpc. For the MW solid line for MW-A1, dotted line for MW-A2 and dashed line for MW-B, for M31 solid grey line represents M31-A1, dotted grey line M31-A2, dashed grey line M31-B and dot-dashed grey line for M31-Bk1.25,  finally for M33 solid line is for M33-A05, dotted line for M33-01 and dashed line for M33-B. See table 2 for more details. } \\label{deut} \\end{figure}   \\subsection{M33}  The chemical evolution of the M33 disk is reproduced with a slower evolution and lower star formation threshold than in the MW and M31  The slope in the abundance gradient is well reproduced, but the oxygen abundances are overestimated by 0.25 dex. This is an indication that the chemical evolution models used for large spiral galaxies need to be adjusted to reproduce the abundances of smaller and lower massive disks.  In any case, the time evolution of the abundance gradient is also very dependent on the chosen efficiency in the SFR, as can be seen for the MW and M31.  The models fail to reproduce the present day gas profile in the inner disk, whereas it is very well reproduced for $R> 5$ kpc (same problem faced by Magrini et al. 2007), indicating a possible bulge-disk interaction in this region despite the small visible bulge of M33.  Compared to the other galaxies in this sample,  M33 presents the steeper stellar distribution along the radius, without taking into account the inner regions of the galaxy.  The SFR predicted for  M33 is in very good agreement with observations and is the one that is best reproduced by our models.  For M33, our best model is the one with the lower efficiency, constant along the galaxy radiusi (M33-A01). This fact suggests that the star formation history in small and low density disks is probably different from the more massive ones such as the MW and M31.  \\par  In conclusion we find that the present day value of the oxygen abundance is more sensitive to the threshold in the SFR than to the efficiency in the star formation and that this latter parameter plays an important role in the time evolution of the gradient. The variable efficiency in the SFR is also important to reproduce the present day gas distribution in the disk of  galaxies with a marked presence of spiral arms. A correlation between the galaxy mass and the star density profile can be seen when observing that the stellar distribution along the galactic radius gets steeper from the most massive (M31) to the lower massive one (M33). Another interesting result is the dependence of the gas distribution along the disks of spirals on the exponent of the Kennicutt law. By varying this exponent by $\\pm0.15$, which corresponds to the observational error,  one can obtain very different gas distributions.   {\\it An important conclusion of this paper is that there should be a downsizing in star formation also in spirals, similar to what applies to ellipticals. A similar conclusion was reached by Boissier et al. (2003).}"
    },
    "1004/1004.1000_arXiv.txt": {
        "abstract": "We investigate the formation of terrestrial planets in the late stage of planetary formation using two-planet model. At that time, the protostar has formed for  about 3 Myr and the gas disk has dissipated. In the model,  the perturbations from Jupiter and Saturn are considered. We also consider variations of the mass of outer planet, and the initial eccentricities and inclinations of embryos and planetesimals. Our results show that, terrestrial planets are formed in 50 Myr, and the accretion rate is about 60\\% - 80\\%. In each simulation, 3 - 4 terrestrial planets are formed inside \"Jupiter\" with masses of  $0.15 - 3.6 M_{\\oplus}$. In the 0.5 - 4AU, when the eccentricities of planetesimals are excited, planetesimals are able to accrete material from wide radial direction. The plenty of water material of the terrestrial planet in the Habitable Zone may be transferred from the farther places by this mechanism. Accretion may also happen a few times between two giant planets only if the outer planet has a moderate mass and the small terrestrial planet could survive at some resonances over time scale of $10^8$ yr. ",
        "introduction": "The discovery of the extrasolar planets \\citep{May95,Lee02,Ji03} around solar-type stars indeed provides substantial clues for the formation and origin of our own solar system. According to standard theory \\citep{Saf69,Wet90,Lis93}, it is generally believed that planet formation may experience such several stages: in the early stage, the dust grains condense to grow km-sized planetesimals; in the middle stage, Moon-to-Mars sized embryos are created by accretion of planetesimals. When the embryos grow up to a core of $\\sim 10M_\\oplus$, runaway accretion may take place. With more gases accreted onto the solid core, the embryos become more massive and eventually collapse to produce giant Jovian planets \\citep{Ida04}. At the end of the stage, it is around that the protostar has formed for about $3$ Myr, the gas disk has dissipated. A few larger bodies with low  $e$ and $i$ are in crowds of planetesimals with certain eccentricities and inclinations. In the late stage, the terrestrial embryos are excited to high eccentricity orbits by mutual gravitational perturbation. Next, the orbital crossings make planets obtain material in wider radial area. In this sense, solid residue is either scattered out of the planetary system or accreted by the massive planet, even being captured \\citep{nag00} at the resonance position of the giant planets.  \\citet{cha01} made a study of terrestrial planet formation in the late stage by numerical simulations, who set $150 - 160$ Moon-to-Mars size planetary embryos in the area of $0.3 - 2.0$ AU under mutual interactions from Jupiter and Saturn. He also examined two initial mass distributions: approximately uniform masses, and a bimodal mass distribution. The results show that $2 - 4$ planets are formed within $50$ Myr, and finally survive over $200$ Myr timescale, and the final planets usually have eccentric orbits with higher eccentricities and inclinations .  \\citet{ray04, ray06} also investigated the formation of terrestrial planets. In the simulations, they simply took into account Jupiter's gravitational perturbation, and the distribution of material are in $0.5 - 4.5$ AU. Their results confirm a leading hypothesis for the origin of Earth's water: they may come from the material in the outer area by impacts in the late stage of planet formation. \\citet{ray06b} explored the planet formation under planetary migration of the giant. In the simulations, super Hot Earth form interior to the migrating giant planet, and water-rich, Earth-size terrestrial planet are present in the Habitable Zone ($0.8 - 1.5$ AU) and can survive over $10^8$ yr timescale.  In our work, we consider two-planet model, in which Jupiter and Saturn are supposed to be already formed, with two swarms of planetesimals distributed in the region among $0.5 - 4.2$ AU and $6.2 - 9.6$ AU respectively. The initial eccentricities and inclinations of planetesimals are considered. We also vary the mass of Saturn to examine how the small bodies evolve. The simulations are performed on longer timescale $400$ Myr in order to check the stability and the dynamical structure evolution of the system. In the following, we briefly summarize our numerical setup and results . ",
        "conclusions": "We simulate the terrestrial planets formation by using two-planet model. In the simulation, the variations of the mass of outer planet, the initial eccentricities and inclinations of embryos and planetesimals are also considered. The results show that, during the terrestrial planets formation, planets can accrete material from different regions inside Jupiter. Among $0.5 - 4.2$ AU, the accretion rate of terrestrial planet is $60\\% - 80\\%$, i.e., about $20\\% - 40\\%$ initial mass is removed during the progress. The planetesimals will improve the efficiency of accretion rate for certain initial eccentricities and inclinations, and this also makes the newly-born terrestrial planets have lower orbital eccentricities. It is maybe a common phenomenon in the planet formation that the water-rich terrestrial planet is formed in the Habitable Zone. The structure, which is similar to that of solar system, may explain the results of disintegration of a terrestrial planet. Most of the planetesimals among Jupiter and Saturn are scattered out of the planetary systems, and this migration caused by scattering \\citep{fer84} or long-term orbital evolution can make planets capture at some mean motion resonance location. Accretion could also happen a few times between two planets if the outer planet has a moderate mass, and the small terrestrial planet could survive at some resonances over $10^8$ yr time scale. Structurally, Saturn has little effect on the architecture inside Jupiter, owing to its protection. However, obviously, a different Saturn mass could play a vital role of the structure outer Jupiter. Jupiter and Saturn in the solar system may form over the same period."
    },
    "1004/1004.5415_arXiv.txt": {
        "abstract": "% We present a set of low resolution empirical SED templates for AGNs and galaxies in the wavelength range from 0.03 to 30$\\mu$m. These templates form a non-negative basis of the color space of such objects and have been derived from a combination 14448 galaxies and 5347 likely AGNs in the NDWFS Bo\\\"otes field. We briefly describe how the templates are derived and discuss some applications of them. In particular, we discuss biases in commonly used AGN mid-IR color selection criteria and the expected distribution of sources in the current WISE satellite mission. ",
        "introduction": "In many current and upcoming extragalactic photometric surveys it will be necessary to use spectral energy distribution (SED) fitting techniques to characterize most objects (e.g., redshifts, K-corrections, stellar masses and bolometric luminosities), because spectroscopic observations are too expensive.  SEDs of galaxies are typically either empirically obtained \\citep[e.g.,][]{cww80} or theoretically created \\citep[e.g.,][]{bc03}. The former have the advantage of giving an accurate representation of the SEDs but generally lack the large wavelength ranges of theoretical ones. Empirical SEDs also have the advantage of easily including non-stellar emission, like active nuclei and dust/PAH emission, which are hard to model from first principles. In particular, most available AGN SEDs are empirically obtained \\citep[e.g.,][]{richards06} by averaging photometric observations of a small number of quasars.  In this proceeding we discuss a non-negative basis of empirically obtained SED templates for galaxies and AGNs that accurately represents the colors of such objects. These templates are derived from a combination of 14448 galaxies and 5347 likely AGNs in the NDWFS Bo\\\"otes field with spectroscopic redshifts and photometry spanning the far-UV to the mid-IR. All results shown in here have been presented and are discussed in detail by \\citet{assef08}, \\citet{assef10a} and \\citet{assef10b}. ",
        "conclusions": ""
    },
    "1004/1004.3996.txt": {
        "abstract": "{We present kinematic data for 211 bright planetary nebulae in eleven Local Group galaxies: M31 (137 PNe), M32 (13), M33 (33), Fornax (1), Sagittarius (3), NGC 147 (2), NGC 185 (5), NGC 205 (9), NGC 6822 (5), Leo A (1), and Sextans A (1).  The data were acquired at the Observatorio Astron\\'omico Nacional in the Sierra de San Pedro M\\'artir using the 2.1m telescope and the Manchester Echelle Spectrometer in the light of [\\ion{O}{3}]$\\lambda$5007 at a resolution of 11\\,km\\,s$^{-1}$.  A few objects were observed in H$\\alpha$.  The internal kinematics of bright planetary nebulae do not depend strongly upon the metallicity or age of their progenitor stellar populations, though small systematic differences exist.  The nebular  kinematics and H$\\beta$ luminosity require that the nebular shells be accelerated during the early evolution of their central stars.  Thus, kinematics provides an additional argument favoring similar stellar progenitors for bright planetary nebulae in all galaxies. % % }  % % \\resumen{Presentamos datos cinem\\'aticos para 211 nebulosas planetarias brillantes en once galaxias del Grupo Local: M31 (137 NPs), M32 (13), M33 (33), Fornax (1), Sagittarius (3), NGC 147 (2), NGC 185 (5), NGC 205 (9), NGC 6822 (6), Leo A (1), y Sextans A (1).  Adquirimos los datos en el Observatorio Astron\\'omico Nacional en la Sierra de San Pedro M\\'artir con el telescopio de 2.1m y el Espectr\\'ometro Manchester Echelle en la l\\'\\i nea de [\\ion{O}{3}]$\\lambda$5007 con una resoluci\\'on de 11\\,km\\,s$^{-1}$. Observamos algunos objetos en H$\\alpha$.  La cinem\\'atica de nebulosas planetarias brillantes no depende fuertemente de la metalicidad o la edad de la poblaci\\'on estelar progenitora en sus galaxias hu\\'espedes, aunque existen peque\\~nas diferencias sistem\\'aticas.  La cinem\\'atica y la luminosidad en H$\\beta$ de las c\\'ascaras nebulares requieren la aceleraci\\'on de las c\\'ascaras durante la evoluci\\'on temprana de las estrellas centrales.  As\\'\\i, la cinem\\'atica representa otro argumento en favor de estrellas progenitoras similares para las nebulosas planeatarias brillantes en todas galaxias. }   % \\addkeyword{ISM: kinematics and dynamics} \\addkeyword{Local Group} \\addkeyword{Planetary Nebulae} \\addkeyword{Stars: Evolution}   % % % \\begin{document} % ",
        "introduction": "Planetary nebulae are the immediate descendants of asymptotic giant branch (AGB) stars of low and intermediate masses ($1\\, M_{\\odot} < M< 8\\,M_{\\odot}$).  The mass lost on the AGB (or part of it) is seen as the ionized nebular shell in planetary nebulae.  The composition of these nebular shells is extremely useful in studying the nucleosynthetic production of their precursor stars.  The cycling of matter through these stars and its transformation is part of the chemical evolution of galaxies, since the progenitors of planetary nebulae are responsible for much of the helium, carbon, nitrogen, and some s-process elements in the universe.  The luminosity function of bright extragalactic planetary nebulae (PNLF) has been used extensively as a distance indicator for about two decades \\citep{jacoby1989, ciardulloetal1989}.  While some progress has been made on understanding the nature of the progenitor stars of bright planetary nebulae \\citep[e.g.,][]{richermccall2008}, the comments made by \\citet{pottasch1990} are still largely valid: \\lq\\lq It may seem rather strange to determine the distance to the galactic centre by calibrating against the much more distant galaxy M31, but it is no stranger than the idea of using PN as standard candles in the first place, since individual distances to PN are so poorly known.\"  Bright extragalactic planetary nebulae have two key advantages for evolutionary studies with respect to their Galactic counterparts.  First, their distances are known, so their absolute luminosities are studied easily.  Second, in galaxies more distant than the Magellanic Clouds, planetary nebulae are unresolved for ground-based observations, making it easy to study their integrated spectral properties, even at high spectral resolution.  The drawback, of course, is the lack (usually) of spatial resolution available for Galactic planetary nebulae that is so useful in studying physical processes.  Over the past decade, a substantial quantity of low resolution spectroscopy has been acquired \\citep[e.g.,][]{jacobyciardullo1999, walshetal1999, rothetal2004, mendezetal2005, penaetal2007, richermccall2008, magrinigoncalves2009}.  On the other hand, the only high resolution spectroscopy of extragalactic planetary nebulae suitable for studying their internal kinematics is that of \\citet{zijlstraetal2006} and \\citet{arnaboldietal2008}, apart from the pioneering efforts of \\citet{dopitaetal1985, dopitaetal1988}.  Here, we present our high resolution spectroscopic observations of 211 planetary nebulae in 11 Local Group galaxies.  For the first time, these data allow redundant comparisons across different stellar populations.  This study is part of a larger effort to understand the systematics of planetary nebula kinematics within our Milky Way and the Local Group.  \\citet{lopezetal2010} present our results for Galactic planetary nebulae.  We present our observations and the data reduction in \\S \\ref{observations}.  We explain the analysis as well as its limitations in \\S \\ref{analysis}.  The results follow in \\S\\ref{results}.  We argue that the [\\ion{O}{3}]$\\lambda 5007$ line width is an adequate description of the kinematics of the majority of the ionized gas.  We demonstrate that our radial velocities are accurate and in agreement with extant data.  We find that the average line widths for bright planetary nebulae in all galaxies are similar, though there is a trend of decreasing line width with nuclear distance in the disc of M31.  In \\S \\ref{discussion}, we consider the implications of the foregoing, with the most important being that the progenitor stars of bright planetary nebulae in all galaxies span a relatively small range in mass and that the central star plays an important role in accelerating the ionized shells of these objects.  \\S \\ref{conclusions} summarizes our conclusions.  Here, we focus on what we can learn of the evolution of bright extragalactic planetary nebulae and their progenitor stars from their [\\ion{O}{3}]$\\lambda 5007$ line widths.  We make no attempt to investigate the kinematics of individual objects nor to interpret the line widths in terms of internal kinematics, as our spatially-integrated line profiles make this very difficult. ",
        "conclusions": "We present kinematic data for 211 bright extragalactic planetary nebulae in 11 Local Group galaxies.  We present line profiles in the line of [\\ion{O}{3}]$\\lambda 5007$ for all objects and in the H$\\alpha$ line for a small minority.  At the signal-to-noise of our data, the line profiles are Gaussian, or nearly so, in all cases.  The intrinsic line widths in [\\ion{O}{3}]$\\lambda 5007$ and H$\\alpha$ are also similar.  Thus, the line widths in [\\ion{O}{3}]$\\lambda 5007$ are an adequate description of the kinematics of most of the matter in the entire ionized shell in these objects.  We find that the average line width for the bright planetary nebulae in all galaxies are similar (where we observed at least three objects).  Given our current understanding, this result implies that the progenitors of bright planetary nebulae at lower metallicity are either slightly more massive or that the higher plasma temperatures at lower metallicities produce larger acceleration of the nebular shell.  The approximate constancy of the line width also implies that there cannot be a large difference in the masses of the progenitors of bright planetary nebulae among galaxies, whether star-forming or not, though small variations are possible.  Within M31, the line width decreases with increasing distance from the nucleus, in agreement with these deductions.  As a general rule, the range of line widths is larger in dwarf galaxies than in M31 or M33.  Presumably, this arises because, at lower metallicity, either a larger range of progenitor masses contribute bright planetary nebulae or that bright planetary nebulae encompass a wider range of evolutionary states. On the other hand, there is no obvious correlation between the line width in [\\ion{O}{3}]$\\lambda 5007$ and either the oxygen abundance or the absolute [\\ion{O}{3}]$\\lambda 5007$ magnitude, $M_{5007}$.  The most significant correlation is between line width and the H$\\beta$ luminosity.  Since [\\ion{O}{3}]$\\lambda 5007$-bright planetary nebulae represent a monotonic fading sequence in the light of H$\\beta$, this correlation implies that the ionized shells of bright planetary nebulae are accelerated during their early evolution, independently of the age or metallicity of the progenitor stellar population.  Theory has long predicted this result, so our data are consistent with its general validity.  More generally, our measurements of the kinematics of bright planetary nebulae are in good accord with the results derived from photometry and low resolution spectroscopy.  All three lines of reasoning require that intrinsically bright planetary nebulae descend from progenitor stellar populations spanning a relatively small range of masses.  We thank the telescope operators during our runs, Gabriel Garc\\'\\i a, Gustavo Melgoza, Salvador Monrroy, and Felipe Montalvo, for their help.  We thank M. Arnaboldi for kindly providing the line widths for the planetary nebulae in the Virgo Cluster that are shown in Fig. \\ref{fig_line_width}.  We thank the referee for helpful comments and suggestions.  We gratefully acknowledge financial support throughout this project from CONACyT grants 37214, 43121, 49447, and 82066 and from DGAPA-UNAM grants 108406-2, 108506-2, 112103, and 116908-3.  We acknowledge the use of NASA's \\emph{SkyView} facility (http://skyview.gsfc.nasa.gov) located at NASA Goddard Space Flight Center in generating Fig. \\ref{fig_M31_pne}."
    },
    "1004/1004.2145_arXiv.txt": {
        "abstract": "We have searched the hybrid BALQSO catalogue of Scaringi et~al.  derived from data release 5 of the Sloan Digital Sky Survey in order to compile the largest sample of objects displaying spectral signatures which may be indicative of radiative line driving.  The feature in question is the ``ghost of Ly-$\\alpha$'', a line-locking feature previously identified in the broad C~{\\sc iv} and Si~{\\sc iv} absorption lines of a small fraction of BALQSOs, and formed via the interaction of Ly-$\\alpha$ photons with N~{\\sc v} ions.  We test, where possible the criteria required to produce an observable ghost feature. These criteria include: significant broad absorption, strong intrinsic Ly-$\\alpha$ emission, narrow Ly-$\\alpha$, strong N~{\\sc v} absorption, and a weak far-UV continuum. No single ghost-candidate meets all of these criteria. Furthermore, we find that these criteria are not met significantly more often in ghost-candidates than in a comparison sample chosen to exhibit relatively featureless broad absorption troughs. Indeed, the only significant differences we find between our ghost-candidate and comparison samples, is that on average, our ghost-candidate sample displays (i) significantly stronger N~{\\sc v} absorption, and (ii) the onset of absorption occurs at lower velocities in our ghost-candidate objects.  Significantly, we find no evidence for an excess of objects whose absorption troughs bracket the location of the Ly-$\\alpha$--N{\\sc v} line-locking region, rather the location of ghost-like features appears to be independent of any systematic velocity, with comparable numbers appearing both redward and blueward of the ghost-zone.  Thus, the majority of objects identified here as strong ghost-candidates are likely multi-trough interlopers whose absorption feature simply bracket the region of interest. ",
        "introduction": "Broad absorption line quasars (BALQSOs) as their name suggests, show strong broad blue-shifted absorption lines in their spectra believed to be indicative of high velocity ($\\sim$0.1c) out-flowing winds. BALQSOs represent approximately 15\\% of quasars in general \\citep{Reich03a,Trump06,Knigge,Sca09}. The majority ($\\sim$ 85\\%) \\citep{Spray,Reich03b} of BALQSOs are known as HiBALs, displaying absorption in lines of high ionisation only (e.g. N~{\\sc v}, Si~{\\sc iv}, and C~{\\sc iv}). The remainder are classified as LoBALs and show in addition broad absorption in lines of low ionisation species, most notably Al~{\\sc iii} and Mg~{\\sc ii}. LoBALs are further sub-classified according to the presence of Fe absorption, the Fe~LoBals. Since the spectra of LoBals are in general redder than HiBals, \\cite{Becker00} suggests that BALQSOs and in particular Fe~LoBALs may be an early phase in the development of emerging or re-fuelled quasars.  There has been much debate about the relationship between BALQSOs and the general quasar population as a whole. Simple unification schemes, suggest that BALQSOs and non-BALQSOs are similar objects and that any observed differences in their spectra arise due to orientation effects \\citep{Ogle99,Wey91,Sch99,Elvis}. In this picture, the relative fraction of BALQSOs to non-BALQSOs has a simple geometric interpretation, representing the fraction of sky, as seen from the source, obscured by gas (ie. the source covering fraction), which in the simplest case, can be related to the flow geometry (e.g. opening angle).  In recent years interest in BALQSOs outflows has risen sharply, principally because of the realisation that such high velocity outflows carry a substantial amount of energy and momentum into the ISM, and may therefore be important in driving AGN feedback as well as providing a mechanism for quenching star formation \\citep{Scan}.  Indeed, the discovery of highly ionised, very high-velocity X-ray outflows (e.g. Pounds et~al. 2003a, 2003b; Pounds and Reeves 2009), for which the energy transport (in terms of mechanical energy) is large enough to interrupt the growth of the host galaxy, may provide the causal link behind the well-known correlation between the mass of the central black hole and the mass of the bulge (e.g. Ferrarese and Merritt 2000; Gebhardt et~al.  2000; Tremaine et~al. 2002).  However, despite the increasing importance of these outflows the precise mechanism responsible for accelerating them to high velocity remains uncertain, with radiative acceleration \\citep{Shlos85,AravI,Murr95,Prog00,Chel01} , Magneto Hydro Dynamic (MHD) driven winds \\citep{blandford,Em92,Kon94,Bot97} and thermally driven winds seen as the main contenders (see e.g. \\cite{Prog07} for a review). A major hindrance to progress in this area is the absence of a clean discriminatory observational signature of what are essentially orthogonal wind geometries. For radio-quiet objects, AGN unification schemes tend to favour equatorial wind geometries (e.g. Elvis 2000). By contrast, observations of radio-loud BALQSOs \\citep{Becker97,Becker00,brotherton98} and in particular polarisation observations of PKS~0040-005 indicate a non-equatorial BAL outflow \\citep{brotherton06}. One possible explanation for the apparent difference in wind geometry between radio-loud and radio-quiet objects, is that different acceleration mechanisms operate in these two classes of objects.  Perhaps the strongest indicator that radiative driving is responsible for accelerating at least some of the high velocity flows, is the appearance of line-locking features in the spectra of a small fraction of BAL quasars (e.g. Turnshek et~al. 1988, Weymann et~al. 1991, Korista et~al. 1993, Arav 1996, Vilkoviskij and Irwin 2001). The most well-studied of these is the so-called ghost of Ly-$\\alpha$, a small hump seen in the absorption troughs of a small fraction (less than a few percent) of BALQSOs formed via the interaction between Ly-$\\alpha$ photons and N~{\\sc v} ions (see e.g. Korista et~al. 1993, Arav et~al.  1995, Arav 1996, and references therein).  However, the reality of ghosts remains very much an open question. Previous studies have been limited to small samples of relatively low quality spectra, with strong ghosts often only becoming apparent in composite spectra (Arav et~al. 1995, North et~al. 2006). Moreover, distinguishing between line-locking features and similar features caused by the chance alignment of multiple absorption systems is difficult with relatively small samples. Indeed, Korista et~al. 1993, showed that in a sample 72 objects, evidence for line-locking features was merely suggestive rather than convincing.  The aim of this paper is twofold: (i) to compile the largest sample of uniformly selected ghost-candidate spectra, and (ii) attempt to determine the origin of their ghost features, by testing on a source by source basis, whether they meet the original criteria as set out by Arav (1995) necessary for ghost-features to be observed.  In \\S2 we describe the mechanism proposed for ghost formation and outline Arav's criteria for the production of strong ghost signatures. Selection of those objects comprising our ghost-candidate spectra and our non-ghost control sample is described in \\S3.  In \\S4 we present the results of testing each of the spectra in our ghost-candidate and comparison samples against each of the criteria in turn necessary for the formation of an observable ghost feature. We discuss the implications of our findings in \\S5.  Our conclusions are summarised in \\S6. ",
        "conclusions": "For the 43 objects in our ghost candidate sample, 21 are at large enough redshift to allow us to test {\\em all\\/} of the criteria necessary for the formation of an observable ghost feature.  Of these 21/21 satisfy the condition for strong N~{\\sc v} absorption. 18/21 also satisfy the condition for strong intrinsic Ly-$\\alpha$ emission. However, less than half of the objects (12/21) satisfy the condition for narrow emission-lines, and even fewer (5/21) show weaker than average He~{\\sc ii} EW. While 11/21 objects satisfy the first 3 criteria, no single object meets all of the requirements for the formation of an observable ghost feature.  For the remaining 22 objects, only 3 objects pass 2 of the first 3 criteria, and of these all show stronger than average He~{\\sc ii}. Of the 5 objects which indicate weaker than average He~{\\sc ii} strengths, only one can be tested for any of the other criteria, in this case, the width of the emission-line, which it fails.   None of the objects in our ghost-candidate sample are present in the samples produced by \\cite{AravV} or \\cite{Kor93} due to differences in the redshift ranges over which they were constructed.  North et~al. (2006) identified 7 strong ghost-candidate spectra in DR3 of the SDSS.  6 of those objects are also in our sample of 43 ghost-candidates. The other, SDSSJ142050.34-002553.1. has an assigned redshift of 2.085 in DR5 compared to the 2.103 used by \\cite{Nor06} and is therefore rejected by our ghost zone cut. Only one of North et~al.'s ghost zone final cut is of sufficient redshift to test all of the criteria for the formation of an observable ghost feature SDSS J110623.52-004326.0. This object fails the narrow emission-line requirement. For the other 5 objects, 3 fail the test for narrow emission-lines, and 4 fail the test for weaker than average He~{\\sc ii} EW.  For our non-ghost comparison sample 12/25 meet the criterion for strong N~{\\sc v} absorption.  20/25 meet the requirement of strong Ly~$\\alpha$, while 12/22 meet the narrow emission-line width constraint. 9/25 objects show evidence for weaker than average He~{\\sc ii} strengths. However, only 3 objects meet 3 out of 4 criteria, and none of these meet the weaker than average He~{\\sc ii} strength. Thus, {\\em none\\/} of the objects in our comparison sample meet all of the criteria necessary for the formation of an observable ghost feature.  A comparison between our ghost-candidate and non-ghost samples suggests that the main difference between them lies in the strength of the absorption, with our ghost candidate sample displaying more objects with strong N~{\\sc v} absorption. Comparison of the geometric mean composite spectra of these two samples, and their ratio (Figure~13 dashed (ghost-candidate) and dotted (non-ghost comparison spectra) lines, indicates that the ghost-candidate composite shows significantly stronger absorption at lower velocities in all of the strong lines Ly-$\\alpha$, N~{\\sc v}, Si~{\\sc iv} and C~{\\sc iv}.   In order to test whether peaks are more common within the ghost-zone than elsewhere, we repeat the ghost selection method using two ``fake'' ghost-zones (red-ward and blue-ward of the original ghost-zone), in a similar fashion to North et~al. (2006).  These zones are created in precisely the same way as the ghost-zone except that the red zone is centred at 4000~km~s$^{-1}$ while the blue zone is centred at 8000~km~s$^{-1}$ blue-ward of the C~{\\sc iv} emission-line.  Of the 258 single peaked objects, 82 have peaks within the ``fake'' red zone and 50 have peaks within the fake ``blue-zone''. Since the original ghost-zone contained 69 objects with single peaks, the evidence for an excess of objects with peaks at a preferred velocity is weak. That is, single peaks within the ghost zone are no more likely than single peaks at other velocities.  We have also examined the link, if any, between the peaks within the C~{\\sc iv} absorption trough and N~{\\sc v} absorption. In order to select N~{\\sc v} BALs we use a modified Balnicity index (Weymann 1991).  We take the continuum to be equal to the mean flux between $\\lambda\\lambda$1315--1330\\AA, and require that the flux drops to below 90\\% of this value continuously for $>1000$~km~s$^{-1}$ between 0 and 7000~km~s$^{-1}$ blue-ward of the N~{\\sc v} emission line.  In order to perform this test we require objects with redshifts in excess of 2.15.  From a sample of 1747 objects with z$>$2.15, we find 1258 N~{\\sc v} BALs and 489 N~{\\sc v} non-BALs.  27\\% (340) of the N~{\\sc v} BALs show multiple troughs in their C~{\\sc iv} absorption compared to only 15.5\\% (76) of the N~{\\sc v} non-BALs.  Similarly, 7.6\\% (95) of the N~{\\sc v} BALs and only 3.1\\% (15) of the N~{\\sc v} non-BALs show a single peak in the C~{\\sc iv} absorption trough. Among the N~{\\sc v} non-BALs we find : i) two objects with single peaks within the ghost-zone, ii) two objects with a single peak in the fake blue zone, and iii) no objects with a single peak in the fake red zone.  For the N~{\\sc v} BALs, we find : i) 26 objects with single peaks within the ghost-zone, ii) 14 with single peaks within the fake blue zone, and iii) 33 with single peaks within the fake red zone. While these results indicate a strong link between the presence of N~{\\sc v} absorption and the mechanism responsible for producing features within the C~{\\sc iv} absorption trough, line-locking between Ly-$\\alpha$ and NV does not appear to be the dominant mechanism.  In summary, N~{\\sc v} BALs are more likely to have multiple troughs within the C~{\\sc iv} absorption than N~{\\sc v} non-BALs (factor of 2). Further, N~{\\sc v} BALs are also more likely to display single-peaks within their C~{\\sc iv} absorption than N~{\\sc v} non-BALs. Approximately 25\\% of the single peaked objects are located within the ghost-zone.  However, similar numbers are found in both the blue and red fake ghost zones. Thus while strong N~{\\sc v} absorption appears to be a strong requirement for the appearance of features within the C~{\\sc iv} absorption trough of BALQSOs, there is no preferred velocity for the location of these features.  \\begin{figure} \\centering \\includegraphics[width=0.5\\textwidth]{fig10} \\caption{Composite spectrum made up of all our ghost candidates fit with a reddened DR5 composite(dotted) and a BALQSO composite (dashed).} \\label{fig:FC} \\end{figure}  \\onecolumn \\begin{figure} \\centering \\includegraphics[width=0.95\\textwidth]{fig11} \\caption{Spectra of ghost candidates from the best-candidate sample for which all criteria can be tested.} \\label{fig:mike_1b} \\end{figure} \\twocolumn  \\onecolumn \\begin{figure} \\centering \\includegraphics[width=0.95\\textwidth]{fig12} \\caption{Spectra of ghost candidates from the best-candidate sample for which all criteria can be tested.} \\label{fig:mike_2b} \\end{figure} \\twocolumn  \\begin{figure} \\centering \\includegraphics[width=0.5\\textwidth]{fig13} \\caption{Upper-panel - composite ghost-candidate (solid line) and comparison (dashed line) spectra. Lower panel - ratio of ghost-candidate to comparison spectra.} \\label{fig:plot_cand} \\end{figure}"
    },
    "1004/1004.0971_arXiv.txt": {
        "abstract": "Extrasolar planet host stars have been found to be enriched in key planet-building elements. These enrichments have the potential to drastically alter the composition of material available for terrestrial planet formation. Here we report on the combination of dynamical models of late-stage terrestrial planet formation within known extrasolar planetary systems with chemical equilibrium models of the composition of solid material within the disk. This allows us to determine the bulk elemental composition of simulated extrasolar terrestrial planets. A wide variety of resulting planetary compositions are found, ranging from those that are essentially \"Earth-like\", containing metallic Fe and Mg-silicates, to those that are dominated by graphite and SiC. This shows that a diverse range of terrestrial planets may exist within extrasolar planetary systems. ",
        "introduction": "Extrasolar terrestrial planets are a tantalizing prospect. Given that the number of planets in the galaxy is expected to correlate inversely with planetary mass, it is expected that Earth-sized terrestrial planets are much more common than giant planets \\citep{marcy2000}. Although still undetectable by current exoplanet searches, the possibility of their existence in extrasolar planetary systems has been examined by several authors. Many such studies have focussed on the long term dynamical stability of regions within the planetary system where such planets could exist for geologic timescales \\citep{br1,br2,asg}. Several systems have been found to posses such regions (e.g. \\citealt{br1}), indicating that if they are able to form, terrestrial planets may still be present within extrasolar planetary systems. Analyses of this nature are of great interest to future planet search missions as they assist in constraining future planet search targets. However, they provide little insight into the formation mechanism and physical and chemical properties of such planets and do not necessarily indicate the presence of a terrestrial planetary companion.  A few other studies have gone one step further and undertaken detailed N-body simulations of terrestrial planet formation. \\cite{ray05} considered terrestrial planet formation in a series of hypothetical `hot Jupiter' simulations and found that terrestrial planets can indeed form in such systems (beyond the orbit of the giant planet) provided the `hot Jupiter' is located within 0.5AU from the host star. Furthermore, such planets may even have water contents comparable to that of the Earth. Terrestrial planets have been found to form even in simulations of systems which have undergone large-scale migration of the giant planet \\citep{raymond:2006,avi}. Terrestrial planets were found to form both exterior and interior to the giant planet after migration has occurred and many were located within the habitable zone of the host star. As many extrasolar planets are believed to have experienced such a migration, it is encouraging that terrestrial planets may still be able to form within these systems. To date, only \\cite{br3} have undertaken terrestrial planet formation simulations for specific planetary systems. They considered four known planetary systems and found that terrestrial planets could form in one of the systems (55Cancri). Small bodies comparable to asteroid sized objects would be stable in another (HD38529).  An even more intriguing question beyond whether or not terrestrial planets could exist within these systems is their potential chemical composition. Extrasolar planetary host stars are already known to be chemically unusual \\citep{g1,g2,butler00,g4,g5,g3,sb,s1,ism04,g6,sm,re,fv,bond1,bond:2008}, displaying systematic enrichments in Fe and smaller, less statistically significant enrichments in other species such as C, Si, Mg and Al \\citep{g3,g6,sb,bo,fv,be,bond1}. Given that these enrichments are likely primordial in origin \\citep{s1,s2,s04,s05,fv,bond1}, it is thus likely that the planet forming material within these systems will be similarly enriched. Hints of such a correlation between transiting giant planets and stellar metallicity have been observed \\citep{metal1,metal2}. Consequently, it is likely that terrestrial extrasolar planets may have compositions reflecting the enrichments observed in the host stars. Furthermore, several known host stars have been found to have C/O values above 0.8 \\citep{bond:2008}. Systems with high C/O ratios will contain large amounts of C phases (such as SiC, TiC and graphite), resulting in any terrestrial planets within these systems being enriched in C and potentially having compositions and mineralogies unlike any body observed within our Solar System.  Despite the likely chemical peculiarities of extrasolar planetary systems and the early successes of extrasolar terrestrial planet formation simulations, no studies of extrasolar terrestrial planet formation completed to date have considered both the dynamics of formation and the detailed chemical compositions of the final terrestrial planets produced. This study addresses this issue by simulating late-stage in-situ terrestrial planet formation within ten extrasolar planetary systems while simultaneously determining the bulk elemental compositions of the planets produced. This is the first such study to consider both the dynamical and chemical nature of potential extrasolar terrestrial planets and it represents a significant step towards understanding the diversity of potential extrasolar terrestrial planets. ",
        "conclusions": "}  \\subsection{Mass Distribution\\label{mass-dist}} Radial midplane mass distributions based on the equilibrium condensation sequence were calculated for each system. As composition is correlated to a specific radial distance within the midplane (via the \\cite{hersant} model), the total mass of solid material present interior to a given radial distance within the disk is given by summing the masses contained within each annulus of the disk for which composition is calculated:  \\begin{equation} {\\rm Mass\\:of\\:solid\\:material} = \\rm {\\Sigma_{i=0.1AU}^{i=5.0AU}}2{\\pi}r_{i}drM_{solid,\\:i} \\end{equation}\\\\  \\noindent where M$_{\\rm solid,\\:i}$ is the mass of solid material determined by the chemical model to be located in an annulus of width dr at r$_{\\rm i}$. The mass of solid material possible is determined by using the minimum mass solar nebula with a gas surface mass density profile varying as r$^{-3/2}$, normalized to 1700 gcm$^{-2}$ at 1 AU. Note that this approach only considers equilibrium driven condensation and neglects other processes that may migrate material and alter the mass distribution.  Based on this calculation, the most carbon-rich systems simulated have significant differences in their mass distributions compared to other systems. The combination of a broad zone of refractory carbon-bearing solids in the inner regions and the relatively small amount of water ice that condenses in the outer regions of these systems suggests that C-rich systems have significantly more solid mass located in the inner regions of the disk than for a Solar-composition disk. This can be seen in Figure \\ref{massdist} which shows the distribution of solid mass within each system normalized to a solar composition disk for disk conditions at t = 5$\\times$10$^{5}$ years.  From Figure \\ref{massdist}, the system with the highest C/O value (HD4203) clearly contains significantly more mass in the innermost regions of the disk than a disk of solar composition does. However, the total mass of solid material produced by the current approximations is only 18.4 M$_{\\bigoplus}$. As such, it is not clear that a giant planet core composed of refractory C-rich species could be produced within several AU of the host star, allowing for giant planet formation to occur much closer to the host star than previously thought. If significantly more mass were present within the disk than is currently estimated, such a scenario may be possible and would obviously alter the extent and nature of planetary migration required within these systems as it would no longer be required that a planet located at 1-2AU originally formed at 5AU and migrated inwards. Alternatively, if insufficient mass is available for Jovian core formation, production of large terrestrial planets in this region may proceed faster and with greater ease, thus increasing the chance of forming detectable terrestrial planets.  It can be seen from Figure \\ref{massdist} that the planetary system with the most Solar-like composition (HD 72659) has a mass distribution similar to that of the Solar System. The enrichment observed over the solar mass distribution within $\\sim$0.5AU from the host star is due to the higher Mg/Si value for HD72659 resulting in the condensation of more Mg silicates. This enhancement is more obvious for the system with the highest Mg/Si value (HD177830). Likewise, the refractory rich system HD27742 also contains more mass in the inner regions of the disk than for a Solar abundance. This is due to the high abundances of several refractory elements ([Al/H]=0.53, [Na/H]=0.41, [Ca/H]=0.12). The variation in mass observed for all systems at $\\sim$4.8AU is due to the condensation of hydrous species in various amounts relative to the solar composition disk.  These results imply that the use of an alternative initial mass distribution may be required for planetary systems with high C/O values. Although the planetary formation simulations presented here utilized the classical Solar-System based mass distribution, such an approach is still valid for the illustrative purposes of the current study. The full implications of these results need to be examined in future work by using alternative mass distributions for extrasolar planetary formation simulations for both gas giant planets and smaller terrestrial planets.  \\subsection{Timing of Formation\\label{EGPvary}} Specific planetary compositions have been found to be highly dependent on the time selected for the disk conditions, especially for those systems containing C-dominated regions. This is primarily due to the low degree of radial mixing encountered within the simulations. As a result, as conditions within the area immediately surrounding the planet evolve, the composition of the solid material and thus the final planet itself drastically change. For disk conditions at later times,the simulated planetary compositions evolve to more closely resemble those of the Solar System. They become dominated by Mg silicate species and metallic Fe. Terrestrial planets in Solar-like systems attain more hydrous material. Variations in composition with disk condition times are most noteworthy for those planets dominated by refractory compositions (such as the inner planets of HD72659 and HD27442). Under later disk conditions, these planets experience a complete shift in their composition, losing the majority of their refractory inventory to be composed primarily of Mg-silicates (olivine and pyroxene). Therefore if solid condensation and initial planetesimal formation occurred significantly later, we would expect to observe predominantly Mg-silicate and metallic Fe planets with enrichments in other elements (such as C) depending on the exact composition of the system. Although disk conditions at 5$\\times$10$^{5}$ years provide the ``best fit'' for Solar System simulations (based on fitting the compositions of Venus, Earth and Mars) and are thus utilized here, it remains to be seen whether or not disk conditions at this time provide an accurate description of the conditions under which planetesimals and embryos formed in other planetary systems. Therefore, we require a more detailed understanding of the timing of condensation and planetesimal and embryo formation within protoplanetary disks to be able to further constrain the predicted elemental abundances.  Similarly, as the disk evolves, the various condensation fronts migrate closer to the host star. For example, the water ice line for Gl777 migrates from 7.29 AU for midplane conditions at 2.5$\\times$10$^{5}$ years to 1.48AU for midplane conditions at 3$\\times$10$^{6}$ years. Similar degrees of migration also occur for the condensation fronts of other species (such as the Mg silicates olivine and pyroxene). In effect, the change in location of the condensation fronts alters the mass distribution within the disk, increasing the mass present in both the very closest regions of the disk ($<$ 1AU) as the refractory species are replaced by the more abundant Mg silicates and in the outer most regions ($>$ 3AU) as water ices appear. The full effects of this change will require formation simulations to be run with alternative mass distributions but it is thought that such conditions will increase the efficiency of forming close-in terrestrial planets and/or the mass of the resulting planets. Additionally, it will also allow for efficient terrestrial planet growth in the outer regions, possibly to the extent of forming gas giant cores. Given that the average disk lifetime is $\\sim$3Myr, not all disks will reach the final chemical compositions modeled here. Thus it remains to be seen not only whether or not sufficient solid mass would be retained during the evolutionary process for Jovian cores to develop but also whether or not core formation can occur before the disk is cleared out.  \\subsection{Detection of Terrestrial Planets} The results of this study are of great importance for the design of terrestrial planet finding surveys. Our simulations, while preliminary, suggest that terrestrial planets can be stable in a wide range of extrasolar planetary systems. Four distinct classes of planetary composition have been produced by the current simulations: Earth-like, Mg-rich Earth-like, refractory (compositions similar to CAI's) and C-rich. These planetary types are primarily a result of the compositional variations of the host stars and thus the system as a whole. Based on their observed photospheric elemental abundances, the majority of known extrasolar planetary systems are expected to produce terrestrial planets with compositions similar to those within our own Solar System. Therefore, systems with elemental abundances and ratios similar to these (e.g. HD72659) are ideal places to focus future ``Earth-like'' planet searches.  Based on our dynamical simulations, the masses of the terrestrial planets produced are too low to be detected by current radial velocity surveys. However, many of the simulated planets are in orbits (assuming the simulated inclinations are correct) that would place them within the prime target space for detection by the Kepler mission. Designed to detect extrasolar planets via transit studies, Kepler is the first mission that has the potential of detecting Earth-mass (and lower) extrasolar planets located within the habitable zone of a planetary system. It has the sensitivity to detect the transit of an Earth mass body within 2AU from the host star and a Mars mass body (0.1Me) within 0.4AU. Single transit events may not be sufficient to positively identify the presence of a planet, although planets at ~1 AU should transit ~3-4 times during the lifetime of the kepler mission, thus reducing the likelihood of contamination in the data. The vast majority of the terrestrial planets formed here (with the exception of the lowest mass, highest semimajor axis planets) are well within this range and thus may be detectable if they are indeed present within these systems (assuming the system is aligned such that transits can be observed from orbit). Only HD4203 produces no potentially detectable planets, based on their predicted masses. Thus it is likely that we will have an independent check of extrasolar terrestrial planet formation simulations within the next 5 years (but not necessarily for these systems, only for a range of systems which may be similar to these). Such information will be vital for further refinement of planetary formation models for both giant and terrestrial planets. Obtaining compositional information about such terrestrial planets, however, will be more difficult as the size and location of the predicted planets will prohibit direct spectroscopic studies. It is also unlikely that the terrestrial planets will contain atmospheres large enough to be detected and characterized by transit surveys. As such, detailed extrasolar terrestrial planetary chemical compositions will remain unknown for the foreseeable future.  In addition to detection via transit surveys, attempts are also being made to obtain direct images of extrasolar planetary systems. One such example is Darwin, a proposed ESA space based mission that would utilize nulling interferometry in the infrared to directly search for terrestrial extrasolar planets. The compositional variations outlined here are likely to influence our ability to successfully detect these planets. Carbon-rich asteroids are known to be highly non-reflective. For example, 624 Hektor (D-type asteroid) has a geometric albedo of 0.025 while 10 Hygiea (C-type asteroid) has a geometric albedo of 0.0717 (compared to a geometric albedo of 0.367 for the Earth and 0.113 for the moon). As both of these asteroids are assumed to be carbon-rich, it is likely that the carbon-rich planets identified here are similarly dark. Thus it is expected that searches for these planets in the visible spectrum will be difficult due to the small amount of light reflected by these bodies. However, a lower albedo results in greater thermal emission from a body, suggesting that the infrared signature from these planets would extend to shorter wavelengths than corresponding silicate planets. As a result, infrared searches (such as that of Darwin) are ideally suited to detect carbon-rich terrestrial planets and thus should be focused on stellar systems with compositions similar to that of the C-rich stars identified here to maximize results. Of course, this conclusion neglects any possible effects of a planetary atmosphere.   \\subsection{Hydrous Species} As one would intuitively expect, hydrous material (water and serpentine in the current simulations) is primarily located in the outer, colder regions of the disks. This corresponds to beyond $\\sim$7.3AU for disk conditions at 2.5$\\times$10$^{5}$ years and beyond $\\sim$1.4AU for disk conditions at 3$\\times$10$^{6}$ years for all compositions examined. As a result of this distribution, terrestrial planets forming in the inner regions of a given planetary system are unlikely to directly accrete significant amounts of hydrated material. In the current in-situ simulations, none of the simulated terrestrial planets directly accrete any hydrous species for disk conditions at 5$\\times$10$^{5}$ years. As the composition of the planetesimals is dictated by the thermodynamic conditions of the disk at the time of condensation only, if any of the simulated planetesimals were to condense/form at later times, they would be more likely to be water-rich given that the `snow line'\\footnote{Note that here `snow line' refers to the location within the disk where the thermodynamic conditions are such that water ice condenses (i.e. where T = 150K).} migrates inwards as the disk cools, producing a greater overlap between their feeding zones and the water-rich region of the disk.  However, a far greater effect is expected to be produced by migration of giant planets within the system \\citep{raymond:2006,avi}. Migration of a giant planet has been shown to be capable of driving a large amount of material from the outer regions of a disk inwards. In the case of hydrous material, this has been found to result in water-rich terrestrial planets being formed both interior and exterior to the giant planet \\citep{raymond:2006,avi}. The full extent of this increase in water content will be examined by a suite of simulations incorporating giant migration that are currently running and will be the focus of future work.  In the Solar System, it has been hypothesized that the Earth's water could have originated from hydrated material in the region where the asteroid belt now lies, or from comets beyond the orbit of Jupiter.  While belts of debris resembling the asteroid belt could exist interior to the giant planets in several of the extrasolar systems and may be incorporated into the terrestrial planets, as noted above, none of that material is expected to be hydrated at early times in the disk. We can not address the issue of cometary delivery of water in these simulations, as bodies beyond the orbits of the giant planets are not presently included in our dynamical models.  Water can also potentially be incorporated into a planetary body via adsorption onto solid grains within the disk \\citep{drake:2005}. While this process has not been considered in our current simulations, it is possible that there will be some water delivered during the formation process to the terrestrial planets produced in the Earth-like systems (HD27442, HD72659 and HD213240) as the solid grains are bathed in water vapor over the entire span of the disk. This same process will likely not be as effective at delivering water to the C-rich systems (55Cnc, Gl777, HD4203, HD17051, HD19994, HD108874 and HD177830) as they only have water vapor present at temperatures below $\\sim$800 K. This temperature range corresponds to beyond a radial distance of $\\sim$1.2AU for \\cite{hersant} midplane conditions at 2.5$\\times$10$^{5}$ years and $\\sim$0.2AU for midplane conditions at 3$\\times$10$^{6}$ years. As few terrestrial planets in our simulations accrete material from beyond 1.2AU, it is expected that C-rich planets forming early in the lifetime of the disk will remain dry without additional water being incorporated into the planets via adsorption. Thus it appears that terrestrial planets are likely to obtain some amount of water (through giant planet migration mixing the disk, variations in composition with time i.e. heterogeneous accretion, adsorption and exogenous delivery), while those within Solar-like systems may receive more water and other hydrous species than terrestrial planets within C-rich systems.  \\subsection{Planetary Interiors and Processes} Given the wide variety of predicted planetary compositions, a similarly diverse range of planetary interiors structures is also expected. To better quantify this, we examined three specific cases: a 1.03M$_{\\bigoplus}$ Earth-like planet located at 0.95AU (HD72659), a 1.22M$_{\\bigoplus}$ Mg-rich Earth-like planet located at 0.43AU (HD177830) and a 0.47M$_{\\bigoplus}$ C-rich planet located at 0.38AU (HD108874). Approximate interior structures for each were calculated using equilibrium mineralogy for a global magma ocean with P = 20GPa and T = 2000$^{\\circ}$C. Equilibrium compositions at these conditions have been found to produce the best agreement between predicted and observed siderophile abundances within the primitive upper mantle of the Earth \\citep{drake}. Elemental abundances were taken from the results discussed in Section \\ref{chem}. Resulting mineral assemblages were sorted by density to define the compositional layers. Note that this assumes that a planet undergoes complete melting and differentiation. Approximate planetary radii were obtained from \\cite{sotin} based on planetary mass. These planetary radii are based on silicate planetary equations of state and as such are unlikely to completely describe the C-rich planets modeled here. However, no studies have considered such assemblages, forcing us to assume a silicate based approximate radius. Density variations at high pressures were not considered in defining the depths of various layers. Although important for planetary evolution processes such as mantle stripping, the effects of large impacts (such as the moon forming impact) are also neglected. The resulting interior structures (shown to scale) can be seen in Figure \\ref{interior}.  The interior mineralogy and structure of one of our model planets orbiting HD72659 is similar to Earth. It contains a pyroxene and feldspathic dominated crust ($\\sim$133 km deep) overlying an olivine mantle ($\\sim$985 km deep) with an Fe-Ni-S core (radius $\\sim$ 4930 km). The crust is thicker than seen on Earth as we are currently neglecting density and phase changes. Given its structure and comparable mineralogy, we would expect to observe planetary processes similar to those seen on Earth. The planets location within the habitable zone of the host star suggests that a liquid water ocean is feasible, provided sufficient hydrous material can be delivered. Melting conditions and magma compositions are expected to be comparable and it is feasible that a liquid core would develop, resulting in the production of a magnetic dynamo. In general, based on their mass and composition, the terrestrial planets of HD72659 are likely to have structures and mineral assemblages similar to those observed in our system.  The simulated planet for HD177830 (the system with the highest Mg/Si value) is depleted in Si, relative to the Earth, resulting in high spinel and olivine content in the mantle (resembling that of type I kimberlites) and a thicker mellite and calcium dominated crust than found for HD72659 ($\\sim$309 km deep). The core would produce a considerable amount of heat via potential energy release during differentiation, potentially producing melts with compositions similar to komatiite (dominated by olivine with trace amounts of pyroxene and plagioclase). Volcanic eruptions would be comparable to basaltic flows observed on Earth due to the low silica content of the melt. However, given the thickness of the crust, extrusive volcanism and plate tectonics are unlikely to occur as high stress levels would be required to fracture the crust. Producing and sustaining such stresses would be challenging. Therefore, it is questionable whether or not a planet with this composition and structure would be tectonically active for long periods of time. Given the similar composition and size of the core compared to Earth, a magnetic dynamo is still expected to be produced within the core.  Finally, carbide planets are expected to form around HD108874. The resulting composition and structure is unlike any known planet. Its small size, refractory composition and possible lack of radioactive elements (due to the potential absence of phosphate species, common hosts for U and Th, and possible lack of feldspar and carbonates, the common host of K) will inhibit long-term geologic activity due to the difficulty of melting the mantle. Only large amounts of heat due to core formation and/or tidal heating would be able to provide the required mantle heating. Once all the primordial heat has been removed, it is unlikely that the mantle would remain molten on geologic timescales. Until that time, given the buoyancy of molten carbon, volcanic eruptions would be expected to be highly enriched in C. The core is also expected to be molten, thus making it likely that a magnetic dynamo would be produced \\citep{gaidos}. Note that this assumes that sufficient heat is initially available to melt the body and allow for differentiation and core formation to occur in the first place. In essence, although initially molten and probably active, old carbide planets of this type would be geologically dead.  Incomplete mixing of material accreted at later times is likely to result in deviations from the equilibrium picture presented here. For example, accretion of oxidized and water-rich material late in the formation process may result in a stratified redox state and water-rich crust as observed for the Earth. Unfortunately, it is not possible to determine these effects with current models as it requires a level of understanding of the impact and accretion process (e.g. mantle mixing, fragmentation) on small planetary bodies that we currently do not have.  These results are also key for super Earth studies such as that of \\cite{plate1} and \\cite{plate2}. Previous simulations have assumed Earth-based compositions and structures. Based on the present simulations, a wide variety of both are possible and will need to be considered.  \\subsection{Planet Habitability\\label{habit}} The habitable zone of a planetary system is defined as being the range of orbital radii for which water may be present on the surface of a planet. For the stars considered here, that corresponds to radii from $\\sim$0.7AU to $\\sim$1.45AU. The vast majority of the planets produced by the current simulations orbit interior to this region (exterior in the case of 55Cnc) and thus are unlikely to be habitable in the classical sense. 10 planets are produced within the classical habitable zone, existing in orbits extending from 0.70AU to 1.19AU. Seven on these planets are formed in Solar-like systems (HD24772 and HD72659) and have compositions loosely comparable to that of Earth. As such, we feel that these systems (and others similar to them) are the ideal place to focus future astrobiological searches as they may not only contain planets with compositions similar to that of Earth but also exist in the biologically favorable region of the planetary system.  Of the seven C-rich systems, only two produced planets within the habitable zone. Gl777 formed two terrestrial planets within the habitable zone while HD19994 formed a terrestrial planet at 0.70AU, just at the inner edge of the habitable zone. All other planets are located well inside the required radii. Both of the habitable planets around Gl777 are C-enriched Earth-like planets, making them potential sites for the development of life. The single habitable planet produced by the HD19994 simulations is dominated by C, along with O, Fe, Si and Mg. It is unclear whether such a composition would be favorable to life. Additionally, the low planet mass (0.06M$_{\\bigoplus}$) further makes it unlikely that this particular simulated planet could ever actually host life. As such, under the current definition of habitable, we conclude that of the seven C-rich systems currently simulated, Gl777 has the best chance of supporting life, but this is by no means guaranteed.  \\subsection{Biologically Important Elements} In addition to water, complex life (as we know it) also requires several key elements to exist. The six essential elements are H, C, N, O, P and S. As was the case for the Solar System simulations discussed in \\cite{bond:2009} none of the planets accreted any N and are also lacking in H. The terrestrial planets formed in the Solar-like systems contained various amounts of O and P but some were deficient in S and all were laking C, as for the Solar System simulations. The most C-rich systems, on the other hand, were lacking in O, P and S.  Thus it is clear that for life as we know it to develop on any of the terrestrial planets formed in the current simulations, significant amounts of several elements must be supplied from exogenous sources within the system. All elements may be supplied from the outer, cooler regions of the disk. Thus it is possible that migration or the radial mixing of cometary-type material into the terrestrial planet region may produce planets with the necessary elements for life to develop. As for the Solar System simulations, all biologically required elements would be introduced in a form that could potentially be utilized by early life. This is especially intriguing for those planets located within the habitable zone. On the other hand, alternative pathways could potentially develop for the formation of an alternative biologic cycle without requiring the same six elements.  \\subsection{Host Star Enrichment} As previously discussed, stellar photospheric pollution has been suggested as a possible explanation for the observed high metallicity of extrasolar planetary host stars \\citep{la,g6,murray}. The current simulations, though, do not support this hypothesis. Enrichments are produced primarily in Al, Ca and Ti, not Fe as is required by the pollution theory. Furthermore, relatively small masses of solid material are accreted by the host stars during planet formation, suggesting that insufficient material is accreted to produce the observed enrichments. Thus unless migration of the giant planets can systematically result in accretion of giant planets by the host star, our results agree with with previous authors (e.g \\citealt{s1,s2,s04,s05,fv}) in finding that the observed host star enrichment is primordial in origin.  Our simulations also imply that enrichments due to stellar pollution are most likely to be observed for the refractory elements in high mass stars with low convective zone masses. This suggests that surveys for pollution effects caused by terrestrial planet formation should focus on Ti, Al and Ca abundances in A-type and high mass F-type stars as they are expected to have the lowest convective zone masses. However, more detailed simulations of the fate of material accreted into radiative zones need to be undertaken to support this hypothesis."
    },
    "1004/1004.0695_arXiv.txt": {
        "abstract": "We update cosmological hot dark matter constraints on neutrinos and hadronic axions. Our most restrictive limits use 7-year data from the Wilkinson Microwave Anisotropy Probe for the cosmic microwave background anisotropies, the halo power spectrum (HPS) from the 7th data release of the Sloan Digital Sky Survey, and the Hubble constant from Hubble Space Telescope observations. We find 95\\% CL upper limits of $\\sum m_\\nu<0.44$~eV (no axions), $m_a<0.91$~eV (assuming $\\sum m_\\nu=0$), and $\\sum m_\\nu<0.41$~eV and $m_a<0.72$~eV for two hot dark matter components after marginalising over the respective other mass. CMB data alone yield $\\sum m_\\nu<1.19$~eV (no axions), while for axions the HPS is crucial for deriving $m_a$ constraints. This difference can be traced to the fact that for a given hot dark matter fraction axions are much more massive than neutrinos. ",
        "introduction": "\\label{sec:introduction}  Cosmological large-scale structure data allow for precise estimates for the parameters of minimal or extended cosmological models. These results have ramifications far beyond cosmology itself, notably in the area of neutrino physics. The well-known hot dark matter constraints provide neutrino mass limits that directly impact neutrino mass searches in single~\\cite{Drexlin:2008zz} and double~\\cite{Schonert:2010zz} beta decay experiments. In a series of papers by our collaboration~\\cite{Hannestad:2003ye,Hannestad:2005df,Hannestad:2007dd,Hannestad:2008js} and another group~\\cite{Melchiorri:2007cd} this approach was extended to hadronic axions where the resulting mass limits are complementary to solar axion searches by the CAST experiment~\\cite{Zioutas:2004hi,Andriamonje:2007ew,Arik:2008mq} and the Tokyo axion helioscope~\\cite{Moriyama:1998kd,Inoue:2002qy,Inoue:2008zp}. In the present work, we use the 7-year data release from the Wilkinson Microwave Anisotropy Probe (WMAP) as an opportunity to update these results, and also modify along the way several other input assumptions as detailed in the main text below.  Within standard cosmological assumptions, the neutrino plus antineutrino number density today, summed over all flavours, is $n_\\nu\\sim336~{\\rm cm}^{-3}$. Currently available cosmological data are not yet sensitive enough to resolve the small mass differences measured in oscillation experiments, so all neutrinos are treated as having the same mass $m_\\nu$, traditionally expressed by the parameter $\\sum m_\\nu=3 m_\\nu$. For axions, on the other hand, the freeze-out temperature and therefore the number density $n_a$ depends on the axion's interaction rate with pions and nucleons via \\begin{eqnarray} \\label{eq:axionthermalisation} a+\\pi &\\leftrightarrow& \\pi +\\pi, \\nonumber \\\\ a+N & \\leftrightarrow& N + \\pi, \\end{eqnarray} where the coupling strength is, in turn, proportional to the axion mass $m_a$ \\cite{Berezhiani:1992rk,Chang:1993gm,Hannestad:2005df}. Figure~\\ref{fig:na} shows the relation between $m_a$ and $n_a$ computed for the thermalisation processes~(\\ref{eq:axionthermalisation}) based on the original calculations of reference~\\cite{Hannestad:2005df}. For small $m_a$, the number density is also small, so assuming $m_a=0$ implies $n_a=0$ which brings us back to standard cosmology. Near the hot dark matter limit of $m_a\\sim 1$~eV one finds a present-day number density of $n_a\\sim 50~{\\rm cm}^{-3}$. Therefore, in the relevant mass range, neutrinos are about 6 times more numerous than axions and one expects a hot dark matter limit on $m_a$ that is roughly twice that on $\\sum m_\\nu$, in agreement with what we find when we use our full range of input data sets.   \\begin{figure}[t] \\hspace{25mm} \\includegraphics[width=12.5cm]{fig1.eps} \\caption{Axion number density $n_a$ as a function of the axion mass $m_a$ assuming the hadronic thermalisation processes~(\\ref{eq:axionthermalisation}) based on the calculations of reference~\\cite{Hannestad:2005df}.\\label{fig:na}} \\end{figure}   In detail, however, the situation is more subtle. Conventional wisdom says that hot dark matter constraints arise primarily from the shape of the measured matter power spectrum, since hot dark matter free-streaming suppresses the growth of matter perturbations and hence the clustering power on small scales. However, recent cosmological data have become so precise that one finds a useful limit on $\\sum m_\\nu$ of order 1~eV already from the cosmic microwave background (CMB) anisotropies alone, notably from the increased amplitude of the first acoustic peak in the temperature auto-correlation spectrum due to the early integrated Sachs--Wolfe (ISW) effect~\\cite{Ichikawa:2004zi}. The same is however not true for axion hot dark matter, since for the same hot dark matter fraction, the axion is necessarily some 6 times heavier than the equivalent neutrino.  Thus while neutrinos with masses near the hot dark matter limit ($m_\\nu \\lwig 0.3~{\\rm eV}$) essentially act like radiation at CMB decoupling and contribute strongly to the early ISW effect, the equivalent axion is already nonrelativistic and thus indistinguishable from cold dark matter as far as the CMB is concerned.  In the latter case, one needs to use the shape of the matter power spectrum from smaller-scale data in order to put any sensible constraint on $m_a$.  In order to derive new hot dark matter limits on neutrinos and axions and to explain their differences, we begin in section~\\ref{sec:model} with a description of our cosmological model and in section~\\ref{sec:data} of the data sets used. In section~\\ref{sec:results} we use standard Bayesian techniques to derive credible intervals for $\\sum m_\\nu$ and $m_a$ separately and for a two-component case based on different combinations of data sets. We discuss and summarise our findings in section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have provided an updated constraint on axion hot dark matter using new cosmological data, most notably CMB data from the WMAP 7-year data release and the final SDSS-DR7 LRG data. We have also pointed out a qualitative difference between neutrino and axion hot dark matter, in that axion masses in the detectable range are large enough that axions are nonrelativistic at recombination. This means they act almost like cold dark matter as far as CMB is concerned and that current CMB data in itself does not provide a useful limit on the axion mass, even when priors from HST or BAO are imposed. In order to properly constrain $m_a$ it is necessary to include information on the shape of the matter power spectrum from, for example, the SDSS halo power spectrum.  An interesting question is whether future CMB data will be sensitive to axion hot dark matter. The Planck mission has an estimated sensitivity of \\hbox{$\\sigma(\\sum m_\\nu) \\sim 0.3$--0.5~eV} \\cite{DeBernardis:2009di,Lesgourgues:2006nd}.  But from figure~\\ref{fig:cls} we can already say that even at high multipoles the effect of axion hot dark matter on the CMB anisotropy spectrum is quite small---generally smaller than the uncertainty due to cosmic variance. This leads to the inevitable conclusion that CMB anisotropy observations will remain poor probes of hot dark matter in the several eV range.  Cosmological bounds on neutrino and axion masses are nicely complementary to experimental searches, but cannot replace them. One caveat concerning our axion results is that we need to assume that the predicted thermal population was actually produced after the QCD epoch. In non-standard cosmologies with low reheating temperature the axion population can be severely suppressed and our bounds would not apply~\\cite{Grin:2007yg}.  In other scenarios a significant cosmic background of low-mass axions can be produced that remain relativistic until today and do not form hot dark matter~\\cite{Chun:2000jr}. For neutrinos such caveats are less relevant because their thermalisation epoch is well probed by big-bang nucleosynthesis and the presence of radiation with roughly the right abundance has been confirmed by precision cosmology."
    },
    "1004/1004.0376_arXiv.txt": {
        "abstract": "High spatial resolution observations of the H$\\alpha$-emitting wind structure associated with the Luminous Blue Variable star P~Cygni were obtained with the Navy Prototype Optical Interferometer (NPOI). These observations represent the most comprehensive interferometric data set on P~Cyg to date. We demonstrate how the apparent size of the H$\\alpha$-emitting region of the wind structure of P~Cyg compares between the 2005, 2007 and 2008 observing seasons and how this relates to the H$\\alpha$ line spectroscopy. Using the data sets from 2005, 2007 and 2008 observing seasons, we fit a circularly symmetric Gaussian model to the interferometric signature from the H$\\alpha$-emitting wind structure of P~Cyg.  Based on our results we conclude that the radial extent of the H$\\alpha$-emitting wind structure around P~Cyg is stable at the 10\\% level. We also show how the radial distribution of the H$\\alpha$ flux from the wind structure deviates from a Gaussian shape, whereas a two-component Gaussian model is sufficient to fully describe the H$\\alpha$-emitting region around P~Cyg. ",
        "introduction": "Luminous Blue Variable (LBV) star P~Cyg (HD~193237, B2pe) is an unusual star with a unique stellar wind structure. LBVs are evolved, very luminous, massive supergiant stars that show some type of instability. As the likely progenitors of Wolf-Rayet stars~\\citep{crowther07} LBV's may provide insight for the ultimate end-stage of these stars, the core-collapse supernovae.  P~Cyg was discovered in 1600 after a violent mass-loss event caused it to quickly brighten to a third magnitude star \\citep{deg69}. The violent mass-loss events that spur these eruptions are not fully understood, but a number of competing models including single and binary star scenarios attempt to explain these eruptions \\citep{hum94}. With a mass-loss between 2$\\times$10$^{-5}$ and 4$\\times$10$^{-4}$ $M_{\\odot}$~$\\rm yr^{-1}$ \\citep{vanb78,abbott80,lei87}, P~Cyg displays characteristics that are suggestive of energetic mass outflows. It is for this reason that P~Cyg has a unique stellar wind structure with possible high density regions in motion within the wind structure. Despite numerous studies published in the literature that discuss P~Cyg or other LBVs, there is still need for observations that can spatially resolve the circumstellar region around P~Cyg.  The spectrum of P~Cyg reveals emission line profiles with their archetypical shapes, which are formed by the circumstellar matter that has been shed from the central star. These profiles are commonly identified by an emission component to the red side of an absorption line. The origin of this Doppler shift can be explained with a spherically symmetric outflowing wind in which the velocity increases with radial distance~\\citep[see, e.g.,][\\S~2.2]{lam99}.  The absorption component corresponds to the continuum light absorbed by the wind structure directly between the observer and the central star. The emission component on the other hand originates in the spherically symmetric halo around the central star.  Because the same characteristics apply to the H$\\alpha$ emission line, which is one of the strongest emission lines in the spectrum of P~Cyg, this makes the H$\\alpha$ emission line an excellent probe of the wind region around P~Cyg.   P~Cyg currently has an estimated mass of 30~$\\pm$~10~$M_{\\odot}$~\\citep{lam83, lam85}. At the start of its lifetime P~Cyg possibly had an initial mass of 50 $M_{\\odot}$ \\citep{deg92}, but has shed much of its mass because of its violent history of mass loss events. The effective temperature and the stellar diameter of P~Cyg are estimated to be $\\rm T_{\\rm eff}$~=~19,300~$\\pm$~2000~K and 76~$\\pm$~14~$R_{\\odot}$, respectively \\citep{lam83}. However, the H$\\alpha$-emitting region extends radially much farther than the central star allowing us to resolve this region using long-baseline interferometric techniques, even though the distance is estimated to be in the range of $1.7~\\pm~0.1$~kpc to $1.8\\pm~0.1$~kpc \\citep{lam83, naj97}.   The circumstellar region of P~Cyg has been spatially resolved in the past using radio and optical interferometry, as well as direct imaging with adaptive optics~(AO).  Radio interferometric observations detect the nebula around P~Cyg at the angular scales from $\\sim$50~mas to almost an arcminute~\\citep{skin97,skin98}, whereas optical long-baseline interferometry is sensitive to structures in the nebula of P~Cyg that are more than an order of magnitude smaller.  Therefore, optical interferometry provides a unique window of opportunity to resolve the inner portion of the wind structure.   Most relevant to our study are the results that probed the H$\\alpha$-emitting region of P~Cyg.  For example, based on near-diffraction limited observations obtained using AO system on a 1.52~m telescope, \\citet{che00} not only resolved the outer H$\\alpha$-emitting region of the extended envelope, but detected signatures of clumping.  Although, the angular scales sampled with a 1.52~m telescope were quite large, in the region of ~600~mas, the angular resolution was limited to 50~mas.  On much smaller angular scales, \\citet{vak97} using a single 17.7~m baseline of the GI2T interferometer have resolved the inner H$\\alpha$-emitting envelope of P~Cyg yielding an angular size of 5.52$\\;\\pm\\;$0.47 mas and reported a marginal detection at the HeI~6678 line.  The spatial scales probed by our study are similar to those sampled by \\citet{vak97}, however, the interferometric observations at the H$\\alpha$ line obtained in our study were acquired using 11 unique interferometric baselines ranging in length from 15.8 to 79.3~m, thus resulting in much higher angular resolution measurements. ",
        "conclusions": "\\subsection{The Radial Intensity Falloff}  The deviation from a Gaussian shape of the spatially resolved wind structure of P~Cyg is one of our most intriguing results.  Examining the squared visibilities in Figure \\ref{fig:all_v2}, it is apparent that the observations at high spatial frequencies cannot be represented by a single-component Gaussian model.  It is worthwhile to mention that the Gaussian intensity distribution is still a closer approximation to the radial intensity distribution from the wind structure than the uniform disk model (shown with dotted line in Fig.~\\ref{fig:all_v2}).  To our knowledge, this is the first time the H$\\alpha$-emitting region of P~Cyg have been observed to clearly deviate from a Gaussian shape.   To reproduce our data at the highest spatial frequencies, we use the two-component Gaussian model as a way of parameterizing the spatially resolved signature of the H$\\alpha$-emitting region of P~Cyg.  This suggests that the wind structure is indeed more complex in P~Cyg as compared to Be stars.  Because P~Cyg possess radiatively driven stellar wind, the velocity structure can be described by the $\\beta$-law of \\citet{castor79}, which implies that the wind is accelerated to its terminal velocity close to the star and then reaches a constant terminal velocity.  This in turn means that the density structure of P~Cyg will change at two different rates, one when the wind is still accelerating and one when the wind has already reached its terminal velocity at which point the density would be expected to drop with radial distance as $r^{-2}$.  Although this cannot be used as a direct explanation for the two component structure because the H$\\alpha$ emission in P~Cyg is optically thick~\\citep{lei88}, by combining the effects of the density distribution with the radiative transfer effects it should be possible to use the data presented in this study to directly constrain a numerical wind model based on a $\\beta$ velocity law.  Lastly, if the two-component wind structure is dominated by optically thick and thin effects, one could argue that the the inner (optically thick) region might be better represented by a uniform disk model and the outer region by a Gaussian model.  This would be in agreement with the results presented in \\S\\ref{observation} that showed that the observations can be equally well represented by a two-component Uniform-Disk and Gaussian model.   P~Cyg has been spatially resolved in the past using interferometric techniques. \\citet{vak97}, using a single 17.7~m interferometric baseline, resolved the extended envelope of P~Cyg at the H$\\alpha$ emission line. Although their data set was limited and did not allow them to fit a model more complex than a uniform disk model, the angular diameter of a uniform disk model they obtained was $\\theta_{\\rm UD}$~=~5.52~$\\pm$~0.47~mas. They also obtained complementary H$\\alpha$ spectrum on P~Cyg on the same night they obtained interferometric observations with a H$\\alpha$ peak signal of 14.6 $F_{\\rm max}/F_{\\rm cont}$. This is somewhat weaker than the $\\approx$20 $F_{\\rm max}/F_{\\rm cont}$ values we obtain based on our observations from 2005--2007~(recall Fig.~\\ref{fig:Pcyg-spectra}), although the effect of lower spectral resolution in the spectra of \\citet{vak97} cannot be ruled out, which could explain the lower peak-to-continuum ratio seen in the observation reported by \\citet{vak97}.   Although a uniform disk model is completely inconsistent with our observations (recall Fig.~\\ref{fig:all_v2}), in order to compare our results to \\citet{vak97} at effectively the same limiting spatial resolution, we fitted a uniform disk model to data from only shortest baselines (up to 18.9 m).  The best-fit uniform disk angular diameters ranged from 8.4 to 10.2 mas depending on which season was used in the fit (see Table~\\ref{modelpar}). Therefore, our UD diameter for the H$\\alpha$-emitting wind structure is not consistent with the reported diameter of 5.52~mas by \\citet{vak97}. This discrepancy might be due to the possibility that \\citet{vak97} has modeled the observational signal that contained both the mostly unresolved central star and H$\\alpha$-emitting envelope with only one uniform disk diameter.  This would be equivalent to our equation~\\ref{eq:model1}, but with the $c_{\\rm p}$ set to 0, which would tend to underestimate the angular extent of the H$\\alpha$ emitting region.  The other possibility is that the H$\\alpha$-emitting region actually grew in size from 1994 to 2007. If we assume that the difference between our values and those reported by \\citet{vak97} to be representative of a real change, then this results in an expansion rate of the H$\\alpha$-emitting region~(due to changes in opacity) over this thirteen year period of $\\sim$3~km~s$^{-1}$.  This is well within the expected wind velocities of stars with radiation driven wind structures, especially P~Cyg, which has a terminal wind velocity almost two magnitudes larger~\\citep{Barlow94,lam85}.  Furthermore, if the H$\\alpha$-emitting region did indeed grow in size, we would expect the H$\\alpha$ emission to be stronger, which is supported by the fact that the $F_{\\rm max}/F_{\\rm c}$ values in Table~\\ref{specvartab} that range between 17 and 24 are larger than the $\\approx$~15 peak value of the spectrum taken by \\citet{vak97}.   \\subsection{Signature of Asymmetry}  To test for the presence of a signature of deviation from point-symmetry in the signal in the H$\\alpha$ channel, we have plotted the weighted mean closure phases in Figure~\\ref{fig:closure_phases}. If the H$\\alpha$ emission originates from an intensity distribution that is not symmetric across the origin of the photocenter, or equivalently the H$\\alpha$ emitting region is not concentric with the central star, we would expect to see a non-zero closure phase in the H$\\alpha$ channel.  At a close inspection of Figure~\\ref{fig:closure_phases} we conclude that the H$\\alpha$ closure phases are generally very close to $0^{\\circ}$.  Even in the cases where the mean closure phase deviates from zero values, it is only at the level of a couple of degrees, which also happens to be only 2 to 3 times the uncertainty of the mean value.  Therefore, we conclude that based on our observations we do not have a strong signature of deviation from point symmetry of the H$\\alpha$-emitting region.  Our conclusions are not necessarily inconsistent with the a very similar detection of a subtle phase variation~(at the level of $\\sim 30^{\\circ}$) across the H$\\alpha$ emission line detected by \\citet{vak97} who interpreted that as being produced by a localized spatial feature~(i.e., a localized blob) within the structure of the wind.  Assuming that the closure phase variations we detect across the H$\\alpha$ spectral channel in Figure~\\ref{fig:closure_phases} are caused by similar spatial feature (which is not necessarily expected since the observations were acquired at different epochs), we would indeed expect to see a much weaker signal across the H$\\alpha$ channel for two reasons.  The first being that we measure the sums of three phases, which already have the tendency to lower the signal when negative and positive phases in a triangle are added together.  The second effect is related to our much wider spectral channel that contains the entire H$\\alpha$ emission line and a significant contribution from the central star, which most likely can be described well by a point-symmetric intensity distribution that contains only real components in the Fourier space (i.e., only real phases).  The net result is that the real components will tend to lower the net complex phase detected in the H$\\alpha$ channel.  Signatures of clumping in the circumstellar region of P~Cyg on much larger spatial scales have also been reported.  For example, based on near-diffraction limited observations using AO system on a 1.52~m telescope, \\citet{che00} reported clumping at the scales of 200--600~mas.  Similarly, P~Cyg was also imaged using the Multi-Element Radio Linked Interferometer Network (MERLIN) array by \\citet{skin97}.  Their 6~cm observations produced images of the circumstellar structure on scales of 100--200 mas with a 50 mas resolution, revealing structural changes on a 40 day timescale along with flux variations at the 20\\% level.  \\citet{skin97} argued that the structural variations could not be attributed to variations in the mass-loss rate. They suggested recombination within the free-free emitting region responsible for the radio emission could explain their observations.  The weak signature of the non-zero closure phase in the NPOI data, if confirmed, could indicate that the clumpiness observed by \\citet{che00} and \\citet{skin97} already originates on scales of less than 10 mas (less than 25 stellar radii).   \\subsection{Photometric Variability}   Lastly, to assess the level of photometric variability during the time frame covered by our interferometric observations, we used photoelectric V-band observations from the American Association of Variable Star Observers.  The $V$-band light curve from 2005 to 2009 is shown in Figure~\\ref{fig:EW_vs_V} where we only include data for which the measured check star~(HD~193369) magnitude was no more than 0.04~mag away from the mean value of 5.573~mag.  We see both a mean seasonal-timescale variability on the order of $\\sim 0.02$ mag, and very short term variability (for observations taken within days to weeks) at the level of $\\pm 0.05$~mag.  Although the short-term variability appears to have large scatter, the photometric quality of the AAVSO PEP data is generally good, with internal errors on the order of 5--10 milli-magnitudes.  Therefore, based on the photometric data we conclude that although the mean photometric level appears to change by less than 0.02~mag, short term variations upwards of 0.1~mag cannot be ruled out.  This implies that the continuum level in P~Cyg could vary up to 10~\\% level, which would affect directly the EW measures at the same level.  This also strongly suggests that the variations in EW measures seen in Table~\\ref{specvartab} are caused by variations in the continuum level and not the emission component itself."
    },
    "1004/1004.2659_arXiv.txt": {
        "abstract": "We present the first view of the magnetic field structure in the OH shell of the extreme OH/IR star OH~26.5+0.6. MERLIN interferometric observations of this object were obtained in December 1993 in full polarisation, at 1612, 1665 and 1667~MHz. The maser spots show a spheroidal distribution both at 1612 and 1667~MHz, while at 1665~MHz emission from the blue-shifted maser peak is concentrated on the stellar position, and the red-shifted peak emission exhibits a filamentary structure oriented on a SE-NW axis. The linear polarisation in both main lines is rather faint, ranging from 9 to 20\\% at 1665~MHz and from 0 to 30\\% at 1667~MHz. At 1612~MHz most maser spots exhibit a similar range of linear polarisation although those in the outermost parts of the envelope reach values as high as 66\\%. This is particularly apparent in the southern part of the shell. The detailed distribution of the polarisation vectors could only be obtained at 1612~MHz. The polarisation vectors show a highly structured distribution indicative of a poloidal magnetic field inclined by 40-60$^\\circ$ to the line of sight. The velocity distribution of the maser spots with respect to the radial distance is well explained by an isotropic outflow at constant velocity in the case of a prolate shaped spheroid envelope, also tilted about 45-65$^\\circ$  to the line of sight. ",
        "introduction": "After leaving the main sequence, low and intermediate mass stars experience a crucial phase in their evolution toward the white dwarf stage: the Asymptotic Giant Branch (AGB) phase. It is at the very end of this phase that the star will shed most of its mass through extensive mass loss (up to a few $10^{-4} {\\rm M_{\\odot} yr^{-1}}$). The exact evolutionary sequence along the AGB to this final stage has not yet been resolved, but OH/IR stars are thought to trace the period before the proto-planetary nebula stage. At that point, the central star is completely obscured in the optical by a thick dust shell built up by mass loss, but the envelope structure can be observed through strong emission in the ground state OH maser lines and at infrared wavelengths.  While AGB stars are fairly spherical objects, asymmetries such as elliptical shapes or bipolar outflows are commonly observed at the planetary nebula stage (Corradi \\& Schwarz 1995).  Recently, a series of papers investigated the polarimetric structure in the intermediate and outermost parts of the circumstellar shells of evolved stars (Bains et al. 2003, Etoka \\& Diamond 2004, Vlemmings et al. 2005, Vlemmings \\& Diamond 2006). Although the origin and evolution of the magnetic field is not well understood and is currently a matter of debate, (cf. Nordhaus et al. 2007 and reference within), this series of papers has shown the importance of the magnetic field in shaping the circumstellar material. \\\\  OH~26.5+0.6 (AFGL~2205; IRAS~18348$-$0526) is an extreme OH/IR star at a distance of 1.37$\\pm$0.30~kpc (van~Langevelde et al. 1990). Its current mass-loss rate has been estimated to be on the order of $5 \\times 10^{-4}$~M$_{\\odot} {\\rm yr^{-1}}$ (Justtanont et al. 1996). It has been classified as a Very-Long Period Variable OH/IR star with a period of 1570~days (le~Bertre 1993). Prior to that work, OH~26.5+0.6 has been imaged several times with the VLA at 1612~MHz with increasing sensitivity (Baud 1981; Bowers et al. 1983; Herman et al. 1985 and Bowers \\& Johnston 1990) where a complete ring-like structure is seen at virtually all velocities. It has also been imaged with MERLIN (Diamond et al. 1985), where the clumpiness of the shell was clearly revealed.  The work presented here is part~II of a series of papers intending to unravel the magnetic structure around extreme OH/IR stars through observations in the ground state OH maser lines at 18~cm. The first paper of the series, Etoka \\& Diamond (2004, hereafter paper~I), presents the magnetic field structure of the red supergiant NML~Cyg at 1612 and 1667~MHz. This first work has shown that a structured polarisation distribution exists for both lines linked with the geometry of the shell itself. This can be explained if the principal driver for the shaping of the shell is the magnetic field.  The details of the observations and data reduction process are given in Section~2. An analysis of the data is presented in Section~3. In Section~4 discussion and interpretation of the results is given, while conclusions are drawn in Section~5. ",
        "conclusions": "The infrared and 18~cm OH maser properties of OH~26.5+0.6 attest to a thick circumstellar envelope, characteristic of a rather evolved star most probably at the tip of the AGB. The 1612~MHz emission reveals an ellipsoidal geometry, while the presumably more central 1665~MHz emission traces a filamentary structure. Both 1612 and 1667~MHz high resolution maps show a lack of maser emission from some parts of the shell, and in the southern and north-eastern part of the shell, 1667~MHz emission extends beyond that at 1612~MHz. All these deviations from the standard spherical model show that the OH/IR stage (i.e., late-AGB phase) is clearly at the stage where asymmetry starts to develop. The presence of acceleration in the shell at the OH maser location may be a secondary factor enhancing the asymmetry observed so far away from the star. The root of this asymmetry is likely to be close to the stellar surface itself, as near infrared results indicate. This latter hypothesis is reinforced by the agreement in orientation of the major axis of the elliptic distribution observed in infrared and at 1612~MHz. Finally, we found that the magnetic field strength, inferred from OH Zeeman splitting, is such that the magnetic field energy density dominates over the thermal and kinetic pressures and that there is a definite correlation between the magnetic field orientation and the main axis of geometrically ellipsoidal maser emission. This suggests that the magnetic field plays a role in the shaping process observed."
    },
    "1004/1004.2684_arXiv.txt": {
        "abstract": "{Disformal transformations have proven to be very useful to devise models of the dark sector. In the present paper we apply such transformation to a single scalar field theory as a way to drive the field into a slow roll phase. The canonical scalar field Lagrangian, when coupled to a disformal metric, turns out to have relations to bimetric dark matter theories and to describe many specific dark energy models at various limits, thus providing a surprisingly simple parametrisation of a wide variety of models including tachyon, Chaplygin gas, K-essence and dilatonic ghost condensate. We investigate the evolution of the background and linear perturbations in disformal quintessence in order to perform a full comparison of the predictions with the cosmological data. % The dynamics of the expansion, in particular the mechanism of the transition to accelerating phase, is described in detail. We then study the effects of disformal quintessence on cosmic microwave background (CMB) anisotropies and large scale structures (LSS). A likelihood analysis using the latest data on wide-ranging SNIa, CMB and LSS observations is performed allowing variations in six cosmological parameters and the two parameters specifying the model. We find that while a large region of parameter space remains compatible with observations, models featuring either too much early dark energy or too slow transition to acceleration are ruled out. } ",
        "introduction": "Observations of the Cosmic Microwave Background temperature anisotropy reveal that a mysterious constituent with negative pressure, so called dark energy, accounts for about $70\\%$ percent of today's mass energy budget and is causing the expansion of the universe to accelerate \\cite{wmap,lss}. These observations are in remarkable concordance with the observations of distant supernovae \\cite{sn,sn1}. A major challenge in present day cosmology is to discover the physical nature of dark energy. Meanwhile, though the evidence for the existence of dark matter has been accumulating for several decades, the quest to find its precise nature, whether as a particle described by an extension of the standard model, a more exotic field or even a modification of General Relativity, is ongoing. Many possibilities have been explored in attempts to explain the workings of the dark sector \\cite{copeland,Durrer:2007re,koi1,koi2,Li1,koi3,aether,koi4,Li2,koi5,li,lli1}.   The phenomenologically simplest but theoretically very problematical dark energy is the Einsteins $\\Lambda$ term \\cite{Nobbenhuis:2006yf}. Quintessence provides a dynamical alternative to the static cosmological constant \\cite{Wetterich:1987fm,Peebles:1987ek}. It tracks or scales with the background energy density, therefore more naturally resulting in similar orders of magnitude for the dark energy and dust-like matter energy densities today. However, one needs some mechanism to end this scaling and to onset the acceleration. Several possibilities have been considered, in particular introducing a suitable bump into the form potential \\cite{Albrecht:1999rm}, taking into account the Gauss-Bonnet term \\cite{Koivisto:2006ai,Koivisto:2006xf}, coupling the field to other matter \\cite{Amendola:1999er,Koivisto:2005nr}, considering non-canonical Lagrangian \\cite{ArmendarizPicon:2000ah}, or introducing many fields \\cite{Copeland:2000vh}. In the following, a disformal relation \\cite{Bekenstein:1992pj} is applied for this purpose.  Disformal relations have recurrently appeared in cosmology in models of alternative dark matter, in particular TeVes \\cite{teves} and other relativistic and covariant formulations of MOND \\cite{mond}, and in bigravity theories \\cite{Banados:2008rm}. This comes about for several reasons: theoretically, it is related to the Born-Infeld type of Lagrangians (as will become clear below), and phenomenologically, it is needed to produce the observed lensing without resorting to particle dark matter \\cite{Skordis:2009bf}. Let us mention that the possible relevance to inflation has also been considered \\cite{Kaloper:2003yf}, and that disformality is a key to varying speed of light, which has been considered as an alternative to inflation \\cite{Clayton:1998hv,Clayton:2001rt}. The usefulness of conformal transformations is well known and appreciated \\cite{Faraoni:2006fx,Kastrup:2008jn,Dabrowski:2008kx}, but many aspects of the wider framework remain uncovered. The conformal transformation relates the simplest scalar-tensor theories (including the $f(R)$ theories \\cite{DeFelice:2010aj}) to general relativity \\cite{Magnano:1993bd}, but in the general case one is prompted to turn into extended classes of transformations.  Here we make a preliminary step along this direction by exploring the perhaps simplest possible set-up: the canonical scalar field coupled to a disformal metric. It turns that already then many relations between seemingly disconnected models emerge. If the scalar field is reduced to a constant term, it appears as a tachyon or a Chaplygin gas \\cite{Kamenshchik:2001cp} in the physical frame. When the kinetic term is included, some other previously considered models can be recovered, as we will discuss in Section \\ref{dm}. In the present study we then focus on one simple possibility, disformal quintessence. We point out that it can act as viable dark energy when the parameters in the Lagrangian are of the Planck scale. Section \\ref{dq} is devoted to study of this model, and contains the main results of this paper. We describe in detail the background evolution of this model in subsection \\ref{background-section}. Particular attention is given to the transition mechanism providing an exit from the scaling era. The details of this transition depend on the two parameters of the model, and thus they can be constrained by the SNeIa data. We also consider the evolution of perturbations in subsection \\ref{perturb-section} in order to compute the CMB and matter power spectra. Armed with these solutions, we perform a Monte Carlo analysis of the model in subsection \\ref{constraints-section}, combining data from many different cosmological experiments. Wide parameter range is found to be compatible with observations, but certain parameter region, corresponding to shallow slope of the exponential potential or of the disformal factor, can be ruled out. We conclude briefly in section \\ref{conclusions-section}. ",
        "conclusions": "\\label{conclusions-section}  Disformal generalizations of the conformal transformation have found several applications in cosmology, particularly in the frameworks of gravitational alternatives to dark matter and varying light speed alternatives to inflation. In this paper we have shown that the disformal relation can be also very prolific in generating alternative explanations for the cosmic acceleration, suggesting new links between models worth further explorations. In the future, it would be interesting to study the disformal relations which may exist between more general field theories. Already the details of motion of point particles in a disformal metric are, to our knowledge, not very well understood. Whilst the nature of dark energy is undisclosed and it is not known to which metric it is coupled, there is in fact no fundamental principle establishing the precise relation between the gravitational metric and the physical metric in which ordinary matter lives \\cite{bi,bi2,bi3,teves,aether}. It is thus useful to investigate the observational signatures and put experimental bounds on it from e.g. classical tests of the equivalence principle.  Here we considered the canonical scalar field action where metric is replaced by the disformal one, i.e. made the substitution $g_{\\mu\\nu}\\to\\bar g_{\\mu\\nu}$ in the Lagrangian. At various limits we can then obtain Chaplygin gas and other models considered in the literature. Disformal quintessence arises in a very simple way as an application of this substitution, which results effectively in a non-standard self-interaction of the scalar field. This self-interaction causes the field to accelerate the universe given exponential functions of the field and under the condition $\\alpha/\\beta<1$ in (\\ref{B},\\ref{V}). The acceleration can begin near the present epoch if both $\\alpha$ and $\\beta$ are of order one and $\\phi_x\\sim M_p$, meaning that {\\it all} the parameters appearing in the Lagrangian can be set to the Planck scale. For the magnitude of the potential and the disformal factor this follows from the special shift symmetry property of the exponential form. Moreover, as well known, in cosmology the exponential potential has the special scaling property, so practically no tuning whatsoever of initial conditions for the field is required. These observations motivated us to consider the phenomenology of the model in detail and to confront it with the latest available data.   We considered evolution both at the level of background and of linear perturbations. The dependence of the background dynamics on model parameters allows us to put constraints on them. These become more stringent if we take into account the evolution of linear structures, though the effect of the perturbations of the field to them is more subtle. A Monte Carlo Markov Chain simulation was used to obtain constraints on the parameter space of the theory by comparing with WMAP 7 year data, SNe from the Union dataset, baryon acoustic peak position and luminous red galaxies power spectrum from SDSS. Six cosmological parameters were allowed to vary together with $\\alpha$ and $\\beta$.  It was then found that small values of $\\alpha$ can be ruled out because of the effect of early dark energy that implies. On the other hand, small values of $\\beta/\\alpha$ result in slow transition to acceleration, which is also disfavored by the data. However, when both $\\alpha$ and $\\beta/\\alpha$ are sufficiently large (the precise limits shown in Figure \\ref{contours}), it becomes very difficult to distinguish the model from $\\Lambda$CDM and thus this parameter region remains compatible with the present data. These constraints might be improved by considering nonlinear scales and gaining a deeper understanding of the effects on large scale structure.  Further research might focus on the role of the disformal relation in dark energy scenarios by the introduction of couplings to other forms of matter and gravity constructed by following the disformal prescription $g_{\\mu\\nu}\\to\\bar g_{\\mu\\nu}$ as well the relation of disformal quintessence with other forms of dark energy and non-minimal derivative coupled version of scalar tensor theories."
    },
    "1004/1004.2960_arXiv.txt": {
        "abstract": "We derived O, Ne, and Mg abundances in the interstellar medium (ISM) of a relatively isolated S0 galaxy, NGC 4382, observed with the Suzaku XIS instruments and compared the O/Ne/Mg/Fe abundance pattern to those of the ISM in elliptical galaxies. The derived temperature and Fe abundance in the ISM are about 0.3 keV and 0.6--2.9 solar, respectively. The abundance ratios are derived with a better accuracy than the  abundances themselves:  O/Fe, Ne/Fe, and Mg/Fe ratios are 0.3, 0.7, and 0.6, respectively, in solar units. The O/Fe ratio is smaller than that of the ISM in elliptical galaxies, NGC 720, NGC 1399, NGC 1404, and NGC 4636, observed with Suzaku. Since O, Ne, and Mg are predominantly synthesized by supernovae (SNe) of type II, the observed abundance pattern indicates that the contribution of SN Ia products is higher in the S0 galaxy than in the elliptical galaxies. Since the hot ISM in early-type galaxies is an accumulation of stellar mass and SN Ia products, the low O/Fe ratio in the ISM of NGC 4382 reflects a higher rate of present SNe Ia, or stars containing more SN Ia products than those in elliptical galaxies. ",
        "introduction": "Early-type galaxies have a hot, X-ray emitting interstellar medium (ISM), which is considered to be gravitationally confined (e.g., \\cite{Form1985,Math2003}).  X-ray observations of the metal abundances in the ISM of early-type galaxies provide a key to understanding the history of star formation and the evolution of galaxies, because the metals in the ISM come from type Ia supernovae (SNe Ia) and stellar mass loss.  A lot of  observational evidence suggests that a significant fraction of present early-type galaxies have transformed from late-type galaxies (e.g., \\cite{Dres1997,Fasa2000,Treu2003,Post2005,Smith2005}). From z$\\sim$0.5, the fraction of S0 galaxies  in clusters of galaxies has increased, whereas, that of spirals has decreased, with no evolution in the fraction of ellipticals \\citep{Koda2004,Desai2007}. \\citet{Pogg2009} found that these changes appear more strongly in less massive clusters with lower velocity dispersion. These results suggest that spiral galaxies changed into S0 galaxies at z$<$0.5, falling into clusters of galaxies. The  metallicity  of stars in galaxies reflects in the star formation history, therefore, it is an important parameter for understanding the evolution of galaxies.  Optical observations indicate that  Mg/Fe and $\\alpha$/Fe ratios of stars are super-solar in the cores of bright early-type galaxies and increases with galactic mass (e.g., \\cite{Fabe1992,Wort1992,Nela2005,Thom2005,Grav2007}). This overabundance of Mg relative to Fe is the key indicator that galaxy formation occurred before a substantial number of SNe Ia could explode and contribute to lowering the these ratios (e.g., \\cite{Bern2003,Nela2005,Smith2006,Grav2007,Pipi2009}). However, absorption-line indices that account for abundance ratios also depend on the age distribution of stars.  Optical spectroscopy is limited within the very center of galaxies.  Using X-ray observations, we can directly determine the metal abundances of the ISM, and  constrain the stellar metallicity of the entire galaxy. The atomic data for lines at X-ray wavelengths are  simpler than for those in optical spectra, and the structure of the hot ISM is also much simpler than stellar population data. Therefore, we can estimate the temperature and metallicity of the hot ISM through X-ray spectra with small systematic uncertainties. XMM-Newton EPIC and RGS provided the means to measure O and Mg abundances in some systems, but reliable results have been obtained only for several central galaxies in groups and clusters (e.g., \\cite{Xu2002,Matsu2003,Tamu2003,Matsu2007a}). The ISM in such regions might be polluted by the gas of clusters or groups (e.g., \\cite{Matsu2001,Matsu2002,Nagi2009}).  With the Suzaku X-ray satellite, O, Ne, and Mg abundances of four elliptical galaxies, NGC 720 \\citep{Tawara2008}, NGC 1399 and NGC 1404 \\citep{Matsu2007b}, and NGC 4636 \\citep{Haya2009}, have been measured. The XIS detector onboard Suzaku \\citep{Mitsu2007} can constrain O, Ne, and Mg abundances well because its energy resolution is better, and its background is lower than any previous X-ray CCD detector \\citep{Koya2007}. According to the new solar abundance table of \\citet{Lodd2003}, the abundance ratios of these elliptical galaxies are close to the solar values. Assuming the SN II abundance pattern of \\citet{Iwa1999}, about 80\\% of Fe is synthesized by SNe Ia \\citep{Matsu2007b}.  The abundance pattern of the hot ISM in S0 galaxies observed with Suzaku has not been reported. In this paper, we present the temperature and abundances of O, Ne, Mg, and Fe in the ISM within 4~$r_e$ of the S0 galaxy NGC 4382, observed with Suzaku. Here, $r_e$ is the effective radius of the galaxy. For NGC 4382, $r_e$ corresponds to 0.91 arcmin \\citep{deVau1991}. The distance to NGC 4382 is 16.8 Mpc \\citep{Tully1988}, and its redshift, z=0.002432, is taken from the NASA/IPAC extragalactic database (NED). NGC 4382 is located in the outskirts of the Virgo cluster, 1.7 Mpc from cD galaxy M87. The central stellar velocity dispersion of the galaxy, 179 km s$^{-1}$ and its $[\\alpha/{\\rm Fe}]$ value, $0.12\\pm0.06$ \\citep{McD2006}, are much smaller than those of the four elliptical galaxies, observed with Suzaku. These values were derived only within two arcsec of the center of NGC 4382, and we do not know the metallicity of stars outside of the central region. So far, NGC 4382 has been observed with ROSAT \\citep{Fabb1994}, ASCA \\citep{Kim1996}, Chandra \\citep{Siva2003,Athey2007} and XMM-Newton \\citep{San2006,Nagi2009}. These X-ray observations suggest that the temperature of the ISM in NGC 4382 is about 0.3--0.4 keV, which is much lower than those of the four early-type galaxies observed with Suzaku. Since no the intracluster medium (ICM) around NGC 4382 has been detected, this galaxy is suitable for investigating heavy elements in the ISM of the galaxy itself.  Throughout this paper, we use a Hubble constant of $H_0 = 70$ ${\\rm km}$ ${\\rm s^{-1}}$ ${\\rm Mpc^{-1}}$. We adopt the new solar abundance table of \\citet{Lodd2003}. The abundances of O and Fe have increased by about 70\\% and 60\\%, respectively. Unless otherwise specified, errors are quoted at 90\\% confidence. ",
        "conclusions": "With the Suzaku observation of S0 galaxy NGC 4382, we measured the ISM temperature, metal abundances of  O, Ne, Mg and Fe, and their abundance ratios for the region within 4~$r_e$ of the galaxy's center. The temperature, 0.3 keV, and  O/Fe ratio, 0.3 in solar units, in the ISM of this galaxy are  smaller than those in elliptical galaxies observed with Suzaku, NGC 720 \\citep{Tawara2008}, NGC 1399 and NGC 1404 \\citep{Matsu2007b}, and NGC 4636 \\citep{Haya2009}. The lower O/Fe ratio of the ISM may reflect a higher rate of present SN Ia or a lower $\\alpha/{\\rm Fe}$ ratio in stars in  NGC 4382 \\citep{McD2006} than in those in the four elliptical galaxies.  \\subsection*{}  We thank to the  referee for a careful reviewing and many helpful comments. We also thank all members of the Suzaku hardware and software teams and the science working group."
    },
    "1004/1004.5401_arXiv.txt": {
        "abstract": "The mass of molecular gas in an interstellar cloud is often measured using line emission from low rotational levels of CO, which are sensitive to the CO mass, and then scaling to the assumed molecular hydrogen H$_2$ mass. However, a significant H$_2$ mass may lie outside the CO region, in the outer regions of the molecular cloud where the gas phase carbon resides in C or C$^+$. Here, H$_2$ self-shields or is shielded by dust from UV photodissociation, whereas CO is photodissociated.  This H$_2$ gas is ``dark\" in molecular transitions because of the absence of CO and other trace molecules, and because H$_2$ emits so weakly at temperatures 10 K $< T < 100$ K typical of this molecular component. This component has been indirectly observed through other tracers of mass such as gamma rays produced in cosmic ray collisions with the gas and far-infrared/submillimeter wavelength dust continuum radiation. In this paper we theoretically model this dark mass and find that the fraction of the molecular mass in this dark component is remarkably constant ($\\sim 0.3$ for average visual extinction through the cloud $\\bar A_V \\simeq 8$) and insensitive to the incident ultraviolet radiation field strength, the internal density distribution, and the mass of the molecular cloud as long as $\\bar A_V$, or equivalently, the product of the average hydrogen nucleus column and the metallicity through the cloud, is constant.   We also find that the dark mass fraction increases with decreasing $\\bar A_V$, since relatively more molecular H$_2$ material lies outside the CO region in this case. ",
        "introduction": "\\label{sec:intro}  Various observations have indicated that a substantial amount of interstellar gas exists in the form of molecular hydrogen (${\\rm H_2}$) along with ionized carbon (${\\rm C^+}$), but little or no carbon monoxide (CO). The total mass in molecular hydrogen has been estimated from gamma ray observations from COS-B \\citep{bloemen1986} and EGRET (Energetic Gamma Ray Experiment Telescope) \\citep{strong1996} and analysis of this data showed more gas mass than can  be accounted for in \\ion{H}{1} and CO alone \\citep{grenier2005}. In addition, the dust column density maps of the Galaxy from DIRBE, and maps of the 2MASS J-K extinction show additional gas not seen in \\ion{H}{1} or CO \\citep{grenier2005} and is presumably molecular hydrogen. In molecular line observations and modeling of low column density molecular clouds the ${\\rm H_2}/{\\rm CO}$ ratio is found to be variable and much larger than $10^{4}$ so that only a small fraction of the C is in CO with the remainder presumably in ${\\rm C^+}$ \\citep{hollenbach2009,goldsmith2008}. \\cite{reach1994} using infrared continuum maps of diffuse clouds along with \\ion{H}{1} and CO observations, found ${\\rm H_2}$ masses comparable to the \\ion{H}{1} masses, with only small amounts of CO. A mixture of ${\\rm H_2}$ and ${\\rm C^+}$ is also inferred to exist in diffuse clouds where the ${\\rm H_2}$ and CO columns are measured by UV absorption spectroscopy and the CO accounts for only a trace amount of C \\citep{sonnentrucker2007,burgh2007,sheffer2008}. This component of the interstellar medium (ISM), termed ``dark gas'' by \\citet{grenier2005}, is also inferred to exist in extragalactic observations comparing far-infrared and CO mass estimates especially in low metallicity galaxies \\citep{israel1997,leroy2007}.  Such an ${\\rm H_2}$ and ${\\rm C^+}$ layer is also predicted from theoretical models of diffuse gas \\citep{vandishoeck1988} and surfaces of molecular clouds \\citep{tielens1985} that indicate the transition from ${\\rm C^+}$ to CO is deeper into the cloud than the transition from H to ${\\rm H_2}$. Essentially, the theoretical models show that  H$_2$ self-shields itself from UV photodissociation more effectively than CO.  This layer is ``dark\" in rotational H$_2$ transitions primarily because the ground state transition 0-0 S(0) at 28 $\\mu$m lies about $\\Delta E/k \\simeq 512$ K above ground; at the temperatures 10 K $< T < 100$ K typical of the dark component the fluxes from the H$_2$ rotational transitions lie below current sensitivities. This layer is dark in CO because of its very low abundance. Although the dark gas layer is ``dark'' in H$_2$ and CO, it does emit mainly in [\\ion{C}{2}] 158 $\\mu$m fine-structure line emission. In preliminary estimates (M.\\ Wolfire et al.\\ 2010, in preparation) the calculated local Galactic [\\ion{C}{2}] emission from the WIM and CNM (diffuse phases that are not associated with molecular clouds) underestimate the COBE observations of the line emission in the plane by a factor of 1/3 to 1/2 and could be another indicator of dark gas. Similarly, \\cite{shibai1991} and \\cite{cubick2008} found the bulk of the [\\ion{C}{2}] emission in the Galactic plane seen by BIRT and COBE arises in neutral gas associated with molecular clouds.  In this paper we present models of molecular cloud surfaces to estimate the mass of gas in the ``dark'' component. In \\S 2 we discuss modifications to existing photodissociation region (PDR) codes, and the modeling procedure. In \\S 3 we define the dark gas mass fraction and other parameters used in our models, discuss the average gas density distribution assumed in our modeling, and derive an expression for the dark gas fraction in terms of parameters found in our numerical modeling procedure.  The modeling results start in \\S 4 with a simple isobaric cloud model in the limit of clumps that are optically thin to the incident radiation field, resulting in an estimate of the dark gas fraction for a typical giant molecular cloud.  We then in \\S 5 enhance this model with the inclusion of turbulent pressure and find the dark gas fraction as a function of incident field strength and cloud mass. We also discuss the variation in the dark gas fraction as a function of metallicity (over a limited range), the average hydrogen column through a cloud, the average visual extinction through a cloud, and the opacity of the clumps. We compare our results with observations.  Finally, in \\S 6 we conclude with a discussion and summary. ",
        "conclusions": "\\label{sec:Discussion}  \\subsection{Model Assumptions}  We have constructed models of molecular clouds to investigate the fraction of gas that is mainly ${\\rm H_2}$  and contains little CO. These conditions exists where the CO is photodissociated into C and C$^+$ but the gas is molecular ${\\rm H_2}$ due to either ${\\rm H_2}$ self-shielding or dust shielding. Such conditions can exist either on the surfaces of molecular clouds or the surfaces of clumps contained within such clouds. Observations indicate that the mass in this ``dark gas'' can be as high as 30\\% of the total molecular mass in the local Galaxy. Previous theoretical plane-parallel models  of individual PDRs have indicated that such a layer should exist \\citep[e.g.,][]{vandishoeck1988}, but here we construct models directed towards molecular clouds as a whole while including observational constraints on cloud mass, radius, average density, and line width and theoretical considerations of the likely strengths of the UV fields impinging on GMCs.  We assume that the surface of each cloud is isotropically illuminated over $2\\pi$ steradians by a soft X-ray/EUV field and an FUV radiation field. We use the standard cosmic-ray ionization rate of $1.8\\times 10^{-17}$ ${\\rm s^{-1}}$ per hydrogen nucleus everywhere in the cloud for all cases but one test  case. There is evidence from observations of ${\\rm H_3^+}$ that the cosmic-ray ionization rate is a factor of 10 higher in some portions of the diffuse ISM \\citep{indriolo2007}, but there is no indication that such rates apply in molecular cloud interiors \\citep{mccall1999}. In preliminary work of M.\\ Wolfire et al.\\ (2010, in preparation) we find that the average FUV field on clouds is $\\sim 20$ times the local Galactic interstellar field. This elevated field arises from the distribution of OB associations around the cloud.  The distribution of temperature, density, and abundances within the \\ion{H}{1}, ${\\rm H_2}$, and CO layers are calculated using the PDR code of \\cite{kaufman2006}. In constructing model clouds from the PDR output, we impose the constraint from \\cite{solomon1987} that the typical column density through the cloud is $\\bar{N}_{22}= 1.5$, independent of cloud mass and independent of radius inside a given cloud (since $\\bar n \\propto r^{-1}$). The locally averaged density $\\bar{n}$  and radius of the CO cloud, $\\rco$, as functions of the CO-cloud mass, $M(\\rco)$ follow from equations (\\ref{eq:barN}) and (\\ref{eq:barnc}). In our notation, $M(\\rco)$ is the molecular mass contained within the $\\tau_{\\rm CO}=1$ surface (of the CO $J=1-0$ transition) including the mass of H$_2$ and He, and $\\rco$ is the radius of this CO photosphere. We impose a volume-averaged density distribution $\\bar{n}(r)$ that behaves as $\\bar{n}\\propto 1/r$ throughout all layers of the cloud. Note that in regions where the thermal pressure $P_{\\rm th}$ lies between $\\pmin$ and $\\pmax$ both warm ($T\\ag 7000$ K) and cold ($T\\al 500$ K) gas solutions are possible. In this regime we use the cold solution from our model results with $n_c$ the density of the cold clumps and $\\bar{n}$ the average over cold and warm gas components.  The warm gas fills the volume, but contains little of the mass.  We test two extreme limits for clumps within the cloud, one in which all clumps are optically thin to the incident FUV field, and one in which all clumps are optically thick to the incident FUV field. In the optically thin approximation, the PDR model output as a function of $A_V$ gives directly the distribution in $A_V$ throughout the cloud. In the optically thin limit we test two different models for the cloud density distribution. First we use constant thermal pressure models and second we include two sources of thermal pressure, radiative heating (which by itself would lead to a two-phase equilibrium) and supersonic turbulence. The distribution of two-phase thermal pressure is calculated as in \\cite{wolfire03} and stored in a look-up table as a function of total column density and molecular fraction (see Fig.\\ \\ref{fig:paveplt}). This pressure drops as one moves into the cloud due to the absorption of the radiation responsible for heating the gas. Supersonic turbulence is characterized by a mass-weighted median density $\\langle n \\rangle_{\\rm med}=\\bar{n}\\exp(\\mu)$, with $\\mu = 0.5\\ln (1+0.25{\\cal M}^2)$ \\citep{padoan1997}. When the two-phase pressure drops below $P_{\\rm th}^{\\rm turb}\\equiv x_t\\langle n \\rangle_{\\rm med} k T$, where $x_t$ is the sum over the abundances of all species relative to hydrogen nuclei, we assume that the turbulence maintains the gas at a thermal pressure $P_{\\rm th}^{\\rm turb}$. The sound speed that enters in the Mach number is calculated from the PDR model output while the turbulent velocity is given by the linewidth-size relation (eq.\\ \\ref{eq:vturb}).   In the limit of very optically thick clumps (with optical depths at least as large as the average optical depth through the cloud), we find that the average radiation field on a clump ranges between about $G_0'/2$ to $G_0'$. In this limit, the dark mass fraction of a clump is the same as the dark mass fraction of the entire cloud.  Since the dark gas mass fraction of a spherical clump or a spherical cloud is insensitive to the incident FUV field (see eqs.\\ \\ref{eq:fDGN} and \\ref{eq:deltaavmainz}), the dark mass fraction of the cloud  does not change significantly compared with the case of a cloud made up of optically thin clumps.  We have tested the steady state assumption for the chemistry by comparing the time to form molecular hydrogen, $t_{\\rm chem}$, with the dynamical time, $t_{\\rm dyn}$, for turbulence to bring molecular gas from the interior to the surface and to bring atomic gas to the interior. We find $t_{\\rm chem}/t_{\\rm dyn} \\sim 0.7$ for a cloud mass of $M(\\rco) = 1\\times 10^5$ $M_{\\odot}$ and $t_{\\rm chem}/t_{\\rm dyn} \\sim 0.5$ for a cloud mass of $M(\\rco)= 1 \\times 10^6$ $M_{\\odot}$. Thus we expect modest affects due to turbulence; mainly to spread out the transition from atomic to molecular hydrogen but not to move $A_V(\\rht)$. We also note that steady state models agree well with observations.  \\subsection{Dark-Gas Fraction}  For our standard cloud mass, $M(\\rco) = 1\\times 10^6$ $M_\\odot$, and incident radiation field, $G_0' = \\zeta_{\\rm XR}' = 10$, we find dark-gas mass fractions of $f_{\\rm DG} = 0.28$ and $f_{\\rm DG} = 0.31$ for constant thermal pressures of $P_{\\rm th} = 10^5$ K ${\\rm cm^{-3}}$ and $P_{\\rm th} = 10^6$ K ${\\rm cm^{-3}}$, respectively. These correspond to a total molecular mass $M(\\rht) \\approx 1.4 M(\\rco)$. For models that include both thermal and turbulent pressures, we find essentially the same results as for the cases with constant thermal pressure. The variation in $f_{\\rm DG}$ ranges from 0.25 to 0.33 over a range in $G_0'$ from 3 to 30 and a range in cloud mass from $10^5$ $M_\\odot$ to $3\\times 10^6$ $M_\\odot$ (Fig.\\ \\ref{fig:ratio}). Figure \\ref{fig:Preslg} shows the distribution in thermal pressures and densities and Figure \\ref{fig:templg} shows the distribution in temperatures and chemical abundances for the standard model.  The constant value of $f_{\\rm DG}$ for fixed $\\bar A_V$  can be understood from the analytic solutions for $A_V(\\rht)$ and $A_V(\\rco)$ in Appendices \\ref{appen:h2form} and \\ref{appen:coform} and the expression for $f_{\\rm DG}$ in equation (\\ref{eq:fDGAV}). In the limit of $G_0'/n > 0.0075 Z'^{0.43}$ ${\\rm cm^3}$, both $A_V(\\rht)$ and $A_V(\\rco)$ increase as $\\ln(G_0'/n)$. Thus, the optical depth through the ${\\rm H_2}$ layer $\\davdg$ is a weak function of $G_0'/n$ and $Z'$ (eq.\\  \\ref{eq:deltaavmainz}) and is nearly constant over our parameter space. Furthermore, we find that $f_{\\rm DG}$ is a function of only $\\davdg$ and the mean extinction through the cloud $\\bar A_V \\equiv 5.26 Z' \\bar N_{22}$ (Eq. \\ref{eq:fDGAV}); thus, for a given $\\bar A_V$, the dark-gas fraction is constant. However,  $f_{\\rm DG}$ increases significantly if $\\bar A_V$ decreases.   Our numerical results compare well with observations of the local Galactic dark-gas fraction, which \\cite{grenier2005} find to be $f_{\\rm DG} \\sim 0.3$, when averaged over their four most massive clouds with masses between $3 \\times 10^4 - 3\\times 10^5$ M$_\\odot$. Lower mass clouds observed by \\cite{grenier2005} are observed to have low $\\bar A_V$ and thus high dark-gas fractions consistent with our prediction. Our \\ion{H}{1} integrated intensities  of 965 K km ${\\rm s^{-1}}$ to 4100  ${\\rm s^{-1}}$ , \\ion{H}{1} cloud-halo thickness of 1 pc to 10 pc, average cloud densities of $\\bar n \\sim 45-150$ cm$^{-3}$, and average \\ion{H}{1}  temperatures of $\\sim 70-80$ K are in good agreement with the observations of \\ion{H}{1} cloud halos observed by \\cite{wannier1983}, \\cite{andersson1991}, and \\cite{andersson1993}.  We have carried out several additional tests to assess the dependence of our results on the cosmic-ray ionization rate, the clump optical depth, the mean cloud column density,  the metallicity, and the mean visual extinction through the cloud. We ran our standard model with the cosmic-ray ionization rate a factor of 10 higher at $A_V< 2$, as suggested by ${\\rm H_3^+}$ observations in some regions of the diffuse ISM, and found that this enhanced cosmic-ray rate only slightly increases the dark-gas fraction. In the limit of very optically thick clumps we find no significant change in $f_{\\rm DG}$.  \\cite{heyer2009} has suggested that the mean column density (and therefore the mean visual extinction $\\bar A_V$) in Galactic CO clouds is about half the value found by \\cite{solomon1987}, i.e.,  $\\bar{N}_{22} \\simeq 0.75$ or $\\bar A_V \\simeq 4$. The results of changing the mean column density, the metallicity, and the mean visual extinction through the cloud are illustrated in Figures \\ref{fig:fdglz2} and \\ref{fig:avmean}. We find that the dark-gas fraction increases with lower extinction since the dark gas occupies a larger fraction of the cloud (see eq.\\ \\ref{eq:fDGN}). There is also a weak dependence of $\\davdg$ on the mean column density that tends to slightly mitigate the dominant effect. Lower mean columns lead to lower mean densities (eq.\\ \\ref{eq:barnc}) and lower $n_c$ (eqs.\\ \\ref{eq:nmed} and \\ref{eq:ncmax}), and thus lower $\\davdg$ (eq.\\ \\ref{eq:deltaavmainz}). We examine metallicities appropriate for the LMC ($Z' = 0.5$), for the local Galaxy ($Z'=1$), and for the molecular ring at $R=4.5$ kpc ($Z'=1.9$). In general, $f_{\\rm DG}$ increases as the metallicity drops for fixed columns because the mean extinction through the cloud decreases, which raises the ratio of the surface dark gas to the interior CO gas (again, see eq.\\ \\ref{eq:fDGN}). There is also a weak dependence of $\\davdg$ on $Z'$ that also slightly increases $\\davdg$ (or $f_{\\rm DG}$)  with decreasing $Z'$ (eq.\\ \\ref{eq:deltaavmainz}) even if the column is changed so that $\\bar A_V$ remains fixed. We note that in the case of varying $\\bar{N}$ and $Z'$, but at constant $\\bar{N}Z'$ or $\\bar A_V$ (see Figure \\ref{fig:avmean}), the change in $f_{\\rm DG}$ is entirely due to the weak dependencies of $\\davdg$ on $\\bar{N}$ and $Z'$ as noted above and shown in equation (\\ref{eq:deltaavmainz}). We also examine the case for dust scaling as $Z'^2$ ($=0.25$) while gas phase metals scale as $Z'$ ($=0.5$). We find $f_{\\rm DG}$ is larger than when both metals and gas scale together.  In Appendices \\ref{appen:h2form} and \\ref{appen:coform} we derive analytic solutions for $A_V(\\rht)$ and $A_V(\\rco)$ as functions of density $n$, FUV field $G_0'$, and metallicity $Z'$. For the case of ${\\rm H_2}$ we find an expression for the abundance of ${\\rm H_2}$ by balancing formation and destruction processes and then solving for the position where $n_{{\\rm H_2}} = 0.25 n$. Similar studies have been carried out by, for example, \\citet{sternberg1988,mckee2010}. We find our fits are good to $\\pm 5$\\% for all models presented in this paper. The fits to CO are found by integrating the expression for the abundance of CO to a column density of $N({\\rm CO})= 2\\times 10^{16}$ ${\\rm cm^{-2}}$, where  $\\tau_{\\rm CO} = 1$ (the optical depth of the CO J= 1-0 transition). The fits are generally good to within $\\pm 15$\\% except for the lowest $Z'$ and $G_0'$ model, where the fit is good to $\\pm 25$\\%. We have also derived analytic expressions for $\\davdg$ and $f_{\\rm DG}$ in the main text (eqs. \\ref{eq:deltaavmainz} and \\ref{eq:fDGN}, respectively).  The overall result of this paper is a theoretical derivation that $f_{\\rm DG}\\simeq 0.3 \\pm 0.08$ for GMCs with $\\bar A_V \\simeq 8$ in our Galaxy.  Therefore, a significant fraction of the molecular gas in our Galaxy lies outside the CO gas.  As discussed in \\S 2.4.2, some calibrations of the CO line intensities to H$_2$ mass take into account this ``dark\" H$_2$ gas.  However, it is important to be aware of this component since it contributes significantly to the gamma ray, infrared/submillimeter continuum, and [CII] 158 $\\mu$m emission from clouds in galaxies. Its contribution to the star formation in a galaxy is as yet undetermined. The importance of this component increases as the metallicity decreases, such as in the outer regions of galaxies or in low metallicity galaxies.  Although the gas in the C$^+$/${\\rm H_2}$ layer is termed ``dark gas,'' it emits [\\ion{C}{2}] 158 $\\mu$m line emission and the dust in the dark layer emits infrared continuum. In a subsequent paper we will estimate the emission from the ``dark gas''. In addition, a future paper will model in detail the lower metallicity clouds found, for example, in the SMC and early universe."
    },
    "1004/1004.2017_arXiv.txt": {
        "abstract": "{ This paper sumarizes research  on abundances in RR Lyrae stars that one of us (GW) has been engaged in with various astronomers.  In addition we report on preliminary analysis of the abundances of C, Si, S and Fe in 24 RR Lyrae stars.  Our model atmosphere analysis, including NLTE effects, are based on the spectra of resolving power 30,000 obtained at the Apache Poing Observatory. } ",
        "introduction": "This paper reports on research by Sergei M. Andrievsky, Valentin V. Kovtyukh, Marcio Catelan, Dana Casetti-Dinescu, and Gisella Clementini as well as ourselves.  The RR Lyrae stars are of great value in studies of the structure of our Galaxy and the history of its composition as modified by stellar nucleosynthesis over the ages. The fact that they show a narrow range of luminosity from about Mv=0.4 to 0.8, as derived from their membership in globular clusters, allows their distances to be derived and their orbits calculated from their proper motions and radial velocities. Their effective temperatures vary from approximately 6000 to 7000 K during their pulsation. Their spectra often show lines of many elements without excessive blending.  The chemical compostion of RR Lyrae stars is a combination of their original composition when they were main sequence stars near Mv=4.5 with important changes induced by nuclear reactions whose products may be convected to the stellar surface. The first mixing event to affect their atmospheres is the deepening of convection as the star leaves the subgiant branch and begins its evolution up the almost vertical giant branch \\citep{hoy55}. CNO-processing converts some carbon into nitrogen and reduces the 12C/13C ratio from its initial value to about 20. As the star's evolution slows down at the red giant clump additional mixing further reduces the carbon and lowers the 12C/13C ratio to about 4-8 \\citep{gil91,cha94,cha98}. These changes have been seen in many globular clusters.  At the tip of the red giant branch important events take place. The triple-alpha reaction starts in the degenerate core producung additional carbon.  At first the tempature rises exponentially because the pressure and density do not respond until the degeneracy is removed by the temperature rise and the reaction rate depends approximately on the temperature raised to the 30th power.  This sets up a grossly superadiabatic temperature gradient and initiates violent convection \\citep{moc09}.  Calculations indicate that the products of helium burning do not reach the surface but observations of the carbon abundance in RR Lyrae stars are useful to confirm the accuracy of the calculations and to describe the initial conditions for stars that will eventually evolve up the AGB, perhaps to become carbon stars.  In addition mass-loss on the red giant branch plus additional mass-loss at the time of the helium-flash reduces the stellar mass from its original value of about 0.8 Msun to about 0.55 Msun as derived from multi-periodic RR Lyrae stars. Such mass-loss may give us a glimpse into the interior of what had once been a red giant.  Our first project to determine the composition of RR Lyrae stars was suggested by M. Catelan who noted that a few RR Lyrae and related stars showed kinematics similar to that of Omega Cen. The first is VY Ser, a genuine RR Lyrae star of period 0.714 days, while 2 more, V716 Oph and XX Vir have periods slightly longer than 1.0 days and are usually called short-period type II cepheids, but might be referred to as long-period RR Lyrae stars. Dana Casetti-Dinescu confirmed that their galactic orbits do indeed relate them to Omega Cen. The period of VY Ser is typical of RR Lyrae stars in Omega Cen, but only 7 short-period cepheids are known in Omega Cen \\citep{cle01}.  In fact it is remarkable that astronomers have defined the break between RR Lyraes and short period cepheid to equal to the period of rotation of the earth. Our observations consisted of 4 echelle spectra of VY Ser, 5 of V716 Oph, and 2 of XX Vir that have been analysed by Andrievsky and Kovtyukh. All 3 stars show [Fe/H] close to $-1.6$ which is typical of Omega Cen but cannot be used to reveal the strong gradient of [s/Fe] when plotted against [Fe/H] in Omega Cen \\citep{van94, nor95}.  Our second group of targets was RR Lyrae stars with period greater than 0.75 days. Such stars are very rare in globular clusters with the conspicuous exception of the two unusual clusters, NGC 6388 and 6441 \\citep{pri01, pri02}.  These relatively metal rich clusters with [Fe/H] near $-0.8$ have numerous long-period RR Lyrae stars. In the general field such stars are rare but some are known and their light curves have been obtained by \\citet{sch02}.  Only a few are sufficiently bright for highres spectroscopy with the 3.5-m telescope of the Apache Point Observatory. They are listed in \\citet{wal09}.  [Fe/H] values were found for 4 stars but their metallicities ranged from $-1.8$ to $+0.2$. Hence our analyses of these stars failed to relate them to NGC 6388 and 6441. KP Cyg remains an almost unique RR Lyrae star with a period of 0.856 days and [Fe/H] $= +0.2$ on the basis of 5 spectra well distributed in phase. It had already been recognized as having delta-S $= 0$ by \\citet{pre59}. Our spectra show a typical velocty curve for an RR Lyrae star with H-alpha emission at phases 0.43, 0.72, and 0.80.  The third aspect of our RR Lyrae analyses was the recognition that the relatively metal-rich stars show an apparant excess of carbon \\citep{wal09}).  The carbon abundance was derived from the lines around 7115\\AA~of multiplets 108 and 109 of \\citet{wie98}.  These lines are  sufficiently weak to be found only in stars with [Fe/H] $> -1.0$. For metal-poor stars stronger lines must be employed which suffer from NLTE \\citep{fab06}. Fabbian has told us (private communication) that the lines around 7115\\AA~should not suffer from significant NLTE effects. The best strong lines are one at 8335\\AA~and from multiplet 62 beyond 9000\\AA.  That spectral region contains many moderately strong atmospheric H$_{\\rm 2}$O lines. ",
        "conclusions": ""
    },
    "1004/1004.3820_arXiv.txt": {
        "abstract": "We present a high-resolution set of adiabatic binary galaxy cluster merger simulations using FLASH. These are the highest-resolution simulations to date of such mergers using an AMR grid-based code with Eulerian hydrodynamics. In this first paper in a series we investigate the effects of merging on the entropy of the hot intracluster gas, specifically with regard to the ability of merging to heat and disrupt cluster ``cool-cores.'' We find, in line with recent works, that the effect of fluid instabilities that are well-resolved in grid-based codes is to significantly mix the gases of the two clusters and to significantly increase the entropy of the gas of the final merger remnant. This result is characteristic of mergers over a range of initial mass ratio and impact parameter. In line with this, we find that the kinetic energy associated with random motions is higher in our merger remnants which have high entropy floors, indicating the motions have efficiently mixed the gas and heated the cluster core with gas of initially high entropy. We examine the implications of this result for the maintenance of high entropy floors in the centers of galaxy clusters and the derivation of the properties of dark matter from the thermal properties of the X-ray emitting gas. ",
        "introduction": "\\label{sec:intro}  The growth of structure in the universe is thought to proceed in a bottom-up fashion, with smaller structures merging into larger structures \\citep[e.g.,][]{whi93}. The largest structures that have formed to date via this process are clusters of galaxies. Clusters of galaxies are important astrophysical objects, both in their own right and for their use as cosmological probes. Galaxy clusters contain a good representation of the different forms and kinds of matter within the universe. The bulk of the mass in clusters of galaxies is comprised of dark matter \\citep{zwy37,bah77}, which is believed to be largely collisionless. The bulk of the baryonic material in clusters of galaxies is comprised of a hot diffuse plasma called the intracluster medium (ICM) \\citep[e.g.,][]{sar88}. The remaining (and smallest) component of mass is that of the galaxies themselves. Galaxy clusters provide an opportunity to witness the interplay between these different kinds and forms of matter in close quarters. Additionally, clusters have become very useful as probes of cosmology. The number density of clusters of a given mass at a given redshift is a sensitive function of the parameters of the standard big bang cosmological model. Accurate measurements of the density and temperature structure of galaxy clusters enable determination of cluster masses (under the assumptions of spherical symmetry and hydrostatic equilibrium) which in turn helps to constrain the values of cosmological parameters \\citep{kita96,voit05}.  Extensive study of both of these aspects of astrophysics with galaxy clusters requires state-of-the-art observations and simulations. The latest generation of X-ray telescopes ({\\it Chandra}, {\\it XMM-Newton}) has provided high-resolution observations of the ICM that have made it possible to probe sensitively the density, temperature, and metallicity structure of the gas. These observations have made possible the discovery of a myriad of features in the ICM such as shocks, cold fronts, and bubbles, to name a few. In addition, such precise measurements have made hydrostatic estimates of cluster masses a useful tool for constraining cosmological parameters on a par with other methods \\citep{vik09b}. In tandem with these advances in observation, astrophysical simulation codes have taken advantage of large-scale computing resources and innovative algorithms to model clusters of galaxies with unprecedented resolution and with the ability to include detailed models of the relevant physical processes occurring in clusters.  One of the most fundamental of these physical processes that shapes the life of clusters is merging. The merging process is responsible for forming the clusters and greatly influences the state of the matter within a cluster throughout its lifetime. Mergers drive shocks into the intracluster medium, heating and compressing the gas and driving turbulent motions. As a result of merging the kinetic properties of the dark matter are also significantly altered. Mergers also drive galaxy activity, producing new bursts of star formation as the gas is compressed, or in some cases possibly inhibiting it \\citep{fuj99,bek03}. Many observed clusters of galaxies show signs of current or recent merging. Two prominent examples of recent note are 1E~0657-56 (the ``Bullet Cluster'') \\citep{mar02,clo06}, and Abell 520 \\citep{mar05,mah07}. Explaining the myriad of observable consequences of mergers using simulations is an active area of research.  From the perspective of simulations, there are two ways to study the merging process. The first is by studying mergers that occur in the context of a cosmological simulation, ``out in the wild.'' The advantage of this approach is that it captures as accurately as possible the merging process as it occurs in the real universe, with many mergers of many different mass ratios and impact parameters going on at the same time, in the context of the overall cosmological expansion. The second computational approach to studying mergers is by focusing on particular merger scenarios in isolation via simulations of idealized cluster mergers. These simulations involve the setup of two or more idealized, spherically symmetric clusters in a localized setting and with an initial orbital trajectory. Though this method is useful for studying individual actual mergers as seen in observations \\citep[e.g.][]{tak06,spr07,aka08,mas08,ran08,zuh09a,zuh09b}, an additional application has been to explore a space over parameters such as the mass ratio and initial impact parameter of the progenitor clusters or by varying the physics of the cluster ICM or dark matter.  Several previous investigations of merger parameter spaces \\citep[e.g.][]{roe97,ric01,gom02,rit02,mcc07,poo06,poo07,poo08,tak10} have emphasized the effects of merging on cluster morphology, the thermodynamics of the gas, and observables. We build on this work with a new set of simulations. The set of simulations presented here is the highest-resolution parameter study of idealized binary cluster mergers performed with an AMR-based PPM code to date, with a resolution 3-4 times higher than the most recent investigation \\citep{ric01} based on such a formalism.  In this first of two papers, we use these simulations to examine the effect of merging on the entropy of the cluster gas. Observational surveys of clusters have shown that although in nearly all clusters the entropy (defined here as $S \\equiv k_BTn_e^{-2/3}$) follows a power-law profile with radius, the central core entropies of clusters vary \\citep{llo00,pon03,voit05c,fal07,mor07,cav09,pra09}. These observed entropy profiles are well-fit by an entropy profile of the form \\begin{equation} S(r) = S_0 + S_1{\\left(r \\over {0.1r_{200}}\\right)}^\\alpha \\label{eqn:entropy} \\end{equation} Where $S_0$ is the value for the core entropy and $\\alpha \\sim 1.0-1.3$. In particular, \\citet{cav09} and \\citet{pra09} showed that clusters exhibit a bimodal distribution of central core entropies. \\citet{cav09} found a low-entropy peak at $S_0 \\sim 15~{\\rm keV~cm}^2$, a high-entropy peak at $S_0 \\sim 150~{\\rm keV~cm}^2$, and very few clusters with core entropies $S_0 \\sim 30-50~{\\rm keV~cm}^2$. Similarly, \\citet{pra09} found peaks at $S_0 \\sim 3$ and $S_0 \\sim 75~{\\rm keV~cm}^2$. These core entropies are higher than what would be expected from a simple ``cooling flow'' scenario \\citep[e.g.][]{fab77,fab94}, where the core gas is allowed to cool and compress unabated until it turns into stars or molecular clouds. High-resolution {\\it XMM-Newton}\\/ spectroscopy \\citep{pet01,pet03} and {\\it Chandra}\\/ spectral imaging \\citep{dav01} have shown that there is indeed little gas below $T\\simeq 1$ keV in the cores of clusters with some of the highest cooling rates.  Since cooling via X-ray radiation is directly observed, but the expected central entropies from unabated cooling is not, clusters with high entropy floors require a heating mechanism.  Proposed candidates for heating of cluster cores include magnetic field reconnection \\citep{sok90}, thermal conduction due to electron collisions (e.g., Narayan \\& Medvedev 2001; Fabian, Voigt, \\& Morris 2002; Zakamska \\& Narayan 2003) and turbulent conduction (e.g., Cho et al.\\ 2003; Voigt \\& Fabian 2004), and heating by cosmic rays (e.g., Inoue \\& Sasaki 2001, Colafrancesco \\& Marchegiani 2008); a recent review can be found in \\citet{petersonfabian06}. The currently favored mechanism is heating by the central AGN (e.g., B\\\"ohringer et al.\\ 1993; Binney \\& Tabor 1995; McNamara et al.\\ 2001, 2005; Fabian et al.\\ 2006; Forman et al.\\ 2007; for a recent review see McNamara \\& Nulsen 2007). The AGN explosions blow the bubbles in the ICM and inject energy into the ICM in the form of relativistic particles as well as mechanical energy.  However, the precise mechanism by which this energy heats the central ICM is still unclear. A fine balance between AGN explosions and cooling is required to avoid the complete blow-up of the cool cores, which gave rise to ``feedback'' models, where the cooling flow itself feeds the AGN.  Thermal conduction is another particularly attractive idea, because it taps the vast reservoir of thermal energy in the gas just outside the cool core, while automatically ensuring that the core will not be overheated, since the heat influx decreases with diminishing temperature gradient. The classic plasma conductivity via Coulomb collisions was shown to be insufficient even at its full Spitzer value (e.g., Zakamska \\& Narayan 2003). It has a strong temperature dependence and decreases right where it is most needed, and tangled magnetic fields should further suppress it (as was indeed observed outside the cool cores, Markevitch et al.\\ 2003a). Additionally, magnetohydrodynamic instabilities such as the heat-flux-driven buoyancy instability (HBI) may suppress heat conduction to cluster cores from the cluster outskirts altogether by aligning magnetic field lines perpendicular to the radial temperature gradients in clusters \\citep{qua08, par08, bog09, par09}, though recent works have shown that modest amounts of turbulence in cluster cores may work to keep the fields randomized and sustain some conduction of heat to the cluster core \\citep{rus10, par10}. Though the effects of AGN and conduction may account for the lack of systems with extremely low central cooling times and entropies, it is also unclear whether or not either effect can explain the highest entropy floors seen in some clusters, which can reach central entropy values of $S_0 \\sim 100$ or higher.  Previous investigations of cluster mergers \\citep[e.g.,][]{rit02, mcc07, poo08} have indicated that so-called cluster ``cool cores'' (characterized by temperature inversions, high gas densities, and low central entropies and cooling times) are largely resilient to the heating that occurs as a result of merging. The central entropies and cooling times increase for a short time after the core passage, but the heating is not enough to offset cooling, which quickly reestablishes a cool core within a few Gyr \\citep{poo08}. \\citet{mcc07}, in a suite of adiabatic idealized merger simulations, determined that the ICM is heated in two major stages: the first is when the shocks are generated at the initial core passage and the second is at a later stage when material is shock-heated as it recollapses onto the merger remnant. Though the overall entropy of the merger remnant was increased with respect to the progenitor systems consistent with the self-similar prediction, the power-law behavior of the entropy profile was maintained all the way down into the central core (see Figure 5 of that paper).  It is notable that these SPH merger simulations also resulted in little mixing of the cluster gases, since such mixing could provide a source of heating to cluster cores if the gases are of very different entropies.  It is known that the various methods for solving the equations of hydrodynamics result in different degrees of mixing. While previous investigations of mixing due to merging \\citep[e.g.,][]{rit02, poo08} have employed the Lagrangian ``smoothed-particle hydrodynamics'' (SPH) formulation for solving the Euler equations, our simulations have been performed using an Eulerian PPM formulation. The two different approaches to hydrodynamics differ in their ability to appropriately model turbulence and the onset of fluid instabilities \\citep[see, e.g.][]{dol05,age07,wad08}. The amount of mixing of the gas is also dependent on the precise nature of the physics of the ICM, as the presence of magnetic fields, turbulence, and dissipative processes will affect the mixing of the gas. The extent to which the ICM mixes as a result of mergers and its effect on the thermodynamic properties of the cluster gas is therefore dependent the physical processes operating in the ICM as well as the chosen algorithm for solving the hydrodynamic equations. Therefore, it is important as a baseline to characterize the mixing process and the effects of such mixing on the thermodynamics of the cluster gas in the simplest model for ICM, that of an inviscid, unmagnetized gas, in the context of an Eulerian, PPM-based hydrodynamics code. Doing so is the goal of this paper.  This paper is organized as follows. In Section \\ref{sec:sims} we describe the simulation method, setup, and rationale for the initial conditions. In Section \\ref{sec:results} we present the results of the simulations, focusing on the degree of mixing of the intracluster gas components and the alteration of the thermodynamic state of the cluster gas as a result. In Section \\ref{sec:disc} we discuss the implications of these results for different approaches to hydrodynamics for galaxy cluster simulations, the survival of cool cores in cluster mergers, and the use of X-ray observables to determine properties of the dark matter. In Section \\ref{sec:conc} we summarize this work and suggest avenues for future investigations. Throughout this paper we assume a spatially flat $\\Lambda$CDM cosmology with $h = 0.7$, $\\Omega_{\\rm m} = 0.3$, and $\\Omega_b = 0.02h^{-2}$. ",
        "conclusions": "\\label{sec:conc}  We have performed a set of fiducial binary galaxy cluster merger simulations that explore a parameter space over initial mass ratio and impact parameter. In this first of a series of papers we focus on the mixing of the intracluster medium between the two clusters and the resulting effect on the thermodynamics of the gas of the merger remnant. In contrast to previous merger parameter space studies but in agreement with recent works focusing on mixing, we find that the gas of the merging clusters mixes very efficiently. The degree of mixing is dependent upon the mass ratio and impact parameter of the cluster merger, but all of the mergers investigated here demonstrate a significant amount of mixing in some regions of the final merger remnant. This is a result of ubiquitous presence of fluid instabilities that occur as the merger generates large shear flows and sharp density gradients, which are resolved in grid-based fluid codes such as the one used in this study but not in particle-based fluid codes used in many previous works. These instabilities result in the mixing of high and low-entropy gas and as a result the thermodynamic state of the primary cluster's core is dramatically affected. In each case the cluster has been changed from one with a high-density, low-entropy, low-temperature core with a short cooling time to a low-density, high-temperature, high-entropy core with a long cooling time. This is the case even for mergers with smaller subclusters. Further evidence that mixing is the culprit is indicated by the correlation between final merged systems with high central gas velocity dispersions and high central entropies, as the greater random motions in these systems allow for more efficient mixing. This result may also pose a difficulty for attempts to derive the kinetic properties of the dark matter from the thermodynamic properties of the X-ray emitting gas, which rely on the hypothesis of the near-equality of the gas temperature and dark matter velocity dispersion profiles. Additionally, such efficiency in heating the cluster core may be a primary factor in the generation of a population of galaxy clusters with high entropy floors and cooling times. Further simulation works including more realistic physics for the ICM (on the one hand, dissipative effects, magnetic fields, which would suppress mixing, and on the other hand, a more accurate characterization of the turbulence via higher resolution or a turbulent subgrid model which would enhance it) are needed to confirm this hypothesis. In spite of this, the main result of this work is that merging is still a serious candidate for the generation of heat to balance against cooling in the cores of galaxy clusters."
    },
    "1004/1004.1963_arXiv.txt": {
        "abstract": "{}{}{}{}{} \\abstract {High-resolution spectroscopy has recently revealed in many low-mass X-ray binaries hosting a neutron  star that the shape of the broad iron  line observed  in  the 6.4-6.97  keV  range is  consistently well-fitted by a relativistically smeared line profile.} {The presence of other broad features, besides the iron line, together with a  high S/N of the  spectra offer the possibility  of testing a self-consistent  approach   to  the  overall   broadband  reflection spectrum and  evaluating the impact  of the reflection  component in the formation of the broadband X-ray spectra.} {We analyzed two XMM-Newton observations of the bright atoll source \\object{4U 1705-44}, which  can be  considered a prototype  of the class  of the persistent NS  LMXBs showing both  hard and soft states.   The first observation was  performed when  the source was  in a hard  low flux state, the second  during a soft, high-flux state.  Both the spectra show  broad  iron  emission  lines.   We fit  the  spectra  using  a two-component model,  together with a  reflection model specifically suited to  the case of a  neutron star, where  the incident spectrum has a blackbody shape.} {In  the   soft  state,  the   reflection  model,  convolved   with  a relativistic  smearing component,  consistently describes  the broad features  present in  the spectrum,  and  we find  a clear  relation between the temperature of the  incident flux and the temperature of the harder X-ray  component that we interpret as  the boundary layer emission.   In  this state  we  find  converging  evidence that  the boundary  layer outer  radius is  $\\sim$  2 times  the neutron  star radius. In the low flux state,  we observe a change in the continuum shape of  the spectrum with respect  to the soft  state.  Still, the broad  local  emission  features  can  be  associated  with  a  disk reflecting  matter, but in  a lower  ionization state,  and possibly produced in an accretion disk  truncated at greater distance. } {Our analysis provides  strong evidence that the  reflection component in soft states of  LMXBs comes from to hard  X-ray thermal irradiation, which we identify with the  boundary layer emission, also present in the continuum model. In the hard  state, the broad iron line if also produced  by reflection,  and  the continuum  disk  emission can  be self-consistently accounted  if the disk  is truncated at  a greater distance than the soft state.} ",
        "introduction": "In  last few  years,  the Epic-pn  instrument  onboard the  XMM-Newton satellite has  allowed deep investigation  of the nature of  the broad emission  lines  observed  in  the  iron K$\\alpha$  region  of  bright accreting   neutron  star   (NS)  low-mass   X-ray   binaries  (LMXBs) \\citep{bhattacharyya07,   pandel08,   cackett08,   dai09,   papitto09, iaria09, disalvo09}.  The authors of these works  have focused their attention on  the shape and  origin of the  broad iron line,  and they agreed on the interpretation of the line broadness being the result of to  special  and general  relativistic  effects  arising  in the  disk reflecting  matter  at a  few  gravitational  radii  from the  compact object.  This interpretation  is supported by theoretical expectations and by  general agreement  between the fitting  model and  X-ray (1-10 keV) data.  However,  if this  interpretation is  correct, the  reflected spectrum should encompass a variety  of other disk reflection features, because there  are other  low-Z,  but  abundant, emitting  ion  metals in  the low-energy band.  The good spectral resolution and the high S/N of the spectra has effectively shown in  many cases a more complex pattern of features,  besides the broad  iron line.   \\citet{dai09} found  in the spectrum  of the bright  Z-source GX  340+0 a  broad emission  line of \\ion{Ca}{xix}  and   an  absorbing   edge  of  highly   ionized  iron; \\citet{iaria09} found in  the spectrum of the Z-source  GX 349+2 three broad  lines,  besides the  iron  line,  identified as  \\ion{Ca}{xix}, \\ion{Ar}{xviii} and  a blending  of L-shell transitions  of moderately ionized  iron.  \\citet{disalvo09}  identified in  the bright  atoll 4U 1705-44 resonant emission  lines of \\ion{Ca}{xix}, \\ion{Ar}{xviii} and \\ion{S}{xvi}, and an \\ion{Fe}{xxv}  iron edge, which appears broad and redshifted  with respect to  the expected  rest-frame energy.   It has been therefore  suggested that  all these features,  and not  only the broad iron  line, originate in  the reflection component.   To support this interpretation, it has been shown that the smearing components of the broad iron  line (i.e., the smearing parameters  of the reflection component, the inner and outer radii, inclination angle of the system, and  the  emissivity  index,  which  measures the  dependence  of  the emissivity power-law  profile from the  distance to the source  of the irradiating photons) consistently describe the shape and the broadness of all the observed emission lines \\citep{dai09,disalvo09,iaria09}.  In  this  work, we  perform  a  detailed  analysis of  two  XMM-Newton observations  of the  bright atoll  source 4U  1705-44.  In  the first observation, the  source was in a  low flux state, while  in the second the  source  was  in a  bright  soft  state.  First results,  using  a phenomenological  approach to  model the  reflection component  of the latter observation  have been presented  in \\citet{disalvo09}.  Here, we  focus on  the spectral  changes  that occurred  between these  two observations  using a self-consistent  reflection model.  The  angular  dependence  of  the  reflected  component  formed  by  a constant-density  partially-ionized medium  was originally  studied by \\citet{zycki94},  \\citet{zycki94b}, and  \\citet{matt03}.  Reprocessing in  a   medium  in  hydrostatic   equilibrium  was  then   modeled  by \\citet{raymond93},      \\citet{nayakshin00},     \\citet{ballantyne01}, \\citet{ballantyne02},     and     \\citet{rozanska02}.     Furthermore, \\citet{nayakshin02}  examined the photoionized  accretion discs  via a novel time-dependent  approach. A  hot layer forms  at the top  of the disk atmosphere, roughly at  the inverse Compton temperature, followed by a steep transition to colder, less ionized layers.  We     use    the    reflection     table    model     described    in \\citet{ballantyne04}.  First  application   of  this  table  model  to consistently fit  the X-ray  spectrum of  an NS LMXB  can be  found in \\citet{ballantyne04b}.      This    reflection     model    (hereafter \\texttt{refbb})  is  calculated  for  an optically  thick  atmosphere, irradiated   by   a  blackbody   incident   spectrum  of   $kT_{ion}$ temperature.  The model gives  the reflected spectrum according to the ionization parameter  $\\xi$, and the  relative abundance of  iron with respect to the other metals. In addition to fully-ionized species, the following     ions    are     included     in    the     calculations: \\ion{C}{iii}-\\ion{C}{vi},                    \\ion{N}{iii}-\\ion{N}{vii}, \\ion{O}{iii}-\\ion{O}{viii},                  \\ion{Ne}{iii}-\\ion{Ne}{x}, \\ion{Mg}{iii}-\\ion{Mg}{xii},                \\ion{Si}{iv}-\\ion{Si}{xiv}, \\ion{S}{iv}-\\ion{S}{xvi},    and   \\ion{Fe}{vi}-\\ion{Fe}{xxvi}.    The ionization parameter,  log $\\xi$, can vary  between 1 and  4, with the density n$_H$ of the illuminated slab constant at 10$^{18}$ cm$^{-3}$. The  constant density prescription  can be  considered a  good diluted approximation of the actual hydrostatic  structure in the disk for the 1.0-10.0 keV energy range \\citep{ballantyne01}.  The $kT_{ion}$ temperature can vary between  1 keV and 5 keV.  All the metals abundances  are fixed at the  solar value, except  for iron for which models  were calculated for 0.1,  0.3, 1.0, 3, and  10 times the solar value.  The space parameters  of this model covers, therefore, a wide range of possible spectral solutions.  In particular, this is one of the few  available reflection models in which  the primary incident spectrum is  a soft thermal  spectrum, and, also  on the basis  of the results already  shown in \\citet{disalvo09},  it is the  most suitable for fitting the soft states of NS LMXBs.  We  show that,  within  the available  energy  range, the  continuum emission  can  be  simply  accounted  for  a  three-component  model composed   of  thermal   disk   emission,  a   saturated/unsaturated Comptonized harder  emission, and  a reflection component.  The last one  arises from  the disk  reflecting matter,  where  the impinging radiation  field \\emph{is} the  hard X-ray  emission.  We  study the chemical  abundances of  the reflecting  disk matter,  the accretion flow in  the two states,  and possible scenarios for  explaining the spectral differences in the two states.  \\subsection{\\object{4U 1705-44}}  The source  is a  persistently bright, accreting  LMXB located  in the direction of  the Galactic bulge  \\citep{forman78}, hosting an  NS. It shows  type-I X-ray  bursts, with  recurrence times  dependent  on the accretion state \\citep{langmeier87, gottwald89, galloway08}.  From the peak  luminosity of  bursts  that exhibited  episodes of  photospheric radius  expansions,  which are  thought  to  happen  at the  Eddington luminosity   and   are,   therefore,   used   as   standard   candles, \\citet{haberl95} derived a  distance of 7.4$_{-1.1}^{+0.8}$ kpc, later confirmed by  \\citet{galloway08}.  The companion source  has still not been identified,  although a near-infrared  counterpart, most probably originated by  X-ray reprocessing by  the outer accretion  disk and/or the companion star, has recently been found by \\citet{homan09}.  \\object{4U 1705-44} shows a secular trend toward alternating high- and low-flux periods (see e.g., Fig.\\ref{asm_countrate}) on a variable timescale of months.   The  spectral variability,  on  the  contrary,  can be  much faster, of days \\citep{barret02}.  Classified by \\citet{hasinger89} as an Atoll  source, it  was later shown  that its spectral  and temporal states are  intermediate between the  classic Atoll and  the Z-sources division \\citep{barret02}.   In particular, broadband  X-ray data have shown that  the switch between hard  and soft states  can be explained using a two-component model of a Comptonized inner emission and a soft thermal emission. \\citet{barret02} and \\citet{olive03} interpreted the alternate  hard and  soft  state transitions  as  caused by  different truncation radii of  the accretion disk.  During hard  states the disk is truncated  at a large  distance from the  compact object and  a hot corona  with high electron  temperatures and  low optical  depth forms around the NS.   During the hard to soft  state transitions, the inner disk  approaches the NS;  this causes  an increase  in its  flux, thus providing a more  efficient Compton cooling for the  hot electrons and softer spectra.   This scenario is supported by  spectral modeling and by  the  timing  analysis  of  the  power-density  spectra  where  the characteristic  frequencies  of  the  band-limited noise  and  of  the low-frequency  noise  components  are  strongly  correlated  with  the position  of  the  source  in the  hardness-intensity  diagram.   High frequency  quasi-periodic  oscillations   (QPOs)  are  also  observed, usually in pairs (so-called kiloHerzt QPOs, kHzQPOs), with the highest reported peak at 1160 Hz \\citep{wijnands98, ford98}.  Spectral analysis with  the Chandra high-resolution gratings revealed, superimposed  to the  continuum, a  set  of local  features, the  most prominent of which was an unambigous, intrinsically broad (FWHM $\\sim$ 1.2 keV) fluorescent iron line \\citep{disalvo05}.  However, it was not possible to  distinguish among different broadening  mechanisms on the basis of the goodness of the spectral fit. This motivated the need for new observations with the  XMM-Newton satellite, given the much larger collecting  area in  the iron  region of  the Epic-pn  CCDs.   A first XMM-Newton observation caught  the source in a low  state, and the S/N was rather poor in the iron range.  A second observation, performed as a target of  opportunity when the source returned  to a high-intensity soft period  was successful  in disclosing the  asymmetry in  the iron line  shape,  which \\citet{disalvo09}  interpreted  as  the result  of reflection on  a disk surface, very  close to the NS,  of hard coronal photons.  A similar scenario has also been proposed for BeppoSAX broad band  data in  \\citet{piraino07}, and  a  claim was  also made,  using INTEGRAL  high  energy data,  of  a signature  of  a  Compton bump  in \\citet{fiocchi07}.  Recently,  \\citet{reis09} using broad  band SUZAKU data also shows  that the asymmetry of the  iron profile are naturally described by a disk reflection scenario. ",
        "conclusions": "We have examined and compared two XMM-Newton observations of the Atoll source 4U  1705-44, in a soft and  in a hard spectral  state. The soft state  is   characterized  by   three  main  spectral   components:  a multicolored thermal disk emission, a harder, blackbody-like, boundary layer emission, and a relativistically smeared reflection component. A self-consistent model  of reflection has been  successfully applied to fit  the  data, where  the  incident spectrum  is  assumed  to have  a blackbody shape.  We have shown  that the temperature of  this thermal irradiating  flux  is  consistent  with  the  thermal  boundary  layer temperature and found from independent constraints that the inner disk radius, coincident with the outer  boundary layer radius is located at a distance of  $\\sim$ 2 R$_{NS}$.  The second  observation, taken when the source  is in  the hard  state, is also  characterized by  a broad emission  feature in  the  iron range,  although  its physical  origin appears less constrained because of the much lower statistics and lack of high-energy response.  We applied different modelizations, and with the constraints obtained  in the analysis of the  soft observation, we showed  and  discussed three  possible  spectral decompositions.   The scenario  that appears  more successful  in  data fitting  and in  its physical implications is  a model of very soft  thermal disk emission, with an inner edge truncated  at relatively large distance from the NS (between 20  R$_g$ and 90  R$_g$) and a thermal  Comptonized emission. The  width of  the iron  line in  this state,  can still  be explained within  the  reflection  scenario,  even with  a  truncated  accretion disk. Although the  line is intrinsically broad, it  does not show any evidence of asymmetry, because of the lower statistics with respect to the  soft  state  observation  and because  relativistic  effects  are strongly reduced if  the disk is truncated at  a greater distance from the NS. To address the nature  of the reflection component in the hard states,  we still  need good  spectral coverage  of the  overall X-ray emission, in the soft (less than  1 keV) range to better constrain the continuum disk emission, in the iron range to resolve the shape of the iron  line, and  in the  hard (above  10 keV)  range to  constrain the reflection    Compton-scattered   continuum.  \\vspace{0.5cm}  During the peer review process of the present manuscript, we become aware of a similar and related work, appeared as an e-print (arXiv:0908.1098), by \\citet{cackett09}."
    },
    "1004/1004.4893_arXiv.txt": {
        "abstract": "We present the results of an intensive photometric and spectroscopic monitoring campaign of the WN4 Wolf-Rayet (WR) star WR\\,1=HD~4004. Our broadband $V$ photometry covering a timespan of 91~days shows variability with a period of P=16.9$^{+0.6}_{-0.3}$~days. The same period is also found in our spectral data.  The light-curve is non-sinusoidal with hints of a gradual change in its shape as a function of time. The photometric variations nevertheless remain coherent over several cycles and we estimate that the coherence timescale of the light-curve is of the order of 60 days. The spectroscopy shows large-scale line-profile variability which can be interpreted as excess emission peaks moving from one side of the profile to the other on a timescale of several days.  Although we cannot unequivocally exclude the unlikely possibility that WR\\,1 is a binary, we propose that the nature of the variability we have found strongly suggests that it is due to the presence in the wind of the WR star of large-scale structures, most likely Co-rotating Interaction Regions (CIRs), which are predicted to arise in inherently unstable radiatively driven winds when they are perturbed at their base. We also suggest that variability observed in WR\\,6, WR\\,134 and WR\\,137 is of the same nature. Finally, assuming that the period of CIRs is related to the rotational period, we estimate the rotation rate of the four stars for which sufficient monitoring has been carried out; i.e. v$_{rot}$=6.5, 40, 70 and 275~km~s$^{-1}$ for WR\\,1, WR\\,6, WR\\,134 and WR\\,137, respectively. ",
        "introduction": "The winds of Wolf-Rayet (WR) stars are well-known to be non-uniform on small physical scales. Because of the inherently unstable nature of radiatively-driven hot stellar winds \\citep{OCR}, instabilities, which reveal themselves in the spectrum as narrow excess emission peaks superposed on the broad emission lines, appear stochastically, propagate in the wind and disappear after several hours \\citep[e.g. ][]{Moffat88}. This clumpiness leads to emission lines that are variable at a level of up to $\\sim$~5~\\% of the line flux \\citep{Stlouis}. These changes are, of course, random and no periodicity is expected.  In certain cases, WR stars also display large-scale line-profile variability (lpv). For some stars this spectroscopic variability has been shown to be periodic and to originate from a different physical processes. Of course, known massive WR+O binaries produce clear periodic radial-velocity (RV) variations from the orbital motion of the stars. It has also been realized in the past two decades that in such systems the winds from the stars collide, forming a shock cone that wraps around the star with the smaller momentum flux (generally the O star). This interaction induces kinematically-characteristic variations in the line {\\it profiles} that are also periodic.  However, large-scale lpv has also been found in WR stars that are not known to be binaries. The two most studied cases, WR\\,6 (WN4) and WR\\,134 (WN6), have been monitored intensively in photometry and spectroscopy by \\citet{Morel97,Morel98,Morel99a}. In both cases, these authors have observed periodic large-scale lpv and a complex light-curve with the same periodicity. To that list of two stars, we may also add WR\\,137, a WC7pd star in a long-period binary system. This star shows large-scale lpv with a period of 1.2~d that is unlikely to be related to its O9 companion nor to the wind-wind collision zone \\citep{Lef}. This behaviour is most likely explained by the presence of large-scale density structures in the wind, such as co-rotating interaction regions (CIRs) \\citep{Cranmer96,Fullerton97}, but the possibility of the presence of a compact companion or of a low-mass main-sequence star cannot be completely excluded.  If the large-scale lpv and the photometric variability of single WR stars can be associated with CIRs, it may imply that the period of these variabilities corresponds directly to the rotational period of the star. Indeed, \\citet{Cranmer96} propose that ``spots'' fixed to the stellar surface, caused either by pulsations or magnetic field activity, are at the origin of the CIRs. Hence, once the radius at which the CIR originates is known, we can determine the rotational velocity of the star at that point. It follows that the rotation rate of at least some WR stars could in principle be determined by carrying out a systematic investigation of the variability of all single WR stars. As a first step, \\citet{Stlouis} and Chen\\'e \\& St-Louis (in prep) set out to identify new candidates for CIR-type variability by conducting a survey of all apparently single Galactic WR stars brighter than $v\\sim$12.5$\\rm ^{th}$ magnitude. For each star in their sample, they obtained 4--5 spectra which allowed them to establish a list of WR stars showing large-scale lpv. The next step is to observe intensively each of the candidates in order to verify and determine the periodicity.  In the above-mentioned survey, WR\\,1 was one of the most striking cases of large-scale lpv with changes reaching 8--10\\% of the line flux and easily distinguishable large scale subpeaks superposed on the broad wind emission profiles. Consequently, it was the first WR star on which we concentrated our efforts. Several previous studies claiming unreproducible periods for the variability of this star have been summarized in \\citet{Morel99b}. Later, \\citet{Nieda,Niedb} demonstrated that the spectrum of WR\\,1 varied greatly from night-to-night reaching 50\\% in the equivalent width of the He{\\sc i}$\\lambda$5876 line. On the other hand, they found that line-profile changes during one night are much smaller with a typical scatter of 0.5\\,\\AA\\, in equivalent width for the He{\\sc ii}$\\lambda$5411, C{\\sc iv}$\\lambda$5808 and He{\\sc i}$\\lambda$5876 lines. The changes, however, were found to be systematic during the course of an entire night and reach a total of 3-4\\,\\AA. Moreover, the search for periods smaller than 2 days failed and only indications of long-term variability could be suggested. The investigation of photometric variability over more than 16~days carried out by \\citet{Morel99b} did not lead to the identification of a period. Consequently, the period search must be done using data taken over a time range of at least twice as long. More recently, \\citet{Flor} claimed a period of P=7.684~days in lpv of this star which they attribute to the ejecta of streams or jets from the stellar surface.  In this paper, we present the results of an intensive monitoring campaign of WR\\,1 extending over several weeks in photometry and spectroscopy with the aim of determining the nature of the variability and, eventually, if found to be associated with CIRs as suspected from our survey observations, the rotation period of WR\\,1. In Section~\\ref{obs}, we present our photometric and spectroscopic observations and the data reduction procedures. In Sections~\\ref{resPhot} and \\ref{ressp} we describe our results, and in Section~\\ref{dis} we discuss the possible interpretations. Finally, our conclusions are presented in Section~\\ref{con}. ",
        "conclusions": "\\label{dis}  \\citet{Stlouis} have already shown that WR\\,1's spectrum shows large-scale lpv with a similar amplitude to what is observed for WR\\,6 and WR\\,134, two presumably single WR stars showing periodic photometric and spectroscopic variability \\citep{Morel97,Morel98,Morel99a,Morel99b}. With the intensive monitoring campaign presented here, we are now able to conclude that, as for those two stars, WR\\,1 presents a unique period in its epoch-dependent light-curve and spectral varibility. The changes consist in large-scale bumps moving periodically from one side of emission lines to the other. The origin of this type of variability is debated in several papers and the two main interpretations put forward are a binary system (with a compact or low-mass companion) and a rotating non-spherically symmetric wind. In what follows, we discuss these two possible interpretations in the context of the variability of WR\\,1 as characterized in this work, as well as that of WR\\,6 and WR\\,134 as detailed in the literature. We also discuss the variability of the WR star member of the binary WR\\,137, which shows large-scale lpv with a 1.2d-period that cannot be caused by the O9 companion nor the wind-wind collision zone, since the orbital period is $\\sim$13~years and the average distance between the two components is high.  \\subsection{The Binarity Scenario}  Like WR\\,6, WR\\,134 and WR\\,137, WR\\,1 shows very small RV amplitude, consistent with null, and its soft X-ray luminosity is within the normal range for single WR stars \\citep[see ][]{Morel99b}. Hence, it is completely excluded that its companion is an OB star. The presence of a compact companion, such as a neutron star or a black hole, in WR\\,1, WR\\,6 and WR\\,134 was not seen as a viable interpretation by \\citet{Ign2}, \\citet{Morel97}, \\citet{Skinner} and \\citet{Morel99a}. Indeed, the strong correlation between the lpv of He{\\sc ii} lines with that of a highly ionized line such as N{\\sc v}$\\lambda$4945 indicates that the ionizing shell around the compact companion would have to be extremely large and would emit a high X-ray flux, which is not observed \\citep{Morel97,Morel99a,Pol}.  Therefore the only remaining possibility for the binary scenario for WR\\,1, WR\\,6 and WR\\,134 is that of a low-mass, non-compact companion. Indeed, assuming circular orbits and taking the published RV variation amplitudes measured for the three stars \\citep[][ and this work]{Morel97,Morel98} as the maximum possible value (recall that RV variations are highly dominated by lpv) the periods lead to an upper limit for the mass of such a companion of 8~M$_\\odot$ in all cases, if the inclination of the orbital plane is higher than 20$^\\circ$. As for WR\\,137, the scatter in RV on a timescale of a few days is $\\sim$10~km~s$^{-1}$ \\citep{Lef} and its measured M~sin$^3$i=3.4$\\pm$1.0~M$_\\odot$. Hence, if the inclination of the orbital plane is higher than 20$^\\circ$, the mass of the companion would have to be as small as 0.8~M$_\\odot$ in an orbit of 7.6~R$_\\odot$, which is very unlikely. But, can a low-mass companion be responsible for such a complex and epoch-dependant pattern of photometric and spectroscopic variability?  \\citet{More1} and \\citet{More2} have shown that tidal interactions involving a relatively low-mass companion in a binary system can produce very small-scale surface oscillations leading to lpv in photospheric absorption features. However, it is unclear how such a process can affect the massive wind of a WR star and lead to large-scale variability of the strong emission lines. It is possible that it could serve as a seed mechanism for CIRs. However, in such a case, it is likely that the period that would be detected would be the rotation period of the WR star rather than the orbital period of the binary, if not a combination of both.  Therefore, although very unlikely, it is not yet possible from the observations to definitively exclude the possibility that WR\\,1, WR\\,6 and WR\\,134 are binaries, and that WR\\,137 has a third low-mass companion. However, the only viable binary scenario involves the formation of massive binaries with a high initial-mass ratio (up to 10), which seems to be extremely unlikely both observationally and theoretically \\citep[e.g.][]{Gar,Kob,Bon}.  \\subsection{The Non-Spherically Symmetric Wind Scenario}  If WR\\,1, WR\\,6, WR\\,134 and the WR star in the WR\\,137 binary system are single stars, the periodic variability very likely originates in an asymmetry in their wind modulated by the stellar rotation. Spectropolarimetric observations of WR\\,6, WR\\,134 and WR\\,137 \\citep{Schulteladbeck91,Robert92,Harries98} have shown the presence of an intrinsic continuum polarization component due to electron scattering, indicating that the winds of these stars are not spherically symmetric. This is revealed by the depolarization of emission lines compared to the neighboring continuum. It was also found that the degree of depolarisation of spectral lines increases with decreasing ionisation state.  This is interpreted to be a consequence of the ionisation stratification of hot stellar winds; as scattering lines with a higher degree of ionization are formed deeper in the wind, their polarisation level is closer to that of the continuum because of the higher density close to the core and therefore their level of depolarisation is lower. If a line has a strong recombination component, wich is not polarised, its polarization level will be even lower and the depolarization with repect to the continuum stronger. Finally, the level of linear polarization of the continuum as measured in braodband light of WR\\,6, WR\\,134 and WR\\,137 is extremely variable and periodic \\citep{Dri,Mof93}. This suggests that the wind density has a varying asymmetric distribution, such as density structures that extend rather far in the wind. Unfortunately, there is no published high signal-to-noise spectropolarimetric nor continuum polarization observation of WR\\,1. The only observation so far was carried out by \\citet{Schmidt88} and does not show a depolarization in its emission lines compared to the underlying continuum. The latter, however, shows a significantly high level of polarization, although it remains to be demonstrated that this is intrinsic to the star rather than interstellar in origin. If the polarization is proved to be interstellar, this would explain why no depolarization is observed in the lines since the continuum is not polarized to start with.  If, on the other hand, the light from the star is confirmed to be highly polarized, a possible explanation for the lack of depolarization in the lines is that for some yet unknown reason, the region in which the lines arise shows a degree of polarization as high as that in which the continuum arises. Also, the epoch-dependent nature of the changes could mean that relatively quiet periods can exist. Therefore, more spectropolarimetric observations are needed to verify wether the polarization of the emission lines of WR\\,1 varies with time.  The possible existence of large-scale density structures, such as CIRs, was first proposed by \\citet{Mul}. Following this idea, \\citet{Cranmer96} modeled the propagation of CIRs in a hot, radiatively, line-driven stellar wind. In that context, CIRs are caused by perturbations at the base of the wind, which in turn could be caused for example by a magnetic field or pulsations. These perturbations propagate through the wind while being carried around by rotation. This generates spiral-like structures in the density distribution that can lead to a characteristic, large-scale, periodic variability pattern in WR-wind emission lines, extremely similar to what is observed in the WR stars we discuss here \\citep{Dessart02}. Thus, taking into account the photometric and spectroscopic periodic variability, the spectropolarimetric observations and the low soft and hard X-ray fluxes, we conclude that CIRs constitute an extremely likely interpretation. As for the epoch-dependency, it can easily be explained by the variable behaviour of perturbations that generate CIRs together with a finite lifetime of the structures in the wind. Indeed, even if the period of the variability caused by the motion of CIRs is always the same, the number and position of the CIRs may change depending on conditions at the surface. The origine of CIRs is still debated and no clear theoritical predictions of the typical lifetime of a CIR have been made yet. However, one can speculate that it will depend on the lifetime of the perturbation at the base of the wind and, to a lesser extent, on the flow speed in the wind. Observationaly, it could be determined from continuous observations during a great number of rotational periods (e.g. several weeks for WR\\,6 and WR\\,134).  \\subsection{Remarks on the putative CIRs in the wind of WR\\,1}\\label{con}  We believe that the best scenario to explain the photometric and spectroscopic variability we have detected in WR\\,1 is the presence in its wind of large-scale structures, most likely CIRs. In that case, each CIR would translate into a bump in the light-curve and a bump over the spectral emission lines.  In theory, CIRs do not suffer from differential rotation and if the recurrence of the changes is caused solely by their rotation, they provide a direct measurement of the rotation period of the underlying star at the position in which they originate, most likely close to the stellar surface. Thus, assuming a value for the radius of WR\\,1 of 2.2~R$_\\odot$ \\citep[a radius that corresponds to a Rosseland optical depth of 20;][]{Ham3}, we obtain an equatorial rotational velocity of 6.5~km~s$^{-1}$. This is an order of magnitude lower than the values obtained for WR\\,6 (40~km~s$^{-1}$) and WR\\,134 (70~km~s$^{-1}$) when assuming a radius of R$_\\ast \\sim 3 R_\\odot$ \\citep{Ham3}. However, all these rotational velocities are in agreement with the very small values predicted for WR stars by massive-star evolutionary models at solar metalicity \\citep{Mey}. The order of magnitude difference between the rotation velocity of WR\\,1 and that of WR\\,6 and WR\\,134 is not completely surprising since the observed rotational velocity of a massive star at the WR evolutionary stage depends on several parameters such as the initial stellar mass, the mass-loss rate and the age of the star (time spent as a WR star). Interestingly, adopting a radius of 4.5~R$_\\odot$ \\citep[corresponding to the radius of the hydrostatic core; ][]{Nug} for WR\\,137, we deduce an equatorial rotational velocity of 275~km~s$^{-1}$, i.e. more than half the breakup velocity. If confirmed, such a fast rotation for a WC star would render WR\\,137 an interesting candidate for an eventual long-term gamma-ray burst.  Inspired by the analysis of \\citet{Dessart02}, we can estimate the inclination of the CIRs in the wind of WR\\,1. Indeed, the maximum Doppler velocity that a bump associated with a CIR reaches during its motion on an emission line is $v_{max}=\\pm v_{lfr} \\rm{cos}(\\theta)$, where $v_{lfr}$ is the velocity of the wind at the radii corresponding to the formation region of the observed emission line over which the bump is observed, and $\\theta$ is the inclination angle of the CIR with the line-of-sight. In Figure~\\ref{montage}, we track two bumps that reach $v_{max1}\\sim$ --1300~km~s$^{-1}$ and $v_{max2}\\sim$ +700~km~s$^{-1}$ on the He{\\sc ii}$\\lambda$5411 line. Assuming that the velocity law of the wind can be described by a $\\beta$-law for which the velocity as a function of radius can be written as follows \\citep{Castor}: \\begin{eqnarray} v(r)=v_\\infty\\left(1-\\frac{R_\\ast}{r}\\right)^\\beta, \\end{eqnarray} where $v_\\infty$ is the terminal velocity of the wind, the value of $v_{lfr}$ for He{\\sc ii}$\\lambda$5411 can be estimated using the emissivity function described by \\citet{Lep}. However, that function depends on the value of $\\beta$. Unfortunately, $\\beta$ cannot be estimated observationally for WR\\,1, since no clumps are observed in the spectra. We can, however, give estimates of $\\theta$ as a function of $\\beta$. If $\\beta$=1, 2 or 3 the first CIR has an inclination angle of $\\theta_1$=$\\pm$40$^\\circ$, $\\pm$38$^\\circ$ or $\\pm$30$^\\circ$, respectively, and the second bump has $\\theta_2$=$\\pm$65$^\\circ$, $\\pm$64$^\\circ$ or $\\pm$62$^\\circ$, respectively. When $\\beta$ is greater than 3, the velocity of the He{\\sc ii}$\\lambda$5411 line formation region is lower than 1300~km~s$^{-1}$. Hence, under our first assumptions, we deduce that the $\\beta$ value for the wind of WR\\,1 should be lower than 3.  Of course, the association of the period of the CIRs with the rotation period of the underlying star remains to be confirmed. In the wind of OB stars, the presence of CIRs is thought to be at the origin of periodically recuring Discrete Absorption Components (DACs) in P~Cygni profile of UV resonance lines \\citep[e.g. ][]{Fullerton97}. In many cases, the period of the outward moving DACs is found to correspond to $v_{eq} \\sin{i}$, where $v_{eq}$ is the rotational velocity at the equator and $i$ the inclinaison of the rotational axis with the line of sight \\citep{Hen,Prin}. But in the case of HD\\,64760, a B0.5 Ib star, the CIRs rotate more slowly than the stellar surface. \\citet{Lob} carried out a hydrodynamical simulation of the CIRs for that star, following \\citet{Cranmer96}, with the difference that they allowed the ``spots'' at the origine of the CIRs to move on the stellar surface. They were able to obtain a fairly good fit to the lpv of the Si{\\sc iv}$\\lambda\\lambda$1394,1403 and concluded that, for that star, the origin of the CIRs must be the interference pattern of a number of non-radial pulsations at the surface of the star.  So far, there is no model for the ``spots'' that are at the origin of the creation of the CIRs and there are only a limited number of datasets available that could potentially be used to confirm either the pulsational or the magnetic origin of CIRs. The range of frequency expected for pulsations in WR stars is not well determined. Up to date, only one detection of a relatively stable 9.8~h period that can be attributed either to g-modes \\citep{Tow} or stange-mode pulsations \\citep{Dor} has been claimed in WR\\,123 \\citep{Lef2}. No pulsational period of a few days is currently known nor predicted. Also, no detection of a magnetic field has been achieved in a WR star so far; only an upper limit of $\\sim$ 25 Gauss for WR\\,6 has been claimed \\citep{Stl}."
    },
    "1004/1004.1364_arXiv.txt": {
        "abstract": "We present results from the \\Spitzer/IRS spectral mapping observations of 15 local luminous infrared galaxies (LIRGs). In this paper we investigate the spatial variations of the mid-IR emission which includes: fine structure lines, molecular hydrogen lines, polycyclic aromatic features (PAHs), continuum emission and the 9.7\\micron\\ silicate feature. We also compare the nuclear and integrated spectra. We find that the star formation takes place in extended regions (several kpc) as probed by the PAH emission as well as the \\Neii\\ and \\Neiii\\ emissions. The behavior of the integrated PAH emission and 9.7\\micron\\ silicate feature is similar to that of local starburst galaxies. We also find that the minima of the \\Neiii\\slash\\Neii\\ ratio tends to be located at the nuclei and its value is lower than that of \\HII\\ regions in our LIRGs and nearby galaxies. It is likely that increased densities in the nuclei of LIRGs are responsible for the smaller nuclear \\Neiii\\slash\\Neii\\ ratios. This includes the possibility that some of the most massive stars in the nuclei are still embedded in ultracompact \\HII\\ regions. In a large fraction of our sample the \\PAHonce\\ emission appears more extended than the dust 5.5\\micron\\ continuum emission. We find a dependency of the \\PAHonce\\slash\\PAHsiete\\ and \\Neii\\slash\\PAHonce\\ ratios with the age of the stellar populations. Smaller and larger ratios respectively indicate recent star formation. The estimated warm (300 K $<$ T $<$ 1000 K) molecular hydrogen masses are of the order of 10$^{8}$M$_\\Sun$, which are similar to those found in ULIRGs, local starbursts and Seyfert galaxies. Finally we find that the \\Neii\\ velocity fields for most of the LIRGs in our sample are compatible with a rotating disk at $\\sim$kpc scales, and they are in a good agreement with H$\\alpha$ velocity fields. ",
        "introduction": "\\label{s:intro}  Infrared (IR) bright galaxies were first identified almost 40 years ago \\citep{Rieke72}. The {\\it IRAS} satellite detected a large number of these galaxies in the local Universe (see \\citealt{Sanders96} for a review). Galaxies were classified according  to their IR ($8-1000\\,\\mu$m) luminosities into Luminous Infrared Galaxies (LIRGs, $10^{11}L_{\\sun}< L_{IR} <10^{12}L_{\\sun}$) and Ultra Luminous Infrared Galaxies (ULIRGs, $L_{IR} > 10^{12}L_{\\sun}$).  Since then, an increasing number of studies showed the important role of (U)LIRGs at cosmological distances. LIRGs and ULIRGs are not common in the local Universe, accounting for just $5\\%$ and $<1\\%$, respectively, of the total IR emission of galaxies. However, at z$\\sim$1, LIRGs dominate the IR background and the co-moving star formation rate density, while ULIRGs are dominant at z$\\sim$2 \\citep{LeFloch2005, PerezGonzalez2005, Caputi2007}.  Thanks to the high sensitivity of the Infrared Spectrograph (IRS) \\citep{HouckIRS} on-board \\Spitzer, we can, for the first time, study systematically the mid-IR properties of (U)LIRGs, both locally and at high-$z$. The majority of IRS local studies have, so far, focused on the properties of starburst galaxies \\citep{Smith07, Brandl06} and ULIRGs \\citep{Armus07, Farrah07}. The latter class of galaxies includes a large variety of physical and excitation conditions. For instance, most of these local ($z<0.3$) ULIRGs are powered predominantly by compact starburst events, although a substantial AGN contribution is present in about half of them \\citep{Imanishi2007, Farrah07, Nardini2008}. In the future the larger sample ($\\sim$200) of local LIRGs of the Spitzer legacy program GOALS will enable a statistical study of this class of galaxies \\citep{Armus09}.  Although ULIRGs at z$\\sim$2 extend to even higher luminosities than local ones, their mid-IR spectra are similar to those of local starbursts, rather than to local ULIRGs \\citep{Rigby2008, Farrah08}. A similar behavior is found for high-z submillimeter galaxies \\citep{Pope2007, MenendezDelmestre2009}. This behavior may indicate that star formation in IR-bright galaxies at z$\\sim$2 is taking place over larger scales (a few kpc) than in local ULIRGs, where most of the IR emission arises from very compact regions (sub-kiloparsec scales, see e.g., \\citealt{Soifer2000, Farrah07}).  As a result, local LIRGs may provide important insights to the behavior of high-redshift ULIRGs. Their star formation is often distributed over a substantial fraction of the galaxy (e.g., \\citealt{AAH06s}), and in any case is usually not embedded at the high optical depths characteristic of the nuclear star forming regions in local ULIRGs. To make good use of this analogy requires observations that map the local LIRGs to get an understanding of their global properties as well as to provide a better comparison with the observations of high-z galaxies, where the entire galaxy is encompassed in the IRS slit. The spectral mapping mode of \\Spitzer/IRS is well suited for this purpose, since it can provide both the nuclear and integrated spectra of local galaxies, enabling study of the spatial distribution of the spectral features that make up the integrated spectrum.   This is the first in a series of papers studying the mid-IR properties of local LIRGs. We present \\Spitzer/IRS spectral mapping observations of a representative sample of fifteen local LIRGs. The first results and the goals of this program were presented by \\citet{AAH09ASR}. In the present paper we focus on the analysis of the spatially resolved measurements. The paper is organized as follows. In Section \\ref{s:observations} we present the sample and the observational details. Section \\ref{s:analysis} describes the analysis of the data. We examine the silicate feature in Section \\ref{s:silicate}. Sections \\ref{s:lines}, \\ref{s:PAH}, \\ref{s:molecular_hydrogen} discuss the spatially resolved properties of the atomic fine structure lines, PAHs, and molecular hydrogen emission, respectively. We explore the velocity fields in Section \\ref{s:velocity_fields}. Finally, in Section \\ref{s:conclusions} we summarize the main results. Throughout this paper we assume a flat cosmology with $H_0 = 70$ km s$^{-1}$Mpc$^{-1}$, $\\Omega_M = 0.3$ and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "\\label{s:conclusions}  We presented the analysis of \\Spitzer/IRS mapping observations of 15 local LIRGs with extended Pa$\\alpha$ emission. We studied the spatial distribution of the mid-IR spectral features and compared the nuclear and integrated spectra. We calculated fine structure line and PAH ratios which trace the physical conditions in the star-forming regions. The main results are the following:  \\begin{itemize}  \\item We used the 9.7\\micron\\ silicate feature strength vs. \\PAHseis\\ equivalent width diagram to classify the activity of the LIRGs. There is a good agreement with the optical classification and this diagram. Most of the LIRGs populate the pure starburst class and only the nuclei classified as Seyfert from optical spectroscopy appear in the diagram in the AGN/SB class. The integrated values show, in general, larger \\PAHseis\\ EW and shallower silicate absorption than the nuclear spectra. That is, the extended star formation partially masks the nuclear activity, resulting in a starburst-like integrated spectra. The silicate feature strength, nuclear and integrated, of our sample of local LIRGs ($S_{\\rm Si} \\sim-0.4 \\, {\\rm to}  \\, -0.9$) is small compared to that found in ULIRGs ($S_{\\rm Si}$ up to -4).  \\item We constructed maps of the spatial distribution of the \\Neiii\\slash\\Neii, which traces the radiation field hardness, and the \\SIIIa\\slash\\Neii\\ that varies with the electron density. In general, the minimum of both ratios is located at the nucleus. The nuclear \\Neiii\\slash\\Neii\\ ratio is below the expected range derived by star-formation models. A possible explanation is that due the high densities in the nuclear regions this ratio is suppressed, including the possibilities that the most massive stars are either missing of buried in ultracompact \\HII\\ regions. Alternatively the star formation rate may have decreased rapidly in all nuclei over the last 10-20 Myr.  \\item We find a positive correlation between the \\Neiii\\slash\\Neii\\ and \\SIIIa\\slash\\Neii\\ ratios for star forming regions in our sample of LIRGs and other starburst galaxies. On the other hand, AGNs, for a given \\SIIIa\\slash\\Neii\\ ratio, show a systematically larger \\Neiii\\slash\\Neii\\ ratio. The few starburst galaxies where we detect the \\SIV\\ emission line follow the correlation between the \\SIV\\slash\\SIIIa\\ and \\Neiii\\slash\\Neii\\ observed in \\HII\\ regions in other galaxies. We observe larger \\SIV\\slash\\SIIIa\\ ratios for those galaxies harboring an AGN.  \\item We find that the \\PAHonce\\ emission is more extended than that of the 5.5\\micron\\ continuum. However, the ratio of the nuclear (2 kpc) \\PAHseis\\ emission with respect to the integrated emission is comparable to that of the 5.5\\micron\\ continuum in most cases. We find no correlation between the \\Neiii\\slash\\Neii\\ ratio and the \\PAHonce\\ EW, thus the effect of the radiation field hardness in the PAH emission, for the \\Neiii\\slash\\Neii\\ ratio range in our sample of LIRGs, may not be important. We propose that the \\PAHonce\\slash\\PAHsiete\\ ratio depends on the age of the stellar population. While the \\PAHsiete\\ comes from young star-forming regions, the \\PAHonce\\ arises also from diffuse medium, thus the \\PAHonce\\slash\\PAHsiete\\ is lower in young \\HII\\ regions. We also found that large \\Neii\\slash\\PAHonce\\ ratios may indicate recent star formation.  \\item We explored the variation of the PAH ratios across the galaxies. In most cases these variations are real, that is, they are not due to extinction). In general the integrated \\PAHonce\\slash\\PAHseis\\ ratios are larger than the nuclear values, probably indicating that the integrated emission includes more neutral conditions. The nuclear and integrated  \\PAHsiete\\slash\\PAHseis\\ ratios are almost constant, except in the nuclei of galaxies classified as AGN where this ratio is higher. Since the \\PAHseis\\ emission is more associated to small molecules, the increased \\PAHsiete\\slash\\PAHseis\\ ratio in AGN could be explained if the small PAH molecules are more easily destroyed in the harsh environments of active nuclei.  \\item Using the \\Hm{0} at 28.2\\micron\\ and the \\Hm{1} at 17.0\\micron\\ lines integrated over 13''$\\times$13'' (3 to 5 kpc) we estimated the mass of the warm (T$\\sim$300 K) molecular hydrogen. It ranges from 0.4 to 2.6$\\times$10$^8$M$_\\sun$, and they are similar to those of ULIRGs, local starbursts and Seyfert galaxies. However these masses are lower limits since it is only included the \\Hmol\\ emission from the central few kpc of the galaxies and the total mass of warm molecular hydrogen is likely to be, at least, a factor of 2 larger.  \\item{The similarity between the PAHs and molecular hydrogen morphologies suggests that the main excitation mechanism of the latter is the UV radiation too. However there are some regions with an excess of \\Hmol\\ emission with respect to the PAH emission, thus the other mechanisms should contribute noticeably to the \\Hmol\\ emission. Some of these regions are associated to interacting systems, where large scale shocks may also play a role in exciting the molecular hydrogen.}  \\item Despite the modest spectral resolution of the SH module, we show that useful velocity information can be obtained from the SH spectra. For most of the galaxies and on the physical scales probed by the IRS spectra ($\\sim$kpc) the velocity fields are comparable with that produced by a rotating disk. For the galaxies with available H$\\alpha$ velocity fields, the \\Neii\\ velocity fields are in good agreement with those of the H$\\alpha$, both in shape and peak-to-peak velocities.   \\end{itemize}"
    },
    "1004/1004.2771_arXiv.txt": {
        "abstract": "A new rapid energization process within a supernova shock transition region (STR) is reported by utilizing numerical simulation. Although the scale of a STR as a main dissipation region is only several hundreds of thousands km, several interesting structures are found relating to generation of a root of the energetic particles. The nonlinear evolution of plasma instabilities lead to a dynamical change in the ion phase space distribution which associates with change of the field properties. As a result, different types of large-amplitude field structures appear. One is the leading wave packet and another is a series of magnetic solitary humps. Each field structure has a microscopic  scale ($\\sim$ the ion inertia length). Through the multiple nonlinear scattering between these large-amplitude field structures, electrons are accelerated directly. Within a STR, quick thermalization realizes energy equipartition between the ion and electron, hot electrons play an important role in keeping these large-amplitude field structures on the ion-acoustic mode. The hot electron shows non-Maxwellian distribution and could be the seed of further non-thermal population. The \"shock system\", where fresh incoming and reflected ions are supplied constantly, play an essential role in our result. With a perpendicular shock geometry, the maximum energy of the electron is estimated by equating a width of the STR to a length of the Larmor radius of the energetic electron. Under some realistic condition of $M_A = 170$ and $\\omega_{pe}/\\Omega_{ce} = 120$, maximum energy is estimated to $\\sim$ 10 MeV at an instant only within the STR. ",
        "introduction": "Recently, the origin of the energetic electron related to astrophysical shocks has been investigated more and more from the point of view of the plasma kinetic processes \\cite{SH00, McC01, HS02, Dieck04, Lem04, Dieck08, Mat06, Tr08, AH09}. Within the theory of well-established classical diffusive shock acceleration (DSA) mechanism \\cite{Bell87, BE87}, a shock wave tends to be treated as a simple discontinuity. At a shock wave, however, we can find that many kind of plasma kinetics play a key role in the energy dissipation and generation of energetic particles. Correct treatment of the plasma kinetic process around the shock may succeed to settle so-called ``injection problem\" to DSA process (e.g. Ref.\\cite{Lev96, Ba03, AH07}) as well as to capture a whole stage of DSA process through nonlinear interaction between the shock structure and energetic particles \\cite{Bere99, Bell04, ZP08, Cap09}. Furthermore, not only as an assistant role to DSA mechanism, investigation of the energy release mechanism in terms of the plasma kinetics can be a powerful tool to search significant electron energization process in much smaller spatial/time scales than that DSA mechanism may need.  In this paper, we focus on a nonrelativistic ion-electron supernova shock with the perpendicular geometry where the shock normal is perpendicular to the direction of the background magnetic field (e.g. Ref.\\cite{Tr08p}). Particle-in-cell (PIC) simulation has been carried out to investigate coupling dynamics among the ion, electron, and field in the shock transition region. The shock transition region (STR) is the front side of the shock wave corresponding to the region flow speed decreases gradually from the upstream toward the shocked downstream region. At the high Mach number shocks, part of the incoming ions are always reflected at the shock front and move into the upstream region. In detail, the STR means here the region from the leading edge of the reflected ion to around the end point of the first macroscopic ion gyration in a shocked region. The upstream electrons are immediately decelerated just after they meet with the reflected ion. The velocity difference between the decelerated electron and the incoming ion is large enough to excite a series of micro-scale potentials through a strong two-stream instability which grows into a packet of large-amplitude spatial-oscillating magnetic field at the leading edge of the reflected ion (hereafter we call this ``leading packet\"). When the incoming ion begins to intermingle with the reflected ion, two ion components together make complex vortices in the velocity phase space, the electric and magnetic field structures change immediately. As a result, just after this dynamical change, some components of the leading packet are converted into magnetic solitary humps almost standing in the electron bulk flow frame. The electrons bouncing around between the leading packet and the standing magnetic humps get energy rapidly by the 1st-order Fermi type process. We define here, the 1st-order Fermi type acceleration as the acceleration of particles by a shock with an extended transition layer, while the original 1st-order Fermi mechanism (DSA) assumes an infinitesimally thin transition layer. Among magnetic humps, the electrons are also bouncing around and sometimes getting energy in a stochastic manner which depends on mainly local motional electric field caused by a combination of magnetic field of the humps and local bulk flow variation due to the macroscopic ion gyro-motion. Within the STR, strong thermalization occurs to attain energy equipartition between the electron and ion. The hot electrons contribute to keep above large-amplitude structures on the ion-acoustic mode. The electron energization mechanisms stated above are quite dynamical process characteristic of a \"shock system\".  In the next section we introduce the simulation setup. Section III. presents simulation results: A. macroscopic overview of shock wave properties, B. the energy re-distribution process in the STR, and C. the structure of the STR and characteristic field generation. In subsection III.D we show electron energy spectra and discuss some examples of electron trajectories in the nonlinear field evolution. Section IV. summarizes and discusses our results. ",
        "conclusions": "The numerical simulation of a collisionless shock wave is reported under high Mach number and high frequency ratio condition ($M_A = 174$ and $\\omega_{pe}/\\Omega_{ce} = 120$) similar to shock waves generated by the strong blast wave of supernovae. We focus on the electron dynamics through the nonlinear particle-field coupling process. We found that the first rapid electron acceleration occurs in the thin shock front region. Nonlinear evolution of the plasma instabilities between the ions and electrons and following the ion-ion instability causes strong particle energization. At the moment of the nonlinear saturation of the instability, strong mixing occurs between the incoming and reflected ion component in the velocity phase space as well as some considerable change in the field structures, for example, from the magnetic leading packet to the magnetic solitary hump. The leading packet consists of a series of large-amplitude negative-positive magnetic fields. (When a 2D or 3D simulation is carried out, we can expect occurrence of the magnetic reconnection there. It may bring about more interesting electron-field dynamics.) Since the each component of the leading packet and the magnetic hump are converging, some electrons scattered and bouncing between these two structures gain energy rapidly by the 1st-order Fermi type process. Some electrons around magnetic humps also gain energy in a rather stochastic manner. In the magnetic hump region, there is still non-zero bulk flow variation due to the ion gyro-motion. The electrons gain energy not only due to the local $E_x$ resulted from nonlinear instability evolution, but also due to the motional electric field $E_y$ accompanied by the magnetic hump $B_z$ and bulk flow. Depending on the gyro-phase, an energetic electron sometimes gains energy when its gyro-motion couples with the $E_y$ variation of the ion-scale.  It is surprising that the electron energy density arises rapidly at just an entrance of the STR and becomes comparable to the ion energy density. This energy equipartition is pushed by the nonlinear saturation process of the plasma instability. So-called ``injection problem\" in DSA mechanism may be settled if our result remain still true under the realistic dimension and mass ratio condition. We showed that in our simulation, the energetic electrons are confined within a STR while they are interacting with the fields and gaining energy. So that we can estimate the maximum energy $\\gamma_{max}$ by equating the electron Larmor radius and characteristic scale of the STR, \\begin{eqnarray} \\displaystyle \\frac{mc^2}{eB} \\sqrt{\\gamma_{max}^2 -1} \\sim \\displaystyle \\frac{1}{2} \\frac{\\Gamma_0 u_0}{\\Omega_{ci}} \\end{eqnarray} where, $\\Gamma_0$ is incoming flow Lorentz factor. Using $u_0 = 2/3 V_A M_A$, above equation results, \\begin{eqnarray} \\gamma_{max} \\sim \\displaystyle \\sqrt{1+\\frac{1}{9} \\frac{M}{m}\\left( \\frac{\\omega_{pe}}{\\Omega_{ce}} \\right) ^{-2} M_A ^{2} \\Gamma_0 ^2} \\end{eqnarray} When we adopt above equation to the shock wave case in this paper, we have $\\gamma_{max} \\sim $5 which is consistent to our result of $\\gamma_{max} \\sim$4. With realistic mass ratio $M/m=1836$, $\\gamma_{max}$ can reach up to 21 (corresponding to $\\sim$ 10 MeV). Under a larger mass ratio condition, the STR becomes wider in terms of the scale of electron dynamics and the ion inertia becomes relatively larger. Since the ion inertia plays an important role to amplify the fields discussed in the current paper, we can expect more electron energization. In fact, some periodic simulations in Ref.\\cite{SH04} shows that when the mass ratio is larger up to the realistic mass ratio, the electric field supported by the ion inertia becomes more important in the electron energization process.  Recent observation shows shocked temperature $T_e < T_i$ for strong shocks (e.g. Ref.\\cite{Rako05, Gh07, He07, Ad08, H09}). Contrary to the observation, our simulation result shows $T_e \\ge T_i$. One reason for that is a scale difference. Our simulation scale is only about several $L_{STR}$ much smaller than observation can resolve, where $L_{STR} \\equiv u_0/\\Omega_{ci} \\sim 4.6\\times10^5$ km (for interstellar density of 0.1 cm$^{-3}$), $\\sim 1.5\\times10^5$ km (for 1.0 cm$^{-3}$) using $\\omega_{pe}/\\Omega_{ci} = 120$ condition. The temperature far outside of the STR is, unfortunately, beyond a scope of the paper. Another reason is a shock configuration treated here, namely, the 1D and perpendicular shock system which is a rather strong restriction. For example, 1D simulation tends to overheat electrons compared to the two-dimensional simulations \\cite{AH09, Umeda09}. Although strong thermalization indeed occurs in our simulation, the hot electrons shows non-Maxwellian distribution and non-thermal part grows as leaving off the STR. The hot electrons can be not only a seed of the future non-thermal population but also contribute to further acceleration by keeping the field structures within a scale $\\sim L_{STR}$.  Unlike the classical diffusive shock acceleration where the electrons are hovering around shocked and unshocked region due to scattering by large scale MHD waves, the electrons in our simulation experience multiple nonlinear interaction, even within one gyro-motion, with the large-amplitude field structures which generated through the nonlinear evolution of the plasma instabilities. These nonlinear transportations may lead a breakout of the classical diffusion model. Recently, without the classical diffusion model, \\citet{MD09} also discuss particle acceleration process under multiple nonlinear interaction among a train of magnetic structure ``shocklets\" of $c/\\omega_{pi}$ scale in the confined region of the STR (shock precursor). As a generator of the structure ``shocklets\", nonlinear evolution of the plasma instability is also considered between cosmic ray flow and shock incident flow. Not only for our rather lower energetic regime, also for higher energetic regime, such microscopic coherent structures due to the nonlinear evolution of the plasma instabilities confined in the STR is expected to play a key role in the particle acceleration process."
    },
    "1004/1004.0890_arXiv.txt": {
        "abstract": " ",
        "introduction": "Topological defects are remnants of spontaneously broken local or global symmetries. The simplest and the most well-known example of the former one is the Abrikosov-Nielsen-Olesen flux tube \\cite{ANO}, which originates from spontaneously broken $U(1)$ gauge symmetry. Most of the attention in the literature has been focused on defects originating from broken gauge symmetries, since grand unified theories have gauge symmetries which are eventually spontaneously broken down to the symmetry of the Standard Model. Cosmic strings are one dimensional topological defects predicted by a large class of unified theories \\cite{kibble76,alexbook,Jeannerot:2003qv}. Cosmic strings were first considered as the seeds of structure formation \\cite{Zeldovich,AV81a}, however, later, it was discovered that cosmic strings were incompatible with the cosmic microwave background (CMB) angular power spectrum. Cosmic strings can still contribute to structure formation, but they cannot be the dominant source.  Cosmic strings are also candidates for the generation of other observable astrophysical phenomena such as high energy cosmic rays, gamma ray burst and gravitational waves \\cite{alexbook, book anderson,khlopov1,khlopov2}. Furthermore, recently it has been shown that in string-theory-inspired cosmological scenarios cosmic strings may also be generated~\\cite{cstrings}. They are referred to as cosmic superstrings. This realization has revitalized interest in cosmic strings and their potential observational signatures. There are some important differences between cosmic strings and cosmic superstrings. The reconnection probability is unity for cosmic strings \\cite{alexbook,Nitta}. Cosmic superstrings, on the other hand, have reconnection probability less than unity. This is a result of the probabilistic nature of their interaction and also the fact that it is less probable for strings to meet since they can live in higher dimensions~\\cite{dvali vilenkin 2004}. The value of $p$ ranges from $10^{-3}$ to $1$ in different theories \\cite{jjp}. Cosmic superstrings could also be unstable, decaying long before the present time. In this case, however, they may also leave behind a detectable gravitational wave signature~\\cite{Leblond:2009fq}.  In the early universe, a network of cosmic strings evolves toward to an attractor solution called  the ``scaling regime\". In the scaling regime the statistical properties of the network, such as the average distance between strings and the size of loops at formation, scale with the cosmic time. In addition, the energy density of the network remains a small constant fraction of the energy density of the universe. For cosmic superstrings in the scaling regime, the density of the network $\\rho$ is inversely proportional to the reconnection probability $p$, that is $\\rho\\propto p^{-\\beta}$. The value of $\\beta$ is still under debate~\\cite{stoica tye,Sakellariadou 2005,avgoustidis shellard}, and as a placeholder in our analysis we assume that $\\beta=1$.  The gravitational interaction of strings is characterized by their tension $\\mu$, or more conveniently by the dimensionless parameter $G \\mu$, where $G$ is Newton's constant. The current CMB bound on the tension is $G \\mu<6.1\\times 10^{-7}$ \\cite{CMB1,CMB2}. It was first believed that gravitational radiation from cosmic strings with $G \\mu\\ll 10^7$ would be too weak to observe. However it was later shown that gravitational radiation produced at cusps, which have large Lorentz boosts, could lead to a detectable signal~\\cite{DV1,DV2,SCMMCR}. Gravitational radiation bursts from (super)strings could be observable by current and planned gravitational wave detectors for values of $G \\mu$ as low as $10^{-13}$, which may provide a test for a certain class of string theories \\cite{pol1}. Indeed, searches for burst signals using ground-based detectors are already underway~\\cite{Abbott:2009rr}.  A gravitational background produced by the incoherent superposition of cusp  bursts from a network of cosmic strings and superstrings was considered in \\cite{Siemens Mandic Creighton}. In this paper we extend this computation to include kinks, long-lived sharp edges on strings that result from intercommutations, and find that kinks contribute at almost the same level as cusps. We investigate the detectability of the total background produced by cusps and kinks by a wide range of current and planned experiments. A similar calculation for the case of infinite strings has been undertaken in the recent paper \\cite{Kawasaki}, see also \\cite{Kawasaki2}.  The organization of the paper is as follows: In Sect. \\ref{Gravitational Radiation} we consider gravitational waves generated by cusps and kinks in the weak field limit \\cite{Weinberg}. In this section we follow the conventions of \\cite{DV1,DV2}, and more details can be found in these references. In Sect. \\ref{Stochastic Background} we derive the expression for the stochastic background, which is a double integral over redshift and loop length. In Sect. \\ref{Analytical app} we evaluate integral analytically with certain approximations, which results in a flat distribution for larger values of the frequency. Finally in Sec. \\ref{parameter scan} we numerically evaluate the background and discuss the observability by various experiment. \\newpage ",
        "conclusions": ""
    },
    "1004/1004.1402_arXiv.txt": {
        "abstract": "The work is an attempt to model a scenario of  inflation in the framework of Anti de Sitter/Conformal Field theory (AdS/CFT) duality, a potentially complete nonperturbative description of quantum gravity via string theory. We look at bubble geometries with de Sitter interiors within an ambient Schwarzschild anti-de Sitter black hole spacetime and obtain a characterization for the states in the dual CFT on boundary of the asymptotic AdS which code the expanding dS bubble. These can then in turn be used to specify initial conditions for cosmology. Our scenario naturally interprets the entropy of de Sitter space as a (logarithm of) subspace of states of the black hole microstates. Consistency checks are performed and a number of implications regarding cosmology are discussed including how the key problems or paradoxes of conventional eternal inflation are overcome. ",
        "introduction": "Making contact with realistic cosmology remains the fundamental challenge for any candidate unified theory of matter and gravity. Precise observations \\cite{Riess:1998dv,Perlmutter:1999rr} indicate that the current epoch of the acceleration of universe is driven by a very mild (in Planck units) negative pressure, positive energy density constituent - {}``dark energy''. Likewise there is strong evidence that the Big Bang phase of the universe was preceded by a exponentially accelerated (inflation) phase \\cite{Guth:1980zm,Linde:1981mu} driven by a negative pressure, positive energy fluid. In most cosmological models this inflationary stage is brought about by a scalar field coupled to gravity, slowly {}``rolling down'' an almost flat potential hill, and upon reaching the bottom of the hill gives rise to the big-bang stage. In most such models inflation is inevitably {}``future eternal'' i.e. there are always some residual regions which keep on rapidly inflating. The steady state picture that emerges is a fractal multiverse structure with many {}``pocket universes'' causally separated by continually inflating sterile regions. These {}``pocket universes'' might possess all possible different {}``fundamental constants'' of nature (including a cosmological constant). So in this scenario one now talks about {}``environmental constants'' of nature instead.  (Super)string theory is the currently leading candidate for an unified theory of fundamental interactions and should be compatible with realistic cosmology if it has to be anything more than an attractive toy model. It is great news that string theory appears to have an enormous {}``landscape'' of vacua including many long-lived metastable states with positive cosmological constants \\cite{Kachru:2003aw} forming a quasi-continuum \\cite{Bousso:2000xa,Feng:2000if,Giddings:2001yu,Susskind:2003kw} which are highly relevant for realistic cosmology. These metastable phases ultimately decay to neighboring stable anti-de Sitter phases via {}``rolling down the hill'' and quantum mechanical tunneling and this is exactly what one needs to attempt a fundamental physics model of the (inflationary) multiverse. So it is important to develop tools in the fundamental theory which describe such transitions.  The AdS/CFT formulation of String theory \\cite{Aharony:1999ti} is a potentially nonperturbative definition of superstring theory in asymptotically AdS spaces and is an perfect setting to embed the string theory landscape. The neighborhood of the landscape that we study is summarized in figure \\ref{fig:The-potential-landscape}. % \\begin{figure} \\includegraphics[scale=0.7]{KKLT}\\caption{The potential landscape showing the dS metastable and AdS stable vacua. The arrows indicate the time evolution of a possible worldline.\\label{fig:The-potential-landscape}}  \\end{figure}   The first step within this approach was taken in \\cite{Alberghi:1999kd} which looked at classical AdS-Reissner-Nordstrom/dS domain wall type geometries (an inflating dS bubble interior patched to an asymptotic AdS black hole spacetime) in the thin wall approximation. Classical bubble spacetimes without charge were further investigated in \\cite{Freivogel:2005qh} and subsequently in \\cite{Lowe:2007ek}. The latter has an added attractive feature that the mysterious Gibbons-Hawking entropy of de Sitter space could be identified as a (logarithm of) subspace of black hole microstates. The aim of the present paper is to build on the work of \\cite{Lowe:2007ek} by attempting to identify the CFT states which code these inflating states in a thermal ensemble, and further explore the consequences of this scenario.  The paper is organized as follows. In section 2, we review the semiclassical consistency condition of \\cite{Lowe:2007ek} which arises from interpreting de Sitter entropy within a unitary framework for quantum gravity. This amounts to the condition that the number of states available must be bounded below by the Gibbons-Hawking entropy if states corresponding to a truly semiclassical region of de Sitter space exist. . Then in section 3, we review the classical bubble spacetimes. Keeping in mind the dual CFT as the underlying quantum description we are forced to rule out the whole class of time-symmetric bubble solutions based on the consistency condition. In particular, this rules out tunneling solutions of the type described in \\cite{Farhi:1989yr}. In section 4, we focus on the CFT picture of the time-asymmetric solutions compatible with the aforementioned consistency condition and their signatures in the dual CFT. We model the effect of the radiation reflecting off the shell by reducing the problem to that of a moving mirror in a black hole background. In section 5 we discuss consistency checks of such a scenario. In section 6 we draw various conclusions regarding the implications for such a picture for cosmology and how several problems/paradoxes that plague the eternal inflation scenario - namely the measure problem, the youngness paradox and the Boltzmann brains paradox are overcome in this setting. ",
        "conclusions": ""
    },
    "1004/1004.1896_arXiv.txt": {
        "abstract": "{Supergiant Fast X-ray Transients (SFXTs) are a new class of High Mass X-ray Binaries, discovered by the \\emph{INTEGRAL} satellite, which display flares lasting from minutes to hours, with peak luminosity of $10^{36} - 10^{37}$~erg~s$^{-1}$. Outside the bright outbursts, they show a frequent long-term flaring activity reaching an X-ray luminosity level of $10^{33} - 10^{34}$~erg~s$^{-1}$, as recently observed with the \\emph{Swift} satellite. Since a few persistent High Mass X-ray Binaries (HMXBs) with supergiant donors show flares with properties similar to those observed in SFXTs, it has been suggested that the flaring activity in both classes could be produced by the same mechanism, probably the accretion of clumps composing the supergiant wind. We have developed a new clumpy wind model for OB supergiants with both a spherical and a non spherical symmetry for the outflow. We have investigated the effects of the accretion of a clumpy wind onto a neutron star in both classes of persistent and transient HMXBs.}  \\FullConference{The Extreme sky: Sampling the Universe above 10 keV - extremesky2009,\\\\ October 13-17, 2009\\\\ Otranto (Lecce) Italy}  \\begin{document} ",
        "introduction": "In the last seven years, the hard X--ray \\emph{INTEGRAL} observatory discovered many new hard X--ray sources \\cite{Bird2007}. In particular, almost 30\\% of the new discovered sources are HMXBs, which were not detected in earlier observations. Among these, \\emph{INTEGRAL} discovered two classes of HMXBs with supergiant companions: the first class is composed of intrinsically highly absorbed hard X--ray sources (e.g. IGR~J16318-4848) \\cite{Filliatre2004}. The members of the second class, called \\emph{Supergiant Fast X-ray Transients} (SFXTs; \\cite{Sguera2005, Negueruela2006}), exhibit outbursts with duration of a few days composed by many flares lasting from minutes to a few hours as discovered by \\cite{Sidoli2008, Romano2008} with the \\emph{Swift} monitoring of 4 SFXTs (IGR~J16479$-$4514, XTE~J1739$-$302, IGR~J17544$-$2619 and AX~J1841.0$-$0536). The behaviour of SFXTs is characterized by a high dynamic range, spanning 3 to 5 orders of magnitude, from a quiescent state at $10^{32}-10^{33}$~erg~s$^{-1}$ up to the peak luminosity during outbursts of $10^{36}-10^{37}$~erg~s$^{-1}$. \\emph{Swift} also discovered that SFXTs display a fainter flaring activity with luminosities of $10^{33}-10^{34}$~erg~s$^{-1}$.  Many different mechanisms have been suggested to explain the SFXT behaviour: \\cite{Grebenev2007, Bozzo2008} proposed that the high dynamic range shown by SFXTs is due to transitions across the neutron star centrifugal barrier produced by a change in the donor wind density. In particular, \\cite{Bozzo2008} proposed that what distinguishes SFXTs from persistent HMXBs with supergiant companions is that SFXTs host magnetars with large spin period ($\\sim 10^3$~s). Another possibility involves the presence of an equatorial wind component denser than the polar wind, and inclined with respect to the orbital plane of the compact object. In this framework, the outburst is produced when the compact object crosses the equatorial wind component and, consequently, accretes more matter \\cite{Sidoli2007, romano09b}. This mechanism has been successfully applied to the SFXT IGR~J11215-5952, which shows periodic outbursts ($P_{orb}\\approx 165$~days, \\cite{Sidoli2007}). \\cite{intZand2005} proposed that the flaring activity in SFXTs is due to the sudden accretion of dense blobs of matter composing the supergiant wind. In the framework of the clumpy wind model proposed by \\cite{Oskinova}, \\cite{Negueruela2008} suggested that different orbital separations could play a role in the different behaviour of SFXTs and persistent HMXBs. Persistent HMXBs have a small orbital period, with a distance supergiant-compact object $<2$ stellar radii, while in SFXTs the compact object orbits the companion at larger distances. ",
        "conclusions": ""
    },
    "1004/1004.1258_arXiv.txt": {
        "abstract": "{The first stars to form in the Universe may be powered by the annihilation of weakly interacting dark matter particles. These so-called dark stars, if observed, may give us a clue about the nature of dark matter. Here we examine which models for particle dark matter satisfy the conditions for the formation of dark stars. We find that in general models with thermal dark matter lead to the formation of dark stars, with few notable exceptions: heavy neutralinos in the presence of coannihilations, annihilations that are resonant at dark matter freeze-out but not in dark stars, some models of neutrinophilic dark matter annihilating into neutrinos only and lighter than about 50 GeV. In particular, we find that a thermal DM candidate in standard Cosmology always forms a dark star as long as its mass is heavier than $\\simeq$ 50 GeV and the thermal average of its annihilation cross section is the same at the decoupling temperature and during the dark star formation, as for instance in the case of an annihilation cross section with a non--vanishing $s$-wave contribution.} ",
        "introduction": "The first stars, also referred to as Population III stars, are the first luminous objects in the Universe. They contribute to the reionization of the interstellar medium, they provide the heavy elements (metals) that eventually become part of the later generations of stars, and they may be the seeds of the very massive black holes observed in quasars.  It was shown in \\cite{dark_stars_prl,Freese:2008hb,Spolyar:2009nt} that the first stars to form in the Universe may be powered by the annihilation of dark matter particles instead of nuclear fusion. These dark-matter powered stars, or dark stars for short, constitute a new phase of stellar evolution. Besides the assumption that dark matter is made of weakly interacting massive particles (WIMPs) that can self-annihilate into ordinary particles, three conditions are necessary for the formation of a dark star.   The first condition is that the density of dark matter at the location of the (proto)star must be high enough for dark matter to efficiently and rapidly annihilate into ordinary particles, releasing a large amount of energy. The first stars are believed to form at the center of dark matter halos when the Universe was young (redshift $z\\sim 10$-50) and denser than today. Not only the dark matter density at the center of those early halos was high, but as the baryonic gas contracted into the first protostars, more dark matter was gathered around the forming object by the deepening of the gravitational potential (gravitational contraction). Cosmological parameters and the evolution of the gas density completely determine the resulting density of dark matter at the location of the first protostars. Analytic and numerical evaluations \\cite{dark_stars_prl,Freese:2008hb,Natarajan:2008db} lead to a resulting density which is high enough to satisfy the first condition for the formation of a dark star.   The second condition is that a large fraction of the energy released in the dark matter annihilation must be absorbed in the gas that constitutes the (proto)star. The fraction $f_Q$ of annihilation energy deposited into the gas depends on the nature of the annihilation products. Typical products of WIMP annihilation are charged leptons, neutrinos, hadrons, photons, W and/or Z bosons, and Higgs bosons. The latter (W, Z, and Higgs) decay rapidly into leptons and hadrons. The hadrons themselves, which are mostly charged and neutral pions) decay rapidly into charged leptons, neutrinos, and photons (although a small number of stable particles like protons can also be produced). After $\\sim 10^{-8}$ seconds, all unstable elementary particles, including the muon, have decayed away, and only protons, electrons, photons and neutrinos survive. Protons have a large scattering cross section with the protostar medium and are quickly absorbed. Electrons and photons can ionize the medium and/or generate electromagnetic showers. For WIMPs with mass $m \\gtrsim 0.5$ GeV, electromagnetic showers are the dominant process. At the time when the following third condition for a dark star is satisfied, the protostar has a diameter of more than 40 radiation lengths, implying that all the energy released in protons, electrons, and photons is absorbed inside the protostar. Only the fraction of energy carried away by the neutrinos is lost for what concerns a dark star.  The third condition for the formation of a dark star is that the heating of the (proto)star gas arising from the dark matter annihilation energy must dominate over any cooling mechanism that affects the evolution of the (proto)star. In \\cite{dark_stars_prl}, it was shown that the dark matter heating rate $Q_{\\rm DM}$, in energy deposited per unit time and unit volume, is given by the expression \\begin{equation} Q_{\\rm DM} = f_Q \\frac{\\svds}{m} \\rho^2 , \\end{equation} where $\\rho$ is the dark matter density inside the (proto)star, which is determined by the cosmological model, and $\\svds$ is the average value of the dark matter annihilation cross section $\\sigma$ times WIMP relative velocity $v$ inside a dark star. To the extent that electromagnetic showers are generated, i.e.\\ $m \\gtrsim 0.5$ GeV, all dark star properties depend on the particle physics model only through the quantity \\begin{equation} f_Q \\frac{\\svds}{m}. \\end{equation}  Ref.~\\cite{dark_stars_prl} fixed the annihilation cross section to $\\svds=3\\times10^{-26}$ cm$^3$/s and examined a range of WIMP masses $m$ from 1 GeV to 10 TeV.  In addition, Ref.~\\cite{dark_stars_prl} assumed $f_Q = 2/3$, based on simulations of neutralino dark matter annihilation in the Minimal Supersymmetric Standard Model (MSSM). For this range of $Q_{\\rm DM}$, they compared the heating and cooling rates along protostar evolution tracks from \\cite{yoshida}, and concluded that there is a time during the evolution of the protostar in which the dark matter heating dominates over all cooling rates. This finding lead to the realization that dark stars may be possible.   In this paper, we examine the possible values of $Q_{\\rm DM}$ for a large selection of particle physics models, and verify if the third condition above is satisfied in these models. We find that not all particle dark matter models lead to the formation of dark stars, although the models that do not form dark stars are either tuned to resonant annihilation or rather artificial.   \\begin{figure} \\begin{center} \\includegraphics[width=0.7\\linewidth]{thermo.eps} \\end{center} \\caption{Condition for the formation of a dark star, in terms of the protostar gas density and temperature. The gray band shows possible evolution tracks of the protostar obtained through numerical simulations in a $\\Lambda$CDM cosmology \\cite{yoshida}. The red lines show critical curves on which the heating rate from dark matter annihilation equals the total cooling rate of the protostar gas. Critical curves are labeled by the value of $f_Q \\svds/m$ in units of cm$^3$ s$^{-1}$ GeV$^{-1}$. A dark star forms at the intersection of a critical line with the gas evolution track. No dark star can form for $f_Q \\svds/m < 10^{-32}$ cm$^3$ s$^{-1}$ GeV$^{-1}$ (solid critical line on the right).} \\label{fig:thermo} \\end{figure}  The restriction imposed by the third condition for dark star formation is best expressed in terms of a condition on the quantity $f_Q\\langle \\sigma v\\rangle_{\\rm ds}/m$. Following \\cite{dark_stars_prl}, we have computed the critical lines in the gas temperature-density plane at which the heating rate from dark matter annihilation equals the total cooling rate. These lines are shown in Figure \\ref{fig:thermo} for a wide range of values of $f_Q\\langle \\sigma v\\rangle_{\\rm ds}/m$, from $10^{-18}$ cm$^3$ s$^{-1}$ GeV$^{-1}$ to $10^{-32}$ cm$^3$ s$^{-1}$ GeV$^{-1}$. Below the latter value, the heating-cooling critical line no longer intersects the thermodynamic track of the protostellar gas, indicated by the gray band obtained through numerical simulations of the formation of the first stars in a $\\Lambda$CDM cosmology \\cite{yoshida}. In other words, for $f_Q\\langle \\sigma v\\rangle_{\\rm ds}/m < 10^{-32}$ cm$^3$ s$^{-1}$ GeV$^{-1}$, the protostar is expected to contract to a regular Population III star powered by nuclear fusion without passing through the dark star phase. At the other side of the $f_Q\\langle \\sigma v\\rangle_{\\rm ds}/m$ range, the critical line reaches a limiting curve given by the vertical line labeled $f_Q\\langle \\sigma v\\rangle_{\\rm ds}/m = 10^{-18}$ cm$^3$ s$^{-1}$ GeV$^{-1}$. Larger values of $f_Q\\langle \\sigma v\\rangle_{\\rm ds}/m$ give the same vertical line. Thus, as expected, if the annihilation rate is large, a protostar passes through the dark star phase. Therefore the third condition for the formation of a dark star is \\begin{equation} f_Q \\frac{\\langle \\sigma v\\rangle_{\\rm ds}}{m} > 1\\times 10^{-32} \\hbox{ cm$^3$ s$^{-1}$ GeV$^{-1}$}. \\label{eq:ds_condition} \\end{equation}  The choice $\\svds=3\\times10^{-26}$ cm$^3$/s in \\cite{dark_stars_prl} was motivated by the assumption that the dark matter WIMPs are produced thermally in the early Universe. That is, that the WIMPs are generated in matter-antimatter collisions at temperatures higher than $T_{\\rm fo} \\sim m/20$, which is the temperature after which WIMP production ``freezes out'' and the comoving WIMP number density remains (approximately) constant. Ref.~\\cite{dark_stars_prl} used the following simple inverse-proportionality relation between the present WIMP density $\\Omega_\\chi$ and the annihilation cross section $\\svfo$ at the time of WIMP freeze-out, \\begin{equation} \\Omega_\\chi h^2 = \\frac{3\\times 10^{-27} {\\rm cm^3/s}}{\\svfo} . \\label{eq:Oh2} \\end{equation} Furthermore, ref.~\\cite{dark_stars_prl} simply assumed that the velocity-averaged annihilation cross section times relative velocities at the time of freeze-out and in a dark star have the same value, $\\svds=\\svfo$. In reality, the relation between $\\Omega_\\chi$ and $\\svfo$ is more complex, and in addition $\\svds$ may differ from $\\svfo$ because $\\sigma v$ may depend sensitively on the WIMP velocity $v$. In this regard, we notice that the average WIMP speed at freeze-out is of the order of \\begin{equation} v_{\\rm fo} \\sim \\sqrt{\\frac{T_{\\rm fo}}{m}}\\sim \\frac{c}{\\sqrt{20}} \\sim7\\times10^{4}{\\rm~km/s}, \\end{equation} while the typical speed of WIMPs in a dark star can be estimated from the orbital velocity \\begin{equation} v_{\\rm ds} \\sim \\sqrt{\\frac{GM}{r}} \\sim {\\rm 30~to~300~km/s}, \\end{equation} namely $\\sim 30$ km/s for a newly-born 1-$M_\\odot$ dark star of 1 AU radius or $\\sim 300$ km/s for a mature 600-$M_\\odot$ dark star of 5 AU radius.  A neutralino in the MSSM provides an example of a more complex relation between $\\Omega_\\chi$ and $\\svfo$. At the same time, it allows the direct evaluation of both $\\svds$ and $f_Q$, and in general it has $\\svds\\ne\\svfo$. Section~\\ref{sec:mssm} explores this case.  Kaluza-Klein dark matter is examined in Section 3, where it is concluded that generically in these models $\\svds$ tends to be larger or comparable to $\\svfo$.  Leptophilic models of dark matter proposed to explain the PAMELA positron excess and the Fermi and HESS cosmic-ray electron-positron data provide another example in which $\\svds$ may not be the same as $\\svfo$. They are examined in Section~\\ref{sec:lepto}.   Finally, we push $f_Q$ down using dark matter particles that annihilate exclusively into neutrinos (``neutrinophilic'' models). In these models, even if annihilation produces predominantly neutrinos that escape the forming star, W- and Z-bremsstrahlung processes may generate enough charged leptons to actually form a dark star. We examine this case in Section~\\ref{sec:neutro}. ",
        "conclusions": "The first stars to form in the Universe may be powered by the annihilation of weakly interacting dark matter particles~\\cite{dark_stars_prl}.  In this paper we explored several popular examples of thermal dark matter models in order to discuss whether they can satisfy the conditions for the formation of a dark star: the neutralino in an effective MSSM scenario; leptophilic models that might explain recent observations in cosmic rays; the KK-photon and the KK-neutrino in UED models; a conservative neutrinophilic model where the dark matter particles annihilate exclusively to neutrinos. We find that in general models with thermal dark matter lead to the formation of dark stars, with few notable exceptions: heavy neutralinos in the presence of coannihilations; annihilations that are resonant at dark matter freeze-out but not in dark stars; neutrinophilic dark matter lighter than about 50 GeV. In particular the discussion of the latter conservative scenario allows us to conclude that a thermal DM candidate in standard Cosmology always forms a dark star as long as its mass is heavier than $\\simeq$ 50 GeV and the thermal average of its annihilation cross section is the same at the decoupling temperature and during the dark star formation, as for instance in the case of a cross section with a non--vanishing $s$-wave contribution.  Therefore, we can conclude that the formation of a first generation of stars powered by dark matter annihilation is an almost inevitable consequence of thermal dark matter when a standard thermal history of the Universe is assumed and if the mechanism of Ref.~\\cite{dark_stars_prl} is at work. So a dark star is always there whenever there is thermal dark matter."
    },
    "1004/1004.2714_arXiv.txt": {
        "abstract": "Numerical simulations of dispersive turbulence in magnetized plasmas based on the Hall-MHD description are presented, assuming spatial variations along a unique direction making a prescribed angle with the ambient magnetic field. Main observations concern the energy transfers among the different scales and the various types of MHD waves, together with the conditions for the establishment of pressure-balanced structures. For parallel propagation, Alfv\\'en-wave transfer to small scales is strongly inhibited and rather feeds magnetosonic modes, unless the effect of dispersion is strong enough at the energy injection scale. In oblique directions, the dominantly compressible character of the turbulence is pointed out with, for quasi-transverse propagation, the presence of conspicuous kinetic Alfv\\'en waves. Preliminary simulations of a Landau fluid model incorporating relevant linear kinetic effects reveal the development of a significant plasma temperature anisotropy leading to recurrent instabilities. ",
        "introduction": "Turbulence in magnetized plasmas remains a main issue in the understanding of the dynamics of media such as the solar corona, the interstellar medium, the solar wind or the planet magnetosheaths. In the solar wind for example the turbulent cascade extends much beyond the ion Larmor radius. One of the questions concerns the spectrum of the magnetic fluctuations that  displays  a power-law behavior on a broad range of wavenumbers,  with a conspicuous change of slope  near  the inverse ion gyroradius \\citep{LSNMW98,GR99,AMM06,SGRK09}. This effect is often  associated with the influence of wave dispersion, induced by the Hall current \\citep{GSRG96,Gal06,ACVS06,GB07,SCPVS07,ShSh09}, but could also result from a superposition of cascades of kinetic Alfv\\'en waves and ion entropy fluctuations,  as suggested by studies based on the gyrokinetic formalism \\citep{HDC08,HCD08,SCD09}.  At scales large compared with the ion inertial length or the ion Larmor radius, the usual MHD description provides a satisfactory description of regimes where, due to the presence of a strong ambient field, a dominant effect is the anisotropic energy transfer to Fourier modes with large transverse wavenumbers (see e.g. \\citet{GG97,OM05} and references therein). This suggests that the  dynamics of transverse small scales may be amenable to a reduced MHD description (\\citep{ZM92} and references therein), possibly including Hall current \\citep{GMD08} or, when retaining scales significantly smaller than the ion Larmor radius, to a  gyrokinetic approach \\citep{HCD06,SCD09}. The latter that appears to be very efficient in describing strongly-magnetized near-equilibrium fusion plasmas is still under discussion concerning its applicability to space and astrophysical plasmas \\citep{MSD08}. In the solar wind for example  magnetic fluctuations may be comparable to the ambient field. Furthermore, longitudinal transfer could a priori be non  negligible in a compressible regime,  at scales where  Hall current  and  kinetic effects play a significant role. A weak turbulence theory performed on the Vlasov-Maxwell system was recently developed \\citep{YF08}, showing the existence of a parallel cascade of low-frequency Alfv\\'en waves through a three-wave decay process mediated by ion-sound turbulence, in a regime where wave-particle interactions are neglected. Addressing this issue by direct numerical simulations of the Vlasov-Maxwell equations being still difficult on the present-day computers, the question arises whether a similar cascade can be observed within a fluid model that retains important ingredients of the above theory, such as compressibility and dispersion. As a first step, we address  the  problem within the simplest description provided by Hall-MHD (HMHD) with Ohmic and viscous dissipations, together with a  large-scale external driving acting on the transverse components of the velocity or magnetic field. We specifically concentrate on a  one-dimensional setting  where the variations of the fields are restricted to a direction making a prescribed angle with the ambient magnetic field, a framework that already reveals a manifold of complex dynamical processes that deserve detailed investigations before including additional physical and multidimensional effects. In the case of quasi-transverse propagation, we also present simulations of a model that extends the HMHD by retaining pressure anisotropy, Landau damping and finite Larmor radius effects up to transverse scales significantly smaller than the ion Larmor radius. This approach developed in \\citet{PS07} extends the so-called Landau fluid model initiated in  \\citet{SHD97} for the MHD scales where Landau damping is the only relevant kinetic effect.  The paper is organized as follows. Section 2 briefly reviews the Hall-MHD description and its one-dimensional reduction.  Section 3  concentrates on the case where the dynamics takes place in the direction of the ambient field. The case of oblique propagation is addressed in Section 4. Landau fluid simulations retaining small-scale kinetic effects are reported in Section 5. Our conclusions are summarized in Section 6. ",
        "conclusions": "Although limited to one space dimension, the present study reveals specific aspects of the turbulent dynamics of magnetized plasmas at scales comparable to the ion inertial length, in a regime where the transverse components of the velocity or the magnetic fields are randomly driven. Special attention was paid to the distribution of the energy among the different MHD waves that can be clearly identified from their linear dispersion relation, in spite of  a possible small shift in the temporal spectrum of the computed fields, in situations where the presence of large-scale coherent structures can be viewed as performing a renormalization of the ambient parameters.  A main observation in the case of parallel propagation is the contrast between a large-scale forcing for which the energy is almost entirely transferred to magnetosonic modes with nevertheless a non-negligible transfer to larger scales, and the regime where the driving takes place at scales where dispersion is more efficient, for which, provided the driving is of kinetic type, a direct Alfv\\'enic transfer establishes,  a result qualitatively consistent with the weak-turbulence analysis performed by \\citet{YF08} in the context of the Vlasov equation. In oblique directions, the combined role of dispersion and compressibility leads to a turbulence dominated by the intermediate modes, with an increasing contribution of low frequency kinetic Alfv\\'en waves, together with a faster establishment of total pressure balance, as the propagation angle is increased. Furthermore, the Landau fluid model shades a light on the development of pressure anisotropy resulting in recurrent instabilities.  To conclude, we would like to stress the complexity of the turbulence problem in magnetized plasmas, even within the strongly simplified description provided by one-dimensional HMHD. The usual picture of inertial ranges where energy ``cascades'' progressively from scale to scale at a constant rate turns out to be strongly affected by the predominance of structures and by a strongly fluctuating transfer, making possibly questionable the usual concepts of the classical turbulence theory,  and leading to a dynamics that turns out to be significantly less universal."
    },
    "1004/1004.5134_arXiv.txt": {
        "abstract": "{The emission  mechanism responsible  for the bulk  of energy from radio to X--rays in low ionization emission line regions (LINERs) and  Low  Luminosity Active  Galactic  Nuclei  (LLAGN)  has been  long debated.  Based  on UV  to X--ray  and radio to  UV flux  ratios, some argue  that  LINERs/LLAGN are  a  scaled-down  version  of their  more luminous predecessors Seyfert galaxies.   Others, based on the lack of X--ray short (hours) time--scale  variability, the non detection of an iron line  at 6.4~keV, and the  faint UV emission  compared to typical AGNs, suggest the  truncation of the classical thin  accretion disk in the inner regions of the AGN where a radiatively inefficient accretion flow  (RIAF)  structure   forms.  We  investigate  the  LINER--Seyfert connection  by studying  the  unabsorbed LINER  galaxy  NGC 4278  that accretes  at  a  low  rate  (L$_{bol/Edd}\\sim$7$\\times$10$^{-6}$)  but exhibits a  broad H$\\alpha$ line,  and a point-like nucleus  in radio, optical,  UV  and  X-rays.  We   analyzed  one  XMM-Newton  and  seven \\chandra\\  X-ray observations  of NGC  4278 spread  over a  three year period,  allowing the  study of  the X--ray  variability  at different time-scales  (hours, months, years).   We also  examined the  radio to X-ray spectral energy distribution  to constrain the accretion mode in the  nucleus  of  \\src.    Long  time-scale  (months)  variability  is observed  where  the  flux increased  by  a  factor  of $\\sim$3  on  a time-scale of a  few months and by a factor of  5 between the faintest and the brightest observation  separated by $\\sim$3 years.  During the XMM-Newton observation,  where the highest flux level  is detected, we found a  10$\\%$ flux increase  on a short  time-scale of a  few hours, while the light curves for the different \\chandra\\ observations do not show short  time-scale (minutes to hours)  variability.  A combination of    an   absorbed   power    law   ($N_{H}\\approx10^{20}$~cm$^{-2}$, $\\Gamma=2.2^{+0.1}_{-0.2}$)      plus     a      thermal     component (kT$\\approx0.6$~keV) were able to fit the \\chandra\\ spectra. The \\xmm\\ spectra, where  the highest X--ray  flux is detected, are  well fitted with an  absorbed power--law with no  need for a  thermal component as the emission from the power--law component is dominant. The power--law photon index  is $\\sim2.1$ and the  hydrogen column density  is of the order  of   $10^{20}$~cm$^{-2}$.   Neither   a  narrow  nor   a  broad Fe~K${\\alpha}$ emission line at 6.4~keV  are detected with a 22~eV and 118~eV  upper limits derived  on their  equivalent widths.   We derive optical  fluxes  from archival  \\hst\\  ACS  observations and  detected optical variability on time--scales of  years.  For the first time for this  source, thanks  to  the  optical/UV monitor  on  board \\xmm,  we obtained simultaneous  UV and X-ray flux  measurements. We constructed SEDs  based on  simultaneous  or quasi  simultaneous observations  and compared them to LINER, radio--loud, and radio--quiet quasar SEDs.  We find that at a low X--ray flux the \\src\\ SED resembles that of typical LINER sources where the radio  to X--ray emission can be considered as originating from  a jet  and/or RIAF, whereas  at a high  X--ray flux, \\src\\ SED  is more like  a low luminosity Seyfert  SED.  Consequently, \\src\\ could exhibit both  LINER and Seyfert nuclear activity depending on the strength of its X--ray emission.} ",
        "introduction": "\\label{sec:intro}  Low-ionization  nuclear  emission  line  regions (LINERs)  were  first identified by \\citet{heckman80aap} as a class of galaxies with optical spectra dominated  by emission lines from low  ionization species. The ionization  mechanism is  yet  poorly known  and  could be  explained, either           in            terms           of           starbursts \\citep{alonso-herrero00ApJ:starburstinliners}  or,  more consistently, being  due  to  a   low  luminosity  active  galactic  nuclei  (LLAGN) \\citep{terashima00ApJ:liners,  ho93ApJ:linerAGN}.   This  latter  idea rises  from  the  detection  of  broad  H$\\alpha$\\  components  in  an important  fraction   of  LINER  sources   \\citep[noted  as  LINER~1.9 (hereinafter LINERI) objects][]{ho97apjs},  and/or a point--like UV or X-ray  source at  the  nucleus.   If the  luminosity  scales with  the accretion  rate, the  study  of  the LINER  nucleus  define a  supreme pattern  for probing  low accretion  rate physics  around supermassive black holes (SMBHs).  How similar are these LINERI sources to classical luminous Seyfert and quasar galaxies?   Based on observational properties,  the weakness or absence   of   a   big   blue   bump   feature   at   UV   wavelengths \\citep{ho08aa:review}       usually        detected       in       AGN \\citep{malkan82apj:uvexcSey1,sanders1989ApJ:bbbseyqua,koratkar99pasp:seybbb}, the   lack  of  a   broad  Fe~K$\\alpha$   emission  line   at  6.4~keV \\citep{terashima02apjs:LLAGNASCA,ptak04apj:ngc3998},  except  for  the peculiar LINER  NGC~1052 \\citep{brenneman09ApJ:ngc1052}, and  last but not  least, the  lack of  short time-scale  (minutes to  hours) X--ray variability  have  been  attributed  to  an  intrinsic  difference  in LINERsI/LLAGN central  engine as  opposed to normal  Seyfert galaxies. One  proposed   scenario  is   that  accretion  in   LINERsI/LLAGN  is radiatively inefficient,  advection dominated compared  to the typical geometrically thin optically thick accretion disks present in luminous AGN \\citep{narayan05apss:adaf}.  However,   more  recently  \\citet{maoz05apj:linervarUV}   revealed  UV variability in the  nuclei of a sample of 17  LINERs observed with the \\hst.   By combining  these results  with non--simultaneous  radio and X--ray   observations,  \\citet{maoz07MNRAS}   demonstrated   that  the UV/X--ray luminosity  ratios are similar to those  of Seyfert~1 nuclei and pointed  out that LLAGN may  be a scaled--down  version of Seyfert galaxies where a  thin accretion disk exists.  This  idea is supported by   \\citet{pianmnras10}  who  studied   simultaneous  UV   to  X--ray observations of four  LINER nuclei observed with the  XRT and the UVOT on--board the  \\textsl{Swift} telescope.  They  discovered short--time scale  (half a day)  X--ray variability  in two  of their  objects and showed that the UV to X--ray  flux ratios are consistent with those of more luminous AGN.   The elliptical  galaxy \\src\\ \\citep[distance of  16.7~Mpc;][ scaled to $H_{0}=70$~km~s$^{-1}$~Mpc$^{-1}$]{tonry01apj:dist}  has  been studied extensively at different wavelengths.   At radio wavelength, a compact non-thermal   radio  source  has   been  detected   at  6   and  18~cm \\citep{jones84apj:4278radobs}.      \\citet{4278nagar05aap}    reported two-sided radio emission on subparsec  scales in the form of twin jets emerging from  a central compact  component ($T_{B}=1.5\\times10^9$~K). \\citet{ho97apjs}  classified  \\src\\  as  a  type 1.9  LINER  from  the definite    detection    of    a    broad   H$\\alpha$    line,    with $log~F(H\\alpha)=-13.07$~ergs~cm$^{-2}$~s$^{-1}$,  supporting  the  AGN nature of the nuclear engine.  Moreover, an unresolved compact nuclear source     has    been    detected     in    \\hst\\     WFPC2    images \\citep{capetti00aa:hst}.  Finally,  \\citet{ho01apjl}, after studying a \\chandra\\ 1.4 ks snapshot taken in  April 2000, gave \\src\\ a class (I) X-ray     morphology,     showing  a  dominant     nuclear     source. \\citet{terashima03apj:rloud}  fit  the  0.5--8 keV  spectrum  obtained during  this same  \\chandra\\ snapshot  with  a power  law modified  by absorption   and   found   a   2--10~keV   corrected   luminosity   of $9.1\\times10^{39}$~ergs~s$^{-1}$.   \\citet{gonzalezmartin09aa}  find a 0.5--10~keV  corrected  luminosity of  $5.6\\times10^{39}$~erg~s$^{-1}$ after  studying a  $\\sim$100~ks \\chandra\\  observation taken  in March 2006.   The mass  of a  black hole in  the nucleus  of \\src, derived  from the M-$\\sigma$        relation        \\citep{termaine02ApJ:Mbh},        is $3.09\\pm0.54\\times10^{8}$~M$_{\\odot}$ \\citep{wang03MNRAS:Mbh,chiaberge05ApJ:Mbh}.     This   leads    to   a $L_{Edd}\\approx3.9\\times10^{46}$~ergs~s$^{-1}$;      assuming     that $L_{Edd}=1.25\\times10^{38}~(M_{BH}/M_{\\odot})$~ergs~s$^{-1}$.      This implies a very low L$_{X}/$L$_{Edd}$ ratio of 2$\\times10^{-7}$.   In this  paper we  report a  timing and X--ray  spectral study  of the LINER galaxy \\src\\  observed with \\xmm\\ and \\chandra.   We describe in Sect.~2  the X-ray observations  and the  data reduction.   Timing and spectral X--ray results as well as optical (\\hst/ACS) and UV (\\xmm/OM) results  are  presented  in  Sect.~3.   In Sect.~4,  we  describe  the construction of the \\src\\ spectral energy distribution.  We discuss in sect.~5 our  results in the context of  the LINER--Seyfert connection. The main results are summarized in Sect.~6. ",
        "conclusions": "We  have studied  in detail  seven  X-ray observations  for \\src,  six ACIS-S  \\chandra\\  observations   plus  one  \\xmm\\  observation.   The observations  covered  a  three  year period  allowing  an  extensive variability study. No short time--scale (minutes to hours) variability is detected during the \\chandra\\ X--ray observations, while during the \\xmm\\ observation, where the highest  flux level is observed, the flux increases by  a factor  of 10\\%\\  in a 1  hour period.   A significant months  time--scale  variability  is  observed between  the  different \\chandra\\  observations and between  \\chandra\\ and  \\xmm\\ observations where  the flux  increased by  a  factor of  $\\sim$3 on  a few  months time--scale  and  by a  factor  of 5  between  the  faintest and  the brightest  observations  separated by  $\\sim$3  years.  We checked  the validity of increasing variability with decreasing luminosity detected in Seyfert  galaxies \\citep{nandra97apj:variance}, \\src,  like most of LINER     sources,    does     not    follow     the     same    trend \\citep{ptak98apj:variance}.   However, based on  the anti--correlation relation between variability and the  mass of the black hole with some exceptions     whenever     the     accretion     rate     is     high \\citep{reeves02mnras:pds456},   no  short  time--scale   variation  is expected in \\src\\ as  it harbors a $\\sim3\\times10^8$~M$_{\\odot}$ black hole  and  accretes  at  low  rate.  However,  the  short--time  scale variability  detected during  the  \\xmm\\ observation  might suggest  a change in the  central engine of \\src\\ at  that point.  No variability is detected in  the hardness ratio during the  observations being soft in all cases.   The best spectral  fit found for the six  \\chandra\\ observations is an absorbed power  law plus a thermal component.   The intrinsic hydrogen column  density  affecting  the  power  law  has  an  upper  limit  of $6.7\\times10^{20}$~cm$^{-2}$, while the  average photon index $\\Gamma$ is $2.2_{-0.2}^{+0.1}$.   The temperature of the  thermal component is $0.63_{-0.04}^{+0.05}$~keV.   Because of  the  \\xmm\\ moderate  angular resolution  compared  to  \\chandra,  we subtracted,  in  a  10\\arcsec\\ radius,  the contribution  from  X-ray binaries  and diffuse  emission calculated  from  the  \\chandra\\   observations.   We  find  that  the power--law  emission is dominant  over the  whole 0.5--8~keV  band and therefore no thermal component is needed.  The hydrogen column density is $(3.8\\pm0.8)\\times10^{20}$~cm$^{-2}$  affecting a power  law with a photon  index  $\\Gamma$  of  2.1.  No Fe~K${\\alpha}$ emission line at 6.4~keV is detected with a 22~eV upper limit on its equivalent width.  We        measured       a        UV        flux       of        about $1.2\\times10^{-15}$~ergs~cm$^{-2}$~s$^{-1}$~\\AA$^{-1}$  from the \\xmm\\ optical  monitor  observation  of  \\src.   The  advantage  of  such  a measurement is  the simultaneity  with the X--ray  observation crucial for  comparing  such  sources  with  normal galaxies  as  \\src\\  shows significant variability on months time--scales.  We calculated optical fluxes  coming from \\hst\\  ACS observations  in two  different filters (F475W and F850LP).  These fluxes are 2 to 4 times higher than optical fluxes, coming from the \\hst\\  WFPC2 F555W and F814W filters, observed $\\sim12$ years  earlier.  Our optical  fluxes are contemporary  to the \\chandra\\  fluxes  coming  from   the  7081  observation  as  the  two observations are separated by less than two months.  In order to determine the origin of the emission mechanism responsible for  the bulk  of  energy from  radio  to X--rays  and  to assess  the geometry of  the central engine in  \\src\\, we plotted  the SED derived from simultaneous  observations (radio\\footnote{We consider  the radio flux constant,  and include it in the  simultaneous observations since it  does   not  show  significant  variability  on   a  several  years time--scale         (see         sect.~5.2).}$+$UV$+$X-rays        and radio$+$optical$+$X--rays) with  the SED of a sample  of LINER sources \\citep{eracleous10:linersed}  along with radio--quiet  and radio--loud quasars \\citep[fig.~8;  adopted from][]{nemmen10:lineradaf}.  As \\src\\ appears to fluctuate in  optical, UV and X--rays on month--timescales, it  appears that  \\src\\ SED  is consistent  with LINER  SEDs  when the X--ray emission is weak and therefore we suppose a jet/RIAF origin for the radio to X--ray bulk of energy. On the other hand, when the X--ray emission is  important, the  \\src\\ SED  is similar to  the SED  of low luminosity Seyfert  galaxies and a thin accretion  disk is responsible for the bulk  of energy from UV to X--rays.   Hence, we consider \\src\\ to  alternate  between LINER-like  and  Seyfert-like nuclear  activity depending on the strength of its X--ray emission."
    },
    "1004/1004.2378_arXiv.txt": {
        "abstract": "There have been a number of attempts to measure the expansion rate of the universe at high redshift using Luminous Red Galaxies (LRGs) as ''chronometers''. The method generally assumes that stars in LRGs are all formed at the same time. In this paper, we quantify the uncertainties on the measurement of H(z) which arise when one considers more realistic, extended star formation histories. In selecting galaxies from the Millennium Simulation for this study, we show that using rest-frame criteria significantly improves the homogeneity of the sample and that $H(z)$ can be recovered to within 3\\% at $z \\sim 0.42$ even when extended star formation histories are considered. We demonstrate explicitly that using Single Stellar Populations to age-date galaxies from the semi-analytical simulations provides insufficient accuracy for this experiment but accurate ages are obtainable if the complex star formation histories extracted from the simulation are used. We note, however, that problems with SSP-fitting might be overestimated since the semi-analytical models tend to over predict the late-time star-formation in LRGs.  Finally, we optimize an observational program to carry out this experiment. ",
        "introduction": "\\label{sec:intro}   An important constraint on cosmological models is the evolution of the Hubble parameter, which is defined as: \\begin{equation} H(z)=-\\frac{1}{1+z} \\frac{dz}{ dt} . \\label{eq:hz} \\end{equation} Most tests of cosmology use measures of the luminosity distance or the angular diameter distance, which include an integral over $H(z)$, but measuring $H(z)$ directly can potentially provide a more direct constraint on cosmological parameters (Jimenez \\& Loeb 2002).  $H(z)$ can be determined at high redshifts by measuring the time-interval, $\\Delta t$, corresponding to a redshift interval, $\\Delta z$, to obtain an approximation of the derivative in equation (1).  A number of authors have attempted to use this method (Ferreras, Mechiorre, \\& Silk 2001, Jimenez \\& Loeb 2002, Jimenez et al. 2003, Capozziello et al. 2004, Ferreras, Melchoirre, \\& Tocchini-Valentini 2003, Simon et al. 2005, Dantas et al. 2007, Verkhodanov, Parijskij, \\& Starobinsky 2007, Samushia et al. 2009) to track the evolution of $H(z)$ as a function of redshift, and place constraints on cosmological parameters.   Most recently, Stern et al. (2010) extended the studies of Jimenez et al. (2003) to a larger, more homogeneous sample to make the measurement of H(z) from LRGs.  Jimenez \\& Loeb (2002) suggested using spectroscopic age-dating of Luminous Red Galaxies (LRGs) to measure $dz/dt$. LRGs are galaxies selected from the Sloan Digital Sky Survey (SDSS, York et al. 2000) by apparent magnitude to trace the evolution of a homogeneous, volume-limited sample of red galaxies (Eisenstein et al. 2001). Jimenez et al. (2003) used LRGs from the SDSS to measure $H_o=69 \\pm 12$ km s$^{-1}$ Mpc$^{-1}$ (the Hubble parameter at z=0). Simon et al. (2005) used the same method to find $H(z)$ as a function of redshift for $0<z<2$. In both cases, the authors assume that LRGs are drawn from the same parent population with the bulk of their stars formed at high redshift and fit single-burst equivalent ages to the galaxies.  Although there is evidence that LRGs make up a fairly homogeneous population of galaxies consisting mostly of old populations, Roseboom et al. (2006) have shown that this is not completely true.  Furthermore, Barber et al. (2007) find that these galaxies have a range of formation ages and stellar histories. Jimenez and Loeb (2002) note that galaxies would not have formed at exactly the same time and explore what effect this would have on one's ability to discriminate between different models for w(z). They do not, however, discuss the consequences of star formation being extended over a period in each galaxy (they assume all stars form at exactly the same time in a galaxy). This extended star formation history (SFH) needs to be considered when one defines and measures an age for a galaxy. In Jimenez et al (2003), the question is addressed in a little more detail: they simulate LRG spectra and add a contribution from a small fraction of stars formed at late times; they then age-date the galaxies and conclude that the age \"edge\" is not significantly different from the theoretical expectation. They do not attempt to quantify what uncertainties arise as a result of the extended star formation.  Given the importance of correct estimates of uncertainties on the measured H(z), we explore, in this paper, the uncertainties that arise from LRGs having extended and varied SFHs. We use LRGs identified in the Millennium Simulation (MS, Springel et al. 2005) since the only way to reliably determine uncertainties is in a simulation where the SFH and input cosmology are known.  The MS is currently one of the most detailed simulations of the evolution of galaxies and large-scale structure and has been tuned to reflect observational data.  Almeida et al. (2008) studied LRGs in simulations and showed that the luminosity function and the clustering of LRGs are reasonably well-modelled in their semi-analytical models of galaxy evolution, which also reproduce many other observables.  The parameters of the de Lucia et al. (2006) semi-analytical model have also been successfully tuned to match the luminosity, colour, and morphology of local elliptical galaxies, the more massive candidates of which would be LRGs.  On the other hand, galaxies are not modelled perfectly in the simulations: there are indications that model galaxies assemble later than observations may suggest (Collins et al. 2009) but more reliable estimates of stellar mass are required.  It should be noted that these models overestimate the small amount of evolution seen in the properties of LRGs (Wake et al. 2006, 2008) and likely over predict the amount of growth occurring in these sources (Masjedi et al. 2008). Specifically for LRGs, Barber et al. (2007) find similar star formation rates, but younger ages than the de Lucia models.  In addition, properties of local galaxies can be reproduced with different models for galaxy formation (e.g. Kere{\\v s} et al. 2009). However, the simulations provide viable models of extended SFHs, allowing us to quantify associated uncertainties and optimize analysis methods before applying them to observational data.   In this paper, we optimize the selection of LRGs based on their properties in the MS and examine how well $H(z)$ can be recovered from their mass-weighted age.  To explore and quantify the uncertainties associated with extended SFHs, we then investigate how well $H(z)$ can be determined, if the mass-weighted age of these LRGs were measurable, using three different methods.  In practice, mass-weighted ages are not directly measurable and accurate age-dating using LRG spectra is potentially a bigger challenge for the cosmic chronometer approach. We carry out an initial study of age-dating, and its implications for $H(z)$ measurements, using the stellar population models of Bruzual \\& Charlot (BC2003, 2003) to synthesize spectra. In contrast to fitting only single stellar population models (SSPs) (e.g. Simon et al. 2005) or dual-burst population models (Jimenez et al. (2003)) to determine the LRG spectral ages, we make a comparison of SSP results to those using more complicated SFHs.  In a companion paper (Olivier et al. 2009), we delve much deeper into issues of age-dating galaxy spectra.  The outline of the paper is as follows: In \\S \\ref{sec:sims} we discuss the extraction of LRGs from the MS. In \\S\\ref{sec:cosmo} we investigate how well $H(z)$ is recovered in the MS using LRGs identified in the simulation. In \\S \\ref{sec:model}, we describe the modelling and age-dating of LRG spectra. In \\S\\ref{sec:obs}, we discuss an observational program to measure $H(z)$. We summarize our results and conclusions in \\S \\ref{sec:conc}.  Throughout this work, we adopt the cosmology used in de Lucia et al. (2006) of $\\Omega_m=0.25$, $\\Omega_{\\Lambda}=0.75$, and $h=0.73$. ",
        "conclusions": "\\label{sec:conc}  In this paper, we have shown that within the self-consistent Universe of the Millennium Simulation that Luminous Red Galaxies, when selected in an appropriate manner, can be used as cosmic chronometers.  By selecting the galaxies from their rest-frame properties, we find that we can create a more homogeneous sample of objects than the apparent magnitude cuts used by Eisenstein et al. (2001).  For galaxies at $z\\sim0.3-0.6$, we find that cuts of $M_V<-23$ and $(B-V)_0>0.81$ select a sample of galaxies with similar star formation histories and formation redshifts in the MS.  The galaxies selected in our final cut do show very similar star formation histories with very few of the galaxies showing any star formation since $z\\sim1.7$.  The ages also show something akin to downsizing (Cowie et al. 1996): the oldest galaxies are in the most massive halos.  The distribution of the ages are very similar with each showing a small tail towards younger age, but with the very strong peak at a single age.  However we do note limitations in the use of the models of de Lucia et al. (2006).  They are able to reproduce some of the major trends seen in local galaxy populations and seem to have similar number density as the SDSS out to $z\\sim0.5$.  However, the model does have some limitations in reproducing some of the observed properties of LRGs while alternative models produce equally acceptable fits to the data.  One such model explored here (Bower et al. 2006) was unsatisfactory in matching the observed number density of LRGs defined by an absolute magnitude cut at a $z\\sim0.5$.  Even though these galaxies have extended SFHs, they can be used as cosmic chronometers to recover the cosmology used in the MS.  Using only galaxies selected from two redshifts, $H(z)$ can be calculated to a precision of less than $3\\%$ using three different methods.  All three methods (averaging the ages, calculating a fit to the distribution, and comparing the ages of pairs of galaxies) were able to recover the cosmology used in de Lucia et al. (2006).  In \\S \\ref{sec:model}, we showed that SSPs were not sufficient to accurately recover the ages for individual galaxies, which has also been shown for other samples (Maraston et al. 2009, Trager \\& Somerville 2009).  Despite the relatively simple nature of the SFHs in Fig. \\ref{fig:civ}, the calculated SSP ages were dominated by the most recent burst of star formation.  If we use the average star formation history from the MS, we are able to replicate the properties of the model spectra, but in the next paper, we look in far greater detail at the modelling and fitting of LRG spectra.  However, the semi-analytic models likely overestimate the extent of star formation in LRGs, and SSP models may provide better fits than indicated here.  Finally, we estimated the required time to complete the project using RSS on SALT.  If systematics in the age can be controlled, $H(z)$ at $z\\approx 0.42$ can be calculated to $3\\%$ in a total of $\\sim 180$ hours, which includes observing overheads.  It is likely that tighter constraints could be put on the evolution of $H(z)$ with additional data that could be obtained while making these observations. In addition to constraining $H(z)$, it would contribute a wealth of detailed information about the evolution of the most massive galaxies at intermediate redshift.  Throughout this work, we have highlighted a number of assumptions that we have made.  Even with our refined selection, the star formation histories of these galaxies are not perfectly homogeneous as assumed by Jimenez \\& Loeb (2002), but $H(z)$ can still be calculated to an accuracy better than $3\\%$.   Other improvements on the measurement can be made by selecting a larger or more homogeneous sample as done by Stern et al. (2010).  Uncertainties in the star formation history of simulated galaxies introduce far smaller errors than the contribution associated with age-dating the galaxies. In our next paper, we explore age-dating in more detail and note that increased uncertainties resulting from age-dating may lead to increased observing requirements."
    },
    "1004/1004.5244_arXiv.txt": {
        "abstract": "Observations and theory both suggest that star clusters form sub-virial (cool) with highly sub-structured distributions.  We perform a large ensemble of $N$-body simulations of moderate-sized ($N=1000$) cool, fractal clusters to investigate their early dynamical evolution.  We find that cool, clumpy clusters dynamically mass segregate on a short timescale, that Trapezium-like massive higher-order multiples are commonly formed, and that massive stars are often ejected from clusters with velocities $> 10$~km~s$^{-1}$ (c.f. the average escape velocity of 2.5~km~s$^{-1}$).  The properties of clusters also change rapidly on very short timescales. Young clusters may also undergo core collapse events, in which a dense core containing massive stars is hardened due to energy losses to a halo of lower-mass stars.  Such events can blow young clusters apart with no need for gas expulsion.  The warmer and less substructured a cluster is initially, the less extreme its evolution. ",
        "introduction": "\\label{sec:intro}  Most stars appear to form in star clusters \\citep{lada03,lada09,portegies_zwart10} and so star formation is inextricably linked with star cluster formation.  Recent advances in observations and theory have allowed us to construct a basic picture of cluster formation in which clusters form dynamically cool (sub-virial), and highly substructured.  Clusters form in highly turbulent molecular clouds.  These clouds are highly substructured, containing dense clumps and filaments \\citep{williams99,williams00,carpenter08} which are presumably formed by the decay of supersonic turbulence \\citep{maclow04,ballesteros-paredes07}. Stars and small stellar groups form in these dense regions and are unsurprisingly observed to have a high degree of substructure when young \\citep{larson95,elmegreen00,testi00,cartwright04,gutermuth05,allen07,schmeja06,schmeja08}, clustering in young stellar groups has even been observed in the SMC \\citep{schmeja09} and LMC \\citep{bastian09}.  The same behaviour is seen in simulations of cluster formation \\citep{klessen00,klessen01,bate03,bonnell03,bate09,offner09}, including in comparison tests between AMR and SPH techniques \\citep{federrath10}.  Clusters are observed to lose their substructure as they evolve, becoming smooth and roughly spherical \\citep{cartwright04,schmeja06,schmeja08}. \\citet{goodwin04} have shown that substructure can only be erased in clusters if the clusters are initially cool (sub-virial) (see also Maschberger et al. 2010). Both observations of pre-stellar cores \\citep{belloche01,andre02,walsh04,peretto06,kirk07} and stars \\citep{peretto06,proszkow09} show that they indeed appear to be sub-virial, a property also found in simulations of cluster formation (Klessen \\& Burkert 2000; Offner et al. 2009; Maschberger et al. 2010).  Following Allison et al. (2009b) we conduct $N$-body simulations of a large number of initially sub-virial, fractal star clusters.  In Section~\\ref{sec:init} we describe our simulations. In Section~\\ref{sec:results} we describe our main results, specifically the early onset of dynamical mass segregation and interesting `post-collapse' evolution. In Section~\\ref{sec:disc} we discuss the results, and we summarise and conclude in Section~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  Observations and theory both strongly suggest that the initial conditions of star clusters are cool and clumpy.  Therefore, we have conducted a large number of $N$-body simulations of the early ($< 4$~Myr) evolution of clusters with virial ratios of $Q = 0.3, 0.4$ and $0.5$ (where $0.5$ is virialised), and fractal dimensions of $D = 1.6, 2.0, 2.6$ and $3.0$ (where $3.0$ is roughly a uniform density sphere) with radii of $1$~pc and $1000$ members (ie. total masses of $\\sim 500$~M$_\\odot$).   In our simulations, all  members were initially single stars selected from a Kroupa IMF, and the simulations lasted for 4~Myr (so not requiring stellar evolution to be included).  We study the evolution of the star clusters with a particular emphasis on the level of mass segregation.  We measure mass segregation using a minimum spanning tree to provide a quantitative measure of mass segregation and which is not biased by clumpy underlying mass distributions (Allison et al. 2009a).  This study follows that of Allison et al. (2009b), in which we showed that clusters with $Q=0.3$ and $D=1.6$ (i.e., very cool and extremely clumpy) undergo collapse to a short-lived but extremely dense core. This core can dynamically mass segregate the most massive stars down to a few M$_\\odot$ via two-body encounters before re-expanding. The re-expansion is partially driven by the increase in the velocity dispersion of the low-mass stars caused by two-body encounters.  Our main results may be summarised as follows:  $\\bullet$  The depth of the collapse, and so the degree of mass segregation depend on both the initial virial ratio, $Q$, and the degree of substructure, $D$.  Low-$Q$ and low-$D$ clusters can collapse to a denser state and so mass segregate more than high-$Q$, high-$D$ clusters.  $\\bullet$ Whilst there is a general trend of increasing mass segregation with lower-$Q$ and lower-$D$ the inherently stochastic nature of fractals means that {\\em statistically} identical clusters may undergo very different evolution.  $\\bullet$  Many features are extremely short-lived, and young clusters can change rapidly and violently.  Observations only provide a snapshot of the evolution of a cluster and it is dangerous to draw conclusions about the past and future state of a cluster from a single snapshot (see also Bastian et al. 2008).  $\\bullet$ Young clusters can undergo core collapse.  Early dynamical mass segregation establishes a core of massive stars (see also Pflamm-Altenburg \\& Kroupa 2006) and a halo of lower-mass stars. Energy loss from the core to the halo can drive the formation of massive, hard binaries and Trapezium-like multiple systems.  The heating of the halo and the hardening and subsequent decay of the central multiple systems can dynamically destroy a cluster within a few~Myr.  $\\bullet$  The interactions of massive stars in the core of a young cluster can cause the ejection of even very massive stars at velocities in excess of $20$~km~s$^{-1}$.  This may help explain ejections from the ONC (see also Pflamm-Altenburg \\& Kroupa 2006).  $\\bullet$ The early evolution of cool, clumpy clusters is rapid, violent, and extreme.  The densities of clusters (and hence their crossing and relaxation times) can change by orders of magnitude during the first few~Myr of their existence.  Thus the currently observed properties of young clusters are just a snapshot in the life of these clusters and extreme care must be taken in inferring the past history or future evolution of clusters from a single snapshot (see also Goodwin \\& Bastian 2006; Bastian et al. 2008; Allison et al. 2009b).  That young clusters are mass segregated down to a few M$_\\odot$, but not below, is due to the short-lived dynamical mass segregation phase which is able to mass segregate only the most massive stars.  That the ONC has an unstable high-order multiple containing four of the most massive stars can be explained by its dynamical formation during the dense phase.  The ejection of high-mass stars from clusters can occur during the dense phase, or afterwards from the decay of higher-order massive multiple systems."
    },
    "1004/1004.3979_arXiv.txt": {
        "abstract": "We study the dynamics of states perturbatively expanded about a harmonic system of loop quantum cosmology, exhibiting a bounce. In particular, the evolution equations for the first and second order moments of the system are analyzed. These moments back-react on the trajectories of the expectation values of the state and hence alter the energy density at the bounce. This analysis is performed for isotropic loop quantum cosmology coupled to a scalar field with a small but non-zero constant potential, hence in a regime in which the kinetic energy of matter dominates.  Analytic restrictions on the existence of dynamical coherent states and the meaning of semi-classicality within these systems are discussed.  A numerical investigation of the trajectories of states that remain semi-classical across the bounce demonstrates that, at least for such states, the bounce persists and that its properties are similar to the standard case, in which the moments of the states are entirely neglected. However the bounce density does change, implying that a quantum bounce may not be guaranteed to happen when the potential is no longer negligible. ",
        "introduction": "\\label{sec:intro}  Loop quantum cosmology \\cite{LivRev} imports mathematical constructions and results obtained in the general setting of quantum geometry realized in loop quantum gravity \\cite{Rov,ThomasRev,ALRev} into quantum cosmology. As a result, the quantum representation is unitarily inequivalent to the one used in Wheeler--DeWitt quantizations of cosmological models, a formal difference entailing also dynamical changes in the behavior of models for the very early universe. Loop quantum cosmology thus allows one to explore the physical implications of quantum geometry and its unconventional space-time structures in a tractable setting.  One of the main changes implied by loop quantum gravity is the use of holonomies \\cite{LoopRep}, objects of the form $\\exp(i {\\rm curv})$ replacing the usual curvature expressions ``${\\rm curv}$'' (extrinsic and intrinsic curvature, or space-time curvature components) in the classical equations. The form of the curvature functional ``${\\rm curv}$'' appearing in holonomies depends on the specific quantization of the Hamiltonian constraint used \\cite{QSDI} as well as on details of the reduction to a symmetric context within quantum gravity \\cite{SymmRed}; however, within isotropic cosmological models it is a function only of the scale factor and its time derivative. Moreover, neither is the full theory of loop quantum gravity uniquely and unambiguously formulated yet, nor is the reduction from general physical states to isotropic ones fully understood at a quantitative level. Despite such ambiguities, several general conclusions can be inferred from mathematical properties of the curvature--holonomy replacement, most importantly related to the boundedness of holonomy expressions in contrast to the curvature components. For instance, the Friedmann equation relates extrinsic curvature $a\\dot{a}$ of spatial slices in a dynamical isotropic universe to the energy density of matter. If extrinsic curvature in the equation is replaced by a bounded function of a certain form, boundedness of the energy density automatically follows in spite of its classical divergence at the big bang singularity.  Replacing curvature expressions by bounded functions is not simply done by hand but is motivated by quantum theory. Quantization also means that one is dealing with states rather than classical geometries subject to the usual evolution equations. Fundamentally, the use of holonomies primarily implies that the dynamics of quantum states of universe models is governed not by the differential Wheeler--DeWitt equation \\cite{DeWitt,QCReview}, but by a difference equation \\cite{cosmoIV,IsoCosmo}. These discrete dynamics turn out to be non-singular \\cite{Sing,BSCG}: wave functions extend uniquely across the classically singular big bang. But while the use of holonomies and the discreteness of the difference equation implied by it are important for this result, detailed properties of the transition depend on various kinds of quantum effects that the use of the wave function introduces. In particular, in addition to quantum geometry effects such as discreteness, there are in general quantum back-reaction effects that characterize deviations of the quantum evolution of expectation values from the classical evolution caused by a changing wave function. Simple intuitive boundedness results on energy densities can be established only if quantum geometry effects can be shown not to be overpowered by potentially adverse quantum back-reaction effects. In loop quantum cosmology, this competition of different quantum effects has not been explored sufficiently in generic regimes. (Examples in Wheeler--DeWitt quantum cosmology have shown the sensitivity of some quantum effects related to singularity avoidance to state properties \\cite{GaussianBohmQC,NonSingBohmQC,WDWBounce}.)  Typically attention has been focused on harmonic models, in which the quantum back-reaction is entirely absent. This has the advantage of being analytically tractable and has been shown to generically lead to bouncing solutions. The robustness of these results for models that deviate from these special cases has only begun to be explored and in this paper we further develop a systematic perturbation theory approach to this question~\\cite{EffAc,Karpacz}. Key results reported here include the interplay of reality conditions (physical normalization of states) and uncertainty relations, which in non-harmonic cases turns out to be important for consistent evolution, numerical or otherwise. By numerical studies, we demonstrate the sensitivity of bounce properties such as the density to quantum effects. Also the high sensitivity of evolution through the bounce on initial values, observed for harmonic models based on an analysis of dynamical coherent states \\cite{BeforeBB,Harmonic}, is qualitatively confirmed in our more general setting. Although the leading-order perturbations used here do not allow quantum dynamical corrections large enough to (potentially) remove the bounce induced by the holonomy modification, the sensitivity we find appears sufficiently strong to warrant caution about extrapolations of the kinetic-dominated behavior to models with significant matter potentials. ",
        "conclusions": "Numerical investigations in loop quantum cosmology have so far only considered semi-classical states of Gaussian type, peaked around some particular value \\cite{APS,APSCurved,NegCosNum}. The expectation value of such states has been shown to closely follow the classical trajectories, except at very high densities where quantum gravity effects, but primarily those of quantum geometry, produce a bounce. Implicit in these analyzes is the idea that a general semi-classical state is well characterized by its expectation values and small spreads alone and that the evolution of the spread and other moments does not affect the evolution of the peak.  Here we have considered a moment expansion of a state around such exactly solvable Gaussian states as a way of probing the consequences of relaxing this assumption. In principle one should consider all orders of the moment expansion, however in order to make the system tractable, we have truncated the series at second order.  Already this expands the configuration space from $4$-dimensional ($(V,\\phi)$ and their momenta) to $18$-dimensional (given by the eight independent variables in Eq.~(\\ref{eq:variables}), $\\phi$ and their momenta).  The inclusion of moments introduces sets of non-trivial uncertainty relations that must be satisfied by the state {\\it at all points along its trajectory}. Dynamical coherent states are those for which all the uncertainty relations are exactly saturated, throughout the state's evolution. We have shown that the conditions necessary to ensure a state is dynamical coherent, effectively reduces the dimension of the state's phase space i.e.\\ the conditions form a constraint surface on which the states, and their trajectories, lie. For the zero potential case, there is only one condition that needs to be met to ensure that a state which initially saturates the uncertainty bounds will be a dynamical coherent state (i.e.\\ will continue to saturate the bounds). For a system with a non-zero potential one has four independent constraints (in addition to the uncertainty bounds) which need to be met, just to ensure that the first derivatives of the uncertainty relations vanish. While requiring that the higher derivatives also vanish is likely to introduce even more constraints. This makes it extremely difficult (and probably impossible) to find exact dynamical coherent states for systems with a non-zero scalar field potential. While conceptually this may be a problem, especially for considering evolution from the infinite past ($\\phi\\rightarrow -\\infty$) or cyclic models, in practice we are used to dealing with quantum states that are only approximately dynamically coherent, such as particle wave-packets in standard quantum mechanics. Provided the states are approximately dynamically coherent for a sufficiently long period of time ($\\phi$) they can, for all practical purposes, be considered semi-classical.  We have investigated the consequences of the moments on the presence of a bounce, for states that remain semi-classical (in the sense that all their moments are small compared to the expectation values and that the states obey the uncertainty conditions). In all cases the bounce was found to be present and the energy density was changed only marginally from the standard case (i.e.\\ the case in which the moments are neglected).  Thus we found that for the (large number of) cases investigated, the bounce occurs at approximately the same density, a feature not at all unexpected due to the setup of our approximations. The fact that there were small deviations from the standard bound is however significant, especially since we show that it is possible for the bounce density to be sometimes larger than the standard case. We restricted ourselves to considering only states that remained highly semi-classical before, after and during the bounce and still found that the moments noticeably back-react on the trajectories of the expectation values. It may be that there are states which are semi-classical at large scales, but which become dominated by the evolution of their moments at small scales (i.e.\\ become highly non-semi-classical) and for such states the presence of a bounce is not guaranteed by current results in the literature."
    },
    "1004/1004.4132_arXiv.txt": {
        "abstract": "We report on the discovery of X-ray--emitting O-Ne-Mg-rich ejecta in the middle-aged Galactic O-rich supernova remnant Puppis~A with {\\it Chandra} and {\\it XMM-Newton}.  We use line ratios to identify a low-ionization filament running parallel to the northeastern edge of the remnant that requires supersolar abundances, particularly for O, Ne, and Mg, which we interpret to be from O-Ne-Mg-rich ejecta. Abundance ratios of Ne/O, Mg/O, and Fe/O are measured to be $\\sim$2, $\\sim$2, and $<$0.3 times the solar values.  Our spatially-resolved spectral analysis from the northeastern rim to the western rim otherwise reveals sub-solar abundances consistent with those in the interstellar medium.  The filament is coincident with several optically emitting O-rich knots with high velocities.  If these are physically related, the filament would be a peculiar fragment of ejecta.  On the other hand, the morphology of the filament suggests that it may trace ejecta heated by a shock reflected strongly off the dense ambient clouds near the northeastern rim. ",
        "introduction": "X-ray emission from evolved supernova remnants (SNRs) is dominated by the interstellar medium (ISM) swept-up by expanding SN ejecta. Therefore, evolved SNRs had been considered to be more suitable for studies of high-Mach number ($> 10$) shock physics, ISM/shock interactions, or the ISM itself, and less suitable for studies of the SN explosion/nucleosynthesis.  However, over the last two decades, the {\\it ASCA}, {\\it ROSAT}, {\\it Chandra}, and {\\it XMM-Newton} X-ray observatories have uncovered a number of ejecta features in evolved SNRs such as the Vela SNR (e.g., Aschenbach 1995; Tsunemi et al.\\ 1999), the Cygnus Loop (e.g., Miyata et al.\\ 1998; Katsuda et al.\\ 2008a), and several LMC (e.g., Hughes et al.\\ 2003; Borkowski, Hendrik, \\& Reynolds 2006) and SMC SNRs (e.g., Park et al.\\ 2003; Hendrick, Reynolds, \\& Borkowski 2005).  Puppis~A is one of the brightest SNRs in the X-ray sky and shows both ISM/shock interactions and SN ejecta.  The large extent ($\\sim$50$^{\\prime}$ in diameter) as well as the high surface brightness allow us to study detailed structures in the remnant.  It has been suggested that the expanding shell of Puppis~A is interacting with dense HI and CO clouds from the eastern (E) rim to the northern (N) rim, based on its asymmetric X-ray surface brightness (e.g., Petre et al.\\ 1982) and the alignment of these dense clouds with the edge of the remnant (e.g., Dubner \\& Arnal 1988).  {\\it Chandra} observations of the most prominent cloud-shock interaction at the E rim can be compared directly to scaled laboratory simulations to infer a mature interaction of age 2000-4000 years (Hwang, Flanagan, \\& Petre 2005; Klein et al.\\ 2003).  Thus, Puppis~A is an excellent astrophysical laboratory where we can study detailed structures of cloud-shock interactions.  On the other hand, evidence of SN ejecta has been detected in Puppis~A.  From optical observations, Winkler \\& Kirshner (1985) discovered an O-rich fast-moving filament (named $\\Omega$ filament) in the northeastern (NE) portion of the remnant.  Subsequently, several O-rich optical knots have been found near the $\\Omega$ filament, with proper-motion vectors suggesting constant expansion from a common center with a dynamical age of 3700$\\pm$300\\,yrs (Winkler et al.\\ 1988).  Recently, signatures of ejecta have been found in X-ray observations as well.  Hwang, Petre, and Flanagan (2008) noticed Si-rich ejecta localized in the NE quadrant, based on a {\\it Suzaku} survey of this SNR [see also Hwang, Petre, \\& Flanagan (2008) for a review of metal-rich ejecta indicated by earlier X-ray observations]. Katsuda et al.\\ (2008b) also reported the discovery of fast-moving metal-rich ejecta knots with blueshifted lines in the NE portion of the remnant, based on {\\it XMM-Newton} observations.  So far, the composition of these ejecta features all indicate that Puppis~A originated from a core-collapse SN.  This is consistent with the presence of a central compact object (CCO) (Petre et al.\\ 1996).  Here, we report the discovery of an O-Ne-Mg-rich ejecta filament in the NE quadrant of the remnant from {\\it Chandra} and {\\it XMM-Newton} observations.  We also discuss the origin of the filament. ",
        "conclusions": "The X-ray surface brightness map and the Ne K line ratio map reveal a distinct filamentary feature (i.e., the NE filament) and two knots in the NE portion of Puppis~A.  Our spectral analysis shows that these features are rich in O, Ne, and Mg, with abundance ratios of Ne/O, Mg/O twice solar and that of Fe/O less than solar.  The super-solar ($>$1000) fitted metal abundances of the NE filament clearly show that its origin is likely O-Ne-Mg-rich ejecta.  The relative abundances of the O-Ne-Mg-rich ejecta can then be used to estimate the progenitor mass.  We compare the relative abundances measured in the filament with those predicted by the theoretical calculations (Rauscher et al.\\ 2002) for different progenitor masses of 15\\,M$_\\odot$, 25\\,M$_\\odot$, 30\\,M$_\\odot$, and 35\\,M$_\\odot$.  We find that the data agree with those calculations for the O-Ne-Mg-rich layers in progenitor stars of $\\lesssim$ 25\\,M$_\\odot$ masses.  We note that the same conclusion can be derived by using the relative abundances of the $\\Omega$ filament (Katsuda et al.\\ 2008b), since the relative abundances of Fil-1/2 are consistent with those of the $\\Omega$ filament.  This is also consistent with the progenitor-mass range of 15\\,M$_\\odot$--25\\,M$_\\odot$ inferred from comparisons of relative abundances of metal-rich (especially Si-rich) ejecta in the NE portion with several nucleosynthesis models (Hwang, Petre, \\& Flanagan 2008).  We caution however that we only measured abundances of just a small fraction of the total ejecta, and so the estimate of the progenitor mass is still not conclusive.  To try to identify the contact discontinuity and the reverse shock, we performed radially-resolved spectral analysis from the NE rim (including the O-Ne-Mg-rich NE filament) to the W rim.  The spectral analysis revealed that all the spectra except for the NE filament show sub-solar metal abundances ($\\sim$0.5 times the solar values) whose relative abundances are consistent with the solar values within a factor of 2---consistent with those from {\\it Suzaku} observations covering the same region (Hwang, Petre, \\& Flanagan 2008).  This result suggests that the X-ray emission is dominated by the ISM {\\it everywhere} except for the NE filament.  Interpretations of the results could be (1) the outer ejecta heated by the reverse shock have already cooled, while the inner ejecta have not yet been heated by the reverse shock, (2) metal-rich ejecta could not be detected (except for the NE filament) because of their faintness compared with the ISM, or (3) most of the ejecta has not yet been heated by the reverse shock, although this last case seems unlikely for a remnant of several thousand years' age such as Puppis A.  With the current data, this remains an open question.  In the area we investigated, the only region that shows evidence of SN ejecta is the small region where the O-Ne-Mg-rich NE filament and knots are located.  This result may indicate that they are peculiar fragments of ejecta inside and/or outside the SNR.  We then note that the NE filament is coincident with three\\footnote{In the vicinity of these knots, recent optical observations identified some more knots whose proper motion vectors are similar with each other (Garber et al.\\ 2010).} optical O-rich fast-moving knots (Winkler et al.\\ 1988).  Figure~\\ref{fig:X_opt} shows the proper motion vectors of the O-rich knots on the X-ray image, in which we see that all the vectors, $\\sim$0$^{\\prime\\prime}$.18\\,yr$^{-1}$ ($\\sim$1900\\,km\\,sec$^{-1}$ at a distance of 2.2\\,kpc) toward the NE, are similar to each other.  If the NE filament is physically associated with these optical knots, we expect the same proper motions for the NE filament as well.  On the other hand, it is interesting to note that the NE filament runs parallel to the NE edge of the remnant (see, Figs.~\\ref{fig:3col} and \\ref{fig:Ne_ratio} (b)).  This suggests that the morphology of the filament is related to the NE edge of the SNR where the forward shock is suggested to be interacting with dense clouds (e.g., Reynoso et al.\\ 1995).  Then, we can speculate that the filament traces ejecta heated by either a strong reflected shock off the dense ambient clouds in the NE rim or a reverse shock strongly developed in this portion of the SNR.  Given that we see evidence of ejecta only in the NE filament, we may interpret the inner side of the filament to be the reflection shock and the outer side to be a boundary between the O-Ne-Mg-rich ejecta and either the ISM or ejecta rich in lighter elements.  A problem for this scenario is that the expected temperature signature at the reflected shock is not seen (Fig.~\\ref{fig:rad_res}); the temperature should be higher in the twice-shocked outer regions than in the singly shocked inner regions (e.g., Hester et al.\\ 1994).  It is also uncomfortable that the filament spans only a portion of the NE edge, whereas it appears that the entire NE edge is interacting with dense clouds (e.g., Reynoso et al.\\ 1995).  However, these signature might be less clear due to projection effects and/or possible density inhomogeneities of the ejecta; the denser ejecta will tend to have the lower temperature, as the temperature is partly determined by the condition that the ram pressure, $\\rho v^2$, be approximately constant.  One way to test the reflection shock scenario would be to measure the proper motion of the filament, as it should then be slower than the freely expanding ejecta.  We tried measuring proper motions of the NE filament, based on X-ray images of {\\it Einstein} in 1979, {\\it ROSAT} in 1992, 1993, and 1994, and {\\it Chandra} in 2006.  Although we indeed found a possible motion, it is difficult to derive conclusive results, because there is no point-like source in the {\\it Einstein} and {\\it ROSAT} images to be used for image registrations, and the brightness of the NE filament seems to change during these observations.  In this respect, future {\\it Chandra} observations of the NE filament are desired to firmly detect the possible proper motion and brightness change, and to conclude the origin of the NE filament."
    },
    "1004/1004.0131_arXiv.txt": {
        "abstract": "Understanding the nature of galactic populations of double compact binaries (where both stars are a neutron star or black hole) has been a topic of interest for many years, particularly the coalescence rate of these binaries. The only observed systems thus far are double neutron star systems containing one or more radio pulsars. However, theorists have postulated that short duration gamma-ray bursts may be evidence of coalescing double neutron star or neutron star-black hole binaries, while long duration gamma-ray bursts are possibly formed by tidally enhanced rapidly rotating massive stars that collapse to form black holes (collapsars). The work presented here examines populations of double compact binary systems and tidally enhanced collapsars. We make use of \\textsc{binpop} and \\textsc{binkin}, two components of a recently developed population synthesis package. Results focus on correlations of both binary and spatial evolutionary population characteristics. Pulsar and long duration gamma-ray burst observations are used in concert with our models to draw the conclusions that: double neutron star binaries can merge rapidly on timescales of a few million years (much less than that found for the observed double neutron star population), common envelope evolution within these models is a very important phase in double neutron star formation, and observations of long gamma-ray burst projected distances are more centrally concentrated than our simulated coalescing double neutron star and collapsar Galactic populations. Better agreement is found with dwarf galaxy models although the outcome is strongly linked to the assumed birth radial distribution. The birth rate of the double neutron star population in our models range from $4-160~$Myr$^{-1}$ and the merger rate ranges from $3-150~$Myr$^{-1}$. The upper and lower limits of the rates results from including electron capture supernova kicks to neutron stars and decreasing the common envelope efficiency respectively. Our double black hole merger rates suggest that black holes should receive an asymmetric kick at birth. ",
        "introduction": "\\label{s:intro}  Pulsars, magnetic oblate spheriods of nuclear densities $20-30$ kilometers in diameter, have been found rotating at speeds of up to almost one thousand times a second (Hessels et al. 2006; see also Galloway 2008) in some of the most exotic settings in the known Universe. For example, some pulsars are found within X-ray binaries (Liu, van Paradijs \\& van den Heuvel 2007) and others with compact object companions in close binaries (Hulse \\& Taylor 1975) thought to be emitting gravitational radiation (see Landau \\& Lifshitz 1951; Paczy\\'{n}ski 1967; Clark \\& Eardley 1977). Recently, the number of known binary pulsars has been rapidly increasing (Lorimer et al. 2006a; Galloway et al. 2008). There are now in excess of $100$ Galactic disk binary pulsars. This includes rotation powered pulsars (radio pulsars: ATNF Pulsar Catalogue Manchester, Hobbs, Teoh \\& Hobbs 2005\\footnote{http://www.atnf.csiro.au/research/pulsar/psrcat/}) and accretion powered pulsars (X-ray binary pulsars: Liu, van Paradijs \\& van den Heuvel 2007; Galloway 2008; Galloway et al. 2008). More than $20$ of these are accreting from a range of stellar masses and companion types (Galloway 2008), $9$ are thought to orbit another neutron star (van den Heuvel 2007; Stairs 2008), while others ($> 70$) dwell in detached systems with white dwarfs companions (ATNF Pulsar Catalogue, Manchester et al. 2005) and possibly main sequence (MS) companions (Champion et al. 2008). From observations of pulsars in such systems it is possible to constrain computational modelling of binary evolutionary phases that are general to many non-pulsar related systems including, but not limited to, tidal evolution, Roche-lobe overflow and common envelope (CE) evolution. It is the most rapidly rotating pulsars that best constrain uncertainties in the theory related to these processes. Such work has been attempted in Kiel et al. (2008) and Kiel \\& Hurley (2009: KH09), where models of the complete Galactic pulsar population have been made. However, we note that these models are yet to include selection effects which are critical when interpreting the results of pulsar surveys (Taylor \\& Manchester 1977; Oslowski et al. 2009).  Short gamma-ray bursts provide an alternative method to study the physics of compact stars. Here compact stars are considered to be the most compact of remnants: neutron stars (NSs) and black holes (BHs). We define a double compact binary (DCB) to be any combination of these compact objects within a binary system (without any limit on the orbital period) and close DCBs are those systems with orbital periods of less than a few days. The DCBs that merge within the assumed $10~$Gyr age of the Galaxy (i.e. during our simulations) are defined as coalescing DCBs. It is postulated that gamma-ray emission is produced during the coalescence of these systems and that radiation of gravitational waves occurs during the preceding in-spiral phase (Clark \\& Eardley 1977). In particular it is thought that short gamma-ray bursts are produced by coalescing double neutron star systems (Paczynski 1986). Gamma-ray burst observations are very interesting in themselves, however, DCB systems offer other observational features of importance. Not only can many tests of general relativity be performed but they offer insights into a host of observable phenomena (e.g. Taylor, Fowler \\& McCulloch 1979; Lyne et al. 2004). The formation of close DCBs requires two stars of sufficient mass to interact gravitationally, triggering mechanisms to decrease the separation between them during their stellar lifetimes. What seems to be the most important binary evolutionary mechanism in forming close DCBs is the common-envelope phase (Paczy\\'{n}ski 1976) where the evolution following the first supernova (SN) generally requires at least one such event. The modelling of the CE phase is associated with much uncertainty (see, for example, Dewi, Podsiadlowski \\& Sena 2006; Belczynski et al. 2007a) and  will be discussed further in the following section. Interestingly, Voss \\& Tauris (2003) found that $1/3$ of BH-NS DCB systems can be formed via the direct-SN mechanism (Kalogera 1998) and therefore a CE phase is not required in all cases. Adding further uncertainty to the mix Brown (1995) suggested that unless the initial mass ratio is very close to unity NSs spiraling-in within a CE should always accrete enough matter to collapse and form BHs.  This work examines population characteristics of DCBs, including predictions of double neutron star (NS-NS) distributions and correlations between orbital properties and location. The results are extended to detail both long and short gamma-ray bursts (GRBs) and their progenitors, modelling tidally enhanced collapsars and coalescing DCBs. In particular, projected distances from the host galaxy of model GRBs are compared to observations. No detailed conclusions from direct comparison to observations are attempted, as in Belczynski, Bulik \\& Rudak (2002) who account for redshift and different galaxy masses. Nor do we consider the evolution of systems in globular clusters (see e.g. Ivanova et al. 2008; Sadowski et al. 2008). Section~\\ref{s:bpbk} briefly outlines the population synthesis tool used in this body of work. The results are spread over Sections~\\ref{s:DCBNS-NS}, \\ref{s:coal} and \\ref{s:grb}, which examine the bound DCB and NS-NS populations (even if they go on to eventually coalesce), coalescing DCB populations and GRB populations, respectively. ",
        "conclusions": "This work investigated the stellar, binary and Galactic kinematical evolutionary features of double compact binary systems and possible short gamma-ray burst (coalescing NS-NSs) and long gamma-ray burst (tidally influenced collapsars) objects. The main conclusions from this investigation are summarised here: \\begin{itemize}  \\item NS-NS systems are typically more kinematically active and have the greatest scale height of all DCBs, eclipsing the BH-BH population especially when BHs are assumed not to receive velocity kicks at birth. Including electron capture supernovae into models alters this somewhat. When assuming electron capture supernovae occur and that BHs receive kicks the BH-BH population has a comparable scale height to NS-NSs.  \\item We find the double neutron-star formation rate ranges between $4~$Myr$^{-1}$ and $162~{\\rm Myr}^{-1}$. The lower value is set by assuming $\\alpha_{\\rm CE} = 1$, the upper value by assuming EC SNe occurs. Our model without EC SNe and with $\\alpha_{\\rm CE} = 3$ gives a double neutron-star formation rate of $38~{\\rm Myr}^{-1}$.  \\item Our model double black-hole formation rate ranges between $42~{\\rm Myr}^{-1}$ and $820~{\\rm Myr}^{-1}$. The lower value arises when we assume BHs receive kicks, the upper value arises when we do not assume BHs receive kicks.  \\item If NS-NSs merge it is likely they do so within a time scale of a few million years and with a merger rate that ranges between $3~$Myr$^{-1}$ and $154~{\\rm Myr}^{-1}$ for our models. The double neutron-star merger rates closely reflect their birth rates.  \\item BH-BH systems typically merge more slowly than NS-NS systems and this difference is enhanced when BHs do not receive kicks. The range of merger rates for double black holes is $25~{\\rm Myr}^{-1}$ to $107~{\\rm Myr}^{-1}$. The most limiting factor is the addition of SN kicks to BH formation, while assuming $\\alpha_{\\rm CE} = 1$ results in a merger rate of $56~{\\rm Myr}^{-1}$. The maximum double BH merger rate arises when we assume they do not feel kicks at birth. Such a high merger rate suggest that at least one merger event should have been detected with LIGO. The null detection of gravitational waves by LIGO gives credence to the latest findings that BHs should receive kicks at birth.  \\item Common envelope evolution is reconfirmed to be a very important process for DCB formation. In particular, merger time scales are sensitive to the number of CE phases that occur within the intervening time between SN events. Decreasing the efficiency of the CE process decreases the typical merger time scale and final number of NS-NS systems. The CE phase is also important in shaping the final NS-NS eccentricity-orbital period distribution.  \\item We find good agreement of the shape of eccentricity-orbital period distribution when compared to observations but poor agreement with observations on the relative number of high eccentricy to low eccentricy systems. Including electron capture SN evolution into population synthesis models does not rectify the situation, as has been suggested previously.  \\item We find poor agreement between the projected distance of long gamma-ray bursts from their host galaxy and the projected distances of NS-NS systems within a Milky Way model.  \\item We find much better agreement between observations and models when the model galaxy mass and size is scaled down by a factor $\\alpha_{\\rm G} = 0.01$. Owing to the short life times of our model systems (both short gamma-ray burst and long gamma-ray burst) the final projected distances depend heavily on the assumed radial birth distribution.   \\end{itemize}"
    },
    "1004/1004.2244_arXiv.txt": {
        "abstract": "We analyze the line-of-sight baryonic acoustic feature in the two-point correlation function $\\xi$ of the Sloan Digital Sky Survey (SDSS) luminous red galaxy (LRG) sample ($0.16<z<0.47$). By defining a narrow line-of-sight region, $r_p<5.5$\\hmpcii, where $r_p$ is the transverse separation component, we measure a strong excess of clustering at $\\sim 110$\\hmpcii, as previously reported in the literature. We also test these results in an alternative coordinate system, by defining the line-of-sight as $\\theta<3^\\circ$, where $\\theta$ is the opening angle. This clustering excess appears much stronger than the feature in the better-measured monopole. A fiducial $\\Lambda$CDM non-linear model in redshift-space predicts a much weaker signature. We use realistic mock catalogs to model the expected signal and noise. We find that the line-of-sight measurements can be explained well by our mocks as well as by a featureless $\\xi=0$. We conclude that there is no convincing evidence that the strong clustering measurement is the line-of-sight \\bafii. We also evaluate how detectable such a signal would be in the upcoming  Baryon Oscillation Spectroscopic Survey LRG volume (BOSS). Mock LRG catalogs ($z<0.6$) suggest that: ({\\it{i}} ) the narrow line-of-sight cylinder and cone defined above probably will not reveal a detectable acoustic feature in BOSS; ({\\it{ii}} ) a clustering measurement as high as that in the current sample can be ruled out (or confirmed) at a high confidence level using a BOSS-sized data set; and ({\\it{iii}} ) an analysis with wider angular cuts, which provide better signal-to-noise ratios, can nevertheless be used to compare line-of-sight and transverse distances, and thereby constrain the expansion rate $H(z)$ and diameter distance $D_\\mathrm{A}(z)$. ",
        "introduction": "The baryonic acoustic feature % serves as an important tool for our understanding of the evolution of the universe  (\\citealt{peebles70a}). Originating from plasma sound-wave residuals that came to a near stop at the end of the baryon drag epoch ($z_{d}\\sim 1010$), it has the potential to serve as a cosmic {\\it{standard ruler}}, which, in turn, can help us measure the expansion of the universe (\\citealt{hubble31}, \\citealt{blake03}, \\citealt{seo03}, \\citealt{hu03a}, \\citealt{linder03a}, \\citealt{glazebrook05a}).   In galaxy clustering measurements, this feature is a peak of over-density at separations of $s\\sim100$\\hmpc in the two point correlation function $\\xi$, which results in an oscillatory feature % in % the power spectrum P$(k)$.  \\Epaper were the first to detect this feature, using the clustering of $\\sim 44,000$ luminous red galaxies (LRGs) from the Sloan Digital Sky Survey (SDSS;  \\citealt{york00a}). Their measurement of the angle-averaged $\\xi$, which is commonly referred to as the {\\it{monopole}}, has the power to constrain a combination of the Hubble expansion rate $H(z)$ and the angular diameter distance $D_\\mathrm{A}(z)$.  To measure $H$ and $D_\\mathrm{A}$ separately, one would like to probe the \\baf independently along the line-of-sight and transverse directions (\\citealt{matsubara04}). Measurements of these potentially promising methods are currently strongly compromised by shot noise and sample variance limitations, due the large scale nature of the feature. For this reason, most studies have focused on measuring and analyzing the baryonic acoustic feature in the angle averaged $\\xi$ (\\citealt{martinez08}, \\citealt{cabre09i}, \\citealt{labini09}, \\citealt{sanchez09a}, \\Kpaperalt), P($k$) (\\citealt{cole05a}, \\citealt{tegmark06a}, \\citealt{hutsi06a}, \\citealt{percival07a}, \\citealt{percival09b}, \\citealt{reid09a}) and the projected two-point function of photo-$z$ samples (\\citealt{padmanabhan07b}, \\citealt{blake07}) in the SDSS and and Two Degree Field Galaxy Redshift Survey (\\citealt{colless03a}) galaxy samples.  \\GpaperA (hereafter referred to as \\Gpaperii), however, claim to have measured the line-of-sight baryonic acoustic feature. Using $\\sim 77,000$ LRGs from the SDSS Data Release 6 (DR6, \\citealt{adelman08a}) and $\\sim 100,000$ from DR7 (\\citealt{abazajian09a}) they report a significant detection of a feature at $\\pi\\sim 110$\\hmpc in the line-of-sight direction, within a projected distance of $r_p<5.5$\\hmpcii.  The clustering excess they focus on appears much stronger than expected from the \\baf according to the concordance $\\Lambda$CDM model. The galaxies are observed in {\\it{redshift-space}}, as opposed to real-comoving-space. In redshift-space ($z$-space) the line-of-sight peak in the non-linear correlation function is expected to smear heavily due to velocity dispersion.  Furthermore, the whole correlation function should appear negative due to the strong squashing effect (\\citealt{kaiser87}) at scales $s>50$\\hmpc in the line-of-sight direction.  These predictions suggest that if the sharp strong positive measurement obtained by \\Gpaper is the real feature, it would require a physical explanation.  Magnification bias has been proposed to increase clustering at the feature scales. This effect results from gravitational lensing modifying the spacial distribution of high redshift objects (\\citealt{turner84a}, \\citealt{hui07a} and references within). \\cite{yoo2009a} and \\Gpaper examine this effect for the redshift-space $\\xi$ at $z=0.35$. Both studies agree that the magnification effect is anisotropic having the strongest impact on the line-of-sight. \\Gpaper show that a model with magnification performs slightly better than without ($2\\sigma$ level).  We do not include an analysis of magnification bias in this study, but show that the line-of-sight clustering does agree well with a fiducial $\\Lambda$CDM model without a magnification bias.     In particular, \\cite{miralda09} argues, using pair-count statistics (based on data analyzed by \\Gpaperii), that the clustering excess is not significant, and should not be regarded as a detection of the \\bafii. We concur with that conclusion here.  The purpose of this study is to revisit the line-of-sight clustering signal in the SDSS LRG sample, examine its reliability and predict the signal and its uncertainties obtainable in the much larger volume and denser Baryonic Oscillation Spectroscopic Survey sample (BOSS; \\citealt{schlegel09a}).  We measure the line-of-sight clustering $\\xi$ at sfcales of $40-200$\\hmpc, finding results similar to that obtained by \\Gpaperii. We predict $z$-space (as well as real-space) signals and uncertainties for SDSS-sized volumes, by using very realistic light-coned mock galaxy catalogs which are based on fiducial $\\Lambda$CDM models.    In \\S\\ref{datamocks} we briefly explain the data and mock catalogs used for analysis. In \\S\\ref{a2pcf_analysis} we present the anisotropic $\\xi$ clustering and the coordinate systems used throughout the study. In \\S\\ref{dr7dim_analysis} we analyze the line-of-sight clustering of \\qvlii, and in \\S\\ref{dr7full_analysis} we perform a similar analysis on the larger DR7-Full, and directly compare results with \\Gpaperii. We examine the significance of the strong line-of-sight clustering signal in  \\S\\ref{interpretation} by applying a Jeffreys scale to compare model fits to data performed here and in \\Gpaperii. In \\S\\ref{boss} we predict the line-of-sight measurement expected from the BOSS sample, along with a detailed comparison of the signal-to-noise ratios of the three volumes discussed here. In \\S\\ref{boss2} we vary the definition of line-of-sight to wider wedges, to show that BOSS may be used to disentangle $H(z)$ and $D_\\mathrm{A}(z)$.  In the following, all calculations assume a flat $\\Lambda$CDM model. When converting data redshifts to comoving distances, we assume a present day matter density $\\Omega_{M0}$=$0.25$, and define $H_{0}=100 h$ \\kms$\\mathrm{Mpc}^{-1}$. ",
        "conclusions": "In this paper we demonstrate that the claim of a {\\it significant detection } in the line-of-sight \\baf in the SDSS LRG sample by \\Gpaper is unjustified. We perform a similar analysis to theirs, and obtain similar results. The main difference in our interpretation, as we elaborate in \\S\\ref{interpretation}, is that we  use a more conservative criterion regarding whether we detect a feature. We find that the data agrees very well with a $\\Lambda$CDM redshift-space non-linear model tested here, which does not contain a clear line-of-sight feature due to its low signal-to-noise ratio. We also find that physical line-of-sight models tested by us and \\Gpaper do not  out-perform a null $\\xi=0$ model, indicating no clear evidence of a line-of-sight \\bafii. The BOSS survey, which has just begun, will have the statistical power to rule out (or confirm) this strong clustering excess at high significance (Figure \\ref{horizon_run}), though not to usefully detect the \\baf in such a narrowly defined line-of-sight measurement. By using broader angular bins, BOSS will be able to independently measure \\baf along the line-of-sight and transverse directions.  We examine two different volumes in the SDSS LRG sample (SDSS-II). In the smaller one ($0.16<z<0.36$; \\qvlii; \\S \\ref{dr7dim_analysis}), for which we use very realistic mock catalogs, we find a good agreement ($1\\sigma-1.5\\sigma$) between the line-of-sight observation and a $\\Lambda$CDM model, which does not have an apparent feature even in $\\sim 160$ \\qvl volumes. For our full sample ($0.16<z<0.47$; DR7-Full) we find that a fiducial model agrees within $1.5\\sigma-2\\sigma$.  Figures \\ref{LOS_SDSS_LD}, \\ref{los_dr7full_plot}, \\ref{sn} clearly show that the line of sight clustering excess at $s\\sim 110$\\hmpc in both volumes is dominated by noise.  We confirm the number of pairs estimate of \\cite{miralda09} which they used on \\Gpaper results. The bin of interest is $106.7<s<111.1$\\hmpc and $0^\\circ<\\theta<3^\\circ$, corresponding to the choice in \\Gpaper of $\\Delta \\pi$ and $r_p<5.5$\\hmpcii.  In our case, we count the number of effective pairs, including the effects of weighting (see \\Kpaperalt \\ for details). \\qvl yields $2259$ pairs and DR7-Full yields $3104$. Examining $160$ very realistic mock catalogs, we find that the high \\qvl value of $\\xi=0.085$ in this bin, although unlikely, is consistent with $\\Lambda$CDM (one mock of $160$ has a higher value).  We point out that the strong clustering excess that {\\it happens to be at the correct scale } is {\\it probably } noise, that will be reduced with increase of volume. We see a hint of this effect when the volume increases from \\qvl ($0.66$\\hgpcii) to DR7-Full ($1.6$\\hgpcii) the line-of-sight clustering excess ($\\theta<3^\\circ$) at  $s \\sim 108.9$\\hmpc is decreased from $\\xi \\sim 0.085$ to $\\xi \\sim 0.045$. This is shown also in \\Gpaper in their Figures 13 and 15.  We interpret the difference between both measurements and the expected value from $\\Lambda$CDM as noise.          Assuming correctness of our fiducial cosmology, the predicted BOSS results, presented in Figure \\ref{horizon_run}, show that the strong line-of-sight clustering excess seen in DR7 should be ruled out at a very high confidence level when measuring $\\xi$ in the BOSS LRG sample.  In the left panels of Figure \\ref{sn} we show the low S/N expected from the line-of-sight measurement in BOSS. We also examine by eye all $32$ line-of-sight results (defined as $r_p<5.5$\\hmpcii), and find that $7$ of $32 \\  (\\sim 22\\%)$ mocks do not show a clear line-of-sight \\bafii, where $12$ do show a clear signal $(\\sim 38 \\%)$. This analysis is, of course, subjective, but suggests that the LRGs alone, might not detect the line-of-sight \\baf in the BOSS volume, if the line-of-sight definition is limited to $\\theta<3^\\circ$ or $r_p<5.5$\\hmpcii.  In Figure \\ref{anglecut} we show, however, reason for optimism in disentangling $H(z)$ and $D_{\\mathrm{A}}(z)$ in the future BOSS LRG data. Unsurprisingly, using wider angular cuts yields better signal-to-noise ratio results. The various $\\Delta\\theta$ slices contain line-of-sight and transverse information that may be used to probe $H(z)$ and $D_\\mathrm{A}(z)$.  Throughout this study, we have ignored the effects of any potential reconstruction techniques that might be employed on the BOSS data (\\citealt{eisenstein07}, \\citealt{huff07a}, \\citealt{padmanabhan09a}, \\citealt{noh09a}). Reconstruction may allow better precision in measuring the peak position than indicated here.  We emphasize that our measurements and uncertainty estimates are similar to those carried out previously by \\Gpaperii. We also confirm that the strong clustering signal at the line-of-sight can not be explained by systematics. We check various binning widths in both $\\pi$ and $r_p$ as well as fiber-collision weightings, and do not find significant changes. We do, however, disagree on interpretation of the results.   The main dispute is regarding the use of the strong line-of-sight clustering  in the data as the \\baf to determine H(z) directly. In their Table 3 \\Gpaper show results for $H(z)$ obtained by two different methods: \\begin{enumerate} \\item ``Shape Method\": using the full shape of the line-of-sight clustering with priors from monopole, quadrupole clustering as well  as CMB temperature fluctuations to determine $H(z$)/$H_0$. \\item ``Peak Method\": using the line-of-sight \\baf peak position to determine $H(z)$ directly. \\end{enumerate} While both methods are correct procedures to perform, the second should be considered valid only if the line-of-sight \\baf is convincingly detected.  To explain the insignificance of a detection, we consider the null $\\xi=0$ test. Using a Jeffreys scale, in \\S\\ref{interpretation} we show that the physical models do not out-perform $\\xi=0$. This does not mean, of course, that $\\xi=0$ is the line-of-sight clustering at these scales, but rather indicates that models can not be distinguished significantly and a clear line-of-sight \\baf can not be declared detected using this data set.   Only future surveys can show definitively what the line-of-sight correlation function is in these bins.  In particular, BOSS will be able to do so. In Figure \\ref{horizon_run} we show that BOSS, which is underway, will be able to rule out this detection or verify it. BOSS is due to cover a comoving volume of ($\\sim8.1$\\hgpcii; $0.16<z<0.80$) by 2014. For a comoving volume of $\\sim3.9$\\hgpcii ($0.16<z<0.60$)  and density expected in BOSS, we find that all our $32$ mock realizations have an apparent {\\it monopole} feature (as opposed to SDSS-II volumes, \\Kpaperalt).  Another survey that is on its way to refine \\baf measurements is the WiggleZ. The WiggleZ Dark Energy Survey (\\citealt{drinkwater10a}) is expected to complete very soon a narrower redshift survey (area $\\sim 1,000$ deg$^2$), using Blue Emission Line Galaxies, in a comoving volume of $\\sim 1$Gpc$^3$ between $0.2<z<1$. % When analyzing a quarter of the predicted final sample, \\cite{blake10a} show in Figure $20$ a hint of a baryonic acoustic wiggle in the P$(k)$ monopole. Although it is not clear whether this survey will detect a significant peak in the line-of-sight $\\xi$ (assuming $\\theta<3^\\circ$), we are hopeful it will give us indication of the true signal.  If the claim for detection by \\Gpaper is correct, the unexpected strong line-of-sight signal would require an explanation.    It is a pleasure to thank Enrique Gazta\\~{n}aga for commenting in detail on our draft,  in depth discussions  on interpretation, uncertainty estimation issues and sharing information regarding analysis. We thank Lam Hui and David Hogg for discussing interpretation of results, as well as Anna Cabre who shared with us her results. We thank Nic Ross for his very useful comments which helped clarify our draft. We thank Daniel Eisenstein for assistance in selecting LRGs and Idit Zehavi for discussions on weighting algorithms and selection function considerations. We thank Nikhil Padmanabhan, Nic Ross and David Schlegel for BOSS related discussions. We also thank Vincent Desjacques, Abraham Loeb, Ariel S$\\acute{\\mathrm{a}}$nchez, and Martin White for useful discussions and insight. We thank the LasDamas collaboration (http://lss.phy.vanderbilt.edu/lasdamas/) for making their mock catalogs publicly available. Thanks are also due to the Horizon team (http://astro.kias.re.kr/Horizon-Run/) for making their mocks public, and in particular  Changbom Park and Juhan Kim for discussions on usage. E.K thanks Jo Bovy and Mulin Ding for their technical help. E.K thanks David Schlegel and Nic Ross at Lawrence Berkeley National Lab for their hospitality during part of the analysis in March.  E.K was partially supported by a Google Research Award and NASA Award NNX09AC85G.  M.B was supported by Spitzer G05-AR-50443 and NASA Award NNX09AC85G. R.S. was partially supported by NSF AST-0607747 and NASA NNG06GH21G.    Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS Web Site is http://www.sdss.org/.  The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington."
    },
    "1004/1004.0288_arXiv.txt": {
        "abstract": "The impact of Type Ia supernova ejecta on a helium-star companion is investigated via high-resolution, two-dimensional hydrodynamic simulations. For a range of helium-star models and initial binary separations it is found that the mass unbound in the interaction, $\\delta M_{\\rm ub}$, is related to the initial binary separation, $a$, by a power law of the form $\\delta M_{\\rm ub} \\propto a^{m}$. This power-law index is found to vary from $-3.1$ to $-4.0$, depending on the mass of the helium star. The small range of this index brackets values found previously for hydrogen-rich companions, suggesting that the dependence of the unbound mass on orbital separation is not strongly sensitive to the nature of the binary companion. The kick velocity is also related to the initial binary separation by a power law with an index in a range from $-2.7$ to $-3.3$, but the power-law index differs from those found in previous studies for hydrogen-rich companions. The space motion of the companion after the supernova is dominated by its orbital velocity in the pre-supernova binary system. The level of Ni/Fe contamination of the companion resulting from the passage of the supernova ejecta is difficult to estimate, but an upper limit on the mass of bound nickel is found to be $\\sim 5\\times 10^{-4}\\ M_\\odot$. ",
        "introduction": "Type Ia supernovae (SNe~Ia) are among the most catastrophic and energetic events in the Universe. Because they are so luminous, and because their light-curve shapes and absolute magnitudes are correlated, SNe~Ia can be used as ``standardizable candles'' in measuring the distances to remote galaxies, allowing constraints to be placed on key cosmological parameters \\citep{Branch:1992p1603, Riess:1996p1623}. SNe~Ia also play a major role in galactic chemical evolution via their energy and metal input to the interstellar medium. Thus, the nature of their progenitor systems and the physical origin of variations in their properties are fundamental problems of great interest \\citep{Dominguez:2001p2574, Howell:2007p2557}.  Most work on progenitor models for SNe~Ia has focused on two general classes of systems. The single-degenerate scenario involves a CO white dwarf (WD) accreting matter from a non-compact stellar binary companion, eventually becoming unstable to explosive nuclear burning \\citep{Whelan:1973p78,Nomoto:1982p1467}. On the other hand, the double-degenerate scenario involves the merger of two CO WDs whose orbital decay results from the loss of angular momentum due to gravitational wave emission \\citep{Iben:1984p43, Webbink:1984p104}.  Single-degenerate models require that the rate of mass accretion be such that the WD avoid a nova explosion.  For a high accumulation efficiency, stable burning on the WD is required, but this burning phase occurs in a fairly narrow range above $10^{-7}\\ M_\\odot\\ {\\rm yr}^{-1}$.  For rates greater than this range, a wind from the WD develops that limits the accretion efficiency \\citep{Nomoto:1982p1467,Hachisu:1996p3228,Ivanova:2004p123}, making it difficult to explain the observed low-redshift SN~Ia rate of $3 \\times 10^{-5}{\\rm \\ Mpc}^{-3}{\\rm\\ yr}^{-1}$ \\citep{Mannucci:2005p3244} using single-degenerate progenitors alone. Although the predicted and observed numbers of double-degenerate systems are sufficient to explain the observed rate \\citep{Nelemans:2001p3347,Napiwotzki:2001p3317, Napiwotzki:2002p3341}, it is more likely that WD-WD mergers lead to the production of an ONeMg WD followed by accretion-induced collapse to a neutron star \\citep{Nomoto:1985p3367, Ivanova:2004p123,Dessart:2006p3379, Wickramasinghe:2009p3391}.  Moreover, the wide range of total WD masses suggests that double-degenerate models should exhibit much more heterogeneity than is observed in SN~Ia light curves \\citep{Goldhaber:2001p3408,Knop:2003p3425}.  \\cite{Maoz:2008p3181} analyzed the observed SN~Ia rate for assumed initial mass functions (IMF) and concluded that almost all intermediate-mass close binary systems in the range $3 -8\\ M_{\\sun}$ should evolve to the SN~Ia stage. This result allows for a wide range of possibilities for SN~Ia progenitor systems. Furthermore, observations of the SN~Ia rate as a function of redshift suggest the need for a two-component model for the delay time distribution (DTD). (The delay time is the interval between a star's arrival on the zero-age main sequence and its destruction in an SN~Ia.) \\cite{Scannapieco:2005p2469} and \\cite{Mannucci:2006p2431} found that the observations can be fit with a short-delay-time population having delays of $\\sim 10^8$~yr and a long-delay-time population having delays of $3 - 4$~Gyr.  To provide an explanation for the two populations suggested by the observed SN~Ia rate, several pre-supernova progenitor models have been investigated.  The long-delay-time population can be understood in terms of progenitor systems characterized by a main-sequence-like companion in the MS-WD channel \\citep{Hachisu:2008p144} and/or by a red giant in the RG-WD channel \\citep{Hachisu:1996p3228, Hachisu:1999p3434, Hachisu:1999p3439,Hachisu:2008p3845}. In contrast, the short-delay-time population may consist of systems characterized by a massive MS star in the MS-WD channel or by a He star in the He-WD channel \\citep{Waldman:2008p628, Wang:2009p3601, Wang:2009p3585, Wang:2009p3598,Meng:2009p3843}.  Numerically, \\cite{Marietta:2000p112} explored the influence of the supernova explosion on the companion star in the single-degenerate channel with hydrogen-rich stars consisting of red giants, subgiants, and main-sequence stars using two-dimensional Eulerian hydrodynamics simulations. They found that significant quantities of hydrogen would be unbound from the companion star in each case (15\\% of the envelope for main-sequence and subgiant cases, and 98\\% of the red giant envelope), in conflict with observational upper limits on the amount of hydrogen inferred from SN~Ia spectra \\citep{Mattila:2005p1442,Leonard:2007p102}.  A more recent hydrodynamical study by \\cite{Pakmor:2008p139} reexamined the main-sequence simulation of  \\cite{Marietta:2000p112} using a three-dimensional smoothed particle hydrodynamics (SPH) simulation. In contrast to Marietta et al., Pakmor et al.\\ adopted the structure for the companion star based on the binary evolutionary models of \\cite{Ivanova:2004p123}, which yielded more compact main-sequence-like companions.  As a consequence, Pakmor et al.\\ found a tenfold reduction in the amount of unbound mass compared with Marietta et al., bringing the prediction of the amount of unbound hydrogen-rich material into agreement with the observational upper limits.  In a complementary analytical study, \\cite{Kasen:2009p2965} investigated the radiation emitted by the collision of SN~Ia ejecta with a red giant, finding that the light curve should depend on the viewing angle. This result was attributed to the fact that the gas is more transparent in the region shadowed by the companion star. This suggests the possibility that the secondary star may be detectable in future observational studies.  These previous simulations of the effect of a supernova impact on a companion star have been carried out only for models applicable to the long-delay-time population of SN~Ia. In contrast, \\cite{Wang:2009p3601, Wang:2009p3585} suggest a progenitor binary model based upon a helium-star channel for the short-delay-time population.  Using a binary evolution model, for this channel they find a SN~Ia birthrate $\\sim 3 \\times 10^{-4}$~yr$^{-1}$ and a corresponding delay time of $\\sim 4.5 \\times 10^7$~yr to $\\sim 1.4 \\times 10^8$~yr \\citep{Wang:2009p3598}. The latter delay time is consistent with that estimated from observations for the short-delay-time population by \\cite{Scannapieco:2005p2469}, \\cite{Mannucci:2006p2431}, and \\cite{Aubourg:2008p2503}.  In this paper we report the results of Eulerian hydrodynamics simulations of the impact of SN~Ia ejecta on companion stars for the single-degenerate helium-star channel. In the next  section, the assumptions underlying our study, the construction of the initial model, and the numerical method are described.  Our numerical results for a range of helium-star models and orbital separations are reported in \\S~3 and their implications are discussed in \\S~4. In the final section, we summarize our results and make some concluding remarks. ",
        "conclusions": "We have investigated the impact of SN~Ia ejecta on a companion star in the single-degenerate helium-star channel for the short-delay-time population of SN~Ia via two-dimensional hydrodynamical simulations using a range of helium-star models and binary orbital separations. Although the general behavior of the temporal evolution is similar to that in previous studies by \\cite{Marietta:2000p112} and \\cite{Pakmor:2008p139} for main-sequence and red-giant companions, the amount of matter unbound from the helium stars is less than for hydrogen-rich companions.  Due to the shorter orbital periods of helium-star progenitor systems, the space motion of the companion star after the explosion is found to be higher.  We find a power-law relation between the unbound mass and initial binary separation that is consistent with the previous studies, suggesting that the power-law behavior is not strongly sensitive to the nature of the companion. The kick velocity can also fitted by a power-law and we conclude that the power-law index may reflect the relative importance of the effect of ablation. An upper limit on the amount of nickel captured by the helium star is found to be $\\sim 5\\times 10^{-4}\\ M_\\odot$. The ratio of nickel to helium abundance may be useful as a diagnostic of such events in future observational studies of SN~Ia stellar remnants.  Future work in this area will include relaxing the assumption of axisymmetry to model the mixing of ejecta with the helium star by including the binary orbital motion in three spatial dimensions.  Additionally, including radiation transfer within these simulations will allow us to determine how much of the helium is ionized, allowing us to make direct contact with spectroscopic constraints on the presence of helium in these systems."
    },
    "1004/1004.5464_arXiv.txt": {
        "abstract": "% The LMC star BI 108 is photometrically variable with the unique light curve: Two strong periods are present in a strict 3:2 resonance, staying coherent over several observing seasons. A spectroscopic data collected at VLT/UVES reveals unexpected and still not fully understood behavior of the system, that has not been observed in any other early-type multiple systems. ",
        "introduction": "An object LMC SC9-127519 was discovered as a variable star with unique photometric properties during our multiperiodicity search in the OGLE-II database (Szyma\\'nski, 2005). Independently this object was announced as two binary systems in 3:2 resonance by Ofir (2008). With the coordinates, the magnitude (V=13.3) and the colour (B-V=-0.142), we identified object as BI 108 in the list of Brunet et al. (1975), for which they give a spectral type of B1-2. Schmidt-Kaler et al. (1999) list B1:II: for the same object. Probably BI 108 is one of the brightest member of a young open cluster, NGC 1881. The OGLE-II data spans almost four years. A time series analysis returns two periods. Of these two periods the stronger one ($P_1$=3.57793 d), with a peak-to-peak amplitude of about 0.18 mag, is clearly of a double wave nature. A light curve of such that particular appearance can, with little doubt, be attributed to an eclipsing binary nature (Fig.~1). The weaker period, with about 0.05 mag peak-to-peak, has the value of $P_2$=5.36654 d. Within the uncertainties this is exactly 3/2 of the stronger period. The second light curve is not sinusoidal (Fig.~1). An analysis of the photometric variability allows two options, including $P_2$ or $P_2$/2, and 3:2 or 3:4 resonance values, respectively. With the MACHO data (ID 79.5378.25) we were able to confirm both periodicities with the same amplitudes in three bands and their long term coherence. Moreover both periods have exactly the same moment of the minimum. It is an additional constraint difficult to explain in the model. Such a clearly resonant and long-term stable period relation is not known for any other early type variable star.  \\begin{figure}[!t] \\plotone{kolaczkowski_fig1.ps} \\caption{ Upper panel: The OGLE I-band light curves for both periods of the BI 108 system. For each period the contribution of the respective other one has been removed. Lower panel: The original light curve folded with superperiod. In all cases $HJD_0$ = 2451163.8915. Dashed lines indicate phases of the spectral observations. } \\end{figure} ",
        "conclusions": ""
    },
    "1004/1004.2843_arXiv.txt": {
        "abstract": "In this note, we explore a pathway to radio--loudness under the hypothesis that {\\it retrograde} accretion onto giant ($M_{\\rm BH}\\sim 10^9 \\, \\msun$) spinning ($a\\gsim 0.7$) black holes (BHs) leads to the launch of powerful jets, as seen in radio--loud QSOs and recently in LAT/Fermi and BAT/Swift Blazars.  Counter--rotation of the accretion disc relative to the BH spin is here associated to {\\it gas--poor} galaxy mergers progenitors of giant (missing--light) ellipticals.  The occurrence of retrograde accretion enters as unifying element that may account for the radio--loudness/galaxy morphology dichotomy observed in AGN. ",
        "introduction": "Thanks to a wealth of observational and theoretical studies of AGN from radio to gamma-ray frequencies, radio-loud AGN are presently thought to contain relativistic jets which propagate from the nucleus out to kpc and, in some cases, Mpc scales.  The power of jets required to at least fuel their radio lobes is huge, up to $10^{47}\\ergs$ corresponding to an energy release of $\\sim 10^{62}$ erg over a time of $10^8$ yr.  In radio-quiet AGN, the radio emission is instead weak and large scale radio jets are absent.  Although the boundary between these two classes is blurred, strong jets (in terms of kinetic relative to accretion power) exist only in a minority of AGN.   Optical studies of {\\it nearby} galaxy nuclei (Capetti \\& Balmaverde 2006) have shown that {\\it all} radio-loud AGN invariably reside in giant {\\it core} (or {\\it missing--light}; Kormendy et al. 2009) ellipticals, i.e. in early-type galaxies whose nuclear surface brightness profile shows a {\\it deficit} in star-light with respect to the outer profile.  By contrast, normal, less massive {\\it coreless} (or {\\it extra--light}; Kormendy et al. 2009) ellipticals {\\it all} host radio-quiet AGN (Capetti \\& Balmaverde 2006), suggesting that a {\\it morphology-related dichotomy} exists among {\\it ellipticals}, likely determined by galaxy evolution. The dichotomy seems also to emerge from a sample of radio-loud and radio-quiet QSOs at redshift $z<0.2$, the former residing in giant ellipticals that show signs of interaction, and the latter residing in merging gas-rich galaxies of intermediate mass (Wolf \\& Sheinis 2008). Evidence exists that the most powerful radio and gamma--loud AGN are associated to the heaviest supermassive BHs in the universe, with $M_{\\rm BH}\\gsim 10^9\\,\\msun$ (Ghisellini et al. 2009a,b) indicating that the BH {\\it mass} can be a key parameter for radio-loudness.   From a theoretical point of view, the launching of jets from accreting BHs have been extensively studied in the past three decades, following the seminal paper by Blandford and Znajek (BZ hereafter; 1978).  The BZ process exploits energy extraction from a rapidly rotating black hole via a purely electromagnetic interaction of a large scale magnetic field which threads the rotating event horizon. In this framework, jets are produced at the expense of the rotational energy of the black hole $E_{\\rm rot}=M_{\\rm BH}[1-2^{-1/2}(1+\\sqrt{1-a^2})^{1/2}]c^2$, where $a$ is the dimensionless spin parameter related to the BH spin ${\\bf J_{\\rm BH}}={\\bf {\\hat l}}\\,(GM^2_{\\rm BH}a/c)$ pointing along ${\\bf {\\hat l}}$ (for a maximally rotating [$a=1$] Kerr BH of $10^9\\,\\msun$, $E_{\\rm rot}=6\\times 10^{62}$ erg).  This may suffice to power a jet and to explain why AGN with comparable optical luminosities can be either radio-loud (large $a$) or radio-quiet (small $a$, as $E_{\\rm rot}$ decays $\\propto\\,M_{\\rm BH}a^2/8$ for $a\\to 0$) depending on $a$ (Wilson \\& Colbert 1995; Sikora et al. 2007).  Advanced fully general-relativistic magneto-hydrodynamical simulations of spinning BHs (McKinney 2005, 2006) predict high collimation of the inner Poynting flow, bulk Lorentz factors of order 10, and jet powers $P_{\\rm jet}\\propto a^2$, or $\\propto a^6$,  implying in this last case high jet luminosities only for close to maximally spinning BHs (Tchekhovskoy et al. 2009). Another essential ingredient determining the BZ power is the value of the magnetic field threading the horizon, which is tied to the accretion rate and disc structure (Ghisellini et al. 2009a).  In two recent works, Garofalo (2009a,b) stressed the importance of the plunging region between the innermost stable circular orbit (ISCO\\footnote{ISCO, expressed in units $GM_{\\rm BH}/c^2$, is at 6 for $a=0$, and at 9 (1) for retrograde accretion with $a=-1$ (prograde accretion with $a=1$).}) of the accretion disc and the BH horizon. In this region the magnetic field can be substantially amplified from its value at ISCO, particularly if the BH has a {\\it retrograde} spin vector with respect to the accretion disc angular momentum (we will use here the convention of negative spin, $a<0$, for retrograde accretion). As a consequence, high jet powers can be extracted even/also from non-maximally rotating BHs (Garofalo 2009b).  Here, we explore the consequences of Garofalo's conjecture in the aim at connecting the BH spin and accretion mode to the structure of the underlying host galaxy. Moderately high values of the spin parameter $a$ ($\\gsim 0.7$) and retrograde accretion appear to enter as unifying ingredients to account for the radio-loud/host morphology dichotomy observed, at low redshifts, among ellipticals.\\footnote{See Sikora et al. (2007) for the discussion on the most general late/early-type versus radio-quiet/loud dichotomy observed in AGN, and Fanidakis et al. (2009) for a different spin dependent cosmological model of jet formation in AGN.} These considerations go in the direction of the very recent proposal that most radio-loud AGN harbor spinning BHs accreting in the retrograde mode (Garofalo et al. 2010). ",
        "conclusions": "In this note, we propose that {\\it retrograde} accretion onto a massive ($M_{\\rm BH}\\gsim 10^9\\,\\msun$) {\\it post-coalescence} spinning BH is conducive to the generation of a powerful large-scale jet, in the gas-poor environment of a giant elliptical with core. This is in agreement with the observation that radio--loud AGN are hosted in bright ellipticals that show a stellar core (whenever optical data are available),  and with the recent finding on the dichotomy in radio jet orientations in elliptical galaxies (Browne \\& Battye 2010).  In less-louder radio-loud AGN, Browne and Battye find a tendency for the axis of the radio emission to align with the minor axis of the starlight of the host, an oblate rotationally supported elliptical.  By contrast they find no preferred radio-optical alignment among the radio-louder objects possibly hosted in triaxial non-rotating ellipticals.  This study appears to support our evolutionary scheme: coreless (extra--light) rotationally supported ellipticals (Kormendy et al. 2009) that likely form in gas-rich mergers are expected to host a BH accreting in the prograde mode with a high spin aligned with the angular momentum of the large-scale disc, and so aligned with the minor axis of the starlight. By contrast, core (missing--light) triaxial ellipticals are expected to host a BH that has a random spin orientation regardless its accretion mode (i.e. whether prograde or retrograde).    The model rests on Garofalo's conjecture that around a retrograde BH, magnetic fields threading the horizon reach maximum amplification, resulting in a dramatic enhancement of the jet power even for non--extreme values of $a$.  The interval of spin values for operation of this process is still weakly constrained by theory. Furthermore {\\it prograde} accretion can also result in large BZ luminosities (Garofalo 2009a,b). Thus, there should exist a {\\it zone of avoidance} delimited by a maximum--negative $a^{-}_{\\rm jet}$ and a minimum--positive $a^+_{\\rm jet}$, outside which the generation of jets is possible. The values of $a^-_{\\rm jet}$ and $a^+_{\\rm jet}$ may not be symmetric as non--symmetric is the underlying Kerr spacetime.  If retrograde accretion is a necessary condition {\\it only} for the most powerful jets hosted in ellipticals, the value of $a^-_{\\rm jet}$ that our model requires is $a^-_{\\rm jet}\\gsim -0.7$, resulting from isotropic BH coalescences in dry mergers.  If prograde accretion vehicles the production of jets as well, a value of $a^+_{\\rm jet}$ close to $+0.4$ (from chaotic accretion) would imply the presence of jets in any type of galaxy, from discs in isolation to spheroids in clusters and groups. The paucity of radio--loud disc hosts and the lack, at low redshift, of radio--loud coreless host--ellipticals suggest that such a low value of $a^+_{\\rm jet}$ is irrealistic.  One would be tempted to require $a^+_{\\rm jet}\\sim +1$, according to the models by McKinney (2005).    The rotational energy stored in giant BH ($M_{\\rm BH}\\gsim 10^9$) may suffice to power a radio--loud AGN at an average level of $\\sim 10^{46}\\,\\ergs$ for a lifetime $\\tau_{\\rm jet}\\sim 10^8$ yr. However retrograde accretion causes the BH to spin down on a time $\\tau_{\\rm down}\\sim$ min($\\tau_{\\rm ac},\\tau_{\\rm jet}$). Spin--down by accretion occurs on $\\tau_{\\rm ac}\\sim J_{\\rm BH}/({\\dot M_{\\rm net}{\\tilde l}_{\\rm ISCO}}) \\sim M_{\\rm BH}/{\\dot M}_{\\rm net}$ (where ${\\dot M}_{\\rm net}$ is the {\\it net} inflow rate onto the BH, and ${\\tilde l}_{\\rm ISCO}$ the angular momentum per unit mass of a test particle at ISCO).  $\\tau_{\\rm ac}$ is close to the Salpeter time for $e$--folding of the BH mass.  The BZ timescale $\\tau_{\\rm jet}$ depends on the magnitude of the magnetic field that threads the horizon and on the thermal/geometrical structure of the accretion disc (Moderski et al. 1998). In Garofalo's scenario, retrograde accretion is a transitory phase during the lifetime of a BH (Garofalo et al. 2010) whose duration and recurrence depend on the way the BH is fed whether through continuous or chaotic  accretion.   Recent observations of extremely bright Blazars with BAT/{\\it Swift} at redshifts $z>2$ (Ajello et al. 2009; Ghisellini et al. 2009a) and of Blazars and FRGs with LAT/{\\it Fermi} (Ghisellini et al. 2009b) indicate that luminous radio-loud AGN host {\\it giant} BHs,  and that a prominent accretion disc co-exists with a powerful jet suggesting that extraction of rotational energy occurs through accretion (Maraschi 2001). We are tempted to associate the extreme BH mass, inferred from the disc luminosity, to the stability of the BH spin orientation and the jet power to its spin modulus and to accretion in the retrograde mode.  Further studies on the {\\it stability} of retrograde accretion will help in disentagling the nature on the AGN radio--loud/quiet dichotomy.  \\vskip 1.2 truecm"
    },
    "1004/1004.4148_arXiv.txt": {
        "abstract": "{In the distant past, astronomy was often intertwined with religion into a unified cosmos. As science became a distinct cultural enterprise, astronomy has witnessed a variety of rich interactions with other fields.  Mathematical statistics was stimulated in the 19th century by astronomical problems, and today astrostatistics is a small but growing cross-disciplinary field advancing methodology to address challenges in astronomical data analysis.  Throughout the 20th century, astronomy became closely allied with physics such that astronomy and astrophysics are now profoundly intertwined.  Physical chemistry played a major role in the identification of molecules in the Milky Way Galaxy, and astrochemistry is now an active subfield giving insights into cosmic molecular processes.  The importance of cross-disciplinary interactions with engineering (for instrumentation), Earth sciences (for planetary studies), computer science (for astroinformatics) and life sciences (for astrobiology) is also growing.  Cross-disciplinary research has been essential both for crucial discoveries in astronomy and for improving the quality of astronomical research.  It should be fostered with increased flexibility in the training of young astronomers and with sufficient funding to nurture these fields. }  \\FullConference{Accelerating the Rate of Astronomical Discovery, sps5\\\\ August 11-14, 2009\\\\ Rio de Janeiro, Brazil}  \\begin{document} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.1247_arXiv.txt": {
        "abstract": "In this paper we discuss the age and spatial distribution of young (age$<$1~Gyr) SMC and LMC clusters using data from the Magellanic Cloud Photometric Surveys. Luminosities are calculated for all age-dated clusters. Ages of 324 and 1193 populous star clusters in the Small and the Large Magellanic Cloud have been determined fitting Padova and Geneva isochrone models to their resolved color-magnitude diagrams. The clusters cover an age range between 10~Myr and 1~Gyr in each galaxy. For the SMC a constant distance modulus of $(m-M)_0$ = 18.90 and a metallicity of Z = 0.004 were adopted. For the LMC, we used a constant distance modulus of $(m-M)_0$ = 18.50 and a metallicity of Z = 0.008. For both galaxies, we used a variable color excess to derive the cluster ages. We find two periods of enhanced cluster formation in both galaxies at 160~Myr and 630~Myr (SMC) and at 125~Myr and 800~Myr (LMC). We present the spatially resolved recent star formation history of both Clouds based on young star clusters. The first peak may have been triggered by a close encounter between the SMC and the LMC. In both galaxies the youngest clusters reside in the supergiant shells, giant shells, the inter-shell regions, and toward regions with a high H$\\alpha$ content, suggesting that their formation is related to expansion and shell-shell interaction. Most of the clusters are older than the dynamical age of the supergiant shells. No evidence for cluster dissolution was found. Computed V band luminosities show a trend for fainter magnitudes with increasing age as well as a trend for brighter magnitudes with increasing apparent cluster radii. ",
        "introduction": "\\label{youngies_intro}  Due to their proximity the Magellanic Clouds (MCs) offer an excellent opportunity to study their spatially resolved star formation (SF) histories. SF can be triggered by several mechanisms such as, e.g., the self-induced gravitational collapse of a molecular cloud, tidal shocking, a turbulent interstellar medium, or cloud-cloud collisions \\citep[e.g., ][]{McKee07}. The MCs and the Milky Way (MW) are interacting with each other. The formation of star clusters younger than $\\lesssim$1~Gyr in the MCs was probably triggered by interactions of the galaxies with each other and with the MW \\citep[e.g., ][]{yoshi03}. Star clusters may be produced through strong shock compressions induced by close encounters of their host galaxies, which causes enhanced star formation. Conversely, the star formation rate decreases again once the galaxies recede from each other. Repeated encounters then lead to episodic cluster formation. In the MCs, a correlation between young star clusters and putative close encounters with each other and MW has been suggested by, e.g., \\citet[][G95]{gir95},~\\citet[][PU00]{piet00}, and~\\citet[][C06]{Chiosi06}.  Strong tidal perturbations induced by the encounters could also have triggered the formation of clusters \\citep*[e.g., ][]{Whitmore99} in the MCs. Possible orbits of the Small Magellanic Cloud (SMC), the Large Magellanic Cloud (LMC), and the MW have been modeled by several authors \\citep[e.g., ][]{Bekki05}. They found that it was difficult to keep the Clouds bound to each other for more than 1~Gyr in the past. The LMC and the SMC have been part of a triple system together with the Milky Way since at least 1~Gyr \\citep[e.g., ][Kallivayalil et al. 2006a/b]{Bekki05}. It is possible that the Clouds are not a bound system and that they are making their first passage close to the MW.  Interestingly, the cluster formation histories of the LMC and SMC show large differences. In the LMC, two main epochs of cluster formation \\citep[e.g., ][]{Bertelli92} have been observed that are separated by an ''age gap'' of about 4-9~Gyr \\citep[e.g., ][]{holtz99,john99,harzar01}, in which no star clusters have formed. The two epochs of pronounced cluster formation occurred $>$9~Gyr ago and $\\sim$3-4~Gyr ago. In the LMC, a few globular clusters are found that are as old as the oldest Galactic globulars \\citep{olsen98}. During the past $\\sim$4~Gyr, star clusters have been forming continuously until the present day. In contrast, the star clusters in the SMC cover a wide range of ages and continued to form over at least the last $\\sim$10.5~Gyr \\citep[e.g., ][]{glatt08a,glatt08b,Parisi08} to the present day. Interestingly, in the SMC the cluster formation history appears to have started with a delay since the SMC formed its first and only globular cluster, NGC\\,121, 2-3~Gyr later than the LMC or the MW \\citep[][and references therein]{glatt08a}. The LMC contains about $\\sim$4200 star clusters, while in the SMC $\\sim$770 star clusters have formed (and survived). The cluster census is probably still incomplete, missing small and faint clusters that are yet to be detected. Ongoing and prospective space-based observations may further increase the number of known star clusters.  The most recent catalog cross-correlating all known objects of the LMC, SMC, and the Magellanic Bridge region was published by \\citet{bica08b} (B08). However, the cluster sample still is highly incomplete as pointed out by the authors. Only for a few clusters in B08's catalog, ages have been determined. For young SMC clusters, \\citet{pietrzynski99} (PU99) used isochrone fitting on data from the Optical Gravitational Lensing Experiment \\citep[OGLE II; ][]{udal98a} to determine ages for 93 clusters. \\citet{Dieball02} compiled ages for 306 binary cluster candidates in the LMC from a variety of literature sources ranging from multiwavelength integrated light studies to isochrone fitting to resolve clusters. \\citet{rafelski05} (RZ05) made use of integrated colors and derived ages for 200 clusters. C06 determined ages of 164 associations and 311 star clusters based on data from the OGLE using isochrone fitting. Their sample is the largest available catalog with SMC cluster ages. Ages for young LMC clusters have been provided by G95 based on integrated colors (624 objects) and by PU00 using isochrone fitting applied to OGLE-II data ($\\sim$600 clusters).  Luminosities have been published for 204 SMC star clusters by RZ05 measuring integrated colors from the Magellanic Clouds Photometric Survey (MCPS). \\citet{bica96} (B96) published integrated photometry of 624 LMC star clusters that was based on observations carried out at the 0.61-m telescope at CTIO in Chile and at the 2.15-m CASLEO telescope in Argentina.  In the present study we increase the number of age-dated young LMC and SMC star clusters and calculate V-band luminosities. We aim at improving the understanding of the cluster age distribution of these two irregular galaxies and present spatial distribution maps of the star clusters in both galaxies. To achieve this goal, we make use of ground-based data of the Magellanic Clouds Photometric Surveys (MCPSs) \\citep{zar02,zar04}. In the next Section the observations and data reduction are described. In $\\S$~\\ref{sec:metunddistmod} the distances, reddenings, and metallicities of both the SMC and the LMC are given. In $\\S$~\\ref{sec:agedist} the clusters' age distribution, spatial distribution, and dissolution effects are discussed and in $\\S$~\\ref{sec:youngies_lum} the correlation between the cluster luminosities and age/radius is derived. ",
        "conclusions": "\\label{sec:youngies_summary}  We have presented ages and luminosities of 324 and 1193 populous SMC and LMC star clusters, respectively. An age range of $\\sim$9~Myr to 1~Gyr was covered based on isochrone fitting to resolved color-magnitude diagrams in both galaxies. Using only cluster ages derived in this study, we find two maxima of enhanced cluster formation for both galaxies which appear to be correlated. In the SMC, the peaks are found at $\\sim$160~Myr and $\\sim$630~Myr, and in the LMC at $\\sim$125~Myr and $\\sim$800~Myr. Model calculations predict that the last close encounter between LMC and SMC occurred around 100-200~Myr ago. During a close encounter, the star formation is expected to be enhanced. Therefore, the first peaks in the cluster age distributions could have been triggered by this tidal interaction. Extending our samples with cluster ages derived by C06 we find a third pronounced period of enhanced cluster formation in the SMC at around 8~Myr. We find the same in the LMC combining our sample with the one of PU00. These peaks are only visible if we extend our sample with objects classified as associations, objects which did not or could not reach higher ages because they dissolve too quickly.  The youngest objects in both galaxies are associated with super giant shells, giant shells, the inter-shell region, and with HII regions. Their formation is probably related to shell expansion and shell interaction. In the spatial distribution of the clusters younger than $\\sim$16~Myr the two SMC shells are clearly visible. The older objects are widely spread across the entire SMC main body, but show a concentration in the western part of the galaxy. In the LMC, the youngest objects are concentrated in 30 Doradus, SGS 11 (LMC~4), and in the giant shells located in the western part and in the bar region. The older LMC clusters are mostly distributed along the bar and along the rim. One can see nicely how star cluster formation propagated along the LMC bar. We find no indication for propagating star cluster formation in the SGSs in either LMC or SMC. Most of the LMC star clusters are older than the dynamical ages of the SGSs and therefore may have formed in shells, which already have dissolved and cannot be detected at the present day.  No obvious dissolution effects were found for MCs star clusters younger than $\\sim$1~Gyr. It is quite difficult to ascertain a real absence of cluster dissolution using this study. Two biases may play a major role: 1. Infant mortality cannot be accounted for, because very young star clusters and OB-associations are not included in our sample; and 2. Cluster dissolution processes for clusters older than $\\sim$1~Gyr, because we did not age-date clusters in this age range. Within the time period considered here - 10~Myr to 1~Gyr - we do not find evidence for cluster dissolution.  In both galaxies, the clusters become fainter with increasing age. The very massive hot stars, which are still present in the young star clusters and contribute most of the light, become fainter and redder with increasing age and so do the star clusters. This trend can be seen in both the LMC and the SMC. The total cluster luminosity increases with increasing radius due to a larger number of stars within the cluster radius."
    },
    "1004/1004.1912_arXiv.txt": {
        "abstract": "We report the discovery of 47 new T~dwarfs in the Fourth Data Release (DR4) from the Large Area Survey (LAS) of the UKIRT Infrared Deep Sky Survey with spectral types ranging from T0 to T8.5. These bring the total sample of LAS T dwarfs to 80 as of DR4. In assigning spectral types to our objects we have identified 8 new spectrally peculiar objects, and divide 7 of them into two classes. H$_2$O-H-early have a H$_2$O-$H$ index that differs with the H$_2$O-$J$ index by at least 2 sub-types. CH$_4$-J-early have a CH$_4$-$J$ index that disagrees with the H$_2$0-$J$ index by at least 2 subtypes. We have ruled out binarity as a sole explanation for both types of peculiarity, and suggest that they may represent hitherto unrecognised tracers of composition and/or gravity. Clear trends in $z'(AB)-J$ and $Y-J$ are apparent for our sample, consistent with weakening absorption in the red wing of the  K{\\sc I} line at 0.77$\\mu$m with decreasing effective temperature. We have used our sample to estimate space densities for T6--T9 dwarfs. By comparing our sample to Monte-Carlo simulations of field T~dwarfs for various mass functions of the form $\\psi(M)~\\propto~M^{-\\alpha}$~pc$^{-3}$ \\Msun$^{-1}$, we have placed weak constraints on the form of the field mass function. Our analysis suggests that the substellar mass function is declining at lower masses, with negative values of $\\alpha$ preferred. This is at odds with results for young clusters that have been generally found to have $\\alpha~>~0$. ",
        "introduction": "\\label{sec:intro}  The study of substellar objects presents a number of important opportunities for extending our understanding of star and planet formation, both through detailed study of individual systems and through statistical population studies. The statistical characteristics of the substellar population, such as binary fraction and distribution, and the form of the initial mass function \\citep[IMF; ][]{salpeter55} across the entire substellar regime provide crucial observational constraints for models of star and planet formation, which currently offer a number of alternative formation scenarios that depend on differing balances of the dominant physics across the low-mass stellar-substellar mass spectrum \\citep[e.g. ][ and references therein]{bate05,padoan02}. The crucial first step for any observational effort to address these issues is the initial identification of a statistically useful sample of brown dwarfs, the first of which were not discovered until the mid-1990s.   The majority of brown dwarfs have been identified via one of two routes: the mining of wide field surveys such as the Sloan Digital Sky Survey \\citep[SDSS; ][]{sdss} and the Two Micron All Sky Survey  \\citep[2MASS; ][]{2mass}  to find nearby L and T dwarf members of the field population and deep optical and near-infrared surveys of the young clusters and OB associations \\citep[e.g.][]{lucas2000,zap2000,caballero07,bihain09}. Most recently, the UKIRT Infrared Deep Sky Survey (UKIDSS) Galactic Clusters Survey has significantly improved the substellar sample across a number of young regions \\citep[e.g.][]{lod06,lod07a,lodieu09}.     To date, the results from cluster and associations have dominated the study of the low-mass extreme of the IMF due largely to the assumption of coevality in clusters that allows the mass-age degeneracy for substellar objects to be broken. A number of determinations have been published that are broadly in agreement across the $\\sim 0.1 - 0.03$~\\Msun\\ range \\citep[e.g. ][]{moraux03,barrado02,lodieu09,moraux07}.   It is the age-mass degeneracy that has hampered efforts to measure the form of the substellar IMF from analysis of the local field population of L and T dwarfs. Although the determination of masses for individual field brown dwarfs is currently prevented by uncertainty about their age, the field mass function can still be constrained through comparison of the observed luminosity function or spectral type distribution with those predicted by Monte Carlo simulations for various star formation histories and underlying mass functions \\citep[e.g. ][]{chabrier2002,burgasser04,dh06}. \\citet{allen05} have applied a different statistical approach to solving this problem, by using Bayesian inference to evaluate the probabilities of different underlying mass functions for space densities of M,~L and~T dwarfs from 2MASS and SDSS, and estimated the age distribution of the field population.   Searches of the SDSS and 2MASS data sets have resulted in the discovery of over 500 L dwarfs and more than 100 T dwarfs in the field (see www.DwarfArchives.org for an up-to-date list of published objects). However, this sample is dominated by objects earlier than type T6, with just 26 objects identified in 2MASS and SDSS with types T6 or later and just a handful of type T8. And it is this population of objects in the $\\geq$T6 range that is most sensitive to the underlying mass function, whilst the spectral type distribution of earlier objects depends more strongly on the Galactic formation history \\citep[see ][ Figure 5]{burgasser04}.   The Large Area Survey (LAS) of the  UKIDSS has now successfully extended the T dwarf sample to types later than those first revealed by the 2MASS and SDSS surveys \\citep{warren07,ben08,ben09} and is now identifying a substantial sample of late-type T dwarfs that will be ideal for constraining the substellar mass function in the field. In this work we extend our searches of earlier data releases of the LAS \\citep[][ data releases 1 and 2 respectively]{lod07,pinfield08} to include all candidates drawn from Data Release 4 (DR4, which incorporates DRs 1-3) which took place on 1$^{st}$ July 2008. Follow-up to confirm spectral types for this sample is now essentially complete for $J \\leq 19.0$ and we report here the discovery of 47 new T~dwarfs, including one T8+ dwarf and a number of spectrally peculiar objects, and use this sample to place improved constraints on the form of the field substellar mass spectrum. ",
        "conclusions": "\\label{sec:summ}  We have reported the discovery of 47 new T~dwarfs in the UKIDSS LAS DR4 with spectral types ranging from T0 to T8.5. These bring the total sample of LAS T dwarfs to 80. In assigning spectral types to our objects we have identified 8 new spectrally peculiar objects, and divide 7 of them into two classes:  \\begin{description} \\item[\\bf{H$_2$O-H-early}] H$_2$O-$H$ index implies an earlier type than that suggested by the H$_2$O-$J$ index by at least 2 subtypes; \\item[\\bf{CH$_4$-J-early:}] CH$_4$-$J$ index implies an earlier type than that suggested by the H$_2$0-$J$ index by at least 2 subtypes; \\end{description}    We have ruled out L-T binarity as a sole explanation for both types of peculiarity, and suggest that they may represent hitherto unrecognised tracers of composition and/or gravity. These objects are ideal candidates for further kinematic and mid-infrared studies.  Clear trends in $z'(AB)-J$ and $Y-J$ are apparent for our sample, consistent with weakening absorption in the red pressure broadened wing of the 0.77$\\mu$m K{\\sc I} line as temperature decreases through the T-sequence.  We have estimated space densities for T6--T9 dwarfs, and by comparing our sample to Monte Carlo simulations have placed weak constraints on the form of the field mass function. Our analysis suggests that negative values of $\\alpha$ (where $\\psi(M)~\\propto~M^{-\\alpha}$~pc$^{-3}$ \\Msun$^{-1}$) are to be preferred. This is at odds with results for young cluster that have generally found $\\alpha~>~0$. We refrain from making a firm estimate for the value of $\\alpha$ in the absence of a more complete examination of the UKIDSS LAS source detection efficiency, and robust $T_{\\rm eff}$ estimates for our sample from mid-infrared photometry. However, it seems unlikely that these factors can fully account for our small number of $\\geq$T6 dwarfs, and a declining underlying mass function across the late T~range seems probable.         \\appendix"
    },
    "1004/1004.1409_arXiv.txt": {
        "abstract": "\\PRE{\\vspace*{.3in}} We outline the expected constraints on non-Gaussianity from the cosmic microwave background (CMB) with current and future experiments, focusing on both the third ($f_{\\rm NL}$) and fourth-order  ($g_{\\rm NL}$ and $\\tau_{\\rm NL}$) amplitudes of the local configuration or non-Gaussianity.  The experimental focus is the skewness (two-to-one) and kurtosis (two-to-two and three-to-one) power spectra from weighted maps. In adition to a measurement of $\\tau_{\\rm NL}$ and $g_{\\rm NL}$ with WMAP 5-year data, our study provides the first forecasts for future constraints on $g_{\\rm NL}$. We describe how these statistics can be corrected for the mask and cut-sky through a window function, bypassing the need to compute linear terms that were  introduced for the previous-generation non-Gaussianity statistics, such as the skewness estimator. We discus the ratio $A_{\\rm NL} =  \\tau_{\\rm NL}/(6f_{\\rm NL}/5)^2$ as an additional test of single-field inflationary models and discuss the physical significance of each statistic.  Using these estimators with WMAP 5-Year V+W-band data out to $l_{\\rm max}=600$ we constrain  the cubic order non-Gaussianity parameters $\\tau_{\\rm NL}$, and $g_{\\rm NL}$ and find  $-7.4 < g_{\\rm NL}/10^5 < 8.2$ and   $-0.6 < \\tau_{\\rm NL}/10^4 < 3.3$ improving the previous COBE-based limit on $\\tau_{\\rm NL} < 10^8$ nearly four orders of magnitude with WMAP. ",
        "introduction": "We have now entered an exciting time in cosmological studies where we are now beginning to constrain simple slow-roll inflationary models with high precision observations of the cosmic microwave background (CMB) and large-scale structure.  In addition to constraining inflationary model parameter space with traditional parameters such as the spectral index $n_s$ and the tensor-to-scalar ratio $r$, we may soon be able to use parameters associated with primordial non-Gaussianity to improve model selection.  In the simplest realistic inflationary models, the field(s) responsible for inflation have minimal interactions.  Such an interaction-less situation should have led to Gaussian primordial curvature perturbations, assuming that pertubations in the inflaton field generates the curvature perturbation. In this case, the two point correlation function contains all the informations on these perturbations.  If the early inflation field(s) have non-trivial interactions, higher-order correlation functions  of the curvature perturbations will contain {\\it connected } pieces encoding information about the primordial inflationary interactions.  This is analogous to the situation encountered in particle physics where correlation functions can be separated into unconnected and connected Feynman diagrams, the later containing information about the underlying interactions (see Fig.~\\ref{fig:Feynman} for an example involving the four-point function).  A detection of non-Gaussianity therefore gives an important window into the nature of the inflation field(s) and their interactions.  To parameterize the non-Gaussianity of a nearly Gaussian field, such as the primordial curvature perturbations $\\zeta({\\bf x})$,  we can expand them perturbatively~\\cite{Kogo:2006kh} to second order as: \\begin{equation} \\zeta({\\bf x}) = \\zeta_g({\\bf x}) + {3 \\over 5}f_{\\rm NL} \\left[\\zeta_g^2({\\bf x}) - \\langle \\zeta^2_g({\\bf x})\\rangle\\right] + {9 \\over 25 } g_{\\rm NL} \\zeta_g^3({\\bf x}), \\label{eq:phi} \\end{equation} where $\\zeta_g({\\bf x})$ is the purely Gaussian part with $f_{\\rm NL}$ and $g_{\\rm NL}$ parametrizing  the first and second order deviations from Gaussianity.   This parameterization of the curvature perturbations is known as the local model as this definition is local in space.  Much effort has already gone into measuring non-Gaussianity at first-order in curvature perturbations using the bispectrum of the CMB anisotropies or large-scale structure galaxy distribution parametrerized by $f_{\\rm NL}$ (see Eq.~\\ref{eq:phi}).  These studies have found $f_{\\rm NL}$ to be consistent with zero~\\cite{Yadav:2007yy,Smith:2009jr,Komatsu:2010fb,Smidt:2009ir}.  However, there is hope that a significant detection may be possible by future surveys that will lead to improved errors~\\cite{Komatsu:2009kd}. \\begin{figure} \\begin{center} \\includegraphics[scale=0.25]{fig1} \\end{center} \\vspace{-0.5cm} \\caption[width=1in]{Four point correlation function for the $\\phi^3$ theory.  The correlation functions breaks up into interaction-less unconnected diagrams and connected diagrams containing information about the interactions. } \\label{fig:Feynman} \\end{figure}  In the trispectrum, two parameters of second-order non-Gaussianity at fourth-order in curvature perturbations, $\\tau_{\\rm NL}$ and $g_{\\rm NL}$, can be measured.  In this paper we also introduce a third parameter, $A_{\\rm NL}$ is an additional parameter that compares $\\tau_{\\rm NL}$ of the trispectrum to $(6f_{\\rm NL}/5)^2$ from the bispectrum as a ratio: \\begin{equation} A_{\\rm NL} = {\\tau_{\\rm NL} \\over (6 f_{\\rm NL}/ 5)^2}. \\end{equation} This ratio can be quite different for many inflationary models~\\cite{Byrnes:2010em,Chen:2009bc} and, as will be shown below, $A_{\\rm NL} \\neq 1$ rules out single-field inflationary models altogether, including the standard curvaton scenario (which neglects perturbations from the inflaton field).  In this paper we discuss the skewness and kurtosis power spectra method for probing primordial non-Gaussianity and give constraints for the first ($f_{\\rm NL}$) and second-order ($g_{\\rm NL}$ and $\\tau_{\\rm NL}$) amplitudes of the local model in addition to their ratio $A_{\\rm NL}$.  Using the bispectrum of CMB anisotropies as seen by WMAP 5-year data, Smidt et al.~(2009) found $-36.4 < f_{\\rm NL} < 58.4$ at 95\\% confidence~\\cite{Smidt:2009ir}. This is to be compared with the most recent WMAP 7 measurement of $-10 < f_{\\rm NL} < 74$~\\cite{Komatsu:2010fb}, where part of the discrepancy is due to a difference in optimization~\\cite{Smith}.   As outlined in Section~VI, using the trispectrum of the same data we find that  $-0.6 < \\tau_{\\rm NL}/10^4 < 3.3$ and $-7.4 < g_{\\rm NL}/10^5 < 8.2$ at 95\\% confidence level showing second order non-Gaussianity is consistent with zero in WMAP.  This paper serves as a guide to the analysis process behind our derived limits on $\\tau_{\\rm NL}$, $g_{\\rm NL}$ and $A_{\\rm NL}$.  Furthermore, in this paper we analyze what to realistically expect when measuring non-Gaussianity from CMB temperature data.  We believe establishing what constraints can be placed upon $f_{\\rm NL}$, $\\tau_{\\rm NL}$, $g_{\\rm NL}$ and $A_{\\rm NL}$ by future experiments is important in determining what models may and may not be tested by future data.  We also highlight several advantages of our work, including ways to correct the cut-sky and mask through a window function without using linear terms which are computationally prohibative~\\cite{Creminelli:2005hu,Yadav:2007ny}.   This paper is organized as follows:  In Section~\\ref{sec:Inf} we review how non-Gaussianity may be used to distinguish between common inflationary models and stress the physical significance of each statistic. In Section~\\ref{sec:Theory} we describe the skewness and kurtosis power spectra and explain how they may be used to extract information about primordial non-Gaussianity from the CMB.  In Section~\\ref{sec:fishertheory}, we describe the signal-to-noise of each estimator, how to add the experimental beam and noise to these calculations and discuss why these power spectra have the advantage for dealing with a cut sky.   In Section~\\ref{sec:fisherresults} we calculate the fisher bounds for upcoming experiments for each statistic. In Section~\\ref{sec:prioranalysis} we discuss the technical details for measuring non-Gaussianity in the trispectrum  and in Section~\\ref{sec:conclusion} we conclude with a discussion. ",
        "conclusions": "\\label{sec:conclusion}  In this paper we discussed the skewness and kurtosis power spectrum approach to probing primordial non-Gaussianity.  We outlined the expected constraints these techniques will place using future experimental data.  These constraints were calculated by computing the signal-to-noise for each estimator, properly taking into account the noise and beam of each experiment.  Optimal error bars for $f_{\\rm NL}$,  $g_{\\rm NL}$ and $\\tau_{\\rm NL}$ are listed as a function of $l_{\\rm max}$.  It was argued that the skewness and kurtosis power spectrum approach to measure non-Gaussianity has several advantages. These advantages include the ability to separate foregrounds and other secondary non-Gaussian signals, the ability to measure the scale dependance of each statistic and an advantage that the cut sky can be corrected from a matrix $M_{l l'}$ without needing to compute extra linear terms.  The physical significance of each non-Gaussian statistic is discussed.  In the bispectrum, different non-Gaussian triangle configurations in Fourier space contributing to $f_{\\rm NL}$ are related to different underlying physics.  By adding a local measurement of the trispectrum, a new statistic $A_{\\rm NL} = \\tau_{\\rm NL}/(6 f_{\\rm NL}/5)^2$ will be a powerful probe to  distinguish between multi-field models.  Single-field models can be ruled out in general if $A_{\\rm NL} \\neq 1$ and we discussed how this may be a real possibility with Planck or EPIC.   Furthermore, for $A_{\\rm NL}$ large enough, the trispectrum becomes a better probe for non-Gaussinity than the bispectrum for analysis utilizing information on very small scales. The parameter $g_{\\rm NL}$ will be the hardest to constrain.  A constraint on this parameter will uncover information on self-interactions."
    },
    "1004/1004.1779_arXiv.txt": {
        "abstract": "We report the detection of a large mass planet orbiting around the K0 metal-rich subgiant HD38801 ($V=8.26$) by precise radial velocity (RV) measurements from the Subaru Telescope and the Keck Telescope. The star has a mass of $1.36M_{\\odot}$ and metallicity of [Fe/H]= +0.26. The RV variations are consistent with a circular orbit with a period of $696.0$ days and a velocity semiamplitude of $200.0\\mps$, which yield a minimum-mass for the companion of $10.7\\mjup$ and semimajor axis of $1.71$ AU. Such super-massive objects with very low-eccentricities and hundreds of days period are uncommon among the ensemble of known exoplanets. ",
        "introduction": "Since 1995, precise RV measurements have unveiled more than 400 extrasolar planets from surveys that include roughly 3000 nearby stars \\citep[e.g.,][]{2007ARA&A..45..397U}. The planets show a great diversity in their physical attributes: masses of the planetary companions range from $4M_{\\oplus} (\\sim 0.013\\mjup)$ to larger than $13\\mjup$, with semimajor axes of 0.02AU - 6AU and orbital eccentricities of 0 - 0.9. A large number of planets have enabled us to discuss correlations between each orbital element and their masses in terms of planet formation and orbital evolution. For example, the diagram of orbital eccentricity against semimajor axis shows that eccentricities of planets with semimajor axis larger than 0.1 AU are almost uniformly distributed from 0 to 1, suggesting orbital evolution due to gravitational interaction between planets, while those inside 0.1 AU tend to be damped below 0.1 due to tidal circularization. The diagram of eccentricity against planet mass also shows another possible hint for planet formation that super-massive planets ($>5 \\mjup$) tend to have relatively high eccentricity ($>0.3$), while less massive planets have a wide range of eccentricities.  As the number of planet surveys has enlarged, diversity of planet-host stars has also been explored. Several surveys have focused on specified stellar properties, particularly stellar mass \\citep{2005A&A...443L..15B, 2010PASP..122..156A, 2008PASJ...60.1317S, 2007ApJ...665..785J} and metallicity \\citep{2005ApJ...620..481F,2009ApJ...697..544S}, in order to investigate correlations between the stellar properties and planetary parameters such as stellar mass and occurrence rate of planets. As a result from these surveys, the statistical correlations between stellar properties and planet parameters have begun to be unveiled. For example, occurrence rate of Jovian planets increases with stellar metallicity from 3\\% for stars with [Fe/H] $< 0$ to 25\\% for those with [Fe/H] $> +0.3$ \\citep{2005ApJ...620..481F}, and also increases with stellar mass from 2\\% around M dwarfs ($< 0.6 M_{\\odot}$) to approximately 9\\% around F and A stars \\citep[1.2 - 1.9 $M_{\\odot}$;][]{2007ApJ...665..785J}. Planetary mass could also be positively correlated with stellar metallicity and mass \\citep{2007A&A...472..657L}. Impacts of stellar properties on characteristics of planets, such as orbital eccentricity and multiplicity, are also expected to emerge as the number of planets grows.  The N2K program is a large scale international exoplanet-search project started in 2004, which consists of a sample of 2000 metal-rich solar-type stars \\citep{2005ApJ...622.1102F}. The search originally targeted short-period planets with a high-cadence observational strategy. However, it has also detected intermediate-period ($18 - 3810$ days) ones thanks to the long-term observations. From the collective N2K surveys, we have discovered 22 planets so far, which have a wide variety of mass (0.22 - 13.1 $\\mjup$) and orbital parameters ($a=$ 0.04 - 4.9AU, $e=$ 0 - 0.7), and orbit around a variety of host stars \\citep[$-0.11 < {\\rm [Fe/H]} < 0.37, 0.88 < M/M_{\\odot} < 1.31$;][]{2007ApJ...657..533W,2006ApJ...647..600J,2009PASP..121..613P}. The planets discovered from our surveys can thus contribute to our general understanding of planet formation and evolution depending on stellar properties.  Here we report the detection of an exoplanet orbiting the metal-rich solar type subgiant HD38801 discovered from the N2K sample at the Subaru Telescope and Keck observatory. The planet has a minimum mass of $10\\mjup$ and a low eccentricity. Such a planet has been rarely discovered so far around solar-type dwarfs and subgiants. We describe the characteristics of HD38801 in \\S \\ref{sec_str}. In \\S \\ref{sec_obs}, we present our observations, and the orbit of HD38801b. We summarize our results in \\S \\ref{sec_sum} and discuss formation scenarios of such a low-eccentricity super-massive planet. ",
        "conclusions": "\\label{sec_sum} We here reported the detection of a large mass planet in an almost circular orbit around a high metallicity K0 IV type star HD38801 from the precise RV observations at the Subaru Telescope and the Keck Telescope. The star has a mass of $1.36M_{\\odot}$ and metallicity of [Fe/H]=0.26. The planet is a new sample around a metal-rich and relatively high-mass star, which makes an important role to investigate dependence of planets on host star's properties.  One remarkable feature of this planet is its low orbital eccentricity ($e\\sim 0$) despite of the large mass ($\\msini =10.7\\mjup$) and the intermediate orbital distance ($a=1.7$AU). Such a planet with low eccentricity ($e<0.1$) and large-mass ($\\msini > 10\\mjup$) has been rarely discovered so far in the regions of $0.1\\ {\\rm AU}<a<5\\ {\\rm AU}$ around FGK dwarfs and subgiants. Only HD16760b ($\\msini=13\\mjup, a=1.08\\ {\\rm AU}, e=0.084$), which was recently discovered from the Subaru Telescope and the Keck Telescope \\citep{2009A&A...505..853B, 2009ApJ...703..671S}, falls into this category except for HD38801b. In the core-accretion model, massive companions are generally formed with circular orbits because of protoplanetary disk's circular motion. The circular orbit of HD38801b is consistent with this scenario. Tidal interaction between the host star and the planet cannot explain the circular orbit because the stellar radius is only $2.5R_{\\odot}$ ($\\sim0.01$AU), which is too small to circularize the planetary orbit at 1.7 AU. Contrary to HD 38801b, most of the intermediate-period ($>18$ days) and large-mass ($>5\\mjup$) planets discovered around FGK dwarfs and subgiants so far tend to reside in eccentric orbits. The origin of such planets still remains to be solved. It is generally thought that eccentric planets are formed by planet-planet scattering in multiple planetary systems. The scenario suggests formation of multi super-massive planets in a system and predicts existence of scattered outer companions as massive as inner ones, which can be good targets for direct imaging. Other scenarios such as giant impact\\citep{2008A&A...482..315B},and disk instability\\citep{1998ApJ...503..923B} and Kozai-mechanism when they are in binary systems \\citep{1962AJ.....67..591K} have also been proposed to explain for formation of such massive eccentric planets. Interestingly, planets with similar properties to HD38801b have been discovered around intermediate-mass giants. Comparing populations of such planets between different types of stars would provide a hint on formation scenario and evolution for the planets.  The high-metallicity and the evolutionary stage of subgiant for HD38801 is another interesting feature of this system. It is well known that the detection rate of giant planets shows positive correlation with the metallicity of the host stars \\citep{2005ApJ...622.1102F}. Two scenarios have been proposed to explain the origin of the correlation; one is that more metal-rich stars tend to form more giant planets and the other is that the stellar surface is polluted by planetary debris \\citep[e.g.,][]{2004ApJ...616..567I}. At the early stage of the main sequence, a stellar surface could be polluted due to the accretion of planetesimals and protoplanets from debris disk and metallicty of the stellar surface could be enhanced. However, if the star has deep convective envelope, the metal-enhanced surface can be mixed with original stellar material and diluted by the convective flow. Thus, the existence of a giant planet around metal-rich evolved star as the subgiant HD38801, which probably has deep convective envelope, favors the former scenario, that is giant planets tend to form in metal-rich environment \\citep{2005ApJ...622.1102F}. It should be noted, however, that many planets have been discovered even in metal-poor (${\\rm [Fe/H]}<0$) giants, more evolved stars than subgiants, and no significant planet-metallicity correlation can be seen among them. \\cite{2007A&A...473..979P} proposed that the lack of metal-rich tendency is due to dilution by deeper convective envelope of giants than that of dwarfs and subgiants. To make sure the origin of the high metallicity of planet-harboring dwarfs, it is important to increase the number of planets discovered around evolved stars with $<1.5M_{\\odot}$, for which we can detect planets in all of the three evolutionary stages, dwarfs, subgiants, and giants, by precise Doppler technique.   We extend our gratitude to Akito Tajitsu and Tae-Soo Pyo for their great expertise and support of HDS observations from the Subaru Telescope. We also gratefully acknowledge the efforts and dedication of all Keck Observatory staff. This research has made use of the Simbad database, operated at CDS, Strasbourg, France. We thank all Hawaiian people for their hospitality and native Hawaiian ancestry on whose sacred mountain of Mauna Kea we are privileged to be guests. Without their generous supports, we would not be able to publish our new planet.  \\clearpage"
    },
    "1004/1004.4640_arXiv.txt": {
        "abstract": "Near future cosmology will see the advent of wide area photometric galaxy surveys, like the Dark Energy Survey (DES), that extent to high redshifts ($z\\sim 1 - 2$) but with poor radial distance resolution. In such cases splitting the data into redshift bins and using the angular correlation function $w(\\theta)$, or the $C_{\\ell}$ power spectrum, will become the standard approach to extract cosmological information or to study the nature of dark energy through the Baryon Acoustic Oscillations (BAO) probe. In this work we present a detailed model for $w(\\theta)$ at large scales as a function of redshift and bin width, including all relevant effects, namely nonlinear gravitational clustering, bias, redshift space distortions and photo-z uncertainties. We also present a model for the full covariance matrix characterizing the angular correlation measurements, that takes into account the same effects as for $w(\\theta)$ and also the possibility of a shot-noise component and partial sky coverage. Provided with a large volume N-body simulation from the MICE collaboration we built several ensembles of mock redshift bins with a sky coverage and depth typical of forthcoming photometric surveys. The model for the angular correlation and the one for the covariance matrix agree remarkably well with the mock measurements in all configurations. The prospects for a full shape analysis of $w(\\theta)$ at BAO scales in forthcoming photometric surveys such as DES are thus very encouraging. ",
        "introduction": "The statistical analysis of the distribution of structure at large astronomical scales has played a key role in advancing the field of Cosmology over the last 20 years. From shaping our understanding of complex processes driving galaxy formation and evolution to constraining the energy density content of the Universe.  The completion of large extra-galactic surveys such at the Sloan Digital Sky Survey (SDSS, \\pcite{york00}) and the 2dF Galaxy Redshift Survey (2dFGRS, \\pcite{colless03}) have bolstered our general knowledge in the field. Particularly more so when combined with the precise measurements of the Cosmic Microwave Background or the increasingly reach data from Supernova data \\cite{sanchez09,percival10,reid10,komatsu10} . One of the most promising, and eventually rewarding, challenges for the field of large scale structure today is the prospect for determining what drives the late time acceleration of the Universe \\cite{riess98,perlmutter99}. This is probed by the presence, in the clustering pattern of galaxies, of remanent features from the coupling of baryon and photons prior to recombination known as the Baryon Acoustic Oscillations (BAO). The BAO have already been detected in the spectroscopic samples of Luminous Red Galaxies (LRGs) in both SDSS and 2dFGRS \\cite{cole05,eisenstein05}, and studied in the early imagining data of SDSS \\cite{padmanabhan07}.  But the observational quest has only started. Several of the next-generation surveys will gain in area and depth, in exchange for a poorer determination of radial positions. In turn this imposes the need for angular clustering analysis in redshift bins of width few times larger that of the photometric error uncertainty at the given redshifts. The difficulty lies in that the projection in redshift bins lowers the clustering amplitude, erasing any particular feature and increasing the noise-to-signal ratio. The achievable precision of our photometrically estimated redshift will play a crucial role. We thus need to understand what affects the angular clustering pattern more severely.  The aim of our work is to tackle this problem, providing a well calibrated model for the clustering signal at large-scales as a function of angle, radial distance and bin width, deepening the available literature in the subject (e.g. \\pcite{padmanabhan07}, \\pcite{blake07} and references therein). We put particular effort in stressing the most relevant effects, redshift distortions and photo-z uncertainties, and how they interplay.  An equally important problem is to have the capability of estimating the full errors of the measurements. We thus provide a well tested description  of the complex error matrix characterizing the measurements of the correlation function in real situations, i.e. including effects of partial sky coverage, photo-z, redshift distortions, bias and shot-noise.  Both, the model for the correlation and the one for the error matrix, will be extensively tested against a very rich set of mocks redshift bins. This work should therefore be relevant for ongoing projects that use photometric redshift estimates like the Dark Energy Survey\\footnote{\\tt www.darkenergysurvey.org} (DES),  the Physics of the Accelerating Universe collaboration\\footnote{\\tt www.pausurvey.org} (PAU) and the the Panoramic Survey Telescope and Rapid Response System\\footnote{\\tt pan-stars.ifa.hawaii.edu} (PanStarrs). But also for upcoming imaging proposals such as the Large Synoptic Survey Telescope\\footnote{\\tt www.lsst.org} (LSST) and the ESA/Euclid\\footnote{\\tt www.euclid-imaging.net} survey.  This paper is organized as follows. In Sec.~\\ref{sec:model} and Sec.~\\ref{sec:errors} we discuss the models proposed in this work for the angular correlation function and its full error matrix respectively. In Sec.~\\ref{sec:simulations} we describe the set of ensembles of mock redshift bins implemented using a large volume N-body simulation. In Sec.~\\ref{sec:mocks.vs.model.I} and \\ref{sec:mocks.vs.model.II} we test the models against the mocks, under different regimes and assumptions. Section~\\ref{sec:conclusions} contains our conclusions and future lines of research. We also include several appendixes. In Appendix~\\ref{Appendix:A} we give a description of our model for the 3-d nonlinear matter correlation function. In Appendix~\\ref{Appendix:B} we study the limitations of the widely used Limber approximation. Finally, Appendix~\\ref{Appendix:C} gives a brief note on the covariance of the angular power spectrum induced by partial sky coverage. ",
        "conclusions": "\\label{sec:conclusions}  The field of large scale cosmological structure will undergo an unprecedented era in the immediate future with several large observational  campaigns proposed or under implementation. Many of these surveys, such as DES, PanStarrs and LSST, will use photometric techniques to estimate the radial position of galaxies instead of measuring their full spectra, which is more time-demanding . This gain allows to survey wider areas and fainter objects, but at the exchange of increasing the uncertainties in the true redshifts and degrading the radial clustering amplitude. The proposal is then to split the data into redshift bins and exploit the information available from angular correlation functions, provided with an accurate determination of the measurement errors. Yet, this goal can only be accomplished if we develop accurate models for the signal and its errors that take into account all relevant effects and are robustly tested in realistic scenarios. In this paper we addresseded this issue in a comprehensive way.   We first developed an extensive set of mock catalogues (in the form of redshift bins) reproducing the angular coverage and radial distribution of a photometric survey like DES. We did it in increasing steps of realism, first in real space, then including redshift distortions or photometric errors and finally altogether. For this we used a large N-body simulation (of $\\sim 450 \\Gpccube$ simulated volume) provided by the MICE collaboration ({\\tt http://www.ice.cat/mice}). These mocks can be regarded as independent realizations as their volume overlap  is minimal and therefore provide a unique statistical framework for model testing (see Table~\\ref{Table:mocks}). We claim they are also equivalent to a light-cone analysis as we are doing narrow redshift bins, which involve negligible evolution.  We next put forward a model for the angular correlation function $w(\\theta)$ accounting for all the relevant effects, namely bin projection, nonlinear gravitational evolution, linear bias, redshift space distortions and photo-z errors.  An exhaustive comparison of our model for $w(\\theta)$ against the mock measurements showed a remarkably good agreement for a wide range of $\\theta$ and scenario (real space, photo-z, RSD and photo-z + RSD), validating the treatment of the different effects and opening the door to the use of this probe for real data analysis. Nonlinear gravitational evolution produces minor distortions in the correlation pattern after the bin projection.  In turn, analysis of halo angular clustering showed a very good consistency with a linear bias assumption. The interplay of photo-z and redshift distortions is the most important consideration regarding the shape of $w(\\theta)$.  Redshift space distortions introduces a large and scale dependent enhancement of $w(\\theta)$, that can reach a factor of a few at BAO scales (see Fig.~\\ref{fig:wtheta_photoz_zspace}). For our widest bin (where the effect should be least important) it still rises the amplitude of $w(\\theta)$ by $\\sim 50\\%$ at $\\theta_{BAO}$ (with respect to real space). Conversely, photo-z effects lower the clustering amplitude by extending the effective bin projection. For example, for the widest bin mentioned before (with $\\Delta z \\sim 4 \\sigma_z$) we find that the two effects counter-act each other at $\\theta_{BAO}$, but leave a scale dependent signal towards smaller angles. For narrower bins photo-z dominates, but redshift distortions is certainly not negligible. These trends were concluded from both, our photo-z + RSD mocks, and the analytical model.  In turn, we showed that the Limber approximation should not be used in precision analysis of large scale clustering as it leads to the incorrect shape of $w(\\theta)$ in the full range of interesting scales, and severely misestimates the amplitude of the $C_\\ell$ spectra for $\\ell \\simlt 40-50$. This is convincingly shown in Figs.~\\ref{fig:Cl},~\\ref{fig:wtheta} and \\ref{fig:w2Limber}. Another interesting issue considered was the impact of uncertainties in the true redshift distribution of objects. This showed an important aspect since it can lead to several percent changes in both the shape of $w(\\theta)$ and the BAO peak position given the accuracy of present photo-z estimators.  We would like to highlight that, in the process of describing $w(\\theta)$, we have also investigated a model for the 3-d matter correlation function that is able to reproduce the clustering signal in a broad range of scales and redshifts with only 2 parameters. This, discussed in detail in Appendix \\ref{Appendix:A}, can be of grand interest for future spectroscopic surveys such as BOSS, Hetdex and WiggleZ.   We have made an equally exhaustive effort in modeling and testing the full error matrix characterizing the measurements of $w(\\theta)$. The covariance matrix is often estimated from the data itself, using internal or re-sampling methods such as Jack-knife or bootstraping. However, their limitation is still a matter of some debate \\cite{norberg09}. Having a full theoretical model is thus very suitable for present and future analysis.  We took into account partial sky coverage by assuming that ${\\rm Cov}(\\theta,\\theta^{\\prime})$ scales as $f_{sky}^{-1}$. The full sky situation is then easily treated by translating errors from harmonic space, where the covariance matrix is diagonal and proportional to $C_\\ell$. Through the angular power spectra we included the same effects considered for $w(\\theta)$ into ${\\rm Cov}(\\theta,\\theta^{\\prime})$ (photo-z, redshift distortions, bias) for a typical survey with $f_{sky}=1/8$.  Our modeling of errors recovers the correct variance in $w(\\theta)$ as measured in the mocks for a wide range of bin configurations , from low to high redshift ($z\\sim 0.3-1.1$) and from thin to thick bins ($100\\Mpc-550\\Mpc$). And this conclusion extend to the more realistic cases where we included photo-z effects and redshift space distortions. In addition we used different halo samples to study cases where the shot-noise was comparable or larger than the sample variance component of the error. This regime was also nicely described analytically by adding a standard Poisson shot-noise contribution to the variance of the $C_\\ell$ spectra.  Moreover, and thanks to the large number of mocks constructed, we measured the full covariance matrix with high precision in different configurations. We find that at least $150-200$ mocks are necessary for a well defined reduced covariance, but this is discussed more properly in \\pcite{cabre10}. Remarkably in all cases tested, the modeling recovers very accurately the true error matrix.  In a parallel line of research we have tested the recovery of cosmological parameters using our model for $w(\\theta)$ and the theoretical expression of the covariance matrix. And compared this with the same analysis but using the {\\it true covariance} as measured in the mock ensembles. Indeed, the best-fit values and errors contours coincide for both approaches \\cite{cabre10} giving very encouraging prospects for the use of our analytical expressions in real data analysis or in realistic forecasts of upcoming photometric surveys."
    },
    "1004/1004.2857_arXiv.txt": {
        "abstract": "We present the results from the {\\it Suzaku} X-ray observations of five flat-spectrum radio quasars (FSRQs), namely PKS\\,0208$-$512, Q\\,0827+243, PKS\\,1127$-$145, PKS\\,1510$-$089 and 3C\\,454.3. All these sources were additionally monitored simultaneously or quasi-simultaneously by the {\\it Fermi} satellite in gamma-rays and the {\\it Swift} UVOT in the UV and optical bands, respectively. We constructed their broad-band spectra covering the frequency range from $10^{14}$\\,Hz up to $10^{25}$\\,Hz, and those reveal the nature of high-energy emission of luminous blazars in their low-activity states. The analyzed X-ray spectra are well fitted by a power-law model with photoelectric absorption. In the case of PKS\\,0208$-$512, PKS\\,1127$-$145, and 3C\\,454.3, the X-ray continuum showed indication of hardening at low-energies. Moreover, when compared with the previous X-ray observations, we see a significantly increasing contribution of low-energy photons to the total X-ray fluxes when the sources are getting fainter. The same behavior can be noted in the {\\it Suzaku} data alone. A likely explanation involves a variable, flat-spectrum component produced via inverse-Compton (IC) emission, plus an additional, possibly steady soft X-ray component prominent when the source gets fainter. This soft X-ray excess is represented either by a steep power-law (photon indices $\\Gamma \\sim 3-5$) or a blackbody-type emission with temperatures $kT\\sim0.1-0.2$\\,keV. We model the broad-band spectra spectra of the five observed FSRQs using synchrotron self-Compton (SSC) and/or external-Compton radiation (ECR) models. Our modeling suggests that the difference between the low- and high-activity states in luminous blazars is due to the different total kinetic power of the jet, most likely related to varying bulk Lorentz factor of the outflow within the blazar emission zone. ",
        "introduction": "Observations with the EGRET instrument ($30$\\,MeV to $30$\\,GeV; Thompson et al. 1993) on board the Compton Gamma-Ray Observatory (CGRO) have resulted in detection of $\\gamma$-ray emission from a few hundred astrophysical sources, 66 of which were securely associated with active galactic nuclei (AGNs; e.g., Hartman et al. 1999). Most of the AGNs detected by EGRET show characteristics of the blazar class. Observationally, this class include flat-spectrum radio quasars (FSRQs) and BL Lac objects. FSRQs have strong and broad optical emission lines, while the lines are weak or absent in BL Lacs. During the first three months of the {\\it Fermi} Large Area Telescope's (LAT) all-sky-survey, 132 bright sources at high Galactic latitudes ($|b|>10^{\\circ}$) were detected at a confidence level greater than $10\\,\\sigma$ (Abdo et al. 2009a). As expected from the EGRET observations, a large fraction (106) of these sources have been associated with known AGNs (Abdo et al. 2009b). This includes two radio galaxies (Centaurus\\,A and NGC\\,1275; Abdo et al. 2009c) and 104 blazars consisting of 58 FSRQs, 42 BL Lac objects, and 4 blazars with unknown classification based on their Spectral Energy Distribution (SED).  The radio-to-optical emission of luminous blazars of the FSRQ type is known to be produced by the synchrotron radiation of relativistic electrons accelerated within the outflow, while the inverse Compton (IC) scattering of low-energy photons by the same relativistic electrons is most likely responsible for the formation of the high energy X-ray-to-$\\gamma$-ray component. In addition, it is widely believed that the IC emission from FSRQs is dominated by the scattering of soft photons external to the jet (external Compton radiation, ECR). These photons, in turn, are produced by the accretion disk, and interact with the jet either directly or indirectly, after being scattered or reprocessed in the broad-line region (BLR) or a dusty torus (DT; see, e.g., Dermer \\& Schlickeiser 1993; Sikora et al. 1994; B\\l a\\.zejowski et al. 2000). Other sources of seed photons can also contribute to the observed IC radiation, and these are in particular jet synchrotron photons through the synchrotron self-Compton process (hereafter SSC; Maraschi et al. 1992; Sokolov \\& Marscher 2005).  In this context, detailed X-ray studies offer a unique possibility for discriminating between different proposed jet emission models, since those scenarios predict distinct components to be prominent in blazar spectra around keV photon energies. For example, in the soft X-ray range a break is expected in the ECR/BLR model, tracking the low-energy end of the electron energy distribution (Tavecchio et al. 2000; Sikora et al. 2009). Indeed, both the {\\it XMM-Newton} and the {\\it Suzaku} X-ray data of RBS 315 show ``convex\" spectra (Tavecchio et al. 2007). Such a curvature, on the other hand, can be alternatively accounted for by an excess absorption below 1 keV over the Galactic value, or by an intrinsic curvature in the electron energy distribution. Furthermore, the situation can be more complex, with the simultaneous presence of yet additional components, such as the high-energy tail of the synchrotron continuum, SSC emission, or the narrow-band spectral feature originating from the ``bulk Comptonization\" of external UV (disk) radiation by cold electrons within the innermost parts of relativistic outflow (Begelman \\& Sikora 1987; Sikora \\& Madejski 2000; Moderski et al. 2004; Celotti et al. 2007).   Ghisellini et al. (1998) have studied the spectral energy distribution of 51 EGRET-detected $\\gamma$-ray loud blazars and have applied the SSC+ECR model to the spectra of these sources. Although most of the broadband data collected by Ghisellini et al (1998) corresponded to non-simultaneous measurements, those authors discovered clear trends and correlations among the physical quantities obtained from the model calculations. In particular, they found an evidence for a well-defined sequence such that the observed spectral properties of different blazar classes (BL Lacs and FSRQs) can be explained by an increasing contribution of an external radiation field towards cooling jet electrons (thus producing the high-energy emission) with the increasing jet power. As a result, while the SSC process alone may account for the entire high-energy emission of low-power sources (BL Lacs), a significant contribution from the ECR is needed to explain the observed spectra of high-power blazars (FSRQs). Meanwhile, when focusing on one particular object, Mukherjee et al. (1999) reported that they found a similar trend in the different spectral states of PKS\\,0528+134. They studied the sequence of flaring and low-flux states of the source and found that the SSC mechanism plays a more important role when the source is in a low state, and the ECR mechanism is the dominant electron cooling mechanism when the source is in a high $\\gamma$-ray state (see in this context also Sambruna et al. 1997).  In order to understand the blazar phenomenon and the differences between BL Lacs and FSRQs, as well as the origin of spectral transitions in a particular object, one has to obtain truly simultaneous coverage across the entire spectrum, during both flaring and low-activity states. However, past $\\gamma$-ray observations in low-activity states have been limited to only a few extremely luminous objects, such as PKS\\,0528-134 or 3C\\,279. Only now, with the successful launch of the {\\it Fermi} satellite and the excellent performance of the {\\it Suzaku} instruments, do we have an opportunity to study high-energy spectra of blazars with substantially improved sensitivity, and therefore can probe the different states of the sources' activity.  In this paper, we report the high-sensitivity, broadband {\\it Suzaku} observations of five FSRQs, namely PKS\\,0208$-$512, Q\\,0827+243, PKS\\,1127$-$145, PKS\\,1510$-$089, and 3C\\,454.3, which were bright gamma-ray sources detected by EGRET. Additionally, all of these sources were monitored simultaneously or quasi-simultaneously by the {\\it Fermi} LAT and {\\it Swift} Ultraviolet/Optical Telescope (UVOT; Roming et al. 2005). These broadband and high-sensitivity observations allow us to reveal the characteristics of the high-energy IC continuum in the low-activity states of luminous blazars. The paper is organized as follows:  in $\\S$2, we describe observation and data reduction in the X-ray ({\\it Suzaku}), UV-optical ({\\it Swift} UVOT) and $\\gamma$-ray ({\\it Fermi} LAT) domains. In $\\S$3, we present the broad-band analysis results. Finally, in $\\S$\\,4 we discuss the constraints on the jet parameters and speculate on the the origin of different activity states in luminous blazars. Throughout the paper we adopt the cosmological parameters $H_0 = 71$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\rm M} = 0.27$, and $\\Omega_{\\Lambda} = 0.73$. ",
        "conclusions": "We have presented the observations and analysis of the data for the $\\gamma$-ray-loud blazars, PKS\\,0208$-$512, Q\\,0827+243, PKS\\,1127$-$145, PKS\\,1510$-$089, and 3C\\,454.3, obtained with the {\\it Suzaku}, {\\it Swift} UVOT and {\\it Fermi} LAT. Observations were conducted between 2008 October and 2009 January. These observations allowed us to construct broadband spectra of the sources in the low $\\gamma$-ray activity state, covering optical to GeV photon energy range. Our results are as follows:  \\begin{enumerate}  \\item The X-ray spectra of five FSRQs are well represented by an absorbed hard power-law model ($\\Gamma\\sim1.4-1.7$). For PKS\\,0208$-$512, PKS\\,1127$-$145, and 3C\\,454.3, the fitted absorption column is larger than the Galactic value (but we note that the ``excess absorption'' is not a unique representation of X-ray spectra of those blazars). Compared with previous X-ray observations, we see a trend of increasing apparent X-ray absorption column with increasing high-energy luminosity of the source.  \\item {\\it Suzaku} observations reveal spectral evolution of the X-ray emission: the X-ray spectrum becomes harder as the source gets brighter. Such spectral changes are most likely due to the underlying and steady low-energy spectral component which becomes prominent when the inverse-Compton emission gets fainter. This soft X-ray excess can be explained as a contribution of the high-energy tail of the synchrotron component, or bulk-Compton radiation.  \\item We adopt the location of the blazar emission region to be outside of the immediate vicinity of the accretion disk but within the BLR, and within the context of this model, we find that the contribution of the synchrotron self-Compton process to the high-energy radiative output of FSRQs is negligible even in their low-activity states.  \\item We argue that the difference between the low- and high-activity states in luminous blazars is due to the different total kinetic power of the jet, most likely related to varying bulk Lorentz factor of the outflow within the blazar emission zone.  \\end{enumerate}"
    },
    "1004/1004.5470_arXiv.txt": {
        "abstract": "{}{We examine the effects of the horizontal turbulence in differentially rotating stars on the GSF instability and apply our results to pre-supernova models.}{We derive the expression for the GSF instability with account  of the thermal transport and  smoothing of the $\\mu$--gradient by the  horizontal turbulence. We apply the new expressions in numerical models of a 20 M$_{\\odot}$ star.} {We show that if $N^2_{\\Omega} < 0$ the Rayleigh--Taylor instability cannot    be killed by the stabilizing thermal and $\\mu$--gradients, so that  the GSF instability is always there and we derive the corresponding diffusion coefficient. The GSF instability grows towards the very latest stages of stellar evolution. Close to the deep convective zones in pre-supernova stages, the transport coefficient of elements and angular momentum  by the GSF instability can very locally be larger than the shear instability and even as large as the thermal diffusivity. However the zones over which the GSF instability is acting are extremely narrow and there is not enough time left before the supernova explosion for a significant mixing to occur.  Thus, even when the inhibiting effects of the $\\mu$--gradient are reduced by the horizontal turbulence, the GSF instability  remains insignificant for the evolution.} {We conclude that the GSF instability in pre-supernova stages cannot be held responsible for the relatively low rotation rate  of pulsars compared to the  predictions of rotating star models.}    \\keywords {stars: massive - evolution - interiors - rotation (instability) - pulsar general (rotation)} \\titlerunning{The GSF instability and turbulence} ",
        "introduction": "The comparison of the observed rotation rate  of pulsars and stellar models in the pre-supernova stages indicate that most stars are losing more angular momentum than currently predicted (\\cite{HegerLW00,HMMXII}). Normally, the conservation of the central angular momentum of a presupernova model would lead to a neutron star spinning with a period of 0.1 ms, which is about two orders of magnitude faster than the estimate for the most rapid pulsars at birth. The question has arisen whether some rotational instabilities may  play a role in dissipating the angular momentum. We can think  in particular of the Golreich-Schubert-Fricke (GSF) instability (\\cite{GoldreichS67,Fricke68}), which has a negligible effect in the Main--Sequence phase and  which may play some role in the He--burning and more advanced phases (\\cite{HegerLW00}), in particular when there is a very steep $\\Omega$--gradient at the edge of the central dense core. This instability is generally not accounted for in stellar modeling.   The aim of this article is to examine whether the  GSF instability is important in the pre-supernova stages, when account is given to the effect of the horizontal turbulence in rotating stars which  reduces the stabilizing effects of the   $\\mu$--gradient.  Sect. 2 recalls the basic properties of the GSF instability, Sect. 3  those  of the horizontal turbulence. The effects of turbulence  on the GSF instability are examined in Sect. 4. Sect. 5 show the results of the numerical models. Sect. 6 gives the conclusion. ",
        "conclusions": "We have examined the effects of the horizontal turbulence on the GSF instability. This instability is present as soon as $N^2_{\\Omega}$ is smaller than zero, whatever the effects of the stabilizing $\\mu$--gradients.  On the whole, the numerical models of rotating stars show that the diffusion coefficient by the GSF instability grows towards the very latest stages of stellar evolution, however the zones over which it is acting are extremely narrow and there is not enough time left before the supernova explosion for a significant mixing to occur. Thus, even when the inhibiting effect of the $\\mu$--gradient is reduced by horizontal turbulence, the GSF instability is unable to  smooth the steep  $\\Omega$--gradients and to significantly transport matter.  We conclude that the amplitude and spatial extension of the GSF instability makes it unable to reduce the angular momentum of the stellar cores in the pre-supernova stages by two orders of magnitude. Therefore, other mechanisms such as magnetic fields (\\cite{Sp02}, \\cite{ROTMII}, \\cite{MZ05}, \\cite{ZBM07}) and gravity waves  (\\cite{TC05}, \\cite{MTPZ08}) must be further investigated.\\\\     {\\bf{Appendix: some approximations} for meridional circulation}\\\\  The coefficient $D_{\\mathrm{GSF}}$ requires, because of the horizontal turbulence, the knowledge of the components $U_2$ and $V_2$ of the meridional circulation. If  the solutions of the 4$^{th}$ order  system of equations governing meridional circulation are not available, some approximations may be considered. We note that the same problem would  occur for Eq.~(\\ref{vgsf}) by Endal and Sofia (\\cite{EndalS78}). As shown by stellar models, the orders of magnitude of $U_2$ and $V_2$ are the same. The numerical models give in general $V_2  \\sim  U_2/3$ and $\\left|2V_2-\\alpha U_2 \\right| \\sim V_2$. Using these orders of magnitude  in Eq.~(\\ref{o2}), we get \\begin{eqnarray} D_{\\mathrm{h}}\\, \\approx \\left(\\frac{\\beta}{10}\\right)^{1/2} \\left(r^2 \\, \\overline{\\Omega}\\right)^{1/2}\\bigg(\\frac{r\\,U_2}{3}\\bigg)^{\\frac{1}{2}} \\label{nuhapprox} \\end{eqnarray} \\noindent For $U_2$, various expressions can be used  taking into account the amount of differential rotation (\\cite{Maeder09}). We can also get an order of magnitude using the approximation for a mixture of perfect gas and radiation with a local angular velocity $\\Omega(r)$, ignoring  the effects of differential rotation on the circulation velocity  and the Gratton-\\\"{O}pik term which is large only  in the outer layers, \\begin{eqnarray} U_2(r) \\, =  \\, \\frac{16}{9} \\, \\frac{\\beta}{(32/3) -8 \\beta -\\beta^2} \\, \\frac{L(r)\\; r^2}{G \\, M^2_r} \\nonumber \\\\ \\frac{1}{\\left(\\nabla_{\\mathrm{ad}}- \\nabla+ \\frac{\\varphi}{\\delta}  \\nabla_{\\mu} \\right)} \\frac{\\Omega^2 r^3}{G \\,M_r}\\, , \\end{eqnarray}  \\noindent where the various quantities have their usual meaning."
    },
    "1004/1004.2120_arXiv.txt": {
        "abstract": "We present observations of a very massive galaxy at $z=1.82$ which show that its morphology, size, velocity dispersion and stellar population properties that are fully consistent with those expected for passively evolving progenitors of today's giant ellipticals. These findings are based on a deep optical rest-frame spectrum obtained with the Multi-Object InfraRed Camera and Spectrograph (MOIRCS) on the Subaru telescope of a high-$z$ passive galaxy candidate (pBzK) from the COSMOS field, for which we accurately measure its redshift of $z=1.8230$ and obtain an upper limit on its velocity dispersion $\\sigma_\\star<326$~\\kms{}.  By detailed stellar population modeling of both the galaxy broad-band SED and the rest-frame optical spectrum we derive a star-formation-weighted age and formation redshift of $t_\\text{sf}\\simeq1$--$2$~Gyr and $z_\\text{form}\\simeq2.5$--$4$, and a stellar mass of $M_\\star\\simeq3$--$4\\times10^{11}\\,M_\\sun$. This is in agreement with a virial mass limit of $M_\\text{vir}<7\\times10^{11}\\,M_\\sun$, derived from the measured $\\sigma_\\star$ value and stellar half-light radius, as well as with the dynamical mass limit based on the Jeans equations.  In contrast with previously reported super-dense passive galaxies at $z\\sim2$, the present galaxy at $z=1.82$ appears to have both size and velocity dispersion similar to early-type galaxies in the local Universe with similar stellar mass. This suggests that $z\\sim2$ massive and passive galaxies may exhibit a wide range of properties, then possibly following quite different evolutionary histories from $z\\sim2$ to $z=0$. ",
        "introduction": "Understanding the formation of massive elliptical galaxies remains a crucial unsolved issue of galaxy evolution. The recent discovery and the first redshift measurements, through deep ultraviolet (UV) rest-frame spectroscopy, of a substantial population of passively evolving galaxies at $z>1.4$ \\citep[e.g.,][]{cimatti:2004,mccarthy:2004,daddi:2005:pbzk} have shown that quenching of star formation in the most massive galaxies was already well under way at $z\\simeq2$.  A puzzling property of such objects has been revealed soon afterwards with some of them being found to have a factor of $\\simeq2-5$ smaller effective radii compared to local early-type galaxies (ETGs) of the same stellar mass \\citep[e.g.,][]{daddi:2005:pbzk,trujillo:2006,longhetti:2007,cimatti:2008,vandokkum:2008}, implying that they are $\\gtrsim 10$ times denser than their possible descendants in the local Universe.  Several alternative mechanisms have been proposed to make such compact ETGs grow in size so to finally meet the properties of local ETGs \\citep[e.g.,][]{khochfar:2006,fan:2008,naab:2009,labarbera:2009,feldmann:2010}, but no general consensus has yet emerged.  On the other hand, ETGs at $z>1.4$ with large effective radii, comparable to the local ETGs, have also been found \\citep[e.g.,][see also \\citealt{daddi:2005:pbzk}]{mancini:2010,saracco:2009}, indicating a diversity of structural properties in the ETG population at $z\\simeq2$. Moreover, possible effects  have also been discussed  that could bias size estimates towards lower values \\citep[e.g.,][]{daddi:2005:pbzk,hopkins:2009,mancini:2010,pannella:2009}.  An independent way to check these issues is by measuring stellar velocity dispersions ($\\sigma_\\star$): if high-$z$ ETGs are really super-dense, their $\\sigma_\\star$ should be much higher than that of local ETGs of the same mass.  \\citet{cappellari:2009:gmass} measured $\\sigma_\\star$ from deep UV rest-frame spectroscopy of a sample of 9 ETGs at $1.4<z<2.0$ from the GMASS survey \\citep{cimatti:2008}, finding two galaxies with similar stellar density and $\\sigma_\\star$ as local ETGs of the same mass.  The remaining galaxies have higher stellar densities and higher $\\sigma_\\star$ from their stacked spectrum, but still overlapping with the densest local ETGs.  In this respect, near-infrared (NIR) spectroscopy offers a great advantage for the galaxies at $z\\gtrsim1.4$ as the optical break at rest-frame $4000$~\\AA{} and \\cahk{} (the strongest spectral features of passively evolving galaxies) are redshifted into NIR, and the rest-frame optical continuum is much brighter than that in the rest-frame UV.  Thus, \\citet{vandokkum:2009} measured a velocity dispersion of $\\sigma_\\star\\simeq500$~\\kms{} for a compact ETG at $z=2.186$ using a deep NIR spectrum from \\citet{kriek:2009}. This value of $\\sigma_\\star$ is much higher than that of the most massive local ETGs, and would be consistent with the small half light radius measured for that galaxy.  In this letter we present a rest-frame optical spectrum of a massive, passively evolving high redshift galaxy candidate taken with the Multi-Object InfraRed Camera and Spectrograph \\citep[MOIRCS;][]{ichikawa:2006:moircs,suzuki:2008:moircs} at Subaru Telescope. A cosmology with $\\Omega_{\\text{M}}=0.3$, $\\Omega_{\\Lambda}=0.7$, and $H_0 = 70$ \\kms~Mpc$^{-1}$ is assumed. ",
        "conclusions": "Figure \\ref{fig:dynprop} compares the properties of the galaxy \\obj and other $z\\simeq2$ galaxies for which the same quantities have been measured \\citep{cappellari:2009:gmass,vandokkum:2009}.  The figure includes ETGs at $z\\simeq0.06$, selected from the Sloan Digital Sky Survey (SDSS) on the base of their red $u-g$ color and high S{\\'e}rsic index $n\\simeq4$ \\citep{blanton:2005}.  The dynamical and stellar masses agree very well for the high-$z$ objects, within a factor of $\\lesssim 2$.  Note that our massive galaxy has physical properties in good agreement with those of local ETGs of similar stellar mass. Our galaxy provides a second example of a very massive passively evolving system for which a stellar velocity dispersion has been measured (the GMASS objects of \\citeauthor{cappellari:2009:gmass}, which in Figure \\ref{fig:dynprop} also lie on the $z=0$ scaling relations, but have stellar masses below $10^{11}\\,M_{\\sun}$).  The ``normal'' size and velocity dispersion of our massive ETG is strikingly in contrast with the extreme properties (i.e., a very high $\\sigma_\\star = 510_{-95}^{+163}$~\\kms{} and a small $\\re=0.78\\pm0.17$~kpc) of the galaxy studied by \\citet{vandokkum:2009} with similar stellar mass ($2\\times10^{11}\\,M_\\sun$).  This suggests a substantial diversity in the physical properties of the \\textit{massive} ETG population at $z\\sim2$ including ``immature'', albeit virialized, systems --- which will have to evolve into normal $z=0$ massive galaxies through some physical processes which decrease their velocity dispersion and increase their sizes --- as well as ``mature'' ETGs, already on the scaling relationships of $z=0$ ETGs.  It is clear that many more observations of similar galaxies are required to establish which kind of ETG is commonest at high redshift: either the compact/high-$\\sigma_\\star$ objects like those found by \\citet{vandokkum:2009}, or the apparently normal, low-$\\sigma_\\star$ objects presented in this paper.  Also, nothing prevents our particular object to evolve further from its present state which mimics that of local ellipticals of the same mass.  For example, it may grow further and become a \\textit{brightest cluster galaxy} (BCG) or a cD galaxy.  For this reason, it would be important to estimate the volume number density of similar objects at high redshifts, and compare it to that of BCGs and cD galaxies.  To conclude, both very compact ETGs and ETGs following the local stellar mass-size and stellar mass-$\\sigma_\\star$ relations appear to co-exist at $z>1.4$ \\citep[see also ][]{mancini:2010}.  However the number of high-$z$ ETGs with individual measurement of the velocity dispersion is still extremely small. Increasing their sample is of great importance to understand the evolution of these galaxies, and in particular how and when they acquire their final structural and dynamical configuration.  This paper demonstrates that with reasonable telescope time several absorption features can be detected in the rest-frame optical spectrum of the high-$z$ ETGs, from which (at least for the most massive ETGs) the velocity dispersion and several stellar population properties can be derived."
    },
    "1004/1004.3689_arXiv.txt": {
        "abstract": "{Among the tracers of the earliest phases in the massive star formation process, methanol masers have gained increasing importance. The phenomenological distinction between Class I and II methanol masers is based on their spatial association with objects such as jets, cores, and ultracompact \\HII\\ regions, but is also believed to correspond to different pumping mechanisms: radiation for Class II masers, collisions for Class I masers. } {In this work, we have surveyed a large sample of massive star forming regions in Class I and II methanol masers. The sample consists of 296 sources, divided into two groups named \\high\\ and \\low\\ according to their [25--12] and [60--12] IRAS colours. Previous studies indicate that the two groups may contain similar sources in different evolutionary stages, with \\high\\ sources being more evolved. Therefore, the sample can be used to assess the existence of a sequence for the occurrence of Class I and II methanol masers during the evolution of a massive star forming region.} {We observed the 6~GHz (Class II) \\METH\\ maser with the Effelsberg 100-m telescope, and the 44~GHz and 95~GHz \\mbox{(Class I)} \\METH\\ masers with the Nobeyama 45-m telescope.} {We have detected: 55 sources in the Class II line (39 \\high\\ and 16 \\low , 12 new detections); 27 sources in the 44~GHz Class I line (19 \\high\\ and 8 \\low , 17 new detections); 11 sources in the 95~GHz Class I line (8 \\high\\ and 3 \\low , all except one are new detections). The detection rate of Class II masers decreases with the distance of the source (as expected), whereas that of Class I masers peaks at $\\sim 5$~kpc. This could be due to the Class~I maser spots being spread over a region $\\la$1~pc, comparable to the telescope beam diameter at a distance of $\\sim 5$~kpc. We also find that the two Class I lines, at 44~GHz and 95~GHz, have similar spectral shapes, confirming that they have the same origin.} {Our statistical analysis shows that the ratio between the detection rates of Class II and Class I methanol masers is basically the same in \\high\\ and \\low\\ sources. Therefore, both maser types seem to be equally associated with each evolutionary phase. In contrast, all maser species (including H$_2$O) have about 3 times higher detection rates in \\high\\ than in \\low\\ sources. This might indicate that the phenomena that originate all masers become progressively more active with time, during the earliest evolutionary phases of a high-mass star forming region.} ",
        "introduction": "\\label{intro}  Maser lines of water (H$_2$O), hydroxyl (OH), and methanol (CH$_3$OH) are commonly detected towards regions of high-mass star formation (e.g. Comoretto et al.~\\citeyear{comoretto}; Menten~\\citeyear{menten}; Kurtz et al.~\\citeyear{kurtz}; Pestalozzi et al.~\\citeyear{pestalozzi}). Among these, methanol masers are the most recently discovered, and their place in the star formation process is still not well understood. Menten~(\\citeyear{menten}) suggested to classify methanol masers into two groups: Class I and Class II. The latter typically coincide in position with hot molecular cores, ultracompact (UC) \\HII\\ regions, OH masers and near-IR sources (e.g. Minier et al.~\\citeyear{minier}; Ellingsen et al.~\\citeyear{ellingsen06}) and are believed to be radiatively pumped (Sobolev et al.~\\citeyear{sobolev} and references therein; Cragg et al.~\\citeyear{cragg05}). In contrast, Class~I masers are also found in massive star forming regions, but usually offset ($\\sim 0.1-1$ pc) from other masers, UC \\HII\\ regions, and bright infrared sources. They seem to be collisionally pumped (Cragg et al.~\\citeyear{cragg}) at the interface between molecular outflows/jets and the quiescent ambient material (Plambeck \\& Menten~\\citeyear{plambeck}). This scenario is supported by observations that reveal a coincidence between Class I methanol masers at 44~GHz and molecular shock tracers (Kurtz et al.~\\citeyear{kurtz04}, Voronkov et al.~\\citeyear{voronkov}).  Even though both species are believed to trace the earliest phases of the massive star formation process, it is still not clear whether any physical relation exists among them. Slysh et al.~(\\citeyear{slysh}) claim an anticorrelation between the intensity of Class I and II methanol masers in the same star forming region. On the other hand, Ellingsen~(\\citeyear{elli}) searched for Class~I masers at 95~GHz towards a sample of known 6~GHz Class~II masers and could not find any (anti)correlation between the two. It is also poorly understood whether a relationship exists between the occurrence of Class I and II masers and the evolution of the corresponding massive star forming region. van der Walt~(\\citeyear{vanderwalt}) estimated a lifetime of the 6~GHz Class II methanol maser of a few $10^4$ yrs, thus covering a consistent part of the early life of a massive star. Ellingsen~(\\citeyear{ellingsen06}) compared the infrared (GLIMPSE) colours of sources containing Class II methanol masers having or not having a Class I methanol maser, and found that those associated with Class I masers have redder GLIMPSE colours, suggesting that the sources hosting Class I masers are less evolved (see also Breen et al.~\\citeyear{breen}). However, until now an evolutionary sequence of methanol masers occurence has not been well established, partly because Class I masers are less studied and models for some Class I lines show that their excitation is very sensitive to the physical properties of the environment (Pratap et al.~\\citeyear{pratap}).  To investigate the existence of a sequence for the occurrence of Class I and II methanol masers along the evolution of a massive star forming region, one needs to study large samples of high-mass young stellar objects (YSOs) believed to be in different evolutionary stages. A large sample of massive YSO candidates with this property was identified by Palla et al.~(\\citeyear{palla}): the sources, all with $\\delta \\geq -30^{o}$, were selected on the basis of their large luminosities ($>10^3~L_\\odot$) and FIR colours typical of dense molecular clumps. These have been divided into two subsamples based on the IRAS colours\\footnote{We define the IRAS colours as [$\\lambda_{1}-\\lambda_{2}$] = $\\log_{10}(F_{\\lambda_1}/F_{\\lambda_2})$}: sources with $[25-12]>0.57$ and $[60-12]>1.3$ were named \\high, the others \\low. The threshold was taken from Wood \\& Churchwell~(\\citeyear{wec}), who suggested that UC \\HII\\ regions have IRAS colours above these limits. Palla et al.~(\\citeyear{palla}) searched for 22~GHz \\WAT\\ masers towards \\high\\ and \\low\\ sources, and found a higher detection rate in \\high\\ sources. In the last decade a series of studies aimed at investigating the environment associated with such sources (association with ammonia cores, centimeter and (sub-)millimeter continuum emission, \\CO\\ outflows: Molinari et al. 1996; Molinari et al. 1998a; Molinari et al. 2000; Brand et al. 2001; Zhang et al. 2001,~2005) indicated that \\low\\ and \\high\\ sources are massive stars in a very early evolutionary stage, with the \\low\\ group being dominated by the youngest sources. Two prototypical examples of \\high\\ and \\low\\ objects are respectively IRAS\\,20126+4104 (Cesaroni et al. 1997; Cesaroni et al. 1999) and IRAS\\,23385+6053 (Molinari et al. 1998b; Fontani et al.~2004; Molinari et al.~\\citeyear{mol08a}).  In this work, we have searched for 6~GHz (Class II), 44~GHz (Class I) and 95~GHz (Class I) \\METH\\ maser emission in the \\high\\ and \\low\\ sources of the Palla et al.~(\\citeyear{palla}) sample, with the aim to find a possible evolutionary sequence for the occurrence of the different masers. The observations were made with the Effelsberg 100-m telescope (at 6~GHz) and Nobeyama 45-m telescope (at 44 and 95~GHz). In Sect.~\\ref{obs} we give an overview of the observations performed and of the data reduction procedure adopted. The results are presented in Sect.~\\ref{res} and discussed in Sect.~\\ref{discu}. A summary of the main findings of this work is given in Sect.~\\ref{conc}. ",
        "conclusions": "\\label{conc}  We have searched for three methanol maser lines (6~GHz Class II, 44~GHz Class I, and 95~GHz Class I) towards a large sample of massive YSOs divided into two groups (\\high\\ and \\low ) on the basis of their IRAS colours (Palla et al.~\\citeyear{palla}). Previous studies indicate that the two groups are in different evolutionary stages, with \\high\\ sources being the more evolved. This work aims to use this sample to test a possible evolutionary sequence in the appearence of Class I and II methanol masers. The main findings of our study are the following: \\begin{itemize} \\item We have detected: 55 sources in the Class II line (39 \\high\\ and 16 \\low , 12 new detections); 27 sources in the 44~GHz Class I line (19 \\high\\ and 8 \\low , 17 new detections); 11 sources in the 95~GHz Class I line (8 \\high\\ and 3 \\low , all new detections). \\item The detection rates of all the masers observed are greater (by a factor $\\sim$3) in \\high\\ than in \\low\\ sources: the \\high/\\low\\ detection rate ratios are 2.9$\\pm1.3$ for the Class I 44~GHz line, $3.3\\pm1.6$ for the 95~GHz Class I line, and $2.5\\pm 1.1$ for the Class II line. All these values are similar to that found for \\WAT\\ masers, i.e. $3.1\\pm 1.1$. Going from \\low\\ to \\high\\ sources, we do not find any statistically significant difference in the {\\it relative} occurrence of Class I masers with respect to Class II masers. An analogous result holds for the ratio between the detection rates of CH$_3$OH and H$_2$O masers. A possible interpretation is that all maser species analysed in this work evolve similarly during the evolutionary phases corresponding to \\high\\ and \\low\\ sources. \\item The detection rate of the Class II masers decreases with the distance of the source, as expected, whereas that of the Class I masers peaks at $\\sim$5~kpc. We interpret this result with Class I maser spots being typically spread over a larger ($\\la$1~pc) region than Class II maser spots. \\item The spectra in the two Class I masers appear to have similar shapes in 7 out of the 11 sources detected in both lines, confirming a common physical origin. Their different detection rates cannot be explained only with the different noise levels at 44 and 95~GHz and we thus conclude that the 95~GHz line is intrinsically fainter.  \\end{itemize}  {\\it Acknowledgments.} We thank the anonymous referee for the constructive criticisms that helped us to shorten, focus, and substantially improve the paper. We are grateful to the Effelsberg and Nobeyama staff for their help during the observations. For part of this work, FF acknowledges support by Swiss National Science Foundation grant (PP002 -- 110504). The research leading to these results has received funding from the European Community's Seventh Framework Programme (FP7/2007--2013) under grant agreement No. 229517. This work is partially supported by a Grant-in-Aid from the Ministry of Education, Culture, Sports, Science and Technology of Japan (No. 20740113)."
    },
    "1004/1004.4256_arXiv.txt": {
        "abstract": "We study  a composition of normal and exotic matter which is required for a flat Emergent Universe scenario permitted by the equation of state (EOS)($p=A\\rho-B\\rho^{\\frac{1}{2}}$) and predict the  range of the permissible values for the parameters $A$ and $B$ to explore a physically viable cosmological model. The permitted values of the parameters are determined taking into account the $H(z)-z$ data  obtained from observations, a model independent BAO peak parameter and CMB shift parameter (WMAP7 data). It is found that although $A$ can be very close to zero, most of the observations favours a small and negative $A$. As a consequence, the effective Equation of State parameter for this class of Emergent Universe solutions remains negative always. We also compared the magnitude ($\\mu (z)$) vrs. redshift($z$) curve obtained in the model with that obtained from the union compilation data. According to our analysis the class of Emergent Universe solutions considered here is not ruled out by the observations. ",
        "introduction": "It is well known from the recent observations that the standard Big Bang model of cosmology fails to describe the present accelerating phase of the universe. The model is also pleagued by a time like singularity in the past. Accelerating phase of the universe can, however, be incorporated in a number of ways. A number of models e.g., models with modified theory of gravity \\citep{b19}, models with unusual matters like Chaplygin gas \\citep{b3,b4}, scalar and tachyon fields \\citep{b18} are taken into account to accommodate present phase of acceleration. There are other models based on mostly non-equilibrium thermodynamics and Boltzmann formulation which do not require dark energy \\citep{b15,b16,b17}.   Only very recently some other models appeared in the literature which discusses cold dark matter (CDM) and CDM interactions as alternative to the $\\Lambda$CDM model \\citep {n1,n2}. However, exploring singularity free cosmological models is an interesting area in cosmology and Emergent Universe scenario (EU) is one of the well known choices. A number of literature appeared which discussed EU model as it was free from initial singularity and the size of the universe was large enough so that quantum gravity effects were not important \\citep{b8,b7}. These models evolve from a static phase in the infinite past into an inflationary phase. The idea is in conformity with Lemaitre-Eddington concepts from early days of modern cosmology. If developed in a consistent manner an emergent universe model is capable of solving  some of the well known conceptual problems  not understood in the Big-Bang model. A model of an ever-existing universe, which eventually enters into the standard Big Bang epoch at some stage and consistent with features known to us today is worth considering. Recently an interesting EU model has been proposed by \\citet{b12} which requires some exotic matter in addition to normal matter as cosmic fluid. The model has been explored in a flat universe as such universe is supported by recent observations. Subsequently the EU model was taken up to examine the suitability of implementing it in the context of various theories \\citep{b1,b5,b2,b13}.  The EOS needed for the model proposed by \\citet{b12} is given by \\begin{equation} p=A\\rho-B\\rho^\\frac{1}{2}, \\end{equation} where $A$ and $B$ are unknown parameters of the theory. We note that different values of $A$ and $B$ pick up different composition of matters which may lead to an EU model. A similar EOS was considered in the literature as a double component dark energy model \\citep{w2} where the model parameters are constrained from Type Ia supernova data. The EOS considered by them is basically a special form of a more general EOS, $p=A\\rho -B\\rho^{\\alpha}$; which may be seen as a realization of Chaplygin gas (with $\\alpha <0$) \\citep{b3,b4}. It may be mentioned here that Chaplygin gas models were introduced in cosmology as an interpolation between a matter dominated era and a de Sitter phase. Later a modified model of Chaplygin gas was proposed \\citep{b10} to describe cosmological evolution. For example models like Modified Chaplygin gas serves well as an interpolation between radiative era and $\\Lambda$CDM era. \\citet{w2} showed that even with $\\alpha>0$ such interpolation is permissible and an EOS like one considered by \\citet{b12} may be considered as a phenomenological realization of string specific configuration.  The model proposed by \\citet{b12} developed the EU scenario in a very general way. Once one considers an EOS given by eq. (1), a class of EU solutions is permitted for $B>0$. The authors however showed that the above EU solution are not permitted with a minimally coupled scalar field. Also, it was found that the EU scenario automatically admitted a composition of three kinds of matter energy density in the universe all having their own way of evolution. This is certainly an interesting issue keeping in mind that the model has the provision for a large class of possible dark energy and dark matter candidate. It is thus worth to investigate the viability of such an EU model with the recent observational data. Nevertheless we intend to explore the allowed range of values of  the parameter $A$, for $B>0$ for a viable cosmological scenario permitted by observations. To determine the range of values for $A$ and $B$ for a viable cosmological model permitted by observations, we adopt here two independent techniques: (i)  Applying $\\chi^2$ minimization technique on $H(z)$ vs. $z$ data \\citep{b14}. Here we use $9$ data points given in Table 1, (ii) using joint analysis of $H(z)$ vs. $z$ data and a  model independent BAO peak parameter and (iii) using joint analysis of $H(z)$ vs. $z$ data, BAO peak parameter and CMB shift parameter. We explore here the suitability  of the model with the help of supernovae data (union compilation data) also. The plan of the paper is as follows : In sec.2 relevant field equations are obtained from Einstein field equation. In sec.3 the constraints on model parameters are determined  from $H(z)$ vs. $z$ data. Subsequently in sec. 4 and sec. 5 we obtain more stringent constraints on model parameters in accordence with the joint analysis with  model independent BAO peak parameter and CMB shift parameter. In sec.5, we draw the $\\mu (z) -z$ curve for our model to compare with that drawn using union compilation data \\citep{b9}. Finally, in sec. 6 we summarize our results and briefly discuss the results.\\\\ \\begin{table} \\begin{minipage}{140mm}  \\caption{$H(z) vs. z$ data} \\begin{tabular}{l|c|r} \\hline {\\it z Data} & $H(z)$ & $\\sigma$ \\\\ \\hline 0.00 & 73  & $ \\pm $ 8.0\t \\\\ 0.10 & 69  & $ \\pm $ 12.0 \\\\ 0.17 & 83  & $ \\pm $ 8.0 \\\\ 0.27 & 77  & $ \\pm $ 14.0 \\\\ 0.40 & 95  & $ \\pm $ 17.4 \\\\ 0.48 & 90  & $ \\pm $ 60.0 \\\\ 0.88 & 97  & $ \\pm $ 40.4 \\\\ 0.90 & 117 & $ \\pm $ 23.0 \\\\ 1.30 & 168 & $ \\pm $ 17.4 \\\\ 1.43 & 177 & $ \\pm $ 18.2 \\\\ 1.53 & 140 & $ \\pm $ 14.0 \\\\ 1.75 & 202 & $ \\pm $ 40.4 \\\\  \\hline \\end{tabular} \\end{minipage} \\end{table} ",
        "conclusions": "\\begin{figure} \\includegraphics[width=240pt,height=200pt]{uz.eps} \\caption{(Colour Online) $\\mu(z) \\: vs. \\: z$ curve comparism with supernovae data} \\end{figure}   \\begin{figure} \\includegraphics[width=240pt,height=200pt]{DP.eps} \\caption{(Colour Online) Plot of density parameter ($\\Omega$) for effective dark energy and effective matter content of Universe.} \\end{figure} Theoretically the class of EU solutions considered here can be realized for a composition of different kinds of matter \\citep{b12} depending on the model parameters $A$ and $B$. In this paper we determine allowed ranges for the model parameters (particularly those involved with the EOS i.e., $A$ and $B$). Considering a EOS required for a flat emergent universe we determine the constraints on the parameters using data available from cosmological observations. We note from the analysis that the EOS that permits a class of EU solution, considered here, should contain exotic matter ($A<0$, $B>0$). This is certainly not ruled out by the theory itself. It may be pointed out here that $H(z)$-$z$ data puts a bound on the model parameters which is further investigated in the light of the other observational data such as BAO peak parameter value and CMB shift parameter. In the above we estimate model parameters ($A$, $B$ and $K$). Most important point to be noted here that the later observations do not permit a positive value for the parameter $A$. Only small negative values seem to be allowed. Although positive $A$ values are permitted when we consider $H(z)$-$z$ data only but the best fit value is found to be negative. However the possibility that $A \\approx 0$ can not be entirely ruled out since our analysis permits values of $A$ which are even very close to zero ($A=0$) and the model may be realized in the presence of dust and dark energy. We also study the evolution of various cosmological parameters of the model. For example density parameter is on such important parameter. We plot density parameter for effective dark energy and effective matter content of the universe with the redshift in fig. 5. We note that almost $80\\%$ of the present matter-energy content is effective dark energy and baryonic and nonbaryonic matter constitutes the remaining part. So, as far as present budget of Dark Energy and Dark matter is concerned, EU differes very little from $\\Lambda$CDM model. However, as mentioned earlier, the class of EU solution considered here has provision for different composition of matter-energy in universe depending on the values of the parametes $A$ and $B$. It can also accommodate a cosmological constant in a special case. The effective equation of state ($\\omega_{eff}$) for EU remains negative always which is evident from fig. 6. The solid line in fig. 6 corresponds to the curve drawn using best fitted values. Dash and dotted curves are drawn with typical model parameters values within 95$\\%$ and 99$\\%$ confidence leves respectively. The transition of the universe from a deceleration phase to an accelerating phase in recent past is depicted from the curve of  deceleration parameter against redshift plotted in fig. 7. The solid curve describes the one drawn with best fitted values and dotted and dash  curves represent curves drawn with values within 99\\% and 95\\% confidence level respectively.  We conclude that the class of EU solutions considered here is not ruled out by the observations. However, this class of EU solutions admits different composition of matter-energies in the universe and the nature of composition depends on the value of parameter $A$ in particular. The observations do in fact severely constrain the nature of allowed composition of matter-energy by constraining the range of the values of the parameters for a physically viable model.  \\begin{figure} \\includegraphics[width=240pt,height=200pt]{weff.eps} \\caption{(Colour Online) Effective EOS parameter for EU ( $\\omega_{eff}$ ) is plotted with redshift. Plot with the best fitted values of model parameters $k=0.0102$, $A=-0.0176$ and $B=0.0103$ (Solid). Plot with values within 95\\% confidence (Dashing) and within 99\\%  (Dotted).} \\end{figure}  \\begin{figure} \\includegraphics[width=240pt,height=200pt]{Dec.eps} \\caption{(Colour Online)Deceleration parameter ($q$) vs. redshift ($z$). Plot with the best fitted values of model parameters(Solid), with values within 95\\% confidence (Dashing) and within 99\\%  (Dotted).} \\end{figure}"
    },
    "1004/1004.1473_arXiv.txt": {
        "abstract": "We use deep V-band surface photometry of five of the brightest elliptical galaxies in the Virgo cluster to search for diffuse tidal streams, shells, and plumes in their outer halos ($r>50$ kpc). We fit and subtract elliptical isophotal models from the galaxy images to reveal a variety of substructure, with surface brightnesses in the range \\muv = 26--29 \\magsec. M49 possesses an extended, interleaved shell system reminiscent of the radial accretion of a satellite companion, while M89's complex system of shells and plumes suggests a more complicated accretion history involving either multiple events or a major merger. M87 has a set of long streamers as might be expected from stripping of low luminosity dwarfs on radial orbits in Virgo. M86 also displays a number of small streams indicative of stripping of dwarf companions, but these comprise much less luminosity than those of M87. Only M84 lacks significant tidal features. We quantify the photometric properties of these structures, and discuss their origins in the context of each galaxy's environment and kinematics within the Virgo cluster. ",
        "introduction": "In a universe where structure grows hierarchically, the assembly of galaxies and galaxy clusters is characterized by an ongoing process of merging and accretion. These processes leave behind signatures which can be used to learn about the dynamical history of galaxies and clusters. On galaxy scales, structural asymmetries (\\eg Zaritsky \\& Rix 1997; Conselice 2006; Coziol \\& Plauchu-Frayn 2007) and diffuse tidal streams and shells (Malin \\& Carter 1980; Schweizer 1982, Tal \\etal 2009) in field galaxies point towards recent accretion events. For galaxy clusters, X-ray substructure (West, Jones, \\& Forman 1995; Mathiesen, Evrard, \\& Mohr 1999), kinematic and spatial asymmetries in galaxy distribution (Dressler \\& Shectman 1988; Knebe \\& M\\\"uller 2000; Oegerle \\& Hill 2001), and structure in the diffuse intracluster light (Feldmeier \\etal 2004; Mihos \\etal 2005, Rudick, Mihos, \\& McBride 2006) are all signatures of dynamically young events.  In field galaxies, the tidal debris formed during galaxy interactions typically remains loosely bound to the galaxy (or merger remnant), forming tidal tails, loops, and shells that spatially mix over many Gyrs (Hernquist \\& Spergel 1992; Hibbard \\& Mihos 1995; Gonz\\'alez-Garc\\'ia \\& Balcells 2005). This tidal material is typically found at very low surface brightnesses (\\muv $> 26.5$ \\magsec), such that very deep imaging is needed to detect it. Such features have been used to reconstruct the dynamical history of nearby galaxies, or determine the role mergers have played in the assembly of specific galaxies.  For example, the tight interleaved shells seen in some elliptical galaxies are indicative of a radial merger with a lower mass companion (Quinn 1984), while long, luminous tidal tails are the hallmark of major mergers (Toomre \\& Toomre 1972). Fainter tidal loops or streams are more ambiguous, and may come either from low mass accretion events (Bullock \\& Johnston 2005) or the delayed fallback of tidal tails from a major merger (Hernquist \\& Spergel 1992, Hibbard \\& Mihos 1995). The variety of tidal morphologies thus provide an important archaeological tool for understanding the formation and growth of galaxies in the local universe.  For galaxies within clusters, the picture is more complicated. The lifetime of diffuse tidal structure depends on two competing processes: galaxy-galaxy and galaxy-cluster interactions which produce tidal debris, and the dynamical heating and mixing of tidal debris in the cluster environment as it is incorporated into the cluster-wide ICL component (Rudick \\etal 2006, 2009). The long-lived shells and loops seen in many field ellipticals may be absent in cluster ellipticals due to rapid tidal stripping of this loosely bound material by the cluster potential (Mihos 2004). Alternatively, the presence of such structure in a cluster elliptical could indicate that the galaxy is only now being accreted into the cluster, so that cluster processes have not yet had an opportunity to strip its diffuse tidal structure.  While the variety of galaxy- and cluster-scale processes simultaneously at work during cluster assembly make it difficult to unambiguously extract the dynamical information held in tidal debris, it is clear that the information content in these features is rich, and provides important constraints on the dynamical history of these galaxies.  With these factors in mind, it is interesting to study the diffuse light around elliptical galaxies in a cluster environment. As part of our ongoing survey for ICL in the Virgo cluster (Mihos \\etal 2005), we have obtained deep, wide-field imaging of the Virgo ellipticals M87 (NGC 4486), M86 (NGC 4406), M84 (NGC 4374), M89 (NGC 4552), and M49 (NGC 4472). These galaxies all occupy different environments within the cluster (see, \\eg Binggeli 1999 for a review of the structure of the Virgo Cluster). As the central dominant elliptical in the Virgo cluster, M87 lives near the center of the cluster potential well (as defined by the X-ray emission; Bohringer \\etal 1994), and represents the center of the most massive subgroup of galaxies in the Virgo Cluster. M86 and M84 lie $\\sim1.3$\\degr\\ (370 kpc)\\footnote{In this work, we adopt a Virgo distance of 16 Mpc (see, \\eg Harris \\etal 1998; Ferrarese \\etal 2000; Mei \\etal 2007); at this distance, 1\\arcsec subtends 77.6 pc.} to the northwest of M87; they are separated from each other by 17\\arcmin\\ (78 kpc) in projection, close enough that, at faint surface brightnesses, their extended halos appear to merge together into a common envelope of light (Mihos \\etal 2005). This is likely a projection effect, however, as distance estimates from surface brightness fluctuations place M84 about 1 Mpc behind M86 (Mei \\etal 2007). Indeed, from a combined optical and X-ray analysis of the cluster, Schindler \\etal (1999) suggest that M86 sits at the center of its own subcluster of galaxies merging with the main body of the cluster. M89 lies $\\sim1.2$\\degr\\ (335 kpc) to the East of M87 and is the least luminous of our selected elliptical galaxies. Finally, lying 4.4\\degr\\ to the south of M87, M49 is the brightest elliptical in Virgo, and defines the center of another distinct Virgo subgroup (cluster B) which is dominated by spiral galaxies (Binggeli, Tammann, \\& Sandage 1987).  The different dynamical environments these ellipticals find themselves in is likely to translate to differences in the structure of their extended luminous halos.  In this work, we study the diffuse outer halos of these ellipticals, searching for tidal structures that may trace the dynamical histories of these galaxies. We use our deep imaging to fit and subtract a smooth elliptical fit to each galaxy's light profile, and identify tidal features in the residual images (for a similar approach, see, \\eg Canalizo \\etal 2007). We then measure the total luminosity and peak surface brightness of each of the cataloged features. We describe the observational dataset and analysis techniques in \\S2, and detail the results for each galaxy in \\S3.  Finally, we end with a discussion of these features in the more general context of the hierarchical assembly of galaxies and galaxy clusters in \\S4. ",
        "conclusions": "In summary, we have used deep, wide-field surface photometry to study the extended envelopes of five luminous Virgo ellipticals. We use the IRAF ELLIPSE task to fit elliptical isophotes to the galaxies' surface brightness profiles. Analytic fits to these isophotal models compare well to fits published elsewhere in the literature. Subtracting the isophotal models from the images, we identify a variety of low surface brightness tidal features -- streams, plumes, and shells -- in the outer envelopes of these elliptical galaxies, which we have quantified in terms of their total luminosity and peak surface brightness. We find that M87 is characterized by an extended diffuse halo with broad plumes and radial streams, but no sharp tidal shells or loops. M86 shows a number of small, relatively high surface streams, while M84 shows no evidence for significant substructure beyond its smooth elliptical isophotes.  In contrast, both M49 and M89 show complex (and in M49, hitherto undiscovered) systems of distinct shells and other tidal features.  The variety of structure we see in these Virgo ellipticals may well reflect differences in their history of accretion of smaller galaxies, their accretion history into the Virgo Cluster itself, or a combination of the two. To begin our discussion of environmental influences, we show in Figure \\ref{virgogeom} three orthogonal projections of the three dimensional positions of our sample galaxies in the Virgo cluster. We use the distances to each galaxy derived by Mei \\etal (2007) using the surface brightness fluctuation technique. We adopt this dataset largely because it has consistently derived distances for all of our galaxies, with very low internal errors. However, there are systematic uncertainties in the absolute distances; other techniques yield somewhat different results. Most notably, the position of M86 is quite uncertain -- while the Mei \\etal distances place M86 in the cluster core, in front of M84, distance determinations derived from planetary nebula put both M86 and M84 at the same distance, 1 Mpc behind M87 (Jacoby \\etal 1990). Finally, Figure \\ref{virgogeom} also plots the line of sight velocity vectors of each galaxy, after subtracting out a mean Virgo velocity of 1064 km/s (Binggeli 1999), and the estimated r200 (1.55 Mpc) for Virgo, determined by McLaughlin (1999).  Without transverse velocities for these galaxies, a full dynamical interpretation is impossible; nonetheless, we can glean useful information from this plot. As Virgo's central galaxy, it is no surprise that M87 has a low line of sight motion ($v_{\\rm rel} = 243$ km/s) with respect to the cluster as a whole. Sitting $\\sim$ 1 Mpc behind the cluster core, M84's low velocity ($-4$ km/s) suggests it is either near apocenter on a low angular momentum orbit -- likely having passed near the cluster center a few Gyr ago -- or moving on a more tangential orbit which keeps it out of the cluster center. As M49 is not projected onto the cluster core, its low velocity ($-67$ km/s) places little constraint on its orbit. However, the presence of an X-ray bowshock to the north of M49 argues that the galaxy is falling into Virgo's hot intracluster medium from the south (Irwin \\& Sarazin 1996). Only M86 and M89 have significant line of sight motion with respect to the cluster center ($-1308$ km/s and $-724$ km/s, respectively). M86 is either moving at high speed through the core, or just about to enter it, depending on the adopted distance, while M89 either passed through the core $\\sim$ 1-2 Gyr ago (if it lacks significant transverse velocity) or is on a more tangential orbit that keeps it further from the core.  \\begin{figure*}[] \\includegraphics[width=6.75in]{fig5.pdf} \\caption{The three dimensional geometry of the Virgo cluster, using line of sight SBF distances from Mei \\etal (2007). In each panel, the dotted line shows the virial radius ($r_{200}$) for Virgo from McLaughlin (1999). Arrows show the line of sight velocity of each galaxy, with the length of the arrow equal to the distance traveled in 1 Gyr.} \\label{virgogeom} \\end{figure*}  \\begin{deluxetable}{cccccc} \\tabletypesize{\\scriptsize} \\tablewidth{0pt} \\tablecaption{Galaxy Properties} \\tablehead{\\colhead{Galaxy} & \\colhead{$R_p$\\tablenotemark{a}} &  \\colhead{$r_{3d}$\\tablenotemark{b}} & \\colhead {$v_{\\rm rel}$\\tablenotemark{c}} & \\colhead{\\Lsub} &  \\colhead{\\fsub} \\\\ \\colhead{ }  & \\colhead{[Mpc]} & \\colhead{[Mpc]} & \\colhead{[km/s]} & \\colhead{[$10^8 L_\\odot$]}  & \\colhead{[\\%]} } \\startdata M49 & 1.23 & 1.32 &   -67 & 7.29 (0.25) & 0.47 (0.02) \\\\ M87 & $\\equiv 0$    & $\\equiv 0$    &   243 & 4.42 (0.46) & 0.39 (0.04) \\\\ M86 & 0.35 & 0.54 & -1308 & 0.91 - 1.27 & 0.10 - 0.14 \\\\ M84 & 0.42 & 1.32 &    -4 & 0.00 - 0.36 & 0.00 - 0.05 \\\\ M89 & 0.33 & 1.41 &  -724 & 9.74 (0.30) & 1.99 (0.06) \\\\ \\enddata \\tablenotetext{a}{Projected clustercentric distance, defined relative to M87} \\tablenotetext{b}{Three dimensional clustercentric distance} \\tablenotetext{c}{Velocity relative to cluster velocity of +1064 km/s (Binggeli \\etal 1999)} \\label{galprop} \\end{deluxetable}   We first examine the environmental question by looking to see if the amount of diffuse substructure in the galaxies correlates with their distance from the center of Virgo (taken to be the position of M87). We quantify the amount of substructure in two ways: the total luminosity in the features (\\Lsub), and their fractional luminosity (\\fsub, measured with respect to the total galaxy light). In the latter case, the galaxy luminosities are calculated analytically from our S\\'ersic surface brightness fits. We then calculate each galaxy's distance from the center of the Virgo Cluster using a combination of their projected distance from M87 on the sky and their line-of-sight distance from Mei \\etal (2007). These quantities are given in Table \\ref{galprop}.  For these five galaxies, we see no obvious correlation between the amount of substructure and cluster-centric distance, as shown in Figure \\ref{environ}. M87, M49, and M89 have significantly more substructure luminosity than the others, although in the case of M87 and M49, the substructures contribute much less to the total luminosity of the galaxies. Neither M84 nor M86 appear to have appreciable substructure, either in total or fractional luminosity. A model where substructure survival depends simply on cluster-centric distance is clearly too simplistic to explain the features seen in these Virgo ellipticals.  \\begin{figure*}[] \\includegraphics[width=6.75in]{fig6.pdf} \\caption{Substructure luminosity (left) and luminosity fraction (right) as a function of clustercentric distance.} \\label{environ} \\end{figure*}   However, M87 is not simply {\\it deep} in the cluster potential, it lives {\\it at the center} of the cluster. In this privileged position, it experiences a constant rain of smaller satellite galaxies falling in at high velocities, which will be tidally stripped and leave behind long streamers such as those seen extending to the NW of M87. Unlike the long-lived structures around isolated galaxies, streamers inside an active cluster environment are dispersed rapidly via interactions with other cluster members, typically within a few crossing times (Rudick \\etal 2009). In this position, M87 has both high creation rates and high destruction rates for its diffuse substructure, so that the window of opportunity for observing cold streams around M87 will be short.  While the amount of substructure does not seem to correlate with cluster-centric radius, there are interesting patterns in the morphological properties of the diffuse light in these ellipticals. M49 and M89 both have complex shell structures out to large radius, where the material would be loosely bound to the host galaxy. Any strong tidal forces that would occur during a galaxy's passage through the cluster would likely be sufficient to strip or at least significantly perturb the shells (see, \\eg Mihos 2004). The sharpness of the shells and the long dynamical timescales in the outskirts of these galaxies ($\\sim$ 0.5 Gyr) argue that these galaxies have not passed through the dense Virgo Core in the recent past.  In contrast, we see little evidence for sharp shell-like structures in M87, M86, or M84. Living at the center of Virgo, M87 experiences repeated encounters with cluster galaxies, and any dynamically delicate feature will quickly be destroyed.  In the radial streams visible to the northwest of M87, we are likely seeing very recent stripping of galaxies falling in on radial orbits.  The nature of the streams observed in M87 and M86 is qualitatively different as well. M87's NW and WNW streams can be traced out to $\\sim$ 1 degree (275 kpc) from the center of M87, while M86's streams are much smaller. The longest, the N Stream, extends 25\\arcmin\\ from M86, while the others are much smaller in length (only a few arcminutes in size) and closer to M86. Moving at such a high velocity near or through the cluster core, M86 simply may not be able to hang on to extended debris. The cluster tidal field may easily strip these streams, or they may lack coherency simply because of the motion of M86. A coherent stream at 300 kpc would have an orbital period of a few Gyr, a timescale over which M86 would have moved a few Mpc through the cluster. The combination of tidal effects and M86's high speed motion through the cluster would make it difficult for M86 to retain highly extended streamers.  Substructure is not the full story of the diffuse light in Virgo ellipticals, of course. Because these features are likely to be short-lived in a dynamically complex environment like a galaxy cluster, they will mix into a more diffuse, extended envelope. Whether or not this envelope is a dynamically or structurally distinct entity from the galaxy itself is a subject of considerable debate. Gonzalez \\etal (2005) argue that in clusters with clear BCGs, the intracluster light settles into a structurally distinct component from the galaxian light of the BCG, and model the two components using distinct $r^{1/4}$ profiles. They find that, in their sample of BCGs, the outer profile typically contains $\\sim$ 80-90\\% of the total luminosity, and has an effective radius 10--40 times larger than the inner component. If we consider M87 in this way, we find a smaller fraction of light in the extended envelope (50\\%) than typical in the Gonzalez \\etal sample. However, Virgo is somewhat different from those BCG clusters, in that it has a number of comparably bright galaxies (M87 at the center; M86 projected 0.5\\degr\\ away; and M49 2\\degr\\ to the south). With multiple bright galaxies and both spatial and kinematic substructure, Virgo may represent a dynamically less-evolved progenitor to the Gonzalez \\etal clusters, in which case the lower fraction of light in the envelope may be a signature of an still-developing ICL (\\eg Rudick \\etal 2006). This inference echoes that of K09, who argue based on the systematics of M87's S\\'ersic fit that the galaxy is at best only a {\\it weak} cD galaxy.  In this context, it is also interesting to look at M49. M49 is the dominant galaxy in the Virgo Southern Extension (VSE), and somewhat more luminous than M87. If clusters build through the hierarchical accretion of sub-clumps, the sub-clumps themselves may have their own diffuse light component, as suggested in the simulations by Rudick \\etal (2006). If the VSE is an evolved group now being accreted into the Virgo cluster, it may have a structure similar to (albeit smaller than) that of the BCG clusters of Gonzalez \\etal (2005). And indeed, the 2dV fit for M49 shows a higher fraction of light (76\\%) in the outer component than does M87, but with a much smaller characteristic scale ($r_{e,out}$(M49) = 20\\% $r_{e,out}$(M87)).  While these results are consistent with a scenario of a gradual, on-going buildup of ICL in Virgo, there are several important caveats. First, the 2dV profiles are not clearly superior to regular S\\'ersic fits, and the justification for dividing luminosity into an inner and outer component based on these fits is not strong. Even the choice of functional form for the division is the source of some debate -- Seigar \\etal (2007) argue that an inner S\\'ersic + outer exponential fit is a better description of the light profile for cD galaxies. However, given the fact that the additional free parameter of the 2dV fit over the S\\'ersic fit did not result in significantly better profile fits for our galaxies, we do not pursue these higher order fits here.  Given our results, we paint a plausible picture of the dynamics of the galaxies within the Virgo Cluster. Sitting at the center of the cluster, M87 experiences a rain of smaller galaxies which are being tidally stripped by the cluster potential, leading to the long diffuse streams seen to the NW of M87. Due to encounters with other galaxies in the cluster core, the lifetime of these streams is short and they mix away to continually build M87's extended envelope and Virgo's intracluster light. The combination of M86's high velocity and its passage through the cluster core makes it difficult to develop or retain very extended streams; instead, it possesses a system of small tidal streams much closer to the galaxy than seen in M87. The lack of tidal structures around M84 may be due to a possible recent passage through the cluster core, or M84 may simply have not experienced much recent accretion.  In contrast, both M49 and M89 display extended ($r \\sim 50-100$ kpc) shell systems, arguing that they have not experienced the strong tidal forces of the Virgo cluster core in the past Gyr or so. As the most luminous galaxy in the Virgo Southern Extension, it may be the dominant galaxy of group falling into the Virgo cluster for the first time, and its shell system reflects its own accretion of a smaller satellite on a radial orbit. As M49 falls into Virgo, its shell system will be disrupted and incorporated into the general ICL of the cluster. The complexity of M89's shell system argues not for an individual accretion event, but rather for multiple satellite accretions or a major merger in its past. Its system of tails, plumes, and shells will likely also be dispersed into Virgo's extended ICL as the galaxy orbits within the cluster environment.  Ultimately, however, morphology alone contains only limited information. A better understanding of these features and how they relate to the dynamical history of Virgo and its galaxies would come through studies of their kinematics and stellar populations. Kinematic information can come from identification and follow-up spectroscopy of planetary nebulae (\\eg Arnaboldi \\etal 2004, Doherty \\etal 2009) and globular clusters (\\eg C\\^ot\\'e \\etal 2003; Hwang \\etal 2008; Lee \\etal 2010) associated with the streams, while stellar populations can be studied using deep multiband surface photometry (Rudick \\etal 2010) or {\\it Hubble Space Telescope} imaging of resolved stars in the streams (\\eg Williams \\etal 2007). Such studies would give an integrated picture of the accretion and stripping processes at work in the Virgo Cluster."
    },
    "1004/1004.5585_arXiv.txt": {
        "abstract": "We study the generalized Chaplygin gas model (GCGM) using Gamma-ray bursts as cosmological probes. In order to avoid the so-called circularity problem we use cosmology-independent data set and Bayesian statistics to impose constraints on the model parameters. We observe that a negative value for the parameter $\\alpha$ is favoured in a flat Universe and the estimated value of the parameter $H_{0}$ is lower than that found in literature. \\\\\\\\ PACS number: 98.80.Es, 98.70.Rz ",
        "introduction": "\\par One of the most important problems of Modern Cosmology is the determination of the matter content of the Universe. The rotation curve of spiral galaxies \\cite{rotation}, the dynamics of galaxy clusters \\cite{dynamics} and structure formation \\cite{struc}, indicate that there is about ten times more pressureless matter in the Universe than can be afforded by the baryonic matter. The nature of this dark matter component remains unknown. Moreover, the Type Ia supernovae (SNe Ia) data indicates that the Universe is accelerating \\cite{super}. Models considering matter content dominated by an exotic fluid whose pressure is negative \\cite{press}, modified gravity theories such as $f(R)$ \\cite{mod} and the evolution of an inhomogeneous Universe model described in terms of spatially averaged scalar variables with matter and backreaction source terms \\cite{back} are some of the proposals to explain this current phase of the Universe. At the same time, the position of the first acoustic peak in the spectrum of CMB anisotropies, as obtained by WMAP, favours a spatially flat Universe\\cite{WM5}. Combining all these data and if we consider the matter content of the Universe dominated by a fluid with negative pressure we have a scenario with a proportion of $\\Omega_{m} \\sim 0.27$ and $\\Omega_{de} \\sim 0.73$, with respect to the critical density, for the fractions of the pressureless matter and dark energy, respectively. This scenario is usually called as the concordance cosmological model. \\par The question is to know what is the nature of the dark matter and dark energy components. For dark matter many candidates have been suggested such as axions, a particle until now undetected which would be a relic of a phase where the grand unified theory was valid \\cite{axion}, the lightest supersymmetric particle (LSP) like neutralinos \\cite{salati1} and the Kaluza-Klein particles \\cite{salati2} that are stable viable Weakly Interacting Massive Particles (WIMPs) and arise in two frameworks: In Universal Extra Dimensions \\cite{UED} and in some warped geometries like Randall-Sundrum \\cite{RS}. For the dark energy, in the hydrodynamical representations of matter, the most natural candidate is a cosmological constant, but there is a discrepancy of some $120$ orders of magnitude between its theoretical and observed values \\cite{cc}. For this reason, other candidates have been suggested like quintessence models that involve canonical kinetic terms of the self-interacting scalar field with the sound speed $c_s^2 = 1$ \\cite{quinte} and k-essence models that employ rather exotic scalar fields with non-canonical (non-linear) kinetic terms which typically lead to a negative pressure \\cite{kesse}. More recently, a string-inspired fluid has been evoked: The Chaplygin gas \\cite{chapl}, that appears as a promising candidate for the dark sector of the Universe. \\par The Chaplygin gas is represented by the equation of state \\begin{equation} p_c = - \\frac{A}{\\rho_c} \\quad , \\end{equation} where $p_c$ represents the pressure, $\\rho_c$ the fluid density and $A$ is a parameter connected with the sound speed. This equation of state is suggested by a brane configuration in the context of string theories \\cite{string}. However, a more general equation of state has been suggested  \\cite{gcg}: \\begin{equation} \\label{chapp} p_c = - \\frac{A}{\\rho^{\\alpha}_c} \\quad , \\end{equation} where again $p_c$ and $\\rho_c$ stand for the generalized Chaplygin gas component and $\\alpha$ is a new parameter, which takes the value $1$ for the traditional Chaplygin gas but values larger than $1$, or even negative may be considered. This is the so-called generalized Chaplygin gas. \\par Much observational data that has been used for comparison with the theoretical cosmological models like the generalized Chaplygin gas model (GCGM). The spectra of anisotropy of cosmic microwave background radiation \\cite{berto1}, baryonic acoustic oscillations \\cite{BAOcha}, the integrated Sachs-Wolfe effect \\cite{SW}, the matter power spectrum \\cite{mass}, gravitational lenses \\cite {lens}, X-ray data \\cite{raiox} and ages estimates of high-$z$ objects \\cite{ages} have been used in this sense. Also, constraints from combined data sources have been obtained in \\cite{combcha}. Another tool used to make this comparison is the Hubble diagram, the plot of redshift $z$ versus luminosity distance $d_L = \\sqrt{\\mathcal{L}/4\\pi\\mathcal{F}}$, where $\\mathcal{L}$ is the luminosity (the energy per time produced by the source in its rest frame) and $\\mathcal{F}$ is the measured flux, i.e., the energy per time per area measured by a detector. Normally, the SNe Ia data are considered good standard candles and are used to construct the Hubble diagram, because their luminosity are well known \\cite{super, SN2}. In particular, constraints on the Generalized Chaplygin gas have been studied in \\cite{SNe}. These assumptions rest on a foundation of photometric and spectroscopic similarities between high- and low-redshift SNe Ia. But this discussion is not yet finished \\cite{SN3}. The other problem comes from the fact that there still does not exist SNe Ia data with $z > 1.8$. To know the properties and behavior of dark energy for high values of $z$ we will have to wait for new data of the SNe Ia or to find other distance indicators.  In this sense, to extend the comparison between observational data and theoretical models at very high redshift we propose to use Gamma-ray bursts (GRBs) due to the fact that they occur in the range of high z beyond the SNe data found today \\cite{Bromm}.  The GRBs are jets that release $\\sim 10^{51} - 10^{53}$ ergs or more for a few seconds and becomes, in this brief period of time, the most bright object in the Universe. They were discovered in the sixties by the Vela satellites in the ``Outer Space Treaty\" that monitored nuclear explosions in space \\cite{hist}. Launched in 1991 The Burst and Transient Source Experiment on the Compton Gamma-Ray Observatory (BATSE on the Compton GRO) \\cite{Costa} observations concluded that the angular distribution of the GRBs on the sky is isotropic within statistical limits. This study ruled out the idea that the GRBs are galactic objects, but it is consistent with the bursts being extra-galactic sources at cosmological distances. More recently, the SWIFT mission (launched in 2004) has provided the most accurate GRB data, available in the Swift BAT Catalog. \\par The search for a self-consistent method to use the GRBs in cosmological problems is intense and promising. But the possibility of using GRBs as standard candles is not a simple question. GRBs are known to have several light curves and spectral properties from which the luminosity of the burst can be calculated once calibrated, and these can make GRBs into standard candles. Just as with SNe Ia, the idea is to measure the luminosity indicators, deduce the source luminosity, measure the observed flux and then use the inverse-square law to derive the luminosity distance. The difficulty arises when these indicators are a priori established through some cosmological model like the concordance one. This means that the parameters of the calibrated relations of luminosity/energy are still coupled to the cosmological parameters derived from a given cosmological model. This is the so called circularity problem. This problem appears in several works that have made use of these GRBs luminosity indicators as standard candles at very high redshift \\cite{circular}. It is possible to treat the circularity problem with a statistical approach \\cite{statist}. On the other hand, many papers have dealt with the use of so called Amati relation, or the Ghirlanda relation for this purpose \\cite{firmani}. However, as argued recently in \\cite{petrosian}, these procedure involve many unjustified assumptions which if not true could invalidate the results. In particular, many evolutionary effects can affect the final outcome. However, recently Liang {\\it et al.} \\cite{liang, liang2010, liang11} made a study considering SNe Ia as first-order standard candles for calibrating GRBs, the second-order standard candles. The sample in reference \\cite{liang} was calibrated from the 192 supernovae obtained in \\cite{davies}. The updated sample used in \\cite{liang2010, liang11} has been obtained and calibrated cosmology-independently from the Union2 (557 data points) compilation \\cite{Union2} released by the Supernova Cosmology Project Collaboration. In these articles the authors found relevant constraints on the Cardassian and Chaplygin gas model by adding to the GRB data the SNe Ia (Union2), the Shift parameter of the Cosmic Microwave Background radiation from the seven-year Wilkinson Microwave Anisotropy Probe and the baryonic acoustic oscillation from the spectroscopic Sloan Digital Sky Survey Data Release galaxy sample. The sample obtained in \\cite{liang2010} will be used in our analysis. These authors obtain the distance moduli $\\mu$ of GRB in the redshift range of SNe Ia and extend this result to very high redshift GRB ($z > 1.4$) in a completely cosmological model-independent way. This approach has been also studied in \\cite{calibration}.  Some analysis have been made with the GCGM and the GRBs as distant markers \\cite{GRB}. In the reference \\cite{bertolami} the authors build a specific distribution of GRB to probe the flat GCGM and the XCDM model. While the GCGM has an equation of state given by expression (\\ref{chapp}) the XCDM model is considered in terms of a constant equation of state $\\omega = p/\\rho < 0$. The main conclusion of this article is that the use of GRBs as a dark energy probe is more limited when compared to SNe Ia. We anticipate that we shall arrive at a similar conclusion. Moreover the XCDM model is better constrained than the GCGM. On the other hand, in \\cite{herman} the GCGM and the $\\Lambda$CDM model are compared by using the GRB and SNe Ia data to build the Hubble diagram. These authors show through the statistical analysis that the Chaplygin gas model (they use $\\alpha = 1$) have the best fit when compared with the data. Also they verify that the transition redshift between the decelerated and the accelerated state of the Universe occurs at $z \\sim 2.5 - 3.5$ rather than $z \\sim 0.5 - 1$ based on the analysis made with the SNe Ia. Here, for our purpose, we will assume the plausible assumption that GRBs are standard candles and we will use the data from Liang {\\it et al.} \\cite{liang2010}, calibrated cosmology-independently from the Union2 compilation of SNe Ia, to constraint the cosmological parameters of the GCGM. We want to show how GRB data could constraint different Chaplygin cosmologies. \\par This paper is organized as follows. In next section, we described a brief review of GCGM. In section $3$ the luminosity distance $d_L$ is obtained for the GCGM and compared with the observational data. Finally, in section $4$ we present our discussion and conclusions. ",
        "conclusions": "\\par  In this study we have analyzed the Chaplygin gas model with a sample of 42 GRBs. Although the use of GRBs as a cosmological tool is a promising way to probe cosmology at high redshifts we have verified that the available data is still insufficient to impose precise constraints in cosmological models. As observed in our analysis, the dispersion is still high when compared with others observational data sets. However we hope that with the future data from the final Swift BAT Catalog we will be able to put strong constraints on the dark energy/matter properties. \\par In our analysis, the unification scenario was not imposed from the beginning. This means that we allow an extra dark matter contribution ($\\Omega_{dm}$) in our calculations in order to probe whether the unification scenario is favoured. In our first analysis the free parameters ($\\bar{A},\\Omega_{dm}$ and $H_{0}$) of the Chaplygin gas ($\\alpha=1$) were well constrained. Our results are in agreement with the Supernova results \\cite{fabris03}. The only difference is that we find a lower value for the Hubble parameter, $H_{0}=51.3^{+9.2}_{-5.7}$ (1$\\sigma$). However, it is possible to find in the literature similar results for the parameter $H_{0}$ \\cite{H0}.  In our second analysis, in order to check the behaviour of the model when $H_{0}=72~km~s^{-1} Mpc^{-1}$ we leave $\\alpha$ free, that is the so called Generalized Chaplygin Gas Model. From Figs. \\ref{GRB1} and \\ref{GRB2} the unification scenario is again favoured. However, the uncertainties are still high. The parameter $\\alpha$ assumes a large negative value. There is no any peak in the parameter $\\alpha$ distribution and the probability remains constant for negative values. For the background dynamics the region ($\\alpha<-1$) represents a behavior different from the matter dominated phase when structures start to form. On the other hand, negative values for $\\alpha$ imply an imaginary sound velocity, leading to small scale instabilities at the perturbative level. Rigourously, the general situation is more complex: such instabilities for fluids with negative pressure may disappear if the hydrodynamical approach is replaced by a more fundamental description using, e.g., scalar fields. However, this is not true for the Chaplygin gas: even in a fundamental approach, using for example the Born-Infeld action, the sound speed may be negative if $\\alpha<0$. Perhaps the restriction $\\alpha\\geq0$ must be imposed for all observational tests.  We work also with a set of four free parameters. Varying all four parameters, the preceding results are confirmed. Leaving the parameter $H_{0}$ free to vary, we confirm that the hypersurface $H_{0}=72~km~s^{-1} ~Mpc^{-1}$ doesn't represent the maximum probability in the 4-D parameters phase space. Chaplygin gas models show values lower than $H_{0}=72~km~s^{-1} Mpc^{-1}$ \\cite{fabris03}. We observe also that there is a significant difference in the final parameter estimation when we consider the prior $\\alpha\\geq0$, instead of $0\\leq\\alpha<1$. For instance, the unification scenario ($\\Omega_{dm}= 0$) is favoured only with the choice $\\alpha\\geq0$. Moreover, there is now a peak in the $\\alpha$ distribution at $\\alpha=-4.3^{+4.8}_{-15.2}$ but with a high dispersion. Again, negative values for $\\alpha$ are favored, despite the two-dimensional PDF ($\\alpha$ x $H_{0}$) in Fig.\\ref{GRB3} indicates a high probability for $\\alpha>6$. Such a contradiction seems to be an artifact of the marginalization process as can be seen also in the two-dimensional PDFs ($\\Omega_{dm}$ x $H_{0}$) and ($\\bar{A}$ x $H_{0}$) in Fig. \\ref{GRB3}. These plots confirm that the final 1D estimation can be very different from the partial 2D ones. This diferrence is due to the integration of the probability function over the adopted prior values of the remmaning parameters.  The analysis with five free parameters confirm some of the previous results. Negative curvature is prefered as well as Sn data data \\cite{fabris03}. Also, the parameter $\\alpha$ is now estimated with a positive value, in constrast with the previous results.  The Chaplygin gas parameters have been estimated in many papers, considering different analysis and several observational data sets. Constraints critically depend on whether one treats the Chaplygin gas as true quartessence (replacing both dark matter and dark energy) or if one allows it to coexist with a normal dark matter component. The former situation is widely considered in the literature. As we leave the density parameter $\\Omega_{dm}$ free to vary in all the cases analysed here, it is not possible to directly compare our results with unified Chaplygin cosmologies  unless we assume the prior $\\Omega_{dm}=0$. This case has been studied using GRBs and other probes in reference \\cite{liang11}. For a comparison with this reference, figure \\ref{UnifGRB} shows the two-dimmensional probability for the free parameter of the unified ($\\Omega_{dm}=0$) GCG model. The best fit occurs at ($\\alpha=0.15, \\bar{A}=0.75$). This result agrees (at 1$\\sigma$) with the joint analysis showed in \\cite{liang11}.  \\begin{figure}[!h] \\begin{center} \\includegraphics[width=0.32\\textwidth]{unificationGCG} \\caption{Constrainst on the free parameters of the unified ($\\Omega_{dm}=0$) GCG model.} \\label{UnifGRB} \\end{center} \\end{figure}  The analysis of section 3 can be compared with \\cite{fabris03}, where the influence of a free $\\Omega_{dm}$ parameter on the final estimations was taken into account. Our results have high confidence with the results obtained in \\cite{fabris03}.  Finally, we remark that, perturbative analysis of Chaplygin models, for instance, reveals a large positive value ($\\alpha>>200$) for the parameter $\\alpha$ \\cite{fabris04} while kinematic tests show values negatives or close to zero. At the background level, the crossing of different data sets (including for example SNe, CMB, BAO, H(z) data and galaxy cluster mass fraction) will provide a more accurate scenario for each Chaplygin-based cosmology studied in this work. We leave this analysis, including the perturbative study, for a future work."
    },
    "1004/1004.0643_arXiv.txt": {
        "abstract": "We present deep high angular resolution observations  of the high-mass protostar NGC\\,7538\\,S, which is in the center of a cold dense cloud core with a radius of 0.5 pc and a mass of $\\sim$ 2,000 \\Msun. These observations show  that NGC\\,7538\\,S is embedded in a compact elliptical core with a mass of 85 - 115 \\Msun. The star is surrounded by a rotating accretion disk, which powers a very young, hot molecular outflow approximately perpendicular to the rotating accretion disk. The  accretion rate is very high, $\\sim$ 1.4 --  2.8 10$^{-3}$ M$_{\\odot}$~yr$^{-1}$. Evidence for rotation of the disk surrounding the star is seen in all  largely optically thin molecular tracers, H$^{13}$CN J = $1 \\to 0$, HN$^{13}$C J = $1 \\to 0$, H$^{13}$CO$^+$ J = $1 \\to 0$, and DCN J = $3 \\to 2$.  Many molecules appear to be affected by the hot molecular outflow, including DCN and H$^{13}$CO$^+$. The emission from CH$_3$CN,  which has often been used to trace disk rotation in young high-mass stars, is dominated by the outflow, especially at higher K-levels. Our new high-angular resolution observations show that the rotationally supported part of the disk is smaller than we previously estimated. The enclosed mass of the inner, rotationally supported part of the disk (D $\\sim$ 5\\arcsec{}, i.e 14,000 AU) is $\\sim$  14 - 24\\Msun. ",
        "introduction": "\\label{Intro}  How high-mass stars are formed is still under debate. High-mass stars are believed to form the same way as low mass stars, i.e. with a rotating accretion disk and driving an outflow, but in denser environments and with much higher accretion rates \\citep{Wolfire87, Stahler00,Norberg00,McKee02,Keto03,Keto07}. However, others argue that high-mass stars are most likely formed by competitive accretion in a clustered environment \\citep{Bonnell06}, or by adiabatic accretion of gas from the surrounding cluster to densities such that stellar collisions are likely to occur \\citep{Clarke08}.  Accretion disks are ubiquitous in young low- and intermediate-mass stars, i.e. T Tauri and Herbig Ae stars \\citep{Simon00,Mannings97,Najita03}, which have sizes of one to a few hundred AU and appear to be in Keplerian rotation. Disks are also common in Class I objects, i.e. younger, more heavily embedded low- and intermediate mass pre-main-sequence stars \\citep{Brown99}, but there are very few examples of confirmed disks for the very earliest stages of  a low-mass protostar, the Class 0 phase \\citep{Jorgensen04,Chandler05}. At the Class 0 phase, the temperature of the protostellar disk is of the same order of that of the collapsing cloud core, and it is therefore an observational challenge  to get enough contrast between the disk and the cloud core.  The situation is even worse for high-mass stars, because they form in much denser environments and always in clusters or small groups. The survival time of disks in high-mass stars is furthermore expected to be much shorter, because once the star is formed it will quickly photo evaporate the disk as soon as the accretion from the surrounding cloud starts to diminish, therefore allowing the uv-photons to reach the disk surface. Yet there are several detections of disks around early B-stars; Lkh$\\alpha$\\,101 (Spectral class B0) \\citep{Tuthill02}, MWC297 (B1.5) \\citep{Manoj07}, MWC\\,349\\,A (B [e]) \\citep{Weintroub08}, clearly confirming that disks do exist around relatively massive stars.  Most young high-mass stars appear to drive outflows \\citep{Shepherd96,Zhang01,Zhang05,Beuther02b}, but whether they also have accretion disks is less clear.  There have been many reports of disks around high-mass stars, but only a few which appear well supported. Two of the most clearcut cases are IRAS $20126+4104$\\citep{Cesaroni97,Cesaroni99, Cesaroni05,Zhang98}, and AFGL\\,490 \\citep{Schreyer06}. IRAS $20126+4104$ has a luminosity of $\\sim$ 1.3 10$^4$ \\Lsun\\ suggesting an early B star. The star drives a parsec scale CO  outflow \\citep{Wilking90,Shepherd00}, and is surrounded by a rotating accretion disk with a diameter of $\\sim$ 15,000 A.U. \\citep{Cesaroni05}. The observed rotation curve for an edge-on Keplerian disk corresponds to $\\sim$ 7 \\Msun. The luminosity of AFGL\\,490, L $\\sim$  2 10$^3$ \\Lsun\\ also suggests an early B-star. It drives a  more compact outflow, 0.3~pc (corrected for inclination) \\citep{Lada81,Snell84} and is surrounded by a nearly face-on disk with a radius  $\\sim$ 1600 AU, which has  a disk mass of 1 \\Msun\\  and appears to be in Keplerian rotation \\citep{Schreyer06}. Another, even younger high-mass protostar, NGC\\,7538\\,S, with a luminosity of $\\sim$ 1.5~10$^4$ \\Lsun,  which  is surrounded by a large, nearly edge-on disk with a radius of 15,000 AU and a disk mass of $\\sim$ 100 \\Msun\\ was reported by \\citet{Sandell03}. The disk drives a compact molecular outflow with a size of 0.1 - 0.4 pc and shows a Keplerian-like rotation.  This disk is the target for our studies. In this paper we present follow-up BIMA spectral line observations of several molecules in the 3 and 1 mm bands, which were chosen to measure the rotation of the disk surrounding NGC\\,7538\\,S. In a  companion paper (Corder et al., in preparation; hereafter Paper II) we discuss high-angular resolution continuum observations of NGC\\,7538\\,S and the Ultra-Compact \\ion{H}{2} regions NGC\\,7538 IRS\\,1 - 3 with BIMA, CARMA, and the VLA, which provide us with better mass estimates of the disk/envelope surrounding the star. ",
        "conclusions": "We have carried out extensive observations of the star forming core NGC\\,7538\\,S with BIMA in mostly optically thin tracers with  spatial resolutions ranging from $\\sim$ 3\\arcsec\\ to 8\\arcsec. Additionally we have acquired complementary observations with FCRAO and JCMT to fill in lacking short spacing in our BIMA observations, which are needed to improve the image fidelity and  reliability of our images of molecules like HCO$^+$, H$_2$CO, and H$^{13}$CN, which are spatially very extended.  We confirm that there is a very young high-mass (proto)star in the compact elliptical core, which has a size of 8\\arcsec\\ $\\times$ 3\\arcsec\\ and a mass of  85 - 115 \\Msun. Recent sub-arcsecond continuum imaging at 110 and 224 GHz with CARMA (Paper II) resolve the elliptical core into three compact sources, all of which are almost certainly protostars.  The strongest one of the three sources agrees within 0\\farcs15 with the adopted position for the high-mass protostar, which is $\\sim$ 2\\arcsec\\ to the northeast from the center of the compact elliptical core. The protostar is seen as a faint, extremely obscured mid-IR source, which coincides with a VLA thermal jet, and an OH and a CH$_3$OH class II methanol maser. The star is surrounded by a massive, rotating accretion disk, which drives a highly collimated thermal jet and powers a very compact,  hot molecular outflow. We see clear evidence for accretion towards the protostar. Almost all molecular transitions that we have observed, show red-shifted self-absorption, which can be explained only by infall motions of the gas in front of the protostar. The accretion signature is very strong in the optically thick HCO$^+$ \\jtra10 transition, which shows infall velocities up to $\\sim$ 15 km~s$^{-1}$. We have used this accretion profile to derive a direct estimate of the accretion rate towards the disk, and find an accretion rate $\\sim$ 10$^{-3}$ M$_{\\odot}$~yr$^{-1}$. We also estimated the accretion rate by assuming that the accretion rate is about three times the outflow rate, and by assuming a steady state infall to the disk, with an  infall velocity equal to the observed rotation velocity of the disk at the radius R.  All three methods give very similar results. Such high accretion rates are sufficient to quench the formation of an \\ion{H}{2} region and allow the central protostar to continue to grow in mass \\citep{Walmsley95,Keto03,Keto07}.  The rotating accretion disk is best seen in H$^{13}$CN \\jtra10. H$^{13}$CN is only marginally affected by emission from the intense hot outflow, while the emission from the outflow may dominate or severely hinder us from seeing the disk in many of the molecular transitions that we have observed. We found that DCN  \\jtra32, which we expected to be an equally good disk traces as  H$^{13}$CN, is strongly affected by the outflow. Due to the high angular resolution, $\\sim$ 2\\ptsec6 (Table \\ref{tbl-2}), we can separate most of the outflow emission from the disk-emission and therefore confirm that NGC\\,7538\\,S is surrounded by a rotating accretion disk. The emission from methylcyanide, CH$_3$CN, however, is largely optically thick and the higher K-levels  are completely dominated by the hot outflow emission near the surface of the disk.   BIMA observations of  H$^{13}$CN \\jtra10 supplemented with FCRAO data show that the cloud core, in which NGC\\,7538\\,S is embedded, has a radius of $\\sim$ 0.5 pc, and a mass $\\sim$ 2000 \\Msun. The size and mass of the core agrees very well with what we derive from analysis of the SCUBA 850 and 450 $\\mu$m data presented by \\citet{Sandell04}, while the mass estimate from a single dish  C$^{18}$O \\jtra21 map is much lower, suggesting that CO is depleted (frozen onto grains) in the cold cloud envelope. The cloud core is  dense and massive enough to provide the necessary conditions for high-mass star formation."
    },
    "1004/1004.2536_arXiv.txt": {
        "abstract": "{Observations of polarized emission are a significant source of information on the magnetic field that pervades the Interstellar Medium of the Galaxy. Despite the acknowledged importance of magnetic field in interstellar processes, our knowledge of field configurations on all scales is seriously limited.} {This paper describes an extensive survey of polarized Galactic emission at 1.4~GHz that provides data with arcminute resolution and complete coverage of all structures from the broadest angular scales to the resolution limit, giving information on the magneto-ionic medium over a wide range of interstellar environments. } {Data from the DRAO Synthesis Telescope, the Effelsberg 100-m Telescope, and the DRAO 26-m Telescope have been combined. Angular resolution is ${\\sim}1'$ and the survey extends from ${\\ell}={66^{\\circ}}$ to ${\\ell}={175^{\\circ}}$ over a range ${-3^{\\circ}} < b < {5^{\\circ}}$ along the northern Galactic plane, with a high-latitude extension from ${\\ell}={101^{\\circ}}$ to ${\\ell}={116^{\\circ}}$ up to ${b}={17\\fdg5}$. This is the first extensive polarization survey to present aperture-synthesis data combined with data from single antennas, and the techniques developed to achieve this combination are described.} {The appearance of the extended polarized emission at 1.4~GHz is dominated by Faraday rotation along the propagation path, and the diffuse polarized sky bears little resemblance to the total-intensity sky.  There is extensive depolarization, arising from vector averaging on long lines of sight, from \\ion{H}{ii} regions, and from diffuse ionized gas seen in H$\\alpha$ images. Preliminary interpretation is presented of selected polarization features on scales from parsecs (the planetary nebula Sh 2-216) to hundreds of parsecs (a superbubble GSH~166$-$01$-$17), to kiloparsecs (polarized emission in the direction of Cygnus~X).} {} ",
        "introduction": "\\label{sec:intro}  The detection of linear polarization in the Galactic radio emission \\citep{west62,wiel62} provided crucial evidence in establishing the synchrotron mechanism as the source of the emission. Assumption of equipartition between relativistic particles and magnetic field then led to estimates of the field strength of a few $\\mu$G \\citep{beck01}, and the best estimates today do not differ substantially \\citep{sun08}. The magnetic field in the Galaxy is a significant reservoir of energy, and the field is likely to play an important role in interstellar processes. Nevertheless, more than four decades after the first detection our knowledge of the field configuration on global and local scales is still quite limited.  The apparent promise of polarization observations has proved difficult to translate into hard information on field configurations. First, the observed fractional polarization is generally far below the 70\\% theoretical maximum.  The radiation is optically thin and superposition of emission contributions along the line of sight can ``depolarize'' the signal through vector averaging. Second, Faraday rotation operates whenever the signal propagates through a magnetized thermal plasma.  Emission and Faraday rotation often occur in the same region, and a wide variety of depolarization effects occur \\citep{burn66,soko98}.  Furthermore, a typical radio telescope operating at low radio frequencies proves to be more sensitive to the Faraday rotation of a plasma region than to its bremsstrahlung (as discussed in Section~\\ref{subsec:small-scale}). The net result of all these effects is that the polarized sky rarely resembles the total-intensity sky and interpretation of polarization images is seldom simple.  Angular resolution of early polarization surveys with single antennas was poor.  A comprehensive set of surveys of the northern sky was published by \\citet{brou76} based on well calibrated (but undersampled) observations at four frequencies between 408 MHz and 1411 MHz made in the 1960s with the Dwingeloo 25-m Telescope; angular resolution ranged from 2$^{\\circ}$ to 36\\arcmin. Here the subject rested for many years until it was revived using the Effelsberg 100-m Telescope in the late 1980s. Extensive surveys made with that telescope had resolutions of 4\\farcm3 at 2695~MHz \\citep{junk87,dunc99} and 9\\farcm4 at 1410~MHz \\citep{uyan98,uyan99,reic04}. A wide-area survey of the Southern Galactic plane was made with the Parkes Telescope at 2417~MHz with resolution 10\\farcm4\\, \\citep{dunc97}. The Northern Galactic plane was surveyed with the Urumqi Telescope with resolution 9\\farcm5 at 4900~MHz \\citep{sun07,gao10}. The WMAP data at 23~GHz and higher frequencies \\citep{hins09} cover the entire sky at an angular resolution smaller than $1^{\\circ}$ and, on simple lines of sight, show almost the intrinsic polarization characteristics because Faraday rotation at these frequencies is extremely low.  A substantial increase in angular resolution has been provided by aperture-synthesis telescopes, and observations of the Galactic polarized emission with angular resolutions from one to a few arcminutes have been made in recent years \\citep{wier93,gray98,gray99,have00,gaen01,uyan02,uyan03,have03a,have03b, have03c,have06a}.  These observations have revealed much about Faraday rotation in the ISM, but interpretation has been hampered by the lack of information on the biggest structures{\\footnote{ Single-antenna data may also suffer from this problem -- see Section~\\ref{subsec:eff_process} and \\citet{reic06}}}.  An interferometer observes the sky through a spatial high-pass filter, so the zero levels for $Q$ and $U$ are lost. Both polarized intensity, ${\\rm{PI}}={\\sqrt{Q^{2}+U^{2}}}$, and polarization angle, ${\\rm{PA}}={{0.5}\\,{\\rm{tan}}^{-1}\\left(\\frac{U}{Q}\\right)}$, change in a very non-linear fashion in response to errors in zero levels. The most serious effect is on angle. Consider an observation of a region several tens of beamwidths in extent comprising broad structure, whose polarization angle, $A$, varies only slowly across the area, and fine structure, contributing rapid changes in PA that add vectorially to $A$. An interferometer can measure only the rapid angle changes, and, if the broad structure is not measured, the apparent PA will vary over a large range, whereas the true distribution of angle is much narrower and is centred on $A$. Examples of this effect are illustrated in \\citet{reic06}, \\citet{sun07}, and \\citet{gao10}. An example in the present data is discussed in Section~\\ref{subsec:w543}. If broad structure is missing then conclusions about PA, and hence rotation measure (RM), are prone to serious error. In contrast, the work described here is the first extensive polarization survey to incorporate single-antenna data with aperture-synthesis data.  In this paper we present a survey of the polarized emission at 1420~MHz covering 1060 square degrees of the northern Galactic plane with an angular resolution of $\\sim$1~arcminute.  With $1.5 \\times 10^7$ independent data points this is the largest polarization survey published to date. The survey combines data from the DRAO Synthesis Telescope, the Effelsberg 100-m Telescope, and the DRAO 26-m Telescope, and we discuss the techniques that we have developed to combine these datasets. Data from two single-antenna telescopes, not just one, were needed to correctly represent broad structure for reasons discussed below (Section~\\ref{subsec:eff_process}).   The new polarization dataset forms part of the Canadian Galactic Plane Survey (the CGPS, described by Taylor et al. 2003). The scientific goal of the CGPS is the study of the Galactic ``ecosystem'', the interplay between the various constituents of the ISM, their role in star formation, the impact of stars on their environments, and the interaction of Galaxy-wide phenomena such as density waves with ISM constituents. The CGPS comprises surveys{\\footnote{The CGPS database is accessible to the astronomy community at {\\tt{http://www3.cadc-ccda.hia-iha.nrc-cnrc.gc.ca/cgps/}}}} of the atomic hydrogen (the 21-cm \\ion{H}{i} line), the ionized gas (seen in radio continuum), the molecular gas (the lines of CO near 115~GHz - \\citep{heye98}), the dust (IRAS data reprocessed with high resolution - \\citep{cao97,kert00}), and, relevant here, the relativistic component and the magnetic field traced by continuum observations at 408 MHz (total intensity only) and 1420 MHz (Stokes parameters $I$, $Q$, and $U$).  The survey described here maps the magneto-ionic component of the ISM with unprecedented detail and precision. ",
        "conclusions": "\\label{sec:conclus}  We have described techniques developed to make a survey of the polarized radio emission from the Galactic plane over a large area, combining data from aperture-synthesis and single-antenna telescopes to provide an accurate portrayal of emission features on all angular scales to the resolution limit of ${\\sim}1'$.  This survey represents a major advance in high-fidelity imaging of the polarized sky.  We have presented data from the survey and have made preliminary interpretations of some features revealed by it.  Mapping the polarized sky opens a new ``window'' on the ISM because the appearance of the sky is dominated by Faraday rotation occurring along the propagation path through the Galaxy, to the point where the polarized sky does not resemble the total-intensity sky.  Some general conclusions can be reached from this work that will be relevant to future polarization imaging.  \\begin{itemize}  \\item {While Faraday rotation tends to break up large emission structures into smaller ones, there is still significant large-scale structure at 1.4~GHz. Consequently, interpretation of aperture-synthesis data will be severely limited unless single-antenna data are accurately incorporated into the images.}  \\item{Faraday rotation is a powerful tool for detecting ionized gas, and polarization observations will lead to the discovery of objects that cannot easily be detected by other means.}  \\item{There are features of the Galactic emission many degrees in extent; despite their large size they are very difficult to recognize without arcminute angular resolution.}  \\item{Polarization features trace structures in the magneto-ionic medium, in electron density or magnetic field or both, and observations of the diffuse polarized emission reveal a diversity of phenomena associated with this component of the ISM. The magneto-ionic medium is in part unstructured and very broadly distributed but it may also be associated with discrete objects such as SNRs, \\ion{H}{ii} regions, planetary nebulae, and stellar-wind bubbles.}  \\item{Some of the features seen in polarization images are the products of propagation and depolarization effects of various kinds, and are not necessarily ``objects'' in the usual astronomical sense.}  \\end{itemize}  This survey has revealed a wealth of structure in the magneto-ionic medium on all scales. Future work on the data presented here will be directed at illuminating the relationship between these structures and other phases of the ISM. Information on other ISM tracers, particularly the CGPS datasets describing the atomic, ionized, and molecular gas and the dust, will be critical to the success of these studies.  The longitude range of the survey has been extended to ${\\ell}{\\approx}{55^{\\circ}}$ towards the inner Galaxy and to ${\\ell}{\\approx}{195^{\\circ}}$ beyond the anticentre.  These new data will be processed using the techniques described in this paper."
    },
    "1004/1004.0705_arXiv.txt": {
        "abstract": "We report results from a survey of Mg\\,II absorbers in the spectra of background QSOs that are within close angular distances to a foreground galaxy at $z<0.5$, using the Magellan Echellette Spectrograph.  We have established a spectroscopic sample of 94 galaxies at a median redshift of $\\langle z\\rangle = 0.24$ in fields around 70 distant background QSOs ($z_{\\rm QSO}>0.6$), 71 of which are in an 'isolated' environment with no known companions and located at $\\rho\\apll 120\\ h^{-1}$ kpc from the line of sight of a background QSO.  The rest-frame absolute $B$-band magnitudes span a range from $M_{B}-5\\log\\,h=-16.4$ to $M_{B}-5\\log\\,h=-21.4$ and rest-frame $B_{AB}-R_{AB}$ colors range from $B_{AB}-R_{AB}\\approx 0$ to $B_{AB}-R_{AB}\\approx 1.5$.  Of these 'isolated' galaxies, we find that 47 have corresponding Mg\\,II absorbers in the spectra of background QSOs and rest-frame absorption equivalent width $W_r(2796)=0.1-2.34$ \\AA, and 24 do not give rise to Mg\\,II absorption to sensitive upper limits.  Our analysis shows that (1) \\ewr\\ declines with increasing distance from 'isolated' galaxies but shows no clear trend in 'group' environments; (2) more luminous galaxies possess more extended Mg\\,II absorbing halos with the gaseous radius scaled by $B$-band luminosity according to $R_{\\rm gas}=75\\times (L_B/L_{B_*})^{(0.35\\pm 0.03)}\\ h^{-1}$ kpc; (3) there is little dependence between the observed absorber strength and galaxy intrinsic colors; and (4) within $R_{\\rm gas}$, we find a mean covering fraction of $\\langle\\kappa_{0.3}\\rangle\\approx 70$\\% for absorbers of $\\ewr\\ge 0.3$ \\AA\\ and $\\langle\\kappa_{0.1}\\rangle\\approx 80$\\% for absorbers of $\\ewr\\ge 0.1$ \\AA.  The results confirm that extended Mg\\,II absorbing halos are a common and generic feature around ordinary galaxies and that the gaseous radius is a fixed fraction of the dark matter halo radius.  The lack of correlation between \\ewr\\ strength and galaxy colors suggests a lack of physical connection between the origin of extended Mg\\,II halos and recent star formation history of the galaxies.  Finally, we discuss the total gas mass in galactic halos as traced by Mg\\,II absorbers.  We also compare our results with previous studies. ",
        "introduction": "The 'forest' of absorption-line systems observed in the spectra of background quasars offers a sensitive probe of otherwise invisible gaseous structures in the universe, where the majority of baryons reside.  Depending on the physical conditions of the absorbing gas, different transitions are expected to display different strengths and kinematic signatures.  Combining QSO absorption-line observations and faint galaxy surveys, in principle, provides a unique and powerful means of establishing a comprehensive picture for understanding the growth of galaxies.  The Mg\\,II $\\lambda\\lambda\\,2796, 2803$ doublets are among the absorption features commonly seen in the spectra of distant quasars. These absorbers are understood to originate in photo-ionized gas of temperature $T \\sim 10^4$ K (Bergeron \\& Stas\\'inska 1986; Charlton et al.\\ 2003) and trace high-column density clouds of neutral hydrogen column density $N(\\hI) \\approx 10^{18}-10^{22}$ \\cmjj\\ (Rao et al.\\ 2006).  At redshifts $z=0.4-2.5$, these doublet features are redshifted into the optical spectral range and are routinely detected in the spectra of background QSOs using optical spectrographs on the ground.  Over the past two decades, a large number of studies have been carried out to characterize the statistical properties of Mg\\,II absorbers, including the frequency distribution function, redshift evolution of the absorber number density, and kinematic signatures (e.g.\\ Lanzetta \\etal\\ 1987; Sargent \\etal\\ 1988; Petitjean \\& Bergeron 1990; Steidel \\& Sargent 1992; Charlton \\& Churchill 1998; Churchill \\etal\\ 2000, 2003; Nestor \\etal\\ 2005; Prochter \\etal\\ 2006).  While the accuracy and precision of these various measurements increase with increasing sample size, the utility of known absorber statistics in advancing our understanding of galaxy evolution has been limited due to an ambiguous origin of these absorbers.  A necessary first step toward the goal of applying QSO absorption-line systems for probing the growth of baryonic structures is to understand and quantify their correlation with galaxies.  The large associated \\hI\\ column density (Rao et al.\\ 2006) suggests that Mg\\,II absorbers are similar to those H\\,I clouds seen around individual galaxies in 21~cm surveys (e.g.\\ Doyle et al.\\ 2005).  A direct connection of these absorbers to galaxies is also supported by the presence of luminous galaxies at projected distances $\\rho = 50-100 \\ h^{-1}$ kpc from known Mg\\,II absorbers (Bergeron 1986; Lanzetta \\& Bowen 1990, 1992; Steidel \\etal\\ 1994; Zibetti \\etal\\ 2005; Nestor \\etal\\ 2007; Kacprzak \\etal\\ 2007).  However, uncertainties remain, because some galaxies found at $\\rho<50\\ h^{-1}$ kpc from a QSO sightline do not produce a corresponding Mg\\,II absorber (e.g.\\ Tripp \\& Bowen 2005; Churchill \\etal\\ 2007), implying that Mg\\,II absorbers may not probe a representative population of galaxies or simply that the gas covering fraction is not unity.  Over the past two years, we have been conducting a program that combines empirical observations and a phenomenological model study to establish a comprehensive description of the correlation between galaxies and cool gas ($T\\sim 10^4$ K) probed by Mg\\,II absorbers.  In Tinker \\& Chen (2008), we have introduced a novel technique that adopts the halo occupation framework (e.g.\\ Seljak 2000; Scoccimarro et al.\\ 2001; Berlind \\& Weinberg 2002) and characterizes the origin of QSO absorption-line systems based on a conditional mass function of dark matter halos in which the absorbers are found. This technique is purely statistical in nature.  It characterizes the cold and warm-hot gas in dark matter halos by comparing the frequency distribution function and clustering amplitude of QSO absorbers with those of dark matter halos.  Our initial halo occupation model is constrained by known statistical properties of Mg\\,II absorbers at $\\langle z\\rangle=0.6$, including an isothermal density profile for describing the spatial distribution of cold gas in individual dark matter halos, the frequency distribution function (e.g.\\ Steidel \\& Sargent 1992; Nestor \\etal\\ 2005; Prochter \\etal\\ 2006), and the clustering amplitude (Bouch\\'e \\etal\\ 2006; Gauthier \\etal\\ 2009; Lundgren \\etal\\ 2009).  The adopted isothermal model is supported by empirical data (Chen \\& Tinker 2008).  The product of the halo occupation analysis is an occupation function (or ``mass function'') that characterizes the fractional contribution of dark matter halos (galaxies) of different masses to the observed Mg\\,II absorbers of different strength.  Our analysis has shown that in order to reproduce the observed overall strong clustering of the absorbers, roughly $5-10$ \\% of the gas in halos up to $10^{14}$ \\hmsol\\ is required to be cold.  The inferred presence of cool gas in massive halos is also supported by (i) the observed strong cross-correlation amplitude of Mg\\,II absorbers and luminous red galaxies (LRG) on projected distance scales of $r_p < 300\\ h^{-1}$ comoving kpc (Gauthier \\etal\\ 2009) and by (ii) direct detections of Mg\\,II absorbers at $\\apll 300\\ h^{-1}$ kpc and $\\apll 320$ \\kms\\ from five LRGs (Gauthier \\etal\\ 2010).  These LRGs are understood to reside in $>10^{13}$ \\hmsol\\ halos (e.g.\\ Blake \\etal\\ 2008; Gauthier \\etal\\ 2009).  For lower-mass halos, our halo occupation analysis has shown that the incidence and covering fraction of extended cool gas is high.  Therefore these halos contribute significantly to the observed Mg\\,II statistics.  The large gas covering fraction is consistent with the empirical findings of Steidel \\etal\\ (1994) and Chen \\& Tinker (2008), though other authors have reported a lower covering fraction from different surveys (e.g.\\ Tripp \\& Bowen 2005; Kacprzak \\etal\\ 2008; Barton \\& Cooke 2009).  In summary, the initial results of our halo occupation analysis demonstrate that combining galaxy and absorber survey data together with a simple semi-analytic model already produces unique empirical constraints for contemporary theoretical models that study the gas content of dark matter halos (e.g.\\ Mo \\& Miralda-Escud\\'e 1996; Birnboim \\& Dekel 2003; Maller \\& Bullock 2004; Kere$\\check{\\rm s}$ \\etal\\ 2005; 2009; Dekel \\& Birnboim 2006; Birnboim \\etal\\ 2007).  In Tinker \\& Chen (2010), we have expanded upon our initial halo occupation analysis to incorporate the observed number density evolution of the absorbers (Nestor \\etal\\ 2005; Prochter \\etal\\ 2006) as an additional constraint to gain insight into the redshift evolution of extended gas around galaxies.  In order to incorporate the expected redshift evolution of the dark matter halo population to explain the observed number density evolution of Mg\\,II absorbers, we have found that the gaseous halos must evolve with respect to their host dark matter halos.  An explicit prediction of our halo occupation model is a more pronounced inverse-correlation between the mean halo mass and absorber strength (e.g.\\ Bouch\\'e \\etal\\ 2006; Gauthier \\etal\\ 2009; Lundgren \\etal\\ 2009) at $z\\apll 0.3$ and a positive correlation between the mean halo mass and absorber strength at $z\\apg 2$.  However, no clustering measurements are available for Mg\\,II absorbers at $z<0.4$ or $z\\apg 1$.  To test this prediction would require a large sample of galaxies and absorbers from these epochs.  At $z\\apg 2$, a large sample of Mg\\,II absorbers is already available (e.g.\\ Nestor \\etal\\ 2005; Prochter \\etal\\ 2006) from searches in the Sloan Digital Sky Survey (SDSS; York \\etal\\ 2000) quasar sample (e.g.\\ Schneider \\etal\\ 2007), but galaxy surveys that cover a cosmological volume are challenging in this redshift range.  In contrast, few Mg\\,II absorbers are known at $z<0.35$ where exquisite details of the galaxy population have been recorded (e.g.\\ the SDSS), due to a lack of spectral sensitivity at wavelength $\\lambda<4000$ \\AA\\ where the low-redshift Mg\\,II absorption features occur.  A new observing window has become available with recently commissioned UV sensitive spectrographs on the ground.  In particular, the Magellan Echellette Spectrograph (MagE; Marshall \\etal\\ 2008) offers high throughput over a contiguous spectral range from $\\lambda=3100$ \\AA\\ to 1 $\\mu$m, allowing searches for Mg\\,II absorbers at redshift as low as $z=0.11$ (see also Barton \\& Cooke 2009 for a similar effort at $z=0.1$).  Building upon the existing SDSS galaxy database, we have initiated a MagE survey of Mg\\,II absorbers in the spectra of background QSOs, whose sightlines intercept the halo of a foreground galaxy at $z\\le 0.5$.  The primary goal of this project is to establish a statistically significant sample ($N\\apg 500$) of low-redshift Mg\\,II absorbers for measuring the clustering amplitude of Mg\\,II absorbers at $\\langle z\\rangle=0.2$ based on the observed cross-correlation amplitude on co-moving scales of $1-30\\ h^{-1}$ Mpc. Combining the detections and non-detections in the vicinity (impact separation of $\\rho\\apll 100\\ h^{-1}$ physical kpc) of known galaxies will also facilitate a comprehensive study of how the properties of extended cool gas (such as the density profile and covering fraction) correlate with known stellar and ISM properties of the host galaxies. Here we introduce the MagE Mg\\,II absorber survey and present initial results from the first year of data.  This paper is organized as follows.  In Section 2, we describe the design of our MagE survey project.  In Section 3, we describe the spectroscopic observations of photometrically selected galaxies in the SDSS data archive and the MagE follow-up of quasars selected from the SDSS spectroscopic QSO catalog.  We provide a summary of the data reduction and analysis procedures.  In Section 4, we present the catalogs of galaxies and Mg\\,II absorbers.  In Section 5, we examine the correlation between Mg\\,II absorption strength and galaxy properties.  Finally, we discuss the properties of extended cool gas in galactic halos and compare our results with previous studies in Section 6.  We adopt a $\\Lambda$CDM cosmology, $\\Omega_{\\rm M}=0.3$ and $\\Omega_\\Lambda = 0.7$, with a dimensionless Hubble constant $h = H_0/(100 \\ {\\rm km} \\ {\\rm s}^{-1}\\ {\\rm Mpc}^{-1})$ throughout the paper. ",
        "conclusions": "We have established a spectroscopic sample of 94 galaxies at a median redshift of $\\langle z\\rangle = 0.2357$ in fields around 70 distant background QSOs ($z_{\\rm QSO}>0.6$).  Our follow-up MagE spectra of the QSOs allow us to examine the extent of cold gas around 88 of these galaxies based on the presence/absence of coincident Mg\\,II absorption features.  In the sample of 88 spectroscopically confirmed galaxies, 17 are found in a 'group' environment with two or more close neighbors at $\\rho<{\\hat R}_{\\rm gas}$ (Equation 1) and velocity separation $\\Delta\\,v<300$ \\kms.  Excluding one galaxy that occurs in the vicinity of the background QSO, we identify seven galaxy 'groups' along the lines of sight of the background QSOs in our sample.  All seven galaxy 'groups' have coincident Mg\\,II absorbers.  Because the association between absorbers and individual galaxies becomes uncertain in a group environment, we have separated these group galaxies from the remaining 'isolated' ones.  Figure 12 shows that while galaxies from a 'group' environment appear to occupy a similar range in the $W_r(2796)$ versus $\\rho$ parameter space with 'isolated' galaxies, no strong inverse correlation is seen in the gaseous profiles of 'group' galaxies.  Because interactions between group members are expected to alter the properties of gaseous halos, such as ram pressure and tidal stripping that could re-distribute cold gas to larger radii (e.g.\\ Gunn \\& Gott 1972; Balogh \\etal\\ 2000; Verdes-Montenegro \\etal\\ 2001), the 'group'-galaxy subsample also presents a unique opportunity to study gas kinematics in overdense galaxy environments beyond the local universe (c.f.\\ Verdes-Montenegro \\etal\\ 2001).  The remaining 71 'isolated' galaxies are located at $\\rho\\apll 120\\ h^{-1}$ kpc from the line of sight of a background QSO.  We identify 47 coincident Mg\\,II absorbers in the QSO spectra with absorber strengths varying from $W_r(2796)=0.1$ \\AA\\ to $W_r(2796)=2.34$ \\AA, and measure a sensitive upper limit of the Mg\\,II absorber strength for the remaining 24 galaxies.  In the absence of a complete spectroscopic survey of galaxies around the 71 'isolated' galaxies, it is likely that some fraction of these galaxies also occur in a group environment.  However, the strong $\\ewr$ versus $\\rho$ anti-correlation after accounting for the luminosity scaling relation seen for the 'isolated' galaxies in Figures 9 \\& 12 indicates that with the exception of a few outliers the majority of the 'isolated' galaxies are indeed different from those 'group' galaxies.  We therefore argue that the majority of the 'isolated' galaxies are in a more quiescent environment than the 'group' ones in our sample.  The results of our likelihood analysis demonstrates that the Mg\\,II absorber strength $W_r(2796)$ scales inversely with galaxy impact parameter.  In addition, the $W_r(2796)$ vs.\\ $\\rho$ anti-correlation is still stronger after accounting for the scaling relation with galaxy $B$-band luminosity, indicating that more luminous galaxies are surrounded by more extended Mg\\,II halos.  However, including intrinsic $B_{AB}-R_{AB}$ color does not improve the observed $W_r(2796)$ vs.\\ $\\rho$ anti-correlation, indicating a lack of physical connection between the origin of extended Mg\\,II halos and recent star formation history of the galaxies (c.f.\\ Zibetti \\etal\\ 2007).  Finally, the $W_r(2796)$ vs.\\ $\\rho$ anti-correlation appears to depend only weakly on galaxy redshift, indicating little evolution of the extended Mg\\,II halos between $z\\approx 0.5$ and $z\\approx 0.1$.  Here we focus on the 'isolated' galaxy sample and discuss the covering fraction and spatial profile of extended Mg\\,II absorbing gas based on the observations and analysis presented in \\S\\ 5.  We also discuss implications for the cool baryon content of galactic halos and compare our results with previous studies.  \\subsection{Incidence and Covering Fraction of Mg\\,II Absorbers}  The analysis presented in \\S\\ 5.2 shows that the strengths of Mg\\,II absorbers depend on the projected distances and intrinsic luminosities of the absorbing galaxies.  In addition to the best-fit scaling relation, we also note that with the exception of two outliers all 33 galaxies at luminosity-scaled impact parameter $\\rho'\\equiv \\rho\\times\\,(L_B/L_{B_*})^{-0.35}<30\\ h^{-1}$ (Equation 11) have a coincident Mg\\,II absorber of $W_r(2796)>0.3$ \\AA.  The observed high incidence of strong Mg\\,II absorbers around galaxies with a broad range of intrinsic colors strongly supports the notion that extended Mg\\,II halos are a common and generic feature of galaxies of all types, from evolved early-type galaxies to late-type star-forming systems.  While extended Mg\\,II absorbing gas reaches out to larger projected distances, our sample shows that both the absorption strength and gas covering fraction declines toward larger radii. Here we examine how the gas covering fraction $\\kappa$ varies with $\\rho$ and $W_r(2796)$.  We perform a maximum-likelihood analysis to determine $\\kappa$.  The probability that a galaxy gives rise to an absorption system of some absorption equivalent width threshold is written as \\begin{equation} P(\\kappa)=\\kappa(\\rho,W_0) \\, B [r_1(L_B),r_2(L_B);\\rho], \\end{equation} where $\\kappa$ is the fraction of galaxies that give rise to Mg\\,II absorption, and $B$ is a boxcar function that defines the impact parameter interval; $B=1$ if $r_1(L_B)\\le \\rho < r_2(L_B)$ and $B=0$ otherwise.  Equation (14) takes into account the scaling relation between the gaseous extent of galaxies and galaxy $B$-band luminosity. The likelihood of detecting an ensemble of galaxies, $n$ of which give rise to Mg\\,II absorption systems and $m$ of which do not, is given by \\begin{eqnarray} {\\cal L}(\\kappa)&=&\\prod_{i=1}^n\\,\\kappa(\\rho_i,W_i) B[r_1(L_{B_i}),r_2(L_{B_i});\\rho_i] \\nonumber \\\\ & & \\times \\prod_{j=1}^m\\,\\{1-\\kappa(\\rho_j,W_j) B[r_1(L_{B_j}),r_2(L_{B_j});\\rho_j]\\} \\nonumber \\\\ &=& \\prod_{i=1}^n\\,\\langle\\kappa\\rangle\\,\\prod_{j=1}^m(1-\\langle\\kappa\\rangle) \\nonumber \\\\ &=& \\langle\\kappa\\rangle^n\\,(1-\\langle\\kappa\\rangle)^m. \\end{eqnarray}  We evaluate $\\kappa$ for three absorption equivalent width thresholds, $W_0=0.5$ \\AA, $W_0=0.3$ \\AA, and $W_0=0.1$ \\AA.  In the 'isolated' galaxy sample, all MagE spectra of the corresponding QSOs have sufficient $S/N$ for detecting an absorber of $W_r(2796)\\ge 0.3$ \\AA\\ and 61 have sufficient $S/N$ spectra for uncovering an absorber of $W_r(2796)\\ge 0.1$ \\AA.  We calculate $\\kappa$ for different impact parameter intervals, using the best-fit scaling relation of Equation (11) and excluding the outliers.  The results are shown in Figure 10. The errorbars indicate the 68\\% confidence interval of the estimated $\\kappa(\\rho,W_0)$.  The impact parameter intervals are chosen to contain ten galaxy--absorber pairs per bin, except for the last bin that includes only the remaining seven pairs at largest separations. Figure 10 shows that the covering fraction of $W_r(2796)\\ge 0.1$ \\AA\\ absorbers varies from 100\\% within $\\rho=30\\ h^{-1}$ kpc of an $L_*$ galaxy to $<20$\\% beyond $\\rho=60\\ h^{-1}$ kpc.  \\begin{figure} \\begin{center} \\includegraphics[scale=0.4]{kappa.eps} \\caption{Incidence and covering fraction of Mg\\,II absorbers versus galaxy impact parameter accounting for the galaxy $B$-band luminosity scaling relation (Equation 11).  Solid points indicate absorbers of $W_r(2796)\\ge 0.1$ \\AA; dotted points indicate absorbers of $W_r(2796)\\ge 0.3$ \\AA; and dot-dashed points indicate absorbers of $W_r(2796)\\ge 0.5$ \\AA.  The impact parameter intervals are chosen to contain ten galaxy--absorber pairs per bin, except for the last bin that includes only the remaining seven pairs at largest separations.  The errorbars represent the 68\\% confidence interval.} \\end{center} \\end{figure}  The large gas covering fraction seen in our galaxy sample at small projected distances is in stark contrast to the finding of Gauthier \\etal\\ (2010), who have compared the incidence of Mg\\,II absorbers around luminous red galaxies (LRGs) of $i'<20$ mag at $z\\approx 0.5$. These authors have found that the covering fraction of extended Mg\\,II absorbers of $W_r(2796)\\ge 0.5$ \\AA\\ is {\\em no more than} $40$\\% at $\\rho<50\\ h^{-1}$ kpc (or $\\rho'<35\\ h^{-1}$ kpc after accounting for the luminosity scaling) and $<30$\\% at $\\rho<100\\ h^{-1}$ kpc (or $\\rho'<70\\ h^{-1}$ kpc) from LRGs.  At $W_r(2796)\\ge 0.5$, our galaxy sample shows $\\kappa_{0.5}\\approx 83$\\% at $\\rho'<30\\ h^{-1}$ kpc.  We note that the LRGs are luminous with $M_B-5\\,\\log\\,h<-21.35$ and are understood to reside in massive halos of $M_h \\apg 10^{13} \\hmsol$, whereas the galaxies in our sample are fainter, with a median rest-frame $B$-band magnitude of $\\langle M_{B}-5\\log\\,h\\rangle=-19.6$, and presumably reside in lower-mass halos of $M_h\\sim 10^{12.3} \\hmsol$ at $z\\approx 0.25$.  To examine how the covering fraction of Mg\\,II absorbing gas varies with galaxy luminosity, we first divide our galaxy sample into luminous ($L_B>L_{B_*}$) and faint ($L_B\\le L_{B_*}$) subsamples and then determine the mean gas covering fraction $\\langle\\kappa_{W_0}\\rangle$ over the entire gaseous halo defined by $R_{\\rm gas}$ (see \\S\\ 6.2) at a given \\ewr\\ threshold, $W_0$.  We consider three different threshold values, $W_0=0.1$ \\AA, $W_0=0.3$ \\AA, and $W_0=0.5$ \\AA.  The results are presented in Figure 11, together with the measurement of Gauthier \\etal\\ (2010) for $W_0=0.5$ \\AA\\ around LRGs.  The decreasing mean covering fraction over a fixed halo radius with increasing absorber strength is consistent with the expectation from a decreasing cloud density profile toward large radii (see \\S\\ 6.2).  In addition, we find little dependence between the mean gas covering fraction and galaxy luminosity within our sample, but a significant reduction in gas covering fraction around LRGs.  The declining gas covering fraction of strong Mg\\,II absorbers with increasing galaxy luminosity (halo mass) is qualitatively consistent with the expectation from the observed clustering properties of Mg\\,II absorbers (Tinker \\& Chen 2008, 2009) and the theoretical expectation of diminishing cool gas fraction in massive halos (Kere\\v{s} \\etal\\ 2005, 2009; Dekel \\& Birnboim 2006).  \\begin{figure} \\begin{center} \\includegraphics[scale=0.4]{kappa_mean.eps} \\caption{Mean covering fraction of Mg\\,II absorbers within $R_{\\rm gas}$ of galaxies in different luminosity intervals.  Circles represent a Mg\\,II absorption threshold of $W_0=0.1$ \\AA; squares represent $W_0=0.3$ \\AA; and triangles represent $W_0=0.5$ \\AA.  The data points are located at the median absolute $B$-band magnitudes of the galaxies with horrizontal errorbars representing the luminosity intervals of the subsamples.  Vertical errorbars represent the 68\\% confidence interval.  For comparison, we have also included the limit of Gauthier \\etal\\ (2010) for Mg\\,II absorbers with $\\ewr\\ge 0.5$ \\AA\\ around LRGs at $z\\approx 0.5$.} \\end{center} \\end{figure}   \\subsection{Density Profile of Extended Mg\\,II Absorbing Gas}  The best-fit power-law model presented in Equation (11) and Figure 9 indicates that the observed absorber strength $W_r(2796)$ scales inversely with increasing projected distance $\\rho$ with a steep slope of $\\Delta\\log\\,W_r(2796)/\\Delta\\log\\,\\rho=-1.9$.  While the power-law model appears to describe the observations well, a physical interpretation of this steep anti-correlation is not straightforward. Using a small sample of 13 galaxy--Mg\\,II absorber pairs and 10 galaxies at $\\rho<100\\ h^{-1}$ kpc that do not give rise to Mg\\,II absorption to a sensitive upper limit ($W_r(2796)\\apll 0.02$ \\AA), Chen \\& Tinker (2008) showed that the anti-correlation between $W_r(2796)$ and $\\rho$, after accounting for $B$-band luminosity distribution, is well described by either an isothermal density profile or a Navarro-Frenk-White (NFW; Navarro \\etal\\ 1996) profile of the absorbing clumps with a finite extent $R_{\\rm gas}$.  An isothermal density profile is motivated by the observed rotation curves of nearby galaxies, while an NFW profile is found to represent the density profiles of dark matter halos in high-resolution numerical simulations.  Using the new sample of galaxy--Mg\\,II absorber pairs that is three times of the Chen \\& Tinker (2008) sample, we examine whether the anti-correlation displayed in Figure 9 is better described by either an isothermal density profile or an NFW profile.  For an isothermal profile of gaseous clumps within a finite extent $R_{\\rm gas}$, the \\ewr\\ versus $\\rho$ relation is characterized following Chen \\& Tinker (2008) as \\begin{equation} \\bar{W}_r^{\\rm iso}(2796) =\\frac{W_0}{\\sqrt{\\rho^2/a_h^2+1}}\\tan^{-1}{\\sqrt{\\frac{R_{\\rm gas}^2-\\rho^2}{\\rho^2+a_h^2}}} \\end{equation} at $\\rho\\le R_{\\rm gas}$ and $\\ewr=0$ otherwise.  The core radius $a_h$ is defined to be $a_h=0.2\\,R_{\\rm gas}$ and does not affect the expected $\\ewr$ at large $\\rho$.  The extent of Mg\\,II absorbing gas scales with the luminosity of the absorbing galaxy according to \\begin{equation} \\frac{R_{\\rm gas}}{R_{{\\rm gas}*}}=\\left(\\frac{L_B}{L_{B_*}}\\right)^{\\beta}. \\end{equation} Following the expectations of an isothermal model in Equations (16) and (17), we perform the likelihood analysis described in Equations (7) through (9) to find the best-fit values of $W_0$, $R_{\\rm gas_*}$, and $\\beta$.  Excluding the outliers, the results of the likelihood analysis show that the observations are best described by $\\log\\,W_0^{\\rm iso}=1.24\\pm 0.03$, \\begin{equation} \\beta^{\\rm iso}=0.35_{-0.04}^{+0.01}, \\end{equation} and \\begin{equation} \\log\\,R_{\\rm gas_*}^{\\rm iso}=1.87\\pm 0.01. \\end{equation} The errors indicate the 95\\% confidence intervals.  The best-fit isothermal profile is also characterized by an intrinsic scatter of $\\sigma_c=0.104$ and an r.m.s.\\ residual between the observed and model Mg\\,II absorber strengths of ${\\rm r.m.s.}(\\log\\,W_r-\\log\\,\\bar{W})=0.233$.  For an NFW profile, the absorber strength is expected to vary with $\\rho$ according to \\begin{equation} \\bar{W}_r^{\\rm NFW}(2796)=\\displaystyle\\int_0^{\\sqrt{R_{\\rm gas}^2-\\rho^2}}\\frac{W_0\\,r_s^3}{(r_s+\\sqrt{l^2+\\rho^2})^2\\sqrt{l^2+\\rho^2}}\\,dl \\end{equation} where $r_s$ is the scale radius.  We adopt $r_s\\equiv R_{200}/15$ that gives a halo concentration index of 15 (e.g.\\ Dolag \\etal\\ 2004).  We perform the likelihood analysis described in Equations (7) through (9) and find $\\log\\,W_0^{\\rm NFW}=-0.56\\pm 0.05$, \\begin{equation} \\beta^{\\rm NFW}=0.35\\pm 0.03, \\end{equation} and \\begin{equation} \\log\\,R_{\\rm gas_*}^{\\rm NFW}=1.89_{-0.12}^{+0.05}. \\end{equation} The errors indicate the 95\\% confidence intervals.  The best-fit NFW profile has associated intrinsic scatter of $\\sigma_c=0.175$ and r.m.s.\\ residual between the observed and model Mg\\,II absorber strengths of ${\\rm r.m.s.}(\\log\\,W_r-\\log\\,\\bar{W})=0.253$.  Figure 12 displays the best-fit models in comparison to observations. While the isothermal model provides a somewhat better characterization of the observations in terms of a minimum derived intrinsic scatter and r.m.s.\\ residual, both models result in a consistent best-fit scaling relation of \\begin{equation} \\frac{R_{\\rm gas}}{R_{\\rm gas_*}}= \\left(\\frac{L_B}{L_{B_*}}\\right)^{0.35}, \\end{equation} where the characteristic gaseous radius of an $L_*$ galaxy is found to be \\begin{equation} R_{\\rm gas_*}\\approx 75\\ h^{-1}\\,{\\rm kpc} \\end{equation} at $z\\sim 0.25$.  \\begin{figure} \\begin{center} \\includegraphics[scale=0.4]{w_rhom_m3.eps} \\caption{Comparison of $W_r(2796)$ versus impact parameter accounting for scaling by $B$-band luminosity and redshift.  Symbols are the same as in Figure 8. The solid curve is the best-fit isothermal model and the dash-dotted curve is the best-fit NFW model, excluding five outliers (according to a 3-$\\sigma$ clipping criterion; see \\S\\ 5.2) marked in dotted circles.  The errorbars in the lower-left corner indicate the intrinsic scatters of the adopted model profiles estimated based on the likelihood analysis.  We note that the residuals between observations and the best-fit model profile do not correlate with galaxy $B$-band luminosity, indicating that $W_0/R_{\\rm gas}$ in Equation (16) does not depend on galaxy luminosity.} \\end{center} \\end{figure}  It is clear that the scatter between the observations and the best-fit isothermal model after accounting for the scaling relation with galaxy luminosity remains large in Figure 12.  A natural expectation of the isothermal model in Equation (16) is that $W_0$ depends on galaxy properties, particularly on $L_B$ or mass, because more luminous (and therefore more massive) galaxies possess more extended gaseous halos and would exhibit stronger $W_0$ (e.g.\\ Tinker \\& Chen 2010). However, we find no correlation between the residuals of $(\\log\\,W_r-\\log\\,\\bar{W})$ in Figure 12 and absolute $B$-band magnitude of the galaxies, indicating that $W_0$ does not depend strongly on galaxy luminosity/mass.  The lack of correlation suggests that the mean absorber strength per halo does not vary with halo mass and that the cool gas fraction declines with increasing halo mass in the mass regime probed by our sample.  The best-fit scaling power in Equation (23) confirms the earlier result of Chen \\& Tinker (2008).  The best-fit characteristic gaseous radius shown in Equation (24) is smaller than the earlier estimate of $91_{-8}^{+3}\\ h^{-1}$ kpc from a smaller galaxy sample, but they are consistent to within 2-$\\sigma$ error uncertainties.  The new galaxy sample presented in this paper, which is three times larger than the initial sample employed by Chen \\& Tinker, has revealed a substantial scatter, particularly at large radii.  The large uncertainty in $R_{\\rm gas_*}$ can be understood as due to the inherent degeneracy between the spatial extent $R_{\\rm gas}$ and diminishing covering fraction $\\kappa(\\rho)$ of Mg\\,II absorbing gas at large radii.  A larger $R_{\\rm gas_*}$ can be accommodated with a smaller gas covering fraction.  For example, within $R_{\\rm gas_*}=75\\ h^{-1}$ kpc, 39 of the 59 galaxies have an associated Mg\\,II absorber of $\\ewr\\ge 0.3$ \\AA.  This yields a gas covering fraction of $\\kappa_{0.3}\\approx 66$\\% at $\\rho' < 75\\ h^{-1}$ kpc.  Within $R_{\\rm gas_*}=91\\ h^{-1}$ kpc, we find $\\kappa_{0.3}\\approx 61$\\% at $\\rho' < 91\\ h^{-1}$ kpc after excluding outliers.  In Chen \\& Tinker (2008), the best-fit $R_{\\rm gas_*}=91\\ h^{-1}$ kpc was close to the maximum impact separation of $r_0=100\\ h^{-1}$ kpc in the galaxy--QSO pair sample.  The selection criterion was imposed by these authors to minimize the incidence of correlated galaxies through large-scale clustering (see their Section 2 for discussion).  Chen \\& Tinker examined the possible bias in the best-fit $R_{{\\rm gas}_*}$ due to the imposed $r_0$ selection, and found that the best-fit $R_{\\rm gas_*}$ was insensitive to the choice of $r_0$.  Namely, they obtained a consistent estimate of $R_{\\rm gas_*}$ for a smaller $r_0$. However, choosing a smaller $r_0$ reduced the already small sample further and the parameters became poorly constrained.  In the present analysis, we have included galaxy--QSO pairs with impact separation as large as $\\rho\\approx 120\\ h^{-1}$ kpc (Figure 3).  The best-fit gaseous radius of $R_{\\rm gas_*}=75\\ h^{-1}$ kpc is significantly below the the maximum separation in the pair sample, although the majority of the data (galaxies with intrinsic colors consistent with present-day disks and irregular's/starbursts) lie at $\\rho<75\\ h^{-1}$ kpc.  To examine whether the best-fit $R_{\\rm gas_*}$ is sensitive to the range of impact parameter covered by the pair sample, we repeat the likelihood analysis adopting only pairs with $\\rho<75\\ h^{-1}$ kpc.  We find that the best-fit results remain the same.  In summary, we find based on a sample of 71 'isolated' galaxies that the extent of cool gas as probed by the Mg\\,II absorption features scales with galaxy $B$-band luminosity according to $R_{\\rm gas}\\propto\\,L_B^{0.35}$, consistent with earlier results obtained using smaller samples.  As discussed in Chen \\& Tinker (2008), the scaling relation, when combined with a fiducial mass-to-light ratio determined from galaxy halo occupation studies of wide-field galaxy survey data, indicates that the gaseous radius is a constant fraction of the halo radius over the mass scale of $M_h=10^{10.6}-10^{13}\\,\\hmsol$ sampled in our observations\\footnote{The halo occupation analysis of 2dFGRS galaxies from Tinker \\etal\\ (2007) showed that the relationship between $M_h$ and luminosity for galaxies of $L_{b_J}\\lesssim L_{*}$ at $z\\sim 0.1$ is $M_h = 10^{12.5}\\,(L_B/L_{B_*})^{1.3}\\,\\hmsol$.  This monotonic relationship is appropriate for galaxies that reside at the {\\it centers} of their dark matter halos and are the brightest galaxy in the halo.  A similar scaling relation was obtained, $M_h = 10^{11.9}\\,(L_B/L_{B_*})^{1.3}\\,\\hmsol$, for DEEP2 galaxies at $z\\sim 1$ by Zheng \\etal\\ (2007).  Interpolating in $\\log\\,(1+z)$, we find that galaxies of $M_{B_*}-5\\log\\,h=-19.8$ on average reside in halos of $10^{12.4}\\,\\hmsol$ at $z\\sim 0.25$.}.  In addition, we find that the residuals of observed and best-fit $W_r(2796)$ do not correlate with galaxy absolute $B$-band magnitude, implying a declined cold gas fraction with increasing halos mass in the mass regime probed by our sample.  \\subsection{The Origin of Extended Cool Gas and Implications for the Baryon Content of Galactic Halos}  The presence of a finite extent $R_{\\rm gas}$ of Mg\\,II absorbing gas is similar to what is found for extended C\\,IV (Chen \\etal\\ 2001) and also for OVI absorbing gas recently published by Chen \\& Mulchaey (2009) around galaxies.  The origin of such a finite boundary for metal-line absorbers can be interpreted as due to a halo fountain phenomenon (c.f.\\ Bregman 1980), according to which the gaseous radius is driven by the finite distance the outflowing material can travel from an early episode of starburst.  We find this scenario unlikely because of the broad range of intrinsic colors covered by our galaxy sample, but a detailed stellar population synthesis to investigate the star formation history will shed more light on this issue.  Alternatively, the finite boundary of absorbing clouds probed by these metal-line absorbers can be understood as a critical radius below which cool clouds can form and survive in an otherwise hot medium. This two-phase model to interpret QSO absorption line systems was formulated in Mo \\& Miralda-Escud\\'e (1996) and later re-visited by Maller \\& Bullock (2004).  This is also shown in recent high-resolution numerical simulations of Milky Way type halos, in which accreted cold streams are disrupted by shocks but the remaining gas overdensities serve as the 'seeds' necessary to form cool gaseous clouds through thermal instabilities within $\\approx 1/3$ of the halo radius (e.g.\\ Kere{\\v s} \\& Hernquist 2009).  The formation and presence of such cool clouds in a hot halo have several important implications for the studies of galaxy formation and evolution.  For example, the confining hot medium around the cool clouds is expected to have a low density and long cooling time, reducing the ``overcooling problem'' in standard galaxy formation models (e.g.\\ Maller \\& Bullock 2004; Kaufmann \\etal\\ 2009).  In addition, these clouds may provide the fuels necessary to support continuous star formation near the center of galactic halos (e.g.\\ Binney \\etal\\ 2000).  Furthermore, the cool clouds together with the confining hot medium offer a reservoir for missing baryons in individual galactic halos (e.g.\\ McGaugh \\etal\\ 2010).  A nominal candidate for these cool clouds in our Milky Way Halo is the high velocity clouds (HVCs) of neutral hydrogen column density $N({\\rm H\\,I})>10^{18}$ \\cmjj\\ seen in all-sky 21~cm observations.  However, unknown distances of these HVCs prohibit an accurate measurement of their total gas mass (e.g.\\ Putman 2006).  High-resolution spectra of strong Mg\\,II absorbers have revealed the multi-component nature of the absorbing gas, with \\ewr\\ roughly proportional to the number of components in the system (e.g.\\ Petitjean \\& Bergeron 1990; Prochter et al. 2006).  These discrete Mg\\,II absorbing components are natural candidates of the condensed cool clouds within $\\sim 75\\ h^{-1}$ kpc radius of known galaxies. Interpreting the intrinsic scatter $\\sigma_c$ as due to Poisson noise in the number of cool clumps intercepted along a line of sight allows us to estimate the number of absorption clumps per galactic halo per sightline, $n^{\\rm clump}$.  In this simple model, each absorber is characterized as $W_r=n^{\\rm clump}\\times k$, where $k$ is the mean absorption equivalent width per absorbing component.  Following Poisson counting statistics, the intrinsic scatter $\\sigma_{c}\\equiv \\sigma_{c,W_r}/(W_r\\,\\ln\\,10)$ is related to $n^{\\rm clump}$ according to \\begin{equation} \\sigma_{c} = \\frac{1}{\\ln\\,10}\\frac{\\sqrt{n^{\\rm clump}+1}}{n^{\\rm clump}} \\end{equation} For the isothermal model, we have determined a best-fit intrinsic scatter of $\\sigma_c=0.104$, which leads to a mean number of absorbing clumps of \\begin{equation} n^{\\rm clump} \\sim 18 \\end{equation} and a mean absorption equivalent width per clump of $k = 0.04$ \\AA\\ at a median projected distance of $\\langle\\rho\\rangle=26\\ h^{-1}$ kpc. The corresponding Mg\\,II absorbing column density per clump is $N({\\rm Mg\\,II})\\approx 10^{12}$ \\cmjj.  We note that if other factors contribute to the observed $\\sigma_c$, then the corresponding Poisson noise is smaller and the inferred $n^{\\rm clump}$ will be bigger.  To estimate the total baryonic mass contained in these cool clouds traced by Mg\\,II absorbers, we first assume a mean clump size of 1 kpc (consistent with the size of H\\,I clouds seen around M31; Westmeier \\etal\\ 2008), a mean metallicity of 1/10 solar (consistent with what is seen in the HVC Complex C; see Thom \\etal\\ 2008 for a list of references), and a mean ionization fraction of $f_{\\rm Mg^+}=0.1$ for photo-ionized clouds (see Figure 6 of Chen \\& Tinker 2008).  We derive a mean gas density of $n_{\\rm H}\\approx 10^{-3}\\ {\\rm cm}^{-3}$. Assuming a spherical shape for the clumps leads to a mean clump mass of $M_b^{\\rm clump}=1.7\\times 10^4\\,{\\rm M}_\\odot$.  Next, we adopt an isothermal density distribution of the clumps (Equation 16) and derive a total number of clumps $N^{\\rm clump}=2\\times 10^5\\, h^{-2}$ within $R_{\\rm gas}=75\\ h^{-1}$ kpc.  The total baryonic mass in these cool clumps is found to be $M_b\\sim 3\\times 10^9\\,h^{-2}\\,{\\rm M}_\\odot$, comparable to the total cold gas content seen in the Milky Way disk (e.g.\\ Flynn \\etal\\ 2006).  The inferred total mass would be still larger if $n^{\\rm clump}$ increases as the Poisson noise decreases.  If we further assume that it takes a free-fall time for the clumps to reach the center of the halo, we infer a cool gas accretion rate of $\\dot{M}_{\\rm gas}\\sim 3\\,{\\rm M}_\\odot\\,{\\rm yr}^{-1}$.  The inferred gas accretion rate is interestingly close to the star formation rate seen in the Milky Way disk (e.g.\\ Robitaille \\& Whitney 2010). However, we note two issues in this simple picture.  The first issue concerns the lifetime of these cool clumps relative to the time it takes for them to reach the inner few kpc of the host halo.  Previous studies have shown that clouds of mass $M_c\\apll 10^5\\,{\\rm M}_\\odot$ cannot survive thermal conduction and ram pressure as they move through a low-density, hot medium (see e.g.\\ Mo \\& Miralda-Escud\\'e 1996; Maller \\& Bullock 2004; Heitsch \\& Putman 2009).  To alleviate this problem, the clumps will need to have lower metallicities and/or bigger sizes than what we have assumed.  For example, if we assume a mean clump size of 2 kpc and a mean metallicity of 0.05 solar, then we derive a mean clump mass of $M_b^{\\rm clump}=1.3\\times 10^5\\,{\\rm M}_\\odot$.  The total number of clumps would be $N^{\\rm clump}=5\\times 10^4\\, h^{-2}$ and the total cool gas mass would be $M_b\\sim 6.5\\times 10^9\\,h^{-2}\\,{\\rm M}_\\odot$ at $r\\le 75\\ h^{-1}$ kpc.  Second, we note that the expected H\\,I column density of these Mg\\,II absorbing clumps is $N({\\rm H\\,I})\\sim 10^{16}$ \\cmjj, below the $N({\\rm H\\,I})$ threshold provided by 21~cm observations.  The H\\,I content is therefore too low for these clumps to be the distant counterparts of the local HVCs.  Nevertheless, the lack of strong ($\\ewr\\ge 0.5$ \\AA) Mg\\,II absorbers beyond $30\\ h^{-1}$ kpc (Figure 10) which have have associated H\\,I of $N({\\rm HI})> 10^{18}$ \\cmjj\\ (Rao \\etal\\ 2006), constrains the distances of the Milky Way HVCs at $d\\apll 50$ kpc.  Within this close distance range, the implied total gas mass in the HVCs is expected to be $M_b^{\\rm HVC}<10^9\\,{\\rm M}_\\odot$ (Putman 2006).  The exercise presented above shows that the cool clouds that give rise to Mg\\,II absorbers and the implied presence of a hot confining medium can provide a substantial reservoir for missing baryons in individual galactic halos (see also Kaufmann \\etal\\ 2009).  The presence of extended Mg\\,II absorbing halos around galaxies of a broad range of intrinsic colors argues against an outflow origin of these absorbers. Under an infall scenario, we show that these cool clouds provide sufficient fuels to support continuous star formation in their host galaxies.  \\subsection{Comparison with Other Studies}  The large covering fraction of Mg\\,II absorbing gas (\\S\\ 6.1) and the scaling relation between gaseous extent and galaxy $B$-band luminosity (\\S\\ 6.2) represent the empirical evidence supporting the halo occupation model developed in Tinker \\& Chen (2008, 2009; see also Gauthier \\etal\\ 2009).  Here we compare the results of our studies with recent surveys by other groups.  First, we have noted the discrepancy between the high gas covering fraction around $\\sim L_*$ galaxies and the low incidence of Mg\\,II absorbers around luminous red galaxies studied in Gauthier \\etal\\ (2010).  We have interpreted the contrast as due to the different halo mass scales probed by the two galaxy samples.  Cosmological simulations have shown that the growth of galaxies progresses primarily through stable accretion of intergalactic gas (e.g.\\ Kere\\v{s} \\etal\\ 2005; 2009).  In low-mass halos ($M_h < M_h^{\\rm crit}=10^{11.5}\\,\\hmsol$), accretion proceeds to the center of the halo through dense cold streams (Dekel \\& Birnboim 2006; Kere\\v{s} \\etal\\ 2005; 2009).  As galaxies grow larger in mass ($M_h > M_h^{\\rm crit}=10^{11.5}\\,\\hmsol$), both numerical simulations and analytic models show that a progressively larger fraction of the accreted gas is shock heated to the virial temperature.  Some fraction of the shock-heated gas may be able to cool in the inner region where gas density is high.  The majority of the galaxies presented in this paper are on the transitional mass scale, $M_h = 10^{11-12.5}\\,\\hmsol$.  In contrast, the luminous red galaxies studied by Gauthier \\etal\\ (2010) are in a higher mass regime of $M_h>10^{13}\\,\\hmsol$, where a diminishing amount of cold gas is expected because the cooling time is long.  The differential gas covering fraction found in our galaxies and in those luminous red galaxies is therefore consistent with theoretical expectations.  Kacprzak \\etal\\ (2008) have also studied the extent and covering fraction of Mg\\,II absorbing gas based on a sample of 37 galaxies at $z=0.3-1$ that are found in the vicinity of a known Mg\\,II absorber. While they did not see a clear correlation between the extent of Mg\\,II absorbing gas and galaxy luminosity, they found that for a scaling power of $\\beta=0.18-0.58$ and a characteristic gaseous extent of $R_*=56-105\\ h^{-1}$ kpc the covering fraction of Mg\\,II absorbers with $\\ewr\\ge 0.3$ \\AA\\ is no more than 70\\%.  Our best-fit isothermal model indicates that the characteristic gaseous radius at $\\ewr=0.3$ \\AA\\ is $R_{\\rm gas_*}(W_r=0.3\\AA)=47\\ h^{-1}$ kpc (Figure 11), less than the gaseous radius considered by Kacprzak \\etal.  Excluding the outliers, we find a mean covering fraction of $\\kappa_{0.3}[\\rho'<47\\ h^{-1}\\,{\\rm kpc}]\\approx 76$\\% for Mg\\,II absorbers with $\\ewr\\ge 0.3$ \\AA.  Including the outliers, we find a mean covering fraction of $\\kappa_{0.3}[\\rho'<47\\ h^{-1}\\,{\\rm kpc}]\\approx 78$\\% for Mg\\,II absorbers with $\\ewr\\ge 0.3$ \\AA.  The gas covering fraction found in our sample is somewhat higher than the parameter space explored by Kacprzak \\etal, but this can be explained by the larger gaseous radius considered by these authors.  In addition, Barton \\& Cooke (2009) have recently carried out a search of coincident Mg\\,II absorbers in the vicinity of 20 luminous galaxies ($M_r-5\\,\\log\\,h\\le -20.5$) identified at $z\\sim 0.1$ in the SDSS galaxy spectroscopic sample.  The authors reported a low gas covering fraction of $\\kappa_{0.3}[\\rho<75\\ h^{-1}\\ {\\rm kpc}]\\apll 0.4$ and $\\kappa_{0.3}[\\rho<35\\ h^{-1}\\ {\\rm kpc}]\\sim 0.25$ for Mg\\,II absorbers of $\\ewr\\ge 0.3$ \\AA.  These results are inconsistent with the high gas covering fraction found in our survey.  The design of the Barton \\& Cooke survey is very similar to ours presented here, with only small differences in the targeted redshift range ($z\\sim 0.1$) and luminosity scale ($>L_*$).  The discrepancy in the observed gas covering fraction therefore requires further perusal.  A fundamental difference in our respective measurements is that Barton \\& Cooke excluded observed strong Mg\\,II absorbers in the covering fraction calculation if the QSO spectrum did not have sufficient $S/N$ for identifying absorbers of $\\ewr=0.3$ \\AA.  Of the six galaxies at $\\rho \\le 35\\ h^{-1}$ kpc, three have associated Mg\\,II absorbers of $\\ewr=1.93\\pm 0.18$ \\AA, $\\ewr=1.91\\pm 0.10$ \\AA, and $\\ewr=0.71\\pm 0.15$ \\AA, respectively, and three do not have corresonding Mg\\,II absorbers to 3-$\\sigma$ limits of $\\ewr<0.3$ \\AA.  The authors excluded from the covering fraction calculation two strong absorbers, $\\ewr=1.93\\pm 0.18$ \\AA\\ and $\\ewr=0.71\\pm 0.15$ \\AA, and derived $\\kappa_{0.3}[\\rho<35\\ h^{-1}\\ {\\rm kpc}]\\sim 0.25$ based on the remaining four galaxies.  We note that the goal of our study is to determine the incidence of Mg\\,II absorbers at the location of a known galaxy.  Absorbers that are observed to have $\\ewr=0.71\\pm0.15$ \\AA\\ and $\\ewr=1.93\\pm 0.18$ \\AA\\ confirm the presence of such absorbing gas in large quantities, despite the fact that the QSO spectrum may not offer the sensitivities required for uncovering weaker absorbers of $\\ewr\\sim 0.3$ \\AA\\ at high confidence levels.  Excluding these sightlines imposes a bias that would result in an underestimate of the gas covering fraction.  We have retrieved from the SDSS data archive the galaxy photometric and spectroscopic data in Barton \\& Cooke (2009), and determined the rest-frame $B$-band absolute magnitude and $B_{AB}-R_{AB}$ color for each of the 20 galaxies according to the procedures described in \\S\\ 4.1.  The rest-frame absolute $B$-band magnitudes of these galaxies span a range from $M_{B}-5\\log\\,h=-19.6$ to $M_{B}-5\\log\\,h=-20.7$ with a median of $\\langle M_{B}-5\\log\\,h\\rangle=-20.3$.  The rest-frame $B_{AB}-R_{AB}$ colors range from $B_{AB}-R_{AB}\\approx 0.6$ to $B_{AB}-R_{AB}\\approx 1.2$ with a median of $\\langle B_{AB}-R_{AB}\\rangle_{\\rm med}\\approx 1.0$.  Figure 13 displays the observed \\ewr\\ versus $\\rho$ relation for the Barton \\& Cooke sample, superimposed on top of our own data.  We have converted the 3-$\\sigma$ upper limits published by these authors to 2-$\\sigma$ upper limits in order to be consistent with our own measurements.  \\begin{figure} \\begin{center} \\includegraphics[scale=0.4]{w_rhom_comp.eps} \\caption{Comparison of the observed $\\ewr$ versus $\\rho$ correlation for galaxies in our sample and those published in Barton \\& Cooke (2009).  We have applied the scaling relation of Equation (23) for both samples.  Symbols and best-fit models for our sample are the same as in Figure 11.  The galaxies of Barton \\& Cooke are shown in star symbols.  We note that the galaxy at $\\log\\,\\rho'=\\log\\,\\rho+0.14\\,(M_B-M_{B_*})\\approx 1.6$ with a strong Mg\\,II absorber of $\\ewr=1.24$ (SDSSJ144033.82$+$044830.9) in the Barton \\& Cooke sample has been found to have a companion galaxy at $\\rho=18.6\\ h^{-1}$ kpc.  The absorber would therefore be considered as in a 'group' environment.} \\end{center} \\end{figure}  Figure 13 shows that with the exception of two outliers at $\\rho'\\approx 25\\ h^{-1}$ kpc, the Barton \\& Cooke sample is in general agreement with what is seen in our sample.  Six of the 12 upper limits in the Barton \\& Cooke sample do not have sufficient sensitivities for constraining the presence/absence of Mg\\,II absorbing gas at the level of $\\ewr\\sim 0.3$ \\AA.  In addition, the galaxy at $\\log\\,\\rho'=\\log\\,\\rho+0.14\\,(M_B-M_{B_*})\\approx 1.6$ with a strong Mg\\,II absorber of $\\ewr=1.24$ (SDSSJ144033.82$+$044830.9) was later found to have a companion galaxy at $\\rho=18.6\\ h^{-1}$ kpc. The absorber would therefore be considered as in a 'group' environment.  The galaxy at $\\rho'\\approx 20\\ h^{-1}$ kpc with a 2-$\\sigma$ upper limit of $\\ewr<0.13$ \\AA\\ (SDSSJ151541.23$+$334739.4) has a luminous companion (SDSSJ151536.18$+$334743.6) at projected distance $< 100 \\ h^{-1}$ kpc and velocity separation $\\Delta\\,v\\approx 27$ \\kms\\ away.  This pair also respresents a 'group' that we have treated separately in our analysis.  Considering only pairs between 'isolated' galaxies and QSOs for which sensitive limits for the corresponding Mg\\,II absorption are available leads to a sample of 12 galaxies with sensitive constraints for the corresponding Mg\\,II absorption features.  We find three of the four galaxies at $\\rho'<35\\ h^{-1}$ kpc have associated Mg\\,II absorbers of $\\ewr>0.3$ \\AA, yielding $\\kappa_{0.3}[\\rho'<35\\ h^{-1}\\ {\\rm kpc}]\\approx 0.75$.  We find seven of the 12 galaxies at $\\rho'<75\\ h^{-1}$ kpc have associated Mg\\,II absorbers of $\\ewr>0.3$ \\AA, yielding $\\kappa_{0.3}[\\rho'<75\\ h^{-1}\\ {\\rm kpc}]\\approx 0.58$. These values are consistent with our estimated gas covering fraction to within 1-$\\sigma$ uncertainties.  We therefore conclude that the observations presented in Barton \\& Cooke (2009) are consistent with the high gas covering fraction seen in our sample.  Finally, Zibetti \\etal\\ (2007) studied the nature of Mg\\,II absorbing galaxies at $z=0.37-1$ based on stacked images of QSOs with known foreground Mg\\,II absorbers of $\\ewr > 0.8$ \\AA.  After subtracting the QSO point spread function in the stacked images, these authors obtained an angular averaged image of Mg\\,II absorbing galaxies out to 100 kpc.  They found that the luminosity-weighted mean colors of the extended emission is consistent with present-day intermediate-type galaxies.  In addition, the authors found that weak absorbers of $\\ewr<1.1$ \\AA\\ originate primarily from red passive galaxies while stronger absorbers display bluer colors consistent with star-forming galaxies.  Based on the differential color distribution, Zibetti \\etal\\ argued that the origin of strong Mg\\,II absorbers might be better explained by models of metal-enriched gas outflows from star-forming/bursting galaxies.  Our analysis in \\S\\ 5.3 shows that the observed Mg\\,II absorber strength does not depend strongly on galaxy intrinsic color, which appears to contradict the observations of Zibetti \\etal.  To understand the discrepant results, we first recall based on the observed $\\ewr$ versus $\\rho$ anti-correlation in Figure 8 that the mean impact parameter of the absorbing galaxies for weak absorbers is expected to be larger than that for stronger absorbers.  In addition, Figure 8 also shows that for a sample of randomly selected QSO--galaxy pairs in our study, the fraction of red galaxies is found to be a factor of two larger at larger impact parameters $\\rho>30\\ h^{-1}$ kpc than at closer distances.  This is not due to selection biases, because we did not preferentially select blue galaxies at smaller impact parameters.  Instead, this can be understood by considering the fact that redder galaxies are on average more luminous and have a lower space density than bluer and on average fainter galaxies (e.g.\\ Faber \\etal\\ 2007).  The increased incidence of red galaxies at larger impact parameter, combined with the luminosity-weighted nature of the image stacking procedure, is expected to skew the stacked images of weaker absorbers toward redder colors.  Whether or not this is sufficient to explain the observed color gradient between strong and weak Mg\\,II absorbers in the Zibetti \\etal\\ sample requires a detailed simulation that takes into account such intrinsic weighting by the space density and luminosity of the galaxies.  Because such analysis has not been performed, it is not clear whether the respective results are inconsistent with each other.  Nevertheless, the conclusion in favor of strong Mg\\,II absorbers being better explained by starburst outflows based on the observed color gradient should be viewed with caution."
    },
    "1004/1004.0901.txt": {
        "abstract": "We present an updated catalog of $113$ X-ray flares detected by \\emph{Swift} in the  $\\sim33\\%$ of the  X-ray afterglows of Gamma-Ray Bursts (GRB). $43$ flares have a measured redshift. For the first time the analysis is performed in $4$ different X-ray energy bands, allowing us to constrain the evolution of the flare temporal properties with energy. We find that flares are  narrower at higher energies: their width follows a power-law relation $w\\propto E^{-0.5}$ reminiscent of the prompt emission. Flares are asymmetric structures, with a decay time which is twice the rise time on average. Both time scales linearly evolve with time, giving rise to a constant rise-to-decay ratio: this implies that both time scales are stretched by the same factor.  As a consequence, the flare width \\emph{linearly} evolves with time to larger values: this is a key point that clearly distinguishes the flare from the GRB prompt emission. The flare $0.3-10$ keV peak luminosity decreases with time, following a power-law behaviour with large scatter: $L_{pk}\\propto t_{pk}^{-2.7\\pm0.5}$. When multiple flares are present, a global softening trend is established: each flare is on average softer than the previous one. The $0.3-10$ keV isotropic energy distribution is a log-normal peaked at $10^{51}$ erg, with a possible excess at low energies. The flare average spectral energy distribution (SED) is found to be a power-law with spectral energy index $\\beta\\sim1.1$. These results confirmed that the flares are tightly linked to the prompt emission. However, after considering various models we conclude that no model is currently able to account for the entire set of observations. ",
        "introduction": "Gamma Ray Bursts (GRB) are short flashes of gamma-rays that during their early lifetime outshine any other source of gamma rays in the sky.  The first event was detected in 1967 and announced in 1973 \\citep{1973ApJ...182L..85K}. Since then and after about 40 years of research our knowledge has increased significantly mainly thanks to three high-energy missions, \\emph{Compton Gamma-Ray Observatory} (CGRO), \\emph{Beppo}-SAX and \\emph{Swift} that, together with related theoretical works, marked fundamental milestones in our knowledge of this phenomenon. These observations characterized the main features of these events.  The timescale of the prompt emission lasts from a few milliseconds \\citep{1981RSPTA.301..645V} to thousands of seconds \\citep{2002astro.ph.11620H}. The distribution of its duration has been shown to be bimodal \\citep{1981Ap&SS..80..119M,1984Natur.308..434N,1989cgrc.conf..337H,1992AIPC..265..304D,1993ApJ...413L.101K}, therefore GRBs can be classified as ``short'' and ``long''. The time profile of the prompt emission may present either multiple spikes of very short duration or relatively broad peaks with no fast variability \\citep{2005ApJ...627..324N}. After the discovery of the isotropic distribution of the BATSE GRBs on the sky \\citep{1992Natur.355..143M,1994ApJS...92..229F,1996ApJ...459...40B,1999ApJS..122..465P} and the detection of the afterglow of GRB970228 \\citep{1997IAUC.6572....1C,1997IAUC.6584....1G,1997Natur.386..686V,1999ApJS..122..465P,2001ApJ...562..654D} it was clearly demonstrated that at least long GRBs were extragalactic and involved the emission of huge amounts of energy in a short time. After the launch of \\emph{Swift} \\citep{2004ApJ...611.1005G} it was firmly established that short GRBs also have an extragalactic origin \\citep{2005Natur.437..851G,2005GCN..3570....1B,2005Natur.437..855V,2005Natur.438..994B} and therefore these bursts involve the emission of a rather large amount of energy as well.  The X-Ray Telescope \\citep[XRT][]{2005SSRv..120..165B} on board the \\emph{Swift} satellite allows the early and well sampled observations of the afterglow. The temporal behaviour of the observed light curve was completely unexpected since, according to the data gathered by \\emph{Beppo}-SAX (the record pointing of this satellite after detection was however of about $4$ hours), the expectation was for a flux decaying smoothly as a power-law, $F \\sim t^{-1.5}$. It was realized rather soon by the \\emph{Swift} team that the light curve of many GRBs was characterized by a more complex temporal behaviour \\citep{2006ApJ...642..389N,2006ApJ...647.1213O,2006ApJ...642..354Z}: a steep early decay, a ``plateau'' and a late decay where the slope observed by \\emph{Beppo}-SAX is essentially recovered. These phases can be either all or in part present \\citep{2009MNRAS.397.1177E}. The significant achievements have been accompanied by substantial theoretical effort to interpret the data. The internal - external shock model (\\citealt{1993ApJ...418L..59M,1992MNRAS.258P..41R,1994ApJ...430L..93R}, see also \\citealt{1999PhR...314..575P,2004RvMP...76.1143P} and references therein) within the fireball scenario \\citep{1978MNRAS.183..359C} explains many of the characteristics of the observed light curve and spectrum of GRBs (see e.g. \\citealt{2006ApJ...642..354Z} and reference therein; for a critical review of this model and possible alternatives see \\citealt{2009arXiv0911.0349L}).  Indeed one of the most intriguing discoveries of the \\emph{Swift}/XRT was the existence of flares in many of the observed GRB afterglows, that released a large emission of energy at later times than the prompt emission \\citep{2005Sci...309.1833B,2006ApJ...641.1010F,2007ApJ...671.1903C,2007ApJ...671.1921F}. The first detection of flares with the \\emph{Swift}/XRT occurred in X-Ray Flash (XRF) 050406 \\citep{2005Sci...309.1833B,2006A&A...450...59R} and GRB050502B \\citep{2005Sci...309.1833B,2006ApJ...641.1010F}. In the first case the afterglow light curve exhibits a rebrightening of a factor of $6$ that decayed quickly to recover the previous temporal behaviour. The flare in GRB050502B was spectacular with a rebrightening of the light curve of a factor $500$; its fluence is comparable to the one of the prompt emission observed by the Burst Alert Telescope \\citep[BAT][]{2005SSRv..120..143B} on board the \\emph{Swift} satellite. Further observations confirmed that flares are quite common events in the light curves of GRBs ($\\sim 33\\%$ of GRB afterglows exhibit flares). The energetics involved as well as their spectral properties, in particular the hard-to-soft evolution, are strong indications that X-ray flares have a common origin with the gamma-ray pulses. Furhermore, the presence of an underlying continuum with the same slope before and after the flaring activity excludes the possibility that flares are related to the afterglow emission by forward external shocks. Therefore their properties can provide an important clue toward the understanding of the mechanism that is at the basis of the GRB phenomenon.  Previous analysis was performed by \\citet[][hereafter Paper I]{2007ApJ...671.1903C} and \\citet[][hereafter Paper II]{2007ApJ...671.1921F}. These authors concluded that: \\begin{enumerate} \\item\tflares occur in all kind of GRBs: short and long, high energy peaked GRBs and X-ray flashes; \\item\tthe flare intensity decreases with time and the flare duration increases with time; \\item\ta sizable fraction of flares cannot be related to the external shock mechanism; \\item\tthe temporal behaviour of flares is very similar to the one of prompt emission pulses; \\item\tthe number of flares of a single event does not correlate with the number of detected prompt pulses; \\item\tthe energy emitted during a bright flare is very large and in some cases is of the order of the prompt emission observed by BAT. Their average fluence, however, is about $10\\%$ of the prompt emission fluence measured by BAT; \\item\tthe peak energy is typically in the soft X-rays, $\\leq 1$ keV; \\item\tthe hardness ratio evolves following closely the evolution of the flare luminosity with a hardening during the rise and a softening during the decay; \\item\ta long lasting activity by the central engine is advocated. \\end{enumerate}  In the present work we expand the statistics of Paper I considering a wider sample of X-ray flares. Moreover, for the first time, we constrain the evolution of the properties of the flares in different X-ray energy bands inside the $0.3-10$ keV bandpass of the XRT. Finally, the flares are fitted with the function proposed in \\citet{2005ApJ...627..324N}: this allows us to study the asymmetry of the flare temporal profiles, to assess the rise time and the decay time evolution with time.  The paper is organized as follows. In Sec. \\ref{data_analysis} we describe the data reduction procedure and in Sec. \\ref{sample} the flare sample and the fitting procedure. In Sec. \\ref{analysis} we describe the analysis of the temporal behaviour of flares (Sec. \\ref{shape}) and their energetic and spectral properties (Sec. \\ref{energy}). In Sec. \\ref{discussion} we discuss the main results of our analysis. Then conclusions follow.  Throughout the paper we follow the convention $f_{\\nu}(t)\\propto \\nu^{-\\beta}t^{-\\alpha}$, where the energy spectral index $\\beta$ is related to the photon index $\\Gamma=\\beta+1$. We have adopted the standard values of the cosmological parameters: $H_\\circ=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.27$ and $\\Omega_{\\Lambda}=0.73$. Errors are given at $1\\, \\sigma$ confidence level unless otherwise stated. ",
        "conclusions": ""
    },
    "1004/1004.1875_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro}  In Chap.~V [Magnetohydrodynamics of Protostellar Disks] it was shown that a sufficiently highly conducting Keplerian disk that is threaded by a weak magnetic field will be subject to the magnetorotational instability (MRI) and may evolve into a turbulent state in which the field is strongly amplified and has a {\\em small-scale}, {\\em disordered} configuration. This turbulence has been proposed as the origin of the effective viscosity that enables matter to accrete by transferring angular momentum {\\em radially} out along the disk plane. In this chapter we focus on an alternative mode of angular momentum transport that can play an important role in protostellar disks, namely, {\\em vertical} transport through the disk surfaces effected by the stresses of a {\\em large-scale}, {\\em ordered} magnetic field that threads the disk.  The possible existence of a comparatively strong, ``open'' magnetic field over much of the extent of at least some circumstellar disks around low- and intermediate-mass protostars is indicated by far-IR and submillimeter polarization measurements, which have discovered an ordered, hourglass-shaped field morphology on sub-parsec scales in several molecular clouds (e.g., Schleuning 1998; Girart et al. 2006; Kirby 2009). The polarized radiation is attributed to thermal emission by spinning dust grains whose short axes are aligned along the magnetic field (e.g., Lazarian 2007). The detected hourglass morphology arises naturally in molecular cloud cores in which a large-scale magnetic field provides dynamical support against the core's self-gravity (see Sect.~\\ref{subsec:forces}). In this picture the field is {\\em interstellar} in origin and is part of the Galactic magnetic field distribution. As discussed in Sect.~\\ref{sec:formation}, the inward gravitational force can become dominant and the core then undergoes dynamical collapse to form a central protostar and a circumstellar disk. The magnetic field is dragged in by the infalling matter and could in principle lead to a large-scale ``open'' field configuration in the disk.  An ordered magnetic field that threads a disk can exert a magnetic torque that removes angular momentum from the interior gas. This angular momentum can be carried away along the field lines either by {\\em torsional Alfv\\'en waves} in a process known as magnetic braking (see Sect.~\\ref{subsec:braking}) or through a rotating outflow in what is known as a {\\em centrifugally driven wind} (CDW; see Sect.~\\ref{subsec:wind}). These mechanisms could supplement or even entirely supplant the radial transport along the disk plane invoked in traditional disk models: by turbulent stresses as in the MRI scenario mentioned above or through gravitational torques in a self-gravitating disk as described in Chap.~IV [Disk Hydrodynamics]. In the case of radial transport, the angular momentum removed from the bulk of the matter is deposited into a small amount of gas at the outer edge of the disk. In the case of vertical transport, this angular momentum is deposited into a small fraction of the disk mass (the tenuous surface layers of the disk) that is removed as a CDW or else (when magnetic braking operates) into the ambient medium through which the torsional Alfv\\'en waves propagate. The introduction of this new transport channel has profound implications to the structure and properties of disks in which it is a major contributor to the angular momentum budget and potentially also to the strong connection that has been found between accretion and outflow phenomena in young stellar objects. This is discussed in Sect.~\\ref{sec:vertical}, where we also consider how to determine which of the two possible angular momentum transport modes (radial or vertical) operates at any given location in a magnetically threaded disk and whether these two modes can coexist.  Of course, large-scale, ordered magnetic fields can also be produced {\\em in situ} by a {\\em dynamo} process; we consider this alternative possibility for the origin of the disk field in Sect.~\\ref{subsec:driving}. In the case of the Sun, high-resolution observations made at extreme ultraviolet and soft X-ray wavelengths and transformed into spectacular false-color images have revealed a complex web of organized structures that appear as loops and prominences near the stellar surface but simplify to a more uniform distribution further out (e.g., Balogh et al. 1995). There is growing evidence that Sun-like stars are already magnetically active in the protostellar phase and, in fact, generate fields that are a thousand times stronger than that of the present-day Sun. The dynamical interaction between such a field and a surrounding accretion disk through which mass is being fed to the nascent star could have important evolutionary and observational consequences. We consider this in Sect.~\\ref{sec:disk-star}. We conclude with a summary and a discussion of future research directions in Sect.~\\ref{sec:conclude}. ",
        "conclusions": "\\label{sec:conclude}  The discussion in this chapter can be summarized as follows:  \\begin{itemize}  \\item There is strong observational evidence for a disk--wind connection in protostars. Large-scale, ordered magnetic fields have been implicated theoretically as the most likely driving mechanism of the observed winds and jets. The ubiquity of the outflows may be related to the fact that centrifugally driven winds (CDWs) are a potentially efficient means of transporting angular momentum from the disk.  \\item Ordered magnetic fields could arise in protostellar disks on account of ({\\it i}) advection of interstellar field by the accretion flow, ({\\it ii}) dynamo action in the disk, and ({\\it iii}) interaction with the stellar magnetic field.  \\item Semianalytic MHD models have been able to account for the basic structure of diffusive disks that drive CDWs from their surfaces as well as for the formation of such systems in the collapse of rotating molecular cloud cores. Some of these models already incorporate a realistic disk ionization and conductivity structure. These studies have established that vertical angular momentum transport by a CDW or through torsional Alfv\\'en waves (magnetic braking) could in principle be the main angular momentum removal mechanism in protostellar disks and determined the parameter regime where wind transport can be expected to dominate radial transport by MRI-induced turbulence. Further progress is being made by increasingly more elaborate numerical simulations (involving nonideal MHD codes) that have started to examine the global properties, time evolution, and dynamical stability of the magnetic disk/wind system.  \\item Robust observational evidence also exists for a magnetic interaction between CTTS disks and their respective protostars, including strong indications of a field-channeled flow onto the stellar surface. This interaction is likely to involve mass ejection and is thought to be responsible for the comparatively low rotation rates of CTTSs.  Since the magnetic field geometry in the interaction region is evidently quite complex and the interaction is likely time dependent, numerical simulations are an indispensable tool in the study of this problem.  \\end{itemize}  Future advances in this area will probably arise from a combination of new observational findings, the refinement of current theoretical approaches, and the incorporation of additional physics into the models. On the observational side, the main challenge is still to demonstrate the existence of CDWs in protostars and to determine their spatial extent (spread out over most of the disk surface or occurring only near its inner edge, and, if the former, is the launching region nearly continuous or is it confined to localized patches?). Recent attempts to measure rotation in the outflows could potentially help to answer this question, but, as noted in Sect.~\\ref{subsec:centrifugal}, the results obtained so far are still inconclusive.  Regarding the further development of theoretical tools, the greatest impact would likely be produced by numerical simulations that study vertical angular momentum transport by either a CDW or magnetic braking with codes that include a realistic conductivity tensor and that have full 3D and mesh-refinement capabilities. Such simulations should be able to clarify the relative roles of vertical and radial angular momentum transport and the possible interplay between them for relevant combinations of the disk model parameters (see Sect.~\\ref{subsubsec:exact_disk}). An interim step might be to solve for the evolution of a vertically integrated disk whose properties, at any radial grid zone, are determined from a vertical integration of a simplified version of the radially localized, steady-state disk model described in Sect.~\\ref{subsec:disk}. Time-dependent models of this type could examine the behavior of a magnetically threaded disk after its mass supply diminishes or stops altogether (corresponding to the protostellar system evolving into the optically revealed phase), which has previously been studied only in the context of ``$\\alpha$-viscosity'' models. As was already noted above, a state-of-the-art, 3D, nonideal-MHD code is also crucial for investigating the stability of disk/wind systems and the various aspects of the star--disk field-mediated interaction. One could, however, also benefit from further development of the semianalytic models, which could include an extension of the ionization/conductivity scheme, a derivation of self-similar disk/wind solutions that allow for a radial drift of the poloidal magnetic field (see Sect.~\\ref{subsubsec:exact_disk}), and a calculation of the predicted observational characteristics of wind-driving disks.  The mass fraction and size distribution of dust grains in the disk have a strong effect on its ionization and conductivity structure and on the degree of field--matter coupling (see Sect.~\\ref{subsubsec:ionize}). Existing models incorporate the effect of dust in a somewhat ad-hoc manner, by adopting an assumed distribution. In reality, the grain distribution is determined by the balance of several processes, including grain collisions due to relative velocities that develop as a result of Brownian motion, differential vertical-settling speeds, and turbulence, which can lead to either coagulation or fragmentation. Grains are also subject to a collisional drag force exerted by the gas and arising from the fact that the gas is subject to thermal and magnetic forces that do not affect the dust. This leads to vertical settling as well as to radial migration, directed either inward or outward depending on whether the gas rotation is sub- or super-Keplerian, respectively. Furthermore, radial or vertical gas motions can affect sufficiently small grains through advection. Yet another effect is evaporation by the ambient radiation field, which could impact dust located at sufficiently high elevations and small radii. Some of these effects have already been incorporated into generic viscous disk models (e.g., Dominik et al. 2007; Brauer et al. 2008), and one could similarly consider them in the context of a wind-driving disk model. In view of the fact that the latter model is characterized by a vertical outflow and by comparatively fast radial inflow speeds, one can expect to find new types of behavior in this case. In particular, grains located near the disk surfaces would either settle to the midplane if they are large enough or be uplifted from the disk if they are sufficiently small, whereas intermediate-size grains would first leave the disk and then re-enter at a potentially much larger radius, from which they could be advected back inward. By including dust dynamics, one could examine whether the effect of dust on the gas motion (through its influence on the field--matter coupling) and the effect of gas on the grain motions (through gas--dust collisions) together place meaningful constraints on the resulting grain distribution. One could also investigate whether the predicted radial transport of intermediate-size grains from small to large radii and their possible thermal processing outside the disk could be relevant to the accumulating evidence for an outward transport of crystalline grains in the protosolar nebula and in other protostellar disks, and whether the implied dust distribution in the disk and the wind might have distinct observational signatures that could be tested by spectral and imaging techniques (see Millan-Gabet et al. 2007, Sect.~\\ref{subsubsec:wind_observe}, and Chaps.~I and~VIII).  Dust particles are thought to be the building blocks of planetesimals and their distribution in the disk is thus a key ingredient of planet formation models. In fact, the general properties of a protostellar disk are evidently relevant to planet formation in light of the growing evidence that the latter is strongly influenced by physical processes that occur when the disk is still predominantly gaseous. A disk threaded by a large-scale, ordered magnetic field could potentially have unique effects on planet growth and migration. One such effect is the generation (through magnetic resonances that are the analogs of Lindblad resonances) of a global torque that may reduce, or even reverse, the secular inward drift (the so-called Type I migration) predicted for low-mass planets, which has posed a conundrum for current theories of planet formation. As was demonstrated by Terquem (2003) and Fromang et al. (2005), a torque of this type could be produced if the disk had a comparatively strong (MRI-stable, but still subthermal) azimuthal field (with a nonzero vertical average of $B_\\phi^2$) that fell off sufficiently fast with radius ($\\propto r^{-1}-r^{-2}$). A poloidal field component could in principle also contribute to this process (Muto et al. 2008). Given that a large-scale field with precisely these properties is expected in wind-driving protostellar disks (see Sects.~\\ref{sec:formation} and~\\ref{sec:vertical}), this possibility clearly merits an explicit investigation in the context of the disk models considered in this chapter. The influence of the vertical channel of angular momentum transport and of the overall effect of an ordered, large-scale field on the disk structure in such systems (e.g., the reduction of the density scale height by magnetic squeezing) may also be worth examining in this connection."
    },
    "1004/1004.2792_arXiv.txt": {
        "abstract": "{Classical novae  are powered by  thermonuclear runaways that  occur on the  white  dwarf component  of  close  binary  systems.  During these violent  stellar events,  whose  energy release  is  only exceeded  by gamma-ray bursts and supernova  explosions, about $10^{-4} - 10^{-5}\\, M_{\\sun}$ of material is ejected into the interstellar medium. Because of the high peak temperatures attained during the explosion, $T_{\\rm peak} \\sim (1 - 4) \\times  10^8$ K,  the  ejecta are  enriched in  nuclear-processed material relative to  solar abundances, containing significant amounts of $^{13}$C, $^{15}$N, and $^{17}$O  and traces of other isotopes. The origin  of these  metal enhancements  observed  in the  ejecta is  not well-known and has puzzled theoreticians for about 40 years. In this paper, we  present  new  2-D  simulations  of  mixing  at  the  core-envelope interface. We  show that Kelvin-Helmholtz  instabilities can naturally lead  to self-enrichment  of  the solar-like  accreted envelopes  with material from the outermost layers of the underlying white dwarf core, at levels that agree with observations.} ",
        "introduction": "The assumption of spherical symmetry  in classical nova models (and in general, in stellar explosions)  excludes an entire sequence of events associated  with  the way that a  thermonuclear  runaway (hereafter,  TNR) initiates (presumably  as a point-like ignition)  and propagates.  The first study of localized TNRs on white dwarfs was carried out by Shara (1982) on the  basis of semianalytical models. He  suggested that heat transport was too inefficient to  spread a localized TNR to the entire white  dwarf surface, concluding  that localized,  {\\it volcanic-like} TNRs were likely to occur.   But his analysis, based only on radiative and conductive transport, ignored  the major role played by convection on the lateral thermalization of a TNR.  The importance  of multidimensional effects  for TNRs in  thin stellar shells was revisited by Fryxell \\& Woosley (1982). In the framework of nova outbursts,  the authors concluded  that the most  likely scenario involves TNRs  propagated by small-scale  turbulence. On the  basis of dimensional  analysis  and  flame  theory,  the  authors  derived  the velocity  of  the deflagration  front  spreading  through the  stellar surface, in the form $v_{\\rm  def} \\sim (h_p v_{\\rm conv} / \\tau_{\\rm burn})^{1/2}$,  where  $h_p$ is  the  pressure  scale height,  $v_{\\rm conv}$ the  characteristic convective velocity,  and $\\tau_{\\rm burn}$ the characteristic timescale for fuel burning. Typical values for nova outbursts  yield $v_{\\rm  def} \\sim  10^4$ cm  s$^{-1}$ (that  is, the flame propagates halfway throughout the stellar surface in about $\\sim 1.3$ days). Shear-driven mixing induced by accretion of matter possessing angular momentum was also investigated by Kutter \\& Sparks (1987), but their numerical simulations failed to obtain a strong enough TNR to power a nova outburst (see Sparks \\& Kutter 1987).  The first  multidimensional hydrodynamic  calculations of this process were performed by Shankar,  Arnett \\& Fryxell  (1992) and Shankar \\&  Arnett (1994). They evolved an accreting,  $1.25 \\, M_{\\sun}$ white dwarf with a 1-D hydro code that was  mapped into a 2-D domain (a spherical-polar grid of 25$\\times$60 km).  The  explosive event was then followed with a 2-D  version of  the Eulerian code  {\\it PROMETHEUS}.   A 12-isotope network, ranging from  H to $^{17}$F, was included  to treat the energetics of  the explosion.   Unfortunately, the subsonic  nature of the problem, coupled with the use of an explicit code (with a timestep limited  by  the   Courant-Friedrichs-Levy  condition),  posed  severe limitations on the  study, which had to be  restricted to very extreme (rare) cases, characterized by huge temperature perturbations of about $\\sim  100  - 600$\\%,  in  small regions  at  the  base of  the envelope. The total computed time was only about 1 second.    The  calculations   revealed  that   instantaneous,  local temperature  fluctuations  cause  Rayleigh-Taylor  instabilities. Their rapid rise  and subsequent expansion (in a  dynamical timescale) cools the hot material  and halts  the lateral  spread of  the burning front, suggesting that such local temperature fluctuations are not important in the initiation or early stages of the TNR.   The study, therefore, favored  the local volcanic-like TNRs proposed by Shara (1982).  Glasner  \\&  Livne  (1995) and  Glasner,  Livne \\&  Truran (1997; GLT97) revisited  these early  attempts using  2-D simulations performed  with  the  code  {\\it VULCAN},  an  arbitrarily  Lagrangian Eulerian (ALE)  hydrocode capable of handling  both explicit and implicit steps. As in Shankar et  al. (1992), a slice of the star (0.1 $\\pi^{rad}$), in spherical-polar  coordinates with reflecting boundary conditions, was adopted.  The  resolution near the envelope base was around 5$\\times$5 km.  As before, the evolution of an accreting, $1 \\, M_{\\sun}$  CO white dwarf  was initially  followed using  a 1-D hydro code (to overcome  the early, computationally challenging phases of  the TNR),  and  then  mapped into  a  2-D domain  as  soon as  the temperature at  the envelope base  reached $T_{\\rm b} \\sim  10^8$ K. As in the previous works, the 2-D runs relied on a 12-isotope network. The  simulations showed  a  good agreement  with  the gross  picture described by 1-D models (for instance, the critical role played by the $\\beta^+$-unstable nuclei $^{13}$N,  $^{14,15}$O, and $^{17}$F, in the ejection  stage, and consequently,  the presence  of large  amounts of $^{13}$C,  $^{15}$N,  and  $^{17}$O  in the  ejecta).   However,  some remarkable  differences  were  also  identified.  The  TNR  was initiated  by  an ensemble of  irregular, localized  eruptions  at  the envelope base caused  by buoyancy-driven temperature fluctuations indicating that  combustion proceeds in a host of many localized flames  -- not  as   a  thin   front --  each   surviving  only   a few seconds. Nevertheless, these  authors concluded that turbulent diffusion efficiently dissipates any local burning around the core, so the  fast stages of  the TNR  cannot be  localized and  the runaway must spread through  the entire envelope. In contrast to 1-D models, the core-envelope  interface was convectively unstable, providing a source for  the metallicity enhancement of the envelope by means of a Kelvin-Helmholtz   instability  ---  resembling the convective overshooting  proposed  by  Woosley (1986).  Efficient  dredge-up of  CO material  from  the outermost white dwarf layers accounts for  $\\sim 30$\\% metal enrichment of the envelope (the accreted envelope  was assumed to be solar-like, without any   pre-enrichment),  in  agreement   with  the  inferred metallicites  in  the  ejecta   from  CO  novae (Gehrz et al. 1998). Finally, larger convective eddies were observed,  extending up to  2/3 of the envelope height with typical velocities  $v_{\\rm conv} \\sim 10^7$  cm s$^{-1}$.  Despite these differences, however, the expansion and  progress  of  the  TNR  towards the  outer  envelope  quickly became almost spherically symmetric, although the initial burning process was not.  The results of  another set of 2-D simulations were published shortly afterward by Kercek,  Hillebrandt \\&  Truran (1998; KHT98), which  aimed to confirm the general behaviors  reported by GLT97, in  this case with a version  of the  Eulerian {\\it  PROMETHEUS} code.  A similar domain (a  box of  about 1800$\\times$1100 km)  was adopted,  but using a cartesian,  plane-parallel  geometry to  allow  the  use of  periodic boundary  conditions.  Two  resolution simulations were performed,  one with a coarser 5$\\times$5 km  grid as in GLT97, and a second with a  finer 1$\\times$1 km  grid.  The calculations used the  same  initial  model  as  GLT97,  and  produced qualitatively  similar but  somewhat  less violent  outbursts. In particular, they obtained longer TNRs  with lower  $T_{\\rm peak}$  and  $v_{\\rm ejec}$, caused by  large differences in the convective  flow patterns.  Whereas GLT97 found that a few, large convective eddies dominated the flow, most  of the early  TNR was now  governed by small,  very stable eddies  (with $l_{\\max} \\sim$  200  km),  which  led to  more  limited dredge-up   and   mixing  episodes.   The   authors  attributed  these discrepancies to  the different  geometry and, more  significantly, to the boundary conditions adopted in both simulations.  The  only 3-D  nova  simulation to  date  was performed  by Kercek,  Hillebrandt  \\&  Truran  (1999),   adopting  a computational  domain  of   1800$\\times$1800$\\times$1000  km  with  a resolution   of  8$\\times$8$\\times$8   km.  It  produced flow patterns that were dramatically different  from those found in the  2-D simulations (much more erratic in the 3-D  case), including mixing by turbulent motions occurring on very small scales (not fully resolved with the adopted resolution) and peak  temperatures being achieved that were slightly lower  than in  the 2-D case (a consequence  of the slower and more  limited dredge-up of core material). The envelope  attained a maximum  velocity that was  a  factor  $\\sim  100$  smaller than  the  escape  velocity  and, presumably, no mass ejection  (except for a possible wind mass-loss phase). In view of these results, the authors concluded that CO mixing must take place prior  to the TNR, in contrast to the main results  of  GLT97\\footnote{Other multidimensional  studies (Rosner  et al.  2001, Alexakis et  al. 2004a,b) focused on  the role of  shear instabilities in the stratified  fluids that form  nova envelopes. They  concluded that mixing can  result from  the resonant interaction  between large-scale shear flows  in the accreted envelope  and gravity waves at the interface between the envelope and the underlying white dwarf. However, to account  for significant mixing, a  very high shear (with a specific velocity profile) had to be assumed.}.  In summary,  two independent studies,   GLT97 and KHT98, based upon the same  1-D initial model, reached nearly opposite  conclusions about  the strength of  the runaway  and its capability to power  a fast nova. The origin  of these differences was carefully analyzed  by Glasner, Livne \\& Truran  (2005), who concluded that the early stages of the explosion, prior to the onset of the TNR -- when the evolution is almost quasi-static --  are extremely sensitive to the outer boundary conditions (see e.g.,  Glasner, Livne \\& Truran (2007), for a  2-D nova simulation  initiated when the temperature  at the envelope  base is  only $5  \\times  10^7$ K).   Several outer  boundary conditions   were  examined.  The   study  showed   that  Lagrangian simulations, where the envelope is allowed to expand and mass is conserved, are consistent  with spherically symmetric solutions. In contrast, in Eulerian schemes with a ``free outflow''  outer boundary condition --- the choice adopted in KHT98 ---, the outburst can be artificially quenched.  In light of these conundrums, a  reanalysis of  the role  of late mixing  at the  core-envelope  interface during  nova outbursts  seems mandatory. To this end, we  performed an independent  2-D simulation, identical to GLT97 and KHT98, with  another multidimensional   hydrodynamic  code to  investigate whether mixing can occur in  an Eulerian framework with an appropriate choice of the  outer boundary conditions. ",
        "conclusions": "We have analyzed the possible self-enrichment  of the solar-composition accreted  envelope with material from  the underlying white dwarf during nova outbursts in a multidimensional framework.  We have found that a shear flow at the  core-envelope interface (which unlike the spherically symmetric case, does not behave like a rigid wall) drives mixing through KH instabilities.  Large convective eddies develop close to the core-envelope interface, of a size comparable to the height  of the envelope (similar to the pressure scale height in 1-D simulations), mixing CO-rich material  from  the outermost  layers  of  the  underlying white dwarf into  the  accreted envelope.  The  metallicity enrichment achieved in  the envelope,  $Z \\sim 0.30$,  is in agreement  with observations  of CO  nova  ejecta.  Our 2-D simulations also show that even for a point-like TNR ignition,  the expansion and progress of the runaway is  almost spherically symmetric  for nova conditions.  We note that  the  adopted  resolution as  well  as  the  size, intensity, and location of the initial perturbation have a very limited impact  on the  results, principally affecting the  timescale for the onset of  the KH instability but not the final, mean  metallicity. Details will be extensively  discussed in a forthcoming publication.   Our results agree  with earlier 2-D  hydrodynamic simulations  (GLT97)  and  solve the controversy raised  by another  2-D study (KHT98) that questioned  the efficiency  of this  mixing mechanism,  and  hence the corresponding strength  of the runaway  and its capability to  power a fast nova outburst."
    },
    "1004/1004.0274_arXiv.txt": {
        "abstract": "{Ultra-high energy (UHE) neutrinos and cosmic rays initiate particle cascades underneath the Moon's surface. These cascades have a negative charge excess and radiate Cherenkov radio emission in a process known as the Askaryan effect. The optimal frequency window for observation of these pulses with radio telescopes on the Earth is around 150 MHz. }{By observing the Moon with the Westerbork Synthesis Radio Telescope array we are able to set a new limit on the UHE neutrino flux. }{The PuMa II backend is used to monitor the Moon in 4 frequency bands between 113 and 175 MHz with a sampling frequency of 40 MHz. The narrowband radio interference is digitally filtered out and the dispersive effect of the Earth's ionosphere is compensated for. A trigger system is implemented to search for short pulses. By inserting simulated pulses in the raw data, the detection efficiency for pulses of various strength is calculated. }{With 47.6 hours of observation time, we are able to set a limit on the UHE neutrino flux. This new limit is an order of magnitude lower than existing limits. In the near future, the digital radio array LOFAR will be used to achieve an even lower limit. } {} ",
        "introduction": "The cosmic ray energy spectrum follows a power law distribution extending up to extremely large energies. At the Pierre Auger Observatory (PAO) cosmic rays (CRs) are observed up to energies around $\\sim 10^{20}$~eV. Above the Greisen-Zatsepin-Kuzmin (GZK) energy of $6\\cdot 10^{19}$~eV, CRs can interact with the cosmic microwave background photons. In the most efficient interaction, a $\\Delta$-resonance is produced which decays into either a proton and a neutral pion or a neutron and a positively charged pion. Charged pions decay and produce neutrinos. The energy loss length for $\\Delta$-resonance production is $\\sim 50$\\, Mpc \\citep{G66,ZK66}.  Recent results of the PAO have confirmed a steepening in the cosmic ray spectrum at the GZK energy \\citep{A08}. This steepening is not necessarily a clear cut-off, as CRs from local sources may arrive at Earth with super-GZK energies. Because of their large energies, these particles will only deflect slightly in the (extra-) Galactic field during their propagation, and their arrival directions correlate with their sources \\citep{A07}.  Sources at distances larger than 50 Mpc can be found by observing neutrinos that are produced in GZK interactions. Since neutrinos are chargeless they will propagate in a straight line from the location where the GZK interaction took place to the observer, thus conserving the directional information. In addition, while CRs from distant sources pile up at the GZK energy, information about the CR spectrum at the source is conserved in the GZK neutrino flux. Other possible sources of UHE neutrinos are decaying supermassive dark matter particles or topological defects. This class of models is refered to as top-down models (see for example \\citet{s04} for a review).  Because of their small interaction cross section and low flux, the detection of cosmic neutrinos calls for extremely large detectors. Assuming the Waxman-Bahcall flux \\citep{wb98,wb01}, even at low energies in the GeV range, the flux is not higher than a few tens of neutrinos per km$^2$ per year. Kilometer-scale detectors are not easily built but can be found in nature. For example, interaction of neutrinos in ice or water can be detected by the Cherenkov light produced by the lepton track or cascade. The nearly completed IceCube detector \\citep{icecube} will cover a km$^3$ volume of South Pole ice with optical modules, while Antares \\citep{ antares} and its successor KM3NET \\citep{KM3NET} exploit the same technique in the Mediterranean sea. Even larger volumes can be covered by observing large detector masses from a distance. The ANITA balloon mission \\citep{anita} monitors an area of a million km$^2$ of South Pole ice from an altitude of $\\sim 37$~km and the FORTE satellite \\citep{forte} can pick up radio signals coming from the Greenland ice mass. Alternatively, cosmic ray experiments like the Pierre Auger Observatory can possibly distinguish cosmic ray induced air showers from neutrino induced cascades at very high zenith angles where the atmosphere is thickest and only neutrinos can interact close to the detector.  The Moon offers an even larger natural detector volume. When CRs or neutrinos hit the Moon they will interact with the medium. CRs will start a particle cascade just below the Lunar surface, while neutrinos will interact deeper inside the Moon, also creating a hadronic shower. The negative charge excess of a particle cascade inside a dense medium will cause the emission of coherent Cherenkov radiation in a process known as the Askaryan effect \\citep{a62}. This emission mechanism has been experimentally verified at accelerators \\citep{s01,g00} and extensive calculations have been performed to quantify the effect \\citep{zhs92,az97}. The idea to observe this type of emission from the Moon with radio telescopes was first proposed by \\citet{dz89} and the first experimental endeavours in this direction were carried out with the Parkes telescope \\citep{parkes}, at Goldstone (GLUE) \\citep{glue}, and with the Kalyazin Radio Telescope \\citep{KALYAZIN}. LUNASKA (Lunar UHE Neutrino Astrphysics with the Square Kilometer Array) is a project that is currently performing lunar Cherenkov measurements with ATCA (the Australia Telescope Compact Array) with a 600 MHz bandwidth at 1.2-1.8 GHz \\citep{LUNASKA}.  \\citet{fg03} suggested to use low-frequency telescopes (like LOFAR) for such an experiment. It is shown by \\citet{scholten} that observing at lower frequencies has the distinct advantage that radio pulses have a much higher chance of reaching the observer, as will be explained in the next section. In this work we use data recorded with the Westerbork Synthesis Radio Telescope (WSRT) in the frequency range of 113-168~MHz to set a new limit on the flux of UHE neutrinos. A first reporting of this limit is made in \\citep{NuMoon-PRL}. ",
        "conclusions": "\\label{sec:peak} We have investigated the nine strongest pulses that survive the applied cuts. Figures \\ref{bigpulse1} and \\ref{bigpulse2} are typical examples of the time traces of such pulses. The pulses are from different observation runs and their $P5$ values are plotted as a function of bin number (bin size is 25 ns) for all frequency bands and both beams. At this stage the RFI has already been mitigated and the signal has been de-dispersed.  \\begin{figure} \\centering \\includegraphics[width=0.45\\linewidth]{14104fg16.eps} \\includegraphics[width=0.45\\linewidth]{14104fg17.eps} \\caption{Time traces of the power (polarizations added) for a typical large pulse seen in the data after applying the cuts. The P5 values are plotted for all bands (top: highest frequency, bottom: lowest frequency) and both beams (left and right). The horizontal axis displays bin number (bin size is 25 ns). The power on the vertical axis is expressed in mean P5 value, with the trigger level at 5 for all bands. A trigger was only found for the right beam.} \\label{bigpulse1} \\end{figure}  \\begin{figure} \\centering \\includegraphics[width=0.45\\linewidth]{14104fg18.eps} \\includegraphics[width=0.45\\linewidth]{14104fg19.eps} \\caption{Like Fig.~\\ref{bigpulse1}, different pulse. The similar features in both beams exclude the pulse as a proper lunar pulse candidate.} \\label{bigpulse2} \\end{figure}  For both example events, the trigger was found in the right-hand beam. For the event in Figure \\ref{bigpulse1} the maximum $P5$ value increases with decreasing frequencies. This could be due to a stronger signal at lower frequencies or an increase of pulsed background at higher frequencies (remember that the $P5$ value is normalized over a 500~$\\mu$s time trace for each individual band). Although the signal is clearly much smaller in the left beam, it should be noted that three out of four bands actually have a pulse that exceeds trigger level ($P5 > 5$). The event displayed in Fig.~\\ref{bigpulse2} has a strong signal in both beams and the only reason this event was not discarded by the anti-coincidence criterion is that the highest frequency band has a very small signal-to-noise ratio. In this band, the signal is suppressed and happens to be just above threshold in the right-hand beam but below threshold in the left-hand beam. This way, strong temporary increases in background radiation are responsible for several of the largest events that pass our criteria.  Although the event in Figure \\ref{bigpulse1} has a curious dependence on frequency, it has the properties of a proper lunar pulse of the type we are looking for, which are: i) present in all frequency bands, ii) strong polarization, iii) short after dispersion correction, and iv) present in one beam only. In order to study the possibility of the pulse to originate from the Moon we can impose an additional condition that the Faraday rotation of the polarization is of the correct magnitude.  The Faraday rotation going through a plasma with STEC=5 in the Earth magnetic field is about $\\pi/4$ radian over 30 MHz (corresponding to the difference in the centroids of bands 1 and 3) at the frequencies of interest for the present study. For a pulse fully polarized in the x-direction in the center of band 1 one would thus expect about equal strength in the x and y polarization for band 3. We have examined whether or not the ratio between the pulse strengths in the x and y polarization in the different frequency bands corresponds to the ratios expected on the basis of the STEC value. This criterion disqualifies the pulse in Figure \\ref{bigpulse1} as originating from outside the ionosphere.  We have studied the nine strongest pulses with $S>62$ and found that all of them are unlikely to come from the Moon, because they either have a strong signal in both beams or do not have the frequency dependent behavior expected from Faraday rotation. As a result we can safely state that we see no pulses originating from a particle cascade in the Moon with a strength larger than $S=62$. Because this analysis is done a posteriori, the threshold used for the determination of the neutrino flux limit is kept at $S=77$. In future studies, additional cut criteria based on temporary power surges in the background and Faraday rotation of the signal in the ionosphere can be implemented to further understand and reduce the background."
    },
    "1004/1004.3511_arXiv.txt": {
        "abstract": "We study the effect of filter zero-point uncertainties on future supernova dark energy missions. Fitting for calibration parameters using simultaneous analysis of all Type Ia supernova standard candles achieves a significant improvement over more traditional fit methods. This conclusion is robust under diverse experimental configurations (number of observed supernovae, maximum survey redshift, inclusion of additional systematics). This approach to supernova fitting considerably eases otherwise stringent mission calibration requirements. As an example we simulate a space-based mission based on the proposed JDEM satellite; however the method and conclusions are general and valid for any future supernova dark energy mission, ground or space-based. ",
        "introduction": "\\label{sec:intro} The discovery of the acceleration of the expansion of the universe \\citep{riess98,perlmutter99} ranks as one of the most significant recent discoveries in cosmology. This acceleration is usually ascribed to a mysterious ``dark energy'' about which almost nothing is known although there are many competing ideas; what is needed to distinguish between them and shed more light on the nature of the acceleration is more and improved data. Observations of type Ia supernovae (SNe Ia) have allowed the discovery of the acceleration of the expansion \\citep{riess98,perlmutter99} and are currently the most established and best \\citep{albrecht06}. The method is described by many authors \\citep{perlmutter97,riess98,perlmutter99,perlmutter03} and is based on the fact that SNe Ia are, to good accuracy, standardizable candles (for a review of SNe Ia as standardizable candles see \\citet{phillips03,branch92}). However current supernova observations are limited by systematic uncertainties; while this was not a problem when the supernova sample was small and statistical uncertainties were the dominant ones, the growing sample size has already reached the point when statistical and systematic uncertainties are of comparable magnitude, as in e.g. the combined sample of $557$ supernovae studied by \\citet[The Union2 compilation:][]{amanullah10}. As more supernovae will be discovered in the future the need to better characterize and reduce systematic uncertainties will become \\emph{the} dominant concern in dark energy experiments. This has been recognized for several years, and the SuperNova/Acceleration Probe (SNAP) satellite\\footnote{\\url{http://www.snap.lbl.gov/}}\\citep[SNAP Collaboration:][]{aldering04}  was proposed as a systematics-controlled space-based experiment that would put much tighter constraints on dark energy than current and near future experiments by following $\\approx 2000$ supernovae out to $z_\\mathrm{max}\\approx 1.7$. More recently NASA and the Department of Energy have announced the Joint Dark Energy Mission (JDEM) \\footnote{\\url{http://jdem.gsfc.nasa.gov/}}$^,$\\footnote{\\url{http://jdem.lbl.gov/}} as a future space-based mission to study the nature of dark energy by employing a combination of techniques including supernovae. Therefore it is important to characterize the sources of systematics of future supernova experiments; studies of this kind have already appeared \\citep{kim04, kim06, nordin08, linder09} and this paper aims at building upon and expanding them. Studies using real data-sets have also appeared: for example \\citet{kilbinger09} use the SNLS supernovae in \\citep{astier06} to evaluate the effect of zero-point uncertainties on the final cosmology. \\par The two most important sources of systematic uncertainty in dark energy experiments that use supernovae are the dimming by dust in the host galaxy and uncertainties in the flux calibration, specifically the filter zero-points, as seen in recent cosmological analyses such as \\citet{astier06}, \\citet{woodvasey07}, \\citet{kowalski08}, \\citet{hicken09}, \\citet{amanullah10}. The problem of host-galaxy dimming is also being aggressively pursued, by e.g. targeting supernovae in rich clusters of galaxies \\citep{dawson09}; we will include it statistically in our analysis but will not go into its systematics. Properly taking into account zero-point uncertainties is nontrivial because their causes are numerous, ranging from any inaccuracy in the response function of telescope, filter, or detector (from now on collectively indicated as ``channel''), or the atmosphere for ground-based experiments, to uncertainties in the calibration procedure. While accurately characterizing all these is obviously an experiment-dependent problem, our aim is to provide a more general way to deal with them. \\par Starting with a simple model of zero-point uncertainty, we perform a complete end-to-end simulation of a supernova dark energy mission, propagating zero-point uncertainties through the simulation chain, and we evaluate its effects on the final cosmology fit. We do not aim at a detailed physical modelling of particular causes of uncertainty such as imperfect knowledge of the standard stars used to calibrate the zero-points or of the filter response functions, but rather at characterizing their overall effect, whatever their underlying reasons, with a set of zero-points representing the contribution of these important sources of systematics to the final error budget. This will serve as a guide to designers of how much specific components (telescope, filters, detectors, calibration procedure and so on) could be imperfectly known and still achieve the mission objectives. \\par Our starting point is the result by \\citet[][hereafter KM]{kim06}. KM introduce a new model of filter zero-point uncertainty and show that, due to the standardizable candle nature of SNe Ia, it is possible to treat these uncertainties as parameters that can be included with other parameters in a cosmology analysis. More precisely, KM model the observed peak magnitude $m$ of a supernova as $m = \\mu + M + \\mathit{Other} + \\mathcal{Z}$ where $\\mu$ is the distance modulus, $M$ is the absolute magnitude after standardization (and therefore the same for all supernovae), ``$\\mathit{Other}$'' indicates all residual effects that influence $m$, such as host galaxy extinction, and $\\mathcal{Z}$ is a new fit parameter for zero-point offsets to be fit with all the other parameters in the model. It is important to note that modelling zero-point offsets as fit parameters would not be possible if SNe Ia were not standardizable candles because $M$ would not be the same for every supernova. KM show that by fitting for all the supernovae distance moduli \\emph{simultaneously} it is possible to achieve a significant reduction in the final uncertainties in the cosmological parameters with respect to the traditional case when supernova distances are fit one by one and calibration uncertainties are then included in the total error budget. In the rest of the paper we will refer to the KM fitting approach as ``simultaneous fit'' and to the traditional approach as ``SN by SN fit''. \\par This work expands KM in several ways: \\begin{enumerate} \\item KM carry on a Fisher matrix analysis of their model; we perform a complete end-to-end simulation of a supernova dark energy mission, with a realistic modelling of all its aspects. \\item KM use a particular $z$ distribution, in which all supernovae are placed at those special redshifts that have zero $K$-correction. At those $z$ the improvement in mission performance is maximum; we study a more realistic $z$ distribution. \\item We include an intrinsic color dispersion. \\item We investigate whether our conclusions are robust with respect to several changes in mission parameters (number of supernovae, maximum redshift, inclusion of additional systematics); our simulation tool allows us to explore a much wider parameter space than KM. \\end{enumerate} It is important to point out that the KM model that we adopt here is applicable to a generic future dark energy mission based on supernovae; however, for concreteness we will present our results by considering a specific example: the supernova survey of the future space-based dark energy mission based on the proposed SNAP satellite. We will also assume that a nearby sample of supernovae, whose characteristics are based on the expected Nearby Supernova Factory sample\\footnote{\\url{http://snfactory.lbl.gov/}} \\citep{aldering02,copin06}, is available: specifically this sample is comprised of $316$ supernovae with $0.03\\leq z\\leq 0.08$. \\par The paper is organized as follows: Section \\ref{sec:snmodel} describes the KM model and its implementation in our simulation tool, Sections \\ref{sec:results} and \\ref{sec:parspace} describe our results and Section \\ref{sec:conclusion} summarizes our conclusions and discusses ways the work can be expanded. In what follows we will use the terms ``mission'' and ``experiment'' interchangeably. ",
        "conclusions": "\\label{sec:conclusion} Adopting the general method of modelling zero-point uncertainties introduced by KM we have carried out simulations of a future space-based supernova dark energy experiment with the main aim of assessing the influence of zero-point uncertainties on its overall performance. We have confirmed KM results for a more realistic experiment: fitting for all supernovae at once results in a greatly improved mission performance over the traditional SN by SN fitting. Whereas this effect may not be evident in today's surveys involving a few hundreds of supernovae and few available bands, it will become very significant for future surveys. We explored a representative section of the mission parameter space paying particular attention to how zero-point requirements can be traded off with other mission parameters; in particular we have shown that in general going to higher redshift results in less stringent zero-point requirements, even without considering other form of systematic. We stress once again that, while our results are for a specific possible space-based mission, the KM model itself is more general. The inclusion of a redshift dependent systematic such as the LH systematic greatly affects the mission performance, both by significantly degrading the FoM and by making the case for higher redshift even stronger; it is therefore extremely important to better characterize other forms of systematic by the time future stage IV experiments get under way. Finally the tools used here can realistically simulate future dark energy supernova experiments. The work can be expanded in many ways, all easily implementable in our simulation tool. The most obvious examples are different mission architectures, both ground and space-based, different redshift distributions, further models of systematics. For the zero-point uncertainties one can explore tighter characterization in the optical vs the near infrared and variation with time. The latter may be especially relevant for ground-based surveys. A more detailed treatment of the nearby supernova sample would introduce a separate set of zero-point parameters; again this can be accommodated by the KM model."
    },
    "1004/1004.1514_arXiv.txt": {
        "abstract": "We report the discovery of a transiting planet orbiting the star TYC~6446-326-1. The star, WASP-22, is a moderately bright (V=12.0) solar-type star ($\\teff=6000\\pm 100$\\,K, [Fe/H]$ = -0.05\\pm 0.08$). The lightcurve of the star obtained with the WASP-South instrument shows periodic transit-like features with a depth of about 1\\% and a duration of 0.14\\,d. The presence of a transit-like feature in the lightcurve is confirmed using z-band photometry obtained with  Faulkes Telescope South. High resolution spectroscopy obtained with the CORALIE  and HARPS spectrographs confirm the presence of a planetary mass companion with an orbital period of 3.533\\,days in a near-circular orbit. From a combined analysis of the spectroscopic and photometric data assuming that the star is a typical main-sequence star we estimate that the planet has a mass $M_{\\rm p} = 0.56\\pm 0.02M_{\\rm Jup}$ and a radius $R_{\\rm p} = 1.12 \\pm 0.04R_{\\rm Jup}$. In addition, there is a linear trend of 40\\,m\\,s$^{-1}$\\,y$^{-1}$ in the radial velocities measured over 16 months, from which we infer the presence of a third body with a long period orbit in this system. The companion may be a low mass M-dwarf, a white dwarf or a second planet. ",
        "introduction": "The WASP project \\citep{2006PASP..118.1407P} is currently one of the most successful wide-area surveys designed to find exoplanets transiting bright stars (V $<$ 12.5).  Other succesful surveys include HATnet \\citep{2004PASP..116..266B}, XO \\citep{2005PASP..117..783M} and TrES \\citep{2006AAS...20922602O}. There is continued interest in finding transiting exoplanets because they can be accurately characterized and studied in some detail, e.g., the mass and radius of the planet can be accurately measured. This gives us the opportunity to explore the relationships between the density of the planet and other properties of the planetary system, e.g., the semi-major axis, the spectral type of the star, the eccentricity of the orbit, etc. Given the wide variety of transiting planets being discovered and the large number of parameters that characterize them, statistical studies will require a large sample of systems to identify and quantify the relationships between these parameters. These relationships can be used to test models of the formation, structure and evolution of short period exoplanets.  A particular puzzle related to the properties of hot Jupiters is the wide range in their densities. Very dense hot Jupiters such as HD~149026 are thought to contain a dense, metallic core \\citep{2005ApJ...633..465S}. There is currently no generally agreed explanation for the existence of hot Jupiters with densities 5\\,--\\,10 times lower than the density of Jupiter, e.g. WASP-17\\,b \\citep{2010ApJ...709..159A}, TrES-4\\,b \\citep{2007ApJ...667L.195M} and WASP-12\\,b \\citep{2009ApJ...693.1920H}. One possibility is that the planets are heated by tidal forces, and that these are driven by the presence of a third body in the system \\citep{2007MNRAS.382.1768M}. Other possibilities include  enhanced opacity in the atmosphere \\citep{2007ApJ...661..502B}, the distribution of heavy elements in the core \\citep{2008A&A...482..315B} and kinetic heating from the irradiated atmosphere into the interior \\citep{2002A&A...385..166S}.  Here we report the discovery of a hot Jupiter system, WASP-22, identified using the WASP-South instrument and present evidence that it is a member of a hierarchical triple system.   \\begin{table} \\caption{Radial velocity measurements.} \\label{rv-data} \\begin{tabular*}{0.5\\textwidth}{@{\\extracolsep{\\fill}}lrrr} \\tableline \\noalign{\\smallskip} BJD--2\\,400\\,000 &\\multicolumn{1}{l}{RV} & \\multicolumn{1}{l}{$\\sigma_{\\rm RV}$} & \\multicolumn{1}{l}{BS}\\\\ & (km s$^{-1}$) & (km s$^{-1}$) & (km s$^{-1}$)\\\\ \\noalign{\\smallskip} \\tableline \\noalign{\\smallskip} CORALIE&&&\\\\ 54704.8563 &$ -7.236 $& 0.013 & $0.005$\\\\ 54706.8761 &$ -7.368 $& 0.026 & $0.030$\\\\ 54708.8417 &$ -7.189 $& 0.011 & $-0.020$\\\\ 54709.7625 &$ -7.320 $& 0.019 & $-0.038$\\\\ 54710.7697 &$ -7.291 $& 0.024 & $0.046$\\\\ 54715.8843 &$ -7.201 $& 0.018 & $0.058$\\\\ 54716.8457 &$ -7.262 $& 0.041 & $-0.127$\\\\ 54717.7794 &$ -7.386 $& 0.041 & $0.034$\\\\ 54720.7520 &$ -7.312 $& 0.017 & $0.002$\\\\ 54721.8184 &$ -7.289 $& 0.019 & $0.028$\\\\ 54722.8233 &$ -7.201 $& 0.012 & $-0.033$\\\\ 54724.7367 &$ -7.314 $& 0.017 & $-0.011$\\\\ 54726.8108 &$ -7.225 $& 0.012 & $0.010$\\\\ 54729.8623 &$ -7.206 $& 0.013 & $-0.005$\\\\ 54731.8658 &$ -7.345 $& 0.012 & $-0.015$\\\\ 54740.7564 &$ -7.199 $& 0.020 & $-0.058$\\\\ 54834.5647 &$ -7.304 $& 0.018 & $0.014$\\\\ 54836.5707 &$ -7.225 $& 0.017 & $-0.019$\\\\ 54853.6313 &$ -7.171 $& 0.014 & $0.046$\\\\ 54854.6322 &$ -7.275 $& 0.013 & $0.018$\\\\ 54855.6091 &$ -7.310 $& 0.012 & $0.001$\\\\ 54860.5404 &$ -7.183 $& 0.013 & $0.051$\\\\ 54862.5589 &$ -7.327 $& 0.015 & $-0.046$\\\\ 54865.5628 &$ -7.313 $& 0.014 & $0.024$\\\\ 54879.5256 &$ -7.280 $& 0.017 & $-0.039$\\\\ 54880.5402 &$ -7.273 $& 0.021 & $0.033$\\\\ 54882.5242 &$ -7.193 $& 0.021 & $0.043$\\\\ 54885.5207 &$ -7.159 $& 0.016 & $-0.014$\\\\ 54886.5443 &$ -7.299 $& 0.017 & $-0.015$\\\\ 55095.8758 &$ -7.308 $& 0.013 & $0.013$\\\\ 55100.8593 &$ -7.154 $& 0.015 & $-0.002$\\\\ 55126.7935 &$ -7.266 $& 0.015 & $0.005$\\\\ 55128.7415 &$ -7.169 $& 0.016 & $0.006$\\\\ 55185.6618 &$ -7.147 $& 0.011 & $-0.007$\\\\ 55186.6153 &$ -7.234 $& 0.011 & $-0.025$\\\\ 55189.6571 &$ -7.164 $& 0.010 & $0.004$\\\\ 55190.6690 &$ -7.284 $& 0.011 & $-0.008$\\\\ HARPS&&&\\\\ 54743.7527 & $-7.1832 $& 0.0029 &$ 0.0090$\\\\ 54746.7683 & $-7.2391 $& 0.0026 &$ 0.0040$\\\\ 54749.7650 & $-7.2941 $& 0.0024 &$ 0.0039$\\\\ 54750.6791 & $-7.1940 $& 0.0029 &$ 0.0109$\\\\ 54754.7257 & $-7.1752 $& 0.0025 &$ 0.0118$\\\\ 54755.7558 & $-7.2726 $& 0.0032 &$ 0.0329$\\\\ \\noalign{\\smallskip} \\tableline \\end{tabular*} \\end{table}     \\begin{figure} \\plotone{wasplc.eps} \\caption{WASP-South photometry of WASP-22 folded on the orbital period $P$~=~3.532759\\,d. Upper panel: all data. Lower panel: data within 0.12 phase units of mid-transit together with the model fit described in Section~\\ref{paramsec} (solid line). \\label{phot-full-and-zoom} } \\end{figure} \\begin{figure} \\plotone{ftslc.eps} \\caption{Faulkes Telescope South z-band photometry of WASP-22 (points) with the model fit described in Section~\\ref{paramsec} (solid line). \\label{ftslc} } \\end{figure} ",
        "conclusions": "The star WASP-22 (TYC~6446-326-1) has a hot Jupiter companion.  A long-term linear trend in the mean value of the radial velocity shows that WASP-22 has a distant companion, i.e., it is a hierarchical triple system. The properties of the third body are poorly constrained by the data available to-date, but it may be an M-dwarf, a white dwarf or a second planet."
    },
    "1004/1004.1178_arXiv.txt": {
        "abstract": "We have carried out a redshift survey using the VIMOS spectrograph on the VLT towards the Cosmic Microwave Background cold spot. A possible cause of the cold spot is the Integrated Sachs-Wolfe effect imprinted by an extremely large void (hundreds of Mpc in linear dimension) at intermediate or low redshifts. The redshift distribution of over seven hundred $z<1$ emission-line galaxies drawn from an $I-$band flux limited sample of galaxies in the direction of the cold spot shows no evidence of a gap on scales of $\\Delta z\\gtsim 0.05$ as would be expected if such a void existed at $0.35<z<1$. There are troughs in the redshift distribution on smaller scales ($\\Delta z\\approx 0.01$) indicating that smaller scale voids may connect regions separated by several degrees towards the cold spot. A comparison of this distribution with that generated from similarly-sized subsamples drawn from widely-spaced pointings of the VVDS survey does not indicate that the redshift distribution towards the cold spot is anomalous or that these small gaps can be uniquely attributed to real voids. ",
        "introduction": "\\label{sec:intro}  A major triumph of modern physics is the measurement and explanation of the Cosmic Microwave Background (CMB) fluctuation spectrum ({\\it e.g.} Bennett et~al., 2003).  Encoded within the pattern of fluctuations from constant temperature on the sky are the signatures of the makeup of the Universe. The amount of baryons, matter, dark energy within the Universe, when and how reionization occurred, and the power spectrum of initial fluctuations can all be inferred either from the properties of these fluctuations alone or in combination with other measurements.   \\begin{figure} \\includegraphics[width=8.5cm]{pointings-sub.eps} \\caption{The six spectroscopic field centres overlayed on the WMAP ILC temperature map of the cold spot region. The pointing centres are: (1) 03:28:51 -21:03:54, (2) 03:10:00 -16:45:44, (3) 03:21:13 -18:43:15, (4) 03:15:57 -20:27:23, (5) 03:09:53 -21:57:24 and (6) 03:04:29 -23:17:10} \\label{fig:1} \\end{figure}    The ``concordance cosmology'', based strongly on $\\Lambda$CDM is extremely successful in explaining CMB anisotropies on small angular scales. However, there are clear anomalies on larger angular scales which pose potential problems for our understanding of the evolution of structure at the most fundamental levels. The most striking anomaly is the presence of an apparent cold spot in the Wilkinson Microwave Anisotropy Probe (WMAP) CMB data in the southern hemisphere (Vielva et al., 2004; Cruz et al., 2005; Spergel et al., 2007). This region of decreased temperature ($\\Delta T/T \\sim -10^{-5}$) extends over an angular scale of 5-10$^\\circ$ (Figure 1). The cold spot has a less than 2 per cent probability of being generated as part of a Gaussian fluctuation spectrum (Cruz et al., 2007a,).   \\begin{figure*}  \\includegraphics[width=18cm]{zdista.eps}  \\vspace{-1cm}  \\includegraphics[width=18cm]{zdistb.eps} \\vspace{-1cm} \\caption{The redshift distribution of the sources in the six fields. Each bin has a width of 0.01 in redshift and the distribution is displayed twice, the second time with the the binning offset by half a bin width. The expected signature of a self-compensating void large enough to produce the cold spot temperature decrement through the ISW effect is an absence of redshifts across several consecutive bins and sharp spikes in the distribution at either edge of the gap. No such feature is seen.} \\label{fig:2} \\end{figure*}   Assuming that the cold spot is not a statistical artifact, or produced as a result of the way the data has been analysed ({\\it e.g.}  Zhang \\& Huterer, 2009), determining the nature of the cold spot is crucial to our understanding of the growth of structure. Either it is intrinsic to the CMB itself or is imposed by a foreground structure along the line of sight to the CMB. Either way, there are important implications for our understanding of structure evolution. In the first instance, an intrinsic origin points to extra physics (such as textures or defects, Cruz et al., 2007b, 2008; Durrer 1999) or non-standard inflation ({\\it e.g.} Peebles 1997) influencing the behavior of the early Universe.   A Galactic origin for a foreground source of the cold spot appears to be ruled out by its spectrum (Cruz et al., 2006). If the cold spot is a secondary CMB anisotropy caused by modulation of the CMB signal as it passes through an intervening massive structure, there is only one possible cause.  An extremely large void of scale $\\sim 200 h^{-1}$Mpc could cause the signal through the Integrated Sachs Wolf (ISW) effect (Inoue \\& Silk 2006).  As the CMB photons pass through such a void, over their transit time the gravitational potential of the void evolves as the void expands leading to a net redshift of the photons and hence cooling of the spectrum at that point in the sky. The size of the void is set not just by the angular size of the cold spot, but also by the light travel time across the void necessary to induce an ISW signal of this magnitude. Any void must be at $z<1$ as at later times the increasing influence of the cosmological constant on the expansion of the Universe enhances the ISW signature of the void.   We have carried out a redshift survey of galaxies selected from six patches of sky separated by $\\sim 1-2^\\circ$ towards the CMB cold spot in order to search for the signature of an extremely large void capable of producing a large enough ISW signal to account for the CMB feature. Assuming galaxies in such a structure have bias factors similar to those found in typical voids, this structure should be largely empty of galaxies, particularly those around $M^*$ which dominate the galactic mass distribution. Given its extraordinary size (2-300 Mpc, of order 5-10 percent the distance back to $z\\sim 1$), it should leave a clear gap of $\\Delta z \\sim$ few $\\times 0.01$ in the redshift distribution towards the cold spot depending upon the exact size and redshift of the void. ",
        "conclusions": "The area of sky towards the cold spot has been examined photometrically in several wave-bands other than in the mm-wave in oder to search for the signature of a super-void. Rudnick et~al (2007) claimed that there is a 25-40 per cent deficit of NRAO VLA Sky Survey (NVSS) sources towards the cold spot and McEwen et~al. (2007) demonstrated that the spatial correlation between the NVSS number counts and fluctuations in the WMAP data are strongly influenced by the cold spot area. This is puzzling given the expected redshift distribution of NVSS sources ({\\it e.g.} Raccanelli et~al. 2008, Brookes et~al. 2008), which is broad and peaks at $z \\sim 1$. Any void at $z<1$ and with a size of $\\Delta z \\sim 0.1$, even if empty of radio sources and non-compensating for these objects, should depress the number counts towards the void by only a few per cent given the volume of the void and the effective volume probed by NVSS (see also Smith \\& Huterer 2008).  Conversely, a deep optical imaging survey carried out by Granett, Szapudi \\& Neyrinck (2009) showed no evidence for a super-void towards the cold spot. Using photometric redshifts derived from $g,r,i,z$ imaging of seven fields within the cold spot region (their pointings A, C and D cover our pointings 6, 2 and 4, and their F \\& G partially contain our 3 and 1), they claimed to rule out any 200 Mpc void between $0.5<z<0.9$. While this agrees with our spectroscopic measurements, we note that attempting to identify a self-compensating void with photometric redshifts which have uncertainties larger than the expected void signal is challenging. This assumes that the uncertainties in photometric redshifts can be fully captured by Gaussian-like error kernels. In reality, while photometric redshifts can be calibrated against the spectroscopic redshifts of objects amenable to spectroscopy, it is not clear that those objects unsuccessfully targeted by spectroscopy (often, as in our case the majority of the targeted objects) follow this calibration. Thus photometric redshift uncertainties could be considerably larger than those characterised from a simple comparison to the successful spectroscopy.  To illustrate the difficulty of using photometric redshifts to probe a putative large self-compensating void, we took the spectroscopic redshifts of all $I_{AB}<22.5$ VVDS sources with reliable or single-line redshifts (over 7300 objects at $z<1.4$), changed the redshifts of sources at $0.675<z<0.725$ to be at the limits of this range, thereby simulating a self-compensating void with $\\Delta z\\sim 0.05$. The redshifts were then randomized by shifting each by a random number drawn from a Gaussian distribution of width $\\sigma=0.05$, simulating the statistical uncertainties of photometric redshifts (but ignoring the systematic effect noted above). Figure 5 shows the three redshift distributions: the original VVDS distribution, the same but with the imposed fake void, and then this distribution with the simulated photometric uncertainties. As can be seen, even a broad, sharp void, one with the strongest possible signature for any cold spot, produces no discernible signal in the ``photometric'' redshift distribution because of the smoothing effect of the uncertainties.  Given the ambiguities of the photometric treatments, the advantage of the spectroscopic approach used here is clear. Our results show no evidence of a void connecting the six pointings towards the cold spot that is large enough to generate the cold spot through the ISW effect. Given the overall redshift distribution of the galaxies with confirmed redshifts, it is possible that such a void could exist at $z<0.35$ and with typically only 1-4 sources per $\\Delta z=0.01$ bin in the absence of a void, we would not be sensitive to it. Any future spectroscopic analysis of this area of sky should concentrate on the region at $z<0.5$ The amplitude of the secondary CMB fluctuation cold spot in the CMB is $7\\pm 3.10^{-5}$ (Cruz et al. 2007b), and if due to the Rees-Sciama effect, varies as the cube of the void radius. Hence the uncertainty in cold spot amplitude at specified angular extent translates into roughly a 15 per cent uncertainty in redshift. Consequently while any void is unlikely to be at $z < 0.5$, this possibility cannot be excluded.  Could we have missed a large void at $0.5<z<1$ through the placement of our fields?  Our fields were chosen based on the WMAP ILC III map of the cold spot region and cover several percent of the region with a decrement larger than 100 $\\mu$K.  If any void had a covering (and filling) factor of less than 100 per cent in this region and our field placings were essentially random within this, then there is a possibility that we could have missed the void. The covering factor cannot be much below 50 per cent as the ISW decrement is directly proportional to the covering factor. For a given signal, any decrease in the covering factor would have to be compensated for by an increased depth to a void in the line-of-sight, leading to potentially implausible geometries for the structure. If an under-density towards the cold spot was made up of several smaller voids with a total covering factor of 50 per cent rather than a single structure, the combined ISW effect from these will produce a far smaller (and therefore undetectable) temperature decrement than a single large void, and so such a scenario is very unlikely. Given the arrangement of our fields across the cold spot region, it is difficult to see how all or even a majority could have missed a single large system.  Our results indicate a large void at $0.5<z<1$ is unlikely to have generated the CMB cold spot temperature decrement through the ISW effect. Either such a void is at considerably lower redshift than expected, or there is some other explanation for the cold spot, one possibly involving new physics or a different analysis of current and future CMB data."
    },
    "1004/1004.3457_arXiv.txt": {
        "abstract": "The temporal instability of stably stratified flow was investigated by analyzing the Taylor-Goldstein equation theoretically. According to this analysis, the stable stratification $N^2\\geq0$ has a destabilization mechanism, and the flow instability is due to the competition of the kinetic energy with the potential energy, which is dominated by the total Froude number $Fr_t^2$. Globally, $Fr_t^2 \\leq 1$ implies that the total kinetic energy is smaller than the total potential energy. So the potential energy might transfer to the kinetic energy after being disturbed, and the flow becomes unstable. On the other hand, when the potential energy is smaller than the kinetic energy ($Fr_t^2>1$), the flow is stable because no potential energy could transfer to the kinetic energy. The flow is more stable with the velocity profile $U'/U'''>0$ than that with $U'/U'''<0$. Besides, the unstable perturbation must be long-wave scale. Locally, the flow is unstable as the gradient Richardson number $Ri>1/4$. These results extend the Rayleigh's, Fj{\\o}rtoft's, Sun's and Arnol'd's criteria for the inviscid homogenous fluid, but they contradict the well-known Miles-Howard theorem. It is argued here that the transform $F=\\phi/(U-c)^n$ is not suitable for temporal stability problem, and that it will lead to contradictions with the results derived from the Taylor-Goldstein equation. However, such transform might be useful for the study of the Orr-Sommerfeld equation in viscous flows. ",
        "introduction": "The instability of the stably stratified shear flow is one of main problems in fluid dynamics, astrophysical fluid dynamics, oceanography, meteorology, etc. Although both pure shear instability without stratification and statical stratification instability without shear have been well studied, the instability of the stably stratified shear flow is still mystery.  On the one hand, the shear instability is known as the instability of vorticity maximum, after a long way of investigations \\cite[]{Rayleigh1880,Fjortoft1950,Arnold1965a,SunL2007ejp,SunL2008cpl}. It is recognized that the resonant waves with special velocity of the concentrated vortex interact with flow for the shear instability \\cite[]{SunL2008cpl}. Other velocity profiles are stable in homogeneous fluid without stratification. On the other hand, \\cite{Rayleigh1883} proved out that buoyancy is a stabilizing effect in the statical case. Thus, it is naturally believed that the stable stratification do favor the stability \\cite[see, e.g.][]{Taylor1931,Chandrasekhar1961}, which finally results in the well known Miles-Howard theorem \\cite[]{Miles1961,Miles1963,Howard1961}. According to this theorem, the flow is stable to perturbations when the Richardson number $Ri$ (ratio of stratification to shear) exceeds a critical value $Ri_c=1/4$ everywhere. In three-dimensional stratified flow, the corresponding criterion is $Ri_c=1$ \\cite[]{Abarbanelt1984}.  However, the stabilization effect of buoyancy is a illusion. In a less known paper, \\cite{Howard1973} had shown with several special examples that stratification effects can be destabilizing due to the vorticity generated by non-homogeneity, and the instability depends on the details of the velocity and density profiles. One instability is called as Holmboe instability \\cite[]{Holmboe1962,Ortiz2002,Alexakis2009}. Then \\cite{Howard1973} stated three main points from the examples without any further proof. (a) Stratification may shift the band of unstable wave numbers so that some which are stable at homogeneous cases become unstable. (b) Conditions ensuring stability in homogeneous flow (such as the absence of a vorticity maximum) do not necessarily carry over to the stratified case, so that 'static stability' can destabilize. (c) New physical mechanisms brought in by the stratification may lead to instability in the form of a pair of growing and propagating waves where in the homogeneous case one had a stationary wave.  Recall the points by \\cite{Howard1973}, and that there is a big gap between Rayleigh's criterion and Miles-Howard' criterion, \\cite{YihCSBook1980} even wrote ``Miles' criterion for stability is not the nature generalization of Rayleigh's well-known sufficient condition for the stability of a homogeneous fluid in shear flow\". The mystery of the instability is still cover for us.  Following the frame work of \\cite{SunL2007ejp,SunL2008cpl}, this study is an attempt to clear the confusion in theories. We find that the flow instability is due to the competition of the kinetic energy with the potential energy, which is dominated by the total Froude number $Fr_t^2$. And the unexpected assumption in Miles-Howard theorem leads the contradiction to other theories. ",
        "conclusions": "In summary, the stably stratification is a destabilization mechanism, and the flow instability is due to the competition of the kinetic energy with the potential energy. Globally, the flow is always unstable when the total Froude number $Fr_t^2\\leq 1$, where the larger potential energy might transfer to the kinetic energy after being disturbed. Locally, the flow is unstable as the gradient Richardson number $Ri>1/4$. The approach is very straightforward and can be used for similar analysis. In the inviscid stratified flow, the unstable perturbation must be long-wave scale. This result extends the Rayleigh's, Fj{\\o}rtoft's, Sun's and Arnol'd's criteria for the inviscid homogenous fluid, but contradicts  the well-known Miles and Howard theorems. It is argued here that the transform $F=\\phi/(U-c)^n$ is not suitable for temporal stability problem, and that it will leads to contradictions with the results derived from Taylor-Goldstein equation.  The author thanks Dr. Yue P-T at Virginia Tech, Prof. Yin X-Y at USTC, Prof. Wang W. at OUC and Prof. Huang R-X at WHOI for their encouragements. This work is supported by the National Basic Research Program of China (No. 2012CB417402), and the Knowledge Innovation Program of the Chinese Academy of Sciences (No. KZCX2-YW-QN514)."
    },
    "1004/1004.4794_arXiv.txt": {
        "abstract": "In this paper we summarise the status of single field models of inflation in light of the WMAP 7 data release. We find little has changed since the $5$ year release, and results are consistent with previous findings. The increase in the upper bound on the running of the spectral index impacts on the status of the production of Primordial Black Holes from single field models. The lower bound on $\\fnleq$ is reduced and thus the bounds on the theoretical parameters of (UV) DBI single brane models are weakened. In the case of multiple coincident branes the bounds are also weakened and the two, three or four brane cases will produce a tensor-signal that could possibly be observed in the future. ",
        "introduction": "We review the status of single-field models of inflation in light of the latest data release from the Wilkinson Microwave Anisotropy Probe \\cite{wmap7}. We utilise the $7$ year $\\rm{WMAP}$ data, combined with the Baryon Acoustic Oscillations ($\\rm{BAO}$) and measurement of the Hubble parameter from the supernovae data, the $\\rm{H}0$ set. This data set combination is considered the best estimate for cosmological parameters at present \\cite{wmap7}.  We categorise our models into `canonical' and `non-canonical' models in concordance with Ref.~\\cite{ALi}. Canonical models have a pressure term $P$ given as $P=X-V(\\phi)$, where $X$ is the kinetic term, $V(\\phi)$ is the potential, and the inflaton ($\\phi$) fluctuations propagate at the speed of light. Non-canonical models on the other hand have a pressure term which is non-linearly dependant on $X$ and the inflaton fluctuations propagate at a different speed to light (see for example Refs.~\\cite{GM1999, Cake}). Our canonical models are then sub-categorised into `small' and `large' field models, where small field models are defined as those with an inflaton variation less than the Planck scale $\\Delta\\phi<\\mpl$.  The observational parameters that we will be utilising in this paper are the spectral index $n_s$, the running of the spectral index $n_s'$, the tensor fraction $r$ and the non-gaussianity parameter for an equilateral configuration $\\fnleq$. The $\\rm{WMAP}7+ \\rm{H}0$ data set gives bounds on these parameters which we list at the $2\\sigma$ confidence limit, \\bea\\label{obs} 0.939<n_s<0.987\\,,\\nonumber\\\\ r<0.24\\,,\\nonumber\\\\ -0.084<n_s'<0.02\\,,\\nonumber\\\\ -214<\\fnleq<266\\,. \\eea We have presented the bounds on $n_s$ for $n_s'=0$ and $r=0$ priors, the bound on $r$ is for an $n_s'=0$ prior and the bounds on $n_s'$ are for an $r=0$ prior.  A key parameter in inflation model discrimination is the number of $e-$folds $N$, which is the logarithmic ratio of the comoving hubble horizon at the end of inflation to its value at the time when scales of cosmological interest left the horizon. Liberal limits on $N$ are taken to be \\be\\label{N-liberal} 10<N<110\\,, \\ee where the lower bound comes from the demand that nucleo-synthesis is well bounded, and the upper bound assumes that the universe underwent a few bouts of ``fast'' roll inflation \\cite{Liddle:2003as}. The following bounds on $N$ are more commonly used in the literature \\be\\label{N-reason} N=54\\pm7\\,, \\ee where the uncertainty comes from our ignorance of reheating. In this paper we assume instant reheating.  In Section (\\ref{can}) we briefly overview models of inflation with canonical kinetic terms, both small and large field, and present the results of their predictions for $n_s$, $r$ and $n_s'$ and how they compare to the latest WMAP data release. We do not give an in depth review of these models or their motivation, for such reviews refer to Refs.~\\cite{Lyth-Riotto,AL1,AL2,ALi,book1,book2}. In Section (\\ref{non-can}) we analyse single field Dirac-Born-Infeld models of inflation in light of WMAP7. ",
        "conclusions": "In this paper we have analysed single field models of inflation with both canonical and non-canonical kinetic terms and find that the results are consistent with those of previous analyses using WMAP 5 data \\cite{ALi,Finelli:2009bs}. In Ref.\\cite{Finelli:2009bs} the authors use WMAP5 data in combination with ACBAR, QUAD, BICEP and SdSS LRG7 data and analyse models both with a zero and non-zero running. They conclude that there is some statistical preference of running models over power law models. In contrast, in this paper we have used the WMAP7 data combined with the BAO and Supernovae data sets and have not analysed models with a non-zero spectral running. We have updated the bounds on the relevant model parameters, and show that for the canonical models of small-field inflation, the power of the self-coupling in the tree-level potential with $p>0$ is more tightly constrained with the $p=3$ model requiring more than $67$ $e-$folds of inflation to satisfy the data at $2\\sigma$, and the logarithmic potential requiring less than $40$ $e-$folds of inflation to satisfy the data at $1\\sigma$. Since the upper bound on the running of $n_s$ has increased, this has expanded the parameter space which allows for the production of PBHs within astrophysical bounds. The hilltop-type model of inflation now leads to the formation of PBHs for $\\{p,q\\}=\\{2,3\\}$ for $N=68$ $e-$folds of inflation. We also find that imposing $N>20$ $e-$folds on the running mass model excludes the parameter space that leads to $n_s'>n_s'|_{WMAP5}$, and as such we find no change from the conclusions of Ref.~\\cite{AK}. We still find that the formation of PBHs is strongly dependent on the allowed upper bound on $N$, consistent with Refs.~\\cite{PE,AK}. However, the fact that PBHs may form after less than $20$ $e-$folds of inflation in this model, raises the question of whether this model leads to the overclosure of the universe, and may be an avenue for further research. The monomial potential with positive power is still consistent with data at $1$ and $2\\sigma$ for even a conservative range of $N$, and we find that we can now rule in the intermediate model with a power much greater than $2$ at the $1\\sigma$ level.  We find that the limits on the theoretical parameters of DBI models, with a single brane falling into a warped throat towards the tip, are weakened, however theoretical expectation is still at odds with observational bounds. In order to `marry' theory to observation we would need to motivate a smaller base volume, a larger Euler number or observe a more negative $\\fnleq$. Ref.~\\cite{Huang:2010up} included the bounds on \\emph{local} $f_{\\rm{NL}}$ in their analysis of the single field DBI model, and as a result exclude it from the $1\\sigma$ regime. We also find that the bounds on $\\fnleq$ and $r$ now allow for up to $4$ branes in multi-brane DBI inflation."
    },
    "1004/1004.1835.txt": {
        "abstract": "We present spectroscopic orbits for the active stars HD\\,82159 (GS\\,Leo), HD\\,89959, BD\\,$+39\\arcdeg$\\,2587 (a visual companion to HD\\,112733), HD\\,138157\\break (OX~Ser), HD\\,143705, and HD\\,160934. This paper is a sequel to one published in this journal in 2006, with similar avowed intention, by \\citeauthor{galvez06}. They showed only graphs, and gave no data, and no orbital elements apart from the periods (only two of which were correct) and in some cases the eccentricities. Here we provide full information and reliable orbital elements for all the stars apart from HD\\,160934, which has not completed a cycle since it was first observed for radial velocity. ",
        "introduction": "In a paper with a title somewhat analogous to ours, \\citet{galvez06} (cited as GMFL in what follows) presented graphs of orbital solutions for the same six stars as we discuss here. We did not set out to follow those authors, but the stars (with one exception, HD\\,160934) came to our attention by featuring in the third edition of \\textit{A catalogue of chromospherically active binary stars} \\citep{eker08} (hereinafter called \u0091CABS3\u0092), of which one of us is a co-author. They were among the minority of entries that lacked proper orbital solutions, and the other of the present authors undertook to make good some of what was missing. Twenty orbits for other such `CABS3' stars have already been presented \\citep{griffin09}, and it was thought useful to discuss the six GMFL stars all together now. Basic data about them are set out in Table~\\ref{table-1}.  \\begin{table*}[t] \\setlength{\\doublerulesep}{\\arrayrulewidth} \\caption{Basic data for the six stars\\label{table-1}\\vspace{2mm}} \\begin{tabular}{@{}crcrcrllcc@{}} \\hline \\\\[-8pt] CABS3 & HD$/$BD &  VS desig. & V~~ & (B\\,$-$\\,V) & Parallax~~~ & ~~$M_{V}$ & Sp. type & N &\\\\ &         &       & m~~ &  m  &  arc ms~~~~  &  ~~ m     &          &   &\\\\[5pt] \\hline\\hline \\\\[-5pt] ~~161&82159&GS Leo&8.85~&0.92&21.11\\hspace{2mm}$\\pm$\\hspace{1.5mm}5.71&+5.5&~~G9\\hspace{1mm}V&30\\\\ ~~177&89959&---&8.30~&0.68&---\\hspace{7mm}~&\\hskip0.6em---&~~K0\\hspace{1mm}V&33\\\\ ~~216&\\llap{+}$39^{\\circ}$2587&---&9.27~&0.84&26.24\\hspace{2mm}$\\pm$\\hspace{1.5mm}1.75&+6.4&~~G6\\hspace{1mm}V\\rlap{(?)}&82\\\\ ~~257&138157&OX Ser&7.14&1.02&5.07\\hspace{2mm}$\\pm$\\hspace{1.5mm}1.00&+0.7&~~K0\\hspace{1mm}III&24\\\\ ~~267&143705&---&7.96~&0.59&16.16\\hspace{2mm}$\\pm$\\hspace{1.5mm}0.98&+4.0&~~G0\\hspace{1mm}V&17\\\\ ~~---\\hspace{2mm}&160934&---&\\llap{1}0.28&1.3:~~&40.75\\hspace{2mm}$\\pm$\\hskip-0.05em 12.06&+8.3\\rlap{:}&~~K8\\hspace{1mm}V&38\\\\[5pt] \\tableline \\end{tabular} \\end{table*}  The new radial-velocity observations presented here were made with the 36-inch (0.91-m) coud\\'{e} reflector at the Cambridge Observatories, England, with a photoelectric radial-velocity spectrometer operating on the cross-correlation principle that was first developed \\citep{griffin67} at the same telescope by one of the present authors; the instrument in current use is largely patterned after the \\textit{`Coravel'} designed by \\citet{bmp79}. The radial-velocity traces (Fig.\\,\\ref{figure-2} below is an example) are cross-correlation functions that exhibit minima at abscissae corresponding to the velocities of the stars observed; in the case of a double-lined star there can be two minima (informally characterized here as `dips'). The dips may be thought of as averaged profiles of the absorption lines in the stellar spectrum.  The traces are routinely reduced by matching to digi\\-tally synthesized models that incorporate rotational broadening, which is computed by starting with an empirical profile whose ordinates are scaled directly from observed traces which have the minimum observed half-width. Many stars share a very sharply defined lower bound to the dip width, which is taken as representing the zero-rotation profile. The spectral lines of active stars, such as those treated in this paper, often show considerable rotational broadening. It is quantified in the model dips by the summation of many elements, into which the stellar disc is conceptually divided, each of which is assigned the appropriate velocity and a brightness which is assessed according to a conventional limb-darkening law. The $v$\\,sin\\,$i$ values (projected rotational velocities) thus determined often repeat very well (to better than $\\pm$~1~km\\,s$^{-1}$ r.m.s.) from one observation to another of the same star, so the mean can easily be precise to a very few tenths, but in view of the neglect of non-rotational sources of broadening the external (true) uncertainty of the mean $v$\\,sin\\,$i$ determined in this fashion is never claimed to be better than 1 km\\,s$^{-1}$.  Inasmuch as the stars treated here are all objects that exhibit chromospheric activity, they may (and at least in some cases do) have starspots.  The rotation of a spotted surface must affect, at some level, the radial velocities measured from the integrated light of the visible hemisphere of a star.  The effect is not large enough in the objects treated here, however, to vitiate the radial-velocity curves significantly, or even to add appreciably to the uncertainties of their determination. ",
        "conclusions": ""
    },
    "1004/1004.3382_arXiv.txt": {
        "abstract": "We measure the UV-optical color dependence of galaxy clustering in the local universe. Using the clean separation of the red and blue sequences made possible by the $NUV - r$ color-magnitude diagram, we segregate the galaxies into red, blue and intermediate ``green'' classes. We explore the clustering as a function of this segregation by removing the dependence on luminosity and by excluding edge-on galaxies as a means of a non-model dependent veto of highly extincted galaxies.   We find that $\\xi(r_p, \\pi)$ for both red and green galaxies shows strong redshift space distortion on small scales -- the ``finger-of-God''  effect, with green galaxies having a lower amplitude than is seen for the red sequence, and the blue sequence showing almost no distortion.  On large scales, $\\xi(r_p,\\pi)$ for all three samples show the effect of large-scale streaming from coherent infall. On scales $1 \\hmpc < r_p < 10 \\hmpc$, the projected auto-correlation function $w_p(r_p)$ for red and green galaxies fits a power-law with slope $\\gamma \\sim 1.93$ and amplitude $r_0 \\sim 7.5$ and $5.3$, compared with $\\gamma \\sim 1.75$ and $r_0 \\sim 3.9 \\hmpc$ for blue sequence galaxies. Compared to the clustering of a fiducial $L^*$ galaxy, the red, green, and blue have a relative bias of $1.5$, $1.1$, and $0.9$ respectively. The $w_p(r_p)$ for blue galaxies display an increase in convexity at $\\sim 1 \\hmpc$, with an excess of large scale clustering. Our results suggest that the majority of blue galaxies are likely central galaxies in less massive halos, while red and green galaxies have larger satellite fractions, and preferentially reside in virialized structures. If blue sequence galaxies migrate to the red sequence via processes like mergers or quenching that take them through the green valley, such a transformation may be accompanied by a change in environment in addition to any change in luminosity and color. ",
        "introduction": "\\label{sec:intro}  With the advent of the Sloan Digital Sky Survey (SDSS;\\citealt{Yor00}) and its value added galaxy catalogs, it has been possible to study the subject of galaxy bimodality and its relationship to fundamental properties, such as stellar mass and star-formation history \\citep[e.g.,][]{Kau04,Sch07,Sal07}. The broad division of galaxies into star forming disks and quiescent early type galaxies is the fundamental principle of Hubble's tuning fork system of classification and is well established. In a plot of optical $g-r$ color vs $M_r$, red galaxies define a clear sequence, while the locus of blue galaxies is broadened into the so-called \"blue cloud\". The red sequence has been shown to maintain its integrity with look-back time \\citep{Bow92} and has grown in mass since redshift $\\sim 1$. Studies by \\cite{Bel04}, \\cite{Bla06}, \\cite{Fab07}, and \\cite{Bro08} argue that the stellar mass contained within the red population has increased by roughly a factor of two in half the Hubble time.  A significant breakthrough in expressing this blue/red dichotomy occurred when photometry from the {\\sl Galaxy Evolution Explorer} (\\galex), notably the near-ultra violet ($NUV$) band, was matched with SDSS photometry \\citep{Mar07,Wyd07,Sch07,Sal07}. When the diagram is plotted using $NUV-r$ as the color, two clear sequences emerge: the familiar red sequence and a new blue sequence in place of the ''blue cloud'' of optical studies. Between the blue and red sequences there are galaxies present in a so-called ``green valley''. Many of these are spectroscopically classified Type II active galactic nuclei \\citep{Ric05,Mar07,Sal07}.  \\cite{Fab07}, \\cite{Mar07}, and \\cite{Sch07} propose several paths by which galaxies might transition from the blue to the red sequence. The presence of AGN in the green valley suggests that AGN activity is associated with a quenching of star formation \\citep{Sil98,Hop06,Hop07}. Other paths from the blue to the red sequence might, hypothetically, involve gas-rich mergers of blue galaxies \\citep{Too72}, or virial shock heating of cold gas streams \\citep{Dek06}. The red sequence might consolidate in luminosity via dissipationless mergers of red galaxies, or low luminosity blue galaxies might acquire bulges through mergers with starbursts, retaining sufficient mass to land the evolved galaxy on the red sequence. However, the green valley might also be populated by casual visitors --- red galaxies that acquire gas and form stars or feed a central AGN. Following this brief burst of star formation, these galaxies might ultimately return to the red sequence from which they started.  In \\cite{Sal07}, a plot of mass against specific star formation rate reveals a clear division between lower mass, star forming, blue sequence galaxies, and more massive AGN, which are not detected in large numbers until stellar mass $M\\sim 3\\times 10^{10}M_\\odot$. The process responsible for populating the green valley and for potentially contributing to evolution from blue to red must bear some relationship to environment and to the dark matter halos in which the galaxies reside. In this study, we investigate the clustering environment of the blue and red sequences, and for the green valley.  The current paradigm of structure formation assumes that galaxies are assembled in dark-matter halos. The dependence of the clustering on galaxy properties may provide clues to the baryonic processes that are important to galaxy formation and evolution. The dependence of galaxy clustering on galaxy type has been known since the earliest studies of extragalactic astronomy \\citep{Hub36, Zwi68}. In the modern era of large-scale galaxy surveys, \\cite{Dav76} showed that the angular auto-correlation of ellipticals has a steeper power-law slope that those of spirals. Recent redshift surveys using the Two-Degree Field Galaxy Survey (2dF) and SDSS confirms these earlier results and the apparent bimodal nature of galaxy clustering \\citep{Mad03,Bud03,Zeh05,Li06,Wan07}.  Studies using SDSS have further revealed that galaxy color is the property most predictive of local environment. \\cite{Bla05b} found that at fixed luminosity and color, density does not correlate with surface brightness nor the Sersic index, and argue that morphological properties of galaxies are less closely related to environment than their star-formation history, and are traced by broadband optical colors. (See \\cite{Par07} for an alternative analysis and point of view.) \\cite{Li06} found that the dependence of clustering on optical $g-r$ color and $D_{4000}$ is much stronger than structural parameters like concentration and surface brightness, and extend to $5 \\hmpc$, beyond what is expected from the localized halo paradigm of structure formation. They concluded that at fixed stellar mass, the clustering properties of the surrounding dark matter haloes are correlated with the color of the selected galaxies. They further argued that different physical processes may be required to explain environmental trends in star formation, distinct from those established by galaxy structure.  In this paper, we will consider the color dependence of the two-point physical correlation function of galaxies using samples constructed from \\galex, augmented with redshift and optical data from SDSS. In particular, we use the natural separation from the $NUV-r$ color to assign galaxies into three subsamples of red, green and blue galaxies. Our study complements recent work by \\cite{Hei08} who investigate the physical clustering of galaxies as a function of star-formation history in the local universe, as well as earlier studies by  \\cite{Mil07}, \\cite{Hei07} and \\cite{Bas08} who investigate the angular-correlation function of rest frame $UV$-selected galaxies and their evolution. We measure the auto-correlation function each of the different subsamples of galaxies, as well as the cross-correlation function between the subsamples. In section \\ref{sec:data}, we describe in detail the data used in this analysis. In section \\ref{sec:methods} we describe the method for estimating correlation functions. We present our results in section \\ref{sec:results}, and discuss their implications for the nature of green valley galaxies and the formation of red sequence galaxies. We summarize our findings in section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary} We have constructed a \\galex and SDSS matched catalog, where we have used the GR3 catalog from \\galex and the SDSS DR5 main spectroscopic galaxy sample. We construct the galaxy distribution of $NUV-r$ vs $M_r$ color-magnitude diagram, and divide the distribution into populations of red sequence, green valley and blue sequence. Since our main goal is to study the color dependence of clustering, we took substantial care in matching the luminosity distribution of each population. For each population, we measure the two-dimensional correlation function $\\xi(\\pi, r_p)$, and the one-dimensional projected correlation function $w_p$. We also perform cross-correlation analyses between each of the sub-populations.  Our principle finding is that the red sequence and green valley appear to show similar clustering properties, as expressed in the finger-of-God effect in the auto-correlation function. The projected correlation function is consistent with red and green galaxies residing as satellites of massive halos, while the blue sequence shows what appears to be a clear two-halo signature, hence primarily serving as central galaxies of less massive halos. The cross-correlation function also shows that green and blue galaxies, on small scales, are not a mere statistical mix, but are spatially segregated from each other.  The findings would appear to place the green valley population with the red sequence. The green valley would largely consist of massive galaxies that reside in massive halos, and which cluster like the red sequence. We note that \\cite{Mar07}, \\cite{Wyd07}, and \\cite{Sal07} show that a large fraction of type II AGN are found in the green valley. \\citeauthor{Sal07} show that in the plot of specific star formation rate vs stellar mass, the AGN tend to be found in massive ($>3\\times 10^{10}M_\\odot$) galaxies. The AGN occupy a region in these plots that strongly resembles that of the reddest class of galaxies, the ``no-H$\\alpha$'' red sequence galaxies. Significantly, the AGN are clearly offset from the locus of the blue sequence, in the plot of specific star-formation rate (SFR) versus mass. The significance is that while a minority of AGN are found with properties that coincide with those of the more massive blue sequence galaxies, green valley galaxies -- the subsample with the largest AGN fraction --- exhibit properties similar to those of the red sequence, but showing mildly elevated star formation.  In this study, we have shown that the green valley population clusters in ways that are characteristic of, but also less strongly than, the red sequence. One may suggest that these studies paint a picture in which both the properties of the green valley and the ``demographics'' are different from those of the blue sequence, at the present epoch. These findings do not necessarily contradict the studies that find an increase in the total mass of red sequence galaxies since $z\\sim 1$. They do suggest, however, that if blue sequence galaxies evolve by some process to the green valley, and ultimately to the red sequence, such evolution must also accompanied by a transition from the field environment to a group/cluster environment. Such a change could conceivably occur if the blue population resides along filaments that infall into clusters, over time. Our cross-correlation results show that green galaxies avoid both red and blue galaxies on small scales is consistent with the change in environment hypothesis. We note that models like ram pressure stripping \\citep{Gun72}, starvation of (cold) gas, and the virial shock heating model of \\citep{Dek06} naturally incorporate environmental factors in their mechanism for color transformation in galaxies.  It is also possible that the downsizing \\citep{Cow96} effect is so strong that most star forming galaxies are evolving rapidly with redshift \\citep[e.g.][]{Tin68}. One must recall that star forming activity at $z\\sim 1$ resides in considerably more massive galaxies, and that a color-based population separation, as we have done, will refer to much higher masses; the rest-frame colors may be similar, but the fundamental nature of the galaxies, not.  One may speculate that the the green valley is occupied by nominally red galaxies that experience the infall of a gas rich system that either induces star formation and/or fuels the AGN, rendering it visible via its emission lines. However, the feedback of an AGN might inhibit star formation and move a blue sequence galaxy to the green valley. Any number of environmental effects (e.g. harassment, starvation) might speed the consumption of gas in a disk, again moving a galaxy to the green valley. A small sample of optically quiescent members of the green valley that nominally have a UV excess show clear star formation signatures (spirals) when imaged in the UV using HST \\citep{Ric09}. There is still the issue of the origin of the low mass red sequence, and the evolution of blue sequence, into the green valley and ultimately the red sequence, might have an important role in the growth of the lower mass portion of the red sequence. This was partially addressed by semi-analytical work of \\cite{Ben03} and \\cite{Bow06}.  In contrast to the construction of color-magnitude diagrams for stellar populations, the environment, for galaxies, is a critical physical variable in their evolution. This is true both the in the sense of their dark matter environment as well as the presence of detectable companion stellar systems. In considering the major processes driving galaxy evolution, it would appear that evolution of both of these observables must be considered, as a function of look-back time."
    },
    "1004/1004.5271_arXiv.txt": {
        "abstract": "{Astronomical observations toward Sagittarius~B2(M) as well as other sources with APEX have recently suggested that the rest frequency of the $J = 1 - 0$ transitions of $^{13}$CH$^+$ is too low by about 80~MHz.} {Improved rest frequencies of isotopologs of methylidynium should be derived to support analyses of spectral recording obtained with the ongoing $Herschel$ mission or the upcoming SOFIA.} {Laboratory electronic spectra of four isotopologs of CH$^+$ have been subjected to one global least-squares fit. Laboratory data for the $J =  1 - 0$ ground state rotational transitions of CH$^+$, $^{13}$CH$^+$, and CD$^+$, which became available during the refereeing process, have been included in the fit as well.} {An accurate set of spectroscopic parameters has been obtained together with equilibrium bond lengths and accurate rest frequencies for six CH$^+$ isotopologs: CH$^+$, $^{13}$CH$^+$, $^{13}$CD$^+$, CD$^+$, $^{14}$CH$^+$, and CT$^+$.} {The present data will be useful for the analyses of $Herschel$ or SOFIA observations of methylidynium isotopic species.} ",
        "introduction": "\\label{intro}  The first three molecules detected in space in the 1930s were methylidynium, CH$^+$, methylidyne, CH, and cyanogen, CN. \\citet{CH+_ident} identified three absorption lines, which had been observed toward several stars, as belonging to the $A\\ ^1\\Pi - X\\ ^1\\Sigma ^+$ electronic transition of CH$^+$. They were assigned to the $R(0)$ ($\\equiv J = 1 - 0$) transitions of the $\\varv = 0 - 0$, $1 - 0$, and $2 - 0$ vibrational bands.  The $J = 2 - 1$, $3 - 2$, and $4 - 3$ rotational transitions were identified by \\citet{CH+_NGC7027} in spectra recorded with the Long Wavelength Spectrometer of the Infrared Space Observatory toward the planetary nebula NGC~7027. The observation of the $J = 1 - 0$ rotational transitions at 835137.5~MHz \\citep{isos-CH+_rot} is hampered by a telluric line of molecular oxygen close by \\citep{O2_rot_2010} and thus requires observations from space, e.g. with the recently launched {\\it Herschel} satellite or with the Stratospheric Observatory For Infrared Astronomy (SOFIA).  Recently, \\citet{det_13CH+} reported on the observation of the $J = 1 - 0$ transition of $^{13}$CH$^+$ in absorption toward the star-forming region G10.6-0.4. They deduced 830132~(3)~MHz as the prefered rest frequency based on scaling from the CH$^+$ laboratory rest frequency. Alternative rest frequencies of 830107~(1) or 830193~(4)~MHz were excluded. Very recently, \\citet{hydrides_2010} carried out an absorption study with the Atacama Pathfinder EXperiment (APEX) of light hydride species toward the evolved massive star-forming region Sagittarius (Sgr for short) B2(M). They detected low-lying rotational transitions of $^{13}$CH$^+$, H$^{35}$Cl, H$^{37}$Cl, and, for the first time, of SH$^+$. Similar studies have been carried out toward additional sources (F. Wyrowski, private communication). A comparison of the line profiles among these species together with observations of additional species at lower frequencies suggested a rest frequency of about 830210~MHz, approximately 80~MHz higher than the prefered rest frequency from \\citet{det_13CH+}, but in reasonable agreement with an alternative rest frequency of 830193~(4)~MHz they had dismissed.  \\citet{CH+_rot} obtained a value of 835078.950~(75)~MHz as the $J = 1 - 0$ transition frequency of CH$^+$, about 5~MHz lower than they deduced from analyses of the $A\\ ^1\\Pi - X\\ ^1\\Sigma ^+$ electronic spectra of four isotopologs \\citep{12CH+_1982,13CH+_UV_1997,CD+_UV_1987,13CD+_UV_1997}.  During the refereeing process of this manuscript I received $J = 1 - 0$ transition frequencies for CH$^+$, $^{13}$CH$^+$, and CD$^+$, very recently obtained by \\citet{isos-CH+_rot}. While the latter two frequencies ($\\sim$830216 and 453522~MHz) were rather close to predictions made from the spectroscopic parameters of the respective isotopolog \\citep{13CH+_UV_1997,CD+_UV_1987}, the CH$^+$ frequency (835137.5~MHz) was almost 40~MHz higher than that deduced from a more recent analysis of the CH$^+$ electronic spectrum \\citep{CH+_UV_2006}, but almost 70~MHz higher than measurements by \\citet{CH+_rot} indicated. Because no further rotational data and no infrared transitions have been reported thus far, determinations of methylidynium rest frequencies still rely quite heavily on the electronic spectra.  In the current work, I present a combined reanalysis of the $A\\ ^1\\Pi - X\\ ^1\\Sigma ^+$ electronic spectra of four isotopic species of CH$^+$ together with the available rotational transitions to derive reliable rest frequencies for CH$^+$ isotopologs. ",
        "conclusions": "\\label{conclusion}  The $A\\ ^1\\Pi - X\\ ^1\\Sigma ^+$ electronic spectra of four isotopologs plus the rotational data for three species have been reproduced rather well with one set of spectroscopic parameters and with only comparatively few transitions omitted from the fit. Considering that the $A$ electronic state has been judged to be perturbed by an unidentified state, the spectroscopic parameters show few peculiarities.  The predicted $J = 1 - 0$ transition frequencies of the three remaining isotopologs as well as higher-$J$ predictions for all species will be useful for future astronomical observations or laboratory spectroscopic investigations. The data for $^{14}$CH$^+$ and CT$^+$ may be useful for determining if $^{14}$C or T play a significant role in certain (circumstellar) environments."
    },
    "1004/1004.2451_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\setcounter{equation}{0}  \\hspace{0.7cm} Investigating field theory in de Sitter space may illuminates deep mysteries surrounding inflation in the early universe and dark energy of the present universe. The past and current exponential expansions of the universe are likely to be driven by the effective cosmological constants of the order of GUT and neutrino mass scales respectively. We are perplexed by the huge disparity of the relevant energy scales. Of course we do not understand why they are small in comparison to the Planck scale in the first place. Phenomenologically it appears that the cosmological constant has evolved with time. Although we may parametrize it by a scalar field with a suitable potential, its microscopic understanding is totally lacking.  Slow roll inflation models possess approximate conformal invariance and the conformal symmetry plays an important role to understand the magnitude of the correlators \\cite{Mald}\\cite{Weinberg}\\cite{Komatsu} and possible dS/CFT correspondence \\cite{Witten}\\cite{Strm}\\cite{BMS} . In string theory, there seems to be no stable de Sitter vacuum as we need to consider brane-antibrane systems to realize it.  Since our understanding is so sparse, we wonder if we are entering a completely new territory. It might very well be the case since the standard Feynman-Dyson perturbation theory breaks down in a time dependent background like de Sitter space. Feynman-Dyson formalism is the backbone not only in relativistic field theory but also in classical statistical mechanics and critical phenomena. In this sense our expertise might be confined in equilibrium physics while our problem belongs to non-equilibrium physics.  In fact we need to use Schwinger-Keldysh formalism to investigate field theory in a time dependent background like de Sitter space. We can derive a Boltzmann equation in this formalism which is a standard tool to investigate non-equilibrium physics \\cite{Polyakov}. In such a setting, it is in principle possible that the effective cosmological constant changes with time. In other words the dynamics may explain deep mysteries of this century if we can demonstrate that the cosmological constant decreases with time in an interacting field theory. Although there exist several proposals along this line of thoughts in the literature, our understanding is still in a preliminary stage \\cite{Polyakov}\\cite{Polyakov1}\\cite{Jackiw}\\cite{TW}\\cite{GB}.  There is a long history of studying Boltzmann equations in Schwinger-Keldysh formalism \\cite{SW} starting from Kadanoff-Baym \\cite{KB}\\cite{Kd}\\cite{kita} . In this paper we derive a Boltzmann equation in de Sitter space from a Schwinger-Dyson equation. This problem has been studied to the leading order of the derivative expansion of the Moyal product in the Wigner representation \\cite{Hohenegger}. However only the energy conserving process has been identified in such a limit. We go beyond the leading order of the expansion to investigate the particle production effects due to energy non-conservation in de Sitter space. We also investigate the energy-momentum tensor of an interacting scalar field theory to estimate the effective cosmological constant.  The organization of this paper is as follows. In section 2, we introduce a scalar field theory in de Sitter space. In section 3, we recall Schwinger-Keldysh formalism. In section 4, we determine the full propagator inside the cosmological horizon to the leading order in perturbation theory. In section 5, we estimate the effective cosmological constant from the energy-momentum tensor of a scalar field. We conclude in section 6 with discussions. ",
        "conclusions": "\\setcounter{equation}{0}  \\hspace{0.7cm} In this paper, we have investigated an interacting scalar field theory in de Sitter space. As Feynman-Dyson perturbation theory breaks down, we need to employ Schwinger-Keldysh formalism. Our problem therefore belong to non-equilibrium physics. It may illuminates great mysteries of our time: inflation in the early universe and dark energy today. Phenomenologically it appears that cosmological constant is not a constant but it changes with cosmic evolution. Since the vacuum changes with time in de Sitter space, it is logically possible that the effective cosmological constant also changes with time. It is because the expectation value of the matter energy-momentum tensor contributes to the effective cosmological constant through the Einstein's equation of motion. We are thus most interested in a possible time dependence of the effective cosmological constant.  We have first investigated the time dependence of the propagator well inside the cosmological horizon. We have derived a Boltzmann equation from a Schwinger-Dyson equation. We have found that the total integral of the spectral weight remains to be unity as the particle creation effects are accompanied by the reduction of the on-shell states. The leading IR contributions cancel between the real and virtual processes. This fact suffices for the complete cancellation of the infra-red divergences at zero temperature as pointed out in section 4. However it is not so at finite temperature. The remaining IR contribution leads to the modification of the particle distribution function as (\\ref{mdf}). This interesting effect vanishes in the zero temperature limit as the Bose distribution function itself vanishes.  Although the propagator is time dependent, explicit time dependence disappears when it is expressed by physical quantities. In another words, time evolution may be identified with the renormalization group evolution in an interacting field theory in de Sitter space. We have thus found that a field theory in de Sitter space shares many common properties with those in Minkowski space. We believe this is due to unitarity of the theory. We also believe it is due to the fact that the degrees of freedom inside the cosmological horizon does not change under cosmic evolution. In any field theory we need to adopt a fixed physical UV cut-off $\\Lambda$ while IR cut-off is provided by the Hubble constant. From this reason, we find that the degrees of freedom inside the cosmological horizon do not lead to diminishing cosmological constant.  On the other hand, the degrees of freedom outside the cosmological horizon increase with time as more and more degrees go out of the horizon. Indeed we find that they may give rise to a desired effect. As they accumulate with cosmic evolution, the screening effect of the cosmological constant grows. Although it is suppressed by $\\kappa\\lambda^2$, the suppression effect can be compensated by large logarithmic factors as arbitrary many degrees of freedom could accumulate outside the horizon.\\footnote{Such a mechanism is proposed by Tsamis and Woodard. See \\cite{TWM2} and references therein.} It is important to generalize our work to more generic situations including quantum effects of gravity. Since our investigation of the propagators is limited to the region well inside the horizon, we also need to understand the behavior of them around and beyond the cosmological horizon."
    },
    "1004/1004.5115_arXiv.txt": {
        "abstract": "We discuss a possible generation of  radio bursts preceding final stages of binary neutron star mergings which can be accompanied by short gamma-ray bursts. Detection of such bursts appear to be advantageous in the low-frequency radio band due to a time delay of ten to several hundred seconds required for radio signal to propagate in the ionized intergalactic medium. This delay makes it possible to use short gamma-ray burst alerts to promptly monitor specific regions on the sky by low-frequency radio facilities, especially by LOFAR. To estimate the strength of the radio signal, we assume a power-law dependence of the radio luminosity on the total energy release in a magnetically dominated outflow, as found in millisecond pulsars. Based on the planned LOFAR sensitivity at 120 MHz, we estimate that the LOFAR detection rate of such radio transients could be about several events per month from redshifts up to $z\\sim1.3$ in the most optimistic scenario. The LOFAR ability to detect such events would crucially  depend on exact efficiency of low-frequency radio emission mechanism. ",
        "introduction": "\\label{SectionI}  Despite the decade of active researches, cosmic gamma-ray bursts (GRBs) remain in the focus of modern astrophysical studies. A huge electromagnetic energy output of $\\sim 10^{48}-10^{53}$~ergs observed in GRBs requires gravitational or rotational energy release possibly mediated by magnetic field under extreme conditions (e.g., core collapse of massive rotating stars, binary neutron star (NS) or NS -- black hole (BH) binary mergings). Among many possibilities, the concept of collapsar \\citep{Woosley1993} for long GRBs (LGRBs) and compact binary mergings for short GRBs (SGRBs) (e.g. \\cite{Blinnikov1984,Eichler1989}, see a review by \\cite{Nakar2007}) appear to be the most viable ones. However, we are apparently still far from full understanding of these most energetic transient natural phenomena (see, e.g., the recent critical discussion by \\cite{Lyutikov2009}).  A magnetic mechanism may be required to explain the rich phenomenology of GRBs (e.g. \\cite{Barkov2008}). A BH surrounded by magnetized torus seems to be the prerequisite condition to form a GRB, since collimated relativistic outflows can not be produced by electromagnetic mechanism without external pressure \\citep{Lyubarsky2009}.  It is very challenging to seek for various messengers from complicated physical processes involved, such as gravitational-wave bursts \\citep{Sengupta2009} expected from compact binary mergings, active neutrino emission \\citep{Ruffert1997} and afterglows in the broad range of electromagnetic frequencies from radio  and optics to x-ray \\citep{Kann2008, Nysewander2009}. The GRB afterglows are associated with the interaction of the relativistic GRB ejecta with the surrounding medium (see e.g. \\cite{Hurley2006} for a review) and will not be considered here.  In the hard electromagnetic domain, the so-called GRB precursors preceding the main burst are found for a sizeable fraction of long GRBs with spectral properties similar to the main GRB emission (see \\cite{Burlon2008}), but none has been detected so far for short GRBs. Yet there are prospects for radio precursors for the merging of magnetized binary NS as well (e.g. \\citep{Lipunov1996,  Hansen2001, Moortgat2005}). If low-frequency radio emission can be generated prior to a SGRB, due to the dispersion  in the intergalactic plasma the radio signal would arrive \\textit{later} than the gamma-ray pulse \\citep{Lipunov_ea97}. So the GRB itself may be used as a trigger to search for such radio transient.  In this paper we would like to consider yet another possibility of the formation of non-thermal low-frequency radio emission that can be generated in relativistic plasma outflow prior to the final collapse of a binary neutron star. In contrast to the previous work by Hansen \\& Lyutikov, in which magnetospheric pair plasma generation was considered before the destruction of merging neutron stars, we shall investigate the next phase of the merging when a single differentially rotating object is formed and when the magnetic field can be amplified to magnetar values.  Our motivations for focusing on low-frequency radio emission are twofold.  First, the lower the frequency, the longer is the intergalactic delay and the more time is accessible for a radio telescope to point to the GRB position.  At frequencies above 1~GHz the delay is about a few seconds, which is insufficient for repointing of a large dish. The second motivation is stimulated by the approaching start of operation  of the LOFAR radio telescope, which will have a record high sensitivity in the low-frequency range and whose design allows one to react to alerts with the necessary rapidity  \\citep{Fenderetal2008,vLeeuwen2009a,vLeeuwen2009b}.  Our further considerations are based on the following three assumptions: (i) SGRBs are produced by binary NS mergers; (ii) an ultra-strong magnetic field ($\\sim10^{15-16}~\\mathrm{G}$) is needed to power the GRB engine; (iii) a rapidly rotating pre-GRB object with strong magnetic field which is formed immediately after the merging can radiate in radio waves very much the same as usual pulsar.  Neutron stars are known to have strong magnetic fields that can survive through long time of the NS life (e.g. internal toroidal field, \\cite{Abdolrahimi2009}). It is quite natural that these fields could give rise to various electromagnetic phenomena during final stages of the binary NS coalescence.  Binary NS population with the time prior to the coalescence due to gravitational wave emission less than the age of the Universe is known from binary pulsar observations \\citep{Regimbau2005}. The rate of binary NS mergings is estimated to be quite high, $\\mathcal{R}_{NS}\\sim 10^{2}-10^3~\\rm Gpc^{-3}yr^{-1}$ (see \\cite{PostnovYungelson06} for a review), which is about two order of magnitudes higher than the estimated rate of SGRBs: $\\mathcal{R}_{SGRBs}\\sim 1-10~\\rm Gpc^{-3}yr^{-1}$ \\citep{Nakar2007}. This discrepancy, of course, is not dangerous for binary NS mergings as the SGRB model considering narrow gamma-ray beaming and the possibility that not every merging ends up with a burst because of lack of required physical conditions (insufficient magnetic field, low mass, etc.).   Bursting electromagnetic emission can be generated  at different stages of the merging process. First, a joint magnetosphere of two coalescing NSs can be restructured to produce strong flares from reconnection of magnetic lines (this process should be considered separately and will be discussed elsewhere). Next, during several last orbital revolutions before the merging the magnetospheric plasma effects come into play (e.g. \\cite{Lipunov1996,Hansen2001}). Simulations show that after the merging  a massive object with rapid differential rotation is formed \\citep{Shibata2005, Rasio2005,Faber2006}, see review by \\cite{Duez2009}. This object can not immediately collapse into BH until it gets rid of excessive  angular momentum via some mechanism. At this stage a significant increase of the seed magnetic fields of NSs can occur: the energy of the differential rotation can be effectively transformed into the energy of magnetic field. Here the situation may be similar to what is thought to occur during the NS formation in the core collapse supernovae \\citep{Spruit2008}.  The growing magnetic field and rapid rotation can lead to the relativistic plasma generation and the formation of the outflow along the open magnetic field lines, in which pulsar-like low-frequency radiation can be produced. As we shall see, such a rapidly rotating strongly magnetized object can have much higher radio luminosity than even the brightest ordinary pulsars. Finally, the merging remnant can collapse into a BH (provided that its mass is above the maximum mass of a NS), possibly surrounded by a magnetized torus; such a configuration is favorable for the launch of a GRB. ",
        "conclusions": "\\label{SectionIV} A binary NS coalescence leads to the formation of a differentially rotating massive object which can eventually collapse to BH possibly surrounded by a magnetized torus. Numerical simulations show \\citep{Duez2006} that the magnetic field is strongly amplified at the pre-collapse stage up to $10^{15}-10^{16}$~G. This configuration is favorable for the formation of a short gamma-ray burst.  At the stage preceding the collapse and GRB, a relativistic plasma outflow can be produced by the differentially rotating strongly magnetized configuration. Some fraction of the total power of this outflow can be converted to electromagnetic radio emission. The main uncertainty in the estimation of the radio luminosity before the collapse of this configuration is the unknown efficiency $\\eta$ of the conversion of the rotational energy losses from the differentially rotating pre-collapse object into radio emission, which we assumed to depend on the total energy loss rate as $\\eta\\propto(\\dot{E})^{\\gamma}$ in analogy with millisecond radio pulsars. We have shown that  in the most optimistic case the planned LOFAR sensitivity at 120 MHz allows the potential detection at the signal-to-noise ratio level $>10$ of short (the proper duration of order of 10-100 ms, smeared by scattering in the intergalactic medium to $\\sim 100$ s) radio bursts associated with SGRBs from redshifts $z<1.3$ at a rate of $\\sim$ 90 events/year. For an assumed dispersion measure of 1000 cm$^{-3}$pc, the low-frequency 120 MHz radio bursts should be delayed by 290 s with respect to the prompt gamma-ray emission, which enables one to use SGRB alerts for rapid redirecting the LOFAR synthesized beam to the SGRB position on the sky.  The detection of non-thernal radio bursts associated with short GRBs will strongly indicate the involvement of high magnetic field in the GRB engine. The dispersion measure of such a burst obtained from radio observations will give an independent direct estimate of the GRB distance. Radio precursors to short GRBs detected by LOFAR will open up an interesting possibility to search for such transients in LOFAR all-sky surveys. Short radio transients with similar characteristics can be a signature of 'orphan' SGRBs. This information can also be used in the analysis of data obtained by modern gravitational wave detectors."
    },
    "1004/1004.2029_arXiv.txt": {
        "abstract": "{Molecular lines in the (sub)millimeter wavelength range can provide important information about the physical and chemical conditions in the circumstellar envelopes around Asymptotic Giant Branch stars.} {The aim of this paper is to study the molecular composition in the circumstellar envelope around the oxygen-rich star \\object{IK~Tau}.} {We observed IK Tau in several (sub)millimeter bands using the APEX telescope during three observing periods. To determine the spatial distribution of the $\\mathrm{^{12}CO(3-2)}$ emission, mapping observations were performed. To constrain the physical conditions in the circumstellar envelope, multiple rotational CO emission lines were modeled using a non local thermodynamic equilibrium radiative transfer code. The rotational temperatures and the abundances of the other molecules were obtained assuming local thermodynamic equilibrium.} {An oxygen-rich Asymptotic Giant Branch star has been surveyed in the submillimeter wavelength range. Thirty four transitions of twelve molecular species, including maser lines, were detected. The kinetic temperature of the envelope was determined and the molecular abundance fractions of the molecules were estimated. The deduced molecular abundances were compared with observations and modeling from the literature and agree within a factor of 10, except for SO$_2$, which is found to be almost a factor 100 stronger than predicted by chemical models.} {From this study, we found that IK Tau is a good laboratory to study the conditions in circumstellar envelopes around oxygen-rich stars with (sub)millimeter-wavelength molecular lines. We could also expect from this study that the molecules in the circumstellar envelope can be explained more faithful by non-LTE analysis with lower and higher transition lines than by simple LTE analysis with only lower transition lines. In particular, the observed CO line profiles could be well reproduced by a simple expanding envelope model with a power law structure.} ",
        "introduction": "\\label{introduction}   Stars with initial masses lower than $\\sim$8  $\\mathrm{M_{\\odot}}$ evolve to a pulsationally unstable red giant star on the Asymptotic Giant Branch (AGB). At this stage, mass loss from the evolved central star produces an expanding envelope. Further on, carbon, C, is fused in the core and then oxygen, O \\citep{Yamamura1996,Fukasaku1994}.  AGB stars are characterized by low surface temperatures, $T_{*}\\mathrm{\\leq 3000}$ K, high luminosities up to several $\\mathrm{10^{4}}$ $\\mathrm{L_{\\odot}}$, % and a very large geometrical size up to several AU \\citep{Habing1996}. In general, these highly evolved stars are surrounded by envelopes with expansion velocities between 5 $\\mathrm{km\\,s^{-1}}$ and 40 $\\mathrm{km\\,s^{-1}}$. They have high mass-loss rates between $\\mathrm{10^{-8}}$ and $\\mathrm{10^{-4}\\,M_{\\odot} \\,yr^{-1}}$. Their atmospheres provide favorable thermodynamic conditions for the formation of simple molecules, due to the low temperatures and, simultaneously, high densities. Due to pulsation, molecules may reach a distance at which the temperature is lower than the condensation temperature and at which the density is still high enough for dust grains to form. Radiation pressure drives the dust away from the star. Molecules surviving dust formation are accelerated  due to dust-grain collisions \\citep{Goldreich1976}.  The chemistry of the atmospheres and, further out, of the circumstellar envelopes (CSEs) around AGB stars is dependent on the chemical class. They are classified either as M stars (C/O abundance ratio $<$ 1), S stars (C/O $\\approx$ 1) or C stars (C/O $>$ 1). The optical and infrared spectra of AGB stars show absorption from the stellar atmosphere. M-type stellar spectra are dominated by lines of oxygen-bearing molecules, e.g., the metal oxides SiO and TiO, and $\\mathrm{H_{2}O}$. In C-star atmospheres carbon-bearing molecules like, a.o., CH, C$_2$, $\\mathrm{C_{2}H_{2}}$ and HCN are detected at optical and infrared wavelenths, and in the microwave regime \\citep[e.g.][]{Loidl2004}. While the \\emph{atmospheric} abundance fractions are nowadays quite well understood in terms of initial chemical composition, which may be altered by nucleosynthetic products which are brought to the surface due to dredge-ups, the main processes determining the \\emph{circumstellar} chemical abundance stratification of many molecules are still largely not understood. In the stellar photosphere, the high gas density ensures thermal equilibrium (TE). Pulsation-driven shocks in the inner wind region suppress TE. This region of strong shock activity is also the locus of grain formation, resulting in the depletion of few molecules as SiO and SiS. Other molecules, as CO and CS, are thought to be inreactive in the dust forming region \\citep{Duari1999}. At larger radii, the so-called outer envelope is penetrated by ultraviolet interstellar photons and cosmic rays resulting in a chemistry governed by photochemical and ion-molecule reactions. This picture on the chemical processes altering the  abundance stratification is generally accepted, but many details on chemical reactions rates, molecular left-overs after the dust formation, shock strengths inducing a fast chemistry zone etc. are not yet known.  Spectroscopical studies of molecular lines in the (sub)millimeter range are a very useful tool for estimating the physical and chemical conditions in CSEs. Due to its proximity, the carbon-rich AGB star \\object{IRC+10216} has attracted lot of attention, resulting in the detection of more than 60 different chemical compounds in its CSE \\citep[e.g.][]{Ridgway1976Natur.264..345R, Cernicharo2000A&AS..142..181C}. Until now, detailed studies of oxygen-rich envelopes have been rare. Recently, \\cite{Ziurys2007} have focused on the chemical analysis of the oxygen-rich peculiar red supergiant \\object{VY~CMa}. VY CMa is, however, not a proto-type of an evolved oxygen-rich star. A complex geometry is deduced from Hubble Space Telescope images \\citep{Smith2001} with a luminosity larger than $\\,10^5$\\,L$_{\\odot}$ and a mass-loss rate of $\\sim$$2 \\times 10^{-4}$\\,M$_{\\odot}$/yr \\citep{Bowers1983, Sopka1985}. VY~CMa is a spectacular object, which because of its extreme evolutionary state can explode as a supernova at any time. Interpreting the molecular emission profiles of VY~CMa is therefore a very complex task, subject to many uncertainties. To enlarge our insight in the chemical structure in the envelopes of oxygen-rich low and intermediate mass stars, we therefore have started a submillimeter survey on the oxygen-rich AGB star \\object{IK~Tau}, which is thought to be (roughly) spherically symmetric \\citep{Lane1987,Marvel2005}. We thereby will advance the understanding on the final stages of stellar evolution of the majority of stars in galaxies as our Milky Way and their resultant impact on the interstellar medium and the cosmic cycle.     \\subsection{IK Tau}  The Mira variable IK Tau, also known as NML Tau, is located at $\\alpha_{2000}$=$3^{\\mathrm{h}}53^{\\mathrm{m}}28^{\\mathrm{s}}$.8, $\\delta_{2000}$=$11^{\\circ}24^{\\prime}23^{\\prime\\prime}$. It was found to be an extremely cool star having large infrared ($J-K$) excess \\citep{Alcolea1999} consistent with a 2000\\,K blackbody. IK Tau shows regular optical variations with an amplitude of $\\sim$ 4.5 mag.  IK Tau is an O-rich star of spectral type ranging from M8.1 to M11.2 \\citep{Wing1973}. Its distance was derived by \\citet{Olofsson1998} to be 250 pc assuming a stellar temperature of 2000\\,K. The pulsation period is $\\sim$470 days \\citep{Hale1997}. The systemic velocity of the star is 33.7 $\\mathrm{km\\,s^{-1}}$. Mass-loss rate estimates range from $\\mathrm{2.4\\times10^{-6}}$ $\\mathrm{M_{\\odot}\\,yr^{-1}}$ \\citep[from the CO(J=1-0) line; ][]{Olofsson1998} to $\\mathrm{3\\times10^{-5}}$ $\\mathrm{M_{\\odot}\\,yr^{-1}}$ \\citep[from an analysis of multiple SiO lines; ][]{Gonzalez2003}.  In the circumstellar envelope of IK Tau maser emission from OH \\citep{Bowers1989}, $\\mathrm{H_{2}O}$ \\citep{Lane1987}, and SiO \\citep{Boboltz2005} and thermal emission of SiO, CO, SiS, SO, $\\mathrm{SO_{2}}$ and HCN have previously been found \\citep{Lindqvist1988,Bujarrabal1994,Omont1993}. Obviously, IK Tau is a prime candidate for circumstellar chemistry studies. ",
        "conclusions": "In this work, we present for the (sub)millimeter survey for an oxygen-rich evolved AGB star, being IK Tau, in order to study the chemical composition in the envelope around the central target.  An extensive non-LTE radiative transfer analysis of circumstellar CO was performed using a model with a power law structure in temperature and density and a constant expansion. The observed line profiles of $\\mathrm{^{12}CO(3-2)}$, $\\mathrm{^{13}CO(3-2)}$, $\\mathrm{^{12}CO(4-3)}$, and $\\mathrm{^{12}CO(7-6)}$ are fit very well by our model, yielding a mass-loss rate of $4.7 \\times 10^{-6}$\\,M$_\\odot$/yr. The line shapes and intensities for all $\\mathrm{^{12}CO}$ transitions are not much influenced by variations of the inner radius, which is understandable since the bulk of the $\\mathrm{^{12}CO}$ emission is produced in the outer envelope. The intensities for the higher excitation CO lines depend strongly on the assumed temperature but not on the value of the outer radius.  For 7 other molecules (SiO, SiS, HCN, CS, CN, SO, and SO$_2$) a fractional abundance study based on the assumption of LTE is performed. A full non-LTE analysis of all molecules is out of the scope of this observational paper, but will be presented in a next paper \\citep{Decin2010}. This study shows that IK Tau is a good laboratory to study the conditions in circumstellar envelopes around oxygen-rich stars with submillimeter-wavelength molecular lines.  The improved abundance estimates of this study will allow refinements of the chemical models in the future.  Molecular line modeling predicts the abundance of each molecule as a function of radial distance from the star, although some ambiguity about an inner or outer wind formation process often exists. To  get a clear picture on the different chemistry processes partaking in the different parts in the envelope, mapping observations for molecules other than CO should be performed. Since most of the submillimeter emission from molecules less abundant than CO  probably arises from  the inner part of the envelope at 2 -- 4$^{\\prime\\prime}$ meaningful observations require interferometers such as the future Atacama Large Millimeter Array (ALMA).  \\space \\space"
    },
    "1004/1004.2503_arXiv.txt": {
        "abstract": "We present a set of cosmological simulations with radiative transfer in order to model the reionization history of the universe from $z = 18$ down to $z=6$. Galaxy formation and the associated star formation are followed self-consistently with gas and dark matter dynamics using the {\\tt RAMSES} code, while radiative transfer is performed as a post-processing step using a moment-based method with M1 closure relation in the {\\tt ATON} code.  The latter has been ported to a multiple Graphical Processing Units (GPU) architecture using the CUDA language together with the MPI library, resulting in an overall acceleration that allows us to tackle radiative transfer problems at a significantly higher resolution than previously reported: $1024^3$ + 2 levels of refinement for the hydrodynamics adaptive grid and $1024^3$ for the radiative transfer Cartesian grid. We reach typical acceleration factor close to $100\\times$ when compared to the CPU version, allowing us to perform 1/4 million time steps in less than 3000 GPU hours.  We observe good convergence properties between our different resolution runs for various volume- and mass-averaged quantities such as neutral fraction, UV background and Thomson optical depth, as long as the effects of finite resolution on the star formation history are properly taken into account. We also show that the neutral fraction depends on the total mass density, in a way close to the predictions of photoionization equilibrium, as long as the effect of self-shielding are included in the background radiation model. Although our simulation suite has reached unprecedented mass and spatial resolution, we still fail at reproducing the $z\\sim 6$ constraints on the neutral fraction of hydrogen and the intensity of the UV background.  In order to account for unresolved density fluctuations, we have modified our chemistry solver with a simple clumping factor model. Using our most spatially resolved simulation (12.5 Mpc/h with 1024$^3$ particles) to calibrate our subgrid model, we have resimulated our largest box (100 Mpc/h with 1024$^3$  particles) with the modified chemistry, successfully reproducing the observed level of neutral Hydrogen in the spectra of high redshift quasars. We however didn't reproduce (by a factor of 2) the average photoionization rate inferred from the same observations.  We argue that this discrepancy could be partly explained by the fact that the average radiation intensity and the average neutral fraction depends on different regions of the gas density distribution, so that one quantity cannot be simply deduced from the other. ",
        "introduction": "After self-gravity, hydrodynamics and radiative cooling ( see e.g. \\citet{1985ApJS...57..241E, 1991ApJS...75..231H, 1992ApJS...78..341C,1996ApJS..105...19K, 1998ARA&A..36..599B} among other historical references), radiative transfer has been included only recently in cosmological simulations of the formation of large scale structure in the Universe (see e.g. \\citet{1999ApJ...523...66A, 2001NewA....6..437G, 2001MNRAS.324..381C, 2002ApJ...572..695R} and more recently \\citet{2006MNRAS.369.1625I,2007ApJ...671....1T,2007MNRAS.377.1043M,2009A&A...495..389B}) . Among many different astrophysical problems that require a proper treatment of light propagation, cosmic reionization stands out as a particularly challenging one, because the ionizing radiation field plays a key role in the transition from the \"dark ages\" to the era of galaxy formation~: the chronometry and the geometry of the process is entirely related to the way matter and radiation interact. The proper numerical modeling of cosmic reionization represents a additional challenge, since it requires to capture a whole set of physical phenomena which are difficult to tackle on their own (see e.g. the review by \\citet{2001PhR...349..125B}). In a nutshell, reionization can be described as \"atoms being dissociated by UV photons emitted by stars formed in collapsed, self-gravitating halos\".  This requires to follow the dynamics of dark matter and gas on large scale, cooling and star formation on galactic scales, the emission of ionizing radiation at microscopic scales and finally, UV light propagation back to the cosmological scales.  Because of this chain of causality involving many different cosmological fluids (dark matter, gas, stars and photons), it is only recently that significant progresses were made in the field of cosmological radiative transfer.  Computer simulations of radiative transfer cover a wide range of techniques, most of them reviewed in \\citet{2009arXiv0906.4348T} and with most implementations gathered in two sets of comparison papers \\citep{2006MNRAS.371.1057I,2009MNRAS.400.1283I}. Current cosmological radiative transfer codes successfully pass these rather academic tests, but it should be noted that only a few observational tests can be used as a probe to calibrate these rather complicate numerical tools. The first major constraint comes from quasars with the detection of Gunn-Peterson troughs and a decrease of the flux transmission in spectra of objects at z$\\sim$6, which can be interpreted as the mark of the transition from a neutral Universe to an ionized one (see e.g. \\citet{2004AJ....127.2598S}, \\citet{2006AJ....132..117F}). From the observed spectra and provided that some assumptions are made on the structure of the density field or the UV background, important quantities such as mean free path, photoionization rate or UV field intensity can be constrained (see e.g. \\citet{2002AJ....123.1247F}, \\citet{2006AJ....132..117F}, \\citet{2007MNRAS.382..325B}). These constraints provide anchor values at $z\\sim 6$ for the calibration of cosmological simulations of reionization and track their ability to simulate the post-overlap era and the overlap itself (see e.g. \\citet{2006ApJ...648....1G}). However, this technique only provide upper/lower boundaries at higher redshifts as complete absorption can be reached with a neutral fraction as low as 0.001. Furthermore, since models are used to infer physical properties from flux transmission, any agreement or disagreement between calculations and quantities derived from observations should be taken with caution (as noted by e.g. \\citet{2007ApJ...671....1T}) and in a reversed role the simulations may happen to be informative about the proper way to interpret data. The second set of constraints comes from the scattering of CMB photons by electrons released during the reionization process. Usually expressed in terms of the Thomson optical depth $\\tau$, current constraints from WMAP set $\\tau=0.084\\pm0.016$ implying a redshift of (instantaneous) reionization of $z\\sim10.9 \\pm 1.4$ \\citet{2009ApJS..180..330K}. This constraint results from the integrated impact of the electrons on the CMB properties and is therefore more sensitive to the complete history of cosmic reionization.  In this paper, our goal is to confront our new radiative transfer code {\\tt ATON} to these observational constraints, using a set of hydrodynamical simulations at different resolutions. This code has already been presented and tested using a standard test suite in \\citet{2008MNRAS.387..295A}. The dynamical simulations include gravity and gas physics with mesh refinement, as well as widely adopted and well tested star formation recipe. The radiative transfer is performed as a post-processing step (full coupling of hydrodynamics with radiation is currently underway). It relies on a moment-based description of the propagation of light in the same spirit as e.g. \\citet{2001NewA....6..437G} or \\citet{2009MNRAS.393.1090F}. The original {\\tt ATON} code has since been fully ported on Graphics Processing Units (GPU hereafter) architecture using CUDA. Thanks to the high acceleration rate ($\\sim 100\\times$ compared to CPU) made possible by such hardware, we have been able to simulate the radiative transfer at the same resolution as the hydrodynamics base grid with $1024^3$ cells.  The current article aims at reaching two objectives~: first, showing the ability of {\\tt ATON} to model properly the reionization process and second, demonstrate the potential of GPU architecture for numerical cosmology. Regarding the ability to model the reionization, we partially recover the observational constraints at $z\\sim 6$ if we include a simple clumping factor model. However we also find that the properties of the radiation field and the neutral fraction distribution are driven by very different regions, making it difficult to relate the average UV intensity to the average fraction of neutral gas. Regarding the adaptation of our code on GPU, we describe in details in the Appendix how such architecture can be used at full power on this type of problems.  This paper is organized as follows:~first we describe the methodology and the simulations. Second we describe a first set of fiducial simulations and assess in particular the issues related to resolution and numerical convergence. Third, we introduce a simple prescription for the subgrid clumping obtained from our most resolved simulation (12.5 Mpc/h with $1024^3$ dark matter particles) and apply it to the largest simulation we have (100 Mpc/h with $1024^3$ dark matter particles). Finally, we discuss our results, forthcoming applications and possible improvements. ",
        "conclusions": "We have presented a set of radiative cosmological simulations in order to model the reionization epoch from  $z \\sim 18$ down to $z \\sim 6$. The gas and dark matter dynamics, as well as the associated star formation have been performed with the {\\tt RAMSES} code, while radiative transfer has been computed by means of a moment--based formalism using the M1 closure relation, implemented in the {\\tt ATON} code. The latter has been ported on a multi-GPU architecture using CUDA, providing an acceleration close to 100x, which allows us to tackle radiative transfer problems at high resolution (a $1024^3$ base grid and 2 levels of refinement for the hydrodynamics and  a $1024^3$ Cartesian grid for the radiative transfer).  A good level of convergence on average quantities (neutral fraction, UV background and Thomson optical depth) has been observed between different simulations of increasing mass and spatial resolution, as long as the effect of finite mass resolution on the simulated star formation history are properly taken in account. We have also shown that the density dependance of the neutral fraction is close to the one predicted by photo-ionization equilibrium, as long as the effect of self-shielding are considered when defining the properties of the UV field. It also appears that without any other ingredients, our simulation fails at reproducing the $z\\sim 6$ constraints on the neutral fraction of hydrogen and the intensity of the UV background, in a similar manner to \\citet{2009MNRAS.400.1049F}.  By combining our best resolved simulation (12.5 Mpc/h and 1024$^3$ particles) with our largest simulated volume (100 Mpc/h with 1024$^3$ particles), we have introduced a subgrid clumping model in our chemistry solver, consistent with the one derived by e.g. \\citet{2007ApJ...657...15K}. We have shown that, although this clumping factor model is quite simplistic, its allowed us to reproduce the level of neutral gas deduced from the spectra of high redshift quasars, as did previously \\citet{2006ApJ...648....1G} or \\citet{2007ApJ...671....1T} among others. However, our estimation of the average photoionization rate is still at least a factor of 2 above the observational constraints. This \"semi-success\" can be explained by the fact that the average radiation intensity and the average neutral fraction depend on different regions of the gas distribution and one cannot simply deduce one from the other using photoionization equilibrium~: in other words, if one constraint is satisfied, the other can't be. However, we have argued that the photoionization rate is probably a more robust observational constraint than the neutral fraction. This suggests that some effort should still be done in our modelisation to reproduce the level of the UV background at $z\\sim 6$.  Among several prospects, one obviously think of increasing the resolution of the calculations. With GPU acceleration $2048^3$ hydro+ radiative transfer calculations are within reach. However, it clearly appears that coupled hydrodynamics and radiative transfer simulations are necessary at this stage, since an increase in resolution will inevitably rise the question of the impact of radiation on mini-haloes or on the star formation history. Also additional physics should be implemented, such as multi-group radiative transfer, where the importance of preheating by X-rays could therefore be fully assessed (see e.g. \\citet{2006MNRAS.371..867F, 2008ApJ...685....1S}), but also population III stars (\\citet{2007ApJ...671....1T}) and  varying star formation efficiencies and escape fractions (\\citet{2008ApJ...672..765G,2009ApJ...693..984W}). Overall, on a final positive note, our current results indicate a satisfying trend of cosmological calculations toward satisfying observational constraints.  \\noindent {\\bf Acknowledgments.} DA is supported by the ANR grant LIDAU and a \\emph{Conseil Scientifique} Grant from the University of Strasbourg. This work was granted access to the HPC resources of CCRT under the \"Grand Challenge Applications\" allocation for 2009."
    },
    "1004/1004.5579_arXiv.txt": {
        "abstract": "\\anta~is currently the largest neutrino telescope operating in the Northern Hemisphere, aiming at the detection of high-energy neutrinos from astrophysical sources. Such observations would provide important clues about the processes at work in those sources, and possibly help solve the puzzle of ultra-high energy cosmic rays. In this context, \\anta~is developing several programs to improve its capabilities of revealing possible spatial and/or temporal correlations of neutrinos with other cosmic messengers: photons, cosmic rays and gravitational waves. The neutrino telescope and its most recent results are presented, together with these multi-messenger programs. ",
        "introduction": "\\label{sec:intro} Astroparticle physics has entered an exciting period with the recent development of experimental techniques that have opened new windows of observation of the cosmic radiation in all its components: photons, cosmic rays, but also gravitational waves and high energy neutrinos, that could be detected both by IceCube \\cite{icecube} and \\anta.   The advantage of using neutrinos as new messengers lies firstly on their weak interaction cross-section ; unlike protons or $\\gamma$, they provide a cosmological-range unaltered information from the very heart of their sources. Secondly, charged particles are deflected by magnetic fields. Neutrinos on the other hand point directly to their sources and exact production site.  The neutrinos \\anta~is aiming at are typically TeV neutrinos from AGNs or galactic sources (microquasars), 30 orders of magnitude lower in flux than solar neutrinos. The detection of those specific neutrinos requires under water/ice instruments, or alternatively acoustic/radio techniques in the PeV-EeV range and air showers arrays above 1 EeV. In spite of efforts in those various energy ranges, since the detection of the MeV neutrino burst from SN 1987A by {\\sc Kamiokande/Baksan/imb/Mont-Blanc} \\cite{sn1987} no astrophysical source for neutrinos above a few GeV has ever been identified.  Sources for TeV $\\nu$ are typically compact objects (neutron stars/black holes), from which often emerge relativistic plasma jets with a still unclear composition - purely leptonic or with some hadronic component. Most of these sources have already been extensively studied from radio wavelengths up to $\\gamma$-rays. These photons can be produced by $e^-$ {\\it via} inverse compton effect (on ambient photon field)/synchrotron radiation, or by protons/nuclei {\\it via} photoproduction of $\\pi^{0}/\\pi^{\\pm}$ :  \\begin{equation} \\begin{array}{l} p / A  +  p / \\gamma \\longrightarrow \\pi^{0}~\\pi^{\\pm} \\\\ \\textrm{where~}\\pi^{0} \\rightarrow \\gamma \\gamma~\\textrm{and}~\\pi^{\\pm} \\rightarrow \\nu_{\\mu}~\\mu, \\textrm{~with~} \\mu \\rightarrow \\nu_{\\mu} \\nu_e e \\\\ \\end{array} \\label{eq:prodmu} \\end{equation}  In the former scenario, no neutrinos are produced, whereas in the latter, the neutrino flux is directly related to the gamma flux: a TeV neutrino detection from gamma sources would then yield a unique way to probe the inner processes of the most powerful events in the universe. Several hints exist which indicates that hadrons could be accelerated up to very high energies. Firstly, the combined radio, X-rays and $\\gamma$-rays observations of the shell-type supernova remnant RX J1713.7-3946 \\cite{rx} favour the production of photons {\\it via} $\\pi^0$ decay (figure \\ref{fig1}, left). Secondly, the correlations between X and $\\gamma$ for the Blazar 1ES1959+650 \\cite{1es} prove the existence of $\\gamma$ flares not visible in X (figure \\ref{fig1}, right), which is difficult to account for in purely leptonic models. \\begin{figure}[h!] \\centerline{\\includegraphics[width=\\linewidth]{fig1.eps}} \\caption{Left: Multiwavelength observations of the SNR RXJ 1713.7-39; the solid curve at energies above $10^7$~eV corresponds to $\\pi^0$-decay $\\gamma$-ray emission, whereas the dashed and dash-dotted curves indicate the inverse Compton (IC) and Nonthermal Bremsstrahlung (NB) emissions, respectively. Right: Whipple vs RXTE flux, for the Blazar 1ES1959+650, which shows the existence of orphan $\\gamma$ flares (crosses). \\label{fig1}} \\end{figure}  The connection between high energy neutrino astronomy and both gamma-ray astronomy, charged cosmic rays and gravitational waves thus emphasizes that not only a multi-wavelength but also a multi-messenger approach, combining data from different observatories, is suited for the study of the most powerful astrophysical sources in the Universe.  Sections \\ref{sec:telescope} and \\ref{sec:results} first describe the \\anta~neutrino telescope and its latest physics results. Section \\ref{sec:grbs} presents the different strategies imagined to detect neutrinos from gamma-ray bursts, while section \\ref{sec:correlations} details the correlations that can be performed with other observatories, such as {\\sc Auger}, {\\sc Hess}, or \\vo/\\lo. ",
        "conclusions": "\\label{sec:conclu}  \\anta~is now taking data since 2008 and has demonstrated the possibility to operate and get competitive physics results from a neutrino telescope under the sea in the Northern Hemisphere, in spite of its reduced size with respect to IceCube. It should also be noted that, because of its location, \\anta~can observe the Galactic Centre and most of {\\sc Hess} sources on the galactic plane.  Through the online and offline multi-messenger programs described hereabove, the \\anta~detector not only enhances its own capabilities as a neutrino telescope, but also contributes to the global effort of understanding the most violent phenomena in our Universe. In addition to offline searches for spatio-temporal correlations with other cosmic messengers (photons, cosmic rays and gravitational waves), \\anta~has the capability to handle external alerts in real time and to trigger follow-up observations with the small latency time required for the study of transient sources. The possibility to store a few minutes of raw data in coincidence with a GCN alert also brings new opportunities for offline analysis. This could be extended in the future to handle alerts involving other messengers, such as gravitational waves, to complete the {\\sc gwhen} project.  Finally, the extension of the follow-up programs to other instruments in different ranges of wavelengths (X, radio) would undoubtely contribute to the development of the astrophysical potential of \\anta.  \\vskip 1cm \\noindent \\textbf{Aknowledgements:} \\textit{Great thanks to the Organizing Committee, especially Fulvio Ricci, for trusting me for this talk.} \\vskip 0.5cm"
    },
    "1004/1004.2390_arXiv.txt": {
        "abstract": "{The Schmidt-Teleskop-Kamera (STK) is a new CCD-imager, which is operated since begin of 2009 at the University Observatory Jena. This article describes the main characteristics of the new camera. The properties of the STK detector, the astrometry and image quality of the STK, as well as its detection limits at the 0.9\\,m telescope of the University Observatory Jena are presented.} ",
        "introduction": "The University Observatory Jena is located close to the small village Gro{\\ss}schwabhausen, west of the city of Jena. The Friedrich Schiller University operates there a 0.9\\,m reflector telescope, which is installed at a fork mount (see, e.g., Pfau 1984)\\nocite{pfau1984}. The telescope, as well as all its instruments are operated from a control room, which is located in the first floor of the observatory building directly beneath the telescope dome.  The 0.9\\,m telescope can be used either as a Schmidt-camera, or as a Nasmyth telescope. In the Schmidt-mode ($f/D=3$) the telescope aperture is limited to \\mbox{$D=0.6$\\,m}, the aperture of the installed Schmidt-plate. In the Nasmyth-mode the full telescope aperture ${D=0.9}$\\,m is used at  ${f/D=15}$. At the Nasmyth port of the telescope the spectrograph \\mbox{FIASCO} (see Mugrauer \\& Avila 2009)\\nocite{mugrauer&avila2009} is installed, which also can be operated together with the CCD-camera CTK (see Mugrauer 2009)\\nocite{mugrauer2009} for simultaneous spectro-photo\\-metric monitoring of targets.  Already during the implementation phase of FIASCO at the 0.9\\,m telescope the design study of a new CCD-imager for the Schmidt-focus of the 0.9\\,m telescope was started. After the specification of its opto-mechanical design the, so called {S}chmidt-{T}eleskop-{K}amera (STK) was then built during 2008 by the company 4pi Systeme GmbH, under contract and after consulting the Astrophysical Institute and University Observatory Jena. First laboratory tests of the STK CCD-detector were then carried out at the end of 2008. Finally, at the begin of February 2009 the STK was installed in the tube of the 0.9\\,m telescope and saw its first light at the University Observatory Jena. Afterwards, the instrument was dismounted again as some modifications of its filter system (in particular software updates and laboratory tests of the filter system hardware) were necessary. End of June 2009 the STK was then reinstalled in the Schmidt-focus of the 0.9\\,m telescope for its first science operation run, which last until mid of November 2009.  During the first months of operation at the 0.9\\,m telescope the properties of the STK detector were determined (dark current, bias level, linearity), and several photo- and astrometric observations were carried out with the new camera to determine its image quality, astrometry, as well as its detection limits. Beside the instrument characterization, the STK was already used in several scientific projects, mainly photometric monitoring programs to study the variability of stars in young stellar clusters, follow-up observations of long-periodical variable stars, as well as of stars with transiting planets.  In this paper we present the main characteristics of the STK. In the second section the opto-mechanical design, as well as all individual components of the new camera are described in detail. Section 3 summarizes all properties of the STK detector. The astrometry and image quality of the new CCD-imager are discussed then in Sect. 4. The STK detection limits at the 0.9\\,m telescope of the University Observatory Jena are presented in Sect. 5. The illumination effect induced by the STK shutter in short integrated images is shown in Sect. 6. Finally, the first light observations of the STK at the 0.9\\,m telescope of the University Observatory Jena are presented in the last section. ",
        "conclusions": ""
    },
    "1004/1004.3443_arXiv.txt": {
        "abstract": "We report the results of our multicolor observations of PG 1115+080 with the 1.5-m telescope of the Maidanak Observatory (Uzbekistan, Central Asia) in 2001-2006. Monitoring data in filter $R$ spanning the 2004, 2005 and 2006 seasons (76 data points) demonstrate distinct brightness variations of the source quasar with the total amplitude of almost 0.4 mag. Our $R$ light curves have shown image C leading B by 16.4d and image (A1+A2) by 12d that is inconsistent with the previous estimates obtained by Schechter et al. in 1997 -- 24.7d between B and C and 9.4d between (A1+A2) and C. The new values of time delays in PG 1115+080 must result in larger values for the Hubble constant, thus reducing difference between its estimates taken from the gravitational lenses and with other methods. Also, we analyzed variability of the A2/A1 flux ratio, as well as color changes in the archetypal \"fold\" lens PG 1115+080. We found the A1/A2 flux ratio to grow during 2001-2006 and to be larger at longer wavelengths. In particular, the A2/A1 flux ratio reached 0.85 in filter $I$ in 2006. We also present evidence that both the A1 and A2 images might have undergone microlensing during 2001-2006, with the descending phase for A1 and initial phase for A2. We find that the A2/A1 flux ratio anomaly in PG 1115 can be well explained both by microlensing and by finite distance of the source quasar from the caustic fold. ",
        "introduction": "Gravitationally lensed quasars are known to potentially provide estimates of the Hubble constant $H_0$ from measurements of the time delays between the quasar intrinsic brightness variations seen in different quasar images (Refsdal 1964).  Since a phenomenon of gravitational lensing is controlled by the surface density of the total matter (dark plus luminous), it provides a unique possibility both to determine the value of $H_0$ and to probe the dark matter content in lensing galaxies and along the light paths in the medium between the quasar and observer.  By now the time delays have been measured in about 20 gravitationally lensed quasars resulting in the values of $H_0$ that are generally noticeably less than the most recent estimate of $H_0$ obtained in the HST Hubble Constant Key Project with the use of Cepheids -- $H_0=72\\pm8$ km s$^{-1}$ Mpc$^{-1}$ (Freedman et al. 2001). This discrepancy is large enough and, if the Hubble constant is really a universal constant, needs to be explained. A detailed analysis of the problem of divergent $H_0$ estimates inherent in the time delay method, and the ways to solve it can be found, e.g., in Keeton \\& Kochanek (1997), Saha \\& Williams (1997), Kochanek (2002), Kochanek \\& Schechter (2004) and Schechter (2005), Read et al. (2007), and in many other works.  The main sources of uncertainties in determining $H_0$ are: \\begin{itemize} \\item low accuracy of the time delay estimates caused by poorly sampled and insufficiently accurate light curves of quasar components, as well as by microlensing events and, as a rule, by low amplitudes of the quasar intrinsic variability; \\item difference in the values of cosmological constants adopted in deriving $H_0$; \\item invalid models of mass distribution in lensing galaxies. \\end{itemize}  The way to reduce the effect of the first source of errors is clear enough: more accurate and better sampled light curves of a sufficient duration are needed. A choice of the cosmological model is usually just indicated - this is mostly a question of agreement. As to the third item, here the problem of estimating the Hubble constant encounters the problem of the dark matter abundance in lensing galaxies.  The problem of determining the Hubble constant from the time delay lenses is known to suffer from the so-called central concentration degeneracy, which means that, given the measured time delay values, the estimates of the Hubble constant turn out to be strongly model-dependent. In particular, models with more centrally concentrated mass distribution (lower dark matter content) provide higher values of $H_0$, more consistent with the results of the local $H_0$ measurements than those with lower mass concentration towards the center (more dark matter). Moreover, it has long been noticed that the time delays are sensitive not only to the total radial mass profiles of lensing galaxies, but also to the small perturbations in the lensing potential (e.g., Blandford \\& Narayan 1986,  Witt et al. 2000, Oguri 2007). It is interesting to note that this effect has been recently proposed as a new approach to detect dark matter substructures in lensing galaxies (Keeton \\& Moustakas 2009).  The Hubble constant -- central concentration degeneracy is a part of the well known total problem of lensing degeneracies: all the lensing observables, even if they were determined with zero errors, are consistent with a variety of the mass distribution laws in lensing galaxies. A strategy for solving this non-uniqueness problem could be a search through a family of lens models that are capable of reproducing the lensing observables (Williams \\& Saha 2000, Oguri et al. 2007). Then many models can be run in order to infer a probability density for a parameter under investigation, e.g. for $H_0$ (Williams \\& Saha 2000). The most recent studies (Saha et al. 2006, Read et al. 2007) have shown that, in such an approach, the discrepancy between the $H_0$ value determined from lensing and with other methods can be substantially reduced if non parametric models for mass reconstruction are used, which can provide much broader range of models as compared to the parametric ones.  In defining priors on the allowed space of lens models, it is naturally to assume that lensing galaxies in the time delay lenses are similar in their mass profiles to other early-type ellipticals, that are presently believed to be close to isothermal and admit the presence of the cold dark matter haloes. The isothermal models are also consistent with stellar dynamics, as well as with the effects of strong and weak lensing.  The quadruply lensed quasars are known to be more promising for solving these problems as compared to the two-image lenses since they provide more observational constraints to fit the lens model. Ten astrometric constraints can be presently regarded as measured accurately enough for most systems. This especially concerns the relative coordinates of quasar images. As to the lensing galaxiy, its less accurate coordinates are often the only reliable information about the lensing object known from observations, with other important characteristics being derived indirectly. This situation is inherent, e.g. in PG 1115+080 with its faint, 0.31- redshift galaxy. Of other observational constraints, the time delays and their ratios are very important. In quadruple lenses, the time delay between one of the image pairs is usually used to determine $H_0$, while the other ones form the $H_0$-independent time delay ratios to constrain the lens model, (Keeton \\& Kochanek 1997).  It has long been known that the observed positions of multiple quasar macroimages are well predicted by smooth regular models of mass distribution in lensing galaxies, while their brightness ratios are reproduced by such models poorly, (e.g., Kent and Falco 1988, Kochanek 1991, Mao \\& Schneider 1998). The first systematic analysis of this problem called \"flux ratio anomalies\" was made by Mao \\& Schneider (1998), who assumed that the anomalies of mutual fluxes of the components in some lenses can be explained by the presence of small-scale structures (substructures) in lensing galaxies or somewhere near the line of sight.  A popular model of forming hierarchical structures in the Universe with a dominant content of dark matter is currently known to poorly explain the observed distribution of matter at small scales. In particular, the expected number of satellite galaxies with masses of the order of $M_G\\approx10^8M_\\odot$ remained after the process of hierarchical formation is completed, is an order of magnitude larger than a number of dwarf galaxies with such masses actually observed within the Local Group (Klypin et al. 1999, Moore et al. 2001). One of the solutions of this contradiction is a suggestion that some substructures, especially those with low masses, are not luminous. \\begin{figure} \\centering \\resizebox{0.7\\hsize}{!}{\\includegraphics{Fig1.eps}} \\caption{PG1115+080 from observations in filter $R$ with the 1.5-m telescope of the Maidanak Observatory. The image was obtained by averaging of six frames from a series obtained in February 24, 2004, with a subsequent Richardson-Lucy processing.  } \\end{figure}  Metcalf and Madau (2001) were the first to note that the dark matter paradigm can naturally explain existence of substructures in galaxies lensing the remote quasars, as proposed by Mao and Schneider (1998) to interpret the anomalies of mutual fluxes of quasar macroimages, and vice versa, confirmation of substructures with masses from $10^6M_\\odot$ to $10^8M_\\odot$ is capable of removing the contradiction between the predicted number of the low-mass satellite galaxies and that one actually observed. The idea turned out to be intriguing and was immediately taken up, (Brada\\v c et al. 2002, Chiba 2002, Dalal \\& Kochanek 2002, Keeton 2001, Metcalf \\& Zhao 2002). Investigation of flux ratio anomalies in gravitationally lensed quasars is presently believed to be a powerful tool in solving the problem of the dark matter abundance in the Universe. It is intensively discussed in numerous recent publications (Congdon \\& Keeton 2005; Keeton et al. 2003, 2005; Kochanek \\& Dalal 2004; Mao et al. 2004; Miranda \\& Jetzer 2007; Pooley et al. 2006, 2007, 2009; Morgan et al. 2008).  \\begin{figure*} \\centering \\resizebox{0.8\\hsize}{!}{\\includegraphics{Fig2.eps}} \\caption{The light curves of PG 1115+080 A1, A2, B, C from observations in filter $R$ with the 1.5-m telescope of the Maidanak Observatory in 2001, 2002, 2004, 2005 and 2006.} \\end{figure*}  In Sec. 3 we report our measurements of the A2/A1 flux ratios in filters $V$, $R$ and $I$ from our data obtained in 2001, 2002 and 2004-2006 at the Maidanak Observatory and analyse their behaviors in time and in wavelength. In Sec. 4, we analyze the new estimates of the time delays between the PG 1115+080 images, obtained from our monitoring in the $R$ filter during 2004-2006 and reported in Vakulik et al. (2009). The new values differ noticeably from those reported by Schecter et al. (1997) and Barkana (1997). In Sec. 5, we discuss our results and their possible effect on selecting an adequate lens model for PG 1115+080 and estimating the value of the Hubble constant. ",
        "conclusions": ""
    },
    "1004/1004.1994_arXiv.txt": {
        "abstract": "{ We analyze the possibility of primordial magnetic field amplification by a stochastic large scale kinematic dynamo during reheating. We consider a charged scalar field minimally coupled to gravity. During inflation this field is assumed to be in its vacuum state. At the transition to reheating the state of the field changes to a many particle/anti-particle state. We characterize that state as a fluid flow of zero mean velocity but with a stochastic velocity field. We compute the scale-dependent Reynolds number $% Re(k)$, and the characteristic times for decay of turbulence, $t_{d}$ and pair annihilation $t_{a}$, finding $t_{a}\\ll t_{d}$. We calculate the rms value of the kinetic helicity of the flow over a scale $\\mathcal{L}$ and show that it does not vanish. We use this result to estimate the amplification factor of a seed field from the stochastic kinematic dynamo equations. Although this effect is weak, it shows that the evolution of the cosmic magnetic field from reheating to galaxy formation may well be more complex than as dictated by simple flux freezing.} ",
        "introduction": "The question of the origin of large scale magnetic fields that permeate almost all structures of the universe is one of the most challenging areas of research in astrophysics. None of the main lines of investigation, namely primordial origin or in situ generation, succeeded up to now to explain both the intensity and the topology of the large scale fields. Local generation mechanisms are mainly based on seed field generation by, e.g., a local battery, amplified by a turbulent dynamo in the interstellar or intergalactic medium (see \\cite{bran-sub-05} and references therein). The primordial origin hypothesis, on the other hand, considers that at least the seed field is generated at some early epoch (inflation, reheating or radiation dominance), and is amplified by flux conservation and/or turbulent dynamo action during gravitational collapse from $z\\approx 100$ on \\cite% {CalKan99}. The seed field must be quite intense for gravitational collapse to produce the detected intensities, and the turbulent dynamo must operate almost since the birth of the galaxy, i.e., during most of the matter dominated era. The recent detection of regular fields in high redshift quasars \\cite{wolfe-08}, \\cite{kronberg-nat-08}, \\cite{kronberg-apj-08} however may challenge the in situ generation, or at least the dynamo mechanism in the form we understand it today, favoring the primordial origin of the fields.  Two obstacles must be overcome by a successful primordial generation mechanism: breaking the conformal symmetry of a massless gauge field in a spatially flat universe and building a large coherence length. Sub-horizon processes, like phase transitions \\cite{hogan-83}, \\cite{qls-89}, \\cite% {co-94}, \\cite{soj-97}, in general produce intense fields, but of very small coherence length (see Refs \\cite{gra-rub-01} and \\cite{widrow-02} for reviews of different magnetogenesis mechanisms). The inflationary epoch of the universe (if ever existed) offers a suitable scenario for large scale field generation as in it sub-horizon scales naturally become super-horizon. Several mechanisms were considered along the years in which conformal invariance is broken either by coupling the magnetic field to curvature in different ways or addressing non-linear electrodynamics \\cite{turn-wid-88}, \\cite{mazz-sped-95}, \\cite{tsa-kan-05}, \\cite{kunze-08}. In general the fields produced are extremely weak, or of marginal intensity, to seed subsequent amplification processes. The reheating period has also been studied as a magnetogenesis scenario (\\cite{cal-kan-mazz}, \\cite% {kan-cal-mazz-wag}, \\cite{giov-shap}, \\cite{cal-kan-02}, \\cite{maroto-01}) but in all scenarios considered so far the obtained fields are too weak to be of astrophysical interest.  Confronted with this situation one wonders if it is possible to have a pre-amplification (or perhaps full amplification) of a seed field created by one of the above mentioned mechanisms already in the early universe. In this sense the reheating epoch offers a good prospect, as it is a period where highly non-linear and out of equilibrium processes take place \\cite% {CalHu08,khleb-tkach-96,khleb-tkach-97,ko-li-sta-96,ko-li-sta-97,fin-bran-99,fin-bran-00,feld-tkach-08,feld-gbell-etal-01,feld-kof-01,gra-cal-02, je-le-ma-10}. This possibility was explored for the first time some years ago by Finelli and Gruppuso \\cite{FiGru-01} and by Bassett et al. \\cite{BaPoTsuVi-01}. In Ref. \\cite{FiGru-01} it is analyzed the amplification of a pre-existing  magnetic field by parametric resonance during the oscillatory regime of a scalar field to which the magnetic field is coupled. In Ref. \\cite{BaPoTsuVi-01} the amplification during preheating is studied considering several different models. Another possibility for such pre-amplification process, and that will be investigated in this paper, could be the operation of a turbulent large scale dynamo \\cite{moffatt}, \\cite% {dan-pab-1}, \\cite{bran-sub-05}, \\cite{mininni-07}, similar to the one that acts in the interstellar plasma.  That the matter fields in reheating can be turbulent was pointed out in Refs. \\cite{khleb-tkach-96}, \\cite{feld-tkach-08}, \\cite{feld-gbell-etal-01}% , \\cite{feld-kof-01} (see \\cite{gra-cal-02} for a theoretical analysis of turbulent reheating). A dynamo requires the presence of a plasma. As the inflaton is a gauge singlet, it will not decay directly into charged species. Therefore to have a plasma we must consider an extra, charged, field. The mechanism by which the plasma is created is particle creation during the transition from inflation to reheating \\cite% {birrel-davies,MukWin07,CalHu08,ParTom09}.  Suppose that the charged species in question was in its vacuum state during inflation. The created particles will generate stochastic currents that on one hand induce a seed field \\cite{cal-kan-mazz}, \\cite{giov-shap} and on the other may constitute the turbulent plasma we are looking for. Creation of spin 1/2 particles such as electrons is suppressed by conformal invariance at the high energies prevailing during inflation \\cite% {cal-kan-mazz}, so the charged species must be a scalar. Suitable candidates can be found in supersymmetric extensions of the standard model \\cite% {kan-cal-mazz-wag}.  The simplest model for a turbulent large scale dynamo is driven by flow velocities and does not take into account the back-reaction of the amplified fields. It is known as a \\textsl{kinematic dynamo} \\cite{moffatt}, \\cite% {dan-pab-1}, \\cite{bran-sub-05}, \\cite{mininni-07}. The sufficient condition for it to be operational is the flow to be helical, i.e., that the volume average of the scalar product of the vorticity (curl of the velocity) and the velocity, known as \\textsl{kinetic helicity} \\cite{lesieur}, does not vanish \\cite{bran-big-sub-01}, \\cite{min-go-mah-05}. Of course this approximation (the neglect of the back reaction of the induced field) is valid for weak magnetic fields and/or very short times of operation. Mathematically speaking, the equation for that dynamo can be written as $% \\partial \\overline{B^{i}}/\\partial t\\simeq -t_{corr}\\mathcal{H}_{c}\\epsilon _{ijk}\\partial B^{k}/\\partial x^{j}$, where $\\overline{B^{i}}$ is the large scale field (or mean field), $\\mathcal{H}_{c}$ the kinetic helicity and $t_{corr}$ a correlation time. If $\\epsilon _{ijk}\\partial B^{k}/\\partial x^{j} \\sim \\overline{B^{i}}/L$, with $L $ the coherence length of the field, then we can estimate $\\overline{B^{i}}% \\left( t\\right) \\sim \\overline{B_{0}^{i}}\\left( 0\\right) \\exp \\left( -t_{corr}\\mathcal{H}_{c}t/L\\right) $.  In this paper we shall investigate the possibility of a dynamo action during reheating. We assume the existence of a charged scalar field, minimally coupled to gravity, that is in its vacuum state during inflation. To simplify the analysis we consider de Sitter inflation and thus a de Sitter invariant vacuum for the field \\cite{allen}. As mentioned above, when the transition from inflation to reheating takes place, the scalar field is amplified, and stochastic currents are generated. The characterization of these particles as a fluid is straightforward. The hydrodynamic energy and pressure are determined by matching the expectation value of the energy-momentum tensor of the scalar field to that of a perfect fluid at rest.  The fluid has a stochastic Gaussian velocity, which is found by matching the self-correlation of the $0i$ components of the energy momentum tensor of the fluid to the symmetric expectation value of the corresponding operator for the field. Finally, the viscosity of the fluid is found by assuming that it is close to saturate the Kovtun, Son and Starinets bound \\cite{son-1}. While initially derived from consideration of the AdS/CFT correspondence, the fact that a similar bound seems to hold for the strongly coupled quark gluon plasma \\cite{LuzRom09} suggests that this bound is a good description of field theories in general.  We characterize the turbulence by finding the momentum dependent Reynolds number $Re\\left( k\\right) $. As for the magnetic Reynolds number, $R_{m}$ we do not need to estimate it because we are interested in the kinematic regime, where magnetic fields are too weak to backreact on the flow. As there are no stirring forces, turbulence will decay eventually. We calculate the decay time of the turbulence for each mode, $t_{d}\\left( k\\right) $. On the other hand, the fluid is made of particles and antiparticles, which are liable to annihilate. We also estimate the characteristic time for pair annihilation, $t_{a}\\left( k\\right) $, finding that $t_{a}\\left( k\\right) <t_{d}\\left( k\\right) $, i.e., the fluid annihilates before turbulence decays. This fact allows us to consider that the turbulence is stationary in the interval $0\\leq t\\leq t_{a} $.  The non-trivial result of our paper is that the rms value of the kinetic helicity of the fluid, $\\mathcal{H}_{c}$, is not zero. This proves that the kinematic dynamo action mentioned above is indeed possible. The key ingredient to have $\\mathcal{H}_{c}\\neq 0$, is the fact that the plasma is made up of two scalar fields, $\\Phi $ and $\\Phi ^{\\dagger }$. We then estimate the amplification factor of the induced field based on the dynamo equation written above, i.e., $\\sim \\exp \\left[ \\mathcal{H}_{c}t_{corr}/L\\right] $.  We work with signature $\\left( -,+,+,+\\right) $ and with natural units, i.e., $\\hbar =c=k_{B}=1$. Greek indexes denote space-time coordinates while latin indexes refer to spatial coordinates. To simplify our analysis we shall consider de Sitter inflation, and define dimensionless variables and fields as $x^{i}=Hr^{i}$, $\\tau =Ht$, and $\\Phi =H^{-1}\\Psi $, where $H$ is the Hubble constant that characterizes the de Sitter phase. The mass of the scalar field will combine with $H$ to produce the dimensionless mass parameter, $m/H$. In section I we make a brief review of dynamo theory. Section II is devoted to the fluid description of a quantum field. In it we find the velocity correlation function and velocity spectrum as well as the kinetic helicity correlation function. In section III\\ we characterize the turbulence by finding the Reynolds number, $Re\\left( k\\right) $, and the characteristic times $t_{d}$ and $t_{a}$. In section IV we find the amplification factor for the magnetic field. Finally in section V we summarize our conclusions. The bulk of the calculations that lead to these results are shown in the four appendices. ",
        "conclusions": "In this paper we have studied the possibility of turbulent dynamo action during reheating. We considered the presence of a charged scalar field minimally coupled to gravity. This field is in its invariant vacuum state during inflation. When the transition to reheating takes place the vacuum state turns into a many-particle state. For sub-horizon modes of the field, the number of created modes depends on the details of the transition. Therefore during reheating, besides the decay products of the inflaton we also have a plasma of scalar particles which is at rest in the comoving frame. We characterize the fluctuating velocities of this plasma giving their spatial two point correlation function and the kinetic energy associated to each Fourier mode of the stochastic velocity field, eq. (\\ref% {energia-k}). We evaluate the Reynolds number associated to each mode, $% Re\\left( \\kappa \\right) $, which turns out to depend on the physical parameters of the problem, namely $m$, $H$ and the duration $\\tau _{0}$ of the transition from inflation to reheating. If $\\tau _{0}$ is small enough, then there is a range of $\\kappa $ for which $Re\\left( \\kappa \\right) >1$ and the flow can be considered as (mildly) turbulent. As there are no stirring forces, the turbulence we refer to decays, each mode doing so in a characteristic time $t_{d}\\left( \\kappa \\right) $ given by eq. (\\ref{t-dec}% ). Besides as the plasma is a particle anti-particle one, each mode of the scalar field (not to be confused with modes of the stochastic velocity field) annihilates in a characteristic time $t_{a}\\left( \\kappa \\right) $ given by eq. (\\ref{t-aniq}). When comparing both times we find that annihilation dominates over decay, eq. (\\ref{2h}) and hence for practical purposes we can consider the turbulence as steady.  The sufficient condition to have a large scale kinematic dynamo is the flow to be endowed with kinetic helicity \\cite{moffatt}. The non-trivial result of this paper is that the scalar plasma does possess a non null rms kinetic helicity,eq. (\\ref{kin-hel-7}) From Figs. (\\ref{lambda-1}) and (\\ref% {lambda-2}) we see that, for the parameters for which $Re\\left( \\kappa \\right) >1$, the characteristic scale of the kinetic helicity is of the order of the particle horizon, thus allowing for kinematic dynamo action.  The existence of an rms helicity is due to the presence of the two scalar fields, $\\Phi $ and $\\Phi ^{\\dagger }$, as is evident from eq. (\\ref% {kin-hel-4}). Moreover, though the helicity may have either sign, in the average the amplification effect dominates. From the simplest model of kinematic dynamo, eq. (\\ref{mft-6}), we compute the amplification factor of an initial seed field, eq. (\\ref{am-2}), and find that for the physical parameters of the scenario considered, it is very small. In spite of this result, we believe our work shows the need for exploring the impact of nonlinear effects in the early universe. These effects offer the most natural answer to the riddle of the survival of the primordial magnetic field until the epoch of structure formation, in spite of the $1/a^{2}$ damping induced by the Hubble expansion."
    },
    "1004/1004.1446_arXiv.txt": {
        "abstract": "We compare the statistics of driven, supersonic turbulence at high Mach number using \\textsc{flash} a widely used Eulerian grid-based code and \\textsc{phantom}, a Lagrangian smoothed particle hydrodynamics (SPH) code at resolutions of up to $512^{3}$ in both grid cells \\emph{and} SPH particles. We find excellent agreement between codes on the basic statistical properties: a slope of $k^{-1.95}$ in the velocity power spectrum for hydrodynamic, Mach~10 turbulence, evidence in both codes for a Kolmogorov-like slope of $k^{-5/3}$ in the variable $\\rho^{1/3}{\\bf v}$ as suggested by \\citet{kritsuketal07} and a log-normal PDF with a width that scales with Mach number and proportionality constant $b=0.33-0.5$ in the density variance--Mach number relation. The measured structure function slopes are not converged in either code at 512$^{3}$ elements.  We find that, for measuring volumetric statistics such as the power spectrum slope and structure function scaling, SPH and grid codes give roughly comparable results when the number of SPH particles is approximately equal to the number of grid cells. In particular, to accurately measure the power spectrum slope in the inertial range, in the absence of sub-grid turbulence models, requires at least $512^{3}$ computational elements in either code. On the other hand the SPH code was found to be better at resolving dense structures, giving maximum densities at a resolution of $128^{3}$ particles that were similar to the maximum densities resolved in the grid code at $512^{3}$ cells, reflected also in the high density tail of the PDF. We find SPH to be more dissipative at comparable numbers of computational elements in statistics of the velocity field, but correspondingly less dissipative than the grid code in the statistics of density weighted quantities such as $\\rho^{1/3}{\\bf v}$.  For SPH simulations of high Mach number turbulence we find it important to use sufficient non-linear $\\beta$-viscosity in order to prevent particle interpenetration in shocks (we require $\\beta_{visc} = 4$ instead of the widely used default value, $\\beta_{visc} = 2$). ",
        "introduction": "\\label{sec:intro} Dense interstellar molecular clouds are ubiquitously observed to have non-thermal line widths implying supersonic internal motions \\citep{ze74}. Furthermore the amplitude of such motions increases with spatial scale in a manner reminiscent of turbulent flows in the laboratory \\citep{larson81,solomonetal87,hb04}. Understanding the nature and origin of such `supersonic turbulence'  is therefore key --- perhaps \\emph{the} key --- to understanding star formation \\citep{es04,mk04,mo07}.  Turbulence provides a natural explanation for the clustered and hierarchical nature of star formation \\citep{ElmegreenFalgarone1996}; the measured fractal dimension of interstellar gas \\citep[][and references therein]{kritsuketal07,FederrathKlessenSchmidt2009}; the few percent efficiency with which gas is converted into stars \\citep*{padoan95,vsbpk03,km05,elmegreen08}; most likely determines in large part the mass distribution of star forming cores (the core mass distribution, CMD) \\citep{bpetal06,dibetal08} and possibly the mass distribution of stars themselves, i.e., the Initial Mass Function (IMF) \\citep{pnj97,pn02,HennebelleChabrier2008}. However, given that a full theory of turbulence is elusive even in the incompressible regime apart from the phenomenology provided by \\citet{kolmogorov41}, one inevitably turns to numerical simulations to glean insight.  Given the importance of numerical simulations in understanding the basic statistics of supersonic turbulence, and the possible implications for star formation theory, it is crucial that results inferred from such simulations are robust with respect to different numerical methods and codes. This has motivated at least two major code comparison projects in the last year or so in which both of the present authors have been involved. The `Potsdam comparison' \\citep{kitsionasetal09} compared simulations of decaying, hydrodynamic turbulence using 7 different codes (3 SPH codes and 4 grid-based codes) at a fixed resolution ($215^{3}$ particles for the SPH codes, and $256^{3}$ grid cells for the grid codes). The results showed generally good agreement on statistics such as the density PDFs and power spectra, similar to earlier studies \\citep{maclowetal98,khm00}. In a similar spirit the KITP07 comparison\\footnote{http://kitpstarformation07.wikispaces.com/Star+Formation+Test+Problems} compared a large number of grid based and SPH calculations of decaying turbulence, for both hydrodynamics and MHD, at a range of resolutions, though the results are yet to be published.  Both of these comparisons are problematic in several respects. The first is that it is difficult to make a statistical comparison using decaying turbulence, since the time evolution is limited and therefore only a few instantaneous snapshots can be compared. Instantaneous snapshots however are subject to intermittent fluctuations that make a head-to-head comparison based on single time slices difficult \\citep{kritsuketal07,FederrathDuvalKlessenSchmidtMacLow2009}. Secondly, both comparisons start from evolved initial conditions, produced from either a previously driven SPH simulation (Potsdam) or a grid-based calculation (KITP) that has to be interpolated and/or downsampled to/from the grid/particles appropriate to the different codes, with an ensuing loss of accuracy and consistency before the comparison has even begun (this problem is much worse for the MHD case where differences in divergence-free representations for the magnetic field between codes is a further issue).  In this paper, we consider only two codes, an SPH code, \\textsc{phantom}, and a grid-based code, \\textsc{flash}, which we take to be broadly representative of the fundamentally different classes of code used for star formation studies (the codes are described in \\S\\ref{sec:numerics}). The turbulence in both codes is driven from stationary, uniform initial conditions with exactly the same energy input and driving pattern over multiple turbulent crossing times. We also consider a range of resolutions ($128^{3}$, $256^{3}$ and $512^{3}$ in both grid elements and the number of SPH particles) in order to estimate resolution requirements and establish where convergence has occurred in one code or the other, or neither. In the present work we limit ourselves to a study of hydrodynamic turbulence, that is, without magnetic fields. This is primarily because the algorithms for Magnetohydrodynamics in SPH currently being used for star formation studies \\citep[e.g.][]{pb08} rely on the Euler potentials formulation of the magnetic field, that cannot be used for turbulence studies over multiple crossing times due to the restricted field representations (see \\citealt{price10} for recent progress).  The goals of this paper are to: i) establish whether or not agreement can be found between SPH and grid codes on the basic statistics of supersonic turbulence; ii) define resolution criteria for various statistical measures of supersonic turbulence such as power spectra, PDFs and structure functions; and iii) establish the relative strengths and weaknesses of each method for turbulence studies. We discuss the numerical methods in \\S\\ref{sec:numerics}, with the Fourier space driving discussed in \\S\\ref{sec:driving}. Results from both codes are presented in \\S\\ref{sec:results} and our findings discussed in \\S\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  In this paper we have performed a detailed comparison of the statistics of supersonic turbulence using high resolution numerical simulations with two fundamentally different codes: \\textsc{flash}, an Eulerian grid-based code and \\textsc{phantom}, a Lagrangian particle-based SPH code. Despite the very different numerical methods we find in general very good agreement between the codes on the many aspects of supersonic turbulence, though it is clear that neither code shows results that are fully converged at $512^{3}$ except for a small scaling range in the velocity power spectrum with both codes and some indication of convergence near the peak and at the high density end of the PDF in the SPH case.  We find good agreement in the fact that hydrodynamic turbulence at Mach$~10$ has a velocity power spectrum with a slope around $E(k) \\propto k^{-1.95}$, very close to Burger's value of $E(k) \\propto k^{-2}$, confirming many previous results using only grid-based codes \\citep[e.g.][]{padoanetal07,FederrathDuvalKlessenSchmidtMacLow2009}. At the highest resolution employed, both codes support the idea that the power spectrum of the mixed quantity $\\rho^{1/3}{\\bf v}$ shows a Kolmogorov-like scaling in the power spectrum of $k^{5/3}$, as proposed by \\citet{kritsuketal07}. Both codes agree that the PDF of the logarithm of the density shows a distribution that is log-normal to good approximation for densities around $3-4$ orders of magnitude either side of the mean density, with the factor $b$ relating the width to the Mach number measured to lie between $b=0.33$ to $b=0.5$, with the converged value expected to lie somewhere around $b\\approx 0.4$. However, both codes also show that there are significant deviations from a log-normal distribution at high density, as also found by \\citet{FederrathDuvalKlessenSchmidtMacLow2009}.  Our conclusions regarding resolution and computational requirements are as follows: \\begin{enumerate} \\item For measuring volumetric statistics such as the power spectrum slope and structure function scaling, we find that SPH and grid codes give roughly comparable results when the number of SPH particles is approximately equal to the number of grid cells. This means that SPH codes are not well suited to measuring such quantities since for this kind of problem an SPH code will be significantly more expensive than a uniform-grid implementation at similar numbers of computational elements (we find the cost of the SPH calculations with \\textsc{phantom} similar to the cost of running the grid code (\\textsc{flash}) at twice the resolution, i.e., with $256^{3}$ particles $\\approx 512^{3}$ grid cells in terms of CPU time). \\item On the other hand the SPH code was found to be better at resolving dense structures, giving maximum densities at a resolution of $128^{3}$ particles that were similar to the maximum densities resolved in the grid code at $512^{3}$ cells, reflected also in the high density tail of the PDF. SPH is therefore a more efficient method in this regard, which is an important reason why it is frequently used for studying star formation (for which a grid based code requires adaptive mesh refinement, adding a significant computational overhead). \\item At comparable resolutions ($N_{part} \\approx N_{cells}$), SPH (\\textsc{phantom}) appears to be more dissipative (steeper dissipative tails in the power spectrum and structure functions) in the velocity field, but the reverse is true for the statistics of the density-weighted quantity $\\rho^{1/3} {\\bf v}$ where the grid (\\textsc{flash}) results appear more dissipative. We attribute the former to the greater sophistication of the shock capturing scheme in the grid-based code, and the latter to the better resolved power in the density field provided by the SPH code. \\item The absolute values of the structure function slopes are not converged in either grid or SPH at $512^{3}$, requiring much higher resolution in order to make a meaningful comparison with scaling models. \\item In order to accurately simulate supersonic turbulence in SPH it is important to ensure that sufficient $\\beta$-viscosity is applied to prevent particle interpenetration in shocks. At Mach~10 we require $\\beta_{visc} = 4$ instead of the usual $\\beta_{visc} = 2$. \\item We find that calculation of the sub-grid density field from the tracer particle distribution using an SPH summation can significantly enhance the resolution of high density structures from the grid-based results. However it is unclear whether or not the statistics of these sub-grid structures can be used in a meaningful way, because of the manner in which tracer particles tend to cluster in high density regions. \\end{enumerate}  The above conclusions mean that the differences between the SPH and grid based results discussed by \\citet{padoanetal07}, at least in terms of the power spectrum slope, can be understood as a consequence of the low resolution employed in the SPH calculations rather than being due to an intrinsic difficulty with the method. Given this, it is also likely that the conclusions drawn by \\citet{bpetal06} regarding the presence or otherwise of an emergent power law slope in the core mass distribution should also be treated with caution. As a follow-up to this paper it would be interesting to examine the properties of dense clumps (or `cores') in our calculations using a clump-finding algorithm in a similar manner to \\citet{bpetal06} and \\citet{padoanetal07} to see whether or not these differences can also be reconciled by improved resolution in the SPH results."
    },
    "1004/1004.4439_arXiv.txt": {
        "abstract": "We present synthetic spectral energy distributions (SEDs) for single-age, single-metallicity stellar populations (SSPs) covering the full optical spectral range at moderately high resolution (FWHM$=$2.3\\AA). These SEDs constitute our base models, as they combine scaled-solar isochrones with a empirical stellar spectral library (MILES), which follows the chemical evolution pattern of the solar neighbourhood. The models rely as much as possible on empirical ingredients, not just on the stellar spectra, but also on extensive photometric libraries, which are used to determine the transformations from the theoretical parameters of the isochrones to observational quantities. The  unprecedent stellar parameter coverage of the MILES stellar library allowed us  to safely extend our optical SSP SEDs predictions from intermediate- to  very-old age regimes, and the metallicity coverage of the SSPs from super-solar  to \\mbox{$\\mbox{[M/H]}=-2.3$}. SSPs with such low metallicities are particularly  useful for globular cluster studies. We have computed SSP SEDs for a suite of  IMF shapes and slopes. We provide a quantitative analysis of the dependence of  the synthesized SSP SEDs on the (in)complete coverage of the stellar parameter  space in the input library that not only shows that our models are of higher  quality than other works, but also in which range of SSP parameters our models  are reliable. The SSP SEDs are a useful tool to perform the stellar populations  analysis in a very flexible manner. Observed spectra can be studied by means of  full spectrum fitting or by using line indices. For the latter we propose a new Line Index System (named LIS) to avoid the intrinsic uncertainties associated with the popular Lick/IDS system and provide more appropriate, uniform, spectral  resolution. Apart from constant resolution as function of wavelength the system  is also based on flux-calibrated spectra. Data can be analyzed at three different  resolutions: 5\\,\\AA, 8.4\\,\\AA\\ and 14\\,\\AA\\ (FWHM), which are appropriate for studying globular cluster, low and intermediate-mass galaxies, and massive galaxies, respectively. Furthermore we provide polynomials to transform current Lick/IDS line index measurements to the new system. We provide line-index tables in the new system for various popular samples of Galactic globular clusters and galaxies. We apply the models to various stellar clusters and galaxies with high-quality spectra, for which independent studies are available, obtaining excellent results. Finally we designed a web page from which not only these models and stellar libraries can be downloaded but which also provides a suite of on-line tools to facilitate the handling and transformation of the spectra. ",
        "introduction": "To obtain information about unresolved stellar populations in galaxies, one can look at colours, ranging from the UV to the infrared. Such a method is popular especially when studying faint objects. Since colours are strongly affected by dust extinction, one prefers to use spectra when studying brighter objects. These have the additional advantage that abundances of various elements can be studied, from the strength of absorption lines (e.g. Rose 1985; S\\'anchez-Bl'azquez et al. 2003; Carretero et al. 2004). Although the UV and the near IR are promising wavelength regions to study, the large majority of spectral studies is done in the optical, mainly because of our better understanding of the region. To obtain information, one generally compares model predictions with measurements of strong absorption line strengths. This method, which is insensitive to the effects of dust extinction (e.g., MacArthur 2005) very often uses the Lick/IDS system of indices, which comprises definitions of 25 absorption line strengths in the optical spectral range, originally defined on a low resolution stellar library (FWHM$>$8-11\\,\\AA) (Gorgas et al. 1993, Worthey et al. 1994, Worthey \\& Ottaviani 1997). There is an extensive list of publications in which this method has been applied, mainly for early-type galaxies (see for example the compilation provided in Trager et al. 1998).   Predictions for the line-strength indices of the integrated light of stellar clusters and galaxies are obtained with stellar population synthesis models. These models calculate galaxy indices by adding the contributions of all possible stars, in proportions prescribed by stellar evolution models. The required indices of all such stars are obtained from observed stellar libraries, but since the spectra in these libraries were often noisy, and did not contain all types of stars present in galaxies, people used the so-called fitting functions, which relate measured line-strength indices to the atmospheric parameters ($T_{\\mbox{\\scriptsize eff}}$, $\\log g$, \\mbox{$\\mbox{[Fe/H]}$}) of library stars. The most widely used fitting functions are those computed on the basis of the Lick/IDS stellar library (Burstein et al. 1984; Gorgas et al. 1993; Worthey et al. 1994) There are, however, alternative fitting functions in the optical range (e.g., Buzzoni 1995; Gorgas et al. 1999; Tantalo, Chiosi \\& Piovan 2007; Schiavon 2007) and in other spectral ranges (e.g., Cenarro et al. 2002; M\\'armol-Queralt\\'o et al. 2008).  The most important parameters that can be obtained using stellar population synthesis are age (as a proxy for the star formation history) and metallicity (the average metal content).  There are several methods to perform stellar population analysis based on these indices. By far the most popular method is to build key diagnostic model grids by plotting an age-sensitive indicator (e.g. H$\\beta$) versus a metallicity-sensitive indicator (e.g., Mg$b$, $\\langle\\mbox{Fe}\\rangle$), to estimate the age and metallicity (e.g. Trager et al. 2000; Kuntschner et al. 2006). There are other methods that, e.g., employ as many Lick indices as possible  and simultaneous fit them in a $\\chi^2$ sense (e.g., Vazdekis et al. 1997; Proctor et al. 2004). An alternative method is the  use of Principal Component Analysis (e.g. Covino et al. 1995; Wild et al. 2009).    In the last decade the appearance of a generation of stellar population models that predict full spectral energy distributions (SEDs) at moderately high spectral resolution has provided new means for improving the stellar population analysis (e.g., Vazdekis 1999, hereafter V99; Bruzual \\& Charlot 2003; Le Borgne et al. 2004). These models employ newly developed extensive empirical stellar spectral libraries with flux-calibrated spectral response and good atmospheric parameter coverage. Among the most popular stellar libraries are those of Jones (1999), CaT (Cenarro et al. 2001a), ELODIE (Prugniel \\& Soubiran 2001), STELIB (Le Borgne et al. 2003), INDO-US (Valdes et al. 2004), and MILES (S\\'anchez-Bl\\'azquez et al. 2006d, hereafter Paper~I). Models that employ such libraries are, e.g., Vazdekis (1999), Vazdekis et al. (2003) (hereafter V03), Bruzual \\& Charlot (2003), Le Borgne et al. (2004). Alternatively, theoretical stellar libraries at high spectral resolutions have also been developed for this purpose (e.g., Murphy \\& Meiksin 2004; Zwitter et al. 2004; Rodr\\'{\\i}guez-Merino et al. 2005; Munari et al. 2005; Coelho et al. 2005; Martins \\& Coelho 2007). Models that use such libraries are, e.g., Schiavon et al. (2000), Gonz\\'alez-Delgado et al. (2005), Coelho et al. (2007).   There are important limitations inherent to the method that prevent us from easily disentangling relevant stellar population parameters. The most important one is the well known age/metallicity degeneracy, which not only affects colours but also absorption line-strength indices (e.g. Worthey 1994). This degeneracy is partly due to  the isochrones and partly due to the fact that line-strength indices change with  both age and metallicity. The effects of this degeneracy are stronger when low resolution indices are used, as the metallicity lines appear blended. Alternative indices that were thought to work at higher spectral resolutions, such as those of Rose (1985, 1994), have been proposed to alleviate the problem.   New indices with greater abilities to lift the age/metallicity degeneracy have been proposed with the aid of the new full-SED models (e.g., Vazdekis \\& Arimoto 1999; Bruzual \\& Charlot 2003; Cervantes \\& Vazdekis 2009). These models allow us to analyze the whole information contained in the observed spectrum at once. In fact, there is a growing body of full spectrum fitting algorithms that are being developed for constraining and recovering in part the Star Formation History (e.g., Panter et al. 2003; Cid Fernandes et al. 2005; Ocvirk et al. 2006ab; Koleva et al. 2008). Furthermore the use of these SSP SEDs, which have sufficiently high spectral resolution, has been shown to significantly improve the analysis of galaxy kinematics both in the optical (e.g., Sarzi et al. 2006) and in the near-IR (e.g., Falc\\'on-Barroso et al. 2003), for both absorption and emission lines.  Here we present single-age, single-metallicity stellar population (SSP) SEDs based on the empirical stellar spectral library MILES that we presented in Paper~I) and Cenarro et al. (2007, hereafter Paper~II). These models represent an extension of the V99 SEDs to the full optical spectral range. These MILES-models are meant to provide better predictions in the optical range for intermediate- and old-stellar populations. In Paper~I we provided all relevant details for MILES, which has been obtained at the 2.5~m Issac Newton Telescope (INT) at the Observatorio del Roque de Los Muchachos, La Palma. The library is composed of 985 stars covering the spectral range $\\lambda\\lambda$ 3540-7410\\,\\AA\\ at 2.3\\,\\AA\\ (FWHM). MILES stars were specifically selected for population synthesis modeling. In Paper~II we present a homogenized compilation of the stellar parameters ($T_{\\mbox{\\scriptsize eff}}$, $\\log g$, \\mbox{$\\mbox{[Fe/H]}$}) for the stars of this library. In fact the parameter coverage of MILES constitutes a significant improvement over previous stellar libraries and allows the models to safely extend the predictions to intermediate-aged stellar populations and to lower and higher metallicities.  The models that we present here are based on an empirical library, i.e. observed stellar spectra, and therefore the synthesized SEDs are imprinted with the chemical composition of the solar neighbourhood, which is the result of the star formation history experienced by our Galaxy. As the stellar isochrones  (Girardi et al. 2000) -- i.e. the other main ingredient feeding the population synthesis code -- are solar-scaled, our models are self-consistent,  and scaled-solar for solar metallicity. In the low metallicity regime, however, our models combine scaled-solar isochrones with stellar spectra that do not show this abundance ratio pattern (e.g., Edvardsson et al. 1993; Schiavon 2007). In a second paper we will present self-consistent models, both scaled-solar and $\\alpha$-enhanced,  for a range in metallicities. For this purpose we have used MILES, together with theoretical model atmospheres, which are coupled to the appropriate stellar isochrones.  We note, however, that the use of our (base) models, which are not truly scaled-solar for low metallicities, does not affect in a significant manner the metallicity and age estimates obtained from galaxy spectra in this metallicity regime (e.g., Michielsen et al. 2008). For the high metallicity regime, where massive galaxies reside and show enhanced \\mbox{$\\mbox{[Mg/Fe]}$} ratios, the observed line-strengths for various popular indices (e.g., Mg$b$, Fe5270) are significantly different from the scaled-solar predictions. It has been shown that by changing the relative element-abundances, and making them $\\alpha$-enhanced, both the stellar models (e.g., Salaris, Groenewegen \\& Weiss 2000) and  the stellar atmospheres (e.g. Tripicco \\& Bell 1995; Korn, Maraston \\& Thomas 2005;  Cohelo et al. 2007) are affected. However it has also been shown that base models, such as the ones we are presenting here, can be used to obtain a good proxy for the \\mbox{$\\mbox{[Mg/Fe]}$} abundance ratio, if appropriate indices are employed for the analysis (e.g. Yamada et al. 2006). In fact we confirm the results of, e.g., S\\'anchez-Bl\\'azquez et al. (2006b), de la Rosa et al. (2007) and Michielsen et al. (2008), who obtain a linear relation between the proxy for \\mbox{$\\mbox{[Mg/Fe]}$}, obtained with scaled-solar models, and the abundance ratio estimated with the aid of models that specifically take into account non-solar element mixtures (e.g., Tantalo et al. 1998; Thomas et al. 2003; Lee \\& Worthey 2005; Graves \\& Schiavon 2008).  In Section \\ref{sec:models} we describe the main model ingredients, the details of the stellar spectral library MILES and the steps that have been followed for its implementation in the stellar population models. In Section \\ref{sec:MILES_SSP_SEDs} we present the new single-age, single-metallicity spectral energy distributions, SSP SEDs, and provide a quantitative analysis for assessing their quality. In Section \\ref{sec:system} we show line-strengths indices as measured on the newly synthesized SSP spectra and propose a new reference system, as an alternative to the Lick/IDS system, to work with these indices. In Section \\ref{sec:colours} we provide a discussion about the colours measured on the new SSP SEDs. In Section \\ref{sec:applications} we apply these models to a set of representative stellar clusters and galaxies to illustrate the capability of the new models. In Section \\ref{sec:web} we present a webpage to facilitate the use and handling of our models. Finally in Section \\ref{sec:conclusions} we provide a summary. ",
        "conclusions": "\\label{sec:conclusions}  Here we present stellar population model SEDs based on the empirical stellar spectral library MILES . In Paper~I we have presented and described this library, whereas in Paper~II we have published the stellar parameter determinations. Here we present synthetic SEDs for single-age, single-metallicity stellar populations (SSPs) covering the full optical spectral range, 3540.5--7409.6\\,\\AA, at moderately high resolution, FWHM$=$2.3\\,\\AA, which is virtually constant as a function of wavelength (see Fig.~4 of Paper~I). The SSP SEDs have a reliable, accurate flux calibration, which is one of the major advantages of the MILES library. Furthermore the redder part of the SSP spectra do not show telluric residuals, as these were carefully cleaned from the stellar spectra as shown in Paper~I.  We consider the SEDs presented here as our base models, as they combine an empirical library, which is imprinted with the chemical composition of the solar neighbourhood, with scaled-solar stellar isochrones (Girardi et al. 2000). Therefore, whereas our models are self-consistent, and scaled-solar for solar metallicity, this is not the case for the low metallicity regime, as the observed stars there do not show this abundance pattern (see Schiavon 2007). In a second paper we will present self-consistent models, both scaled-solar and $\\alpha$-enhanced, for a range in metallicities. For this purpose we have used MILES, together with theoretical model atmospheres, which are coupled to the appropriate stellar isochrones.  An interesting feature of our base models is that they rely, as much as possible, on empirical ingredients. Apart from using MILES stellar spectra, we also use transformations of the theoretical parameters of the isochrones to observational, measurable, quantities that are based on extensive empirical, photometric libraries, with marginal dependence on model atmospheres. The implementation of MILES in the population synthesis code has been done very carefully: each star of the library has been compared with other MILES stars with similar atmospheric parameters, using the same interpolation algorithm used in the synthesis code. As a result of these tests and other considerations, we have removed or decreased the weight with which a given star contributes to the synthesis of a stellar spectrum for 135 (out of 985) stars of MILES.  The unprecedented stellar parameter coverage of the MILES library has allowed us to safely extend our optical SSP SEDs from intermediate to very old age regimes (0.06--18\\,Gyr) and the metallicity coverage of the SSPs from supersolar (\\mbox{$\\mbox{[M/H]}=+0.22$}) to \\mbox{$\\mbox{[M/H]}=-2.32$}. The very low metallicity SSPs should be of particular interest for globular cluster studies. In addition, we have computed SSP SEDs for a suite of IMF shapes and slopes: unimodal and bimodal IMF shapes as defined in Vazdekis et al. (1996) with slopes varying in the range 0.3-3.3 (Salpeter: unimodal with slope 1.3), and the segmented IMFs of Kroupa (2001) (Universal and Revised cases).  In addition to our reference models, which adopt the MILES temperatures published in Paper~II, we provide an alternative set of models that adopt the recently published temperature scale of Gonz\\'alez-Hern\\'andez \\& Bonifacio (2009) (i.e., about $\\sim$51\\,K hotter). The most relevant feature of these models is that the well known model zero-point problem affecting the age estimates obtained from the Balmer lines, i.e. the obtained ages are larger than the CMD-derived ages, is partially alleviated.  To constrain the range of ages, metallicities and IMFs where the use of these SEDs are safe, we provide a quantitative analysis. This analysis, which basically takes into account the parameter coverage of the stellar library feeding the models, clearly shows that our SSP SEDs can be safely used in the age range 0.06--18\\,Gyr, and for all the metallicities, except for \\mbox{$\\mbox{[M/H]}=-2.32$}, for which only stellar populations with ages above $\\sim$10\\,Gyr can be considered safe. This applies to the standard IMF shapes. However our quantitative analysis has shown that the quality of the SSP SEDs decreases as we increase the IMF slope, especially for the Unimodal IMF. For example, for slopes above 1.8, the synthesized SSP SEDs for \\mbox{$\\mbox{[M/H]}=-2.32$} are no longer safe for a Unimodal IMF. According to this analysis our predictions are of significantly higher quality than any previous model predictions in the literature  for all the ages and metallicities that have been computed in this work. For example, any predictions based on the standard fitting functions of the Lick/IDS system (Worthey et al. 1994), e.g. line-strength indices, with which most stellar population studies have been performed so far, should not be considered safe for \\mbox{$\\mbox{[M/H]}<-1.3$} (for standard IMF shapes). This limit is set at higher metallicity, \\mbox{$\\mbox{[M/H]}\\sim-0.7$}, for predictions based on STELIB (e.g., Bruzual \\& Charlot 2003). The comparison of our new models, based on MILES, with the SSP SEDs published in V99, reveals systematic differences due to limitations in the flux calibration of the Jones (1999) stellar library that feeds the latter models. As for the line-strengths the most important difference between these two models is found for the H$\\beta$ index, which shows a lower metallicity sensitivity in the V99 models. We show an example where this effect has a non-negligible impact on our interpretation of galaxy age/metallicity gradients.  The SSP SEDs increase our ability to perform the stellar populations analysis, which can be done in a very flexible manner. Observed spectra can be studied by fitting the full spectrum or by focussing on selected line indices, either standard or newly defined. An example of the previous method can be found in Koleva et al. (2008), where various stellar populations models were compared following this approach. Our favourite approach is to analyze the spectrum of a galaxy at the resolution imposed by its velocity dispersion (and its instrumental resolution),  which require us to smooth the SSP SEDs to match the observed total resolution. Therefore, the line index analysis to obtain relevant stellar population parameters is performed in the system of the galaxy, either through full spectrum fitting or line index approaches.  As many authors prefer to plot the line-strength measurements of their whole galaxy sample in a single index-index diagnostic diagram, or to compare their data to tabulated indices in the literature, we propose a new Line Index System (LIS), which makes this method straightforward to apply. With LIS we avoid the well known uncertainties inherent to the widely employed Lick/IDS system. Unlike the latter, LIS has two main advantages, i.e., a constant resolution as function of wavelength and a universal flux-calibrated spectral response. Data can be analyzed in this system at three different resolutions: 5\\,\\AA, 8.4\\,\\AA\\ and 14\\,\\AA\\ (FWHM), i.e. $\\sigma=$127, 214 and 357\\,${\\rm km\\,s}^{-1}$ at 5000\\,\\AA, respectively. These resolutions are appropriate for studying globular cluster, low and intermediate-mass galaxies, and massive galaxies, respectively. Line-strength measurements given at a higher LIS resolution can be smoothed to match a lower LIS resolution.  Furthermore we provide polynomials to transform current Lick/IDS line-index measurements in the literature to the new system. We provide LIS line-index tables for various popular samples of Galactic globular clusters and galaxies. We also show various popular index-index diagnostic diagrams for these samples in the LIS system.  As an application, we have fitted a number of representative stellar clusters of varying ages and metallicities with our models, obtaining good agreement with CMD determinations. Unlike for the open cluster M\\,67, our SED fits for a sample of Galactic globular clusters show non negligible residuals blueward of 4300\\,\\AA, which mostly reflect the characteristic CN-strong features of these clusters among other deviations from the scaled-solar pattern. We also applied our models to representative galaxies with high quality spectra, for which independent studies are available, obtaining good results. We show that our base models can be used for studying line-strength indices of galaxies with $\\alpha$-enhanced element partitions. Examples of such a use can be found in, e.g., Vazdekis et al. (2001b), Kuntschner et al. (2002), Carretero et al. (2004), Yamada et al. (2006). The method consists in plotting a highly sensitive age indicator, such as those of Vazdekis \\& Arimoto (1999), or the recently defined H${\\beta_o}$ index of Cervantes \\& Vazdekis (2009) versus Mg (e.g., Mg$b$) and versus Fe (e.g., Fe4383, $\\langle\\mbox{Fe}\\rangle$), obtaining a virtually orthogonal model grid where the estimated age does not depend on the metallicity indicator in use. Unlike the age, for a \\mbox{$\\mbox{[Mg/Fe]}$} enhanced galaxy the obtained metallicities differ when plotted against, for example,  Mg$b$ and $\\langle\\mbox{Fe}\\rangle$ indices. This metallicity difference, [Z$_{{\\rm Mg}b}$/Z$_{<{\\rm Fe}>}$], can be used as a good proxy for the abundance ratio determined with the aid of stellar population models that specifically take into account non-solar element partitions. Although this proxy yields larger \\mbox{$\\mbox{[Mg/Fe]}$} values, there is a linear relation between these two ways of estimating the abundance ratios (e.g., S\\'anchez-Bl\\'azquez et al. 2006b; de la Rosa et al. 2007; Michielsen et al. 2008). In the second paper of this series we will present both $\\alpha$-enhanced and scaled-solar self consistent SSP SEDs, all based on MILES, but modified with the aid of model atmospheres.  By applying these models to galaxies with varying abundance ratios, for which high S/N spectra are available, we confirm the linear relation mentioned above. There is one advantage  in using this proxy: the abundance ratios determined in this way do not depend  on the specific details of modeling of $\\alpha$-enhanced stellar populations such as those that are implicit to the computations of the atmosphere, or the adopted element mixtures.  Our fits to the stellar cluster and galaxy spectra of high quality have confirmed the high precision of the flux calibration of our SSP SEDs. Furthermore the colours that we obtained from our SSP SEDs are consistent with the ones that we derived via the employed photometric stellar libraries, within typical zero point uncertaintes. Such agreement shows the reliability of the temperatures adopted for the MILES stars (Paper~II), which are consistent with the temperature scales of Alonso et al. (1996,1999). Such accuracy opens new applications for these models. For example, the model SEDs can be used as templates for codes aiming at determining spectrophotometric redshifts, particularly for those based on narrow-band filters (e.g. ALHAMBRA, PAU).  Finally we present a webpage from which not only these models and stellar libraries can be downloaded, but where we also provide a suite of user-friendly webtools to facilitate the handling and transformation of the spectra. For example, once the user enters the details of the instrumental setup employed in the observations (e.g., PSF, sampling), the users can obtain their favourite line strength indices and diagnostic diagrams, ready to plot their observational measurements.  Apart of the model SEDs shown here, the webpage also provides predictions for a suite of observable quantities."
    },
    "1004/1004.2053_arXiv.txt": {
        "abstract": "We update constraints on cosmic opacity by combining recent SN Type Ia data compilation with  the latest measurements of the Hubble expansion at redshifts between 0 and 2.  The new constraint on the parameter $\\epsilon$ parametrising deviations from the luminosity-angular diameter distance relation ($d_L=d_A(1+z)^{2+\\epsilon}$), is $\\epsilon=-0.04_{-0.07}^{+0.08}$ (2-$\\sigma$). For the redshift range between $0.2$ and $0.35$ this corresponds to an opacity $\\Delta\\tau<0.012$ ($95\\%$ C.L.), a factor of 2 stronger than  the previous constraint. Various models of beyond the standard model physics that predict violation of photon number conservation contribute to the opacity and can be equally constrained. In this paper we put new limits on axion-like particles, including chameleons, and mini-charged particles. ",
        "introduction": "Introduction}  Cosmological observations provide constraints on different distance measures: luminosity distance (as provided e.g., by supernovae), angular diameter distance (as provided e.g., by baryon acoustic oscillations) and even on the expansion rate or the Hubble parameter as a function of redshift $z$. Both luminosity distance and angular diameter distance are functions of the Hubble parameter. While combining these measurements helps to break parameter degeneracies and constrain cosmological parameters, comparing them helps to constrain possible deviations from the assumptions underlying the standard cosmological model (e.g. isotropy), or to directly constrain physics beyond the standard model of particle physics (e.g. couplings of photons to scalar or pseudo-scalar matter).  The Etherington relation \\cite{Etherington1} implies that,  in a cosmology based on a metric theory of gravity,  distance measures are unique: the luminosity distance is $(1+z)^2$ times the  angular diameter distance. This is valid in any cosmological background where photons travel on null geodesics and  where, crucially, photon number is conserved.  There are several scenarios in which the Etherington relation would be violated: for instance we can have  deviations from a metric theory of gravity, photons not traveling along unique null geodesics,  variations of fundamental constants, etc. In this paper we want to restrict our attention on violations of the Etherington relation arising from the violation of photon conservation.  A change in the photon flux during propagation towards the Earth  will affect the Supernovae (SNe) luminosity distance measures but not the determinations of the angular diameter distance. Photon conservation can be violated by simple astrophysical effects or by exotic physics. Amongst the former we find, for instance, attenuation due to interstellar dust, gas and/or plasmas. Most known sources of attenuation are expected to be clustered and can be typically constrained down to the 0.1\\% level \\cite{Menard, Bovy}. Unclustered sources of attenuation are however much more difficult to constrain.  For example, gray dust \\cite{Aguirre}  has been invoked to explain the observed dimming of Type Ia Supernovae without  resorting to cosmic acceleration.  More exotic sources of photon conservation violation involve a coupling of photons to particles  beyond the standard model of particle physics.  Such couplings would mean that, while passing through the intergalactic medium, a photon could disappear --or even (re)appear!-- interacting with such  exotic particles, modifying the apparent luminosity of sources. Here we consider the mixing of photons with scalars, known as axion-like particles, and the possibility of mini-charged particles which have a tiny, and unquantised electric charge. A recent review \\cite{Jaeckel:2010ni} highlights the rich phenomenology of these weakly-interacting-sub-eV-particles (WISPs), whose effects have been searched for in a number of laboratory experiments and astronomical observations. In particular, the implications of this particles on the SN luminosity have been described in a number of publications~\\cite{Csaki2,Mortsellaxions,Mirizzi:2006zy,Burrage:2007ew,Ahlers:2009kh}.  One of the most interesting features of these models is that the exotic opacity involved could  in principle ``mimic\" the value of a non-zero cosmological constant inferred from SNe measurements. However, this possibility can already be excluded (at least in the simplest WISP models) by the absence of distortions in the CMB or the spectra of quasars for axion-like-particles, and by arguments of stellar evolution in the case of mini-charged particles.  In this paper we use improved bounds on cosmic opacity to further constrain the existence of exotic particles which can couple to the photon. The rest of the paper is organised as follows. In section 2 we update constraints on transparency from the latest available data. In section 3 we discuss the implications of this for axion-like particles and chameleons, and in section 4 we consider mini-charged particles. We then forecast, in section 5, how the constraints will improve with distance measures from future, planned and proposed, surveys. We conclude in section 6.  Sections 3 and 4 discuss in detail the motivation, modelling  and  regime of applicability of the beyond the standard model physics we consider. Readers with a more focused interest on cosmology may concentrate on the beginning of section 3, sub-sections 3.4, 3.5 and figures 2, 3, 4, 5, 6.  Appendix A summarises the cosmologically-relevant results of sections 3 and 4.  Precursors of this paper can be found in~\\cite{bassettkunz1,bassettkunz2,Song:2005af,More,AVJ}. ",
        "conclusions": "If new particles from physics beyond the standard model couple to photons then the propagation of light may be altered.  In this paper we have focused on two scenarios for exotic particles  which can significantly modify the propagation of photons as they pass through magnetic fields.    Measurements of cosmic opacity are a strong tool to constrain such scenarios, as interactions between photons and exotic particles in the magnetic fields of the intergalactic medium leads to a new source of cosmic opacity. Uniform deviations from cosmic transparency (i.e. opacity) can be constrained through their effects on distance duality, by parameterizing possible deviations from the Etherington relation. The Etherington relation  implies that, in a cosmology based on a metric theory of gravity, distance measures are unique: the luminosity distance is $(1 + z)^2$ times the angular diameter distance. Both luminosity distance and angular diameter distance depend on the Hubble parameter $H(z)$, but this relation is valid in any cosmological background where photons travel on null geodesics and where, crucially, photon number is conserved. We have   restricted our attention on violations of the Etherington relation arising from the  violation of photon conservation.  We have combined direct measurements of cosmic expansion (from the latest determinations of the Hubble parameter)  at redshifts $0<z<2$ and recent SN data yielding the luminosity distance.  SN-inferred luminosity distances are affected by violation of photon conservation, but the $H(z)$ measurements  we use are not. Assuming an underlying flat $\\Lambda$CDM model, we have placed tight limits on possible deviations from photon-conservation. Photon-conservation can be violated by simple astrophysical effects which give uniform attenuation such as gray dust. We have reported updated constraints on this effect.  More exotic sources of photon-conservation violation involve a coupling of photons to particles beyond the standard model of particle physics. We have focused on axion-like particles, new scalar or pseudo scalar fields which couple to the kinetic terms of photons, and mini-charged particles which are hidden sector particles with  a tiny electric charge.  Photons passing through intergalactic magnetic fields may be lost by pair production of light mini-charged particles.  If the mixing between axion-like particles and photons is significant, then interactions in the intergalactic magnetic fields will also lead to a loss of photons due to conversion into ALPs.  However if the coupling between photons and ALPs is sufficiently strong, one-third of any initial flux will be converted into ALPs, and two-thirds into photons, resulting in a redshift-independent dimming of supernovae which we cannot constrain or exclude with cosmic opacity bounds.  The improved measurement of the cosmic opacity found here leads to improved bounds on these exotic physics scenarios which are summarised in Fig. \\ref{sumario}.  Future measurements of baryon acoustic oscillations, and an increase in the number of observations of high redshift supernovae will lead to further improvements in the constraints on physics beyond the standard model.  \\begin{figure}[h] \\begin{center} \\includegraphics[width=3in]{boundALPpro.eps} \\includegraphics[height=3in,width=3in]{boundMCPpro.eps} \\caption{\\label{sumario} Prospects of the future bounds on ALPs (left) and MCPs (right). ALPs: The dark region corresponds to the current bounds while gray and light grey regions cover the possible future bounds using EUCLID and EUCLID+SNAP data. Taking into account the redshift dependence of the background improves these bounds in the region covered by the lines. MCPs: Current bounds (solid) and prospects from EUCLID (dashed) and EUCLID+SNAP (dotted).} \\end{center} \\end{figure}"
    },
    "1004/1004.0783_arXiv.txt": {
        "abstract": "{} { We probe the dependence of $\\alpha^2/\\mu$ on the ambient matter density by means of submm- and mm-wave bands spectral observations in the Milky Way. } {A procedure is suggested to explore the value of $F = \\alpha^2/\\mu$, where $\\mu = m_{\\rm e}/m_{\\rm p}$ is the electron-to-proton mass ratio, and $\\alpha = e^2/(\\hbar c)$ is the fine-structure constant. The fundamental physical constants, which are measured in different physical environments of high (terrestrial) and low (interstellar) densities of baryonic matter are supposed to vary in chameleon-like scalar field models, which predict that both masses and coupling constant may depend on the local matter density. The parameter $\\Delta F/F = (F_{\\rm obs} - F_{\\rm lab})/F_{\\rm lab}$ can be estimated from the radial velocity offset, $\\Delta V = V_{\\rm rot} - V_{\\rm fs}$, between the low-laying rotational transitions in carbon monoxide $^{13}$CO and the fine-structure transitions in atomic carbon [\\ion{C}{i}]. A model-dependent constraint on \\daa\\ can be obtained from \\dff\\ using \\dmm\\ independently measured from the ammonia method. } { Currently available radio astronomical datasets provide an upper limit on $|\\Delta V| < 110$ \\ms\\ ($1\\sigma$). When interpreted in terms of the spatial variation of $F$, this gives $|\\Delta F/F| < 3.7\\times10^{-7}$. An order of magnitude improvement of this limit will allow us to test independently a non-zero value of \\dmm = $(2.2\\pm0.4_{\\rm stat}\\pm0.3_{\\rm sys})\\times10^{-8}$ recently found with the ammonia method. Taking into account that the ammonia method restricts the spatial variation of $\\mu$ at the level of $|\\Delta\\mu/\\mu| \\leq 3\\times10^{-8}$ and assuming that \\dff\\ is the same in the entire interstellar medium, one obtains that the spatial variation of $\\alpha$ does not exceed the value $|\\Delta\\alpha/\\alpha| < 2\\times10^{-7}$. Since extragalactic gas clouds have densities similar to those in the interstellar medium, the bound on \\daa\\ is also expected to be less than $2\\times10^{-7}$ at high redshift if no significant temporal dependence of $\\alpha$ is present. } {} ",
        "introduction": "\\label{sect-1}  The dimensionless physical constants like the electron-to-proton mass ratio, $\\mu = m_{\\rm e}/m_{\\rm p}$, or the fine-structure constant, $\\alpha = e^2/(\\hbar c)$, are expected to be dynamical quantities in modern extensions of the standard model of particle physics (Uzan 2003;  Garcia-Berro \\etal\\ 2007; Martins 2008; Kanekar 2008; Chin \\etal\\ 2009). Exploring these predictions is a subject of many high precision measurements in contemporary laboratory and astrophysical experiments. The most accurate laboratory constraints on temporal $\\alpha$- and $\\mu$-variations of $\\dot{\\alpha}/\\alpha = (-1.6\\pm2.3)\\times10^{-17}$ yr$^{-1}$, and $\\dot{\\mu}/\\mu = (1.6\\pm1.7)\\times10^{-15}$ yr$^{-1}$ were obtained by Rosenband \\etal\\ (2008), and Blatt \\etal\\ (2008), respectively.  In case of monotonic dependence of $\\alpha(t)$ and $\\mu(t)$ on cosmic time, at redshift $z \\sim 2$ (corresponding look-back time is $\\Delta t \\sim 10^{10}$ yr) the changes of $\\alpha$ and $\\mu$ would be restricted at the level of $|\\Delta\\alpha/\\alpha| < 4\\times10^{-7}$ and $|\\Delta\\mu/\\mu| < 3\\times10^{-5}$. Here \\daa\\ (or \\dmm) is a fractional change in $\\alpha$ between a reference value $\\alpha_1$ and a given measurement $\\alpha_2$ obtained at different epochs or at different spatial coordinates: $\\Delta\\alpha/\\alpha = (\\alpha_2 - \\alpha_1)/\\alpha_1$.  These constraints are in line with geological measurements of relative isotopic abundances in the Oklo natural fission reactor which allows us to probe $\\alpha(t)$ at $\\Delta t \\sim 2\\times10^9$ yr ($z \\sim 0.4$). Assuming possible changes only in the electromagnetic coupling constant, Gould \\etal\\ (2006) obtained a model dependent constraint on $|\\Delta\\alpha/\\alpha| < 2\\times10^{-8}$. However, when the strength of the strong interaction,~-- the parameter $\\Lambda_{QCD}$,~-- is also suggested to be variable, the Oklo data does not provide any bound on the variation of $\\alpha$ (Flambaum \\& Shuryak 2002; Chin \\etal\\ 2009).  Current astrophysical measurements at higher redshifts are as follows. There was a claim for a variability of $\\alpha$ at the 5$\\sigma$ confidence level: $\\Delta\\alpha/\\alpha = -5.7\\pm1.1$ ppm (Murphy \\etal\\ 2004)\\footnote{Hereafter, 1 ppm = $10^{-6}$.}, but this was not confirmed in other measurements which led to the upper bound $|\\Delta\\alpha/\\alpha| < 2$ ppm (Quast \\etal\\ 2004; Levshakov \\etal\\ 2005; Srianand \\etal\\ 2008; Molaro \\etal\\ 2008a).  Measurements of the cosmological $\\mu$-variation exhibit a similar tendency. Non-zero values of $\\Delta\\mu/\\mu = -30.5\\pm7.5$ ppm, $\\Delta\\mu/\\mu =  -16.5\\pm7.4$ ppm (Ivanchik \\etal\\ 2005), and $\\Delta\\mu/\\mu = -24\\pm6$ ppm (Reinhold \\etal\\ 2006) found at $z = 2.595$ (Q 0405--443) and $z = 3.025$ (Q 0347--383) from the Werner and Lyman bands of H$_2$ were later refuted by Wendt \\& Reimers (2008), King \\etal\\ (2008) and Thompson \\etal\\ (2009) who used the same optical absorption-line spectra of quasars and restricted changes in $\\mu$ at the level of $|\\Delta\\mu/\\mu| < 6$ ppm. The third H$_2$ system at $z = 2.059$ towards the quasar J2123--0050 also does not show any evidence for cosmological variation in $\\mu$: $\\Delta\\mu/\\mu = -5.6\\pm5.5_{\\rm stat}\\pm2.9_{\\rm sys}$ ppm (Malec \\etal\\ 2010). More stringent constraints were obtained at lower redshifts from radio observations of the absorption lines of NH$_3$ and other molecules: $|\\Delta \\mu/\\mu| < 1.8$ ppm at $z = 0.68$ (Murphy \\etal\\ 2008), and $|\\Delta \\mu/\\mu| < 0.6$ ppm at $z = 0.89$ (Henkel \\etal\\ 2009). Two cool gas absorbers at $z = 1.36$ (Q 2337--011) and $z = 1.56$ (Q 0458--020) were recently studied in the \\ion{H}{i} 21cm and \\ion{C}{i}$\\lambda\\lambda1560, 1657$ absorption lines providing a constraint on the variation of the product $X = g_{\\rm p}\\alpha^2\\mu$ (here $g_{\\rm p}$ is the proton gyromagnetic ratio): $\\Delta X/X = -6.8\\pm1.0_{\\rm stat}\\pm6.7_{\\rm sys}$ ppm (Kanekar \\etal\\ 2010). Thus, the most accurate astronomical estimates restrict cosmological variations of the fundamental physical constants at the level of $\\sim$1-2 ppm.  The estimate of fractional changes in \\daa\\ and \\dmm\\ by spectral methods is always a measurement of the relative Doppler shifts between the line centers of different atoms/molecules and their comparison with corresponding laboratory values (Savedoff 1956; Bahcall \\etal 1967; Wolfe \\etal\\ 1976; Dzuba 1999, 2002; Levshakov 2004; Kanekar \\& Chengalur 2004). To distinguish the line shifts due to radial motion of the object from those caused by the variability of constants, lines with different sensitivity coefficients, ${\\cal Q}$, to the variations of $\\mu$ and/or $\\alpha$ are to be used\\footnote{${\\cal Q}$ is a dimensionless coefficient showing a relative change of the atomic transition frequency $\\omega_i$ in response to a change of the physical constant $F$: $\\Delta \\omega_i/\\omega_i = {\\cal Q}_i \\Delta F/F$.}. It is clear that the larger the difference $|\\Delta {\\cal Q}|$ between two transitions, the higher the accuracy of such estimates.  Optical and UV transitions in atoms, ions and molecular hydrogen H$_2$ have similar sensitivity coefficients with $|\\Delta Q|$ not exceeding 0.05 (Varshalovich \\& Levshakov 1993; Dzuba 1999, 2002; Porsev \\etal\\ 2007). For atomic spectra, the estimate of \\daa\\ is given in linear approximation ($|\\Delta\\alpha/\\alpha| \\ll 1$) by (e.g., Levshakov \\etal\\ 2006): \\begin{equation} \\frac{\\Delta \\alpha}{\\alpha} \\approx \\frac{(V_2 - V_1)}{2c({\\cal Q}_1 - {\\cal Q}_2)} \\equiv \\frac{\\Delta V}{2c\\Delta {\\cal Q}}\\, , \\label{eq1} \\end{equation} where $V_1, V_2$ are the radial velocities of two atomic lines, and $c$ is the speed of light. It was shown in Molaro \\etal\\ (2008b) that the limiting accuracy of the wavelength scale calibration for the VLT/UVES quasar spectra at any point within the whole optical domain is about 30 \\ms, which corresponds to the limiting relative accuracy between two lines measured in different parts of the same spectrum of about 50 \\ms. Taking into account that $|\\Delta Q| \\simeq 0.05$, it follows from Eq.(\\ref{eq1}) that the limiting accuracy of $\\Delta\\alpha/\\alpha$ is 2 ppm, which is the utmost value that can be achieved in observations of extragalactic objects with present optical facilities.  A considerably higher sensitivity to the variation of physical constants is observed in radio range. For example, van Veldhoven \\etal\\ (2004) first showed that the inversion frequency of the $(J,K) = (1,1)$ level of the ammonia isotopologue $^{15}$ND$_3$ has the sensitivity coefficient ${\\cal Q}_\\mu = 5.6$. Compared to optical and UV transitions, the ammonia method proposed by Flambaum \\& Kozlov (2007) provides 35 times more sensitive estimate of \\dmm\\ from measurements of the radial velocity offset between the NH$_3$ $(J,K) = (1,1)$ inversion transition at 23.7 GHz and low-lying rotational transitions of other molecules co-spatially distributed with NH$_3$: \\begin{equation} \\frac{\\Delta \\mu}{\\mu} \\approx 0.289\\frac{\\Delta V}{c}\\ . \\label{eq2} \\end{equation} The ammonia method was recently used to explore possible spatial variations\\footnote{Hereafter, the term `spatial variation' means a possible change in $\\mu$ between its terrestrial and interstellar values.} of physical constants from observations of prestellar molecular cores in the Taurus giant molecular cloud (Levshakov \\etal\\ 2010, hereafter L10), the Perseus cloud, the Pipe Nebula, and Infrared dark clouds (Levshakov \\etal\\ 2008b; Molaro \\etal\\ 2009). In contrast to the mentioned above laboratory constraints on {\\it temporal} variations, this method reveals a tentative {\\it spatial} variation of \\dmm\\ at the level of \\dmm = $(2.2\\pm0.4_{\\rm stat}\\pm0.3_{\\rm sys})\\times10^{-8}$ (L10). The corresponding conservative upper limit in this case is equal to $|\\Delta \\mu/\\mu| \\leq 3\\times10^{-8}$.   In the present paper we consider fractional changes of a combination of two constants $\\alpha^2$ and $\\mu$, $F = \\alpha^2/\\mu$, which are estimated from the comparison of transition frequencies measured in different physical environments of {\\it high} (terrestrial) and {\\it low} (interstellar) densities of baryonic matter. The idea behind this experiment is that some class of scalar field models~--- so-called chameleon-like fields~--- predict the dependence of both masses and coupling constant on the local matter density (Olive \\& Pospelov 2008; Upadhye \\etal\\ 2010). Chameleon-like scalar fields were introduced by Khoury \\& Weltman (2004a,b) and by Brax \\etal\\ (2004) to explain negative results on laboratory searches for the fifth force which should arise inevitably from couplings between scalar fields and standard model particles. The chameleon models assume that a light scalar field acquires both an effective potential and effective mass because of its coupling to matter that depends on the ambient matter density. In this way, the chameleon scalar field may evade local tests of the equivalence principle and fifth force experiments since the range of the scalar-mediated fifth force for the terrestrial matter densities is too small to be detected. Similarly, laboratory tests with atomic clocks for $\\alpha$-variations are performed under conditions of constant local density and, hence, they are not sensitive to the presence of the chameleon scalar field (Upadhye \\etal\\ 2010). This is not the case for space-based tests, where the matter density is considerably lower, an effective mass of the scalar field is negligible, and an effective range for the scalar-mediated force is large. Light scalar fields are usually attributed to a negative pressure substance permeating the entire visible Universe and known as dark energy (Caldwell \\etal\\ 1998). This substance is thought to be responsible for a cosmic acceleration at low redshifts, $z \\la 1$ (Peebles \\& Rata 2003; Brax 2009). ",
        "conclusions": "\\label{sect-4}   The level of 0.2 ppm  represents a model-dependent upper limit on the {\\it spatial} variations of $\\alpha$. Under model-dependence we assume here that both \\dff\\ and \\dmm\\ do not change significantly from cloud to cloud, since astrophysical measurements of these parameters are made in low density regions of the interstellar medium with $\\rho_{\\rm cloud} \\ll \\rho_{\\rm  terrestrial}$.  For comparison, the upper limit on the {\\it temporal} $\\alpha$-variation obtained from high-redshift quasar absorbers is $|\\Delta \\alpha/\\alpha| < 2$ ppm (Sect.~1). If dependence of constants on the ambient matter density dominates over temporal (cosmological), as suggested in chameleon-like scalar field models, then one may expect that $|\\Delta\\alpha/\\alpha| < 0.2$ ppm at high redshifts as well, since quasar absorbers have gas densities similar to those in the interstellar clouds. Taking into account that the predicted changes in $\\alpha$ and $\\mu$ are not independent and that $\\mu$-variations may exceed variations in $\\alpha$ (e.g., Calmet \\& Fritzsch 2002; Langacker \\etal\\ 2002; Dine \\etal\\ 2003; Flambaum \\etal\\ 2004), even a lower bound of $|\\Delta\\alpha/\\alpha| \\leq 0.03$ ppm is conceivable within the framework of the chameleon models.  We note that if a theoretical prediction $|\\Delta\\alpha/\\alpha| \\ll |\\Delta\\mu/\\mu|$ is valid, then $\\Delta F/F \\approx -\\Delta\\mu/\\mu$, and, hence, the $F$-estimate with a further order of magnitude improvement in sensitivity will provide an independent test of the tentative change of $\\mu$.  The factors limiting accuracy of the current estimate of \\dff\\ at $z = 0$ are a relatively low spectral resolution of the available observations in submm- and mm-wave bands, a rather large uncertainty of the rest frequencies of the [C\\,{\\sc i}] FS lines, and a small number of objects observed in both [C\\,{\\sc i}] and $^{13}$CO transitions.  Modern telescopes like the recently launched Herschel Space Observatory can provide for Galactic objects the spectral resolution as high as 30 \\ms\\ (e.g., the Heterodyne Instrument for the Far Infrared, HIFI, has resolving power $R = 10^7$). This means that the positions of the [\\ion{C}{i}] FS lines can be measured with the uncertainty of $\\sim$3 \\ms. In the near future, high precision measurements will be also available with the Atacama Large Millimeter/submillimeter Array (ALMA), the Stratospheric Observatory For Infrared Astronomy (SOFIA), the Cornell Caltech Atacama Telescope (CCAT) and others. Thus, any further advances in exploring $\\Delta F/F$ depend crucially on new laboratory measurements of the [C\\,{\\sc i}] FS frequencies. If these frequencies will be known with uncertainties of a few \\ms, then the parameter $\\Delta F/F$ can be probed at the level of $10^{-8}$ which would be comparable with the non-zero signal in the spatial variation of the electron-to-proton mass ratio $\\mu$."
    },
    "1004/1004.5359_arXiv.txt": {
        "abstract": "We make use of a 500~ks Chandra HRC-S/LETG spectrum of the blazar H~2356-309, combined with a lower S/N (100 ks) pilot LETG spectrum of the same target, to search for the presence of warm-hot absorbing gas associated with two Large--Scale Structures (LSSs) crossed by this sightline, and to constrain its physical state and geometry. Strong (logN$_{OVII} \\ge 10^{16}$ cm$^{-2}$) OVII K$\\alpha$ absorption associated with a third LSS crossed by this line of sight (the Sculptor Wall, SW), at $z=0.03$, has already been detected in a previous work. Here we focus on two additional prominent filamentary LSSs along the same line of sight, at z=0.062 (the Pisces-Cetus Supercluster, PCS) and at z=0.128 (the ``Farther Sculptor Wall'', FSW).  The combined LETG spectrum has a S/N of $\\sim 11.6-12.6$  per resolution element in the $20-25$ \\AA, and an average 3$\\sigma$ sensitivity to intervening OVII K$\\alpha$ absorption line equivalent widths of EW$_{OVII} \\gs 14$ m\\AA\\ in the available redshift range ($z<0.165$). No statistically significant (i.e. $\\ge 3\\sigma$) individual absorption is detected from any of the strong He- or H-like transitions of C, O and Ne (the most abundant metals in gas with solar-like composition) at the redshifts of the PCS and FSW structures, and down to the above EW thresholds. However we are still able to constrain the physical and geometrical parameters of the putative absorbing gas associated with these structures, by performing joint spectral fit of various marginal detections and upper limits of the strongest expected lines with our self-consistent hybrid ionization WHIM spectral model.  At the redshift of the PCS we identify a warm phase with $\\rm \\log{T} =5.35_{-0.13}^{+0.07} ~K$ and $\\rm \\log{N_H} = (19.1 \\pm 0.2) ~cm^{-2}$ possibly coexisting with a much hotter and statistically less significant phase with $\\rm \\log{T} = 6.9^{+0.1}_{-0.8} ~K$ and $\\rm \\log{N_H} = 20.1^{+0.3}_{-1.7} ~cm^{-2}$ ($1\\sigma$ errors). These two separate physical phases are identified through, and mainly constrained by, CV K$\\alpha$ (warm phase) and OVIII K$\\alpha$ (hot phase) absorption, with single line significances of $1.5\\,\\sigma$ each.  For the second LSS, at $z\\simeq 0.128$, only one hot component is hinted in the data, through OVIII K$\\alpha$ ($1.6\\,\\sigma$) and NeIX K$\\alpha$ ($1.2\\,\\sigma$). For this system, we estimate $\\rm \\log{T} =6.6_{-0.2}^{+0.1} ~K$ and $\\rm \\log{N_H} = 19.8_{-0.8}^{+0.4} ~cm^{-2}$.  Our column density and temperature constraints on the warm-hot gaseous content of these two LSSs, combined with the measurements obtained for the hot gas permeating the SW, allow us to estimate the cumulative number density per unit redshifts of OVII WHIM absorbers at 3 different equivalent width thresholds of $0.4 \\rm m\\AA$, $7 \\rm m\\AA$ and $25.8 \\rm m\\AA$. This is consistent with expectations only at the very low end of EW thresholds, but exceed predictions at 7m\\AA\\ and 25.8 m\\AA (by more than $2 \\sigma$). We also estimate the cosmological mass density of the WHIM based on the 4 absorbers we tentatively detect along this line of sight, obtaining $\\Omega_b^{WHIM} = (0.021^{+0.031}_{-0.018}) (Z/Z_{\\odot})^{-1}$, consistent with the cosmological mass density of the intergalactic 'missing baryons' only if we assume high metallicities ($Z\\sim Z_{\\odot}$). ",
        "introduction": "Cosmological hydrodynamical simulations predict the gradual formation of a local ($\\rm{z}<1$) web of low density ($n_b=10^{-6}-10^{-5}\\,cm^{-3}$), warm-hot ($\\rm{T}=10^5-10^7\\, \\rm{K}$) intergalactic gas, connecting virialized halos (i.e. galaxies, galaxy groups and clusters of galaxies), permeating the large--scale structures (LSS) of which these systems are constituents, and ultimately providing the necessary fuel for their continuous growth \\citep{cen95,cen,dave,cen06}. This low-redshift intergalactic gas is largely the same primordial gas present at redshift higher than $\\sim 2$ in a cool photo-ionized phase (the so called 'Lyman-$\\alpha$ Forest)', but in a much hotter and metal-enriched phase, because of efficient shock-heating during the continuous process of LSS assembly and growth in a non-linear Universe, and of strong feedback with the same structures for which it provides building blocks. Due to its high temperatures this IGM phase has been dubbed Warm-Hot Intergalactic Medium (WHIM).  Electrons and baryons in the WHIM are shock-heated during their infall in the dark matter LSS potential well, and settle in filamentary/sheet-like structures surrounding LSSs. Such matter is predicted to account for a sizable fraction ($\\sim50\\%$) of all the baryons in the local ($z<1$) universe, and it is thus considered the best candidate to host the baryons seen at high redshift and missing from the low redshift census \\citep{fukugita98,fukugita}.  Given its high temperature the WHIM can only emit or absorb in the Far UV and soft X-rays, mainly through Li- through H-like metal transitions and bremsstrahlung continuum emission. However, at WHIM densities both line and continuum emission are highly depressed (due to the dependency of these mechanisms on the square power of the emitters volume densities), and on average well below the sensitivity of current instrumentations. Nonetheless, statistical techniques based on cross correlation of regions with excess diffuse X-ray emission in ROSAT, {\\em Chandra} and XMM-{\\em Newton} deep exposures, with large--scale galaxy distribution, have probably already allowed the marginal detection of the density peak of the WHIM distribution \\citep[e.g.][]{scharf,zappacosta,zappacosta2,werner}.  A far more promising way to detect the WHIM is through resonant absorption by highly ionized metals. Intervening WHIM filaments should imprint a 'forest', i.e. the so called X-ray Forest \\citep{hellsten,perna}, by analogy with the HI Ly~$\\alpha$ Forest copiously seen at $z\\gtrsim 2$, of metal absorption lines onto the spectra of bright background sources, whose strength depends only linearly on the absorber density, and is therefore less suppressed than the corresponding emission. Predicted equivalent widths (EWs) from the most abundant ions in WHIM, range from 1 m\\AA\\ to $\\lesssim 20$ m\\AA\\ for the densest filaments. The detection of even the strongest of such absorption lines (probing only the very high density tail of the WHIM distribution: overdensity $ \\delta = n_b/<n_b> \\gtrsim 300$, compared to the average density of the Universe: $<n_b>= 2 \\times 10^{-7}  (1+z)^3 (\\Omega_bh^2/0.02)$, where $\\Omega_b$ is the baryonic density parameter and $h$ is the Hubble constant in units of 100~km/s/Mpc) is extremely challenging with the limited sensitivity ($A_{eff}<60cm^2$) and resolution (R=$E/\\Delta E=400-800$) of the current {\\em Chandra} and XMM-{\\em Newton} high resolution X-ray gratings. A $\\ge 3\\sigma$ detection of an EW=20 m\\AA\\ absorption line, requires spectra of the background targets with S/N$\\gtrsim 8$ per 50 m\\AA\\ resolution element (i.e. $\\gtrsim 70$ net counts per bin). These can only be obtained, for the brightest ($\\sim 10^{-11} erg\\,s^{-1}\\,cm^{-2}$) soft X-ray  ($\\rm0.5-2\\,keV$) targets in the extragalactic sky (preferably blazars, because of their intrinsically featureless spectra), with $\\gtrsim 0.5$ Ms  exposures, while in quiescence, or $\\gtrsim 100$~ks exposures in outburst.  Such dense WHIM filaments are rare. According to hydrodynamical simulations, at $z\\simeq 0$ one expects $\\le 0.05$ WHIM filaments with EW(OVII$_{K\\alpha}$)$\\ge 20$ m\\AA, per unit redshift. The probability of having one of such filaments along a random line of sight up to $z=0.3$, is therefore only $1.5$\\% \\citep{gehrels}, and only few $z\\ge 0.3$ targets with quiescent F(0.5-2 keV)$\\ge 10^{-11}$ erg s$^{-1}$ cm$^{-2}$ are available (e.g. Conciatore et al., 2009, priv. comm.). The number density of OVII WHIM absorbers per unit redshift increases by almost two orders of magnitude by lowering the EW(OVII$_{K\\alpha})$ detection threshold down to $\\ge 2$ m\\AA\\, so dramatically increasing the probability  of randomly detecting one of such filaments to P($z=0.3$)=97.5\\%. This observational strategy has the advantage to probe, in theory, the bulk of the WHIM mass distribution, but requires incredibly high S/N spectra ($\\ge 75$ per resolution element with the {\\em Chandra} LETG, and $\\ge 180$ per resolution element with the XMM-{\\em Newton} RGS, for a $\\ge 3\\sigma$ detection), obtainable only with several Ms exposures on the brightest possible $z\\ge 0.3$ targets in their quiescent states. Moreover, even when such high S/N spectra are obtained \\citep[for example by observing background target during exceptionally high outburst:][]{nicastro05apj,nicastro05} the clear assessment of the statistical significance of such weak 2-3 m\\AA\\ absorption lines is hampered by our limited knowledge of the instrument systematics, which is comparable to the statistical uncertainties on the lines EWs \\citep{kaastra,rasmussen07,nicastro08}   An alternative observational strategy is to select lines of sight where the probability to cross dense WHIM filament is enhanced \\citep[e.g.][]{zapproposal,zapproposal2}. WHIM gas density and metallicity is predicted \\citep[and possibly in the UV observed][]{stocke} to correlate with LSS galaxy overdensity \\citep{viel}. Thus, chances of intercepting a dense WHIM filament may be improved by carefully selecting bright sources in the background of extreme filamentary LSS concentrations (Fig.~\\ref{LSS}). This observational strategy has been successfully exploited recently in \\citet{buote09}, who detected a strong (EW$\\sim30 m\\AA$) absorption line in two XMM-{\\em Newton} RGS and {\\em Chandra} LETG spectra of the blazar H~2356-309 (z=0.165), and tentatively identified it with OVII K$\\alpha$ at a redshift consistent with that spanned by the intervening Sculptor Wall \\citep[SW;][]{daCosta}.  \\begin{figure*}[!] \\begin{center} \\includegraphics[scale=0.7]{f1_wa.ps} \\end{center} \\caption{Sky map and wedge diagram of the region of the Sculptor Wall where the blazar H~2356-309 is located. The upper sky map refer to the FSW and the lower to the PCS. Galaxies (black point) clusters and groups of galaxies (magenta circles) in the wall are taken from NED. The wedge diagram show galaxies inside the dashed blue box drawn in the sky map and report the sightline to the blazar as red dashed line. The galaxy and cluster catalogs contain objects belonging to different parent catalogs. Hence we point out that some holes in the projected galaxy distribution are artificially caused by this non homogeneity like the hole at north-east of the blazar position visible at $\\rm RA=0.1 hours$ and $\\rm DEC=-29.5 deg$ in the PCS sky map.} \\label{LSS} \\end{figure*}  Here, and in our companion paper \\citep{fang10}, we report on the successful extension of this observational program. H~2356-309 has been recently re-observed with the {\\em Chandra} LETG for 500 ks, as part of an approved cycle 10 GO observational program. The main goal of this deep observation was to confirm, at higher significance, the SW OVII detection \\citep{fang10}. Secondary objectives, were to confirm the presence of other (lower significance) lines from the same absorber (Buote et al., 2010 in prep.). Here instead we focus on constraining the physical parameters of the putative WHIM gas content of two additional galaxy LSSs present along this line of sight, at $z=0.062$ (the Pisces-Cetus Supercluster; PCS) and $z=0.128$ (a farther wall which we will call Farther Sculptor Wall; FSW). In this paper we use all the existing {\\em Chandra} LETG data of H~2356-309, to characterize the physical properties of the WHIM permeating these additional structures, and conclude by estimating the contribution of such dense gaseous component of LSSs, to the WHIM cosmological mass density. In \\S\\ref{sightline} we describe the richness of LSSs along the line of sight to H~2356-309. In \\S\\ref{datareduction} and \\S\\ref{analysis} we present the data and describe their reduction and analysis. \\S\\ref{discussion} is devoted to a critical discussion of our findings. In \\S\\ref{conclusions} we summarize our conclusions. Throughout the paper we adopt a $\\Lambda$-CDM cosmology, with $h=0.71$, $\\Omega_M=0.27$, $\\Omega_{\\Lambda}=0.73$. ",
        "conclusions": "We have analyzed a long integration (600 ksec) Chandra HRC-S/LETG spectrum of the blazar H~2356-309 to search and characterize the WHIM in two large--scale galaxy structures along this line of sight and, more specifically, the ``Pisces-Cetus Supercluster'' (PCS) at z=0.062 and the ``Farther Sculptor Wall'' (FSW) at z=0.128. These are structures more prominent and more distant with respect to the ``Sculptor Wall'', at z=0.03, where WHIM was identified in previous works through the detection of the OVII absorption line.  Although we do not detect significant individual absorption lines in the PCS nor in the FSW, the joint analysis of the marginally detected lines (as well as of stringent upper limits) through a self-consistent hybrid ionization spectral model, allow us to constrain the physics of the WHIM in the two farther superstructures. The main results are summarized in the following:  \\begin{itemize}  \\item At the redshift of the PCS we identify two distinct phases: a warm phase, with $\\rm \\log{T}= 5.35^{+0.07}_{\u22120.13}~K$ and $\\rm log{N_H} = (19.1 \\pm 0.2)~cm^{\u22122}$, and a much hotter less significant phase, with $\\rm log{T} = 6.9^{+0.1}_{\u22120.8}~ K$ and $\\rm log{N_H} = 20.1^{+0.3}_{-1.7}~cm^{-2}$ (1$\\sigma$ errors).  \\item At the redshift of the FSW only one hot component is hinted in the data, with $\\rm log{T} = 6.6^{+0.1}_{\u22120.2}~ K$ and $\\rm log{N_H}= 19.8^{+0.4}_{-0.8}~ cm^{-2}$.  \\item Under the assumption that the absorbers are embedded in their galaxy superstructures having baryonic column densities $N_b$ equal to the $-1\\sigma$ $N_H$ value, and have dense cores extending $<$4~Mpc along the line of sight, we can infer conservative lower limits on the baryons densities in the two systems for the three absorbers. More specifically, for the two PCS WHIM phases we get $\\delta > 2.7 (Z/Z_{\\odot})^{-1}$ and $\\delta > 0.9 (Z/Z_{\\odot})^{-1}$, while for the hot FSW phase we obtain $\\delta > 2.8 (Z/Z_{\\odot})^{-1}$, all consistent with predicted WHIM overdensities.  \\item By combining the constraints on the OVII absorbers in the PCS and in the FSW, with the previous detection in the SW \\citep{fang10}, we derive the cumulative number density of OVII absorbers per unit redshift, as a function of the EW(OVII). While at low equivalent widths (EW(OVII)$>$0.4~m\\AA) the absorbers number density is fully consistent with the predictions of hydrodynamical simulations, at EW(OVII)$>$10~m\\AA \\ the inferred absorbers number density exceeds significantly the theoretical predictions. The latter finding may result from a line of sight selection bias, or may reflect the possibility that one or more of the detected absorptions do not arise in WHIM filaments but in galaxy halos. We considered also the possibility that the two most marginally detected systems (Hot-PCS and FSW) may be just statistical fluctuations and not absorbers. In this case the number density of OVII absorbers per unit redshift would agree with the predictions within the large uncertainties.  \\item Finally, by combining the measurements in all absorbers we derive a cosmological mass density of the WHIM of $\\rm \\Omega _b^{WHIM}=(0.021^{+0.031}_{-0.018})(Z/Z_{\\odot})^{-1}$, consistent with the cosmological mass density of intergalactic missing baryons in the local universe. Yet, if the WHIM metallicities are $<$10\\% Solar, then the resulting $\\rm \\Omega _b^{WHIM}$ is significantly in excess of the missing baryon density, again possibly indicating a line of sight selection bias, or reflecting the possibility that one or more absorbers are not associated with WHIM. The exclusion of the two most marginally detected systems may bring in agreement the $\\rm \\Omega _b^{WHIM}$ value if we assume $0.1\\rm Z_\\odot$.    \\end{itemize}"
    },
    "1004/1004.3749_arXiv.txt": {
        "abstract": "{In galaxy clusters the entropy distribution of the IntraCluster Plasma modulates the latter's equilibrium within the Dark Matter gravitational wells, as rendered by our Supermodel. We argue the entropy production at the boundary shocks to be reduced or terminated as the accretion rates of DM and intergalactic gas peter out; this behavior is enforced by the slowdown in the outskirt development at late times, when the Dark Energy dominates the cosmology while the outer wings of the initial perturbation drive the growth. In such conditions, we predict the ICP temperature profiles to steepen into the cluster outskirts. The detailed expectations from our simple formalism agree with the X-ray data concerning five clusters whose temperature profiles have been recently measured out to the virial radius. We predict steep temperature declines to prevail in clusters at low $z$, tempered only by rich environs including adjacent filamentary structures.} ",
        "introduction": "Galaxy clusters constitute the largest bound structures in the Universe, with their masses up to $M\\sim 10^{15}\\, M_{\\odot}$ and outskirts extending out to sizes $R\\sim$ a few Mpcs. These set the \\emph{interface} between the intergalactic environment keyed to the cosmology at large, and the confined intracluster plasma (ICP). The latter pervades the clusters at temperatures $k_B T\\propto GM/R\\sim 5$ keV and number densities $n\\sim 10^{-3}$ cm$^{-3}$, and so emits copious X-ray powers mainly by thermal Bremsstrahlung \\citep[see][]{Sarazin1988}. The ICP coexists with the gravitationally dominant dark matter (DM) component in the baryonic fraction $m/M$ close to the cosmic value $0.16$, and the two build up together from accretion across the cluster boundary.  The build up comprises an early collapse of the cluster body, tailing off into a secular development of the outskirts by smooth accretion and minor mergers \\citep[see][]{Zhao2003, Diemand2007, Vass2008, Navarro2010}. In radius, the body ranges out to $r\\sim r_{-2}$ where the slope of the DM density run $n(r)$ equals $-2$; the adjoining outskirts extend out to the current virial radius $R$ with steepening density.  In time, the transition is marked by the redshift $z_t$; thereafter $r_{-2}$ stays put while $R$ grows larger in a quasi-static, self-gravitating DM equilibrium \\citep[described through the Jeans equation, see ][]{Lapi2009a}, to imply for the standard concentration parameter $c\\equiv R/r_{-2}$ observed values $c\\approx 3.5\\, H(z_t)/H(z_{\\rm obs})$ in terms of the Hubble parameter $H(z)$. In the following we adopt the standard flat cosmology \\citep[see][]{Dunkley2009}. So values of $c$ ranging from $3$ to $10$ correspond for $z_{\\rm obs}\\approx 0-0.2$ to \\emph{young} or to \\emph{old} dynamical cluster ages $z_t\\sim 0.2-3$. The concentration can be directly if laboriously probed with gravitational lensing \\citep[see][]{Broadhurst2008, Lapi2009b}.  Secular accretion of DM goes along with inflow of intergalactic gas. The ensuing ICP equilibrium is amenable to the powerful yet simple description provided by the Supermodel\\footnote{\\textsl{IDL} and \\textsl{FORTRAN} algorithms which implement the Supermodel and run in a fraction of a second on a standard laptop, can be found at \\textsl{http://people.sissa.it/$\\sim$lapi/Supermodel/}.} \\citep[SM; see][hereafter CLFF09]{Cavaliere2009}.  Clearly, inflows into the cluster outskirts are exposed to the cosmological grip. This is the focus of the present paper. ",
        "conclusions": "These data clearly bear out our expectations of steep runs for $T(r)$, stemming from outer entropy production being reduced as the accretion rates $\\dot M$ of DM and intergalactic gas peter out (see \\S~2). We trace back such a behavior to two concurring sources: (i) the cosmological slowdown in the growth of outskirts developing at late cosmic times; (ii) shallow perturbation wings scantily feeding the outskirts growth in average or poor environs.  Under such conditions, we show that on average the entropy profiles are to \\emph{flatten} out or even \\emph{decline} into the cluster outskirts as represented by Eq.~(4). Then our Supermodel yields \\emph{steep} outer profiles of projected temperatures (and flatter densities), in close agreement with current data out to $R$.  We also expect the cosmological decrease of the accretion rates to be locally offset in rich environs \\emph{biased} high, or in cluster \\emph{sectors} adjacent to filamentary large-scale structures; more standard runs of temperature are predicted there. These loom out in a sector of A1689 as observed by \\citet{Kawaharada2010}.  Compared with numerical simulations like the ones by \\citet{Nagai2007} and \\citet{Roncarelli2006}, our temperature declines and entropy dearths are comparable as for clusters like A2204 and A1413; but they are quite sharper for low $z$ clusters like A1795 and PKS0745-191, well beyond possible contaminations due to \\textsl{Suzaku} PSF smearing \\citep[see][]{Reiprich2009}. Low-noise, high-resolution simulations addressing the issue of entropy production related to richness of surrounding environment or adjoining filamentary structures will be most fruitful.  We have checked that with the SM the gas mass $m(<r)$ grows monotonically outwards, and actually faster than the DM's, so as to produce an increasing baryonic fraction by a factor $10$ from $R/20$ to $R$ \\citep[cf.][]{Zhang2010}. So sharply declining $T(r)$ and increasing masses are in fact consistent with thermal HE. Note that as $\\dot{M}$ decreases to the point that the infall velocities decrease to transonic values, the shocks weaken, thermalization becomes inefficient (see CLFF09), the entropy production is terminated, and thermal pressure alone cannot support HE any longer.  Equilibrium in the outskirts can be helped by bulk or turbulent, merger-induced motions; these contribute up to $10-20\\%$ of the total support in relaxed clusters, as gauged in terms of pressure and X-ray masses both observationally \\citep[][adding to the refs. in \\S~3]{Mahdavi2008, Zhang2010} and numerically \\citep[][]{Nagai2007, Piffaretti2008, Bode2009, Lau2010, Meneghetti2010}. A similar accuracy is intended for our predictions based on thermal HE, while a finer approximation with non-thermal contributions included in the SM will be provided elsewhere.  On the other hand, even extended HE constitutes just a useful approximation, due to fail even in relaxed clusters: locally, when minor lumps of cold gas fall into the cluster from an adjacent filament, along with smooth accretion; globally, beyond $R$ where the DM equilibrium is in jeopardy.  In sum, from the SM with thermal HE and optimal energy conversion in strong shocks, we find temperature profiles declining \\emph{sharply} outwards. These stem from progressive \\emph{exhaustion} of mass inflow, especially for clusters in average or poor environments at \\emph{low} $z_{\\rm obs}$; such a predicted trend is consistent with the current evidence for the handful of clusters in Figs.~2 and 3.  Our punchline is that cosmology - besides affecting the cluster statistics like mass or temperature distributions \\citep[see][]{Vikhlinin2009} - concurs with the perturbation shape to set the outer structure of \\emph{individual} clusters and their development.  All these rich phenomena expected at the \\emph{interface} between the ICP and the intergalactic medium call for extensive probing even at $z\\ga 0.2$ with the next generation of X-ray telescopes planned to detect at high resolutions low-surface brightness plasma \\citep[see][]{Giacconi2009}, like \\textsl{WFXT} (see \\textsl{http://wfxt.pha.jhu.edu/}) and eventually \\textsl{IXO} (see \\textsl{http://ixo.gsfc.nasa.gov/})."
    },
    "1004/1004.4675_arXiv.txt": {
        "abstract": "We present 1D numerical simulations, which consider the effects of radiative cooling and gravity on the hydrodynamics of the matter reinserted by stellar winds and supernovae within young nuclear starbursts with a central supermassive black hole (SMBH). The simulations confirm our previous semi-analytic results for low energetic starbursts, evolving in a quasi-adiabatic regime, and extend them to more powerful starbursts evolving in the catastrophic cooling regime. The simulations show  a bimodal hydrodynamic solution in all cases. They  present a  quasi-stationary accretion flow onto the black hole, defined by the matter reinserted  by massive stars within the stagnation volume and a stationary starburst wind, driven by the high thermal pressure acquired in the region between the stagnation and the starburst radii. In the catastrophic cooling regime, the stagnation radius rapidly approaches the surface of the starburst region, as one considers more massive starbursts. This leads to larger accretion rates onto the SMBH and concurrently to powerful winds able to inhibit interstellar matter from approaching the nuclear starburst. Our self-consistent model thus  establishes a direct physical link between the SMBH accretion  rate and the nuclear star formation activity of the host galaxy  and provides a good upper limit to the accretion rate onto the central black hole. ",
        "introduction": "Powerful starbursts have been conclusively detected in the  nuclear regions of galaxies with active galactic nuclei (AGN; see for a review \\citet{2005ARA&A..43..769V} and \\citet{2009and..book..335H}). These include quasars \\citep{2005ApJ...625...78H, 2008ChJAA...8...12H}, Seyferts \\citep{2003ApJ...599..918I, 2007ApJ...671.1388D, 2008ApJ...677..895W, 2009ApJ...695L.130C}, submillimeter galaxies with an extreme star formation rate \\citep{2005ApJ...632..736A, 2009Natur.457..699W} and even low luminosity AGNs \\citep{1995MNRAS.272..423C, 2004ApJ...605..127G}. Some  supermassive black holes (SMBHs) coincide with massive and compact nuclear star clusters \\citep{2008ApJ...678..116S}, some of them with a complex history of star formation are found in a number of nearby spiral \\citep{2006AJ....132.1074R} and elliptical galaxies \\citep{2006ApJ...644L..17W, 2006ApJS..165...57C}.  The interplay between young nuclear starbursts and their central SMBHs is not well understood and remains one of the central issue in  the theory of AGN galaxies, and in the process of cosmological growth of SMBHs \\citep{2006MNRAS.368.1001L, 2009MNRAS.398...53B} and  their co-evolution with the bulges of their host galaxies \\citep{2005MNRAS.361.1387B, 2005Natur.433..604D, 2008MNRAS.391..481S}. In fact, even basic issues regarding, for example, the impact of type II supernovae (SN) on the matter left over from star formation, seem to be  still undecided. One can find in the literature massive starbursts with a SN rate of 1 yr$^{-1}$ structuring a gaseous disk of just 10$^7$ M$_\\odot$ \\citep{2002ApJ...566L..21W} while other calculations assume that type II supernova might evacuate most of the nuclear region from gas and dust \\citep{2009MNRAS.393..759S} and have considered only the evolution after the end of the type II SN epoch.  \\citet{2008ApJ...686..172S} presented a self-consistent, spherically symmetric stationary solution for the gaseous flow around a SMBH at the center of a young ($\\le 40$ Myr) starburst region. They also found a threshold mechanical luminosity ($L_{crit}$) which separates systems evolving in the quasi-adiabatic regime with a small stagnation radius, $R_{st}$  (at which the velocity is zero km s$^{-1}$) and consequently a low accretion rate,  from those with a large $R_{st}$ and a high accretion rate  whose hydrodynamics are strongly affected by  radiative cooling (see their Figure 6). The self-consistent semi-analytic solution found by \\citet{2008ApJ...686..172S} is only valid  for nuclear starbursts with a mechanical luminosity $L_{NSB}<L_{crit}$. Above the threshold luminosity, strong radiative cooling drives the gas thermally unstable.   This has a strong impact on the thermal pressure gradient and thus on the  location of the stagnation radius, facts which inhibit a complete semi-analytic solution. Thus cases with a mechanical luminosity $L_{NSB}>L_{crit}$ demand of a numerical integration of the flow equations. Nevertheless, from the semi-analytic results of \\citet{2008ApJ...686..172S}  it was suggested  that the thermalization of the kinetic energy released within the volume occupied by a young nuclear starburst (NSB) should  lead to a central well regulated accretion flow onto the massive object, while  a powerful wind, capable of  preventing  the accretion of the ambient interstellar gas (ISM) onto the SMBH, should emanate from the outer regions. Here we follow these suggestions and measure the power of the resultant winds, as well as the accretion rates onto the SMBH and show how their luminosity  correlates with the starburst parameters. Following \\citet{2000MNRAS.311..346N},  \\citet{2009ApJ...699...89C} we have assumed that the angular momentum of the thermalized gas is not too large and here we present results from our 1D numerical simulations as a first approximation. The simulations confirm Silich's \\etal (2008) semi-analytic results  and extend them to starbursts with a mechanical energy input rate, $L_{NSB} > L_{crit}$. In this paper  the input physics account for the  mechanical feedback that a young nuclear starburst may provide to the accretion flow while allowing  for  radiative cooling and the gravitational pull from both the central SMBH and the  nuclear starburst. The feedback provided by the central SMBH  \\citep{1998A&A...331L...1S, 2001ApJ...551..131C, 2009ApJ...699...89C} will be the subject of a forthcoming communication.  The paper is organized as follows: the numerical scheme used and the initial and boundary conditions are presented  in section 2. Section 3 deals with the results from the simulations and compares them  with the semi-analytic results. In section 4  the accretion rates and luminosities of the central SMBHs are obtained, as well as a measure of the possible impact of the resultant winds onto the host galaxy ISM. Section 5 gives our conclusions. ",
        "conclusions": "By means of 1D numerical simulations and  semi-analytic estimates, we  have worked out the stationary hydrodynamic solution for  the matter  reinserted by stellar winds and type II supernovae from a young, massive and compact starburst in presence of a  central SMBH.  The solution is bimodal in all cases, below and above the threshold line ($L_{crit}$), with a stagnation radius ($R_{st}$) which defines the outer boundary of the  accretion flow onto the SMBH as well as the inner boundary of the  starburst wind.  We have shown that at the stagnation radius, the force of gravity perfectly balances the outward pressure gradient  acquired by the thermalized reinserted matter.   We have also shown  that radiative cooling becomes an important issue for massive  starbursts with a mechanical luminosity above the threshold line ($L_{NSB} > L_{crit}$). In all these cases,  radiative cooling depletes the pressure established through  thermalization of the injected matter and this leads to the development of  a thermal instability in the accretion flow. The stagnation radius then moves rapidly, towards the starburst boundary, with the mass of the considered starburst. In all simulations with $L_{NSB} > L_{crit}$, strong radiative cooling occurs at a radius interior to the stagnation  radius. Radiative cooling re-structures the inner accretion flow lowering the temperature to the lowest allowed value, $T = T_{min}$. From then onwards and despite the continuous input of energy, the rapid velocity increase leading to a rapid density enhancement keeps and sustains the in-falling gas temperature at  $T = T_{min}$. It is the larger mass deposition rate, provided by more massive starbursts, what triggers the onset of strong radiative cooling and with it the shift of the stagnation radius  towards the starburst boundary.  This results in a rapid increase  of the central SMBH luminosity for starbursts further above the  critical threshold  in the $L_{NSB}-R_{NSB}-M_{BH}$ parameter space.   The larger mass deposition rates provided by more massive starbursts also leads to more powerful starburst winds  and an estimate of their mechanical power rules out  the possibility  of the ISM feeding the SMBH, at least during the type II supernova era.  Clearly, spherically symmetric calculations, as the ones presented  here, cannot account for the redistribution of the net angular  momentum in the accretion flow. Nevertheless, they provide a good  estimate of the upper limit to the accretion rate onto the central  black hole, while pointing to a direct  physical link between nuclear  starbursts and the central SMBH luminosity. Our calculations do  realistically account for the symmetric deposition of mass and energy from massive  stars around the central object. This suggests  that in a more  realistic 2D or 3D geometry, able to account for the redistribution of the net angular momentum, the hydrodynamics would still lead to a bimodal solution with an accretion flow and an outward wind. In such a case however, the residual angular momentum could favor the formation of a gaseous disk well contained within the nuclear starburst region. We shall consider some of these issues in a future communication."
    },
    "1004/1004.1099_arXiv.txt": {
        "abstract": "We present the results of our investigation into the stellar populations of 24 radio galaxies at $z \\simeq 0.5$ drawn from four complete, low-frequency selected radio surveys. We use the strength of the 4000\\ang\\ break as an indicator of recent star formation, and compare this with radio luminosity, optical spectral classification and morphological classification. We find evidence of different star formation histories for high- and low-luminosity radio sources; our group of low radio luminosity sources (typically FRI-type sources) has systematically older stellar populations than the higher radio luminosity group. Our sample is also fairly well divided by optical spectral classification. We find that galaxies classified as having low excitation spectra (LEGs) possess older stellar populations than high excitation line objects (HEGs), with the HEGs showing evidence for recent star formation. We also investigate the link between radio morphology, as used by  Owen \\& Laing (1989), and the stellar populations. We find that there is a preference for the ``fat-double'' sources to have older stellar populations than the ``classical double'' sources, although this is also linked to these sources lying predominantly in the LEG and HEG categories respectively. These results are consistent with the hypothesis that HEGs are powered by accretion of cold gas, which could be supplied, for example, by recent mergers, secular instabilities, or filamentary cold flows. These processes could also trigger star formation in the host galaxy. The host galaxies of the LEGs do not show evidence for recent star formation and an influx of cold gas, and are consistent with being powered by the accretion of the hot phase of the inter-stellar medium. ",
        "introduction": "There is increasing evidence that active galactic nuclei (AGNs) have an important role to play in the formation and evolution of galaxies via AGN-driven feedback. Various relationships exist between the central black hole of an AGN and the stellar population of the galaxy: for example, the relation between black hole mass and stellar velocity dispersion ($M_{BH}$-$\\sigma$) (Gebhardt et al. 2000; Ferrarese \\& Merritt 2000; Merritt \\& Ferrarese 2001; G\\\"ultekin et al. 2009) and the relation between black hole mass and bulge luminosity (e.g. McLure \\& Dunlop 2001; McLure \\& Dunlop 2002; G\\\"ultekin et al. 2009). The latter may also be expressed as the Magorrian relation between black hole mass and bulge stellar mass (Magorrian et al. 1998). These tight correlations provide compelling evidence for links between galaxy formation and evolution and the growth of black holes through AGN (accretion) activity. One suggested mechanism is that of quasar outflows limiting black hole masses, dependent on the depth of the potential wells of dark matter halos (Silk \\& Rees 1998).  The current models of galaxy formation and evolution also invoke AGN-driven feedback to halt the overproduction of stars in the most massive galaxies. For example, Croton et al. (2006) include AGN-driven feedback phenomena in their semi-analytic model, separated into ``quasar-mode'' (associated with the efficient accretion of cold gas) and ``radio-mode'' (associated with the less efficient accretion of warm gas) feedback. Sijacki et al. (2007) also include AGN-driven feedback in their full hydrodynamical model, once again separated into two modes: the ``quasar regime'' (corresponding to central black holes with high accretion rates) and mechanical feedback (corresponding to central black holes with low accretion rates). Studying the host galaxies of radio-loud AGN should therefore enable us to place important constraints on the validity of the feedback mechanisms in the models.  One method of investigating the cold- and hot-mode accretion in AGNs is with radio samples. Selecting AGNs on the basis of their low-frequency radio emission enables the  selection of type-2, or obscured AGNs, in the same way as radio-loud quasars. This is because the radio waves are unaffected by dust obscuration, and the radio emission from the extended lobes is optically thin and therefore orientation independent. Furthermore, selecting the type-2 radio-loud quasars, or radio galaxies, allows the study of the stellar populations in the AGN host galaxy, as the central nuclear emission is obscured by the putative dusty torus invoked in unified schemes (e.g. Antonucci 1993; Hill, Goodrich \\& DePoy 1996). We also note that in the lower power AGNs this optical faintness may arise not merely from obscuration, but also from an inherent lack of bright optical-UV continua due to the absence of an accretion disc (e.g. Chiaberge, Capetti \\& Celotti 1999). The obscuration-related optical faintness in powerful AGN has enabled many authors to study the evolution of their host galaxies, with most work based on the near-infrared ($K-z$) Hubble diagram (Lilly \\& Longair 1982; Jarvis et al. 2001; Willott et al. 2003) and broad-band colours (e.g. Lilly 1989) suggesting that the hosts of powerful radio galaxies are comprised of an old stellar population which forms at high redshift ($z>2$) and then passively evolves.  Radio galaxies are commonly classified by their radio morphologies into Fanaroff-Riley type I (FRI) and Fanaroff-Riley type II (FRII) classes (Fanaroff \\& Riley 1974). An alternative division is that introduced by Hine \\& Longair (1979) based on optical spectra into high excitation galaxies (HEGs) and low excitation galaxies (LEGs); this was further refined by Laing et al. (1994) and Jackson \\& Rawlings (1997) who classified 3CR radio galaxies on this basis (see section~\\ref{specclass}).  There is increasing evidence that this emission line classification scheme, rather than the FR-class, has a direct link to the accretion mode of the AGN. Hardcastle, Evans \\& Croston (2007) showed that Bondi accretion of the hot phase of the intergalactic medium (IGM) is sufficient to power all low excitation radio sources, whereas the HEGs are powered by the accretion of cold gas typically thought to be driven towards the central engine during a galaxy merger. There are, however, other theories for the origin of the accreted cold gas in these galaxies. For example, in the hydrodynamical model of Ciotti \\& Ostriker (2007), gas from dying stars can be recycled by cooling radiatively and falling to the nuclear region where a small fraction ($\\approx 1 \\%$ or less) is accreted onto the central black hole. Alternatively, various authors (Kere\\v s et al. 2005; Dekel \\& Birnboim 2006; Ocvirk, Pichon \\& Teyssier 2008; Dekel et al. 2009; Brooks et al. 2009; Kere\\v s et al. 2009) have used hydrodynamic simulations to show that cold mode accretion via streams or filaments of cold gas could be responsible for the majority of the gas supplied to massive galaxies at high redshifts. For lower mass galaxies this cold mode accretion could be the dominant source of cold gas right up to the present day.   Following on from earlier work by Lilly \\& Longair (1984), Baldi \\& Capetti (2008) studied a sample of nearby 3CR radio galaxies and compared optical and UV images in order to detect evidence of recent star formation.  They found evidence of recent star formation in the HEGs in their sample, but not in the LEGs. They suggest that the HEGs have undergone recent major mergers which both trigger star formation and power the AGN by providing a supply of cold gas for accretion. The LEGs in their picture, which supports the one put forward by Hardcastle et al. (2007), have experienced no such mergers.  In this case no bursts of star formation are triggered and the AGN are powered by accretion of the hot interstellar medium (ISM).  In complementary work, Emonts et al. (2008) find no evidence for large-scale H{\\sc i} structures in the host galaxies of FRI-type sources, but find that the host galaxies of most of the FRII-type sources in their sample contain significant amounts of H{\\sc i}. They suggest that this dichotomy is a result of different formation histories for FRI and FRII-type radio sources. This is again consistent with the picture outlined above in which high- and  low-luminosity radio galaxies have undergone different star formation histories and this also influences the mode of accretion on to the central supermassive black hole.  In recent work using the vast data set available from the Sloan Digital Sky Survey, Kauffmann, Heckman \\& Best (2008) study a sample of radio-loud AGNs with emission lines. They find a strong correlation between the presence of emission lines and the presence of a young stellar population, consistent with the findings of Baldi \\& Capetti (2008). They also find a correlation between the age of the stellar population and the radio luminosity normalised by the black hole mass. They note that strong optical AGNs have a significantly enhanced probability of hosting radio jets, and thus the radio and optical phenomena are not independent. However, although the SDSS is a wide-area survey, it does not probe the necessary volume to obtain a significant sample of the more luminous FRII-type or HEG populations. Therefore, in order to study these effects further, we need to use a complete sample of radio sources selected by other methods.  In this paper we present the results of our study into the star formation histories (SFHs) of 24 radio galaxies at $0.4 < z < 0.6$. We compare our 4000\\ang\\ break indices with radio luminosity, spectral classification and morphological classification. We begin in section \\ref{thesample} of the paper by describing our sample, our observations and our data reduction. Section \\ref{analysis} describes our 4000\\ang\\ break strength measurements and our spectral and morphological classifications. We present our findings in section \\ref{discussion} before concluding in section \\ref{conclusions}. Throughout the paper we assume a standard cosmology in which $H_{0}$ = 70 km s$^{-1}$, $\\Omega_{M}$ = 0.3 and $\\Omega_{\\Lambda}$ = 0.7. ",
        "conclusions": "In this paper we have used deep spectroscopic observations to determine how the age of the stellar populations in the host galaxies of powerful radio sources is related to the structure of the radio emission and the ionizing power of the AGN.  \\begin{itemize}  \\item We have shown that our sample of $z \\simeq 0.5$ radio galaxies forms two distinct populations in the D$_{n}$(4000) - L$_{151}$ plane. The population at lower radio luminosity is composed entirely of LEGs, the one at higher radio luminosity predominantly of HEGs (although a few LEGs fall within this population).  \\item Our lower radio luminosity population has systematically higher D$_{n}$(4000) indices than the higher radio luminosity population. We find that this is most likely due to the presence of a younger stellar population in the higher radio luminosity (or HEG) population after considering the possibility that scattered quasar light may significantly alter the D$_{n}$(4000). We find that scattered quasar light could only decrease D$_{n}$(4000) by a maximum value of 0.1, which is not sufficient to explain the dichotomy that we observe. However, we are unable to measure the H$\\delta$ for our objects, which does leave some uncertainty in determining the actual extent of the star formation in our galaxies.  \\item We find that there is a preference for ``fat-double'' sources to have older stellar populations than the ``classical double'' sources. However, this may simply reflect the underlying links between morphology and excitation state or radio morphology, which appear to be more strongly correlated with D$_{n}$(4000) than is the morphology.  \\item Our results are consistent with the hypothesis that HEGs are powered by the accretion of cold gas, the influx of which could be due to mergers -- which would also trigger star formation -- or other processes, such as cooled, recycled gas from dying stars or cold mode accretion. On the other hand, we suggest that LEGs are powered by the accretion of the hot phase of the inter-stellar medium (as they have undergone no recent influx of cold gas).   \\end{itemize}  The results presented here show that deep optical spectroscopy can provide important information on the host galaxy properties of moderate redshift, powerful radio galaxies. However, this also needs to be linked to other indicators of star formation and merger activity to obtain a comprehensive and consistent explanation for the observed differences between AGN and their corresponding effect on the evolution of massive galaxies.  In future work (Herbert et al. in preparation) we will use this sample to investigate where these radio sources fall on the Fundamental Plane and how they relate to the general massive galaxy population. Furthermore, we also have deep multi-band imaging data to assess the environmental richness of the full $z \\simeq 0.5$ sample which will enable us to link how the power of the AGN, its host galaxy mass and star formation history are related to the environmental density of the AGN."
    },
    "1004/1004.3325_arXiv.txt": {
        "abstract": "We attempt to measure possible miscalibration of the wavelength scale of the VLT-UVES spectrograph. We take spectra of QSO HE0515-4414 through the UVES iodine cell which contains thousands of well-calibrated iodine lines and compare these lines to the wavelength scale from the standard thorium-argon pipeline calibration. Analyzing three exposures of this $z=1.71$ QSO, we find two distinct types of calibration shifts needed to correct the Th/Ar wavelength scale.  First, there is an overall average velocity shift of between 100 $\\ms$ and 500 $\\ms$ depending upon the exposure. Second, within a given exposure, we find intra-order velocity distortions of 100 $\\ms$ up to more than 200 $\\ms$. These calibration errors are similar to, but smaller than, those found earlier in the Keck HIRES spectrometer. We discuss the possible origins of these two types of miscalibration. We also explore the implications of these calibration errors on the systematic error in measurements of $\\delalpha$, the change in the fine-structure constant derived from measurement of the relative redshifts of absorption lines in QSO absorption systems. The overall average, exposure-dependent shifts should be less relevant for fine-structure work, but the intra-order shifts have the potential to affect these results. Using either our measured calibration offsets or a Gaussian model with sigma of around 90 $\\ms$, Monte Carlo mock experiments find errors in $\\delalpha$ of between $1\\ten{-6} \\Nsys^{-1/2}$ and $3\\ten{-6} \\Nsys^{-1/2}$, where $\\Nsys$ is the number of systems used and the range is due to dependence on how many metallic absorption lines in each system are compared. \\bigskip ",
        "introduction": "In a recent paper (Griest et al. 2010) it was shown that Keck HIRES spectrograph measurements of absorption features in high redshift QSOs calibrated with the normal thorium-argon (Th/Ar) technique contained overall average velocity shifts between exposures of up to 2000 $\\ms$ over several nights and calibration errors of up to 500 $\\ms$ even within a single echelle order in a single exposure. Similar wavelength calibration problems were found previously by Osterbrock et al. (2000) and Suzuki et al. (2003). These results were found by re-calibrating the wavelength scale using spectra taken through the Keck HIRES iodine cell.  The iodine cell, which has been used extensively in the Doppler method search for extrasolar planets, puts several sharp lines per Angstrom on top of the QSO spectrum, and its reproducible absorption spectrum allows extrasolar planet workers to attain a relative velocity precision of around 5 $\\ms$ between exposures taken up to several years apart (Johnson et al. 2006; Konacki 2009). The analysis in Griest et al. (2010) compared the iodine cell wavelength calibration to Th/Ar calibration with two separate Keck HIRES analysis pipelines (XIDL and MAKEE) and so is probably measuring calibration problems with the Th/Ar calibration of the spectrograph itself rather than problems just with the standard calibration software.  For most astronomical work, wavelength calibration errors of around 1 $\\kms$ are not important, but over the past few years several groups have used absorption lines in the spectra of high redshift QSOs to measure or place limits on possible changes in the fine-structure constant over cosmological time. There has been considerable controversy in experimental measurements of $\\delalpha$ using high redshift absorption systems, with claims of detection and also limits on variation inconsistent with those claims.  See Murphy et al. (2008) for a summary.  If the fine-structure constant was different 10 billion years ago by the claimed detection of $\\delalpha = (-5.7 \\pm 1.1)\\ten{-6}$ (Murphy et al. 2003, 2004), relative atomic transition wavelengths are expected to differ from their lab values by up to $\\sim$200 $\\ms$. For such experimental uses, it is important to pay careful attention to calibration errors of the size reported above.  However, it is important to understand that while individual measurement errors may be larger than the expected signal, if many atomic transitions are used in many QSO absorption systems, it may be possible to average away these calibration uncertainties. See, for example, Murphy et al. (2009).  The results of Griest et al. (2010) apply only to the Keck HIRES instrument, but there is another instrument, the VLT UVES spectrograph, that is playing a key role in the search for possible changes in the fine-structure constant using absorption lines in high redshift QSOs. In this paper, we perform a similar recalibration of the standard UVES Th/Ar wavelength calibration pipeline using the VLT iodine cell. We find similar, but smaller, wavelength calibration errors than found in HIRES. We discuss the possible origin of these offsets, in particular whether they arise from the UVES pipeline software or systematic errors within the telescope and/or spectrograph. We also make a first attempt at calculating whether these calibration errors can give rise to important systematic errors in the measurements to date of $\\delalpha$. ",
        "conclusions": "We tested the accuracy of the wavelength calibration of the VLT UVES spectrograph by recalibrating the standard Th/Ar calibration pipeline output using spectra taken though the UVES iodine cell. We found several types of miscalibration:  first, an exposure-dependent overall average velocity shift that is probably mostly due to the position of the QSO within the spectrograph slit. Repositioning of the spectrograph gratings and pressure/temperature drifts inside the spectrograph may also contribute to these velocity shifts.  These overall average shifts range up to 500--600 $\\ms$, but do not affect fine-structure constant measurements as long as comparisons between absorption features are only performed within the same spectrum.  Second, and more importantly, we found intra-order calibration shifts of up to several hundred meters/second that occur within the same spectrum and within each echelle order.  These can affect measurement of the fine-structure constant. We investigated several possible causes for these shifts.  We explored the effect of the degree of the polynomial used in the Th/Ar calibration and found a velocity dispersion of 20--40 $\\ms$ due to this. This is several times smaller than the total dispersion we find in the velocity measurements, so the main effect we are finding is not due to the degree of the Th/Ar calibration polynomial. In fact, we found that the iodine cell recalibration of extractions using different degree polynomials allowed us to regain the same absolute wavelength scale to an accuracy of better than 5 $\\ms$, thus showing the robustness of our iodine recalibration method.  We considered several other possible causes for the intra-order shifts and showed that spectrograph blaze function and related possibilities were unlikely. Thus, we conclude that most likely there are unknown hardware-related systematic errors within UVES and/or the VLT responsible for these shifts. We noticed some weak patterns in the shift as a function of CCD pixel and attempted a correction based upon these patterns.  This correction was not very successful, so we conclude that additional data and analysis are required to discover the cause and potential fixes for these wavelength scale shifts. We note that it is important look for such a correction, since if the intra-order distortion pattern for a given object exposure could be well approximated from a subsequent I$_2$ exposure, either of the QSO or of a nearby bright star, then one could correct the Th/Ar wavelength scale of the former, i.e., a sort of calibration transfer function could be established and applied.  We next explored the effect of these shifts on determination of or limits on the value of the fine-structure constant.  Using our measured shifts we performed Monte Carlo experiments and found that the effect of the intra-order shifts is well modeled by a Gaussian dispersion in wavelength calibration with a magnitude of around 80--120 $\\ms$.  We found that the effect on $\\delalpha$ depended strongly on the number of lines being compared within the spectrum, with the use of a small number of lines resulting in larger than expected errors.  Thus, until a correction to these intra-order wavelength miscalibrations is found, we recommend against focusing on comparison between pairs of lines or the use of spectra which contain only a few lines. We summarized the results of our Monte Carlo experiments in a fitting formula Equation~(\\ref{eqn:sigmada}), which should be useful in estimating the effect of these (or similar) wavelength calibration problems on future fine-structure work."
    },
    "1004/1004.4505_arXiv.txt": {
        "abstract": "We describe and use two different statistical approaches to try and detect low-frequency solar oscillations in Sun-as-a-star data: a frequentist approach and a Bayesian approach. We have used frequentist statistics to search contemporaneous Sun-as-a-star data for coincident, statistically-prominent features. However, we find that this approach leads to numerous false detections. We have also used Bayesian statistics to search for evidence of low-frequency p modes and g~modes in Sun-as-a-star data. We describe how Bayesian statistics can be used to search near-contemporaneous data for coincident prominent features. Near-contemporaneous data were used to circumvent the difficulties in deriving probabilities that occur when common noise is present in the data. We find that the Bayesian approach, which is reliant on the assumptions made when determining the posterior probability, leads to significantly fewer false detections and those that are observed can be discredited using a priori knowledge. Therefore, we have more confidence in the mode candidates found with Bayesian statistics. ",
        "introduction": "\\citet{Appourchaux2008} describes how Bayesian statistics can be used to determine the frequencies of modes in solar data and \\citet{Appourchaux2009} uses Bayesian statistics to search asteroseismic data observed by the COnvection ROtation and planetary Transits (CoRoT) spacecraft for evidence of oscillations. We now summarize the methods used by \\citet{Appourchaux2009} and adapt them so that they can be used to search contemporaneous solar data.  Bayes theorem can be used to determine the `posterior probability', $p(\\textrm{H}0|x)$, that the H0 hypothesis is true given the observed data, $x$. \\citet{Appourchaux2009} shows that the posterior probability is given by \\begin{equation}\\label{equation[Bayes reduced]} p(\\textrm{H}0|x)=\\left[1+\\frac{(1-p_0)}{p_0}\\frac{p(x|\\textrm{H}1)}{p(x|\\textrm{ H}0)}\\right]^{-1}, \\end{equation} where $p_0$ is the prior probability, which is defined as the subjective probability, before any observations have been made, that H0 is true. The H0 hypothesis assumes that the data contains only noise. Therefore, in equation (\\ref{equation[Bayes reduced]}), $p(x|\\textrm{H}0)$ is the probability that a spike with a power $x$ is observed in a frequency-power spectrum that contains only noise. Similarly, $p(x|\\textrm{H}1)$ is the probability that a spike with a power, $x$, is observed if the $\\textrm{H}1$ hypothesis is true. Here, the $\\textrm{H}1$ hypothesis assumes that there is a signal in the data with a maximum power spectral density per unit bin, $H$. We are assuming that H0 and H1 cover all possibilities. In other words there is no third hypothesis, H2 say, that is true. \\citet{Appourchaux2009} derives the expressions for $p(x|\\textrm{H}0)$ and $p(x|\\textrm{H}1)$ that can be used to search frequency-power spectra. We now quote these expressions for reference.  We define a spike as the power spectral density contained in a single bin of the frequency-power spectrum. Frequency-power spectra have a $\\chi^2$, 2 d.o.f. distribution and so the probability of observing a spike with a power spectral density of $x$ if the data contain noise only is given by \\begin{equation}\\label{equation[p(x|h0)]} p(x|\\textrm{H}0)=\\textrm{e}^{-x}, \\end{equation} where we have taken the mean power spectral density of the spectrum to be 1. For the alternative H1 hypothesis we assume that the spike contains the signal from a mode which has been stochastically excited, as is the case for solar oscillations. However, as we do not know the power of the mode we assume, a priori, that the maximum power spectral density per unit bin of the mode in a frequency-power spectrum is distributed uniformly between $0$ and $H$. The probability of observing a spike with a power spectral density, $x$, is then given by \\begin{equation}\\label{equation[p(x|h1)]} p(x|\\textrm{H}1)=\\frac{1}{H}\\int_0^H\\frac{1}{1+H'}\\textrm{e}^{-x/(1+H')}\\textrm{ d}H'. \\end{equation} Equations (\\ref{equation[p(x|h0)]}) and (\\ref{equation[p(x|h1)]}) can be substituted into equation (\\ref{equation[Bayes reduced]}) to give the posterior probability.  \\subsection{Range of heights}\\label{section[amplitude range]} In order to calculate $p(x|\\textrm{H}1)$ we have to make assumptions about the maximum power spectral density per unit bin of a mode profile, $H$, with a frequency, $\\nu$. Fig. \\ref{figure[post prob vs a]} shows the variation with $H$ of the posterior probability, which is related to $p(x|\\textrm{H}1)$ by equation (\\ref{equation[Bayes reduced]}), for power spectral densities, $x$. The choice of $H$ is most influential at low $H$ and high $x$ and so it is important to choose the value of $H$ carefully to enable the significance of any prominent feature to be accurately assessed. As we are searching for oscillations that have not yet been observed we need to use models to estimate the value of $H$.  Notice that for each $x$ plotted in Fig. \\ref{figure[post prob vs a]} a minimum posterior probability is observed. Such minima can be understood as a manifestation of the fact that the prior on $H$ being large is contradicted by the low power spectral density level reached by the peak \\citep{Appourchaux2009}. In other words, if the mode power spectral densities were truly so high the observed peak would also be prominent.  \\begin{figure} \\centering \\includegraphics[width=0.4\\textwidth, clip]{post_prob_vs_h}\\\\ \\caption{The variation of the posterior probability with the maximum power, $H$, for different threshold powers, $x$.} \\label{figure[post prob vs a]} \\end{figure}  Before we discuss the values used to constrain $H$ it is important to define exactly what we mean by the maximum power spectral density per unit bin of a mode. In this paper we define the maximum power spectral density as the peak height of a mode's profile in a frequency-power spectrum. Solar oscillations are damped and, if the lifetime of a mode is significantly shorter than the length of the observations, the mode will appear in a frequency-power spectrum in the form of a Lorentzian. The width of the Lorentzian, $\\Delta$, is related to a mode's lifetime, $\\tau$, by the following equation: \\begin{equation}\\label{equation[width-lifetime]} \\Delta=\\frac{1}{\\pi\\tau}. \\end{equation} The maximum power spectral density per unit bin (or height) can be described by the following equation \\citep{Fletcher2006}: \\begin{equation}\\label{equation[intermediate H]} H=\\frac{2V^2}{\\pi T\\Delta+2}. \\end{equation} where $V^2$ is the total power of the mode, $T$ is the length of the time series and $\\Delta$ is in units of frequency. Consequently, $T\\Delta$ gives the width of the Lorentzian in units of bins.  For p~modes the value of $H$ was estimated in two different ways, each of which will now be described. \\citet{Broomhall2008} extrapolated to low frequencies the total mode powers, $V^2$, and widths, $\\Delta$, of well-defined, higher-frequency modes, whose profiles in a frequency-power spectrum can be fitted accurately. The total power, $V^2$, was extrapolated in preference to $H$ as it is not dependent on $\\Delta$. \\citeauthor{Broomhall2008} did this in two different ways and here we have used the method that produced the largest $H$ values.  Even so, it is possible that this method under-estimates the power of the signal from a mode. For example, \\citet{Broomhall2008} found that the $l=0$, $n=6$ mode at $\\sim972\\,\\rm\\mu Hz$ is excited to a larger power than is expected from the extrapolations. In fact, the simulations performed by \\citeauthor{Broomhall2008} indicated that, if the actual power was equal to the value given by the extrapolation, the probability that the $l=0$, $n=6$ mode would be excited to the observed power is approximately 0.01. Furthermore, because solar oscillations are excited stochastically, the total power of a mode, $V^2$, varies with time. If the value of the power given by the extrapolation is regarded as a mean, the observed power of a mode is equally likely to be greater than this value as it is to be lower. Consequently, by integrating over the range 0 to $H$ we could be underestimating the true height of the mode. More realistic probabilities may be obtained if we overestimate the range of heights over which the integration in equation (\\ref{equation[p(x|h1)]}) is made. Therefore, we have also performed a search where the maximum power spectral density per unit bin, $H$, was assumed to be constant at all frequencies and to have the height predicted by \\citet{Broomhall2008} for a mode with a frequency of $1500\\,\\rm\\mu Hz$. A larger value of $H$ should allow signals to be detected more easily as when $H$ tends to 0, $p(x|\\textrm{H}1)$ tends to $p(x|\\textrm{H}0)$ and so it is very difficult to determine which hypothesis is most likely to be correct (see Fig.~ \\ref{figure[post prob vs a]}).  As for low-frequency p~modes we are reliant on models to predict the powers of g~modes. Two groups have produced predictions of g-mode powers. The first group uses the Cambridge stochastic excitation model, as applicable to g~modes \\citep{Houdek1999}. The second group \\citep{Belkacem2009} uses a similar model with a different temporal-correlation between the modes and the turbulent eddies: \\citeauthor{Belkacem2009} take the eddy-time correlation function to be a Lorentzian, while \\citeauthor{Houdek1999} assume that the eddy-time correlation is a Gaussian. The authors of both papers feel that their respective models are only applicable to limited ranges in frequency and these ranges do not coincide. However, the work of \\citeauthor{Belkacem2009} predicts significantly larger g-mode powers than those predicted by \\citeauthor{Houdek1999}. Therefore, we have taken the upper limit for $H$ to be the maximum power predicted by the work of \\citet{Belkacem2009} as the larger the mode powers the more likely we are to be able to detect any oscillations. We have taken the total power, $V^2$, to be the constant value of $9\\,\\rm mm^2\\, s^{-2}$ at all frequencies in the range $(50-350)\\,\\rm\\mu Hz$. We have then assumed that the lifetime of the g~modes is significantly longer than the length of the observations and so all of the modes power will be confined to a single bin. Therefore the height of the mode, $H$, is equal to the total power, $V^2$ i.e. $H=9\\,\\rm mm^2\\, s^{-2}\\,bin^{-1}$. \\citet{Belkacem2009} note that, given the uncertainties in their model, this maximum power could, potentially, be a factor of four larger. Therefore, we have also determined the posterior probability given the observed prominent features when the maximum height, $H$, was $36\\,\\rm mm^2\\,s^{-2}\\,bin^{-1}$.  \\subsection{Searching a range of frequencies} The statistics described above can be used to search individual bins of a frequency-power spectrum. However, there are $\\sim400,000$ bins in the total range of frequencies searched here $(50-1500\\,\\rm\\mu Hz)$. Consequently the chances of making a false detection are significant when searching each bin individually. We therefore consider a range of frequencies containing $M$ bins. In the manner described by \\citet{Sturrock2008} we define the following null and alternative hypotheses: \\newline \\newline \\noindent H00: The principal peak in $M$ bins is due to noise,\\newline \\noindent H11: The principal peak is due to an oscillatory signal with a frequency, $\\nu$.\\newline \\newline \\noindent Then \\begin{equation}\\label{equation[p(h00)]} P(x|\\textrm{H}00)=1-[1-P(x|\\textrm{H}0)]^M\\approx MP(x|\\textrm{H}0), \\end{equation} when $P(x|\\textrm{H}0)$ is small. We assume only one frequency bin out of the $M$ of interest contains a signal and so \\begin{equation}\\label{equation[p(h11)]} P(x|\\textrm{H}11)=P(x|\\textrm{H}1). \\end{equation} Therefore, \\begin{eqnarray}\\label{equation[p(h00|x)]} p(\\textrm{H}00|x)&=&\\left[1+\\frac{(1-p_0)}{p_0}\\frac{p(x|\\textrm{H}11)}{ p(x|\\textrm{H}00)}\\right]^ {-1}\\nonumber\\\\ &=&\\left[1+\\frac{(1-p_0)}{p_0}\\frac{p(x|\\textrm{H}1)}{Mp(x|\\textrm{H}0)}\\right]^ {-1}. \\end{eqnarray} To ensure that only one mode component is present in the $M$ bins we have taken $M$ to be the number of bins in $0.4\\,\\rm\\mu Hz=106$, as this is the expected splitting between adjacent $m$ components. Note that we have erred on the side of caution when setting this value of $M$ as in Sun-as-a-star data only $(l+m)$-even components can be detected. Therefore, the average splitting between adjacent visible components is closer to $0.8\\,\\mu\\rm Hz$. We have assumed that the components of g~modes are also rotationally split by $0.4\\,\\rm \\mu Hz$. It should be noted that the core rotation rate is very uncertain and so this assumption may not be valid. For example, \\citet{Garcia2007} suggest that the core might be rotating between three and five times faster than the radiative zone. However, even if this is the case there will still be only one component present in the $M=106$ bins under consideration. One possible extension to this work would be to test the data taking into account a faster core rotation rate.  \\subsection{Searching smoothed spectra using Bayesian techniques}\\label{section[smoothed bayes]} The lifetimes of p~modes are expected to be short enough for their profiles to have a resolved width in the frequency-power spectrum examined here. Until now we have assumed that g~modes have very long lifetimes but this is not necessarily correct. \\citet{Garcia2007} have found a very interesting peak in a periodogram of a periodogram of GOLF data. They claim that this peak is due to the superposition of g~modes that lie in the asymptotic range where the separation in period is approximately constant ($\\sim24\\,\\rm min$). The peak is positioned close to the period predicted by standard solar models. \\citeauthor{Garcia2007} use artificial data to show that, if this peak is evidence of g~modes, the modes might have widths in a frequency-power spectrum that are commensurate with damping times of several months. However, this is not the only explanation for the observed mode widths. The observed results could also be due to an internal magnetic splitting, changes with time of the cavities in which the oscillations were trapped or a noise effect in the GOLF data. If the damping times are of the order of months and given that the data searched here are $3071\\,\\rm d$ in length, the width of these modes will be resolved in frequency-power spectra.  \\citet{Appourchaux2009} describe how Bayesian statistics can be used to search for modes with resolved widths. This is done by summing the frequency-power spectrum over a given number of bins. Smoothing a spectrum alters the underlying statistics of the noise and so it is necessary to derive new equations for $p(x|\\textrm{H}0)$ and $p(x|\\textrm{H}1)$. If we smooth a frequency-power spectrum over $R$ bins the underlying statistics of the noise becomes $\\chi^2$ with $2R$ d.o.f.. We have used the equations for $p(x|\\textrm{H}0)$ and $p(x|\\textrm{H}1)$, which are described in Section 3.2 of \\citet{Appourchaux2009}, to search solar data for modes with lifetimes that are shorter than the length of the time series. We now give the equations for reference.  The probability that, in a spectrum that has been smoothed over $R$ bins, a spike with a power spectral density of $x$ is observed if the $\\textrm{H}0$ hypothesis is true is given by \\begin{equation}\\label{equation[smoothed p(x|h0)]} p(x|\\textrm{H}0)=\\frac{x^{R-1}\\textrm{e}^{-x}}{\\Gamma(R)}, \\end{equation} where $\\Gamma(R)$ is the gamma function and we have assumed that the mean of the unsmoothed frequency-power spectrum is unity. For the alternative hypothesis we again assume that the spike contains the signal from a stochastically excited mode. However, when dealing with smoothed spectra, instead of making assumptions about $H$ we must make assumptions about the total mode power, $V^2$, and the mode linewidth, $\\Delta$. Here we assumed, a priori, that the mode amplitude is distributed uniformly between 0 and $V$ and that the mode linewidth is distributed uniformly between 0 and $\\Delta$. Then the probability of observing a spike with a power spectral density $x$ in a smoothed spectrum, if the $\\textrm{H}1$ hypothesis is true, is given by \\begin{equation}\\label{equation[smoothed p(x|h1)]} p(x|\\textrm{H}1)=\\frac{1}{V\\Delta}\\int^V_0\\int^\\Delta_0\\frac{\\lambda^\\nu} {\\Gamma(\\nu)}x^{\\nu-1}\\textrm{e}^{-\\lambda x}\\textrm{d}V'\\textrm{d}\\Delta', \\end{equation} where $\\lambda$ and $\\nu$ are given in \\citet{Appourchaux2004} and are related to the line profile of the mode, i.e. they are related to $V$ and $\\Delta$. Equations (\\ref{equation[smoothed p(x|h0)]}) and (\\ref{equation[smoothed p(x|h1)]}) can be substituted into equation (\\ref{equation[p(h00|x)]}) to give the posterior probability.  When searching for p~modes, we have used the values for $V^2$ and $\\Delta$ extrapolated by \\citet{Broomhall2008} as the upper limits on the mode powers and lifetimes. When searching for g~modes we took $V^2$ to be both $9\\,\\rm mm^2\\,s^2$ and $36\\,\\rm mm^2\\,s^2$ (see Section \\ref{section[amplitude range]} for details). As the lifetimes of the g~modes are not known we used a range of different values: 1, 2, 3, and 4 months. Equation (\\ref{equation[p(h00|x)]}) was then used to determine the posterior probability. The posterior probability, given by Equation (\\ref{equation[p(h00|x)]}), is also dependent on the prior probability, $p_0$, and we now discuss the value given to $p_0$. ",
        "conclusions": "\\label{section[discussion]} We begin by comparing the results observed in the unsmoothed spectra using frequentist and Bayesian search methods. The frequencies of candidates that are detected by both approaches show very good agreement with each other and with previously claimed detections \\citep[see e.g.][]{Toutain1998, Bertello2000, Garcia2001, Chaplin2002, Broomhall2007, Salabert2009}. When the frequentist approach was applied, more false detections (i.e. detections that were found to lie away from the predicted mode frequencies) were observed than expected. We believe this is significant and highlights the misleading nature of the frequentist approach. When the Bayesian approach was applied significantly fewer false detections were observed, even before the uneven prior probability was applied. Furthermore, the Bayesian approach allows us to use our knowledge of the Sun in a statistically rigourous manner to guide the frequencies at which we consider detections to be mode candidates (with the uneven prior probability). Although Bayesian methods may lead to more stringent threshold levels and consequently fewer detections we feel that this is a small price to pay for the improved accuracy and confidence one achieves in any detections. It is possible that using the uneven prior probability may discount detections that are in fact real mode signatures (type~II errors). However, we feel that use of the uneven prior probability is justified as it also helps to minimize the number of noise signatures that are considered as mode candidates (type~I errors). The Bayesian approach is also less open to misinterpretation than the frequentist approach. However, once again we note that the Bayesian approach is reliant on the assumptions made when determining the posterior probability. Different assumptions may lead to different mode candidates being uncovered. For example, the assumption that the rotational splitting of g~modes is $0.4\\,\\mu Hz$ may be incorrect \\citep{Garcia2007}.  Mode candidates were also detected in smoothed spectra. This method was particularly successful at higher frequencies, where the modes have broader widths. It is worth noting that as the p-mode powers and widths were determined by extrapolation it is likely that the values used here are more accurate higher in the frequency search range. This may have influenced the results especially if the mode widths were overestimated. We were able to eliminate mode candidates that were positioned away from the predicted mode frequencies by employing the non-uniform prior probability. This cannot be done for the frequentist approach as we are effectively using a priori knowledge to guide the regions in frequency we search. The use of a priori knowledge in a statistically rigourous manner can only be done through Bayesian statistics.  The g-mode candidate that was observed in the frequentist approach was not detected when Bayesian statistics were used. One g~mode candidate was detected using the Bayesian approach. However, each detection was only made in one combination of the data sets and neither candidate was detected with both the frequentist and Bayesian approaches. Therefore, we regard these two candidates with suspicion and feel that each candidate still requires further investigation and independent confirmation. The candidate that was detected with the Bayesian approach was not detected when the spectra were smoothed. This could indicate that, if it is a g~mode, it has a lifetime that is longer than the lifetimes assumed here."
    },
    "1004/1004.3529_arXiv.txt": {
        "abstract": "Major stellar-wind emission features in the spectrum of $\\eta$ Car have recently decreased by factors of order 2 relative to the continuum. This is unprecedented in the modern observational record. The simplest, but unproven, explanation is a rapid decrease in the wind density. ",
        "introduction": "%  Today, 150 years after the close of its Great Eruption, $\\eta$ Car has not yet returned to thermal and rotational equilibrium \\citep{2005ASPC..332....3M,2005ASPC..332..101D}.  This fact is important because the ``SN impostor'' phenomenon and its aftermath constitute a major gap in the theory of massive stars, and $\\eta$ Car is the only example that can be studied in detail; see reviews by many authors in \\citet{2005ASPC..332.....H}.  Its recovery has been {\\it unsteady\\/},  with unexplained photometric and spectral changes in the 1890s and 1940s \\citep{2008AJ....135.1249H}.   This object may have entered a phase of accelerated development 12-15 years ago.  From 1953 to the mid-1990's, ground-based `V' photometry of star plus ejecta brightened at a rate of 0.024 mag/yr, with brief deviations smaller than $\\pm$0.3 mag (Fig.\\ 2 in \\citealt{1999AJ....118.1777D}). In the past decade, however, it has risen 0.6 mag above that earlier trend line (Fig.\\ 3 in \\citealt{2009A&A...493.1093F}).   The central star shows a more dramatic increase,  a factor of more than 3 in UV-to-visual {\\it Hubble Space Telescope (HST))\\/} data since 1998 \\citep{2004AJ....127.2352M, 2006AJ....132.2717M, 2009IAUC.9094....1D}.   A decrease in the amount of circumstellar dust may be responsible, but that requires some change in the wind and/or radiation field.  Meanwhile the periodic ``spectroscopic events'' of 1998.0, 2003.5, and 2009.0, defined in \\S2 below, differed in major respects \\citep{2005AJ....129..900D,2010AJ....139.1534R,Corcoran:2009:Online}. Very likely the mass-loss rate has been decreasing at an inconstant pace, while rotational spin-up may play a role \\citep{2008AJ....135.1249H, 2006AJ....132.2717M, 2005AJ....129..900D, 2003ApJ...586..432S}.  All those discussions, however, seemed to face an embarrassing observational contradiction.  From the first {\\it HST\\/} spectroscopy in 1991 until the Space Telescope Imaging Spectrograph (STIS) failed in 2004, {\\it $\\eta$ Car's spectrum showed no major change\\/} except during the temporary spectroscopic events.  One might have expected some sort of spectral evolution to accompany the rapid brightening after 1998.  In this letter we report a novel development:  Observations in 2007--2010 with Gemini/GMOS and {\\it HST\\/}/STIS reveal major spectral changes. They are not subtle;  evidently the wind has been altered, at least temporarily and perhaps for the indefinite future. ",
        "conclusions": ""
    },
    "1004/1004.4415_arXiv.txt": {
        "abstract": "% {After several generations of interferometers in radioastronomy, wide-field imaging at high angular resolution is today a major goal for trying to match optical wide-field performances.} % {All the radio-interferometric, wide-field imaging methods currently belong to the mosaicking family. Based on a 30 years old, original idea from Ekers \\& Rots, we aim at proposing an alternate formalism.} % {Starting from their ideal case, we successively evaluate the impact of the standard ingredients of interferometric imaging, \\ie{} the sampling function, the visibility gridding, the data weighting, and the processing of the short spacings either from single-dish antennas or from heterogeneous arrays. After a comparison with standard nonlinear mosaicking, we assess the compatibility of the proposed processing with 1) a method of dealing with the effect of celestial projection and 2) the elongation of the primary beam along the scanning direction when using the on-the-fly observing mode.} % {The dirty image resulting from the proposed scheme can be expressed as a convolution of the sky brightness distribution with a set of wide-field dirty beams varying with the sky coordinates.  The wide-field dirty beams are locally shift-invariant as they do not depend strongly on position on the sky: Their shapes vary on angular scales typically larger or equal to the primary beamwidth. A comparison with standard nonlinear mosaicking shows that both processing schemes are not mathematically equivalent, though they both recover the sky brightness. In particular, the weighting scheme is very different in both methods. Moreover, the proposed scheme naturally processes the short spacings from both single-dish antennas and heterogeneous arrays.  Finally, the sky gridding of the measured visibilities, required by the proposed scheme, may potentially save large amounts of hard-disk space and cpu processing power over mosaicking when handling data sets acquired with the on-the-fly observing mode.} % {We propose to call this promising family of imaging methods \\emph{wide-field synthesis} because it explicitly synthesizes visibilities at a much finer spatial frequency resolution than the one set by the diameter of the interferometer antennas.} % ",
        "introduction": "The instantaneous field of view of an interferometer is naturally limited by the primary beam size of the individual antennas. For the ALMA 12m-antennas, this field of view is $\\about9''$ at 690~GHz and $\\about 27''$ at 230~GHz.  The astrophysical sources in the (sub)-millimeter domain are often much larger than this, but still structured on much smaller angular scales.  Interferometric wide-field techniques enable us to fully image these sources at high angular resolution.  These techniques first require an observing mode that in one way or another scans the sky on spatial scales larger than the primary beam.  The most common observing mode in use today, known as stop-and-go mosaicking, consists in repeatedly observing sky positions typically separated by half the primary beam size. The improvement of the tracking behavior of modern antennas now leads astronomers to consider on-the-fly observations, with the antennas slewing continuously across the sky. The improvements in correlator and receiver technologies are also leading to techniques that could potentially sample the antenna focal planes with multi-beam receivers instead of the single-pixel receivers installed on current interferometers.  The ideal measurement equation of interferometric wide-field imaging is \\begin{equation} \\V{}(\\up,\\as) = \\DFTint{\\ap}{\\B{}(\\ap-\\as)\\,\\I{}(\\ap)}{\\up}, \\label{eq:measurement:otf} \\end{equation} where \\V{} is the visibility function of 1) \\up{} (the spatial frequency with respect to the fixed phase center) and 2) \\as{} (the scanned sky angle), \\I{} is the sky brightness, and \\B{} the antenna power pattern or primary beam of an antenna of the interferometer~\\citep[][chapter 2]{thompson86}. For simplicity, 1) we assume that the primary beam is independent of azimuth and elevation, and 2) we use one-dimensional notation without loss of generality. We do not deal with polarimetry~\\citep[see \\eg][]{hamaker96,sault96a,sault99} because it is adds another level of complexity over our first goal here, \\ie{} wide-field considerations. Several aspects make Eq.~\\ref{eq:measurement:otf} peculiar with respect to the ideal measurement equation for single-field observations. First, the visibility is a function not only of the $uv$ spatial frequency (\\up{}) but also of the scanned sky coordinate (\\as{}). Second, Eq.~\\ref{eq:measurement:otf} is a mix between a Fourier transform and a convolution equation. It can be regarded, for example, as the Fourier transform along the \\ap{} dimension of the function, $B{}(\\ap-\\as)\\,\\I{}(\\ap)$, of the (\\ap,\\as) variables. But Eq.~\\ref{eq:measurement:otf} can also be written as the convolution: \\begin{equation} \\V{}(\\up,\\as) = \\sint{\\ap}{\\Bg{}(\\as-\\ap)\\,\\Ig{}(\\ap,\\up)}, \\end{equation} \\begin{equation} \\mbox{where} \\quad \\Bg{}(\\as-\\ap) \\Definition \\B{}(\\ap-\\as) \\end{equation} \\begin{equation} \\mbox{and} \\quad \\Ig{}(\\ap,\\up) \\Definition \\I{}(\\ap)\\,\\dexp{\\ap\\up}. \\end{equation} For each \\up{} kept constant, \\V{}(\\up,\\as) is the convolution of \\Bg{} and \\Ig{}. Indeed, $\\Ig{}(\\ap,\\up = 0) = \\I{}(\\ap)$, so we derive \\begin{equation} \\V{}(\\up=0,\\as) = \\conv{\\Bg{}}{\\I{}}{\\ap}{\\as}, \\end{equation} \\ie{}, the convolution equation for single-dish observations.  \\citet{ekers79} were the first to show that the measurement equation (Eq.~\\ref{eq:measurement:otf}) enables us to recover spatial frequencies of the sky brightness at a much finer $uv$ resolution than the $uv$ resolution set by the diameter of the interferometer antennas.  Interestingly enough, the goal of \\citet{ekers79} was ``just'' to find a way to produce the missing short spacings of a multiplying interferometer.  However, \\citet{cornwell88} realized that Ekers \\& Rots' scheme has a much stronger impact, because it explains why an interferometer is able to do wide-field imaging. \\citet{cornwell88} also demonstrated that on-the-fly scanning is not absolutely necessary to interferometric wide-field imaging. Indeed, the large-scale information can be retrieved in mosaics of single-field observations, provided that the sampling of the single fields follows the sky-plane Nyquist sampling theorem.  As a result, all the information about the sky brightness is coded in the visibility function. From a data-processing viewpoint, all the current radio-interferometric wide-field imaging methods~\\citep[see, \\eg{},][]{gueth95,sault96b,cornwell93,bhatnagar04,bhatnagar08,cotton08} belong to the mosaicking family\\footnote{In the rest of this paper, stop-and-go mosaicking refer to the observing mode, while mosaicking alone refer to the imaging family.} pioneered by \\citet{cornwell88}.  In this family, the processing starts with Fourier transforming \\V{}(\\up{},\\as{}) along the \\up{} dimension (\\ie{} at constant \\as{}) to produce a set of single-field dirty images before linearly combining them and forming a wide-field dirty image. In this paper, we propose an alternate processing, which starts with a Fourier transform of \\V{}(\\up{},\\as{}) along the \\as{} dimension (\\ie{} at constant \\up{}). We show how this explicitely synthesizes the spatial frequencies needed to do wide-field imaging, which are linearly combined to form a ``wide-field $uv$ plane'', \\ie{}, one $uv$-plane containing all the spatial frequency information measured during the wide-field observation.  Inverse Fourier transform will produce a dirty image, which can then be deconvolved using standard methods. The existence of two different ways to extract the wide-field information from the visibility function raises several questions: Are they equivalent? What are their relative merits?  We thus aim at revisiting the mathematical foundations of wide-field imaging and deconvolution. Sections~\\ref{sec:notations} to~\\ref{sec:mosaicking} propose the new algorithm, which we call wide-field synthesis: Section~\\ref{sec:notations} first defines the notations and it then lays out the basic concepts used throughout the paper. Section~\\ref{sec:ekers-rots} states the steps needed to go beyond the Ekers \\& Rots scheme and explores the consequences of incomplete sampling of both the $uv$ and sky planes.  Section~\\ref{sec:gridding} discusses the effects of gridding by convolution and regular resampling. Section~\\ref{sec:deconvolution} describes how to influence the dirty beam shapes and thus the deconvolution.  Section~\\ref{sec:short-spacings} states how to introduce short spacings measured either from a single-dish antenna or from heterogeneous interferometers. Section~\\ref{sec:mosaicking} compares the proposed wide-field synthesis algorithm with standard nonlinear mosaicking. Some detailed demonstrations are factored out in Appendix~\\ref{sec:demo} to enable an easier reading of the main paper, while ensuring that interested readers can follow the demonstrations. Appendices~\\ref{sec:projection} and~\\ref{sec:otf} then explain how the wide-field synthesis algorithm can cope with non-ideal effects: Appendix~\\ref{sec:projection} discusses how at least one standard way to cope with sky projection problems is compatible with the wide-field synthesis algorithm.  Appendix~\\ref{sec:otf} explores the consequences of using the on-the-fly observing mode. Finally, we assume good familiarity with single-field imaging in various places. We refer the reader to well-known references: \\eg{} chapter 6 of \\citet{thompson86} and~\\citet{sramek89}.  \\vspace{-\\medskipamount} ",
        "conclusions": "Interferometric wide-field imaging implies scanning the sky in one way or another (\\eg{} stop-and-go mosaicking, on-the-fly scanning, sampling of the focal plane by multi-beams). This produces sampled visibilities \\SV{}, which depends both on the $uv$-plane and sky coordinates (\\eg{}, \\up{} and \\as{}).  Based on a basic idea by~\\citet{ekers79}, we proposed a new way to image the interferometric wide-field sampled visibilities: \\SV{}(\\up,\\as).  After gridding the measured visibilities both in the $uv$ and sky planes, the gridded visibilities \\SV[\\G{}]{} are Fourier-transformed along the \\as{} sky dimension, yielding synthesized visibilities \\SVft[\\G{}]{} sampled on a $uv$ grid whose cell size is related to the total field of view; \\ie{}, it is much finer than the diameter of the interferometer antennas. We thus proposed calling this processing scheme ``wide-field synthesis''.  The Fourier transform is performed for each constant \\up{} value. As many independent estimates of the $uv$ plane are produced as independent values of \\up{} measured.  A shift-and-average operator is then used to build a final, wide-field $uv$ plane, which translates into a wide-field dirty image after inverse Fourier transform, \\ie{}, \\begin{equation} \\Iwfft[\\G{}](\\uf) \\Definition \\sint{\\up}{\\Wu{}(\\up,\\uf-\\up)\\,\\SVft[\\G{}]{}(\\up,\\uf-\\up)}, \\label{eq:shift-and-average} \\end{equation} where \\Wu{} is a normalized weighting function. Using these tools, we demonstrated that: \\begin{enumerate} \\item The dirty image (\\Iwf[\\G{}]{}) is a convolution of the sky brightness distribution (\\I{}) with a set of wide-field dirty beams (\\Dw[\\G{}]{}) varying with the sky coordinate \\af{}, \\ie{}, \\begin{equation} \\Iwf[\\G{}](\\af) = \\sint{\\aff}{\\Dw[\\G{}]{}(\\af-\\aff,\\af)\\,\\I{}(\\aff)}. \\end{equation} Compared to single-field imaging, the dependency on the primary beam is transferred from a product of the sky brightness distribution into the definition of the set of wide-field dirty beams. \\item The set of gridded dirty beams (\\Dw[\\G{}]{}) can be computed from the ungridded sampling function (\\S{}), the transfer function (\\Bft{}, the inverse Fourier transform of the primary beam), and the gridding convolution kernel (See Eq.~\\ref{eq:samp:def:1}, \\ref{eq:grid:aver} and~\\ref{eq:grid:samp}). \\item The dependence of the wide-field dirty beams on the sky position is slowly-varying, with their shape varying on an angular scale typically larger than or equal to the primary beamwidth. \\end{enumerate} Adaptations of the existing deconvolution algorithms should be straightforward.  A comparison with standard nonlinear mosaicking shows that it is not mathematically equivalent to the wide-field synthesis proposed here, though both methods do recover the sky brightness. The main advantages of wide-field synthesis over standard nonlinear mosaicking are \\begin{enumerate} \\item Weighting is at the heart of the wide-field synthesis because it is an essential part of the shift-and-average operation. Indeed, not only can a multiplicative weight be attributed to each visibility sample before any processing, but the $uv$-plane weighting function (\\Wu{}, see Eq.~\\ref{eq:shift-and-average}) is also a degree of freedom, which should be set according to the conditions of the observation and the imaging goals, \\eg{} highest signal-to-noise ratio, highest resolution, or most uniform resolution over the field of view.  The \\Wu{} weighting function thus enables us to modify the wide-field response of the instrument. On the other hand, mosaicking requires a precise weighting function in the image plane, which freezes the wide-field response of the interferometer. \\item Wide-field synthesis naturally processes the short spacings from both single-dish antennas and heterogeneous arrays along with the long spacings. Both of them can then be jointly deconvolved. \\item The gridding of the sky plane dimension of the measured visibilities, required by the wide-field synthesis, may potentially save large amounts of hard-disk space and cpu processing power relative to mosaicking when handling data sets acquired with the on-the-fly observing mode. Wide-field synthesis could thus be particularly well suited to process on-the-fly observations. \\end{enumerate}  The wide-field synthesis algorithm is compatible with the $uvw$-unfaceting technique devised by~\\citet{sault96a} to deal with the celestial projection effect, known as non-coplanar baselines (see appendix~\\ref{sec:projection}). Finally, on-the-fly observations imply an elongation of the primary beam along the scanning direction. These effects can be decreased by an increase in the primary beam sampling rate. However, it may limit the dynamic range of the image brightness if the primary beam sampling rate is too coarse (see appendix~\\ref{sec:otf})."
    },
    "1004/1004.4765_arXiv.txt": {
        "abstract": "We examine the temperature structure of the intergalactic medium (IGM) surounding a hard radiation source, such as a Quasi-Stellar Object (QSO), as it responds to the onset of helium reionization by the source. We model the reionization using a radiative transfer (RT) code coupled to a particle-mesh (PM) $N$-body code. Neutral hydrogen and helium are initially ionized by a starburst spectrum, which is allowed to gradually evolve into a power law spectrum ($\\propto \\nu^{-0.5}$). Multiple simulations were performed with different times for the onset and dominance of the hard spectrum, with onset redshifts ranging from $z=3.5$ to 5.5. The source is placed in a high-density region to mimic the expected local environment of a QSO. Simulations with the source placed in a low-density environment were also performed as control cases to explore the role of the environment on the properties of the surrounding IGM. We find in both cases that the IGM temperature within the \\HeIII\\ region produced exceeds the IGM temperature before full helium reionization, resulting in a ``thermal proximity effect,'' but that the temperature in the \\HeIII\\ region increases systematically with distance from the source. With time the temperature relaxes with a reduced spread as a function of impact parameter along neighbouring lines of sight, although the trend continues to persist until $z=2$. Such a trend could be detected using the widths of intervening metal absorption systems using high resolution, high signal-to-noise ratio spectra. By contrast, the Doppler widths of \\HI\\ absorption features in mock spectra along neighbouring lines of sight show a weak trend with impact parameter prior to full helium reionization reflecting the behaviour of the underlying density and peculiar velocity fields, but they take on a near constant value after helium reionization, with a median value near $25\\kms$ at $z=3$, in good agreement with observations. ",
        "introduction": "\\label{sec:Intro}  An abundance of observational evidence shows that subsequent to the Recombination Era, during which both hydrogen and helium recombined to their neutral forms, the Universe was reionized. The best constraints on the epoch of hydrogen reionization are derived from measurements of the intergalactic Lyman resonance line transitions in the spectra of high redshift Quasi-Stellar Objects (QSOs) and polarization measurements of the Cosmic Microwave Background (CMB). The measured optical depths of intergalactic \\Lya\\ and \\Lyb\\ absorption in high redshift QSOs show that the Universe was reionized by $z\\gsim 6$ \\citep{Becker01, 2002AJ....123.1247F}. This is consistent with CMB measurements by the {\\it Wilkinson Microwave Anisotropy Probe} ({\\it WMAP}), which place the hydrogen reionization epoch, if a sudden process, at $z=11.0\\pm1.4$ \\citep{2009ApJS..180..306D}. The sources of hydrogen reionization are currently unknown, but are strongly suspected of being young star-forming galaxies. More speculative possibilities include pockets of Population III stars, as may arise in isolated star clusters \\citep{1999ApJ...514..648M, CF06}, or miniquasars \\citep{Madau04}. Hydrogen reionization by QSOs is untenable based on current estimates of the QSO luminosity function, unless there is a sharp rise at the faint end at high redshifts \\citep{Meiksin05, 2009ApJ...692.1476C}.  The reionization epoch of helium, by contrast, is far less well constrained. This is primarily because of the need to go to space to measure the intergalactic Lyman series absorption of intergalactic gas in the spectra of high redshift QSOs. The current constraints place the helium reionization epoch, when helium becomes nearly fully doubly ionized, at $z\\gsim3$ \\citep{Reimers05}. Because of the high photoelectic threshold energy of \\HeII, it is expected that helium was reionized by QSOs rather than stars, most likely in the redshift interval $z=3-4$ \\citep{MM94, Meiksin05, 2009ApJ...694..842M}.  Measurements of the \\HeII\\ \\Lya\\ optical depth show the optical depth rises rapidly approaching $z\\lsim3$ with an increasing amount of patchiness \\citep{Zheng04, 2006A&A...455...91F}. By $z\\approx3.4$, the optical depth is immeasurably large \\citep{2008ApJ...686..195Z}. The limits are not yet strong enough to demonstrate the epoch of helium has been detected, but the rapid rise in optical depth and degree of patchiness are suggestive of an approach to the helium reionization epoch. This interpretation, however, is not unique:\\ it may also indicate a diminishing attenuation length of \\HeII-ionising photons, resulting in large fluctuations in the \\HeII-ionising metagalactic background due to large local Poisson fluctuations in the numbers of QSOs contributing to the background \\citep{2007arXiv0711.3358M}. The uncertainty in QSO counts and the spectral shape of high redshift QSOs still does not allow much earlier reionization, up to $z\\lsim5.5$, to be excluded \\citep{Meiksin05}.  In this paper, we examine an alternative means of probing the epoch when helium in the intergalactic medium (IGM) surrounding a QSO was reionized:\\ detections of a heightened IGM temperature above the post-reionization equilibrium value. The temperature following full helium reionization will reach high values as the helium ionization front (I-front) passes, eventually establishing a lower temperature in overdense regions as the gas achieves thermal equilibrium \\citep{Meiksin94, 2007MNRAS.380.1369T, 2009ApJ...694..842M}.  If the IGM near the QSO is polluted with metal lines, then a direct determination of the gas temperature may be made by combining the measured linewidths of the metals, provided they are adequately resolved \\citep{1996ApJ...467L...5R, 2007arXiv0711.3358M}. Because of the low abundances of intergalactic metals, however, measurements have so far been restricted to high \\HI\\ column density systems. The advent of large telescopes, like the Thirty Meter Telescope\\footnote{\\tt http://www.tmt.org} or an Extremely Large Telescope\\footnote{\\tt http://www.eso.org/sci/facilities/eelt/}, may permit temperature measurements to be extended to much smaller column densities.  Whilst such direct temperature estimates are currently limited to high column density absorbers, the reionization of helium by a QSO will induce a ``thermal proximity effect'' in which the increased temperature near the QSO will be reflected by broadened hydrogen, as well as helium and metallic, absorption lines. The effect has been exploited as a means of probing the epoch of full helium reionization through the global statistical effect on the distribution of the \\HI\\ Doppler parameters for systems throughout the IGM. Rises in the lower cutoff of the Doppler parameter distribution have been reported by \\citet{2000ApJ...534...41R} and \\citet{Schaye00}, and interpreted, based on calibrations with numerical simulations, as evidence for a sharp temperature jump and reduction in the effective polytropic index towards isothermality near $z=3.0$, due to the onset of \\HeII\\ reionization. A similar conclusion was reached by \\citet{2002ApJ...564..153Z} in a wavelet-based analysis of QSO spectra. Global reionization simulations and semi-analytic estimates allowing for distributed sources, however, are not clearly in accord with such sharp changes in the absorption line parameters. The results suggest only a gradual rise in the mean IGM temperature should occur during helium reionization \\citep{2009ApJ...694..842M, 2009MNRAS.395..736B}. Searches for isolated regions of increased temperature, as may arise around a helium-ionising QSO, have been only partially successful. \\citet{2002ApJ...564..153Z} argues at least ten QSO sitelines are required to probe temperature fluctuations at a statistically significant level. Using a similar method, \\citet{2002MNRAS.332..367T} find a region $10^4\\kms$ across suggestive of a statistically significant elevated temperature in a single QSO siteline out of eight examined.  A more direct means of searching for the reionization signature is to target the region around specific QSOs. The temperature of the IGM gas surrounding a QSO may be probed both using lines of sight towards the QSO and nearby on the sky towards background QSOs. Surveys of QSO neighbours with follow-up spectroscopy along multiple lines of sight are currently underway in an attempt to measure the transverse hydrogen ionization proximity effect and to use transverse spatial correlations in the hydrogen \\Lya\\ forest as a constraint on the vacuum energy density of the Universe \\citep{2008ApJS..175...29M}. The detection of a thermal proximity effect would provide additional support for unification models of Active Galactic Nuclei (AGN). In these models, the variety of AGN activity observed, from radio galaxies to QSOs and from Seyfert 1s to Seyfert 2s, is a consequence of geometry:\\ an obscuring dusty torus or other obstacle blocks the optical/UV radiation from the central engine along some lines of sight, possibly re-radiating it in the infra-red. The observed character of the object then depends on the viewing angle. While good evidence exists for elements of the model \\citep{2005SSRv..119..355V, 1993ARA&A..31..473A}, direct evidence for optical/UV radiation transverse to the line-of-sight of a putative obscured QSO or Seyfert 1 is generally lacking. Prior to the completion of helium reionization in the IGM, the broadening of \\Lya\\ absorption features transverse to the line of sight toward bright radio or infra-red galaxies, Type II QSOs, or Seyfert 2 galaxies would provide evidence for intense photoionising UV radiation beamed transverse to the line of sight, as predicted by the unification models.  The purpose of this paper is to use numerical reionization simulations to quantify the impact of \\HeII\\ reionization by a QSO on the temperature and absorption line widths produced in the surrounding gas. The simulations are performed using a radiative transfer code coupled to an $N$-body code to model the evolution of the IGM. The method was previously described in \\citet{2007MNRAS.380.1369T}. The details of the simulation volume and sources are given in \\sect{sec:Simulations}. Results of the simulations are provided in \\sect{sec:Results} and discussed in \\sect{sec:Discussion}. Our conclusions are summarised in \\sect{sec:Conclusions}. In an Appendix, we describe modifications to the radiative tranfer code implemented to simulate reionization by a central source. ",
        "conclusions": "\\label{sec:Conclusions}  \\begin{figure} \\scalebox{0.5}{\\includegraphics{fig12.eps}}  \\caption{The energy per \\HeII\\ ionization as a function of the \\HeI\\ and \\HeII\\ ionization fractions, $x_{\\rm HeI}$ and $x_{\\rm HeII}$, respectively, forming narrowly-defined contours. The dependence is well-parametrised as a function of $x_{\\rm HeII} + 2x_{\\rm HeI}$ (thin line). The contour levels are spaced by a decade. The left contours correspond to $x_{\\rm HeI}=0$. The data are for $z=4.0$. } \\label{fig:E_per_ion_parametrised_contour} \\end{figure}  Using a coupled radiative-transfer and Particle-Mesh $N$-body code, we study the effects of helium reionization by a QSO source on the surrounding IGM. We consider QSO source turn-on redshifts ranging from $z=3.5$ to $z=5.5$, with the QSO spectrum modelled as a power law. In the overdense environment expected for a QSO, the crossing of the \\HeIII\\ ionization front boosts the temperature behind the front to median values of $T\\approx20-30\\times10^3$~K, gradually relaxing to $T\\approx10\\times10^3$~K by $z=2$ as thermal equilibrium between photoionization heating and atomic radiative cooling processes and adiabatic expansion cooling is established. The peak temperature tends to be higher for earlier turn-on redshifts. A ``thermal proximity effect'' results with a heightened IGM temperature within the \\HeIII\\ region produced. An identical source placed in an underdense environment shows similar behaviour.  The median gas temperature surrounding the QSO, however, is not uniform; it shows an increasing trend towards the \\HeIII\\ ionization front, where spectral filtering by the intervening gas still hardens the radiation field from the source in patches and where helium was most recently ionized. We generate synthetic spectra to assess the possibility of detecting the trend in the absorption lines along the lines of sight to background neighbouring QSOs. The spectral signature of the post-helium ionization temperature trends nearly vanishes in the \\HI\\ Doppler parameters, which show both little evolution and little dependence on distance from the source. For a source placed in an overdense region, a median Doppler parameter of $25\\kms$ at $z=3$ is found, independent of impact parameter and the time of \\HeII\\ reionization. A somewhat lower median Doppler parameter of $23\\kms$ is found for a source placed in the centre of a void. The median values are in good agreement with the measured median Doppler parameter. This is in contrast to simulations without radiative transfer, for which the value is underpredicted by a few to several $\\kms$ \\citep{MBM01}. Late \\HeII\\ reionization appears to account for the higher values. The absence of a dependence of the median Doppler parameter on redshift and impact parameter may be attributed to the absence of any temperature trend with redshift and impact parameter of the denser regions which dominate the \\Lya\\ absorption spectra. By contrast, a trend {\\it is} found prior to full helium reionization, with an increase in the Doppler parameters for a smaller impact parameter of the line of sight, reflecting similar behaviour in the peculiar velocity field. Thus the disappearance of a trend with impact parameter may indicate the onset of full helium reionization.  A near polytropic relation is obtained between the gas temperature and density with a slope that is close to adiabatic, extending to overdensities of $\\sim10^2$ in the vicinity of the source placed in a massive halo. The temperatures found, approaching $10^5$~K, are those expected for accretion onto halos and sheets, with radiative cooling lowering the temperature at higher overdensities. In regions without a massive halo, the turnover in the temperature occurs at overdensities of only a few, again as expected from radiative cooling losses. No inverted temperature--density relation is found in the underdense gas. Whilst the temperature--density relation flattens soon after \\HeII\\ reionization, the temperature in the underdense gas follows a general trend of increasing with density, although with a wide spread in values.  The detection of metal absorption systems would allow a more direct probe of the gas temperature. Because the peculiar velocity contribution to the velocity width of an absorption line would be identical for features arising from different elements, while the thermal contribution (in quadrature) decreases like the mass of the ion, it is possible to separate these two contributions provided absorption features due to two or more elements from the same structure are detected. Such temperatures have been measured for a few systems, although the corresponding \\HI\\ column densities tend to be high, typically exceeding $10^{14}\\,{\\rm cm^{-2}}$, corresponding to overdensities of a few to several (eg, \\cite{2007arXiv0711.3358M}). The temperature trend as a function of impact parameter is less clearly defined for such overdensities, with a weak trend of higher values for smaller impact parameters. None the less, a spread with impact parameter could indicate the onset of full helium reionization, particularly if very high temperatures exceeding $30\\times10^3$~K are found at $z<4$.  It would be very interesting to use metal absorption features arising in the very diffuse IGM, corresponding to \\HI\\ column densities of $10^{13}\\,{\\rm cm^{-2}}$ and lower, to probe the temperature structure of the IGM, both before and after full helium reionization. There are currently no detections of individual metal absorption features in such low density diffuse systems, but neither do observations preclude them. Exploiting them for temperature measurements, however, may need to await spectroscopy on a much larger future generation of telescopes such as the Thirty Meter Telescope or an Extremely Large Telescope."
    },
    "1004/1004.4553_arXiv.txt": {
        "abstract": "In this paper, we present an automated system, which has the capability to catch and track solar limb prominences based on observations from EUV 304~\\AA\\ passband.  The characteristic parameters and their evolution, including height, position angle, area, length and brightness, are obtained without manual interventions. By applying the system to the STEREO-B/SECCHI/EUVI 304~\\AA\\ data during 2007 April -- 2009 October, we obtain a total of 9477 well-tracked prominences and a catalog of these events available online at \\url{http://space.ustc.edu.cn/dreams/slipcat/}. A detailed analysis of these prominences suggests that the system has a rather good performance. We have obtained several interesting statistical results based on the catalog. Most prominences appear below the latitude of 60 degrees and at the height of about 26 Mm above the solar surface. Most of them are quite stable during the period they are tracked. Nevertheless, some prominences have an upward speed of more than 100 km/s, and some others show significant downward and/or azimuthal speeds. There are strong correlations among the brightness, area and height. The expansion of a prominence is probably one major cause of its fading during the rising or erupting process. ",
        "introduction": "Solar prominences, also called filaments when they are viewed on-disk, are long-observed but still not well-known structures in the solar atmosphere. Since they are outstanding features in multiple-wavelength observations of the Sun and have close relationships with various solar eruptive phenomena, prominences are always one of major topics in solar and space physics. The key issues in prominence/filament studies are their formation, maintenance, dynamic processes, and their roles in other related solar activities, e.g., coronal mass ejections (CMEs) and flares. With the aid of modern sensor technology, many facts of prominences are revealed \\citep[e.g.,][and the references therein]{Poland_1986, Tandberg-Hanssen_1995, Martin_1998, Patsourakos_Vial_2002}. Prominences are dense (electron density $\\sim10^{9}-10^{11}$ cm$^{-3}$) and cool ($\\sim5000-8000$ K) plasmas  floating in hot and diluted solar corona \\citep[e.g.,][]{Engvold_Brynildsen_1986, Hiei_etal_1986, Hirayama_1986, Madjarska_etal_1999}. They can appear anywhere from active regions to polar regions, and live for days to months. They have spines and barbs, and always straddle above polarity inversion lines. There are sometimes strong counterstreamings along spines and very dynamic vertical flows. The chirality of prominences/filaments obeys the pattern that most prominences on northern hemisphere are dextral while most ones on southern hemisphere are sinistral. The association rate of eruptive prominences with CMEs is more than about 70\\% \\citep[e.g.,][]{Gilbert_etal_2000, Gopalswamy_etal_2003}.  Many of the above findings are made through statistical investigations combined with case studies. A continuously updated catalog of prominences with unbiased parameters is undoubtedly helpful for such researches, especially in the age of the explosive growth of observational data. For instance, the successful launch of STEREO (Solar Terrestrial Relationship Observatory) spacecraft (A and B) in 2006 led the amount of solar observations explosively growing to more than 12 GB a day, and now it has increased to about 2 TB a day from SDO (Solar Dynamic Observatory), which was just launched in February 2010. NOAA/SWPC\\footnote{http://www.swpc.noaa.gov/Data/index.html} routinely compiles a list of solar events, in which on-disk filaments and limb eruptive prominences are included; but the list is far from complete. Thanks to the unique properties of prominences/filaments, they can be clearly observed at multiple wavelengths, such as H$\\alpha$, He I 10830~\\AA, He II 304~\\AA, radio waves, etc \\citep[e.g.,][]{Schmahl_etal_1974, Hanaoka_etal_1994, Penn_etal_1994, ChiuderiDrago_etal_2001, Labrosse_Gouttebroze_2001}, and therefore  it is possible to  extract them from the vast amount of data automatically and consistently.  Recognitions of on-disk filaments and limb prominences are different. The former is mainly accomplished by studying H$\\alpha$ data. For example, \\citet{Gao_etal_2002}, \\citet{Shih_Kowalski_2003}, \\citet{Fuller_etal_2005} and \\citet{Zharkova_etal_2005} developed codes to automatically detect filaments in full-disk H$\\alpha$ images. The automated system developed by \\citet{Bernasconi_etal_2005} is able to detect, classify and track H$\\alpha$ filaments efficiently. EUV observations are much more difficult to be used in detecting on-disk filaments due to the low contrast of filaments and the involvement of coronal holes. However, EUV observations are suitable for limb prominence detection. Through the usage of Fe IX/X 171~\\AA, Fe XII 195~\\AA, Fe XV 284~\\AA\\ and  He II 304~\\AA\\ images from SOHO/EIT (\\citet{Delaboudiniere_etal_1995}) instrument, \\citet{Foullon_Verwichte_2006} developed algorithms to recognize limb prominences. In their method, He II 304~\\AA\\ data  provide the basic criteria for the selection of candidate prominence regions, and other emission lines are used to remove active regions, which also appear brightly in EUV 304~\\AA. Most recently, \\citet{Labrosse_etal_2010} also developed an efficient code to detect limb prominences in EUV 304~\\AA\\ images.  In this paper, we present an automated system of detecting and tracking solar limb prominences based on only He II 304~\\AA\\ data, and a resultant on-line catalog as well, which can be continuously updated. The performance and limitations of the system are presented in section \\ref{sec_performance}. Based on our catalog, some preliminary statistical results of solar limb prominences are also presented. The reasons we choose He II 304~\\AA\\ emission line rather than H$\\alpha$ are the following. First, for prominence/filament observations, He II 304~\\AA\\ line is the only one uninterruptedly imaging the Sun with high cadence (operated by space-borne instruments SECCHI/EUVI on board STEREO twins, and AIA on board SDO). A complete database of limb prominences is therefore possible to be established. Secondly, the high time resolution of the data allows us to track their evolution, even small changes. Thirdly, the projection effect can be minimized for certain  parameters, such as height, radial speed, etc. Fourth, they are complementary to the catalogs of on-disk filaments. Further, there is so far no well-established on-line catalog for limb prominences. ",
        "conclusions": "We have developed an automated system of catching and tracking solar limb prominences in EUV 304~\\AA\\ images. The system, called SLIPCAT, is able to generate (1) a catalog of solar limb prominences and (2) some characteristic parameters of each detected prominence, including the height, position angle, area, length, brightness and their first and second derivatives with respect to time. SLIPCAT is composed by five modules, (1) prominence candidate selection, (2) parameter extraction, (3) non-prominence feature removal, (4) prominence tracking, and (5) catalog generating. At present, an online catalog for STEREO-B/EUVI 304~\\AA\\ data has been generated (refer to \\url{http://space.ustc.edu.cn/dreams/slipcat/}), and catalogs for STEREO-A/SECCHI/EUVI and SDO/AIA data are in preparation.  Based on the STEREO-B/EUVI 304~\\AA\\ data, SLIPCAT proved to perform well in detecting limb prominences. \\begin{enumerate} % \\item It can distinguish real prominences from non-prominence features, e.g., active regions, without observations in other wavelengths by using the technique of linear discriminant analysis. The goodness of the overall classification is at about 86\\% successful rate.  \\item It detects as many as 9477 well-tracked prominences during 2007 April -- 2009 October, which means in average about 10 events per day. Compared to H$\\alpha$ data, it is found that SLIPCAT is sensitive enough to recognize almost all prominences, even those invisible in H$\\alpha$ images or very small ones.  \\item Thanks to the high-cadence EUV 304~\\AA\\ data, SLIPCAT is able to provide the detailed evolution processes of prominences quantitatively without manual interventions. The upper-right panel of Figure \\ref{fg_duration} implies that a well-tracked prominence is at least detected in 28 images in average; and the case in Figure \\ref{fg_20081113_kin} shows that such high-cadence detection allows us to make a detailed analysis of  its evolution, including its eruption, oscillation, etc. \\end{enumerate} However, not all the parameters extracted by SLIPCAT can precisely reveal the real behavior of prominences. The limitations have been addressed in the last paragraph of Sec.\\ref{sec_performance}. Summarized here, they are that (1) the parameters characterize the properties of prominences during the period they are detected, not over the whole life-time, (2) they suffer from the projection effect, and (3) the speeds, change rates of area, length and brightness, which are derived from linear and quadratic fittings, may not be accurate.  By applying SLIPCAT to the STEREO-B/EUVI 304~\\AA\\ data from 2007 April to 2009 October, we obtain the following preliminary statistical results of solar limb prominences. \\begin{enumerate} % \\item In average, there are about 10 prominences standing above the solar limb per day during the solar minimum. For most days, about 14 prominences are expected to be detected, and sometimes the number could be as large as 32 or as small as zero.  \\item Most (99\\%) prominences appear below latitude of 60 degrees, and the long extended prominences tend to arise between latitudes of 30 and 60 degrees.  \\item Most (82\\%) prominences have a height of about 26 Mm from the solar surface.  \\item The projected area of a prominence on the plane-of-sky is about 1072 Mm$^2$ in average, and nearly 60\\% of prominences have a smaller area.   \\item Most prominences are quite stable during the period they are detected; no obvious change in position, area or brightness can be found.  \\item Particularly, more than 80\\% of prominences do not show obvious motion in either radial or azimuthal direction.  However, some prominences have an upward speed of more than 100 km/s, and a few prominences present a significant downward or azimuthal speed.  \\item The brightness of prominences is anti-correlated with the height. The prominences at higher altitude look dimmer.  \\item The area of prominences is positively correlated with the height. The prominences at higher altitude are generally larger.  \\item From the statistical point of view, the expansion of prominences is probably one of the major causes of the fading of prominences during their rise or eruption. \\end{enumerate}   \\paragraph{Acknowledgments.} We acknowledge the use of the data from STEREO/SECCHI. This research is supported by grants from NSFC 40525014, 973 key project 2006CB806304, FANEDD 200530, and the fundamental research funds for the central universities."
    },
    "1004/1004.1189_arXiv.txt": {
        "abstract": "Analyses have found a ``haze'' of anomalous microwave emission surrounding the Galactic Center in the WMAP sky maps.  A recent study using Fermi data detected a similar haze in the $\\gamma$-ray.   Several studies have modeled these hazes as radiation from the leptonic byproducts of dark matter annihilations, and arguably no convincing astrophysical alternative has been suggested.  We discuss the characteristics of astrophysical cosmic ray sources that could potentially explain this microwave and $\\gamma$-ray emission.  The most promising astrophysical scenarios involve cosmic ray sources that are clustered such that many fall within $\\sim 1$~kpc of the Galactic Center.  For example, we show that several hundred Galactic Center supernovae in the last million years plus a diffusion-hardened electron spectrum may be consistent with present constraints on this emission.  Alternatively, it could be due to a burst of activity probably associated with Sagittarius A$^*$ occurring $\\sim 1$~Myr ago and producing $>10^{51}$~erg in cosmic ray electrons.  Different models predict contrasting trends for the spectral index of the microwave and $\\gamma$-ray spectrum as a function of angle from the Galactic Center that we show should be robust to cosmic ray propagation uncertainties.  In particular, if the haze is from dark matter annihilations, it should have a very hard microwave and $\\gamma$-ray spectrum for which the spectral shape does not change significantly with angle, which we argue would be difficult to achieve with any astrophysical mechanism.  Observations with the Planck and Fermi satellites can distinguish between viable haze models using these signatures. ",
        "introduction": "Using data from the Wilkinson Anisotropy Probe (WMAP), \\citet{finkbeiner04a} discovered an approximately circular excess of diffuse emission centered around the Galactic Center of radial extent $\\sim 20^\\circ$ and luminosity $L_{\\rm haze} \\sim 10^{37} \\; {\\rm erg \\, s^{-1}}$ between $20$ and $60$ GHz. Nicknamed the ``WMAP haze,'' this radiation was the residual emission after \\citet{finkbeiner04a} subtracted templates for Galactic H$\\alpha$, synchrotron, and dust emission from the WMAP foreground maps.  \\citet{finkbeiner04a} argued that the spectrum of the haze was too hard to be due to synchrotron from cosmic ray electrons produced in supernovae.  Therefore, this emission was initially attributed to free-free emission from $\\gtrsim  10^5$ K gas \\citep{finkbeiner04a}.  However, in a follow-up study \\citet{finkbeiner04} argued that the lack of diffuse X-ray emission from the Galactic Center ruled out this hypothesis.  If the haze is not due to free-free emission, it must be synchrotron.  Interestingly, \\citet{finkbeiner04} showed that the energetics and profile of the haze emission could be explained by synchrotron emission from the byproducts of dark matter annihilations for a generic weakly interacting $\\sim 100$ GeV particle with a standard dark matter halo profile.   This idea has been pursued further in several subsequent studies \\citep{hooper07, hoopergray, zhang08, cholis08, cumberbatch09,  harding09}.  However, cosmic rays from astrophysical sources may also be able to create the haze emission. The total energy in synchrotron electrons required to produce the haze is \\begin{equation} E_{\\rm syn}^{\\rm haze} \\sim 10^{51}\\, \\left(\\frac{L_{\\rm haze}}{10^{37} {\\rm erg \\, s^{-1}}} \\right) \\, \\left(\\frac{10 \\, {\\rm \\mu G}}{B} \\right)^{\\frac{3}{2}}  \\left(\\frac{50 \\,  {\\rm GHz}}{\\nu}\\right)^{\\frac{1}{2}} ~{\\rm erg}. % \\label{eqn:energy} \\end{equation} Equation (\\ref{eqn:energy}) is derived by equating $L_{\\rm haze}$ with the synchrotron loss rate for an electron that emits at $\\nu$ in a magnetic field of strength $B$, noting the relationship between electron energy $E$ and the synchrotron peak frequency $\\nu$: \\begin{equation} E \\approx  30  \\; \\left( \\frac{\\nu}{50 \\; {\\rm GHz}} \\right)^{1/2} \\; \\left(\\frac{B_\\perp}{10 \\; {\\rm \\mu G}}\\right)^{-{1}/{2}} ~ {\\rm GeV},  % \\end{equation} where $B_\\perp$ is the magnetic field component orthogonal to the particle motion. For comparison, $\\sim 10^{51}~$erg is roughly equal to the kinetic energy of a single supernova explosion.   This energy must be output within the time it takes an electron to radiate its energy via synchrotron and inverse Compton or \\begin{equation} \\Delta t = 1.0 \\;  \\left( \\frac{B_\\perp}{10 \\; {\\rm \\mu G}} \\; \\frac{50 \\, {\\rm GHz}}{\\nu} \\right)^{{1}/{2}} \\left( \\frac{U_{\\rm r} + U_{\\rm B}}{10 \\, {\\rm eV \\,  cm^{-3} }} \\right)^{-1}  ~ {\\rm Myr}, \\label{eqn:cooling} \\end{equation} where $U_{\\rm r}$ and $U_{\\rm B}$ are respectively the energy densities in radiation and magnetic fields.   While it is conceivable that astrophysical sources could satisfy these energetics, the difficulties with creating the haze with astrophysical sources are threefold: \\begin{enumerate} \\item \\citet{dobler08} find the haze spectrum in the microwave to be hard, measuring a power-law electron spectrum with index $\\gamma \\approx 2$.  Such a spectrum is much harder than the galactic synchrotron towards other directions (where $\\gamma \\approx 3$). \\item The distribution of known galactic cosmic ray sources appears to not be sufficiently concentrated toward the Galactic Center to create the haze \\citep{finkbeiner04, dobler08, zhang08}.  The distribution of supernovae and pulsars is estimated to peak at $\\approx 4$ kpc from the Galactic Center, corresponding to an angular separation of $\\approx 30^\\circ$ \\citep{lorimer06}.  In contrast, the haze intensity is strongly increasing with decreasing angle at $<20^\\circ$ from the Galactic Center \\citep{dobler08}. \\item By construction, a similar excess towards the Galactic Center is not present at $408$ MHz.  If the haze is from synchrotron, this suggests that the source of the haze does not contribute significantly to the population of cosmic ray electrons at $\\sim 5$~GeV.  In addition, the sky at $408~$MHz is dominated by emission from the galactic disk, without a significant enhancement in the haze region. \\end{enumerate}  No study that has proposed an astrophysical explanation for the haze has demonstrated that it can satisfy all of these criteria.  Pulsars are the primary astrophysical solution that has been discussed in the literature \\citep{zhang08, kaplinghat09, harding09}.  We argue here that a disk population of pulsars likely cannot satisfy the third criterion.  %  If dark matter annihilations from the Milky Way's hypothetical dark matter density cusp are the source of the haze, \\citet{finkbeiner04} found that these three criteria can easily be satisfied.  Furthermore, for standard dark matter models, large clumping factors or boost factors in the annihilation cross section above theoretically preferred values may not be required to explain this radiation \\citep{finkbeiner04, hooper07}.  If large boost factors are not required, this would be in contrast to other recent astrophysical anomalies that have been interpreted as arising from dark matter annihilations.  Recently, a $\\gamma$-ray ``haze'' was detected by \\citet{doblerFermi} using data from the Fermi Space Telescope, which has a similar spatial morphology to the WMAP haze.  This haze was interpreted as inverse Compton emission off of starlight from the same hard population of electrons that generate the haze in the microwave \\citep{doblerFermi}.  The $\\gamma$-ray haze may be another difficulty for astrophysical models of the haze that likely can be accommodated in dark matter models \\citep{cholis09}.  However, we argue that present measurements of this excess in the $\\gamma$-ray do not yet provide a stringent test for models of the haze.  The bulk of this paper investigates different explanations for the haze emission and their observational signatures.  Section \\ref{sec:cosmic_rays} summarizes the standard diffusive model for galactic cosmic ray propagation.  Section \\ref{sec:sources} details several candidate sources for the haze.  Section \\ref{sec:results} contrasts the microwave properties of the candidate haze models, and Section \\ref{sec:IC} contrasts these models' inverse Compton signatures.  Throughout this article we refer to the radiating particles as ``electrons'', even though some fraction may be positrons.% ",
        "conclusions": "We have investigated several models for the WMAP haze emission.  We argued that free-free emission is unlikely to be the source of this radiation, in agreement with previous studies.  Therefore, it is most likely that the haze is synchrotron.  If it is synchrotron, the haze indicates that the spectrum of electrons between a few and $20~$GeV is harder towards the Galactic Center.  We showed that it is difficult to produce such a hardened source of Galactic Center electrons with the standard disk population of supernovae.  In addition, we argued that a disk population of pulsars (which has been cast by the literature as the astrophysical alternative to the dark matter annihilation model) is also unlikely to produce a hardened spectrum \\emph{just} towards the Galactic Center.  The bulk of this paper investigated different models that may achieve sufficient hardening.  We showed that a Galactic Center population of supernovae and a diffusion-hardened electron distribution may be able to create the haze emission.  We found that $\\sim 1000$ Galactic Center supernovae per Myr could be consistent with present constraints on the energetics, profile, and (marginally) the  spectral hardness of the haze if a relatively large value for the diffusion length is assumed.  This supernovae rate is consistent with predictions for its value in the central couple hundred parsecs.  %  A transient event in the Galactic Center that releases $\\sim 5\\times 10^{51}$~erg in cosmic ray electrons in the past few Myr is another possibility that can reproduce the properties of the haze, even with a relatively soft spectrum ($\\gamma_{\\rm int} \\approx 2- 2.5$).  It is unlikely that a single stellar event can source such energies, and, thus, such an event would be almost certainly associated with Sagittarius~A$^*$.  Intriguingly, there is evidence for enhanced star formation in the Galactic Center $\\approx 2$~Myr ago.  This activity could have been accompanied by a brief episode of accretion onto the SMBH, potentially producing an accelerated population of electrons.  %  Synchrotron radiation from the byproducts of dark matter annihilations is the most exciting explanation for the WMAP haze \\citep{finkbeiner04, hooper08}. This explanation is naturally able to account for the measured microwave hardness of the WMAP haze emission and its radial extent.  This is true even if the diffusion length of $\\sim 30$~GeV electrons is much smaller than the radius of the haze, in contrast to the other haze models.   A relatively light particle ($\\lesssim 100$~GeV) is required in order not to invoke large boost factors in the annihilation rate beyond theoretically expected values.  However, even for a $100~$GeV WIMP, we require a boost factor of $\\sim 10$ unless the inner halo density profile is steeper than a power-law with index $-1.2$.   Electrons with energies that radiate in the microwave and $\\gamma$-ray are expected to lose significant energy as they diffuse a distance comparable to the extent of the haze. Because of how these losses occur in the different models, all of the considered models produce differing trends in the hardness of the microwave and $\\gamma$-ray spectra as a function of angle from the Galactic Center.  These trends do not rely on the specific details of galactic cosmic ray propagation model.  Instead, all of these trends owe to the theoretically expected and empirically confirmed fact that more energetic particles diffuse more quickly and cool faster than less energetic ones.  Previous observations in the microwave and $\\gamma$-ray are not yet able to constrain radial trends in the haze hardness, but future observations are projected to constrain them with high precision.  We illustrated these trends using simple models that were tuned to be consistent with the \\citet{dobler08} haze measurement.   In particular, dark matter annihilation models for the haze emission differ from astrophysical models in that a lot of the energy is input even at several kpc from the Galactic Center.  Because of this difference, the hardness of the synchrotron and inverse Compton emission in this model is not as strong a function of galactic latitude as in the astrophysical models.  This observable may allow one to distinguish this model from the others.   Future observations have the potential to significantly constrain potential explanations for the WMAP haze.  In particular: \\begin{itemize} \\item Radio maps at frequencies higher than $408$~MHz would probe electrons that are more akin to the population that emits in the WMAP bands, and such maps could test the two electron population model for the microwave synchrotron emission.  In fact, such data already exists at $1420~$MHz \\citep{2001A&A...376..861R}. \\item Planck temperature data (and possibly more sophisticated techniques applied to the WMAP data) can provide a better measurement of the haze profile as well as a measurement of the hardness as a function of angle. Planck's improved measurement relative to WMAP stems from the added channels at higher frequencies, which break degeneracies between the haze and other foregrounds \\citep{dobler08}. \\item Future CMB polarization measurements can search for a polarized haze.  While a recent analysis of WMAP polarization data did not detect a hard electron component \\citep{gold10}, the haze should be present at some level if it owes to a hard population of synchrotron electrons.% \\item $\\gamma$-ray observations with Fermi are sensitive to slightly more energetic electrons than those that emit in the microwave.  Future $\\gamma$-ray measurements have the potential to constrain the hardness of these electrons as a function of angle and thereby haze models. \\end{itemize}  We thank Greg Dobler for providing the data for the \\citet{dobler08} measurement of the haze, and Lars Hernquist, Philip Hopkins, Eliot Quataert, Avi Loeb, Ryan O'Leary, and Jennifer Yeh for useful discussions.  We would especially like to thank Greg Dobler  and David Spergel for useful comments on the manuscript.  MM is supported by the NASA Einstein fellowship."
    },
    "1004/1004.4881_arXiv.txt": {
        "abstract": "In a density-stratified turbulent medium the cross helicity $\\langle \\mathrm{\\bf u}'\\cdot \\mathrm{\\bf B}'\\rangle$ is considered as a result of the interaction of the velocity fluctuations and a large-scale magnetic field. By means of a quasilinear theory and by numerical simulations we find the cross helicity and the mean vertical magnetic field anti-correlated. In the high-conductivity limit the ratio of the helicity and the mean magnetic field equals the ratio of the magnetic eddy diffusivity and the (known) density scale height. The result can be used to predict that the cross helicity at the solar surface exceeds the value of $1$ gauss\\,km\\,s$^{-1}$. Its sign is anti-correlated with that of the radial mean magnetic field. Alternatively, we can use our result to determine the value of the turbulent magnetic diffusivity from observations of the cross helicity. ",
        "introduction": "Dynamo theory for convective zones needs to know both the values of the $\\alpha$-effect and the eddy diffusivity. The $\\alpha$-effect is strongly related to the kinetic helicity which has opposite signs in the two hemispheres. Almost all of the theoretical calculations for rotating stratified turbulence lead to negative helicity ({\\it i.e.} positive $\\alpha$-effect) for the northern hemisphere and positive helicity ({\\it i.e.} negative $\\alpha$-effect) for the southern hemisphere. Despite all of the complications to measure the helicity on the solar surface, a new result has recently been presented by Komm, Hill, and Howe (2008). They do indeed find negative (positive) values for the kinetic helicity in the northern (southern) hemisphere. This result is based on a ring-diagram analysis of GONG data. The values remained constant as long as the magnetic field did not exceed 10\\,gauss (G). Using the observation of the DIV-CURL correlation, ${\\cal C}=\\langle (u_{,x}+v_{,y})(v_{,x}-u_{,y})\\rangle$, which is proportional to the kinetic helicity (R\\\"udiger, Brandenburg, and Pipin, 1999; see R\\\"udiger and Hollerbach (2004) for more details), Duvall and Gizon (2000) found $\\cal C$ to be negative (positive) in the northern (southern) hemisphere, as derived from the horizontal velocity components of mesogranulation patterns. Egorov, R\\\"udiger, and Ziegler (2004) simulated these observations with the NIRVANA code and reproduced them with Taylor numbers as small as $10^3$.  By use of the finding of Keinigs (1983) the current helicity and the $\\alpha$-effect are anti-correlated, so one also can derive the sign of the $\\alpha$-effect by observation of the current helicity $\\langle \\mathrm{\\bf J}'\\cdot  \\mathrm{\\bf B}'\\rangle$. Seehafer (1990) started to observe the current helicity at the solar surface showing that it is negative (positive) in the northern (southern) hemisphere. Again the $\\alpha$-effect is found to be positive (negative) in the northern (southern) hemisphere.  The numerical value of the helicity derived by Komm {\\it et al.} (2008) is of the order of $10^{-7}$ cm\\,s$^{-2}$. This turns out to be very small, because the resulting $\\alpha$-effect is below 1 cm\\,s$^{-1}$. By comparison, K\\\"apyl\\\"a, Korpi, and Brandenburg (2009) find $\\alpha \\simeq 0.03 u_{\\rm rms}$ near the surface from their convection simulations. With $u_{\\rm rms}\\simeq 300$ m\\,s$^{-1}$, this corresponds to 10 m\\,s$^{-1}$. The maximal $\\alpha$-value in their box center is of the order of 0.3$u_{\\rm rms}$. This highlights a major discrepancy between theory and observations or, at least, a difficulty in determining $\\alpha$ from observations.  The empirical definition of the turbulent magnetic diffusion seems to be more straightforward. The decay of non-permanent magnetic structures such as sunspots or larger active regions lead to numerical values of the turbulent magnetic diffusivity. One finds $\\eta_{\\rm T}\\simeq 10^{11}$ cm$^2$\\,s$^{-1}$ from sunspot decay (Stix, 1989) or $\\eta_{\\rm T}\\simeq 10^{12}$ cm$^2$\\,s$^{-1}$ from the decay of active regions (Schrijver and Martin, 1990). These values are smaller than the value of $3\\times 10^{12}$ cm$^2$\\,s$^{-1}$, which results from the widely used formula $\\eta_{\\rm T}\\approx u_{\\rm rms} \\ell /3$ with correlation or mixing length $\\ell$ and parameter values taken close to the surface. There is no possibility until now to observe the turbulent diffusivity on the solar surface for the quiet Sun where the magnetic quenching of this quantity by magnetic fields is negligible. We shall demonstrate in the present paper that there is a rather simple possibility to observe the magnetic diffusivity even in the presence of very weak magnetic fields ($< 10$\\,G), for which quenching should be negligible.  We shall show that in the presence of a mean magnetic field along the direction of density stratification, hydromagnetic turbulence will attain cross helicity, $\\langle \\mathrm{\\bf u}'\\cdot \\mathrm{\\bf B}'\\rangle$, whose value is proportional to the turbulent magnetic diffusivity. Indeed, our work is an extension of that by Kleeorin {\\it et al.} (2003), who considered the effect of stratification of turbulent intensity. More recently, Kuzanyan, Pipin, and Zhang (2007) emphasized the importance of cross-helicity for estimating internal solar parameters important for the dynamo. We mention in this connection that cross helicity itself may constitute a potentially important dynamo effect \\cite{Yos90,Yok96}. In this paper we propose a measurement of cross-helicity in the Sun for estimating the turbulent magnetic diffusivity in quiet regions. We argue that this is more accurate than measuring, for example, the mean electromotive force. ",
        "conclusions": "We  have shown that nonrotating turbulence at the top of the solar convection zone under the influence of a vertical magnetic field yields a finite cross helicity. The only requirement is the existence of density stratification which enters the induction equation via the anelastic relation $\\mathrm{div}(\\rho\\mathrm{\\bf u}') = 0$. The Boussinesq approximation cannot be used. The effect exists mainly in the high-conductivity limit, {\\it i.e.} for sufficiently large magnetic Reynolds numbers (see Figure \\ref{pRm_dep}). The radial magnetic field, on the other hand, must be weak enough to remain passive so that it does not dominate the flow. Figure \\ref{pB_dep} shows that the maximum field is given by $B_{\\rm eq}$ which is {\\em much} higher than the mean vertical field, which is of the order of a few gauss on the Sun.  To estimate the value of the cross helicity at the solar surface we shall assume a density scale height of 100 km. Then one finds from Equation (\\ref{12}) that \\begin{equation} {\\langle u'_r B'_r\\rangle}  \\simeq -  \\frac{\\langle B_r\\rangle}{1 {\\rm G}}\\frac{\\eta_{12}}{H_7} \\;\\mbox{G km s}^{-1}. \\label{15} \\end{equation} The average is here to be applied over many turbulent cells which, in the Sun, might correspond to 30--100 Mm. The magnetic diffusivity in Equation (\\ref{15}) has been used in the form $\\eta_{\\rm T}= 10^{12} \\eta_{12}$\\,cm$^2$\\,s$^{-1}$ and the density scale height as $H_\\rho= 100 H_7$ km. We thus predict the existence of a cross helicity of more than 1 G\\,km\\,s$^{-1}$. We also emphasize that the cross helicity is anti-correlated to the mean radial magnetic field, {\\it i.e.} \\begin{equation} \\langle u'_r B'_r\\rangle \\   \\langle B_r \\rangle < 0. \\label{16} \\end{equation} For a dipolar background field the sign of the cross helicity will be opposite in the two hemispheres.  Relation (\\ref{12}) can also be used to measure the magnetic diffusivity if the cross helicity is known by observations. In order to find the cross helicity one only has to correlate the observed flow fluctuations with observed magnetic fluctuations. Together with the calculated mean value of the radial magnetic field, Equation (\\ref{12}) provides the unknown quantity $\\eta_{\\rm T}$. We hope that such an analysis of the observations using, for example, data from the Hinode satellite will soon provide supporting evidence for an anti-correlation between $\\langle u'_r B'_r\\rangle$ and $\\langle B_r \\rangle$, and that a meaningful value of $\\eta_{\\rm T}$ can be obtained in that way.      \\begin{acks} The authors are thankful to Hongqi Zhang for attracting attention to the cross-helicity problem. This work was supported by the Deutsche Fo\\-r\\-schungs\\-ge\\-mein\\-schaft and by the Russian Foundation for Basic Research (projects 10-02-00148, 10-02-00391). We acknowledge the allocation of computing resources provided by the Swedish National Allocations Committee at the Center for Parallel Computers at the Royal Institute of Technology in Stockholm and the National Supercomputer Centers in Link\\\"oping. This work was supported in part by the European Research Council under the AstroDyn Research Project No.\\ 227952 and the Swedish Research Council Grant No.\\ 621-2007-4064. \\end{acks}"
    },
    "1004/1004.0146_arXiv.txt": {
        "abstract": "The radio-quiet quasar SDSS J1536+0441A shows two broad-line emission systems, recently interpreted as a binary black hole (BBH) system with a subparsec separation; as a double-peaked emitter; or as both types of systems.  The NRAO VLBA was used to search for 8.4 GHz emission from SDSS J1536+0441A, focusing on the optical localization region for the broad-line emission, of area 5400 mas$^2$ (0.15 kpc$^2$).  One source was detected, with a diameter of less than 1.63 mas (8.5 pc) and a brightness temperature $T_b > 1.2 \\times 10^7$ K.  New NRAO VLA photometry at 22.5 GHz, and earlier photometry at 8.5 GHz, gives a rising spectral slope of $\\alpha = 0.35\\pm0.08$.  The slope implies an optically thick synchrotron source, with a radius of about 0.04 pc, and thus $T_b \\sim 5 \\times 10^{10}$ K.  The implied radio-sphere at rest frame 31.2~GHz has a radius of 800 gravitational radii, just below the size of the broad line region in this object.  Observations at higher frequencies can probe whether or not the radio-sphere is as compact as expected from the coronal framework for the radio emission of radio-quiet quasars. ",
        "introduction": "\\label{motivation} Binary black hole (BBH) systems with subparsec scales are predicted in merging scenarios for galaxy evolution and also factor prominently in predictions for the gravitational wave background \\citep[e.g.,][]{col09}.  But can such BBH systems be found?  A promising search method - identifying candidate BBH systems through their optical broad-line emission properties - recently yielded three such candidates: SDSS J153636.22+044127.0 \\citep{bor09}, SDSS J092712.65+294344.0 \\citep{bog09,dot09}, and SDSS J105041.35+345631.3 \\citep{shi09}.  In this Letter we focus on SDSS J153636.22+044127.0 (SDSS J1536+0441 hereafter), a quasar at a redshift $z = 0.388$ \\citep{bor09}.  For the assumed flat cosmology\\footnote{$H_0 = 71$~km~s$^{-1}$~Mpc$^{-1}$ and $\\Omega_m = 0.27$, implying a luminosity distance of 2.1 Gpc, an angular size distance of 1.1 Gpc and a scale of 5.2 pc per mas.}, the inferred 0.1 pc separation subtends 0.02 mas \\citep{bor09}.  This quasar has also been interpreted as a lone double-peaked emitter \\citep[DPE;][]{gas10,cho10} or as twin DPEs with a subparsec separation \\citep{tan09}.  The emission lines from a DPE are thought to arise from rotational motion in a relativistic accretion disk. SDSS J1536+0441 is radio quiet and our imaging of it at 8.5 GHz \\citep{wro09} revealed two faint sources, SDSS J1536+0441A and J1536+0441B, separated by 0.97\\arcsec\\, (5.1 kpc) with each source being unresolved with a diameter of less than 0.37\\arcsec\\, (1.9 kpc). It is now clear that SDSS J1536+0441A is hosted by the radio-quiet quasar (RQQ), while SDSS J1536+0441B is hosted by a companion radio-loud elliptical galaxy \\citep{dec09,lau09}.  In the BBH scenarios for SDSS J1536+0441A \\citep{bor09,tan09}, each of the two broad-line emission systems is itself a quasar.  Recent optical spectroscopy finds that the broad-line emission systems have relative positions that localize them to 3 $\\sigma$ bands of full width 90 mas (470 pc) along a position angle (PA) of 48\\arcdeg\\, \\citep{cho10} and 60 mas (310 pc) along PA 90\\arcdeg\\, (T. Boroson, 2010, private communication).  The resulting parallelogram localizes SDSS J1536+0441A to an area of 5400 mas$^2$ (0.15 kpc$^2$). Improvements in the localization of the emission from SDSS J1536+0441A would further test its candidacy as a BBH system.  In this Letter we use measurement techniques at radio frequencies to investigate the compactness of SDSS J1536+0441A, first by seeking evidence for a synchrotron self-absorbed spectrum and then by direct imaging with mas resolution to localize the emission on parsec scales.  Our new imaging is presented in \\S~\\ref{imaging} and its implications are explored in \\S~\\ref{implications}.  A summary appears in \\S~\\ref{summary}. ",
        "conclusions": "\\label{summary}  Our VLBA search for 8.4 GHz emission from the RQQ SDSS J1536+0441A found only one source within the localization region, of area 0.15 kpc$^2$, for the broad-line emission.  The VLBA detection has a diameter of less than 8.5 pc and a $T_b > 1.2 \\times 10^7$ K.  This detection of SDSS J1536+0441A represents the first parsec-scale localization of a candidate BBH system identified through its broad-line properties.  The VLA photometry at a rest frequency of 31.2 GHz yields a rising spectrum, consistent with synchrotron self-absorption, which implies a radius of about 0.04 pc for the radio-sphere, and $T_b \\sim 5 \\times 10^{10}$ K.  The observed compact flat spectrum radio sphere is consistent with the trait predicted in the coronal framework for RQQs.  The radio-sphere at 31.2 GHz happens to be just inside the estimated inner boundary for the BLR, of 800 gravitational radii, in this object.  It would be useful to investigate the spectrum at higher frequencies, corresponding to the mm and sub-mm range, to determine where the spectral slope steepens as the source becomes optically thin.  This will establish the size of the most compact synchrotron-emitting region, and potentially allow a direct exploration of relativistic electrons in the accretion disk corona.  We plan such investigations using the Expanded VLA \\citep{per09} and the Atacama Large Millimeter/submillimeter Array \\citep{woo09}.  Concerning the BBH scenarios for SDSS J1536+0441A, the VLBA detection is consistent with a quasar separation of less than 8.5 pc.  No additional sources with $T_b = 4.5 \\times 10^6$ K or more are found in within the localization region for the broad-line emission.  Concerning the lone DPE scenario for SDSS J1536+0441A, its emission line profiles do make it an unusual DPE.  But as no other radio-quiet DPEs have been imaged on parsec scales, it is impossible to say whether or not the VLBA detection of SDSS J1536+0441A is in any way unusual."
    },
    "1004/1004.5139_arXiv.txt": {
        "abstract": "The issue of giant planet formation by core instability (CI) far from the central star is rather controversial because the growth of massive solid core necessary for triggering the CI can take longer than the lifetime of the protoplanetary disk. In this work we assess the range of separations at which the CI may operate by (1) allowing for arbitrary (physically meaningful) rate of planetesimal accretion by the core and (2) properly taking into account the dependence of the critical mass for the CI on the planetesimal accretion luminosity. This self-consistent approach distinguishes our work from similar studies in which only a specific planetesimal accretion regime was explored and/or the critical core mass was fixed at some arbitrary level. We demonstrate that the largest separation at which the CI can occur within $3$ Myr corresponds to the surface density of solids in the disk $\\gtrsim 0.1$ g cm$^{-2}$ and is $40-50$ AU in the minimum mass Solar nebula. This limiting separation is achieved when the planetesimal accretion proceeds at the fastest possible rate, even though the high associated accretion luminosity increases the critical core mass delaying the onset of the CI. Our constraints are independent of the mass of the central star and vary only weakly with the core density and its atmospheric opacity. We also discuss various factors which can strengthen or weaken our limits on the operation of the CI. ",
        "introduction": "\\label{sect:intro}  Recent discoveries of distant planetary companions (separations of at least tens of AU) around several nearby stars by direct imaging (Marois \\etal 2008; Thalmann \\etal 2009) have stimulated investigation of their formation mechanisms. Two contending theories of the planet formation --- the core instability (Perri \\& Cameron 1974; Harris 1978; Mizuno 1980; hereafter CI) and the gravitational instability (Cameron 1978; Boss 1998) --- differ quite dramatically in their ability to form planets at various distances from the star.  It is generally thought that CI should be capable of forming giant planets close to the star, at separations $\\lesssim 10$ AU. Both theoretical modelling (Mizuno 1980; Stevenson 1982; Pollack \\etal 1996) and the existence of two gas giants at semi-major axes of $5.2$ AU and $9.5$ AU in our own Solar System (coupled with some knowledge about the evolution of the planetary orbital architecture in the early Solar System) attest to this statement. At the same time it seems unlikely that giant planets can be formed by the gravitational instability in the inner parts of protoplanetary disks because of the long local cooling time of gas (Matzner \\& Levin 2005; Rafikov 2005, 2007). On the contrary, giant planet formation by the direct gravitational instability appears feasible at large distances from the star (Boley 2009; Rafikov 2009; Clarke 2009; but see Boss (2006) for an opposite view) where the disk cooling time is short enough to permit fragmentation of the gravitationally unstable disk (Gammie 2001). Whether the CI can operate beyond several tens of AU from the star is not so clear.  The well known problem faced by the CI far from the star is that the buildup of massive refractory core (necessary for triggering the vigorous gas accretion) is thought to take very long time, longer than the $1-10$ Myr lifetime of the gaseous component of protoplanetary nebula. Studies illustrating this problem are usually based on two key assumptions (Dodson-Robinson \\etal 2009): first, that the CI is triggered whenever a $M_{crit}\\sim 10$ M$_\\oplus$ core is built by planetesimal accretion, and, second, that the planetesimals accrete onto the growing core at rather modest rates (Ida \\& Lin 2004; we will explain later what do we mean by that). When these conditions are imposed it is generally found that the CI does not commence at $a\\gtrsim 20-30$ AU prior to the nebular gas removal.  In this note we show that aforementioned assumptions are too restrictive and neither of them need to be adopted in determining the feasibility of the CI. First, planetesimal accretion by the growing core can (at least potentially) be much faster than what has been previously assumed. Second, the critical core mass itself strongly depends on the planetesimal accretion rate. Based on these observations we formulate in this work a novel constraint on the operation of the CI in protoplanetary disks. ",
        "conclusions": "\\label{sect:disc}  Despite the robustness of our arguments it is not inconceivable that some additional factors can weaken them and make giant planet formation by the CI possible even beyond the limits represented by equations (\\ref{eq:Sig_lim}) and (\\ref{eq:a_lim}). Alternatively, it is quite possible that some of the assumptions used in deriving these results are too extreme and one can get even better constraints by focussing on less dramatic assumptions. Below we review factors that can work one way or another.   \\subsection{Extending CI to larger radii.} \\label{subsect:weaken}  One possible way to  facilitate CI and increase $a_{lim}$ is to consider possibility of planetesimal accretion at rates {\\it exceeding} $\\dot M_{max}$. This is very difficult (since there are many factors that tend to reduce $\\dot M$ compared to $\\dot M_{max}$, see \\S\\ref{subsect:strengthen}) but may be possible if e.g. one takes into account the increase of planetesimal capture cross-section by the core caused by its extended, dense atmosphere. This effect has been previously investigated by Inaba \\& Ikoma (2003) who demonstrated that an increase of $\\dot M$ by a factor of $\\sim 10$ compared to the value computed without atmosphere is possible. According to equation (\\ref{eq:a_lim}) such an enhancement of $\\dot M$ (incorporated by increasing $\\chi$) would boost $a_{lim}^{MMSN}$ by a factor of $\\sim 4$.  Our present calculations assume that the core is accreting planetesimals continuously until the protoplanetary nebula dissipates --- this is important at large $a$ since massive core requires long time to be built. But one may wonder if building smaller core in shorter time and then cutting off subsequent planetesimal accretion (and energy release at the core surface, which supports atmosphere against going unstable) completely may still lead to the CI and potentially extend it to larger semi-major axes. Such accretion scenario has been adopted by e.g. Pollack \\etal (1996). The problem in this case is that even if $\\dot M=0$ it still takes long time for the atmosphere around the core to grow to the mass comparable to $M_c$. INE0 show that this process occurs on thermal timescale of the atmosphere and typically takes millions of years.  Similar problem is also encountered in a scenario where the core grows rapidly by planetesimal accretion in the inner regions of protoplanetary disk and then gets scattered out to large radii by some massive perturber. One might expect that after the orbit of the scattered core circularizes by dynamical friction the core would gradually accrete massive atmosphere and undergo CI at some point. Given that both the orbit circularization and envelope accretion are likely to take long time it is not at all obvious whether the CI could be achieved in this scenario within several Myrs.   \\subsection{Limiting CI to smaller radii.} \\label{subsect:strengthen}   There are many factors that can potentially reduce $a_{lim}$ compared to 44 AU estimated in equation (\\ref{eq:a_lim}). For example, there are several reasons why it may not be possible for $\\dot M$ to reach the maximum rate $\\dot M_{max}$.  First, the growing core may clear out a gap in planetesimal disk around its orbit, thus significantly reducing $\\dot M$ (Tanaka \\& Ida 1997; Rafikov 2001; 2003a). In our previous calculations we implicitly assumed this not to happen e.g. because of the core migration through the disk, which allows fresh planetesimal material to be constantly supplied for core accretion (Alibert \\etal 2005).  Second, as we mentioned in \\S\\ref{sect:accr}, a known pathway to $\\dot M\\approx\\dot M_{max}$ is via the growth of the core to the size at which it starts dominating dynamical evolution of nearby planetesimals and triggers their efficient collisional fragmentation (Rafikov 2004). However, there is a strong implicit assumption in this scenario --- that the core can reach this critical size within the nebula lifetime. Rafikov (2003b) has shown that at $a\\sim 30-40$ AU a dynamically dominant core would need to have mass of order $10^{24}$ g and would require on the order of $10-100$ Myr to grow in the MMSN. This time scale is apparently in conflict with the typical dissipation times of protoplanetary disks. Thus, one may need to either require a more massive disk at these radii or to find other pathways for accretion at $\\dot M_{max}$. Formation of massive solid bodies by direct gravitational instability facilitated by various streaming instabilities (Johansen \\etal 2009) may be quite relevant for the latter option. Regardless of this (arguably rather serious) issue our estimate (\\ref{eq:a_lim}) still remains a useful upper limit on $a_{lim}$.  Among other factors hindering the onset of CI and reducing $a_{lim}$ we should mention the possibility of high opacity in the protoplanetary atmosphere. It was suggested in \\S\\ref{subsect:opacity} that $\\kappa$ may be very low because of the dust sedimentation and growth. However, infalling planetesimals which feed core accretion likely get partly disrupted in the atmosphere leaving behind large amount of refractory material. This may actually {\\it increase} $\\kappa$ compared to the value of $0.1$ cm$^2$ g$^{-1}$ assumed in equation (\\ref{eq:a_lim}). Nevertheless, given the weak sensitivity of $a_{lim}$ to $\\kappa$ the potential increase of atmospheric opacity is unlikely to have huge effect on $a_{lim}$."
    },
    "1004/1004.0236_arXiv.txt": {
        "abstract": "We compare current and forecasted constraints on dynamical dark energy models from Type Ia supernovae and the cosmic microwave background using figures of merit based on the volume of the allowed dark energy parameter space.  For a two-parameter dark energy equation of state that varies linearly with the scale factor, and assuming a flat universe, the area of the error ellipse can be reduced by a factor of $\\sim 10$ relative to current constraints by future space-based supernova data and CMB measurements from the Planck satellite. If the dark energy equation of state is described by a more general basis of principal components, the expected improvement in volume-based figures of merit is much greater.  While the forecasted precision for any single parameter is only a factor of 2--5 smaller than current uncertainties, the constraints on dark energy models bounded by $-1\\leq w\\leq 1$ improve for approximately 6 independent dark energy parameters resulting in a reduction of the total allowed volume of principal component parameter space by a factor of $\\sim 100$.  Typical quintessence models can be adequately described by just 2--3 of these parameters even given the precision of future data, leading to a more modest but still significant improvement.  In addition to advances in supernova and CMB data, percent-level measurement of absolute distance and/or the expansion rate is required to ensure that dark energy constraints remain robust to variations in spatial curvature. ",
        "introduction": "\\label{sec:intro}   In the absence of physically compelling models for dark energy, figures of merit (FoMs) are a useful tool that encapsulate the constraining power of cosmological data.  FoMs combine various constraints on the expansion history of the universe (e.g.\\ the distance-redshift relation) in a single number, or at most a handful of numbers, which serve as simple and quantifiable metrics by which to evaluate current and proposed experiments \\cite{DETF,Albrecht_Bernstein,FoMSWG}.  The simplest schemes adopt fixed functional forms for the evolution of the dark energy equation of state and define the FoM as the inverse of the allowed parameter volume \\cite{Huterer_Turner}.  One widely used version is the two dimensional $w_0$--$w_a$ parametrization \\cite{DETF} but other higher dimensional versions have also been considered \\cite{Huterer_Turner,Wang_FoM,Crittenden_Pogosian_Zhao}.  Any simple parametrization of the expansion history risks biasing the FoM in favor of or against certain types of data by choosing a fixed functional form \\cite{Albrecht_Bernstein}. To avoid this problem, one can use more complicated schemes that parametrize the whole functional freedom in the dark energy equation of state evolution and separate the expansion history and growth of structure information.  For example, uncorrelated modes of piecewise-constant discretizations of the equation of state that are local in redshift \\cite{Hu_PC,Huterer_Cooray,Wang_Tegmark_2005,Shapiro_Turner,Dick,Suletal07,Zhao_Huterer_Zhang,Zhao_Zhang:2009,Serra:2009} or constructed from principal components (PCs) \\cite{Hu_PC,Huterer_Starkman,PaperI,PaperII,Huterer_Peiris,Tang:2008hm,Kitching:2009yr} have been employed to characterize both current and future data.  In particular, the inverse parameter volume of the PC amplitudes, defined separately for each experiment, has been advocated as a FoM \\cite{FoMSWG}.  Besides avoiding biasing results towards a particular functional form, a more model-independent FoM has the advantage of being able to identify improvements in dark energy constraints that might be missed by FoMs with fewer parameters. On the other hand, not all improvements in a multidimensional PC FoM reflect improvements in constraining the space of known dark energy models \\cite{dePutter_Linder2,Barnard:2008mn}. For example, stronger constraints may exclude regions of the parameter space not occupied by typical models.   In this paper, we study the FoMs defined both with the commonly-used $w_0$--$w_a$ parametrization and with PCs based on forecasts for Type Ia supernova (SN) and cosmic microwave background (CMB) data. Previous studies of PC-based FoMs have generally relied on the Fisher matrix approximation, and the implementation and utility of these FoMs for real data has not been addressed. Here we define straightforward generalizations of PC FoMs and apply them to both actual data from current measurements and forecasts for future data. While we still construct the PCs using the Fisher matrix approach, we then treat the PC amplitudes as free parameters and constrain them using full Markov Chain Monte Carlo (MCMC) likelihood explorations with current or future data. This methodology follows that employed in previous papers where we studied generalized predictions of classes of dark energy models based on forecasts \\cite{PaperI} and current data \\cite{PaperII}.  The new element here is the application of these methods to the study of dark energy FoMs.   The standard approach for constructing the principal component-based FoM is to use the PCs specific to the experiment and the cosmological probe(s) considered; see e.g.~Ref.~\\cite{FoMSWG}. However, this approach makes it difficult to directly compare PC-based FoMs for different experiments and to assess improvements in specific regions of parameter space.  To facilitate such direct comparison between FoMs for current and future data, we choose instead to fix the set of PCs based on a specific projection of future data rather than computing separate PCs for future and current data.   This paper is organized as follows. In Sec.~\\ref{sec:methods}, we describe the current and forecasted SN and CMB data sets as well as additional priors from baryon acoustic oscillations (BAO) and measurements of the Hubble constant. We also briefly review the MCMC analysis methods used to infer dark energy constraints. In Sec.~\\ref{sec:results}, we compare FoMs from the current and future SN and CMB constraints for both the commonly-used $w_0$--$w_a$ model (Sec.~\\ref{sec:fomw0wa}) and a more general PC-based parametrization (Sec.~\\ref{sec:fompc}). We summarize and discuss these results in Sec.~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}   We have compared current and forecasted figures of merit (FoMs) for dark energy using both a simple $w_0$--$w_a$ description of the equation of state, and a more complicated but more complete principal component (PC) parametrization.  By consistently using a fixed set of PCs based on future data, and by generalizing the FoM definition to include possible parameter covariance, we showed how PC FoMs can be applied to likelihood analyses of both existing data sets and forecasts.  We have also shown that the covariance-based FoMs accurately represent relative changes in the parameter volume, even in the presence of strongly non-Gaussian posterior distributions such as those caused by imposing top hat priors.  Traditional variance-based PC FoMs are consistent with those that include the full covariance for the cases we have tested, but we do not expect such agreement to hold in general. For example, if our PC eigenfunctions from SN and CMB data are applied to qualitatively different data, e.g.\\ weak lensing or BAO, the FoMs that account for the parameter covariances should be employed.   For the $w_0$--$w_a$ FoM, future space-based supernova data and CMB measurements from the Planck satellite can improve on current measurements by a factor of $\\sim 10$.  For the PC FoM, the expected improvement is much greater still.  While the forecasted precision for any single principal component is only a factor of 2--5 smaller than current uncertainties, the constraints on general quintessence models bounded by $-1\\leq w\\leq 1$ improve for approximately 6 components resulting in a reduction of the total allowed volume of dark energy parameter space by a factor of $\\sim 60$--$140$. The FoM improvement depends mainly on whether or not variations in spatial curvature are allowed, with the maximum ratio of $\\sim 140$ requiring either fixing the curvature with a theoretical prior ($\\ok=0$) or combining future SN and CMB data with a measurement of absolute distance and/or the expansion rate with at least $\\sim 1\\%$ accuracy.  Although improvements from future SN and CMB data extend to 6 components of $w(z)$, many commonly-considered quintessence models are adequately described by the first three PCs for which the allowed volume only decreases by a factor of $\\sim 20$--$40$ relative to current measurements. While figures of merit provide a convenient model-independent metric to evaluate experiments, they must be interpreted with caution when science goals include discriminating amongst specific dark energy models.       \\vspace{1cm} {\\it Acknowledgments:} MJM was supported by CCAPP at Ohio State; DH by the DOE OJI grant under contract DE-FG02-95ER40899, NSF under contract AST-0807564, and NASA under contract NNX09AC89G; WH by the KICP under NSF contract PHY-0114422, DOE contract DE-FG02-90ER-40560 and the Packard Foundation.    \\vfill"
    },
    "1004/1004.4003_arXiv.txt": {
        "abstract": "We study resolution effects in numerical simulations of gas-rich and gas-poor major mergers, and show that the formation of slowly-rotating elliptical galaxies often requires a resolution that is beyond the present-day standards to be properly modelled. Our sample of equal-mass merger models encompasses various masses and spatial resolutions, ranging from about 200~pc and $10^5$ particles per component (stars, gas and dark matter), i.e. a gas mass resolution of $\\sim 10^5$~M$_{\\sun}$, typical of some recently published major merger simulations, to up to 32~pc and $\\sim 10^3$~M$_{\\sun}$ in simulations using $2.4 \\times 10^7$ collisionless particles and $1.2 \\times 10^7$ gas particles, among the highest resolutions reached so far for gas-rich major merger of massive disc galaxies. We find that the formation of fast-rotating early-type galaxies, that are flattened by a significant residual rotation, is overall correctly reproduced at all such resolutions. However, the formation of slow-rotating early-type galaxies, which have a low residual angular momentum and are supported mostly by anisotropic velocity dispersions, is strongly resolution-dependent. The evacuation of angular momentum from the main stellar body is largely missed at standard resolution, and systems that should be slow rotators are then found to be fast rotators. The effect is most important for gas-rich mergers, but is also witnessed in mergers with an absent or modest gas component (0--10\\%\\ in mass). The effect is robust with respect to our initial conditions and interaction orbits, and originates in the physical treatment of the relaxation process during the coalescence of the galaxies. Our findings show that a high-enough resolution is required to accurately model the global properties of merger remnants and the evolution of their angular momentum. The role of gas-rich mergers of spiral galaxies in the formation of slow-rotating ellipticals may therefore have been underestimated. Moreover, the effect of gas in a galaxy merger is not limited to helping the survival/rebuilding of rotating disc components: at high resolution, gas actively participates in the relaxation process and the formation of slowly-rotating stellar systems. ",
        "introduction": "Numerical simulations have been intensively used for more than two decades to study the properties of the remnants of galaxy mergers and the role of hierarchical merging in the formation of elliptical-like early-type galaxies \\citep{HB91,Barnes92, mihos95}. With the increasing resolution and large statistical samples \\citep[e.g.,][]{naab03, bour05, dimatteo1,dimatteo2}, modern work tends to quantify in details the properties of major and minor merger remnants, and accurate comparisons with observed properties of early-type galaxies can now be envisioned \\citep[e.g.,][]{burkert-sauron}.  A general concern, though, is that the impact of the spatial and the mass resolutions on the detailed properties of the systems under scrutiny remains largely overlooked, and whether or not simulations of mergers have converged with today's typical resolution remains unexplored. Obviously, increasing resolution enables simulations to directly resolve cold gas clouds and clustered star formation \\citep[e.g.][]{B08,kim09}, but whether these additional small-scale ingredients can significantly impact the global, large-scale morphology and kinematics of merger remnants has not been studied in detail. In cosmological simulations, an increase in resolution (i.e., an increase in the number of particles and/or decrease of the softening length) can affect the baryonic density and circular velocity profiles of individual galaxies in a halo (\\citealt{naabresol}). \\citet{navresol} also studied numerical convergence via a suite of $\\Lambda$CDM simulations and confirmed that the halo mass distributions were better described by Einasto profiles that are not, stricly speaking, universal.  While many resolution studies have been made in cosmological simulations, few have focused on galaxy merger simulations. \\citet{coxresol}, \\citet{hopkinsresol} and e.g., \\citet{dimatteo2} included some checks of the effect of resolution on the star-formation activity of on-going mergers. But a resolution study aimed at examining the detailed morphology and kinematics of relaxed merger remnants (i.e., galaxies which tend to be roughly S0- or elliptical-like) has not yet been conducted.  Models of galaxy mergers have reached particularly high resolution with the work of \\citet{wetzstein07} (70~pc softening with $4\\times 10^6$ particles in total -- but only 45,000 for the gas component), \\citet{li04} (10 to 100~pc and $5 \\times 10^5$ gas particles per galaxy), \\citet{naabresol} ($8\\times 10^6$ particles with a 125~pc resolution in a cosmological re-simulation of an individual galaxy halo). The highest resolution for gas-rich mergers have been achieved recently by \\citet{B08} for mergers of bright spiral galaxies, with a total of $3.6 \\times 10^7$ particles including more than $10^7$ gas particle, and a 32~pc softening size, and \\citet{kim09} with a spatial resolution of 3.8~pc and a mass resolution of $2\\times 10^3$~M$_{\\sun}$ (for dwarf or low-mass spirals, though). But such high-resolution studies have focused on small-scale gas physics and structure formation, without studying the impact of high resolution on the global properties of the elliptical-like galaxies formed in major mergers.  Large samples of simulations of idealized galaxy mergers remain typically limited to softening lengths of about 100--300~pc, and $\\sim 10^5$ particles per galaxy \\citep[see samples in][]{naab03,NJB06,bour04,bour07,coxal06,coxal08,dimatteo1}. Whether or not the relatively limited numerical resolution used in such studies affects the global properties of merger remnants is still a largely open question: for instance, the detailed comparison of major merger remnants with the observed anisotropy-flattening relation by \\citet{burkert-sauron} relies on simulations with $2 \\times 10^4$ gas particles per galaxy, a gas particle mass $\\sim 3 \\times 10^5$~M$_{\\sun}$, and a spatial resolution (softening) of about 200~pc.  Within the context of the \\atlas{} project (\\href{http://purl.org/atlas3d}{http://purl.org/atlas3d}), an extensive set of numerical simulations is being conducted to support the multi-wavelength survey of a complete sample of early-type galaxies within the local (40Mpc) volume, in terms of various formation mechanisms of early-type galaxies: binary mergers, multiple mergers, disc instabilities, etc. An ambitious series of simulations of mergers are being specifically performed and analysed for this purpose (Bois et al. in preparation). To properly interpret the results from these simulation efforts, as well as to understand the robustness of existing and past studies of galaxy mergers, we first probe the effect of spatial and mass resolutions on the global structure of binary disc merger remnants.  In this paper, we study the effect of numerical resolution on the global morphology and the kinematics of the simulated remnants of binary, equal-mass major mergers. We wish to examine resolutions ranging from the typical resolutions used in recent, large simulations samples, to some of the highest merger simulations ever performed. We study both dry (collisionless) and wet (gas-rich) mergers of disc galaxies. The modeled interaction orbits lead to the formation of both fast rotators, i.e., early-type galaxies flattened by  significant rotational support, and slow rotators, i.e., early-type galaxies with low residual rotation, supported (and flattened) by (anisotropic) velocity dispersions, following the classification detailed in \\citet[][see also Section~\\ref{sec:nbody}]{ems07}. We find that the formation of fast rotators is overall correctly reproduced with numerical simulations at modest resolutions. In contrast, the formation of slow-rotating systems is correctly reproduced only at high resolution (Section~\\ref{sec:results}), above the resolution of most of the recently published merger simulations. The influence of gas on the structure of merger remnants, compared to dry mergers, also differs at high resolution, and is not limited to easing the survival and/or rebuilding of rotating disc components. In Section~\\ref{sec:origins}, we further examine the origin of this observed resolution effect in the formation of slow-rotating systems. We show that it is not an artefact from different initial conditions or interaction orbits, but that the physical treatment of the merging process is actually biased when the resolution is too low. The effect of the resolution has been tested on other simulations producing slow rotators and we find that it is a systematic one (Section~\\ref{sec:systematic}). We summarize our results, discuss the required resolution for accurate studies and the general implications for the formation of elliptical galaxies in Section~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} In this paper, we have studied the effect of numerical resolutions (spatial and mass resolution) on the global properties of merger remnants. Our simulations at ``standard'' resolution, are comparable to the majority of merger simulation samples published in the last few years: the spatial resolution (gravitational softening and typical hydrodynamical smoothing lengths) is 180~pc, and the number of particles $\\sim$ 10$^{5}$ per galaxy and per component (gas, stars, and dark halo). These simulations have been compared to models of the same mergers with increased resolution, up to 32~pc and almost 10$^{7}$ particle per galaxy and per component.  We have analyzed the morphology and kinematics of the relaxed merger remnants. In particular, we have studied whether they are ``fast rotators'', with an apparent spin parameter $\\lambda_R > 0.1$ and have small misalignments between the morphological and kinematic axes, i.e. in broad terms early-type galaxies with significant flattening and rotational support. At the opposite end, ``slow rotators'' are systems with a low $\\lambda_R \\leq 0.1$ (at one effective radius), large kinematic misalignments, i.e. early-type galaxies dominated by (anisotropic) pressure support and low residual rotation. Such slow rotators usually have central KDCs in our high-enough resolution simulations.  Our mains findings can be summarized as follows: \\begin{itemize} \\item The formation of \\fasts{} is not significantly affected by numerical resolution. Models that produce fast rotators at the highest resolution also result in fast rotators at lower resolution, with some random fluctuation of their properties but no sign of systematic variation in the morphology or angular momentum profile against resolution.  \\item The formation of \\slows{} is greatly affected by numerical resolution. Models that produce slow rotators at the highest resolution  result in much faster rotators at lower, standard resolution. The effect is present, but relatively minor, in purely collisionless dry mergers. Discrepancies become major in wet mergers, even in cases with modest gas fractions like 10\\% of the baryonic mass.  \\item These effects cannot be attributed to our choice of initial conditions or interaction orbits, but actually relate to the physical treatment of the merging process itself. In particular, small-scale density fluctuations increase at high resolution, and they participate to scattering stellar orbits and largely influence the final degree of relaxation and orbital structure in the merger remnants.  \\item The effect of gas on the properties of merger remnants is generally considered to consist in preserving a higher angular momentum, in particular through enhancing the survival/rebuilding of disc components in merger remnants. We find that this picture is incomplete: at high resolution, gas still reforms discy components, but also forms a large number of dense substructures (massive star clusters, dense nuclei, tails and filaments, etc) that trigger rapid variations of the gravitational potential and the degree of relaxation of the final system. This effect is missed with a too low resolution. At high-enough resolution, adding gas to a given merger does not necessarily increase the rotational support of the final merger remnant; we even find a case of a wet merger with 20\\% of gas that has a final angular momentum parameter $\\lambda_R$ slightly lower than the corresponding dry merger. \\end{itemize}  At the present stage, our results do not indicate how frequently real slowly-rotating ellipticals were formed by binary (wet) mergers of disc galaxies, but they show that this can be a robust pathway for their formation. In the course of the \\atlas{}\\ project, we are conducting, analyzing, and comparing a large set of numerical simulations for various formation mechanisms, in order to derive which is (are) the main formation mode(s) of real slow rotators in the nearby Universe. Our present results already indicate the limitations of existing samples of galaxy merger simulations, and will then serve to estimate the required resolution, the limitations of numerical models and their possible biases.  More generally, the immediate implications of these findings on our understanding of early-type galaxy formation are: \\begin{itemize} \\item High resolution in simulations of major mergers does not just allow to resolve small-scale structures like nuclear systems and star clusters, but impacts the whole global properties of the elliptical-like merger remnants, at least for the slow-rotating ones. \\item The formation of slow-rotating elliptical galaxies can be achieved through a major merger relatively more easily than previously believed. It can be frequent even in wet mergers with relatively high gas fraction, and with late-type, disc-dominated progenitor galaxies. \\item Repeated mergers and/or dry mergers of galaxies that are already early-type systems are thus not the only theoretical path to produce slow-rotating galaxies. Major mergers of two dic galaxies, including wet mergers, can produce slow-rotating early-type galaxies. Further studies are needed to determine how common this formation mechanism is for slow-rotators. \\item Quantitative comparisons of major merger simulation results with the observed properties of real early-type galaxies require high resolution models. A typical requirement, according to our study, would be a spatial resolution better than 100~pc for both the gravitational N-body aspects (i.e., softening length) and the hydrodynamical ones (for instance, the size of groups of particles other with quantities are smoothed in SPH models). The mass resolution should correspond to at least $\\sim 10^6$ particle per galaxy per component, which typically corresponds to a mass resolution $\\sim 10^4$~M$_{\\sun}$ for the gas discs of bright spiral galaxies. We find reasonable convergence above this resolution, but cannot rule out that some systematic effects still exist; in any case simulations below this resolution level show clear and strong resolution effects. Unfortunately, most published samples of major merger simulations are below this typical resolution limit. \\item The small-scale properties of interstellar gas and clustered star formation are important for the global, large-scale properties of merger remnants. Simulations directly resolving gas cooling down to low temperatures, the formation of cold (molecular) gas clouds and star formation therein, are highly desirable to understand the whole process of early-type galaxy formation. Modern hydrodynamic codes are promising in this respect \\citep[e.g.][]{B09,kim09}. \\end{itemize}"
    },
    "1004/1004.1410_arXiv.txt": {
        "abstract": "{ We compare the spectral properties of 227 Gamma Ray Bursts (GRBs) detected by the \\fe\\ Gamma Ray Burst Monitor (GBM) up to February 2010 with those of bursts detected by the CGRO/BATSE instrument. Out of 227 \\fe\\ GRBs, 166 have a measured peak energy \\epo\\ of their $\\nu F_{\\nu}$ spectrum: of these 146 and 20 belong the long and short class, respectively. \\fe\\ long bursts follow the correlations defined by BATSE bursts between their  \\epo\\  vs fluence and peak flux: as already shown for the latter ones, these correlations and their slopes do not originate from instrumental selection effects. \\fe/GBM bursts extend such correlations toward lower fluence/peak energy values with respect to BATSE ones whereas no GBM long burst with  \\epo\\ exceeding a few MeV is found,  despite the possibility of detecting them. Again as for BATSE, $\\sim$ 5 \\% of long and almost all short GRBs detected by \\fe/GBM  are outliers of  the \\ep--isotropic equivalent energy (''Amati'') correlation while no outlier (neither long nor short) of the \\ep--isotropic equivalent luminosity (``Yonetoku\") correlation is found. \\fe\\ long bursts have similar typical values of  \\epo\\ but a harder low energy spectral index with respect to all BATSE events, exacerbating the inconsistency with the limiting slopes of the simplest synchrotron emission models. Although the short GRBs detected by \\fe\\ are still only a few, we confirm that their  \\epo\\ is greater and the low energy spectrum is harder than those of long ones. We discuss the robustness of these results with respect to observational biases induced by the differences between the GBM and BATSE instruments. } ",
        "introduction": "Our current knowledge of the spectral properties of the prompt emission in GRBs mainly relies on the data collected in almost 10 years by the Burst And Transient Source Experiment (BATSE) onboard the {\\it Compton Gamma--Ray Observatory (CGRO)} satellite.  BATSE was composed by eight identical modules each containing two NaI(Tl) scintillation detectors: a Large Area Detector (LAD), optimised for sensitivity and directional response,  and a Spectroscopy Detector (SD) optimised for energy band and resolution. Thanks to its wide energy coverage ($\\sim$25 keV - $\\sim$1 MeV) and all--sky viewing, BATSE detected more than 2700 GRBs. The published spectral catalogs comprise relatively small sub-samples of bright GRBs, selected on the basis of the burst fluence and/or peak flux \\citep{preece00,kaneko06}.  The analysis of such samples revealed two main results about the spectral and energetics properties of these GRBs.  The first one is that the prompt spectra can be typically well described by a curved function showing a peak - in a $\\nu F_\\nu$ representation - at a typical energy \\epo\\ of a few hundreds of keV but whose distribution spans nearly three orders of magnitude. Large dispersions characterise also the distributions of the low-- and high--energy photon indices, whose characteristic values are $\\alpha\\sim -1$ and $\\beta\\sim -2.3$, respectively \\citep{preece00,kaneko06}.   The second one is that the rest frame peak energy \\ep\\ of long bursts correlates both with the bolometric isotropic energy \\eiso\\ emitted during the prompt \\citep{amati02} and with the bolometric isotropic luminosity \\liso\\ estimated at the peak of the light curve \\citep{yonetoku04}. Such correlations are clearly an intriguing powerful clue about the dominant emission mechanism during the prompt phase. Furthermore, as by correcting by the presumed jet opening angle the dispersion of the correlations even reduces \\citep{ghirlanda04a}, GRBs become potential standard candles \\citep{ghirlanda04b}.  \\cite{nava08} (N08 hereafter) performed the spectral analysis on a larger fainter sample of BATSE bursts, selected by setting a limiting fluence to $F=10^{-6}$ erg/cm$^2$. They found that the spectral parameter \\epo\\ shows a strong correlation with the fluence $F$ (or the peak flux $P$). These imply that the derived distribution of \\epo\\ is strongly affected by a selection based on cuts in fluence (or peak flux). The correlations in the observer frames (\\epo--\\F and \\epo--\\Pf) may be just the consequence of the rest frame correlations mentioned above. Alternatively, it has been claimed that the rest frame correlations are the result of instrumental selection effects \\citep{band05,nakar05}.  \\citet[hereafter G08]{ghirlanda08} and N08 examined two instrumental selection effects which may affect the relation between \\epof\\ and \\epop\\ and found that - although they do affect the burst sample properties - they are not responsible for the correlations found in the observational planes. Moreover \\citet{ghirlanda10} recently showed that the correlation \\ep-- \\eiso\\  and \\ep--\\liso\\ holds for the time--resolved quantities within individual \\fe\\ bursts (see also \\citealt{firmani09} for Swift bursts and \\citealt{krimm09} for Swift--Suzaku GRBs) and that this ``time--resolved\" correlation is similar to that defined by the time integrated properties of different GRBs. These results support the hypothesis that the \"Amati\" and \"Yonetoku\" correlations do have a physical origin and that the trends seen in the \"observational planes\" are just their outcome.  The same analysis has been applied to short bursts \\citep{ghirlanda09}, though their lower fluence limits the possibility to determine with reasonable accuracy the properties of a statistically significant sample of events. Also in this case spectral information mainly arise from BATSE data.  Evidence of a spectral diversity between long and short bursts comes from their different hardness--ratios (HR) \\citep{kouveliotou93}. The larger HR of short bursts might be ascribed to a larger \\epo. However, due to the relation between \\epo\\ and the bolometric fluence and peak flux, a direct comparison between the \\epo\\ distributions of the two different burst classes must take the different fluence/peak flux selection criteria into account. \\citealt{ghirlanda09} (G09 hereafter) performed a detailed spectral analysis of samples of short and long BATSE bursts selected on the same limit on the peak flux. They found that the peak energy distributions of the two classes are similar, while the most significant difference concerns the low--energy power--law indices, with short bursts having typically a harder $\\alpha\\sim -0.4$.  The existence of intrinsic spectral--energy correlations for short bursts is still an open question, due to the paucity of short GRBs with measured redshift. However, the findings by G09 show that short events do not follow the \\ep--\\eiso\\ correlation defined by long ones, while their are consistent with a similar \\ep--\\liso\\ correlation \\citep{amati06,ohno08,krimm09}.   The \\fe\\ satellite, launched in June 2008, opened a new and powerful opportunity to shed light on the origin of the GRB prompt emission thanks to its two high energy instruments: the Large Area Telescope (LAT) and the Gamma--ray Burst Monitor (GBM). LAT detected in 1.5 year very high energy emission ($>$ 100 MeV) from 15 GRBs, out of which 12 are discussed in \\citet{ghisellini10}.  The GBM extends the BATSE energy range at both low and high energies. Its twelve NaI detectors assure a good spectral resolution between $\\sim$8 keV and $\\sim$1 MeV and two BGO detectors extend the instrument sensitivity up to $\\sim$40 MeV. While - as for BATSE - the NaI detectors guarantee full--sky coverage, the smaller geometric area (16 times lower than that of the LADs) result in a lower sensitivity.  Until now (beginning of March 2010), about 280 triggers of the GBM are reliably classified as GRBs\\footnote{http://heasarc.gsfc.nasa.gov/W3Browse/fermi/fermigbrst.html}. While a spectral catalogue has not yet been published, preliminary results of the spectral analysis performed by the GBM team have been distributed to the community through the Galactic Coordinates Network (GCNs) Circulars. For 166 events the spectral properties (\\epo) are well constrained.  Their statistically significant number allows to start a meaningful comparison between BATSE and GBM results. The fact that the two instruments have a different energy range and effective area sensitivities could be in turn exploited to infer the properties of GRBs prompt emission. This is the aim of this work.  In other words we try to assess whether the GBM is confirming any/which ones of the BATSE results or is revealing something new about prompt spectra in the hard X/$\\gamma$--ray energy range and the existence of the above mentioned spectra vs energetics correlations.  In particular the GBM wider energy window should allow to examine two important aspects: on one hand the low energy photon index - as first inferred from BATSE data - indicate (\\citealt{preece98a,preece02,ghisellini00}) an inadequacy of a simple synchrotron model to account for the GRB prompt emission; on the other hand better constraints on the high energy spectral index may help to discriminate between quasi--thermal and non--thermal processes and would give clues on the energy distribution of the emitting electrons. Moreover, due to its wider energy range, the GBM could more easily detect soft and hard bursts, i.e. those events whose peak energy is outside the BATSE energy range.  In this paper we study these issues with \\fe/GBM bursts. In \\S 2 we present the samples of \\fe\\ bursts (both long and short) and their spectral parameters, as reported in the GCNs. In \\S 3 the properties of \\fe/GBM bursts are compared with the BATSE ones in terms of \\epo, \\F\\ and \\Pf, i.e. in the \\epof\\ and \\epop\\ planes. We compute the relevant instrumental selection effects introduced by the GBM characteristics on these quantities, considering separately short and long GRBs. We compare the spectral parameters of BATSE and \\fe\\ bursts in \\S 4 and discuss in \\S 5 the similarities and differences of the spectral properties of GRBs between the two instruments and, for each of them, between the population of short and long events. ",
        "conclusions": "\\begin{figure*} \\vskip -2 true cm \\includegraphics[scale=0.6]{figurella4.ps} \\vskip -2 true cm \\caption{ Schematic view of the distribution of long and short GRBs in the \\epof\\ and \\epop\\ planes. The horizontal dashed line at $\\sim$ 30 keV represents the lower limit for the GBM instrument: the simulations performed in this work show that \\epo\\ can be hardly determined below this value. For the BATSE instrument this limit corresponds to $\\sim$ 50 keV. The upper limit for BATSE is at $\\sim$ 1 MeV, while for the GBM there is no upper limit in this plane. The vertical dashed line (left panel) shows an example of fluence selection, while the dashed curve (right panel) refers to the photon flux selection criterion adopted by G09.}  \\label{figurella} \\end{figure*}  In this work we study the spectral properties of \\fe\\ bursts detected by the GBM instrument up to February 2010. We collect a sample of 227 GRBs (37 of which are classified as short) from the GCN circulars. For 146 long and 20 short bursts the spectrum could be fitted by a curved model so that the peak energy of the $\\nu F_{\\nu}$ spectrum could be constrained.  The population of long bursts is large enough to allow a statistical comparison with the BATSE results. We consider a complete sample of BATSE bursts down to a limiting fluence of $F\\sim10^{-6}$ erg cm$^{-2}$. Fig. \\ref{fluence_tutti} shows that the two samples have very similar fluence distributions. For both instruments the \\epo\\ distribution peaks around \\epo$\\sim 150$ keV (Fig. \\ref{ep_tutti}). Despite of its larger energy range, the GBM extends the \\epo\\ distribution of long bursts only at low energies with respect to BATSE. The $\\alpha$ distribution, instead, reveals some difference: GBM bursts have on average a harder low energy photon index ($\\langle\\alpha_{\\rm GBM}\\rangle=-0.9$ and $\\langle\\alpha_{\\rm BATSE}\\rangle=-1.1$, KS probability $P=6\\times10^{-9}$). Also for GBM short bursts we can draw some conclusions about their spectral properties. Their \\epo\\ distribution is shifted towards higher energies  compared both to long bursts from the same instrument and to short bursts seen by BATSE. The lack of low energy \\epo\\ (below $\\sim 250$ keV) can be accounted for by the spectral threshold we derived for the GBM instrument (see Fig. \\ref{piani correle}), while the larger energy coverage allows the detection of \\epo\\ up to $\\sim$ 4 MeV. \\fe\\ data confirm that short bursts have on average a harder $\\alpha$ compared to long bursts ($\\langle\\alpha_{\\rm GBM,short}\\rangle\\sim -0.46$), as already found in the BATSE sample by G09.  Similarly to what has been done for long and short GRBs detected by BATSE (N08; G09) we examine the distribution of \\fe\\ bursts in the \\epo--Fluence and \\epo--Peak Flux planes in order to study instrumental selection effects and test their consistency with the rest frame correlations (i.e. \\ama\\ and \\yone, respectively) defined by GRBs with measured redshifts.  Our main results are: \\begin{itemize}  \\item {\\it Long GRBs} detected by \\fe\\ follow the same \\epof\\ and \\epop\\ correlations defined by BATSE GRBs (Fig. \\ref{piani correle}). We computed the instrumental selection effects of \\fe/GBM -- as already done for BATSE (G08; N08): the trigger threshold and the spectral threshold are not responsible for the correlations defined by long GRBs in both planes (see Fig. \\ref{piani correle}). The \\fe/GBM spectral range down to 8 keV with respect to the limit of 30 keV of BATSE,  allows to extend the correlations to lower peak energies/fluences. Instead, despite of the higher energy sensitivity of \\fe/GBM (up to 40 MeV) no long GRBs with \\epo\\ larger than a few MeV are detected. This can be due to a real absence of bursts with such high \\epo\\ or to the fact that they have large fluences, thus being too rare to be detected during 1.5 year of \\fe\\ observations.  We conclude that long GRBs detected by \\fe\\ confirm what found with BATSE bursts, in particular that they follow a correlation both in the \\epof\\ and in the \\epop\\ plane. Moreover, the fraction of bursts detected by \\fe\\ which are outliers at more than 3$\\sigma$ with respect to the \\ama\\ correlation is $\\sim$5\\%, similarly to what found in the BATSE sample, while there are no outliers (at more than 3$\\sigma$) of the \\yone\\ correlation among \\fe\\ long GRBs.  \\item {\\it Short GRBs} detected by \\fe\\ populate a different region in the \\epof\\ plane with respect to long events, the firmer having a larger peak energy and lower fluences compared to the latter ones. This is consistent with what found by BATSE and confirms  that short GRBs do not follow the \"Amati\" correlation but they obey the \"Yonetoku\" correlation defined by long events. \\end{itemize}   The comparison of short and a representative sample of long BATSE GRBs (selected with a similar peak flux threshold) led G09 to conclude that their main spectral diversity is due to a harder low energy spectral index in short bursts while their \\epo\\ of BATSE is similarly distributed. \\fe\\ bursts provide the the opportunity of re-examining this result for the population of short and long GRBs detected by the GBM and also compare their spectral properties with those of the  BATSE ones. We find that:  \\begin{itemize} \\item \\epo\\ of short \\fe\\ GRBs is larger and \\al\\ smaller that those of long one, indicating that short events are harder, both in terms of their peak energy and low energy spectral index.  \\item A comparison between GBM and BATSE short bursts reveals that they have similar \\al\\ while the \\epo\\ of short GBM bursts is larger than that of short BATSE events (see Fig. \\ref{piani correle}, bottom left panel). This information is allowed by the higher energies which can be detected by the GBM. Moreover, the different \\epo\\ distribution of BATSE and GBM short bursts is affected by the lower sensitivity of the GBM instrument, which misses short bursts at low fluence (and therefore low \\epo).  \\item \\fe\\ and BATSE long bursts have a similar \\epo\\  while \\fe\\ events tend to have a harder low energy spectral index (Fig. \\ref{a_fermi}).  \\end{itemize}  Fig. \\ref{figurella} shows a schematic representation of the current information about the distribution of short and long bursts in the \\epof\\ and \\epop\\ planes. With respect to BATSE, the GBM reveals that long bursts extend to lower \\epo, consistently with what previously found with other instruments (mainly Hete--II and Swift).  Despite of the high energy sensitivity, also the \\epo\\ distribution of \\fe\\ long events extends only up to $\\sim$ 1 MeV. The situation is different for short GRBs whose \\epo\\ reach up to $\\sim$ 4 MeV in the present sample. These high \\epo\\ were not detectable by BATSE, whose sensitivity drops at $\\sim$ 1 MeV (upper horizontal dashed line in Fig. \\ref{figurella}). Therefore, the \\fe/GBM shows that short GRBs have larger \\epo\\ with respect to long ones, contrary to what found with BATSE (G09).  When comparing the \\epo\\ distribution of short and long bursts, different conclusions can be drawn, according to the selection criterion of the samples. The left panel in Fig. \\ref{figurella} shows that a given cut in fluence (represented by the vertical dashed line) would result in different \\epo\\ distributions between short and long bursts, resulting from their different location in the \\epof\\ plane. The right panel in Fig. \\ref{figurella} illustrates, instead, what happens for a selection in {\\it photon flux}. This translates into a curve in {\\it energy flux}: the dashed curve represents the cut applied by G09 to select both short and long bursts, corresponding to a photon flux larger than 3 ph cm$^{-2}$s$^{-1}$. This criterion applied to BATSE bursts produces similar \\epo\\ distributions of long and short events, as indeed found by G09. The very same criterion applied to \\fe/GBM bursts results, instead, in different distributions, since short bursts can have very high \\epo\\ values, not detected in the sample of long bursts."
    },
    "1004/1004.3286_arXiv.txt": {
        "abstract": "We used spectral observations of Fe I line profiles with a 200 000 resolution to determine micro and macroturbulent velocities in the atmospheres of the Sun as a star, $\\alpha$ Cen A, Procyon ($\\alpha$~CMi), Arcturus ($\\alpha$ Boo), and Canopus ($\\alpha$ Car). Isotropic microturturbulent velocities ($V_{mi}$) and radial-tangential macroturbulent velocities ($V_{ma}^{RT}$) were found to be a quite suitable approximation to the velocity field in the atmospheres of all stars studied except Canopus. The average velocities $V_{mi}$ and $V_{ma}^{RT}$ are $0.8\\pm0.1$ and $2.6\\pm0.3$~km/s for the Sun as a star, $0.8\\pm0.2$ and $2.9\\pm0.4$~km/s for $\\alpha$ Cen A, $0.8\\pm0.3$ and $5.9\\pm0.2$~km/s for Procyon, $1.0\\pm0.2$ and $4.6\\pm0.3$~km/s for Arcturus. The velocity field in the atmosphere of Canopus can be described  by an anisotropic radial-tangential distribution of microturbulence  with $V_{mi}^{RT} = 2.1$~km/s and anisotropic distribution of macroturbulence with $V_{ma}^{rad} =17\\pm2$~km/s and $V_{ma}^{tan}=1.3\\pm1.0$~km/s. From Fourier analysis of broadening and shapes of three spectral lines of Fe I, we have  derived the rotation velocity $V_e \\sin i=3.5 \\pm 0.2$~km/s for Canopus. ",
        "introduction": "\\label{S-Introduction} An advantage of the Fourier analysis of stellar spectra is that it allows the separation of the contributions to the Fourier transforms from various processes which are responsible for line broadening. Some examples of application of the Fourier analysis to the investigations of stellar rotation and astrospectroscopic turbulence as well as the description of its details and problems can be found in Gray's book \\cite{5}. However, it is not always possible to realize all the advantages of this technique in actual practice, as the Fourier transforms of the principal broadening processes are reliably separated in the high frequency domain, which can be studied only in observations with high signal-to-noise ratios and high spectral resolutions. Nevertheless, the Fourier analysis extended the possibilities for interpreting the broadening of spectral lines produced by mechanical motions in stellar atmospheres. While it was found for solar type stars that line broadening and shifts are caused by the photospheric overshooting convection and oscillatory motions  \\cite{11,19,40,41}, the question of the nature of motions in the atmospheres of nonsolar type stars remains open  \\cite{14,17,31}.  Our analysis is based on unique spectral observations made by Dravins  \\cite{16}, and its main objective is to determine the rotation velocity and study the distribution of micro-macroturbulent velocities in the atmospheres of stars of various types. With the stars appropriately selected, we are able to solve the following problems.  1. Follow variations in the velocity field when going from the Sun to other stars in the H--R diagram. The Sun can be used as a lest star.  2. Models are available for all stars studied. They are either theoretical three-dimensional (3D) inhomogeneous models (the Sun and Procyon  \\cite{11,40,41}, $\\alpha$ Cen A  \\cite{41}) or semiempirical multicomponent models (Arcturus and Canopus  \\cite{17}). Thus the analysis results can be interpreted more or less reliably and used to explain the granulation boundary in the H--R diagram found by Gray and Nagel  \\cite{29}. Lines in the spectra of the stars located to the right of this boundary (cooler stars) have a red C-shaped asymmetry of the solar type, while lines in hotter stars display a violet inversion asymmetry. The position of the boundary approximately corresponds to those stars of luminosity classes III--V in which convective envelopes adjacent to atmospheric layers can already form. Thus, a red asymmetry of spectral lines may suggest that granulation exists on the stellar surface and the velocity field of the thermal convection (or photospheric overshooting convection) strongly affects the line profiles. This conclusion is at variance, however, with the data for stars of luminosity classes I--II, for which Gray's granulation boundary and the boundary of convection-induced envelopes (according to the mixing-length theory  \\cite{29}) do not coincide. At the same time a violet asymmetry of stellar spectral lines may be interpreted as a prevalence of mechanical motions other than convection in stellar atmospheres (oscillations, mass loss, etc.).  3. Four stars in our list (the Sun, $\\alpha$ Cen A, Procyon, Arcturus) have convective envelopes. The photospheric overshooting convection is of different character in each star, however, and it is of interest to relate the distinctions to the astrospectroscopic turbulence parameters. The fifth star, the supergiant Canopus, lies to the left of the granulation boundary, it has no convective envelope, and the lines in its spectrum have a violet inversion asymmetry  \\cite{16}. Analysis of its atmospheric velocity field is of special interest. ",
        "conclusions": "\\label{Conclusion} The isotropic distribution of microturbulent velocities and the radial-tangential distribution of macroturbulent velocities seem to represent most adequately the velocity field in the atmospheres of the Sun, $\\alpha$ Cen A, Procyon, and Arcturus within the scope of their homogeneous models. For Canopus, the microturbulence is anisotropic and the macroturbulence displays a pronounced asymmetry between its radial and tangential components. The rotation velocity found by us for Canopus turned out to be low.  \\vspace{1.cm} {\\bf Acknowledgements.} We thank Prof. D. Dravins for furnishing the observational data. The study was partially financed by the State Fundamental Research Foundation of Ukraine (Grant No. 6.4/192)."
    },
    "1004/1004.4974_arXiv.txt": {
        "abstract": "We model the formation of  the Galactic stellar halo via the accretion of satellite galaxies onto a time-dependent semi-cosmological galactic potential. Our goal is to  characterize the substructure left by these accretion events  in a close manner  to what may be  possible with the {\\it  Gaia} mission.   We have  created a  synthetic {\\it  Gaia} Solar Neighbourhood catalogue  by convolving the  6D phase-space coordinates of  stellar particles from  our disrupted  satellites with  the latest estimates of the {\\it Gaia} measurement errors, and included realistic background  contamination due to  the Galactic  disc(s) and  bulge. We find  that,  even  after  accounting for  the  expected  observational errors,  the resulting phase-space  is full  of substructure.   We are able to successfully isolate  roughly 50\\% of the different satellites contributing to  the `Solar Neighbourhood' by  applying the Mean-Shift clustering algorithm in  energy-angular momentum space. Furthermore, a Fourier  analysis of  the space  of orbital  frequencies allows  us to obtain accurate  estimates of  time since accretion  for approximately 30\\% of the recovered satellites. ",
        "introduction": "How galaxies form and evolve remains one of the most interesting and challenging puzzles  in astronomy.  Although a great  deal of progress has  been  made  in  the  last  decade  many  questions  await  to  be addressed. Both,  theory and observations  are now converging  about a key ingredient  of this formation  process; galaxies as our  own Milky Way seem  to have  experienced the accretion  of smaller  objects that have  come   together  thanks  to  the  relentless   pull  of  gravity \\citep[e.g.][]{wr}.  From  the observational side,  several new  studies have  revealed the presence  of a  large amount  of stellar  streams in  galactic stellar haloes,  echoes of  ancient as  well as  recent and  ongoing accretion events. These  stellar streams have  been preferentially found  in the outer   regions   of   galaxies.    The  Sagittarius   Tidal   Streams \\citep{ibata94,ibata01b,majewski} and the Orphan Stream \\citep{belu07} are just two examples of  satellite debris in the Milky Way \\citep[see also][]{new02,ibata03,yanny03,belu06,  grill,  else09}.  The  abundant substructures    found    not    only    in   the    halo    of    M31 \\citep[e.g.][]{ibata01a,McCo09}, but  also in several  haloes of other galaxies \\citep[e.g.][]{Mart08,Mart09}  are an unambiguous  proof that accretion is inherent to the process of galaxy formation.  On the theoretical side, models  of the formation of stellar haloes in the  $\\Lambda$CDM cosmogony  have  been able  to  explain their  gross structural properties \\citep[e.g.][]{bj05,lh,cooper}. In these models, the inner  regions of the  haloes (including the  Solar Neighbourhood) typically formed  first and hence  contain information about  the most ancient   accretion   events   that   the   galaxy   has   experienced \\citep[e.g.][]   {ws2000,hws03,bj05,tumli}.   Previous   studies  have predicted  that several  hundreds kinematically  cold  stellar streams should  be present  in  the Solar  Vicinity  \\citep{hw}. However,  the identification  of  these  streams  is  quite  challenging  especially because of the  small size of the samples of  halo stars with accurate 3-D velocities  currently available.  Full  phase-space information is necessary because  of the very  short mixing time-scales in  the inner regions of the halo. Some progress has been made towards building such a  catalogue but  only a  few detections  have been  reported  to date \\citep{hwzz99,klement09,smith09}.  These have  made use  of catalogues such   as  SEGUE   \\citep{yanny09}  and   RAVE   \\citep{zwitter08}  in combination  with  Tycho   \\citep{thyco}  and  Hipparcos  \\citep{hip}. Clearly, this field will only advance significantly with the launch of the astrometric satellite {\\it  Gaia} \\citep{gaia}.  This mission from the European  Space Agency will  provide accurate measurements  of the 6-D  phase-space coordinates  of  an extraordinarily  large number  of stars\\footnote{Details about the latest {\\it Gaia} performance numbers may  be   found  at:  http://www.rssd.esa.int/gaia}.    Together  with positions, proper  motions and parallaxes  of all stars  brighter than $V=20$,  {\\it Gaia}  will also  measure  radial velocities  down to  a magnitude of $V=17$.  To  unravel  the accretion  history  of  the  Milky Way  requires  the development  of  theoretical  tools  that will  ultimately  allow  the identification and characterization of the substructure present in the {\\it Gaia} data  set. Several authors have studied  the suitability of various  spaces to  isolate debris  from accretion  events \\citep[see, e.g.][]{hz00,knebe05,ari_fu,font,h06,sj09}.  Recently \\citet[][]{mcm}, followed  by  \\citet[][hereafter   GH10]{gh10},  showed  that  orbital frequencies form  a very  suitable space in  which to  identify debris from past merger events. In frequency space individual streams from an accreted  satellite can  be  easily identified,  and their  separation provides a  direct measurement of its time  of accretion. Furthermore, GH10 showed  for a few  idealized gravitational potentials  that these features are preserved  also in systems that have  evolved strongly in time.  While  very promising, these  studies have focused on  a single accretion  event  onto a  host  galaxy.   In  reality, we  expect  the Galactic stellar  halo to have formed  as a result  of multiple merger events  \\citep[where most  of its  mass  should have  originated in  a handful  of significant  contributors; see  e.g.][]{lh,cooper}.   As a consequence, it  is likely that debris from  different satellites will overlap in frequency space, complicating their detection.  In  this  paper  we  follow   multiple  accretion  events  in  a  live (cosmologically motivated) galactic  gravitational potential. We focus on how the latest estimates of the measurement errors expected for the {\\it Gaia}  mission will affect  our ability to recover  these events. In  particular  we  study  the  distribution  of  stellar  streams  in frequency  space  and  the  determination of  satellite's  time  since accretion from  stellar particles located  in a region like  the Solar Neighbourhood.  Our  paper is organized  as follows.  In Section  2 we present the details  of the $N$-body simulations carried  out, as well as the  steps taken to generate  a Mock {\\it  Gaia} stellar catalogue. In Section  3 we  characterize the distribution  of debris in  a Solar Neighbourhood-like sphere in the  absence of measurement errors, while in  Section  4  we focus  on  the  analysis  of  the Mock  {\\it  Gaia} catalogue. We discuss and summarize our results in Section 5. ",
        "conclusions": "\\begin{figure} \\centering \\includegraphics[width=83mm,clip]{fig11.eps} \\caption{Distribution of  stellar particles  inside a sphere  of 4~kpc radius at 8~kpc from the galactic centre in $E$ vs. $L_{z}$ space as would  be observed  by {\\it  Gaia}. The  different colours  show the groups identified  by the Mean  Shift algorithm. Black  open circles denote  those for which  the time  since accretion  was successfully derived.} \\label{fig:mstest} \\end{figure}  We  have studied  the characteristics  of merger  debris in  the Solar neighbourhood as  may be observed by  ESA's {\\it Gaia}  mission in the near  future. We  have  run a  suite  of $N$-body  simulations of  the formation  of the  Milky Way  stellar  halo set  up to  match, at  the present  time, its known  properties such  as the  velocity ellipsoid, density  profile and  total  luminosity.  The  simulations follow  the accretion of a  set of 42 satellite galaxies  onto a semi-cosmological time dependent  Galactic potential.  These satellites  are evolved for 10~Gyr,  and  we  use  the  final  positions  and  velocities  of  the constituent particles to generate a Mock {\\it Gaia} catalogue.  Using synthetic  CMDs, we have resampled the  satellite's particles to represent  stars down to  $M_V \\approx  4.5$. This  absolute magnitude corresponds to an apparent magnitude  $V \\approx 16$ at 2.5~kpc, which is  comparable  to the  {\\it  Gaia}  magnitude  limit for  which  full phase-space information  will be  available.  Our Mock  catalogue also includes a  Galactic background population  of stars represented  by a Monte   Carlo  model   of  the   Galactic  disc   and  bulge,   as  in \\citet{brown05}.  At 8~kpc from  the Galactic  centre, stars  from the disc  largely outnumber  those  of  the stellar  halo,  however it  is possible  to  reduce  their  impact   by  applying  a  simple  cut  on metallicity.  We  have estimated,  using the latest  determinations of the MDF of the Galactic  disc(s) \\citep{gc,ive08}, that only $10\\ 000$ stars out  of the  estimated $4.1 \\times  10^{7}$ disc  stars brighter than $V=17$ in the Solar neighbourhood should have [${\\rm Fe}/{\\rm H}] \\leq -1.1$.  A smaller number  of bulge stars ($\\sim 800$) with [${\\rm Fe}/{\\rm H}]  \\leq -1.1$ is  expected to contaminate our  stellar halo sample,  down to  the faintest  $M_V$. This  fraction  represents only $2.6$ per  cent of the whole  Mock {\\it Gaia}  stellar halo catalogue, and therefore does not constitute an important obstacle to our ability to characterize this component.  Finally, we have convolved the positions and velocities of all `observable stars' in our Solar Neighbourhood sphere with the latest models of the {\\it Gaia} observational errors, according to performances given by ESA (http://www.rssd.esa.int/gaia).  The  analysis presented  here confirms  previous results,  namely that satellites can  be identified as coherent clumps  in phase-space, e.g. in  the $E$-$L_z$  projection \\citep[see][]{hz00,  knebe05,  font}. We find    that    a   clustering    algorithm    such   as    Mean-Shift \\citep[][]{fukuhost,comameer,derpa} is able to recover roughly 50\\% of all   satellites   contributing  stellar   particles   to  the   Solar neighbourhood sphere.  We have also demonstrated that  even after accounting for the Galactic background contamination  and the expected measurement  errors for the {\\it Gaia} mission, the space of orbital frequencies is also very rich in substructure. In this space streams from a given satellite define a regular  pattern  whose  characteristic  scale is  determined  by  the satellite's time  since accretion.  We find  that reasonable estimates of the time  since accretion may be provided when  the number of stars with accurate parallaxes ($\\sigma_{\\varpi} / \\varpi \\leq 0.02$) from a given  satellite is at  least $\\sim  100$.  This  was possible  in our simulations for 4 cases.  Together  with  {\\it  Gaia},  ground  based  follow-up  campaigns  are currently  being planned  by European  astronomical  community.  Their purpose  is   to  complement  a  future  {\\it   Gaia}  catalogue  with information that either will not be obtained or for which the accuracy will  be  very low.   As  an  example,  high resolution  multi  object spectroscopy  in combination  with ground  based photometry  could not only push the limiting magnitude  of the phase-space catalogue down to $V  \\approx 20$, but  could also  provide detailed  chemical abundance patterns of individual stars.  The former will require the development of accurate and precise  photometric distances indicators which can be calibrated  using the large  number of  stars near  the Sun  with very precise trigonometric parallaxes  from {\\it Gaia}.  The identification of satellites in, e.g., $E-L_{z}$ space could be considerably improved by having such an  extended full phase-space catalogue.  Firstly, very faint satellites could now be  observed and, secondly, the much larger sample of stars could allow us to apply a clustering algorithm in very small volumes  around the  Sun with enough  resolution to  avoid large overfragmentation.  The  later is  important because by  extending the volume probed  we expect to  reduce the overlap between  satellites in $E-L_{z}$ space.  In addition, chemical tagging will be very important to characterize the satellite's  star formation and chemical histories \\citep{FBH}.   Note, however,  that in  our simulations  we  find that stellar particles in a given  stream do not originate from a localized region  in  physical  space   (such  as  a  single  molecular  cloud). Therefore even individual  streams are likely to reflect  the full MDF present in the object at the time it was accreted.   Although  our satellites  were evolved  in a  cosmologically motivated time  dependent host  potential, our  simulations do  not  contain the fully   hierarchical   and  often   chaotic   build-up  of   structure characteristic of  the $\\Lambda$ cold dark matter  model.  The violent variation  of the  host potential  during merger  events,  the chaotic behavior induced by a triaxial dark matter halo \\citep[e.g.][]{vol08}, and the orbital evolution  due to baryonic condensation \\citep{vallu}, are potentially  important effects which we have  neglected and should be taken into account in future work.  To assess the impact of some of these effects on the distribution of debris in the Solar Neighbourhood we  are  currently  analysing  a fully  cosmological  high  resolution simulation of the formation of  the Galactic stellar halo based on the Aquarius project \\citep{cooper}."
    },
    "1004/1004.2555_arXiv.txt": {
        "abstract": "The properties of the -- presumably -- youngest Galactic supernova remnant (SNR) G1.9+0.3 are investigated within the framework of nonlinear kinetic theory of cosmic ray acceleration in SNRs. The observed angular size and expansion speed as well as the radio and X-ray emission measurements are used to determine relevant physical parameters of this SNR. Under the assumption that SNR G1.9+0.3 is the result of a Type Ia supernova near the Galactic center (at the distance $d=8.5$~kpc) the nonthermal properties are calculated. In particular, the expected TeV gamma-ray spectral energy density is predicted to be as low as $\\epsilon_{\\gamma}F_{\\gamma} \\approx 5\\times10^{-15}$~erg~cm$^{-2}$~s$^{-1}$, strongly dependent ($F_{\\gamma}\\propto d^{-11}$) upon the source distance $d$. ",
        "introduction": "G1.9+0.3 has been known as a potentially young shell type Galactic supernova remnant (SNR) of very small angular size \\citep{gg84}. Recently, the interest in this SNR was revived by \\citet{rey08, rey09} who analyzed the expansion rate of the object and deduced an age $t_\\mathrm{SN}$ of about 100~yr, which makes it the youngest known SNR in the Galaxy. Although the expansion rate was derived by a comparison of radio observations in 1985 and {\\em Chandra} observations in 2007, this rate has been confirmed very soon thereafter by independent radio observations \\citep{green08, murphy08}.  According to \\citet{rey08}, the line-free X-ray emission has a pure synchrotron origin which clearly indicates that effective particle acceleration takes place, at least for electrons. There are also arguments, like the bilateral symmetry of the  X-ray synchrotron emission suggesting a roughly uniform ambient magnetic field, that favor a type Ia origin for G1.9+0.3. Finally, the distance estimate $d =8.5$~kpc is based on an analysis of the absorption toward G1.9+0.3.  During the survey of the inner Galaxy by H.E.S.S. in very high energy $\\gamma$-rays, no emission was reported from the direction to G1.9+0.3 \\citep{aha06}. Therefore, one can derive an upper limit at the level of 2\\% of the Crab flux above 200 GeV.  For the purpose of a more general study of such an unusual object regarding its nonthermal properties, it is of interest to describe it by a kinetic theory of cosmic ray (CR) acceleration in SNRs, coupled with the gas dynamics of the thermal plasma, as given by \\citet{bek96} and \\citet{bv97}. This model assumes spherical symmetry, although the assumption is later relaxed.  Similar models on almost the same physical basis have recently been developed by two other groups \\citep{KangJ06,zp09}, whose calculations very well confirm these earlier results.  The kinetic description allows a corresponding analysis of the nonthermal evolution of an SNR at a very early phase. This assumes that the plasma physics underlying especially the temporal dependence of the magnetic field amplification process can be extrapolated to such an early evolutionary phase, where the dynamical behavior of the ejecta plays an essential role. Such a study is reported here. It is combined with a discussion about the influence of the assumption of a smaller distance on the TeV \\gr flux. ",
        "conclusions": "The existing data for G1.9+0.3, when analyzed within the framework of the nonlinear kinetic theory of CR production in SNRs described above, are consistent with a type Ia explosion in a rarefied medium at a distance of $d=8.5$~kpc, whose nuclear CR spectrum reaches the energy of the ``knee'' in the observed Galactic CR spectrum at the present epoch. This conclusion also concerns the derived strong magnetic field amplification. A test particle, purely leptonic gamma-ray scenario, is inconsistent with existing TeV gamma-ray observations with the H.E.S.S. telescope array. However, the data set is not complete enough to unequivocally determine the value of the source distance. Since the expected \\gr flux is very sensitive to the distance $d$, $F_{\\gamma} \\propto d^{-11}$, a detection of the \\gr flux from G1.9+0.3, as improbable as it may be, would yield the distance. However, in the case when G1.9+0.3 is located near the Galactic center ($d=8.5$~kpc), the expected TeV \\gr energy flux is so low, $\\epsilon_{\\gamma}F_{\\gamma}\\approx 5\\times10^{-15}$~erg~cm$^{-2}$~s$^{-1}$, that it is not detectable with present instruments. It is clear that for a distance that was not larger than $d=5.6$~kpc G1.9+0.3 could be visible with future instruments like CTA which can be assumed to have a sensitivity of $\\sim 1$~mcrab at a few 100 GeV \\citep{bernloehr09}."
    },
    "1004/1004.2249_arXiv.txt": {
        "abstract": "{  A previously reported literature search suggested that the AGB stars in Galactic globular clusters may be showing different distributions of CN-strong and CN-weak stars as compared to their RGB stars. In most cases the second giant branches of GCs appeared to be deficient in stars with strong CN bands. However the sample sizes of AGB stars at that time were too small to give a definitive picture. Thus an observing campaign targeting GC AGB stars was proposed. We now have medium resolution spectral observations of about 250 GC AGB stars across 9 globular clusters, obtained with the 2dF/AAOmega instrument on the Anglo-Australian Telescope. In this paper we report some preliminary results regarding the distributions of CN-strong and CN-weak stars on the two giant branches of a selection of globular clusters. We find that some GCs show a total lack of CN-strong stars on the AGB, whilst some show a reduction in CN-strong stars as compared to the RGB. Standard stellar evolution does not predict this change in surface abundance between the two giant branches. We discuss some possible causes of this unexpected phenomenon. ",
        "introduction": "Although Galactic globular clusters (GCs) are chemically homogeneous with respect to Fe and most other heavy elements (see e.g. \\citealt{KSL92}), it has long been known that they show inhomogeneities in many lighter elements (e.g. C, N, O, Mg, Al). These inhomogeneities are considered anomalous because they are not seen in halo field stars of similar metallicity (see e.g. \\citealt{GSC00}).  One of the first inhomogeneities discovered was that of the molecule cyanogen (CN, often used as a proxy for nitrogen). A picture of `CN-bimodality' emerged in the early 1980s whereby there appears to be two distinct chemical populations of stars in most, if not all, GCs. One population is known as `CN-strong', the other `CN-weak' (we note that the CN-weak population might be more informatively called `CN-normal' -- as these stars show CN abundances similar to the Halo field stars). Originally, observations of CN were mainly made in stars on the giant branches but more recently there have been observations on the main sequence (MS) and sub-giant branch (SGB) of some clusters (e.g. \\citealt{CCB98}). These observations show that there is little difference in the bimodal CN pattern on the MS and SGB as compared with the giants --- indicating a primordial origin for the differing populations. Figure 6 in \\citet{CCB98} exemplifies this situation.  Due to the paucity of asymptotic giant branch (AGB) stars in GCs (a result of their short lifetimes) there have been very few systematic observational studies of the CN anomaly on the AGB in globular clusters (\\citealt{M78} is one of which we are aware).  What little that has been done has been an aside in more general papers (e.g. \\citealt{NCF81}, \\citealt{BSH93}, \\citealt{ISK99}). However these studies have hinted at a tantalising characteristic: most (observed) GCs show a lack of CN-strong stars on the AGB. If this is true then it is in stark contrast to the red giant branch (RGB) and earlier phases of evolution, where the ratio of CN-Strong to CN-Weak stars is roughly unity in many clusters.  This \\textit{possible} discrepancy was noted by \\citet{NCF81} in their paper about abundances in giant stars in NGC 6752. They state that ``The behaviour of the CN bands in the AGB stars is... quite difficult to understand... not one of the stars studied here has enhanced CN... yet on the [first] giant branch there are more CN strong stars than CN weak ones.''  (also see Figure 3 of that paper).  Motivated by this observation we conducted a literature search to ascertain whether this was true of other GCs (\\citealt{Paper1}). The literature search revealed that the AGB star counts for all studies (which are not, in general, studies about AGB stars in particular) are low, usually being $\\leq 10$ (see Table 1 of \\citealt{Paper1}). The search also revealed that the picture may not be consistent between clusters. Although most clusters appear to have CN-weak AGBs, at least two seem to have CN-strong AGBs (M5 \\& 47 Tuc).  To further complicate the picture, clusters often appear to have a combination of both CN-strong and CN-weak stars on their AGBs. Again, all these assertions are however based on small sample sizes.  We note that \\citet{SIK00} also presented a conference paper on this exact topic. Compiling the data in the literature at the time they discussed the relative amounts of CN in AGB and RGB stars in the GCs. They also discuss Na abundance variations in M13.  In their closing remarks they suggest observations with larger sample sizes are needed --- which may be done using wide-field multi-object spectroscopes. This is the same conclusion the present authors also came to.  Here we present some preliminary results of a study that increases the GC AGB sample sizes substantially. With this new information we hope to confirm or disprove the existence of these abundance differences. ",
        "conclusions": "The preliminary results of the current study strongly support the unexpected phenomenon in which CN-rich stars seem to `disappear' between the RGB and AGB, leaving CN-weak dominated AGBs.  There is the possibility that the measurements of CN in AGB stars are biased in some way, either by gravity or temperature differences as compared to the RGB stars. However, as discussed by \\cite{NCF81}, the fact that we see stars at the same temperature on the RGB with varying CN band strengths indicates that this is not affecting the results. The higher luminosity of the AGB stars at the same temperatures would tend to increase the CN band strengths, indicating that gravity also is not affecting the results. Furthermore, in some of our observations we see CN-rich AGBs as well as CN-weak AGBs in the same cluster, which again indicates that the AGB measurements are not affected by their slightly different properties. We do however plan to perform spectral synthesis calculations to double-check. A benefit of the current study is that our spectra data set is homogeneous (since all the spectra were taken on the same instrument).  If we accept the above arguments and that the paucity (or total lack of) CN-strong stars on the AGB is real, then we need to look at evolutionary affects. There is however no known reason why the surface abundance of CN (which is a proxy for N) should reduce when a star evolves from the RGB to the AGB, at least in standard stellar evolution theory. Moreover, due to the well-known `deep mixing' effect that increases the N abundances (at the expense of C) in stars as they ascend the RGB (e.g. \\citealt{S03}), this is actually the opposite to what we would expect on the AGB -- we would predict, if anything, more CN-strong stars on the AGB!  \\cite{NCF81} proposed two possible explanations to explain the (apparent) lack of CN-strong stars on the AGB:  \\begin{enumerate}  \\item The two populations in NGC 6752 have different He abundances (they suggest $\\bigtriangleup Y\\sim0.05$). This may have come about through a merger of two proto-cluster clouds with differing chemical histories or through successive generations of stars (i.e. self-pollution). The He-rich material would also be N-rich. The He-rich stars would then evolve to populate the blue end of the HB -- and not ascend the AGB -- leaving only CN-normal stars to evolve to the AGB.  \\item Mixing in about half of the RGB stars pollutes their surfaces (increasing N) and also increases mass-loss rates, again leading to two separate mass groups on the HB. As before, the CN-strong, low mass group does not ascend the AGB.  \\end{enumerate}  A constraint on the first explanation (for NGC 6752) is that about \\textit{half the mass of the cluster} must be polluted, as the number of CN-strong and CN-normal stars is (usually) roughly equal. As Norris et al. state, this is not a serious problem for the merger scenario, as the merging clouds/proto-clusters may very well have had similar masses. However, due to the constancy of Fe group elements, the chemical histories of the two clouds/proto-clusters would have to have been identical with respect to these heavy elements. This is more difficult to explain since we require differing chemical histories for the light elements.  In light of recent observations on the MS and SGBs of some clusters, the second explanation by Norris et al. may require some clarification.  As N appears to have a preformation source (as evidenced by MS observations), the extra mixing is not required (although it does still exist). However, the general suggestion that the differing compositions may affect mass-loss rates and lead to different mass populations on the HB may be a valid one. We are calculating stellar models that include this effect.  An important point apparent in the current results is that there appears to be a variation in the number of CN-strong stars that `disappear' -- some GCs show a total lack of CN-strong stars on their AGBs, whilst some show a mix of CN-strong and CN-weak, although they are still dominated by CN-weak stars.  Theories such as those of \\cite{NCF81} will have to account for this point also.  Finally, the substantial abundance differences between the GC RGBs and AGBs reported here may reveal other clues to the GC abundance anomaly problems (i.e. those of the heavier p-capture products - see \\citealt{SIK00} for a discussion), and even the GC formation problem. The rest of the results from this study will be published soon, including a more in-depth analysis and discussion."
    },
    "1004/1004.2413_arXiv.txt": {
        "abstract": "{} {Molecules that trace the high-density regions of the interstellar medium may be used to evaluate the changing physical and chemical environment during the ongoing nuclear activity in (Ultra-)Luminous Infrared Galaxies.}  {The changing ratios of the HCN\\,(1$-$0), HNC\\,(1$-$0), \\HCOP\\,(1$-$0), CN\\,(1$-$0) and CN\\,(2$-$1), and CS\\,(3$-$2) transitions were compared with the HCN\\,(1$-$0)/CO\\,(1$-$0) ratio, which is proposed to represent the progression time scale of the starburst. These diagnostic diagrams were interpreted using the results of theoretical modeling using a large physical and chemical network to describe the state of the nuclear ISM in the evolving galaxies.}  {Systematic changes are seen in the line ratios as the sources evolve from early stage for the nuclear starburst (ULIRGs) to later stages. These changes result from changing environmental conditions and particularly from the lowering of the average density of the medium. A temperature rise due to mechanical heating of the medium by feedback explains the lowering of the ratios at later evolutionary stages.  Infrared pumping may affect the CN and HNC line ratios during early evolutionary stages.}  {Molecular transitions display a behavior that relates to changes of the environment during an evolving nuclear starburst. Molecular properties may be used to designate the evolutionary stage of the nuclear starburst. The HCN(1$-$0)/CO(1$-$0) and HCO$^{+}$(1$-$0)/HCN(1$-$0) ratios serve as indicators of the time evolution of the outburst.  } ",
        "introduction": "The (Ultra-) Luminous Infrared Galaxies (ULIRGs) are powered by a (circum-)nuclear starburst and/or an Active Galactic Nucleus (AGN). ULIRGs belong to a much larger population of Sub-Millimeter Galaxies (SMGs) that peaks at redshift 2$-$3 and extends to z = $\\sim$6.  The relatively short episodes of intense nuclear activity in ULIRGs and SMGs have likely been triggered by galaxy mergers or collisions, which produce the most luminous galaxies in the universe.  Among other spectral observations, the emissions from the high-density molecular material in the nuclear region of external galaxies allow a clear view of the ISM and are a preferred tool for diagnosing the physical and chemical environment of the starforming and AGN-excited ISM. A strong relation has been found between the far-infrared luminosity and the luminosity of molecular emissions, that are indicators of the high-density (\\numd$\\ga$10$^{4-6}$\\,\\percc) component of the interstellar medium \\citep{GaoS2004a,CurranAB2000, BaanEA2008}. Since these integrated emissions predominantly arise from the nuclear region, they are closely associated with the nuclear starburst/AGN environment and the production of the FIR and sub-millimeter luminosities.  Each molecular species responds differently to the changing physical and chemical environment of the nuclear region. The collective behavior of molecular tracers may thus be used to characterize the evolution of the nuclear environment during a nuclear activity.  In this paper, we consider evolutionary changes of the ISM as seen by the molecular tracer emissions and resulting from the ongoing star-formation, the AGN activity, and any feedback processes. Diagnostic interpretation of the behavior of molecular line ratios is based on theoretical modeling using a large physical and chemical network to describe the state of the nuclear ISM in the evolving galaxies.  This study consider mostly single molecular transitions for the multi-molecular modeling.  The inclusion of multiple higher transitions will augment the diagnostic accuracy of the modeling analysis. Modeling of the observed ground state transitions already predicts the behavior of higher transitions and their ratios and will be used in future work to incorporate newly obtained observational data and data expected from Herschel and ALMA. ",
        "conclusions": "\\label{sec:phenomenon}  The evolutionary stages of a nuclear starburst represent a well-determined sequence of events that would affect the ISM in a complex but predictable manner.  As a result, the characteristics of the ISM and the subsequent star-formation process will change systematically during the course of the outburst and ultimately leads to the termination of the process. The dominant effects describing the nuclear ISM are the steady depletion of the high-density molecular component, the decreasing average density of the medium, the increasing temperature resulting from the heating due to feedback.  A  simple  model has  been  discussed  that  describes the  decreasing high-density component of the ISM during an outburst. It has also been argued that HCN is the molecular species that is least affected by the chemical  and  physical changes  in  the ISM  and  best  serves as  an indicator of the  high-density component in the ISM.  For this reason, all molecular  variations in this  study are measured relative  to the strength  of   the  HCN\\,(1$-$0).  Molecular species will react differently to  the changing environment of the nuclear  ISM and these molecules will  reveal different aspects of the changing physical and chemical environment. In this example,  the  decreasing HCN\\,(1$-$0)/CO\\,(1$-$0)   ratio   serves  as   an   initial  measure   of evolutionary  time for the  starburst. Systematic changes may be seen for all characteristic line ratios with evolutionary time. The observed ratios with HCN\\,(1$-$0) for the two transitions with lower critical densities, CN\\,(1$-$0) and \\HCOP\\,(1$-$0), display a systematic increase with evolutionary time of the starburst. Furthermore, the CN\\,(2$-$1)/CN\\,(1$-$0) ratio for starbursts and LIRGs (tentatively) decreases systematically with evolutionary time. All these ratios may be used as indicators of evolutionary time of the starburst.  The two other transitions with higher critical densities, HNC\\,(1$-$0) and CS\\,(3$-$2) display a more complicated picture because it appears that additional physical processes affect the ratios, particularly for the highest-luminosity ULRGs. Both transition display weak evolutionary variation with evolutionary age. While FIR pumping may account for the anomalous HNC/HCN ratio in these ULIRGs, a pumping mechanism for CS has yet to be identified. Besides HNC and CS, CN also displays anomalous excitation during the early ULIRG stage. Here the CN\\,(1$-$0) line is enhanced and the CN\\,(2$-$1) line is reduced relative to HCN\\,(1$-$0), which may reduce the CN\\,(2$-$1)/CN\\,(1$-$0) ratio during early stages of evolution.  In addition to the main group of sources that display a systematic increase in the observed ratios with a lowering of the HCN/CO ratio, a group of evolved sources has been identified that would display a different dependence on evolutionary time. Besides the systematic lowering of density in the PDR-dominated environment (from 10$^{5.5}$ to 10$^{4}$ cm$^{-3}$), the feedback from SNe and SNRs during later stages of evolution provides mechanical heating leads to lower ratios for all line ratio with respect to HCN\\,(1$-$0) \\citep{LoenenSBM2008}.  The observed variation of characteristic line ratios may be used to diagnose the evolutionary state of the nuclear activity in galaxies. This would allow to establish an evolutionary sequence for ULIRGs and (future) SMGs and a way to understand and diagnose their physical processes.  The systematic changes of the characteristic molecular line ratios already provide a first view of evolutionary sequence for the nuclear activity. While the relation between the HCN\\,(1$-$0)/CO\\,(1$-$0) line ratio and evolutionary time only provides relative time during the evolution of the starburst, further modeling of the observed changes would provide a more accurate translation to evolutionary time. In addition, sequential modeling to reconstruct the changing nuclear environment will provide a better understanding and identification of the different physical and chemical processes that determine the evolution of the molecular environment during a nuclear outburst. Additional observational data of multiple molecular transitions will further improve the understanding of the complex myriad of competing processes."
    },
    "1004/1004.0291_arXiv.txt": {
        "abstract": "The detection of neutrinos from massive stellar collapses can teach us a lot not only about source objects but also about microphysics working deep inside them. In this study we discuss quantitatively the possibility to extract information on the properties of dense and hot hadronic matter from neutrino signals coming out of black-hole-forming collapses of non-rotational massive stars. Based on our detailed numerical simulations we evaluate the event numbers for SuperKamiokande with neutrino oscillations being fully taken into account. We demonstrate that the event numbers from a Galactic event are large enough not only to detect it but also to distinguish one hadronic equation of state from another by our statistical method assuming the same progenitor model and non-rotation. This means that the massive stellar collapse can be a unique probe into hadron physics and will be a promising target of the nascent neutrino astronomy. ",
        "introduction": "\\label{intro}  One of the important roles of astrophysics is to explore physics under extreme conditions that are difficult to realize in terrestrial experiments. In this sense, hadron physics at supra-nuclear densities and non-zero temperatures (e.g., hyperon appearance, quark deconfinement and so on) is a natural target of astrophysics and the gravitational collapse of massive stars at the end of their lives will set the stage \\cite{sumi09, self08a}. In particular, black-hole-forming collapses expected for very massive stars with a mass larger than $\\sim \\!\\! 30$ solar masses ($M_\\odot$) will be the most promising site. Although such an event has not been observed yet, a black hole candidate with an estimated mass of 24-33$M_{\\odot}$ was discovered \\cite{prest07} and this might be a remnant of the collapse of such a massive star. Recently, a regular monitoring of $\\sim \\!\\! 10^6$ supergiants within a distance of 10~Mpc is proposed to detect their silent disappearances \\cite{kocha08} and the black-hole-forming collapse, which would be invisible optically, might be observed that way.  In our previous studies \\cite{sumi06, sumi07, sumi08}, we showed that the event, which we refer to as the ``failed supernova'' hereafter, is as bright in neutrino emissions as ordinary core-collapse supernovae. We showed also that its time evolutions of luminosities and spectra are qualitatively different from those of the ordinary supernova explosion and  the ensuing proto-neutron star cooling \\cite{totani98}, which may lead to the delayed black hole formation for some reasons \\cite{refb, refb2}. Our numerical data were adopted as a reliable basis to predict the relic neutrino background from stellar collapses \\cite{luna09}. More importantly, however, we also demonstrated by employing different hadronic equations of state (EOS) that the duration of neutrino emissions from the failed supernova is sensitive to the stiffness of EOS at supra-nuclear densities and, therefore, that the observation of neutrinos from such an event will provide us with valuable information on the properties of dense and hot hadronic matter as well as on the maximum mass of proto-neutron stars.  Although this approach is simple and robust, valid irrespective of neutrino oscillations, it can not distinguish EOS's with a similar duration of neutrino emissions: a soft nucleonic EOS and a hyperonic EOS, for example. In this study, we attempt to break this degeneracy by analyzing more in detail the time variation of neutrino numbers observed at a terrestrial detector, which we refer to as the ``light curve'' hereafter. While we have studied the detection of failed supernova neutrinos taking fully into account the neutrino oscillation and its parameter dependence so far \\cite{self08b}, we innovate a new method here by employing the Kolmogolov-Smirnov (KS) test, which is free from the ambiguity of the distance to the progenitor. We adopt the results of our detailed numerical simulations and evaluate the neutrino event number for SuperKamiokande (SK) as a representative of currently operating neutrino detectors. This is the first ever serious self-contained attempt to demonstrate that for Galactic events it is indeed possible to break the degeneracy for hadronic EOS's by the statistical analysis.  We arrange this paper as follows. A brief review of the neutrino detection and a description of the newly proposed statistical method are given in Sec.~\\ref{setup}. The main results of our study are reported in Sec.~\\ref{result}. In Sec.~\\ref{disc}, we mention the possible uncertainties and observational issues. Sec.~\\ref{sum} is devoted to a summary. ",
        "conclusions": "\\label{disc}  Certainly, there are some issues remaining unaddressed. First of all, investigating other EOS's constructed by different manners is important future work. In particular, hyperon mixing in neutron star matter has been studied by the microscopic approach \\cite{refa, refa2, refa3} and it is desired to be applied to EOS tables including finite temperature. If neutrinos have masses of $\\sim$~eV, the arrival time depends on the neutrino energy and the difference becomes $\\sim \\!\\! 0.01$~sec between the neutrinos with an energy of $E \\sim 10$~MeV and those with $\\gtrsim \\!\\! 30$~MeV \\cite{beacom00} for Galactic sources. Recently, the core collapse simulations for the same progenitor model are done utilizing a spherically symmetric $\\nu$-radiation-hydrodynamics code by another grope \\cite{fischer08}. While they conclude that their results are in qualitative agreement with our study, there are some quantitative differences. For instance, in the case with LS180-EOS, they found that the interval time between the bounce and black hole formation is 0.436~sec, which is $\\sim \\!\\! 20$\\% shorter than our result. Further numerical investigation is important.  Multi-dimensional hydrodynamical effects such as rotations and magnetic fields, which are not taken into account in our model, may also affect the signals. In the context of supernovae explosions, Ref.~\\cite{marek09} compared the neutrino signal by an one-dimensional simulation with those by two-dimensional simulations (with and without rotation) for the collapse of a $15M_\\odot$ star adopting LS180-EOS (see Figs.~1 and 14 of this paper). According to this paper, when the initial rotation of progenitor core is a constant angular frequency of 0.5~rad~s$^{-1}$, the stellar rotation does not affect the luminosity and mean energy of the emitted neutrinos very much. Incidentally, the neutrino emission after the collapse to a black hole is estimated to be negligible for low-angular momentum cases \\cite{fryer09}.  Uncertainties with progenitor structures are another concern. While a certain class of progenitors provides similarly short neutrino burst, it has been shown that the density profile of outer layer may affect the duration \\cite{sumi08}. Especially, in the case with lower matter density outside the core \\cite{woosley02}, the interval time between the bounce and black hole formation becomes longer as 1.477~sec even for the model with $40M_\\odot$ and LS180-EOS \\cite{fischer08}. However, if a neutrino event is actually detected, we will be able to determine the direction of the progenitor to some extent by the neutrino detection itself \\cite{ando02} and the progenitor is highly likely to be identified or at least constrained by earlier records of optical observations as in the case of SN1987A. Then we can study the progenitor dependence much more efficiently.  As discussed in Refs.~\\cite{kocha08, self08b}, the event rate of the black-hole-forming failed supernovae is estimated to be somewhat low $\\sim \\!\\! 0.008$/yr. This problem may be circumvented, however, by deploying a large detector such as Deep-TITAND with a fiducial volume of 5~Mton \\cite{kist08} proposed currently. With this large facility, the event rate goes up to $\\sim \\!\\! 0.02$/yr since we should be able to detect $\\sim \\!\\! 400$ neutrinos for the first 0.5~sec of a failed supernova in galaxies as far away as the Andromeda galaxy (M31) at 780~kpc from us. This is large enough to distinguish the EOS's we studied in this letter for the inverted mass hierarchy with $\\sin^2\\theta_{13}=10^{-2}$. Note that the cumulative core-collapse supernova rate within this range is 10\\% of that within 10~Mpc \\cite{ando05}, where a regular monitoring of $\\sim \\!\\! 10^6$ supergiants is proposed in Ref.~\\cite{kocha08}."
    },
    "1004/1004.1542_arXiv.txt": {
        "abstract": "{We report the discovery of WASP-26b, a moderately over-sized Jupiter-mass exoplanet transiting its 11.3-magnitude early-G-type host star (1SWASP\\,J001824.70-151602.3; TYC 5839-876-1) every 2.7566~days. A simultaneous fit to transit photometry and radial-velocity measurements yields a planetary mass of $1.02 \\pm 0.03$~$M_{\\rm Jup}$ and radius of $1.32 \\pm 0.08$~$R_{\\rm Jup}$. The host star, WASP-26, has a mass of $1.12 \\pm 0.03$~$M_{\\sun}$ and a radius of $1.34 \\pm 0.06$~$R_{\\sun}$ and is in a visual double with a fainter K-type star. The two stars are at least a common-proper motion pair with a common distance of around $250 \\pm 15$~pc and an age of $6 \\pm 2$~Gy.} ",
        "introduction": "Most of the known exoplanets have been discovered using the radial velocity technique \\citep{1995Natur.378..355M}. However, in recent years an increasing number have been discovered using the transit technique, via ground-based and space-based survey projects. Transiting exoplanets allow parameters such as the mass, radius, and density to be accurately determined, as well as their atmospheric properties to be studied during their transits and occultations \\citep{2005ApJ...626..523C,2009MNRAS.394..272S,2009IAUS..253...99W}.  The SuperWASP project has a robotic observatory on La Palma in the Canary Islands and another in Sutherland in South Africa. The wide angle survey is designed to find planets around relatively bright stars in the $V$-magnitude range $9\\sim13$. A detailed description is given in \\citet{2006PASP..118.1407P}.  In this paper we report the discovery of WASP-26b, a Jupiter-mass planet in orbit around its $V = 11.3$~{mag.} host star 1SWASP\\,J001824.70-151602.3 in the constellation Cetus. We present the SuperWASP-South discovery photometry, together with follow-up optical photometry and radial velocity measurements. ",
        "conclusions": "WASP-26b is a moderately over-sized Jupiter-mass exoplanet transiting a G0 host star every 2.7566~days. A simultaneous fit to transit photometry and radial-velocity measurements gave a planetary mass of $1.02 \\pm 0.03$~$M_{\\rm Jup}$ and radius of $1.32 \\pm 0.08$~$R_{\\rm Jup}$. The mass and radius of WASP-26b place it within the group of bloated hot Jupiters. The incident flux received by WASP-26b is $1.8 \\times 10^9$~erg\\,s$^{-1}$\\,cm$^{-2}$, which is clearly within the theoretical `pM' planetary class proposed by \\citet{2008ApJ...678.1419F}.  The host star, WASP-26, and its K-type companion are at least a common-proper motion pair with a common distance of around $250 \\pm 15$~pc and an age of approximately $6 \\pm 2$~Gy. With a physical separation of at least 3800~AU, the companion star is unlikely to have any significant influence on the planetary system's dynamics \\citep{2007A&A...462..345D}. Spectroscopic confirmation of the stellar parameters of the K-type companion should further improve the distance and age determination of this system."
    },
    "1004/1004.4367_arXiv.txt": {
        "abstract": "A Bayesian approach to the determination of stellar distances from photometric and spectroscopic data is presented and tested both on pseudodata, designed to mimic data for stars observed by the RAVE survey, and on the real stars from the Geneva-Copenhagen survey. It is argued that this method is optimal in the sense that it brings to bear all available information and that its results are limited only by observational errors and the underlying physics of stars. The method simultaneously returns the metallicities, ages and masses of programme stars. Remarkably, the uncertainty in the output metallicity is typically 44 per cent smaller than the uncertainty in the input metallicity. ",
        "introduction": "Present attempts to reconcile models of the Galaxy with observations are hampered by a lack of full 6d phase space data for a large number of stars. While the majority of recent stellar surveys -- for example the Geneva-Copenhagen survey (\\citealt{Nordstrom}) and the Radial Velocity Experiment \\citep[RAVE;][]{Steinmetz:RAVE} -- provide sky positions, proper motions and often radial velocities, the persistent difficulty is the determination of distances, without which proper-motion measurements are of limited value. Since the fitting of models to observational constraints is arguably the method with the most potential for understanding the true nature of the Galaxy, the lack of these distance data is a key obstacle to our understanding of the Milky Way, and thus to our understanding of galaxies in general.  Several methods of distance determination are known. Perhaps most famously and successfully, the Hipparcos satellite measured trigonometric parallaxes for $\\sim\\! 10^5$ stars down to a \\textit{V}-band magnitude of around 12 (\\citealt{Hipparcos}). Trigonometric parallax is conceptually the most fundamental distance measurement technique for stars; however the Hipparcos measurements were accurate only out to a distance of around 200\\,pc. Within the next decade, the Gaia mission (\\citealt{Gaia}) should return parallaxes for around $10^9$ stars, dramatically increasing the size of the available data set for investigation; however in the meantime trigonometric parallaxes to a vast number of otherwise well-observed stars are lacking. Furthermore, even after the Gaia mission is complete, other estimation methods will remain vital for more distant stars.  Two other trigonometric methods of distance measurement are known: so-called `Galactic parallax' (\\citealt{Galpar,Galpar2}) and `geometrodynamical' techniques (\\citealt{LyndenBelldistance}).  However, both these methods require strong constraints on the orbits of target stars, so they are restricted to a small subset of all stars.  The next most reliable distance determination technique is the use of so-called `standard candles' -- stars with particular properties such as RR~Lyrae and Cepheid variables, whose luminosities are a reasonably sharply-defined function of observables such as period and colour. While the accuracies attainable for such stars may be high, standard candles form an extremely small subset of any  stellar population. Hence although they can be used to give a broad-brush picture of Galactic structure (see \\citealt{Gautschy}), the technique cannot be applied to the majority of stars of interest.  The other major technique for distance estimation is the determination of `photometric distances'. This involves deducing the luminosity of a star from its colours and perhaps metallicity, allowing its distance to be inferred from its apparent magnitude. \\cite{Juric_cut} have applied this technique to millions of stars measured by the Sloan Digital Sky Survey (hereafter SDSS). The technique works well for SDSS stars because they are apparently faint and therefore overwhelmingly likely to lie on the main sequence, and the luminosity of such a star is a well-defined function of colour and metallicity. However for samples at brighter apparent magnitudes it is not safe to assume that stars lie on the main sequence, and a more sophisticated technique is required.  An example of such a sample is the RAVE survey, which is in many ways complementary to the SDSS sample. It covers a nominal $I$-band magnitude range of $9 < I < 12$ and is expected to provide spectrophotometric data on up to a million stars by 2011 (\\citealt{RaveDR1_cut}). On account of its magnitude range, the RAVE sample contains a non-negligible number of giants. Such a heterogeneous sample therefore requires a more sophisticated approach to parameter determination than a direct colour-to-absolute magnitude mapping. Furthermore it would be much more satisfactory to obtain a methodology that can give a definitive distance estimation (and corresponding individual uncertainty) for specific stars rather than only for sizeable statistical samples.  Recently \\cite{Breddels} estimated photometric distances to $\\sim\\! 25\\,000$ stars in the RAVE survey. These distances were obtained by repeatedly seeking the stellar model that provides the best fit to first the data and then $5\\,000$ pseudo-data obtained by scattering the observables of the first stellar model by the observational errors. The logical justification of this procedure is unclear. In this paper we argue that the principles of Bayesian inference lead unambiguously to a different procedure that has much in common with the procedures for determining stellar ages introduced by \\cite{PontEyer} and \\cite{Jorgensen}. In fact, our procedure yields not only a distance to each star but also estimates of its mass, age and metallicity. We argue that our procedure is optimal in the sense that it exploits the available data in their entirety -- including information about the survey in question. From these facts it constructs a pdf for each star in model space, so estimated distances, masses, ages and metallicities are accompanied by error estimates.  The method is applicable to any survey that provides more than one non-degenerate observable for each star and is limited only by degeneracies in the underlying stellar physics and any inherent inaccuracies in the stellar models.  The paper is structured as follows: Section~\\ref{sec:theory} covers the theoretical basis of the method. Section~\\ref{sec:pseudodata} explores the application of the method to a fake data set. Section~\\ref{sec:GCS} then looks at the results of its application to the Geneva-Copenhagen sample. Section~\\ref{sec:previouswork} details the relationship of this paper to previous. Section~\\ref{sec:conclusions} presents a discussion of the results and highlights the potential for future applicability.  \\section[]{Theory} \\label{sec:theory}  Let $\\mathbf{y}$ represent an $n$-tuple of a star's observable quantities (e.g.\\ effective temperature $T_{\\rm{eff}}$, surface gravity $\\log g$, observed metallicity $\\mh_{\\rm{obs}}$, colours, apparent magnitudes, sky position) and let $\\mathbf{x}$ represent its six `intrinsic' variables (true initial metallicity $\\mh$, age $\\tau$, initial mass $\\cM$, distance $s$, and true sky position $(l,b)$) -- i.e.\\ \\begin{equation} \\mathbf{x} = \\left( \\mh, \\tau, \\cM, s, l, b \\right). \\end{equation} We will assume that a stellar model can be used to provide a direct mapping from $\\mathbf{x}$-space to $\\mathbf{y}$-space. For the purpose of concision, let $G(\\mathbf{w},\\bsigma)$ represent the multivariate Gaussian function \\begin{equation} G(\\mathbf{w},\\bsigma) \\equiv \\prod_{i=1}^n \\left( \\sigma_i \\sqrt{2\\pi}\\right)^{-1} \\exp \\left( - w_i^2 / 2 \\sigma_i^2 \\right) \\end{equation} for an $n$-tuple $\\mathbf{w}$.  In order to estimate a value and an uncertainty for each stellar parameter, the natural approach to take is probabilistic. We know a set of facts: the actual measured values of a star's observables ($\\bar{\\mathbf{y}}$), the quoted errors thereon ($\\bsigma_y$) and, subtly, the fact that the star is in the given sample (let this fact be denoted by $S$). Hence the logical distribution to consider is the pdf of a star's parameters $x_i$ given these three facts. If we can find this distribution, then we have all the parametric information that can logically be inferred from the known facts -- most importantly, we can give an expectation value for each parameter and an estimated error on this value.  Explicitly, for each component $x_i$ of $\\mathbf{x}$ we seek the three moments $\\mathcal{I}_{ik}$ defined by \\begin{equation} \\label{eq:firstmoments} \\mathcal{I}_{ik} = \\int \\! x_i^k \\,\\, p(\\mathbf{x} | \\bar{\\mathbf{y}}, \\bsigma_y , S ) \\, \\dd^6 \\mathbf{x} , \\end{equation} where $k \\in \\{0,1,2\\}$. (The $k=0$ moment should properly be unity, but it is useful if the pdf is left unnormalized.) From these moments we can then infer the expectation and variance of each stellar parameter.  In order to express the pdf in terms of known distributions, it is useful to consider the full pdf \\begin{eqnarray} \\label{eq:fullpdf} p(\\mathbf{x}, \\bar{\\mathbf{y}}, \\bsigma_y , S) &\\!\\! =& \\!\\! p(\\mathbf{x} | \\bar{\\mathbf{y}}, \\bsigma_y , S) \\, p(\\bar{\\mathbf{y}}, \\bsigma_y , S) \\\\ &\\!\\! =& \\!\\! P(S | \\bar{\\mathbf{y}}, \\mathbf{x}, \\bsigma_y) \\, p(\\bar{\\mathbf{y}} | \\mathbf{x}, \\bsigma_y ) \\, p(\\bsigma_y | \\mathbf{x}) \\, p(\\mathbf{x}) .\\nonumber \\end{eqnarray} Hence we can expand the pdf in a form analogous to Bayes' theorem to give \\begin{equation} p(\\mathbf{x} | \\bar{\\mathbf{y}}, \\bsigma_y , S) = \\frac{ P(S | \\bar{\\mathbf{y}}, \\mathbf{x}, \\bsigma_y) \\, p(\\bar{\\mathbf{y}} | \\mathbf{x}, \\bsigma_y ) \\, p(\\bsigma_y | \\mathbf{x}) \\, p(\\mathbf{x}) }{ p(\\bar{\\mathbf{y}}, \\bsigma_y , S) } , \\end{equation} which can be simplified to \\begin{equation} \\label{eq:firsteq} p(\\mathbf{x} | \\bar{\\mathbf{y}}, \\bsigma_y , S)  \\propto P(S | \\bar{\\mathbf{y}}, \\mathbf{x}, \\bsigma_y) \\, p(\\bar{\\mathbf{y}} | \\mathbf{x}, \\bsigma_y ) \\, p(\\bsigma_y | \\mathbf{x}) \\, p(\\mathbf{x}) , \\end{equation} wherein factors independent of $\\mathbf{x}$ have been neglected as irrelevant to the problem we aim to solve, since we will require only ratios of different $\\mathcal{I}_{ik}$'s.  Each of the factors in equation~(\\ref{eq:firsteq}) can now be analysed independently: \\begin{enumerate} \\item $\\selfn$, which we term the \\textit{selection function}, expresses the probability of a star being in the sample given its inherent values and the measurement errors. This factor therefore reflects our selection criteria, whether they be, for instance, a magnitude cut or any cut on errors used to `clean' a sample.  \\item For many types of observation the \\textit{likelihood} $p(\\bar{\\mathbf{y}} | \\mathbf{x}, \\bsigma_y )$ can be approximated by a Gaussian $G(\\bar{\\mathbf{y}} - \\mathbf{y}(\\mathbf{x}), \\bsigma_y)$. More generally it takes a functional form determined purely by the measuring instrument, in which the different components of $\\bar{\\mathbf{y}}$ may or may not be independent.  \\item $p(\\bsigma_y | \\mathbf{x})$ is the probability of the quoted errors in $\\mathbf{y}$ given the  object's underlying characteristics $\\mathbf{x}$.  \\item $p(\\mathbf{x})$ is our \\textit{prior}. This will describe as many stellar populations as we wish to take into account, in terms of an initial mass function (IMF) and a spatial distribution. It describes the relative abundance of stars of various types, without regard to their observability.  \\end{enumerate}  The uncertainties on the observed values of $(l,b)$ are assumed to be sufficiently small for us to regard the corresponding pdfs as delta functions and drop the integrals over $l$ and $b$, leaving a 4-dimensional integral in $\\mathbf{x}$. However, when this is done, any locational factors in the prior must gain a multiplicative factor of $s^2$ in order to take account of the conical shape of the volume surveyed.  In this work we assume a uniform distribution for $p(\\bsigma_y | \\mathbf{x})$; although in reality this will not be the case (for instance, more distant stars will tend to have higher errors), the dependence of $\\bsigma_y$ on $\\mathbf{x}$ will generally be sufficiently weak not to affect our results greatly. In this case, equation~(\\ref{eq:firsteq}) simplifies to \\begin{equation}\\label{eq:final} p(\\mathbf{x} | \\bar{\\mathbf{y}}, \\bsigma_y , S) \\propto P(S | \\bar{\\mathbf{y}}, \\mathbf{x}, \\bsigma_y) \\, p(\\bar{\\mathbf{y}} | \\mathbf{x}, \\bsigma_y ) \\, p(\\mathbf{x}) . \\end{equation} This can be substituted into the simplified form of equation~(\\ref{eq:firstmoments}), \\begin{equation} \\mathcal{I}_{ik} = \\int \\! x_i^k \\, p(\\mathbf{x} | \\bar{\\mathbf{y}}, \\bsigma_y , S ) \\, \\dd^4 \\mathbf{x} , \\end{equation} giving a set of three integrals that together will give us an estimated parameter value for each star ($\\left< x_i \\right> = \\mathcal{I}_{i1} /\\mathcal{I}_{i0}$) and an uncertainty thereon ($\\sigma_i = \\sqrt{\\left(\\mathcal{I}_{i2} /\\mathcal{I}_{i0}\\right) - \\left< x_i \\right>^2}$), taking into account all of the known information. Thus in a single pass we can infer all these values for each star's metallicity, age, mass and distance.  \\subsection{Selection function}  The nature of the selection function $P(S | \\bar\\mathbf{y}, \\mathbf{x}, \\bsigma_y)$ bears comment, as it is by no means trivial. We must distinguish carefully what is meant by the symbol $\\bar\\mathbf{y}$ in the formalism above: it refers to the actual observed parameters of the star, but only those we actually use in the analysis. There may be other parameters that have been observed and enter into the selection function. If this is the case, they can be expressed only as probabilities dependent on $\\mathbf{x}$, not on $\\bar\\mathbf{y}$. Therefore it becomes important to split the selection function into two parts, one dependent on $\\bar\\mathbf{y}$ and the other on $\\mathbf{x}$: \\begin{equation} P(S | \\bar\\mathbf{y}, \\mathbf{x}, \\bsigma_y) = \\psi(\\bar\\mathbf{y}, \\bsigma_y) \\,\\phi(\\mathbf{x}). \\end{equation} In this product $\\psi$ will generally be the dominant factor for two main reasons:  \\begin{itemize} \\item[(i)] If the sample is based on an input catalogue, it will frequently be a selection on observables that are not then reobserved -- for example, the Geneva-Copenhagen survey (\\citealt{Nordstrom}) was selected on the basis of the Str\\\"omgren photometry of \\cite{Olsen1,Olsen2,Olsen3,Olsen4}, which is simply transcribed to the Geneva-Copenhagen catalogue.  \\item[(ii)] The fundamental limitations on stellar photometry also tend to make $\\psi$ dominant: saturation leads to a bright-end cut being imposed on the survey; at the faint end a cutoff is imposed to eliminate objects too faint for the detector's source extraction algorithm to distinguish true signals from noise. Both cuts are generally imposed on the basis of the observed magnitude and therefore are functions of the observed properties of a star \\textit{after scattering by observational errors}. That is, they are functions of $\\bar\\mathbf{y}$ rather than $\\mathbf{x}$.  \\end{itemize} However, there can be contributions of the nature of $\\phi(\\mathbf{x})$, and these are of much greater importance to our analysis. For example, we will later be interested in comparison of the results provided by our method on a real stellar sample with those of the Hipparcos satellite. In such cases, it is common (e.g.\\ \\citealt{Breddels}) to consider only those stars for which Hipparcos provides reasonably well-constrained distances, by restricting the sample to stars with fractional parallax errors below a certain threshold. It is important to recognize that this induces a bias in the sample: by considering only those stars with, say, $\\sigma_\\varpi / \\overline\\varpi < 20\\%$ (where $\\overline\\varpi$ represents the observed parallax and $\\sigma_\\varpi$ the error thereon), one will preferentially cut away stars with low parallaxes and hence large distances, leaving a sample biased towards small distances.  This effect can, however, be at least mitigated by a consideration of the r\\^{o}le of $\\phi(\\mathbf{x})$. For any value of $\\mathbf{x}$, and hence a `true' model distance $s$, it is possible to define a pdf for a star's observed parallax $\\overline\\varpi$ from a knowledge of $\\sigma_\\varpi$, assuming Gaussian errors: indeed \\begin{equation} p\\left( \\overline\\varpi | s \\right) = G \\left( \\overline\\varpi -1/s , \\sigma_\\varpi \\right) . \\end{equation} Therefore one can incorporate this bias into one's analysis in the form \\begin{equation} \\phi(\\mathbf{x}) \\propto \\int_{5 \\sigma_\\varpi}^\\infty \\!\\! G \\left( \\overline\\varpi - 1/s , \\sigma_\\varpi \\right) \\, \\dd \\overline\\varpi , \\end{equation} providing an element of balance in the analysis that maximum-likelihood techniques neglect.  If one can, as will generally be the case at least approximately, decompose the selection function into the form $\\psi(\\bar\\mathbf{y}, \\bsigma_y) \\,\\phi(\\mathbf{x})$, then since we are interested only in terms dependent on $\\mathbf{x}$, we can ignore any occurrence of $\\psi(\\bar\\mathbf{y}, \\bsigma_y)$. Hence our final formula for the moments of a star's pdf collapses to \\begin{equation}\\label{eq:integrate} \\mathcal{I}_{ik} = \\int \\! x_i^k \\, \\phi(\\mathbf{x}) \\,p(\\bar{\\mathbf{y}} | \\mathbf{x}, \\bsigma_y ) \\, p(\\mathbf{x}) \\, \\dd^4 \\mathbf{x} , \\end{equation} which can be readily calculated.  \\subsection{Implementation}  Since for many surveys the errors on apparent magnitude are very small, the integration of equation~(\\ref{eq:integrate}) need only cover a small range in distance for each set of $(\\mh,\\tau,\\cM)$ values. Consequently, after experimenting with both fixed-grid and Markov Chain Monte Carlo techniques, it was decided that the best compromise between speed and reliability was provided by a simple fixed-grid Newton-Cotes rectangle integration method, where one defines a grid of points in metallicity, the logarithm of age, and mass. For each point in this three-dimensional space one then integrates over a range in distance corresponding to an apparent magnitude spread of say ten times the magnitude error either side of the observed value -- any distances beyond this range will give a negligible contribution to the integral due to the Gaussian factor in the likelihood term $p(\\bar{\\mathbf{y}} | \\mathbf{x}, \\bsigma_y )$. For the purposes of calculating the integral it was found to be sufficient simply to multiply the integrand's value at each grid point by an effective volume determined by the distance between the given grid point and its nearest neighbours.  No two numbers completely characterize a general probability distribution. Medians, means and modes can all be used to characterize the centre of the distribution. The median is the most stable measure and generally to be preferred. Unfortunately, it is in general hard to calculate. The mode, which was favoured by \\cite{Jorgensen}, is susceptible to noise when the distribution is at all flat-topped. Therefore we favour the mean as a robust and readily calculated measure of the pdf. Similarly, we favour the variance as an estimate of the width of the pdf, rather than the more complex confidence intervals used by \\citeauthor{Jorgensen}.  Another advantage of using these two numbers to characterize the distribution is that their calculation, via the integration technique outlined above, avoids the thorny issue of interpolation on the isochrones by considering only the values of observables at tabulated grid points. Interpolation using isochrones is notoriously difficult, and thus we avoid it as far as is possible.  To summarize, the method we propose runs as follows: \\begin{enumerate} \\item Define a grid of points in $(\\mh,\\log\\tau,\\cM)$ space. We found a suitable grid spacing to be given by the metallicity values given in Table~\\ref{table:Zs} (although a greater range would do no harm -- one should include the lowest metallicity that is available from the isochrones); a spacing in $\\log\\tau$ of 0.025; and the mass spacing provided automatically by the Padova isochrones, which ranges from $10^{-6}$\\,M$_\\odot$ to 0.255\\,M$_\\odot$. \\item At each grid point, define a further set of points in distance $s$, covering the range \\begin{equation}\\label{eq:distance} s \\in [r(m-n\\sigma_m), r(m+n\\sigma_m)], \\end{equation} where $m$ is the apparent magnitude used for fitting to observations, $\\sigma_m$ is the observational error thereon, and $r(m)$ represents the distance that produces an apparent magnitude of $m$ given the metallicity, age and mass at the grid point in question. $n$ is some positive number chosen by the user -- we found $n=5$ to be ample for high precision.  \\item Assign to each point a 4-volume corresponding to a hypercuboid with boundary planes halfway to each of the point's nearest neighbours (at the edges of parameter space, take the length in the limiting dimension to be that between the final two data points).  \\item Perform the integral in equation~(\\ref{eq:integrate}) by finding the value of $\\phi(\\mathbf{x}) \\, p(\\bar{\\mathbf{y}} | \\mathbf{x}, \\bsigma_y ) \\, p(\\mathbf{x})$ at each grid point, multiplied by the grid point's associated 4-volume. We term this product of probability density and volume the `weight' of each grid point. For each of the first four dimensions of $\\mathbf{x}$, one can then define an array that keeps track of the sum of these weights multiplied by powers of each parameter $x_i$, and also keep a running total of the sum of the weights themselves, which we term $\\mathcal{I}_0$.  \\item Once one has run over the entire grid (or some subset thereof, since one may be able to disregard metallicities too far from the observed value), one can then calculate the moments $\\mathcal{I}_{ik}$ by dividing each running total by $\\mathcal{I}_0$. From the $\\mathcal{I}_{ik}$ one then has directly an expectation and uncertainty for each dimension of $\\mathbf{x}$.  \\item It is of course imperative to check that one has sampled $\\mathbf{x}$-space sufficiently finely to achieve sufficient precision in the work. For these purposes one should redo the analysis of a sample of stars using a finer sampling in each dimension and check that there is minimal alteration in the outputs for the new sampling. If time is not a concern, one would ideally do this for every star, refining the grid iteratively to a point where there is minimal change in the outputs. This may require interpolation if the isochrones cannot be provided with a fine enough sampling. \\end{enumerate}  \\subsection{Terminology\\label{sec:terminology}}  In what follows, we shall refer to the technique described above as the `Bayesian method', in order to distinguish it from what we shall call the `maximum-likelihood method'. The distinction here essentially concerns the nature of the prior and selection function. Maximum-likelihood techniques for fitting a model to data consist of finding the model (from some specified range) that maximizes the probability of the data in question -- i.e.\\ in our case one would seek to maximize the likelihood $p(\\bar{\\mathbf{y}} | \\mathbf{x}, \\bsigma_y )$. This is mathematically equivalent to setting both the selection function and the prior in the above formalism to uniform distributions, but conceptually it has an important difference from our Bayesian approach, which seeks to find the model with the maximum probability of being correct given the data. Maximum-likelihood techniques are popular due to their simplicity -- they involve no consideration of the nature of the prior and they generally do not take into account the selection function. Furthermore, from a conceptual viewpoint, the Bayesian technique is a more justified approach to the problem of model selection than maximum likelihood: since the data are given, it is logical to seek the model with the maximum probability of being correct given the data. For the case of determining stellar ages, \\cite{Jorgensen} give a nice example of the advantage of a Bayesian approach over maximum likelihood, while \\cite{PontEyer} provide an extensive discussion of the relative merits of a Bayesian technique.  One more point bears comment with regard to maximum likelihood (and na{\\\"i}ve Bayesian approaches) when the prior is set to be uniform in order to represent an unprejudiced starting assumption. For a uniform continuous pdf, the prior can only be defined to be uniform in a specific coordinate system: a transformation of those coordinates will in general not leave it so. Thus, for example, while a uniform prior in age may seem a reasonable starting point, it is difficult to justify such an assumption as opposed to a prior uniform in e.g.\\ the logarithm of age. Thus the assumption of a uniform prior in some space is nonetheless an assumption, and cannot be considered safer than making an explicit choice of prior after considering all the circumstances of the particular case. In general one hopes for the likelihood function to be sufficiently strongly peaked at some value to render the exact form of the prior unimportant for the posterior distribution; however in such complex, degenerate cases as those addressed here involving stellar evolution, such hopes are not always well-founded. For this reason a prior that is based on what we do know of stars in the Galaxy is an important factor in any calculation. ",
        "conclusions": "\\label{sec:conclusions}  We have presented a Bayesian technique for determining stellar parameters from photometric and spectroscopic data. It has been demonstrated that the technique outperforms maximum-likelihood techniques, and the mathematical and physical basis of the system ensures that all available information can be exploited in the calculation. The resultant uncertainties, assuming a reasonable prior is chosen, are therefore a consequence purely of the underlying physics, and by virtue of this the technique is optimal -- given a set of data and some level of understanding of the underlying physics, one cannot do better.  Since the uncertainties derived from our technique simply reflect the physics, so long as one employs a reasonable prior, the technique will provide the most accurate possible estimates of the true uncertainties -- smaller estimated uncertainties could only come from an overly restrictive prior, undersampling the pdf in some manner or defining the uncertainties in a different manner, as for example in \\cite{Jorgensen}.  The problem of fitting stellar models to observations is an ideal setting for the application of Bayesian probability theory. We have a wealth of advance knowledge of the underlying physics with which to construct a prior, and the complex relationship between input parameters and observables for stellar models renders a simple maximum-likelihood approach suboptimal for providing reliable parameter determinations. There \\textit{are} favoured regions of observable space; our approach exploits this by virtue of the input prior. The Bayesian method also enables a valuable synthesis of different areas of astrophysical research: large-scale analysis and mapping of stellar populations in the Galaxy feeds us information for the prior, whilst the specifications of observations give us the form of the likelihood and selection functions.  As with all photometric techniques, obscuration, which we have neglected, is an important issue. It would be simple to include reddening and extinction models, but inevitably the results would then be vulnerable to weaknesses in the adopted models. When a large body of trigonometric parallaxes for relatively distant stars is at hand, it will be possible to adjust such models by making photometric distances compatible with trigonometric parallaxes. Such work will be a key project to be undertaken with the Gaia Catalogue.  A simplification we have made is to consider all metallicity information to be encoded in the value of $\\mh$ for each star. The chemistry of stars is in reality more complex and both helium abundance and alpha enhancement play significant roles in stellar evolution. We have not explicitly addressed this issue because to do so we would require both further sets of isochrones and additional spectroscopic observables. However, such refinements will be possible within a decade; most importantly, their incorporation would represent no fundamental change to the method presented here.  While our method is strictly only applicable to single stars, dividing the GCS sample into probable binaries and single stars reveals that the parallaxes of binaries are under-estimated by at most $\\sim\\! 0.2\\sigma$ in the mean, and probably by only half this figure (Section~\\ref{sec:binary}). The effect of binarity on the estimation of other stellar parameters is more subtle, however \\cite{Jorgensen} provide a demonstration that its effect on Bayesian age determination is quite limited.  The effect of binaries could be incorporated into our formalism by including a probability of binarity in the prior, along with its associated effect on the various observables; this is the approach taken in the work of \\cite{PontEyer}. Taking full account of such contamination would require observational simulations to find the effect of different mass-ratio binaries on each observable.  In this work we have introduced the selection function $\\selfn$ as an explicit term in our calculations.  Its inclusion provides an intuitive and logically consistent means of taking account of all selection effects introduced by instrumental apparatus, observing strategies and later choices (such as cuts on errors to `clean' a sample). The presence of this factor in our calculations in the form of $\\phi(\\mathbf{x})$ is key, as it truly brings the model to the data rather than vice versa, and thus permits a theoretically much more satisfactory basis for all our calculations.  The distribution defined by the product of the selection function and the prior has another use. This distribution can be used as in equation~(\\ref{eq:fake}) to generate a sample of pseudodata, which can then be scattered by Gaussian `observational errors'. The distributions of the various observables of this pseudodata can then be compared with those of the true data set that one wishes to mimic, in order to assess the accuracy of the different factors one has included. Thus, if one is confident of the selection function (as will often be the case), one can adapt the prior until a good fit is found between the pseudodata and the real data. Once this has been achieved, this optimized prior can then be used in the analysis of the true data with a high degree of confidence.  In this work the form of the prior has been kept intentionally simple: we have dealt with a case in which the Galaxy is described by three stellar populations. Depending on the nature of the sample being analysed, the prior could take a large range of forms. An interesting case that we have not considered is when the data include kinematic quantities. Then the prior would include a full phase-space distribution function for each stellar population, and this enrichment of the prior information that is brought to bear on the data would further reduce the uncertainties in estimated distances, metallicities, etc. Indeed the power of the Bayesian method described in this paper lies in its generality; although it was applied only to spectrophotometric data in Sections~\\ref{sec:pseudodata} and \\ref{sec:GCS}, nothing restricts the inclusion of other observables.  Of course the correct form for the prior is unlikely to be known absolutely. One may however be able to specify it with a reasonable degree of certainty -- one might, for example, be able to give the prior in a parametric form with a certain pdf over the parameters that specify it (such as disc scale lengths). The integration described in Section~\\ref{sec:theory} can then also be performed over these parameters, providing slightly less assumption-dependent values for the method's outputs.  We have not explored the question of the joint probability distributions of different parameters conditional upon the data for each star. Using the same technique as we describe but marginalizing only over selected dimensions of $\\mathbf{x}$, it would be possible to examine such distributions and analyse covariances between different parameters for selected stars. This extension would be simple and one can envisage cases in which it could prove fruitful.  Another consideration regarding the pdf for each star is that of bimodality. We have characterized each star's pdf in distance by its mean and dispersion; however an extension of this approach would be to search through each star's pdf in distance to identify any instances of multimodality.  Some stars essentially present two distinct solutions for the given observables: one being a dwarf lying comparatively nearby, the other a more distant giant (as in the cases explored for the GCS in Section~\\ref{sec:GCS}). It could be fruitful to identify such degeneracies explicitly, and provide a mean and dispersion for each peak in the pdf.  Starting from the algorithm described in this paper, one could identify any points in the assigned grid whose probabilities (by equation~\\ref{eq:final}) lie above some chosen threshold value, and then seek to identify sizeable islands in this subset of the space by a friends-of-friends type algorithm (\\citealt{Huchra}). This would provide a method of partitioning the pdf into discrete distributions, each of which could then be assigned an average value and associated spread. Thus significant degeneracies in the pdf resulting from certain values of the observables could be identified and handled suitably. We hope to explore this avenue in a future paper.  We have shown that by using all of the available information one can constrain the metallicities of stars to greater accuracy than the observations themselves. It is reasonable to expect that this same effect can be achieved for other observable quantities such as the surface gravity, for which the nominal observational errors are often sizeable. In this manner one can use the small errors on certain observables (such as apparent magnitudes) to shrink conservative errors on less well-constrained values. So long as one has confidence in the stellar models used, this promises to be a powerful technique for survey analysis. It should be noted that this technique can only realistically be applied once -- `looping', by rerunning the analysis using the new values and errors, is prohibited -- since its effectiveness comes from the application of independent prior knowledge to data with mildly overquoted errors. The mathematical formalism does not permit the subsequent application of a prior that is not independent of the data.  One more point deserves mention. If one knows the distances to the stars in some specific sample (of a generic type that can be described by model isochrones), then one can in theory explore different forms for the prior in order to optimize one's estimation of their distances. This could then be fed back in to the algorithm in order to provide a more accurate estimation of distances to other stars.  We are currently using this method to obtain distances to in excess of $250\\,000$ RAVE stars and hope shortly to report the results of this work."
    },
    "1004/1004.1318_arXiv.txt": {
        "abstract": "{} {A strong decline of magnetic activity towards hotter stars occurs in the regime of mid/late A-type stars due to the vanishing of the outer convection zone. X-ray emission is an important diagnostic of studying possible activity in intermediate-mass stars.} {We present a Chandra observation of the exceptional planet bearing A5\\,V star HR~8799, more precisely classified as a kA5\\,hF0\\,mA5 star and search for intrinsic X-ray emission.} {We clearly detect HR~8799 at soft X-ray energies with the ACIS-S detector in a 10~ks exposure; minor X-ray brightness variability is present during the observation. The coronal plasma is described well by a model with a temperature of around 3\\,MK and an X-ray luminosity of about $L_{\\rm X} = 1.3 \\times 10^{28}$~erg/s in the 0.2\\,--\\,2.0~keV band, corresponding to an activity level of log~$L_{\\rm X}$/$L_{\\rm bol}\\approx -6.2$. Altogether, these findings point to a rather weakly active and given a RASS detection, long-term stable X-ray emitting star.} {The X-ray emission from HR~8799 resembles those of a late A/early F-type stars, in agreement with its classification from hydrogen lines and effective temperature determination and thus resolving the apparent discrepancy with the standard picture of magnetic activity that predicts mid A-type stars to be virtually X-ray dark.} ",
        "introduction": "The star \\object{HR~8799} (V\\,342 Peg, HD~218396) is usually classified as spectral type A5\\,V, located at a distance of 39.9~pc and has a visible magnitude of V\\,=\\,5.96~mag with B-V\\,=\\,0.23. It is an exceptional and so far unique star, it is a $\\lambda$~Bootis, $\\gamma$\\,Doradus and Vega-like at the same time. $\\gamma$\\,Doradus stars are a class of pulsating stars that typically reside at the cool edge of the Cepheid instability strip with spectral types mid-A to mid-F. HR~8799 was linked to the class of $\\lambda$~Bootis stars by \\cite{gray99}; $\\lambda$~Bootis stars are chemically peculiar, metal-poor (in particular the Fe-peak element,s but with the exception of C, N, O) A-type stars that do not exhibit ordered large-scale magnetic fields as observed in Ap stars \\citep{boh90}. \\cite{gray99} derived a metallicity of [M/H]\\,=\\,-0.47 for HR~8799 and determined its stellar parameters to $M\\,=\\,1.47\\pm 0.30$\\,M$_{\\odot}$, $R\\,=\\,1.34\\pm 0.05$\\,R$_{\\odot}$, $L\\,=4.92\\pm 0.41$\\,L$_{\\odot}$ ($1.9 \\times 10^{34}$~erg/s) and $T_{\\rm eff}=7430 \\pm 75$~K. They assigned it the spectral type kA5\\,hF0\\,mA5\\,v\\,$\\lambda$\\,Boo to indicate that HR~8799 is a mild $\\lambda$\\,Bootis star that exhibits the A5 standard in the CaII\\,K line as well as in other metallic lines, however the hydrogen profiles (and thus effective temperature) are in better agreement with the F0 standard. The $\\lambda$\\,Bootis phenomenon is thought to be caused by the accretion of metal depleted material, i.e. external processes, in line with the finding that HR~8799 is also a 'Vega-like' star that shows a far-IR excess flux at 60~$\\mu$m from circumstellar dust in IRAS data \\citep{sad86}. Additionally, there is strong evidence for a dusty debris disk from IR observations \\citep{zuck04}.  Recently, HR~8799 gained even more interest when multiple orbiting planets were detected by direct imaging \\citep{mar08}. The three planets are located at distances of several tens (24, 38, 68) AU and have masses around 10~$M_{\\rm J}$ each. The authors derived an age of about 60\\,Myr (30\\,--\\,150 Myr) from various lines of evidence, thus HR~8799 is likely relatively young. While stellar radiation, especially energetic ones like UV and X-rays, influences the evolution of circumstellar material and thus planet formation, the planets themselves are likely not important in the generation of X-ray emission in such stellar systems. Although its $Vsini=38 \\pm2$~km/s \\citep{kaye98} is only moderate for an A-type star and the exact orientation of the rotation axis is unknown, astrometric analysis of the detected planetary system suggests a rather low inclination and consequently fast rotating star, e.g. $V_{\\rm rot} \\gtrsim$~100 (200)~km/s for $i \\lesssim 20~(10)^{\\circ}$.   Hints for X-ray emission from HR~8799 came from RASS (ROSAT All-Sky-Survey) data via a cross-correlation search between X-ray positions and bright A-type stars \\citep{hue98}. They found that the position of the soft X-ray source 1RXS~J230729.0+210802 matches quite well with the one of HR~8799; however, the positional uncertainty and the low number of detected counts (10 photons) prevent a detailed investigation. On the other hand, \\cite{kaye98} attributed the chromospheric \\ion{Ca}{ii}~K flux of HR~8799 to be of basal origin. From the X-ray point of view HR~8799 is remarkable, since an A5 star should be virtually X-ray dark in the standard paradigm that is valid for main sequence stars. In this scheme X-ray emission is expected from magnetic activity in the regime of 'cool stars' (spectral type around A7 and later) and from wind-shocks in the regime of 'hot stars' (B2 and earlier); intermediate spectral types are virtually X-ray dark. Exceptions are found among young Herbig AeBe stars and peculiar Ap/Bp stars, where fossil magnetic fields are thought to play a mayor role in the generation of X-rays. This picture is well established by X-ray observations with {\\it Einstein} and ROSAT, showing that stellar activity develops in late A-type stars and increases strongly towards cooler stars, with activity levels being in the range of $\\log L_{\\rm X}$/$L_{\\rm bol}= -3~...~-7$ \\citep[e.g.][]{schmitt85, schmitt97}. A recent {\\it XMM-Newton} observation of the fast rotating A7 star Altair confirmed the presence of weak magnetic activity and coronal X-ray emission \\citep{rob09a}. The vanishing of magnetic activity in mid A-type stars is expected theoretically and confirmed with other activity indicators; e.g. FUSE observations of main sequence stars have shown the disappearance of chromospheric emission lines at effective temperatures above $T_{\\rm eff} \\approx 8200$\\,K \\citep{sim02}.  In this paper we present results from a {\\it Chandra} ACIS-S observation of HR~8799. Our paper is structured as follows: in Sect.\\,\\ref{ana} we describe the observation and data analysis, in Sect.\\,\\ref{res} we present our results and summarize our findings in Sect.\\,\\ref{sum}. ",
        "conclusions": "\\label{sum}  \\begin{enumerate} \\item We have detected soft X-ray emission from the extraordinary A5/F0 V star HR~8799. The derived X-ray luminosity of about $L_{\\rm X} = 1.3 \\times 10^{28}$~erg/s in the 0.2\\,--\\,2.0~keV band corresponds to an activity level of log~$L_{\\rm X}$/$L_{\\rm bol}\\approx -6.2$, pointing to a weakly active star.  \\item The ACIS data confirms the ROSAT detection, significantly improves the positional accuracy and additionally enables a deeper study of the X-ray properties of HR~8799. We find that its X-ray emitting coronal plasma has an average temperature of about 3.0~MK, typical for a star with a rather low activity level. Minor X-ray brightness variations are present in our observation; the accompanying spectral changes point to variable magnetic activity.  \\item Overall, the X-ray properties of HR~8799 resemble those of mildly active late A/early F stars, rather than those of mid/late A-type stars. Keeping in mind that the attributed spectral type for HR~8799 depends on the classification criterion, the one based on hydrogen lines (F0) most suitable reflects its X-ray properties.  \\end{enumerate}"
    },
    "1004/1004.0963_arXiv.txt": {
        "abstract": "We investigate the stellar populations of Ly$\\alpha$ emitters (LAEs) at $z=5.7$ and $6.6$ in a $0.65$ deg$^2$ sky of the Subaru/\\textit{XMM-Newton} Deep Survey (SXDS) Field, using deep images taken with Subaru/Suprime-Cam, UKIRT/WFCAM, and Spitzer/IRAC. We produce stacked multiband images at each redshift from $165$ ($z=5.7$) and $91$ ($z=6.6$) IRAC-undetected objects, to derive typical spectral energy distributions (SEDs) of $z\\sim 6$ -- $7$ LAEs for the first time. The stacked LAEs have as blue UV continua as the HST/WFC3 $z$-dropout galaxies of similar $M_{\\rm UV}$, with a spectral slope $\\beta \\sim -3$, but at the same time they have red UV-to-optical colors with detection in the $3.6\\mu$m band. Using SED fitting we find that the stacked LAEs have low stellar masses of $\\sim \\left( 3-10 \\right) \\times 10^{7}M_\\odot$, very young ages of $\\sim 1-3$ Myr, negligible dust extinction, and strong nebular emission from the ionized interstellar medium, although the $z=6.6$ object is fitted similarly well with high-mass models without nebular emission; inclusion of nebular emission reproduces the red UV-to-optical colors while keeping the UV colors sufficiently blue. We infer that typical LAEs at $z\\sim 6-7$ are building blocks of galaxies seen at lower redshifts. We find a tentative decrease in the Ly$\\alpha$ escape fraction from $z=5.7$ to $6.6$, which may imply an increase in the intergalactic medium neutral fraction. From the minimum contribution of nebular emission required to fit the observed SEDs, we place an upper limit on the escape fraction of ionizing photons to be $f_{\\rm esc}^{\\rm ion} \\sim 0.6$ at $z=5.7$ and $\\sim 0.9$ at $z=6.6$. We also compare the stellar populations of our LAEs with that of stacked HST/WFC3 $z$-dropout galaxies. ",
        "introduction": "\\label{sec:introduction}  Ly$\\alpha$ emitters (LAEs) are a galaxy population commonly seen at high redshift. The high number density of LAEs indicates their importance in the evolutionary history of galaxies. Since galaxies with strong Ly$\\alpha$ emission can be identified as LAEs even when their continuum emission is too faint to be detected in broadband imaging, LAEs provide an opportunity to probe low-mass galaxies with active star-formation. Some of them may be building blocks of more evolved galaxies.     Surveys of LAEs have been made primarily with narrow-band imaging to isolate Ly$\\alpha$ emission \\citep[e.g.,][]{hu1998,rhoads2000,iye2006}, and until now over a thousand LAEs have been photometrically selected and/or spectroscopically identified \\citep[e.g.,][]{hu2002,ouchi2003,malhotra2004,taniguchi2005,kashikawa2006,shimasaku2006,dawson2007,gronwall2007,murayama2007,ouchi2008,shioya2009,nilsson2009,guaita2009,hayes2010b}.  Studying the stellar population of LAEs is essential to understand their physical nature and to reveal the relationship between LAEs and other high-redshift galaxies such as Lyman break galaxies \\citep[e.g.,][]{steidel1996a}. At $z \\simeq 3$ -- $5$, much progress has been made recently, and it has been revealed from large samples with multiband photometry that most LAEs are small galaxies with masses $10^8$ -- $10^9 M_\\odot$ and young ages $<10^8$ yr \\citep[e.g.,][]{gawiser2007,nilsson2007,pirzkal2007}, while some have large stellar masses of $\\simeq 10^{10} M_\\odot$ \\citep[e.g.,][]{lai2008,finkelstein2009,ono2009}.            At higher redshifts, a few studies have reported on stellar populations of LAEs, but they are all based on a very small sample and consensus has not been reached. \\cite{lai2007} have studied stellar populations of three bright LAEs at $z = 5.7$ detected in rest-frame optical wavelengths, and found that they have high stellar masses of $\\sim 10^{10} M_\\odot$, as old ages as the Universe at their redshifts, and some amount of dust. \\cite{ouchi2009} have reported the discovery of a giant LAE at $z=6.595$, Himiko, with a Spitzer/Infrared Array Camera \\citep[IRAC;][]{fazio2004} counterpart, and estimated its stellar mass to be ($0.9$ -- $5.0$) $\\times 10^{10} M_\\odot$. It should be noted that these four LAEs are all detected in at least one IRAC band, which means that they are not typical LAEs but rare, massive-end objects among the overall LAE population. On the other hand, \\cite{pirzkal2007} have studied three faint LAEs at $5.2 \\leq z \\leq 5.8$ % found by Hubble Space Telescope (HST)/Advanced Camera for Survey (ACS) slitless spectroscopy in the Hubble Ultra Deep Field (HUDF), to show that they are all very young ($\\simeq$ several $\\times 10^6$ yr) with low masses ($\\simeq 10^6$ -- $10^8 M_\\odot$) and small dust extinctions ($A_V = 0$ -- $0.1$). \\cite{chary2005} have studied a gravitationally-lensed LAE at $z=6.56$, HCM6A, and reported that it has a stellar mass of $\\sim 10^9 M_\\odot$ and very young age ($\\sim 5$ Myr). Although these two studies may be picking up typical LAEs at each redshift, it is hard to draw robust conclusions from such small numbers of objects.            Recently, \\cite{ouchi2008} and \\cite{ouchi2009} have constructed the largest available sample of $z = 5.7$ and $6.6$ LAEs in an about $1$ deg$^2$ of the Subaru/\\textit{XMM-Newton} Deep Field (SXDF) from deep optical broadband and narrowband data. About $65${\\%} of the field is also covered by deep $JHK$ images taken with the United Kingdom Infrared Telescope (UKIRT)/ Wide Field Infrared Camera \\citep[WFCAM;][]{casali2007} from UKIRT Infrared Deep Sky Survey (UKIDSS) Ultra Deep Survey \\citep[UDS;][]{lawrence2007,warren2007}, and $3.6$ -- $8.0\\mu$m images taken with the Spitzer/IRAC from the Spitzer legacy survey of the UDS field (SpUDS; Spitzer Proposal ID {\\#}40021; PI: J. Dunlop). In this paper we use these wide-field, multiwaveband survey data to study the stellar population of typical LAEs at $z=5.7$ and 6.6 from stacking a large number of faint objects (165 at $z=5.7$ and 91 at $z=6.6$). At each redshift the stacked object is detected in several broadbands including the $3.6\\mu$m band, enabling us to place meaningful constraints on stellar population parameters. We also examine the stellar population of $z$-dropout galaxies recently discovered in the deep HST/Wide Field Camera 3 (WFC3) data of the HUDF, using the stacked spectral energy distribution (SED) obtained by \\cite{labbe2010} from $14$ objects of the \\cite{oesch2010} sample.         As noted by \\cite{schaerer2009}, careful SED fittings including nebular emission are required to accurately determine stellar population parameters for high-$z$ galaxies \\citep[see also][]{zackrisson2008,raiter2010,schaerer2010}. Nebular emission is produced in the interstellar medium (ISM) ionized by hot stars. Hence, the nebular emission of a galaxy increases as the fraction of ionizing photons absorbed by HI gas in the ISM increases, or equivalently, as the fraction of ionizing photons escaping from the galaxy, $f_{\\rm esc}^{\\rm ion}$, decreases. To see the influence of nebular emission on the determination of stellar population parameters, we examine SED models for two extreme $f_{\\rm esc}^{\\rm ion}$ values: $f_{\\rm esc}^{\\rm ion}=1$ where the SED is determined solely by stellar emission as most previous studies assumed, and $f_{\\rm esc}^{\\rm ion}=0$ where the contribution of nebular emission to the SED is largest. We find that the latter models generally give a very good fit to the observed SEDs, especially the $z=5.7$ one which is fit by the latter significantly better.        In addition, we treat $f_{\\rm esc}^{\\rm ion}$ as a free parameter in our SED fitting to place rough constraints on $f_{\\rm esc}^{\\rm ion}$, because redshifts of $z\\sim 6$ -- $7$ are close to the end of cosmic reionization \\citep[e.g.,][]{fan2006,becker2007} and the $f_{\\rm esc}^{\\rm ion}$ of galaxies is a key parameter which determines the ionizing photon budget. At redshifts up to $z \\sim 4$, the $f_{\\rm esc}^{\\rm ion}$ of galaxies has been estimated or constrained by a variety of methods using, e.g., FUV spectra and narrow-band images of star-forming galaxies \\citep[e.g.,][]{steidel2001,shapley2006,iwata2009} and the distribution of neutral hydrogen column densities in the after-glow spectra of long duration GRBs \\citep[e.g.,][]{chen2007,fynbo2009}, although the results do not agree well with each other. For example, \\cite{shapley2006} have derived $f_{\\rm esc}^{\\rm ion} \\geq 0.65$ for two $z \\sim 3$ LBGs with detected Lyman continua and $f_{\\rm esc}^{\\rm ion} = 0.14$ from a composite spectrum of $14$ $z \\sim 3$ LBGs\\footnote{They have assumed the intrinsic UV to Lyman continuum flux density ratio to be $3.0$, and corrected for IGM absorption.}, while \\cite{chen2007} have placed an upper limit of $f_{\\rm esc}^{\\rm ion} \\leq 0.075$ using a compiled sample of $28$ GRBs \\citep[see also][]{fynbo2009}.       At redshifts as high as $z \\sim 6$ -- $7$, most constraints are based on the UV luminosity density of galaxies. For example, \\cite{ouchi2009b} and \\cite{finkelstein2009f} inferred the production rate of ionizing photons in galaxies from the observed UV luminosity density of $z$-dropout galaxies and obtained lower limits to $f_{\\rm esc}^{\\rm ion}$ to keep the intergalactic medium (IGM) ionized. However, this method has a number of uncertainties such as the UV luminosity function of galaxies and the clumpiness of the IGM. Recently, \\cite{bouwens2010b} proposed that the very blue UV color of $z$-dropout galaxies they found may be due to weak nebular emission and hence high $f_{\\rm esc}^{\\rm ion}$, because strong nebular emission makes the UV color too red. This is interesting in relating the SED to $f_{\\rm esc}^{\\rm ion}$. However, as they already state, $f_{\\rm esc}^{\\rm ion}$ is not uniquely determined from the UV color but it also depends on the age of the stellar population. In this paper we show that inclusion of a UV-to-optical color reduces this degeneracy and provides constraints on $f_{\\rm esc}^{\\rm ion}$ for our LAEs.       The outline of this paper is as follows. After describing the imaging data used in this study in Section 2, we produce stacked images in Section 3. Our SED fitting method is presented in Section 4. In Section 5, we present and discuss our SED fitting results and constraints on $f_{\\rm esc}^{\\rm ion}$. A summary is given in Section 6. Throughout this paper, we use magnitudes in the AB system \\citep{oke1983} and assume a flat universe with ($\\Omega_m$, $\\Omega_\\Lambda$, $h$) $=$ ($0.3$, $0.7$, $0.7$). The age of the universe is $\\simeq 0.98$ Gyr at $z = 5.70$ and $\\simeq 0.82$ Gyr for $z = 6.56$.                    \\begin{figure*}./ \\begin{center} \\includegraphics[scale=1.0]{./faces.eps} \\end{center} \\caption[] { Median-stacked images of LAEs at $z=5.7$ (top) and $z=6.6$ (bottom) taken by the Subaru Suprime-Cam ($B$, $V$, $R$, $i'$, $NB$, $z'$), the UKIRT WFCAM ($J$, $H$, $K$), and the Spitzer IRAC ($3.6\\mu$m, $4.5\\mu$m, $5.8\\mu$m, $8.0\\mu$m). The $z=5.7$ LAE is detected (brighter than the $3\\sigma$ magnitude) in $i'$, $NB816$, $z'$, and $3.6\\mu$m, and nearly detected in $J$. The $z=6.6$ LAE is detected in $NB921$, $z'$, and $3.6\\mu$m. Each panel is $6'' \\times 6''$ in size, or $\\simeq 35$ $(33)$ kpc at $z = 5.70$ $(6.56)$. } \\label{fig:faces} \\end{figure*}   \\begin{deluxetable}{cccccc} \\tablecolumns{3} \\tablewidth{0pt} \\tablecaption{Median-Stacked Lyman Alpha Emitters \\\\ at $z=5.7$ and $6.6$ \\label{tab:photo_prop}} \\tablehead{ \\colhead{}    & \\colhead{$z=5.7$ LAE}    & \\colhead{$z=6.6$ LAE} } \\startdata $NB$\\tablenotemark{a}  &  $25.5$ ($27.6$)  &  $25.6$ ($27.0$) \\\\ $m_{\\rm con}$\\tablenotemark{b}  &  $27.2$ ($28.2$)  &  $27.4$ ($26.6$) \\\\ $B$  &  $99.9$ ($30.4$)  &  $99.9$ ($30.1$) \\\\ $V$  &  $99.9$ ($29.8$)  &  $99.9$ ($29.5$) \\\\ $R$  &  $99.9$ ($29.9$)  &  $99.9$ ($29.7$) \\\\ $i'$  &  $27.9$ ($30.0$)  &  $29.9$ ($29.6$) \\\\ $z'$  &  $27.5$ ($29.0$)  &  $27.7$ ($28.6$) \\\\ $J$  &  $27.6$ ($27.5$)  &  $27.7$ ($27.1$) \\\\ $H$  &  $99.9$ ($26.6$)  &  $99.9$ ($26.3$) \\\\ $K$  &  $99.9$ ($27.0$)  &  $99.9$ ($26.7$) \\\\ $3.6 \\mu$m  &  $26.7$ ($27.0$)  &  $26.6$ ($26.7$) \\\\ $4.5 \\mu$m  &  $32.0$ ($26.6$)  &  $26.7$ ($26.3$) \\\\ $5.8 \\mu$m  &  $99.9$ ($25.4$)  &  $99.9$ ($24.4$) \\\\ $8.0 \\mu$m  &  $99.9$ ($25.3$)  &  $99.9$ ($24.1$) \\\\ $\\beta$\\tablenotemark{c}  &  $-2.9 \\pm 1.0$  &  $-3.0 \\pm 2.7$ \\\\ $f({\\rm Ly}\\alpha)$ [$10^{-18} $erg s$^{-1}$cm$^{-2}$]\\tablenotemark{d}  &  $11.0 \\pm 0.77$  &  $8.93 \\pm 1.25$ \\\\ $L({\\rm Ly}\\alpha)$ [$10^{42} $erg s$^{-1}$]\\tablenotemark{d}  &  $3.90 \\pm 0.27$  &  $4.37 \\pm 0.61$ \\\\ ${\\rm EW}({\\rm Ly}\\alpha)_{\\rm rest}$ [{\\AA}]\\tablenotemark{d}  &  $78.1^{+96.3}_{-65.4}$  &  $84.2^{+170.8}_{-49.3}$ \\\\ redshift\\tablenotemark{e} & $5.70$ & $6.56$ \\enddata \\tablecomments{ All magnitudes are total magnitudes. $99.9$ mag means negative flux densities. Magnitudes in parentheses are $3\\sigma$ limiting magnitudes. } \\tablenotetext{a}{% $NB816$ for $z=5.7$ LAE, $NB921$ for $z=6.6$ LAE. } \\tablenotetext{b}{% Rest-frame UV continuum magnitudes at (observed) effective wavelengths of $8303${\\AA} ($z=5.7$) and $9445${\\AA} ($z=6.6$) after correction for Ly$\\alpha$ emission and the IGM absorption, derived from $NB816$ and $i'$ magnitudes ($z=5.7$) and $NB921$ and $z'$ magnitudes ($z=6.6$). } \\tablenotetext{c}{% Derived from $m_{\\rm con}$ and $J$ magnitude using eq. (\\ref{eq:beta}) } \\tablenotetext{d}{% Derived assuming that the redshifted Ly$\\alpha$ line is at the center of the $NB$ filter. } \\tablenotetext{e}{% Corresponding to the central wavelengths of $NB816$ and $NB921$. } \\end{deluxetable} ",
        "conclusions": "\\label{sec:conclusion}  In this paper, we investigated the stellar populations of LAEs at $z = 5.7$ and $6.6$ found in $0.65$ deg$^2$ of the SXDF, based on deep rest-frame UV-to-optical photometry obtained from the three surveys: SXDS, UKIDSS/UDS, and SpUDS. We made composite images from $165$ and $91$ LAEs at $z=5.7$ and $6.6$, respectively, which are fainter than the $3\\sigma$ magnitude in the IRAC $3.6\\mu$m band, and derived typical SEDs of $z\\sim 6-7$ LAEs for the first time. We found that their UV continua are as blue as those of dropout galaxies at similar redshifts, with UV spectral slopes $\\beta\\sim -3$, albeit with large photometric uncertainties. Fitting stellar population synthesis models with and without nebular emission, which is parameterized by $f_{\\rm esc}^{\\rm ion}$, to the multiband data of the stacked objects at $z=5.7$ and $6.6$, we derived their stellar masses, ages, and dust extinction.   Our main results are as follows:  \\begin{enumerate} \\item[(i)] We find that the stacked LAEs at both redshifts are fitted well by $f_{\\rm esc}^{\\rm ion}=0$ models. The best-fit $f_{\\rm esc}^{\\rm ion}=0$ models have low stellar masses of $\\sim \\left( 3-10 \\right) \\times 10^{7}M_\\odot$, very young ages of $\\sim 1-3$ Myr, and negligible dust extinction. In these models, young stellar populations reproduce the observed blue UV continua, and strong nebular emission redward of 4000 {\\AA} makes the UV-to-optical color as red as observed.  While we find that the $z=6.6$ LAE is also fitted similarly well by an old, massive model without nebular emission, we do not take this models as the best-fit model, since typical LAEs up to $z \\sim 6$, including our $z=5.7$ LAEs, have been consistently found to be very young and low-mass galaxies. We propose that typical $z\\sim 6-7$ LAEs are candidates of galaxy building blocks at the early stage of galaxy formation.     \\item[(ii)] We estimate the Ly$\\alpha$ escape fraction to be $f_{\\rm esc}^{{\\rm Ly}\\alpha} \\simeq 0.36$ for the $z=5.7$ LAE and $f_{\\rm esc}^{{\\rm Ly}\\alpha} \\simeq 0.04$ for the $z=6.6$ LAE with large errors. This decrease from $z=5.7$ to $6.6$, if real, might be due to an increase in the neutral fraction of the IGM.     \\item[(iii)] We also apply SED fitting to a stacked object from $z \\sim 7$ $z$-dropout galaxy candidates found by a recent HST/WFC3 deep survey, and find that it is either a young, low-mass galaxy with strong nebular emission similar to the $z=5.7$ and 6.6 LAEs, or a very old and massive galaxy with little nebular emission.     \\item[(iv)] From the constraints on nebular emission models, we estimate the upper limit of the Ly-continuum escape fraction to be $f_{\\rm esc}^{\\rm ion} \\sim 0.6$ for the $z=5.7$ LAE and $\\sim 0.9$ for the $z=6.6$ LAE. We apply this technique to the $z$-dropout galaxies, but obtain no meaningful constraints on $f_{\\rm esc}^{\\rm ion}$. \\end{enumerate}"
    },
    "1004/1004.3017_arXiv.txt": {
        "abstract": "The Sun is the only star that we can spatially resolve and it can be regarded as a fundamental plasma laboratory of astrophysics. The solar transition region (TR), the layer between the solar chromosphere and corona, plays an important role in solar wind origin and coronal heating. Recent high-resolution observations made by \\textit{SOHO}, \\textit{TRACE}, and \\textit{Hinode} indicate that the TR is highly nonuniform and magnetically structured. Through a combination of spectroscopic observations and magnetic field extrapolations, the TR magnetic structures and plasma properties have been found to be different in coronal holes and in the quiet Sun. In active regions, the TR density and temperature structures also differ in sunspots and the surrounding plage regions. Although the TR is believed to be a dynamic layer, quasi-steady flows lasting from several hours to several days are often present in the quiet Sun, coronal holes, and active regions, indicating some kind of plasma circulation/convection in the TR and corona. The emission of hydrogen Lyman lines, which originates from the lower TR, has also been intensively investigated in the recent past. Observations show clearly that the flows and dynamics in the middle and upper TR can greatly modify the Lyman line profiles. ",
        "introduction": "\\label{}  The Sun is the only star that we can spatially resolve. Observations suggest that many phenomena such as flares, particle acceleration, and collimated outflows not only occur in the solar atmosphere, but also occur in other astrophysical systems in similar forms \\citep[e.g.,][]{Zhang2000,Zhang2007}. Thus, the Sun can be regarded as a fundamental plasma laboratory of astrophysics.  The solar transition region (TR), the temperature regime from $\\sim$10$^{4}$~K to 10$^{6}$~K, is traditionally believed to be a very thin layer between the solar chromosphere and corona. Figure~\\ref{fig.1} illustrates the structure of the solar interior and solar atmosphere, and shows the height variations of the temperature and density of the average solar atmosphere following the model of \\citet{Vernazza1981}. Abrupt changes of the temperature and density are clearly present through the TR. As an interface between the cool chromosphere and the hot corona, the TR plays an important role in coronal heating and solar wind origin \\citep[e.g.,][]{Chae1998b,Peter2003,Tu2005a}.  \\begin{figure} \\centering \\includegraphics[width=.48\\textwidth]{f1_1.eps} \\includegraphics[width=.48\\textwidth]{f1_2.eps} \\caption{Upper: Structure of the solar interior and solar atmosphere. Adapted from the magazine of Scientific American (2003) and the thesis of \\citet{McAteer2004} (Credit: Don Dixon.). Lower: Height variations of the temperature and density of the average solar atmosphere. Adapted from \\citet{Peter2004}.} \\label{fig.1} \\end{figure}  Most of the TR emission falls into the FUV/EUV range ($\\sim$50~{\\AA}-2000~{\\AA}). And thus spectroscopic observations in this spectral range are the most important way to obtain information about the physics of the TR. Since the plasma is optically thin for most FUV/EUV emission lines, we can extract information about the physical conditions prevailing in their source regions from the line profiles \\citep[see reviews by][]{Mariska1992,Xia2003,Wilhelm2007}. Wavelengths and associated ions of several typical TR emission lines (the Mg~{\\sc{x}}~625~{\\AA} line is formed in the lower corona) are marked in Figure~\\ref{fig.1} according to the formation temperatures of the ions (temperature of maximum fractional ionization). The TR can be divided into three parts according to the temperature: lower TR ($10^{4}$~K-$3\\times10^{4}$~K), middle TR ($3\\times10^{4}$~K-$5\\times10^{5}$~K), and upper TR ($5\\times10^{5}$~K-$10^{6}$~K). The lines of \\lya~1216~{\\AA}, He~{\\sc{i}}~584~{\\AA} and Si~{\\sc{ii}}~1533~{\\AA} are all typical emission lines formed in the lower TR. The Ne~{\\sc{viii}}~770~{\\AA} line is a typical upper-TR line and has been widely used in observations in the last decade. In the middle TR, thousands of emission lines such as C~{\\sc{ii}}~1037~{\\AA} and C~{\\sc{iv}}~1548~{\\AA} have been identified \\citep{Curdt2001}, and many of them have been used in observations to probe the physical properties of the TR.  Early observations made by instruments such as the S082B EUV spectrograph onboard the Skylab space station \\citep{Bartoe1977}, the UV spectrometer onboard the OSO 8 (Orbiting Solar Observatory) satellite \\citep{Bonnet1981}, and the High-Resolution Telescope Spectrograph (HRTS) flown on some rockets and Spacelab 2 \\citep{Bartoe1975,Brueckner1986} have provided much valuable information about the solar TR. Early instruments and results of these early observations have been reviewed by \\citet{Wilhelm2004} and \\citet{Mariska1992}, respectively. However, our knowledge of the TR was limited due to the lack of long-duration high-resolution observations before middle 1990s.  With the launch of SOHO (Solar and Heliospheric Observatory), TRACE (Transition Region and Coronal Explorer), and Hinode, our understanding of the TR has been significantly improved since 1996. Huge amount of data acquired by these missions reveal clearly that the TR is a highly dynamic and nonuniform layer. High spatial ($\\sim$1$^{\\prime\\prime}$), temporal ($\\sim$1~s), and spectral (1st order $\\sim$44~m{\\AA}/pixel, 2nd order $\\sim$22~m{\\AA}/pixel) observations of the SUMER (Solar Ultraviolet Measurements of Emitted Radiation) instrument \\citep{Wilhelm1995,Lemaire1997} onboard SOHO have greatly enhanced our knowledge of the TR. Thousands of emission lines with wavelengths from 660~{\\AA} to 1610~{\\AA} and formation temperatures from $10^{4}$~K to $2\\times10^{6}$~K have been identified and used in extensive studies of the upper solar atmosphere \\citep{Curdt2001,Curdt2004}. Moreover, the Doppler shifts of these FUV/EUV spectral lines could be measured with an accuracy of about 1-2~km/s \\citep{Brekke1997,PeterJudge1999,Popescu2004}, and thus the flow velocities in the TR could be determined with a high accuracy. Besides, the Coronal Diagnostic Spectrometer (CDS) onboard SOHO \\citep{Harrison1995} and the EUV Imaging Spectrometer (EIS) onboard Hinode \\citep{Culhane2007} have the capability of simultaneous multi-wavelength observations using a number of TR and coronal lines, which are very suitable for plasma diagnostics in the upper solar atmosphere.  Disk observations reveal that the bright network pattern is the dominant emission structure in intensity images of TR lines \\citep[e.g.,][]{Brueckner1974,Reeves1976}. The TR network is cospatial with the underlying chromospheric network, and is believed to be the upward extension of the supergranular boundary to higher layers \\citep[e.g.,][]{Brueckner1974,Patsourakos1999}. \\citet{Reeves1976} found that in the quiet Sun the contrast between the network and internetwork emission is largest for lines formed at $\\sim$2$\\times$10$^{5}$~K and decreases at higher and lower temperatures. \\citet{Patsourakos1999} investigated the height variation of the network size (width of the network lane) in the upper atmosphere of the quiet Sun, and found that the network size is almost constant in the middle TR and increases strongly at the upper TR. The TR is a low-$\\beta$ plasma region and the magnetic field largely controls the distribution of the plasma.  The first magnetic network model was proposed by \\citet{Gabriel1976}. In this model, magnetic fields concentrated in the supergranular boundary expand with height and open into the corona. The TR radiation is balanced by the downward thermal conduction. However, the emission measure reproduced by this model is far less than the observed value at temperatures below $10^5$~K \\citep{Athay1982}. Ten years later, \\citet{Dowdy1986} proposed a modified model in which only part of the network magnetic fields open into the corona as a funnel shape, while the rest of the network is occupied by a population of low-lying loops with lengths less than 10~Mm (1~Mm=10$^6$~m). In this model, the small loops are heated internally. \\citet{Peter2001} further pointed out that magnetic funnels originating from the network could be connected to the solar wind or be legs of large-scale coronal loops. It might also be possible that the fundamental structure of the TR has not yet been resolved due to limitations of current instruments \\citep{Feldman1983,Feldman1987,FeldmanLaming1994,Doschek2004}.  The TR materials are often flowing. By analyzing the spectra obtained by the S082B instrument, \\citet{Doschek1976} found that TR lines are often redshifted. The persistent redshifts of emission lines formed in the middle TR were confirmed in later observations. With SUMER observations, \\citet{PeterJudge1999}, \\citet{Teriaca1999}, \\citet{Stucki2000} and \\citet{Xia2004} studied the variation of the average Doppler shift with the line formation temperature. They found that the redshift is most significant ($\\sim$10~km~s$^{-1}$) at $\\sim$2$\\times$10$^{5}$~K and decreases at lower and higher temperatures. Emission lines formed in the lower TR reveal no obvious Doppler shift on average, while emission lines formed in the upper TR are clearly blue shifted. The temperature variation of the TR line shift shows a similar trend in both the quiet Sun and coronal holes. In both regions, the value of the redshift is positively correlated with the line radiance, and is more prominent in the network than in cell interiors \\citep[e.g.,][]{Curdt2008a,Aiouaz2008}. The blue shifts of upper TR lines are predominant in coronal holes and are suggested to be signatures of the fast solar wind \\citep{Dammasch1999,Hassler1999,Wilhelm2000,XiaEtal2003,Tu2005a,Aiouaz2005}, while blue shifts are localized at network junctions in the upper TR of the quiet Sun \\citep{Hassler1999}. In active regions, the redshifts of middle TR lines are about two times larger than those in the quiet Sun \\citep[e.g.,][]{Teriaca1999}. Different mechanisms such as spicule activities \\citep[e.g.,][]{Athay1984}, siphon flows from one loop leg to the other \\citep[e.g.,][]{McClymont1987,Mariska1988}, downward propagating MHD waves produced by nano-flares \\citep[e.g.,][]{Hansteen1993,Hansteen1997} and solar wind related reconnections in the network \\citep{Tu2005a,He2008} have been suggested to explain these prominent Doppler shifts of TR lines \\citep[see reviews and discussions in][]{Mariska1992,Brekke1997,PeterJudge1999,Xia2003}.  Small-scale transient phenomena such as explosive events, coronal bright points, blinkers and macrospicules are frequently observed in the TR. Explosive events (EEs), which were first identified through HRTS observations \\citep{Brueckner1983}, are characterized by a small spatial scale of about 1800 km, a short lifetime of about 60~s, and non-Gaussian line profiles showing Doppler shifts of $\\sim$100~km~s$^{-1}$ \\citep[e.g.,][]{Dere1989,Innes1997a,Doyle2006}. EEs tend to occur along boundaries of the magnetic network, where weak mixed-polarity magnetic features are present \\citep[e.g.,][]{Chae1998a,Teriaca2004}. They are believed to be produced by magnetic reconnections in the TR \\citep{Innes1997b,Chen2006}. Coronal bright points (BPs), which are characterized by locally enhanced emission in X-ray and EUV, are associated with the underlying bipolar magnetic field at network boundaries \\citep[e.g.][]{Habbal1990,Webb1993,Falconer1998,Brown2001,Madjarska2003,Tian2007}. A BP consists of several miniature dynamic loops \\citep[e.g.,][]{Sheeley1979,Ugarte2004a} and is typically $30^{\\prime\\prime}-40^{\\prime\\prime}$ in size. The average lifetime of BPs is 20 hours in EUV \\citep{Zhang2001} and 8 hours in X-ray observations \\citep{Golub1974}. The energization of BPs may be associated with the interaction between two magnetic fragments of opposite polarities \\citep{Priest1994,Parnell1994}, magnetic reconnection along separator field lines \\citep{Longcope1998}, or current sheets induced by photospheric motion \\citep{Buchner2006,Santos2007}. Blinkers are transient TR network brightenings with a typical duration of 5 to 30 minutes \\citep[e.g.,][]{Harrison1997,Doyle2004}. Macrospicules are jet-like structures extending from the chromosphere into the corona, which could be regarded as the EUV counterpart of the H$\\alpha$ spicules \\citep[e.g.,][]{Bohlin1975,Xia2005}.  The hydrogen Lyman lines, especially the \\lya~line, play a dominant role in the radiative energy transport in the lower TR \\citep{Fontenla1988}. Due to the effect of radiative transfer, a central reversal is usually present in \\lya~and \\lyb~line profiles. The Lyman line profiles can be used to reveal information on fine structures and physical properties of quiescent prominences \\citep[e.g.,][]{Heinzel2005,Vial2007,Schmieder2007,Gunar2008}. They can also be used to diagnose nonthermal effects in solar flares \\citep{Henoux1995,Fang1995,Xu2005}. Full \\lya~and \\lyb~line profiles were acquired through early observations more than two decades ago \\citep[e.g.][]{Nicolas1976,Kneer1981,Lemaire1978,Vial1982,Bocchialini1996,Basri1979,Fontenla1988}. These observations revealed valuable information on the stratification of the solar upper atmosphere, although the obtained Lyman line profiles were hampered by the geocoronal absorption at the center. The SUMER observations at the first Lagrangian point are free from geocoronal absorptions. By analyzing spectra obtained by SUMER, \\citet{Warren1998} found that the average profiles for \\lyb~through Ly$\\epsilon$ are all reversed at the center and stronger in the red wing. It has been suggested that the asymmetries of Lyman line profiles are related to the differential flows in the upper solar atmosphere \\citep{Gouttebroze1978,Fontenla2002,Gunar2008}. \\citet{Warren1998} also found larger peak separations of the profiles at limb than at disk center. The full high-quality \\lya~line profiles were not recorded on the detectors of SUMER for a long time, because the high radiance of the line usually leads to a saturation of the detector microchannel plates. Attempts have been made to record \\lya~on the bare part of detector A. However, the signal determination turned out to be highly uncertain due to the local-gain depression correction \\citep{Teriaca2005a,Teriaca2005b,Teriaca2006}.  Although much has been known about the emission and structures of the TR, there are still many important things that have not or rarely been investigated. For example, almost all studies of the TR network structures are concentrated on the average Sun or quiet Sun \\citep[e.g.,][]{Gabriel1976,Dowdy1986,Patsourakos1999,Peter2001}. But the TR is in fact highly nonuniform \\citep[e.g.,][]{Marsch2006} so that differences between different solar regions (quiet Sun, coronal holes, active regions) should be expected. Another aspect is concerning the nature of the flow field in the TR. High-cadence observations reveal that the TR is highly dynamic and transient flows are often detected \\citep[e.g.,][]{Innes1997a}, while quasi-steady flows lasting from several hours to several days are also observed to be present in the TR \\citep[e.g.,][]{Marsch2004,Dammasch2008}. It is hard to imagine how these two kinds of flows could coexist in the TR. Moreover, upflows in the upper TR outside coronal holes are not well understood \\citep{Hassler1999,Marsch2004}. The TR above sunspots is also an interesting topic. Although sunspots and plage have been intensively studied in the photosphere and chromosphere \\citep{Solanki2003}, their possible different properties in the TR are not well understood. Finally, the lower TR line \\lya~has been suggested to be used in proposed future solar missions such as Solar Orbiter \\citep{Marsch2002}, Kuafu \\citep{Tu2008} and SMESE \\citep{Vial2008}. Despite its importance, high-quality undisturbed \\lya~profiles have not been obtained, and signatures of TR flows and transient phenomena in Lyman line profiles have not been identified in observations before 2008. These open questions are summarized below:  (1) Are there any differences between the TR properties in coronal holes and in the quiet Sun?  (2) Why are quasi-steady flows frequently observed in the TR and corona?  (3) Are there any differences between the TR properties above sunspots and the surrounding plage regions?  (4) What are the characteristics of the origin and initial acceleration of the solar wind in the TR and corona?  (5) What does the genuine solar hydrogen \\lya~radiance profile (that unaffected by the atmosphere of the earth) look like?  In the past several years, a lot of work has been done to understand the problems mentioned above. The magnetically structured TR was found to have different properties in different regions through a combination of spectroscopic observations and magnetic field extrapolations. With the help of magnetic field extrapolations, the TR flows in different regions were better explained, and the persistent and transient flows can now be understood in a unified picture. Comprehensive comparisons of TR properties between sunspots and plage regions have also been done in the recent past, and distinct differences were found. A new technique of recording \\lya~profiles has been employed by the SUMER team since June 2008, and high-quality \\lya~profiles without geocoronal absorption were obtained for the first time. It has been found that TR flows and EEs can modify the observed Lyman line profiles. This article reviews these recent progresses and provides a new view on the TR emission and structures. ",
        "conclusions": "Recent advances in the investigation of the solar TR, resulting especially from spectroscopic observations in the FUV/EUV spectral range, have been reviewed in this article. Significant progress was achieved in many aspects of the TR research, concerning the magnetic field structures and plasma properties, behaviors of the nascent solar wind and other systematic flows, and characteristics of the hydrogen Lyman line profiles in different parts of the Sun (the quiet Sun, coronal holes, and ARs).  Based on studies combining spectroscopic observations with magnetic field extrapolations, a new scenario of the TR consisting of different magnetic field and emission structures in the quiet Sun and coronal holes has been proposed. In this new picture (shown in Figure~\\ref{fig.12}), the TR is higher and more extended in coronal holes than in the quiet Sun. The dominant magnetic field structures in CHs are open magnetic funnels which are related to the fast solar wind, whereas there are only low-lying magnetic loops in coronal holes. In contrast, in the quiet Sun magnetic loops with different spatial scales are crowded. Occasionally, some large loops might become locally open through magnetic reconnection, and then possibly release material to the solar wind. These different scenarios naturally lead to a strong expansion of the magnetic funnels in coronal holes and a relatively weak expansion of magnetic field structures in the quiet Sun through the upper TR.  Although the TR is believed to be a dynamic layer, quasi-steady TR flows lasting from several hours to several days have also been observed to be present in the quiet Sun, coronal holes, and ARs. In coronal holes, the predominant blue shift of emission lines was found to increase steadily from the upper TR to the lower corona, indicating the steady acceleration of the nascent fast solar wind, while in the quiet Sun the blue shifts of upper-TR lines are localized in network junctions and considered to be signatures of mass supply to large coronal loops. The persistent redshift of middle-TR lines in the network could correspond to the downflowing plasma occurring after continuous magnetic reconnections or cooling of the upflows. Persistent redshifts of emission lines formed in the TR and lower corona are usually found in the legs of AR loops, and steady blue shifts lasting for several days are often present at the boundaries of some ARs. All these flows seem to be guided by the dominant magnetic field structures, and their quasi-steadyness suggests the presence of some kind of convection/circulation in the TR and corona. Sporadically, small-scale transient flows will locally occur and be superposed on the background quasi-steady flow field.  Based on spectroscopic observations in EUV/FUV, properties of the TR above sunspots and the surrounding plage regions have also been comprehensively investigated. By comparing the Lyman line profiles, electron densities, DEM curves, and filling factors of the sunspot plume, umbra, penumbra, and the surrounding plage regions, the TR above sunspots was found to be higher and probably more extended (shown in Figure~\\ref{fig.18}). The opacity of the hydrogen lines is apparently much smaller above sunspots than above plage regions. Theoretical calculations and observational inferences also indicate that the enhanced TR emission of the sunspot plume is probably caused by a large filling factor.  As the dominant emission line in the lower TR, the hydrogen \\lya~line provides a good tool to diagnose through its profile the coupling between the lower-TR emission and the flows in upper-TR layers. With SUMER observations, high-quality solar \\lya~profiles without geocoronal absorption were recently obtained for the first time. As an unexpected result, this strongly self-reversed line is often asymmetric with a stronger blue peak, and there is a clear correspondence between this asymmetry and downflows in the middle TR. Profiles with deeper self-reversals were found statistically in the internetwork rather than the network. A larger peak separation indicative of a larger opacity was found in coronal holes than in the quiet Sun. The \\lyb~line is also a useful line to study the coupling between the lower TR and higher layer, because of its opposite asymmetry in the quiet Sun and polar coronal holes. Signatures of TR explosive events are often clearly present in the \\lyb~line."
    },
    "1004/1004.0938_arXiv.txt": {
        "abstract": "In radio astronomy, reference signals from auxiliary antennae, receiving only the radio frequency interference (RFI), can be modified to model the RFI environment at the astronomy receivers. The RFI can then be cancelled from the astronomy signal paths. However, astronomers typically only require signal statistics. If the RFI statistics are changing slowly, the cancellation can be applied to the signal correlations at a much lower rate than required for standard adaptive filters. In this paper we describe five canceller setups; pre- and post-correlation cancellers that use one or two reference signals in different ways. The theoretical residual RFI and added noise levels are examined, and demonstrated using microwave television RFI at the Australia Telescope Compact Array. The RFI is attenuated to below the system noise, a reduction of at least 20 dB. While dual-reference cancellers add more reference noise than single-reference cancellers, this noise is zero-mean and only adds to the system noise, decreasing the sensitivity. The residual RFI that remains in the output of single-reference cancellers (but not dual-reference cancellers) sets a non-zero noise floor that does not act like random system noise and may limit the achievable sensitivity. Thus dual-reference cancellers often result in superior cancellation. Dual-reference pre-correlation cancellers require a double-canceller setup to be useful and to give equivalent results to dual-reference post-correlation cancellers. ",
        "introduction": "\\label{INTRODUCTION}  Radio astronomy is poised to move ahead with a suite of new instruments such as the Square Kilometre Array (SKA). These instruments will improve sensitivity, resolution, bandwidth, and many other instrument and observational parameters by more than an order of magnitude. At present, many radio astronomy observations are corrupted to at least some extent by radio frequency interference (RFI). This interference comes from ground-based communication transmitters, satellites, and the observatory equipment itself. With the increase in sensitivity and frequency coverage of radio astronomy instruments, and with telecommunication signals occupying more of the spectrum, it is essential to develop ways of removing or suppressing this RFI.  In radio astronomy, real-time adaptive filters can be used to modify an auxiliary voltage time series (the reference signal) so that it cancels RFI from an astronomical voltage time series \\citep{Barnbaum1998, Bower2001, Mitchell2005}. For each voltage sample the filters are allowed to slightly vary their internal coefficients in order to adapt to any changes taking place in the RFI. If one is interested in the power spectrum of the astronomy signal, and the filter coefficients stay fairly constant over the interval in which the power spectrum is estimated, cancellation in the post-correlation domain can give better results, since a second reference antenna can be used to give complete suppression of the RFI (only zero-mean random receiver noise is added to the complex correlations and it will average away, see \\citealt{Briggs2000, Mitchell2001}). However, many applications (particularly in the communications field, but also some astronomy applications), require the recovery of the actual symbol stream (i.e., a transmitted sequence of symbols such as bits or words) from the noisy RF environment, which is not retained in post-correlation. Following a suggestion of \\citet{Briggs2000}, we have devised a modified approach that can give improved RFI attenuation in the voltage domain. In an attempt to minimise any RFI in an astronomical voltage series the standard approach is to minimise the canceller's output power, which it will be shown means that some residual RFI always remains. The modified approach that we give here forces the RFI in the output power to zero. It results in residual power that is always greater than that of the standard approach (output power is no longer minimised), but which does not contain RFI. That is, superior RFI cancellation is obtained, but at the expense of somewhat increased thermal noise.  In the following sections the standard adaptive canceller is discussed, and the new approach is introduced. This is followed by an overview of how cancellation can be applied after correlations are formed. Residual RFI power and added receiver noise are investigated and an example from the Australia Telescope Compact Array is given. ",
        "conclusions": "\\label{SUMMARY}  Interference cancelling using a single reference antenna can give excellent results when the reference signal interference-to-noise ratio is large, and there is more gain towards the interfering signal for the reference antenna than for the astronomy antennae. However, receiver noise in the reference signal means that a fraction of the interference will always remain after cancelling. A second reference signal can be used to remove the noise bias and give infinite interference attenuation, but a larger amount of reference receiver noise is added during cancelling. For pre-correlation systems, a dual canceller setup can be used to average the (zero-mean) receiver noise away, a process that comes automatically with post-correlation cancellers. A breakdown of the main properties for the different mitigation techniques is given in table \\ref{SUMMARY TABLE}.  It is important to note that even though the single-reference cancellers leave residual RFI, the residual may be extremely small and well below the RMS noise. This occurs when the interference-to-noise ratios of the reference signals are very large, and the use of two references (in pre-correlation systems) might just add complexity to the system with little or no benefit. However, if maximum sensitivity is required, one should be aware that they will eventually reach a non-zero residual signal.  Using two reference signals to remove the reference receiver noise bias removes the inherent stability of the algorithms in situations where some or all of the frequency channels are interference-free. Although there are applications in which the passband will always be entirely filled with RFI (such as observations in the GPS L1 and L2 bands), many interfering signals will only take up a part of the band. In these cases the algorithms need a mechanism to turn themselves off in the vacant frequency channels.    \\begin{table*} \\center \\begin{tabular}{|c|c|c|c|c|c|c|} \\hline & \\parbox[t]{1.0cm}{Fig. no.} & \\parbox[t]{0.8cm}{Ref. Ants.} & Filters & Configuration & Behaviour & \\parbox[t]{1cm}{Eqn. no.} \\\\ &&&&&&\\\\ \\hline MK1a & \\ref{mk1 adaptive cancellers} & 1 & 1 & \\parbox[t]{4.2cm}{One reference antenna feeds both filter and weight generator.} & \\parbox[t]{4.2cm}{Minimises residual power, but includes some RFI. Turns off gracefully as $INR_r\\rightarrow0$.} & \\ref{mk1 residual power}, \\ref{mk1 residuals a}, \\ref{mk1 residuals b} \\\\ &&&&&&\\\\ \\hline MK2a & \\ref{mk2 adaptive cancellers} & 2 & 1 & \\parbox[t]{4.2cm}{Independent reference antennae for filter and weight generator.} & \\parbox[t]{4.2cm}{Totally cancels RFI, but adds extra noise. Must be controlled as $INR_r\\rightarrow0$.} & \\ref{mk2 residual power} \\\\ &&&&&&\\\\ \\hline MK1b & \\ref{indepRX MK1} & 2 & 2 & \\parbox[t]{4.2cm}{Two MK1a cancellers, operating on copies of main signal.} & \\parbox[t]{4.2cm}{Sensitivity decrease, and non-zero RFI floor.} & \\ref{mk1 indepRX} \\\\ &&&&&&\\\\ \\hline MK2b & \\ref{indepRX MK2} & 2 & 2 & \\parbox[t]{4.2cm}{Two MK2a cancellers, operating on copies of main signal.} & \\parbox[t]{4.2cm}{Sensitivity decrease, but no non-zero RFI floor.} & \\ref{mk2 indepRX} \\\\ &&&&&&\\\\ \\hline PC2 & \\parbox[t]{1cm}{not shown} & 2 & 1 & \\parbox[t]{4.2cm}{Two reference correlations with each astronomy signal.} & \\parbox[t]{4.2cm}{Sensitivity decrease, but no non-zero RFI floor. Must be controlled as $INR_r\\rightarrow0$.} & \\ref{pc2 residual power}, \\ref{stablised pc2 resid} \\\\ &&&&&&\\\\ \\hline \\end{tabular} \\caption{Summary of the main properties for various interference mitigation techniques\\label{SUMMARY TABLE}} \\end{table*}"
    },
    "1004/1004.0109_arXiv.txt": {
        "abstract": "Simulations of the magnetorotational instability (MRI) in `unstratified' shearing boxes exhibit powerful coherent flows, whereby the fluid vertically splits into countermoving planar jets or `channels'. Channel flows correspond to certain axisymmetric linear MRI modes, and their preponderance follows from the remarkable fact that they are approximate nonlinear solutions of the MHD equations in the limit of weak magnetic fields. We show in this paper, analytically and with one-dimensional numerical simulations, that this property is also shared by certain axisymmetric MRI modes in \\emph{vertically-stratified} shearing boxes. These channel flows rapidly capture significant amounts of magnetic and kinetic energy, and thus are vulnerable to secondary shear instabilities. We examine these parasites in the vertically stratified context, and estimate the maximum amplitudes that channels attain before they are destroyed. These estimates suggest that a dominant channel flow will usually drive the disk's magnetic field to thermal strengths. The prominence of these flows and their destruction place enormous demands on simulations, but channels in their initial stages also offer a useful check on numerical codes. These benchmarks are especially valuable given the increasing interest in the saturation of the stratified MRI. Lastly we speculate on the potential connection between `run-away' channel flows and outburst behaviour in protostellar and dwarf nova disks. ",
        "introduction": "Differential rotation and weak magnetic fields conspire to destabilise laminar flow in accretion disks (and other astrophysical settings) via the magnetorotational instability (MRI) (Balbus and Hawley 1991). The MHD turbulence that this linear instability sustains can transport significant angular momentum, and thus drive the mass accretion that is observed in many disk systems (Hawley et al.~1995, Stone at al.~1996, Balbus and Hawley 1998, Hawley 2000).  The characteristics of MRI turbulence have been studied extensively with local three-dimensional simulations in the `unstratified' shearing box\\footnote[1]{Throughout the paper we adhere to the common (mis-) use of the term `stratified', by which we mean a background variation in density and pressure but not necessarily in entropy. Normally, however, `stratified' explicitly means that the background entropy does vary.}, which is a model that mimics a small block of gas deeply embedded in the disk, ignorant of its large-scale vertical and radial structure (Balbus and Hawley 1998). When a net vertical magnetic flux passes through such a box, the fluid can be dominated by coherent structures (often recurring) called channel flows. Akin to Kolmogorov shear flow, these consist of several countermoving planar streams of fluid situated at different vertical levels (Meshalkin and Sinai 1961, Hawley et al.~1995, Sano and Inutsuka 2001, Sano 2007, Goodman and Xu 1994, hereafter GX94). Previous work has investigated the role these structures, and their destruction, play in controlling MRI saturation in the shearing box model (GX94, Latter et al.~2009, hereafter LLB09, Pessah and Goodman 2009, Pessah 2009). In this paper we generalise these analyses to the case of a shearing box in which the \\emph{vertical structure} of the disk has been incorporated.   The preponderance of channel behaviour in unstratified boxes is connected to the fact that the channels correspond not only to linear MRI modes but also to approximate nonlinear solutions of the MHD equations, when in the incompressible, weak field limit (GX94). In this paper, we show explicitly that analogous linear modes in a vertically-stratified box possess the exact same property --- they are also approximate nonlinear solutions\\footnote[2]{Alternatively one could state that the Goodman and Xu channels are exact nonlinear solutions of the equations of incompressible MHD, and the solutions we present in this paper are exact nonlinear solutions of the equations of \\emph{anelastic} MHD. But, since both classes of solution leave each regime so rapidly, we prefer to call them approximate solutions to the full equations.}. This property has important consequences. It means that channel flows also dominate stratified simulations with a net vertical flux, growing exponentially and free of nonlinear interactions for a time sufficient to accumulate significant amounts of kinetic and magnetic energy. But the ultimate destruction of such a large amplitude channel can be so violent that the computational domain is evacuated of mass, and the simulation compromised (Stone et al.~1996, Miller and Stone 2000). In contrast, behaviour elicited by other magnetic configurations (a purely azimuthal field, for example) is not nearly so striking, nor numerically problematic (Miller and Stone 2000). As a result, published simulations with a net vertical field are comparatively rare, even if the magnetic configuration itself is of astrophysical importance.  The agents responsible for the disruption of isolated channel flows are nonaxisymmetric parasitic instabilities, which feed on their velocity and magnetic shear (GX94, LLB09, Pessah and Goodman 2009, Pessah 2009). In general, these parasitic modes are ineffectual until large channel amplitudes are attained because of the combined action of the background shear and the exponential growth of the channel itself. In this paper, we compute directly some of the ideal MHD parasitic modes that might be present in stratified boxes in a convenient intermediate limit. Using these results, we estimate the minimum channel amplitude that permits a successful parasitic attack. In a later article a general survey with three-dimensional simulations will be presented.  The paper is organised as follows. First, in Section 2 we present linear MRI channel modes in a vertically-stratified box and show that they are also nonlinear solutions in the limit of weak magnetic fields. These are directly computed numerically and then checked with one-dimensional stratified simulations in Section 3. Following that, in Section 4, the ideal MHD parasitic instabilities which could disrupt these channels are characterised and discussed. Our conclusions are drawn in Section 5, where we also briefly discuss the potential role of these solutions in real accretion disks. ",
        "conclusions": "In summary, we have shown that channel flows, a subset of the linear axisymmetric MRI modes, are approximate nonlinear solutions in the vertically-stratified shearing box, when in the limit of large midplane $\\beta$. This generalises the better known unstratified result (GX94), but only holds when the equilibrium consists of a net vertical magnetic flux. We compute stratified channel flows numerically for an isothermal model, and confirm the basic result with 1D simulations. The MRI eigenfunctions and growth rates, conversely, offer a valuable numerical benchmark on the performance of codes operating in the vertically-stratified shearing box. In addition, these simulations provide estimates (for future 3D simulations) of the resolution required to resolve the fastest growing channels when using second order numerical schemes: 25 cells per scaleheight are needed when $\\beta=100$, 50 cells when $\\beta=10^3$, and more than 200 cells per scale height are mandatory when $\\beta=10^4$.  Stratified channel flows, like their unstratified counterparts, are subject to a variety of magnetised shear flow instabilities, or `parasitic modes'. Some of these we directly compute in an intermediate limit of channel amplitude. As in the unstratified setting, the fastest growing parasite is the `unmagnetised' kink mode. But it is unlikely to destroy its host before equipartition fields have been reached.  At present we are extending these findings with 3D numerical simulations in stratified boxes, where the parasitic modes' behaviour can be confirmed, not only in the initial stages of a simulation but once it has relaxed into a saturated turbulent state. Because of the inevitably of large magnetic fields, and also finite resolution, these simulations will permit a set of interesting physical effects which we have neglected here but which deserve further study. These include, most importantly, magnetic buoyancy and magnetic reconnection. Also, these simulations will allow us to test the recurrent formation of channels once the MRI has saturated in turbulence (Suzuki and Inutsuka 2009). In particular, they can establish whether larger radial domains impede recurrent channel formation, by analogy with unstratified boxes (Bodo et al.~2008, LLB09).  Lastly, we speculate on the potential role dominant channel flows play in real magnetised accretion disks. Global disk models support analogues of the stratified MRI channels, localised in radius to narrow rings. Such modes might also exhibit a weaker version of the nonlinear property examined in this paper, permitting one mode to dominate the initial evolution of small amplitude initial conditions. Indeed, simulations with a poloidal field piercing the disk report such behaviour (Hawley 2000, 2001). In fully-developed turbulent disks, on the other hand, the dominance of a single channel will be rare. But occasional `run-away' channel flows might occur in the marginally stable fluid that divides regions of stable and unstable MRI --- dead-zone interfaces in protostellar disks, for instance. It may take only a small fluctuation in the ionisation and thermal properties of the gas to ignite instability in this previously quiescent critical layer, and hence launch a powerful MRI channel flow. Disruptive events of this kind may control the integral-scale intermittency that should characterise the dead-zone interface, and may even be related to the initial phases of outburst behaviour (Zhu et al.~2009). Similarly, channel flows may also feature in the onset of episodic accretion in dwarf nova disks (Gammie and Menou 1998)."
    },
    "1004/1004.0615_arXiv.txt": {
        "abstract": "We use the first compilation of 72 core-collapse supernovae (SNe) from the Palomar Transient Factory (PTF) to study their observed subtype distribution in dwarf galaxies compared to giant galaxies. Our sample is the largest single-survey, untargeted, spectroscopically classified, homogeneous collection of core-collapse events ever assembled, spanning a wide host-galaxy luminosity range (down to $M_r \\approx -14$ mag) and including a substantial fraction ($>20\\%$) of dwarf ($M_r \\geq -18$ mag) hosts. We find more core-collapse SNe in dwarf galaxies than expected and several interesting trends emerge. We use detailed subclassifications of stripped-envelope core-collapse SNe and find that all Type I core-collapse events occurring in dwarf galaxies are either SNe~Ib or broad-lined SNe~Ic (SNe~Ic-BL), while ``normal'' SNe~Ic dominate in giant galaxies. We also see a significant excess of SNe~IIb in dwarf hosts. We hypothesize that in lower metallicity hosts, metallicity-driven mass loss is reduced, allowing massive stars that would have appeared as ``normal'' SNe~Ic in metal-rich galaxies to retain some He and H, exploding as Ib/IIb events. At the same time, another mechanism allows some stars to undergo extensive stripping and explode as SNe~Ic-BL (and presumably also as long-duration gamma-ray bursts). Our results are still limited by small-number statistics, and our measurements of the observed $N(\\textrm{Ib/c})/N(\\textrm{II})$ number ratio in dwarf and giant hosts ($0.25_{-0.15}^{+0.3}$ and $0.23_{-0.08}^{+0.11}$, respectively; 1$\\sigma$ uncertainties) are consistent with previous studies and theoretical predictions. As additional PTF data accumulate, more robust statistical analyses will be possible, allowing the evolution of massive stars to be probed via the dwarf-galaxy SN population. ",
        "introduction": "Understanding the formation, evolution, and ultimate fate of massive stars is a key open question in astrophysics. A particularly interesting problem concerns the formation of highly stripped objects known as Wolf-Rayet (W-R) stars. It is still unclear whether W-R stars are formed through wind driven mass loss from single very massive stars, or if binary interactions play an important role in which case the minimal initial mass of W-R stars may be much lower. The final stages in the evolution of a massive star influence the type of explosion it will undergo at the end of its life; studying the resulting core-collapse supernovae (SNe) is thus an attractive way to learn about these massive progenitors. The hope is to reveal how the physical properties of the star (such as its mass and metallicity) determine its pre-explosion mass loss, as reflected in the properties of the ensuing SN. An additional benefit is that luminous SN explosions are visible to great distances and hence provide access to a wide range of stellar populations, while normal individual stars can be effectively studied only within the Local Group.  Core-collapse SN explosions of massive stars are observationally classified according to their spectroscopic properties (see Filippenko 1997 for a review). H-rich Type IIP (``plateau'') SNe have been associated with red supergiant progenitors (Smartt 2009, and references therein). The progenitors of Type IIb SNe appear to have lost most (but not all) of their hydrogen envelope before they exploded. Further mass loss leads to Type Ib explosions showing strong He lines but lacking H, while SNe~Ic, lacking both H and He, probably arise from the most highly stripped progenitors. A particular subclass of SNe~Ic, displaying very broad lines (SNe~Ic-BL) indicating high expansion velocities, is notable for the fact that its members are the only events to have been observationally linked with gamma-ray bursts (GRBs) and X-ray flashes (XRFs; see Woosley \\& Bloom 2006 for a review).  While it stands to reason that H-poor SNe arise from H-poor W-R stars, attempts to directly detect such stars in pre-explosion images of the locations of SNe~Ib and Ic have so far failed (Gal-Yam et al. 2005; Maund et al. 2005; Crockett et al. 2008b; Smartt 2009). The exact identity of the progenitor stars of SNe~IIb/Ib/Ic, as well as those of GRBs and XRFs, remains an unsolved puzzle.  Given the limited efficacy of direct pre-explosion imaging studies, one can attempt to study massive-star populations using statistical analyses of the resulting SNe. Two main methods have been employed for this purpose thus far, as follows.  The number ratio of SNe~Ib/c to SNe~II as a function of host-galaxy metallicity has been measured using host luminosity as a metallicity proxy for 280 and 497 core-collapse SNe (Prantzos \\& Boissier 2003; Boissier \\& Prantzos 2009, respectively), as well as via direct metallicity measurements for 159 core-collapse SNe (Prieto et al. 2008). Both methods show that the ratio of SNe~Ib/c to SNe~II increases with increasing host metallicity. This suggests that wind-driven mass loss plays an important role in the formation of stripped W-R stars exploding as SNe~Ib/c, and that at lower metallicities massive stars that would have ended as SNe~Ib/c explode instead as Type II events (``failed Ib/c''). Heger et al. (2003) predicted this behavior, but their decline rate for the SN~Ib/c ratio with decreasing metallicity, as well as that of Fryer et al. (2007), does not agree with the Boissier \\& Prantzos (2009) data.  The location of a SN in its host galaxy can also be indicative of the mass and metallicity of its progenitor. Anderson \\& James (2009) examined 173 core-collapse SN locations and confirmed previous results by van den Bergh (1997), Tsvetkov et al. (2004), and Hakobyan (2008) which indicated that SNe~Ib/c are preferentially found in more central regions of their host galaxies compared to SNe~II (see also Boissier \\& Prantzos 2009; Hakobyan 2009), with SNe~Ic the most centrally concentrated. This implies a progenitor-metallicity sequence (II-Ib-Ic) with SNe~Ic having the highest metallicity progenitors. Like the population-ratio studies, these results also suggest that SN~Ib/c progenitor analogs (in terms of mass) would explode as SNe~II in low-metallicity regions. Other location studies, however, indicate otherwise. Fruchter et al. (2006) found that long-duration GRBs occur preferentially in faint (metal-poor) galaxies (see also Stanek et al. 2006) and in brighter regions than the general core-collapse population (based on 16 core-collapse SN locations). Kelly et al. (2008) showed that SNe~Ic (26 in their sample) occur in areas within their hosts similar to those of long-duration GRBs. Taken together, these results, as well as those of Leloudas et al. (2010), would suggest that some SN~Ic progenitor analogs in low-metallicity environments tend to explode as SNe~Ic-BL (accompanied by long GRBs) rather than as SNe~II.  While the results from the studies described above are intriguing and hint at the power of the statistical analysis of large SN samples, a clear picture has yet to emerge. Perhaps this is due to the small number of events in dwarf, low-metallicity galaxies (leading Prantzos \\& Boissier 2003 and Boissier \\& Prantzos 2009, for example, to extend the ``dwarf'' galaxy bin to luminosities as bright as $-19$ mag). It should also be noted that many studies to date were based on heterogeneous compilations of SN data, inheriting their strong and unknown selection effects, and many included SNe discovered decades ago, the exact classification of which was unclear (Anderson \\& James 2009). Here we present a population-ratio analysis for the first sample of 72 core-collapse SNe from the Palomar Transient Factory (PTF; Rau et al. 2009; Law et al. 2009), 15 of which are in dwarf galaxies having $r$-band magnitudes $M_r \\geq -18$. The observational data are described in $\\S~2$ and the results in $\\S~3$. We discuss our findings in $\\S~4$ and summarize in $\\S~5$. ",
        "conclusions": "While small, the number of core-collapse SNe discovered in dwarf galaxies is higher than expected. Using rough scaling relations to compare PTF with the Pan-STARRS1 predictions presented by Young et al. (2008), it would appear that the number of such SNe is larger than predicted by a factor of a few. Two of these events, PTF10bfz and PTF10bzf (SN 2010ah; Ofek et al. 2010), are SNe~Ic-BL having no obvious association with a GRB. Along with the recent discovery of SN 2007bg (Moretti et al. 2007; Young et al. 2009) as a dwarf-host (Prieto et al. 2008) SN~Ic-BL without an associated GRB\\footnote{Neither SN 2007bg nor PTF10bzf show radio emissions indicative of the relativistic outflows present in GRBs (Soderberg 2009; Chomiuk \\& Soderberg 2010).}, these events suggest that not all SNe~Ic-BL in low-metallicity galaxies are associated with long-duration GRBs (cf. Young et al. 2008; Modjaz et al. 2008). However, without direct metallicity measurements, we cannot determine whether they violate the dividing line presented by Modjaz et al. (2008), as these hosts may be outliers to the metallicity vs. luminosity relation.  We find that the SN~Ib/c to SN~II number ratio is $N(\\textrm{Ib/c})/N(\\textrm{II})=0.23_{-0.08}^{+0.11}$ in giant galaxies and $N(\\textrm{Ib/c})/N(\\textrm{II})=0.25_{-0.15}^{+0.3}$ in dwarfs ($1\\sigma$ uncertainties\\footnote{These confidence levels have been calculated using the Clopper-Pearson method, discussed, for example, by Gehrels (1986).}). While this does not recover the reduction in the SN~Ib/c fraction in dwarf or low-luminosity hosts reported by Prantzos \\& Boissier (2003) and Prieto et al. (2008), the difference is not significant. Theoretical studies (Fryer et al. 2007) predict a fraction of $N(\\textrm{Ib/c})/N(\\textrm{II}) \\lesssim 0.1$ in dwarf hosts. Distinguishing between SN~Ib/c to SN~II ratios of 0.25 and 0.1 at the 95\\% confidence level requires a dwarf-host SN sample size of $\\geq57$ (assuming binomial statistics), which is expected to be attained by PTF within the next few years. The data to be gathered by PTF (as well as other surveys) will thus allow strong conclusions to be drawn in the near future.  Our distinction among SNe~Ib, Ic, and Ic-BL subtypes, as well as separating Types II and IIb, demonstrates the power of a detailed spectroscopic classification. We find that the only SNe~Ib/c in dwarf galaxies are either Ib or Ic-BL. ``Normal'' SNe~Ic, abundant in large galaxies, are not present, or are rare, in dwarf hosts. This is in agreement with past trends: all SNe~Ib/c in dwarfs studied by Prieto et al. (2008) were either SNe~Ib (SN 2005hm) or SNe~Ic-BL (SN 2006qk, SN 2007I, and SN 2007bg), while all GRB-associated SNe in dwarf hosts are also Ic-BL events (Woosley \\& Bloom 2006). Discoveries from the ROTSE-III surveys (Quimby 2006; Yuan et al. 2007) may include one normal SN~Ic hosted by a galaxy meeting our definition of a dwarf (SN 2004gk; Quimby et al. 2004), but the sample otherwise shows trends similar to the PTF results.  We further identify an overabundance of SNe~IIb and SNe~Ic-BL in dwarf galaxies. Assuming the fraction $p$ of SNe~IIb is independent of host luminosity, we can calculate the probability of finding $\\leq 2$ SNe~IIb in 55 bright ($M_r < -18$ mag) hosts, and $\\geq 3$ SNe~IIb in 15 dwarf ($M_r \\geq -18$ mag) hosts. This probability as a function of $p$ is given by \\[\\sum_{x=0}^{2}\\left(\\begin{array}{c}55\\\\x\\end{array}\\right)p^{x} \\left(1-p\\right)^{55-x} \\sum_{y=3}^{15}\\left(\\begin{array}{c}15\\\\y \\end{array}\\right)p^{y}\\left(1-p\\right)^{15-y}.\\] \\noindent Varying $p$ between 0 and 1, we find that this probability is $<5\\%$. A similar result is found for the overabundance of SNe~Ic-BL in dwarf hosts. It is therefore unlikely that the observed trends are produced by chance from the same underlying population.  Our data suggest that ``normal'' SNe~Ic occurring in giant galaxies are replaced by IIb and Ic-BL events in dwarf galaxies. This seems to confirm both trends suggested by previous location and population-ratio studies, implying two distinct mechanisms at work. Metallicity might be the main process responsible for the evolution of stars which explode as ``normal'' SNe~Ic. In lower metallicity environments, stars that would otherwise explode as SNe~Ic end their lives as Ib or IIb events, while many SN~Ib progenitors are also driven toward Type IIb explosions following a decrease in their mass loss. However, an additional process which allows for strong stripping in low-metallicity environments must be present to explain the SNe~Ic-BL and long GRB events observed in dwarf hosts. Binarity and rapid stellar rotation are possible candidates for such a stripping mechanism. Finally, we note an apparent excess of stripped (IIb+Ib+Ic+Ic-BL) SNe in dwarf galaxies, compared to ``normal'' SNe~II (see Figures \\ref{fighostmag} and \\ref{figpie}). If this is indeed the case, it is unlikely to be caused by mass-loss mechanisms, which are expected to be stronger in high-metallicity environments. A top-heavy initial mass function (IMF) in dwarf galaxies might be required to explain this effect (see also Hakobyan et al. 2009; Habergham et al. 2010). Such an IMF is also consistent with pair-production SNe (Gal-Yam et al. 2009) and highly luminous SNe (Quimby et al. 2009), which are associated with very massive progenitors, occurring only in dwarf hosts.  Our analysis is limited by small-number statistics. Continued observations by PTF are expected to rapidly increase the sample size during the coming few years. An attractive complementary study is to analyze the locations of our SNe within their host galaxies (e.g., as done by Fruchter et al. 2006 and Kelly et al. 2008 on similar-sized samples). A combined analysis should provide a powerful constraint on the evolution and mass-loss properties of massive stars as a function of their metallicity. The distance and angular size of our host sample require the high spatial resolution of the {\\it Hubble Space Telescope} to carry out this test."
    },
    "1004/1004.0753_arXiv.txt": {
        "abstract": "We report on observations of the superbubbles (SBs) N11 and N51D in the Large Magellanic Cloud (LMC) with {\\it Suzaku} and {\\it XMM-Newton}. The interior of both SBs exhibits diffuse X-ray emission, which is well represented by thin thermal plasma models with a temperature of 0.2--0.3~keV. The presence of nonthermal emission, claimed in previous works, is much less evident in our careful investigation. The 3$\\sigma$ upper limits of 2--10~keV flux are $3.6 \\times 10^{-14}$~\\flux\\ and $4.7 \\times 10^{-14}$~\\flux\\ for N11 and N51D, respectively. The previous claims of the detection of nonthermal emission are probably due to the inaccurate estimation of the non X-ray background. We conclude that no credible nonthermal emission has been detected from the SBs in the LMC, with the exception of 30~Dor~C. ",
        "introduction": "\\label{sec:introduction}   \\begin{table*} \\begin{center} \\caption{Log of observations appearing in this paper. \\label{tab:log}} \\begin{tabular}{llccccc} \\tableline\\tableline Target & Telescope/ & Observation & \\multicolumn{2}{c}{Aim point} & Observation & Effective  \\\\ name   & Instrument & ID  & $\\alpha$(J2000.0) & $\\delta$(J2000.0) & start date & exposure (ks) \\\\ \\tableline N11 & Suzaku/XIS  & 501091010 & $04^{\\rm h}56^{\\rm m}48^{\\rm s}$ & $-66^{\\circ}24'38''$ & 2006 Nov 7 & 30.5  \\\\ ~ & XMM-Newton/EPIC  & 0109260201 & $04^{\\rm h}55^{\\rm m}58^{\\rm s}$ & $-66^{\\circ}25'53''$ & 2001 Sep 17 & 32.5 (MOS)/25.1 (pn)  \\\\ N51D & Suzaku/XIS  & 804009010 & $05^{\\rm h}25^{\\rm m}37^{\\rm s}$ & $-67^{\\circ}29'57''$ & 2009 May 17 & 85.5/87.1/87.3\\tablenotemark{a}  \\\\ ~ & XMM-Newton/EPIC  & 0071940101 & $05^{\\rm h}26^{\\rm m}03^{\\rm s}$ & $-67^{\\circ}28'52''$ & 2001 Oct 31 & 30.5 (MOS)/23.1 (pn)  \\\\ NEP & Suzaku/XIS  & 500026010 & $18^{\\rm h}11^{\\rm m}18^{\\rm s}$ & $+66^{\\circ}00'45''$ & 2006 Feb 10 & 88.5  \\\\ \\tableline \\end{tabular} \\tablecomments{ $^{a}$The exposure times of the XIS0/1/3. The XIS2 was out of operation during this observation. } \\end{center} \\end{table*}  Massive stars are often found clustered in OB associations. Their high-speed stellar winds and the blast wave from core-collapsed supernova (SN) explosions collectively sweep up the ambient interstellar medium (ISM) to generate large (10--100-pc-scale) shell-like structures known as superbubbles (SBs) (e.g., McCray \\& Snow 1979; Mac Low \\& McCray 1988). A tenuous cavity created inside an SB shell allows the blast wave of a subsequent supernova remnant (SNR) to propagate rapidly without deceleration for a long time. Furthermore, wind-wind and/or shock-cloud interactions, expected inside an SB, can maintain significant turbulence and magnetic inhomogeneities. Thus, SBs are believed to be very efficient site for the acceleration of relativistic cosmic rays (e.g., Bykov \\& Fleishman 1992; Parizot et al.\\ 2004).   The Large Magellanic Cloud (LMC) is particularly suitable for the study of SBs, because of its low foreground extinction \\citep{dl} and its known distance of $\\sim$50~kpc with a small uncertainty (e.g., Persson et al.\\ 2004). Using {\\it ROSAT}, Chu \\& Mac Low (1990) and \\cite{dunne01} systematically studied several SBs in the LMC. They found that the cavities inside the SBs exhibit diffuse X-ray emissions from rarefied hot gas with a temperature on the order of 0.1~keV. They also showed that the X-ray surface brightnesses are significantly higher than those expected from the standard pressure-driven SB model (Weaver et al.\\ 1977). This suggests the presence of interior SNRs colliding with the inner walls of the SB shell.   The detection of nonthermal X-rays has been reported in three LMC SBs, 30~Dor~C \\citep{bamba04}, N11 \\citep{m09}, and N51D \\citep{c04}. 30~Dor~C exhibits the nonthermal emission with a clear shell-like structure. The spectra obtained with {\\it Chandra} and {\\it XMM-Newton} were represented by a power-law with a photon index of $\\Gamma$ = 2--3 \\citep{bamba04}, typical for synchrotron X-rays from nonthermal SNRs. Using {\\it Suzaku}, \\cite{yama09} discovered the spectral cutoff in the power-law emission. They concluded that the synchrotron X-rays originate from the rapidly-expanding shell of the SNR with an age of (4--20) $\\times 10^3$~yr.   On the other hand, the validities of the other two results may be doubtful. \\cite{m09} have recently observed N11 with {\\it Suzaku} and have reported the detection of nonthermal X-rays above 2~keV. However, the claimed flux is comparable to that of 30~Dor~C despite the fact that no nonthermal component was detected by the previous {\\it XMM-Newton} observation \\citep{naze04}. Moreover, a morphology of the emission has been unclear. The {\\it Suzaku} data should be therefore carefully reanalyzed.   The nonthermal emission in N51D is claimed to be distributed entirely in the SB cavity with a centrally peaked morphology in contrast to 30~Dor~C \\citep{c04}. The photon index of $\\Gamma \\sim 1.3$ is much lower than those expected for synchrotron X-rays. On the other hand, in the preceding work by \\cite{bomans03}, the same {\\it XMM-Newton} spectrum was well represented by a thermal plasma model without any additional nonthermal components. Thus, the existence of nonthermal emission in N51D is still controversial. Motivated by this issue, we have obtained a long-time exposure of this SB with {\\it Suzaku}.   The aim of this paper is to investigate for the existence of diffuse hard X-ray emission. Therefore, other properties of SBs and their stellar contents are not focused on. For the analyses, the {\\it Suzaku} data are mainly utilized, since these provide a low and stable non-X-ray background (NXB), particularly for diffuse sources (Mitsuda et al.\\ 2007; Tawa et al.\\ 2008). In addition, we also refer to {\\it XMM-Newton} archival data to compensate the limited angular resolution of {\\it Suzaku}.   First, we briefly introduce the characteristics of N11 (\\S\\ref{ssec:intro_n11}) and N51D (\\S\\ref{ssec:intro_n51d}). In Section~\\ref{sec:observation}, we outline our observation and data reduction procedure. The results for N11 and N51D are presented in Sections~\\ref{sec:n11} and \\ref{sec:n51d}, respectively. We discuss our results in Section~\\ref{sec:discussion}, and give a summary in Section~\\ref{sec:summary}. The errors quoted in the text and tables refer to the 90\\% confidence intervals for one parameter, and the error bars in the figures are for 1$\\sigma$ confidence, unless otherwise stated.     \\subsection{Characteristics of N11} \\label{ssec:intro_n11}  N11 (DEM L~34), located at the northwest edge of the LMC \\citep{henize56}, contains four OB associations, LH9, LH10, LH13, and LH14 \\citep{lh}. The \\Ha\\ emission exhibits a complex structure consisting of several shells and filaments (e.g., Mac~Low et al.\\ 1998). The largest \\Ha\\ shell with a diameter of $\\sim$120~pc is situated around the central cluster LH9. The age of LH9 is estimated to be $\\sim$3.5~Myr (Walborn et al.\\ 1999), the oldest among the four clusters. It is suggested that the stellar activity of LH9 triggered a starburst at the periphery, leading to the formation of the other OB associations (e.g., Walborn \\& Parker 1992; Hatano et al.\\ 2006).   \\subsection{Characteristics of N51D} \\label{ssec:intro_n51d}  N51D (DEM L~192) is located at the southeast corner of the supergiant shell LMC4. Two OB associations, LH51 and LH54 \\citep{lh}, are surrounded by a filamentary SB shell with a size of about 140~pc $\\times$ 120~pc. Both the OB clusters are estimated to have an age of $\\sim$3~Myr (Oey \\& Smedley 1998). ",
        "conclusions": "\\label{sec:discussion}   \\begin{figure*} \\begin{center} \\rotatebox{270}{ \\includegraphics[scale=.31]{fig11a.eps} } \\rotatebox{270}{ \\includegraphics[scale=.31]{fig11b.eps} } \\caption{XIS-FI unfolded spectra (black crosses) of N11 (a) and N51D (b). The EPIC-pn spectra of the point sources are also shown with gray diamonds. Solid and dashed gray lines are the model components of the diffuse and point sources, respectively. The spectra of ``nonthermal emission'' claimed by \\cite{m09} and \\cite{c04} (for N11 and N51D, respectively) are indicated with the black thick lines. The data points in the energies above $\\sim 2$~keV are obviously far below the claimed fluxes although the point sources are not subtracted from the XIS spectra. \\label{fig:unfold}} \\end{center} \\end{figure*}    The X-ray spectra of the SBs N11 and N51D are well reproduced by single-component thin-thermal plasmas (with the models for the group of point sources). Figure~\\ref{fig:unfold} shows the unfolded spectra of both SBs. The presence of additional nonthermal emission, argued in previous works (N11: Maddox et al.\\ 2009; N51D: Cooper et al.\\ 2004), is much less evident in our detailed analyses, supporting the earlier results by \\cite{naze04} and \\cite{bomans03}. The 2--10~keV flux of the diffuse component $F_{\\rm D}$ ($\\pm~\\sigma_{\\rm D}$) for each SB can be calculated as \\begin{eqnarray} F_{\\rm D} &=& \\left[ F_{\\rm S} -  F_{\\rm PS} - \\frac{\\Omega_{\\rm S}}{\\Omega_{\\rm B}} \\cdot F_{\\rm B} \\right] \\pm \\sigma_{\\rm D}  \\\\ \\sigma_{\\rm D} &=& \\left[ \\sigma_{\\rm S}^2 +  \\sigma_{\\rm PS}^2 + \\left( \\frac{\\Omega _{\\rm S}}{\\Omega _{\\rm B}} \\right)^2 \\sigma_{\\rm B}^2 \\right]^{1/2}, \\end{eqnarray} where $F_{\\rm S}$, $F_{\\rm B}$, and $F_{\\rm PS}$ (and $\\sigma_{\\rm S}$, $\\sigma_{\\rm B}$, $\\sigma_{\\rm PS}$) are the 2--10~keV fluxes (and uncertainties) of the SRC and BGD regions and that of the point sources. Using the values given in Tables~\\ref{tab:sbn_n11} and \\ref{tab:sbn_n51d}, we obtain 3$\\sigma$ upper limits of the fluxes to be $3.6 \\times 10^{-14}$~\\flux\\ and $4.7 \\times 10^{-14}$~\\flux\\ for N11 and N51D, respectively.   For N11, \\cite{m09} claimed that the hard X-rays, detected at energies of up to $\\sim$7~keV, required a power-law component with a photon index of $\\Gamma \\sim 1.7$ and an observed flux of $\\sim$$3.0\\times 10^{-13}$~\\flux\\ (in the 0.2--10~keV band)\\footnote{The energy range from which the flux was derived is not given in \\cite{m09}. Therefore, we independently estimated it from reported values of the absorption column density, photon index, and the ratio of the observed to unabsorbed fluxes.}$\\!\\!$. For N51D, the photon index and observed flux in the 0.3--3.0~keV band were reported to be $\\Gamma \\sim 1.3$ and $\\sim$$3.3\\times 10^{-13}$~\\flux, respectively \\citep{c04}. To directly compare these values with our results, these fluxes are modified to values in the 2--10~keV band, giving $2.3\\times 10^{-13}$~\\flux\\ for N11 and $7.1\\times 10^{-13}$~\\flux\\ for N51D. These values are about one order higher than the upper limits that we have obtained. The spectral models for the claimed hard X-ray emissions are shown in Figure~\\ref{fig:unfold} with black solid lines. Note that the point sources had been excluded from the EPIC spectrum of N51D \\citep{c04}. Nevertheless, both models give fluxes clearly exceeding those given by the XIS spectrum in which the point-source fluxes are contained.   Given that the point-source flux in N11 is negligibly small (Fig.~\\ref{fig:unfold}), the detection of hard X-rays claimed in the previous report was probably due to incorrect background subtraction. In fact, \\cite{m09} commented that the subtraction of Ni lines originating from the NXB was unsuccessful. The NXB spectrum of the XIS consists of emission-line and continuum components, and their intensities are correlated with each other \\citep{tawa08}. Therefore, this also caused the incomplete subtraction of the continuum component in the previous reports. In our analyses, on the other hand, the fluorescent NXB lines are accurately subtracted (see Section~\\ref{sssec:n11_suzaku1}), and hence the background reduction is considered to be successful. The problem concerning N51D is analogous to the case of N11, because the residual of the Al-\\Ka\\ line due to incomplete background subtraction has been reported \\citep{c04}. We point out that the analysis of faint extended sources, such as SBs in the LMC, is very sensitive to the estimation of the NXB and the vignetting effect, and thus should be performed with great care.   \\begin{table} \\begin{center} \\caption{Comparison of nonthermal luminosities with other sources. \\label{tab:comparison}} \\begin{tabular}{lcc} \\tableline\\tableline Object &  $L_{\\rm X}$~(\\lumi)\\tablenotemark{a}  &  Reference  \\\\ \\tableline N11  & $<$ $1 \\times 10^{34}$  & This work. \\\\ N51D & $<$ $1 \\times 10^{34}$  & This work. \\\\ 30~Dor~C & $2 \\times 10^{35}$ & Yamaguchi et al.\\ (2009) \\\\ Arches & $1 \\times 10^{34}$ & Tsujimoto et al.\\ (2007) \\\\ NGC~6334 & $8 \\times 10^{32}$ & Ezoe et al.\\ (2006) \\\\ RCW38 & $5 \\times 10^{32}$ & Wolk et al.\\ (2002) \\\\ Westerlund~1 & $3 \\times 10^{34}$ & Muno et al.\\ (2006) \\\\ Westerlund~2 & $2 \\times 10^{34}$ & Fujita et al.\\ (2009) \\\\ \\tableline \\end{tabular} \\tablecomments{ $^{a}$Luminosity in the 2--10~keV band. } \\end{center} \\end{table}    Nonthermal X-ray emissions have so far been discovered from several Galactic star-forming regions, the Arches cluster (Tsujimoto et al.\\ 2007), NGC~6334 (Ezoe et al.\\ 2006), RCW~38 (Wolk et al.\\ 2002), Westerlund~1 (Muno et al.\\ 2006), and Westerlund~2 (Fujita et al.\\ 2009). In these cases, however, the emissions are correlated with relatively compact (1--10-pc-scale) regions around an OB association or in a molecular cloud core. 30~Dor~C is another example of a star-forming region that exhibits nonthermal (synchrotron) X-rays \\citep{bamba04}. However, the X-ray shell can be interpreted as a single middle-aged SNR expanding rapidly inside the SB cavity \\citep{yama09}. Therefore, at present, there is no evidence for large-scale ($\\sim$100~pc) nonthermal emission distributed entirely in an SB. As given in Table~\\ref{tab:comparison}, the luminosities of the nonthermal X-ray emissions in N11 and N51D are, if any, less than those of 30~Dor~C and some of the other star-forming regions.    It is theoretically expected that efficient cosmic-ray acceleration can take place in SBs due to repeated SN shocks and/or magnetic turbulence (e.g., Bykov \\& Fleishman 1992; Parizot et al.\\ 2004). Moreover, Butt \\& Bykov (2008) suggested that more than 10\\% of the energy stored in an SB can be constantly transferred to accelerating cosmic-ray particles at the evolution stage of N11 and N51D (a few Myr). If this is the case, accelerated electrons would quickly lose their energies via radiation so that no nonthermal X-rays would be detected. It is still an open issue whether or not SBs can efficiently accelerate cosmic-rays, but this is beyond the scope of this paper. Future detailed observations and the development of theoretical models are required to solve this problem.    \\newpage    We have analyzed the {\\it Suzaku} XIS data of the extended X-ray emission from the SBs N11 and N51D. The analysis of the {\\it XMM-Newton} EPIC data has also been performed to quantify the contribution from the point-like sources. In both SBs, the spectra of the diffuse components are well represented by thin thermal plasma at collisional ionization equilibrium. The electron temperatures derived for N11 and N51D are $\\sim$0.3~keV and $\\sim$0.2~keV, respectively. The elemental abundances are comparable with average LMC values, suggesting that the thermal emission mainly originates from the shock-heated ISM with only slight SN enrichment.   Neither N11 nor N51D show any evidence for hard X-ray emission, contrary to the previous claims by \\cite{m09} and \\cite{c04}. The 3$\\sigma$ upper limits of the 2--10~keV flux are $3.6 \\times 10^{-14}$~\\flux\\ for N11 and $4.7 \\times 10^{-14}$~\\flux\\ for N51D. The published claims of the detection of nonthermal X-rays are thought to be simply due to inadequate background subtraction. Careful analyses (i.e., NXB and point-source subtraction, vignetting correction, and so on) are required to obtain an accurate flux for faint extended sources, such as N11 and N51D. As the result of this work, no observational evidence for the nonthermal X-ray emission widely associated with SBs remains, with the exception of 30~Dor~C (however, this is likely to be a single remnant of a recent SN explosion). It is still an open issue whether or not SBs can efficiently accelerate cosmic-rays. Observations of more SBs with longer exposures as well as the development of theoretical calculations are required."
    },
    "1004/1004.0279_arXiv.txt": {
        "abstract": "The gyrokinetic simulation code \\T{AstroGK} is developed to study fundamental aspects of kinetic plasmas and for applications mainly to astrophysical problems. \\T{AstroGK} is an Eulerian slab code that solves the electromagnetic gyrokinetic-Maxwell equations in five-dimensional phase space, and is derived from the existing gyrokinetics code \\T{GS2} by removing magnetic geometry effects. Algorithms used in the code are described.  The code is benchmarked using linear and nonlinear problems. Serial and parallel performance scalings are also presented. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.5389_arXiv.txt": {
        "abstract": "In this paper we present simulated observations of massive self-gravitating circumstellar discs using the Atacama Large Millimetre/sub-millimetre Array (ALMA).  Using a smoothed particle hydrodynamics model of a $0.2\\msolar$ disc orbiting a $1\\msolar$ protostar, with a cooling model appropriate for discs at temperatures below $\\sim 160$K and representative dust opacities, we   have constructed maps of the expected emission at sub-mm wavelengths. We have then used the CASA ALMA simulator to generate simulated images and visibilities with various array configurations and observation frequencies, taking into account the expected thermal noise and atmospheric opacities.  We find that at 345 GHz (870 $\\mu$m) spiral structures at a resolution of a few AU should be readily detectable in approximately face-on discs out to distances of the Taurus-Auriga star-forming complex. ",
        "introduction": "\\label{intro} Circumstellar discs play an important role in the formation and evolution of both stars and planets, and as such have been the object of much study and observation in recent years.  Millimetre wavelength surveys of star forming regions in Taurus \\citep{Beckwithetal90,Kitamuraetal02,AndrewsWilliams05}, Orion \\citep{Eisneretal08} and $\\rho$ Ophiucus \\citep{AndreMontmerle94,AndrewsWilliams07a} have provided extensive evidence of circumstellar discs, while optical images have been provided by the Hubble Space Telescope (HST) of HH30 \\citep{Burrowsetal96} and the Fomalhaut system \\citep{Kalasetal08}.  However, relatively few systems have been imaged with great enough resolution to determine the disc structure on scales of less than a few tens of $AU$.  This may be about to change however, with the Atacama Large Millimetre/sub-millimetre Array (ALMA) due to come on line in the near future. With a minimum beam diameter of approximately 5 milli-arcseconds at 900 GHz, ALMA should ideally provide resolution down to $\\approx 2 AU$ for discs observed in Orion, with sub-$AU$ resolution for systems in Taurus-Auriga.  It will therefore allow observations of disc sub-structure in unprecedented detail.  The discovery of structures within discs could have important implications for our understanding of the evolution of both the discs themselves \\citep{RiceA09} and the protostars they orbit.  Gaseous circumstellar discs are expected to be turbulent \\citep{Ebert94,Gammie96}, with the internal stresses that this turbulence induces being the principal driver of angular momentum transport, and thus accretion.  This turbulence will lead to sub-structure being present at all scales within the disc, with those on smaller scales (due to the magneto-rotational instability) remaining undetectable, while the larger scale structure induced by the gravitational instability should be resolvable with ALMA.  It is expected that protostellar discs about Class 0/Class I objects will go through a self-gravitating phase \\citep{LodatoBertin01a,VorobyovBasu05c,Hartmannbook} as the infall rate from the gaseous envelope is much greater than the accretion rate onto the protostar \\citep{VorobyovBasu06}.  In this phase large amplitude spiral structures will form, driving accretion onto the protostar \\citep{LinPringle87,LaughlinBodenheimer94,ArmitageLP01,LodatoR04,LodatoR05} and potentially leading to the formation of companion objects.  These companions may vary in size from stellar mass \\citep{Bonnell94,BonnellBate94,Whitworthetal95} to brown dwarf (BD) or planetary mass \\citep{RiceLA05,StamatellosHW07,Clarke09,Cossinsetal10}.  It is possible that protoplanetary discs around Classic T-Tauri (Class II) stars may also exhibit spiral patterns due to the presence of gravitational instabilities \\citep{LodatoBertin01a,Boleyetal06,VorobyovBasu08,Cossinsetal10}, and indeed there are already possible detections of spiral structures in the discs of GSS 39 in Ophiuchus \\citep{Andrewsetal09} and in IRAS 16293-2422B \\citep{Rodriguezetal05}.  At radii greater than approximately $50 AU$, planet formation through direct fragmentation of these spiral overdensities into bound objects is possible \\citep{Boss97,Boss98,Clarke09,Rafikov09,Cossinsetal10}.  While \\citet{WolfD'Angelo05} have indicated that the forming giant (proto-)planets themselves may be detectable using ALMA, the observability of the large-scale spiral structure itself within protostellar and protoplanetary discs, though implied \\citep{TestiS08}, has remained undemonstrated.  In this paper we therefore present a simple self-gravitating disc simulation, and from it we derive mock observations of disc systems at the resolutions and sensitivities that should be possible with ALMA.  Hence in \\sref{simulations} we detail the simulation, and in \\sref{generation} we discuss how the mock observations are generated from it, taking into account telescope effects and sensitivities.  In \\sref{results} we present the observations for various system and telescope parameters, and finally in \\sref{discussion} we discuss the significance of our results. ",
        "conclusions": "\\label{discussion}  In this paper, we have used a 3D, global SPH simulation of a massive ($0.2\\msolar$) compact ($R_{\\mathrm{out}} = 25$ AU) self-gravitating disc about a young star ($1.0\\msolar$) to demonstrate that the spiral modes excited by the gravitational instability should be detectable in face-on circumstellar discs using ALMA.  At distances comparable to the TW Hydrae association ($\\sim 50$ pc), such spiral density waves are readily apparent with observation times of 2 hours, whereas at Taurus distances of $\\sim 140$ pc a careful choice of the observing frequency and excellent observing conditions  may be required for significant detections.  Our results suggest that structure in such discs in Orion ($\\sim 400$pc) will most likely not be resolvable.  In order to generate these predicted observations we have used temperature and density maps provided by numerical simulations, together with an empirical relationship for the dust opacity of circumstellar material to obtain the optical depth at each part of the disc face.  In a precisely similar manner to that used to obtain disc properties from sub-mm observations \\citep{Beckwithetal90} we have then been able to determine the specific intensity of the disc emission across the disc face, which has then been used as input for the ALMA simulator in CASA to produce simulated observations.  We find that observation frequencies of around 345 GHz (870 $\\mu$m) are ideal for this kind of observation, as we find the spiral arms to be optically thick, whereas the inter-arm regions are optically thin, maximising the contrast in emission between the regions.  It should be noted however that this `ideal' frequency range is dependent on our assumptions concerning the dust opacity, in which there is considerable scatter.  There are certain limitations to our model which should be noted however.  Our stellar and disc masses are both at the upper end of the expected distributions \\citep{AndrewsWilliams05,AndrewsWilliams07a,Beckwithetal90}, although in the early stages of star formation (Class I objects) such high disc masses are not unreasonable due to the infalling envelope.  Of necessity, in order to be self-gravitating our discs are cold ($20 - 40$ K), which implies both a relatively low ambient temperature (not unreasonable, since giant molecular clouds tend to have temperatures of $\\sim 10K$, \\citealt{Myersetal83,MyersBenson83}) and for heating from the protostar to be negligible.  While this latter assumption is clearly unlikely to be valid for the surface layers of the disc that are irradiated directly by the star, the disc midplane (which dominates the emission) is likely to be cold enough to justify this assumption \\citep{AndrewsWilliams05,Dullemondprpl,D'Alessioetal98}.  In a similar manner, this assumption of a colder inner layer allows us to ignore the effects of the magneto-rotational instability \\citep{Gammie96}.  Given that our disc is quite compact, and many discs are observed to extend out to much larger radii ($\\sim 10^{2} - 10^{3}$ AU, \\citealt{Andrewsetal09,Eisneretal08,AndrewsWilliams07a,Kitamuraetal02}) the colder regions outside of $\\sim 50$ AU are if anything more likely to show evidence of gravitational instability that at the radii we simulate, and indeed any spiral structures forming at large radii will be more easily resolvable.  In this sense therefore our predictions may even be conservative in terms of the maximum distance at which spiral structures may be detectable.  Likewise our cooling prescription, although simplistic, is valid for regions at large radii where the temperature is below the ice sublimation temperature, and is therefore not an unreasonable simplification. It should be noted that the ratio of the cooling time to the dynamical time $\\beta=t_{\\rm cool}\\Omega$ determines the amplitude of the spiral perturbation and hence the contrast in the simulated images. The value $\\beta=7$ adopted in this paper is in the right range for discs at a few tens of AU. However, a larger value of $\\beta$ (that is, a less efficient cooling) would provide a relatively smaller contrast in the ALMA images.  Finally, note that in this paper we have considered the contribution to the sub-millimetre emission due solely by the disc. In the earliest phases of star formation, the system might show a substantial emission on larger scales, produced by the infalling envelope feeding the disc. This larger scale contribution has been neglected in the present paper.  As noted in the Introduction there are already possible detections of spiral structures in the discs of GSS 39 in Ophiuchus \\citep{Andrewsetal09} and in IRAS 16293-2422B \\citep{Rodriguezetal05}.  However the structures in GSS 39 are not robust at the 3$\\sigma$ level \\citep{Andrewsetal09}, and those in IRAS 16293-2422B, whilst appearing to be genuine, may plausibly be due to interactions with a companion.  Furthermore, the presence of planetary mass objects within the disc may also lead to the excitation of spiral density waves \\citep[see for instance][]{GoldreichTremaine80,OgilvieLubow02,ArmitageRice05}. However in this latter case the one-armed spiral wakes are in general tightly wound (dependent on the disc thickness, \\citealt{OgilvieLubow02}) and would therefore be difficult to detect directly.  Furthermore, analysis by \\citet{WolfD'Angelo05} on the observability of embedded gas giant planets using ALMA suggests that the spirals would only be visible very close to the planet itself.  Any detections of large scale spirals in isolated protostellar discs are therefore likely to be due principally to the action of the gravitational instability.  Confirmed, unambiguous observations of gravitationally induced spiral structures within protostellar discs would be valuable for a number of reasons.  As the gravitational instability is expected to operate during the early phases of star formation, processing the infalling envelope and allowing rapid accretion onto the protostar, such detections would validate this mechanism for growing the masses of protostars. More controversially, it may enable models of brown dwarf/low mass stellar companion formation through disc fragmentation \\citep{BateBB02,StamatellosHW07,StamatellosW09,Clarke09} to be validated, as the presence (or otherwise) and amplitudes of spiral arms would allow constraints to be placed on the numbers of such companions that may be expected.  In a similar vein, detections of spiral features may enable us to determine the dominant mode of planet formation at various radii, about which there is some debate.  While the standard core-accretion model is favoured at low radii \\citep{Boley09,Klahr08,Bodenheimeretal00,Lissauer93}, direct fragmentation of gravitationally induced spirals remains a candidate mechanism at radii above $\\sim 50$ AU \\citep{Boss97,Boleyetal06,Rafikov09,Cossinsetal10}. As the presence of a peak in the 1.3 cm emission around HL Tau has been put forward as a promising candidate for a planetary mass companion formed through gravitational instability, detections of the spiral wave progenitors of such companions would provide significant backing for this mechanism.  Moreover, the presence of large amplitude spiral density perturbations may be important for the formation and growth of planetesimals, both due to the concentration of the dust fraction within the arms \\citep{RiceLPAB04,RiceLPAB06,ClarkeL09} and further due to the possible scattering of planetesimals by the spiral potential \\citep{BritschCL08}.  In either case, observations of the spiral arms themselves would place constraints on, and therefore allow us to discriminate between, the two planet formation modes at large radii.  \\begin{figure} \\centering \\includegraphics[width=15pc]{f_345ghz_410pc.eps} \\caption{Simulated 345 GHz aperture synthesis image of a disc at 410~pc, with a transit duration of 2 hours.  Contours start at and are spaced by 60 mJy/beam.  Clear asymmetries are present, but the underlying spiral structure is not well resolved.} \\label{fsimul410} \\end{figure}"
    },
    "1004/1004.4988_arXiv.txt": {
        "abstract": "Acetic acid (CH$_3$COOH) has been detected mainly in hot molecular cores where the distribution between oxygen (O) and nitrogen (N) containing molecular species is co-spatial within the telescope beam. Previous work has presumed that similar cores with co-spatial O and N species may be an indicator for detecting acetic acid. However, does this presumption hold as higher spatial resolution observations become available of large O and N-containing molecules? As the number of detected acetic acid sources is still low, more observations are needed to support this postulate. In this paper, we report the first acetic acid survey conducted with the Combined Array for Research in Millimeter-wave Astronomy (CARMA) at 3 mm wavelengths towards G19.61-0.23, G29.96-0.02 and IRAS 16293-2422. We have successfully detected CH$_3$COOH via two transitions toward G19.61-0.23 and tentatively confirmed the detection toward IRAS 16293-2422 A. The determined column density of CH$_3$COOH is 2.0(1.0)$\\times$10$^{16}$ cm$^{-2}$ and the abundance ratio of CH$_3$COOH to methyl formate (HCOOCH$_3$) is 2.2(0.1)$\\times$10$^{-1}$ toward G19.61-0.23. Toward IRAS 16293 A, the determined column density of CH$_3$COOH is $\\sim$1.6$\\times$10$^{15}$ cm$^{-2}$ and the abundance ratio of CH$_3$COOH to methyl formate (HCOOCH$_3$) is $\\sim$1.0$\\times$10$^{-1}$ both of which are consistent with abundance ratios determined toward other hot cores. Finally, we model all known line emission in our passband to determine physical conditions in the regions and introduce a new metric to better reveal weak spectral features that are blended with stronger lines or that may be near the 1-2$\\sigma$ detection limit. ",
        "introduction": "With high resolution observations, Sgr B2 has been resolved into many regions of  compact molecular/continuum emission \\citep{gaume95}; one of these regions, Sgr B2(N-LMH) (i.e., the ``Large Molecule Heimat'') \\citep{snyder94,miao95}, is a  molecule-rich hot core that has historically been the best place to search for large and complex molecules. Interstellar acetic acid (CH$_{3}$COOH) was first detected in the hot  molecular core (HMC) of Sgr B2(N-LMH) with the Berkeley-Illinois-Maryland Association (BIMA) and Owens Valley Radio Observatory (OVRO) millimeter-wavelength arrays by \\citet{meh97}. The compactness of the CH$_{3}$COOH emission in Sgr B2(N-LMH) ($<$3\\arcsec) \\citep[e.g.,][]{remi02} highlights the importance of interferometers. To clearly detect the compact emission features from CH$_{3}$COOH and  other large molecules toward hot cores embedded in regions of more extended molecular emission, high resolution interferometric observations are necessary. Thus, following the first CH$_{3}$COOH detections in Sgr B2(N-LMH) and using the BIMA array, \\citet{remi02} confirmed CH$_{3}$COOH in Sgr B2(N-LMH) and discovered the second CH$_{3}$COOH source, the hot core toward W51e2.  \\citet{remi03} launched an extensive CH$_{3}$COOH survey of 12 galactic HMCs, including high-mass ($\\geq$ 10 M$_\\sun$) and low-mass ($<$ 10 M$_\\sun$) sources, and they found the third CH$_3$COOH source toward the high-mass HMC G34.3+0.15. Using the IRAM 30-m telescope, \\citet{cazaux03} reported the first detection of CH$_{3}$COOH in IRAS 16293-2422 via the $9_{*,9}-8_{*,8}$ $E$ line at 100.855 GHz. This is the first low-mass star forming region with detected CH$_{3}$COOH to date, making only a total of four known interstellar sources of CH$_{3}$COOH. At the conclusion of the search for new sources of CH$_{3}$COOH, it was found that the three high mass regions containing CH$_{3}$COOH emission share certain chemical properties. For example, the abundances of CH$_{3}$COOH and its isomer methyl formate (HCOOCH$_3$) are very different; in Sgr B2(N-LMH), the abundance of HCOOCH$_3$ is 26 times higher than that of CH$_3$COOH \\citep{snyder06} assuming the emission from CH$_3$COOH and HCOOCH$_3$ are co-spatial within the telescope beam. In addition, it has been suggested by \\citet{remi04} that there is a relationship between the CH$_3$COOH sources and the spatial separation of complex oxygen (O) (e.g.\\ HCOOCH$_3$ and CH$_3$COCH$_3$) and nitrogen (N) containing species  (e.g.\\ CH$_3$CH$_2$CN and NH$_2$CH$_2$CN). At the limit of the current observational resolution of these surveys, CH$_3$COOH emission is only detected where the emission from O and N-containing species is co-spatial. In contrast, there have not been any CH$_{3}$COOH detections in  HMCs with reported spatial separation of complex O and N-containing species, e.g. Orion KL \\citep[e.g.][]{sutt95} and W3(OH) \\citep[e.g.][]{wyro99}, even though the observed separation is very dependent on source distance.  Gas phase and grain surface chemistry may both play critical roles in the formation of CH$_{3}$COOH and other large molecules in the ISM \\citep{hase92,garr06,garr08}.  However, the exact pathways of CH$_{3}$COOH formation in gas phase or on grain  surfaces in astronomical environments are still not clear. The apparent spatial coincidence of the emission of complex O and N-containing species with the CH$_3$COOH emission may suggest a strong correlation with CH$_3$COOH formation  \\citep[e.g.,][]{remi04}. Another possibility is that CH$_{3}$COOH  and large N-bearing species share approximately the same time scale to form in hot cores. Nevertheless, it is important to investigate the spatial relationship between CH$_{3}$COOH and the distribution of other complex molecular species. Thus, we also observed methyl formate (HCOOCH$_3$) and ethyl cyanide (CH$_3$CH$_2$CN) as proxies to determine the spatial distribution of O and N-bearing species toward our sources in this survey.  To search for regions of CH$_{3}$COOH emission and continue to determine its spatial relationship with  the distribution of other complex molecular species, we have used the high  sensitivity of the Combined Array for Research in Millimeter-wave Astronomy (CARMA)  to carry out a CH$_{3}$COOH survey toward three more hot cores. The observed sources are  G19.61-0.23, G29.96-0.02 and IRAS 16293-2422 and are described in more detail in $\\S$2. A complete description of the CARMA observations are described in $\\S$3; the results are presented in $\\S$4; and in $\\S$5 we describe the analysis routine and discuss the implication of these observations with future surveys and chemical formation models. ",
        "conclusions": "\\subsection{Molecular Column Densities}  Molecular column density is one essential physical parameter needed to constrain chemical models (see e.g. Garrod et al. 2008 and references therein). In order to  determine the column densities of each molecular  species of interest, we assume that each region has uniform physical conditions, that the populations of the energy levels can be characterized by a Boltzmann distribution, and finally, that the emission is optically thin.  Assuming that the  molecular species is in local thermodynamic equilibrium (LTE) and low  optical depth, the total beam averaged column density is \\begin{equation} N_T (\\mathrm{cm^{-2}})=2.04\\times10^{20} \\it \\frac{\\int I(\\mathrm{Jy~beam^{-1}}) d\\upsilon(\\mathrm{km~s^{-1}})Q_re^{E_u(K)/T_r}}{\\Omega_b(\\mathrm{arcsec^2}) \\nu^{3}(\\mathrm{GHz^3})<S\\mu^{2}>(\\mathrm{debye^2})}, \\label{eqn_Ntot} \\end{equation} where $\\Omega_b$ is the solid angle of the beam , $\\int I_{\\upsilon}d\\upsilon$ is the integral of the line intensity over velocity, $\\nu$ is the spectral line frequency, $<S\\mu^{2}>$ is the line strength parameter, $Q_r$ is the rotational partition function, $T_r$ is the rotational temperature, and $E_u$ is the upper level energy of the transition \\citep{miao95}. In each case, we assume the source emission fills the synthesized beam of the observations.  All of the line parameters used for the analysis are given in Table 2.  A least-square fitting routine with Gaussian functions was used to determine the linewidths ($\\Delta v$), peak intensities ($\\Delta$I) and $\\int I d\\upsilon$ of each observed spectral line toward each source.  All least-square fitting was done using the  standard packages contained in the $Mathematica$ software package.  Using the spectroscopic  parameters from Table 2 and by varying the temperature and column density over a wide  range of values (e.g. the temperature range investigated was from 10-500 K and the column  density range was from 10$^{12}$ - 10$^{18}$ cm$^{-2}$ for each molecular species), Equation (1) gave a predicted integrated line intensity for each spectral line observed in our data. These predictions of integrated line intensity were then compared to the measured values and, using a least-squares fitting routine, a best fit temperature and column density was found for each molecule toward each source. However, there were some sources where an independent fit of the temperature did not converge.  These sources are described in more detail in the  following sections. Columns 3 and 4 of Table 3, 4 and 5 give the fitted line parameters of the  observed peak intensity and line width. Columns 5 and 6 give the modeled peak intensity  based on the least square fitting routine and the fitted line width from the observed data,  respectively.  Once all the relevant line parameters have been determined, a model spectrum  with the appropriate amount of random noise added and convolved with the spectral resolution of the observation can then be produced and be directly compared with the observed dataset (see $\\S$5.5).  The rotational temperature $T_r$ generally is representative of the physical kinetic temperature in regions where the spatial density is larger than the critical density of the molecule in question. The assumption in this work is that  the densities of the HMCs investigated are expected to be high enough for this to be true for these molecules.  The rotational temperature is often found by observing several different transitions of a molecular species over a range of $E_u$ values.  However, given the limited bandwidth and low number of transitions observed in our data, it was often difficult to accurately constrain the temperature and column density of our three observed molecules toward these regions \\citep[e.g.,][]{snyder05}. Therefore, when the least squares fit failed to give a reasonable temperature for a given molecular species, we adopted a rotational temperature from previous observations by other investigators toward these sources at similar spatial resolution and were then able to fit for the total beam averaged column densities or associated upper limits if no transitions were detected beyond the statistically significant 1$\\sigma$ limit.  It was also the case that if the fitting procedure was given too many free parameters, a constraint was needed on one or more or else it would not converge (see the G29 discussion below).  To properly constrain the temperature of these regions over  a variety of molecular species, observations of additional transitions are  necessary and complete spectral line surveys may be required \\citep[see e.g.][and references therein]{font07,frie04}.  For the G19 region, the least square fitting routine determined a rotational temperature of CH$_3$CH$_2$CN of T$_{rot}$=161(58) K, and a beam averaged column density of  $N_T$=6(3)$\\times$10$^{16}$ cm$^{-2}$.  The errors on all fits are 1$\\sigma$. The routine could not find a reasonable fit for the temperature and column density  for either HCOOCH$_3$ or CH$_3$COOH presumably because the detected transitions  were too close in energy.  Therefore, because the emission regions of each molecular species  around G19 are nearly co-spatial (Figure \\ref{g19map}), we adopted the same rotational temperature for HCOOCH$_3$ and CH$_3$COOH as found from the CH$_3$CH$_2$CN fit.  Assuming a  T$_{rot}$=161 K for HCOOCH$_3$, we find a best fit to the column density of $N_T$=9(2)$\\times$10$^{16}$ cm$^{-2}$ and for CH$_3$COOH, we find a best fit to the  column density of $N_T$=2(1)$\\times$10$^{16}$ cm$^{-2}$.  Given the minimum and maximum values of the measured column density, we determine an abundance ratio range of N$_{HCOOCH_3}$/N$_{CH_3COOH}$=3-11. Table \\ref{tbl-g19} gives the fitted intensities of each of the detected spectral features compared to the predicted intensities determined from Equation (1) for G19.  For the G29 region, the least square fitting routine determined a rotational temperature of CH$_3$CH$_2$CN of T$_{rot}$=107(49) K, and a beam averaged column density of  $N_T$=1.1(7)$\\times$10$^{16}$ cm$^{-2}$ only if we adopted a rotational temperature of 150 K, determined from \\citet{olmi03} and supported by the observations of \\cite{font07} and \\cite{beut07}, for the transitions of HCOOCH$_3$.  This temperature was also set to determine the upper limit for the column density of CH$_3$COOH. Assuming a  T$_{rot}$=150 K for HCOOCH$_3$, we find a best fit to the column density of $N_T$=4(1)$\\times$10$^{16}$ cm$^{-2}$ and for CH$_3$COOH, we find an upper limit to the column density of $N_T$$<$9$\\times$10$^{14}$ cm$^{-2}$ for a relative abundance ratio upper limit of N$_{HCOOCH_3}$/N$_{CH_3COOH}$$>$50.  Table \\ref{tbl-g29} gives the fitted intensities of each of the detected spectral features compared to the predicted intensities determined from equation (1) for G29.  We also note the detection of another feature shown in Figure \\ref{g19spec} near 107.540 GHz that is most likely due to two high energy ($>$110 K) spectral lines of CH$_3$CH$_2$CN.  The discussion of weak spectral features  near the 1$\\sigma$ noise level will be given in section $\\S$5.5.  A recent observation of the G19 and G29 regions with the IRAM 30-m telescope \\citep{font07} covered our two acetic acid line frequencies (Figure \\ref{spec_30m}). However, with a 28\\arcsec~beam, the lines detected are under the noise level of 26 mK. These lines are brighter than those observed  by \\cite{cazaux03}, but the beam dilution is severe. The high resolution capability of CARMA is indeed critical for searching for acetic acid and large molecule research.  Finally, toward IRAS 16293-2422, the least square fitting routine could not determine a rotational temperature for the detected transitions of HCOOCH$_3$ presumably for the same reason as described toward G19. Thus, we adopted a rotational temperature of 62 K determined from \\citet{bott04}. In this case, we find a best fit to  the column density of  HCOOCH$_{3}$ of $N_T$=1.5(3)$\\times$10$^{16}$ cm$^{-2}$  toward region A and 6(1)$\\times$10$^{15}$ cm$^{-2}$ toward region B. Since there was no clear detection of CH$_3$CH$_2$CN  beyond the 1$\\sigma$ detection limit toward either region A or B, based on the noise level in each of the  passbands containing those lines, we determined an upper limit to the total beam averaged column density  of CH$_3$CH$_2$CN to be $<$6$\\times$10$^{13}$ cm$^{-2}$ for region A; and $<$3$\\times$10$^{14}$ cm$^{-2}$ toward region B.  For CH$_3$COOH, the 2$\\sigma$ detection toward region A gives a total beam averaged column density of $\\sim$1.6$\\times$10$^{15}$ cm$^{-2}$ and for region B, we determine and upper  limit of $<$6$\\times$10$^{14}$ cm$^{-2}$.  Table \\ref{tbl-i16293} summarizes the  fitting results of the HCOOCH$_{3}$ transitions compared to the model intensities  determined from Equation (1).  \\subsection{A Proxy for the Detection of CH$_{3}$COOH Sources}  We have successfully detected CH$_{3}$COOH toward G19. Following Sgr B2(N-LMH), W51e2 and G34.3+0.15, G19.61-0.23 is the fourth high mass hot core that is a source of CH$_3$COOH. We compare the CH$_3$COOH column densities and the abundance ratios to HCOOCH$_3$ of all detected CH$_3$COOH hot cores in Table \\ref{tbl-aca_cmp}. The values are generally consistent. Therefore, the column densities and ratios can be used to constrain hot core chemical models assuming the co-spatial dependence of HCOOCH$_3$ and CH$_{3}$COOH remains.   This has yet to be tested at extremely high spatial resolution. In addition, a consistency can be found between all the sources containing CH$_3$COOH emission by comparing the relative line strengths between the detected transitions of HCOOCH$_3$ and CH$_3$CH$_2$CN to the CH$_3$COOH emission.  Taking a ratio between measured (or predicted intensities) of the 107.594 GHz line of CH$_3$CH$_2$CN to the 111.507 GHz line of CH$_3$COOH, we find that this ratio is between 5.5 and 6 for those sources where CH$_3$COOH is detected. Therefore, given the line intensity of 0.21 Jy/beam for the 107.594 GHz line of CH$_3$CH$_2$CN toward G29.96-0.02, the expected line intensity for the 111.507 GHz line of CH$_3$COOH is $\\leq$0.04 Jy/beam.  This intensity is right at the 1$\\sigma$ detection limit of the current CARMA observations.  Therefore, in addition  to the criteria for the detection of CH$_3$COOH outlined in \\citet{remi03}, we now have a good proxy for the detectability of the 111 GHz CH$_3$COOH transitions relative to the detected line strengths of the HCOOCH$_3$ and CH$_3$CH$_2$CN transitions near 107 GHz. Further observations of new sources of  HCOOCH$_3$ and CH$_3$CH$_2$CN that fit the criteria set by \\citet{remi03} and higher S/N observations of sources like G29, are necessary to validate this hypothesis.  \\subsection{CH$_{3}$COOH in IRAS 16293-2422?}  High spatial resolution maps and spectra of IRAS 16293-2422 clearly reveal that the morphologies of the emission of complex molecules are clearly not co-spatial and that there is no simple separation of O and N-bearing ``chemical differentiation''  \\citep{guelin08}.  \\citet{blak87} and \\citet{rodg01} have suggested the difference in the morphologies of the emission of complex molecules in these regions reflects the timescales of chemical evolution. Predominantly N-bearing species (e.g.\\ CH$_3$CH$_2$CN) represent the later stage of chemical evolution while  predominantly O-bearing species (e.g.\\ HCOOCH$_3$) represent the earlier stage.  Their models can estimate the ages of hot cores by comparing this difference between the O and N  morphology of emission.  However, to date, no direct observation can accurately  measure hot core ages, so the uncertainty of these estimates is large.  IRAS 16293-2422 also shows clear distinctions in the morphology of molecular emission \\citep{chan05,remi06}.  In fact, the relative distributions of molecular species will undoubtedly give insight into their formation chemistry. It may also turn out that the relative abundances of molecular species and the overall population of  energy levels of the detected transitions of a given molecule may be determined  more by the formation chemistry and not on the local thermodynamic environment.  To test this hypothesis will require high spectral and spatial resolution observations  of a group of species that may form via a common formation pathway.  IRAS 16293-2422 has been reported as the first low-mass CH$_3$COOH source by \\citet{cazaux03} using the IRAM 30-m telescope, but the one detected line was likely blended. While IRAS 16293-2422 is not resolved in their work (beamsize of 28\\arcsec),  we have resolved it into I16293A and B with a separation of less than 5\\arcsec. Even though their observation used different  transitions, we should detect our CH$_3$COOH lines in our synthesized beam of 3\\farcs1$\\times$1\\farcs9 using their flux estimates.  Although at a lower flux than expected from their results, we do detect two CH$_3$COOH transitions toward region A but only at the 2$\\sigma$ limit.  We estimate a column density of $\\sim$ 1.6 $\\times$ 10$^{15}$ cm$^{-2}$, and \\citet{cazaux03} estimated an upper limit of 2.5 $\\times$ 10$^{15}$ cm$^{-2}$, which is similar.  Furthermore, there is no evidence beyond the 1$\\sigma$ levels of any emission of CH$_3$COOH toward region B. In addition, while we have observed  HCOOCH$_3$ in I16293A and B, we did not detect the corresponding CH$_3$CH$_2$CN emission features. The detection of CH$_3$COOH toward G19 (and possibly toward I16293A) and the other hot cores have demonstrated that a diversity of abundant large O and N-bearing species are strongly correlated with CH$_3$COOH. One possibility is that the mix of O and N-bearing species may serve as an efficient chemical network for CH$_3$COOH and other large molecules. Thus, they need to wait for large N-bearing molecules to form first before they can be efficiently produced. Another possibility is that the formation of large O and N-bearing species is simply time dependent, and the gas-phase CH$_3$COOH concentration  just by coincidence starts to increase along with large N-bearing species.  Given the criteria set above, it is intriguing that CH$_3$COOH is detected toward I16293A given there was no clear evidence of any emission from CH$_3$CH$_2$CN.  Further higher sensitivity observations are once again necessary to confirm this detection and determine the evolutionary stage of the I16293A and B cores.  \\subsection{CH$_{3}$COOH Formation}  Both gas phase and grain surface chemistry undoubtedly play important roles in hot cores \\citep[e.g.][]{hase92,garr06,garr08}. Large molecule abundances can be enhanced enormously by grain surface chemistry while current models of gas phase chemistry alone are not sufficient to match observations.  \\citet{garr06} have demonstrated that radicals in the icy mantles on grain surfaces can move around as the temperature gradually increases from 10 to 200 K in star forming regions. Grain surfaces become efficient environments for large molecules to form. However, the details of CH$_3$COOH formation in hot cores still remain unclear. The correlation between CH$_3$COOH and large N-bearing molecules may suggest that some N-bearing molecules act as catalysts to the CH$_3$COOH formation. The chemical and physical evolution timescales of star formation lacks direct observational evidence to confirm chemical models. Without the support of observation, it is difficult to establish solid chemical models. \\citet{garr08} have reported their new chemical model for CH$_3$COOH in hot cores. Grain surface reactions have been enhanced to match observations. They show the time scales of hot cores warming up affect the abundances of the secondary radical CH$_3$CO, which can react with OH to form CH$_{3}$COOH. This effect can also cause the dissimilar abundances of structural isomers, which is consistent with our survey.  \\subsection{Identifying Weak Spectral Features Using Model Spectra.}  When CH$_3$COOH was first detected \\citep{meh97}, it emphasized the importance of interferometry for the study of astrochemistry in hot cores. Star forming regions generally contain many regions of compact molecular emission.  Although G19 has a complex structure, including extended and compact components \\citep[see][]{furu05}, the CH$_3$COOH detection toward G19 was conducted with a high resolution beam, 1\\farcs9$\\times$1\\farcs6, which is the smallest used in a CH$_{3}$COOH survey to date. The three observed molecules in G19, CH$_3$COOH, HCOOCH$_3$ and CH$_3$CH$_2$CN, are mainly located at the same component, which agrees with the CH$_3$CH$_2$CN observation by \\citet{furu05}. As we continue extending CH$_{3}$COOH surveys to smaller hot cores, the required resolution of observations will be higher.  Higher resolution observations also have the advantage of reducing the effect of line confusion present in a spectrum.  Also, if the source size of the emitting region is well coupled to the synthesized beam of the array, weak spectral features that are lost due to beam dilution effects from single dish observations can be significantly resolved out of the noise. Thus, many new molecule detections over the last several years have used both a combination of single dish and array observations to confirm the detection \\citep[e.g.,][]{bell08,bell09}. In addition, these detections routinely use model spectra based on the laboratory and calculated rotational spectroscopy of energy levels and line strengths for molecular transitions, the physical conditions of the emitting regions including temperature and density and finally, the characteristics of the telescope including forward beam efficiency and field of view  (i.e. beam size).  Given the model spectra, the weak transitions of molecular species appear statistically significant in the observed spectrum. However, the model data have near infinite spectral resolution compared to the observed data and do not include an estimate of the noise level in the passband.  These types of models are shown in the spectral line data presented in this work (red Gaussian traces in the first row of spectra in Figures 2, 4, 6 and 7).  In order to correctly ascertain whether a weak feature in the observed spectrum is due to  a calculated line of a molecular species, we advocate taking the model data and convolving it with the spectral resolution of the instrument and adding the appropriate amount of random noise to the resultant spectrum to give a more accurate representation of  the data collected by the telescope.   The noise level is based on the measured noise of line free channels in the observed data.  These modeled data are then subtracted from the observed data to form a residual  spectrum.  From these residuals, with the caveat of an accurate model of the physical conditions of the source, one can identify the remaining weak spectral features and determine if features close to the 1-2$\\sigma$ level are in fact statistically significant spectral line features.  Based on the residual passbands shown in the last row of Figures 2, 4, 6 and 7, we include in Tables 3, 4 and 5 the possibility of unidentified lines (i.e. U-lines) in the data toward G19, G29 and IRAS 16293-2422. In addition, given the spectral line fit (red Gaussian traces) toward IRAS 16293-2422, there is a suggestion that the weak emission features shown toward I16293B may in fact be due to CH$_3$COOH.  However, when this model spectrum is convolved with the noise level in the passband, these weak features completely disappear and there are definitely no statistically significant feature seen in the residual passbands. As a result, while we claim that the  emission features that are coincident with the CH$_3$COOH transitions toward I16293A are in fact from CH$_3$COOH, these features are only  barely above the noise level so further high sensitivity observations are  needed for a definitive detection.  It is our recommendation that if model spectra are going to be the new proxy for identifying weak spectral features in observational data, the model data should be first convolved to the spectral resolution of the telescope and the appropriate amount of noise be added to the model in order to make an accurate representation as possible to the observed data.  Only then can the 1-2$\\sigma$ spectral features be corrected assessed as statistically significant  features in the observed data and not an artifact of the noise and spectral resolution."
    },
    "1004/1004.4493_arXiv.txt": {
        "abstract": "We report sub-arcsecond scale radio continuum observations of a field of 35 by 22\\,arcmin centred in NGC\\,3351 obtained with the Multi-Element Radio Linked Interferometer Network (MERLIN). We found 23 radio sources in this field, 6 of which are projected within the D$_{25}$ isophote of the galaxy, and 3 are located inside the central 100\\,arcsec in radius. Two of these three are significantly extended, while the third one is relatively compact. This one is the only source with a previously detected counterpart at other wavelengths and could constitute the radio counterpart of a young supernova remnant. The other two are probably related to jets from a background AGN. We are not able to detect individual supernovae or SNRs in the central region ($r<600$ pc) of the galaxy. This could imply that the ionising populations of the circumnuclear star-forming regions (CNSFRs) are too young (less than a few Myr) to host supernovae. Also the diffusion length of the relativistic electrons in the ISM associated with the SN from previous events of star formation seems to be larger than our maximum resolution of 50 pc in radius. Detecting the thermal bremsstrahlung emission from the circumnuclear \\HII regions probably requires deeper observations. ",
        "introduction": "\\label{secIntro}   The radio emission of a normal spiral galaxy consists of a thermal bremsstrahlung component, arising from inelastic Coulomb collisions between the free electrons and ions of the interstellar medium (ISM), and a synchrotron component emitted by relativistic cosmic-ray electrons from supernova remnants (SNRs) and active galactic nuclei (AGN) spiralling in the galactic magnetic field \\citep[see e.g.][]{Condon92}. Star-forming regions may host a large number of supernovae, and the \\HII regions created by young massive stars can be detected as thermal sources \\citep[e.g.][]{KobulnickyJohnson99,Turner+00,Johnson+01,Johnson+03,Johnson+09,CannonSkillman04,Tsai+06,Tsai+09,Rosa-Gonzalez+07a}. In particular, the inner ($\\sim1$~kpc) parts of some spiral galaxies show higher star formation rates (SFRs) than usual, frequently arranged in a ring pattern around the nucleus. Circumnuclear star-forming regions (CNSFRs), also referred to as ``hot-spots'' \\citep{Sersic+65,Sersic+67}, are similar to luminous and large disk \\HII regions, but they are more compact and show higher peak surface brightness \\citep{Kennicutt+89}. Their massive stars can dominate the observed ultraviolet (UV) emission even in the presence of AGN \\citep{Gonzalez-Delgado+98,Colina+02}, and the ionised gas can contribute up to 50 per cent to the total H$\\beta$ luminosity in some Seyfert~2 nuclei \\citep{Cid-Fernandes+01}. The \\Ha luminosities, typically higher than $10^{39}$~erg~s$^{-1}$, overlap with those of \\HII galaxies \\citep[see e.g.][and references therein]{Melnick+88,Diaz+00,HoyosDiaz06}. Due to their location near the centre of the galaxy, CNSFRs display high metal abundance \\citep[see e.g.\\ ][and references therein]{Diaz+07}, which makes them ideal laboratories to investigate star formation in environments of large metallicity.  NGC\\,3351, also known as M95 and UGC5850, is a nearby \\citep[10.05\\,Mpc, 49 pc/arcsec,][]{Graham+07} SBb(r)II early-type barred spiral galaxy \\citep{Sandage+87} in the Leo group. Its coordinates are $\\alpha_{\\rm 2000}=10^{\\rm h}\\,43^{\\rm m}\\,57\\fs7$, $\\delta_{\\rm 2000}=+11^{\\circ}\\,42\\arcmin\\,12\\farcs7$, and its major- and minor-axis diameters at the 25 magnitude isophote (D$_{25}$) are $7.4 \\times 5.0$~arcmin, with a position angle (PA) of 13$^{\\circ}$ \\citep[RC3 catalogue,][]{deVaucouleurs+91}. NGC\\,3351 has a star-forming circumnuclear ring, and has been classified as a ``hot-spot'' galaxy by \\cite{Sersic+67}. Early detailed studies of its nuclear regions by \\cite{Alloin+82} showed that NGC\\,3351 harbours high-mass star formation along a ring of about 20\\,arcsec in diameter in its central zone. In fact, it can be considered a nuclear starburst galaxy, since the SFR per unit area is significantly higher in the nuclear region than in the disc \\citep{Devereux+92}. HST UV images show that the present star formation in the ring of NGC\\,3351 is arranged in complexes of diameters between 1.4 and 2.0~arcsec, made up of several high-surface brightness knots a few parsec in size embedded in a more diffuse component \\citep{Colina+97}. A circumnuclear SFR of 0.38\\,M$_\\odot$\\,yr$^{-1}$ was inferred by \\cite{Elmegreen+97} from near infrared (IR) photometry in  the J and K bands, and a value of 0.24\\,M$_\\odot$\\,yr$^{-1}$ was derived by \\cite{Planesas+97} from the H$\\alpha$ emission. From CO observations, these authors estimated a mass of molecular gas of 3.5\\,$\\times$\\,10$^{8}$\\,M$_\\odot$ inside a circle of 1.4~kpc in diameter. According to \\cite{Jogee+05}, NGC\\,3351 hosts 5.3\\,$\\times$\\,10$^8$\\Msun of molecular hydrogen, with most of the emission coming from the inner $\\sim600$\\,pc. Within this radius, a SFR of 0.5\\,M$_\\odot$\\,yr$^{-1}$ can be estimated from the Br$\\gamma$ luminosity given by \\cite{Puxley+90}, whereas the non-thermal component of the radio continuum emission favours a lower value of 0.3\\,M$_\\odot$\\,yr$^{-1}$ \\citep{Jogee+05}. \\cite{vandeVen+09} estimated that the molecular mass inside the circumnuclear ring is about 5\\,$\\times$\\,10$^7$\\Msun, and \\cite{Leroy+09} found a well-defined molecular ring with a diameter of about 50\\,arcsec, located at the end of the stellar bar \\citep[with a semimajor axis of 47\\,arcsec; see e.g.][]{Martin95}. These CO structures are compatible with the excesses observed by \\cite{Regan+01} and \\cite{Helfer+03}. A small CO bar perpendicular to the large-scale stellar bar was also observed by \\cite{Devereux+92} and \\cite{Helfer+03}.  NGC\\,3351 is part of the SINGS Legacy Science Program sample \\citep{Kennicutt+03}, and it has been included in a wide variety of statistical studies. The SINGS data set provides broadband imaging, as well as low- and high-resolution spectral maps in the IR. The integrated spectral energy distribution (SED) of NGC\\,3351 from 0.15 to 850~$\\mu$m has been derived by \\cite{Dale+07} by combining the {\\it Spitzer} data with those from 2MASS and {\\it IRAS} in the IR, GALEX in the UV, and other data in the optical, modelling the dust and stellar contributions. They concluded that most of the sample's spectral variations stem from two underlying components, one representative of a galaxy with a low IR-to-UV ratio and one representative of a galaxy with a high IR-to-UV ratio, as given by a principal component analysis of the sample. \\cite{Draine+07} used the IRAC and MIPS observations to estimate the total dust mass (10$^{7.46}$\\,\\Msun) and the fraction of the dust mass contributed by PAHs (3.2 per cent). They found that the dust properties of spiral galaxies resemble those in the local environment of the Milky Way, with similar dust-to-gas ratio and similar PAH abundance. \\cite{Munoz+09} found that NGC\\,3351 shows significantly larger IR concentration indices than the remaining galaxies of the same morphological type of the SINGS sample.   We have studied the circumnuclear region of NGC\\,3351 in two previous works. In \\cite{Diaz+07}, the physical conditions of the gas were determined for 7 CNSFRs. We developed a semi-empirical method for the derivation of chemical abundances in the high-metallicity regime, obtaining values consistent with solar within the errors and comparable to those found in high-metallicity disc \\HII regions using direct measurements of electron temperatures. An analysis of the kinematics of the gas and stars in 5 CNSFRs and the nucleus of NGC\\,3351 is presented in \\cite{Hagele+07}. We derived the dynamical masses (between $4.9\\times10^6$ and $4.3\\times10^7$\\,M$_\\odot$) for the stellar clusters in these regions, and found that the CNSFR complexes, with sizes of about 100 to 150\\,pc in diameter, are seen to be composed of several individual star clusters (with sizes between 1.7 and 4.9\\,pc on a HST image). The radial velocity curve shows deviation from circular motions for the ionised hydrogen, consistent with infall towards the central regions of the galaxy at a velocity of about 25\\,km\\,s$^{-1}$ \\citep[in agreement with][]{Rubin+75,Dicaire+08}. These last authors also point out that the H$\\alpha$ velocity field is fairly regular outside the  central bar, and \\cite{Jogee+05} found that the velocity field of the molecular material in the inner 500\\,pc is generally dominated by circular motions, with some evidence of weaker non-circular streaming.  In this paper, we will look for radio-continuum signatures of massive star formation in NGC\\,3351, both in the nucleus (in particular, the CNSFRs) and the galaxy disk. Section~\\ref{secObs} describes the observations and the data reduction. The spatial distribution of radio sources and their observed properties are discussed in Section~\\ref{secResults}, and Section~\\ref{secDiscus} is devoted to the interpretation of these results in terms of star formation activity. Our main conclusions are briefly summarised in Section~\\ref{secConclus}.    \\begin{table*} \\caption{Sources around NGC\\,3351 detected by MERLIN. The first column gives the adopted reference name for each source (in brackets the adopted names for the three closest sources to the centre of NGC\\,3351), the second column gives the FIRST counterpart name (within 1\\,arcsec of the MERLIN source in all cases) and the third column gives the radial offset of the MERLIN sources from the pointing centre. The remaining columns give the MERLIN source properties: positions, peak flux densities, integrated flux densities, and their corresponding errors. These measurements for R1, R2 and R3 were obtained by fitting a single component; see Table \\ref{ext-sources}\\ for further details of these sources. } \\label{tab:brights} \\begin{tabular}{lcrccrrrrr} \\hline Source & FIRST & Radius& R.A. & Dec. & $\\sigma_{\\rm {RA}}$ & $\\sigma_{\\rm {Dec}}$ & Peak & Int. &$\\sigma_{\\rm {rms}}$\\\\ & name & (arcmin) & \\multicolumn{2}{c}{(J2000)}& (mas)& (mas)&(mJy b$^{-1}$)&(mJy)&(mJy b$^{-1}$)\\\\ \\hline 1    &  J104252.6+114853 &17.283 & 10 42 52.5564 & +11 48 54.157 & 47 & 76 &  0.31 &   1.78 & 0.08 \\\\% C-FIRST    \\\\ 2    &  J104253.2+114848 &17.076 & 10 42 53.3084 & +11 48 48.482 & 27 & 44 &  0.55 &   1.91 & 0.08 \\\\% B-FIRST    \\\\ 3    &  J104254.6+114841 &16.727 & 10 42 54.6482 & +11 48 41.204 & 24 & 38 &  0.63 &   5.33 & 0.08 \\\\% A-FIRST    \\\\ 4    &  J104301.1+115230 &17.228 & 10 43 01.1845 & +11 52 30.364 & 27 & 44 &  0.54 &   1.28 & 0.08 \\\\% D-FIRST    \\\\ 5    &                   &15.086 & 10 43 12.8952 & +11 52 35.685 & 48 & 78 &  0.31 &   0.73 & 0.08 \\\\% new3\t     \\\\ 6    &  J104320.5+114301 &9.1344 & 10 43 20.5287 & +11 43 01.933 & 10 & 16 &  1.48 &   5.00 & 0.08 \\\\% 2 7    &  J104328.4+113723 &8.6480 & 10 43 28.4280 & +11 37 23.613 &  5 & 10 &  2.43 &   5.53 & 0.08 \\\\% 3 8    &                   &5.2228 & 10 43 36.7801 & +11 43 15.556 & 21 & 34 &  0.72 &   1.65 & 0.08 \\\\% new1\t     \\\\ 9    &  J104342.7+115209 &10.575 & 10 43 42.7584 & +11 52 09.373 &  2 &  4 &  4.33 &  11.82 & 0.08 \\\\% 4 10    &                   &2.8629 & 10 43 47.3236 & +11 40 54.775 & 35 & 57 &  0.42 &   0.81 & 0.08 \\\\% 3J\t     \\\\ 11    &  J104349.1+113628 &6.1266 & 10 43 49.1642 & +11 36 28.452 &  3 &  5 &  2.74 &   4.00 & 0.09 \\\\% 5 12    &                   &1.9497 & 10 43 49.8033 & +11 41 58.791 & 43 & 70 &  0.43 &   0.85 & 0.08 \\\\% 9J\t     \\\\ 13    &                   &2.8251 & 10 43 50.6611 & +11 39 59.677 & 35 & 56 &  0.42 &   0.78 & 0.08 \\\\% 2J\t     \\\\ 14(R1)&  J104355.2+114129 &0.9541 & 10 43 55.2618 & +11 41 29.337 &  2 &  4 &  4.11 &   4.08 & 0.07 \\\\% 6(R1) 15(R2)&  J104359.5+114332 &1.3864 & 10 43 59.5464 & +11 43 32.638 &  3 &  5 &  3.53 &  22.04 & 0.07 \\\\% 7(R1) 16(R3)&  J104401.4+114323 &1.4793 & 10 44 01.4482 & +11 43 23.623 & 21 & 18 &  1.38 &  15.93 & 0.06 \\\\% 8(R3) 17    &                   &2.9359 & 10 44 08.4135 & +11 40 54.837 & 38 & 58 &  0.39 &   1.09 & 0.08 \\\\% 1J\t     \\\\ 18    &                   &3.0888 & 10 44 10.3120 & +11 42 08.643 & 35 & 57 &  0.42 &   0.50 & 0.08 \\\\% 5J\t     \\\\ 19    &  J104417.4+114614 &6.2824 & 10 44 17.4533 & +11 46 14.678 &  1 &  1 & 62.77 & 125.40 & 0.21 \\\\% 9 20    &                   &6.6425 & 10 44 17.5254 & +11 46 46.157 & 15 & 24 &  0.96 &   1.79 & 0.09 \\\\% new2\t     \\\\ 21    &  J104423.4+114458 &6.8677 & 10 44 23.4286 & +11 44 58.340 &  3 &  3 & 14.05 &  32.41 & 0.32 \\\\% 10 22    &  J104429.2+113811 &8.7224 & 10 44 29.2655 & +11 38 11.349 &  7 &  6 &  4.05 &  13.05 & 0.14 \\\\% 11 23    &  J104430.6+113811 &9.0252 & 10 44 30.6493 & +11 38 11.184 &  1 &  1 & 21.12 &  88.18 & 0.13 \\\\% 12 \\hline \\end{tabular} \\end{table*}        \\begin{figure*} \\centering \\includegraphics[width=1\\textwidth]{figures/all-sources-Halpha-2.png} \\caption[]{Distribution of the MERLIN sources in a 35\\,$\\times$\\,22\\,arcmin$^2$ field centred at $\\alpha_{\\rm 2000}$\\,=\\,10$^{\\rm h}$\\,43$^{\\rm m}$\\,57\\fs7, $\\delta_{\\rm 2000}$\\,=\\,+11\\degr 42\\arcmin 14\\farcs0 superimposed in an H$\\alpha$ image from Palomar 48-inch Schmidt telescope. Crosses and small circles indicate the MERLIN and FIRST sources, respectively. The MERLIN sources are labelled according to Table~\\ref{tab:brights}. The big circle indicates the central 100\\,arcsec (in radius) of NGC\\,3351, and the ellipse is the D$_{25}$ isophote of this galaxy as defined by \\citet{deVaucouleurs+91}. The orientation is north up, east to the left. [{\\it See the electronic edition of the Journal for a colour version of this figure.}]} \\label{all-source-dist} \\end{figure*} ",
        "conclusions": "\\label{secDiscus}  Early-type spiral galaxies such as NGC\\,3351 provide an ideal laboratory to study the very different environments in which star formation may take place, from the peculiar conditions of the CNSFRs \\citep[e.g.][and references therein]{Hagele+09,Hagele+10} to the outskirts of the galaxy disk. Maps of radio-continuum emission can reveal the presence of SNRs and \\HII regions through their synchrotron and thermal emission, respectively. These are closely related to the evolution of massive stars and are, therefore, excellent tracers of the recent star formation activity within the galaxy \\citep[see e.g.][]{Muxlow+94}. Moreover, young massive star-forming regions have very characteristic signatures in the radio band, providing strong constraints on the age of the most recent starbursts \\citep[see e.g.][]{Rosa-Gonzalez+07a}.  We have divided the galaxy in several regions, corresponding to different star formation regimes: the centre, which we define as the central 600~pc (12~arcsec radius) features the highest densities and metallicities, $Z \\sim Z_\\odot$. The CNSFR complexes identified in \\citet{Planesas+97} have been considered separately from the diffuse nuclear star formation. Finally, the galaxy disk ($r>600$ pc, $r>12$~arcsec) represents a more typical, low-density environment with low to moderate metallicities. The main properties (diameter, \\Ha luminosity, electron density and temperature, filling factor, thermal optical depth and flux density, star formation rate, total number of SNRs, expected number above our detection threshold, and total synchrotron flux density) of each region are summarised in Table~\\ref{tab:regions}. The eight CNSFRs are listed in the same order as in \\citet{Planesas+97}.  Diameters ($D$) of the CNSFR complexes are estimated from the HST F606W image, and \\Ha luminosities [$L(\\Ha)$] are taken from \\citet{Planesas+97}, corrected for internal extinction using the colour excess [E(B-V)] estimated by \\citet{tesisdiego} and assuming the Galactic extinction law of \\citet{Miller+72} with $R_v$\\,=\\,3.2 \\citep{HagelePhD}. For the rest of the nuclear region (excluding the CNSFRs), the star formation rate derived from the Br$\\gamma$ luminosity \\citep{Puxley+90} has been converted to an \\Ha luminosity by means of the empirical relation given by \\citet{Rosa-Gonzalez+02}. For the galactic disk, the SFR measured by \\citet{Leroy+08} has been used.  Electron densities ($n_{\\rm e}$) of the emitting gas in the CNSFRs have been computed from the [S{\\sc ii}]$\\lambda\\lambda$6717/6731\\,\\AA\\ line ratio, and electron temperatures refer to $T_{\\rm e}$([O{\\sc iii}]) \\citep{Diaz+07}. For the eighth region, values of 400~cm$^{-3}$ and 5000~K have been assumed. The ambient ISM densities of the central diffuse region and the disk are 10 and 1~cm$^{-3}$, respectively, with a temperature of $10^4$~K. The associated filling factors ($f$) are given by expression~(\\ref{eqFillFactor}). As explained in Appendix~\\ref{secThermal}, the predicted optical depth ($\\tau_{\\rm 1.42\\,GHz}^{\\rm brems}$) and flux density ($S_{\\rm 1.42\\,GHz}^{\\rm brems}$) at 1.42~GHz arising from thermal bremsstrahlung take into account self-absorption, although the effect is found to be mild at this frequency.  Star formation rates have been estimated from the \\Ha luminosity \\citep{Rosa-Gonzalez+02}, and the expected number of SNRs ($N_{\\rm SNR}$) has been computed according to equation~(\\ref{eqNumberSNR}), assuming an ambient ISM density of 10~cm$^{-3}$ for the CNSFRs. The number of SNRs brighter than 0.3~mJy ($N_{\\rm SNR}^{( 0.3~{\\rm mJy} )}$) is given by expression~(\\ref{eqN_SNR_thr}), and the total synchrotron emission from the whole population of unresolved supernovae ($S_{\\rm 1.42\\,GHz}^{\\rm synchr}$) has been estimated from the empirical relation between the star formation rate and the 1.4~GHz luminosity density reported by \\cite{Yun+01}.   As shown in the previous section, our analysis of the sub-arcsecond spatial resolution radio map of NGC\\,3351 shows the presence of 23 sources in the L-band. Three of them (R1, R2 and R3) are located within the central 100~arcsec in radius. They have clear counterparts in the FIRST survey \\citep{White+97} and the VLA data analysed by \\citet{Paladino06}, but only R1 seems to be visible at other wavelengths. It is unresolved by MERLIN, putting an upper limit of $\\sim10$ pc on its size, but its variability (Section~\\ref{secResults}) suggests that it is $<3$ pc.  As shown in Figure~\\ref{R2+R3-multiw}, neither R2 nor R3 display significant emission in X-ray, UV, optical, or IR images.  These sources are significantly extended (see Figure~\\ref{extended-sources}). They are separated by about half an arcmin and could possibly be associated with jets from a background AGN. FIRST contains about 7 sources deg$^{-2}$ brighter than 20 mJy; the probability of finding one of them within a region of 100~arcsec radius chosen at random is only $\\sim 2$ per cent, but their observed fluxes and extents are difficult to reconcile with a stellar origin. Moreover, the chance of two unrelated sources of such high brightness occurring so close together is extremely low, favouring the background AGN scenario. Thus, the only radio point source above 0.3~mJy that may be related to star formation in NGC\\,3351 is R1.  In particular, we have not detected any point source in the nuclear region. The thermal emission predicted for each of the CNSFR complexes ranges from 0.03 to 0.36\\,mJy, slightly below our detection limit\\footnote{For the requested observing time of 48 hours, the expected noise level would have been $\\sigma_{\\rm rms}=0.03$~mJy under good weather conditions.}. The estimated fluxes are fairly robust with respect to the adopted sizes. Smaller diameters, appropriate for the individual knots \\citep[see e.g.][]{Hagele+07}, yield higher filling factors and optical depths, smaller areas, and, then, similar fluxes. On the other hand, a number of supernovae is expected to be present in these regions. The expected probability of finding a SNR brighter than 0.3\\,mJy in the circumnuclear ring is relatively high ($\\sim 50$~per cent), and the synchrotron flux associated to the unresolved population within each CNSFR (last column of Table~\\ref{tab:regions}) would have been easily detected by MERLIN. Therefore, the observed lack of emission suggests that these systems are younger than the lifetimes of the most massive stars \\citep[see e.g.][]{Rosa-Gonzalez+07a}. This seems also to be the case in similar CNSFRs observed in other galaxies \\cite[e.g.][in NGC\\,6951]{Dors+08}.  Regarding the star formation activity in the rest of the central region, the predicted number of bright SNR is also $\\sim 0.5$, and the total synchrotron flux is estimated to be about 36~mJy. Indeed, VLA observations at a minimum resolution of $4.5 \\times 4.1$~arcsec reveal a large central core, about 20~arcsec in size, featuring a double emission peak with maximum surface brightness about 10$^{10}$ mJy sr$^{-1}$ \\citep{Paladino06}, $\\sim 0.05$ mJy per MERLIN beam (close to the rms noise of our observations) if the emission was uniformly distributed. If star formation takes place in compact knots, similar in size to those in the CNSFRs, this emission ought to display conspicuous small-scale structure that should be detected by MERLIN. However, MERLIN is only sensitive to scales smaller than $\\sim2$ arcsec, which is $\\sim 100$ pc at the distance of NGC\\,3351. Therefore, the relativistic electrons generated by the supernova explosions must have diffused over distances longer than our maximum resolution radius of $\\sim 50$~pc, and the true diffusion scale would thus be somewhere in between this value and the 300~pc (6~arcsec) where the radio-IR correlation is observed by \\citet{Paladino06}.  The total flux density at 1.42\\,GHz, integrated over the whole galaxy, is predicted to be $\\sim 0.14$~Jy. According to equation~(\\ref{eqN_SNR_thr}), the expected number of radio sources above 0.3\\,mJy in the disk of NGC\\,3351 is of the order of one, consistent with the interpretation that R1 is a young SNR. However, an increase in flux of a factor of two between the FIRST observations and our data seems difficult to justify in this scenario. Two possible explanations would be that FIRST caught the supernova near the time of the explosion, where the radio continuum was still rising, and we are now observing a very young remnant, as suggested by its high luminosity. An alternative possibility would be that the radio flux has increased due to density variations in the ambient medium. Some Galactic nebulae, such as Cas A \\citep{Erickson+75,Read+77,Chevalier+78} and GK Per \\citep{Anupama+05} also show anomalous increases in flux density, albeit at lower frequencies. Ideally SNRs could be identified based on the radio spectral index, and future observations will be sought based on the results of the current work."
    },
    "1004/1004.4036_arXiv.txt": {
        "abstract": "We present results on the Soft Gamma Repeater (SGR)~0501+4516, discovered by the \\textit{Swift} Burst Alert Telescope (BAT) on 2008 August 22. More than 50 bursts were identified from this source, out of which 18 bursts had enough counts to carry out spectral analysis. We performed time-averaged spectral analysis on these 18 bursts using 8 models, among which the cut-off powerlaw and the two-blackbody models provided the best fit in the 15--150 keV energy range. The cut-off powerlaw model fit yields a mean photon index $\\Gamma_{CPL}$ = 0.54$\\pm$0.11 and a cut-off energy $E_C$ = 19.1$\\pm$1.8 keV for the bursts. The mean hard and soft blackbody temperatures are found to be $kT_{BB_h}$ = 12.8$\\pm$0.7 keV and $kT_{BB_s}$ = 4.6$\\pm$0.5 keV, respectively, and are anti-correlated with the square of the radii of the hard and soft emitting regions ($R_{BB_h}$ and $R_{BB_s}$) as $R_{BB_h}^2$ $\\propto$ $kT^{-5.8}$ and $R_{BB_s}^2$ $\\propto$ $kT^{-2.7}$, respectively. The soft and hard component temperatures with different indices support the idea of two distinct emitting regions with the hard component corresponding to a smaller radius and the soft component corresponding to a larger radius, which further corroborate the idea of the propagation of extraordinary (E) and ordinary (O) mode photons across the photosphere, as predicted in the magnetar model. We notice strong burst fluence--duration correlation as well as hardness ratio--duration and hardness ratio--fluence anti-correlations for the SGR 0501+4516 bursts. The burst fluences range from $\\sim$ 4.4$\\times$10$^{-9}$ ergs~cm$^{-2}$ to $\\sim$ 2.7$\\times$10$^{-6}$ ergs~cm$^{-2}$, consistent with those observed for typical short SGR bursts. ",
        "introduction": "The \\textit{Swift} mission (Gehrels et al. 2004), launched in November 2004 to explore the gamma-ray bursts (GRBs), has provided an excellent opportunity to detect and study the $\\gamma$-ray activities from known Soft Gamma Repeaters (SGRs) as well as to discover new SGRs. SGRs, belonging to the class of `magnetars' (Duncan \\& Thompson 1992), are highly magnetized ($B$ $\\sim$ 10$^{14-15}$ G) and slowly rotating ($P$ $\\sim$ 2$-$12 s) neutron stars characterized by short, bright bursts of hard X-rays and soft $\\gamma$-rays (see Mereghetti 2008 for a recent review). During its active state, an SGR can go through periods of intense bursting activity lasting from a few days to months, however, it can also remain dormant for many years. The bursts, often varying in duration, are classified into three main categories: short, intermediate and giant flares. The short bursts are the typical kind observed during an outburst, marked by timescales $\\sim$ 0.1--0.5~s and luminosities $\\sim$ 10$^{38}$--10$^{41}$ ergs s$^{-1}$, whereas the intermediate bursts have timescales $\\sim$ 1--60~s and luminosities $\\sim$ 10$^{41}$--10$^{43}$~ergs~s$^{-1}$ (Mereghetti 2008). Giant flares, the rare and unique events, are distinguished by their extreme energies (10$^{44}$--10$^{47}$~ergs~s$^{-1}$), long durations ($\\sim$ hundreds of seconds), and the presence of a coherent pulsating decaying tail, consistent with the spin period of the neutron star. Persistent X-ray emission, interpreted as originating due to the magnetospheric currents driven by twists in the evolving ultra-high magnetic field (Thompson \\& Duncan 1995), has also been observed from SGRs in the 0.1--10 keV band with a typical X-ray luminosity of $\\sim$ 10$^{35}$ ergs~s$^{-1}$, and the spectrum is generally described by an absorbed powerlaw (PL) plus a blackbody component (Mereghetti 2008).   According to the magnetar model (Thompson \\& Duncan 1995), the dominant form of energy powering an SGR is its decaying ultra-strong magnetic field. The surface of the neutron star is heated and fractured by instabilities generating Alfv\\'{e}n waves that accelerate electrons, and in turn give away their energy in short bursts. The model also suggests that the high-energy dissipated remains trapped in the magnetosphere as the `trapped fireball', and it shrinks in size with time. Alternatively, the bursts can arise from heating of the corona by magnetic reconnection in the stellar magnetosphere resulting in intermediate-type flares (Lyubarsky 2002). The giant flares are caused by the sudden rearrangement of the star's magnetic field producing global crustal fractures (Thompson \\& Duncan 1995).  \\begin{table*}[th] \\caption{Summary of SGR 0501+4516 bursts observed with \\textit{Swift-BAT}.} \\begin{tabular}{l l l l l l} \\hline Triggers  & Date of Observation & Exposure$^{a}$  & Bursts & T$_{100}$  & Count-rate$^b$ \\\\ & & (ks) & & (s) & (counts s$^{-1}$) \\\\ \\hline  00321174000 & 2008 August 22 & 34.8 & 321174A  & 0.132 & 1.22$\\pm$0.06 \\\\ & & &  321174B & 0.036 & 0.58$\\pm$0.07 \\\\ \\hline 00321177000 & 2008 August 22 & 3.0 & 321177  & 0.150 & 1.33$\\pm$0.06  \\\\ \\hline 00321174001 & 2008 August 22 & 12.8 & 321174001 & 0.200 & 0.48$\\pm$0.03 \\\\ \\hline 00321252000 & 2008 August 22 & 2.6 & 321252  & 0.078 & 0.43$\\pm$0.04 \\\\ \\hline 00321481000 & 2008 August 23 & 3.2 & 321481A  & 0.419 & 7.35$\\pm$0.21 \\\\ & & &  321481B  & 0.263 & 1.75$\\pm$0.05\\\\ \\hline 00321551000 & 2008 August 23 & 4.2 & 321551  & 0.487 & 15.2$\\pm$0.40\\\\ \\hline 00321574000 & 2008 August 23 & 3.4 & 321574  & 0.044 & 0.44$\\pm$0.06\\\\ \\hline 00321583000 & 2008 August 23 & 4.0 & 321583A & 0.043 & 0.61$\\pm$0.06\\\\ & & & 321583B  & 0.068 & 0.22$\\pm$0.04\\\\ & & & 321583C  & 0.092 & 0.43$\\pm$0.04\\\\ && &  321583D  & 0.119 & 0.62$\\pm$0.04\\\\ & & & 321583E  & 0.332 & 2.38$\\pm$0.06\\\\ & & & 321583F  & 0.140 & 1.11$\\pm$0.04\\\\ & & & 321583G  & 0.059 & 0.33$\\pm$0.04\\\\ \\hline 00323192000 & 2008 September 01 & 3.3 & 323192  & 0.044 & 0.33$\\pm$0.05\\\\ \\hline 00323650000 & 2008 September 03 & 4.3 & 323650  & 0.151 & 0.48$\\pm$0.06\\\\ \\hline \\end{tabular} \\tablecomments{T100 is the total burst duration obtained by noting the start and the end time of the burst, manually. See \\S\\ref{2} for details. \\\\ $^{a}$ Total BAT exposure for the trigger numbers. \\\\ $^b$ Count-rates are in the 15--150 keV energy range.} \\end{table*}   SGR~0501+4516 is a recently discovered SGR by the \\textit{Swift} $\\gamma$-ray observatory (Holland et al. 2008; Barthelmy et al. 2008). \\textit{Rossi X-ray Timing Explorer} (RXTE) observations discovered a 5.7~s X-ray pulsating counterpart (Gogus et al. 2008),  with a dipole magnetic field $B$ = 3$\\times$10$^{14}$ G  estimated from its period and spin-down rate (Woods et al. 2008), thus confirming the magnetar nature of the source. Soon after, the GLAST Burst Monitor (GBM) onboard \\textit{Fermi} satellite triggered and located the bursts from SGR 0501+4516 (van der Horst \\& Connaughton 2008; McBreen et al. 2008).  Multiwavelength observations reported the detection of its infrared (Tanvir et al. 2008; Rea et al. 2008) and optical (Fatkhullin et al. 2008; Rol et al. 2008) counterparts. However, no radio emission was detected from this source during the outburst (Kulkarni \\& Frail 2008; Gelfand et al. 2008). $Suzaku$ observations of the SGR performed from 2008 August 26--27 detected 32 bursts with the X-ray imaging spectrometer (XIS) and the hard X-ray detector (HXD). The persistent X-ray  emission obtained with the XIS in the 0.4--10 keV range was best fitted by the combination of a blackbody ($kT$ = 0.69$\\pm$0.01 keV) plus  a PL ($\\Gamma$ = 2.8$\\pm$0.1) component, modified by an interstellar absorption $N_H$ = (0.89$\\pm$0.08)$\\times$10$^{22}$~cm$^{-2}$ (Enoto et al. 2009). The 1--200 keV combined XIS+HXD burst spectrum was described by a two-component blackbody model having soft and hard temperatures of  $kT$ = 3.3$_{-0.4}^{+0.5}$ keV and 15.1$_{-1.9}^{+2.5}$ keV,  respectively, and was interpreted as the  population of ordinary (O) and extraordinary (E) mode photons propagating across the magnetosphere (Enoto et al. 2009).  \\textit{Konus-Wind} $\\gamma$-ray burst spectrometer observed the SGR bursts from 2008 August 23--26 and the 20--200 keV spectra  were best fitted by either a powerlaw with an exponential cut-off (CPL) model or an optically-thin thermal bremsstrahlung model (OTTB), both of  which gave a peak energy ($E_p$) or temperature ($kT$) in the range of 20--45 keV (Aptekar et al. 2009). Subsequently, five \\textit{XMM-Newton} and two \\textit{INTEGRAL} observations of the source were reported, with a hard X-ray variable source detected by \\textit{INTEGRAL} only during the first pointing (Rea et al. 2009).  The phase-coherent timing analysis performed with \\textit{XMM-Newton}, \\textit{Suzaku}-XIS and \\textit{Swift}-XRT data further refined the period $P$ = 5.7620695(1)~s, period derivative $\\dot{P}$ = 6.7(1)$\\times$10$^{-12}$~s~s$^{-1}$, and  magnetic field $B$ $\\sim$ 7$\\times$10$^{13}$ -- 2$\\times$10$^{14}$ G, with the evidence of a second period derivative  $\\ddot{P}$ = $-$1.6(4)$\\times$10$^{-19}$ s~s$^{-2}$ (Rea et al. 2009). Moreover, the spectrum of the outburst indicates a  trend of spectral softening, with the blackbody component decaying slower than the PL component and the spectral evolution changing  from phase to phase (Rea et al. 2009).   In this paper, we report the \\textit{Swift}-BAT observations of the recently discovered magnetar SGR~0501+4516. The paper is organized as follows: \\S\\ref{2} and \\S\\ref{3} describe the observations and the burst spectroscopic analysis, respectively. In \\S\\ref{4}, we present the results of the time-averaged spectral and statistical analyses of the SGR bursts, which are then discussed in \\S\\ref{5}. The conclusions are summarized in \\S\\ref{6}. ",
        "conclusions": "\\label{6}  We have presented a quantitative analysis of the newly discovered SGR 0501+4516 observed by \\textit{Swift}-BAT .  The spectra of the bursts from SGR 0501+4516 in the 15-150 keV energy range are best-fitted by both the CPL and BB+BB models. In the BB+BB model, the hotter temperature represents the smaller trapped fireball regions and the colder temperature corresponds to the regions consistent more or less with the radius of the neutron star. These temperatures are also  associated with the propagation of E-mode and O-mode photons across the photosphere as predicted in the magnetar model. The nonthermal emission observed from the short bursts seems to originate from the magnetosphere. All the above findings are consistent with the typical short bursts observed from other SGRs and with the magnetar model prediction of SGRs. However, we believe that more burst observations will help in further constraining the SGR emission mechanisms."
    },
    "1004/1004.1425_arXiv.txt": {
        "abstract": "We provide an analytical study of the coupling of short and long wavelength fluctuation modes during the initial phase of reheating in two field models like hybrid inflation. In these models, there is - at linear order in perturbation theory - an instability in the entropy modes of cosmological perturbations which, if not cut off, could lead to curvature fluctuations which exceed the current observational values. Here, we demonstrate that the back-reaction of short wavelength fluctuations is too weak to lead to a truncation of the instability for the long wavelength modes on time scales comparable to the typical instability time scale of the long wavelength entropy modes. Hence, unless there are other mechanisms which truncate the instability, then in models such as hybrid inflation the curvature perturbations produced during reheating on scales of current observational interest may be very important. ",
        "introduction": "The inflationary scenario \\cite{Guth} is the current paradigm of early universe cosmology. It addresses several conceptual problems which Standard Big Bang cosmology was unable to solve, and - perhaps more importantly - provides a causal mechanism to generate the primordial cosmological fluctuations which lead to the large-scale structure and cosmic microwave background anisotropies which are currently observed \\cite{ChibMukh}. Reheating after the period of inflation is a key aspect of inflationary cosmology (see the recent review of \\cite{Reheatingrev}). During reheating the energy is transformed from the inflaton, the scalar matter field which is responsible for providing the inflationary expansion of space, to the matter we see today. Without reheating, the inflationary universe would leave behind a large universe empty of any matter, in obvious contradiction with observations. Reheating requires some coupling between the inflaton and regular matter.  The initial studies of reheating were based on first order perturbation theory applied to the inflaton decay \\cite{initial}. This approach, however, fails to take into account the coherent nature of the inflaton field. At the end of the period of inflation, the inflaton begins coherent homogeneous oscillations about the ground state of its potential. In \\cite{TB} (see also \\cite{DK,STB}) it was realized that this gives rise to a potential parametric instability in any matter field which couples to the inflaton. This instability can take various forms depending on the nature of the coupling between the inflaton and the matter field. Rather generically, the process leads to a rapid energy transfer from the inflaton to the matter fields \\footnote{Rapid meaning on a time scale less than the Hubble time at the end of the period of inflation.}. In a wide class of models, the resonance is of ``broad resonance\" \\cite{KLS1} type and affects all fluctuations with wavenumbers smaller than a characteristic mass set by the particle physics Lagrangian. In some cases - in particular in the hybrid inflation model which we study in this paper - there is a tachyonic instability in the matter sector at the end of inflation, in which case the reheating instability is of ``tachyonic resonance\" \\cite{tachyonic} type and hence extremely efficient.  The inflaton field obviously couples to gravity, and therefore, as first conjectured in \\cite{BKM1}, it is possible that its oscillations will lead to a parametric resonance instability of the cosmological fluctuation modes. Due to the exponential growth of the causal horizon during the inflationary phase with constant Hubble radius, even super-Hubble modes can be causally affected by this instability \\cite{FB1}. This will inherently cause a long wavelength curvature mode to develop which potentially exhibits a tachyonic resonance. These large scale modes are of prime importance, since they correspond to scales observable today. If an inflationary model fails to control this possible instability of the curvature mode, the observed smallness of the magnitude of curvature perturbations today might rule out this model.  If matter consists of a single scalar field, then - as studied in \\cite{FB1} and \\cite{Afshordi} - there is no parametric instability of super-Hubble scale modes. However, if entropy fluctuations are present, then an instability of entropy modes can occur which will lead to rapid growth of the curvature fluctuations. For this instability to be effective, the entropy mode cannot have been exponentially redshifted during the period of inflation. Models which satisfy the conditions for parametric instability of the metric fluctuations were first discussed in \\cite{BaVi} and \\cite{FB2}.  If the duration of the parametric instability of long wavelength curvature fluctuations were long, then the curvature fluctuations induced by the resonant entropy modes could become larger than the primordial curvature perturbations and - in fact - larger than the observed amplitude of the inhomogeneities. Thus, one might potentially be able to use the parametric resonant instability of the entropy modes of the cosmological fluctuations to rule out large classes of multi-field inflationary models.  One class of inflationary models which at linear order in perturbation theory leads to an instability of the entropy fluctuations is hybrid inflation \\cite{hybrid}. In this scenario, two scalar fields $\\phi$ and $\\psi$, are involved in the inflationary process, the slowly rolling inflaton field $\\phi$ and the ``waterfall\" field $\\psi$. The key point of this class of models is that while $\\phi$ is slowly rolling during inflation, the energy density of the Universe is dominated by the potential of $\\psi$. Over the past few years, a lot of interest has been devoted to these models, mainly since they provide a framework for the realization of inflation in the context of both supersymmetry \\cite{Rachel}  and string theory (see \\cite{stringinflrev} for reviews and \\cite{Tye} for an original reference). Among the promising approaches stand the D-brane / antibrane inflation models (e.g. the D3/D7 brane inflation model \\cite{Keshav} and the KKLMMT model \\cite{KKLMMT}). Though these models may resolve some of the conceptual problems from which simple single scalar-field  driven inflation models suffer, the large number of light moduli fields present in string compactifications can give rise to entropy fluctuation modes, which could in turn enter a parametric resonance phase at early stages of reheating \\cite{BaVi,FB2}.  Hybrid inflation models are characterised by symmetry breaking along the $\\psi$ direction: The symmetry $\\psi~\\rightarrow-\\psi$ is spontaneously broken for field values $\\phi$ smaller than a critical value $\\phi_c$. Inflation takes place while $\\phi$ is slowly rolling at field values $\\phi > \\phi_c$. Once the inflaton crosses the critical value, the effective potential for $\\psi$ develops a tachyonic instability which triggers the rapid rolling of $\\psi$ towards one of the ground states of its potential. This causes inflation to end.  Reheating in hybrid inflation models on a fixed Friedmann background cosmology (i.e. no cosmological perturbations) was studied in detail in \\cite{Felder1,Felder2} (see also \\cite{hybridpreh}). By means of numerical simulations it was observed that within a very short time, non-linearities on a length scale given by the mass of the waterfall field develop and dominate the subsequent stages of the reheating process. It is important to study how large-scale metric perturbations evolve during reheating in these models. In previous work, linear evolution of the entropy and curvature modes was studied \\cite{BDD,BFL} \\footnote{Other work on fluctuations beyond linear analysis in hybrid inflation model focused on non-Gaussianities \\cite{FIN,Barnaby}}, and it was shown that it is possible that an important entropy perturbation mode develops. In these works it was assumed that non-linear effects, e.g. the back-reaction of small-scale fluctuations, would unlikely have a dominant effect on large-scale fluctuation. Due to the large range of scales (the small wavelength modes we are interested in have wavelengths comparable to the inverse Hubble radius at the end of inflation whereas the wavelengths of modes we are interested in for current cosmological observations are of the order of $1 {\\rm mm}$, assuming that the scale of inflation is about $10^{15} {\\rm GeV}$) numerical studies do not have the dynamical range to study this question. Instead, an analytical understanding of the dynamics is required.  Our work demonstrates, by means of an analytical analysis, that the back-reaction of short wavelength modes with random phases is too weak to truncate the instability of the long wavelength entropy modes. We have studied the effects of short wavelength fluctuations both on the background inflaton and on long wavelength modes of the inflaton and waterfall fields. Most importantly, we have shown that the back-reaction on large scale fluctuations of the waterfall field is too weak to shut off the resonance of this entropy mode one the relevant time scale of the instability. This implies that hybrid inflation models of the type analyzed in this paper may suffer from a potential ``entropy mode problem\", unless there are other ways of truncating the resonance (e.g. the exponential decrease in the initial value of the fluctuations modes during the period of inflation - see e.g. \\cite{BFL} for a discussion) \\footnote{There is another way to test these models observationally: hybrid inflation models typically produce topological defects such as cosmic strings. These strings, if stable on cosmological scales as they are in field theory models of hybrid inflation, lead to a scaling solution of strings at all late times (see e.g. \\cite{CSrevs} for reviews on cosmic strings and structure formation). The strings, in turn, produce line discontinuities in cosmic microwave temperature anisotropy maps \\cite{KS} which can be searched for in observational maps using edge detection algorithms such as the Canny algorithm, as recently studied in \\cite{Canny}.}.  Our study of the back-reaction of fluctuations in hybrid inflation models is not the first analytical study. For an analytical study of the back-reaction effect of fluctuations on the background using very different techniques than the ones we use see \\cite{Katrin}. For a study of the effects of metric fluctuations on the observable local expansion rate of the universe due to fluctuations see \\cite{Ghazal}.  The structure of this paper is as follows. In Section 2, we review the class of hybrid inflation models which are studied here, and introduce the perturbative approach employed to study non-linear interactions that arise  during preheating. In Section 3, we solve the background theory, and in Section 4 we review the evolution of quantum fluctuations to linear order. In Section 5, we analyze the generation of long wavelength second order perturbation modes sourced by shorter wavelength first order fluctuations (summing over the contribution from all frequencies and assuming random phases), for both the inflaton and the waterfall field. This is the leading order back-reaction effect between short and long wavelength modes. We also study the back-reaction of the fluctuations on the background. We find that small-scale first-order modes with random phases have a contribution on the evolution of large-scale modes which does not start to dominate until long after the instability of the first order long wavelength modes has had a chance to develop. Hence, in spite of their large phase space, the non-linear effects of the short wavelength perturbations do not become dominant for the evolution of long wavelength modes during the  early stages of preheating. ",
        "conclusions": "We have considered classical dynamics in a two field inflation model like hybrid inflation in which one field - $\\phi$ - is slowly rolling during inflation and a second field, the so-called ``waterfall field\" $\\psi$, develops a tachyonic instability once $\\phi$ decreases below some critical value $\\phi_c$. We focus on the dynamics right after the instability develops, which corresponds to the initial stages of preheating at the end of inflation.  We have provided a perturbative analytical analysis of the back-reaction effects of linear fluctuations on the background fields and on the perturbations themselves during the initial phase of the tachyonic instability which arises at the end of a period of hybrid inflation. We conclude that the tachyonic growth of the short wavelength fluctuation modes leads - at second order in perturbation theory, the lowest order at which mode mixing appears - to contributions which could in principle shut off the tachyonic growth of long wavelength entropy modes. However, in the case of random phases for the linear fluctuations considered here, we find that time scale when the back-reaction terms become important is much longer than the typical instability time scale of the tachyonic field. This process cannot shut off the instability of entropy modes before the latter have had time to become important. Note that there are other mechanisms which could dramatically weaken the strength of the instability of the long wavelength modes of the waterfall field, e.g. the decrease in amplitude of the linear fluctuations on super-Hubble scales during inflation.  It is important to stress the crucial role which the assumption of random phases has played in our analysis. Had we worked with correlated phases, then the back-reaction effects would have been larger by a factor comparable to $L^{3/2}$ (in Planck units), and back-reaction would then be able to cut off the tachyonic instability of modes of the waterfall field long before they have time to develop.  Our work complements the numerical studies of this model \\cite{Felder1,Felder2}. Our analytical approach has as a big advantage that we can access the large dynamical range of scales which is required in order to draw conclusions for modes of cosmological interest today from dynamics taking place on microphysical scales at the end of inflation, the scales which are studied in the numerical works. Disadvantages of our approach based on a discrete Fourier space analysis are that it is purely perturbative and that we cannot study position space phenomena expected in the model such as the formation of topological defects.  The reader may worry that the back-reaction effect which we find on the adiabatic mode of the fluctuations, the $\\phi$ fluctuations, might affect the predicted curvature fluctuations of the model. However, this is not the case. On scales larger than the Hubble radius, the curvature fluctuations are dominated by gravity. Microphysical effects which occur at the time of reheating cannot affect the adiabatic modes of the curvature perturbations. This follows in full generality from the Traschen integral constraints \\cite{Traschen} which hold as long as we operate in the framework of General Relativity. In the case of scalar field models of inflation, it was shown in \\cite{FB1} that effects at preheating cannot affect the adiabatic mode of the curvature fluctuations.  The reader should also keep in mind that our analysis does not concern other sources of entropy fluctuations in models like the one considered here. Specifically, if topological defects are produced during preheating, then the active evolution of the defects between the time of formation and the present time will lead to a large contribution to the entropy fluctuations. In fact, the specific model considered in this paper, where both $\\phi$ and $\\psi$ are single component real fields, would lead to domain walls and would in fact be ruled out because of the domain wall problem \\cite{DW}, the over-abundance of energy in the walls."
    },
    "1004/1004.0804_arXiv.txt": {
        "abstract": "This is a summary of presentations delivered at the OC1 parallel session ``Primordial Gravitational Waves and the CMB\" of the $12^{\\rm th}$ Marcel Grossmann meeting in Paris, July 2009. The reports and discussions demonstrated significant progress that was achieved in theory and observations. It appears that the existing data provide some indications of the presence of gravitational wave contribution to the CMB anisotropies, while ongoing and planned observational efforts are likely to convert these indications into more confident statements about the actual detection. ",
        "introduction": "\\label{sec1}  Here, we report on presentations and discussions that took place at the OC1 parallel session entitled ``Primordial Gravitational Waves and the CMB\". The programme was designed to include both theoretical and observational results. The electronic versions of some of the talks delivered at this session are available at the MG12 website \\cite{mg12website}. We shall start from the overall framework of this session and motivations for studying the primordial (relic) gravitational waves.  Primordial gravitational waves (PGWs) are necessarily generated by a strong variable gravitational field of the very early Universe \\cite{gr74}. The existence of relic grtavitational waves relies only on the validity of basic laws of general relativity and quantum mechanics. Specifically, the generating mechanism is the superadiabatic (parametric) amplification of the waves' zero-point quantum oscillations. In contrast to other known massless particles, the coupling of gravitational waves to the external (``pump\") gravitational field is such that they could be classically amplified or quantum-mechanically generated by the gravitational field of a homogeneous isotropic FLRW (Friedmann-Lemaitre-Robertson-Walker) universe. Under certain extra conditions the same applies to the primordial density perturbations. The PGWs are the cleanest probe of the physical conditions in the early Universe right down to the limits of applicability of currently available theories, i.e. the Planck density $\\rho_{\\rm Pl} = c^5/G^2 \\hbar \\approx 10^{94}{\\rm g}/{\\rm cm}^{3}$ and the Planck size $l_{\\rm Pl} = (G \\hbar/c^3)^{1/2} \\approx 10^{-33}{\\rm cm}$.  The amount and spectral content of the PGWs field depend on the evolution of the cosmological scale factor $a(\\eta)$ representing the gravitational pump field. The theory was applied to a variety of $a(\\eta)$, including those that are now called inflationary models \\cite{gr74,a12a,starobinski}. If the exact $a(\\eta)$ were known in advance from some fundamental ``theory-of-everything\", we would have derived the properties of the today's signal with no ambiguity. In the absence of such a theory, we have to use the available partial information in order to reduce the number of options. The prize is very high - the actual detection of a particular background of PGWs will provide us with a unique clue to the birth of the Universe and its very early dynamical behaviour.  To be more specific, let us put PGWs in the context of a complete cosmological theory hypothesizing that the observed Universe has come to the existence with near-Planckian energy density and size (see papers \\cite{birthpapers,Grishchuk2009} and references therein). It seems reasonable to conjecture that the embryo Universe was created by a quantum-gravity or by a ``theory-of-everything\" process in a near-Planckian state and then started to expand. (If you think that the development of a big universe from a tiny embryo is arrant nonsense, you should recollect that you have also developed from a single cell of microscopic size. Analogy proposed by the biophysicist E. Grishchuk.) The total energy, including gravity, of the emerging classical configuration was likely to be zero then and remains zero now.  In order for the natural hypothesis of spontaneous birth of the observed Universe to bring us anywhere near our present state characterized by $\\rho_p =3 H_0^2/8 \\pi G \\approx 10^{-29} {\\rm g}/{\\rm cm}^3$ and $l_p = 500 l_H$ (which is the minimum size of the present-day patch of homogeneity and isotropy, as follows from observations \\cite{GrZeld}) the newly-born Universe needs a significant `primordial kick'. During the kick, the size of the Universe (or, better to say, the size of our patch of homogeneity and isotropy) should increase by about 33 orders of magnitude without losing too much of the energy density of whatever substance that was there, or maybe even slightly increasing this energy density at the expence of the energy density of the gravitational field. This process is graphically depicted in Fig.~\\ref{birth} (adopted from the paper \\cite{Grishchuk2009}). The present state of the accessible Universe is marked by the point P, the birth of the Universe is marked by the point B. If the configuration starts at the point B and then expands according to the usual laws of radiation-dominated and matter-dominated evolution (blue curve), it completely misses the desired point P. By the time the Universe has reached the size $l_p$, the energy density of its matter content would have dropped to the level many orders of magnitude lower than the required $\\rho_p$. The only way to reach P from B is to assume that the newly-born Universe has experienced a primordial kick allowing the point of evolution to jump over from the blue curve to the black curve.  If we were interested only in the zero-order approximation of homogeneity and isotropy, there would be many evolutionary paths equally good for connecting the points B and P. However, in the next-order approximations, which take into account the inevitable quantum-mechanical generation of cosmological perturbations, the positioning and form of the transition curve in Fig.~\\ref{birth} become crucial. The numerical value of the Hubble parameter $H$ (related to the energy density of matter driving the kick, as shown on the vertical axis of the figure) determines the numerical level of amplitudes of the generated cosmological perturbations, while the shape of the transition curve determines the shape of the primordial power spectrum.   \\begin{figure} \\includegraphics[width=12cm,height=10cm]{g3.eps} \\caption{A primordial kick is required in order to reach the present state of the Universe P from the birth event B. Red lines describe possible transitions that would be accompanied by the generated cosmological perturbations of observationally necessary level and spectral shape \\cite{Grishchuk2009}. The ``legitimate transition in inflationary theory\" is an evolution allowed by the incorrect (inflationary) formula for density perturbations. See explanations in Sec.{\\ref{section3}} below.} \\label{birth} \\end{figure}  The simplest assumption about the initial kick is that its entire duration was characterized by a single power-law scale factor \\cite{gr74} \\begin{eqnarray} \\label{scfactor} a(\\eta) = l_o|\\eta|^{1+\\beta}, \\end{eqnarray} where $l_o$ and $\\beta$ are constants, and $\\beta < -1$. In this power-law case, the gravitational pump field is such that the generated primordial metric power spectra (primordial means considered for wavelengths longer than the Hubble radius at the given moment of time), for both gravitational waves and density perturbations, have the universal power-law dependence on the wavenumber $n$: \\begin{equation} \\label{primsp} h^{2}(n) \\propto n^{2(\\beta+2)}. \\end{equation} It is common to write these metric power spectra separately for gravitational waves (gw) and density perturbations (dp): \\begin{equation} \\label{primsp1} h^2(n)~({\\rm gw}) = B_t^2 n^{n_t}, ~~~~~h^2(n)~({\\rm dp})=B_s^2 n^{n_s -1}. \\end{equation}  According to the theory of quantum-mechanical generation of cosmological perturbations (for a review, see \\cite{a12a}), the spectral indices are approximately equal, $n_s-1 = n_t = 2(\\beta+2)$, and the amplitudes $B_t$ and $B_s$ are of the order of magnitude of the ratio $H_i/H_{\\rm Pl}$, where $H_i\\sim c/l_o$ is the characteristic value of the Hubble parameter $H$ during the kick. The straight lines in Fig.~\\ref{birth} symbolize the power-law kicks (\\ref{scfactor}). They generate primordial spectra with constant spectral indices throughout all wavelenghts. In particular, any horizontal line describes an interval of de Sitter evolution, $\\beta = -2$, ${\\dot H} =0$, $H = const$. (Initial kick driven by a scalar field is appropriately called inflation: dramatic increase in size with no real change in purchasing power, i.e. in matter energy density.) The gravitational pump field of a de Sitter kick transition generates perturbations with flat (scale-independent) spectra $n_s-1 = n_t = 0$. The red horizontal line shown in Fig.~\\ref{birth} corresponds to $H_i/H_{\\rm Pl} \\approx 10^{-5}$  and the generated primordial amplitudes $B_t \\approx B_s \\approx 10^{-5}$. In numerical calculations, the primordial metric power spectra are usually parameterized by \\begin{eqnarray} \\label{PsPt} P_{t}(k)=A_t(k/k_0)^{n_t},~~P_{s}(k)=A_s(k/k_0)^{n_s-1}, \\end{eqnarray} where $k_0=0.002$Mpc$^{-1}$.  Although the assumption of a single piece of power-law evolution (\\ref{scfactor}) is simple and easy to analyze, the reality could be more complicated and probably was more complicated (see the discussion of WMAP data in Sec.\\ref{section3}). A less simplistic kick can be approximated by a sequence of power-law evolutions, and then the primordial power spectra will consist of a sequence of power-law intervals.  The amplitudes of generated cosmological perturbations are large at long wavelengths. According to the widely accepted assumption the observed anisotropies in the cosmic microwave background radiation (CMB) are caused by perturbations of quantum-mechanical origin. This assumption is partially supported by the observed ``peak and dip\" structure of the CMB angular power spectra. Presumably, this structure reflects the phenomenon of quantum-mechanical phase squeezing and standing-wave pattern of the generated metric fields \\cite{a12a}. The search for the relic gravitational waves is a goal of a number of current and future space-borne, sub-orbital and ground-based CMB experiments \\cite{Planck,WMAP5,BICEP,quad,Clover,QUITE,EBEX,SPIDER,cmbpol}.  The CMB anisotropies are usually characterized by the four angular power spectra $C_{\\ell}^{TT}$, $C_{\\ell}^{EE}$, $C_{\\ell}^{BB}$ and $C_{\\ell}^{TE}$ as functions of the multipole number $\\ell$. The contribution of gravitational waves to these power spectra has been studied, both analytically and numerically, in a number of papers \\cite{Polnarev1,a8,a11,a12,a13}. The derivation of today's CMB power spectra brings us to approximate formulas of the following structure \\cite{a12}: \\begin{eqnarray} \\label{exact-clxx'} \\begin{array}{l} C_{\\ell}^{TT}= \\int \\frac{dn}{n} [h(n, \\eta_{rec})]^2 \\left[F^T_{\\ell}(n)\\right]^2, \\\\ C_{\\ell}^{TE}=\\int \\frac{dn}{n} h(n, \\eta_{rec}) h^{\\prime}(n, \\eta_{rec}) \\left[F^T_{\\ell}(n) F^E_{\\ell}(n)\\right], \\\\ C_{\\ell}^{YY}=\\int \\frac{dn}{n} [h^{\\prime} (n, \\eta_{rec})]^2 \\left[F^Y_{\\ell}(n)\\right]^2, ~~~~{\\rm where}~Y=E,B. \\end{array} \\end{eqnarray} In the above expressions, $[h(n, \\eta_{rec})]^2$ and $[h^{\\prime} (n, \\eta_{rec})]^2$ are power spectra of the gravitational wave field and its first time-derivative. The spectra are taken at the recombination (decoupling) time $\\eta_{rec}$. The functions $F^X_{\\ell}(n)$ ($X=T, E, B$) take care of the radiative transfer of CMB photons in the presence of metric perturbations. To a good approximation, the power residing in the metric fluctuations at wavenumber $n$ translates into the CMB $TT$ power at the angular scales characterized by the multipoles $\\ell \\approx n$. Similar results hold for the CMB power spectra induced by density perturbations. Actual numerical calculations use equations more accurate than Eq.~(\\ref{exact-clxx'}).   There are several differences between the CMB signals arising from density perturbations and gravitational waves. For example, gravitational waves produce B-component of polarization, while density perturbations do not \\cite{a8}; gravitational waves produce negative TE-correlation at lower multipoles, while density perturbations produce positive correlation \\cite{a12a,a12,polnarev,zbg,zbg2}, and so on. However, it is important to realize that it is not simply the difference between zero and non-zero or between positive and negative that matters. (In any case, various noises and systematics will make this division much less clear cut.) What really matters is that the gw and dp sources produce different CMB outcomes, and they are in principle distinguishable, even if they both are non-zero and of the same sign. For example, if the parameters of density perturbations were precisely known from other observations, any observed deviation from the expected $TT$, $TE$ and $EE$ correlation functions could be attributed (in conditions of negligible noises) to gravitational waves. From this perspective, the identification of the PGWs signal goes well beyond the often stated goal of detecting the B-mode of polarization. In fact, as was argued in the talk by D.~Baskaran, delivered on behalf of the group including also L.~P.~Grishchuk and W.~Zhao, the $TT$, $TE$, and $EE$ observational channels could be much more informative than the $BB$ channel alone. Specifically for the Planck mission, the inclusion of other correlation functions, in addition to $BB$, will significantly increase the expected signal-to-noise ratio in the PGWs detection.  It is convenient to compare the gravitational wave signal in CMB with that induced by density perturbation. A useful measure, directly related to observations, is the quadrupole ratio $R$ defined by \\begin{eqnarray} \\label{defineR} R \\equiv \\frac{C_{\\ell=2}^{TT}({\\rm gw})}{C_{\\ell=2}^{TT}({\\rm dp})}, \\end{eqnarray} i.e. the ratio of contributions of gw and dp to the CMB temperature quadrupole. Another measure is the so-called tensor-to-scalar ratio $r$. This quantity is built from primordial power spectra (\\ref{PsPt}): \\begin{eqnarray} \\label{definer} r \\equiv \\frac{A_t}{A_s}. \\end{eqnarray} Usually, one finds this parameter linked to incorrect (inflationary) statements. Concretely, inflationary theory substitutes (for ``consistency\") its prediction of arbitrarily large amplitudes of density perturbations in the limit of models where spectral index $n_s$ approaches $n_s=1$ and $n_t$ approaches $n_t=0$ (horizontal transition lines in Fig.~\\ref{birth}) by the claim that it is the amount of relic gravitational waves, expressed in terms of $r$, that should be arbitrarily small. (For more details, see Sec.~{\\ref{section3}}.) However, if $r$ is defined by Eq.(\\ref{definer}) without implying inflationary claims, one can use this parameter. Moreover, one can derive a relation between $R$ and $r$ which depends on the background cosmological model and spectral indices. For a rough comparison of results one can use $r\\approx2R$. ",
        "conclusions": "In general, the OC1 parallel session was balanced and covered both theoretical and experimental aspects of primordial gravitational waves and the CMB. Thanks to excellent contributions of all participants, this subject received a new momentum. Hopefully, new observations and theoretical work will bring conclusive results in the near future."
    },
    "1004/1004.0942_arXiv.txt": {
        "abstract": "An extension of the MSSM called the $\\mu\\nu$SSM does not allow a conventional thermal leptogenesis scenario because of the low scale seesaw that it utilizes.  Hence, we investigate the possibility of electroweak baryogenesis.  Specifically, we identify a parameter region for which the electroweak phase transition is sufficiently strongly first order to realize electroweak baryogenesis.  In addition to transitions that are similar to those in the NMSSM, we find a novel class of phase transitions in which there is a rotation in the singlet vector space. ",
        "introduction": "An extension of the MSSM called the $\\mu\\nu$SSM \\cite{LopezFogliani:2005yw} is a model similar to the NMSSM \\cite{Ellis:1988er} (with the usual $\\mathbb{Z}_{3}$ charge assignment) except that the singlet whose vacuum expectation value (VEV) gives rise to the $\\mu$ term also serves the role of a right-handed neutrino, thereby violating R-parity. Because the VEV generates the $\\mu$-term and the right handed neutrino mass, the right-handed neutrino masses are of order TeV, leading to a low scale type I seesaw. Given the absence of a high scale seesaw, thermal leptogenesis is difficult in the $\\mu\\nu$SSM.  Hence, it is interesting to consider whether or not electroweak baryogenesis (EWBG) \\cite{Kuzmin:1985mm} can occur in this class of models.   One of the most stringent constraints of EWBG on the $\\mu\\nu$SSM is the requirement of a sufficiently strongly first order phase transition (SFOPT) such that the created baryons are not washed out \\cite{Bochkarev:1990gb}.  Because the $\\mu\\nu$SSM contains 3 singlet chiral superfields (right handed neutrinos), mainly motivated by generality, standard model generation replication pattern, and phenomenological convenience \\cite{LopezFogliani:2005yw,Escudero:2008jg}, there is a ``larger'' SFOPT parameter space for EWBG when compared to the NMSSM.  More precisely, there can be SFOPT where the singlet VEVs rotate in the singlet vector space during the electroweak phase transition.  The price paid for this is a more complicated global minimum analysis at both finite and zero temperatures.  The aim of this paper is not to uncover the most general parameter space consistent with EWBG, but is to simply give a couple of parametric regions to show the existence of possibilities.  Depending on the path of the phase transition, the exact $\\mu\\nu$SSM parametric dependence of the phase transition strength $v(T_c)/T_c$ is complicated.  Nonetheless, we find that it is typically true that to achieve SFOPT, the parameters are close to satisfying the following condition: \\begin{equation} \\frac{E_{\\eff}}{\\lambda_{\\eff} v(0)} \\approx \\frac{1}{2} \\end{equation} where $E_{\\eff}$ is the effective cubic coupling, $\\lambda_{\\eff}$ is the effective quartic coupling, and $v(0)$ is the magnitude of the scalar field space VEV (including both the Higgs and singlets) at zero temperature.  Physically, this corresponds to the parametric region where the critical temperature $T_c$ is small compared to $v(0)$ during the electroweak phase transition.  In the examples provided in this paper, whether or not the SFOPT proceeds from the origin, the leading nonvanishing value of $E_{\\eff}$ in the $\\mu\\nu$SSM arises from the soft terms \\begin{equation} \\sum_{i}^{3}\\left[-a_{\\lambda}H_{1}H_{2}\\tilde{\\nu}_{i}^{c}+\\frac{1}{3}a_{\\kappa}\\left(\\tilde{\\nu}_{i}^{c}\\right)^{3}+\\hc \\right] \\end{equation} where $\\tilde{\\nu}_{i}^{c}$ are singlet fields.  The dimensionful coupling $a_\\lambda$ is distinguished from $a_\\kappa$ in that $a_\\lambda$ also enhances the mixing between the Higgs sector and the singlet sector.  The leading contribution to $\\lambda_{\\eff}$ comes from the superpotential and D-terms.  Beyond these general results, we find a somewhat interesting feature because we focus on the parametric region analyzed by \\cite{Escudero:2008jg}.  In this parametric region, an approximate $\\mathbb{S}_3$ symmetry (permutation symmetry) arises due to the right handed neutrino generation independence of the non-Yukawa couplings and the smallness of the neutrino Yukawa couplings.  Hence, to avoid any extra complications associated with domain wall formations, one might naively try to avoid $\\mathbb{S}_3$ symmetry breaking phase transitions by considering parameters which yield zero temperature vacua preserving $\\mathbb{S}_3$.  Hence, this is the boundary condition that we impose in this paper.  Interestingly, we find that despite this boundary condition, $\\mathbb{S}_3$ is typically spontaneously broken multiply at finite temperatures in a way that is sensitive to quantum radiative corrections.  As the temperature is lowered from high temperatures, this leads to multistep phase transitions starting from the trivially $\\mathbb{S}_3$ symmetric vacuum in which all VEVs vanish. The electroweak symmetry breaking phase transition occurs with $\\mathbb{S}_3$ symmetry {\\em restoration} to a vacuum in which all sneutrino VEVs are identical and nonvanishing.  We also find one step SFOPTs in which the scalar fields (including the singlet fields) make a transition from the origin to the electroweak symmetry breaking minimum.  The numerical values of the parametric regions uncovered in this paper is in the paragraph containing Eq.~(\\ref{eq:sigmadef}) and regions IIIa and IIIb depicted in Fig.~\\ref{fig:Scan1}.  There have been many studies of EWBG and the electroweak phase transition in models with gauge singlets; some of these are Refs.~\\cite{McDonald:1993ey, Ham:2004nv, Huber:2006wf, Ham:2007xz, Ham:2007wu, Ham:2007wc, Chiang:2009fs, Kang:2009rd, Kang:2004pp}, and we will discuss others throughout the remainder of this section. Since our work is most closely related to previous work on SFOPT in the NMSSM, we give here a little preview of some of the differences between our work and select previous works, in addition to the multidimensional aspect stressed above. In Ref.~\\cite{Pietroni:1992in}, SFOPT in the context of the NMSSM is first analyzed and the author points out that the tree level cubic term coming from the soft SUSY-breaking sector is important. Note that Ref.~\\cite{Pietroni:1992in} uses the definition of critical temperature in which scalar mass squared matrix develops a vanishing eigenvalue. We take a more robust definition of $T_{c}$ being the temperature at which a new coexistence phase occurs even though this definition is harder to implement in practice.  The authors of Ref.~\\cite{Davies:1996qn} also analyze the NMSSM, but they include a $\\mu$-term on the basis that it is more general and its nonzero value eliminates the $\\mathbb{Z}_{3}$ symmetry which can be cosmologically dangerous with respect to the problem of domain wall formation \\cite{Abel:1995wk}.  The nonzero $\\mu$-term leads to false vacuum not being at the origin.  In this case the critical temperature criterion used by Ref.~\\cite{Pietroni:1992in} is invalid. Therefore, the authors of Ref.~\\cite{Davies:1996qn} take the coexistence phase definition of critical temperature as we do in this paper.  They also include a bilinear soft term in the Higgs which breaks the $\\mathbb{Z}_{3}$ symmetry.  Although we do not include such $\\mathbb{Z}_{3}$ breaking terms directly, we will assume that nonrenormalizable terms can be included to obtain acceptable phenomenology with respect to any possible domain wall formation. However, it is to be noted that $\\mathbb{Z}_3$ breaking can often lead to UV instabilities in the singlet tadpoles, making the UV stability of these theories (including the one considered in this paper) a model building challenge as noted by \\cite{Abel:1995wk}.  The analysis \\cite{Huber:2000mg} considers the generalized NMSSM similar to \\cite{Davies:1996qn}. They run 9 parameters with a popular choice of ``universal'' boundary conditions from the GUT scale down to the electroweak scale to generate their model. They do not reject metastable vacua based on the intuition that longevity of the false vacuum on the horizon scale today is not difficult to attain.  To be conservative and to avoid potentially complicated discussions of metastability, we accept only stable vacua in this paper.  A model related to the NMSSM and the $\\mu\\nu$SSM is the nMSSM in which the discrete charge assignment is modified as to eliminate the singlet cubic term in the superpotential. This model was analyzed by \\cite{Menon:2004wv} for SFOPT. For a significant portion of the parameter space in which SFOPT occurs, a linear tadpole term in the superpotential plays a significant role in contrast to our scenario.   The analysis of \\cite{Profumo:2007wc} considers the EWPT in an extension of the SM which adds a real singlet $S$.  These authors find a large region of the parameter space of their model that is consistent with SFOPT and LEP Higgs search bounds.  They argue that the strength of the phase transition can be enhanced by 1) choosing a large negative value for the $S H^2$ coupling, 2) choosing a negative value for the $S^2 H^2$ coupling, or 3) allowing the singlet to have a nonzero VEV before the electroweak symmetry is broken.  In the language of this paper, the first two points correspond to increasing $E_{\\mathrm{eff}}$ and decreasing $\\lambda_{\\mathrm{eff}}$, respectively.  Before we begin the main body of the work, let us list here all the caveats to our analysis.  We do not take into account explicitly the high energy Landau pole constraint (i.e. perturbativity up to the GUT scale) because we will take the attitude that the $\\mu\\nu$SSM is well motivated mainly by its ability to have all fields participate at low energy and thereby have potential measurability. Nonetheless, the parametric region that we uncover lies at the border of perturbativity up to the GUT scale (inferring from the work of Refs.~\\cite{Escudero:2008jg,Miller:2003ay}), which means that the UV cutoff for our theory can be taken to be far above the TeV scale.  We do not take into account explicit $\\mathbb{Z}_{3}$ breaking effects because a small amount of breaking can address most cosmological domain wall problems, as we later demonstrate.  We do not take into account explicit CP violation effects in the phase transitions as this will typically lead to less than order 10\\% effects since CP violating phases compatible with phenomenology are typically order $0.1$ or smaller.  For robustness, we accept in this paper as phenomenological possibility only absolutely stable global zero temperature vacua instead of analyzing the phenomenological possibilities of metastable vacua.  Finally, all of our numerical work is kept in control to only order 10\\% accuracy.  The order of presentation is as follows.  In the next section, we present the Lagrangian including its discrete symmetry properties and radiative/thermal corrections.  The section concludes by highlighting the $\\mu\\nu$SSM differences from the NMSSM scenario.  In Sec.~\\ref{sec:qualitativedesc}, we describe the parametric region relevant for SFOPT in terms of one-dimensional field space slice parametrization.  There we also qualitatively describe how the multidimensional paths of the phase transition and discrete symmetries play a role.  Next, in Sec.~\\ref{sec:sphaleron}, we explicitly show that singlets do not play a significant role in terms of numerical value of the sphaleron action controlling the $B+L$ violating rate in the broken phase.  The main numerical results are presented in Sec.~\\ref{sec:paramscan} where the explicit existence of the SFOPT parameter region is demonstrated.  Details of the transition paths organized in terms of discrete symmetries, phenomenological bounds placed, and explicit mass spectra for a sample parametric point are given.  In Sec.~\\ref{sec:domainwall}, we demonstrate that the cosmological domain wall problem is easily evaded with an inclusion of a weak $\\mathbb{Z}_3$ symmetry breaking operator in our scenario.  We then conclude with a summary of the results.  Several appendices then follow giving useful technical details.  In Appendix A, we list the field-dependent mass matrices used for computing the effective potential.  In the next appendix, we give details regarding the approximate thermal masses used in the paper.  In Appendix C, we describe analytically the boundaries of in Fig. \\ref{fig:Scan1} which is one of our main results.  Finally, in Appendix D, we show that it is generically possible to construct a nonrenormalizable $\\mathbb{Z}_3$ superpotential to obtain a CP conserving global minimum in the absence of any explicit CP violating parameters. ",
        "conclusions": "We have uncovered a $\\mu\\nu$SSM parametric region giving rise to a first order phase transition sufficiently strong to be useful for the electroweak baryogenesis scenarios involving electroweak symmetry breaking bubbles as the source of out of equilibrium and $SU(2)_L$ $F\\tilde{F}$ operators as a source of baryon number violation.  The parametric region corresponds to tuning the soft terms in the Lagrangian $a_{\\lambda}H_{1} \\! \\cdot \\!  H_{2}\\tilde{\\nu}_{i}^{c}$ and $-m_{\\tilde{\\nu}^{c}}^{2}|\\tilde{\\nu}_{i}^{c}|^{2}$ to achieve Eq.~(\\ref{eq:crucialsaturate}).  The numerical values of the uncovered parametric region is in the paragraph containing Eq.~(\\ref{eq:sigmadef}) and regions IIIa and IIIb depicted in Fig.~\\ref{fig:Scan1}.  As expected, the Yukawa coupling of the singlets to the leptonic sector does not play a role in determining the strength of the phase transitions because of the weakness of the coupling tied to the smallness of the neutrino masses.  The region IIIa transitions are two-step transitions in which the electroweak symmetry breaking is the second transition that starts from a phase in which the singlet scalars of the $\\mu\\nu$SSM have a nonzero vacuum expectation value (e.g. starts from a vacuum which spontaneously breaks the approximate $S_3$ symmetry in the singlet sector).  These transitions contain a rotation in the singlet field space and do not have an analog in the NMSSM transitions because of the different dimensionality in the singlet complex vector space.  The region IIIb transitions are the ones in which electroweak symmetry breaking transition starts from the origin of the scalar field space. All these transitions have useful descriptions in terms of the representations of the approximate discrete symmetries in the system.  Our phenomenological bounds were rather minimal and placed using Ref.~\\cite{PDG}, but in many parametric regions, the observables are sufficiently far away from the bounds that the plausibility of the phenomenological self-consistency is strong.  Follow-up possibilities include a more complete collider related phenomenological investigation in this parametric regime, studies of domain wall histories due to weak global symmetry breaking operators, and a complete computation of CP asymmetry creation and transport leading to baryon asymmetry.  Given that the $\\mu\\nu$SSM had to give up the popular thermal leptogenesis scenario due to its low scale implementation of the type I seesaw, this work is of interest as it shows that electroweak baryogenesis may be a promising avenue to create baryon asymmetry in this class of models.  Given that the $\\mu\\nu$SSM is one of the few supersymmetric models in which all dynamical degrees of freedom responsible for the neutrino mass may be accessible at TeV-scale colliders, it is encouraging that the model has a good chance at being consistent with the observed baryon asymmetry in the universe."
    },
    "1004/1004.3953.txt": {
        "abstract": "We present a hidden abelian extension of the Standard Model including a complex scalar as a dark matter candidate and a light scalar acting as a long range force carrier between dark matter particles. The Sommerfeld enhanced annihilation cross-section of the dark matter explains the observed cosmic ray excesses. The light scalar field also gives rise to potentially large cross-sections of dark matter on the nucleon, therefore providing an interesting way to probe this model simultaneously at direct and indirect dark matter search experiments. We constrain the parameter space of the model by taking into account CDMS-II exclusion limit as well as PAMELA and FermiLAT data. ",
        "introduction": "The existence of dark matter (DM) in the present Universe has been firmly supported by a range of evidences~\\cite{dm_review}. The prime among them are galaxy rotation curves, large scale structure, cosmic microwave background and gravitational lensing. However, the identity of DM within the Standard Model (SM) of particle physics is still missing. Therefore, its experimental verification is expected to be a new discovery and a strong indication for physics beyond the SM.  Recently the Cryogenic DM Search (CDMS) Collaboration~\\cite{cdms} in the Soudan mine reported the observation of two events compatible with a positive dark matter detection at $1.64\\  \\sigma$ confidence level (C.L.). The DAMA/LIBRA~\\cite{dama} experiment claims an evidence of DM in its modulated signal at $8.8\\ \\sigma$ C.L. Several other direct detection experiments are running and setting upper bounds, including EDELWEISS-II~\\cite{edelweiss}, ZEPLIN-II~\\cite{zeplin}, XENON-10~\\cite{xenon2007} and Cresst-III~\\cite{cresst}. The Xenon collaboration will soon release the data from the first run of the Xenon100 experiment~\\cite{xenonup} and superCDMS is planned.  A significant amount of effort has also been devoted to detection of DM through indirect searches. For example, the PAMELA collaboration~\\cite{pamela} reported an unexpected rise of the positron fraction compared to that of the galactic background at energies above 10 GeV, while confirming the earlier results of AMS~\\cite{ams} and HEAT~\\cite{heat}. Similarly the HESS~\\cite{hess} and Fermi Large Area Telescope (FermiLAT)~\\cite{fermilat} collaborations also reported an excess of electron plus positron flux with respect to the galactic background at energies above 100 GeV, but without confirming the spectral features observed by the balloon-based experiments ATIC~\\cite{atic} and PPB-BETS~\\cite{ppbbets}. It has been widely interpreted that DM could be a viable candidate for the observed cosmic ray anomalies, although they could be explained by astrophysical sources~\\cite{astrophysics}.  In light of the above experimental results, several models have been considered in the literature, which leave signatures at either direct\\cite{direct_detection} or indirect \\cite{indirect_scattering,indirect_decay} DM searches. Typically, for a given model, the predictions for direct and indirect signatures of DM depend on different parts of the parameter-space, and the derived constraints thus do not overlap. However, in some models the same couplings are responsible for both the scattering of DM on the nucleon and large annihilation cross sections, in which case an interesting complementarity between direct and indirect searches exists~\\cite{cfs,cfs2}.  A simple possibility is to consider singlet extensions of the SM, in which the DM is a singlet scalar~\\cite{Singlet}, coupled to the SM Higgs particle through the so-called Higgs portal~\\cite{wilzeck}, namely the Higgs to DM quartic coupling (see also doublet extensions of the SM~\\cite{doublet_models}). Depending on the strength of this portal, the singlet can account for the observed relic density, $\\Omega_{\\rm DM} h^{2}\\sim 0.1$~\\cite{Komatsu:2010fb}. Furthermore, through the Higgs portal coupling, DM scatters with the nucleon and is thereby constrained by direct searches, as well as annihilates into SM fermions which can be observed at indirect detection experiments~\\cite{cfs,cfs2}.  The anomalous positron and electron fluxes observed by PAMELA and FermiLAT require a large enhancement of current DM annihilations. In their minimal versions~\\cite{Singlet}, the singlet extensions cannot reproduce such features. Furthermore, in these models, the stability of DM is supported by an {\\textit{ad hoc}} discrete symmetry. In this paper, we study a model which naturally solves these issues. We introduce a hidden sector gauged under an Abelian $U(1)_{H}~$\\cite{HiddenU1}-\\cite{HU1_2}, containing two complex scalars $\\Phi$ and $S$. While all SM fields are hidden sector singlet, the extra scalars are singlet under the SM but charged under $U(1)_{H}$. This model provides all the ingredients for a viable DM model with potentially large direct and indirect detection signals.  The paper is organized as follows. In Sec.~\\ref{sec:model} we present the model, with particular attention to the mass spectrum and mixings in the scalar sector. The model parameter space is then examined in Sec.~\\ref{sec:omegah2} and constrained by requiring the relic density of the DM candidate to be in the WMAP7 range. Sections~\\ref{sec:dd} and \\ref{sec:ind} describe the phenomenology of the model in light of the present DM searches: CDMS-II, XENON10, PAMELA and FemiLAT. First, in Sec.~\\ref{sec:dd}, we discuss the direct detection bounds constraining the hidden sector parameter space and the interplay with the indirect detection. Then, in Sec.~\\ref{sec:ind}, we investigate the indirect detection bounds, and give the results in terms of positron and electron excesses. The conclusions are presented in Sec.~\\ref{sec:conclusion}. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "\\label{sec:conclusion} In this paper we studied a hidden Abelian extension of the standard model. The DM is a complex scalar, singlet under the SM gauge group but charged under the hidden sector. We also introduce a light scalar $\\Phi$, whose non-zero vev $u$ breaks $U(1)_{H}$ to a $Z_2$ symmetry under which all fields but $S$ are even. As a result $S$ is a stable DM candidate, with mass ranging from the GeV to the TeV scale. The three fields $H$, $\\Phi$ and $S$ couple together via three Higgs portal couplings. The physical scalars $h$ and $\\phi$ mix together with the mixing angle $\\theta_{m}\\propto f_{H\\phi}$.  The relic density of $S$ mainly results from the annihilation channel $S^{\\dagger}S\\to h_{2} h_{2}$ through $f_{S \\phi}$. All three portal couplings enter in the spin-independent DM-nucleon cross section. While the usual Higgs-mediated channel depends on $f_{HS}$, the mixing between $h$ and $\\phi$ provides an additional channel, mediated by the light scalar $h_{2}$, which is $\\propto u\\times f_{S\\phi}\\times \\theta_{m}$. Given the mass scales we consider, the main contribution to the direct detection signal actually comes from this mixing term. For the parameter space we scanned over, solutions are found saturating or exceeding current experimental bounds from XENON10 and CDMS-II, or are in the reach of sensitivity of XENON100. The model also provides indirect signatures of $S$ through cosmic ray flux measurements. The main annihilation channel $S^{\\dagger}S\\to h_{2} h_{2}$, followed by fast $h_{2}\\to \\mu^{+}\\mu^{-}$ decays, ends up in high-energy electron and positron fluxes. The suppression of $S$ annihilations into standard Higgs compared to our dominant channel, together with the light mass of $h_{2}$, entail that no antiproton flux is expected at a significant level. A Sommerfeld enhancement of the current DM annihilation cross section occurs through the light $h_{2}$ exchange, which explains cosmic ray excess observations. This enhancement depends on $u\\times f_{S\\phi}$.  Interestingly, in this model, direct and indirect DM searches constrain same part of the parameter space. More precisely, in order to fully account for the anomalous positron fraction observed by PAMELA, large values of $u\\times f_{S\\phi}$ are required. Such large values give rise to large direct detection signals, saturating the current experimental exclusion limits. From this, an upper bound on the mixing angle is inferred. As an example, for $M_S=400\\GeV$ and $u\\times f_{S\\phi}\\sim 410 \\GeV$, $\\theta_{m}$ should be less than $1.9\\times 10^{-6}$ in order to satisfy simultaneously the current direct and indirect detection limits. We stress that in all models where the Sommerfeld enhancement occurs thanks to a light scalar that mixes with the Higgs particle, direct and indirect detection of dark matter are tightly connected."
    },
    "1004/1004.4799_arXiv.txt": {
        "abstract": " ",
        "introduction": "Planetary nebulae (PNe) evolve from  Asymptotic Giant Branch (AGB) stars after a short post-AGB transition and represent the last evolutionary stage of solar-type stars (M $\\leq$ 8\\,M$_{\\odot}$) before the white dwarf phase \\citep{BLO95}. The transformation of an AGB star into a PN involves a dramatic change in all stellar and circumstellar properties \\citep{BYF02}. In particular, while mass ejection in the AGB phase is spherical, the PN phase is dominated by axisymmetric shells, peculiar envelope geometries and collimated outflows \\citep{MAN96,MRG10}. Direct images allow us to study the morphology of PNe and to identify the structural components present in these objects. However, the study of the internal kinematics of PNe  is crucial to understand the nature of the components observed in the images, to determine their physical structure, and to investigate the ejection processes involved in their formation \\citep[e.g.,][]{MGT99}.  We have started an intermediate-resolution spectroscopic program to establish the true nature of sources classified as post-AGB star and PN candidates on the basis of their IRAS colours. Among other results \\citep[see][]{PYM07}, we have confirmed the PN nature of several objects and analyzed their morphology \\citep[hereafter MPG09]{PYM05,MPG09}.  In this paper we present high-resolution long-slit spectra of three of these confirmed PNe: PM\\,1-242, PM\\,1-318 and PM\\,1-322. The spectra allow us to analyze their internal kinematics and to investigate their physical structure. At least in two objects, the real structure derived from the long-slit spectra differs from that suggested by the morphology observed in direct images.  \\begin{figure}[h] \\begin{center} \\includegraphics[scale=1.00, angle=0]{Miranda2010fig1.eps} \\caption{({\\it top)} [N\\,{\\sc ii}]$\\lambda$6584 image of PM\\,1-242 (adapted from MPG09). The slit position used for high resolution spectroscopy is indicated. {\\it (bottom)} Grey-scale and contours position-velocity map of the [N\\,{\\sc ii}] emission line derived from the long-slit spectrum. The orientation (E-W) and the spatial and velocity scales are indicated.}\\label{fig1} \\end{center} \\end{figure} ",
        "conclusions": "In the following, we will present and discuss the results obtained for the three PNe. A brief description of their morphology and spectral properties is also provided from the previous results by \\cite{PYM05} and MPG09.    \\subsection{PM\\,1-242 (IRAS\\,18320+0005)}  PM\\,1-242 is a medium-to-high excitation PN with relatively strong He\\,{\\sc ii} emission and relatively weak [N\\,{\\sc ii}] and [S\\,{\\sc ii}] emissions (see MPG09). Figure\\,1 shows the [N\\,{\\sc ii}] image of PM\\,1-242 presented by MPG09. Morphologically, PM\\,1-242 can be classified as an elliptical PN with its major axis along the north-south direction and a size of $\\simeq$ 8$''$$\\times$5$''$. Two bright low-excitation, point-symmetric arcs trace the ellipse. Faint protrusions and/or knots are observed along PAs 50$^{\\circ}$, 225$^{\\circ}$ and  230$^{\\circ}$, particularly in H$\\alpha$ and [O\\,{\\sc iii}] (see MPG09). As mentioned by MPG09, the images do not allow us to determine whether PM\\,1-242 is an ellipsoid containing two point-symmetric arcs, with the major axis along the north-south direction, or a bipolar PN with a bright ring, faint bipolar extensions represented by the protrusions/knots and with the major axis along the east-west direction.  A position-velocity (PV) map of the [N\\,{\\sc ii}] line derived from the long-slit spectrum is also shown in Figure\\,1. The PV map of the H$\\alpha$ line (not shown here) is qualitatively similar to the [N\\,{\\sc ii}] one. The PV maps show two velocity components separated $\\simeq$ 36 and $\\simeq$ 42 km\\,s$^{-1}$ in H$\\alpha$ and [N\\,{\\sc ii}], respectively. The two components correspond to emission from the arcs at the minor axis of the ellipse (Figure\\,1). In particular, the blueshifted bright component corresponds to a bright knot at the eastern egde of the ellipse while the redshifed faint component corresponds to a faint knot at the western edge. From the center of the H$\\alpha$ and [N\\,{\\sc ii}] emission features in the PV maps we derive a heliocentric systemic velocity of $\\simeq$ $+$72.6 {\\kms}.  The presence of two velocity components at different velocities is not compatible with the slit being oriented along the minor axis of an ellipsoid. If this were the case, one would expect the two components to be located at the systemic velocity. On the contrary, the results strongly suggest that the observed ellipse is a tilted ring with its polar axis oriented along the east-west direction. The ring is tilted so that its eastern side points toward the observer and its western side recedes away. Assuming that the ring is circular, we derive an inclination angle of $\\simeq$ 30$^{\\circ}$ for its polar axis with respect to the observer. Taking into account this angle, the expansion velocity of the ring is $\\simeq$ 20 and $\\simeq$ 25 km\\,s$^{-1}$ in H$\\alpha$ and [N\\,{\\sc ii}], respectively.  These results show that the main structure of PM\\,1-242 is a tilted ring. The faint protrusions and knots (see MPG09) suggest that faint bipolar lobes may also be present and, hence, PM\\,1-242 could be a bipolar PN, although the presence of faint lobes needs to be confirmed by deeper observations. In any case, PM\\,1-242 is probably related to ring-like PNe like, e.g., IC\\,2149 \\citep{VAZ02}, Me\\,1-1 \\citep{PER08}, WeBo\\,1 \\citep{BPW03} and SuWt\\,2 \\citep{JON10}. In these PNe, a bright ring is the dominant structure in the nebula, which is, in some cases, accompanied by very faint bipolar lobes. Particularly interesting in PM\\,1-242 is the point-symmetric brightness distribution in the ring. Point-symmetry is usually observed in the lobes of bipolar PNe or in collimated structures usually associated to these objects, but to the best of our knowledge PM\\,1-242 is the only PN that shows this characteristic in the equatorial structure.   \\begin{figure}[h] \\begin{center} \\includegraphics[scale=0.78, angle=0]{Miranda2010fig2.eps} \\caption{{\\it (top)} [O\\,{\\sc iii}]$\\lambda$5007 image of PM\\,1-318 (adapted from MPG09). The slit position used for high resolution spectroscopy is indicated. {\\it (bottom)} Grey-scale position-velocity map of the [O\\,{\\sc iii}] emission line derived from the long-slit spectrum. The orientation (NW-SE) and the spatial and velocity scales are indicated.}\\label{fig2} \\end{center} \\end{figure}   \\begin{figure}[h] \\begin{center} \\includegraphics[scale=0.7, angle=0]{Miranda2010fig3.eps} \\caption{Line profiles of the H$\\alpha$, [N\\,{\\sc ii}] and [O\\,{\\sc iii}] emission lines in PM\\,1-322. An identification chart obtained from the POSS is also shown (upper left).}\\label{fig4} \\end{center} \\end{figure}    \\subsection{PM\\,1-318 (IRAS\\,20077+3722)}  PM\\,1-318 is a high-excitation PN with strong He\\,{\\sc ii} lines (e.g., He\\,{\\sc ii}$\\lambda$4686/H$\\beta$ $\\simeq$ 1) and extremely weak or absent low-excitation emission lines (MPG09). The [O\\,{\\sc iii}] image of PM\\,1-318 by MPG09, reproduced in Figure\\,2, shows that the nebula consists of an inner ring-like shell with two enhanced opposite regions, surrounded by a faint, round attached shell. The morphology observed in the images classifies PM\\,1-318 as a round PN.  A PV map of the [O\\,{\\sc iii}] emission line is also presented in Figure\\,2. The [O\\,{\\sc iii}] emission feature corresponds to emission from the inner shell because the spectrum is not deep enough to detect the faint outer shell. The [O\\,{\\sc iii}] emission feature appears as a tilted velocity ellipse on the PV map. The southeastern nebular regions are blueshifted while the northwestern nebular regions are redshifted. A velocity splitting of $\\simeq$ 60 km\\,s$^{-1}$ is measured at the center of the emission feature and a heliocentric systemic velocity of $\\simeq$ $-$99 {\\kms} is derived from that center. We note that radial velocities up to $\\simeq$ $\\pm$ 40 km\\,s$^{-1}$ from the systemic velocity are observed. The inner regions of the velocity ellipse show a decreasing [O\\,{\\sc iii}] intensity, which suggests that the [O\\,{\\sc iii}] emission mainly originates in a relatively thin layer of the inner shell, as also suggested by the direct image.  The tilt of the velocity ellipse on the PV map rules out that the inner shell of PM\\,1-318 is spherical (or round) because, in this case, the velocity ellipse would not be tilted on the PV map. This result strongly suggests that the inner shell is an ellipsoid with the major axis tilted with respect to the observer and oriented near PA 315$^{\\circ}$. Moreover, the largest velocities in the velocity ellipse are observed almost along the line of sight to the center of the nebula. This implies a relatively small inclination angle of the polar axis with respect to the observer and, hence, a polar expansion velocity of the order of $\\sim$ 30--40 km\\,s$^{-1}$. From these results we conclude that PM\\,1-318 is an (almost) pole-on elliptical PN. The enhanced regions in the inner shell probably represent a bright equatorial region in the ellipsoid.   \\subsection{PM\\,1-322 (IRAS\\,20124+5844)}  PM\\,1-322 was identified as a new PN by \\citet{PYM05}. It exhibits very strong [O\\,{\\sc iii}]$\\lambda$4363 line emission ([O\\,{\\sc iii}]$\\lambda$4363/H$\\beta$ $\\simeq$ 1.2) indicating a high electron density ($\\geq$ 10$^6$ cm$^{-3}$). The nebula is unresolved in our direct images. The compactness and high electron density suggest that PM\\,1-322 is a very young PN. On the other hand, these properties and the position of PM\\,1-322 in the [O\\,{\\sc  iii}]$\\lambda$4363/H$\\gamma$ versus [O\\,{\\sc iii}]$\\lambda$5007/H$\\beta$ diagram indicate that PM\\,1-322 could be a symbiotic star \\citep{PYM05}.  PM\\,1-322 is also spatially unresolved in the long-slit spectra. Therefore, we show in Figure\\,3 the H$\\alpha$, [N\\,{\\sc ii}] and [O\\,{\\sc iii}] emission line profiles derived from the spectra. The  [N\\,{\\sc ii}] and [O\\,{\\sc iii}] profiles presents a single peak at a heliocentric radial velocity of $\\simeq$ $+$27.2\\,{\\kms} that can be considered as the systemic velocity of the nebula. The FWHM (corrected of instrumental resolution) of the [N\\,{\\sc ii}] and [O\\,{\\sc iii}] profiles is $\\simeq$ 30 and $\\simeq$ 18 km\\,s$^{-1}$, respectively. The H$\\alpha$ emission line shows a double-peaked profile. The heliocentric radial velocity derived from the centroid of the H$\\alpha$ profile is $\\simeq$ $+$27.2\\,{\\kms}, in excelent agreement with the systemic velocity of PM\\,1-322 derived from the forbidden lines. The two emission peaks of the H$\\alpha$ profile are observed at $-$17.5 and $+$9\\,{\\kms} with respect to the systemic velocity and the red peak is stronger than the blue one. The emission peaks are separated by an apparent ``absorption reversal'' that is blueshifted by $-$9.6\\,{\\kms} with respect to the systemic velocity. Finally, the wings of the H$\\alpha$ line can be traced up to $\\simeq$ 325\\,km\\,s$^{-1}$ (FWZI) while the [N\\,{\\sc ii}] and [O\\,{\\sc iii}] lines present a much smaller FWZI of $\\simeq$ 90\\,km\\,s$^{-1}$.  The large differences between the forbidden lines and H$\\alpha$ profiles suggest a different origin for these emissions. Given the high density in the object, the regions probed via the H$\\alpha$ emission may critically depend upon the density and its variation within the object. Therefore, the H$\\alpha$ emission may be produced in regions that do not contribute to [N\\,{\\sc ii}] and [O\\,{\\sc iii}] emissions. In particular, the relatively large wings of the H$\\alpha$ emission could be due to Rayleigh-Raman scattering in a very dense region close to the central star \\citep[see][]{LEE00}. The spectral properties of PM\\,1-322 are remarkably similar to these found in other PNe suspected to host a symbiotic central star. These PNe are characterized by differences between the forbidden lines and H$\\alpha$ profiles, a double-peaked H$\\alpha$ profile, with very similar properties to these observed in PM\\,1-322, and extended H$\\alpha$ wings \\citep[e.g.,][]{BAL89,MTE96,GUE01,AYT03}. These results reinforce the idea that the central star of PM\\,1-322 is a symbiotic star.   \\subsection{Conclusions}  We have presented high-resolution long-slit spectra of three new confirmed PNe: PM\\,1-242, PM\\,1-318 and PM\\,1-322 with the aim of investigating their internal kinematics and physical structrure. The main conclusion of this work can be summarized as follows.  1) PM\\,1-242 is a ring-like PN, but not an elliptical PN as suggested by direct images.Interestingly, the ring displays a point-symmetric brightness distribution which is unusual among ring-like PNe and bipolar PNe with central rings.  2) PM\\,1-318 is an almost pole-on elliptical PN with an enhanced equatorial region, but not a round PN as suggested by direct images.  3) PM\\,1-322 is spatially unresolved and presents large differences between the forbidden lines and H$\\alpha$ profiles, a double-peaked H$\\alpha$ profile and relatively extended H$\\alpha$ wings. These characteriscs are also found in other PNe believed to host symbiotic central stars."
    },
    "1004/1004.3493_arXiv.txt": {
        "abstract": "% {% Weak gravitational lensing is a powerful probe of large-scale structure and cosmology. Most commonly, second-order correlations of observed galaxy ellipticities are expressed as a projection of the matter power spectrum, corresponding to the lowest-order approximation between the projected and the three-dimensional power spectrum. }% {% The dominant lensing-only contribution beyond the zero-order approximation is the reduced shear, which takes into account not only lensing-induced distortions but also the isotropic magnification of galaxy images. This involves an integral over the matter bispectrum. We provide a fast and general way to calculate this correction term. }% {% Using a model for the matter bispectrum, we fit elementary functions to the reduced-shear contribution and its derivatives with respect to cosmological parameters. The dependence on cosmology is encompassed in a Taylor-expansion around a fiducial model. }% {% Within a region in parameter space comprising the WMAP7 68\\% error ellipsoid, the total reduced-shear power spectrum (shear plus fitted reduced-shear correction) is accurate to 1\\% (2\\%) for $\\ell<10^4$ ($\\ell<2 \\times 10^5$). This corresponds to a factor of four reduction in the bias compared to the case where no correction is used. This precision is necessary to match the accuracy of current non-linear power spectrum predictions from numerical simulations. }% {% }% ",
        "introduction": "\\label{sec:intro}  Cosmic shear surveys exploit the distortions of distant galaxy images induced by the matter structures in the Universe on large scales. Weak cosmological lensing has become an important tool to measure cosmological parameters. The current precision in the measured shapes of high-redshift galaxies is at the few percent level providing important constraints on the matter density $\\Omegam$ and the power-spectrum normalisation $\\sigma_8$. Future surveys aim to achieve sub-percent level accuracy in measuring dark-energy properties or deviations from general relativity.  To reach this precision for cosmological parameters, efforts not only on the observational but also on the theoretical side have to be made. To predict the shear power spectrum from theoretical models, several approximations are usually made to facilitate calculations. When comparing and fitting models to observations, these simplifications can significantly bias the resulting cosmological parameters. One of these approximations is to neglect the correction due to the so-called reduced shear.  The observed ellipticity of galaxies is an estimator of the reduced shear \\begin{equation} g = \\frac{\\gamma}{1 - \\kappa}, \\label{g_red_1} \\end{equation} where $\\gamma$ is the shear and $\\kappa$ is the convergence. Here and in the following, both shear and reduced shear are represented as complex quantities, $\\gamma = \\gamma_1 + \\ii \\gamma_2$, $g = g_1 + \\ii g_2$. Since in weak lensing the convergence is small, $|\\kappa| \\ll 1$, one usually simplifies the above relation to linear order, $g \\approx \\gamma$.  The effect of reduced shear was examined using $N$-body simulations by \\citet{W05} and compared with theoretical predictions in \\cite{2006PhRvD..73b3009D}. Corrections to the lensing skewness were discussed in \\cite{1998MNRAS.296..873S} and \\cite{2005PhRvD..72h3001D}. An extensive study of higher-order corrections to the weak-lensing power spectrum, including reduced shear, was presented by \\cite{2009arXiv0910.3786K}. Compared with corrections due to light-path integration (relaxation of the Born approximation) and lens-lens coupling \\citep[e.g.,][]{1993LIACo..31..579S,1997A&A...322....1B}, reduced-shear corrections represent the dominant contribution to the weak-lensing power spectrum. The lowest-order correction term, which is of order ${\\cal O}(\\phi^3)$ in the gravitational potential $\\phi$, reaches 10\\% of the total lensing spectrum. This is the term that we consider in this work. The next-order correction at ${\\cal O}(\\phi^4)$ is one magnitude below the ${\\cal O}(\\phi^3)$-term \\citep{2009arXiv0910.3786K}. Ignoring reduced shear will cause a bias in cosmological parameters, which is comparable to the statistical error from future surveys like DES\\footnote{http://www.darkenergysurvey.org}, LSST\\footnote{http://www.lsst.org}, or Euclid\\footnote{http://www.ias.u-psud.fr/imEuclid} \\citep{2009ApJ...696..775S}. For the weak-lensing results of the COSMOS survey\\footnote{http://cosmos.astro.caltech.edu}, an estimate from numerical simulations yielded a 1\\% underestimation of $\\sigma_8$ \\citep{SHJKS09}.   Another correction to weak lensing from a magnitude-limited galaxy sample is the so-called \\emph{magnification bias}. \\cite{2001MNRAS.326..326H} found its contribution to the lensing power spectrum to be small. The corresponding third-order correction terms were also calculated and compared to numerical simulations by \\cite{2003AA...403..817M}. However, if not corrected for, future surveys will yield significantly biased cosmological parameter constraints \\citep{2009ApJ...702..593S}.  The lowest-order correction term has the same functional form as the reduced-shear one, and therefore both can be dealt with simultaneously. Higher-order terms contribute by smaller than one percent to the lensing power spectrum \\citep{2009arXiv0910.3786K}.  There are additional important sources of contamination to the weak lensing observables, some of which may be larger than the reduced-shear corrections. These involve galaxy clustering, the intrinsic alignment of galaxies with each other or with the surrounding dark matter haloes, and the modelling of the non-linear dark-matter power spectrum and baryonic processes. The modelling of these effects requires detailed understanding of baryonic physics, galaxy formation, the interplay between galaxies and dark matter, and massive numerical simulations in the case of the non-linear power spectrum. In contrast, taking into account the reduced shear is straightforward and can be done with sufficient accuracy to render it a sub-dominant contamination.  In the next section, we review the expressions for the lowest-order reduced-shear term, and introduce a fitting scheme to efficiently model this term. In Sect.~\\ref{sec:res}, we present our results, demonstrating the accuracy of the fits. We conclude this paper in Sect.~\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:disc}   The lowest-order correction term for the reduced-shear power spectrum is the dominant contribution from lensing-only effects \\citep{2009arXiv0910.3786K}. This term is proportional to the third power of the gravitational potential $\\phi$, and includes an integral over the lensing bispectrum. In this paper, we have presented fitting formulae for this integral and its derivative with respect to cosmological parameters. This has permitted a more efficient calculation of the reduced-shear correction power spectrum within a parameter range comprising the 68\\% confidence region of WMAP7. The fitting scheme reduces the computational effort from tens of minutes for the full integration to a fraction of a second.  For currently available weak-lensing data, the uncertainty in the non-linear power spectrum on small scales is still larger than the bias due to the reduced-shear. For example, the widely-used fitting prescription by \\citet{2003MNRAS.341.1311S} strongly underestimates the power on small scales. More recent numerical simulations however provide fitting formulae that improve the accuracy by a factor 5 to 10 \\citep{CoyoteI, CoyoteII, CoyoteIII}. Moreover, their emulation scheme can be applied to a wide range of cosmological parameters. This range is comprised in the hypercube for which our fits are valid. In combination, these results provide us with predictions for the weak lensing power spectrum that are accurate to the few percent level."
    },
    "1004/1004.3988_arXiv.txt": {
        "abstract": "Gravitational wave sources are a promising cosmological standard candle because their intrinsic luminosities are determined by fundamental physics (and are insensitive to dust extinction).  They are, however, affected by weak lensing magnification due to the gravitational lensing from structures along the line of sight. This lensing is a source of uncertainty in the distance determination, even in the limit of perfect standard candle measurements.  It is commonly believed that the uncertainty in the distance to an ensemble of gravitational wave sources is limited by the standard deviation of the lensing magnification distribution divided by the square root of the number of sources.  Here we show that by exploiting the non-Gaussian nature of the lensing magnification distribution, we can improve this distance determination, typically by a factor of 2--3; we provide a fitting formula for the effective distance accuracy as a function of redshift for sources where the lensing noise dominates. ",
        "introduction": "In 1998 two groups measuring the luminosity distance-redshift relation (the Hubble diagram) from Type Ia supernovae (SNe) reported that the expansion of the Universe was accelerating \\cite{1998AJ....116.1009R, 1999ApJ...517..565P}.  This discovery has stimulated a range of efforts to measure the cosmic expansion history and assess whether it is consistent with a cosmological constant or if alternatives such as quintessence are required.  The Type Ia SNe continue to provide one of the most valuable constraints \\cite{2007ApJ...666..694W}, due to quality data at a range of redshifts, and the lack of cosmic variance limitations that plague alternatives such as weak gravitational lensing (WL) and baryon-acoustic oscillations (BAO) at low redshift.  The advent of gravitational wave astronomy has prompted interest in gravitational wave (GW) sources as a standard candle.  Schutz \\cite{1986Natur.323..310S} showed that a gravitational waveform from merging compact objects can be used to measure the distance to the source; a redshift obtained of an electromagnetic counterpart or host galaxy would then allow one to place the GW source on a luminosity distance-redshift relation.  The GW source method has the key advantage over other standard candles that its luminosity can be determined from fundamental physics, thus alleviating the common concern with standard candles that they could evolve with cosmic time.  They are also insensitive to dust opacity.  Finally, many proposed space-based gravitational wave detectors measure test mass separations directly in units of the laser wavelength, as opposed to supernovae, which measure relative distances and require an independent calibration ladder in order to measure absolute distances.  Thus GW sources offer the potential for a low-systematics probe of the expansion history of the Universe.  Examples of such sources would include mergers of massive black holes, observable with the Laser Interferometer Space Antenna (LISA \\footnote{{\\tt http://lisa.jpl.nasa.gov/}}), and binary neutron star or stellar mass black hole binaries, observable with a more futuristic space-based detector in the $\\sim 1$ Hz band such as the Big Bang Observer (BBO).  The latter, in particular, could potentially observe hundreds of thousands of neutron star-neutron star binary inspirals \\cite{2009PhRvD..80j4009C}.  A final advantage is that the distance determination to a GW source is limited by the signal-to-noise of the measurement (and by partial degeneracies with binary parameters), as opposed to Type Ia SNe, which have a seemingly random $\\sim 15$\\% scatter in their luminosities even after the light curve stretch correction.  This means that the statistical power of GW sources may be interesting even if the number of usable sources is far less than the number of usable SNe.  In fact, for high signal-to-noise detections of GW sources, the distance determination should be limited not by the intrinsic dispersion of the source nor by the measurement error, but rather by weak lensing magnification: the intervening matter between us and the source will magnify the GW source and affect our measurement of the distance.  The apparent flux from the source is increased by some factor $\\mu$, which is usually $\\sim 1$, and the apparent distance $D_{\\rm app}$ thus differs from the true distance $D$ according to $D_{\\rm app}=D\\mu^{-1/2}$.  This phenomenon of course occurs for all standard candles, and has long been recognized as an issue for SNe \\cite{1997ApJ...475L..81W, 1998PhRvD..58f3501H, 2005ApJ...631..678H}, but its importance relative to intrinsic dispersion is much greater for gravitational waves.  (Gravitational wave measurements with nearby sources or with lower signal-to-noise per source, such as the proposed binary progenitors of short gamma ray bursts \\cite{2006PhRvD..74f3006D, 2009arXiv0904.1017N}, are much less affected since the lensing scatter is subdominant.)  The usual way of accounting for the magnification effect is to suppose that it adds in quadrature with the intrinsic-luminosity and apparent flux measurement contributions to the distance error.  That is, \\begin{equation} \\sigma_{\\ln D^2} = \\sqrt{\\sigma_{\\ln L}^2 + \\sigma_{\\ln F}^2 + \\sigma_{\\mu}^2}, \\label{eq:basic} \\end{equation} where $D$ is the distance, $L$ is the intrinsic luminosity, $F$ is the measured flux, and $\\mu$ is the magnification; with $N$ sources, this uncertainty is reduced by a factor of $\\sqrt N$ \\cite{2005ApJ...629...15H, 2005ApJ...631..678H}.  Since the last term dominates for GW sources and is significant for high-$z$ SNe, there is great motivation to try to reduce it.  One possibility would be to try to construct an estimated magnification $\\hat\\mu$ from external data sets; the last term should then be replaced by $\\sigma^2(\\mu-\\hat\\mu)$.  Unfortunately, WL shear maps have too much high spatial frequency noise to be useful as a magnification estimator for point sources \\cite{2003ApJ...585L..11D}, but galaxy maps are highly correlated with the mass distribution and may be able to reduce the lensing dispersion term by a factor of $\\sim\\sqrt3$ \\cite{2006ApJ...640..417G, 2007ApJ...658...52J, 2009A&A...493..331J}.  Maps of flexion (i.e. the gradient of the shear measured using the banana- or trefoil-shaped distortion of a galaxy \\cite{2005ApJ...619..741G}) could also be useful if very high source densities ($>100\\,$arcmin$^{-2}$) can be obtained \\cite{2010MNRAS.tmp..245S}.  The purpose of this paper is to explore yet another method of reducing the lensing dispersion in Eq.~(\\ref{eq:basic}).  Because the probability density function (PDF) of $\\mu$ is highly non-Gaussian (technically non-$\\Gamma$, as we show in Appendix~\\ref{app:ineq}), our ability to centroid it using $N$ sources can be much better than $\\sigma_{\\mu}/\\sqrt N$.  This is in fact a familiar result: as an extreme case, many distributions such as the Airy diffraction pattern in optical astronomy can be centroided even though their variances are formally infinite. In the case of lensing magnification, the variance of $\\mu$ is often dominated by a long tail to positive values, corresponding to lines of sight that pass through a galaxy or its halo, whereas the information content is dominated by relatively empty lines of sight with $\\mu-1$ sharply peaked around a slightly negative value. In such cases, the use of outlier-rejecting statistics not only removes sources with misidentified redshifts \\cite{2009PhRvD..80j4009C} but also reduces the uncertainty in $D(z)$ for correctly identified sources.  (A similar point has been made in the recent paper by Shang \\& Haiman \\cite{2010arXiv1004.3562S}.)  This paper is organized as follows.  In Section~\\ref{sec:centroid}, we discuss the information-theoretic limits on centroiding a distribution. Numerical results and simulations are presented in Section~\\ref{sec:numerics}. Section~\\ref{sec:constraints} gives cosmological constraints obtainable with the reduced centroid errors for BBO and for LISA. We conclude and briefly discuss systematics in Section~\\ref{sec:discussion}.  We focus here on the problem of measuring $D(z)$ from GW sources.  However, the formalism is applicable to any standard candle, and we briefly discuss the implications for Type Ia supernovae. ",
        "conclusions": "\\label{sec:discussion}  The luminosity distance-redshift relation is one of the oldest and most fundamental probes of cosmology, and future gravitational wave detectors offer the possibility of measuring it accurately using binary inspirals whose luminosity can be calculated directly from measured parameters and fundamental physics.  These sources are however affected by weak gravitational lensing by intervening inhomogeneities in the cosmic mass distribution.  This introduces changes of typically a few percent (but occasionally much larger) in the flux, while not significantly affecting the redshift, and thus provides a source of noise in the $D(z)$ relation.  We have shown in this paper that exploiting the full power of the likelihood function can reduce this noise: the noise in the $D(z)$ relation is not limited by the lensing dispersion divided by the square root of the number of sources, but rather can be less by a factor of 2--3 if one centroids the distribution of apparent distances using the known non-Gaussian form of the lensing magnification PDF.  We have not discussed here the systematic errors associated with using large numbers of gravitational wave sources for precision low-redshift cosmology, as suggested for BBO.  While the signal itself is expected to be clean, there are potential sources of systematic error.  Some of these, such as strain calibration, coherent peculiar velocities at low redshift \\cite{2006PhRvD..73j3002C, 2006PhRvD..73l3526H}, and host redshift misidentification, exist irrespective of the method used to estimate the true $D(z)$ from a collection of weakly lensed GW sources with their apparent fluxes and redshifts.  However, the issue of uncertainties in the magnification PDF is worth discussing here.  It may seem at first glance that the flux-averaging technique is more robust than the centroiding technique described here, because it relies only on the flux conservation constraint $\\int_0^\\infty P(\\mu)d\\mu=1$.  This may not be the case for three reasons.  One is that in order to remove misidentified host galaxies, it is likely that a BBO-type mission would reject outliers from the $D(z)$ relation \\cite{2009PhRvD..80j4009C}.  This outlier rejection would eliminate the the tails of the magnification distribution with $|\\ln\\mu|>(0.4\\ln10)\\Delta M$, where $\\Delta M$ is the half-width of the cut in magnitudes.  Since $P(\\mu)$ is highly asymmetric, with the large-$\\mu$ tail much stronger than the small-$\\mu$ tail, it follows that outlier rejection will result in the conditional probability $\\int_0^\\infty P(\\mu|{\\rm accept})d\\mu<1$ and hence give a positive bias in $D(\\mu)$.  This can be corrected, but it requires knowledge of the contribution to $\\int_0^\\infty P(\\mu)d\\mu$ from the strong-magnification tail.  A second reason has to do with strong lensing: flux conservation implies that the magnification satisfying $\\int_0^\\infty P(\\mu)d\\mu=1$ is the {\\em total} magnification of all of the images.  However, since strong-lens time delays are often measured in months (and even longer if the strongly de-magnified central image is significant), it is likely that for many BBO sources there will be additional images whose time delay places them outside the BBO mission lifetime.  Finally, in obtaining a successful host redshift (or identifying a source with the correct host rather than another object nearby on the sky), there is likely to be a bias in favor of brighter sources, which results in a success probability that has some dependence on the magnification, $P({\\rm success})\\sim\\mu^\\beta$.  The presence of lensing dispersion then results in a bias in the mean flux of $\\sim\\beta{\\rm Var}\\,\\mu$; since Var$\\,\\mu$ is of order $10^{-2}$ at $z\\sim 1$, such biases will likely be far above the BBO statistical errors, and will have to be corrected using knowledge of $P(\\mu)$.  Therefore, even the flux-averaging method is sensitive to the particular distribution $P(\\mu)$.  The problem of how well $P(\\mu)$ can be determined via theory (particularly in the presence of baryonic effects), or reduced to a parametric model whose parameters can be simultaneously fit using BBO, is left to future work.  In this paper, we have not attempted to reduce the lensing dispersion by using external information, i.e. only the shape of $P(\\mu)$ was assumed.  The main external sources of information that are commonly considered are weak lensing shear maps and galaxy maps.  In principle the way one would incorporate this information would be to write a conditional probability density, e.g. $P(\\mu|g)$, where $g$ represents the galaxy density field.  Galaxy maps have been shown to be helpful because they contain information about the small scales that dominate the lensing variance \\cite{2006ApJ...640..417G, 2007ApJ...658...52J, 2009A&A...493..331J}; however the conditional probability distribution $P(\\mu|g)$ may be very hard to determine theoretically at the required precision.  For example, despite recent advances in determining the relation between galaxy luminosity and host halo mass (a key quantity of interest if one is trying to infer the matter density field from galaxy observations) using clustering and lensing data \\cite{2005ApJ...630....1Z, 2006MNRAS.368..715M}, a measurement of the scatter in this relation is not yet possible, and the full distribution of this scatter -- required if one is going to compute $P(\\mu|g)$ -- is woefully underconstrained.  Nevertheless, for a mission such as LISA that may have only a limited number of usable sources and hence may be dominated by statistical errors due to weak lensing, this may be a useful approach.  The weak lensing field is sensitive only to the matter distribution, and so one could imagine that it would be profitable to utilize the smoothed convergence field, $\\kappa_{\\rm sm}$, and attempt to centroid the conditional density $P(\\mu|\\hat\\kappa_{\\rm sm})$. WL has traditionally been viewed as not useful for de-lensing of GW sources because most of the lensing variance comes from small scales where weak lensing measurements are noisy \\cite{2003ApJ...585L..11D}.  This conclusion should be revisited in future work using the centroiding technique; in particular, this small-scale structure contributes strongly to the high-magnification tail of $P(\\mu|\\hat\\kappa_{\\rm sm})$, and it is not yet known what happens to the Fisher information, which depends largely on the width of the peak of $P(\\mu|\\hat\\kappa_{\\rm sm})$."
    },
    "1004/1004.4602_arXiv.txt": {
        "abstract": "A new harassment model is presented that models the complex, and dynamical tidal field of a Virgo like galaxy cluster. The model is applied to small, late-type dwarf disc galaxies (of substantially lower mass than in previous harassment simulations) as they infall into the cluster from the outskirts. These dwarf galaxies are only mildly affected by high speed tidal encounters with little or no observable consequences; typical stellar losses are $<10\\%$, producing very low surface brightness streams ($\\mu_B > 31$ mag arcsec$^{-2}$), and a factor of two drop in dynamical mass-to-light ratio. Final stellar discs remain disc-like, and dominated by rotation although often with tidally induced spiral structure. By means of Monte-Carlo simulations, the statistically likely influences of harassment on infalling dwarf galaxies are determined. The effects of harassment are found to be highly dependent on the orbit of the galaxy within the cluster, such that newly accreted dwarf galaxies typically suffer only mild harassment. Strong tidal encounters, that can morphologically transform discs into spheroidals, are rare occurring in $<15 \\%$ of dwarf galaxy infalls for typical orbits of sub-structure within $\\Lambda$CDM cluster mass halos. For orbits with small apocentric distances ($<$250 kpc), harassment is significantly stronger resulting in complete disruption or heavy mass loss ($>90 \\%$ dark matter and $> 50 \\%$ stellar), however, such orbits are expected to be highly improbable for newly infalling galaxies due to the deep potential well of the cluster. ",
        "introduction": "Dwarf galaxies are easily the most common type of galaxy in the Universe, and yet their very origin remains a subject of much debate. They can be found in vast numbers in the cluster environment - our nearest large cluster, the Virgo cluster, contains over a thousand dwarfs. The understanding of the formation and evolution of cluster dwarfs is hampered by their intrinsic properties of low luminosity and low surface brightness causing them to be challenging objects to observe, and study. Intuitively, it might be expected that their low mass, and hence shallow potential wells, cause them to be extremely sensitive to their local environment. As such, cluster dwarfs aught to provide sensitive indicators to the effects of mechanisms occurring in the high density environment.  The Virgo cluster catalogue (\\citealp{VCC}), continues to provide an invaluable tool for the study of dwarf galaxies in the Virgo cluster. Complete down to apparent magnitude $B_T>$18, and with a suprisingly accurate morphological classification that has, on the whole, stood the test of time, the Virgo cluster catalogue (VCC) clearly demonstrates that dwarf galaxies are far from a homogenous family of galaxies. Dwarf ellipticals are the most common type of cluster dwarf - typically containing no current star formation, and an old stellar population. Dwarf ellipticals can be additionally split into further categories, with some containing a compact, central stellar nucleus - the nucleated dwarf ellipticals, and dwarf spheroidals - generally considered to be low luminosity analogues of dwarf ellipticals. Not all the dwarfs are `red and dead' with many containing active star formation, with knotty HII regions clearly visible in the Las Campanas observatory photographic plates utilised to create the VCC. These star forming dwarfs are further split into sub-categories; dwarf irregulars containing atomic gas, and with evidence of rotation (\\citealp{Zee2004}), and blue compact dwarfs whose gas and star formation is centrally concentrated and also show significant rotation (\\citealp{Zee2001b}). A further classification is made for `transition-types', objects whose properties are somewhere mid-way between dwarf ellipticals and dwarf irregulars, and may form a possible evolutionary link between the two classes. The large variety in dwarf galaxy properties makes a single theory for their creation highly challenging.  With the release of SDSS data set, covering the Virgo region, further progress was made in the observation of cluster dwarfs. A series of papers (\\citealp{Lisker2006},\\cite{Lisker2006b},\\cite{Lisker2007}) have provided a quantitive study of over 400 Virgo dwarf galaxies, suggesting that the VCC dwarf classifications can be further sub-categorised. The improved imaging quality, and additional colour information have shown that many dwarf ellipticals thought to be typically `red and dead', infact contain blue centres indicating recent or current central star-formation. Furthermore, dwarf ellipticals that appear smooth and featureless, present intricate spiral patterns and disks when an un-sharp mask technique is applied. \\cite{Lisker2007} stress the importance of recognising these sub-categories and considering them separately when attempting to un-ravel dwarf galaxy origins, and conclude that multiple environmental mechanisms are required to explain the diverse properties of cluster dwarfs.  Current suggested dwarf galaxy formation and evolution scenarios are numerous and varied. The $\\Lambda$CDM paradigm suggests that current day dwarf galaxies are the left over debris from hierarchical merging of haloes. While $\\Lambda$CDM simulations predicts the presence of an extended dark matter halo about each dwarf galaxy, they make less detailed predictions about properties of the baryonic content of a galaxy. It is necessary to include physical mechanisms that influence the visible components of dwarf galaxies in order to bring the simulations into agreement with observations (\\citealp{Bullock2000}). These physical mechanisms can be broadly split into two categories; global mechanisms that affect all dwarf galaxies regardless of their location, such as supernovae feedback (\\citealp{Dekel1986}), and environmentally dependent mechanisms. One such environmental mechanism, known as harassment, is believed to be of importance in high density environments (\\citealp*{Moore1998}, \\citealp{Mastropietro2005}). Numerous high speed tidal encounters within a cluster act to morphologically transform medium mass, low surface brightness disc galaxies into smaller galaxies - cluster dwarfs. A second scenario, ram pressure stripping, acts on in-falling galaxies that are already dwarfs - late-type dwarfs - and that this can produce galaxies that resemble cluster dwarf ellipticals (\\citealp{Zee2004}, \\citealp{Boselli2008}). These two scenarios have one thing in common - a newly infalling disc galaxy is subjected to cluster environmental mechanisms causing transformation into an object resembling a cluster dwarf elliptical.  Indeed, radial velocity measurements of cluster dwarf ellipticals in the Virgo cluster do indicate that they are far from a virialised population within the cluster potential, bearing signatures of a recent in-fall (\\citealp{Conselice2001}). Meanwhile dwarf galaxies obey their own morphology-density relationship; early -type dwarfs have been said to be the most strongly clustered of all galaxies, whereas gas-rich dwarfs are the most weakly clustered (\\cite{Ferguson1994}). As cluster-centric distance increases, the ratio of dwarf ellipticals to dwarf irregulars decreases (\\citealp{Sabatini2005}). These two observations combined may suggest that many cluster dwarf ellipticals have their origin in infalling late-type dwarfs.  Further evidence for galaxy transformation of discs is present in the form of hidden disc features, in some cases even including fine spiral arm structure (\\citealp{Lisker2006b}, \\citealp{Lisker2007}) in Virgo cluster dwarf ellipticals. It is as if the original discy nature of the galaxies has yet to be completely removed by the cluster environmental mechanisms. This is additionally supported by evidence of significant rotation in a sub-sample of the dwarf ellipticals in \\cite{Zee2004b}. Remarkably the rotating dwarf ellipticals, and dwarf irregulars, both appear to obey the Tully-Fisher relation for bright late-type spirals. If we assume that a substantial fraction of present-day dwarf ellipticals are indeed accreted as dwarf discs into the cluster, the next question is whether it is harassment or ram pressure stripping that is the key environmental influence driving the tranformation. As ram pressure stripping has been shown to remove significant amounts of atomic hydrogen from {\\it giant} spiral galaxies within the Virgo cluster (\\citealp{Chung2008}, \\citealp{Vollmer2002}), it should be of little surprise that ram pressure can strip small late-type dwarfs of their HI gas content, causing a cessation of star-formation, and evolving the dwarfs from blue to red colours, typical of cluster dwarf ellipticals (\\citealp{Boselli2008}). However, some form of physical morphological transformation may additionally be required to transform dwarfs from thin discs to spheroidals. One suggestion is that both mechanisms are required - ram pressure halts the star formation, while harassment thickens up the disc.  In the literature (\\citealp{Moore1998}), harassment is the effects of repeated and numerous long range tidal encounters in the potential well of the cluster. The tidal forces experienced by a galaxy within a cluster can be broadly split into two categories - tidal heating, and tidal shocking. The key difference between these categories is the time-scale on which the galaxy experiences the tidal force. For tidal shocking the time-scale is short - for example in a high speed tidal encounter between two cluster galaxies, a strong but short-lived tidal force may be experienced by the two galaxies. Tidal heating occurs when the time-scale is far longer - for example a cluster galaxy that passes close to the cluster centre spends a significant length of time under the influence of the deep potential well of the cluster, and may be entirely dismantled if the resulting tidal forces are in excess of the galaxy's self-gravity. Harassment is the combined effects of both tidal heating (from the tidal forces associated with the deep potential well of the cluster) and tidal shocking (from the tidal forces associated with the high speed galaxy-galaxy encounters) on a cluster galaxy.  The aim of this work is to study the effects of harassment on newly infalling, late-type dwarf irregulars. This is accomplished using a new, and fast algorithm to model a galaxy clusters tidal fields that produces similar results to far more complex models set in a cosmological context. In concordance with cosmological models, both dwarf galaxy models and the cluster tidal field utilise cuspy cold dark matter haloes.The dwarf galaxy models are significantly lower mass ($\\sim$10$^9$ - $ 1 \\times 10^{10}$ M$_\\odot$) than previous studies of harassment on dwarf galaxies. \\cite{Mastropietro2005} dwarf models were of mass $7 \\times 10^{10}$ M$_\\odot$ in comparison. Studying the effect of harassment on dwarfs in a lower mass regime is interesting for a number of reasons - obviously it extends harassment studies to cover a wider range of the parameter space. Additionally while larger dwarfs may be harassed into smaller objects resembling dwarf ellipticals (\\citealp{Mastropietro2005}), the fate of smaller mass dwarfs is less clear. A weakly harassed low mass dwarf may be able to complete a morphological transformation into an object resembling a dwarf elliptical without significant mass-loss. Alternatively, a heavily harassed low mass dwarf may be entirely destroyed in the process.  For these low mass dwarfs, the significance of tidal heating from the potential well of the cluster alone can be estimated using a Roche limit analysis. When the tidal forces of the cluster potential well exceeds the dwarf galaxy's own self-gravity, the dwarf will be entirely dismantled. For our standard dwarf model, and for the cluster potential utilised (see Section \\ref{harassmodel}), this can be expected to occur at a radius of less than $\\sim 50$ kpc within the cluster. Hence only dwarfs on extremely plunging orbits are significantly influenced by tidal heating. Hence for the orbits we consider (see Section \\ref{simeffects}) whose peri-cluster distance is $\\sim 200$ kpc, the effects of tidal heating alone are expected to be weak.  We can also analytically estimate the effects of tidal shocking from galaxy-galaxy encounters for our low-mass dwarf models. Although it seems intuitive, that the lowest mass galaxies might be expected to be the most significantly affected by tidal encounters, this is not necessarily the case. The strength of the effects of a high-speed encounter can be expected to depend on the dynamical time at the scale-radius, $r_s$ of the dwarf ($t_{dyn} \\sim \\left[(r_s^3/{G M(<r_s)}\\right]^{\\frac{1}{2}}$, where $G$ is the gravitational constant. When this dynamical time is long compared to the passage-time of a passing harasser galaxy, then the the dwarf can respond effectively to the encounter. It is interesting to note that in the dwarf models considered here, for material surrounding the disk of the galaxy, $t_{dyn}$ is of the same order as the passing time of a harasser galaxy. Despite the lowered mass, they are significantly more compact, and hence resilient to harassment. These properties are not peculiar to our models - in this mass regime, this is to be expected.  Although tidal stripping and shocking are expected to be weak for the dwarfs we consider, strong tidal shocking between galaxies in combination with tidal heating can still produce significant effects. These effects are studied in detail on model dwarf galaxies in Section \\ref{simeffects}. We attempt to quantify the likelihood of significant harassment in section \\ref{Monty}, and test how the likelihood varies for a variety of orbits in Section \\ref{orbitsection}.  \\newline \\noindent The paper is organised as follows. In section 2, the numerical code utilised in all simulations is discussed, and in Section 3 the methods used to build model dwarf galaxies are shown. The harassment model is presented in section 4. In section 5, the influence of harassment on 11 in-falling dwarf galaxy models are shown. The statistical influences of harassment are found by means of a Monte-Carlo simulation in section 6, and tested for a variety of orbits.    \\section[]{The code} In this study we make use of `gf' (\\citealp{Williams2001},\\citealp{Williams1998}), which is a Treecode-SPH algorithm that operates primarily using the techniques described in \\cite{Hernquist1989}. `gf' has been parallelised to operate simultaneously on multiple processors to decrease simulation run-times. While the Treecode allows for rapid calculation of gravitational accelerations, the SPH code allows us to include a HI gas component to our dwarf galaxy models. In all simulations, the gravitational softening length, $\\epsilon$, is fixed for all particles at a value of 100 pc, in common with the harassment simulations of \\cite{Mastropietro2005}. Gravitational accelerations are evaluated to quadropole order, using an opening angle $\\theta_c=0.7$. A second order individual particle timestep scheme was utilised to improve efficiency following the methodology of \\cite{Hernquist1989}. Each particle was assigned a time-step that is a power of two division of the simulation block timestep, with a minimum timestep of $\\sim$0.5 yrs. Assignment of time-steps for collisionless particles is controlled by the criteria of \\cite{Katz1991}, whereas SPH particle timesteps are assigned using the minimum of the gravitational time-step and the SPH Courant conditions with a Courant constant, $C$=0.1 (\\citealp{Hernquist1989}). As discussed in \\cite{Williams2004}, the kernel radius $h$ of each SPH paricle was allowed to vary such that at all times it maintains between 30 and 40 neighbours within 2$h$. In order to realistically simulate shocks within the SPH model, the artificial viscosity prescription of \\cite{Gingold1983} is used with viscosity parameters $(\\alpha,\\beta)$ = (1,2). The equation of state for the gas component of the galaxies is isothermal with a bulk velocity dispersion of 7.5 km s$^{-1}$, based on the measured velocity dispersion of molecular clouds in the local interstellar medium (\\citealp{Stark1989}). By choosing an isothermal equation of state, we are intrisically assuming that stellar feedback processes are balanced by radiative cooling producing a constant velocity dispersion. `gf' contains a simple star formation prescription, where each SPH particle converts it's mass into stellar mass at a rate controlled by a Schmidt law with a power of 1.5. When the global stellar mass formed in this manner reaches a specified mass, a specified number of new N-body star particles are formed. The new particles have their position chosen statistically such that recent star formation history of individual SPH particles is taken into account. For details of code testing, please refer to \\cite{Williams1998}. ",
        "conclusions": "Harassment of small late-type dwarf galaxies {\\it that have recently been accreted into the cluster} is not an effective mechanism for forming the thickened discs of dwarf ellipticals. Despite their low mass, they are additionally very concentrated in mass, and hence are robust against the effects of harassment. In previous studies (\\citealp{Moore1998}), harassment is the effect of repeated long-range fast encounters. For dwarfs in the low-mass but more concentrated regime studied here, harassment is only signicant in repeated, or even individual, close-range high speed encounters where tidal forces are the strongest, and these are rare for objects infalling into the cluster at the current epoch. Currently it would be a challenge to observe the typical effects of harassment on dwarfs that have recently infallen into the cluster (very low surface brightness extragalactic features, mild disc thickening, and a $\\sim 50 \\%$ reduction in dynamical mass-to-light ratios). However harassment appears far more effective on dwarf galaxies that have formed in situ within the cluster as it formed, possibly forming nucleated cluster dwarfs. Alternatively harassment of larger, low surface brightness spirals into dwarf ellipticals in the manner suggested by \\cite{Moore1998} is another possibility. However, the simulations of \\cite{Moore1999} demonstrate the harassment is not effective on all disc galaxies - while a giant low surface brightness disc is effectively tranformed, a high surface brightness giant like our own galaxy is only mildly affected. Therefore if cluster dwarf ellipticals have their origins in larger disc galaxies, a very large population of specifically low surface brightness field disc galaxies is required, in order to form the vast numbers of dwarf ellipticals seen in clusters today. The formation of dwarf ellipticals from dwarf irregulars via ram pressure stripping is another convincing possibility as discussed in \\cite{Boselli2008}. Ram-pressure stripping need only terminate their star formation to produce objects that appear very much like dwarf ellipticals today (\\citealp{Zee2004})."
    },
    "1004/1004.3487_arXiv.txt": {
        "abstract": "{ Recent sub-millimetric observations at the Plateau de Bure interferometer evidenced a cavity at $\\sim 46$ AU in radius into the proto-planetary disk around the T Tauri star LkCa15 (V1079 Tau), located in the Taurus molecular cloud. Additional Spitzer observations have corroborated this result possibly explained by the presence of a massive ($\\ge 5 M_{\\rm{Jup}}$) planetary mass, a brown dwarf or a low mass star companion at about 30 AU from the star.} { We used the most recent developments of high angular resolution and high contrast imaging to search directly for the existence of this putative companion, and to bring new constraints on its physical and orbital properties.} { The NACO adaptive optics instrument at VLT was used to observe LkCa15 using a four quadrant phase mask coronagraph to access small angular separations at relatively high contrast. A reference star at the same parallactic angle was carefully observed to optimize the quasi-static speckles subtraction (limiting our sensitivity at less than 1.0~$\\!''$).} { Although we do not report any positive detection of a faint companion that would be responsible for the observed gap in LkCa15's disk (25-30 AU), our detection limits start constraining its probable mass, semi-major axis and eccentricity. Using evolutionary model predictions, Monte Carlo simulations exclude the presence of low eccentric companions with masses $M \\ge 6~M_{Jup}$ and orbiting at $a\\ge100$~AU with significant level of confidence. For closer orbits, brown dwarf companions can be rejected with a detection probability of 90\\% down to 80~AU (at 80\\% down to 60~AU). Our detection limits do not access the star environment close enough to fully exclude the presence of a brown dwarf or a massive planet within the disk inner activity (i.e at less than 30~AU). Only, further and higher contrast observations should unveil the existence of this putative companion inside the LkCa15 disk.} {}  \\keywords {Stars: pre-main sequence - Planetary systems: proto planetary disks - Instrumentation: adaptive optics, coronagraphy} ",
        "introduction": "\\label{sec:intro}  Since more than two decades, the first glimpse of planetary formation has been directly observed (Smith \\& Terrile 1984). Then, following the discovery of the first extrasolar planet around a solar-type star \\citep[51 Peg, see][]{1995Natur.378..355M} , more than 400 planets have been found, mainly using the radial velocity technique (RV). However, the explored time span limits the radial velocity study (as well as pulsar timing, micro-lensing, photometric transit, astrometric techniques) to the close ($\\le 4-5$ AU) circumstellar environment. To understand the way exo-planetary systems form and evolve at wider orbits, it is then particularly interesting to use a complementary technique, such as direct imaging.  Around young solar analogues the current deep imaging performances enable the detection of companions down to the brown dwarf (BD) and planetary mass regimes.  Typical separations larger than 50-100 mas (i.e. $\\ge 5-10$ AU for a star at 100 pc) can be explored for BD companions and for planetary mass objects, typical separations in the range 200-400 mas (ie $\\ge 20-40$ AU for a star at 100 pc).  The discoveries of the planetary mass companions like 2M1207\\,b (Chauvin et al. 2004) and, more recently, around Fomalhaut \\citep{2008Sci...322.1345K}, a triple system around HR 8799 \\citep{2009AAS...21331907M} and a close candidate around $\\beta$ Pic (Lagrange et al. 2009) confirm the high contrast and high angular resolution imaging capabilities of current instruments. To resolve ultra-cool and light companions, large telescopes equipped with adaptive optics (AO) are required, as well as dedicated focal devices, such as coronagraph and differential imagers.  Of special interest is the detection of planets still embedded in their parental proto-planetary disk, as this may give us clues on the timescales for planet formation and on the disk dissipation mechanisms.  Indirect information can be given by gaps and asymmetries observed in the disk structures as detected through scattered light images in near infrared (NIR) \\citep[e.g. for HD 141569, see][]{1999A&A...350L..51A} or, through millimeter interferometric measurements as recently done for AB Aur \\citep{2005A&A...443..945P} and TW Hydra \\citep{2007ApJ...664..536H}. Two controversial discoveries have been recently obtained one with RV for the giant planet ($M \\sim 10 M_{\\rm{Jup}}$) at 0.04 AU in the disk that surrounds TW Hya \\citep[See][]{2008Natur.451...38S,2008A&A...489L...9H}, the other by differential polarimetric and coronagraphic imaging for a point source (companion or disk over-density due to dust accretion on an unseen companion) in the outer annulus of dust of the AB Aur disk (Oppenheimer et al. 2008). More recently, a planetary candidate has been detected in the $\\beta$ Pic circumstellar disk (Lagrange et al. 2009).  If confirmed, these companions will directly evidence the link between planetary formation and the main morphological and dynamical peculiarities of transition disks.  In this paper, we present the outcomes of  NACO observations, taken using a four quadrant phase mask coronagraph, which aims at offering enhanced detection performances at small angular separations. Our goal was to detect the companion responsible for the large gap, observed at mm wavelengths, in the disk that surrounds the young star LkCa~15 \\citep[see][]{2006A&A...460L..43P}.  After summarizing the properties of this star and its proto-planetary disk in section~2, we present the NACO observations and the data reduction and analysis in section~3 and 4 respectively. We finally report our detection limits in section~5 and discuss them in regards of the domain of mass and orbital parameters explored for a putative companion. ",
        "conclusions": "\\label{sec:conclusions}  The T-Tauri star LkCa 15 was observed with VLT/NACO using the 4QPM coronagraph, reaching a contrast lower than 9.5 in $K_s$ band, at separations higher than 0.5~$\\!''$.  Our goal was the detection of a low-mass companion, with a mass spanning from 0.2~$M_{\\odot}$ down to 5~$M_{Jup}$, which presence has been suggested as an explanation for the large cavity, evidenced by sub-millimeter observations, in the disk surrounding the star.  We do not report any positive detection of a close companion to the star LkCa15. Based on our detection limits and Hot Start evolutionary model predictions, we ran simulations to take into account that the presence of a putative companion enable us to constrain its mass and semi-major axis given reasonable assumptions from the disk geometry.  The innovative approach in our case is the combined use of either the available information on the disk configuration (inclination, ascending node, gap prediction) and the complete 2D information of the coronagraphic image, to actually constrain the probability of existence of a stellar or sub-stellar companion around LkCa15. As a main result of our statistical analysis we can exclude the existence of a low mass star and brown dwarf companion  with semi-major axis larger than 55 and 80~AU respectively. The planetary mass regime is only partially constrained and at large semi-major axis ($a \\ge 100$~AU) the existence of massive ($M \\ge 6~M_{Jup}$) planetary mass companions seem unlikely.  The limitations of our study is also discussed in regards of the LkCa15 age uncertainty that does not affect much our conclusions in the case of low mass star and brown dwarf companions, and of the uncertainty related to the initial conditions adopted for the evolutionary model, particularly critical in the planetary mass regime. Finally, in case of a true companion orbiting LkCa15 in the inner disk cavity, our observations would favor a planetary mass or low-mass brown dwarf companion, although more massive companions cannot be completely excluded with high detection probability.  Further deep imaging studies at the 5--50~AU scale at a new epoch should provide complementary information to completely reject the existence of close stellar or brown dwarf companion to LkCa15 and pursue the search for the putative companion that would be responsible for the disk geometry and inner cavity within 46~AU."
    },
    "1004/1004.2448_arXiv.txt": {
        "abstract": "The Ultraviolet/Optical Telescope (UVOT) is one of three instruments onboard the {\\it Swift} observatory. The photometric calibration has been published, and this paper follows up with  details on other aspects of the calibration including a measurement of the point spread function with an assessment of the orbital variation and the effect on photometry. A correction for large scale variations in sensitivity over the field of view is described, as well as a model of the coincidence loss which is used to assess the coincidence correction in extended regions. We have provided a correction for the detector distortion and measured the resulting internal astrometric accuracy of the UVOT, also giving the absolute accuracy with respect to the International Celestial Reference System. We have compiled statistics on the background count rates, and discuss the sources of the background, including instrumental scattered light.  In each case we describe any impact on UVOT measurements, whether any correction is applied in the standard pipeline data processing or whether further steps are recommended. ",
        "introduction": "The Ultraviolet/Optical Telescope \\citep[UVOT;][]{Roming05} is one of three instruments which make up the {\\it Swift} observatory \\citep{GN2004} along with the 15--150~keV Burst Alert Telescope \\citep[BAT;][]{BS2005} and the 0.2--10~keV X-Ray Telescope \\citep[XRT;][]{BDN2005}. {\\em Swift} is primarily a gamma-ray burst (GRB) observatory, but is increasingly being used to look at a wide range of targets including active galactic nuclei, supernovae and X-ray transients.  The UVOT obtains ultraviolet and optical data in parallel with the BAT and the XRT. It incorporates a modified Ritchey--Chr\\'{e}tien telescope with a $17 \\times 17$ arcmin field-of-view and covers the wavelength range 1600--8000\\AA. A filter wheel carries seven broadband filters: three in the optical range ({\\em v, b {\\em and} u}), three in the UV ({\\em uvw1, uvm2 {\\em and} uvw2}) and one clear filter ({\\it white}) covering the whole wavelength range. In addition there are two grisms, a magnifier and a blocked filter in the filter wheel.  The UVOT uses a fast readout, micro-channel-plate (MCP) intensified, photon-counting CCD detector with $256 \\times 256$ active pixels. Each pixel is subdivided by use of an onboard centroiding algorithm into $8 \\times 8$ subpixels, giving a full field of $2048 \\times 2048$ subpixels, each of which subtends 0.502 arcsec on the sky after correcting for distortion. Throughout this paper the word `pixel' refers to one of these 0.5 arcsec subpixels, unless it is specifically called a CCD pixel. The full field is read out in 11.0329 ms and each photon may be time-tagged with this timing resolution; this is known as event mode. Alternatively, and more usually, image mode is used, where the photons are built up into an image on board to reduce telemetry. The UVOT is suitable for viewing sources from $\\sim~10.5$ to $\\sim~23.5$ mag, depending on the filter used and the nature of the source. Further details on the UVOT detector may be found in Section 2 of Poole et al. (2008, hereafter Paper I).  The Swift Data Center\\footnote{http://swift.gsfc.nasa.gov/docs/swift/sdc/} at Goddard Space Flight Center  is responsible for processing raw telemetry data from the {\\it Swift} satellite and making them available to the HEASARC\\footnote{http://heasarc.gsfc.nasa.gov/} for distribution to the public as reduced data products. The standard processing is known as `the pipeline'; the current release is 3.13.17 (August 3rd, 2009).  The Swift Science Center\\footnote{http://swift.gsfc.nasa.gov/docs/swift/swiftsc.html} provides a suite of software tools to the HEASARC suitable for analysing or reprocessing the data and applying the calibration (UVOT specific \\texttt{FTOOLS} are released as part of HEAsoft {\\it Swift} software package\\footnote{HEAsoft software can be found at \\newline http://heasarc.gsfc.nasa.gov/docs/software/lheasoft/}), as well as maintaining the Calibration database (CalDB\\footnote{CalDB files and associated documentation can be found at \\newline http://swift.gsfc.nasa.gov/docs/heasarc/caldb/swift/}; the current release of the UVOT CalDB is version 20090930).  The preliminary in-orbit calibration of UVOT was described in \\citet{spie1} and \\citet{spie2}. With this paper and Paper I we present a more detailed analysis supported by more data and superceding the earlier work. In Paper I we presented the photometry calibration. Here in Paper II we  include the following: measurements of the point spread function and the effect of the orbit (Section~\\ref{SecPSF}); a model of how coincidence loss is affected by high backgrounds (Section~\\ref{SecCoi}); mod-8 noise and how it can be removed from images (Section~\\ref{SecMod8}); small and large scale positional uniformity (Sections~\\ref{SecSSS} and \\ref{SecLSS}); the accuracy of UVOT astrometry both relative and absolute (Section~\\ref{SecAstrometry})  and  background count rates (Section~\\ref{SecBackground}), including scattered light. In each case we assess any impact on UVOT science, describe any correction that is applied in the standard pipeline data processing (and how good this correction is), or whether further steps are recommended under certain circumstances.  The grism calibration will be published in a separate paper. ",
        "conclusions": "This paper substantially improves the calibration of the UVOT and compliments the photometric calibration already covered in Paper~I. This paper also discusses more fully some issues raised in that paper.  We have measured the Point Spread Function out to a large radius (30 arcsec) for each filter to enable, for example, reliable surface photometry. We have also described how binning, rotation and orbital variations affect the PSF FWHM, and that photometry is not affected by the variations provided a big enough aperture is used.  Coincidence loss was described for a point source in Paper~I. Here we have extended the study to extended sources, and sources in regions of high background, presenting a model which not only shows us the limitations of the current coincidence correction, but also gives us the tools to improve it in the future. We also described the effect of mod-8 noise and how it can be removed without degrading photometric accuracy.  We have measured the sensitivity variation over the detector and have constructed an effective correction for each filter that enables us to achieve a photometric response uniform to 1--2 per cent over the entire detector area.  The positional accuracy of UVOT astrometry (after aspect correction) and distortion has been compared with the International Celestial Reference System and shown to be accurate to better than 0.42 arcsec (90 per cent confidence) including systematic and random errors.  Lastly, we have measured the range of observed background count rates in all filters and the dependencies on the Earth limb and Galactic latitude. We have discussed the sources of the background from both astrophysical and instrumental causes.  This extended calibration enables a more consistent use of UVOT data over the whole field of view. The work on the PSF, coincidence loss and astrometry are of particular use in the analysis of extended objects."
    },
    "1004/1004.2212_arXiv.txt": {
        "abstract": "We present a study of optical \\ion{Fe}{2} emission in 302 AGNs selected from the SDSS. We group the strongest \\ion{Fe}{2} multiplets into three groups according to the lower term of the transition (b${\\ }^4F$, a${\\ }^6S$ and a${\\ }^4G$ terms). These correspond approximately to the blue, central, and red part respectively of the ``iron shelf'' around H$\\beta$. We calculate an \\ion{Fe}{2} template which takes into account transitions into these three terms and an additional group of lines, based on a reconstruction of the spectrum of I Zw 1. This \\ion{Fe}{2} template gives a more precise fit of the \\ion{Fe}{2} lines in broad-line AGNs than other templates. We extract \\ion{Fe}{2}, H$\\alpha$, H$\\beta$, [\\ion{O}{3}] and [\\ion{N}{2}] emission parameters and investigate correlations between them. We find that \\ion{Fe}{2} lines probably originate in an Intermediate Line Region. We notice that the blue, red, and central parts of the iron shelf have different relative intensities in different objects. Their ratios depend on continuum luminosity, FWHM H$\\beta$, the velocity shift of \\ion{Fe}{2}, and the $\\mathrm{H\\alpha}$/$\\mathrm{H\\beta}$ flux ratio. We examine the dependence of the well-known anti-correlation between the equivalent widths of \\ion{Fe}{2} and [\\ion{O}{3}] on continuum luminosity. We find that there is a Baldwin effect for [\\ion{O}{3}] but an inverse Baldwin effect for the \\ion{Fe}{2} emission. The [\\ion{O}{3}]/\\ion{Fe}{2} ratio thus decreases with L$_{\\lambda5100}$. Since the ratio is a major component of the Boroson and Green eigenvector 1, this implies a connection between the Baldwin effect and eigenvector 1, and could be connected with AGN evolution. We find that spectra are different for H$\\beta$ FWHMs greater and less than $\\sim$3000 $\\mathrm{kms^{-1}}$, and that there are different correlation coefficients between the parameters. ",
        "introduction": "Optical \\ion{Fe}{2} ($\\lambda\\lambda$ 4400-5400 \\AA) emission is one of the most interesting features in AGN spectra. The emission arises from numerous transitions of the complex \\ion{Fe}{2} ion. \\ion{Fe}{2} emission is seen in almost all type-1 AGN spectra and it is especially strong in narrow-line Seyfert 1s (NLS1s). It can also appear in the polarized flux of type-2 AGNs when a hidden BLR can be seen. Origin of the optical \\ion{Fe}{2} lines, the mechanisms of their excitation and location of the \\ion{Fe}{2} emission region in AGN, are still open questions. There are also many correlations between \\ion{Fe}{2} emission and other AGN properties which need a physical explanation.  Some of the problems connected with \\ion{Fe}{2} emission are:  1.  There have been suggestions that \\ion{Fe}{2} emission cannot be explained with standard photoionization models \\citep[see][and references therein]{b60,bb}. To explain the \\ion{Fe}{2} emission, additional mechanisms of excitation have been proposed:  a) continuum fluorescence \\citep{b1,b2},  b) collisional excitation \\citep{b3,b4,b5,b6}, c) self-fluorescence among \\ion{Fe}{2} transitions \\citep{b7},  d) fluorescent excitation by the Ly$\\alpha$ and Ly$\\beta$ lines \\citep{b8,b300}.  2. There are uncertainties about where the \\ion{Fe}{2} emission line region is located in an AGN. A few proposed solutions are that the \\ion{Fe}{2} lines arise: a) in the same emission region as the broad H$\\beta$ line \\citep{b9}, b) in the accretion disk near the central black hole, which produces the double-peaked broad Balmer emission lines  \\citep{bb, b10}, c) in a region which can be heated by shocks or from an inflow, located between BLR and NLR -- the so-called the Intermediate Line Region \\citep{b001,b31,b66,b222,b51,b151} and d) and in the shielded, neutral, outer region of a flattened BLR \\citep{baa,ba}.  3. It has long been established that the \\ion{Fe}{2} emission is  correlated with the radio, X and IR continuum. \\ion{Fe}{2} lines are stronger in spectra of Radio Quiet (RQ) AGNs, than in Radio Loud (RL) objects \\citep{b11,b12}. But, if we consider RL AGNs, \\ion{Fe}{2} emission is stronger in Core Dominant (CD) AGNs than in Lobe Dominant (LD) AGNs \\citep{b13,b5}.  The relative strength of the \\ion{Fe}{2} lines correlates with the soft X-ray slope, but anti-correlates with X-ray luminosity \\citep{b14,bc,b15}. Strong optical \\ion{Fe}{2} emission is usually associated with strong IR luminosity \\citep{b17,b19}, but it is anti-correlated with IR color index $\\alpha$(60,25) \\citep{be}.  4.  Many correlations have been observed between \\ion{Fe}{2} and other emission lines. \\citet{b9} give a number of correlations between the  equivalent width (EW) of \\ion{Fe}{2} and properties of the [\\ion{O}{3}] and H$\\beta$ lines. The most interesting are the anti-correlations of EW \\ion{Fe}{2} vs. EW [\\ion{O}{3}]/EW H$\\beta$, EW \\ion{Fe}{2} vs. peak [\\ion{O}{3}], EW \\ion{Fe}{2}/EW H$\\beta$ vs. peak [\\ion{O}{3}], and the correlation of EW \\ion{Fe}{2}/EW H$\\beta$ vs. H$\\beta$ asymmetry. An anti-correlation between EW \\ion{Fe}{2} and FWHM of H$\\beta$ has been also found \\citep{b20,b21,b9}.  These correlations dominate the first eigenvector of the principal component analysis of  \\citet{b9} (hereinafter EV1). Correlations between EV1 and some properties of C IV $\\lambda1549$ \\AA \\ lines have also been pointed out \\citep{b22}. With increasing continuum luminosity, the EW of the C IV lines decreases, \\citep[the so-called ``Baldwin effect'',][]{b35,b23}, while the blueshift of the C IV line and the EW of \\ion{Fe}{2} increase.  Recently, \\citet{bf} have found the Baldwin effect to be related to EV1. \\ion{Fe}{2} lines are usually strong in AGN spectra with low-ionization Broad Absorption Lines (BALs) \\citep{b24,b18}.  Searching for a physical cause of these correlations may help in understanding of the origin of \\ion{Fe}{2} emission. Some of physical properties which may influence these correlations are: Eddington ratio (L/$\\mathrm{L_{Edd}}$) \\citep{b9,b22}, black hole mass ($M_{BH}$) and inclination angle \\citep{b25,b26,b27}.  In addition to these properties, \\ion{Fe}{2} emission strength may vary with evolution of AGNs. For example, \\citet{b29} have explained some properties of AGN (as BALs, \\ion{Fe}{2} strength, radio properties, emission line width, luminosity of narrow lines, etc.) by an ``evolutive unification model''. In this model, accretion arises from the interaction between nuclear starbursts and the supermassive black hole.  This would mean that not only the orientation, but also the evolutionary state of an AGN, influence its spectral properties. Thus, young AGN have strong \\ion{Fe}{2}, BALs, weak radio emission, the NLR is compact and faint, and broad lines are relatively narrow. In contrast with this, old AGNs have weak \\ion{Fe}{2}, no BALs, strong radio emission with extended radio lobes, the NLR is extended and bright, and the broad lines have greater velocity widths.  In this paper, we investigate the \\ion{Fe}{2} emitting region by analyzing the correlations between the optical \\ion{Fe}{2} lines and the other emission lines in a sample of 302 AGN from the SDSS. To do this we construct an \\ion{Fe}{2} template. The strongest \\ion{Fe}{2} multiplets within the $\\lambda\\lambda$4400-5500 \\AA \\ range are sorted into three groups, according to the lower terms of the transitions: $3d^6({\\ }^3F2)4s{\\ }^4F$, $3d^54s^2{\\ }^6S$ and $3d^6({\\ }^3G)4s{\\ }^4G$ (which we will hereinafter refer to as the $F$, $S$ and $G$ groups of lines). The $F$ group mainly contains the lines from the \\ion{Fe}{2} multiplets 37 and 38 and describes the blue part of the \\ion{Fe}{2} shelf relative to H$\\beta$. The $S$ group of lines describes the part of \\ion{Fe}{2} emission under H$\\beta$ and [\\ion{O}{3}] and contains lines from multiples 41, 42 and 43 \\ion{Fe}{2}. Finally, the $G$ group contains lines from multiplets 48 and 49, and describes the red part of the \\ion{Fe}{2} shelf. A simplified Grotrian diagram of these transitions is shown in Figure \\ref{1}. We analyze separately their relationships with other lines in AGN spectra (H$\\beta$, [\\ion{O}{3}], H$\\alpha$, [\\ion{N}{2}]). In this way, we try to connect details of the \\ion{Fe}{2} emission with physical properties of the AGN emission regions in which these lines arise. Each of the \\ion{Fe}{2} multiplet groups has specific characteristics, which may be reflected in a different percent of correlations with other lines, and give us more information about the complex \\ion{Fe}{2} emission region.  The paper is organized as follows: In Sect. 2 we describe the procedure of sample selection and details of analysis; our results are presented in Sect. 3; in Sect. 4 we give discussion of obtained results, and finally in Sect. 5 we outline our conclusions. ",
        "conclusions": "  \\begin{enumerate}  \\item  We have proposed here an optical \\ion{Fe}{2} template for the $\\lambda\\lambda$4400-5500 \\AA \\ range, which consists of three groups of  \\ion{Fe}{2} multiplets, grouped according to the lower terms of transitions ($F$, $S$ and $G$), and an additional group of lines reconstructed from the I Zw 1 spectrum. We found that template can satisfactorily fit the \\ion{Fe}{2} lines. In spectra in which \\ion{Fe}{2} emission lines have different relative intensities than in I Zw 1, this template fit better than empirical and theoretical templates based on I Zw 1 spectrum (see more in Appendix B). Using this template, we are able to consider different groups of transitions which contribute to the blue, central and red part of \\ion{Fe}{2} shelf around the H$\\beta$+[\\ion{O}{3}] lines.  \\item We find that the ratios of different parts of the iron shelf ($F/G$, $F/S$, and $G/S$) depend of some spectral properties such as: continuum luminosity, H$\\beta$ FWHM, shift of \\ion{Fe}{2}, and $\\mathrm{H\\alpha}$/$\\mathrm{H\\beta}$ flux ratio. Also, it is noticed that spectra with H$\\beta$ FWHM greater and less than $\\sim$3000 $\\mathrm{kms^{-1}}$ have different properties which is reflected in significantly different coefficients of correlation between the parameters.  \\item We found that the \\ion{Fe}{2} emission is mainly characterized with a random velocity of $\\sim$1400 $\\mathrm{kms^{-1}}$ that corresponds to the ILR origin, which is also supported by the significant correlation between the width of \\ion{Fe}{2} and H$\\alpha$, H$\\beta$ ILR widths. This is in agreement with the previous investigations \\citep{b31,b51,b151}, but unlike the earlier investigations, we found a slight correlation with the width of VBLR.  Therefore, it is possible that \\ion{Fe}{2} is partly produced in the VBLR and cannot be resolved from continuum luminosity. Also, we found that the \\ion{Fe}{2} lines  are slightly redshifted relative to the narrow lines ($\\sim$ 100 $\\mathrm{kms^{-1}}$) (see Figure \\ref{18}).  \\item  The Balmer lines were decomposed into NLR, ILR and VBLR components and relationships among H$\\beta$ components and some spectral properties were investigated. We found a positive correlation between the width of narrow lines and EW \\ion{Fe}{2} ($r=0.30$, $P < 0.0001$),  while the width of the broad H$\\beta$ components anti-correlates with EW \\ion{Fe}{2}. Also, we found a Baldwin effect in the case of H$\\beta$ NLR component, while the ILR component showed no correlation with continuum luminosity, and VBLR component shows an inverse Baldwin effect for the FWHM $<$ 3000 $\\mathrm{kms^{-1}}$ subsample.   \\item  We confirm in our sample the anti-correlation between EW \\ion{Fe}{2} and EW [\\ion{O}{3}]  which is related to Eigenvector 1 (EV1) in \\citet{b9} and we examined its dependence on the continuum luminosity and redshift. We found an inverse Baldwin effect for \\ion{Fe}{2} lines from the central and red part of the \\ion{Fe}{2} shelf ($S$ and $G$), but for \\ion{Fe}{2} lines from blue part ($F$) no correlation was observed. A Baldwin effect was confirmed for the [\\ion{O}{3}] lines. Since EW \\ion{Fe}{2} increases, and EW [\\ion{O}{3}] decreases with increases of continuum luminosity (and also the trend of decreasing of EW [\\ion{O}{3}]/EW \\ion{Fe}{2} ratio with luminosity is obvious), the observed EW \\ion{Fe}{2} vs. EW [\\ion{O}{3}] anti-correlation is probably due to the same physical reason which causes the Baldwin effect.  Moreover, it is observed that the coefficients of correlation due to Baldwin effect depend on H$\\beta$ FWHM range of a sub-sample, which also implies the connection between the Baldwin effect and EV1. A decreasing trend is observed for EW [\\ion{O}{3}]/EW H$\\beta$ vs. continuum luminosity.  \\item  We found that the equivalent width of the \\ion{Fe}{2} lines, as well as the EW H$\\beta_{\\mathrm{total}}$/EW \\ion{Fe}{2} ratio have a stronger correlation with redshift, than with continuum luminosity. This implies that the inverse Baldwin effect of \\ion{Fe}{2} may be primarily caused by evolution.  \\end{enumerate}"
    },
    "1004/1004.5432_arXiv.txt": {
        "abstract": "Detecting dark matter as it streams through detectors on Earth relies on knowledge of its phase space density on a scale comparable to the size of our solar system. Numerical simulations predict that our Galactic halo contains an enormous hierarchy of substructures, streams and caustics, the remnants of the merging hierarchy that began with tiny Earth mass microhalos. If these bound or coherent structures persist until the present time, they could dramatically alter signatures for the detection of weakly interacting elementary particle dark matter (WIMP). Using numerical simulations that follow the coarse grained tidal disruption within the Galactic potential and fine grained heating from stellar encounters, we find that microhalos, streams and caustics have a negligible likelihood of impacting direct detection signatures implying that dark matter constraints derived using simple smooth halo models are relatively robust. We also find that many dense central cusps survive, yielding a small enhancement in the signal for indirect detection experiments. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.4214_arXiv.txt": {
        "abstract": "We present new measurements of dark matter distributions in 25 X-ray luminous clusters by making a full use of the two-dimensional (2D) weak lensing signals obtained from high-quality Subaru/Suprime-Cam imaging data. Our approach to directly compare the measured lensing shear pattern with elliptical model predictions allows us to extract new information on the mass distributions of individual clusters, such as the halo ellipticity and mass centroid. We find that these parameters on the cluster shape are little degenerate with cluster mass and concentration parameters. By combining the 2D fitting results for a subsample of 18 clusters, the elliptical shape of dark matter haloes is detected at $7\\sigma$ significance level. The mean ellipticity is found to be $\\langle e \\rangle = \\langle 1-b/a \\rangle =0.46\\pm 0.04$ ($1\\sigma$), which is in excellent agreement with a theoretical prediction based on the standard collisionless cold dark matter model. The mass centroid can be constrained with a typical accuracy of  $\\sim 20$~arcseconds ($\\sim 50h^{-1}$~kpc) in radius for each cluster. The mass centroid position fairly well matches the position of the brightest cluster galaxy, with some clusters showing significant offsets. Thus the 2D shear fitting method enables to assess one of the most important systematic errors inherent in the stacked cluster weak lensing technique, the mass centroid uncertainty.  In addition, the shape of the dark mass distribution is found to be only weakly correlated with that of the member galaxy distribution or the brightest cluster galaxy. We carefully examine possible sources of systematic errors in our measurements including the effect of substructures, the cosmic shear contamination, fitting regions, and the dilution effect, and find none of them to be significant. Our results demonstrate the power of high-quality imaging data for exploring the detailed spatial distribution of dark matter, which should improve the ability of future surveys to conduct cluster cosmology experiments. ",
        "introduction": "\\label{sec:intro}  The internal structure of clusters predicted by $N$-body simulations of the collisionless Cold Dark Matter (CDM) universe exhibits several important features, including the universal radial density profile with progressively shallower slopes toward the centre \\citep[e.g.,][]{navarro96,navarro97} and the highly non-spherical structure fitted well by a triaxial density profile \\citep[e.g.,][]{jing02}. These features are closely related to the cold and collisionless nature of hypothetical dark matter particles \\citep[e.g.,][]{ostriker03}, which suggests that detailed observational tests of cluster density profiles may provide a clue to the nature of dark mater. In addition, an accurate determination of the mass distribution of individual clusters is also important in calibrating observable-mass relations which are necessary for constraining cosmological parameters from cluster abundance \\citep[e.g.,][]{bahcall93,eke96,kitayama97,henry00,borgani01, bahcall03,dahle06,mantz08,henry09,vikhlinin09,rozo10}.  Gravitational lensing is one of the most powerful methods to constrain the cluster mass distribution. This is because gravitational lensing probes the matter distribution directly, regardless of whether it is luminous or dark. This is in marked contrast with other probes such as X-ray and Sunyaev-Zel'dovich effect for which assumptions on gas state and other pressure components are crucial in extracting the mass distribution. In particular the cluster mass distribution has extensively been studied using weak lensing technique which makes use of small distortions of background galaxies produced by a foreground cluster \\citep[see][for a review]{bartelmann01}. Indeed the weak lens technique has been quite successful in constraining dark matter distributions in clusters independently of the distribution of the hot gas component \\citep{kneib03,clowe06,jee07,mahdavi07,mahdavi08,okabe08,bradac08}.  While the weak lensing technique allows the non-parametric reconstruction of the mass map \\citep{kaiser93}, azimuthally-averaged one-dimensional (1D) tangential shear profiles have been used in most of quantitative studies so far \\citep[e.g.,][]{bardeau07,broadhurst08,hamana09,oguri09b,okabe10}. The main reason for this is that the 1D profile fitting provides an easy and efficient way to extract several important cluster properties, such as mass, concentration, and radial density profiles of clusters. However, the highly non-spherical nature of CDM haloes predicted by $N$-body simulations indicates that the full two-dimensional (2D) study of weak lensing data is crucial for further tests of the cluster mass distribution. Moreover, an important disadvantage of the 1D profile fitting is that we have to assume the centre of the cluster a priori. It has often been assumed that the mass centroid coincides with the location of the brightest cluster galaxy (BCG), despite the fact that offsets of BCGs from the mass centroids appear to be common \\citep[e.g.,][]{katayama03,lin04,koester07,ho09,sanderson09}.  The 2D weak lensing studies with non-spherical mass models have been attempted, but only for a handful of lensing clusters \\citep[e.g.,][]{cypriano04,deb10}. Fitting of weak lensing data by taking fully into account a 3D (triaxial) cluster shape and the projection along arbitrary line-of-sight has been performed by \\citet{oguri05}, and later by \\citet{corless09}, but their interests lied in the determination of masses and concentration parameters rather than the shape of the projected mass distribution. Given the strong selection bias of the apparent 2D cluster ellipticity for lensing clusters \\citep{oguri09a}, more systematic 2D weak lensing studies of massive clusters should be conducted using well-defined statistical cluster samples.  Cluster mass distributions have also been studied in details using stacked weak lensing data \\citep{sheldon07a,sheldon07b,johnston07, mandelbaum06b,mandelbaum08}. Although the stacked weak lensing technique is quite successful in extracting average properties of clusters, its main disadvantage comes from the centroid problem \\citep[e.g.,][]{johnston07,mandelbaum10}; in stacking clusters one has to assume a priori the mass centroid for each cluster. The misidentification of the mass centroid results in the suppression of weak lensing signals, particularly near the cluster centre. Recently, \\citet{evans09} measured the ellipticity of isolated clusters with stacked weak lensing technique to obtain the average cluster ellipticity of $e\\sim 0.5$.  However, their result is built on the assumption that the cluster mass distributions is aligned with the spatial distribution of member galaxies, and also on the assumption about the cluster centre, which  makes the interpretation somewhat difficult.  In this paper, we present a systematic 2D study of weak lensing maps. We analyse 25 clusters at $0.15<z<0.3$ presented by \\citet{okabe10}, who examined azimuthally-averaged 1D tangential profiles to study the radial mass profiles of the clusters.  The cluster sample is essentially X-ray flux-limited statistical sample from the Local Cluster Substructure Survey (LoCuSS; G. Smith et al., in preparation). In this paper we extend the lensing analysis of \\citet{okabe10} and perform a full 2D fitting of the shear fields with particular emphasis on the projected 2D ellipticities of the clusters. All the weak lensing data are based on optical images taken with Subaru/Suprime-cam \\citep{miyazaki02} which is known to be the current best ground facility for weak lensing studies because of its exquisite image quality and wide field-of-view \\citep[e.g.,][]{kasliwal08}. In our analysis we leave the centre of the mass profile as a free parameter and fit it simultaneously in order to see how the 2D analysis of shear maps constrains the cluster mass centroid which has often been assumed to coincide with the location of the BCG. Our approach to fit individual clusters with the centroid and orientation of the cluster as free parameters not only overcomes the problems inherent to stacked weak lensing technique but also allows us to even test the assumptions on the centroid and orientation from weak lensing data alone.  The structure of this paper is as follows. We describe our fitting procedure in \\S\\ref{sec:fit}, and present results in \\S\\ref{sec:result}. In \\S\\ref{sec:discussion} we discuss possible systematic effects, and in \\S\\ref{sec:conclusion} we draw our conclusions. Throughout the paper we assume a flat universe with the matter density $\\Omega_M=1-\\Omega_\\Lambda=0.274$, the dimensionless Hubble constant $h=0.704$, the baryon density $\\Omega_b h^2=0.02267$, the spectral index $n_s=0.96$, and the normalisation of the matter power spectrum $\\sigma_8=0.812$ \\citep{komatsu09}. ",
        "conclusions": "\\label{sec:conclusion}  In this paper, we have presented a systematic study to explore the ellipticity of cluster mass distribution directly from the 2D pattern of weak lensing distortion measured from background galaxy images. Our sample of high quality Subaru/Suprime-cam images for 25 X-ray selected clusters \\citep{okabe10} has enabled reliable statistical studies of dark matter distributions in massive cluster in great details. We developed a method of comparing the measured distortion patter with the prediction computed from an elliptical NFW model, where the halo ellipticity is modeled in the projected mass density and the resulting shear field is computed by solving the 2D Poisson equation (see \\S~\\ref{sec:model}).  Also notable is our method properly includes the error covariance matrix arising from the shot noise due to intrinsic ellipticities as well as the cosmic shear contamination due to large scale structures that are at different redshifts, but along the same line of sight to the cluster (see \\S~\\ref{sec:2dfit}).  By conducting the 2D shear fitting, we can extract information on the shape of dark matter distribution in the cluster region in addition to the cluster mass and the profile parameters (halo concentration). We have shown several interesting results, and our findings are summarised as follows.  From a sample of 18 clusters, we have determined the mean ellipticity of the dark matter distribution to be $\\langle e \\rangle=\\langle 1-b/a \\rangle =0.46\\pm 0.04$ (1$\\sigma$), in excellent agreement with the prediction of the triaxial halo model by \\citet{jing02} (see Figure~\\ref{fig:ell}). By comparing the results with the fitting results obtained from the spherical NFW model ($e=0$), we concluded that the ellipticity is detected at $7\\sigma$ confidence level, which has been achieved without assuming the location of mass peaks and the alignment between mass and light. Interestingly, we found that the constraints on halo ellipticity parameters are not largely degenerated with halo mass and concentration, implying that the halo ellipticity is constrained mainly from the 2D shear pattern rather than the shear amplitudes (see Figure~\\ref{fig:cont_a2390}).  According to the CDM predictions, dynamically young and more massive haloes tend to have larger ellipticities. That is, the halo ellipticity should depend on redshift and mass of clusters as well as on cosmological models. However, our results are not sufficient to discriminate the model differences due to the limited size of our cluster sample (see Figure~\\ref{fig:ell_paras}).  It appears that our results are robust against several possible systematics, such as the dilution effect, substructures in clusters, and projections of other haloes along the line-of-sight. For masses and concentrations, we find no significant systematic differences between the 1D and 2D analysis (see Figure~\\ref{fig:mc_1d2d}). The mass centroid positions constrained from weak lensing is typically consistent with the BCG positions with a typical accuracy of $\\sim 20$~arcseconds ($\\sim 50h^{-1}$~kpc), although there are a few clusters that show significant offsets (see Figure~\\ref{fig:dis}). We have found that the correlation between the orientations of the dark matter distribution and cluster member galaxy distributions is not very strong, possibly because of the limitation of our member galaxy sample which is defined photometrically using only two-band images. Therefore it is of interest to repeat the analysis using spectroscopically confirmed cluster members.  Overall, our results demonstrate the importance of analysing the full 2D distributions of cluster weak lensing data. The 2D shear map contains much richer information than azimuthally-averaged 1D tangential shear profiles which have been a main focus in most of previous studies. Direct tests of dark matter distributions should grow its importance in the era of wide-field imaging surveys which will bring an exciting opportunity to explore the full use of the method presented in this paper in order to obtain stringent, quantitative constraints on the collisionless CDM scenario on small non-linear scales."
    },
    "1004/1004.3602_arXiv.txt": {
        "abstract": "Pulsars are very stable clocks in space which have many applications to problems in physics and astrophysics. Observations of double-neutron-star binary systems have given the first observational evidence for the existence of gravitational waves (GWs) and shown that Einstein's general theory of relativity is an accurate description of gravitational interactions in the regime of strong gravity. Observations of a large sample of pulsars spread across the celestial sphere forming a ``Pulsar Timing Array'' (PTA), can in principle enable a positive detection of the GW background in the Galaxy. The Parkes Pulsar Timing Array (PPTA) is making precise timing measurements of 20 millisecond pulsars at three radio frequencies and is approaching the level of timing precision and data spans which are needed for GW detection. These observations will also allow us to establish a ``Pulsar Timescale'' and to detect or limit errors in the Solar System ephemerides used in pulsar timing analyses. Combination of PPTA data with that of other groups to form an International Pulsar Timing Array (IPTA) will enhance the sensitivity to GWs and facilitate reaching other PTA goals. The principal source of GWs at the nanoHertz frequencies to which PTAs are sensitive is believed to be super-massive binary black holes in the cores of distant galaxies. Current results do not signficantly limit models for formation of such black-hole binary systems, but in a few years we expect that PTAs will either detect GWs or seriously constrain current ideas about black-hole formation and galaxy mergers. Future instruments such as the Square Kilometre Array (SKA) should not only detect GWs from astrophysical sources but also enable detailed studies of the sources and the gravitational theories used to account for the GW emission. ",
        "introduction": "\\label{sec:intro} A pulsar is a rapidly rotating neutron star with very strong magnetic fields both inside and outside the star, most likely formed in a supernova explosion at the end of the life of a massive star. The combination of strong magnetic field and rapid rotation generates enormous electric fields which accelerate charged particles in the surrounding magnetosphere to ultra-relativistic energies. These particles radiate beams of emission which sweep across the sky as the star rotates. If one or more of these beams cross the Earth each rotation we can detect the neutron star as a pulsar. The interval between pulses, known as the pulse period, is equal to the rotation period of the star. There are currently more than 1800 pulsars known.\\footnote{See the ATNF Pulsar Catalogue \\cite{mhth05} at www.atnf.csiro.au/research/pulsar/psrcat.}  Almost all of these lie within our Galaxy, typically at distances of a few kiloparsec.  There are two main classes of pulsar: ``normal'' pulsars and ``millisecond'' pulsars (MSPs). Most normal pulsars have periods of between 0.1 and 5 seconds, whereas MSPs have much shorter periods, generally between 2 and 50 ms, hence the name. Most MSPs are binary, that is, in orbit with another star. Orbital periods are typically days or months, but can be as short as a few hours. Only a small proportion of normal pulsars, just over 1\\%, are binary. This gives a pointer to the mechanism by which MSPs get their very short period -- it is believed that they were old, slow and probably dead pulsars which were spun up by accretion of mass and angular momentum from an evolving binary companion star (see, e.g., Ref.~\\refcite{bv91}). As well as increasing the spin frequency, the accretion evidently reduces the effective magnetic field by several orders of magnitude. Despite this, the increased spin rate is generally sufficient to reactivate the emission beams -- hence the resulting short-period pulsars are often known as ``recycled'' pulsars. Most of these are MSPs.  In general, pulsars are most easily detected at radio wavelengths and hence most pulsar surveys are carried out in the radio band (see, e.g., Ref.~\\refcite{mlc+01}). However, with the recent launch of the {\\em Fermi} Gamma-ray Space Telescope, an increasing number of pulsars, mostly young but also including many MSPs, have been found or detected at gamma-ray wavelengths \\cite{aaa+09c,aaa+09f}. At the time of writing, the number of reported gamma-ray detected pulsars is 46.  Because neutron stars are tiny, with a radius of about 15 km, but massive, typically about 1.4 M$_\\odot$, their rotation period is extremely stable -- they are very difficult to slow down or spin up! It is this property which makes pulsars extraordinarily useful as probes. This is especially true for MSPs as they have narrow pulses (in time units) and hence measurements of pulse delays and periodicities have higher precision compared to normal pulsars. Having an array of precise clocks spread through the Galaxy enables many different investigations, for example, studies of the interstellar medium, pulsar astrometric measurements, including proper motions and even annual parallaxes in some cases, and potentially the direct detection of gravitational waves (GWs). The fact that most of the MSPs are in orbit with another star is a bonus as it allows many additional properties relating to binary motion to be investigated.  For double-neutron-star binary systems measurement of relativistic perturbations to the orbital parameters is generally possible, thereby allowing precise determinations of neutron star masses and making possible stringent tests of general relativity (GR) and other gravitational theories in the strong-field regime. The famous Hulse-Taylor binary pulsar, PSR B1913+16, was the first binary pulsar to be discovered \\cite{ht75a}. Precise timing measurements over 30 years have given significant measurements of three relativistic or ``Post-Keplerian'' parameters, leading to accurate mass determinations for the two stars which confirm that it is a double-neutron-star system. They have also shown that the observed orbit decay is consistent with the predictions of GR for energy loss from the system in GWs, thereby providing the first observational evidence for the existence of these waves \\cite{wt05}. The Double Pulsar, PSR J0737-3039A/B, discovered at Parkes in 2003/4 \\cite{bdp+03,lbk+04}, is an even more relativistic binary system than the Hulse-Taylor binary. Relativistic precession of the longitude of periastron is an amazing 17$^{\\rm o}$~yr$^{-1}$ and was detected in just a few days of observation. In less than three years of observation five Post-Keplerian parameters have been observed, which together with the mass ratio of the two stars (measurable since both stars are detected as pulsars), has shown that GR accurately predicts the motions of the two stars at a level of precision exceeding 0.05\\%, the most precise test so far in this strong-field regime \\cite{ksm+06}.  The main topic of this presentation is the current effort to make a direct detection of GWs using pulsar timing. The concept of a ``Pulsar Timing Array'' (PTA) is described in Section~\\ref{sec:pta} and Section~\\ref{sec:ppta} describes the Parkes Pulsar Timing Array (PPTA) in some detail. Section~\\ref{sec:discn} discusses some of the implications of current results and prospects for the future. Finally, the main points are summarised in Section~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:discn} Current PTAs are most sensitive to GWs in the frequency range 3 -- 10 nHz. The astrophysical GW signal in this frequency range most likely to be detected is a stochastic background from super-massive black-hole binary systems in the cores of distant galaxies \\cite{jb03,wl03a}. Simulations of the expected GW signal by Sesana, Vecchio \\& Colacino (Ref.~\\refcite{svc08}) show that it is dominated by radiation from galaxies with intermediate redshifts $z\\sim 1$ (Fig.~\\ref{fg:gw_sim}) and that the strongest signals come from systems with black-hole masses of $10^9$ -- $10^{10}$ M$_\\odot$.  \\begin{figure}[ht] \\begin{center} \\psfig{file=svc08_f4.eps,width=90mm} \\end{center} \\caption{Contribution of different redshift intervals to the stochastic GW background for two GW frequencies, $8\\times 10^{-9}$~Hz (upper curve in each subplot) and $10^{-7}$~Hz (lower curve in each subplot). The upper panel shows the signal amplitude and the lower panel shows the number of galaxies contributing. The solid lines are from a Monte Carlo method and the dashed lines are from a semi-analytic approach. (From Ref.~\\refcite{svc08}.)} \\label{fg:gw_sim} \\end{figure}  Fig.~\\ref{fg:gw_amp} shows the range of model predictions for the strain amplitude as a function of GW frequency. Current published limits (e.g., Ref.~\\refcite{jhv+06}) do not constrain the model, but the ``ideal'' PPTA or, more realistically, longer data spans with current instrumentation and combined IPTA data sets will either detect the stochastic GW background or significantly modify our current understanding of black-hole formation in galaxy cores and/or galaxy merger processes. Future instruments such as the Square Kilometre Array (SKA) will have enormously improved sensitivity and are expected to obtain high-precision TOAs for hundreds of MSPs \\cite{kbc+04}. With a data span of five years or more, this will certainly give a significant detection of GWs unless the current predictions are totally wrong. It will enable not just a detection of GWs, but also detailed study of the signal. For example, if one of several non-Einsteinian theories of gravitation is correct, then GWs will have polarization properties which differ from those predicted by GR and this will have an effect on the observed correlations \\cite{ljp08,cjp09}. Even current results are constraining some proposed sources of the GW background. For example, the results presented by Jenet et al. (Ref.~\\refcite{jhv+06}) limit the equation of state of matter at the epoch of inflation, $w = P/\\epsilon > -1.3$ (cf. Ref.~\\refcite{gri05}) and the tension in cosmic strings (cf. Ref.~\\refcite{dv05}).  \\begin{figure}[ht] \\begin{center} \\psfig{file=svc08_f13.eps,width=90mm} \\end{center} \\caption{The range of predictions of the strain amplitude as a function of GW frequency from different models is represented by the hashed area. The lines within this region are predictions from particular models. The long-dashed lines are the approximate sensitivity limits for PTAs, with the upper one being the current best published limit, the middle one being an ``ideal'' PPTA with 20 pulsars, 100 ns rms timing residuals and a 5-year data span, and the lower one is a prediction of the limit obtainable with the proposed Square Kilometer Array radio telescope after 10 years. (From Ref.~\\refcite{svc08}.) } \\label{fg:gw_amp} \\end{figure}  Intermediate-mass black holes (IMBHs), with masses in the range of a few 100 to a few 1000 M$_\\odot$, have been invoked as the basis of the so-called ultra-luminous X-ray sources (see Ref.~\\refcite{mc04a} for a review). IMBHs are difficult to form by aggregation of stellar-mass objects. It has been suggested by Saito \\& Yokoyama (Ref.~\\refcite{sy09}) that the most likely formation path is by collapse of overdense regions at the end of the inflation era. Saito \\& Yokoyama show that formation of such IMBHs at this time will lead to a potentially detectable background of GWs. Fig.~\\ref{fg:imbh} shows the spectrum of the emitted GWs. The GW frequency is dependent on the mass of the formed black hole and is in the PTA range for IMBHs. In fact, the predicted amplitude of the GW background from IMBH collapse is greater than the limits already placed by pulsar timing, making primordial formation of IMBHs unlikely. This raises the issues of whether or not such IMBHs exist, what the energy source of ultra-luminous X-ray sources is and maybe even whether or not these X-ray sources are actually ultra-luminous.  \\begin{figure}[ht] \\begin{center} \\psfig{file=imbh_sy09.eps,width=100mm} \\end{center} \\caption{Spectrum of GWs emitted by formation of black holes at the end of the inflation era (Ref.~\\refcite{sy09}). For a given formation event, the emitted spectrum is concentrated at a GW frequency which is inversely related to the collapsing mass as indicated on the upper x-scale. Spectra are shown for two specific black-hole masses: 600 M$_\\odot$ (IMBH, peak around $10^{-8}$~Hz in the PTA frequency range) and $1.5\\times 10^{-11}$~M$_\\odot$ (Dark Matter, peak around 0.1~Hz in the LISA frequency range). These spectra peak near the dotted and dashed lines which are envelopes for the GW emssion from this mechanism with two different sets of assumptions.  Sensitivity limits for other proposed or existing instruments are also shown.  Current pulsar limits, indicated by the thick black line, already severely limit IMBH formation by this process. (See Ref.~\\refcite{sy09} for more details.)} \\label{fg:imbh} \\end{figure}   It is possible that future PTAs will have sufficient sensitivity to detect individual sources of GWs.  Sesana, Vecchio \\& Volonteri (Ref.~\\refcite{svv09}) show that current models for the evolution of super-massive binary black holes in galaxies predict that at least one such system will produce a sinusoidal signal of period comparable to the data span with an amplitude of between 5 and 50 ns in timing residuals. A signal of this size is potentially detectable with the IPTA and certainly should be detectable by the SKA. Such a detection would open up the field of GW astronomy, allowing investigation of the source characteristics and the properties of the GW emission itself. For example, it will be possible to localise the source on the sky. Fig.~\\ref{fg:sky_posn} shows a simulation of the ability of the PPTA to determine source positions \\cite{abc+09}. Since the PPTA pulsars are predominantly in the southern hemisphere, source positions are better determined in the south with a point-spread function of half-power width about 15$^{\\rm o}$ or 20$^{\\rm o}$. An improved sky distribution of pulsars and higher sensitivity from the IPTA or SKA will improve the accuracy to which a source can be localised on the sky.  \\begin{figure}[ht] \\begin{center} \\psfig{file=gw_posn_abc+09.eps,width=90mm} \\end{center} \\caption{Localisation of a source of GWs in the northern sky (upper) and in the southern sky (lower) using simulated data for the pulsars observed by the PPTA. The actual assumed source positions are 06$^{\\rm h}$, 45$^{\\rm o}$ and 18$^{\\rm h}$, $-45^{\\rm o}$. (Ref.~\\refcite{abc+09})} \\label{fg:sky_posn} \\end{figure}   Pulsars are extremely stable clocks distributed throughout our Galaxy. The high precision to which it is possible to measure pulse arrival times at the Earth and the long data spans now available leads to many interesting applications. For example, pulsar timing has given the first observational evidence for the existence of gravitational waves and shown that Einstein's general theory of relativity is accurate in the regime of strong gravitational fields. With observations of many pulsars distributed across the celestial sphere -- a ``Pulsar Timing Array'' -- we can in principle make a positive detection of gravitational waves from astrophysical sources. Such an array can also be used to establish a pulsar-based timescale and to identify (or limit) errors in Solar-system ephemerides.  We are now approaching the level of timing precision and data spans which are needed to achieve the main goals of PTA projects. Specifically, if current predictions of the strength of the stochastic GW background in the Galaxy from binary super-massive black holes in the cores of distant galaxies are realistic, within a few years we should be able to make a significant detection of this signal. Other possible sources of a stochastic GW background include fluctuations in the inflation era and oscillations of cosmic strings in the early Universe. Fig.~\\ref{fg:gw_spec} summarises the signals expected from different astrophysical sources and places the PTA in the context of ground- and space-based laser interferometer systems, specifically LIGO and LISA.  \\begin{figure}[ht] \\begin{center} \\psfig{file=GW_spectrum_hsens.eps,width=120mm} \\end{center} \\caption{Spectrum of potentially detectable GW sources and sensitivity curves for PTA systems, the space-based laser interferometer LISA and the ground-based laser interferometer LIGO. } \\label{fg:gw_spec} \\end{figure}  Current efforts are concentrating on eliminating systematic errors from the timing data and improving signal-processing and GW-detection algorithms. While intrinsic period irregularities are significant in a few PTA pulsars, so far they are not a limiting factor for GW detection. Several groups around the world are collaborating to combine data sets to form an International Pulsar Timing Array, which will improve our sensitivity to GWs and help us to reach the other PTA goals. Future instruments such as the Square Kilometre Array will greatly enhance the sensitivity of these efforts and will surely allow the detection and detailed study of GWs and their sources, both stochastic and individual. These studies have the potential to give new information on the properties of gravitational waves and gravitational interactions in the near and distant Universe."
    },
    "1004/1004.4487_arXiv.txt": {
        "abstract": "We examine how the uncertainties involved in supernova dynamics as well as in nuclear data inputs affect the $\\nu$p-process in the neutrino-driven winds. For the supernova dynamics, we find that the wind termination by the preceding dense ejecta shell, as well as the electron fraction ($Y_\\mathrm{e, 3}$; at $3\\times 10^9$~K) play a crucial role. A wind termination within the temperature range of $(1.5-3) \\times 10^9$~K greatly enhances the efficiency of the $\\nu $p-process. This implies that the early wind phase, when the innermost layer of the preceding supernova ejecta is still $\\sim 200-1000$~km from the center, is most relevant to the $\\nu$p-process. The outflows with $Y_\\mathrm{e, 3} = 0.52-0.60$ result in the production of the p-nuclei up to $A=108$ with interesting amounts. Furthermore, the p-nuclei up to $A=152$ can be produced if $Y_\\mathrm{e, 3} = 0.65$ is achieved. For the nuclear data inputs, we test the sensitivity to the rates relevant to the breakout from the pp-chain region ($A < 12$), to the $(n, p)$ rates on heavy nuclei, and to the nuclear masses along the $\\nu$p-process pathway. We find that a small variation of the rates of triple-$\\alpha$ and of the $(n, p)$ reaction on $^{56}$Ni leads to a substantial change in the p-nuclei production. We also find that $^{96}$Pd ($N=50$) on the $\\nu$p-process path plays a role as a second seed nucleus for the production of heavier p-nuclei. The uncertainty in the nuclear mass of $^{82}$Zr can lead to a factor of two reduction in the abundance of the p-isotope $^{84}$Sr. ",
        "introduction": "The astrophysical origin of the proton-rich isotopes of heavy elements (p-nuclei) is not fully understood. The most successful model to date, the photo-dissociation of pre-existing neutron-rich isotopes ($\\gamma$-process) in the oxygen-neon layer of core-collapse supernovae (or in their pre-collapse stages), cannot explain the production of some light p-nuclei including $^{92, 94}$Mo and $^{96, 98}$Ru \\citep{Woos1978, Pran1990, Raye1995, Raus2002, Haya2008}. The recent discovery of a new nucleosynthetic process, the $\\nu $p-process, has dramatically changed this difficult situation \\citep{Froe2006, Froe2006b, Prue2006, Wana2006}. In the early neutrino-driven winds of core-collapse supernovae, $\\bar{\\nu}_e$ capture on free protons gives rise to a tiny amount of free neutrons in the proton-rich matter. These neutrons induce the $(n, p)$ reactions on the $\\beta^+$-waiting point nuclei along the classical rp-process path ($^{64}$Ge, $^{68}$Se, and $^{72}$Kr), which bypass these nuclei (with the $\\beta^+$-decay half-lives of 1.06~min, 35.5~s, and 17.1~s, respectively). \\citet{Wana2006} has shown that the p-nuclei up to $A\\sim 110$, including $^{92, 94}$Mo and $^{96, 98}$Ru, can be produced by the $\\nu$p-process in the neutrino-driven winds within reasonable ranges of the model parameters.  All the recent hydrodynamic studies of core-collapse supernovae with neutrino transport taken into account suggest that the bulk of early supernova ejecta is proton rich \\citep{Jank2003, Lieb2003, Bura2006, Kita2006, Fisc2010, Hued2010}. This supports the $\\nu$p-process taking place in the neutrino-driven winds of core-collapse supernovae. However, different works end up with somewhat different outcomes. \\citet{Froe2006} showed that the p-nuclei up to $A\\sim 80$ were produced with the one-dimensional, \\textit{artificially} induced explosion model of a $20 M_\\odot$ star, while \\citet{Prue2006} obtained up to $A\\sim 100$ with the two-dimensional, \\textit{artificially} induced explosion model of a $15 M_\\odot$ star. On the contrary, \\citet[][also S. Wanajo et al., in preparation]{Wana2009} found negligible contribution of the $\\nu$p-process to the production of p-nuclei with the one-dimensional, \\textit{self-consistently} exploding model of a $9 M_\\odot$ star \\citep[electron-capture supernova,][]{Kita2006, Hued2010}. These diverse outcomes indicate that the $\\nu$p-process is highly sensitive to the physical conditions of neutrino-driven winds.  Besides the supernova conditions, there could be also uncertainties in some key nuclear rates, in particular of $(n, p)$ reactions, because no attention was paid to neutron capture reactions on proton-rich nuclei before the discovery of the $\\nu $p-process. Uncertainties in some reactions relevant to the breakout from the pp-chain region ($A < 12$), which affect the proton-to-seed ratio at the onset of $\\nu$p-processing, might also influence the nucleosynthetic outcomes. There are still a number of isotopes without experimental nuclear masses on the $\\nu$p-process pathway \\citep{Webe2008}.  Our goal in this paper is to examine how the variations of supernova conditions as well as of nuclear data inputs influence the global trend of the $\\nu$p-process. The paper is organized as follows. In \\S~2, a basic picture of the $\\nu$p-process is outlined. A semi-analytic neutrino-driven wind model and an up-to-date reaction network code are described, which are used in this study (\\S~3). We take the wind-termination radius (or temperature), the neutrino luminosity, the neutron-star mass, and the electron fraction as the key parameters of supernova conditions (\\S~4). In previous studies \\citep{Froe2006, Froe2006b, Prue2006, Wana2006}, some of these parameters were varied to test their sensitivities, but only for limited cases. In particular, the effect of wind termination has not been discussed at all in previous studies. As the key nuclear reactions, we take triple-$\\alpha$, $^7$Be$(\\alpha, \\gamma)^{11}$C, $^{10}$B$(\\alpha, p)^{13}$C (all relevant to the breakout from the pp-chain region), and the $(n, p)$ reactions on $^{56}$Ni, $^{60}$Zn, and $^{64}$Ge (\\S~5). Sensitivities of the masses of the nuclei along the $\\nu$p-process path are also discussed. We then discuss the possible role of the $\\nu$p-process as the astrophysical origin of the p-nuclei (\\S~6). A summary of our results follows in \\S~7. ",
        "conclusions": "We investigated the effects of uncertainties in supernova dynamics as well as in nuclear data inputs on the $\\nu$p-process in the neutrino-driven outflows of core-collapse supernovae. The former includes the wind-termination radius $r_\\mathrm{wt}$ (or temperature $T_\\mathrm{wt}$), neutrino luminosity $L_\\nu$, neutron-star mass $M_\\mathrm{ns}$, and electron fraction $Y_\\mathrm{e, 9}$ (or $Y_\\mathrm{e, 3}$, at $T_9 = 9$ and 3, respectively). The latter includes the reactions relevant to the breakout from the pp-chain region ($A < 12$), the $(n, p)$ reactions on heavy nuclei ($Z \\ge 56$), and the nuclear masses ($40 \\le Z \\le 50$) on the $\\nu$p-process pathway. Our result is summarized as follows.  1. Wind termination of the neutrino-driven outflow by colliding with the preceding supernova ejecta causes a slowdown of the temperature decrease and thus plays a crucial role on the $\\nu$p-process. The termination within the temperature range of $T_9 = 1.5-3$ (relevant to the $\\nu$p-process) substantially increases the number of neutrons captured by the seed nuclei ($\\Delta_\\mathrm{n}$) and thus enhances efficiency of the p-nuclei production. In the current case, the efficiency is maximal at $T_\\mathrm{wt, 9} = 2.65$ ($r_\\mathrm{wt} = 231$~km for $L_\\nu = 10^{52}$~erg~s$^{-1}$). This implies that the early wind phase with the termination radius close to the proto-neutron star surface is favored for the $\\nu$p-process.  2. A lower $L_\\nu$ (with the other parameters $T_\\mathrm{wt}$, $M_\\mathrm{ns}$, and $Y_\\mathrm{e, 9}$ unchanged) leads to more efficient $\\nu$p-processing. This is due to the larger entropy per nucleon for a lower $L_\\nu$, which increases the proton-to-seed ratio $Y_\\mathrm{p}/Y_\\mathrm{h}$ and thus $\\Delta_\\mathrm{n}$. However, the role of the wind termination is more crucial and thus we presume that the maximum efficiency is obtained during the early phase with $L_\\nu \\sim 10^{52}$~erg~s$^{-1}$.  3. A larger $M_\\mathrm{ns}$ (with the other parameters $T_\\mathrm{wt}$, $L_\\nu$, and $Y_\\mathrm{e, 9}$ unchanged) results in a larger efficiency of the $\\nu$p-process. This is a consequence of the larger entropy per nucleon and the faster expansion of the neutrino-driven outflow for a larger $M_\\mathrm{ns}$, both of which help to increase $Y_\\mathrm{p}/Y_\\mathrm{h}$ and thus $\\Delta_\\mathrm{n}$. This implies that a more massive progenitor is favored for more efficient $\\nu$p-processing, if other parameters remain unchanged. In reality, however, the evolutions of $L_\\nu$, $r_\\mathrm{wt}$, and $Y_\\mathrm{e}$ will be dependent on the progenitor mass, making it difficult to draw definitive conclusions.  4. The $\\nu$p-process is highly sensitive to the electron fraction $Y_\\mathrm{e, 3}$ that controls $Y_\\mathrm{p}/Y_\\mathrm{h}$ at the onset of the $\\nu$p-process and thus $\\Delta_\\mathrm{n}$. An increase of only $\\Delta Y_\\mathrm{e, 3} \\sim 0.03$ results in $\\Delta A_\\mathrm{max} \\sim 10$. The models with $Y_\\mathrm{e, 3} = 0.52-0.60$ (with the other parameters unchanged) produce sufficient amounts of the light p-nuclei up to $A = 108$. Furthermore, the models with $Y_\\mathrm{e, 3} = 0.60-0.65$ produce the p-nuclei up to $A = 152$. Note that this is a combined effect of the high $Y_\\mathrm{e, 3}$ and the wind termination at sufficiently high temperature ($T_\\mathrm{wt, 9} = 2.19$ in the standard model). Our result shows no substantial enhancement of the p-nuclei with $A > 152$, since the nuclear flow reaches the $\\beta$-stability line and enters to the neutron-rich region at $A \\sim 130-160$. This is a consequence that a large $\\Delta_\\mathrm{n}$ leads to the strong $(n, \\gamma)$ flows that compete with those by $(p, \\gamma)$ for $Z > 50$.  5. Variations on the nuclear reactions relevant to the breakout from the pp-chain region ($A < 12$), namely of triple-$\\alpha$, $^7$Be$(\\alpha, \\gamma)^{11}$C, and $^{10}$B$(\\alpha, p)^{13}$C affect the $\\nu $p-process by changing $Y_\\mathrm{p}/Y_\\mathrm{h}$ (and $\\Delta_\\mathrm{n}$) or producing intermediate-mass nuclei (proton poison) during $\\nu $p-processing. Among these reactions, triple-$\\alpha$ has the largest impact, although the other two show non-negligible effects, on the production of the p-nuclei at $A \\sim 100-110$.  6. Variations on the $(n, p)$ reactions on $^{56}$Ni (seed nuclei), $^{60}$Zn, and $^{64}$Ge (first $\\beta^+$-waiting point on the classical rp-process) show great impact on efficiency of the $\\nu $p-process for heavy element synthesis. Only a factor of two variation leads to a factor of 10 or more changes in the production of the p-nuclei with $A \\sim 100-110$ for the first reaction (but somewhat smaller changes for the latter two reactions). This is a consequence that these reactions control the strength of the nuclear flow passing through the $(n, p)$-waiting points ($^{56}$Ni, $^{60}$Zn, and $^{64}$Ge) on the $\\nu$p-process path. We also find that the $N=50$ nucleus $^{96}$Pd plays a role of the ``second seed'' for production of heavier nuclei.  7. Application of the new experimental masses of \\citet[][for $39 \\le Z \\le 46$]{Webe2008} exhibits a suppression of the flow $^{87}$Mo$(p, \\gamma)^{88}$Tc ($N=45$), which however do not affect the nucleosynthetic p-abundances. Our sensitivity tests for unmeasured nuclear masses indicate that a future measurement of the $^{82}$Zr mass (and of $^{93}$Nb with a large estimated error) on $N=42$ could reduce the abundance of $^{84}$Sr by a factor of two.  8. Our result implies that, within possible ranges of uncertainties in supernova dynamics as well as in nuclear data inputs, the solar inventory of the light p-nuclei up to $A = 108$ ($^{108}$Cd) can be attributed to the $\\nu$p-process, including the most mysterious ones $^{92, 94}$Mo and $^{96, 98}$Ru. The molybdenum isotopes are, however, tend to be underproduced compared to the neighboring p-isotopes. If highly proton-rich conditions with $Y_\\mathrm{e, 3} = 0.60-0.65$ are realized in neutrino-driven ejecta, the solar origin of the p-nuclei up to $A = 152$ ($^{152}$Gd) can be explained by the $\\nu$p-process.  Our explorations in this study suggest that more refinements both in supernova conditions and in nuclear data inputs are needed to elucidate the role of the $\\nu$p-process as the astrophysical origin of the p-nuclei. In particular, multi-dimensional studies of core-collapse simulations with full neutrino transport, as well as experiment-based rates of triple-$\\alpha$ and the $(n, p)$ reactions on heavy nuclei will be important in the future works."
    },
    "1004/1004.3869_arXiv.txt": {
        "abstract": "The presence of Dark Matter (DM) is required in the universe regulated by the standard general relativistic theory of gravitation. The nature of DM is however still elusive to any experimental search. We discuss here the process of accumulation of evidence for the presence of DM in the universe, the astrophysical probes for the leading DM scenarios that can be obtained through a multi-frequency analysis of cosmic structures on large scales, and the strategies related to the multi-messenger and multi-experiment astrophysical search for the nature of the DM. ",
        "introduction": "\\label{sect:intro}  There is overwhelming evidence that we live in a geometrically flat ($\\Omega_0 \\approx 1$) universe which is dominated - in the standard general relativistic cosmological paradigm -  by a dark form of matter (Dark Matter - DM) and by an obscure form of energy (Dark Energy - DE). In such standard cosmological paradigm, DM provides a substantial fraction $\\Omega_{DM} \\approx 0.227$ of the overall matter-energy content (the rest being provided by Dark Energy with $\\Omega_{DE} \\approx 0.728$ with the baryonic contribution limited to $\\Omega_b \\approx 0.045$, see, e.g., Komatsu et al. 2010) and amounts to $\\sim 83 \\%$ of the total mass content of the universe.\\\\ How did we arrive to such a conclusion? Is there a general and definite consensus on it?   \\subsection{The Dark Matter timeline}  The discovery of the presence of DM can be considered as a typical scientific revolution (Kuhn, 1970) that induced a change in the reference (cosmological) paradigm. However, as often occurs in a paradigm shift, there was no single discovery, but new concepts were developed and integrated step-by-step \\footnote{In this context, it is necessary, since the beginning of this discussion, to acknowledge many reviews on the Dark Matter and related problems, as given by Faber \\& Gallagher (1979), Trimble (1987), Kormendy \\& Knapp (1987), Srednicki (1990), Turner (1991), Silk (1992), Ashman (1992), van den Bergh (2001), Ostriker \\& Steinhardt (2003), Rees (2003), Turner (2003), Colafrancesco (2006, 2007), Einasto (2004, 2009), among others.}.  To begin, there are actually two dark matter problems, the {\\it local} Dark Matter close to the plane of our Galaxy, and the {\\it global} Dark Matter surrounding galaxies, clusters of galaxies and large scale structures.  The local Dark Matter (Oort 1932) in the Galactic disk is baryonic [faint stars or planet (jupiter)-like objects], since a collection of matter close to the galactic plane is possible if it has formed by the contraction of pre-stellar matter towards the plane accompanied by dissipation of the extra energy, so to conserve the flat shape of the population. It is now clear that the amount of local Dark Matter is low, and it depends on the mass boundary between luminous stars and faint invisible stars or planet-like objects (see Einasto 2009 for a discussion and references therein).  The global Dark Matter (to which we refer in this paper) is the dominating mass component in the universe; it is mainly concentrated in galaxies, clusters and superclusters of galaxies, and populates all other large-scale structures in the universe.\\\\ The first evidence for such global DM was obtained by F. Zwicky (1933). From observations of the radial velocities of eight galaxies in the Coma cluster, Zwicky found an unexpectedly large velocity dispersion $\\sigma_v = 1019  \\pm 360$ km s$^{-1}$ [we note in passing that Zwicky's velocity dispersion from only eight galaxies agrees well with the modern value $\\sigma_v = 1082$ km s$^{-1}$, as obtained by, e.g., Colless \\& Dunn (1996)]. Zwicky concluded from these observations that, for a velocity dispersion of $\\sim 1000$ km s$^{-1}$, the mean density of the Coma cluster would have to be $\\sim 400$ times greater than that which is derived from luminous matter. \\footnote{Zwicky actually overestimated the mass-to-light ratio of the Coma cluster because he assumed a Hubble parameter $H_0 = 558 km s^{-1} Mpc^{-1}$. His value for the overdensity of Coma should therefore be reduced from 400 to $\\sim$ 50 by using the actual value of the Hubble parameter. It is interesting to note that Hubble's prestige was so great at that time, that none of the early authors thought of reducing Hubble's constant as a way of lowering their mass-to-light ratios.} Zwicky concluded, therefore, that in order to hold galaxies together in the cluster, the Coma cluster must contain huge amounts of some dark (i.e. not visible), cold matter (Zwicky used actually the words {\\it \"Dunkle Kalte Materie\"} which might be regarded as the first reference to Cold Dark Matter ... even though not in the modern sense).\\\\ Not until three years later it was found (Smith 1936) that the Virgo cluster also appears to exhibit an unexpectedly high mass. Smith made the speculation that the excess mass of Virgo \"{\\it represents a great mass of internebular material within the cluster}\" (see van den Bergh 2001).  Six years later than Zwicky's 1933 paper, Babcock (1939) obtained spectra of the Andromeda galaxy (M31) and found that in its outer regions the galaxy is rotating with an unexpectedly high velocity, far above the expected Keplerian velocity. He interpreted this result either as a high mass-to-light ratio in the periphery or as a strong dust absorption. Dark Matter in galaxies was also envisaged.\\\\ One year later, Oort (1940) studied the rotation and the surface brightness of the edge-on S0 galaxy NGC 3115. He found that {\\it \"The distribution of mass in this system appears to bear almost no relation to that of light.\"} He concluded that a value $M/L \\sim 250$ should be present in the outer regions of NGC 3115 (note that this value is reduced by almost an order of magnitude if the modern distance to this galaxy is adopted). Oort ended his paper by writing that {\\it \"There cannot be any doubt that an extension of the measures of rotation to greater distances from the nucleus would be of exceptional interest.\"} However, no connection was made between the missing mass in this S0 galaxy, and the Zwicky and Smith missing mass problem in clusters of galaxies.  The DM problem remained as a kind of \"anomaly\" for roughly a quarter of century after the initial Zwicky's paper until when Kahn \\& Woltjer (1959) pointed out that M31 and our Galaxy were moving towards each other, so that they must have completed most of a (very elongated) orbit around each other in a Hubble time. These authors found that if M31 and our Galaxy started to move apart $\\sim 15$ Gyr ago, the mass of the Local Group had to be $\\simgt 1.8 \\times 10^{12} M_{\\odot}$. Assuming that the combined mass of the Andromeda galaxy and the Milky Way system was $0.5 \\times 10^{12} M_{\\odot}$, Kahn \\& Woltjer (1959) concluded that most of the mass of the Local Group existed in some invisible form. They opined that it was most likely that this missing mass was in the form of hot (with temperature of $\\sim 5 \\times 10^5$ K) gas. From a historical perspective, it is interesting to note that also Kahn \\& Woltjer did not seem to have been aware of the earlier papers by Zwicky (1933) and Smith (1936) on the missing mass problem in clusters of galaxies.  In the 1960's and 1970's, Babcock's optical rotation curve of M31, and that of Rubin \\& Ford (1970), were extended to even larger radii by Roberts \\& Whitehurst (1975) using 21-cm line observations that reached a radial distance of $\\sim 30$ kpc. These observations clearly showed that the rotation curve of M31 did not exhibit any Keplerian drop-off. From these observations Roberts \\& Whitehurst (1975) concluded that the mass-to-light ratio had to be $\\simgt 200$ in the outermost regions of M31. Again, it is interesting to note that neither Babcock, nor Roberts \\& Whitehurst, cited the 1933 paper by Zwicky. In other words, no connection was made between the missing mass in the outer region of spiral galaxies and the missing mass in galaxy clusters such as Coma (Zwicky 1933) and Virgo (Smith 1936). Regarding the flat outer rotation curves of galaxies Roberts (1999, as quoted in van den Bergh 2001) recalls that this result {\\it \"was, at best, received with skepticism in many colloquia and meeting presentations.\"} Nonetheless, the paper by Roberts \\& Whitehurst (1975) was important because it, together with papers on the stability of galactic disks (e.g. Ostriker \\& Peebles 1973), and on the apparent increase of galaxy mass with increasing radius (e.g. Ostriker, Peebles \\& Yahil 1973), first convinced the majority of astronomers that a missing mass problem existed.  By the mid of 1970's enough research was done for the community to see that the missing mass ''anomaly'' was not going to go away, and the majority of astronomers had become convinced that missing mass existed in cosmologically significant amounts. {\\it ''When... an anomaly comes to seem more than just another puzzle of normal science, the transition to crisis and to extraordinary science has begun''} (Kuhn 1970). In other words, it began to be clear in the mid 1970's that a paradigm shift would be required to interpret the observations that seemed to support the ubiquitousness of missing (dark) matter in the universe.  At the beginning of the 1980's the presence of DM was directly confirmed by many independent sources: the dynamics of galaxies and of stars within galaxies (Rubin et al. 1978, 1980), the mass determinations of galaxy clusters based on gravitational lensing \\footnote{It is noted, in passing, that Zwicky (1937) also pointed out that gravitational lensing might provide useful information on the total masses of galaxies.} (see, e.g., Bartelmann \\& Schneider 2001 for a review and references therein), X-ray studies of clusters of galaxies (see, e.g., Forman et al. 1985), the indications of the first N-body simulations of large scale structure formation and of galaxies and cluster formation (see, e.g., White et al. 1987).\\\\ At this time it was also already clear that structures in the universe form by gravitational clustering, started from initially small fluctuations of the density of matter (Peebles 1980). In order to form the presently observed structures, the amplitude of the density fluctuations must be at least one thousandth of the density itself at the epoch of recombination, when the universe started to be transparent. The relic emission coming from this epoch was first detected in 1965 (Penzias \\& Wilson 1965) as an (almost) isotropic cosmic microwave background (CMB). When finally the fluctuations (anisotropies) of the CMB were measured by the COBE satellite, first an large angular scales (Smooth et al. 1992), and by the Boomerang experiment, also on smaller angular scales (de Bernardis et al. 2000), they appeared to be two orders of magnitude lower than expected from the density evolution of the luminous (baryonic) mass. The solution of the problem was suggested independently by several theorists. If we suppose that the dominating mass component of the universe - Dark Matter - is not made of ordinary matter but of some sort of non-baryonic matter, then density fluctuations can start to grow much earlier, and have at the time of recombination the amplitudes needed to form the structures observed today. The interaction of this non-baryonic matter with radiation is much weaker than that of ordinary matter, and radiation pressure does not slow the early growth of fluctuations.  However, beyond the indubitable success of the DM-dominated cosmological framework, the fundamental uncertainty on the DM nature remained.\\\\ The first suggestion for a non-baryonic DM was made referring to particles well known at that time to physicists, i.e. neutrinos ($\\nu$) (see, e.g., Pontecorvo 1967), for which ''theoretical and practical considerations'' push progressively some physicists to agree on a non zero mass neutrino (Quantum mechanics allow then the oscillation of neutrinos: a $\\nu_e$ can, along its travel in the universe, become a $\\nu_{\\mu}$ and vice-versa). This can be considered as the first input from particle physics to the search for the nature of DM. Rees (1977) was one of the first astrophysicists to tackle with the idea of non-baryonic nature for the cosmologically relevant DM, and suggested that: {\\it ''There are other possibilities of more exotic character - for instance the idea of neutrinos with small ($\\sim$ few eV) rest mass has been taken surprisingly seriously by some authors''.} However, this scenario of hot DM soon led to major problems, because neutrinos move with very high velocities which prevents the formation of small structures as galaxies (Silk 1968).  Thus some other hypothetical non-baryonic particles were suggested, such as axions (Abbott \\& Sikivie 1983, Preskill, Wise \\& Wilczek, 1983; Dine \\& Fischler, 1983) or Weakly Interacting Massive Particles (e.g. WIMPs, Goldberg 1983, Ellis et al. 1983). The essential character of these particles to make them good for cosmology is that they have much lower velocities than neutrinos. Because of this, the new version of Dark Matter scenario was called Cold, in contrast to neutrino-dominated Hot Dark Matter scenario.\\\\ Numerical simulations of the evolution of the structure of the universe confirmed the formation of filamentary superclusters and voids in the Cold Dark Matter scenario of structure formation (e.g. Ostriker 1993 and references therein). The suggestion of the Cold Dark Matter has solved most problems of the new cosmological paradigm.  One unsolved problem remained. Estimates of the total matter density (ordinary plus DM) yield values $\\Omega_m \\sim 0.27$ of the critical density. This value - not far from unity, but definitely smaller than unity - is neither favoured by theory nor by the data, including the measurements of the cosmic microwave background, the galaxy dynamics and the expansion rate of the universe obtained from the study of distant type Ia supernovae (SNe). To fill the matter/energy density gap between unity and the observed matter density it was assumed that some sort of Dark Energy (DE) exists. \\footnote{This assumption was not new: already Einstein added to his cosmological equations a cosmological constant term $\\Lambda$ (corresponding to a vacuum energy).} By the early 2000's, refined CMB anisotropy experiments (e.g., Boomerang, Maxima, DASI, and then WMAP) demonstrated that the CMB anisotropy spectrum is the largest detector for the presence of Dark Matter in their fluctuation spectrum, and by combining their measurements with independent cosmological distance measurement using type Ia SNe (Perlmutter 2000), a first determination of the overall matter-energy composition of the universe was possible (see Komatsu et al. 2010 for a recent determination of the cosmological parameter set). The inclusion of the DE term in the general relativistic cosmological scenario has filled the last gap in the modern cosmological paradigm.  \\subsection{False alarms and diversionary manoeuvres}  No good detective story is complete without at least one false clue or some diversionary manoeuvres.  Oort (1960, 1965) believed that he had found some dynamical evidence for the presence of missing mass in the disk of the Galaxy. If true, this would have indicated that some of the dark matter was dissipative in nature. However, late in his life, Jan Oort confessed (as reported by van den Bergh 2001) that the existence of missing mass in the Galactic plane was never one of his most firmly held scientific beliefs. Detailed observations, that have been reviewed by Tinney (1999), show that brown dwarfs cannot make a significant contribution to the density of the Galactic disk near the Sun.  The presence of large amounts of matter of unknown origin has given rise to speculations on the validity of the Newton law of gravity at very large distances. One of such attempts is the Modified Newton Dynamics (MOND) model, suggested by Milgrom (1983) and Milgrom \\& Bekenstein (1987) (see, e.g., Sanders 1990 for a discussion). Indeed, MOND is able to explain spiral galaxies quite well imposing a minimum acceleration scale $a_0 \\sim c H_0$, without assuming the presence of some hidden matter. If the MOND scale $a_0$ is allowed to run, galaxy clusters might be explained as well. However, there exist several arguments which make this model unrealistic (see Einasto 2009 for a discussion). In addition, a full relativistic theory is needed in any case to construct a consistent cosmology. A tensor-vector-scalar theory (TeVeS, see Bekenstein 2004) has been then proposed and it is successful in reproducing MOND in the proper limit even for photons, and respecting also the classical tests of general relativity (GR).\\\\ Whether a consistent cosmology can be constructed with this theory remains to be seen (in fact, both MOND and TEVES models are not covariant formalisms of a general relativistic theory of gravitation, see Capozziello et al. 2008).  Extended Gravity scenarios have been therefore suggested to explain the large-scale cosmological problem of the accelerated expansion of the universe and have been also worked out to describe the morphology and dynamics of very peculiar systems, like the bullet cluster (Brownstein \\& Moffat 2007). Such a scenario is today the best alternative to a pure particle DM scenario.   \\subsection{The Astro-Particle connection}  All the available information indicates that the standard scenario for structure formation in the GR framework requires that the global Cold Dark Matter must be non-baryonic, its density fluctuations start to grow much earlier than perturbations in the baryonic matter, and have at the recombination epoch large enough amplitudes to form all structures seen in the universe.\\\\ However, the actual nature of the CDM particles still remains unknown. Physicists have attempted to discover particles which have the properties needed to explain the structure of the universe, but so far without success. This means that a true Astro-Particle connection should be developed in the search for the nature of DM. In conclusion, even though the direct information on the dark components of the universe (DM and DE) comes solely from astronomical observations, it is clear that a definite understanding of the nature of DM will come through the discovery and the multi-messenger study of the fundamental particles of which DM consists of.  \\subsection{A last remark}  The discovery of DM has the general character of a typical scientific revolution connected with changes of paradigms, as discussed by Kuhn (1970) in his book \"The Structure of Scientific Revolutions\". There are not so many areas in modern astronomy and cosmology where the development of ideas can be described in these terms. ",
        "conclusions": ""
    },
    "1004/1004.5448_arXiv.txt": {
        "abstract": "We present CCD photometric observations of the W UMa type contact binary EK Comae Berenices using  the 2 metre telescope of $IUCAA$ Girawali Observatory, India. The star was classified as a W UMa type binary of subtype-W by \\citet{sam1996}. The new V band photometric observations of the star reveal that shape of the light curve has changed significantly from the one observed by \\citet{sam1996}. A detailed analysis of the  light curve obtained from the high-precision CCD photometric observations of the star indicates that EK Comae Berenices is not a W-type but an A-type totally eclipsing W UMa contact binary. The photometric mass ratio is determined to be 0.349 $\\pm$ 0.005. A temperature difference of $\\Delta T = 141 \\pm 10 $ K between the components and an orbital inclination of $i [^{o}] = 89.800 \\pm 0.075$ were obtained for the binary system. Absolute values of masses, radii and luminosities are estimated by means of the standard mass-luminosity relation for zero age main-sequence stars. The star shows O'Connell effect, asymmetries in the light curve shape around the primary and secondary maximum. The observed O'Connell effect is explained by the presence of a hot spot on the primary component. ",
        "introduction": "The eclipsing binary EK Comae Berenices [hereafter EK Com] was discovered by \\citet{kin1966}.~They classified the star as a W UMa type binary.~The star was also observed visually by \\citet{bor1990} who determined the period of $0.2666864$ days.~The first photometric solution of EK Com was obtained by \\citet [hereafter SA96]{sam1996} based on the observations of 1994 with the help of the Wilson-Devinney [hereafter WD] code. These authors determined a period of 0.2668726 days and derived other photometric elements from their observed light curves and found the evidence for a period decrease at a rate of $\\sim 1.5 \\times 10^{-10} \\rm cycle^{-1}$. From the light curve solutions using the WD code, they classified the star as a W UMa type binary of subtype W with a mass ratio $q = 0.304$, a temperature difference of $\\Delta T = 310$ K between the components and a fill-out factor of 15\\%. They estimated the temperatures of the components as $T_{1} = 5000$ K and $T_{2} = 5310$ K respectively, where the subscripts 1 and 2 refer to the more massive (primary) and the less massive (secondary) component respectively. The change in the orbital period of EK Com has been analysed by \\citet{liz2006} and found the evidence of continuous decrease of orbital period $dP/dt = -5.96~\\times 10^{-10}$ days cycle$^{-1} = -0.071$ s yr$^{-1}$.  Since there is no photometric measurement of the star available since 1994, we have obtained a complete and well-defined light curve of EK Com through high-precision CCD photometry using the $IUCAA$ Girawali Observatory (IGO) 2-m telescope. The newly obtained light curve is quite different in shape than that obtained by SA96. SA96 classified it as a W UMa binary of subtype-W by pointing out that the light curve showed an interval of constant light in the primary eclipse. However, it is difficult to make out whether a minimum is curved or flat and SA96 modelled the light curve assuming a flat-bottom primary minimum. Using the new observations, we are able to present highly accurate and more precise light curve of EK Com than SA96. Based on the appearance of the eclipses in our data, we conclude that the system is clearly an A-type W UMa binary, i.e, the deeper eclipse is the transit of the more massive star by the less massive one. Our new photometric observations, in fact, clearly suggest that the primary minimum is rounded and the secondary minimum is flat indicating its A-type nature.  In this paper, we report the complete V band light curve analysis of the star EK Com from new photometric observations and employ the WD light curve modelling technique.~This technique is used extensively to characterize the observed photometric light curves of eclipsing binary stars with the computed synthetic light curves thus yielding the parameters of components of the binary system. It has been seen from the literature that determination of photometric mass ratio $q$ obtained from the WD light curve modelling is reliable only for a total-eclipse configuration of the binary system as such light curves show characteristic inner contacts with duration of totality setting a strong constraint on the $(q, i)$ pair \\citep{ruc2001}. Since EK Com is found to be a total-eclipsing over-contact binary, the various geometrical and physical parameters obtained from the light curve modelling are expected to be reliable. Although we have only two nights of data observed only in the V band, the phase coverage is excellent enough to reveal the accurate shapes of the primary and secondary minima. In Section 2, we present the observations and data reduction procedures. Section 3 deals with the period analysis and determination of ephemeris from new observations. In Section 4, we describe detailed photometric analysis of the star in the V band using the WD code. In Section 5, we describe determination of distance to EK Com. Finally, in Section 6, we summarise the results and main conclusions of our study. ",
        "conclusions": "We have presented a photometric solution for the eclipsing binary EK Com obtained from the new high-precision time-series CCD photometric observations with complete phase coverage in the V band. The new observations indicate that the star may be of A-type with a flat secondary minimum. By using the WD code, we have found that the system is a total-eclipsing binary star with an orbital inclination of $i[^{o}] = 89.800 \\pm 0.075$ and a high degree of over-contact configuration 33.0\\%. The difference in the mean temperatures between the two components is found to be $141 \\pm 10~K$. The observed light curve is quite asymmetric, especially at quadrature levels, because of surface activity on either of the two components. This feature of the light curve is explained assuming a hot spot on the  primary component.  It has been found that the photometric mass ratio determined from the WD light curve modelling approach is reliable for a totally eclipsing over-contact binary. Keeping in view of this fact, the absolute physical parameters of the two components are determined based on the results of the light curve solution and by the MV96 method. The luminosity of the star is calculated using the temperature and radius of each of the individual components. The photometric spot solutions given in Table~\\ref{table3} are tentative, as it has been observed that with the spot solutions included in the WD code, the whole light curve can be easily fitted, but a serious problem of uniqueness of the solution persists \\citep{mac1993} unless other means of investigation such as Doppler Imaging techniques are applied \\citep{mac1994}.  The spectroscopic radial velocity observations and long term photometric monitoring of the star is important to determine its actual type and to answer whether the deeper rounded shape of the primary eclipse seen in our observations is either due to the actual configuration of the system or due to its higher level of spot magnetic activity at different phases or both."
    },
    "1004/1004.2929_arXiv.txt": {
        "abstract": "We investigate the effect of long-range scalar interactions in dark matter (DM) models of cosmic structure formation with a particular focus on the formation times of haloes. Utilising $N$-body simulations with $512^3$ DM particles we show that in our models dark matter haloes form substantially earlier: tracing objects up to redshift $z\\sim6$ we find that the formation time, as characterised by the redshift $z_{1/2}$ at which the halo has assembled half of its final mass, is gradually shifted from $z_{1/2}\\approx 1.83$ in the fiducial \\lcdm\\ model to $z_{1/2}\\approx 2.54$ in the most extreme scalar-interaction model. This is accompanied by a shift of the redshift that marks the transition between merger and steady accretion epochs from $z_{*}\\approx 4.32$ in the \\lcdm\\ halos to $z_{*}\\approx 6.39$ in our strongest interaction model. In other words, the scalar-interacting model employed in this work produces more structures at high redshifts, prolonging at the same time the steady accretion phases. These effects taken together can help the \\lcdm\\ model to account for a high redshift reionisation as indicated by the WMAP data and can alleviate issues related to the survival of the thin-disk dominated galaxies at low redshifts. ",
        "introduction": "The $\\Lambda$ cold dark matter (\\lcdm) model has proven to be capable of explaining a tremendous amount of observational data. As we entered the era of precision cosmology --- both observationally and from a modellers perspective --- we are left with a more and more detailed picture of what actually took place when the Universe cooled down and formed structures starting from $z\\sim1100$ to the present day. But there are still issues related to the structure formation on galactic scales to be resolved. For instance, observations of high redshift quasars \\citep{quazars_z6_reionisation} and galaxies \\citep{early_reionisation} imply that the epoch of the reionisation of the Universe ended at $z\\sim 6$.  On the other hand observations of the polarisation of the cosmic microwave background radiation \\citep{WMAP07,reion2,reion3} suggest that the Universe was already reionised at $z{\\stackrel{>}{{}_\\sim}}10-11$.  However, galaxy formation within the \\lcdm\\ model using the suggested moderate normalisation of the power spectrum of primordial density perturbations $\\sigma_8=0.8$ \\citep[as favoured by recent WMAP observations; ][]{WMAP07} may not happen early enough to account for such a high redshift of reionisation.  Furthermore the astonishing array of recent observations of ultra-faint distant objects \\citep{high_z_evolution_LF,z_10_galaxies} as well as advancements in the understanding of the high-$z$ physics \\citep{star_formation_z8,supermassive_BH_growth1,supermassive_BH_growth2} have shown that the young Universe was a busy and already structure-rich place.  In addition, there are issues related to the survival of thin-disk dominated galaxies.  It was shown that in a pure CDM model of galaxy-halo formation, the majority of Milky-Way-sized haloes with $M_0\\sim10^{12}h^{-1}M_{\\odot}$ have accreted at least one object with greater total mass than the Milky Way disk ($m>5\\times 10^{10}h^{-1}M_{\\odot}$) over the redshifts $z\\leq 3$ \\citep{thin_disk_prob,disk_surv}.  It has been shown recently \\citep{disk_surv2, disk_surv3}, that the disk may reform after mergers or even survive minor mergers, somewhat easing the problem of disk survival.  Still, late-time bombardment of Milky-Way like haloes, predicted within the \\lcdm\\ model, could still be considered as a concern about the survival of thin-disk-dominated galaxies and the existence of galaxies with unusually quiet merger history like our Galaxy. Moreover, recent observations of the distant Universe suggest that much of the stellar mass of bright galaxies was already present in place at $z>1$, this presents the ``downsizing'' challenge for models of galaxy formation because massive haloes are assembled late in the hierarchical clustering process intrinsic to the CDM cosmology \\citep{downsizing,downsizing_problem, hierarchy_problem_star_mass_z_gt_1}. \\sout{Finally} In addition there were also reports on massive superclusters found at intermediate redshifts ($z\\sim1$), which seem to be very rare and extreme objects within standard \\lcdm\\ model \\citep{2sFGRS_superclusters_problem, 2dFGRS_final_data, UKIDSS_z09_supercluster}. Finally there is a recent claim of possible observational evidence for a ``fifth force'' in DM sector \\citep{lee}.  Even though we would not go as far as to call this a ``galaxy formation crisis'', there are noticeable tensions between what observations suggest and what (cosmological simulations of) the \\lcdm\\ model seems to predict. The ReBEL (\\textit{daRk Breaking Equivalence principLe}) model discussed in recent years in the literature \\citep{GP1,GP2,PeeblesIDMDE,Farrar2007,rebel_nature} has the potential to alleviate these tensions and to help the theoretical \\lcdm\\ model to better match the observations. The ReBEL model is a slight modification of the standard cold dark matter model. It introduces an additional long-range scalar mediated force only in the dark sector. This has been studied in numerical simulations \\citep{NGP,LRSI1,LRSI2,Rebel,LRSI3} and these studies have shown that the model can match observational properties of the large-scale structures (like the one-dimensional Ly-$\\alpha$ power spectrum and the power spectrum of the SDSS\\footnote{Sloan Digital Sky Survey - \\url{http://www.sdss.org/}} galaxies) as well as the standard \\lcdm\\ does, providing at the same time interesting features with respects to the formation of galactic structures. We also acknowledge that recently other authors performed structure formation tests in similar scalar-interacting DM scenarios \\citep{Coupled_DE1,Coupled_DE2,Coupled_DE3,Coupled_DE4}, but they focused on models containing scalar-field-like Dark Energy interacting with DM. Elementary considerations related to the ReBEL model and the reionisation mechanism showed that scalar-interacting DM could accommodate early reionisation of the intergalactic medium \\citep{reion1}. The low-resolution simulations presented in the ReBEL literature have indicated that this model can foster a picture in which structure formation starts at higher redshifts. In this letter we will show that this is indeed the case. We will present, for the first time in the literature, results of analysis of the high-resolution N-body simulations of hierarchical DM halo formation in the ReBEL framework. ",
        "conclusions": "The main conclusion of this \\textit{Letter} is that a slight modification to the CDM model in the form of scalar interactions can accommodate more gravitationally bound structures at high redshifts (e.g. fig.~\\ref{fig:cmfs}) which in turn promotes an early reionisation of the Universe. This can be expressed by examining the $\\eta$ factor that describes the efficiency of reionisation of the Inter Galactic Medium (IGM) \\citep{reion1,Cen2003}: \\beq \\label{eqn:eta_factor} \\eta\\equiv{c_{\\star}f_{esc}(\\textrm{d}f_M/\\textrm{d}t)e_{UV}\\over C(1+z)^3}\\,, \\eeq where $c_{\\star}$ is the star formation factor efficiency, $f_{esc}$ is the ionising photon escape fraction, d$f_M/$d$t$ is the halo formation rate, $e_{UV}$ is the ionising photon production efficiency and $C$ is the gas clumping factor. Here the numerator reflects the ionising photon production rate while the denominator is the ionising photon destruction rate. When $\\eta$ exceeds a certain threshold, the IGM becomes ionised. As the ReBEL models exhibit a richer structure at $z=6$, their formation rate must have been higher than the one of the \\emph{\\lcdm} run.  This would make it easier to reach the threshold of reionisation at high reshifts in the ReBEL framework.  We further have to acknowledge that earlier formation of the DM haloes fostered by the ReBEL model, results in shifting a significant part of the violent merger events to higher redshifts. This indicates that there are many more haloes with extended steady accretion histories than expected in the standard \\lcdm\\ structure formation scenarios. This would help thin-disk dominated galaxies to preserve their disks at low redshifts and would also imply that the exceptionally long period of peace experienced by the Milky-Way is no more so extraordinary \\citep{Milky-Way}.  Finally, we have also shown that long-range scalar interactions between DM particles increase to some extent the masses of DM haloes. That is also an important effect that may help to understand the existence of objects like massive superclusters at moderate redshifts. The simulations presented in this \\textit{Letter} however do not encompass a large enough volume to study the formation of these massive objects in detail.  In conclusion, we have shown that the ReBEL model allows to alleviate some of the tension between predictions of the \\lcdm\\ and observations \\citep{reion2,reion3,thin_disk_prob,disk_surv,2sFGRS_superclusters_problem}, however, as we only focused on the issue of formation time, it still needs to be investigated if and how internal properties of dark matter haloes are affected by the ReBEL formalism.  One interesting quantity is the halo concentration parameter, as it correlated with the halo formation time \\citep[eg.][]{Wechsler2002}.   Preliminary work shows this to indeed be the case, with the  mean concentration $\\langle c_{vir}\\rangle$ being $9$ ($11$, $11$, $13$) for \\lcdm\\ (\\textit{b05rs05}, \\textit{b05rs1}, \\textit{b1rs1}) for haloes with masses $M\\sim2\\times 10^{12} h^{-1}M_{\\odot}$, respectively.  However the issue of the halo concentration parameters require deeper studies and, since a more detailed study of the haloes internal properties is in progress, we will cover the concentration parameter issue in a forth-coming work."
    },
    "1004/1004.0437_arXiv.txt": {
        "abstract": "We study a class of early dark energy (EDE) models, in which, unlike in standard dark energy models, a substantial amount of dark energy exists in the matter-dominated era. We self-consistently include dark energy perturbations, and show that these models may be successfully constrained using future observations of galaxy clusters, in particular the redshift abundance, and the Sunyaev-Zel'dovich (SZ) power spectrum. We make predictions for EDE models, as well as $\\ld$CDM for incoming X-ray (eROSITA) and microwave (South Pole Telescope) observations.  We show that galaxy clusters' mass function and the SZ power spectrum will put strong constraints both on the equation of state of \\de~today and the redshift at which EDE transits to present-day $\\ld$CDM like behavior for these models, thus providing complementary information to the geometric probes of dark energy.  Not including perturbations in EDE models leads to those models being practically indistinguishable from $\\ld$CDM. An MCMC analysis of future galaxy cluster surveys provides constraints for EDE parameters that are competitive with and complementary to background expansion observations such as supernovae. ",
        "introduction": "Over the last decade, observational evidence has mounted in favor of \\de, the mysterious component which dominates the energy content of the universe at present, and causes the expansion of the universe to accelerate \\citep{sne, wmap7}.  The nature of \\de~is one of the most tantalizing mysteries of present day cosmology. The simplest model for \\de, the cosmological constant, fits the current data well \\citep{sne}, however, there are no strong constraints on the time evolution of \\de~at present. Thus, evolving models of \\de~remain as alternative candidates for dark energy. Many non-cosmological constant models have been suggested for \\de, including scalar field quintessence models, modifications of the Einstein framework of gravity, \\etc \\citep[see reviews][and references therein]{rev1, rev2, rev3, rev4, rev5, rev6, rev7}, but as of now there is no clear consensus on the nature on \\de.  An interesting class of models which have been suggested in the literature are early dark energy (EDE) models, in which the early universe contained a substantial amount of dark energy. Geometric probes of \\de~put strong constraints on the present-day nature of \\de, and these constraints are expected to improve with future surveys. However, very little is known as to the nature of \\de~at early times due to paucity of data beyond redshifts of few. It has thus been hypothesized that even if \\de~at present behaves like the cosmological constant, in the past it could have had completely different behavior, leading to the idea of the EDE models. Different facets of these models have been studied in recent works \\citep{ede1,ede2,ede3} and references therein, and have been analyzed with respect to observations extensively in recent times in \\citep{ede4,ede5,ede6,ede7,ede8, ede9, ede_cmb, ede10}. Current data places some constraints on these models, but does not rule them out. In this work we use a parameterization of the equation of state of \\de~to study the possibility of constraining EDE models using future large scale structure surveys.  Ground and space-based telescopes targeting clusters in microwave [Atacama Cosmology Telescope (ACT) \\citep{act}, the South Pole Telescope (SPT) \\citep{spt} and Planck \\footnote{http://www.rssd.esa.int/index.php?project=Planck}] have begun operation, or will return data shortly. Also current (Ebeling et al.~2010) and future missions have been planned for detecting clusters in X-ray waveband (Predehl et al.~2007). It has been recognized for some time now, that galaxy cluster surveys in X-ray or microwave via Sunyaev-Zel'dovich (SZ) effect could be precision probes of cosmological parameters, in particular the dark energy density and its equation of state parameter \\citep{haiman2001}. This is because redshift distribution of massive clusters is exponentially sensitive to the growth of structure history of the universe, which in turn, bears the signature of dark energy. The presence of early dark energy reduce structure formation, consequently the cluster abundances reduce. Thus cluster counts can provide a smoking gun probe for detecting early dark energy. For example, \\cite{mantz09a} have used a flux limited X-ray cluster data to constrain early dark energy component assuming the redshift of transition between 0 to 1. The current generation of SZ surveys will also measure the power spectrum of the cosmic microwave background with an unprecedented accuracy down to scales of 1 arc-min. The SZ surveys have started measuring the SZ power spectrum which is the dominant signal at scales of few $arc-min$ \\citep{sptcl, actcl}. The amplitude of the SZ power spectrum is proportional to the total number of objects that have formed and hence is extremely sensitive to the amount of dark energy present at early epoch.  The paper is organized as follows. In Section~\\ref{ede} we explain the early dark energy formalism and model used for this analysis. Section~\\ref{obs} expounds the nature of X-ray and microwave observations and mass-observable scaling relations. In section~\\ref{res} we show predictions from current and future observations, comparing our EDE models to a fiducial $\\Lambda$CDM cosmology. Finally, section~\\ref{concl} is devoted to conclusions and discussion.  \\begin{figure*}[t] \\includegraphics[height=8cm, width= 16cm]{fig1.eps} \\caption{\\small Equation of state (panel (a)) and matter power spectrum at $z=0$ (panel (b)) for EDE models with $\\w = -1.0, \\wm = -0.05, \\at =0.17, \\dt =0.17$ (EDE1, red lines), and with $\\w = -0.9, \\wm = -0.05, \\at =0.1, \\dt -=0.1$ (EDE2, blue lines). The black line in each panel represents the corresponding $\\ld$CDM model, in panel (b), the dashed lines represent EDE without DE perturbations, while the solid lines represent EDE with perturbations. } \\label{fig:ede} \\end{figure*} ",
        "conclusions": "In this work, we have studied early dark energy models in light of future galaxy cluster data. We study two EDE models, both of which are allowed within the current observational constraints. The first model (EDE1) transits from EDE behavior at redshift $z \\ggeq 5$ to $\\ld$CDM-like behavior at present, while the second model (EDE2) transits at a higher redshift ($z \\simeq 9$) to an equation of state $w_0 = -0.9$, which is close to but not identical to $\\ld$CDM. For both the models, as well as for the fiducial $\\ld$CDM cosmology, we show the predictions for cluster abundance and SZ power spectrum. For that purpose, we consider two instruments -- eROSITA for X-ray surveys, and ACT/SPT for SZ effect. We also do a likelihood analysis of data simulated replicating the SPT and eRosita survey specifications, along with the simulated Planck CMB data, to obtain the constraints on the EDE parameters from future surveys, and find that future galaxy cluster counts and SZ power spectrum can put competitive constraints on these parameters.  It is worth noting some apparent differences between our work here, and some other recent works on EDE signatures. For example, Fedeli et al.~(2009) also considered several EDE models, but their cluster counts are higher than for the fiducial $\\Lambda$CDM model, although $\\sigma_8$ for EDE models are commonly lower in their work (as it is the case here as well).  That is because for their choice of EDE models Hubble parameters, $E(z)$ differs significantly from the corresponding $\\ld$CDM model. This affects minimum observable mass as a function of redshift; their EDE models have much lower mass threshold than $\\Lambda$CDM, thus resulting in higher observable cluster counts. In contrast, we have chosen models which have $E(z)$ similar to that of the corresponding $\\ld$CDM model, to highlight the difference coming from the perturbations rather than the background expansion.  Similar difference from this work, due to the radical departure of $E(z)$ from $\\Lambda$CDM, can be seen in Waizmann \\& Bartelmann (2009), for the effects of EDE on SZ power spectrum. In some other cases differences can arise depending whether one uses low-k, CMB normalization of the power spectrum as done here, or $\\sigma_8$ normalization as in Sadeh et al.~(2007), although it is clear that a successful model has to fulfill both constraints.  We show that the inclusion of dark energy perturbations has a major effect on the matter power spectrum, therefore increasing the possibility of discriminating EDE models from $\\ld$CDM using large scale structure probes. Neglecting the \\de~\\perts~leads to severe underestimation of the imprint which the sharp transition in dark energy equation of state leaves on the \\de~\\perts~and therefore on the matter power spectrum.  We show that the models considered here which are allowed by the current observations can be ruled out using the future galaxy cluster probes. It is also interesting that both the cluster counts and the SZ power spectrum are sensitive to the redshift at which the transition between early and present day dark energy occurs. We expect to put strong constraints on the equation of state of dark energy at present using low redshift geometric observables (such as the luminosity distance of type Ia SNe and Baryon Acoustic Oscillations peaks). These probes of geometry of the Universe, however, are insensitive to the high redshift behavior of dark energy. The galaxy cluster observables such as their redshift abundance and the SZ power spectrum, although at low redshift, are sensitive to the perturbations in the early dark energy models, and hence would be able to put constraints on the redshift of transition from EDE behavior to present-day, $\\ld$CDM-like behavior. With the ongoing and future cluster surveys in microwave, optical and X-ray waveband, galaxy clusters will be able to provide strong constraints on the dynamical dark energy sector."
    },
    "1004/1004.5024_arXiv.txt": {
        "abstract": "{The three most iron poor stars presently known ([Fe/H] equal to -5.96, -5.4 and -4.75) are carbon-rich, they are called C-Rich Ultra-Metal Poor Stars (CRUMPS). The origin of their peculiar surface abundances is not understood.} {We propose a synthetic view of the different models so far proposed to explain the peculiar abundances observed at the surface of the CRUMP stars. We deduce some expected trends based on nucleosynthetic arguments and look for signatures allowing to discriminate among models. We discuss the conditions for having CRUMP stars which are He-rich,  i.e. with a mass fraction of helium greater than 0.30 and up to 0.60.} {We discuss the chemical composition of stars made of interstellar medium mixed with wind material of very metal poor massive stars, with wind plus supernova ejecta and with material extracted from the envelope of early AGB stars. Rotating and non-rotating models are considered.} {The high nitrogen abundances observed in CRUMP stars imply that the material which is responsible for their peculiar abundance pattern must be heavily enriched in primary nitrogen. We show that rotating stars (both massive and intermediate mass stars) can produce the required amount of primary nitrogen, and can also account for the observed enhancements in C, O, Na, Mg and Al. CRUMP stars formed from wind material of massive stars mixed with small amounts of pristine interstellar medium are He-rich (helium mass fraction between 0.30 and 0.60), Li-depleted and  present low  $^{12}$C/$^{13}$C ratios (inferior to 10 in number). Such He-rich stars, if discovered, would confirm that the most metal poor CRUMPs formed from essentially pure wind/envelope material. They would provide the most direct way to probe the nucleosynthetic outputs of the first generations of stars.} {We show that rotation is a key ingredient to explain the abundance patterns of CRUMPS stars and probably also of at least some Carbon-Enhanced Metal Poor (CEMP) stars, in particuliar the CEMP-no stars. Similar non-rotating models, without any extra-mixing, do not succeed to explain the enhancements in the three CNO elements.}  \\keywords {Stars: interiors, rotation, mass loss -- Stars: early--types, AGB -- Stars: chemically peculiar } \\titlerunning{Are CRUMPS He-rich?} ",
        "introduction": "The most iron-poor objects found to date in the Universe are halo field stars with [Fe/H] as low as -5.96 \\citep{Frebel2008}. Surprisingly, the three most iron poor  stars present strong overabundances of carbon with respect to iron. More generally, \\citet{Lucatello2006} find a lower limit of 21$\\pm$2\\% for the number of stars with [Fe/H]$\\le$-2.0 and [C/Fe]$\\ge$+1.0. These stars are collectively named C-Enhanced Metal Poor (CEMP) Stars. The C-rich Ultra Metal Poor Stars (CRUMPS) are a subset of CEMP stars, those with [Fe/H]$\\le -4$. A more correct name should be C-Rich Ultra  {\\it Iron} Poor Stars. The overall metallicity, where the metals would not only contain iron but all the elements heavier than helium, would amount to a [M/H] above -1, which is, by far, not ultra metal poor!  For some of these stars, strong overabundances of nitrogen, oxygen, sodium, magnesium and aluminium with respect to iron have been found. Compared to solar values, the abundance ratios may be one to four orders of magnitudes greater depending on the element and/or star considered.  The CEMP stars can be classified into different categories  according to the presence or not of {\\it s-} and/or {\\it r-}process elements  \\citep[see the discussion in][]{Masseron2009}.  In the present work, when comparisons are made with observations, we shall focus on the ``CEMP-no'' category, {\\it i.e.} on those stars which present no evidence for enhancements in the {\\it s-} and {\\it r-}process elements. The absence of {\\it s-}process elements  makes less favorable the possibility that the peculiar surface abundances of these stars result from accretion of material from an AGB companion, which is in general rich in {\\it s-}process elements. As can be seen in  Fig.  21 of \\citet{Masseron2009}, the ``CEMP-no'' category contains the most iron poor objects and are thus the best candidates to study the enrichment processes due to short lived massive stars. Despite concentrating here on  ``CEMP-no''  stars, the model studied in the present work may also be relevant for at least some CEMP stars of other categories.   The main purpose of the present work is to discuss in details the consequences of the  ``spinstar'' model proposed by \\citet{Meynet2006} and \\citet{Hirschi2007} where CRUMPS are formed from material ejected by metal poor rotating stars (hereafter called source material) mixed with some amount of interstellar material. We present the various outputs which can be obtained from such models, distinguishing the following three possibilities for the origin of the source material: \\begin{enumerate} \\item The  wind material ejected by one massive spinstar (spinstar wind model). \\item Material ejected under the form of wind material during the whole massive spinstar lifetime together with some amount of matter ejected at the time of the supernova event (spinstar wind and supernova model). \\item The envelope of intermediate mass spinstar at the early AGB phase (spinstar E-AGB model). \\end{enumerate}  We shall discuss the main characteristics which may distinguish these various origins for the source material.  In particular, we present the conditions required  for producing CRUMP/CEMP-no stars, which in addition to be C-rich, are also He-rich.  The fact that some CEMP or CRUMP stars may be He-rich have important consequences for the determination of their physical characteristics: \\begin{itemize} \\item The mass determinations of the C-rich stars will be different depending on the abundance of helium considered. Typically, a star, at a given position in the HR diagram, has a smaller mass if it is He-rich. \\item Adopting a He-rich model atmosphere modifies the temperature and density structure of the outer layers where the absorption lines are formed. This may have an impact on the abundances deduced from spectral synthesis. \\end{itemize}  The discovery of very iron-poor He-rich stars in the field of the halo would be interpreted in the frame of the present model as the existence of stars made up of nearly ``pure'' source material. Such objects would therefore represent the most direct opportunities for studying the nucleosynthetic outputs of the first stellar generations.  He-rich stars are already indirectly found in some globular clusters \\citep[see e.g.][]{Piotto2007}. They may also be present in elliptical galaxies. Let us recall that all ellipticals show the UV upturn phenomenon, i.e. an UV excess, observed in their spectra. These excesses are caused by a population of stars on the blue end of the horizontal branch, {\\it i.e.} by an old population of hot helium-burning stars \\citep[see the review by][]{OConnell1999}. Different models have been proposed to explain the origin of this population  \\citep[see e.g.][]{Bressan1994, Han2007}. A possibility, which still remains to be explored, would be that these stars are the descendants of He-rich stars \\citep{Meynet2008c}. He-rich small mass stars are indeed expected to populate the blue end of the horizontal branch \\citep{DecressinIAU2582009}.  The paper is organized as follows: in Sect. 2, we recall the main characteristics of the ``spinstar'' model. Sections 3  to 5 discuss the chemical abundance of CRUMPS expected in the frame of the three possibilities listed above. In each case, we compare the theoretical expectations with the observed abundances in the most iron poor stars known today and in CEMP-no stars. In Sect. 6 we briefly discuss consequences of alternative scenarios found in the literature. The isotopic ratio $^{12}$C/$^{13}$C can be used to discriminate between different models. This is presented in Sect. 7. Conclusions and perspectives are given in Sect. 8. ",
        "conclusions": " \\begin{enumerate} \\item CRUMP stars originating from the winds of spinstars present strong signs of material processed by H-burning, {\\it i.e.}, high [N/C] and [N/O] ratios, and small  $^{12}$C/$^{13}$C ratios. This is independent of the dilution factor. If the dilution factor is small, these stars are He-rich and Li-poor. \\item The spinstar wind model does appear in the present study as the only model which can explain CEMP stars with low $^{12}$C/$^{13}$C ratio. It might be that new rotating AGB star models could also be potential interesting candidates but this has to be confirmed by further computations. \\item Some CRUMPS can also be formed of a mixture of both wind and SN ejecta from massive fast rotating stars. Supernovae from non-rotating models cannot reproduce the abundance pattern observed in CRUMPS. \\item Stars made of the envelope of E-AGB stars resulting from rotating models present also abundance pattern in qualitative agreement with CEMP-no stars. Similar non-rotating models do not fit the observed abundances. \\item Both the wind+supernova model and the spinstar AGB model requires some dilution with pristine interstellar medium to obtain a reasonable fit. This means that such stars would not be He-rich and would present high $^{12}$C/$^{13}$C ratios. \\end{enumerate}  \\subsection{Why are there no analogs of CRUMPS at higher metallicities?}  One may wonder why the scenario leading to the formation of CRUMPS no longer works in the present universe. Two reasons can be invoked: 1) At high metallicity, rotation cannot produce primary nitrogen and thus the peculiar simultaneous overabundance of CNO elements. 2) At high metallicity, the chemically enriched winds of massive stars propagate in a medium which has already been pushed away by the strong and fast O-type star stellar winds. This may prevent the formation of stars made almost completely of pure chemically enriched ejecta. At very low metallicity the situation can be very different, the line driven stellar winds are first very weak preventing the star from pushing out the circumstellar material. Only when the stellar surface is enriched in heavy elements,  does the star lose large amount of material by stellar winds. Therefore, these enriched, slow winds (slow because ejected when the star is a supergiant), will have to go through a dense circumstellar medium. This may favor the formation of clumps. Further detailed hydrodynamical modelings are needed to confirm or reject such scenarios.            \\subsection{How to show that some CRUMPS are He-rich?}  Let us end by saying a few words about the possibility to test the hypothesis of CRUMPS  (or maybe also some CEMP) stars being He-rich. Although detection of He will remain problematic in  cool stars, it is not hopeless. Among the possibilities let us cite the following ones: \\begin{itemize} \\item Helium lines may be observable in cool stars, whose strength might be related to the He abundance. The line He I 10 830 \\AA\\    was used by \\citet{Moehler2000}  mainly for velocity diagnostics. It is not a straightforward line to interpret because it arises from a very excited metastable level in He I.  However, according to Dupree (private communication)  if stars of similar $T_{\\rm eff}$ are observed,   a relative measure could be done. It might also be possible to get abundances  but that would  need probably  non LTE modeling\\footnote{In the hotter HB stars (T$_{\\rm eff}$ $>$10000) other Helium lines become visible and those have been used to determine abundances \\citep[see][]{Moehler2000}.}. An additional difficulty with this approach comes from the fact that atomic diffusion would deplete the surface in Helium \\citep[see the models by][]{Korn2009}. \\item Asteroseismology may help in revealing the He-abundance of these stars. \\item To observe an eclipsing binary system whose components are CRUMPS or CEMP stars and to determine the positions of the two components in the HR diagram as well as their masses. The comparison with the stellar models may put in evidence the need to enhance the He-abundance in order to simultaneously reproduce all the observed features. \\item The presence of a significant population of extremely blue horizontal branch stars in the halo field may be an indirect hint of He-rich stars, since such stars can be the descendants of He-rich stars. \\end{itemize}   \\subsection{What to conclude if He-enhancement is detected? If He abundance is normal?}  If He-enhancement is detected and found to be a feature concerning the whole star and not only its surface, then this would definitively imply that the material from which the CRUMPS are made do contain relatively small amount of interstellar medium. This would discard models explaining CRUMPS as formed by the accretion of a small amount of mass at the top of the secondary in a binary system. It would also discard the normal supernovae as a possible source for the material from which the CRUMPS are made of, because the dilution factors which are used to bring the composition of normal supernova ejecta to the level of the  observed abundances in CRUMPS are much too high to make the star He-rich. Faint supernovae with large fall back would remain interesting candidates together with the effect of rotation to explain the high nitrogen enhancements. As discussed in the present work, the wind of spinstars would also be a favored candidate in that case. Let us stress here that this model is  not incompatible with the faint supernova model and both may be linked in a same evolutionary history, rotation allowing to enrich the envelope in H- and He-burning products and to eject it as a wind before the faint supernova event which will add its contribution. As indicated in the introduction, the finding of He-enhancement in  the CRUMPS will also modify the estimates of the mass of these objects. It may be that the most massive of these He-rich stars now populate the blue end of the horizontal branch.  If He-abundance is found to be normal, namely near the primordial helium abundance, then it favors either models invoking small amount of accreted mass (as in the AGB model described above), or models with large dilution factors. In that last case the very low abundance of Li observed in some CEMP stars should be due to depletion mechanism in the star itself. Many of the spinstar models discussed here are compatible with normal He-stars (see Table \\ref{models}). In particular, the wind of spinstars model would remain an interesting model for normal He-stars showing small  $^{12}$C/$^{13}$C number ratios, since combinations of dilution factor and of the initial mass of the spinstar may yield very modest helium enrichments, while still allowing very low $^{12}$C/$^{13}$C number ratios."
    },
    "1004/1004.2923_arXiv.txt": {
        "abstract": "The deposition of mechanical feedback from a supermassive black hole (SMBH) in an active galactic nucleus (AGN) into the surrounding galaxy occurs via broad-line winds which must carry mass and radial momentum as well as energy. The effect can be summarized by the dimensionless parameter $\\eta={\\Mdotoutf}/{\\Mdotacc}= {2 \\epsw c^2}/{\\vw^2}$ where $\\epsw$ ($\\equiv \\Edotw/(\\Mdotacc c^2)$) is the efficiency by which accreted matter is turned into wind energy in the disc surrounding the central SMBH. The outflowing mass and momentum are proportional to $\\eta$, and many prior treatments have essentially assumed that $\\eta=0$. We perform one- and two-dimensional simulations and find that the growth of the central SMBH is very sensitive to the inclusion of the mass and momentum driving but is insensitive to the assumed mechanical efficiency. For example in representative calculations, the omission of momentum and mass feedback leads to an hundred fold increase in the mass of the SMBH to over $10^{10} \\Msun$. When allowance is made for momentum driving, the final SMBH mass is much lower and the wind efficiencies which lead to the most observationally acceptable results are relatively low with $\\epsw \\lesssim 10^{-4}$. ",
        "introduction": "Feedback from active galactic nuclei (AGNs) at the centers of galaxies is believed to have a significant effect on the evolution of those galaxies. However, the precise physical mechanisms by which this feedback occurs are greatly uncertain---perhaps more so than is commonly acknowledged. While much path-breaking and insightful work has been done, it is also true that some of the most basic requirements, such as the necessity that mass, energy, and momentum be conserved, have not been imposed in several of the popular treatments of this subject. And the inclusion of the presently known and observed feedback processes is often treated selectively. The purpose of this paper is to attempt to lay out the physical framework for discussing the issues and to provide illustrative examples of the results obtained primarily from one-dimensional computations that include or exclude specific processes. We also include a treatment of the two-dimensional, axisymmetric case, presented in less detail, to show how the qualitative features carry over to this more realistic case. Definitive solutions are beyond present art in this field, so the focus will be on the qualitative features of the physical solutions rather than the detailed comparison with observations.  In outline, there are three phases to the overall phenomenon: 1) the provision of fuel for the central supermassive black hole (hereafter SMBH); 2) the generation of the outflowing stream of energy, mass and momentum from the vicinity of the SMBH; and 3) the absorption and transmission of this energy, mass and momentum by the ambient gas in the galaxy and the subsequent reactions of the ambient gas to that input.   1. The fueling is generally believed to be via infalling gas, and typically, two origins for that gas have been proposed; at high redshift ambient gas in discs liberated by the non-axisymmetric forces released during mergers is certainly important \\citep{bar91}, while at lower redshift mergers fail and probably do not fuel AGNs \\citep{li08}, but the processed gas released via normal stellar evolution provides an ample source \\citep[][hereafter CO07]{mat83,shu83,cio91,pad93,cio07}. A primary clue as to which of these sources dominates in a specific case is provided by details of the metallicity distribution, since the reprocessed gas probably has super-solar metal abundance and will also show signs of stellar evolution such as higher nitrogen or S-process abundances. The clue to fueling by infalling globular clusters might be relatively low abundances of elements made in SN I such as Fe. In almost all treatments a central disc mediates between the inflowing material and the SMBH. Other sources, such as small stellar systems dragged in by dynamical friction have been considered from time to time. These stars, or others added to the central regions via loss-cone processes  \\citep{beg80} can be shredded during tidal interactions with the central SMBH or by collisions with one another or with a central disc---the debris collecting in the disc and feeding the central SMBH via conventional mechanisms \\citep{ost83,dai10}.  2. The outflows fall into three categories. The signature of AGNs is, of course, the enormous electromagnetic, luminous output, with major contributions from the IR bands to the gamma ray region. The bulk of the flux is typically in the ``UV bump\" and the flux from this region thus dominates for the ``momentum driven winds\" \\citep[e.g., see][]{pro00, kin03,pro04,deb09}. The region where this driving occurs is fairly close to the quasar  (50$\\rbh \\lesssim r \\lesssim 500 \\rbh$, where $\\rbh$ is the SMBH Schwarzschild radius). However, the moderately hard X-rays determine the average photon energy: $\\left< h \\nu \\right>={h \\int \\lnu \\nu {\\rm d}\\nu}/{\\int \\lnu {\\rm d}\\nu}$ when integrated over the spectrum. This region of the spectrum dominates the photon heating, causing the heated gas to approach the mean photon energy: $1.5 k T_X = \\left< h \\nu \\right>$ with $T_X \\sim 2 \\times 10^7$K \\citep{saz04}. The resultant heating occurs over an extended range of radii: 100 pc$ <r< $3 kpc \\citep[][hereafter Papers I and III]{cio09,cio10}, and it can be significant for $r \\gtrsim$ 0.1 pc \\citep{pro07}. It can efficiently drive outflows as shown in a series of papers by Ciotti, Ostriker and collaborators (cf.CO07 and references therein). For electromagnetic output there is, of course, no rest-mass component. The total energy emitted in this form has been established fairly accurately via the \\citet{sol82} argument to be $\\Delta E = \\epsrad \\Delta \\Macc c^2$ with $\\epsrad \\sim 0.1-0.15$ \\citep{yu02}. The momentum output, of course, is $\\Delta p = \\Delta E / c$.  In optically thick cases ($\\tau \\gg 1$), the total momentum absorbed by the fluid can approach $\\Delta p = \\tau \\Delta E / c$ \\citep{deb09}. \\citet{sil10} also consider the importance of radiative momentum driven winds on galactic and cluster scales but limit the input to $L/c$, which can be considerably less than allowed in the optically thick case by \\citet{cio07} or \\citet{deb09}.  Next, let us turn to mechanical output. Both broad- and narrow-line regions inject mass, energy and momentum into the surrounding gas, with the broad-line winds probably dominant. Since these are material flows with velocity in the vicinity of the SMBH, $\\vw$, the mass outflow can be considerable. If we let the inflowing and outflowing mass rates be ($\\Mdotinf$, $\\Mdotoutf$), then conservation of mass, energy and momentum can be summed up with the following simple equations: \\beq \\Mdotacc = \\Mdotinf - \\Mdotoutf, \\label{eq:Mdot} \\eeq where $\\Mdotacc$ is the mass rate actually accreted by the SMBH, and \\begin{subequations} \\begin{eqnarray} \\Edotw &=&  \\frac{1}{2} \\Mdotoutf \\vw^2\\label{eq:edotw2}\\\\ &=& \\epsw \\Mdotacc c^2, \\label{eq:edotw1} \\end{eqnarray} \\end{subequations} \\beq \\pdot = \\Mdotoutf \\vw, \\label{eq:pdot} \\eeq are the wind energy and momentum, respectively. We have oversimplified matters by allowing only one wind velocity, when in fact Equation~\\ref{eq:edotw2} requires $\\left<\\vw^2\\right>$ and Equation~\\ref{eq:pdot} requires $\\left<\\vw\\right>$. Also, it is important to specify exactly where and when the quantities in equations \\ref{eq:Mdot}--\\ref{eq:pdot} are to be measured. In the conventional treatment of this subject, the SMBH is surrounded by a disc or torus to which matter has fallen from larger radius. Then, placing a sphere around this disc or torus (at $r \\sim 1$ pc), the instantaneous spherically averaged infall through the sphere is $\\Mdotinf(t)$ and the spherically averaged outflow is $\\Mdotoutf(t)$. The difference will be accreted onto the SMBH unless driven out in disc originating winds; the latter of course contributes to $\\Mdotoutf$ and so the remainder $\\Mdotinf - \\Mdotoutf$ will be accreted. Two further complications are allowed for in some detailed treatments: (i) the actual instantaneous value of $\\Mdotacc$ is a time-lagged convolution of the quantity in Equation \\ref{eq:Mdot} since a finite time elapses as material is transported through the disc to the central SMBH and (ii) star formation may (in fact frequently will) occur in the disc, removing mass that would otherwise have accreted onto the SMBH. Both of these complications are allowed for in CO07 and other work, neither is of dominant importance.  Now, defining the dimensionless ratio from Equation~\\ref{eq:edotw2} and \\ref{eq:edotw1} to be, \\beq \\eta \\equiv \\frac{\\Mdotoutf}{\\Mdotacc}= \\frac{2 \\epsw c^2}{\\vw^2}, \\label{eq:eta} \\eeq we can now rewrite Equations~\\ref{eq:Mdot}--\\ref{eq:pdot} as \\begin{subequations} \\begin{eqnarray} \\Mdotacc &=& \\Mdotinf \\frac{1}{1+\\eta}\\label{eq:Mdot_sol} ,\\\\ \\Mdotoutf &=& \\Mdotinf\\frac{\\eta}{1+\\eta}, \\\\ \\Edotw     &=& \\frac{1}{2} \\Mdotinf \\vw^2 \\frac{\\eta}{1+\\eta} =\\epsw \\Mdotinf c^2 \\frac{1}{1+\\eta}, \\\\ \\pdot        &=& \\Mdotinf \\vw \\frac{\\eta}{1+\\eta}. \\end{eqnarray} \\label{eq:sol} \\end{subequations} These Equations, \\ref{eq:sol}a\\--d, are, in fact, the ones that most authors have adopted who treat AGN feedback as a unified process comprising both infall and outflow. However, they typically adopt $\\eta=0$, implicitly assuming $\\vw \\to \\infty$, so that $\\Mdotoutf$ and $\\pdot$ are neglected and the two terms that are included, $\\Edotw$ and $\\Mdotacc$, may be overestimated. If it eventuated that $\\eta$ really is a very small number, then not much error would be induced and one would be justified in neglecting the out-flowing mass and momentum and in setting $\\Edotw \\sim \\epsw \\Mdotinf c^2 $, as most authors assume. If we adopt for the efficiency of generating mechanical energy the value $\\epsw = 5 \\times 10^{-3}$, as done by \\citet{spr05,joh09}, hereafter SDMH05 and JNB09 respectively, \\citet{mcc10} and other authors, and we take $\\vw = 10^4$ km/s ($\\vwten$) \\citep{moe09}, then we have from Equation~\\ref{eq:eta}, $\\eta = 9 \\vwten^{-2}$ and all of the neglected effects may in fact be dominant; the bulk of the inflowing mass may be ejected in a broad-line disc wind, and the mass and momentum input deposited in the ambient gas may dominate over the energy input, which may be largely radiated away. Papers I and Paper III do include these effects, but do not spell out their significance. The principal purpose of the present paper is to do just that---to show, in specially simple one- and two-dimensional calculations the effects of including or excluding mass, energy and momentum conservation when $\\eta > 0$. In addition to the papers referred to above which attempt to compute both the infall to the central SMBH and the outflow from it in a unified fashion, there are many others which postulate a central source and then, after estimating the mass, momentum and energy flowing at of that source (and some angular and temporal distribution thereof) do effectively compute the effects of that injection of energy, mass and momentum onto the surrounding fluid. Space does not permit a comprehensive description of this related subject of research, but important papers include the following: cf. \\citet{met94,ste08,fab09,ree09,ari10,gas10}.  The wind efficiency, $\\epsw$, is not known very well---neither from observations nor from detailed physical simulations. But the best estimates from either of these sources might be in the range $1\\times10^{-3} > \\epsw > 3\\times10^{-4}$ \\citep{pro00,pro04, kro07,sto09,kur09}, a factor of 5 to 17 smaller than the commonly adopted values and in a range where $\\eta \\lesssim 1$ if $\\vwten \\approx 1$.  A specific example may be useful.  \\citet{moe09} study the quasar SDSS J0838+2955. They find a mechanical energy output of $4.5 \\times 10^{45}$ erg/s, a mass outflow rate 10 times the accretion rate and a mechanical efficiency of $1 \\times 10^{-3}$, and they quote other observational studies which indicate similar numbers.  From analyses of the ionization parameters in the broad-line winds, estimates of the radial extent of the winds can be made; the above paper, and those quoted within indicate radii measured in kpc---consistent with the one-dimensional numerical work in Paper III.  As shown in Papers I and III, an additional important question asks ``what fraction of the sky is covered with the broad-line winds?\" Again two approaches are possible. Empirically, on the order of 20-25\\% of bright quasars show broad-line winds; this translates to $\\sim \\pi$ steradians or $\\pi/2$ steradians in each conical outflow, if we assume that the wind is emitted symmetrically above and below the inner AGN disc. On the theory side, the radiation driven winds found by \\citet{pro04}, via detailed hydro radiation-transfer calculations, cover $\\sim \\pi$ steradians, roughly consistent with the observational estimates.  Finally, let us turn to the narrow jets, the outflow observed from AGN in ``radio mode\", when the electromagnetic luminosity is considerably below the Eddington limit. M87 is an excellent nearby example of such a system. These are standard FRI radio sources. Here the jets are quite narrow and appear to be comprised primarily of a relativistic fluid. The same type of calculation as presented in the last section would indicate that the out-flowing mass is of negligible importance and the energy output greatly dominates over the momentum output. The total energy output from these phases is considerable, but the accretion rates are thought to be low in these phases so the efficiencies of energy generation may be very high \\citep[cf. for a computation][]{gam04}. Since so little mass is accreted in radio mode, the Soltan argument cannot be used to empirically estimate efficiencies, but, from the observational estimates of the energies available in the giant radio lobes, it may be that the AGN emits in radio mode considerably more energy than it does in wind mode. However the deposition from the intense but extremely narrow streams appears to be inefficient, and the jet drills through the gas in the surrounding galaxy, dumping most the energy into the intergalactic medium.  Thus, while it may act as the dominant feedback mechanism for the IGM (and we will return to this in a subsequent paper), it is of lesser importance than the radiative or wind components in heating and driving out the ambient gas from within a galaxy.  3. The interactions between the out-flowing energy, mass and momentum with the ambient fluids are complex and are just beginning to be studied with the needed detail. We focus here on the relatively gas poor elliptical systems, since it is in these that the bulk of the mass in SMBHs is found, The radiative interactions are perhaps easiest to describe. Since the mechanical momentum is conserved and cannot be radiated away, it can be a dominant effect. The minimum level of interaction is provided by electron scattering and, since the most luminous quasars are found to be clustered near the Eddington luminosity limit (at which level the momentum absorbed by electron scattering balances the gravitational force on the fluid from the central SMBH), we know that this effect is significant in many cases. Absorption of the out-flowing radiation will not, in general, reduce this effect, since typically the radiation is simply re-emitted in another band and electron scattering opacity is wavelength independent until the Klein-Nishina limit is reached at very high energies. In fact, in the optically thick limit, the radiation is transformed by dust absorption into the infra-red, but the effects in this case are even greater than in the simple case, since the scattering opacity of the dust to infra-red is, per atom, larger (by roughly a factor of 5) than the electron scattering cross section. For the bright ULIRGs, which may contain both an active AGN and a brighter starburst, there will be a near balance between the inward gravitational forces and the outward radiative momentum transfer on the dust \\citep[cf.][CO07]{tho05}. Under these circumstances, the inner several hundred parsecs of the galaxy are analogous in their equilibrium structure to a very massive star in so far as there is a nearly equilibrium balance between radiative and gravitational forces.   The effects of heating from the AGN are, for quite different reasons, also likely to be independent of absorption (so long as it is not excessive, i.e. not Compton thick). \\citet{saz05} present a simple analytical exploration of the effects and Paper I presents a more detailed one-dimensional treatment. The photons which dominate the heating process are in the moderately hard region ($\\sim$50keV), and we know from X-ray absorption studies that AGN are typically optically thin to such radiation. Metal line resonance absorptions dominate the absorption unless the spectrum is extremely hard, and in those cases Compton absorption would be dominant. If we consider the issue on a per atom basis, all that matters is the heating per atom, which scales as $r^{-2}$ (assuming that the fluid is optically thin to hard X-rays), and the cooling rate per atom which scales as the density. Since the latter can also scale as $r^{-2}$ or even falls off at a steeper rate in some circumstances, the heating can balance or exceed cooling over an extended range of radii. If that happens, the gas temperature will rise towards the radiation temperature, $T_X \\sim 2 \\times 10^7$ K. Then, since this exceeds the virial temperature in almost all galaxies, the heated gas, having thermal energy higher than its gravitational energy, will be accelerated outwards and tend to drive a wind into the surrounding fluid. Of course, since this will shut off the accretion flow and the fuel to the central source, the result will be a burst of energy output followed by much slower cooling of the shocked gas and a repeated burst at a much later time. Thus, episodic accretion is expected.  The mechanical energy input is more localized to the vicinity of the SMBH and would be efficient in ``protecting\" the SMBH from very high rates of accretion, except for one important caveat. It necessarily happens that such episodes of high rates of energy deposition will occur when the central gas densities are high, and under such circumstances the gas will tend to radiate away the input energy unless forbidden to do so as has occurred in some calculations \\citep{boo09}. This, as we shall see, makes the energy input rather inefficient in driving outflows and in protecting the SMBH from excessive accreation. But the momentum input cannot be radiated away, and, as we shall see in the remainder of the paper, it is very efficient in limiting the infall and accretion onto the central SMBH. Mechanical input, via either thermal or momentum based mechanisms, will also tend to produces episodic accretion.   The broad-line gas outflow must drive a strong shock into the ambient gas, and that, in turn, given standard physics, should accelerate charged particles efficiently via a variant of the first order Fermi process \\citep[cf.][]{bla78,bel78,bla87}. Then this relativistic fluid will further drive the outflow and, since thermal radiation is suppressed for this component, the conversion may somewhat enhance the effects of feedback. But, overall, this process simply transforms internal energy from one form to another and so, whereas it may be observationally quite significant, it will have a relatively small global effect. Two recent papers that have explored these processes are \\citet{fuj07} and \\citet{jia10}; see also \\citet{sir10}. ",
        "conclusions": "\\label{Discuss} The primary purpose of this paper is to quantitatively show, based on one- and two-dimensional computations, exactly which processes are most important during AGN feedback episodes; which processes are most useful in protecting the central SMBH from excessive mass growth and which have most effect on the ambient galaxy.  After a central outburst, the mechanical energy must be communicated to the ambient gaseous fluid by a wind and we, in fact, see these winds in luminous galaxies labeling them ``broad-line regions\" with outflow velocities observed to be $\\sim 10,000$ km/s covering conical regions subtending 20-25 \\% of the sky. These winds {\\it must} carry mass and radial momentum to the ambient fluid, reducing thereby the mass deposited on the central SMBH and adding a driving component which cannot  be reradiated away by thermal processes. Equations 5a-d summarize the physics, with the dimensionless parameter $\\eta$, indicating the importance of mass and momentum outflows.  In the case of an assumed high mechanical efficiency ($\\epsw=0.005$), we find that, if we suppress the mass and momentum input, then the SMBH grows by over a factor of 100 more than if momentum and mass flux were properly included in the calculation, and reaches masses~$> 10^{10} \\Msun$ in both one- and two-dimensional calculations. Turning on or off the energy input has relatively much less effect, altering the SMBH growth by roughly a factor of 2. Ignoring the mass and momentum feedback inputs also leaves the galaxy with a central optical luminosity from the AGN which is orders of magnitude brighter than is seen in nearby elliptical systems.  Compared to these dramatic effects, the uncertainty due to not knowing accurately the wind efficiencies has a relatively minor effect. Reducing the efficiency by a factor of 20 from $5\\times10^{-3}$ to $2.5 \\times 10^{-4}$ reduces the wind energy output by only a factor of 2 (to $10^{58.7}$ erg) and reducing the efficiency by another factor of 10 reduces the wind energy output by only another factor of 2.  We also found that redirecting much of the inflowing mass into a BAL wind has, by itself, an important effect on models with only radiative feedback. In those computations, which do not allow for the redirection, the central SMBH again grows far too much in both the one- and two-dimensional computations.  In summary, it eventuates that enforcing the conservation of mass, momentum and energy provides extremely useful constraints in estimating the growth of central SMBHs and the feedback effects on the surrounding galaxies."
    },
    "1004/1004.0117_arXiv.txt": {
        "abstract": "The formation of stars is inextricably linked to the structure of their parental molecular clouds. Here we take a number of nearby giant molecular clouds (GMCs) and analyse their column density and mass distributions. This investigation is based on four new all-sky median colour excess extinction maps determined from 2MASS. The four maps span a range of spatial resolution of a factor of eight. This allows us to determine cloud properties at a common spatial scale of 0.1\\,pc, as well as to study the scale dependence of the cloud properties.  We find that the low column density and turbulence dominated part of the clouds can be well fit by a log-normal distribution. However, above a universal extinction threshold of $6.0 \\pm 1.5$\\,mag $A_V$ there is excess material compared to the log-normal distribution in all investigated clouds. This material represents the part of the cloud that is currently involved in star formation, and thus dominated by gravity. Its contribution to the total mass of the clouds ranges over two orders of magnitude from 0.1 to 10\\,\\%. This implies that our clouds sample various stages in the evolution of GMCs. Furthermore, we find that the column density and mass distributions are extremely similar between clouds if we analyse only the high extinction material. On the other hand, there are significant differences between the distributions if only the low extinction, turbulence dominated regions are considered. This shows that the turbulent properties differ between clouds depending on their environment. However, no significant influence on the predominant mode of star formation (clustered or isolated) could be found. Furthermore, the fraction of the cloud actively involved in star formation is only governed by gravity, with the column density and mass distributions not significantly altered by local feedback processes. ",
        "introduction": "\\label{intro}  The study of Giant Molecular Clouds (GMCs) is of great importance to understand star formation. Their properties (density distribution, dynamics, temperature, etc.) are thought to be vital in determining whether, where and how stars are forming within them. Turbulent motions inside a cloud cause fragmentation and hence determine the density distribution (Padoan et al. \\cite{1997ApJ...474..730P}, Ballesteros-Paredes et al. \\cite{1999ApJ...527..285B}). This in turn should have an influence on which mode of star formation (isolated or clustered) is occurring (Klessen, Heitch and Mac Low \\cite{2000ApJ...535..887K}). It is, however, unclear if this is the sole determinant of the star formation mode. Or are there other important causes such as the environment, feedback and/or magnetic fields?  One way to try to answer these questions is to investigate the column  density structure of a large number of GMCs in a systematic, comparable and unbiased fashion. This means ideally one should use the same tracer of column density for all clouds. Such a tracer should be as bias free as possible and all clouds should be investigated at the same physical spatial resolution. There are a number of techniques to determine the column density of clouds. These include molecular line emission, dust continuum emission, scattered infrared light or extinction measurements. The latter on the bases of star counts, colour excess or combined methods. See Froebrich \\& Rowles \\cite{2010BASI..........F} for a brief summary and discussion of the advantages and disadvantages of the various techniques. It has been shown by Goodman et al. \\cite{2009ApJ...692...91G} that near infrared extinction mapping is the best tracer of the `real' column density of a cloud, based only on the assumption of a constant gas to dust ratio.  A number of large scale maps of the column density distribution of clouds have been made to date, which could be used to compare the structure of a number of GMCs. These are for example the maps by Cambr\\'{e}sy \\cite{1999A&A...345..965C}, Schlegel et al. \\cite{1998ApJ...500..525S}, Dobashi et al. \\cite{2005PASJ...57S...1D}, Froebrich et al. \\cite{2007MNRAS.378.1447F} and Rowles and Froebrich \\cite{2009MNRAS.395.1640R}. Some of them do not use the ideal tracer (extinction) for the column density. Others use star counts, which result in a distance dependent bias in the column density distribution (Froebrich \\& del Burgo \\cite{2006MNRAS.369.1901F}), similar to using the mean colour excess instead of the median. One bias they all have in common is that clouds at different distances are mapped at different spatial resolutions. This will naturally lead to a change in the observed column density distribution, since for more distant clouds one averages over larger physical scales. This will lead to the effect that small scale high extinction regions are averaged out and the column density distributions are skewed. In order to compare the structure of GMCs with each other, one has to determine the column density distribution at similar physical scales and using the same `good' tracer for all investigated clouds.  Only then can we attempt to draw conclusions about similarities and differences in their structures. This can be done by analysing the column density and mass distribution (the focus of this paper) or determining structure functions (the focus of the next paper in the series) and comparing it to model predictions such as published by  Kolmogorov \\citep{1941DoSSR..30..301K}, She \\&  Leveque \\citep{1994PhRvL..72..336S} and Boldyrev \\citep{2002ApJ...569..841B}. In our previous paper (Rowles \\& Froebrich \\cite{2009MNRAS.395.1640R}, Paper\\,I hereafter) we presented all-sky  extinction maps determined using data from 2MASS (Skrutskie et al. \\cite{2006AJ....131.1163S}). These were created using an adaptation of the NICE method (Lada et al. \\cite{1994ApJ...429..694L}), median colour excess determination and variable spatial resolution to obtain a constant signal to noise ratio. Here we present and analyse additional maps, determined using a range of constant spatial resolutions to map different clouds at similar spatial scales.  In Sect.\\,\\ref{method} we briefly describe the new maps created for the project, and explain the data analysis methods. In Sect.\\,\\ref{results} we present the results for our selection of nearby GMCs. We discuss these results and draw conclusions in Sect.\\,\\ref{dissc}.  \\begin{table} \\centering \\caption{\\label{cloudlst} Summary table of the clouds analysed in this paper. We list the cloud name, the range in galactic coordinates ($l$;$b$) covered by the region and the distance (d) adopted in this paper. $^*$The cloud Auriga\\,1 is referred to as the California Molecular Cloud in Lada et al. (2009). The references for the adopted distances are the following: $^1$ Lada et al. (2009); $^2$ Kenyon et al. (2008); $^3$ Kun (2008); $^4$ Knude \\& Hog (1998); $^5$ Bally et al. (1999); $^6$ Muench (2008); $^7$ Lombardi et al. (2008); $^8$ Carpenter (2008); $^9$ Motte (2001); $^{10}$ Straizys (1996).} \\begin{tabular}{llll} Name & $l$ range & $b$ range & d [pc] \\\\ \\hline Auriga\\,1$^*$    &      156$\\degr$ --           173$\\degr$ &      $-$12.0$\\degr$ -- \\,\\,\\,$-$3.0$\\degr$ & 450$\\pm$\\,\\,\\,23$^1$     \\\\ Auriga\\,2        &      175$\\degr$ --           186$\\degr$ &      $-$11.0$\\degr$ -- \\,\\,\\,$-$2.0$\\degr$ & 140$\\pm$\\,\\,\\,28$^2$     \\\\ Cepheus          &      100$\\degr$ --           120$\\degr$ & \\,\\,\\,$+$5.0$\\degr$ --      $+$20.0$\\degr$ & 390$\\pm$125$^3$          \\\\ Chamaeleon       &      296$\\degr$ --           304$\\degr$ &      $-$18.0$\\degr$ --      $-$13.0$\\degr$ & 150$\\pm$\\,\\,\\,30$^4$     \\\\ Circinus         &      316$\\degr$ --           319$\\degr$ & \\,\\,\\,$-$6.0$\\degr$ -- \\,\\,\\,$-$3.0$\\degr$ & 700$\\pm$350$^5$          \\\\ Corona Australis &      359$\\degr$ -- \\,\\,\\,\\,\\,\\,1$\\degr$ &      $-$20.0$\\degr$ --      $-$17.0$\\degr$ & 170$\\pm$\\,\\,\\,34$^4$     \\\\ $\\lambda$-Ori    &      188$\\degr$ --           201$\\degr$ &      $-$18.0$\\degr$ -- \\,\\,\\,$-$6.0$\\degr$ & 400$\\pm$\\,\\,\\,80$^6$     \\\\ Lupus\\,1, 2      &      334$\\degr$ --           342$\\degr$ &      $+$11.0$\\degr$ --      $+$19.0$\\degr$ & 155$\\pm$\\,\\,\\,\\,\\,\\,8$^7$\\\\ Lupus\\,3, 4, 5, 6&      335$\\degr$ --           344$\\degr$ & \\,\\,\\,$+$5.0$\\degr$ --      $+$11.0$\\degr$ & 155$\\pm$\\,\\,\\,\\,\\,\\,8$^7$\\\\ Monoceros        &      212$\\degr$ --           222$\\degr$ &      $-$13.0$\\degr$ -- \\,\\,\\,$-$7.0$\\degr$ & 830$\\pm$\\,\\,\\,50$^8$     \\\\ Ophiuchus        &      350$\\degr$ --           360$\\degr$ &      $+$12.0$\\degr$ --      $+$19.0$\\degr$ & 119$\\pm$\\,\\,\\,\\,\\,\\,6$^7$\\\\ Orion\\,A         &      208$\\degr$ --           219$\\degr$ &      $-$21.0$\\degr$ --      $-$16.0$\\degr$ & 410$\\pm$\\,\\,\\,80$^6$     \\\\ Orion\\,B         &      201$\\degr$ --           211$\\degr$ &      $-$17.0$\\degr$ -- \\,\\,\\,$-$8.0$\\degr$ & 410$\\pm$\\,\\,\\,80$^6$     \\\\ Perseus          &      156$\\degr$ --           163$\\degr$ &      $-$25.0$\\degr$ --      $-$15.0$\\degr$ & 310$\\pm$\\,\\,\\,65$^9$     \\\\ Serpens          & \\,\\,\\,30$\\degr$ --      \\,\\,\\,32$\\degr$ & \\,\\,\\,$+$4.0$\\degr$ -- \\,\\,\\,$+$6.0$\\degr$ & 260$\\pm$\\,\\,\\,10$^{10}$  \\\\ Taurus           &      164$\\degr$ --           178$\\degr$ &      $-$19.0$\\degr$ --      $-$10.0$\\degr$ & 140$\\pm$\\,\\,\\,28$^2$     \\\\ \\end{tabular} \\end{table} ",
        "conclusions": "\\label{dissc}  We investigated the column density and mass distribution of a number of GMCs based on near infrared median colour excess extinction maps determined from 2MASS. Our multi scale extinction mapping approach enabled us to determine the cloud properties homogeneously at a common scale of 0.1\\,pc for all clouds. Furthermore, it allowed us to investigate how the cloud properties depend on the spatial scale.  We find that for about half the investigated clouds the slope of the column density and mass distribution changes to more negative (steeper) values with increasing spatial scale. For the other half no change is found over a scale range of almost an order of magnitude. The former can be understood by the fact that at larger scales, small scale high extinction cores are not detected anymore, naturally leading to a steeper mass and column density distribution. In the latter case this does not happen. This could mean that there are no small scale high extinction regions. But this does not seem plausible, in particular since the clouds in this group are e.g. Taurus, Perseus, Ophiuchus and Orion. Another possibility is that there is no significant structure in those clouds at the smallest scales, and hence no change over the observed range of scales. We will analyse this in more detail in a forthcoming paper where we will determine the structure functions of all clouds.  Our fits to the column density distributions using a log-normal function resulted in variable outcomes. Some clouds, in particular Auriga\\,2, Corona Australis, Lupus\\,3, 4, 5, 6 and Serpens can be fit very well ($rms$ lower than 3\\,$\\sigma$). Other clouds (Cepheus, Monoceros, Ophiuchus) cannot be fit properly by a log-normal function. This is in agreement with earlier works on some of these clouds e.g. by Lombardi et al. \\cite{2008A&A...489..143L} and more recently by Kainulainen et al. \\cite{2009A&A...........K}. All investigated clouds show an excess of column density compared to a log-normal distribution at higher $A_V$ values. The quality of the fit is hence a measure of how much excess there is in a particular cloud. This excess material has decoupled from the general turbulent field and is dominated by gravity, in other words more likely actively involved in star formation. However, in our sample we could not find any significant correlations of the achieved quality of the log-normal fit with other cloud properties or the amount of young stars in the cloud.  The analysis of the mass distribution enables us to draw further conclusions about the structure of the clouds and their star formation properties. For all clouds the change in the slope of the mass distribution is more pronounced than in the column density distribution. It is hence easier to determine the threshold at which gravity becomes the dominant force in shaping the structure. This value ranges from 3.4\\,mag in Lupus\\,1, 2 to 8.6\\,mag of optical extinction in Ophiuchus. We associate this with the extinction threshold for star formation, originally found to be about 7\\,mag of $A_V$ in Ophiuchus by Johnstone et al. \\cite{2004ApJ...611L..45J}, in agreement with our value. On average for all clouds we find $\\left< A_{\\rm V, SF} \\right> = 6.0 \\pm 1.5$\\,mag optical extinction, a rather small range.  In contrast, the mass fraction of each cloud which is currently involved in star formation varies by almost two orders of magnitude. It ranges from 0.12\\,\\% in $\\lambda$-Ori to almost 10\\,\\% in Corona Australis. We note that these values are not an estimate of the star formation efficiency in these clouds. They rather should be seen as the potential of the cloud to form new stars in the next few 10$^6$\\,yrs. The wide range of values indicates that our sample contains clouds at different stages in their evolution. We see clouds that currently only form a small number of stars (e.g. Auriga\\,1 and Corona Australis) but with a completely different fraction of mass involved in star formation. In Auriga\\,1 only a small fraction of the total material seems to be available for future star formation. On the contrary in Corona Australis a much larger fraction of material is expected to be forming stars. While other regions that currently form stars (Orion, Taurus, Perseus) still possess a significant fraction of material in a state which is expected to continue star formation. One example for a cloud which might have reached the end of star formation seems to be Ophiuchus. It is currently forming a large number of stars, but has only less than one percent of material in a gravity dominated form.  We have searched amongst our determined cloud properties for further correlations. In particular we hoped to find a link between the predominant mode of star formation in the cloud (clustered vs. isolated) with one or several cloud properties. No such correlation could be found. This is partly expected since our extinction maps show only how much material is potentially involved in star formation, and not how much gas and dust (and its properties) have lead to the star formation mode we currently see in these clouds. One could expect that the properties of the turbulence dominated part of the clouds (out of which the denser parts are formed) show some dependence on the star formation mode. Either because different turbulent properties cause different modes of star formation, or that feedback from young star clusters and isolated YSOs causes different turbulent properties. But again, no such dependence could be found in our data.  However, there are clear differences in the properties of the clouds at low column densities compared to high $A_V$ values. The scatter of the $\\gamma_{\\rm low}$ and $\\delta_{\\rm low}$ values between clouds is a factor of 2.5 (for $\\gamma$) and 4.5 (for $\\delta$) larger than for the high extinction regions. Essentially this paints the following picture of the cloud structure: The low $A_V$ and turbulence dominated regions differ from cloud to cloud. Their column density and mass distributions are determined by the environment of the cloud and the feedback they have experienced from the star formation processes within them. Above a (column) density threshold of about 6\\,mag $A_V$, which is independent of the cloud, gravity becomes the dominant force in shaping the structure. This part of the cloud is then more and more decoupled from the influences of its surrounding turbulent field, and thus the column density and mass distributions for all clouds are virtually identical above the extinction threshold. This also implies that local feedback from young clusters or stars has no significant influence in shaping the column density and mass distributions."
    },
    "1004/1004.0926_arXiv.txt": {
        "abstract": "We calculate the screening of the ion-ion potential due to electrons in the presence of a large background magnetic field, at densities of relevance to neutron star crusts. Using the standard approach to incorporate electron screening through the one-loop polarization function, we show that the magnetic field produces important corrections both at short and long distances.  In extreme fields, realized in highly magnetized neutron stars called magnetars, electrons occupy only the lowest Landau levels in the relatively low density region of the crust. Here our results show that the screening length for Coulomb interactions between ions can be smaller than the inter-ion spacing. More interestingly, we find that the screening is anisotropic and the screened potential between two static charges exhibits long range Friedel oscillations parallel to the magnetic field. This long-range oscillatory behavior is likely to affect the lattice structure of ions, and can possibly create rod-like structures in the magnetar crusts. We also calculate the imaginary part of the electron polarization function which determines the spectrum of electron-hole excitations and plays a role in damping lattice phonon excitations. We demonstrate that even for modest magnetic fields this damping is highly anisotropic and will likely lead to anisotropic phonon heat transport in the outer neutron star crust. ",
        "introduction": "Highly magnetized neutron stars called magnetars that feature extreme magnetic fields ($B$) --- as large as $10^{15}$G on the surface and perhaps even larger fields inside --- have been detected in recent years.  Currently there are 13 magnetar candidates which are classified as anomalous x-ray pulsars (AXPs) or soft gamma repeaters (SGRs) (http://www.physics.mcgill.ca/~pulsar/magnetar/main.html). Such large magnetic fields  can strongly influence the state of matter in neutron stars, especially in the lower density region were the characteristic energy scales associated with the matter fields is small compared to the magnetic ``perturbation''. Several decades ago, it was realized that even modest magnetic fields could influence the structure of atoms in the atmospheres of magnetized neutron stars and affect the spectral features of the emitted radiation \\cite{Ruderman:1974} .  At the surface, where electrons are localized around nuclei and the system consists of neutral atoms, large magnetic fields distort the atomic structure resulting in rod-like or cigar shapes elongated parallel to $B$. (For a recent review see \\cite{Ruder:1994}.) It was realized quite early that these cigar shaped atoms would bind with each other  along the magnetic field direction forming polymer chains, and that these chains would interact in the perpendicular plane, resulting in a condensed phase \\cite{Ruderman:1971}.  More recently the physics of these condensed phases for hydrogen has been analyzed and its implications for neutron atmospheres is discussed in Ref.  \\cite{Lai:1997}. Here we will consider the effects of the magnetic fields at higher density, in the crusts of magnetars where the electrons are not localized in atomic states but rather form a degenerate Fermi system. This will be important in determining the structural and transport properties of the magnetar crust. Our finding will have implication for crust oscillations (excited in explosive phenomena such as giant flares) and for heat and electrical transport which are important ingredients in the study of magnetic and thermal evolution of magnetars \\cite{Geppert:2004,Aguilera:2007}.  In this work we will typically restrict the study to strong fields where only a few Landau levels are occupied and this will necessarily restrict us to relatively low density region in the outer crust of the neutron star.  ($\\rho \\lsim 10^{10}$ g/cm$^3$ for $B$ $\\lsim$ 10$^{15}$ G.) The key microscopic quantity required to characterize the response in our approach is the polarization function, which is the Fourier transform of the density-density correlation function. The expression for the polarization function for a magnetized electron gas is calculated and is found in agreement with previous work \\cite{Danielsson:1995}.  Our new findings in this work are related to identifying and understanding the implications of the electronic response in strong magnetic fields for neutron star structure and transport. The real part of the electron polarization function is related to the screening of the ion-ion potential in the crust. In an unmagnetized crust, screening is dominated by the usual Debye screening. In neutron stars where the typical magnetic field is small ($B$ $\\lsim10^{11}$ G),  the electron screening length is larger than the inter-ion distance and this implies that the effects of electron screening are usually unimportant in determining the crystal structure and thermodynamics. In contrast, we find that for large fields when only one or a few Landau levels are filled, the screening length can be either much smaller or much larger than in the unmagnetized case depending on the density  and magnetic field. At relatively low density the screening is enhanced resulting in a screening length that is generically smaller than the inter-ion distance. Our finding of enhanced screening for large $B$ is not new, this was already noted in early work which explored some its implications for nuclear fusion reactions \\cite{Shalybkov:1987}. The qualitatively new finding of our work is the realization that the sharp Fermi surface in the electron momentum distribution along the $B$ field will give rise to long-range oscillatory behavior called Friedel oscillations in the ion-ion potential --- a well studied feature in the condensed matter context for $B=0$. The period of this spatial oscillation $\\lambda_F = \\pi/(k_{f}^z)$, where $k_{f}^z$ the Fermi momentum of the electrons in the $z$-direction, where we have assumed that $B$ is parallel to $z$. This potential falls off slowly as $1/z$ on a scale of several inter-ion distances and consequently can be very important in determining the phase structure of the crust. We speculate that this will lead to the formation of an anisotropic crystal with rod-like structures akin to the polymer chains expected in the atomic regime.  The imaginary part of the polarization is related to dissipation and gives rise to the usual Landau damping observed in Fermi systems.  In the crust of a neutron star, excitations such as lattice phonons couple to the electrons, and their mean free paths are largely determined by the efficiency of the electronic Landau damping where the phonon decays by producing an electron-hole pair. We show that this decay rate has a non-trivial dependence on the angle between their propagation direction and the magnetic field because the electron-hole excitation for large $B$ is highly anisotropic. This anisotropy leads to an anisotropy in heat transport properties of lattice phonons. We find that the lattice phonons can be scattered more easily if they are moving perpendicular to the field, compared to when they are moving parallel to it.  The paper is organized as follows. In Section~\\ref{Section:PolarizationMagnetic} we derive the polarization function of electrons in the presence of a magnetic field. In the Sections~\\ref{Section:Screening} and \\ref{Section:HeatTransport}, we apply the result to two physical problems. In Section~\\ref{Section:Screening} we present numerical results for Friedel oscillations. In Section~\\ref{Section:HeatTransport} we discuss the effect of anisotropic polarization on the decay rates of the lattice phonons as a function of the angle between the propagation direction and the magnetic field.  We discuss the implications for the structure of the magnetar crust in Section~\\ref{Section:Conclusions}. Details of calculations of the expression for the screened potential are given in Appendix~\\ref{Section:LLL dominates} and Appendix~\\ref{Section:Friedel derivation}. ",
        "conclusions": "}  By calculating the one-loop electron-hole polarization function in a strong magnetic field we have found that the screening of the ion-ion potential is significantly different from those obtained in earlier studies where the Debye approximation was assumed. We show for the first time that the Fermi surface of electrons in the direction parallel to the magnetic field leads to Friedel oscillations in the ion-ion potential. These oscillations are large and could fundamentally change the structure of the ionic solid at large field and relatively low density when only the lowest Landau levels are occupied. For typical magnetar field strengths of order $10^{15}$ G, matter up to densities of the order $10^{10}$gm$/$cc will be affected by our finding here. To evaluate how Friedel oscillations and modifications to the average screening length affects the structure and melting properties of the solid we will need to perform either Path Integral Monte Carlo or classical Monte Carlo simulations because the ion-ion interaction is highly non-perturbative. This is beyond the scope of this work and will be investigated separately. Here, we present plausible implications based on qualitative arguments to explore how our findings will modify the ionic structure.  Considering the parameters used in Fig.~\\ref{fig:g_separate}, $eB=100eB_c$ and $\\mu=0.7$(MeV) and assuming $Z=26$, the average number density of ions in the system is $n_I=3.2\\times10^{-9}$ (fm)$^{-3}$. If one assumes that the ions form a regular bcc lattice, then the separation between nearest ions is $(\\sqrt{3}/2)\\times (2/n_I)^{1/3}=743$(fm).  This is comparable to the screening length $\\lambda_D$, and smaller than the wavelength for Friedel oscillations $\\lambda_F$. If significant screening in the $x-y$ plane prevents the formation of a regular bcc structure, it may be favorable to form an anisotropic crystal structure. One possibility is that the ions fit into the troughs in the potential formed by superposition of the $V(z)$ separated in the $z$ axis by $\\lambda_F$. To maintain neutrality, the ions should arrange themselves in the $x-y$ plane (in a regular or irregular structure) with an average separation $a_\\perp=(\\pi n_I\\lambda_F)^{-1/2} = 278$(fm).  If indeed the ions in the $z$ direction arrange themselves with separation $\\lambda_F$, then the following picture arises for the structure of the ion-electron system with changing density, or as we move deeper into the neutron star crust. A single chain of ions arranged along the $z$ axis will look like a rod. $\\lambda_F\\propto eB/n_I$, and for fixed $eB$ decreases linearly with increasing density. The transverse separation between the rods, $a_\\perp\\propto (eB)^{1/2}$, on the other hand remains a constant. At the surface, it is known that the elongated atoms form chains parallel to the $B$ field. At low densities (near the surface of the crust), ions will form a plasma of rods with a small charge per unit length (since $\\lambda_F$ is large). These rods will interact with each other by a two dimensional Debye screened Coulomb interaction. As we move deeper into the star, the charge per unit length will keep on increasing, and once $\\lambda_F$ becomes comparable to $a_\\perp$, the system will look more isotropic. This picture will break down when higher Landau levels start being occupied. It is interesting to note that for $Z=26$, $a_\\perp\\sim \\lambda_F$ implies that $\\mu^2 \\gtrsim 2 eB$, meaning that the next Landau level is occupied. Hence, the point where the screened potential becomes more isotropic coincides with the point where the lattice structure becomes more isotropic. Hence, in our picture, the region where the electrons occupy only the lowest Landau level interpolates smoothly between the surface and the inner crust of the neutron star.  We have also showed that the heat conduction due lattice phonons can become anisotropic because their damping rate due to electron-hole excitations is anisotropic. It is well known that the electronic heat conductivity perpendicular to the magnetic field is suppressed~\\cite{Urpin:1980}, and this in turn results in a temperature anisotropy at the surface of the neutron star~\\cite{Geppert:2004}. Our finding that the heat conductivity due to phonons in the direction perpendicular to the magnetic field is also suppressed, is new. This effect may modify the temperature anisotropy in the outer regions of the magnetars during their early thermal evolution."
    },
    "1004/1004.0051_arXiv.txt": {
        "abstract": "{} {We perform a comprehensive multiwavelength analysis of a sample of 20 starburst galaxies that show a substantial population of Wolf-Rayet (WR) stars. In this paper, the third of the series, we present the analysis of the O and WR star populations within these galaxies.} {We study the spatial localization of the WR-rich clusters via the detection of the blue \\WRBUMP, mainly composed by the broad \\ion{He}{ii} $\\lambda$4686 line and attributed to WN stars, and the red \\WRBUMP, composed by the broad \\ion{C}{iv} $\\lambda$5808 line observed in WC stars. We perform a detailed fitting of the nebular and broad emission lines within these broad features and derive the numbers of WN, WC and O stars using (i) the standard assumption of constant WR luminosities and (ii) considering metallicity-dependent WR luminosities. We then compare our results with the predictions given by evolutionary synthesis models and with previous empirical results.} {We report the detection of blue \\WRBUMP\\ in 20 regions, but the red \\WRBUMP\\ is only detected in six. Aperture effects and the exact positioning of the slit onto the WR-rich bursts play a fundamental role in their detection. The nebular \\ion{He}{ii} $\\lambda$4686 line is detected in 21 regions; its intensity clearly decreases with increasing metallicity. We derive an empirical estimation of the WNL/(WNL+O) ratio using the intensity of the broad \\ion{He}{ii} line assuming metallicity-dependent WR luminosities. As expected, the total number of WR stars increases with increasing metallicity, but objects with \\abox$<$8.2 show a rather constant WR/(WR+O) ratio. The computed WCE/WNL ratios are different than those empirically found in nearby star-forming galaxies, indicating that the observed galaxies are experiencing a strong and very short burst. Considering metallicity-dependent WR luminosities, our data agree with a Salpeter-like IMF in all regimes.} {We consider that the contribution of the WCE stars is not negligible at low metallicities, but deeper observations are needed to detect the red \\WRBUMP\\ because of the decreasing of the WR luminosities with decreasing metallicity. Although available models reproduce the WR properties at high metallicities fairly well, new evolutionary synthesis models for young starbursts including all involved parameters (age, metallicity, star-formation history, IMF and WR stars properties such as metallicity-dependent WR luminosities, stellar rotation and the WR binnary channel) are absolutely needed to perform an appropriate comparison with the observational data.}   \\titlerunning{Massive star formation in Wolf-Rayet galaxies III: O and WR populations}  \\authorrunning{L\\'opez-S\\'anchez \\& Esteban} ",
        "introduction": "Despite their relatively low number and short lifetime in terms of evolutio\\-nary timescales, massive stars have a fundamental influence on the interstellar medium (ISM) and galaxy evolution: they generate most of the ultraviolet ionizing radiation in galaxies, powering the far-infrared luminosities through the heating of dust; they enrich the interstellar medium, not only by returning nuclear processed material during their whole lifetime \\citep{Maeder81} but also in supernova explosions; finally, they deposit mechanical energy, first via strong stellar winds and later as supernovae \\citep{Abbott82}, which are a triggering mechanism of star formation \\citep{Woodward78}. Massive stars are the progenitors of the most energetic phenomena known today, the long gamma-ray bursts (GRB), as they collapse after supernova explotions into black holes \\citep{WoosleyHeger06}. The most massive ($M\\geq$ \\mbox{25 \\Mo}\\ for \\Zo) luminous (10$^5$ to 10$^6$ \\Lo) and hot (temperatures $\\sim$50~000~K) O stars evolve to the Wolf-Rayet phase between 2 and 5 \\Myr\\ since their birth. A WR star is interpreted as a central He-burning core that has lost the main part of its H-rich envelope via strong winds \\citep{Conti76,Maeder90,Maeder91}. This stage of the most massive stars lasts only some few hundreds of thousands of years ($t_{WR}\\leq5\\times10^5$~yr) before they explode as Type Ib/Ic supernovae  \\citep{MeynetMaeder05}.   The broad emission features that characterize the spectra of WR stars are often observed in extragalactic \\HII regions. Actually, these Wolf-Rayet galaxies make up a very inhomogeneous class of star-forming objects, all having in common ongoing or recent star formation which has produced stars massive enough to evolve to the WR stage \\citep{SCP99}. The blend of the broad \\HeII\\ $\\lambda$4686, \\ion{C}{iii}/\\ion{C}{iv} $\\lambda$4650 and \\ion{N}{iii} $\\lambda$4640 emission lines constitutes the {\\bf blue \\WRBUMP}; it mainly originates in WN stars with a minor contribution of WC stars. The original blue \\WRBUMP\\ detection was made in the blue compact dwarf galaxy He~2-10 by \\citet*{Allen76}. Later, \\citet{KS81} recognized the importance of this \\WRBUMP\\ to characterize the mode of star formation in starbursts (burst versus continuous mode) and developed the first method to empirically derive the WR/O ratio. The blend of the \\ion{C}{iii} $\\lambda$5698 and \\ion{C}{iv} $\\lambda$5808 broad emission lines constitutes the {\\bf red \\WRBUMP}. \\ion{C}{iv} $\\lambda$5808 is the strongest emission line in WC stars, but it is barely seen in WN stars. The first detections of the red \\WRBUMP\\ in integrated galaxy spectra were reported by \\citet{KS86} and \\citet{Dinerstein86}, but it seems to be more difficult to detect than the blue \\WRBUMP\\ and it is always weaker than the blue \\WRBUMP\\ \\citep{GIT00,FCCG04}. Actually, the red \\WRBUMP\\ is rarely observed at low metallicities. However, both WN and WC stars have been detected in the lowest metallicity ($Z\\sim$1/50~\\Zo) galaxy known, I~Zw~18 \\citep{Izotov97,Legrand97,DeMello98}.  Although the main feature of WR galaxies is a broad \\HeII\\ $\\lambda$4686 emission line, a considerable fraction of objects also show the nebular (narrow) \\HeII\\ $\\lambda$4686 line. This emission line is very rarely found in Galactic \\HII regions \\citep{Garnett91,Schaerer97} but common in planetary nebulae, and has been sometimes observed in giant \\HII regions where no WR stars have been detected. Sources with sufficient photons of energy $>$54 eV are needed to produce the nebular \\HeII\\ emission line, so its origin is still puzzling \\citep{Garnett91,G04}. Some observations concluded that this nebular line is intimately linked with the appearance of hot WR stars \\citep{Schaerer96,SV98,CH06} but recent analyses suggest that O stars may also contribute to the \\ion{He}{ii} ionizing flux at low metallicities (Brichmann, Kunth \\& Durret 2008, hereafter BKD08). In one way or in another, studies of the nebular \\HeII\\ emission line give clues about the physical processes in \\HII regions, the ionizing fluxes of starbursts and their contribution to the ionization of the intergalactic medium (Garnett et al. 1991; Schaerer et al. 1998; Stasi\\'nska 1998). This is the reason why \\citet*{SCP99} also compiled all the extragalactic objects (54) showing the nebular \\HeII\\ $\\lambda$4686 emission line in their catalog of WR galaxies. The analysis performed by \\citet{TI05} listed 465 Blue Compact Dwarf galaxies (\\BCDG s) showing nebular \\ion{He}{ii} $\\lambda$4686 emission.   Wolf-Rayet stars can be individually observed in our Galaxy, the Magellanic Clouds, in some galaxies of the Local Group (M~31, M~33, IC~10, NGC~6822, IC~1613) and in the brightest galaxies of the Sculptor group (NGC~300 and NGC~55). A recent review of the physical properties of WR stars was presented by \\citet{Crowther07}. Thanks to the analysis of these individual WR stars a quantitative classification in WN, WC and WO stars could be established. This led the developing of stellar population synthesis models tailored to the analysis of massive star populations in young starburst galaxies. The best example of these models are those presented by Schaerer \\& Vacca (1998; hereafter SV98), which provide detailed predictions of many stellar and nebular UV/optical features. Famous \\STARBURST\\ models \\citep{L99,VL05} also include predictions for the WR and O stars populations. Except for some very low metallicity objects, a good agreement between the observations and the synthesis models of SV98 is usually found. For objects with subsolar metallicity (mostly \\BCDG s), this comparison indicates fairly short timescales of star formation, with burst ages lower than 2--4 Myr, which is the initial mass function (IMF) compatible with a Salpeter slope, but requires the existence of high-mass stars (Guseva et al. 2000; De Mello et al. 1998; Mas-Hesse \\& Kunth 1999, Fernandes et al. 2004; Buckalew, Kobulnicky \\& Dufour 2005). However, some authors \\citep{Zhang07} suggested a dependence of the slope of the IMF on metallicity. The timescales of star formation in high-metallicity enviroments are more extended than those found at low metallicity, with burst durations of \\mbox{4--10~Myr} \\citep{SGIT00,FCCG04}, although observations are also explained assuming a superposition of several bursts.   We performed a detailed photometric and spectroscopic analysis of a sample of 20 WR galaxies. Our main aim is the study of the formation of massive stars in starburst galaxies and the role that the interactions with or between dwarf galaxies and/or low surface brightness objects have in its triggering mechanism. In Paper~I \\citep{LSE08} we explained the motivation of this work, compiled the list of the analyzed WR galaxies (Table~1 of Paper~I) and presented the results of optical/\\NIR\\ broad-band and \\Ha\\ photometry. In Paper~II \\citep{LSE09} we presented the results of the analysis of intermediate-resolution long-slit spectroscopy of 16 objects of our sample of WR galaxies --the results for the other four objects were published separately. In many cases, two or more slit positions were used to analyze the most interesting zones, knots or morphological structures belonging to each galaxy or even surrounding objects. In this paper we analyze the localization and the properties of the O and WR populations within those galaxies. Papers~IV and V will present the global multiwavelength analysis of our WR galaxy sample considering all available data.  In Sect.~2 we describe the detection of the WR features. The procedure used to fit the broad and nebular emission lines in the blue and red \\WRBUMP s is explained in Sect.~3. We analyze the detection of the nebular \\ion{He}{ii} $\\lambda$4686 line in Sect.~4. The quantitative analysis of the WNL and WCE star populations in described in Sect.~5. In Sect.~6 we analyze the metallicity dependence of the WR/(WR+O) ratio. The comparison with \\citet{SV98} and \\STARBURST\\ \\citep{L99,VL05} models is shown in Sect.~7. We analyze the derived WCE/WNL ratio in Sect.~8, while the dependence of the IMF on the WR/(WR+O) ratio is discussed in Sect.~9. Finally, we compile our main conclusions in Sect.~10. ",
        "conclusions": "We presented a detailed analysis of the broad stellar features originated by winds of Wolf-Rayet stars in a sample of starbursting galaxies previously classified as WR galaxies. The photometric properties of these galaxies were compiled in Paper~I, while their spectroscopic analysis was shown in Paper~II. The metallicity of these galaxies lies between 7.58 and 8.75 --in units of \\abox--. Our goals are to locate the WR-rich stellar clusters, derive the number of O, WN and WC stars and compare the results with previous observations and the predictions of evolutionary synthesis models. Our main results are the following:  \\begin{enumerate} \\item The blue \\WRBUMP, mainly composed by the broad, stellar \\ion{He}{ii} $\\lambda$4686 emission line and attributed to WN stars, is unambiguously detected in 12 regions and it is probably observed in eight other regions. Aperture effects and/or the position of the slit within the starburst are clearly playing a fundamental role in the detection of the WR features. We consider that multi-aperture spectroscopy or narrow-band imagery using filters sensitive to WR features and the adjacent continuum later followed by slit spectroscopy may be more appropriate for this kind of analysis. The advent of the optical 3D spectroscopy is already showing the localization of WR features with star-forming regions (i.e. Kehrig et al. 2008; Garc\\'{\\i}a-Lorenzo et al. 2008; James et al. 2009,2010; L\\'opez-S\\'anchez et al. 2010). \\item The broad \\ion{C}{iv} $\\lambda$5808 emission line (red \\WRBUMP), attributed to WC stars, is clearly detected in four regions, and it is probably observed in two other objects. \\item We have performed a detailed fitting of the spectra considering the broad stellar and narrow nebular emission lines in both the blue and the red \\WRBUMP s. In the blue \\WRBUMP\\ we usually found a contribution of the nebular and broad \\ion{He}{ii} $\\lambda$4686 components. \\item The nebular \\ion{He}{ii} $\\lambda$4686 is detected in 17 regions and it is probably observed in six other knots. We confirm the trend that galaxies with lower metallicities have a higher $I$(nebular \\ion{He}{ii} $\\lambda$4686)/$I$(\\Hb) ratio. Some of our lowest-metallicity regions show the nebular but not the broad \\ion{He}{ii} emission line. The SV98 models seem to underpredict the intensity of this emission line in low metallicity objects.  Both findings agree with the conclusions reached by BKD08, who suggested that O stars and not only WR stars contribute to the ionization of the He$^+$ at low metallicities. \\item We used the intensity ratios of the broad \\ion{He}{ii} $\\lambda$4686 and \\ion{C}{iv} $\\lambda$5808 lines to derive the number of O, WN and WC stars within each region. We have considered two methods: (i) assuming constant luminosities for the WR features, as it is commonly done and it used in theoretical models, and (ii) assuming metallicity-dependent luminosities for the WR features, as both recent observations and new WR stars models indicate (CH06). This second method gives higher number of WR stars at lower metallicities. We derive an empirical estimation of the WNL/(WNL+O) ratio using the intensity of the broad \\ion{He}{ii} $\\lambda$4686 line assuming the metallicity-dependent of the WR luminosities. \\item As expected following theoretical models, the total number of WR stars increases with increasing metallicity. However, we observe a possible flattening of the WR/(WR+O) ratio for objects with \\abox$<$8.2, because all the low-metallicity galaxies show a constant value of $\\sim$0.03 -- 0.04, but more data are needed to confirm this trend. We also conclude that the contribution of the WCE stars is not negligible at low metallicities. \\item The comparison of the observational data with the theoretical predictions given by the most accurate models available to date, SV98 and SB99, is difficult because of all the parameters involved, namely, age, metallicity, star-formation history, assumption of the IMF and WR stars properties (variation of the WR luminosities with the metallicity, effect of star rotation, contribution of the WR binnary channel). New models including all these factors are absolutely needed to perform an appropriate comparison with the observational data. The available SV98 and SB99 models seem to work better at higher metallicities, but more complex star-formation histories are needed to explain the observed WR luminosities and ratios in this regime. At low metallicities the SV98 and SB99 models fail to reproduce the observed WR intensities because these models do not consider the decreasing of the WR luminosities with decreasing metallicity. Our data agree with the SV98 models considering a Salpeter-like IMF for all metallicites. \\item Finally, the WCE/WNL ratios observed in our starbursting galaxies are quite different than those empirically found in nearby star-forming and irregular galaxies. We consider that this effect is a consequence of the very different star-formation histories that these objects are experiencing. \\end{enumerate}"
    },
    "1004/1004.2262_arXiv.txt": {
        "abstract": "We present the spectroscopic discovery of a broad-lined Type Ic supernova (SN~2010bh) associated with the nearby long-duration gamma-ray burst (GRB) 100316D.  At $z$ = 0.0593, this is the third-nearest GRB-SN.  Nightly optical spectra obtained with the Magellan telescopes during the first week after explosion reveal the gradual emergence of very broad spectral features superposed on a blue continuum.  The supernova features are typical of broad-lined SNe~Ic and are generally consistent with previous supernovae associated with low-redshift GRBs.  However, the inferred velocities of \\bh\\ at 21 days after explosion are a factor of $\\sim2$ times larger than those of the prototypical \\bw\\ at similar epochs, with $v \\approx$ 26,000~\\kms, indicating a larger explosion energy or a different ejecta structure.  A near-infrared spectrum taken 13.8 days after explosion shows no strong evidence for \\ion{He}{1} at 1.083~\\micron, implying that the progenitor was largely stripped of its helium envelope.  The host galaxy is of low luminosity (M$_R \\approx -18.5$~mag) and low metallicity ($Z \\lesssim 0.4$ Z$_{\\odot}$), similar to the hosts of other low-redshift GRB-SNe. ",
        "introduction": "While the connection between long-duration gamma-ray bursts (GRBs) and Type Ic supernovae (SNe Ic) has been established in broad terms, the details remain poorly understood \\citepeg{wb06}.  Of particular importance is identifying the physical parameter(s) that distinguish the $0.1-1\\%$ of SNe Ic which give rise to GRBs from those that do not.  Related to this question is the diversity within the GRB-SN sample itself.  Progress requires detailed spectroscopic observations, which limit in-depth studies of GRB-SNe to the nearest events ($z\\lesssim 0.2$).  The low detection rate of such events, just $\\sim 0.3$ yr$^{-1}$, has hampered observational progress.  Over the past 12 years, detailed spectroscopy has been obtained for just four GRB-SNe: 980425/1998bw \\citep{galama98,patat01}, 030329/2003dh \\citep{hjorth03,stanek03,matheson03}, 031203/2003lw \\citep{malesani04,avishay04}, and 060218/2006aj \\citep{modjaz06,mirabal06,pian06}. In each case, the spectra revealed unusually broad features, indicative of large ejecta velocities that are faster than 90-95\\% of all ordinary SNe Ic \\citep{pod04}.  Recently, the broad-lined SN\\,2009bb \\citep{max09} was also shown to produce relativistic ejecta typical of GRBs \\citep{sod09bb}.  Joining this small set of events, \\grb\\ was detected as an image trigger with the \\swift\\ Burst Alert Telescope (BAT) on 2010 March 16.531 UT \\citep{gcn10496}. The combined pre-trigger BAT survey data and BAT triggered observations revealed a peculiar GRB with a long duration ($\\sim2300$~s), unusual light curve shape, and a soft spectrum \\citep{gcn10524}.  A slowly-varying X-ray counterpart was detected with the {\\it Swift} X-ray Telescope (XRT; Starling et al. 2010), but no variable source was detected with the UV/Optical Telescope in early observations \\citep{gcn10496,gcn10520}.  These peculiarities resembled the unusual behavior of the low-redshift XRF 060218 \\citep{gcn10511} and indeed a galaxy at $z=0.059$ was found inside the XRT error circle \\citep{gcn10511,gcn10512}.  Follow-up observations with the Gemini-South 8-m telescope revealed a brightening optical counterpart at $\\Delta t\\approx 1.5-2.5$ d after the burst, suggestive of an emerging SN \\citep{levan_gcn,wiersema_gcn}.  In this Letter, we present our spectroscopic discovery of a SN associated with GRB\\,100316D \\citep{chornockgcn}, subsequently designated SN\\,2010bh \\citep{bufano_cbet,chornockcbet}, and describe its spectral evolution using extensive optical/near-IR observations spanning $1.47-22$ days after the burst.  As in the case of previous GRB-SNe, our spectra reveal that SN\\,2010bh is a broad-lined SN~Ic.  We also find no evidence for helium.  The photospheric velocity at late time is significantly higher than even SN\\,1998bw. Finally, we show that the metallicity of the explosion site is $Z\\lesssim 0.4$ Z$_\\odot$, in line with the other nearby GRB-SNe host galaxies. ",
        "conclusions": "The spectra presented in this Letter provide conclusive evidence that the $z=0.0593$ GRB 100316D was accompanied by the broad-lined SN~Ic 2010bh.  We highlight the following basic properties:  \\begin{enumerate}  \\item Spectral features from \\bh\\ were detected beginning 2.33 days after the GRB, one of the earliest such detections in a GRB-SN.  \\item The \\ion{Si}{2} velocities on day 21 were twice as large as \\bw.  \\item The velocity evolution was slower in \\bh\\ than in other GRB-SNe, possibly indicating diversity in the ejecta structure of these events.  \\item We find no evidence for helium, indicating a highly-stripped progenitor star.  \\item The SN is superposed on a star-forming region with a low-metallicity ($Z \\lesssim 0.4 Z_{\\odot}$) in a low-luminosity host galaxy.  \\end{enumerate}  Continued observations to the nebular phase will shed further light on the ejecta velocity and geometry."
    },
    "1004/1004.2748_arXiv.txt": {
        "abstract": "{Short-period high-amplitude pulsating stars of Population I ($\\delta$~Sct stars) and II (SX~Phe variables) exist in the lower part of the classical (Cepheid) instability strip. Most of them have very simple pulsational behaviours, only one or two radial modes being excited. Nevertheless, BL~Cam is a unique object among them, being an extreme metal-deficient field high-amplitude SX~Phe variable with a large number of frequencies. Based on a frequency analysis, a pulsational interpretation was previously given.} {We attempt to interpret the long-term behaviour of the residuals that were not taken into account in the previous Observed-Calculated (O-C) short-term analyses.} {An investigation of the O-C times has been carried out, using a data set based on the previous published times of light maxima, largely enriched by those obtained during an intensive multisite photometric campaign of BL~Cam lasting several months.} {In addition to a positive (161 $\\pm$ 3) x 10$^{-9}$~yr$^{-1}$ secular relative increase in the main pulsation period of BL~Cam, we detected in the O-C data short- (144.2 d) and long-term ($\\sim$ 3400 d) variations, both incompatible with a scenario of stellar evolution.} {Interpreted as a light travel-time effect, the short-term O-C variation is indicative of a massive stellar component (0.46 to 1 M$_{\\sun}$) with a short period orbit (144.2 d), within a distance of 0.7 AU from the primary. More observations are needed to confirm the long-term O-C variations: if they were also to be caused by a light travel-time effect, they could be interpreted in terms of a third component, in this case probably a brown dwarf star ($\\geq$ 0.03 \\ M$_{\\sun}$), orbiting in $\\sim$ 3400 d at a distance of 4.5 AU from the primary.} ",
        "introduction": "SX~Phe-type stars of the Galactic field are high-amplitude Population II pulsators located in the lower part of the classical (Cepheid) instability strip, close to the main sequence, among the high-amplitude Population I $\\delta$ Scuti variables (e.g., Nemec \\& Mateo 1990; Breger 2000; Rodr\\'{i}guez 2003). Only a few (14) field SX~Phe-type stars have been found to date. However, more than one hundred SX~Phe variables have been discovered among the blue straggler population of globular clusters, and in nearby dwarf galaxies (e.g., Rodr\\'{i}guez \\& L\\'opez-Gonz\\'alez 2000; Jeon et al. 2003, 2004; Mazur et al. 2003).   Multiperiodicity with more than two independent modes is not a common feature among the high-amplitude pulsators of the low instability strip. However, BL~Cam (= GD 428, $\\alpha_{2000}$ = 3 h 47 m 19 s, $\\delta_{2000} = + 63\\degr 22\\arcmin  07\\arcsec$, $<$V$>$ $\\sim$ 13.1 mag, $\\Delta V=0.33$ mag) is one of these very few exceptions. In addition to its main pulsation frequency $f_0 = 25.5769$ cycles per day (d$^{-1}$), the multiperiodicity of this extreme metal-deficient ([Fe/H] = -2.4 dex; McNamara 1997) field SX~Phe-type star has been claimed by different authors (cf. Fauvaud et al. 2006  for a review, hereafter F06). Analysis of the main period temporal changes of BL~Cam has also revealed a secular increase in its main period and the possibility that this star is a member of a detached binary system (F06).  Since BL~Cam is a very attractive target for observational asteroseismology, we initiated the first multisite photometric campaign for this object to investigate in detail its pulsational behaviour. The results of a frequency analysis (Rodr\\'{i}guez et al. 2007, hereafter R07) has largely allowed the reduction of the daily alias and provided the true value of the first overtone frequency (32.6~d$^{-1}$). R07 confirmed the existence of a very dense pulsational frequency spectrum in this star, in addition to the already known high-amplitude main periodicity.  In this paper, we present a new analysis of the main period variation of BL~Cam by means of the Observed-Calculated (O-C) diagram of the O-C times. From a general point of view, the O-C method (see, e.g., Sterken 2005a) allows us to assess the consistency between the measurement of an observable phenomenon and its prediction, and the possible importance of the light travel-time (LTT) effect in the O-C diagram. A historical example of this method is the discovery made by Ola\\\"{u}s R\\\"{o}mer, in 1676, of the finite velocity of light, obtained observing the eclipse phenomena of Jupiter's Io satellite (D\\'{e}barbat 1978; Sterken 2005b). The variable LTT effect, caused by the different positions of the Earth on its orbit, makes it necessary to convert Julian dates into Heliocentric Julian dates. In variable star studies, the analysis of the O-C diagram is a widely used tool.  The paper is organized as follows. The observations and data reduction are presented in Sect. 2. A new ephemeris and a description of the O-C diagram are given in Sect. 3. An interpretation of the O-C values is proposed in Sect. 4, and Sect. 5 is devoted to a discussion. A conclusion and perspectives are given in Sect. 6. ",
        "conclusions": "From 2005 to 2009, an intensive photometric campaign has been dedicated to the SX~Phe-type star BL~Cam. This campaign has allowed a large amount of observations to be gathered by amateurs and professionals from Europe and North-America. This multisite campaign has led to a new pulsation analysis of BL~Cam. An analysis of the O-C diagram has been performed to investigate the main pulsation period of the star. A short period (144.19 $\\pm$ 0.02 d) has been detected in the O-C variations, as well as a possible long-term periodic (3390 $\\pm$ 26 d) variation. We therefore invoked a light travel-time effect to interpret the short-term variation, and proposed the presence of a rather massive stellar companion (from $\\sim$0.5 to 1~M$_{\\sun}$), closer than $\\sim$0.7 AU to BL~Cam, with a short period orbit (144.2 d). In addition, about two long cycles of quasi-periodic fluctuations have been observed over a period of about two decades. This long timescale O-C modulation could also be interpreted as a light travel-time effect, which is indicative of an additional object orbiting BL~Cam. If real, it should be a low mass object ($\\geq$ 0.03 M$_{\\sun}$), very probably a brown dwarf, orbiting in $\\sim$9.3~yr within a $\\sim$4.5 AU distance.  The analysis of the O-C diagram has also allowed us to measure a positive secular increase in the main period at the rate (4.41 $\\pm$ 0.06) x $10^{-10}$ d$^{-1}$ = (161 $\\pm$ 3) x $10^{-9}$ yr$^{-1}$. A continuous photometric monitoring is essential for obtaining a superior knowledge of the oscillations and behaviour of the main pulsation period of BL~Cam, and confirming the three-body system hypothesis.  BL~Cam is quite distant ($\\sim$1000 pc; McNamara 1997), and its companion(s) are too faint and too closely orbiting to be observed by direct imaging techniques. The spectrum of BL~Cam has only a few lines -- mostly of hydrogen (McNamara \\& Feltz 1978; Alvarez \\& Fox-Machado, private comm). From the above orbital parameters, we can infer that the radial velocity curve should have a semi-amplitude of $\\sim$15 km s$^{-1}$. This could be within the reach of modern spectrographs, provided we achieve a high enough signal-to-noise ratio to unambiguously detect and identify lines."
    },
    "1004/1004.0267_arXiv.txt": {
        "abstract": "Motivated by recent simulations of galaxy formation in which protogalaxies acquire their baryonic content through cold accretion, we study the gravitational fragmentation of cold streams flowing into a typical first galaxy. We use a one-zone hydrodynamical model to examine the thermal evolution of the gas flowing into a $10^{8}\\,M_{\\odot}$ DM halo at redshift $z=10$. The goal is to gain an understanding of the expected fragmentation mass scale and thus the characteristic mass of the first population of stars to form by shock fragmentation. Our model accurately describes the chemical and thermal evolution of the gas  as we are specifically concerned with how the chemical abundances and initial conditions of the low density, metal enriched, cold accretion streams that pass an accretion shock alter the cooling properties and tendency to fragment in the post-shock gas. Cold accretion flows are not shock heated at the virial radius but instead flow along high-baryonic-density filaments of the cosmic web and penetrate deep into the host halo of the protogalaxy. In this physical regime, if molecular cooling is absent because of a strong LW background, we find there to be a sharp drop in the fragmentation mass at a metallicity of $Z \\sim10^{-4}\\,Z_{\\odot}$. If, however, $\\text{H}_2$ and HD molecules are present, they dominate the cooling at $T < 10^{4}\\, \\text{K}$, and metallicitiy then has no effect on the fragmentation properties of the cold stream.  For a solar abundance pattern of metallicity, O is the most effective metal coolant throughout the evolution, while for a pair instability supernova (PISN) metallicity yield, $\\text{Si}^{+}$ is the most effective coolant. PISN abundance patterns also exhibit a slightly smaller critical metallicity. Dust grains are not included in our chemical model, but we argue their inclusion would not significantly alter the results. We also find that this physical scenario allows for the formation of stellar clusters and large, $10^{4}\\, M_{\\odot}$ bound fragments, possibly the precursors to globular clusters and supermassive black holes. Lastly, we conclude that the usual assumption of isobaricity for galactic shocks breaks down in gas of a sufficiently high metallicity, suggesting that metal cooling leads to thermal instabilities. ",
        "introduction": "An important open question in cosmology is to identify the conditions under which the first Population II (Pop II) stars could have formed. It has been argued that clusters of predominantly low-mass stars could have arisen in the first galaxies at $z=10$ \\citep{BC02, GB06}. Their emergence straddles the end of the so-called cosmic dark ages, arguably the most uncharted time in the history of the universe \\citep{BarLoeb01, MirES03}. Between the release of the cosmic microwave background (CMB) at $z\\simeq 1,100$ and the highest redshift gamma ray burst (GRB) at $z\\simeq8.2$ \\citep[][]{Salvaterra:09,Tanvir:09}, no direct observations have been made yet, thus rendering theoretical models our best tools for exploring this formative cosmic epoch \\citep[for indirect probes of this era, see][]{GIS00, CHR02, CEN03, FRE05, FRE09}.   The standard cosmological model ($\\Lambda$CDM) predicts the first stars (Pop III) formed at redshifts $z\\simeq20-30$ in $\\sim10^{6}\\,M_{\\odot}$ dark matter `minihalos' \\citep[e.g.,][]{CR86, HAI96, TEG97} and this brought about an end to the cosmic dark ages. The cooling properties of $\\text{H}_{2}$, the most effective primordial gas coolant at temperatures below $10^{4} \\text{ K}$, set a characteristic mass scale of $100~M_\\odot$ for Pop III stars \\citep{ABN00, BCL02}. The primordial initial mass function (IMF) would then have been top-heavy. This result has been confirmed by subsequent studies \\citep[e.g.,][]{OP2003, ONor0700, Y06, Y08}. Further investigations were directed at elucidating their properties and key implications, such as the chemical signature \\citep{HW02, Iwamoto05}, the detailed IMF \\citep{NU01}, the effect on reionization \\citep{WL03, SOK04}, the formation rate \\citep{TS09}, and the tendency toward binarity \\cite[][]{TUR09, STA09}. These massive Pop~III stars had extremely hot photospheres \\citep[$T_{\\rm eff} \\sim 10^5\\,\\text{K}$,  see][]{BK01, SCH02}, and very short lifetimes of $\\sim$ 3 Myr \\citep[for the case against this standard picture see][]{TUM04, SL06, CLA10}.  It is still not known with any certainty when and how the universe shifted its star formation mode from the top-heavy IMF for Pop III to the roughly Salpeter distribution seen today. \\citet{BRO01} suggested it is the chemical input from Pop III stars that allowed gas in minihalos halos to cool more efficiently and fragment at a smaller mass scale. They found that when molecular cooling is excluded by a soft-UV photodissociating background, gas enriched to $Z_{\\rm crit} \\simeq 10^{-3.5}\\,Z_{\\odot}$, the so-called critical metallicity, fragments into smaller mass clumps. The idea of metallicity being the main driver of the mode transition was advanced by many other studies. \\citet{SS06} found individual critical metallicities for C, Fe, Si, and O and suggested that for densities and temperatures conducive to Pop III star formation, these species contribute more to the cooling rate than atomic hydrogen or $\\text{H}_2$. \\citet{SMI09} found that there exists a critical metallicity window, due in part to the CMB temperature floor, which regulates the transition. In a series of papers \\citep{OMU00, OMU05,OMU08} dust was shown to have a significant impact on the Pop III/II transition. Its inclusion lowers the critical metallicity derived from pure metal line cooling by several orders of magnitude. Dust grains act as a catalyst for molecule formation and radiate thermally, however these effects generally become important late in the pre-stellar gas contraction phase when density is high. \\citet{SO09} argued that the interplay of dust and metal cooling is the key driver, together with the CMB temperature floor. A different view was developed in a series of studies \\citep{JAP07, Jappsen:09a,JAP09}, now taking into account molecular cooling. According to these authors, metallicity plays no role in the fragmentation of collapsing clouds while initial conditions, set by the details of galaxy formation, the CMB temperature, turbulence, and rotation, dominate the Pop. III/II transition.  This star formation mode transition is clearly a complex, gradual process, regulated in part by all of these factors.  Most studies of the critical metallicity to date have focused on $\\sim 10^{6} \\,M_\\odot$ DM halos, or minihalos, as the primordial star formation environment. However, the majority of the first Pop II stars likely formed in deeper potential wells, possibly in halos that can cool via atomic hydrogen cooling \\citep[e.g.,][]{Betal09}. Whereas the quasi-hydrostatic, roughly spherical nature of baryonic collapse inside of minihalos admits modeling of low metallicity cooling in a rather straightforward fashion, the hydrodynamics inside the first galaxies is much more complex. Simulations with realistic cosmological initial conditions have recently demonstrated how cold accretion streams feed dense, turbulent gas to the center of the galaxy, where second generation star formation will take place \\citep[e.g.,][]{GB08,WA07,WA08}. Since detailed simulations of the fragmentation properties in such a cloud with different levels of pre-enrichment are still lacking, it is timely to carry out an exploratory survey of the crucial physics involved in the fragmentation of such shocked, stream-fed gas.  These cosmic filaments play a pivotal role in galaxy formation by adding a third accretion mode to the two processes that have traditionally been thought to partake in galaxy formation. Spheroidal galaxy components were thought to have been built by gas channeling in galactic mergers \\citep{TOO77, WR78} and disk components by accretion from a shock-heated intergalactic medium (IGM) during virialization \\citep{RO77, SIL97}. Hot accretion occurs when gas is shock heated to approximately the virial temperature close to the virial radius of the DM halo; the gas then cools, falls inward, and circularizes in the galactic disk. Recently, however, a new mode of accretion has been found to dominate in high-$\\sigma$ cosmological peak halos at high redshift \\citep{DB06,DEK09}. This mode, coined `cold flow accretion,' occurs along cosmic filaments that feed gas directly to the inner parts of a galaxy, leading to shocks at radii much smaller than the halo virial radius \\citep[see, also,][]{WA07,Binney:04}. This differs qualitatively from hot accretion in that the infalling gas passes through the virial shock unaffected, has a much higher density, and is deposited much closer to the center of the galaxy. Cold accretion is theoretically very appealing. It presents a solution to the observed blue-red galaxy bimodality \\citep{BAL04}. It provides a natural explanation for high redshift galaxies with large star formation rates (SFRs) but no evidence of a major merger \\citep[e.g.,][]{Genzel:06}.  Also, cold accretion naturally arises in many simulations of galaxy formation \\citep{KER05, DB06, WA07, GB08, HAR08,Regan:09}, suggesting this type of accretion is natural to high-$\\sigma$ peak halos at high redshifts within $\\Lambda$CDM models.  DM halos with masses $\\sim 10^{8} \\,M_\\odot$ represent the smallest structures that can be classified as `first galaxies' \\citep{KY05, REA06, WA07, GB08, Betal09}. Halos of this mass have virial temperatures $T_{\\text{vir}} >  10^{4} \\,\\text{K}$ which allows efficient cooling via excitation of atomic hydrogen \\citep[e.g.,][]{Haiman:97,TEG97,MiraldaEscude:98,Barkana:99,Oh:99}. This is the halo mass scale where we first expect to see sustained, possibly self-regulated, star formation, similar to the present-day case. Finally, in $10^{8}\\,M_\\odot$ halos, gas accreting along cosmic filaments will drive supersonic turbulence \\citep{WA07, GB08,Regan:09,KLE10}, fundamentally changing the star formation dynamics \\citep{PN02,MacLow:04, MOst07}. In this work, as in many others, we define a first galaxy to be one of the first DM halos to form with mass $\\geq5\\times10^{7}\\,M_{\\odot}$, thus meeting the requirement that gas heated to the virial temperature is able to cool through hydrogen excitation (or Lyman-$\\alpha$) emission.   A further motivation for this work is understanding the formation of low-luminosity dwarf spheriodal satellite galaixes (dSph) and globular clusters (GCs). These structures have similar luminosities, stellar masses, generally no gas nor recent star formation, and contain metal-poor, old stellar populations. However, GCs are small ($\\sim 1-10\\,\\text{pc} $) and do not seem to contain much DM while dSph galaxies are much larger ($>100\\,\\text{pc} $) and are heavily dominated by DM as indicated by their large mass-to-light ratios \\citep{MAT98,STR08}. It is clear these are two structurally distinct populations \\citep[e.g.,][]{KOR85,BEL07}, but it is necessary to understand the cosmological circumstances of their formation. To answer whether the first galaxies are progenitors to dSph galaxies, GCs, or both requires an understanding of star formation in these objects.   With this in mind, we investigate the fragmentation properties of the metal-pre-enriched cold flow accretion in protogalaxies at high redshift, and explore implications for star formation in these objects. Dense, metal-enriched baryonic streams flow along the filaments of the cosmic web and penetrate deep into a protogalaxy.  The multiple (e.g., ternary) streams collide with each other, or collide with a turbulently-supported gas accumulation that has already collected in the center of the protogalaxy.  Since the stream inflow is supersonic, the colliding gas must pass a strong shock, but then cools rapidly under approximate isobaric conditions.  We identify the characteristic fragmentation mass scale in the shocked, isobarically cooling gas.  We find that if a strong photodissociating Lyman-Werner (LW) background has significantly suppressed molecule formation, the characteristic mass scale depends strongly on the metallicity content of the gas.  The outline of this work is as follows. In Section \\ref{sec:cooling} we describe our methodology for the cooling of shocked, metal-enriched gas.  Section \\ref{sec:initial_conditions} motivates the initial conditions for our model, including metal abundances and the physics of cold accretion. In Section \\ref{sec:thermal_evolution} we discuss our fragmentation criteria and the results from our numerical integrations, and summarize our findings and broader implications in Section \\ref{sec:discussion}.  Throughout the paper we assume the $\\Lambda\\text{CDM}$ cosmological parameters: $\\Omega_{\\rm m} = 0.27$, $\\Omega_{\\Lambda} = 0.73$, and  $h=0.70$, consistent with the 5 yr. WMAP results  \\citep{KOM09}. ",
        "conclusions": "\\label{sec:discussion}  In this work, we have examined the fate of metal-enriched, cold, dense, and dust-free accretion streams into DM halos of mass $10^{8}\\, M_\\odot$ at redshift $z=10$; we regard these objects as the first galaxies. From simulations, we know these cold streams penetrate very close to the center of the protogalaxy, where they shock and fragment. Our goal was to determine the fragmentation properties of these cold streams and to identify the key physical parameters governing this first step in the high redshift star formation process. We find that metallicity (due to metals being an effective low temperature coolant) and LW background strength (as it is capable of destroying $\\text{H}_{2}$ and HD)  are the dominant factors in determining the final fate of the post-shock region. Using a fit to the baryonic growth rate of halos, and assuming, as recent studies suggest, that these galaxies are fed entirely by cold streams and that their star formation rate closely matches the mass accretion rate, we can estimate the number of fragments produced by these flows. We either see one large ($\\sim10^{4}M_{\\odot}$) fragment being produced, or roughly $10^{4}$ solar-mass fragments, depending on the initial conditions of the flow. The CMB temperature floor is also a key factor in determining the mass of fragments, as it represents a strict temperature lower limit of radiatively cooling gas. Roughly speaking, this scenario occuring at higher redshifts would still demonstrate similar behavior in terms of a critical metallicity, but would produce fragments of slightly higher mass. Overall, this effect is of little importance as even low-redshift star forming molecular clouds do not generally cool below $10\\, \\text{K}$, wheres as the CMB temperature at $z=10$ is $\\sim 30\\text\\,{K}$. Any observational effects from this temperature floor wold be difficult to extract from our simulations. \\citet{BAI10} studied CMB regulated star formation in some detail (at lower redshifts than we examine) and find that observations currently cannot confirm or refute CMB regulation on star formation.  We did not include dust, even though it can have a major impact in that it allows very efficent $\\textrm{H}_2$ formation on grain surfaces, though we do argue its presence will not significantly alter our findings. The presence of dust and its effect on cooling has been discussed as a possible catalyst for the star formation mode transition. \\citet{SCH06} and \\citet{SO09} found that dust induced fragmentation can decrease the Pop II fragmentation mass scale by about four orders of magnitude over that accessible with metal fine structure lines only. \\citet{JAP09} suggested that metallicity does not affect the fragmentation of gas with hot initial conditions, similar to the environment probed in this study. We attribute the difference in our conclusions to our examination of a slightly different physical regime and to our inclusion of a soft UV background that stifles $\\textrm{H}_2$ cooling \\citep[as supported, e.g., by][]{SHA09}. Taking guidance from galaxy formation simulations, we follow isobarically evolving gas in an expanding shocked slab to study its thermal evolution and fragmentation. Our model is designed to reflect the filament-dominated inflow topology of cold accretion flows, whereas \\citet{JAP09} examine the collapse into a spherically-symmetric, centrally concentrated halo potential. In our isobarically evolving case, the tight coupling between density and temperature will increase the importance of the role of the cooling rate in determining the final outcome of fragmentation. This aside, if we include molecule formation and molecular cooling ($\\textrm{H}_2$ and HD) by setting $J_{21} = 10$, our results agree with \\citet{Jappsen:09a} in that metal line cooling does not affect the fragmentation properties of the evolving gas. This emphasizes the importance of a LW soft UV background in modeling the formation of Pop~II stars in the first galaxies.  The first galaxies are suspected to be the first structures where turbulence begins to play a major role, driven by high Mach-number cold streams penetrating to the protogalactic core \\citep{GB08, WA07}. It is the interplay of turbulence and self gravity which is suspected to be responsible for the shape of the present-day IMF, especially the high mass power law extension (Padoan \\& Nordlund 2002; Mac Low \\& Klessen 2004; McKee \\& Ostriker 2007). Thus, it is likely that star formation inside the first galaxies will be characterized by a stellar mass distribution that for the first time in the history of the universe resembles already the present-day IMF. Specifically, we can expect turbulence to introduce density perturbations on the order of ${\\cal M}^{2}$, suggesting within our above model that fragments as low as $\\sim0.3 \\,M_\\odot$ might form. Such stars, below the mass limit of $0.8 \\,M_\\odot$ required to survive for the age of the universe, would still exist today, and could be probed by Stellar Archaeology \\citep[][and references therein]{FRE09}. To predict the detailed shape of the Pop. II IMF inside the first galaxies, the source and impact of the turbulence needs to be studied with realistic numerical simulations.   Here, we have constrained some of the key physical processes that govern star formation in high-redshift protogalaxies within an idealized analytical framework that allowed us to explore the relevant parameter space. This work needs to be followed up by detailed simulations, including a number of ingredients ignored here, such as high optical depth effects, H and He ionizing radiation \\citep{RIC02}, turbulence, feedback effects (AGN, stellar outflows, etc.), and possibly cosmic rays from the first SNe \\citep{JCE07, SB07}. It needs to be emphasized that the estimates presented in this paper are derived with the Bonnor-Ebert mass and a simplified treatment of the post-shock dynamics. To truly probe the fragmentation mass scale in the first galaxies one must use multidimensional numerical simulations. The goal of understanding the complex first stages of galaxy formation seems now within reach, using state-of-the-art three dimensional radiation-hydrodynamical simulations, together with realistic cosmological initial conditions. Such calculations are crucial to make useful predictions for the {\\it James Webb Space Telescope (JWST)}."
    },
    "1004/1004.4672_arXiv.txt": {
        "abstract": "Expressing explicitly the parameters of the standard Skyrme interaction in terms of the macroscopic properties of asymmetric nuclear matter, we show in the Skyrme-Hartree-Fock approach that unambiguous correlations exist between observables of finite nuclei and nuclear matter properties. We find that existing data on neutron skin thickness $\\Delta r_{np}$ of Sn isotopes give an important constraint on the symmetry energy $E_{\\text{\\textrm{sym}}}({% \\rho _{0}})$ and its density slope $L$ at saturation density ${\\rho _{0}}$. Combining these constraints with those from recent analyses of isospin diffusion and double neutron/proton ratio in heavy-ion collisions at intermediate energies leads to a more stringent limit on $L$ approximately independent of $E_{\\text{\\textrm{sym}}}({\\rho _{0}})$. The implication of these new constraints on the $\\Delta r_{np}$ of $^{208}$Pb as well as the core-crust transition density and pressure in neutron stars is discussed. ",
        "introduction": "The nuclear symmetry energy $E_{\\text{\\textrm{sym}}}(\\rho )$ that encodes the energy related to the neutron-proton asymmetry in the equation of state (EOS) of isospin asymmetric nuclear matter (ANM) plays a crucial role in both nuclear physics and astrophysics \\cite% {LiBA98,Dan02,Lat04,Ste05,Bar05,LCK08}. It is also relevant to some interesting issues regarding possible new physics beyond the standard model \\cite{Hor01b,Sil05,Kra07,Wen09}. Although significant progress has been made in recent years in determining the density dependence of $E_{\\text{\\textrm{% sym}}}(\\rho )$ \\cite{Bar05,LCK08}, large uncertainties still exist even around the normal density $\\rho _{0}$. For instance, while the value of $E_{% \\text{\\textrm{sym}}}(\\rho_{0} )$ is determined to be around $30\\pm 4$ MeV mostly from analyzing nuclear masses, the extracted density slope $L$ of $E_{% \\text{\\textrm{sym}}}(\\rho )$ at $\\rho _{0}$ scatters in a very large range from about $20$ to $115$ MeV depending on the observables and methods used in the studies \\cite{Tsa09,Car10,She10}. Since many observables in terrestrial laboratory experiments intrinsically depend on both $E_{\\text{% \\textrm{sym}}}(\\rho_{0} )$ and $L$, the extraction of $L$ at an accuracy required for understanding more precisely many important properties of neutron stars \\cite{Lat04,XuJ09} is still severely prohibited, although the uncertainty of $E_{\\text{\\textrm{sym}}}(\\rho_{0} )$ is relatively small. To extract $L$ with higher accuracy is thus of crucial importance.  Theoretically, studies based on both mean-field theories \\cite% {Bro00,Hor01,Fur02,Yos04,Che05b,Tod05,Rei10} and droplet-type models \\cite% {Mye69,Oya03,Dan03,Cen09} have shown that the neutron skin thickness $\\Delta r_{np}=\\langle r_{n}^{2}\\rangle ^{1/2}-\\langle r_{p}^{2}\\rangle ^{1/2} $ of heavy nuclei, given by the difference of their neutron and proton root-mean-squared radii, provides a good probe of $E_{\\text{\\textrm{sym}}% }(\\rho )$. In particular, $\\Delta r_{np}$ has been found to correlate strongly with both $E_{\\text{\\textrm{sym}}}(\\rho_0 )$ and $L$ in microscopic mean-field calculations \\cite{Bro00,Hor01,Fur02,Yos04,Che05b,Tod05,Rei10} using different parameter sets for the nuclear effective interactions, which all fit the binding energies and charge radii of finite nuclei but correspond to different $E_{\\text{\\textrm{sym}}}(\\rho)$ and give different $% \\Delta r_{np}$. It is, however, difficult to extract an accurate value for $% L $ from comparing calculated $\\Delta r_{np}$ of heavy nuclei with experimental data as $\\Delta r_{np}$ depends on a number of nuclear interaction parameters in a highly correlated manner \\cite{Fur02,Rei10} and the calculations have been usually carried out by varying simultaneously the interaction parameters. Similar difficulties also exist when one tries to extract other physical quantities from observables of finite nuclei within mean-field theories or density functional theories~\\cite{Col04}. A well-known example is the Skyrme-Hartree-Fock (SHF) approach using normally $% 9$ interaction parameters. Although experimental data on nucleon-nucleon scatterings and properties of both finite nuclei and infinite nuclear matter would in principle put strong constraints on the combinations of these parameters~\\cite{Vau72,Ton84,Mar02,Yos04,Agr05,Rei06,Cao06}, there is generally no constraint on most of the individual interaction parameters. Instead of varying directly the $9$ interaction parameters within the SHF, we propose here an alternative approach based on a modified Skyrme-Like (MSL) model \\cite{Che09} to express them explicitly in terms of $9$ macroscopic observables that are either experimentally well constrained or empirically well known. This opens the possibility to explore transparently the correlations between properties of finite nuclei and the macroscopic properties of nuclear matter within the SHF approach. In the present work, we use this method to study the correlation between $\\Delta r_{np}$ and various macroscopic observables of infinite nuclear matter by varying individually the values of the latter within their known ranges. We then demonstrate that existing $\\Delta r_{np}$ data on Sn isotopes can give important constraints on $L$ and $E_{\\text{\\textrm{sym}}}(\\rho_0 )$. Combining these constraints with those from recent analyses of isospin diffusion and double neutron/proton ratio in heavy-ion collisions at intermediate energies~\\cite{Tsa09}, we further show that a more stringent limit on $L$ is obtained approximately independent of the value of $E_{\\text{% \\textrm{sym}}}({\\rho _{0}})$. Finally, we discuss the implication of these new constraints on both the $\\Delta r_{np}$ of $^{208}$Pb and the core-crust transition density and pressure in neutron stars. ",
        "conclusions": "\\label{Summary}  We have proposed to analyze the correlation between observables of finite nuclei and some macroscopic properties of asymmetric nuclear matter by expressing explicitly the parameters of the nuclear effective interaction in terms of the macroscopic properties of asymmetric nuclear matter. This would allow us to extract information on some important physical quantities from data on finite nuclei in a more transparent way.  Using such a correlation analysis within the standard SHF approach, we have demonstrated that the neutron skin thickness of heavy nuclei can provide reliable information on the symmetry energy, and the existing neutron skin data on Sn isotopes can give important constraints on the symmetry energy parameters $E_{\\text{\\textrm{sym}}}({\\rho _{0}})$ and $L$. In particular, combining the obtained $L$-$E_{\\text{\\textrm{sym}}}({\\rho _{0}})$ constraints with that from recent analyses of isospin diffusion and double $% n/p$ ratio in heavy-ion collisions has led to a quite accurate value of $% L=58\\pm 18$ MeV approximately independent of the value of $E_{\\text{\\textrm{% sym}}}({\\rho _{0}})$. The obtained $L$-$E_{\\text{\\textrm{sym}}}({\\rho _{0}})$ constraints also put a stringent limit of $\\Delta r_{np}=0.175\\pm 0.02$ fm for the neutron skin thickness of $^{208}$Pb.  Furthermore, we have explored how the core-crust transition density $\\rho _{t}$ and the corresponding pressure $P_{t}$ in neutron stars correlate with the macroscopic properties of asymmetric nuclear matter. Our results have indicated that the $\\rho _{t}$ displays a particularly strong correlation with $L$, a weak dependence on $E_{\\text{% \\textrm{sym}}}({\\rho _{0}})$ and $K_{0}$, but almost no sensitivity to other macroscopic parameters. On the other hand, the $P_{t} $ exhibits a strong correlation with both $L$ and $E_{\\text{\\textrm{sym}}}({\\rho _{0}})$, and a weak dependence on $K_{0} $. The $L$-$E_{\\text{\\textrm{sym}}}({\\rho _{0}})$ constraints obtained in the present work leads to $\\rho _{t}=0.069\\pm 0.018$ fm$^{-3}$ and $P_{t}=0.33\\pm 0.21$ MeV/fm$^{3}$.  Although we have mainly analyzed in the present work the correlation between the neutron skin thickness of finite nuclei and the symmetry energy within the standard SHF approach, our method can be generalized to other correlation analyses or mean-field models. In particular, it will be interesting to see how our results will change if different energy density functions are used. These studies are in progress."
    },
    "1004/1004.1782_arXiv.txt": {
        "abstract": "In  the  Hamiltonian formulation of General Relativity the energy associated to an asymptotically flat  space-time with metric $\\gmn$ is related to the Hamiltonian $H_{\\rm GR}$ by $E=H_{\\rm GR}[\\gmn]-H_{\\rm GR}[\\emn]$, where the subtraction of the flat-space contribution is necessary to get rid of an otherwise divergent boundary term. This classic result indicates that the energy associated to flat space does not gravitate. We apply the same principle to study the effect of zero-point fluctuations of quantum fields in cosmology, proposing that their contribution to the cosmic expansion is obtained computing the vacuum energy of quantum fields in a FRW space-time with Hubble parameter $H(t)$ and subtracting from it the flat-space contribution. Then the term proportional to $\\lc^4$ (where $\\lc$ is the UV cutoff) cancels and  the remaining (bare) value of the vacuum energy density is proportional to $\\Lambda_c^2 H^2(t)$. After renormalization, this produces a renormalized  vacuum energy density $\\sim M^2H^2(t)$, where $M$ is the scale where quantum gravity sets is, so for $M$ of order of the Planck mass a vacuum energy density of the order of the critical density can  be obtained without any fine tuning. The counterterms can be chosen so that the renormalized energy density and pressure satisfy $p=w\\rho$, with $w$ a parameter that can be fixed by comparison to the observed value, so in particular one can chose $w=-1$. An energy density evolving in time as $H^2(t)$  is however observationally excluded as an explanation for the dominant dark energy component which is responsible for the observed acceleration of the universe. We rather propose that  zero-point vacuum fluctuations  provide a new subdominant  ``dark\" contribution to the cosmic expansion that, for a UV scale  $M$ slightly smaller than  the Planck mass, is consistent with existing limits and potentially detectable. ",
        "introduction": "The cosmological constant problem is a basic issue of modern cosmology (see e.g. \\cite {Weinberg:1988cp,Peebles:2002gy,Padmanabhan:2002ji,Copeland:2006wr}). One aspect of the problem is to understand what is the effect of the zero-point fluctuations of quantum fields on the cosmic expansion. Zero-point quantum fluctuations seem to give a contribution\tto the vacuum energy density of order $\\Lambda_c^4$, where $\\lc$ is the UV cutoff.  Even for a cutoff as low as $\\lc\\sim 1\\, {\\rm TeV}$, corresponding to scales where quantum field theory is well tested, this exceeds by many orders of magnitude the value of the critical density of  the universe, which is of order $(10^{-3}\\, {\\rm eV})^4$. This aspect of the problem has by now a long history, which even goes back to works of Nernst and of Pauli in the 1920s~\\cite{Peebles:2002gy}.  It is clear that this $\\Lambda_c^4$ contribution should somehow be subtracted. Renormalization in quantum field theory (QFT) offers a logically viable possibility. One should not forget, in fact, that this cutoff-dependent contribution, as any similar quantity in QFT, is just a bare quantity rather than  a physical observable. Working for instance with a cutoff $\\lc$ in momentum space one finds that, in Minkowski spacetime, the bare, cut-off dependent, contribution to the vacuum energy density is \\be\\label{bare} \\rho_{\\rm bare}(\\lc)=c\\lc^4\\, , \\ee with $c$ some constant that depends on the number and type of fields  of the theory. The standard procedure in QFT is to add to it a cutoff-dependent counterterm, chosen so that it cancels the divergent part and leaves us with the desired finite part (for an elementary discussion of this point in the cosmological constant case, see e.g. the textbook~\\cite{Maggiore:2005qv}), i.e. \\be\\label{rhocount} \\rho_{\\rm count}(\\lc)=-c\\lc^4+\\rho_{\\rm vac}\\, , \\ee where $\\rho_{\\rm vac}\\sim (10^{-3}\\rm eV)^4$ is independent of the cutoff and equal to the observed value of the vacuum energy density (assuming that vacuum energy is indeed responsible for the observed acceleration of the universe). The physical, renormalized, vacuum energy density \\be\\label{eq3v3b} \\rho_{\\rm ren}\\equiv \\rho_{\\rm bare}(\\lc)+\\rho_{\\rm count}(\\lc) \\ee is therefore by definition  independent of the cutoff and equal to the observed value $\\rho_{\\rm vac}$. In principle this procedure is not different from what is done when one renormalizes, say, the electron mass or the electron charge in QED, where again with suitable counterterms one removes the divergent, cut-off dependent, parts and fixes the finite parts so that they agree with the observed values (in this sense the statement, often heard, that QFT gives a {\\em wrong} prediction for the cosmological constant, is not correct. Strictly speaking QFT makes no prediction for the cosmological constant, just as it does not predict the electron mass nor the fine structure constant).    Still, for the vacuum energy this procedure is not really satisfying. The source of uneasiness  is partly due to the fact that the counterterm must be fine tuned to a huge precision in order to cancel the $\\Lambda_c^4$ term and leave us with a much smaller result. In fact, if in \\eq{bare} we take $\\lc$ at least larger than a few TeV, where  quantum field theory has been successfully tested,  $\\rho_{\\rm bare}$ is at least of order $({\\rm TeV})^4=(\\rm 10^{12}\\, eV)^4$, so $\\rho_{\\rm count}$ in \\eq{eq3v3b} must be fine tuned so that it cancels something of order $(\\rm 10^{12}\\, eV)^4$, leaving a result of order $(\\rm 10^{-3}\\, eV)^4$. This fine tuning becomes even much worse if one dares to take $\\lc$ of the order of the Planck mass $\\mpl\\sim 10^{19}$~GeV. The fact that the counterterm must be tuned to such a precision creates a naturalness problem. Here, however,  one might argue that neither the bare term  nor the counterterm  have any physical meaning and only their sum is physical, so this fine tuning is  different from an unplausible cancellation between observable quantities. The same kind of cancellation appears, for instance, when one computes the Casimir effect. The crucial point, however, is that this renormalization procedure leaves us with no clue as to the physical value that emerges from this cancellation, so it gives no explanation of why the energy density associated to the cosmological constant appears to have just a value of the order of the critical density of the universe at the present epoch.  The point of view that we develop in this paper is that, even if renormalization must be an ingredient for understanding the physical effects of vacuum fluctuations, it is not the end of the story. The Casimir effect mentioned above gives indeed a first hint of what could be the missing ingredient for a correct treatment of vacuum energies in cosmology. In the Casimir effect the quantity that gives rise to observable effects,  which have indeed been detected experimentally, is the difference between the vacuum energy in a given geometry (e.g., for the electromagnetic field, between two parallel conducting plates) and the vacuum energy computed in  a reference geometry, which is just  flat space-time in an infinite volume. Both terms are separately divergent as $\\lc^4$, but their difference is finite and observable. This  can suggest that, to obtain the physical effect of the vacuum energy density on the expansion of the universe, one should similarly compute the vacuum energy density in a FRW space-time and subtract from it the value of a reference space-time, which is naturally taken as Minkowski space.  A possible objection to this procedure could be that General Relativity requires that any form of energy should be a source for the gravitational field, which seems to imply that even the vacuum energy associated to flat space should contribute. A more careful look at the formalism of GR shows however that the issue is not so simple and that, in a sense, the subtraction that we are advocating is in fact already part of the standard tenets of classical GR. To define carefully the energy associated to a field configuration in GR, it is convenient to use the Hamiltonian formulation, which goes back to the classic paper by Arnowitt,  Deser and Misner \\cite{ADM} (ADM). As we will review in more detail in Section~\\ref{sect:ADM}, in order to define the Hamiltonian of GR one must at first work in a finite three-dimensional volume, and then the Hamiltonian takes the form \\be\\label{HHH} H_{\\rm GR}=H_{\\rm bulk}+H_{\\rm boundary}\\, , \\ee where $H_{\\rm bulk}$ is given by an integral over the three-dimensional finite spatial volume at fixed time, and $H_{\\rm boundary}$ by an integral over its (two-dimensional) boundary. At this point one would like to define the energy of a classical field configuration as the value of this Hamiltonian, evaluated on the classical solution, but one finds both a surprise and a difficulty. The ``surprise\" (which actually is just a consequence of the invariance under diffeomorphisms) is that $H_{\\rm bulk}$, evaluated on any classical solution of the equations of motion, vanishes, so the whole contribution comes from the boundary term. The difficulty is that the boundary term, evaluated on any asymptotically flat metric (including flat space-time) diverges when the boundary is taken to infinity. The solution proposed by ADM is to subtract from this boundary term, evaluated on an asymptotically flat space-time with metric $\\gmn$, the same term computed in flat space-time. The corresponding energy (or mass) is finite and is known as the ADM mass, and provides the standard definition of mass in GR. For instance, when applied to the \\Sch space-time, the ADM mass computed in this way turns out to be equal to the mass  $M$ that appears in the \\Sch metric.  The ADM prescription can be summarized by saying that, in GR, the energy $E$ associated to an asymptotically flat space-time with metric $\\gmn$ can be obtained from the Hamiltonian $H_{\\rm GR}$  by \\be\\label{EHH} E=H_{\\rm GR}[\\gmn]-H_{\\rm GR}[\\emn]\\, , \\ee where $\\emn$ is the flat metric. Even if the context in which this formula is valid, namely asymptotically flat space-times, is  different from the cosmological context in which we are interested here, still \\eq{EHH} suffices to make the point that the intuitive idea that GR requires that any form of energy acts as a source for the gravitational field is not really correct. \\Eq{EHH} tells us that the energy associated to Minkowski space does not gravitate.  The idea of this paper is to generalize \\eq{EHH} to the case of zero-point quantum fluctuation in curved space, proposing that the effect of zero-point fluctuations  on the cosmic expansion should be obtained by computing the vacuum energy of quantum fields in a FRW space-time with Hubble parameter $H(t)$ and subtracting from it the flat-space contribution. Computed in this way, the physical effect of zero-point fluctuations on the cosmic expansion can be seen as a sort of ``cosmological Casimir effect\": while in the standard Casimir effect one computes the vacuum energy in a given geometry (say, for the electromagnetic field, between two infinite parallel conducting planes) and subtracts from it the value computed in  a reference geometry (flat space-time in  an infinite volume), here we compute the vacuum energies of the various fields in a given curved space-time, e.g. in FRW space-time, and we subtract from it the value computed in a reference space-time, i.e. Minkowski.  A possible objection to the idea of applying \\eq{EHH} to vacuum fluctuations could be that one might think that, after all, the structure of UV divergences is determined by local properties of the theory, while whether a space-time is Minkowski is a global question. However one should keep in mind that, even if $\\Tmn(x)$ is a local quantity, the vacuum expectation value $\\cav\\Tmn\\vac$ in a curved background  involves the global aspects of the space-time in which it is computed. This is due to the fact  that $\\cav\\Tmn\\vac$ requires a definition of the vacuum state $\\vac$. The vacuum is  defined from the condition $a_{\\bf k}\\vac=0$, where the annihilation operators $a_{\\bf k}$ are defined with respect to a set of mode functions $\\phi_{\\bf k}(t)$. The mode functions, in turn, are obtained solving a wave equation over the whole space-time, and therefore are sensitive to global properties of the space-time itself. In particular, in a FRW background the mode functions   depend on the scale factor so their time derivatives (which enter in the computation of $\\cav\\Tmn\\vac$) depend on the Hubble parameter $H(t)$. As we will review below,  this results in the fact that  the quadratic divergence in $\\cav\\Tmn\\vac$ depends on the expansion rate $H(t)$ of the FRW background.   The paper is organized as follows. In Section~\\ref{sect:ADM} we recall how \\eq{EHH} is derived in the ADM formalism. The reader familiar with the subject, or not interested in the derivation, might simply wish to move directly to Section~\\ref{sect:appl}, where we apply this classical subtraction, together with standard renormalization theory, to zero-point vacuum fluctuations. Some cosmological consequences of our proposal are discussed in Section~\\ref{sect:impl}, while Section~\\ref{sect:Concl} contains our conclusions. ",
        "conclusions": "\\label{sect:Concl}  One aspect of the cosmological constant problem, or more generally of the problem of understanding the origin of dark energy, is to understand why zero-point fluctuations of quantum fields do not produce an energy density of the order of $M^4$, where $M$ is the UV mass scale  of the quantum field theory (e.g. the Planck mass, or the string mass scale), despite the fact that this seems to be the natural value suggested by quantum field theory. We have proposed that the solutions  to this long-standing puzzle has a purely classical origin, and is related to the correct definition of energy in classical General Relativity, which already involves the subtraction of the flat-space contribution, see \\eq{EHH}.  We have applied this subtraction procedure to a FRW space-time with Hubble parameter $H(t)$ and we have found that the remaining energy density, after renormalization, has a ``natural\" value proportional to $M^2H^2(t)$ (and a sign that could in principle be either positive or negative, just as in the Casimir effect). For $M\\simeq \\mpl$ this gives an energy density just of the order of the critical density of the universe. As we have discussed, however, such an energy density has a time dependence that is not compatible with present observations, if we identify it with the dark energy component with $\\Omega_{\\Lambda}\\simeq 0.7$ which  in the standard ${\\Lambda}$CDM cosmology is responsible for the observed acceleration of the universe. It is however possible that it represents a new form of dark energy, whose normalized energy density today, $\\Omega_Z$, is smaller than $\\Omega_{\\Lambda}$. Values of $|\\Omega_Z|\\lsim {\\,\\,\\rm a\\,\\, few}\\times 10^{-3}$  are compatible with the observations that we have discussed, but could  give observable effects in more detailed studies that make use of the specific signature of zero-point fluctuations, namely an energy density with a time dependence proportional to $H^2(t)$, as well as in future more accurate cosmological observations.   \\vspace{5mm} \\noindent {\\bf Acknowledgments}. I thank Ruth Durrer,  Juan Garcia-Bellido, Alberto Nicolis, Massimiliano Rinaldi, Antonio Riotto, Joan Sol\\`a and the referees for very useful comments on the manuscript. This work is supported by the Fonds National Suisse.   \\appendix"
    },
    "1004/1004.3322_arXiv.txt": {
        "abstract": "We investigate the impact of mergers on the mass estimation of galaxy clusters using $N$-body + hydrodynamical simulation data. We estimate virial mass from these data and compare it with real mass. When the smaller subcluster's mass is larger than a quarter of that of the larger one, virial mass can be larger than twice of the real mass. The results strongly depend on the observational directions, because of anisotropic velocity distribution of the member galaxies. We also make the X-ray surface brightness and spectroscopic-like temperature maps from the simulation data. The mass profile is estimated from these data on the assumption of hydrostatic equilibrium. In general, mass estimation with X-ray data gives us better results than virial mass estimation. The dependence upon observational directions is weaker than in case of virial mass estimation. When the system is observed along the collision axis, the projected mass tends to be underestimated. This fact should be noted especially when the virial and/or X-ray mass are compared with gravitational lensing results. ",
        "introduction": "Mass is one of the most important physical parameters to characterize astrophysical objects. This is especially true in many kinds of self-gravitating objects such as star clusters, galaxies, and clusters of galaxies, where gravity plays a crucial role in their evolution. In addition, mass distribution in large scales in the universe have some information about dark matter (DM). For example, mass distribution in the very central part of relaxed dark halos could be altered by the self-interaction of DM particles \\citep{Yosh00}. Investigating spatial distributions of DM, galaxies, and intracluster medium (ICM) also gives us important clues of DM properties \\citep{Clow04, Brad06, Jee07, Okab08}.  There are several methods to estimate mass distribution of galaxy clusters in an observational way. Assuming dynamical equilibrium, we can estimate cluster mass from line-of-sight velocity distribution of member galaxies through virial theorem or Jeans equation \\citep{Kent82, Oege94, Gira98, Barr07, Maur08}. Cluster mass distribution can also be estimated from the X-ray observational data on the assumption of hydrostatic equilibrium of the ICM \\citep{Reip02, Taki03, Ota04, Vikh06, Gast07}. Weak and strong gravitational lensing techniques enable us to determine gravitational mass directly without any assumptions of dynamical status of the system \\citep{Kais93, Mira95}.  We have multiple observational ways of cluster mass estimation as mentioned above. However, these methods do not always give us consistent results \\citep{Mahd08}. In case of well-known gravitational lensing cluster CL 0024+17, for example, mass derived from X-ray data \\citep{Ota04} and those from gravitational lensing \\citep{Tyso98, Broa00, Jee07} are significantly different, where the latter results are larger than the former by a factor of two or three. Recently, \\citet{Zhan10} made a comparison of X-ray and lensing mass for 12 clusters, and showed that mass of disturbed clusters could be both overestimated and underestimated.   Some assumptions are necessary in the mass estimation method presented above. In the method with line-of-sight galaxy velocities, we assume that the system is in dynamical equilibrium and often assume isotropic velocity distribution of the member galaxies and spherical symmetric structure. In addition, another critical aspect in this method lies in the difficulty of removing foreground and background galaxies, which significantly affect the estimate of the velocity dispersion and mass \\citep{Bivi06}. In the X-ray method we assume that the systems are in hydrostatic equilibrium and spherically symmetric. Simple methods with strong lensing often assume axial symmetry. Even more elaborate methods based on strong lensing technique are essentially model-dependent. Although the weak gravitational lensing method seems to be free from these assumptions about dynamical state and geometry, it does not provides us with three-dimensional mass density distribution but two-dimensional surface-mass density one. This fact should be noted especially when the lensing results are compared with those derived from other methods.  A lot of $N$-body + hydrodynamical simulations with cosmological initial conditions have been done so far to investigate accuracy of cluster mass estimation \\citep{Evra96, Rasi06, Kay07, Naga07, Piff08, Vent08, Ameg09}. It is true that such studies are useful in order to investigate this issue in realistic situations. However, such simulations are not suitable to model a single merger event because it is difficult to control initial conditions and achieve very high resolution in particular interesting regions efficiently. Therefore, studies based on binary merger simulations are complementary and worthwhile, in which it is relatively easy to keep initial conditions under control, improve numerical resolution, and derive clear physical interpretations. The assumptions for mass estimation listed above are not very good in clusters during or a few Gyr after mergers. Binary merger cluster simulations tell us that mergers generate bulk flow motion and complex temperature structures in the ICM, and anisotropic velocity distribution in DM and galaxies \\citep{Roet96, Taki99, Taki00, Rick01, Ritc02, Rowl04, Asca06, Pool06, Taki06, McCa07, Pool07, Spri07, Akah08, Mast08, ZuHo09}. Of course, asymmetric spatial structures are also often seen in ICM, DM, and galaxies. However, it is not trivial how these systems will be overestimated or underestimated. This depends on the phase of mergers, geometry of the system and observational direction, and which mass estimation method we use.   In this paper, we adopt the latter approach to investigate the impact of mergers on the cluster mass estimation using $N$-body+hydrodynamical simulation data. We make mock observational data such as line-of-sight velocity of member galaxies, X-ray brightness distribution, and spectroscopic-like temperature \\citep{Mazz04} maps of ICM from the simulation data. We perform ``simulations of mass estimation'' for these mock data, and compare the results with actual mass distribution in the simulation data.  The rest of this paper is organized as follows. In \\S \\ref{s:simu} we describe the adopted numerical methods and initial conditions for our $N$-body + hydrodynamical simulations. In \\S \\ref{s:resu} we present the results of mass estimation. In \\S \\ref{s:summ} we summarize the results and discuss their implications. ",
        "conclusions": "\\label{s:summ} We investigate the impact of mergers on the mass estimation of galaxy clusters using $N$-body + hydrodynamical simulation data. We randomly select a part of $N$-body particles and recognize them as galaxies whose line-of-sight velocity is observed. We estimate virial mass from these data and compare it with actual mass. The results strongly depend on the observational directions mainly because of anisotropic velocity distribution. When observed along the collision axis, the mass of the system tends to be overestimated. When the smaller cluster mass is larger than a quarter of the larger one, the virial mass can be larger than twice of the real mass. We also make the X-ray surface brightness and spectroscopic-like temperature maps from the simulation data. Mass profiles are calculated from these ``X-ray observational data'' under the assumption of hydrostatic equilibrium. Again, the estimated mass profiles are compared with those of actual mass. In general, mass estimation with X-ray data gives us fairly better results than Virial mass estimation. The dependence upon the observational directions is weaker, because the gas pressure is isotropic, and because anisotropic temperature fluctuations are smoothed out in azimuthal direction. When the systems are observed along the collision axis, the projected mass tends to be underestimated. This fact should be noted when the Virial and/or X-ray mass are compared with gravitational lensing results.   \\subsection{Mass Estimation Uncertainty and the Morphology of X-ray Image}   It is useful to discuss the relationship between the mass estimation uncertainty and the morphology of X-ray image. One may easily guess that clusters with an extremely elongated X-ray image (e.g., the snapshot at t=3.67 Gyr in figure \\ref{fig5}) are not in hydrostatic equilibrium. Indeed, the mass is systematically underestimated in this system (see ``t=3.67Gyr, vertical'' panel in Fig8). Roughly speaking, the X-ray mass estimation results in underestimation of 20-40 \\%. In other word, mass can be estimated as accurately as within a factor of two even for such an irregular cluster. The estimation accuracy becomes better for the projected mass comparison (e.g., the snapshot at t=3.67 Gyr in figure \\ref{fig10}). This is because the underestimation trend caused by irregular morphology and bulk flow motion is partly compensated by the overestimation trend caused by the projection effect.   To quantify the image morphology, we use the axial ratio of the surface brightness image defined through the second moments of the surface brightness, \\begin{eqnarray} M_{ij} = \\sum I_X x_i x_j, \\end{eqnarray} where the summation is conducted over the coordinates whose origin is at the X-ray peak and within a certain aperture centered on the origin \\citep{Ohar06}. Our models have axial symmetry for the collision axis ($x$-axis). Therefore, we only need to calculate $M_{11}$ and $M_{22}$ and do not have to diagonalize the matrix $M$ for the X-ray images seen from the direction perpendicular to the collision axis, when we choose the $x_1$ and $x_2$ axes to be parallel to the $x$ and $y$, respectively. We actually calculate $M_{12}(=M_{21})$ and confirm that $|M_{12}/M_{ii}|$ is typically $\\sim  0.005$. Then, the axial ratio $\\eta$ is defined as, \\begin{eqnarray} \\eta = \\Bigg\\{ \\begin{array}{cc} M_{11}/M_{22}, & M_{11} \\le  M_{22}, \\\\ M_{22}/M_{11}, & M_{11} > M_{22}. \\end{array} \\end{eqnarray} Thus, the axial ratio lies between 0 and 1 by definition, and $\\eta=1$ means a circular cluster. In some clusters with complex morphology, such as ones with three or more peaks, the axial ratio is not a good indicator \\citep{Vent08}. However, because our model clusters tends to have a relatively simple elliptical shape, it does not matter. Figure \\ref{fig12} shows the $\\eta$ as a function of the aperture radius (solid lines) and the ratio of estimated and real mass (dashed lines) for the standard run seen from the direction perpendicular to the collision axis. It is interesting that the inner region is nearly circular ($\\eta \\simeq 1$) but mass estimation error is relatively large at $t=3.67$ and $4.67$ Gyr. This means that the morphology determined from only inner region is not a good indicator of mass estimation error, which should be cared especially for the distant faint clusters. \\begin{figure} \\begin{center} \\FigureFile(160mm,160mm){fig12.eps} \\end{center} \\caption{The axial ratio of the X-ray surface brightness as a function of the aperture radius (solid line) and ratio of estimated mass with X-ray data to the real mass (dashed lines) for the standard run seen from the direction perpendicular to the collision axis.}\\label{fig12} \\end{figure}     On the other hand, it should be noted that clusters with spherical and regular X-ray morphology do not always result in good mass estimation. This is especially true for head-on merger clusters observed nearly along the collision axis. Spherical mass can be both overestimated and underestimated, depending on the radial distance and the phase of merger, although the results tend to be more accurate than in case of virial mass (see upper panels in figure \\ref{fig8}). About $\\sim 20$ \\% errors are typically seen. Even nearly 50 \\% errors can be found in the ``t=3.67Gyr, parallel'' panel in figure \\ref{fig8}. As for the projected mass results, however, the situation is a little bit different, where the X-ray mass is systematically smaller than the real mass (see the upper panels in figure \\ref{fig10}). In these cases, $\\sim$ 50 \\% underestimations can occur. In both cases, anyway, it should be noted that X-ray and virial mass estimations give us inconsistent results, and that both are significantly different than the real mass, although the X-ray image is quite symmetric. This implies that these clusters do not follow the correlation between X-ray temperature and galaxy velocity dispersion for relaxed clusters. In other words, these clusters show a peculiar $\\beta_{\\rm spec}$ value, which has been already recognized as a good indicator for merging clusters \\citep{Ishi96, Taki00}. These results can qualitatively explain the discrepancy between X-ray and lensing mass reported in CL 0024+17 \\citep{Ota04, Jee07}.  Recently, \\citet{ZuHo09} studied a line-of-sight galaxy cluster collision using numerical simulations, and compared their results with Cl 0024+17 in detail. Their results are basically consistent with ours, but there are several differences. As seen in figure \\ref{fig10}, systematic underestimation trends are clearly seen in the projected mass when observed along the collision axis, which is qualitatively consistent with theirs. As written in subsection 2.2 of \\citet{ZuHo09}, their collision velocity is approximately twice of the free fall velocity. This is an extremely rare case in the CDM universe though it might be good for reproducing the Cl 0024+17 results. On the other hand, the collision velocities in our simulations are comparable to but a little bit smaller than the free fall one, which means a more usual case. Thus, it is not surprising that \\citet{ZuHo09} shows more serious underestimation, and that their results are quantitatively different than ours. In addition, though both \\citet{ZuHo09} and our models used the double-beta model fitting, its interpretation and resultant procedure of mass estimation are different. \\citet{ZuHo09} assumes that double-beta surface brightness distribution means two individual spherical clusters aligned along the line-of-sight. Therefore, they calculate hydrostatic mass for two spherical clusters and sum them in deriving projected mass. On the other hand, we assume that the double-beta surface brightness distribution means a single spherical cluster with a relatively complex density profile that cannot be well represented by a single beta-model \\citep{Ikeb04}. Thus, we calculate hydrostatic mass for a single spherical cluster.    \\subsection{ICM Internal Velocity} It is quite natural that ICM internal velocity as well as member galaxies' velocities can be another probe of merger dynamics and mass estimation uncertainty. In addition, although subclusters are in hydrostatic equilibrium in the initial conditions in our merger simulations, it is most likely that clusters before major mergers have already undergone some mergers in the CDM universe. As a result, the merger progenitors could have significant bulk and/or turbulent flow motions in the ICM. This causes additional underestimation in mass measurement with X-ray data \\citep{Fang09}. Anyway, firm direct detection of ICM motion is highly desired in this regard though this is not an easy task in the current status of X-ray astronomy (e.g. \\cite{Ota07,Fuji08,Suga09}). The ASTRO-H satellite (former NeXT; e.g. \\cite{Taka06}), which is planned to be launched around 2013, will enable us to measure the line-of-sight velocity directly with the X-ray microcalorimeters and provide us with useful information on the dynamical status of the ICM. Another possible observational approach to the ICM motion is the Sunyaev-Zel'dovich (SZ) effect. In principle, the line-of-sight and tangential velocity components can be measured with the kinematic SZ effect and SZ polarization, respectively \\citep{Suny80}, though they are still challenging in the present status of the observations.   \\subsection{Member Galaxy Distribution} As mentioned above, we sampled particles corresponding to the member galaxies from $N$-body particles in virial mass estimation. This means that the phase space distribution of the galaxies are assumed to be essentially the same as that of DM. However, it is likely that this assumption is not so good in realistic situations. Using cosmological simulations that include cooling and star formation, \\citet{Naga05} showed that the concentration of the radial distribution of simulated galaxies is lower than that of dark matter even for relaxed systems, thereby implying errors in the mass estimation through the virial theorem. On the other hand, our aim is to evaluate errors caused by dynamical motions during mergers, therefore, for systems that are far from dynamical equilibrium. We conclude that our findings about virial mass estimates are not significantly affected by such biases \\citep{Naga05}.    \\bigskip M. T. would like to thank N. Okabe, M. Takada, and M. Nagashima for helpful comments. The authors are grateful to an anonymous referee for his/her useful comments and suggestions which improved the manuscript. Numerical computations were carried out on VPP5000 and XT4 at the Center for Computational Astrophysics, CfCA, of the National Astronomical Observatory of Japan. M. T. was supported in part by a Grant-in-Aid from the Ministry of Education, Science, Sports, and Culture of Japan (16740105, 19740096).       \\appendix"
    },
    "1004/1004.2790.txt": {
        "abstract": "We report the final optical identifications of the medium-depth ($\\sim 60$ ksec), contiguous (2 deg$^{2}$) XMM-Newton survey of the COSMOS field. XMM-{\\it Newton} has detected $\\sim1800$ X-ray sources down to limiting fluxes of $\\sim 5\\times10^{-16}$, $\\sim 3\\times10^{-15}$, and $\\sim 7\\times10^{-15}$ \\cgs\\ in the 0.5-2 keV, 2-10 keV and 5-10 keV bands, respectively ($\\sim 1\\times10^{-15}$, $\\sim 6\\times10^{-15}$, and $\\sim 1\\times10^{-14}$ \\cgs, in the three bands, respectively, over 50\\% of the area). The work is complemented by an extensive collection of multi-wavelength data from 24$\\mu$m to UV, available from the COSMOS survey, for each of the X--ray sources, including spectroscopic redshifts for $\\gs50$\\% of the sample, and high-quality photometric redshifts for the rest. The XMM and multiwavelength  flux limits are well matched: 1760 (98\\%) of the X-ray sources have optical counterparts, 1711 ($\\sim95$\\%) have IRAC counterparts, and 1394 ($\\sim78$\\%) have MIPS 24 $\\mu$m detections. Thanks to the redshift completeness (almost 100\\%) we were able to constrain the high-luminosity tail of the X--ray luminosity function confirming that the peak of the number density of logL$_{X}>44.5$ AGN is at z$\\sim2$. Spectroscopically-identified obscured and unobscured AGN, as well as normal and starforming galaxies, present well-defined optical and infrared properties. We devised a robust method to identify a sample of $\\sim150$ high redshift (z$>1$), obscured AGN candidates for which optical spectroscopy is not available. We were able to determine that the fraction of the obscured AGN population at the highest (L$_{X}>10^{44}$ erg s$^{-1}$) X--ray luminosity is $\\sim 15-30$\\% when selection effects are taken into account, providing an important observational constraint for X--ray background synthesis. We studied in detail the optical spectrum and the overall spectral energy distribution of a prototypical Type 2 QSO, caught in a stage transitioning from being starburst dominated to AGN dominated, which was possible to isolate only thanks to the combination of X-ray and infrared observations. ",
        "introduction": "First introduced to explain the properties we observe today in {\\it normal} galaxies such as their color bimodality (Blanton et al. 2003, Bell et al. 2004, Faber et al. 2007) and the so-called local scaling relations (Ferrarese \\& Merrit 2000, Gebahrdt et al. 2000, G\\\"ultekin et al. 2009), the existence of a ``feedback\" between the accreting Super Massive Black Hole (SMBH) and the host galaxy in which it resides ultimately challenged our understanding of {\\it Active Galactic Nuclei} (AGN). Indeed, the widely accepted scenario of galaxy-AGN coevolution (e.g. Silk \\& Rees 1998, Fabian \\& Iwasawa 1999, Granato et al. 2004, Di Matteo et al. 2005, Menci et al. 2008) points towards a physical coupling responsible for the  self-regulated SMBH and galaxy growth. In particular, according to the theoretical scenario emerging from extensive semi-analytical models and hydrodynamics simulations, a key event in the SMBH vs. host galaxy co-evolution is represented by the highly obscured AGN phase, when large quantities of gas driven to the center in the merger of two disk galaxies were available to efficiently feed and obscure the growing black hole and build galaxy stellar mass through significant episodes of star formation. This phase ends when strong winds and shocks from the central AGN heat the interstellar medium, blowing away the dust and gas, thus cleaning the line of sight and inhibiting further star-formation (see e.g. Granato et al. 2004, Di Matteo et al. 2005; Croton et al. 2006; Sijacki et al. 2007; Menci et al. 2008; Hopkins et al. 2008).  In this general framework for merger-induced AGN feedback, capable of passing numerous observational tests (e.g., Sanders et al. 1988,  Hopkins et al. 2006, Veilleux et al. 2009), the differences between ``obscured\" and ``unobscured\" AGN are no longer and uniquely described under a geometrical unification model (in which they are simply related to orientation effects, Antonucci \\& Miller 1985, Antonucci 1993, Urry \\& Padovani 1995), but can be interpreted as due to the fact that the {\\it same} objects are observed in {\\it different evolutionary} phases. This hypothesis is consistent with the finding, mainly from X--ray surveys, that absorption is much more common at low luminosities (see e.g. Ueda et al. 2003, la Franca et al. 2005, Maiolino et al. 2007, Hasinger 2008) and, possibly, at high redshift (La Franca et al. 2005; Ballantyne et al. 2006; Treister \\& Urry 2006; Hasinger 2008). The luminosity and redshift dependence of the obscuring fraction may be naturally linked to the AGN radiative power (related to the intrinsic X--ray luminosity, see e.g. Lawrence \\& Elvis 1982) which is able to ionize and expel gas (more common at high-z) and dust from the nuclear regions, nicely fitting the current framework of AGN formation and evolution sketched above (see, for example, Hopkins et al. 2006). The complete picture is likely to be more complex, depending on many other parameters such as the BH mass, the Eddington ratio, the star formation activity in the host galaxy and, in particular, the timescales associated with the AGN duty cycle activity.  A correct and complete identification of unobscured, obscured and highly obscured AGN at all redshifts (and especially in the z=1-3 interval, where most of the feedback is expected to happen) is therefore crucial for a comprehensive understanding of the still little explored phase of the common growth of SMBHs and their host galaxies. While it is relatively straightforward to select unobscured AGN from optical multicolor surveys (e.g. Richards et al. 2002) and/or from spectroscopic samples (e.g. the VVDS survey, Gavignaud et al. 2006, Bongiorno et al. 2007), the most efficient, reddening independent method to select obscured, type--2 AGN is the presence of luminous X-ray emission (L$_{[2-10keV]}>10^{42}$ erg s$^{-1}$) and hard X-ray colors (see Brandt \\& Hasinger 2005 for a review). Indeed, hard (2-10 keV) X-ray nuclear emission is an almost unambiguous mark of the presence of an AGN, given the extremely small contamination from star-formation induced emission at these luminosities and frequencies. Combined with the fact that hard X-rays are not seriously affected by obscuration due to neutral gas along the line of sight (up to N$_{\\rm H}\\sim10^{23}$ cm$^{-2}$), it is clear that hard X-ray selection is the most effective method to uncover unobscured to moderate obscured AGN and study in detail their demographics (see Comastri \\& Brusa 2008; Brandt \\& Alexander 2010 for a review).   Since the launch in 1999 of both the XMM--{\\it Newton} and {\\it Chandra} satellites, a large ($>$ 30) number of surveys covering a wide fraction of the area vs. depth plane (see Fig.~1 in Brandt \\& Hasinger 2005; see also Hickox 2009) have been performed, and our understanding of AGN properties and evolution has received a major boost. Thanks to vigorous programs of multiwavelength follow-up campaigns, sensitive X--ray observations turned out to be highly efficient in unveiling  weak and/or \"elusive\" accreting black holes, in a variety of otherwise {\\it non-active} galaxies (i.e. not recognized as AGN from the optical spectra or continuum emission), such as (among others): X--ray Bright Optically Normal Galaxies (XBONG, e.g. Comastri et al. 2002, Severgnini et al. 2003, Civano et al. 2007), Extremely Red Objects (e.g. Alexander et al. 2002, Brusa et al. 2005, Severgnini et al. 2005), Sub Millimeter Galaxies (e.g. Alexander et al. 2005, Laird et al. 2009), high-z starforming systems (e.g. Daddi et al. 2007, Fiore et al. 2008, 2009, Treister et al. 2009b), Lyman Break Galaxies (LBG, e.g. Brandt et al. 2002, Nandra et al. 2002, Aird et al. 2008). In many of these cases, the AGN responsible for the X-ray emission is overwhelmed at longer wavelength by the host galaxy light and/or the obscuration might be connected to processes within the host galaxy itself, such as the presence of dust lanes or starburst disks (see e.g. Martinez-Sansigre et al. 2005, Ballantyne et al. 2006, Ballantyne 2008, Hopkins et al. 2009). This suggests that the accretion activity (especially in high-redshift sources) can be unambiguously revealed thanks to the presence of a strong X--ray emission (see e.g. discussion in Brusa et al. 2009) and, therefore, the combination of both X-ray and optical classifications, coupled with the multiwavelength analysis, can be crucial to fully assess the nature of the candidate AGN.  The high level of completeness in redshift determination for a large number of X--ray selected AGN (up to a few thousands) has made possible a robust determination of the luminosity function  and evolution of unobscured and mildly obscured AGN which turned out to be luminosity dependent: the space density of bright QSOs ($L_X > 10^{44}$ erg s$^{-1}$) peaks at z$\\sim$ 2--3, to be compared with the z$\\sim$0.7--1 peak of lower luminosity  Seyfert galaxies (Ueda et al. 2003; Hasinger et al. 2005;  La Franca et al. 2005, Silverman et al. 2008; Ebrero et al. 2009, Yencho et al. 2009). Based on the Ueda et al. (2003) work, Marconi et al. (2004) and Merloni (2004) were the first to propose that SMBH undergo a ``anti--hierarchical\" evolution, in the form of a differential growth (earlier and faster for more massive black holes). This anti--hierarchical behavior observed in AGN evolution (similar to that observed in normal galaxies, e.g. Cowie et al. 1996) provided an important and independent confirmation that the formation and evolution of SMBH and their host galaxies are likely different aspects of the same astrophysical process. It should be noted that the Luminosity Dependent Density Evolution (LDDE) parameterization, corresponding to a strong downsizing, has been recently questioned by Aird et al.(2010). Their preferred model for the evolution of the XLF is a luminosity and density evolution model (LADE), where the shift in the redshift peak of AGN space density as function of X--ray luminosity is much weaker than in LDDE models.  A full characterization of the AGN bolometric luminosity function can be obtained only exploring the entire area vs. flux plane (e.g. combining samples from deep and large area surveys) and adopting different AGN selection methods (X--ray vs. optical vs. infrared). Nevertheless, it seems plausible that the blow-out phase associated with the obscured growth represents a relatively short, but very powerful, episode in the QSO lifetime. For this reason, the probability to detect rare and luminous X--ray events is maximized by large area surveys and, in particular, by hard X--ray observations (to cope with obscuration effects) and associated deep multiwavelength coverage.  The  XMM--{\\it Newton} wide-field survey in the COSMOS field (hereinafter: XMM-COSMOS, Hasinger et al. 2007) is an important step forward in addressing the topics described above. The $\\sim2$ deg$^2$ area of the HST/ACS COSMOS Treasury program (Scoville et al. 2007a,b; Koekemoer et al. 2007) has been surveyed with \\xmm\\ for a total of $\\sim$1.55 Ms during AO3, AO4 and AO6 cycles of XMM observations (Cappelluti et al. 2007, Cappelluti et al. 2009, hereafter C09). XMM-COSMOS provides an unprecedently large sample of point-like X-ray sources ($\\gsimeq1800$), detected over a large, contiguous area, with {\\it complete} ultraviolet to mid-infrared (including Spitzer data) and radio coverage, and extensive spectroscopic follow--up granted through the zCOSMOS (Lilly et al. 2007; 2009) and Magellan/IMACS (Trump et al. 2007; 2009) projects. The excellent multiband photometry available in this area allows a robust photometric redshift estimate for the faint sources not reachable by optical spectroscopy, thus allowing a virtually complete sample of X-ray sources. The XMM-COSMOS project is described in Hasinger et al. (2007), while the X--ray point source catalog and counts from the complete XMM-COSMOS survey are presented in a companion paper (C09). The present paper, which extends the work presented by Brusa et al. 2007 (hereiafter B07) on the optical identifications of the X--ray point sources in the XMM-COSMOS survey, discusses the multiwavelength properties of this large sample of X-ray selected AGN. Several works have already appeared in the literature and are based on this catalog, or on previous versions of it, e.g.: the derivation of AGN photometric redshifts (Salvato et al. 2009); the IMACS AGN spectroscopic campaign (Trump et al. 2009); the clustering properties of spectroscopically confirmed AGN (Gilli et al. 2009); the space density of high-redshift QSOs (Brusa et al. 2009); the host galaxies properties of AGN in COSMOS (Gabor et al. 2009); the on-going and co-evolving starformation and  AGN activity at z$<0.8$ (Silverman et al. 2009a); the environments of AGN within the galaxies density field (Silverman et al. 2009b); the relation of the X--ray and optical emission in BL AGN (Lusso et al. 2010).  The paper is organized as follows: Section~2 presents the multiwavelength datasets drawn from the COSMOS survey and used in the paper; Section 3 describes the method used to identify the X-ray sources and its statistical reliability; the XMM-COSMOS multiwavelength catalog is presented in Section~4, while Section~5 reports the redshift information for the X-ray sources. The number density evolution of luminous XMM-COSMOS  sources is presented in Section~6. Section~7 describes the X--ray to optical and near infrared properties of the identified population and of the obscured AGN candidates, while in Section~8 we discuss the obscured AGN fraction as a function of the X--ray luminosity. Section~9 presents the spectrum and SED of a prototype obscured QSO at z$\\sim1.6$, and in Section 10 we summarize the most important results. Throughout the paper, we adopt the cosmological parameters $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m$=0.3 and $\\Omega_{\\Lambda}$=0.7 (Spergel et al. 2003). In quoting magnitudes, the AB system will be used, unless otherwise stated. ",
        "conclusions": "We presented the catalog of optical and infrared identifications for $\\sim$1800 X--ray sources detected in the XMM-COSMOS survey. This catalog comprises the totality of the XMM sources presented in C09, with the exception of 65 faint sources ($\\sim$3\\%) detected in two additional fields obtained in AO6. The optical/IR matches with the X--ray positions were driven by the {\\it Chandra} subarcsec centroids, when available, or estimated through statistical methods already successfully tested in different XMM and {\\it Chandra} fields (see B07, Cardamomone et al. 2008). For each source, a flag on the reliability of the X--ray to optical/infrared association is provided, classifying the counterparts as ``reliable\" (87.7\\%), ``ambiguous\" (11.3\\%) and ``not identified\" (1\\%; see Table 1). For sources with two possible associations both entries are reported in the associated catalog (Table 2). The availability of \\chandra\\ data allowed us also to quantify the reliability of the likelihood ratio technique when applied to XMM sources matched to deep optical and infrared images, which turned out to be very high ($\\gs98.5$\\%, see Section 3.1).  Along with the positions of the optical counterparts we publish a wealth of multiwavelength information, from 24$\\mu$m to X--ray (see Section 4 and Table 2), most notably  a collection of all the available spectroscopic redshifts ($\\sim 900$) obtained through different programs ongoing in the COSMOS field (see Section 5 for details). To maximize the completeness over a well defined large area and, at the same time, keep selection effects under control, for the multiwavelength analysis presented in this paper we further considered the subsample of 1640 identified sources detected above the limiting fluxes corresponding to a coverage of more than 1 deg$^{2}$ in at least one  of the three X--ray energy ranges considered (see Section 3.1). The multiwavelength properties of the fainter X--ray sources are investigated in a much greater detail in the framework of the C-COSMOS survey (Elvis et al. 2009, Civano et al. in prep.).  The spectroscopic completeness in the soft (0.5-2 keV) and hard (2-10 keV) subsamples is $\\sim$50\\% and $\\sim$60\\%, respectively, remarkably high for identification campaigns of X--ray sources, to be  compared with the spectroscopic completeness in the AEGIS survey (40\\% in the 2-10 keV band, see Aird et al. 2010); in the X-Bootes survey ($\\sim$50\\%, Ryan Hickox, private communication; see also Murray et al. 2005; Kenter et al. 2005); and in the CDFS 2Ms sample ($<50$\\%, Luo et al. 2008, Luo et al. 2010). By far, the highest spectroscopic completeness is for the sources detected in the ultra-hard (5-10 keV) band, where it reaches $\\sim80$\\%: for the first time, XMM-COSMOS provides a statistically significant sample (221) of sources detected at $>5$ keV with a very high spectroscopic breakdown at fluxes. %%%%%%%%%%%%%%%%%%%%%%%%%%%% As a comparison, the Beppo-SAX HELLAS survey had a 50\\% completeness (71/147) at higher X-ray fluxes ($>3\\times10^{-14}$ \\cgs, Vignali PhD thesis), while the XMM-HBSS has an almost complete (97\\%) spectroscopic information, but it is limited to 62 sources at fluxes $>7\\times10^{-14}$ \\cgs\\ (Della Ceca et al. 2008).  For the sources without spectroscopic redshifts, accurate photometric redshifts (down to $\\sigma(\\Delta z/(1+z))\\sim0.01$ at I$<22.5$) are available (Salvato et al. 2009), allowing an almost 100\\% complete sample of X--ray sources with redshift information. This, coupled with the large number of objects, allowed us to sample with an unprecedented statistic the high luminosity tail  of the X--ray luminosity function. The present data allow us to derive the AGN space density as obtained by the COSMOS survey over a relatively broad redshift range (z$\\sim$ 1--3) and luminosity interval ($Log L_{\\rm X}=44.5-47$). We compared the COSMOS results with those discussed in recent papers (Ebrero et al. 2009; Yencho et al. 2009; Aird et al. 2010) combining data from different surveys and covering in a fairly homogeneous way the luminosity redshift plane. Our results suggest that at high luminosities ($log L_{\\rm X}>44.5$) the AGN number density peaks at z$\\sim 2-2.5$, in good agreement with LDDE parameterization of the above mentioned XLF.  We investigated the dependence of the X--ray to optical flux ratio versus the X--ray luminosity for BL AGN and NL AGN, and sources with photometric redshifts only. The normalization of the correlation between X/O and X--ray luminosity for the obscured AGN increases going from the spectroscopically identified sample to the objects with only photometric redshifts available (i.e., towards fainter magnitudes; see Figure~\\ref{xolx}). The well-known relation between the X/O and the X--ray luminosity (F03; Barger et al. 2005) for obscured AGN is mostly driven by the fact that, while the nuclear AGN X-ray luminosity can span several decades (depending on the BH mass {\\it and} on the accretion rate), the host galaxy R-band luminosity (which dominates the optical emission in obscured sources) has a moderate scatter, around the M$_{\\rm BH}-M_{\\rm Bulge}$ relation, i.e., depending only on the BH mass.   We then studied  the dependence of the $R-K$ color on the X--ray luminosity and we found an opposite behavior for BL AGN and NL AGN in our sample. For NL AGN there is evidence for a trend where high-luminosities sources have redder colors than low-luminosity sources (see Figure~\\ref{rklx}). This correlation is most likely due to redshift effects, given that the $R-K$ color of normal and star forming galaxies increases up to redshift z$\\sim 2$. On the contrary, while high-luminosity BL AGN show a median $R-K$ color consistent with that observed in optically selected samples (Barkhouse \\& Hall 2001), lower luminosity BL AGN show a median $R-K$ color statistically different (redder) from the one observed for the higher-luminosity ones (see Figure~\\ref{rklx}). This difference in the $R-K$ colors for high and low luminosity sources is likely due to the contribution of the host galaxies in the latter sample. We conclude that X--ray surveys are effective in isolating also {\\it unobscured} objects otherwise missed by optical multicolor surveys, with preselection on the basis of optical colors and pointlike emission (see also similar conclusions obtained on BL AGN samples from pure spectroscopic selection, Bongiorno et al. 2007). However, the effect of host galaxy contribution  may play an important role in getting the correct classification of the nuclear source.  The large body of COSMOS  multiwavelength data allowed us to devise a robust method to build a sizable sample of luminous obscured AGN candidates at $z>1$ among the still spectroscopically unidentified population, and control the selection effects (see Section 7.1.1 and 7.1.2). In total, we ended up with a robust sample of $\\sim 150$ objects with photometric redshifts information, and multicolor properties consistent with those of luminous, obscured quasars at z$>1$. This sample should not be considered complete with respect to the X--ray selected, high-z obscured AGN population,  but can be considered as representative of the average properties of this class of sources. We compared the average X--ray hardness ratio of the high-z obscured AGN candidate sample with the ones derived for the spectroscopically identified samples (Section 7.2). We found that the average HR for this sample, HR$=-0.17$, is considerably higher than that derived for spectroscopically classified NL AGN (HR$=-0.29$), further confirming that a red observed  $R-K$ color ($R-K\\gsimeq5$) coupled with high X/O or $\\nu_{24}F_{24}/\\nu_{R}F_{R}$ ratios, is a good indicator for both optical and moderate X--ray obscuration in high-luminosity sources (see also Maiolino et al. 2006).  For the vast majority (75-80\\%) of the high-z obscured AGN candidates the AGN component is visible in the near-IR (roughly corresponding to the rest frame $\\sim 1-3\\mu$m range), where the dust reprocessed emission is expected to dominate the observed SED. These sources would have been selected on the basis of IRAC colors only (see Fig.~\\ref{sternflux} and Fig.~\\ref{iraccm}, and Section 7.3). However, at the limiting fluxes of the IRAC COSMOS survey the contamination from non-AGN objects within the selection wedge starts to become important, and the selection criterion would become efficient only when combined with an additional criterion (in this case, the X--ray emission).  We studied the fraction of obscured AGN in the XMM-COSMOS sample as a function of the 2-10 keV X--ray luminosity, and discussed the differences between optical (spectroscopic and photometric) and X--ray classifications. We confirmed that the fraction decreases with increasing luminosity (see Section~8 and Fig.~\\ref{fracl}), for both the spectroscopically identified and the full sample. At L$_{\\rm X}\\gs10^{44}$ erg s$^{-1}$, the obscured AGN fraction is of the order of 25-30\\% in the logL$_{\\rm X}=44-44.5$ bin and $\\sim 15-20$\\% in the  logL$_{\\rm X}=44.5-45$ bin, about a factor of 2-3 lower than the predictions from the Gilli et al. (2007) and Treister \\& Urry (2006) models. At lower-luminosities, the optical and X--ray classifications largely differ (see green triangles and squares at logL$_{\\rm X}<43$ erg s$^{-1}$ in Figure~\\ref{fracl}). This may be related, at least in part, to host galaxy dilution of the intrinsic AGN spectrum, which is expected to be important at low Eddington ratios and at low luminosities (Hopkins et al., 2009).  We studied in detail the SED and the spectral properties of XID 2028, that we consider the prototype of an obscured QSO at high-z (z=1.592) caught in a transition stage from being starburst dominated to AGN dominated, which was possible to isolate only thanks to the combination of X-ray and infrared observations. XID 2028 is one of the brightest XMM-COSMOS sources, with an unabsorbed X--ray luminosity L$_{\\rm X}\\sim10^{45}$ erg s$^{-1}$ and significant X--ray absorption (N$_{H}\\sim10^{22}$), i.e. a Type 2 QSO. The optical spectrum of this source presents strong AGN emission lines ([NeV] doublet) consistent with our selection method. Moreover,  evidence of outflowing material at a velocity of $\\sim 300$ km s$^{-1}$ (see Figure~\\ref{spe2028}) is also clearly detected from the blueshifted MgI and MgII complex absorption. The X--ray, optical and infrared properties and band ratios of XID 2028 are very similar to those predicted by the above mentioned models, and the overall SED is best represented by numerical templates resulting from the models published by Narayanan et al. (2009) to describe z$\\sim2$ DOG (see Fig.~\\ref{sed2028}, right panel).  XID 2028 is the brightest object for which such a detailed analysis can be conducted in the XMM-COSMOS catalog. However, as demonstrated in Section 7, it is reasonable to assume that a large fraction of the sources among the high-z obscured AGN candidates share the same properties, although at lower X--ray (L$_{\\rm X}\\gs10^{44}$ erg s$^{-1}$) and bolometric luminosities. Most of these sources are well within the reach of the {\\it Herschel} (Pilbratt 2005) instruments between 75 and 500 $\\mu$m, which in the GT project in the COSMOS field will reach limiting fluxes of few mJy. Such longer wavelength observations can greatly help in separating nuclear activity and star formation, and assessing the real bolometric output of obscured AGN. Further investigations, as for example, direct infrared spectroscopy with X-Shooter@VLT (D'Odorico et al. 2004) or LUCIFER@LBT (Mandel et al. 2007) may be crucial in confirming the nature of these candidates."
    },
    "1004/1004.1920_arXiv.txt": {
        "abstract": "By relaxing the conventional assumption of a purely gravitational interaction between dark energy and dark matter, substantial alterations to the growth of cosmological structure can occur. In this work we focus on the homogeneous transfer of energy from a decaying form of dark energy.  We present simple analytic solutions to the modified growth rates of matter fluctuations in these models, and demonstrate that neglecting physics within the dark sector may induce a significant bias in the inferred growth rate, potentially offering a false signature of modified gravity. ",
        "introduction": "One of the major goals of future cosmological studies is to categorise dark energy as one of three candidates: a cosmological constant, a physical fluid, or simply some manifestation of new gravitational physics. Geometric measurements such as supernovae and baryon acoustic oscillations will scrutinise the null hypothesis of a cosmological constant. However, a physical fluid and modified gravity could both reproduce  almost arbitrary expansion histories, so distinguishing them requires a study of the growth of structure. It is the validity of this diagnostic step that provides the focus of the present work.  Anisotropic stress in the dark energy fluid may mimic metric theories of gravity, as demonstrated by Kunz \\& Sapone \\cite{2007PhRvL..98l1301K} and more generally by Hu \\cite{2009NuPhS.194..230H}. Here we take this one step further: even \\emph{without} anisotropic stress, it can be shown that significant deviations in structure growth are achievable. The only requirement is that some interaction should exist between dark matter and the smooth dark energy component, permitting the transfer of energy.  Any change in the dark energy density is conventionally attributed to the equation of state $w(z)$, which dictates the adiabatic behaviour of a physical fluid. But density evolution may also arise from non-gravitational interactions with other fluids, thereby violating adiabaticity. Cosmologies with energy exchange  have been extensively studied in the literature \\cite{2000PhRvD..62d3511A, 2006PhRvD..73j3520B, valvi, valvi2, 2009PhRvD..79d3522C, jackson, 2001PhLB..521..133Z, baldi2010, LiBarrow}, and generate an expansion history that is fully reproducible by a single inert scalar field whose evolution matches the effective equation of state $w_{\\rm{eff}}(z)$. The degeneracy between these two models may be broken by studying the growth of structure, which is disrupted by the evolving matter density \\cite{HeWang}. Yet it is this same test that would conventionally be used to identify modifications to gravity.     In \\S \\ref{sec:coupling}  we begin by reviewing a simple case of dark energy decaying into a form of dark matter.   The evolution of perturbations are quantified in \\S \\ref{sec:growth}, before exploring the observational consequences in \\S \\ref{sec:observables}. ",
        "conclusions": "In the case of a decaying cosmological `constant', we have demonstrated that the growth rate of large scale structure is subject to a constant decrement $f=\\Omega_m^{\\gamma} - c$. For the simple model we consider, this can largely be attributed to the dragging effect on bulk motions induced by the production of stationary matter. Smaller contributions arise from both the background dynamics, and the diluting effect of gradually introducing a homogeneous density field. It is also important to note that in more general interacting models the presence of dark energy perturbations may also influence the growth of structure. These act to enhance the growth rate, and may overpower the mechanisms considered here. Indeed in many cases the dark energy perturbations in coupled models lead to pathological instabilities.  Energy exchange within the dark sector leads to a change in the comoving matter density, which in turn invokes a number of changes to cosmological observations. If these changes are not taken into consideration, then suppression to the growth of large scale structure leads to a na\\\"{i}vely inferred value of the growth index $\\gamma$ rising above the conventional value $\\gamma > 0.55$. While the physical motivation for such models remains unclear, we believe this is no less true for current approaches to modified gravity.  As noted by Blandford et al. \\cite{2005ASPC..339...27B}, a blind cosmologist living in the radiation dominated era might measure the evolution of the scale factor and erroneously conclude that the expansion is driven by a scalar field with an exponential potential. This thought experiment may be extended further: a ``dark\" astronomer only capable of studying the dark matter would notice a strange behaviour of non-linear structure. They might attempt to construct a modified theory of gravity which accounts for this behaviour. However the true source of this discrepancy is the momentum exchanged between the baryons and photons, inducing new features to the growth of structure.  This modification is apparent both at the era of recombination, and the present day. We are at risk of falling into a similar trap here: the baryocentric assumption that neither dark energy nor dark matter exhibit any complex behaviour may lead to dark physics being mistaken for modified gravity.  \\noindent{\\bf Acknowledgements} \\\\ FRGS and BMJ are grateful for support from the STFC, and we thank the anonymous referee for helpful comments."
    },
    "1004/1004.0671_arXiv.txt": {
        "abstract": "Observations made during the last ten years with the Chandra X-ray Observatory have shed much light on the cooling gas in the centers of clusters of galaxies and the role of active galactic nucleus (AGN) heating.  Cooling of the hot intracluster medium in cluster centers can feed the supermassive black holes found in the nuclei of the dominant cluster galaxies leading to AGN outbursts which can reheat the gas, suppressing cooling and large amounts of star formation.  AGN heating can come in the form of shocks, buoyantly rising bubbles that have been inflated by radio lobes, and the dissipation of sound waves. ",
        "introduction": "\\dropcap{C}handra observations of cavities in cluster centers in the last decade have brought the issue of feedback to the center stage in astrophysics. Feedback from AGN driven by supermassive black holes in the cores of massive galaxies is now seen as a necessary ingredient to adequately describe the formation and evolution of galaxies and the clusters in which they sometimes reside.  The energy injected by AGN is required to produce a galaxy luminosity function in simulations (e.g. \\cite{croton06}) that matches observations.  This feedback may set the upper limit to the observed masses of galaxies and contributes to cluster preheating and the observed ``entropy floors'' seen in cluster centers. It may affect scaling relations, such as the relationship between X-ray luminosity and temperature and potentially affects cluster properties that are used for constraining cosmological models, such as the gas mass fraction (e.g. \\cite{allen04}).  Data from the {\\it{Chandra X-ray Observatory}} over the last ten years reveal AGN at work in the centers of clusters, inflating bubbles that rise buoyantly through the intracluster medium (ICM), sometimes producing shocks and sound waves.  Many of the most spectacular {\\it{Chandra}} images of AGN feedback are of interactions in the centers of cooling flow (sometimes referred to as ``cool core'') clusters of galaxies.  Under the assumption of hydrostatic equilibrium, the X-ray-emitting gas in a cluster is the most dense at the cluster center.  The gas will cool as it radiates, and the cooling time is inversely proportional to the density of the gas.  Therefore, gas will cool at cluster centers before it cools at larger radii.  If the cooling time is shorter than the cluster age, the central gas will cool, lose pressure, and outer gas must flow in to maintain hydrostatic equilibrium (see \\cite{fabian94} for a review). Data from earlier X-ray observatories such as {\\it{Einstein}} \\cite{white97}, the ROentgen SATellite ({\\it{ROSAT}}) \\cite{peres98}, and the Advanced Satellite for Cosmology and Astrophysics ({\\it{ASCA}}) \\cite{white00} were used to give estimates of the gas inflow rates of typically hundreds or thousands of solar masses per year.  The classic ``cooling flow problem'' has been that sufficient quantities of cool gas and/or star formation were not detected to match the predictions from the X-ray data.  {\\it{Chandra}} and {\\it{XMM-Newton}} observations have shown that while the gas temperature drops in the centers of cooling flow clusters, the temperature reaches a threshold value below which there is little cooling. This is seen in the temperature profiles as well as high resolution spectroscopy.  An important early result from {\\it{XMM-Newton}} high resolution spectra from cooling flows was that the emission lines from cool gas were not present at the expected levels, and the spectra were well-fitted by a cooling flow model with a low temperature cutoff \\cite{peterson01,peterson03}.  This cutoff is typically one-half to one-third of the average cluster temperature, and is usually not much below 1 keV.  The gas is either then cooling non-radiatively or being heated to higher temperatures. Throughout this paper, we use the term ``cooling flow'' to indicate clusters with central drops in temperature and large central density peaks above that of a typical $\\beta$-model.  There is a clear association with the presence of a cooling flow in a cluster and the central galaxy hosting a radio source. Using {\\it{Einstein}} data, \\cite{burns90} found $\\approx70\\%$ of cooling flow clusters had giant elliptical galaxies that emitted in the radio, as compared to only $\\approx20\\%$ of central galaxies in non-cooling flows.  Recent studies, depending on the sample, find a fraction of up to 100\\% for cooling flows and 45\\% for non-cooling flows \\cite{mittal09} using {\\it{Chandra}} data.  This association is consistent with the cooling gas supplying fuel to the central black hole, leading to the AGN outburst observed in the radio.  As seen in {\\it{Chandra}} images, the radio jets and lobes associated with the AGN significantly affect the X-ray gas, inflating ``cavities'' or ``bubbles'' in the cluster centers.  Clear correspondence between regions of radio emission and deficits in the X-ray are seen in a number of cases (e.g. Hydra A \\cite{mcnamara00,wise07}, Perseus \\cite{fabian00}, Abell 2052, \\cite{blanton01}, Abell 2597 \\cite{mcnamara01,clarke05}, Abell 2199 \\cite{johnstone02}, Centaurus \\cite{fabian05}, Abell 4059 \\cite{heinz02,choi04}, Abell 262 \\cite{blanton04,clarke09}, MS0735.6+7321 \\cite{mcnamara05}, M87/Virgo \\cite{forman07}). The X-ray gas, in turn, affects the radio lobes by confining and distorting them.  Energy can be transported to the ICM through bubbles that rise buoyantly to larger cluster radii, as well as through weak shocks and sound waves generated by AGN outbursts.  The energy transferred to the cluster gas can diminish the cooling, giving an explanation for the lack of copious quantities of cool gas and star formation seen in the centers of cooling flow clusters. Composite images of A2052 and A262 are shown in Figs.\\ 1 and 2, respectively, where blue is the X-ray emission as seen with {\\it Chandra}, red is optical emission, and yellow/green (in the cluster centers) is 1.4 GHz radio emission. ",
        "conclusions": ""
    },
    "1004/1004.1441_arXiv.txt": {
        "abstract": "Using infrared photometry from the \\emph{Spitzer Space Telescope}, we perform the first inventory of aromatic feature emission (AFE, but also commonly referred to as PAH emission) for a statistically complete sample of star-forming galaxies in the local volume. The photometric methodology involved is calibrated and demonstrated to recover the aromatic fraction of the IRAC $8\\mum$ flux with a standard deviation of 6\\% for a training set of 40 SINGS galaxies (ranging from stellar to dust dominated) with both suitable mid-infrared \\emph{Spitzer} IRS spectra and equivalent photometry. A potential factor of two improvement could be realized with suitable $5.5\\mum$ and $10\\mum$ photometry, such as what may be provided in the future by \\emph{JWST}. The resulting technique is then applied to mid-infrared photometry for the 258 galaxies from the Local Volume Legacy (LVL) survey, a large sample dominated in number by low-luminosity dwarf galaxies for which obtaining comparable mid-infrared spectroscopy is not feasible. We find the total LVL luminosity due to five strong aromatic features in the $8\\mum$ complex to be $2.47\\times10^{10}\\,\\Lsun$ with a mean volume density of $8.8\\times10^6\\,\\Lsun\\,\\mathrm{Mpc}^{-3}$. Twenty-four of the LVL galaxies, corresponding to a luminosity cut at $\\mathrm{M}_{\\mathrm{B}} = -18.22$, account for 90\\% of the aromatic luminosity. Using oxygen abundances compiled from the literature for 129 of the 258 LVL galaxies, we find a correlation between metallicity and the aromatic to total infrared emission ratio but not the aromatic to total $8\\mum$ dust emission ratio.  A possible explanation is that metallicity plays a role in the abundance of aromatic molecules relative to the total dust content, but other factors such as star formation and/or the local radiation field affect the excitation of those molecules. ",
        "introduction": "It is well established that the infrared spectra of galaxies tend to be dominated by emission from dust \\citep[see, e.g., the review by][]{1979rie17araa477}. This light is a significant fraction of the bolometric emission from galaxies, making up roughly half the extragalactic background luminosity \\citep[e.g.,][]{2005lag43araa727}.  Investigation into the composition of the responsible dust grains is hindered by the small number of known associated emission and absorption lines.  The strongest features, by far, are a set of emission bands in the $3-17\\,\\mum$ range that are commonly observed in galaxies.  In luminous galaxies where the mid-infrared emission is dominated by star formation, they appear to be ubiquitous \\citep{1991roc248mnras606} with little \\citep{2003lu588apj199}, but not negligible \\citep{2007smi656apj770}, relative variation in feature strength.  These features have been referred to as the mid-infrared emission bands, the unidentified infrared bands, aromatic emission features, and variations thereof.  They have been tentatively identified as arising from polycyclic aromatic hydrocarbons \\citep[PAHs;][]{1984leg137aap5} and are often cited as resulting from a mixture of such particles. Regardless of the exact composition, it is widely agreed that the features in question arise from various bending and stretching modes of aromatic molecules largely composed of carbon and hydrogen; therefore, we refer to the resultant emission simply as aromatic emission or aromatic feature emission (AFE).  While present in nearly all star formation dominated galaxies, aromatic features do exhibit some dependence on galaxy host properties. It has been known for years that they can be weaker, or absent altogether, in systems characterized by low metallicity and high star formation intensity \\citep{1991roc248mnras606,1999thu516apj783}.  \\citet{2000mad44newar249} suggested that metallicity might be an important parameter, and this was shown to be true using the \\emph{Spitzer Space Telescope} \\citep{2004wer154apjs1} by \\citet{2005eng628apj29} and others \\citep{2006wu639apj157,2006jac646apj192,2008ros674apj814} and by detailed examination of data from the \\emph{Infrared Space Observatory} \\citep[ISO;][]{2006mad446aap877}. Further, the weakening of aromatic features in the low-metallicity outskirts of the galaxy M101 was attributed by \\citet{2008gor682apj336} to the hardness of the radiation field, which was also suggested to be a factor by the work of \\citet{2006mad446aap877} and \\citet{2008eng678apj804}.  \\begin{figure*}[t] \\begin{center} \\includegraphics[width=7.1in]{fig1.eps} \\end{center} \\caption{ (a) Spectral response of the $J$, $K_s$, IRAC 3.6/4.5/5.8/8.0 $\\mum$, and MIPS $24\\mum$ broad bands (the lighter regions below indicate wavelength ranges where the response is greater than half of the maximum). (b) IRS spectrum of NGC~4536 (solid black line) and its constituent flux contributions.  The stellar model (blue) is an interpolation of the 100 Myr Starburst99 SEDs (for the metallicity of NGC~4536), extinction-corrected and scaled to match the $J-K_s$ color and $3.6\\mum$ flux of NGC~4536.  The aromatic features (green), atomic/molecular lines (purple), and dust continuum (red) were simultaneously fit using PAHFIT.} \\label{fig_sed} \\end{figure*}  Investigations of mid-infrared aromatic feature emission are generally based on spectroscopic observations made from above the Earth's atmosphere; for example, with \\emph{Spitzer}'s Infrared Spectrograph \\citep[IRS;][]{2004hou154apjs18}. As such, they are largely relegated to samples of galaxies that are relatively small and/or bright.  Taking advantage of the fact that several of the strongest features fall within the rest-frame $8\\mum$ band (see Figure~\\ref{fig_sed}) of the Infrared Array Camera \\citep[IRAC;][]{2004faz154apjs10} on {\\it Spitzer}, \\citet{2005eng628apj29,2008eng678apj804} introduced a purely photometric approach to measuring the strengths of aromatic features. However, the agreement between photometrically and spectroscopically determined AFE equivalent widths for 27 starburst galaxies in \\citet{2008eng678apj804} was characterized by a Spearman rank-order correlation coefficient of only $r_s = 0.46$.  In this paper, we refine the photometric prescription for measuring aromatic feature emission at $8\\mum$, demonstrate its improved reliability, and apply it to the statistically complete sample of 258 local star-forming galaxies comprising the Local Volume Legacy survey \\citep[LVL;][]{2009dal703apj517}.  By doing so, we hope to advance future AFE investigations with a larger, more diverse, and less biased data set. In \\S~\\ref{sec_galaxies}, we introduce the LVL galaxies as well as an additional spectroscopic sample used for diagnostic purposes.  The corresponding datasets are described in \\S~\\ref{sec_data}, while in \\S~\\ref{sec_methods} our methodology is detailed and tested. Finally, we present an inventory of LVL aromatic feature emission measurements as well as a brief characterization of them in \\S~\\ref{sec_lvlafe} before summarizing our findings in \\S~\\ref{sec_concl}. ",
        "conclusions": "We have presented a purely photometric methodology for measuring the strength of the emission from mid-infrared aromatic features in nearby galaxies that are either too faint, too numerous, or too spatially extended for spectroscopy to be feasible. \\emph{Spitzer} IRAC (3.6, 4.5, and 5.8 $\\mum$) and MIPS $24\\mum$ photometry are used to disentangle contributions to the IRAC $8\\mum$ flux density from starlight, the dust continuum, and the $8\\mum$ complex of aromatic features. The latter is then calibrated to approximate the integrated flux one would obtain spectroscopically for the individual features at 7.42, 7.60, 7.85, 8.33, and 8.61 $\\mum$. With carefully matched photometry, we demonstrate that the photometric technique recovers the spectroscopically derived aromatic fraction of the total $8\\mum$ luminosity in a sample of 40 SINGS galaxies with a maximal relative difference of 12\\% and an rms of 6\\%.  Given suitable 5.5 and 10 $\\mum$ photometry, the rms could decrease by a factor of two in future studies.  The photometric approach was then used to measure aromatic emission from galaxies in the Local Volume Legacy sample, a \\emph{Spitzer} IRAC and MIPS survey of a statistically complete and nearly volume-limited sample of 258 nearby galaxies within 11 Mpc. In doing so, we have presented the first inventory of aromatic feature emission for a statistically complete sample of star-forming galaxies in the local universe. The total luminosity in the LVL sample resulting from the $8\\mum$ complex of five strong aromatic features is $2.47\\times10^{10}\\,\\Lsun$ with a mean volume density of $8.8\\times10^6\\,\\Lsun\\,\\mathrm{Mpc}^{-3}$. The corresponding values for all mid-infrared aromatic features in the wavelength range $5.5\\mum<\\lambda<20\\mum$ are larger by a factor of approximately 1.9. Twenty-four of the LVL galaxies, corresponding to $\\mathrm{M}_\\mathrm{B} > -18.22 \\pm 0.87$ and $\\LTir > 10^{9.25 \\pm 0.17}\\,\\Lsun$, are responsible for 90\\% of the total $8\\mum$ aromatic luminosity.  For a given LVL galaxy, the $8\\mum$ complex features contribute, on average, 2.5\\% of the total infrared luminosity, or roughly half that of the more luminous SINGS galaxies.  As a whole, the combined contribution for the LVL sample is 5.1\\%.  With respect to the total dust emission in the IRAC $8\\mum$ band, the aromatic contribution ranges from 0\\% to 90\\%.  Similarly, the aromatic features are responsible for between 0\\% and 80\\% of all the light in the IRAC $8\\mum$ band.  Oxygen abundances compiled from the literature were presented for 129 of the 258 LVL galaxies.  These are shown to follow the well-known luminosity-metallicity relationship with a best fit for the IRAC $4.5\\mum$ band of $12+\\log\\mathrm{(O/H)} = 5.06\\pm0.04 - (0.164\\pm0.002)\\,\\mathrm{M}_{4.5}$. We find that the total infrared luminosity and oxygen abundance correlate weakly with the aromatic to total infrared flux ratio and tentatively identify the latter relationship as being more fundamental.  The trend with oxygen abundance breaks down when the total dust emission at $8\\mum$ is used to scale the aromatic emission instead of the total infrared emission from $3-1100\\,\\mum$.  One possible explanation for this is that metallicity affects the abundance of aromatic molecules, while other factors such as star formation and/or the local radiation field determine their excitation."
    },
    "1004/1004.3734_arXiv.txt": {
        "abstract": "Stochastic effects during inflation can be addressed by averaging the quantum inflaton field over Hubble-patch sized domains. The averaged field then obeys a Langevin-type equation into which short-scale fluctuations enter as a noise term. We solve the Langevin equation for a inflaton field with Dirac Born Infeld (DBI) kinetic term perturbatively in the noise and use the result to determine the field value's Probability Density Function (PDF).  In this calculation, both the shape of the potential and the warp factor are arbitrary functions, and the PDF is obtained with and without volume effects due to the finite size of the averaging domain. DBI kinetic terms typically arise in string-inspired inflationary scenarios in which the scalar field is associated with some distance within the (compact) extra dimensions. The inflaton's accessible range of field values therefore is limited because of the extra dimensions' finite size. We argue that in a consistent stochastic approach the distance-inflaton's PDF must vanish for geometrically forbidden field values. We propose to implement these extra-dimensional spatial restrictions into the PDF by installing absorbing (or reflecting) walls at the respective boundaries in field space. As a toy model, we consider a DBI inflaton between two absorbing walls and use the method of images to determine its most general PDF.  The resulting PDF is studied in detail for the example of a quartic warp factor and a chaotic inflaton potential. The presence of the walls is shown to affect the inflaton trajectory for a given set of parameters. ",
        "introduction": "\\label{sec:intro}  Putting the successful inflationary scenario on the firm footing of a fundamental theory is one of the remaining challenges in cosmology. Recent years have seen considerable progress towards this goal with the construction of several concrete string inflation models; for recent reviews, see \\eg~\\cite{McAllister:2007bg,Kinney:2009vz,Baumann:2009ni,Lorenz:2010nx} and references therein. A top-level distinction among these models is the either closed or open string mode character of the inflaton field. A typical example of the second class are brane inflation scenarios~\\cite{Dvali:1998pa,Burgess:2001fx,Kachru:2003sx}: the inflaton field $\\phi$ corresponds (up to renormalization) to the distance between two branes embedded in a higher-dimensional background. While there is an ongoing debate about the form of the inflationary potential $V(\\phi)$~\\cite{Kachru:2003sx,Baumann:2006th,Baumann:2007np,Baumann:2007ah, Krause:2007jk,Pajer:2008uy,Baumann:2008kq}, a generic feature of these models is the field's kinetic term, which is of Dirac Born Infeld (DBI) type rather than canonic~\\cite{Silverstein:2003hf,Alishahiha:2004eh}. One can understand these DBI dynamics as a geometry-imposed upper limit on the field's velocity $\\dot{\\phi}$, characterized by the so-called warp factor $T(\\phi)$. This relativistic speed limit acts like a brake on the inflaton, forcing $\\phi$ to ``slow-roll'' even in regions where the potential is not flat. Hence, open string mode inflaton models provide an additional mechanism to generate quasi-exponential expansion.  \\par  It is interesting to investigate the DBI analogues of standard inflationary calculations, such as the field perturbations' evolution and spectra~\\cite{Lorenz:2008je,Lorenz:2008et}, but also the effects of stochastic inflation~\\cite{Starobinsky:1986fx,Starobinsky:1994bd,Martin:2005ir, Martin:2005hb,Kuhnel:2008yk,Finelli:2008zg,Kuhnel:2008wr,Seery:2009hs, Kuhnel:2010pp,Finelli:2010sh}. Stochastic inflation provides a technique to assess quantum effects on the inflaton $\\phi$'s trajectory averaged over a scale beyond the Hubble patch. In this way, one can define a coarse-grained field $\\varphi$ which consists exclusively of the large-scale Fourier components. To zeroth order, $\\varphi$ obeys the slow-roll Klein Gordon equation, but its full evolution is subject to stochastic noise $\\xi$ from small-scale Fourier modes. For standard inflation with arbitrary potentials, this equation was solved perturbatively up to $\\mathcal{O}(\\xi^{2})$ in Ref.~\\cite{Martin:2005ir}. Using this perturbative solution, Ref.~\\cite{Martin:2005ir} showed how to obtain the Probability Density Function (PDF) for $\\mean{\\varphi}$ in the Gaussian approximation, with and without the volume effects (\\ie the size of each averaging domain) taken into account. The reliability of this treatment was further studied in Ref.~\\cite{Martin:2005hb}. In this paper, we generalize both of these results to the case of DBI inflation with arbitrary warp factors and potentials.  \\par  The Langevin equation for DBI models was given for the first time in Ref.~\\cite{Chen:2006hs}, which aimed at studying ``eternal inflation''~\\cite{Vilenkin:1983xq,Linde:1986fd} in the brane inflation context. Since then, stochastic effects in DBI (or general $k$-inflation) models have been studied by several authors~\\cite{Helmer:2006tz,Tolley:2008na}. In this paper, we use the approach of Refs.~\\cite{Martin:2005ir, Martin:2005hb} to calculate the PDFs of a DBI inflaton field (with and without volume effects) for arbitrary functional form of the potential and warp factor. We illustrate our results by applying them to the example model of ``chaotic Klebanov Strassler (CKS) inflation'', where $T(\\phi)$ and $V(\\phi)$ are known functions. The integrals encountered are, within certain limits, exactly calculable, and we discuss the behavior and reliability of the resulting DBI inflaton probability densities.  \\par  However, we argue that all PDFs used so far in the literature suffer from a serious problem: they predict a non-vanishing probability for the moving $D3$-brane to find itself outside the so-called ``Klebanov Strassler (KS) throat'', \\ie outside the (part of the) extra-dimensional geometry with warp factor $T(\\phi)$. In fact, this problem is twofold. Firstly, at the bottom end of the throat this means that there is a non-vanishing probability to literally find the brane ``outside'' the extra dimensions, in other words ``out of space'', which is clearly meaningless. Secondly, since the metric of the 6d bulk space [and hence the continuation of $T(\\phi)$] is typically unknown beyond the KS region, a string-inflationary scenario based on the brane's motion inside the throat becomes inconsistent beyond its top end. Note that, while the latter question is of a more technical nature and we may hope to resolve it as our understanding of string geometries improves, the former issue is rather severe as the inflaton's PDF does not respect the fundamentally geometric origin of the scenario.  \\par Put a different way, the compact character of the stringy extra dimensions (for the purposes of concrete model building, this means the well-known KS ``corner'' of the 6d manifold) can be translated directly into a restricted field range for the DBI inflaton. For consistency, any modifications of the classical trajectory induced by stochastic effects should still respect these geometry-imposed boundaries in field space. Hence, studies of stochastic DBI inflation so far were missing a tool to ensure the consequences of a stringy inflaton's geometric interpretation at the effective field theory level.  \\par  To amend this problem, we propose to install ``walls'' at the boundaries (\\ie the bottom $r_{0}$ and the edge $r_{\\UV}$) of the inflationary KS throat (whose radial coordinate is denoted by $r$). As a consequence, the stochastically corrected inflaton field value should remain within its allowed range $\\phi_{0}<\\phi<\\phi_{\\UV}$ at all times. This requires the calculation of a new PDF respecting the boundary conditions imposed by the presence of the reflecting or absorbing walls. For the first time, we then determine this PDF in the presence of two absorbing walls at $\\phi_{0}$ and $\\phi_{\\UV}$, using the method of images, see \\eg \\cite{Sommerfeld:ED, Sommerfeld:PDEs}.  As a by-product, we obtain the modified stochastic trajectory of the mobile brane within the KS throat and show that, in some cases, the presence of the walls has a significant effect.  \\par  This paper is organized as follows. In the next Section, we start from the underlying background equations of DBI inflation and discuss how they can be used to formulate their stochastic counterpart, \\ie the DBI Langevin equation. We pay special attention to the normalization factor of the noise term. We then solve the DBI Langevin equation up to second order in the noise and calculate the corresponding PDF along with its volume correction. In Sec.~\\ref{sec:applications}, we use these expressions to calculate the stochastic effects in the case of chaotic Klebanov Strassler inflation, both for a potential with and without a constant term. Where applicable, we calculate the domain of validity of the perturbative approach and consider the existence of a regime of eternal inflation~\\cite{Vilenkin:1983xq,Linde:1986fd}.  In Sec.~\\ref{sec:finitesize}, we use the additional information on the inflaton's geometric bounds as an argument to implement two absorbing walls into the calculation of the PDF. In this way, it can be assured that even quantum effects do not violate these field space limits. Finally, we summarize our main findings in Sec.~\\ref{sec:conclusions}, commenting on eternal inflation as well as on an overall picture for a realistic brane trajectory across the entire compact 6d geometry. ",
        "conclusions": "\\label{sec:conclusions}  We now conclude our investigation by revisiting our main results. In this paper, we have generalized the approach of Refs.~\\cite{Martin:2005ir,Martin:2005hb} to DBI inflation models, in which the kinetic term of the inflaton is modified through an geometry-imposed upper limit on the field velocity. To this end, we solved the DBI Langevin equation using a perturbative expansion in the noise. It turns out that the results of Refs.~\\cite{Martin:2005ir,Martin:2005hb} essentially only change by additional factors of $\\gamma$, reflecting the fact that the distinction between long and short wavelength fluctuations, which is at the heart of the stochastic inflationary approach, now is defined with respect to the sound (instead of the Hubble) radius scale.  \\par  Our calculation yields easy expressions for the PDFs describing the probability for the patch-averaged inflaton to assume a given value in one Hubble domain, or in the entire Universe taking into account the size of each averaging domain. We calculated and plotted these PDFs for the example of chaotic Klebanov Strassler inflation, both for the cases with and without a constant term in the potential. In the absence of a constant term $V_{0}$, stochastic effects are found not to alter the classical field trajectory in a significant way. However, in both the UV ($\\varepsilon=+1$) and IR ($\\varepsilon=-1$) formulation of a potential including a constant term, the quantum behavior of the field plays a dominant r\\^ole for small field values, \\ie at the bottom of the Klebanov Strassler throat. Hence, in the UV model, where the brane moves downwards in the throat, these effects occur at the end of inflation, whereas in the IR case, with the brane climbing up the throat, they come into play at the very beginning.  \\par  The main result of this paper is our demonstration that an additional subtlety arises due to the geometric restriction which confines the field value of the coarse-grained field between $\\phi_{0}<\\varphi<\\phi_{_{\\rm UV}}$: to prevent the PDFs from extending below $\\phi_{0}$ (where the throat --and, consequently, the stringy spacetime itself-- ends), we introduced two absorbing walls located at $\\phi_{0}$ and $\\phi_{_{\\rm UV}}$.  This changes the PDF's shape, and within the bounds of validity it is the ``absorbing walls probability density functions'' that must be used to describe the stochastic behavior of the inflaton field. It should be interesting to put this technique to use for the ``DBI version'' of quintessence proposed in Ref.~\\cite{Martin:2008xw}.  \\par  The question of existence of an eternally self-reproducing regime in the DBI case is more complicated. We have argued that such a regime does not exist if the potential is of the pure $m^{2}\\phi^2$ type, but that in the case of an additional constant term $V_{0}$, eternal inflation seems possible.  Given the intrinsic validity limitations of our present approach, a fully numerical treatment of the DBI Langevin equation would be necessary to answer this question.  \\par  Finally, we would like comment on possible brane trajectories beyond the single throat model. Globally, the string geometric 6d background is given by a compactified bulk Calabi Yau space (where the metric is unknown) whose ``corners'' may comprise several (possibly generalized) Klebanov Strassler throats with different parameters. Some of these throats can contain additional anti-branes at their bottom, like in the original scenario of Ref.~\\cite{Kachru:2003sx}.  One could therefore imagine that a brane starts out at the bottom of an IR throat, moving upwards towards the bulk, then transverses a short distance within the bulk before dropping down into a UV throat.  \\par  Let us consider qualitatively the quantum effects experienced by a brane following such a trajectory: While in the IR throat, the impact of stochastic inflation is strongest at the very bottom, \\ie quantum effects will push the brane upwards even faster than its classical evolution demands. Having made its way across the bulk, the brane will descend into the UV throat first on a purely classical trajectory. However, as it nears the bottom of the second throat, the stochastic influence grows again, pushing the brane away from the bottom and possibly upwards into the bulk again. At the limit, one may therefore imagine a scenario where the brane is propelled out of a UV throat back into the bulk every time it comes sufficiently close to a throat's bottom.  \\par  Since we ignore the formulation of the metric in the bulk, it seems difficult to carry out the full calculation for the brane trajectory sketched above. However, we do know how to describe the metric within each individual throat. As a first toy model, we therefore considered the case of a brane trapped between two absorbing walls located at one throat's bottom and edge, respectively. This successfully prevents the brane from meandering ``outside of spacetime'' (below the bottom of the throat at $\\phi_{0}$), and we have therefore solved the severe conceptual issue stated in the introduction of this paper. Near the edge of the throat at $\\phi_{_{\\rm UV}}$, the technical difficulty of describing the stochastic effects remains because to-date the 6d bulk metric remains unknown. The next step would be to write down a joint warp factor for two throats of different depth which are ``glued together'' at their edges. Given the quantum behavior found in this paper, our expectations might point us towards an increased probability to find the brane near the ``throat matching point''."
    },
    "1004/1004.3698.txt": {
        "abstract": "We present a catalog of 93 very-well-observed nova light curves.  The light curves were constructed from 229,796 individual measured magnitudes, with the median coverage extending to 8.0 mag below peak and 26\\% of the light curves following the eruption all the way to quiescence.  Our time-binned light curves are presented in figures and as complete tabulations.  We also calculate and tabulate many properties about the light curves, including peak magnitudes and dates, times to decline by 2, 3, 6, and 9 magnitudes from maximum, the time until the brightness returns to quiescence, the quiescent magnitude, power law indices of the decline rates throughout the eruption, the break times in this decline, plus many more properties specific to each nova class.  We present a classification system for nova light curves based on the shape and the time to decline by 3 magnitudes from peak ($t_3$).  The designations are \u00d4S\u00d5 for smooth light curves (38\\% of the novae), \u00d4P\u00d5 for plateaus (21\\%), \u00d4D\u00d5 for dust dips (18\\%), \u00d4C\u00d5 for cusp-shaped secondary maxima (1\\%), \u00d4O\u00d5 for quasi-sinusoidal oscillations superposed on an otherwise smooth decline (4\\%), \u00d4F\u00d5 for flat-topped light curves (2\\%), and \u00d4J\u00d5 for jitters or flares superposed on the decline (16\\%).  Our classification consists of this single letter followed by the $t_3$ value in parentheses; so for example V1500 Cyg  is S(4), GK Per is O(13), DQ Her is D(100), and U Sco is P(3). ",
        "introduction": "Nova eruptions (Payne-Gaposhkin 1964) occur on the surface of white dwarf stars in interacting binaries, generically called `cataclysmic variables' (Warner 1995), usually with a red dwarf companion star, where material accumulates on the surface until the pressure and temperature are high enough that a thermonuclear runaway reaction occurs (Bode \\& Evans 2008).  During a nova eruption, a normally faint star is observed to brighten by 8-20 magnitudes to peak within a few days, to fade initially fairly quickly over a timescale of weeks to months, and then slowly fade back to the quiescent level.  The light curves of novae are usually the best-measured feature of the events, and are seen to vary substantially from system to system.  The light curves carry much information about each nova.  The literature is full of nova magnitudes and fragmentary light curves.  Only a handful of the brightest novae (e.g., V1500 Cyg, T CrB, RS Oph, V603 Aql, DQ Her) have well observed published light curves.  Of these five examples of the best, only two were followed to quiescence in the professional literature, while all were followed to quiescence by amateur observers reporting to the American Association of Variable Star Observers (AAVSO).  Even amongst the 13 classical novae presented as ``First Class'' by Payne-Gaposhkin (1964), that is the best-of-the-best up until 1964, only five have professional light curves going to quiescence and only two of those appear in the book.  In modern times, little has changed, with RS Oph having wonderful coverage in 2006 but most novae having only a handful of magnitudes near peak published.  In one of the better cases, V5114 Sgr, one of us helped push for many professional observations, yet still our omnibus paper (Ederoclite et al. 2006) only has 46 V-band magnitudes not even going to 5.5 mag below the peak, and the position is mis-identified in the Downes \\& Shara Catalog (Downes \\& Shara 1993; Downes et al. 1997; 2001; henceforth referred to collectively as D\\&SCat).  For another recent example, Figure 1a displays all the V-band magnitudes published in the professional literature for V705 Cas, and we see only 220 magnitudes with no coverage after the end of the prominent dip in the light curve.  And this is for one of the {\\it best} observed events in recent decades.  For a set of light curves with the best coverage of all novae, one of us (Schaefer 2010, henceforth S2010) has made an exhaustive compilation of essentially all photometry of all ten galactic recurrent novae (RNe, RN singular).  This comprises 37 known eruptions.  Of these, 16 are known with fewer than 20 magnitudes, while 8 have 20-100 points on their light curves.  Such coverage is largely useless for classifying or understanding the light curves.  Only 6 of these nova light curves (for T CrB and RS Oph) have been followed to quiescence.  Excepting these six eruptions, the median magnitude below peak of this coverage is just six magnitudes.  Even the famous U Sco has never been followed past 30 days after maximum.  Of the 13 eruptions with more than 100 magnitudes, all of them have the majority (typically 80-98\\%) of their data from the AAVSO data base.  This sampling of nova light curves shows a wide range of coverage quality, mostly poor yet with a useful fraction of good coverage provided only by the AAVSO.  These RNe have attracted much attention (at least after their first eruption), so this sampling of light curve coverage is better than for ordinary novae.  Several prior works have collected large numbers of nova light curves (McLaughlin 1939; Payne-Gaposhkin 1964; van den Bergh \\& Younger 1987).  McLaughlin (1939) displayed only 17 light curves, with median coverage going to 7 magnitudes below peak and none going to quiescence.  Payne-Gaposhkin (1964) presented 43 light curves, 15 of which are copied from McLaughlin; in total, the coverage goes to a median of 5 magnitudes below peak.  van den Bergh \\& Younger (1987) collected all nova light curves with UBV photometry and could only find 32 novae, for which their light curves have a median coverage of only 26 measures extending just 2.7 mag below peak.  The classification of classical nova light curves is in a similar primitive state.  McLaughlin (1939) merely divided the light curves into `fast' and `slow'.  Payne-Gaposhkin (1964) made five finer divisions (`very fast', `fast', \\ldots) based solely on the time it took the light curve to decline by two magnitudes from peak ($t_2$).  Payne-Gaposhkin never used her classification system for any purpose, nor has anyone else.  Duerbeck (1981) largely kept the speed classes as the primary classification parameter, but added a new class for those novae that showed `dust dips' (the sharp drop to a faint level followed by a fast recovery of brightness like seen for DQ Her) caused by the formation of dust in the expanding nova shell that hides the interior light.  His five classes are fast/smooth, irregulars, dust dip events, slows, and very slows.  The only later attempt to classify nova light curves that we are aware of is by Rosenbush (1999a; 1999b) who plots \u00d2outburst amplitude\u00d3 versus the logarithm of the main shell radius.  The primary reason why none of these classification schemes has been usefully adopted is that the classes are so broad that greatly different physical settings are lumped together and they have not been found to correlate with anything else.  The measure of nova light curve properties is still fairly primitive.  The only properties that are known for a significant number of novae are the peak magnitude and the time it takes to decline by two or three magnitudes below peak ($t_2$ and $t_3$ respectively).  The most complete compilation from the literature, a heroic task in itself, is given in Shafter (1997).  Burlak \\& Henden (2008) used the AAVSO data base to derive the peak magnitude, $t_2$, and $t_3$ for 64 novae, with no other properties being measured.  But this work was seriously flawed because no attempt was made to include observations from before the AAVSO data started, so the compilation missed the peak of about half the novae, and the decline times are correspondingly wrong.  Recently, theorists and observers have been starting to ask questions about nova light curves that need measures from many events.  Hachisu \\& Kato (2006; H\\&K06) have proposed a `universal decline law', where they make specific predictions as to the power law slopes at late times and early times, and the time when the break occurs is related closely to the mass of the white dwarf.  They have been systematically collecting many light curves and fitting them to broken power laws to find white dwarf masses for many individual systems (Hachisu \\& Kato 2007; 2008; Hachisu et al. 2009a).  Independently, they have also been pointing to the existence of plateau phases as being a hallmark that a nova is recurrent (Hachisu et al. 2000; 2002; 2003; 2009a).  This has caused one of us to measure plateau parameters for the known RNe, and to ask what fraction of classical novae show similar plateaus (S2010).  Pejcha (2009) has been looking at a mixture of oscillations and flares (jitters) in seven novae trying to find whether the peak spacing is uniform in logarithmic or linear time, so as to test his theoretical model.  In this paper, we present a catalog of 93 well-sampled nova light curves.  Almost all our light curves are all for the V-band as taken from the AAVSO database.  This sample solves the problems described above, in that we have a large number of light curves with extensive time coverage (often to quiescence) that go deep and have many observations each day.  With this large sample, we have created a new classification for nova light curves, with the categories more in line with modern ideas.  We also have measured many light curve properties for all of the novae, with our measures hopefully being useful to theorists.  We present plots and tabulated light curves for all novae.  Finally, we have sought and found a variety of correlations between our nova classes and our measured properties.  These new correlations and fits now form a challenge to theorists to explain their physical origin. ",
        "conclusions": ""
    },
    "1004/1004.4608_arXiv.txt": {
        "abstract": "Precision radial velocities from Keck/HIRES reveal a Saturn-mass planet orbiting the nearby M4V star HIP 57050. The planet has a minimum mass of $M \\sin i \\sim 0.3 M_{\\rm J}$, an orbital period of 41.4 days, and an orbital eccentricity of 0.31. V-band photometry reveals a clear stellar rotation signature of the host star with a period of 98 days, well separated from the period of the radial velocity variations and reinforcing a Keplerian origin for the observed velocity variations. The orbital period of this planet corresponds to an orbit in the habitable zone of HIP 57050, with an expected planetary temperature of $\\sim$230 K. The star has a metallicity of [Fe/H] = $0.32\\pm0.06$ dex, of order twice solar and among the highest metallicity stars in the immediate solar neighborhood. This newly discovered planet provides further support that the well-known planet-metallicity correlation for F, G, and K stars also extends down into the M-dwarf regime. The {\\it a priori} geometric probability for transits of this planet is only about 1\\%. However, the expected eclipse depth is $\\sim 7\\%$, considerably larger than that yet observed for any transiting planet. Though long on the odds, such a transit is worth pursuing as it would allow for high quality studies of the atmosphere via transmission spectroscopy with HST. At the expected planetary effective temperature, the atmosphere may contain water clouds. ",
        "introduction": "Due to their low masses and surface temperatures, M dwarfs present the most promising targets for searching for terrestrial-mass and potentially habitable planets. As the least massive stars, these objects experience the greatest reflex accelerations in response to an orbiting planet. This advantage was first realized with the detection of a Neptune-mass extrasolar planet around the star GJ 436 \\citep{Butler04}, and the first super-Earth around the star GL 876 \\citep{Rivera05}. The low surface temperatures of M dwarfs place their (liquid water) habitable zones at conservative distances of approximately 0.1 AU to 0.2 AU. These distances correspond to orbital periods of 20 to 50 days, implying another advantage of M dwarfs as potential targets for detecting habitable planets relatively quickly.  As precision Doppler surveys are optimally sensitive to small orbits, it is not surprising that terrestrial-mass planets around M dwarfs, in particular those in the habitable zone, have been the subject of research for more than a decade \\citep{Joshi97,Segura05,Boss06,Scalo07,Grenfell07,Tarter07}. During the past few years, such research resulted in the detection of 17 extrasolar planets around 12 M dwarfs\\footnote{We refer the reader to exoplanet.eu for more details.}. Slightly more than half of these planets are Neptune-mass or smaller, consistent with the fact that M dwarfs have smaller circumstellar disks, and experience has shown that they are less frequently accompanied by readily detectable planets, and/or their planets are less massive compared to those of G stars.   While the majority of the currently known extrasolar planets have been detected around nearby F,G, and K stars, more than 70\\% of the nearest stars are M dwarfs. For the past decade, we have had a sample of $\\sim$300 nearby quiet stars under precision radial velocity survey (P.I. Butler's NASA M-dwarf Exoplanet Survey) with the Keck telescope and its HIRES spectrometer. Here, we present 9.9 years of precision radial velocities for the nearby M4 dwarf HIP 57050 and report the detection of the exoplanet they imply. ",
        "conclusions": "In the quest for potentially habitable planets, the nearest stars are of special importance. They have accurate distances and precisely determined stellar parameters, and are the only stars for which follow up by astrometry and direct imaging is possible. Within the Sun's immediate neighborhood, M-dwarfs constitute the majority of nearby stars. As such, these stars have the special properties (distances, masses, and habitable zones) that drive exoplanetary science, astrobiology, and the next generation of interferometry and direct imaging missions. The habitable (liquid water) zones of nearby M-dwarfs are typically between 0.1 AU and 0.2 AU which corresponds to orbits with periods of 20 to 50 days. Establishing (by direct detection) the prevalence and nature of low-mass planets, such as HIP 57050 b, in these orbits informs us greatly about the possibility for potentially habitable planets (and/or moons) in the solar neighborhood.  A zeroth-order prediction of the core-accretion paradigm for giant planet formation is that the frequency of readily detectable giant planets should increase with both increasing stellar metallicity and with increasing stellar mass \\citep{Laughlin04,Ida05}. During the past decade, both of these trends have been established observationally \\citep[see, e.g.][for a discussion of the metallicity trend and Johnson et al. 2009 for a discussion of the mass trend]{Fischer05}. Until recently, however, there appeared to be little evidence for the strong expected planet-metallicity correlation among the handful of M-dwarf stars that are known to harbor giant planets. Attempts to determine accurate metallicities of M-dwarfs have largely been stymied by ambiguity in the continuum levels of their heavily line-blanketed spectra and by the profusion of molecular features in their spectra. Conventional estimates for the metallicities of 7 of the currently known planet-hosting M dwarfs as given by \\citet{Bailey09}, and a comparison between these estimates and those of \\citet{Schiavon97} and \\citet{Bean06}, suggest a spread of metallicity among these M-dwarfs (four are metal poor, one has high metallicity, and the metallicities of the remaining two are solar).  One would naively expect that a low-mass disk will need all the help it can get in order to build giant planet cores before the gas is gone. If anything, the planet-metallicity correlation should be strongest among the M-dwarfs. If observations show that the planet-metallicity correlation breaks down for M-dwarfs, then one is naturally led to speculate that the infrequent giant planets in a systems like Gliese 876 might be the outcome of gravitational instability \\citep[e.g.][]{Boss00} rather than core accretion.  \\citet{Bonfils05} pioneered a new approach to the determination of M-dwarf metallicities. The long evolutionary time scales for M-dwarfs imply that age-related $L$ and $T_{\\rm eff}$ changes should be minimal once a low-mass star has landed on the zero-age main sequence. M-dwarf positions on the color-magnitude diagram, therefore, should be parameterized only by mass and metallicity, opening the possibility of a metallicity determination based on ${\\rm M}_{\\rm K}$ and V-K alone. \\citet{Bonfils05} developed such a calibration by assuming that M-dwarf binary companions to F, G, and K stars share the readily determined metallicities of their primaries. \\citet{Johnson09} have recently provided an update to the \\citet{Bonfils05} calibration. The \\citet{Johnson09} calibration indicates that the planet-bearing M-dwarfs do appear to be systematically metal-rich, suggesting that there is no breakdown of the planet-metallicity correlation as one progresses into the red dwarf regime.   HIP 57050 b appears to offer further support for the emerging M-dwarf planet-metallicity correlation. Using HIP 57050's values, $V=11.88$, $K=6.822$, and $d=11.03\\,{\\rm pc}$, the \\citet{Johnson09} calibration yields $[{\\rm Fe/H}]=0.32\\pm0.06\\,{\\rm dex}$, indicating that HIP 57050 has a metallicity of order twice solar, which places it among the highest metallicity stars in the immediate solar neighborhood.   The {\\it a priori} geometric transit probability for HIP 57050 b is $\\sim 1\\% $. The small size of the primary star and the planet's unfavorable orbital alignment ($\\omega=238^{\\circ}$) conspire to diminish the odds that transits can be observed. An analysis of our photometry data also shows no signs of a transit. However, because the orbital elements can change if a second planet emerges, it is premature to conclude at this point that transits do not occur. The eclipse depth in this system is expected to be $\\Delta F/F \\sim 7\\%$, which is considerably larger than that yet observed for any transiting planet. Such a large depth makes this system suitable for small-telescope observers to check. We therefore suggest that small-telescope observers carry out photometric monitoring of HIP 57050 during the predicted transit windows centered on HJD 2455201.400239. The large planet-to-star ratio would allow for detailed study of the atmosphere via transmission spectroscopy with HST. The expected planetary effective temperature, $T_{\\rm eff}\\sim230\\,{\\rm K}$, suggests that the atmosphere may contain water clouds.  It is interesting to speculate about the possible presence of a habitable moon around HIP 57050 b. By analogy with our own solar system, whose gas giants all have dozens of moons, one might expect HIP 57050 b to also harbor such moons. In our solar system, $\\sim0.02\\%$ of the masses of the gas giants are assigned to their satellites. This would translate to a satellite with $\\sim2\\%$ of Earth's mass (similar to Titan) orbiting HIP 57050 b. While it is not out of the question that HIP 57050 b could harbor a moon, and that moon would thus be in the liquid water habitable zone of the parent star, an object with only 1/5th of the mass of Mars in the liquid water habitable zone, from various standpoints is probably not a particularly good prospect for habitability. In any case, direct detection of such a moon would be extremely challenging.  We conclude this study by noting that the Doppler radial velocity method continues to be the most productive and cost-effective way to find those extrasolar planets that impart the greatest scientific returns \\citep{Butler04,Rivera05,Lovis06,Udry07_HARPS,Mayor08,vogt10,Rivera10}. During the past several years, the threshold $M\\sin(i)$ for radial velocity planets has rapidly approached the 1$M_{\\oplus}$ regime. Radial velocity surveys, furthermore, have led to the discovery of all but one of the most readily characterizable transiting planets, and the rapidly growing catalog of Doppler-detected planets has been instrumental in providing our best current view of the nearby planetary population\\footnote {see the planet lists and correlation diagrams at exoplanet.eu}. The future looks bright!"
    },
    "1004/1004.3502_arXiv.txt": {
        "abstract": "{The southern \"infrared companion\" of T\\,Tau is known to show strong photometric variations of several magnitudes on timescales of years, as well as more modest $\\lesssim$1\\,mag variations on timescales as short as one week. The physical mechanism driving these variations is debated, intrinsic luminosity variations due to a variable accretion rate were initially proposed, but later challenged in favor of apparent fluctuations due to time-variable foreground extinction.} {We seek to investigate the nature of the observed photometric variability. Based on simple geometric arguments and basic physics laws, a minimum variability timescale can be derived for which variable extinction is a viable mechanism. Because this timescale increases rapidly with wavelength, observations at long wavelengths provide the strongest constraints.} {We used VISIR at the VLT to image the T\\,Tau system at two epochs in February 2008, separated by 3.94~days. In addition we compiled an extensive set of near- and mid-infrared photometric data from the literature, supplemented by a number of previously unpublished measurements, and constructed light curves for the various system components. We constructed a 2D radiative transfer model for the disk of T\\,Tau\\,Sa, consisting of a passively irradiated dusty outer part and a central, actively accreting component.} {Our VISIR data reveal a $+$26$\\pm$2\\% change in the T\\,Tau\\,S/T\\,Tau\\,N flux ratio at 12.8\\,$\\mu$m within four days, which can be attributed to a brightening of T\\,Tau\\,Sa. Variable extinction can be excluded as a viable mechanism for the observed flux variation based on the short timescale and the long observing wavelength. We show that also the large long-term photometric variability and its associated color-magnitude behavior can be plausibly explained with variable accretion. However, variable extinction is also a viable mechanism for the long-term variability, and a combination of both mechanisms may be required to explain the collective photometric variability of Sa.} {We conclude that the observed short-term variability is caused by a variable accretion luminosity in T\\,Tau\\,Sa, which leads to substantial fluctuations in the irradiation of the disk surface and thus induces rapid variations in the disk surface temperature and IR brightness. Both variable accretion and variable foreground extinction can plausibly explain the long-term color and brightness variations. We suggest that the periods of high and variable brightness of Sa that we witnessed in the early and late 1990s were due to enhanced accretion induced by the periastron passage of Sb, which gravitationally perturbed the Sa disk.} ",
        "introduction": "Six decades ago \\cite{ambartsumian1947,ambartsumian1949} first proposed T\\,Tauri stars to be newly formed low-mass stars. They had been defined earlier as a class of optically variable stars with emission lines in the vicinity of bright or dark nebulosity by \\cite{1945ApJ...102..168J}. Ever since, the name-giving object of the class, T\\,Tau, has been regarded as the proto-type of young low-mass stars. It is a K0 star with a mass of approximately 2\\,M$_{\\odot}$ and an age of 1-2\\,Myr, that has a modest and time-variable line of sight extinction of \\Av\\simil1\\,mag \\citep{1974A&AS...15...47K,2005AstL...31..427M,2007ApJ...671..546L}. It resides in the Taurus molecular cloud at a distance of 148\\,pc \\citep{2007ApJ...671..546L}.  An \"infrared companion\", discovered by \\cite{1982ApJ...255L.103D}, is located approximately 0\\farcs7 south of T\\,Tau. It was soon realized that the bolometric luminosity of this source (T\\,Tau\\,S) rivals that of the optically visible star (T\\,Tau\\,N), and that it must be a self-luminous object rather than a dust cloud heated by the northern component. Even though T\\,Tau\\,S is at least 10\\,mag fainter than the northern component in the optical \\citep{1998ApJ...508..736S}, it dominates the system flux at infrared wavelengths beyond \\simil3\\,\\mum \\ \\citep{1991AJ....102.2066G,1997AJ....114..744H,2004ApJ...614..235B}, but it contributes only a minor fraction to the total flux at millimeter wavelengths \\citep[][]{1997ApJ...490L..99H,1998ApJ...505..358A}. The southern component was itself found to be a close binary with components Sa and Sb at a projected separation of \\simil50\\,mas at the time of discovery \\citep{2000ApJ...531L.147K}, corresponding to only 7$-$8\\,AU at the system distance.  T\\,Tau\\,S was shown to be strongly variable at near- and mid infrared wavelengths by \\cite{1991AJ....102.2066G}, who detected a \\simil2\\,mag increase in brightness at 2.2\\,\\mum \\ within $\\sim$1\\,yr, and a similarly large increase at 10\\,\\mum \\ between two observations separated by about five years. They attributed the brightness fluctuations to variable accretion and derived an accretion rate of 3.6\\,$\\times$\\,$10^{-6}$\\,\\msunyr.  \\cite{2004ApJ...614..235B} presented a large set of (near-) simultaneous K and \\Lp \\ photometry obtained over the course of about seven years starting 1995, and found T\\,Tau\\,S to vary with a total amplitude of nearly 3\\,mag at K-band and more than 2\\,mag at \\Lp \\ during this period. The fastest brightness variation they detected occurred in December 1997, when T\\,Tau\\,S brightened by 0.9\\,mag in the K~band over only a seven day period. They observed the system during three consecutive nights in September 1998 and another three nights in November 1999 in order to detect even faster variations, but did not detect any. The brightness variations in T\\,Tau\\,S were shown to be color-dependent in a K-\\Lp \\ vs. K color-magnitude diagram, in a \"redder when fainter\" fashion that closely resembles the ISM extinction curve \\citep{1985ApJ...288..618R}. This behavior was not expected if the observed brightness fluctuations were due to variable accretion, in this case models predicted the opposite behavior, i.e. \"bluer when fainter\" \\citep{1997ApJ...481..912C}, mostly because of the changing contrast ratio between the blue stellar spectrum and the much redder disk spectrum as the accretion rate varies. Therefore \\cite{2004ApJ...614..235B} concluded that variable obscuration is the main responsible factor for the observed large brightness variations in T\\,Tau\\,S, although these authors did also find evidence for variable accretion in T\\,Tau\\,S \\citep{2001ApJ...551.1031B}.  A decade of spatially resolved observations of the southern binary pair have now covered a substantial fraction of the Sa-Sb orbit and have allowed direct determination of the orbital parameters \\citep{2000ApJ...531L.147K,2003ApJ...596L..87F,2004ApJ...614..235B,2006A&A...457L...9D,2008A&A...482..929K}. Current estimates still allow a substantial range for some of the orbital parameters, of which the orbital period and the mass of Sb are the most relevant for this paper. Continued observations over the coming years are expected to yield a converged solution. The mass of the primary Sa is \\simil2.2\\,\\Msun, the most recent estimates for the mass of Sb range from 0.4 to 0.8~\\Msun, and for the orbital period from \\simil25 to \\simil100 years \\citep{2008A&A...482..929K,2008JPhCS.131a2028K}. The spatially resolved observations have also revealed Sb to show only minor brightness variations at 2.2\\,\\mum \\ ($RMS$\\simil0.2\\,mag), while Sa has varied by more than 3\\,mag during the last decade at this wavelength.  \\vspace{0.1cm} Here we revisit the physical mechanism causing the observed photometric variations of T\\,Tau\\,S. This work was motivated by short-term mid-infrared variability that was serendipitously detected during an observing campaign with VLT/VISIR in February 2008, primarily aimed at characterizing the [Ne\\,{\\sc ii}] emission in the T\\,Tau system. We reported on the [Ne\\,{\\sc ii}] observations in a separate paper \\citep{2009A&A...497..137V} and here focus on the \\simil26\\% increase in continuum brightness of T\\,Tau\\,S at 12.8\\,\\mum \\ that we detected in two observations separated by only four days. In this paper we will distinguish between the very fast variations with modest amplitude as detected in our VISIR observations, and the long-term variations of much larger amplitude that are seen in the entire body of photometry taken over the past decades. These are not necessarily caused by the same physical mechanism.  The paper is organized as follows. In \\sek~\\ref{sec:observations} we describe the VISIR observations as well as additional previously unpublished infrared measurements performed with a range of facilities. The results of the infrared photometry are reported in \\sek~\\ref{sec:results}, in which we also construct updated near- and mid-infrared light curves of the various components of the system. In \\sek~\\ref{sec:variable_extinction} we explain why the short-term variability revealed in the VISIR observations cannot be due to variable extinction and must instead be due to intrinsic luminosity variations, i.e. variable accretion. In \\sek~\\ref{sec:radiativeTransfer} we address the nature of the long-term variability, which has previously been attributed to time-variable extinction. We present radiative transfer calculations of the disk of Sa that include variable accretion, and show that these models qualitatively reproduce the observed long-term color-dependent near- and mid infrared brightness variations. In \\sek~\\ref{sec:TTauS_scenario} we outline a tentative scenario for T\\,Tau\\,S in which the periods of high and variable brightness in which Sa has prevailed from \\simil1990 to \\simil2003 were due to enhanced accretion activity, quantitatively intermediate between an EXOR and a FUOR outburst, induced by gravitational perturbation of the Sa disk during the periastron passage of Sb in \\simil1995. We summarize our results in \\sek~\\ref{sec:summary}. ",
        "conclusions": "\\label{sec:discussion} In this section we will first discuss why the fast mid-infrared variability that we detected excludes variable extinction as a viable mechanism for the observed short-term mid-IR brightness fluctuations in T\\,Tau\\,S, {and argue that these must instead be due to variations in intrinsic luminosity.} Then we will present a radiative transfer disk model of Sa, including both an active accretion and a passive reprocessing component, that qualitatively reproduces the observed long-term color-magnitude variations when the accretion rate is varied. Thus we will show that variable accretion may also explain the long-term variability, but do not exclude variable extinction as the prime mechanism responsible for the long-term variations. Lastly, we sketch a tentative scenario for T\\,Tau\\,S, in which the \\simil15\\,yr period of irregular and enhanced brightness that we witnessed in the recent past is induced by the periastron passage of Sb, gravitationally disturbing the Sa disk and triggering an accretion outburst that may be reminiscent of EXOR variables.   \\subsection{Can the observed short-term brightness fluctuations be caused by variable extinction?} \\label{sec:variable_extinction}  For variable extinction to be a viable cause of the observed brightness fluctuations, it is a necessary condition that the opaque medium causing the extinction can (un-) cover the emitting region \\emph{within the timescale of the variations}. Because only Sa shows large brightness variations whereas T\\,Tau\\,N and Sb do not, one may assume the absorbing \"screen\" to be local to the system, i.e. at a distance of \\simil148\\,pc. Thus we can estimate the \\emph{minimum speed} at which the screen must travel to cause the observed variations. To this purpose, we will first calculate the minimum size of the emitting region.  Our sole constraint is that T\\,Tau\\,S brightened by +3.3\\,Jy within 94.5~hours at 12.8\\,\\mum. Because at 12.8\\,\\mum \\ the emission is completely dominated by thermal emission from dust grains, the maximum intensity the emitting region can have is that of a \\simil1500\\,K black body (at higher temperatures, the dust evaporates and the material loses its IR opacity). Thus the minimum solid angle of the region that needs to be (un-) covered is: $$ \\Omega = F_{\\nu}/B_{\\nu}(1500\\,\\rm{K}) \\approx 1.9\\times10^{-16} \\ {\\rm sr}, $$ where $F_{\\nu}$\\,=\\,3.3\\,Jy, and $B_{\\nu}$ denotes the Planck function (evaluated at 12.8\\,\\mum). At the system distance of 148\\,pc, and approximating the emitting region by a circle\\footnote{The argument is qualitatively independent of the exact shape of the emitting region and we adopt the simplest possible geometry.}, this yields a diameter of \\simil0.48\\,AU. The minimum required speed for the absorbing screen to (un-) cover this region within the available four days then becomes \\simil210\\,km\\,s$^{-1}$.  This velocity is much higher than the velocities one may expect in the close environment of T\\,Tau\\,Sa. The Kepler speed at 0.24\\,AU, the absolute minimum distance from the star at which the screen may be positioned in order to cover a region of 0.48\\,AU in diameter, is \\simil90\\,km\\,s$^{-1}$, assuming a mass of 2.2\\,\\Msun \\ for Sa \\citep{2008A&A...482..929K,2008JPhCS.131a2028K}. This velocity corresponds to a circular orbit, screens orbiting on eccentric orbits may reach velocities of up to \\simil130\\,km\\,s$^{-1}$ at this distance. If the absorbing screen would be on an eccentric orbit, though, it would itself become warmer and brighter at 12.8\\,\\mum \\ on closest approach, thus at least partially compensating for its own dimming effect. Thus, extinction caused by a \"structure\" existing within the disk, e.g. a warp or spiral arm, cannot be responsible for the observed IR variability. Moreover, the velocity dispersion within a molecular cloud is only a few km\\,s$^{-1}$ \\citep[e.g.][]{1981MNRAS.194..809L}, which excludes dust clouds in relatively close vicinity of the star, but not directly related to it, as viable absorbing screens.  In the above analysis we have made several simplifications, which were all chosen to \\emph{lower} the required velocities for the observing screen, i.e. to favor the variable extinction scenario. The goal was to show that variable extinction can be ruled out, even with these unrealistically conservative assumptions:  \\begin{itemize} \\item[$\\bullet$] The required velocity of the obscuring screen that we derived is the absolute minimum value that does not violate fundamental laws of physics (Planck's or Kepler's law), under the assumption that the observed IR continuum radiation is thermal dust emission. However, in reality, the region of the disk that emits the bulk of the flux at 12.8\\,\\mum \\ is likely much larger, leading to a correspondingly higher minimum required velocity. For a typical disk model around an object with approximately the stellar parameters of T\\,Tau\\,Sa, the \"size\" of the disk (here defined as the region within which 75\\% of the flux at 12.8\\,\\mum \\ is emitted) is about 15\\,AU in radius \\citep[e.g.][]{2005A&A...441..563V}. In the special case of T\\,Tau\\,Sa, where the disk is likely truncated on the outside at a radius of 3$-$5\\,AU due to tidal interactions with T\\,Tau\\,Sb, the region we see at 12.8\\,$\\mu$m probably encompasses the entire disk. Conservatively assuming an outer radius of 3\\,AU for the disk, an absorbing screen would need to travel at \\simil700\\,km\\,s$^{-1}$ in order to (un-) cover 26\\% of the emitting region within four days, whereas the Kepler velocity at this radius is only \\simil26\\,km\\,s$^{-1}$.  \\item[$\\bullet$] We assumed the absorbing screen to be 100\\% opaque at 12.8\\,\\mum. If the screen's transparency at this wavelength is considerable, a larger area needs to be (un-) covered within the available time and the required velocity becomes correspondingly higher. \\end{itemize} We conclude that variable extinction \\emph{cannot} cause the short-term brightness fluctuations observed at 12.8\\,\\mum. The absolute minimum velocity required for an absorbing screen to (un-) cover the emitting region within the available time is about a factor of 2 higher than the highest possible speed such a screen could have at the relevant location. A more realistic estimate yields a discrepancy of a factor of \\simil20-30 between the required and the physically possible velocity.  Note that if we apply the same simple model to the fastest variations detected by \\cite{2004ApJ...614..235B} \\emph{at 2.2\\,\\mum} \\ (\\simil0.9\\,Jy in seven days)\\footnote{The brightness of T\\,Tau\\,S increases from K=7.4\\,mag on 1997 December 6 to K=6.5\\,mag seven days later \\citep{2004ApJ...614..235B}. Assuming Sb was constant at K=8.6\\,mag, this implies that Sa went from K=7.84\\,mag to K=6.67\\,mag, i.e. it brightened by \\simil1.2\\,mag, or \\simil0.9\\,Jy.}, we cannot exclude variable extinction: a 1500\\,K blackbody emitting the corresponding flux would have a diameter of \\simil0.15\\,AU and the minimum speed needed for an opaque screen to (un-) cover this region within 7 days is only \\simil37\\,\\kms. This is substantially below the corresponding Kepler speed of \\simil160\\,\\kms, and thus variable extinction is a perfectly viable explanation for the observed variations. Considering that the stellar photosphere may contribute significantly to the total flux at 2.2\\,\\mum, the case for variable extinction becomes even stronger. \\emph{Only the long wavelength of our VISIR observations allows rejection of the variable extinction hypothesis}.   \\begin{figure}[t] \\hspace{-0.7cm} \\includegraphics[width=10.0cm,angle=0]{TTau_Sa_KminLp_vs_K.pdf} \\caption{\\label{fig:KminLp} K-\\Lp \\ vs. K color-magnitude diagram similar to that presented by \\cite{2004ApJ...614..235B}, except that we plot only Sa (\\cite{2004ApJ...614..235B} plot the sum of Sa and Sb). The diamonds represent \\emph{indirect measurements} calculated from the data by \\cite{2004ApJ...614..235B}, in which Sa-Sb remained spatially unresolved, assuming that Sb was of constant brightness at K=8.6\\,mag and \\Lp=6.25\\,mag, respectively. The diamonds in the upper left part of the diagram, in a darker shade, are less reliable because Sa was comparatively faint during the epochs of these measurements. The cross and triangle are from \\cite{2005ApJ...628..832D} and \\cite{2007AJ....134..359H}, respectively, whose measurements \\emph{did} spatially resolve the Sa-Sb pair. Over-plotted are radiative transfer model magnitudes, labeled with their respective accretion rates in \\msunyr, see \\sek\\,\\ref{sec:radiativeTransfer} and \\tab\\,\\ref{tab:model_parameters}.} \\end{figure}  \\begin{figure}[t] \\hspace{-0.7cm} \\includegraphics[width=10.0cm,angle=0]{TTau_Sa_KminN_vs_K.pdf} \\caption{\\label{fig:KminN} K-12.8 vs. K color-magnitude diagram of T\\,Tau\\,Sa. The diamonds represent measurements in which the Sa-Sb pair was spatially resolved in the K-band, crosses denote measurements in which we calculate the K-band magnitude of Sa from the total flux of T\\,Tau\\,S assuming Sb has K=8.6. The N-S pair was spatially resolved at 12.8\\,\\mum, except for the two points encircled in grey, for which the 12.8\\,\\mum \\ flux was calculated from the total flux of T\\,Tau assuming T\\,Tau\\,N has a flux of 8.2\\,Jy. The point surrounded by a grey box represents a broad band measurement centered at 12.5\\,\\mum \\ by \\cite{1991AJ....102.2066G}. The uncertainties arise mostly because the near- and mid-infrared observations were not performed simultaneously (see \\sek~\\ref{sec:color_magnitude_behavior}). Over-plotted are radiative transfer model magnitudes, labeled with their respective accretion rates in \\msunyr, see \\sek\\,\\ref{sec:radiativeTransfer} and \\tab\\,\\ref{tab:model_parameters}.} \\end{figure}   \\begin{table}[b] \\caption{\\label{tab:model_parameters} Summary of the parameters of the radiative transfer model described in \\sek\\,\\ref{sec:radiativeTransfer} and \\figs~\\ref{fig:KminLp}, \\ref{fig:KminN}, and \\ref{fig:Sa_SEDs}.}  \\begin{tabular}{lccc}  parameter      &  symbol        &  value                           & unit    \\\\ \\hline  stellar radius         & $R_*$  &  3.5  &  \\Rsun \\\\ stellar mass           & $M_*$  &  2.2  &  \\Msun \\\\ effective temperature  & $T_{\\rm{eff}}$  &  5500 &  K  \\\\ stellar luminosity     & $L_*$  &  10   & \\Lsun \\\\  \\vspace{-0.15cm} \\\\  disk mass      & $M_{\\rm{disk}}$ &  10$^{-3}$ & \\Msun \\\\ exponent surface density  & $p$          &  -1  &   \\\\ inner radius passive disk & $R_{\\rm{in}}$ &  0.5 & AU \\\\ outer radius passive disk & $R_{\\rm{out}}$ &  5.0 & AU \\\\ scale height at inner edge & $H_0$    &  0.021 & AU \\\\ flaring index              & $\\gamma$ &  2/7 &    \\\\  \\vspace{-0.15cm} \\\\  accretion rate &       \\Mdotacc &  10$^{-8}$, 10$^{-6}$, 10$^{-5}$ & \\msunyr \\\\ inner radius active disk  &  $R_{\\rm{in, SS}}$ &  0.036 & AU \\\\ outer radius variable \\Mdotacc  &  $R_{\\rm{out, acc}}$ &  0.5 & AU \\\\  \\vspace{-0.15cm} \\\\  inclination               & $i$       &  74     &  deg \\\\ foreground extinction     & $A_{\\rm{V}}$ & 15      &  mag \\\\  \\hline \\end{tabular}  \\end{table}   \\subsection{Can the observed long-term brightness fluctuations be caused by variable accretion?} \\label{sec:radiativeTransfer}  In the previous section we showed that apparent brightness fluctuations due to variable extinction cannot be the cause of the observed short-term mid-IR photometric variations of T\\,Tau\\,Sa. Therefore the variability must be due to intrinsic changes in the luminosity of Sa. A high and strongly variable accretion rate would naturally lead to such luminosity changes in a young star like Sa, as already argued by \\cite{1991AJ....102.2066G}. Thus, variable accretion is the prime candidate mechanism to explain the observed short-term variability. One may ask whether it can explain the observed long-term photometric variations, which are of much larger amplitude, as well. The variable accretion mechanism has previously been deemed unfit to explain the ``bluer when brighter'' color-magnitude behavior of the long-term brightness fluctuations \\citep{2004ApJ...614..235B}. In this section we will show that constructing a disk model that qualitatively reproduces the observed color-magnitude behavior \\emph{is} in fact possible, with parameters that remain in a realistic range.  Our model of the Sa disk consists of two components: an ``active'' inner disk in which luminosity is generated by the release of gravitational energy from accreting material, and a ``passive'' outer disk that is irradiated by the central star and the inner disk, and reprocesses the absorbed optical/NIR photons into longer wavelength infrared radiation. The model parameters are briefly discussed here and are also listed in \\tab~\\ref{tab:model_parameters}. They are chosen so that they represent Sa and its disk to the best of current knowledge. In particular the small outer radius of the disk and the near edge-on orientation are key features \\citep[e.g.][]{2005ApJ...628..832D}.  The active, inner part is modeled with an optically thick, geometrically thin accretion disk with a radial temperature profile following the \\cite{1973A&A....24..337S} model. In addition, a ``hot-spot'' covering 1\\% of the star, representing the accretion shock of infalling material hitting the stellar surface, is included by adding a blackbody of the corresponding temperature and luminosity to the stellar spectrum. For the passive, outer disk we used the 2D radiative transfer code {\\sc radmc} \\citep[][]{2004A&A...417..159D} to calculate the temperature structure and emergent infrared SED. The passive disk part has an inner radius of $R_{\\rm{in}}$\\,$=$\\,0.5\\,AU, an outer radius of $R_{\\rm{out}}$\\,$=$\\,5\\,AU, a mass of $M_{\\rm{disk}}$\\,$=$\\,$10^{-3}$\\,\\Msun \\ (assuming a gas to dust ratio of 100) with a radial surface density profile of $\\Sigma$\\,$\\propto$\\,$R^{-1}$, a flared geometry with a flaring index\\footnote{The flaring index $\\gamma$ is defined so that the ratio of the disk scale height $H$ and the distance to the central star $R$ scales like $H/R$\\,$\\propto$\\,$R^{\\gamma}$.} of $\\gamma$\\,$=$\\,2/7 and a scale height of $H_0$\\,$=$\\,0.021\\,AU at the inner disk edge of the passive, re-processing disk ($H/R$\\,$=$\\,0.042, chosen to be the value corresponding to hydrostatic equilibrium in the \"high\" state). The outer radius is much smaller than that of typical circumstellar disks because of tidal truncation by Sb, whose highly eccentric orbit brings it to within \\simil12\\,AU from Sa at closest approach. This naturally explains the non-detection of the Sa disk at mm~wavelengths by \\cite{1997ApJ...490L..99H} and \\cite{1998ApJ...505..358A}, as argued before by \\cite{2000ApJ...531L.147K}.  By varying the accretion rate, we change the amount of energy that is released in the center of the system, heating the circumstellar material. The luminosity of the stellar photosphere dominates at low accretion rates, providing a \"base level\" for the energy output. At accretion rates above \\simil7$\\times$10$^{-7}$\\,\\msunyr \\ accretion becomes the main energy source and the heating rate of the circumstellar material scales roughly linearly with the accretion rate. Variations in the central luminosity lead to essentially instantaneous variations in the temperature of the circumstellar material and to corresponding changes in the infrared emission because the dust grains are small and attain their new equilibrium temperature practically instantly if the radiation field changes, and the light travel time to the outer disk edge is only about 30 minutes.  For the luminosity of the central star itself, we adopt a value of 10\\,\\Lsun. This is lower by a factor of a few compared to the estimate by \\cite{2005ApJ...628..832D}, who base their value on a mass estimate of 2.5-3\\,\\Msun \\ for Sa. However, dynamical modeling of the Sa-Sb orbit including more recent data yields a mass estimate of $\\sim$2.2\\,\\Msun \\ for Sa \\citep{2008A&A...482..929K,2008JPhCS.131a2028K}, which is essentially identical to the mass of T\\,Tau\\,N \\citep[2.1\\,\\Msun,][]{2001ApJ...556..265W}. The bolometric luminosity of T\\,Tau\\,N is relatively well determined to be 7.3$^{+1.3}_{-1.1}$\\,\\Lsun \\ \\citep{2001ApJ...556..265W}, and assuming the stars are co-eval therefore yields a luminosity of $\\lesssim$10\\,\\Lsun \\ for Sa. A lower stellar luminosity is favorable for a scenario in which the IR variability is explained by variable accretion: the required extra luminosity in ``high states'' is always relative to the stellar photosphere and thus directly scales with the stellar luminosity. The stellar parameters that we chose correspond to a 2\\,Myr-old star of 2.2\\,\\Msun \\ according to the pre-main sequence evolutionary tracks by \\cite{2008ApJS..178...89D}.   \\begin{figure}[t] \\hspace{-0.7cm} \\includegraphics[width=10.0cm,angle=0]{TTau_Sa_SEDs.pdf} \\caption{\\label{fig:Sa_SEDs} Observed infrared photometry of Sa, with radiative transfer model SEDs over-plotted. The photometry is the same as in \\figs~\\ref{fig:KminLp} and \\ref{fig:KminN}. The model SED curves are for accretion rates of 10$^{-5}$, 10$^{-6}$, and 10$^{-8}$\\,\\msunyr \\ (top to bottom black curve), with the other model parameters as given in \\sek~\\ref{sec:radiativeTransfer} and \\tab~\\ref{tab:model_parameters}. Also shown, with a grey curve, is the 10$^{-5}$\\,\\msunyr-model without foreground extinction.} \\end{figure}  We performed model calculations for mass accretion rates of 10$^{-8}$, 10$^{-6}$, and 10$^{-5}$~\\msunyr. For the lowest accretion rate, the accretion luminosity amounts to $\\lesssim$2\\% of the stellar luminosity and the resulting fluxes are essentially identical to those obtained with a zero accretion rate. We calculated models for a range of nearly edge-on inclinations. Because our objective is to demonstrate that a variable accretion model is in principle capable of reproducing the observed color-magnitude behavior, we searched for a model that gives a fair match to the observations but did not comprehensively explore parameter space. A model with parameters as listed earlier in this section and in \\tab~\\ref{tab:model_parameters}, seen at an inclination of 74\\degree \\ and with a foreground extinction of 15\\,mag agrees fairly well with the observations (the extinction matches the value of the foreground extinction observed toward Sb \\citep{2005ApJ...628..832D}). This particular inclination yields the best fit for the set of model parameters that we chose, but other choices of disk parameters will result in slightly different best-fit inclinations. Thus, our model does not accurately constrain the inclination of the Sa disk, but we stress that a near edge-on inclination is required to obtain the correct color-magnitude behavior of the variations, regardless of the specific choice of disk parameters.  In \\fig\\,\\ref{fig:Sa_SEDs} we show the model SEDs for accretion rates of 10$^{-8}$, 10$^{-6}$, and 10$^{-5}$~\\msunyr, as well as the observed infrared photometry. We also show the model for the highest accretion rate as it would appear if the foreground extinction is omitted (grey curve). In \\figs~\\ref{fig:KminLp} and \\ref{fig:KminN} we showed the observed color-magnitude behavior of the photometric variations of Sa. The model magnitudes are over-plotted and labeled with their respective accretion rates. The results of our modeling efforts can be summarized as follows:  \\begin{itemize} \\item[$\\bullet$] the near-infrared color and brightness variations are qualitatively reproduced: we obtain a \"redder when fainter\" behavior in a K-\\Lp \\ vs. K color-magnitude plot (see \\fig\\,\\ref{fig:KminLp}). The slope of the relation in the model is somewhat shallower (i.e. closer to grey) than observed. \\item[$\\bullet$] the observed K-[12.8] \\ vs. K color-magnitude behavior is qualitatively reproduced. The slope of the modeled relation fits the observations well. The considerable observed scatter can at least in part be attributed to the time lag between the near- and mid-infrared observations (see \\fig\\,\\ref{fig:KminN}). \\item[$\\bullet$] the disk model always produces a 10\\,\\mum \\ silicate in emission, even at 90\\degree \\ (edge-on) inclination. Additional foreground extinction is needed to explain the observed silicate absorption feature. \\end{itemize}  Why does our model reproduce the observed ``bluer when brighter'' behavior, whereas the previous model considered in the context of T\\,Tau\\,S \\citep[][]{1997ApJ...481..912C,2004ApJ...614..235B} did not? In the Calvet model the central star is not highly extincted. At low accretion rates, the very blue stellar photospheric emission dominates the near-infrared spectrum. As the accretion rate is increased, the comparatively red disk spectrum becomes more prominent and dominates the near-infrared emission above 10$^{-7}$\\,\\msunyr. The changing star-disk contrast causes a ``redder when brighter'' color-magnitude behavior with varying mass accretion rate. In contrast, in our highly reddened model of the Sa star and disk system, the photosphere never substantially contributes directly to the near-infrared flux, because it is obscured by the Sa disk. At low accretion rates, the emission is dominated by the far side of the inner edge of the passive disk part, irradiated by the central star. With increasing accretion rate, the inner edge of the passive disk becomes hotter and, in addition, the comparatively blue active Shakura-Sunyaev disk contributes stronger to the near-infrared spectrum. The combined effect is a ``bluer when brighter'' color-magnitude behavior.  The variable accretion mechanism has been proposed previously, based on the apparent grey behavior of the observed near- to mid-infrared flux variations \\citep{1991AJ....102.2066G,2002ApJ...568..771D}. However, based on a much larger data set we show that the large, long-term flux variations at 2.2 and 12.8\\,\\mum \\ are accompanied by substantial color variations (\\fig\\,\\ref{fig:KminN}), which can be reproduced by our variable accretion model. However, these data can also be explained with variable foreground extinction (see the reddening vectors drawn in \\figs~\\ref{fig:KminLp} and \\ref{fig:KminN}). The total amplitude of the foreground extinction required to cover the observed brightness range is $\\Delta$\\Av\\simil40\\,mag.  \\subsubsection{Variable extinction causing variable disk-illumination?} In the previous paragraphs we showed that variable extinction in its simplest form, i.e. absorbing material passing in front of a source of approximately constant intrinsic luminosity, cannot cause the observed fast mid-IR brightness variations. However, the color-magnitude behavior of the large, multi-magnitude variations that have been observed during the last two decades, \\emph{can} be explained with variable extinction (see \\sek~\\ref{sec:radiativeTransfer}). The variable accretion model is in principle capable of explaining both the detected sub-magnitude fast variations, as well as the presumably slower multi-magnitude variations. Here we will consider a conceptually different realization of the variable extinction scheme, in which we attempt to explain both the short- and long-term variations with extinction alone.  To this purpose, we will introduce ``central extinction events'': variable absorption columns close to the central star that cause substantial changes in the illumination of the disk surface at larger radii with optical and near-infrared radiation. This could in principle lead to fast variations at mid-IR wavelengths because the material causing the variable disk illumination resides much closer to the central star than the material emitting the bulk of the mid-IR radiation, thereby circumventing the dynamical timescale problem encountered in the ``classical'' variable extinction scenario.  In order to explore this scheme in a more quantitative sense, we constructed a passive disk model for Sa, in which we introduce variable absorption columns by increasing the scale height at the inner disk edge. The disk parameters\\footnote{$R_{\\rm{in}}$\\,$=$\\,0.5\\,AU, $R_{\\rm{out}}$\\,$=$\\,5\\,AU, $M_{\\rm{disk}}$\\,$=$\\,$10^{-3}$\\,\\Msun, $\\Sigma$\\,$\\propto$\\,$R^{-1}$, $\\gamma$\\,$=$\\,2/7, $H_0$\\,$=$\\,0.021\\,AU at 0.5\\,AU.} are the same as those of the model presented in \\sek\\,\\ref{sec:radiativeTransfer}, but as the central illuminating source we replaced the 10\\,\\Lsun \\ star plus variable accretion component with a star-only configuration with a constant luminosity of 40\\,\\Lsun \\ \\citep{2005ApJ...628..832D}.  We considered both a scenario in which the scale height at the inner disk edge is time variable, and one in which a fixed ``structure'' rotates at the inner disk edge with the Keplerian speed, casting a shadow over the outer disk that rotates along with the same angular velocity. For the former experiment, we artificially increased the scale height of the disk between $R_{\\rm{in}}$ and 2$R_{\\rm{in}}$, at all azimuthal angles. If we increase the scale height by 55\\%, with respect to the disk parameters given in the previous paragraph, the flux at 12.8\\,\\mum \\ decreases by 0.26\\,mag because the inner disk somewhat shadows the regions at larger radii. The \\simil0.26\\,mag brightening at 12.8\\,\\mum \\ that we observed in February 2008 may conversely be explained by the opposite change in inner rim structure: a reduction of the pressure scale height from \\simil0.034\\,AU to \\simil0.022\\,AU within $\\lesssim$4\\,days. The height above the midplane where the disk becomes optically thin to the stellar irradiation\\footnote{Here this is taken to be the height where the optical depth in the radial direction at $\\lambda$\\,$=$\\,0.55\\,\\mum \\ falls below unity.} then decreases from \\simil0.16\\,AU to \\simil0.11\\,AU, i.e. by more than twice the pressure scale height corresponding to hydrostatic equilibrium. The orbital period at 0.5\\,AU from a 2.2\\,\\Msun \\ star is \\simil87~days, and thus vertical movements on the order of one scale height require roughly 20 days, which is much longer than our observational constraint. Thus we conclude that changes in the inner disk structure cannot cause the observed fast mid-IR brightness variations.  A somewhat different approach is to take a fixed enhancement of the disk scale-height at the inner rim over only part of the azimuthal range, causing a ``shadow'' to be cast over the disk behind it. This shadow will move over the outer disk with an angular velocity corresponding to the Kepler speed at the inner edge, and thus cause a ``cool region'' at larger radii that moves at locally vastly super-keplerian velocities over the disk surface. For near edge-on inclinations, where the mid-IR flux we see is dominated by the far side of the disk, this can cause substantial dimming. We used the 3D version of {\\sc radmc} (Dullemond et al. in prep.) to calculate light curves of a system in which the scale height at the inner rim is increased with respect to hydrostatic equilibrium over a limited azimuthal range\\footnote{Over $\\pi$ radians in azimuth the scale height is increased with a sinusoidal shape, so that at the central azimuthal angle the scale height is twice the hydrostatic value, and at angles of $-\\pi$/2 and $+\\pi$/2 with respect to the maximum it reaches the hydrostatic value.}. With this we can achieve brightness variations of $\\lesssim$0.15\\,mag on timescales of \\simil20~days. Because the enhancement of the scale height is unrealistically large and yet the brightness variations are substantially smaller and slower than observed, we do not consider this scenario to be a viable explanation for our observations.  \\vspace{0.15cm}  We conclude that variable accretion is the prime mechanism responsible for the observed \\emph{short-term} mid-IR variability, and exclude a significant role for variable extinction in this matter, as explained in \\sek~\\ref{sec:variable_extinction}. The short-term near-IR photometric variations are naturally also explained by variable accretion, though the existing data do not strictly exclude variable extinction as their cause. Based on the hitherto available data, both variable accretion and variable extinction are viable mechanisms to explain the \\emph{long-term} photometric variations and the associated color-magnitude behavior. Data presented in past studies, in particular a time-variable optical depth in the 3.05\\,\\mum \\ ice absorption feature \\citep[][]{2001ApJ...551.1031B}, do suggest that at least part of the variability can be attributed to variable extinction.  \\subsubsection{The 10\\,\\mum \\ silicate feature in Sa}  The fact that T\\,Tau\\,S shows a deep silicate absorption feature in the 10\\,\\mum \\ spectral region \\citep{1991AJ....102.2066G}, which was later shown to arise in Sa only \\citep{2008ApJ...676.1082S,2009A&A...502..623R}, has been regarded as evidence that the 10\\,\\mum \\ emission is dominated by a disk that is seen nearly edge-on from our vantage point \\citep[e.g.][]{2005ApJ...628..832D}. Our radiative transfer model of the Sa disk, however, always shows a silicate feature in emission, even for an edge-on inclination (see the unreddened model curve in \\fig\\,\\ref{fig:Sa_SEDs}, grey line). This is due to the \"artificial\" truncation of the Sa disk by the tidal influence of Sb, with an outer radius of 5\\,AU in our model. In \"normal \" T\\,Tauri disks, which are much larger, cold material in the outer disk regions obscures the warm inner disk when viewed at near edge-on inclinations, causing a silicate feature in absorption. In the special case of Sa, the entire disk is warm enough to emit in the 10\\,\\mum \\ spectral region. The radiation that we receive is dominated by the optically thin upper disk layers, yielding a silicate feature in emission, regardless of the viewing angle. \\cite{2009A&A...502..623R} reach the same conclusion based on their  radiative transfer model for Sa using the MC3D code. Note that the shape and peak to continuum ratio of the silicate band depend on the grain size used. The small, sub-micron sized grains used in our model cause a relatively strong emission band peaking just short-ward of 10\\,\\mum. The use of larger grains results in a weaker, broader silicate feature, and one can make the silicate band essentially disappear from the spectrum by using very large ($\\gtrsim$10\\,\\mum) grains. However, one never gets a silicate feature in absorption. Thus, an additional, external absorption component needs to be invoked. This is consistent with the fact that Sb, whose disk is thought to be seen more or less face-on \\citep[e.g.][]{2005ApJ...628..832D}, suffers an extinction of \\Av\\simil15\\,mag. The shape of the observed silicate absorption feature implies that the material causing the absorption consists predominantly of small ($\\lesssim$1\\,\\mum) grains.    \\subsection{A tentative scenario for T\\,Tau\\,Sa} \\label{sec:TTauS_scenario} The light curves shown in \\fig\\,\\ref{fig:light_curve} can qualitatively be described as follows\\footnote{We will describe the K-band light curve because it has the best temporal sampling. The \\Lp \\ and 12.8\\,\\mum \\ light curves follow the K-band (see also Figs.\\,\\ref{fig:KminLp} and \\ref{fig:KminN}).}. From the beginning of the IR observations, in 1982, until 1989 T\\,Tau\\,S appears to have been in a ``low state''. Because there are only three measurements in this period, it is by no means certain that it has been faint during this whole period, but the available measurements all show T\\,Tau\\,S to be of essentially the same (low) brightness at 2.2\\,\\mum. Sometime in 1989, the brightness of the southern component started to increase strongly. Based on data from later epochs this increase, or at least the vast majority of it, is attributed to Sa only. During the period from \\simil1991 to \\simil1995 the temporal sampling was very poor, but the few available measurements suggest Sa to be relatively faint. From \\simil1996 to late 2001 Sa has been in a ``high and variable'' state. This period has been very well sampled by \\cite{2004ApJ...614..235B} and the K-band brightness varied between K\\simil7.5\\,mag and K\\simil6\\,mag. From late 2001 to late 2003 the brightness of Sa decreased to K\\simil10.5\\,mag. The temporal coverage of the observations is poor and thus we cannot assess how ``smooth'' this decrease was. Sometime between late 2004 and late 2006 the brightness of Sa rose again to K\\simil9\\,mag, and is currently varying around this value.  In view of the variable accretion scenario, the dramatic brightening events in the early and late 1990s and the interspersed episodes of irregular flux variations, of which we may be seeing the last twitch at the current time, could be interpreted as an ``outburst'' of accretion. In this respect, it is particularly interesting to compare the timing of the outburst with that of the periastron passage of the companion Sb. In the current best-fit solution of the Sa-Sb orbit, the closest approach occurred in June 1995, with an uncertainty of somewhat more than a year (1$\\sigma$), as we have indicated in \\fig\\,\\ref{fig:light_curve} with a grey vertical band. Thus, we have a period of enhanced activity around the time of closest approach of the companion, with the middle of the active period timed somewhat after periastron passage, and the strongest activity occurring in the second half of the active period. We will now argue that this behavior provides a good qualitative match to a scenario in which an accretion outburst in the center of the Sa disk is induced by the gravitational influence of Sb, tidally disturbing the Sa disk during periastron passage.  When a young star and disk system has a low-mass companion, the latter will tidally truncate the disk around the primary to $\\sim$1/3 times the semi-major axis of the orbit \\citep[e.g.][]{1993ApJ...419..166A}. If the companion is on an eccentric orbit, its gravitational influence on the primary disk will be strongly time-dependent, being relatively weak during most of the orbit but culminating near periastron passage. \\cite{2008A&A...486..617K} modeled the close binary system $\\gamma$~Cephei in its youth, which is thought to have been very similar to the Sa-Sb system today. It has primary and secondary masses of \\simil1.59\\,\\Msun \\ and \\simil0.38\\,\\Msun, respectively, a semi-major axis of \\simil18\\,AU, and an eccentricity of \\simil0.36 \\citep{2003ApJ...599.1383H}. In addition, the primary has a planetary companion with a mass of $M$\\,sin\\,$i$\\,\\simil\\,1.7\\,M$_{\\rm{Jup}}$ \\citep{2003ApJ...599.1383H}, which is, however, not of central importance to the current discussion. \\cite{2008A&A...486..617K} performed hydro-dynamical simulations of the disk around the young $\\gamma$\\,Cep primary and considered the tidal effects of the low-mass secondary. The simulations reveal a nearly undisturbed disk at most orbital phases, when the companion is at relatively large distances. As the companion approaches the primary, however, two strong spiral shock waves are excited which propagate all the way to the inner disk (see \\fig\\,\\ref{fig:kley_simulation}). This takes some time, and therefore the gravitational disturbance of the inner disk is not symmetrical in time with respect to the periastron passage, but lags somewhat. The strongest disturbances at the inner disk edge occur substantially after closest approach.  We speculate that such shock waves induce enhanced accretion in the innermost disk region of Sa when Sb is near periastron. For example, the mechanical energy dissipated in the shocks may cause a normally essentially neutral disk medium to be sufficiently ionized for MRI to become effective, naturally leading to strongly enhanced accretion. We have re-run the Kley \\& Nelson model including viscous heating and radiative cooling for the parameters of the Sa-Sb system, and find the same behavior as seen in the $\\gamma$\\,Cep calculations\\footnote{Note that the model discussed here and shown in \\fig\\,\\ref{fig:kley_simulation} is a hydrodynamical calculation and is independent of the radiative transfer model discussed in \\sek\\,\\ref{sec:radiativeTransfer}.}. The simulations provide a good qualitative match to the observations: a period of enhanced accretion activity near every periastron passage, whose peak has a substantial time lag with respect to the closest approach of the secondary. Quantitatively, the model shows much more modest variations in the accretion rate than inferred by the observations, and the time between periastron passage and maximum accretion activity is longer in the simulations by a factor of $\\sim$3. The Kley \\& Nelson model focuses on the dynamical evolution and mass loss of the disk at its outer edge, but does not attempt to treat the accretion physics at the inner disk edge in detail. A model with a more detailed treatment of the disk viscosity and accretion physics in the innermost disk regions is needed, which allows the spiral shock waves to propagate inward more quickly and can account for the required high accretion rates. This is beyond the scope of the current work, and whether our qualitative model can be refined to provide a quantitative match to the observations once all the relevant physics are included remains to be confirmed.  Here we pose the companion-induced accretion outburst as an idea to explain the observed behavior of T\\,Tau\\,S during the past decades, and provide a clear prediction that can be easily verified observationally: if this scenario is correct, Sa should return to its ``low state'' some time in the near future, and stay there until the next periastron passage of Sb. Exactly how long it will take for Sa to return to ``quiescence'' in detail depends on the actual accretion physics and inner disk parameters, neither of which we know. Thus, while it is technically trivial to test the proposed scenario with currently available observing facilities, it may in practice take a substantial number of years to convincingly falsify our scheme. Supportive evidence could be built up over the years if Sa is observed to remain faint, the final prove may come only many years from now if the next outburst coincides with the next periastron passage. The current astrometric data still allow a wide range of orbital periods, from \\simil25 to \\simil100 years, but as new observations accumulate over the coming years, we should be able to fairly accurately predict the onset of the next outburst long before it would occur.   \\begin{figure}[t] \\includegraphics[width=9.7cm,angle=0]{temp2d-140c.jpg} \\caption{\\label{fig:kley_simulation} Two-dimensional temperature distribution of the Sa disk 10.8 yrs after the periastron passage of Sb, for a hydrodynamic simulation using the physical parameters of the T\\,Tau\\,S system. X and Y are spatial coordinates, and the disk is viewed face-on in this representation. To better bring out the entire range of temperatures we show $T[K]^{1/4}$, the colors cover the temperature range of 10 to 2000\\,K.} \\end{figure}   \\subsection{Issues} Here we briefly discuss some issues that are relevant to the current discussion.  \\subsubsection{The mass and lifetime of the Sa disk} \\label{sec:Sa_disk_mass} The amount of material that we can reasonable assume to be present in the Sa disk is at most a few times 10$^{-3}$\\,\\Msun. If the average accretion rate during an outburst is 5\\,$\\times$\\,10$^{-6}$\\,\\msunyr \\ and an outburst lasts for 10 years, we can have only on the order of 10$^{2}$ these events before the disk is all gone, unless the Sa disk can be supplied with fresh material. If the orbital period of Sa-Sb is 50\\,yrs, this means the current situation can exist for only $\\lesssim$5000\\,yrs. Possibly the Sa-Sb pair became a close binary only within the last several thousand years due to gravitational interaction with a third object. It may also be possible that the Sa disk is still supplied with fresh material coming in at high latitudes with respect to the orbital plane as the Sa-Sb pair moves through ambient cloud material. Because the disks of Sa and Sb appear strongly misaligned \\citep{2005ApJ...628..832D,2009A&A...502..623R}, the Sa disk and the Sa-Sb orbit are misaligned, and cloud material is present close to the Sa-Sb system \\citep[\\sek\\,\\ref{sec:naco_results},][]{2007AJ....134..359H,2008ApJ...676..472B,2008A&A...488..235G}, the notion that cloud material is still being freshly supplied to the Sa disk is realistic.   \\subsubsection{A note on the Sb disk} \\label{sec:Sb_disk} Like the primary star Sa, the low-mass companion Sb also shows a substantial infrared excess, indicating that some circumstellar material is still present, presumably also in a small circumstellar disk. As our calculations show, Sb will strongly disturb the Sa disk during periastron passage. Conversely, the gravitational pull of Sa will affect the disk around Sb even more strongly, and accretion onto Sb should also be enhanced during periastron passage. Therefore one may wonder why the observed brightness variations are much smaller in Sb (maximum amplitude \\simil0.4\\,mag $RMS$$=$0.2\\,mag) than in Sa (amplitude $\\gtrsim$3\\,mag).  \\cite{2002ApJ...568..771D,2005ApJ...628..832D} measured the Br\\,$\\gamma$ line flux of Sb to be 7.2$\\times$10$^{-17}$\\,Wm$^{-2}$ on 2000 November 19, and 9.4$\\times$10$^{-17}$\\,Wm$^{-2}$ on 2003 December 12, with corresponding K-band continuum magnitudes of 8.77 and 8.50, respectively. the spectral type of Sb corresponds to $T_{\\rm{eff}}$\\simil3800\\,K \\citep{2005ApJ...628..832D}, and assuming an age of 2\\,Myr this yields a mass of approximately 0.6\\,\\Msun \\ and a luminosity of \\simil0.45\\,\\Lsun \\ by comparison to the pre-main sequence evolutionary tracks of \\cite{2008ApJS..178...89D}, which seem reasonable. The corresponding stellar radius of Sb is $R_*$$=$1.55\\,\\Rsun. Assuming a foreground extinction of $A_{\\rm{V}}$$=$15\\,mag, a standard extinction law ($A_{\\rm{K}}$=0.112\\,$A_{\\rm{V}}$), the relation between Br\\,$\\gamma$ luminosity and accretion luminosity given in equation~3 of \\cite{2004AJ....128.1294C}, and using the approximate relation $L_{\\rm{acc}}$$=$G$M_*$$\\dot M_{\\rm{acc}}$/$R_*$, one derives mass accretion rates of 3.5$\\times$10$^{-8}$\\,\\msunyr \\ and 4.4$\\times$10$^{-8}$\\,\\msunyr \\ for the respective epochs. The corresponding accretion luminosities are approximately equal to the stellar luminosity. Therefore, variable accretion is a viable mechanism for the observed modest near-infrared variations in Sb as well. How variations in the accretion luminosity translate to variations in the near-infrared continuum excess emission is somewhat dependent on the disk structure: the extra luminosity is released in the UV which cannot be observed due to the high extinction, and we only see the reprocessed radiation from the hot inner disk regions.  We note, however, that the existing data also perfectly agree with a time variable extinction towards Sb. Between the aforementioned epochs of Br\\,$\\gamma$ spectroscopy by Duch\\^ene et al. the extinction would have decreased by $\\Delta$$A_{\\rm{V}}$\\simil2.4\\,mag, assuming a standard extinction law.  \\subsubsection{Mid-infrared brightness during the peak of the outburst} We are not aware of mid-IR measurements performe between roughly 1998 January and 2001 November, but we argue Sa must have been very bright during that period, peaking at \\simil40\\,Jy at 12.8\\,\\mum. Should any mid-IR observations have been performed during this period, they are highly relevant to the current discussion. The variable extinction scenario predicts a lower peak-brightness in mid-IR than the variable accretion scenario. However, only if mid-IR observations have been performed (near-) simultaneously with the existing K-band data they may provide supportive evidenced for either scenario. In case of a substantial time lag the potential short-term variations will inhibit drawing firm conclusions.    The southern infrared companion of T\\,Tau is a close binary system of which the more massive member, T\\,Tau\\,Sa, shows strong long-term photometric variability, with an amplitude of $\\gtrsim$3\\,mag at 2.2\\,\\mum \\ and \\simil2\\,mag in the 10\\,\\mum \\ region on timescales of years. Short-term variability of \\simil1\\,mag within one week has been detected at 2.2\\,\\mum. The physical mechanism driving the observed flux variations has been under debate, the originally proposed intrinsic luminosity variations due to a variable accretion rate have later been challenged in favor of apparent brightness fluctuations due to variable foreground extinction.  We revisited the nature of the photometric variability of T\\,Tau\\,Sa using new observations in the mid-infrared. In February 2008 we detected a very fast increase of the mid-IR brightness of T\\,Tau\\,S which can be attributed to Sa: $+$26$\\pm$2\\% at 12.8\\,$\\mu$m in 3.94 days. Using simple geometric arguments and basic physic laws we can exclude that this \\emph{short-term} variability is due to time-variable extinction. We show that the fast brightness fluctuations must instead reflect changes in intrinsic luminosity, and attribute these to variable accretion. The key aspect of our argument is the short timescale of the variations in combination with the long observing wavelength.  We show that variable accretion can plausibly explain also the much larger \\emph{long-term} variability and its associated near-infrared ``bluer when brighter'' color-magnitude behavior, by creating a 2D radiative transfer disk model of the Sa disk that qualitatively reproduces the observed properties. However, the hitherto observed long-term variability can also be explained within the variable foreground-extinction scenario. A combination of both mechanisms, variable accretion and variable foreground extinction, may be required to explain the collective photometric variability of T\\,Tau\\,Sa.  In view of the tentative scenario that both the short- and long-term variability of Sa are caused by variable accretion, we reviewed the infrared light curve in relation to the parameters of the Sa-Sb orbit. We propose that the dramatic brightening events in the early and late 1990s and the interspersed episodes of irregular flux variations were an accretion outburst, induced by gravitational perturbation of the Sa disk during the periastron passage of Sb, with closest approach in \\simil1995. The induced maximum accretion rate is on the order of 10$^{-5}$\\,\\msunyr. While we do not model the physics of this outburst in detail, we make a clear observational prediction: if this scheme is correct, Sa should return to ``quiescence'' some time in the near future, and remain in a low state until the next periastron passage of Sb."
    },
    "1004/1004.1507_arXiv.txt": {
        "abstract": "We have derived disk scale lengths for 30374 non-interacting disk galaxies in all five SDSS bands. Virtual Observatory methods and tools were used to define, retrieve, and analyse the images for this unprecedentedly large sample classified as disk/spiral galaxies in the LEDA catalogue. Cross correlation of the SDSS sample with the LEDA catalogue allowed us to investigate the variation of the scale lengths for different types of disk/spiral galaxies. We further investigat asymmetry, concentration, and central velocity dispersion as indicators of morphological type, and are able to assess how the scale length varies with respect to galaxy type. We note however, that the concentration and asymmetry parameters have to be used with caution when investigating type dependence of structural parameters in galaxies. Here, we present the scale length derivation method and numerous tests that we have carried out to investigate the reliability of our results. The average $r$-band disk scale length is $3.79$ kpc, with an RMS dispersion of $2.05$ kpc, and this is a typical value irrespective of passband and galaxy morphology, concentration, and asymmetry. The derived scale lengths presented here are representative for a typical galaxy mass of $10^{10.8\\pm 0.54} \\rm{~M}_\\odot$, and the RMS dispersion is larger for more massive galaxies. Separating the derived scale lengths for different galaxy masses, the $r$-band scale length is $1.52 \\pm 0.65$ kpc for galaxies with total stellar mass $10^{9}$--$10^{10}\\rm{~M}_\\odot$ and $5.73 \\pm 1.94$ kpc for galaxies with total stellar mass between $10^{11}$ and $10^{12}\\rm{~M}_\\odot$. Distributions and typical trends of scale lengths have also been derived in all the other SDSS bands with linear relations that indicate the relation that connect scale lengths in one passband to another. Such transformations could be used to test the results of forthcoming cosmological simulations of galaxy formation and evolution of the Hubble sequence. ",
        "introduction": "\\label{sec:intro}  The exponential scale length of a galaxy disk is one of the most fundamental parameters to determine its morphological structure as well as to model its dynamics, and the fact that the light distributions are exponential makes it possible to constrain the formation mechanisms \\citep{Freeman70}. The scale length determines how the stars are distributed throughout a disk, and can be used to derive its mass distribution, assuming a specific M/L ratio. Ultimately, this mass distribution is the primary constraint for determining the formation scenario \\citep[e.g.,][ and referenced therein]{LP87,Dutton09}, which dictates the galaxy's evolution. As the galaxy evolves, substructures such as bulges, pseudo-bulges, bars, rings, and spiral arms may build up, and this will then considerably change the morphology of the host disks \\citep{CE93,Eetal05,Bournaud07}. The scale length value is intimately connected to the circular velocity of the galaxy halo, which in turn relates closely to the angular momentum of the halo in which the disk is formed \\citep{DSS97, Mo98}. Up to the last few years, cosmological simulations were limited to rather low resolution, were disks and spheroids were barely resolved, and generally limited to high redshifts, so reproducing realistic disk scale lengths for modern galaxies was clearly out of reach. The current simulations reach resolutions that allow resolving the disks from high redshift down to redshift zero, and subtle mechanisms changing the disk masses and scale lengths can be studied \\citep[e.g., ][ listed in alphabetic order]{Ceverinoetal10,Governatoetal10,MartigBournaud10,Schayeetal10}, thus calling for a comprehensive observational determination of these parameters to test the state of the art cosmological simulations.   Previous observations of NGC, UGC, and low surface brightness galaxies have shown that scale lengths span over a range of three orders of magnitudes \\citep[e.g., ][]{Boroson81, Romanishin83, vdKruit87, Schombertetal92, Knezek93, deJong96}. Any physical galaxy formation scenario should be able to explain this wide range of values while simultaneously explaining the similarities among disk galaxies throughout this range.  Analytic disk formation scenarios predict that, in cases where angular momentum is conserved, the disk scale length is determined by the angular momentum profile of the initial cloud  \\citep{LP87}, and the scale length in a viscous disk is set by the interplay between star formation and dynamical friction \\citep{Silk01}. These processes form the basis of a galaxy's gravitational potential, and determine the strength of gravitational perturbations, the location of resonances in the disk, the formation and evolution of spiral arms and bars, kinematically decoupled components in centres of galaxies, and the dynamical feeding of circumnuclear starbursts and nuclear activity \\citep[e.g.,][]{Eetal96, Knapen04, Fathi04, KK04}.  Photometrically, one generally derived this scale length by azimuthally averaging profiles of the surface brightness which is in turn decomposed into a central bulge and an exponential disk, and when spatial resolution allows other components such as one or several bars and rings can be taken into account.  \\begin{figure} \\begin{center} \\includegraphics[width=.49\\textwidth]{radec_densityplot.jpg} \\caption{Top: right ascension and declination distribution of the galaxies for which the LEDA services provide a Hubble classification number $\\ge 1$ (irrespective of classification error). Middle: same  for the 56096 SDSS galaxies fulfilling the first sorting criteria described in section~\\ref{sec:sample}, and Bottom: the final sample of 30374 selected SDSS disk galaxies for which we have reliable scale $r$-band scale lengths (regardless of uncertainty in the morphological classification, see also section~\\ref{sec:reliable}).} \\label{fig:radec} \\end{center} \\end{figure} As images in different bands probe different optical depths and/or stellar populations, it is likely that a derived scale length value should depend on waveband, and these effects may vary as a function of galaxy type where different amounts of dust and star formation are expected. Dusty disks are more opaque, resulting in larger scale length values in bluer bands when compared with red and/or infrared images. Similar effects can also be caused by differences in the stellar populations. Differences in scale length as a function of passband can therefore be used to derive information about the stellar structure and contents of galactic disks. Both the effects of stellar populations and dust extinction have been subject to much discussion over the years \\citep[e.g.,][]{SdeV83, Kent85, Val90, vDetal95, Petal94, Petal95, Betal96, Courteau96, Cunow98, BBA98, Petal01, Graham01, GdeB01, Cunow01, GH02, MacA03, Cunow04, GW08}. A detailed and extensive analysis of the dust effects has also been presented for a few tens of galaxies in \\citet{Holwerda05} and subsequent papers by this author, however, as noted by \\citet{Petal94} and \\citet{vDetal95}, the scale length alone in different wavelengths in small sample cannot be used to break the age/metallicity and dust degeneracies. Investigating the scale length variation as a function of inclination for large numbers of galaxies is necessary to distinguish between the effects of dust and stellar populations.  The common denominator in all the previous studies is the roughly comparable sample sizes (at most few hundred galaxies). Most studies have so far analysed individual galaxies, or samples containing a few tens, and in unique cases a few hundred \\citep[e.g.,][]{KvK91, Courteau96}, galaxies. Although a number of great results from studies with the Sloan Digital Sky Survey \\citep[SDSS; ][]{Yorketal00} in the last years have appeared, these works have not studied the astrophysical parameters targeted here.   As a part of a European Virtual Observatory\\footnote{http://www.euro-vo.org} Astronomical Infrastructure for Data Access (Euro-VO AIDA) research initiative, we have undertaken a comprehensive analysis of the scale length in disk galaxies using an unprecedentedly large sample of disk galaxies. We have used the Virtual Observatory (VO) tools to retrieve data in all ($u$, $g$, $r$, $i$, and $z$) bands from the sixth SDSS major data release \\citep[DR6; ][]{AMcetal08} which includes imaging catalogues, spectra, and redshifts freely available. We use the LEDA\\footnote{http://leda.univ-lyon1.fr} catalogue \\citep{Petal03} to retrieve morphological classification information about our sample galaxies, and those with types defined as Sa or later are hereafter refereed to as disk galaxies (distribution of both samples are presented in Fig.~\\ref{fig:radec}).   In the present paper, we present the data retrieval and analysis method used to automatically derive the scale lengths for a sample of disk galaxies which contains 56096 objects (described in section~\\ref{sec:sdsssample}), and after rigourous tests described in section~\\ref{sec:scalelengths}, we find that a subset of 30374 of these can be called reliable following these criteria. The scale lengths presented here relate only to the disk components, and we have tried to avoid the regions that could be dominated by the bulge component, in order to avoid complications related to the uncertainties of bulge-disk decomposition procedure \\citep[as demonstrated in, e.g., ][]{KvK91}. In section~\\ref{sec:results} we present the first results based on our unprecedentedly large sample of galaxies and finally discuss their implications in section~\\ref{sec:discussion}.  \\begin{figure*} \\begin{center} \\includegraphics[width=.99\\textwidth]{distplots.jpg} \\caption{Distribution of some key parameters retrieved from the SDSS database and morphological types from LEDA (bottom right). The filled histograms are the 56096 disk galaxies described in section~\\ref{sec:sample}, and the open histograms show the distributions of the final 30374 galaxies for which we derive reliable scale lengths. The distribution of the sample remains unchanged.} \\label{fig:design} \\end{center} \\end{figure*} ",
        "conclusions": "\\label{sec:discussion} We have derived reliable scale lengths for 30374 disk/spiral galaxies, with no sign of ongoing interaction or disturbed morphology, in all five $u,g,r,i,$ and $z$-bands from SDSS DR6 images. Cross-correlation of the SDSS sample with the LEDA catalogue has enabled us to investigate the variation of the scale lengths for different types of disk/spiral galaxies. Although the typical scale length in $u$-band is 35\\% larger than that in the $r$-band, the scale lengths in the $g,r,i,$ and $z$-bands are similar and only become smaller on the average for late morphological types. This result remains consistent when using by-eye morphological classification or when using asymmetry parameter, concentration parameter, or velocity dispersion as an indicator for galaxy morphological type. Our sample spans a range of total stellar masses between $10^{6.6}$ and $10^{12.2} \\rm{~M}_\\odot$ with a typical galaxy mass of $10^{10.8\\pm0.54} \\rm{~M}_\\odot$, and shows that the while scale length increases for more massive galaxies, the scale length spread also increases with galaxy mass. Overall, these results are in full agreement with the recent work by \\citet{Courteauetal07}.  Scale length variations between bands are commonly studied to better understand the content and distribution of different stellar populations, metals, and/or dust. A colour gradient is expected to increase from early-type spirals to late-type spirals,  mainly due to extinction, which increases to later types. However, for Scd galaxies or later the colour gradient becomes smaller, because of decreasing amount of extinction \\citep[see, e.g., ][]{PB96,deGrijs98}.  Changes in galaxy scale length in different wavelengths can be attributed to extinction by moderate amounts of dust, with radial metallicity and age gradients as other contributing factors \\citep[][]{EE84,Petal94, Betal96}. All these parameters will probably change as a function of redshift, enabling us to measure the variation of intrinsic scale length with cosmological epoch. It is expected that the opacity of disk galaxies is expected to have been systematically higher in the past \\citep[e.g., ][]{Dwek98, Pei99}. The fact that the observed radial colour vary little suggests strongly that stellar population effects are not important here. Dust effects are studied by using radiative transfer models which take into account scattering as well as absorption by dust, and observationally by investigating the scale length ratio in different bands as a function of inclination, i.e., optical depth. However, there are some degeneracies. Any tendency of the stars in the outer parts of disks to be bluer would tend to result in underestimated dust content. Any tendency of the dust to concentrate towards the centre would result in an overestimate of the bluer scale lengths, and would not be distinguishable photometrically from a tendency of the stars in the outer disc to be bluer \\citep[c.f.][]{PT06,Erwinetal08,Azzolinietal08}. \\citet{Petal94,Petal95} found that scale length ratios could change, due to stellar population changes, by a factor of approximately 1.1-1.2 from blue to near-infrared ($K$-band). This would correspond to a factor of about 1.03-1.06 from $g$ to $z$. Since these numbers are very small, a very accurate analysis is needed to derive conclusions from the SDSS database. Our results seem to first order in agreement with these numbers.  The derived scale lenghts and our presentation of the transformation coefficients for converting observed scale lengths from one SDSS band to another, furthermore, are meant to be useful tools to test the results of cosmological galaxy formation models, whether numerical, or semi-analytical.   In the future, we plan to add the stored parameters form the Galaxy ZOO\\footnote{http://www.galaxyzoo.org} project to obtain a larger sample with morphological classifications (at this stage, detailed morphologies are not available), but also compare scale length as a function environment, nuclear activity, and colour gradients \\citep[e.g., comparing the sample with that of ][]{Chatzimi05}. Many more parameters can be further investigated, and with the data and the derived parameters at hand, we now have the capability to continue this project in various and potentially unforeseeable directions."
    },
    "1004/1004.3819_arXiv.txt": {
        "abstract": "In this work, we investigate the accuracy of various approximate expressions for the transit duration of a detached binary against the exact solution, found through solving a quartic equation.  Additionally, a new concise approximation is derived, which offers more accurate results than those currently in the literature. Numerical simulations are performed to test the accuracy of the various expressions.  We find that our proposed expression show yields a $>200$\\% improvement in accuracy relative to the most previously employed expression.  We derive a new set of equations for retrieving the lightcurve parameters and consider the effect of falsely using circular expressions for eccentric orbits, with particularly important consequences for transit surveys. The new expression also allows us to propose a new lightcurve fitting parameter set, which minimizes the mutual correlations and thus improves computational efficiency.  The equation is also readily differentiated to provide analytic expressions for the transit duration variation (TDV) due to secular variations in the system parameters, for example due to apsidal precession induced by perturbing planets. ",
        "introduction": "Transiting exoplanets and eclipsing binaries produce familiar U- and V-shaped lightcurves with several defining quantities, such as the mid-eclipse time, eclipse depth and duration. Out of these, the transit duration is undoubtedly the most difficult observeable to express in terms of the physical parameters of the system. Kipping (2008) showed that the duration is found by solving a quartic equation, to which exists a well-known solution. In general, two roots correspond to the primary eclipse and two to the secondary but this correspondence determination has an intricate dependency on the input parameters for which there currently exists no proposed rules. As a consequence, there currently exists no single exact expression for the transit duration.  In many applications, the process of discarding unwanted roots may be performed by a computer, but naturally this can only be accomplished for case-by-case examples. The benefits of a concise, accurate and general expression for the transit duration, as we will refer to it from now on, are manifold.  The solution provides lower computation times, deeper insight into the functional dependence of the duration and a decorrelated parameter set for fitting eclipse lightcurves (see \\S6).  Such a solution may also be readily differentiated to investigate the effects of secular and periodic changes in the system (see \\S7). With the changes in transit duration recently being proposed as a method of detecting additional exoplanets in the system (Miralda-Escud\\'{e} 2002; Heyl \\& Gladman 2007) and companion exomoons (Kipping 2009a,b), there is a strong motivation to ensure an accurate, elegant equation is available.  In this work, we will first propose a new approximate expression for the transit duration in \\S2. In \\S3, we will derive two new approximate expressions for the transit duration and discuss others found in the literature; exploring their respective physical assumptions and derivations. In \\S4, we present the results of numerical tests of the various formulae.  We find that one of our new proposed expressions offers the greatest accuracy out of the candidates.  In \\S5 and \\S6, we utilize the new equation to derive mappings between the observeable transit durations and the physical model parameters.  These mappings may be used to obtain a decorrelated transit lightcurve fitting parameter set, which enhances the compuational efficiency of Monte Carlo based methods. Finally, in \\S7, we differentiate the favoured solution to predict how the transit duration will change due to apsidal precession, nodal precession, in-fall and eccentricity variation. ",
        "conclusions": "We have derived and tested a new approximate expression for the transit duration of an extrasolar planet with non-zero orbital eccentricity (equation 15). The expression has been shown to analytically reduce down to the exact expressions for a circular orbit, unlike the most previously utilized equation.  In numerical tests, the new equation is shown to be more accurate than the other candidate expressions considered in this work, in particular for highly eccentric systems.  Quantitvely, the new expression yields a $>200$\\% improvement in accuracy over the previously most utilized expression.  Manipulation of the new expression provides for a new lightcurve fitting parameter set which is based upon a modification of a previously proposed set by \\citet{bak07}.  The new parameter set is shown to demonstrate the lowest mutual correlations compared to other most commonly used parameter sets and therefore yields the most efficient algorithm for fitting lightcurves.  Additionally, we have shown that the effect of even mild eccentricity can cause very large biases in the lightcurve derived stellar density, which is often used a method for discriminating between planets and blends in transit surveys.  Consequently, planetary candidates can be either falsely rejected or accepted for systems with poor contraints on the eccentricity.  Finally, we have used the new equation to derive the rates of secular transit duration variation (TDV) as a result of apsidal precession, nodal precession (e.g. due to a perturbing planet), in-falling extrasolar planets and eccentricity variation (e.g. Kozai mechanism).  These derivatives will provide for a more accurate interpretation of secular TDV."
    },
    "1004/1004.4336_arXiv.txt": {
        "abstract": "{} {We study the star formation histories (SFH) and stellar populations of 213 red and 226 blue nearly face-on low surface brightness disk galaxies (LSBGs), which are selected from the main galaxy sample of Sloan Digital Sky Survey (SDSS) Data Release Seven (DR7). We also want to compare the stellar populations and SFH between the two groups. } {The sample of both red and blue LSBGs have sufficient signal-to-noise ratio in the spectral continua. We obtain their absorption-line indices (e.g. $Mg_2$, $H\\delta_A$), $D_n$(4000) and stellar masses from the MPA/JHU catalogs to study their stellar populations and SFH. Moreover we fit their optical spectra (stellar absorption lines and continua) by using the spectral synthesis code STARLIGHT on the basis of the templates of Simple Stellar Populations (SSPs). } {We find that red LSBGs tend to be relatively older, higher metallicity, more massive and have higher surface mass density than blue LSBGs. The $D_n$(4000)-$H\\delta_A$ plane shows that perhaps red and blue LSBGs have different SFH: blue LSBGs are more likely to be experiencing a sporadic star formation events at the present day, whereas red LSBGs are more likely to form stars continuously over the past 1-2 Gyr. Moreover, the fraction of galaxies that experienced recent sporadic formation events decreases with increasing stellar mass. Furthermore, two sub-samples are defined for both red and blue LSBGs: the sub-sample within the same stellar mass range of 9.5 $\\leq$ log$(M_\\star/M_\\odot)$ $\\leq$ 10.3, and the surface brightness limiting sub-sample with $\\mu_0(R)$ $\\geq$ 20.7 mag arcsec$^{-2}$. They show consistent results with the total sample in the corresponding relationships, which confirm that our results to compare the blue and red LSBGs are robust. } {} ",
        "introduction": "Galaxies with central surface brightness much fainter than the value of $\\mu_0(B) = 21.65 \\pm 0.3$ mag arcsec$^{-2}$ are well known as low surface brightness galaxies (LSBGs). Although they are faint compared to the night sky and hard to find, they represent a significant fraction of the number density of galaxies in the universe (O'Neil \\& Bothun 2000; Trachternach et al. 2006) and may comprise up to half of the local galaxy population (McGaugh et al. 1995).  The most widely studied LSBGs are blue (e.g. Zackrisson et al. 2005; Vorobyov et al. 2009), which showed that they appear to have lower metallicity (Burkholder et al. 2001, Haberzettl et al. 2007), lower star formation rate (O'Neil et al. 2007), are evolving much more slowly (van den Hoek et al. 2000), have larger gas fraction (McGaugh \\& de Blok 1997; Schomert et al. 2001), lower galaxy density (Rosehanm et al. 2009) and larger amounts of dark matter (de Blok \\& McGaugh 1997) than what is typically found in normal galaxies.  The wide-field CCD survey of O'Neil et al. (1997) firstly found several red LSBGs with optical colors compatible with those seen in old stellar populations. In the following, some work compared the properties of the two groups. Based on the optical-near infrared color-color diagrams of 2 red and 3 blue LSBGs, Bell et al. (1999) found that red and blue LSBGs have different star formation histories (SFH): blue LSBGs are well described by models with low, roughly constant star formation rate, whereas red LSBGs are better described by a `faded disk' scenario. Furthermore, with 5 red LSBGs, Bell et al. (2000) suggested that the red LSBGs cataloged by O'Neil et al. (1997) are a heterogeneous group, which seem to have relatively few common traits.  Due to the considerable uncertainty regarding the SFH of LSBGs, further studies on the properties of red and blue LSBGs, especially their spectroscopic properties, are important for understanding their formation and evolution. Moreover, the comparison of their SFH presents a good opportunity for understanding the global properties of low surface brightness systems.  Nevertheless, the previous studies of red and blue LSBGs have been traditionally carried out with very small samples. With the advent of the large sky survey of the Sloan Digital Sky Survey (SDSS), it is now possible to extend such studies dramatically in size. Moreover, this enormous amount of high-quality data will be undoubtedly important to allow us to study the photometric and spectroscopic properties of those galaxies more carefully. We have selected a large sample of LSBGs (Zhong et al. 2008) from SDSS-DR4 (Adelman-McCarthy et al. 2006), which consists of much more red and blue LSBGs than before. In this work, we select red and blue LSBGs from the latest data release of SDSS, the Data Release Seven (DR7)\\footnote{http://www.sdss.org/DR7} (Abazjian et al. 2009), which greatly extends the sample. This large sample of galaxies will be helpful to explore and compares the SFH of red and blue LSBGs through their photometric properties and spectral features, such as the relations of $Mg_2$ vs. log$M_*$, $D_n$(4000) vs. $H\\delta_A$, $D_n$(4000) vs. log$M_*$, surface density and spectral synthesis etc.  This paper is organized as follows. In Section \\ref{sec.2}, we describe the selection of the sample. In Section \\ref{sec.3}, we study the SFH of red and blue LSBGs from some property parameters, including stellar absorption indices, $D_n$(4000), stellar mass, and surface mass density. The stellar populations, studied by using spectral synthesis through the STARLIGHT\\footnote{http://www.starlight.ufsc.br} code and the simple stellar populations (SSPs) of Bruzual \\& Charlot (2003), are given in Section \\ref{sec.4}. The properties of two sub-samples (surface brightness limiting ($\\mu_0(R)$) and stellar mass limiting) are given in Section \\ref{sec.5}. Then we discuss our results in Section \\ref{sec.6}, and summarize this work in Section \\ref{sec.7}.  Throughout the paper, a cosmological model with $H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M$ = 0.3 and $\\Omega_\\lambda$ = 0.7 is adopted. All the magnitudes, in Petrosian magnitudes, and colors presented here are corrected for Galactic extinction and $K$-correction by using the reddening maps of Schlegel et al. (1998) and the code provided by Blanton et al. (2003), respectively. ",
        "conclusions": "\\label{sec.6}  We discuss the uncertainties of spectral synthesis and the aperture effects in this section.  \\subsection{Uncertainties of Spectral Synthesis}   We discuss the uncertainties of the stellar populations of galaxies calculated by using STARLIGHT code and the SSP templets from Bruzual \\& Charlot (2003). In our earlier work, Chen et al. (2010), by using STARLIGHT to compare six popular evolutionary stellar population synthesis models, we have also discussed the uncertainties of the resulted stellar populations.  As we know, STARLIGHT group has used this code to analyze large sample of SDSS galaxies as shown in their series of work (Cid Fernandes et al. 2005, 2007; Mateus et al. 2006; Asari et al. 2007). And they have tested the uncertainties of the resulted stellar populations. In the study of Cid Fernandes et al. (2005), they found uncertainties are smaller than 0.05, 0.1, and 0.1 for young ($t < 10^8$), intermediate (10$^8 < t < 10^9$), and old ($t > 10^9$) populations for S/N$\\geq$10, respectively. In our fittings on the sample galaxies, the code provides the values (e.g. the contributed light fractions of SSPs) of last-chain-values for 7 Markov chains, and we find that most of the discrepancies in these adopted values are less than 1\\%. Thus we believe the uncertainties of the resulted light fractions are small, and we consider the uncertainties of the resulted stellar populations as insignificant.  In the fittings, most of the age sensitivity comes from the continuum shape and $D_n$(4000) break, thus degeneracy with dust should be discussed carefully. The dust extinction has been considered in STARLIGHT as a variable parameter, and we adopt the extinction law of Cardelli et al. (1989) for the code. The resulted dust extinction $A_V$ from the code is around 0.4 for the sample galaxies as given in Fig.~\\ref{fig.2} and Fig.~\\ref{fig.4}. We have also measured the $A_V$ value from H${\\alpha}$/H${\\beta}$ for the blue LSBGs from the combined spectra as given in Fig.~\\ref{fig.2}a. It is just about 0.4, consistent with the resulted one from STARLIGHT program. Thus, the dust extinction effect could have been reliably considered in the the stellar population analyses here.  Furthermore, we use a simpler method to test the age and metallicity degeneracy in the resulted stellar populations from STARLIGHT. We use 15 SSPs with 15 ages at $Z = Z_\\odot$ to do the spectral synthesis on the blue and red LSBGs, and then use another 15 SSPs with same ages but at $Z = 0.2 Z_\\odot$ to re-do the spectral synthesis. Comparing these results with those from all three $Z$ case (Table 1), for blue LSBGs, the $Z = Z_\\odot$ results show $\\sim$12\\% more young population, and $\\sim$7\\% less intermediate population; and the $Z = 0.2 Z_\\odot$ results show $\\sim$9\\% less young population, and  $\\sim$10\\% more old population; for red LSBGs, the light fractions have little changes. Thus, from these analysis the metallicity will cause uncertainties about 10\\% for the stellar populations of blue LSBGs, and very small effect on red LSBGs. But here we prefer the 3 metallicity results. In any case, the dominant population of the sample galaxies are unchanged, i.e. the young population.   \\subsection{Aperture effects}  The SDSS is a fiber based survey. Therefore, we consider the aperture effects in this section. Tremonti et al. (2004) and Kewley et al. (2005) have discussed the weak effect of the 3{$^{\\prime \\prime }$}  aperture of SDSS spectroscopy. They believe that redshifts $z >$ 0.03 and 0.04 are required for SDSS galaxies, respectively to get reliable metallicities. Given that the fiber mag is a measurement of the light going down the fiber and the petrosian mag is a good estimate of the total magnitude. Thus, to check how much the light of the galaxies was covered by the fiber observation, one simple and accurate way is to compare the ``fiber\" and ``petrosian\" magnitudes of the SDSS galaxies (Liang et al. 2010). We choose the formula: \\begin{equation} \\label{eq.fiber} light\\_fraction = 10^{(-0.4\\times(fiber\\_mag - petro\\_mag)_r)}, \\end{equation} to estimate how much light was covered by the fiber observations. Fig. \\ref{fig.5} shows the light fractions for red and blue LSBGs of our total sample. It shows that the light fractions of red (bottom) and blue (top) LSBGs are 0.16 and 0.13, respectively. Therefore, the properties of red and blue LSBGs are main from the central regions, which may be redder, older and more metal-rich than the outer regions. However, it could not affect much our results of the differences between red and blue LSBGs, because the light fractions are almost same. Moreover, all our sample galaxies are disk-dominated galaxies with very small bulges, we assume the resulted stellar populations could be good representatives of their disk populations.  \\begin{figure} \\begin{center} \\includegraphics [width=8.0cm, height=7.2cm] {./figure/histgram_aperture.eps} \\caption{The histogram distributions of the light fraction (Eq. \\ref{eq.fiber}) with bins of 0.01 for both red (bottom) and blue (top) LSBGs of our total sample. The median values of the light fractions are 0.16 for red LSBGs and 0.13 for blue LSBGs.} \\label{fig.5} \\end{center} \\end{figure}"
    },
    "1004/1004.3499_arXiv.txt": {
        "abstract": "Efforts to detect gravitational waves by timing an array of pulsars have focused traditionally on stationary gravitational waves: e.g., stochastic or periodic signals. Gravitational wave bursts --- signals whose duration is much shorter than the observation period --- will also arise in the pulsar timing array waveband. Sources that give rise to detectable bursts include the formation or coalescence of supermassive black holes (SMBHs), the periapsis passage of compact objects in highly elliptic or unbound orbits about a SMBH, or cusps on cosmic strings. Here we describe how pulsar timing array data may be analyzed to detect and characterize these bursts. Our analysis addresses, in a mutually consistent manner, a hierarchy of three questions: \\emph{i}) What are the odds that a dataset includes the signal from a gravitational wave burst? \\emph{ii}) Assuming the presence of a burst, what is the direction to its source? and \\emph{iii}) Assuming the burst propagation direction, what is the burst waveform's time dependence in each of its polarization states? Applying our analysis to synthetic data sets we find that we can \\emph{detect} gravitational waves even when the radiation is too weak to either localize the source of infer the waveform, and \\emph{detect} and \\emph{localize} sources even when the radiation amplitude is too weak to permit the waveform to be determined. While the context of our discussion is gravitational wave detection via pulsar timing arrays, the analysis itself is directly applicable to gravitational wave detection using either ground or space-based detector data. ",
        "introduction": "It has been just over thirty years since \\citet{sazhin:1978:ofd} and \\citet{detweiler:1979:ptm} showed how gravitational waves could be detected by correlating the timing residuals of a collection of pulsars, and twenty years since \\citet{foster:1990:cpt} proposed using a collection of pulsars --- i.e., a pulsar timing array --- to achieve greater sensitivity. Over the ensuing years telescope collecting area has increased, antenna temperature has decreased, pulsar timing electronics and methodology has improved, and pulsars with exceptionally low intrinsic timing noise have been discovered. As a result of these advances the near-future detection of a stochastic gravitational wave signal through pulsar timing observations is a strong possibility.  Analyses aimed at detecting gravitational waves using pulsar timing array observations have traditionally focused on stationary signals (i.e., stochastic or periodic gravitational waves). More generally, analyses aimed at detecting gravitational waves have merged the questions of detection and characterization, overlooking the possibility of detecting a signal that is too weak to be characterized. Here we describe how pulsar timing array data may be analyzed to search for gravitational wave \\emph{bursts}, demonstrating that \\emph{i}) pulsar timing array data is sufficiently rich to allow the detection of gravitational wave bursts, the localization of the burst source, and the time-dependent waveform of the radiation in its (two) polarization states; and, \\emph{ii}) that gravitational wave signals too weak to be characterized, or too weak to allow their source to be localized, may still be strong enough to be unambiguously detected.  The first detections of gravitational waves will be important for confirming their existence and testing whether general relativity correctly predicts their properties (e.g., polarization modes, propagation speed). Of perhaps greater long-term significance will be the use of gravitational waves as a tool of observational astronomy that gives us direct insight into phenomena that we can now observe only indirectly, if at all. For example, \\citet{jaffe:2003:gwp,wyithe:2003:lgw} and \\citet{jenet:2006:ubo} have shown that root-mean-square (rms) amplitude of a stochastic signal arising from the confusion limit of a large number of supermassive black hole binary coalescences is within an order of magnitude of the current sensitivity of the most advanced pulsar timing array. Since the signals that contribute to this background arise from a population of discrete sources distributed throughout space we quite reasonably expect that some of the individual sources may be observable as gravitational wave bursts rising above this background. Indeed, recent work by \\citet{sesana:2008:sgb} shows that at frequencies greater than a few times $10^{-8}$ nanohertz the gravitational ``background'' arising from supermassive binary black hole coalescense should be dominated by a few bright sources. Other potential burst gravitational wave sources in the pulsar timing array band include cosmic (super)string cusps and kinks \\citep{damour:2001:gwb, siemens:2007:gsb,leblond:2009:gwf}.  Pulsar timing array observations are sensitive to gravitational waves of periods ranging from the interval  between timing observations (days to months) and the duration of the observational data sets (years). The corresponding wavelengths are much greater than those explored in existing or proposed human-built ground or space-based detectors. Ground-based detectors, whether of the  acoustic \\citep{astone:2010:i1s} or interferometric variety \\citep{accadia:2010:sap,riles:2010:pgw}, are currently sensitive to waves in the $\\sim$ 100~Hz -- 1~kHz band with proposed advances opening-up the 10--100~Hz band \\citep{smith:2009:pte,kuroda:2006:sol}. Space-based detectors, which have been the subject of extensive design studies over the last thirty years, would be sensitive to gravitational waves in the $\\sim3\\times10^{-5}$~Hz -- 10~Hz band \\citep{stebbins:2006:lmt,jennrich:2009:trl,kawamura:2008:jsg}. Over this broad band --- $10^{-9}$ -- $10^3$~Hz --- the scale and character of the sources varies dramatically: e.g., ground-based detectors will be sensitive to gravitational waves from neutron star or solar mass black hole binaries, supernovae and gamma-ray burst progenitors; space-based detectors to gravitational waves from white dwarf binaries, stellar disruptions about intermediate mass black holes and the inspiral of solar mass compact objects or intermediate mass black holes about $10^{4.5}--10^{7.5}\\,\\mathrm{M}_\\odot$ black holes; and pulsar timing arrays to the formation, interaction and evolution of supermassive black holes. Pulsar timing array observations thus offer their own, unique perspective on the gravitational wave universe, distinct from that provided by the either ground- and space-based detectors.  Analyses aimed at detecting gravitational waves using pulsar timing array observations have traditionally focused on stationary signals: i.e., an isotropic stochastic gravitational wave background \\citep{hellings:1983:ulo,mchugh:1996:pta,thorsett:1996:ptl,lommen:2001:pmt,lommen:2003:nlo,jenet:2005:dsg,jenet:2006:ubo,demorest:2007:mgw,hobbs:2008:upt,haasteren:2009:omg,anholm:2009:osf} or gravitational waves from discrete periodic sources \\citep{lommen:2001:upt,jenet:2004:cpo,jenet:2005:pta,jenet:2005:cpo}. More recent work \\citep{haasteren:2009:gma} has investigated the detection of gravitational wave ``memory'' \\citep{christodoulou:1991:nno} associated with sources that radiate a significant amount of energy in gravitational waves \\citep{wiseman:1991:cng} or that become unbound \\citep{thorne:1992:gbw}.  Gravitational wave detection using a pulsar timing array begins with the collection of timing residuals from the several array pulsars. These timing residuals are the difference between the expected pulse arrival times (taking into account all non-gravitational-\\emph{wave} propagation effects) and the actual pulse arrival times at each pulsar observational epoch. For pulsars used in current timing arrays the timing precision is in the 50~ns -- 5~$\\mu$s range. In \\S\\ref{sec:response} we summarize how these timing residuals reflect the passage of a plane gravitational wave through the pulsar-Earth baseline. In \\S\\ref{sec:methodology} we describe our analysis for gravitational wave bursts, which  takes advantage of the correlation of the timing residuals measured for different pulsars. In \\S\\ref{sec:examples} we demonstrate the effectiveness of the analysis by applying it to simulated data arising from a thirty pulsar timing array and including a gravitational wave burst such as would be expected from a parabolic encounter of two supermassive black holes. Finally, in \\S\\ref{sec:conclusions} we summarize our findings and describe planned future work. ",
        "conclusions": "\\label{sec:conclusions}  In the history of astronomy few (if any) new observational windows have been as eagerly anticipated as the gravitational wave window, whose opening will provide us with a novel and direct view of astronomical phenomena that can now be inferred at best dimly and indirectly.  Making sense of what we see through this new window requires analysis tools and techniques adapted to the unique nature of our new ``telescopes'' and the sources they enable us to study. Here we have described an intra-related suite of analysis techniques for gravitational wave astronomy designed to address quantitatively three specific questions: \\begin{enumerate} \\item What are the ``odds'' that a gravitational wave detector data set includes the signal from a gravitational wave burst? \\item Assuming that a gravitational wave burst is present in a data set, what is the probability that the wave is propagating in direction $\\hat{k}$? \\item Assuming the presence of a burst propagating in direction $\\hat{k}$, what is the probability that the wave at Earth is characterized by the functions $h_{+}(u)$ and $h_{\\times}(u)$, $u = t-\\hat{k}\\cdot\\vec{x}$, representing the $+$ and $\\times$ polarization state waveforms? \\end{enumerate}  We address these questions in the specific context of gravitational wave detection using pulsar timing array data. Until recently, analyses for gravitational wave detection using timing data from an array of pulsars has focused on stationary sources: e.g., a stochastic gravitational wave signal or the signal from a binary system. By addressing burst sources we also add to the very recent literature examining how pulsar timing data can be used to detect gravitational wave bursts \\citep{haasteren:2009:gma} such as might arise a close fly-by or collision of two supermassive black holes or from a cosmic string cusp \\citep{binetruy:2009:gwb,key:2009:cgw}.  To demonstrate the efficacy of our analysis we applied it to four synthetic timing residual data sets  representative of observations using the International Pulsar Timing Array \\citep[IPTA]{hobbs:2009:ipt}. Each data set included simulated timing noise, constructed to be characteristic of actual IPTA timing noise. Three of the datasets included the timing residual signature of a gravitational wave burst characteristic of a parabolic fly-by of two supermassive black holes; the fourth did not. The three ``signal-present'' cases varied only by the gravitational wave signal amplitude. In the case of the strongest signal the burst was unambiguously detected, localized to much better than a $\\deg^2$, and the waveform in the individual polarization states recovered. In the moderate signal amplitude case the signal was, again, unambiguously detected and the general direction to the source clearly determined; however, the signal amplitude was too low to infer the waveform characteristics. In the third case the signal was strong enough to be detected but too weak to be characterized or to allow the source to be localized. Finally, in analyzing noise alone the calculated odds were, as they should be, unambiguously against the presence of a gravitational wave burst.  At present, pulsar timing array data sets are constructed by fitting a timing mode to the time of arrival (TOA) data for each pulsar. This timing model includes, in parameterized form, all the non-gravitational-wave contributions that affect the pulse arrival times. The residual differences between the timing model predictions and the actual arrival times are then analyzed for the signature of a passing gravitational wave. This procedure has the disadvantage of being incapable of identifying any gravitational wave whose effect on the arrival time of individual pulsars is degenerate with any of the non-gravitational-wave effects that are part of the timing model. Our analysis may be extended to infer, in addition to the gravitational radiation waveform, the other timing model parameters. This extended analysis may be applied to TOA data directly, avoiding entirely the problem of ``fitting-out'' gravitational wave contributions whose character is similar to other timing model contributions. We intend to investigate this extension in future work.  While our presentation and discussion have focused on pulsar timing array observations the analysis methodology that we describe applies equally well and without modification to gravitational wave data taken from ground-based detector networks like the LIGO-Virgo network \\citep{accadia:2010:sap,riles:2010:pgw}, or for the analysis of LISA \\citep{jennrich:2009:trl,merkowitz:2009:cls} data."
    },
    "1004/1004.3666_arXiv.txt": {
        "abstract": "By means of Monte Carlo simulations of extensive air showers (EAS), we have performed a comprehensive study of the shower to shower fluctuations affecting the longitudinal and lateral development of EAS. We split the fluctuations into physical fluctuations and those induced by the thinning procedure customarily applied to simulate showers at EeV energies and above. We study the influence of thinning on the calculation of the shower to shower fluctuations in the simulations. For thinning levels larger than $R_{\\rm thin}=10^{-5} - 10^{-6}$, the determination of the shower to shower fluctuations is hampered by the artificial fluctuations induced by the thinning procedure. However, we show that shower to shower fluctuations can still be approximately estimated, and we provide expressions to calculate them. The influence of fluctuations of the depth of first interaction on the determination of shower to shower fluctuations is also addressed. ",
        "introduction": "Extensive air showers (EAS) have been studied over the last 70 years \\cite{watson-nagano}.  They result from the interaction in the atmosphere of high-energy protons and nuclei arriving from space.  The product of these collisions are a set of secondary particles carrying a fraction of the primary energy.  These secondaries move through the atmosphere and interact again generating new secondaries.  The process continues, increasing the number of secondary particles, until their energies are too low to contribute to the generation of new particles. Particles reaching ground are sampled with arrays of detectors, and their properties are used to infer the properties of the primary initiating the shower. Measurements of the electron and muon density, of the arrival time of the particles at ground, and of the depth at which the shower has the maximum number of particles (X$_{\\rm max}$), give information on the arrival direction, primary energy, and on the mass of the primaries \\cite{watson-nagano}.  The complexity of the cascade phenomena, and the poor knowledge of the hadronic interactions at very high energy \\cite{hadronic_models}, make the experimental determination of the properties of the primaries very difficult. Moreover, primary particles with the same energy, mass and direction produce secondary particles with parameters that vary from shower to shower. This feature is called ``shower to shower fluctuations\".  An understanding of the shower to shower fluctuations will help to improve the interpretation of cosmic-ray data.  The calculation of shower to shower fluctuations can in principle be addressed with Monte Carlo simulations of extensive air showers. However, the number of particles that are produced in an air shower at ultra high energy (above $\\sim 10^{18}$ eV) is so large ($\\sim 10^{10}$), that it is almost impossible to follow the propagation to ground level of all the secondaries in the Monte Carlo in a reasonable amount of time, or even to store the large amount of information produced. For this reason, a statistical sampling procedure called ``thinning\" \\cite{Hillas_thin} is used in the simulations. Thinning algorithms typically consist on propagating only a small, representative fraction of the total number of particles in the shower, assigning statistical weights to the sampled particles to compensate for the rejected ones. However, thinning algorithms introduce artificial fluctuations in the simulated showers, hampering the determination of the intrinsic, physical shower to shower fluctuations with Monte Carlo simulations. For this reason the study of fluctuations using Monte Carlo simulations is quite difficult and uncertain. This is of utmost importance in cosmic-ray physics, since an incorrect assumption on the shower to shower fluctuations can lead to systematic errors on the determination of the parameters of the primary particles.  In this work we address the problem of determining the true, physical shower to shower fluctuations in Monte Carlo simulations, and quantify the effect of thinning on their determination. We give expressions that allow the estimation of physical fluctuations from Monte Carlo simulations, even in the case of relatively strongly thinned showers.  The paper is organized as follows: In Section \\ref{S:TheSimulations} we describe the simulations performed in this work, and the thinning algorithm adopted. In Section \\ref{S:Fluctuations} we identify the different sources of fluctuations in shower simulations. In Section \\ref{S:Results} we perform a comprehensive study of the fluctuations in the longitudinal and lateral shower development, and give expressions that allow to separate physical shower to shower fluctuations from the artificial fluctuations induced by the thinning procedure. In Section \\ref{S:Composition}, we quantify the influence on the shower to shower fluctuations of the fluctuations of the depth of first interaction of the primary initiating the shower. Finally, we summarize our conclusions in section \\ref{S:Conclusions}. In the Appendix we give an explicit mathematical derivation of the expressions presented in Section~\\ref{S:Results}. ",
        "conclusions": "\\label{S:Conclusions}  In this work we have performed a comprehensive study of shower to shower fluctuations by means of Monte Carlo simulations of extensive air showers. An understanding of the shower to shower fluctuations will help to improve the interpretation of cosmic-ray data.  We have shown that the determination of the true, physical shower to shower fluctuations is hampered by the thinning procedure necessary to simulate in a practical manner air showers at EeV energies and above. However, we also show that the artificial fluctuations induced by thinning ($\\sigma_{\\rm thin}$) can be identified and splitted from the physical fluctuations ($\\sigma_{\\rm phys}$) with the aid of Eq.~(\\ref{eq:sigma}), which we have shown to account for all, true and artificial fluctuations appearing in the simulations. Eq.~(\\ref{eq:sigma}) reproduces the expectation that as the thinning level decreases $R_{\\rm thin}\\rightarrow 0$, and showers are less thinned, then the artificial fluctuations decrease, the physical ones become dominant, and they do not depend on $R_{\\rm thin}$.  Our simulations also indicate that the physical shower to shower fluctuations of the number of particles at ground behave proportionally to the number of particles $N$, while the artificial fluctuations are Poissonian, i.e., behave as $\\sqrt N$.  Besides, we have shown that the size of the relative fluctuations due to the depth at which the first interaction initiating the shower occurs, is smaller or of the same order as the fluctuations occuring in the subsequent secondary interactions in the shower."
    },
    "1004/1004.2383_arXiv.txt": {
        "abstract": "s{WIMP direct detection experiments probe the ultra-local dark matter density and velocity distribution. We review how uncertainties in these quantities affect the accuracy with which the WIMP mass and cross-section can be constrained or determined.} ",
        "introduction": "If Weakly Interacting Massive Particles (WIMPs) are detected then measuring the WIMP mass and cross-section will shed light on the particle nature of the WIMP. It will allow us to constrain models of particle physics beyond the standard model, for instance supersymmetry and universal extra dimensions. Furthermore a convincing WIMP discovery may require consistent signals (i.e. same inferred mass) in multiple experiments in multiple channels (direct, indirect, collider). We discuss the effect of astrophysical uncertainties on the accuracy with which the WIMP mass and cross-section could be measured from direct detection experiments.  Direct detection experiments aim to detect WIMPs in the Milky Way halo via the energy deposited when they elastically scatter off detector nuclei~\\cite{ddtheory,book}. Assuming spin-independent coupling the differential event rate (number of events per unit energy, per unit time, per unit detector mass) is given by~\\cite{jkg,ls} \\be \\label{drde} \\frac{{\\rm d} R}{{\\rm d}E} = \\frac{\\sigma_{{\\rm p}} \\rho_{\\chi}}{2 \\mu_{{\\rm p} \\chi}^2 m_{\\chi}} A^2 F^2(E)   \\int^{\\infty}_{v_{{\\rm min}}} \\frac{f(v)}{v} {\\rm d}v     \\,, \\ee where $\\rho_{\\chi}$ is the ultra-local WIMP density, $f(v)$ the ultra-local WIMP speed distribution in the rest frame of the detector, $\\sigma_{{\\rm p}}$ the WIMP scattering cross section on the proton, $\\mu_{{\\rm p} \\chi} = (m_{\\rm p} m_{\\chi})/(m_{{\\rm p}}+ m_{{\\chi}})$ the WIMP-proton reduced mass, $A$ and $F(E)$ the mass number and form factor of the target nuclei respectively and $E$ is the recoil energy. The lower limit of the integral, $v_{{\\rm min}}$, is the minimum WIMP speed that can cause a recoil of energy $E$: \\begin{equation} \\label{vmin} v_{{\\rm min}}=\\left( \\frac{ E m_{A}}{2 \\mu_{{\\rm A} \\chi}^2} \\right)^{1/2} \\,, \\end{equation} where $m_{A}$ is the atomic mass of the detector nuclei and $\\mu_{{\\rm A} \\chi}$ the WIMP-nucleon reduced mass.  The energy, time and direction dependence of the event rate provide potential signals for discriminating WIMP events from backgrounds. If we assume, as a starting point, the standard halo model with an isotropic Gaussian speed distribution \\be \\label{max} f(v) \\propto \\exp{ \\left( - \\frac{v^2}{v_{\\rm c}^2} \\right)} \\,, \\ee where $v_{\\rm c}$ is the local circular speed (the speed with which stars in the solar neighbourhood orbit the Galactic centre), then the differential event rate has a characteristic energy~\\cite{ls} \\be \\label{er} E_{\\rm R} =  \\frac{2 \\mu_{{\\rm A} \\chi}^2 v_{\\rm c}^2}{m_{\\rm A}} \\,. \\ee If $m_{\\chi} \\ll m_{\\rm A}$ then $E_{\\rm R} \\propto m_{\\chi}^2$, while if $m_{\\chi} \\gg m_{\\rm A}$ then $E_{\\rm R} \\sim {\\rm const}$. Therefore, provided the WIMP mass is not too large or not small~\\footnote{If the WIMP mass is too small the event rate above the detector energy threshold becomes small.}, a measurement of the WIMP mass can be extracted from the energy spectrum of events.  The motion of the Earth with respect to the Galactic rest frame provides two potential WIMP `smoking guns'. Due to the Earth's orbit about the Sun the net velocity of the lab with respect to the Galactic rest frame varies over the course of a year. The net speed is largest in the Summer and hence there are more  high speed WIMPs, and less low speed WIMPs, in the lab frame. This produces an energy dependent, ${\\cal O}(5\\%)$, annual modulation in the differential event rate~\\cite{dfs}. To detect this a large exposure (detector mass times time) is required.  Our motion with respect to the Galactic rest frame also produces a directional signal~\\cite{spergel}. The WIMP flux in the lab frame is sharply peaked in the direction of motion of the Sun (towards the constellation CYGNUS).  The recoil spectrum is then peaked in the direction opposite to this. This is a large signal; with an ideal detector only 10 events would be needed to demonstrate that the recoil rate was isotropic~\\cite{mgs} and $\\sim 30$ to confirm that the median recoil direction is opposite to the direction of solar motion~\\cite{billard,gm}. A detector which can measure the recoil directions is required though (e.g. DMTPC, DRIFT, MIMAC, NEWAGE). See Ref.~\\cite{ahlen} for a review of the current status of directional detection experiments. ",
        "conclusions": "The energy, time and direction dependence of the direct detection event rate all depend on the ultra-local dark matter distribution. Deviations from the standard halo model which affect these signals are expected. The uncertainty in the ultra-local dark matter density translates directly into an uncertainty in the event rate and hence measurements of (or constraints on) the WIMP cross-section. The uncertainty in the circular speed (which is related to the WIMP velocity dispersion) leads to an uncertainty in characteristic scale of energy spectrum and hence measurements of the WIMP mass. The uncertainty in the shape of WIMP velocity dispersion produces an uncertainty in the amplitude and phase of the annual modulation signal and hence the WIMP parameters.  Several approaches to dealing with the impact of astrophysical uncertainties have recently been proposed. Strigari and Trotta~\\cite{strigari} have suggested using data (such as measurements of the kinematics of Milky Way halo stars and the local escape speed) and a model for the Milky Way mass distribution in a Monte Carlo Markov Chain analysis of direct detection data. Peter~\\cite{peter} has presented an approach which involves combining data sets from different direct detection experiments and jointly constraining  a parametrisation of the WIMP speed distribution and the WIMP parameters (mass and cross-section). Both of these approaches are contingent on formulating a reliable, robust model of the Milky Way.  In the long term the dependence of the direct detection event rate on the ultra-local dark matter distribution is an opportunity rather than an inconvenience. With sufficient data direct detection experiments (in particular directional ones) could measure the ultra-local dark matter distribution and shed light on the formation of the Milky Way."
    },
    "1004/1004.0386_arXiv.txt": {
        "abstract": "Data from the SDSS ($\\sim 300,000$ galaxies) indicates that recent star formation (within the last 1 billion years) is bimodal: half the stars form from gas with high amounts of metals (solar metallicity), and the other half form with small contribution of elements heavier than Helium ($\\sim 10-30\\%$ solar). Theoretical studies of mass loss from the brightest stars derive significantly higher stellar-origin BH masses ($\\sim 30-80 \\msun$) than previously estimated for sub-solar compositions. We combine these findings to estimate the probability of detecting gravitational waves (GWs) arising from the inspiral of double compact objects. Our results show that a low metallicity environment significantly boosts the formation of double compact object binaries with at least one BH. In particular, we find the GW detection rate is increased by a factor of $20$ if the metallicity is decreased from solar (as in all previous estimates) to a 50--50 mixture of solar and $10\\%$ solar metallicity. The current sensitivity of the two largest instruments to NS-NS binary inspirals (VIRGO: $\\sim 9$ Mpc; LIGO: $\\sim 18$) is not high enough to ensure a first detection. However, our results indicate that if a future instrument increased the sensitivity to $\\sim 50-100\\,\\mbox{Mpc}$, a detection of GWs would be expected within the first year of observation. It was previously thought that NS-NS inspirals were the most likely source for GW detection. Our results indicate that BH-BH binaries are $\\sim 25$-times more likely sources than NS-NS systems and that we are on the cusp of GW detection. ",
        "introduction": "Gravitational waves (GWs) are a consequence of Einstein's (1918) theory of general relativity. The first indirect evidence for the existence of GWs was presented by Hulse and Tylor (1974) through measurement of the orbital decay of a double binary pulsar. After four decades of effort, we have yet to directly detect GWs. In the past decade, two large interferometric observatories have been built to search for gravitational waves: LIGO and VIRGO. The first observations have been collected, and although no positive detection has been made, some useful upper limits have been placed (Abbott et al. 2008, 2009a). Both observatories will soon undergo major upgrades aimed to increase their sensitivity roughly ten-fold. In this {\\em Letter} we show that existing detectors are at the cusp of detecting gravitational waves, and even modest improvements are likely to lead to the first direct detection of gravitational waves.  We utilize the {\\tt StarTrack} population synthesis code (Belczynski et al. 2002) to perform a suite of Monte Carlo simulations of the stellar evolution of stars in environments with a range of metallicity. Our calculations are guided by recent results from the Sloan Digital Sky Survey observations (Panter et al. 2008), combined with recent estimates of mass loss rates (Belczynski et al. 2010a). Utilizing these latest results, we calculate a population of 2 million massive binary stars, tracking the ensuing formation of relativistic binary compact objects: double neutron stars (NS-NS), double black hole binaries (BH-BH), and mixed systems (BH-NS). Our modeling utilizes updated stellar and binary physics, including results from supernova simulations (Fryer \\& Kalogera 2001) and compact object formation (Timmes et al. 1996), incorporating elaborate mechanisms for treating stellar interactions like mass transfer episodes (Belczynski et al. 2008) and tidal synchronization and circularization (Hut 1981). We put special emphasis on the common envelope evolution phase (Webbink 1984), which is crucial for close double compact object formation as the attendant mass transfer allows for efficient hardening of the binary. This orbital contraction can be sufficiently efficient to cause the individual stars in the binary to coalesce and form a  single highly rotating object, thereby aborting further binary evolution and preventing the formation of a double compact object. Due to significant radial expansion, stars crossing the Hertzsprung gap (HG) very frequently initiate a common envelope phase. HG stars do not have a clear entropy jump at the core-envelope transition (Ivanova \\& Taam 2004); if such a star overflows its Roche lobe and initiates a common envelope phase, the inspiral is expected to lead to a coalescence (Taam \\& Sandquist 2000). In particular, it has been estimated that for a solar metallicity environment (e.g., our Galaxy), properly accounting for the HG gap may lead to a reduction in the merger rates of BH-BH binaries by $\\sim 2-3$ orders of magnitude (Belczynski et al. 2007). The details of the common envelope phase are not yet fully understood, and thus in what follows we consider two models, one which does not take into account the suppression (optimistic model: A), and one that assumes the maximum suppression (pessimistic model: B). ",
        "conclusions": "Recent observational data indicates that about half of recent star formation arises in low metallicity galaxies (Panter et al. 2008). We have created a synthetic model of the local Universe that assumes half of the stars are forming at solar metallicity, and the other half are forming at $10\\%$ solar. This range of metallicity has profound effects on the detectable rate of gravitational waves from binary systems. The predicted detection rates are presented in Table~2, as a function of gravitational detector sensitivity. LIGO (with its sensitivity of about 18 Mpc) has yet to detect a binary inspiral (Abbott et al. 2009b). Our optimistic model (A), with a predicted detection rate of $5$ yr$^{-1}$ for $d_{\\rm 0,nsns}=18$ Mpc, is thus already excluded by LIGO. Our pessimistic model (B), with a much lower detection rate of $0.05$ yr$^{-1}$ for $d_{\\rm 0,nsns}=18$ Mpc, is consistent with observations. This hints that some fraction of massive binaries cannot be surviving the common envelope phase.  We employ the results of our pessimistic model (B) to estimate the sensitivity at which instruments like LIGO or VIRGO would detect 1 and 10 inspirals per year. Assuming a local density of Milky Way-size galaxies of $\\rho_{\\rm gal} = 0.01$ Mpc$^{-3}$ (O'Shaughnessy et al. 2008), we extrapolate to a distance sufficient to contain enough star forming mass to generate the desired merger rate (our calculated rate of mergers per Milky Way-size galaxy are listed in Table~1). It is found that instruments need to reach a sensitivity of $d_{\\rm 0,nsns}=45\\,\\mbox{Mpc}$ and $97\\,\\mbox{Mpc}$ to provide $1$ and $10$ detections per year, respectively. We also note that the detections will be most likely dominated by BH-BH binaries ($\\sim 80\\%$), with a smaller contribution from BH-NS systems ($\\sim 15\\%$), and a negligible contribution from NS-NS inspirals ($\\sim 5\\%$).  Our results are subject to a number of uncertainties. First, the star formation rates may be different from those presented in Panter et al. (2008). Stars which formed in the remote past but are merging at present day may manifest themselves as an additional and significant component to the detection rates. In-depth study of the overall cosmic star formation rate, and the contribution of mergers from stars born at high redshift, using existing tools (Belczynski et al. 2010b) is underway. Second, the process of BH formation is still somewhat uncertain. Evolutionary effects like wind mass loss rates may alter the final BH mass, while natal kicks (or lack thereof) or the disputed mass delineation between NS and BH formation may substantially affect the rates. Although we have chosen the best available models to describe these processes, future theoretical and observational constraints may significantly affect our estimates. Finally, we can exclude 100\\% of stars following model A (ignoring suppression through the common envelope phase), since this produces rates that should have already been detected by LIGO. However, mixtures of model A and model B (e.g., 50--50 of each model) are still consistent with observational limits. The physical interpretation of this sort of mixture model is that massive stars begin developing core-envelope structure halfway through the Hertzsprung gap, and the survival through common envelope phase depends not only on the donor type, but also on its evolutionary state. Therefore a fraction of binaries survive the common envelope phase.  Our results suggest a very high potential for the detection of gravitational radiation from a binary inspiral in the near future. Our merger rates for the most likely source to be detected, a BH-BH binary, are significantly higher than previous estimates in the literature. For example, Abadie et al. (2010) estimated the realistic level of BH-BH detection at the level of $20$ yr$^{-1}$ for the advanced LIGO, while our realistic prediction is of the order of $\\gtrsim 200$ yr$^{-1}$. Note that we have not included the potentially significant contribution of dynamically formed BH-BH binaries in globular clusters (e.g., Sadowski et al. 2008; Downing et al. 2010). Our prediction is based on the most recent observational and theoretical results, and employs a standard stellar evolutionary model. Further support for our enhanced theoretical predictions is the existence of high mass X-ray binaries, such as IC10 X-1 and NGC300 X-1, that are expected to form close BH-BH binaries on a very short timescale (set by the evolution of the massive Wolf-Rayet companion). These systems have formed massive BHs in low metallicity environment (consistent with our results), but also using their nearby locations leads to a semi-empirical estimate of merger rates that are comparable to the rates presented here (Bulik et al. 2008). Extending the sensitivity of existing instruments by factors of only a few, to about $50\\mbox{--}100$ Mpc, should result in the first detection of gravitational waves. The lack of detections at increased sensitivity will put stringent constraints on the BH-BH formation (BH natal kicks) and/or on the star formation (metallicity) in the local Universe."
    },
    "1004/1004.0948_arXiv.txt": {
        "abstract": "We present a statistical characterization of the carbon-star to M-giant (C/M) ratio in the halo of M31. Based on application of pseudo-filter band passes to our Keck/DEIMOS spectra we measure the \\col-color index of 1288 stars in the giant stellar stream and in halo fields out to large distances. From this well-established narrow-band system, supplemented by V$-$I colors, we find only a low number (five in total) of C-star candidates. The resulting low C/M ratio of 10\\% is consistent  with the values in the M31 disk and inner halo from the literature. Although our analysis is challenged by small number statistics and our sample selection, there is an indication that the oxygen-rich M-giants occur in similar number throughout the entire halo. We also find no difference in the C-star population of the halo fields compared to the giant stream. The very low C/M ratio is at odds with the observed low metallicities and the presence of intermediate-age stars at large radii. Our observed absence of a substantial carbon star population in the these regions indicates that the (outer) M31 halo cannot be dominated by the debris of disk-like or SMC-type galaxies, but rather resemble the dwarf elliptical NGC~147. ",
        "introduction": "As more observations accumulate, it has become increasingly clear that the field population of massive galaxy halos consists largely of the debris of accreted satellites, along the lines of  Searle \\& Zinn (1978).  In the case of the Milky Way, the orbital stream of the Sagittarius (Sgr) dwarf spheroidal (dSph) is so complicated that it is easy to imagine its members being dispersed throughout much of the halo over a few Gyr timescale (e.g., Law et al. 2005; Fellhauer et al. 2006).  Additional observations find coherent structures and satellites over much of the Galactic halo (e.g., Belokurov 2007a,b).    Evidence of  mergers  in the form of the giant stream and other structures, is clearly even more striking in M31 (Ibata et al. 2001, 2007; Ferguson et al. 2002).  N-body models of these events show that even a single minor merger is capable of filling much of the halo with debris after only 2 Gyr (Mori \\& Rich 2008).   One is also struck by both the surface brightness and spatial extent of the giant stream and its shell feature, all of which are attributable to one relatively recent minor merger, that did not disrupt the disk (Fardal et al. 2006, 2008; Mori \\& Rich 2008).    The importance of the giant stream event emphasizes that, as halo histories are stochastic, the populations of halos may be dominated by the debris of just a few events involving massive satellites -- certainly not the picture that one may infer from the study of their globular cluster systems.  In the case of M31, direct measurements of the halo star formation history from the main sequence turnoff revealed a complex range of age and abundance (e.g., Brown et al. 2003, 2006, 2008).  In all of the M31 halo fields, there exist significant populations of intermediate-age stars, with metallicities ranging from 1/10 to twice Solar.   In addition to a contingent of stars ranging from 5--10 Gyr, nearly all M31 halo fields show signs of a blue extension of $\\sim 1$ Gyr old stars (e.g. Brown et al. 2006; Richardson et al. 2008).   If stellar systems are being ingested then  their properties should, in principle, reflect those of similar galaxies in the Local Group.   Although star formation likely ceases as a satellite is tidally disrupted, it is clear that much of the complexity can be observationally reconstructed, as has been accomplished for the Sagittarius dSph galaxy (e.g., Chou et al. 2007).  An efficient tracer of more recent accretion events through intermediate-age populations is the number ratio, C/M, of carbon-rich (C) stars to the oxygen-rich (M) stars. C- and M-type stars are the most luminous single stars found in intermediate-age populations, and their spectral signatures are distinctive.  In complex stellar systems, the presence of carbon stars in significant numbers guarantees the presence of substantial intermediate-age populations. The primary correlation in  the C/M ratio is with metallicity (Iben \\& Renzini 1983; Brewer et al. 1995); in more metal rich systems, the balance tilts toward the oxygen-rich M-stars.   For example, the Galactic bulge is dominated by M giants (Blanco et al. 1984).   Intermediate-age populations are carbon star dominated, with this tendency to increase at lower metallicity (e.g. Groenewegen 2002; Mouhcine \\& Lan\\c con 2003; Batinelli \\& Demers 2005). Likewise, radial age and/or metallicity gradients in a given stellar system have been confirmed to be accompanied by  gradients in the C/M ratio (e.g., Cioni et al. 2008). Furthermore, the Small Magellanic Cloud is well known to contrast with the LMC in having a high C/M ratio (Blanco et al. 1984).   Although carbon stars are known in star clusters of the Magellanic Clouds (see e.g. Mould \\& Aaronson 1979), no late-type, luminous carbon stars are found in Galactic globular clusters despite there being a large contingent of massive, metal poor clusters. It may be inferred that the deep mixing on the AGB necessary to produce a carbon star is only possible for intermediate age stars.  Most studies of the C/M ratio in the Local Group rely on imaging through narrow band filters that measure the relative strength of CN and TiO bands (e.g., Cook et al. 1986).  In this context, the \\col~-color index is a measure for the relative strengths of CN and TiO bands and thus efficiently segregates M-stars with strong TiO and C-stars with strong CN bands. Quantitative analyses require additional color information (such as V$-$I), in order to construct distinctive two-color diagrams (TCDs). Specifically, C-Stars in M31 have been recorded (Brewer et al. 1995; Battinelli et al. 2003; Demers \\& Battinelli 2005; Battinelli \\& Demers 2005). The inferred relatively low C/M ratios of the order of 0.1 in the M31 disk and the inner halo (within 10 kpc) are then in agreement with the metal rich character of these populations.  For  the present work, we have examined spectra from a large survey of the M31 halo (Koch et al. 2008; hereafter K08), and we have synthesized the CN and TiO filters for our spectra.   From those we explore the C/M ratio in M31 halo fields, some of which are further than 100 kpc from the nucleus, but we do so without the completeness of a magnitude-limited photometric survey.  Even considering these limitations, we can statistically investigate the late-type stellar content in the Andromeda halo and connect occurrences of either population to the prevailing ages and metallicities out to large radii. This Paper is organized as follows: In \\textsection 2 we briefly summarize our observations and the standard reductions, while \\textsection 3 explains our definitions of the CN$-$TiO pseudo-filters. The  separation into C- and M-stars from these filters, possible biases, and the implications for M31's halo populations are discussed in \\textsection 4.  Finally, \\textsection 5  summarizes our findings. ",
        "conclusions": "We have analyzed a total of 1288 spectra of  giant stars in the M31 halo, with projected distances ranging from 9 to 160 kpc.  Using mock narrow-band filter convolutions, we have been able to plot these giants' populations in the CN$-$TiO vs. V$-$I diagram; from this we find that the number ratio C/M5+ is less than 0.1 in {\\em all halo fields}. We detect a total of only five C-star candidates. Therefore,  the outer halo exhibits the low C/M ratio that typifies other fields that have been targeted in M31 (Brewer et al. 1995; Nowotny et al. 2001;  Batinelli \\& Demers 2005). Moreover, a few of the M31 dSph satellites only show  null-detections or upper limits of C-star numbers (Groenewegen 2002; Harbeck et al. 2004). In contrast, disk fields closer than 30 kpc (Batinelli \\& Demers 2003) host far greater absolute numbers of carbon stars  that roughly trace the disk surface brightness.  If we assume the mean metallicity (and 1$\\sigma$ spread) of those fields, which harbor these candidates (at $<$[Fe/H]$>$=$-1.5\\pm0.41$ dex for the halo fields and $-1.24\\pm0.30$ dex for the Stream fields; K08), this low C/M5+ ratio in the  M31 halo fields falls $\\sim 1$ order of magnitude below the well known correlation of Local Group galaxies  (Fig.~8), which is defined over a broad range of morphological types by systems like the SMC, M33, and NGC 205 (Groenewegen 2002; Mouhcine \\& Lan\\c con 2003; Cioni \\& Habing 2003).  It is noteworthy, though, that the M31 dwarf galaxy satellites And~VII  and NGC~147 (Battinelli \\& Demers 2004a; Harbeck et al. 2004) fall near our M31 halo and Stream fields in this plot\\footnote{Nowotny et al. (2003) measure a C/M5+ ratio for NGC~147 that is higher by an order of magnitude than the value in Harbeck et al. (2004), which may be due to the different areas between these studies.}. \\begin{figure}[tb] \\begin{center} \\includegraphics[angle=0,width=1.1\\hsize]{f8.eps} \\end{center} \\caption{C/M ratios of Local Group galaxies from the compilation of Groenewegen 2002, supplemented with data for And VII, NGC~147, NGC~185, WLM, and AndII by Nowotny et al. (2003), Harbeck et al. (2004),  Battinelli \\& Demers (2004b), and Kerschbaum et al. (2004). Open circles indicate that only C/M0+ ratios are available in the literature, while filled symbols refer to C/M5+. Open triangles indicate C-star null-detections, or lower limits in the M-star population, in the Groenewegen, or  Harbeck et al.  data, respectively. Our measurement for the M31 halo fields is shown as a red filled circle at the mean metallicity of the fields surrounding the C-star candidates, while the solid square is the value from the two Stream fields. Black lines show the models of Mouhcine \\& Lan\\c con (2003) for different star formation histories.} \\end{figure} NGC~147 is an M31 dwarf elliptical satellite, devoid of dust and gas (Grebel et al. 2003), that exhibits a star formation history similar to the Magellanic Clouds or the Galactic dSphs Fornax and Leo~I (Harbeck et al. 2004) in that it is massive enough to form stars over most of its life time (M$_V$=15.1). The M31 dSph companion And~VII (Grebel \\& Guhathakurta 1999), on the other hand, is different from the remainder of the M31 dSphs in that it experienced prolonged star formation and it shows a predominant intermediate-age population of the order of 3--5 Gyr (Harbeck et al. 2004). These galaxies and our halo value also lie along lines of similar star formation histories, as indicated by the models by Mouhcine \\& Lan\\c con (2003; black lines in Fig.~8), in particular those that are characterized by a higher dominance of the younger populations (solid and long-dashed lines). The agreement of our C/M ratio in the M31 halo with those of NGC 147 and And VII is then consistent with the interpretation that such (dSph and dE) galaxies are likely connected to some of the distinct merger features in M31;  the structures and mean stellar metallicites of these features (Ferguson et al. 2002; Harbeck et al. 2004) are further consistent with this picture.  Owing to the pre-selection criteria for the present spectroscopic sample, we cannot exclude the possibility that we have underestimated the number of M-type AGB stars (Sect.~4.3.3). Under realistic assumptions, however, we showed that the C/M5+ ratio is unlikely to be significantly larger than 1.0, unless we missed a substantial C-star population in our sample. This value is then still lower than expected from a SMC-type population and rather similar to the other M31 dwarf elliptical satellite, NGC~185. In order to reach a C/M ratio as high as  predicted by the models of Mouhcine \\& Lan\\c con (2003) at the given halo metallicity, M31's AGB-fraction, $f_{\\rm AGB}$, would have to be remarkably low, of the order of $\\la2$\\%. Moreover, only if our survey was 100\\% incomplete at the bright end could we attain SMC-like C/M ratios of $\\sim$3 (following our methodology in Sect.~4.3.3.).  Prior surveys of the M31 disk, extending to 40 kpc where the halo likely dominates, find compelling evidence that the C/M5+ fraction actually drops at the boundary of the disk (Demers \\& Battinelli 2005).  The G1 cloud of Ferguson et al. (2002), once thought to be the likely remnant of a dwarf spheroidal satellite, exhibits a C/M5+ ratio typical of our halo fields. Deep HST observations (Brown et al. 2003, 2006, 2007, 2008; Richardson et al. 2009), find a widespread intermediate age population in the M31 halo, while Kalirai et al. (2006) and K08 both find evidence of a metal poor outer halo in M31.  Low metallicity and a strong intermediate age population would appear to be strong factors favoring a large carbon star population. Indeed, this appears to be the case for the Galactic dwarf spheroidals (e.g., Azzopardi et al. 1985,1986; Groenewegen 1999; Demers \\& Battinelli 2002). However, the Galactic halo does not appear to be comprised of the debris of such systems with major C-star populations.  Even if the parent systems disrupted, the carbon-rich stars should have survived in the halo field (see Ibata et al. 2001b for the Sgr dwarf).  We are left to conclude that two factors may decrease the expected numbers of carbon stars.  First, the age distribution in the outer halo may be tilted toward older stars.   However, in the G1 clump (at a projected distance of 34 kpc), there is evidence that the field has a strong intermediate-age component (Rich et al. 2004; Faria et al. 2007) and also an HST field at 35 kpc was shown to contain $\\sim$1/3 of stars younger than 10 Gyr (Brown et al. 2008). The second explanation would argue that the halo is too metal rich to support a substantial population of C-stars. However, both Kalirai et al. (2006) and K08 agree that [Fe/H]$\\sim -1.5$ and lower at large distances, a value that is low enough to support carbon star production; moreover, there is a dispersion in metallicity with a metal rich tail high enough to spawn M giants. Considering that the origin of this metal rich tail must arise in mergers, the absence of a strong carbon star population in the outer halo must be considered to be noteworthy. Therefore, the low C/M ratio in the M31 field is inconsistent with a presumed disk-like or dwarf galaxy origin (with properties like those indicated in Fig.~8) of those halo fields over a broad radial extent. Although the total numbers of stars are low in the halo, the C/M {\\em ratio} should be preserved.  The cool M-giants with detectable TiO, on the other hand, are found throughout the halo, further strengthening the case for a widespread population with [Fe/H]$>-1$, even for the most distant giants. Further, stars with measureable TiO follow the same kinematics as is found for all giants with measurable radial velocities. Both regarding all M stars and with respect to the C/M ratio (as measured by our very small numbers) the giant stream appears to harbor no special stellar population, in contrast to what has been found for the Sgr dwarf spheroidal galaxy (Ibata et al. 2001b), which does have a concentration of carbon stars.  With the explosive discovery of numerous streams and stellar systems in the halo, the picture of the halo population comprised nearly entirely of stellar subsystems has become the canon.  Yet well constructed surveys of our Milky Way halo (e.g. Mauron et al. 2008) find only a handful of luminous carbon stars.  If indeed halos are dominated by debris, as cannot ruled out from a dynamical standpoint, they are not similar in population to the SMC, NGC 6822, or M33.  If we appeal to the disruption of a disk-like galaxy to have created the giant stream (e.g., bearing on the model of Fardal et al. 2008), that progenitor was especially devoid of carbon stars and more closely resembled NGC 147 rather than M33. In the broader picture, the late-type stellar content of the M31 halo must be judged to differ from that of the dwarf galaxies that one might propose to have contributed its stars."
    },
    "1004/1004.4720_arXiv.txt": {
        "abstract": "{The pulse profiles of the transient Be/X-ray binary \\exo\\ show strong dependence on energy, as well as on its luminosity state, and are asymmetric in shape.} {We want to identify the emission components of the two magnetic poles in the pulsed emission to understand the geometry of the neutron star and its beam pattern.} {We utilize a pulse-profile decomposition method that enables us to find two symmetric pulse profiles from the magnetic poles of the neutron star. The symmetry characteristics of these single-pole pulse profiles give information about the position of the magnetic poles of the neutron star relative to its rotation axis.} {We find a possible geometry for the neutron star in \\exo\\ through the decomposition of the pulse profiles, which suggests that one pole gets closer to the line of sight than the other and that, during the revolution of the neutron star, both poles disappear behind the horizon for a short period of time. A considerable fraction of the emission arises from a halo while the pole is facing the observer and from the accretion stream of the other pole while it is behind the neutron star, but the gravitational line bending makes the emission visible to us. } {} ",
        "introduction": "\\object{EXO 2030+375} is an accreting X-ray pulsar with a pulsation period of $\\sim$42~s, which was discovered with \\exosat\\ in 1985 during a giant outburst \\citep{1989ApJ...338..359P}. A B0 Ve star was found as its counterpart in follow-up observations in the optical and infrared bands \\citep{1988A&A...202...81J,1987A&A...182L..55M,1988MNRAS.232..865C}. During the giant outburst, \\exo\\ showed a spin-up of $-P/\\dot{P} \\approx 30$~yr \\citep{1989ApJ...338..359P} and quasi-periodic oscillations with a frequency of $\\sim$0.2~Hz \\citep{1989ApJ...346..906A} interpreted as caused by the formation of an accretion disk. Detailed analyses have shown that its rate of pulse-period change $\\dot{P}$, energy spectrum, and pulse profile are strongly luminosity dependent \\citep{1989ApJ...338..373P,1989ApJ...338..359P,1993ApJ...414..302R}. The orbital period is 46 days \\citep{2002ApJ...570..287W}, and a normal outburst has been detected for nearly every periastron passage since 1991 \\citep{2005ApJ...620L..99W}. In 2006, \\exo\\ underwent the first giant outburst since its discovery in 1985 \\citep{2006ATel..843....1C,2006ATel..861....1K,2006ATel..868....1M}, during which it reached a maximum luminosity of $L_{1 - 20~{\\rm keV}} \\approx 1.2 \\times 10^{38}$~erg~s$^{-1}$ \\citep{2008A&A...491..833K} and again showed a strong spin-up. Rossi X-ray Timing Explorer (\\rxte) monitored \\exo\\ extensively during the 2006 giant outburst \\citep{2008ApJ...678.1263W}. The source was also observed by the INTErnational Gamma Ray Astrophysics Laboratory \\citep[\\integral,][]{2003A&A...411L...1W} and \\swift\\ \\citep{2004ApJ...611.1005G}. The spectra indicate a cyclotron absorption line \\citep{2007A&A...464L..45K,2008ApJ...678.1263W}. \\citet{2008A&A...491..833K} have shown that the spectrum of \\exo\\ changes with pulse phase, suggesting a fan beam geometry during the maximum, while towards the end of the giant outburst, it changes to a combination of a fan beam and a pencil beam.  In X-ray pulsars, a neutron star accretes matter from a companion star via stellar wind or Roche lobe overflow. The accreted matter is channeled along the field lines of the strong magnetic field of the neutron star onto the magnetic poles. X-ray emission from the neutron star is produced in regions around the two magnetic poles. As the magnetic dipole axis is most likely inclined against the rotation axis of the neutron star, a distant observer sees pulsed emission. X-ray pulsars exhibit a wide variety of pulse shapes that differ from source to source. Generally, high-energy pulses have simpler shapes than low-energy pulses \\citep[][and references therein]{1983ApJ...270..711W,1989ESASP.296...57F,1997ApJS..113..367B}. If one assumes an axially symmetric geometry for the two emission regions of the neutron star in a dipole configuration, the observed pulse profile should be symmetric. However, the observed pulse profiles typically show an asymmetry. To explain the asymmetric shape of the total pulse profile, a distorted magnetic dipole field in which the two magnetic poles are not located opposite each other have been discussed \\citep{1989ApJ...338..373P,1991MNRAS.251..203L,1993ApJ...406..185R,1995ApJ...444..405B}. \\citet{1995ApJ...450..763K} shows that, starting from the observed, asymmetric pulse profile, it is possible to disentangle the contribution of the two emission regions of the neutron star. Once the pulsed emission from each of the poles has been obtained, one can derive the geometry of the neutron star. This again allows us to construct the beam pattern, i.e., the flux distribution from one emission region. Using this pulse-profile decomposition method, \\citet{1996ApJ...467..794K} have analyzed the pulse profiles of Cen~X-3 and find indications of both pencil and fan beam. In the case of Her~X-1, the results of the pulse-profile decomposition by \\citet{2000ApJ...529..968B} have not only shed light on the beam pattern of the magnetic poles, but have also confirmed that a warped and tilted accretion disk attenuates the emission from one pole of the neutron star. For A~0535+26, the reconstructed beam pattern suggests that the emission comes from a hollow column plus a halo of scattered radiation on the neutron star surface \\citep{Caballero2010}.  In this paper we present the analysis of the energy-resolved pulse profiles of \\exo\\ utilizing the decomposition method developed by \\citet{1995ApJ...450..763K}. Section \\ref{data} gives an overview of the data used for our analysis and Sect.\\,\\ref{analysis} describes the analysis and the results obtained with the pulse-profile decomposition method. The results are discussed in Sect.\\,\\ref{discussion}. Section \\ref{summary} summarizes the possible geometry of the neutron star and the origin of the observed emission. ",
        "conclusions": "\\label{discussion}  In this section we want to further examine the results of the decomposition method and present the possible geometry of the neutron star.  \\subsection{Disentangling the emission components}  \\begin{figure} \\centering \\includegraphics[width=0.49\\textwidth,clip=]{onespp.A1.comp.minus.eps} \\caption{ Decomposed single-pole pulse profiles of the data taken near the maximum of the giant outburst for the lowest energy band (same as the upper left diagram of the upper [solution 1] panels in Fig.\\,\\ref{sppbp}) with proposed emission components derived from the beam patterns of the {\\tt minus} solution (Fig.\\,\\ref{bptheta1}, lower panel). We suggest that the emission from the first pole consists of emission from the halo (diagonally striped component) and from the accretion column (vertically striped). The dominant emission from the second pole seems to have its origin in the scattered photons in the upper accretion stream, which is visible while the second pole is behind the horizon of the neutron star due to gravitational light bending, while there is also a contribution from the accretion column (vertically striped). In this solution, the first pole gets closer to the line of sight than the second. } \\label{sppcomp} \\end{figure}  For solution 1, there is no overlap between the two beam patterns obtained from the single-pole pulse profiles as shown in Sect.\\,\\ref{asymbp}. It means that only a part of the emission is seen from each pole during the revolution of the neutron star. By putting the two visible parts together, we obtain the total beam pattern of the emission around one magnetic pole. In doing so, we make the assumption that the two magnetic poles have the same emission pattern. The {\\tt plus} solution (Fig.\\,\\ref{bptheta1}, upper panel) can be described as a composition of a forward directed emission (towards 0\\degr) that is more pronounced at higher energies and an extended, relatively soft emission ($\\sim$80\\degr\\ -- 180\\degr). The harder, forward-directed emission indicates a pencil beam. The relative flux of the softer emission component is higher in the data from the observation performed near the maximum compared to the one during the decay and can be interpreted as a fan beam, in agreement with the largely accepted picture that an optically thick accretion column is formed during the giant outburst \\citep[][and references therein]{1983ApJ...270..711W}. The {\\tt minus} solution has a soft emission at $\\theta \\la$ 60\\degr and a harder emission at $\\theta \\ga$ 120\\degr\\ that increases for larger $\\theta$. As newest calculations by Kraus et al.\\ (in prep.) have shown, reprocessing of photons in the upper accretion stream creates a significant emission component that dominates the beam pattern at higher energies and can be observed while the pole, hence the accretion column, is on the other side of the neutron star (`anti-pencil'). In addition, the emission from a halo that is formed by scattered photons at the bottom of the accretion column dominates the beam pattern at lower energies and at lower $\\theta$ and is stronger in the data near the maximum of the giant outburst.  Figure \\ref{sppcomp} shows how the emission from the two poles contributes to the pulse profiles for the decompositions of the low energy band data taken at the maximum of the giant outburst in the case of the {\\tt minus} solution. Let us assume that the first magnetic pole with the polar angle $\\Theta_{1}$ gets closer to the line of sight than the second magnetic pole, i.e., $\\Theta_{0}-\\Theta_{1} \\le \\Theta_{2}-\\Theta_{0}$. At phase 0.0, both poles are right behind the horizon of the neutron star: the first pole is going to reappear, the second pole is turning farther away from the observer. At phase $\\sim0.1$ the first pole becomes visible, and from then on the emission from the halo of the first pole makes the largest contribution until the pole disappears behind the horizon at $\\sim0.6$. When the first pole, which is closer to the observer's line of sight, is right at the horizon of the neutron star and its accretion column is seen from the side, a minimum is likely to be observed in the pulse profile (see also Sect.\\,\\ref{emorigin}). At the major maximum at phase 0.15 -- 0.25, the second magnetic pole is behind the horizon of the neutron star and the scattered and gravitationally bent photons from the upper accretion stream cause the pronounced increase in flux. The second pole that is rotating on a circle farther away from the line of sight than the first pole comes back to the front side at about phase 0.45 and is closest to the observer at about phase 0.7. However, since the line of sight is closer in latitude to the first pole than to the second pole, the observer never gets as close to the surface normal of the second pole as to that of the first pole. The main emission seen from the second pole while it is on this side of the horizon comes directly from the accretion column.  \\citet{1989ApJ...338..373P} have modeled the luminosity dependent pulse profiles from the first observed giant outburst in 1985 by assuming a fan beam and a pencil beam component for the two magnetic poles based on a model by \\citet{1981A&A...102...97W}. They obtain a fit for all ten pulse profiles for different luminosities ranging from 0.1 -- 10.0 $\\times 10^{37}$~erg~s$^{-1}$ with some residuals and find that the fan beam mainly produces the major peak at phase 0.2, whereas the peak at phase 0.95 can be ascribed to a pencil beam. The way they have chosen phase 0.0 is different than in our work, resulting in a shift by about 0.4. The best-fit parameters of \\citet{1989ApJ...338..373P} correspond to $\\Theta_{1}$ = 70\\degr, $\\Theta_{2}$ = 110\\degr, and $\\delta$ = 70\\degr, assuming $\\Theta_{0}$ = 25\\degr. This result is consistent with our solution 1 for the case of $\\Theta_{0}$ = 25\\degr\\ (see Fig.\\,\\ref{geometry}).  \\subsection{Comparison to phase-resolved spectral analysis}  \\citet{2008A&A...491..833K} analyzed the same \\integral\\ data as used in this work and performed pulse phase resolved spectroscopy. For the data of the maximum, they find that the spectrum is harder at the main peak (phase 0.2). They conclude that the observer might be seeing Compton scattered photons from the optically thick accretion column. In the data taken at the end of the decay, a new peak appears at about phase 0.95. At this point, the spectrum is again harder, leading to the conclusion that the line of sight of the observer is closest to the magnetic field lines. This picture is in good agreement with the solution 1 of the decomposition, in which the emission of the upper accretion stream of the second pole, which is visible owing to gravitational light bending, has its maximum at phase $\\sim$0.2 while the halo of the first pole also has substantial emission, thus contributing significantly to the main peak. At phase $\\sim$0.95 the scattered emission from the upper accretion stream of the second pole starts to increase, while the emission from the accretion column of the first pole also contributes to the observed emission.  \\subsection{Intrinsic beam patterns}  \\begin{figure*} \\centering \\includegraphics[width=0.4\\textwidth,angle=270,clip=]{polar.A1.bw.eps_pages3} \\includegraphics[width=0.4\\textwidth,angle=270,clip=]{polar.A1.bw.eps_pages4} \\caption{Polar diagrams of the intrinsic beam pattern for the {\\tt minus} solution of solution 1 for the same observations as shown in Fig.\\,\\ref{bptheta1}, for the softest and the hardest bands for an assumed $\\Theta_{0}$ = 50\\degr. Dotted lines are drawn to clearly show the range in which the emission is visible for the observer. } \\label{bpintrA1} \\end{figure*}   The observed asymptotic beam pattern differs from the intrinsic local beam pattern of the emission region because of relativistic light deflection near the neutron star. Making assumptions on the radius $r$ and mass $m$ of the neutron star, the asymptotic angle $\\theta$ can be transformed into the intrinsic angle $\\vartheta$ relative to the surface normal at the first magnetic pole to describe the local emission pattern. We use the canonical values $r = 10$~km and $m = 1.4~M_{\\sun}$, thus the ratio between $r$ and the Schwarzschild radius $r_{\\rm S}$ is $r/r_{\\rm S}$ = 2.4. The intrinsic beam patterns of the {\\tt minus} solution are shown in Fig.\\,\\ref{bpintrA1} in polar diagrams. As can be seen in these polar diagrams, the range for the visible angle is $\\sim$10\\degr\\ -- 105\\degr\\ if $\\Theta_{0}$ = 50\\degr\\ is assumed.  \\subsection{Possible origin of the emission}\\label{emorigin}  In X-ray pulsars, a strong magnetic field funnels matter accreted from the companion star onto the polar caps. The infalling particles deposit their energy in the atmosphere of the neutron star. Heating due to accretion is balanced by radiative cooling through bremsstrahlung emission, Compton scattering, and the so-called cyclotron radiation after collisional excitation of electrons. In low-luminosity states, the accreted matter is decelerated in the atmosphere of the neutron star by Coulomb scattering. Emission can be seen as pencil beams from the hot spots at polar caps \\citep{1981A&A...102...97W}. In high-luminosity states, however, the accretion rate is higher, and an accretion column is believed to form. Plasma is decelerated by radiation pressure in the column, and a radiative shock forms above the neutron star surface \\citep{1973Natur.246....1D,1976MNRAS.175..395B}. Above the radiative shock, plasma is in free fall, while below the discontinuity, there is a region of nearly stagnant plasma from which photons escape from the sides of the column in a fan beam.  In addition, a luminous halo might form around the accretion column \\citep{1973ApJ...179..585D,1988SvAL...14..390L} and radiation from the polar cap and the halo can also be scattered in the upper accretion stream \\citep{1985A&A...144..485S,1991ApJ...369..179B}. As the accretion stream is delimited by magnetic field lines of the neutron star, it opens up wide far from the neutron star. Therefore, the emission from the upper part of the accretion stream can dominate the observed flux and can also screen the polar caps and the halo. Due to relativistic light deflection, emission from one pole can be deflected to the antipodal direction \\citep{1983ApJ...274..846P,1988ApJ...325..207R,1995MNRAS.277.1177L,2001ApJ...563..289K}. All these effects can modify the local beam pattern and thus have an effect on the pulse profile of the X-ray pulsar. \\citet{1981A&A...102...97W} have modeled the pulse profiles of X-ray pulsars by assuming hot spots and fan beams, although \\exo\\ is not included in their list of selected X-ray pulsars. They show that, in general, sharp minima can be seen in the pulse profiles at the moment when the fan beam rotates behind the horizon of the neutron star and when it reappears on the side of the neutron star facing the observer. This is exactly what seems to be happening in \\exo\\ at phases 0.1 and 0.6 as the decomposition method has shown.  Kraus et al.\\ (\\citeyear{2003ApJ...590..424K}; in prep.) have modeled beam patterns and pulse profiles for medium-luminosity X-ray pulsars assuming an accretion column with energy-dependent local beaming of radiation, a luminous halo formed by illumination of the neutron star surface, and magnetic scattering in the upper accretion stream. They show that the emission from the upper accretion stream can dominate the local beam pattern for $\\theta >$ 120\\degr. Halo emission has a maximum at $\\theta \\approx$ 30\\degr\\ -- 60\\degr, with significant contribution to the beam pattern for photon energies below $\\sim$5~keV. If one assumes isotropic emission from the accretion column below the discontinuity, the halo emission becomes comparable or negligible relative to the column emission at $\\ga$10~keV. However, if the radiation below the shock is beamed downwards, the halo remains more luminous than the accretion column and the upper stream up to $\\sim$30~keV. The decomposition has shown that the beam patterns show dominant soft emission for $\\theta <$ 60\\degr, which can be interpreted as halo emission, while the emission $\\theta >$ 150\\degr\\ most likely arises from the upper accretion stream (see Fig.\\,\\ref{bptheta1}, lower panel)."
    },
    "1004/1004.3877_arXiv.txt": {
        "abstract": "The pair decay width of the first excited $0^+$ state in $^{12}$C (the Hoyle state) is deduced from a novel analysis of the world data on inelastic electron scattering covering a wide momentum transfer range, thereby resolving previous discrepancies. The extracted value $\\Gamma_{\\pi} = (62.3 \\pm 2.0)$~$\\mu$eV is independently confirmed by new data at low momentum transfers measured at the \\mbox{S-DALINAC} and reduces the uncertainty of the literature values by more than a factor of three. A precise knowledge of $\\Gamma_\\pi$ is mandatory for quantitative studies of some key issues in the modeling of supernovae and of asymptotic giant branch stars, the most likely site of the slow-neutron nucleosynthesis process. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.2633.txt": {
        "abstract": "Results of modeling the spectra of two supernovae SN\\,2008D and SN\\,2006aj related to the \\linebreak X-ray flash XRF\\,080109 and gamma-ray burst GRB\\,/\\,XRF\\,060218, respectively, are studied. The spectra were obtained with the 6-meter BTA telescope of the Special Astrophysical Observatory of the Russian Academy of Sciences in 6.48 and 27.61 days after the explosion of SN\\,2008D, and in 2.55 and 3.55 days after the explosion of SN\\,2006aj. The spectra were interpreted in the Sobolev approximation with the \\texttt{SYNOW} code. An assumption about the presence of envelopes around the progenitor stars is confirmed by an agreement between the velocities of lines interpreted as hydrogen and helium, and the empiric power-law velocity drop with time for the envelopes of classic core-collapse supernovae. Detection of a P Cyg profile of the H$\\beta$ line in the spectra of optical afterglows of GRBs can be a determinative argument in favor of \\mbox{this hypothesis.} ",
        "introduction": "The modern classification of core-collapse supernovae includes several subtypes. If the spectra demonstrate hydrogen lines, then the supernova belongs to the II type. Its light curve shape is determined by the hydrogen recombination features in the envelopes. If a star loses its hydrogen envelope, but the helium shell is still present in the spectrum, then the star is classified into the Ib type. The impacts of hydrogen recombination are absent in its light curve, and its shape is determined by the heating of the emitting matter as a result of the \\linebreak $^{56}{\\rm Ni} \\rightarrow ~^{56}{\\rm Co} \\rightarrow ~^{56}{\\rm Fe}$ decay, depending on the explosion energy. If the helium envelope is lost and the spectra show only the traces of heavier elements \\mbox{(O, Mg, Si, S, Ca, Fe),} then the supernova is classified into the Ic type. Different subtypes of supernovae vary by different explosion energies, and, consequently, by the velocities of matter expansion in the envelopes.  Since in some cases the classification depends on the phase of the supernova at the moment of observation, some intermediate types are introduced: Ib/c (appearance of weak helium lines in late spectra), IIb (disappearance of hydrogen lines in late spectra). The classification of core-collapse supernovae can be conceived as a manifold of intermediate subtypes, depending continuously on the initial parameters of the progenitor star.  %some events show ambiguous behavior at different stages of %evolution forming intermediate classes Ib/c (appearance of weak %helium lines in later spectra), IIb (disappearance of helium lines %in later spectra), the supernovae classification could be %conceived with more diffuse, smooth borders depending continuously %on initial parameters of a progenitor star.  This paper discusses the manifestations of stellar-wind envelopes around the progenitor stars of an \\mbox{X-ray} flash XRF\\,080109 and a gamma-ray burst with a strong X-ray component GRB / XRF\\,060218. Some signs of supernovae (SNe) were noticed in the spectra and light curves of both events, which allowed studying these phenomena from the very start of the explosion (unlike most SNe that are detected at the moments close to the ``classical'' maximum of the light curve). The study of this kind of phenomena makes it possible to approach the solution to the supernova explosion mechanism and the origin of gamma-ray bursts.   On January 9.57, 2009 (UT) the X-ray telescope XRT on board of the \\textit{Swift} space observatory registered an X-ray flash XRF\\,080109. The X-ray component was detectable for about 15 \\mbox{minutes \\cite{Chevalier_Fransson2008:Moskvitin_n, Soderberg2008:Moskvitin_n}.} The Burst Alert Telescope (BAT) did not detect any gamma-ray quanta. The passage of a shock wave through the envelope was observed in multi-color light curves during several days \\cite{Soderberg2008:Moskvitin_n, Modjaz2008:Moskvitin_n}. The afterglow light curves and spectra distinctly showed signs of the supernova, designated as SN\\,2008D.  Almost two years earlier (on February 18.149, 2006) the same space platform \\textit{Swift} detected a peculiar gamma-ray burst with the SN signs in the afterglow spectra and light curves of GRB\\,/\\,XRF\\,060218\\,/ \\,SN\\,2006aj. It is the first phenomenon for which we observed the passage of a shock wave through the wind envelope surrounding a massive SN progenitor star, which manifested itself as a thermal component in the X-ray spectrum observed during the first 2 hours, and then as a powerful UV burst with the maximum observed in 11 hours after the burst ~\\mbox{\\cite{Blustin2007:Moskvitin_n, Campana2006:Moskvitin_n}}. The effect of the shock wave passage is seen in the optical light curve as a relatively short peak several days of duration \\cite{Campana2006:Moskvitin_n, Sonbas2008:Moskvitin_n}.  %(see Fig.2 in \\cite{Campana2006} and Fig.2 in our previous paper %\\cite{Sonbas2008}).  In both cases the moment of the X-ray flash and/or gamma-ray burst registration was taken as the beginning of the SN explosion.  The structure of the paper is as follows: spectral observations and preliminary data processing are described in section \\ref{SpecObs:Moskvitin_n}; the interpretation of the obtained spectra with the \\texttt{SYNOW} code is laid out in \\ref{SYNOW:Moskvitin_n}; a comparison of the obtained measurements of the photosphere and envelope expansion velocities with the analogous velocities of SNe that did not show any relation with gamma-ray bursts or X-ray flashes is in section \\ref{Vel_phot_env:Moskvitin_n}; our conclusions are discussed in section \\ref{results:Moskvitin_n}. ",
        "conclusions": "\\label{results:Moskvitin_n}  An agreement between the velocities of photosphere expansion and those of the helium/hydrogen envelopes of the SNe associated with the X-ray flash XRF\\,080109 and the gamma-ray burst GBR\\,060218 argues for similar explosion dynamics of gamma-ray bursts and X-ray flashes.  A comparison with analogous velocities of classical type Ib supernovae not accompanied by gamma- or X-ray bursts may indicate common properties of the progenitors of these two classes of phenomena, and similar explosion mechanisms.  Previously, an idea about the presence of hydrogen in the spectra of type Ib and Ic supernovae was suggested ~\\cite{Branch2006:Moskvitin_n, Elmhamdi2006:Moskvitin_n}. The idea was confirmed in our paper ~\\cite{Sonbas2008:Moskvitin_n} dedicated to the study of early spectra of SN\\,2006aj.  Branch et al. \\cite{Branch2006:Moskvitin_n} list the main candidates for the description of a spectral feature near 6200\\,\\AA{} observed in the early spectra of supernovae. Apart from the hydrogen line H$\\alpha~ \\lambda6563$ these are: Si\\,II $\\lambda6355$ spaced from H$\\alpha$ at 9500 km/s, Ne\\,I $\\lambda6402$ at 7360 km/s, and only C\\,II $\\lambda6580$ is spaced from the hydrogen line at 777 km/s, which can be within velocity measurement errors. Candidate velocities must agree with the photosphere velocity determined from the heaviest ions (e.g., from the ionized iron). Reasoning from the results of modeling the SN\\,2006aj and SN\\,2008D spectra, it can be concluded that the most possible sources of absorption near $\\lambda$6200 are H$\\alpha$~ $\\lambda6563$ and C\\,II $\\lambda6580$. To choose between these candidates we compared the velocity variability of hydrogen envelopes of the two SNe under study with the classical supernovae. The results of comparison of the velocity evolutions of different supernovae argue for the existence of hydrogen-containing envelopes around massive progenitor stars that appear before the explosion due to more or less powerful stellar winds.  Detection of P Cyg profile absorption components of other Balmer lines, in particular H$\\beta$ $\\lambda4861$, can become a decisive argument in favor of hydrogen presence in the envelope. Since the optical depth of H$\\beta$ line is smaller than that of H$\\alpha$, there arises a problem of its detection in a noisy spectrum among numerous lines of heavy metals, particularly, Fe\\,II. Detailed long-term observations of the evolution of subtle details in the spectra obtained with a high signal/noise ratio can solve this problem.   % Nevertheless, the identi?cation of the 6200 A absorption is plagued by ambiguities. The % strongest line of Ne I, ?6402, is about 2200 km s?1 to the red of Si II ?6355, and H-alpha % ?6563 is only about 800 km s?1 to the blue of C II ?6580. Because these four ions, each of % which could appear in supernova spectra (although Ne I probably would require nonthermal % excitation), have their strongest optical lines to the red, but not too far to the red, of the"
    },
    "1004/1004.5383_arXiv.txt": {
        "abstract": "We report Westerbork Synthesis Radio Telescope and Arecibo Telescope observations of the redshifted satellite OH~18cm lines at $z \\sim 0.247$ towards PKS~1413+135. The ``conjugate'' nature of these lines, with one line in emission and the other in absorption, but with the same shape, implies that the lines arise in the same gas. The satellite OH~18cm line frequencies also have different dependences on the fine structure constant $\\alpha$, the proton-electron mass ratio $\\mu = m_p/m_e$, and the proton gyromagnetic ratio $g_p$.  Comparisons between the satellite line redshifts in conjugate systems can hence be used to probe changes in $\\alpha$, $\\mu$, and $g_p$, with few systematic effects.  The technique yields the expected null result when applied to Cen.A, a nearby conjugate satellite system.  For the $z \\sim 0.247$ system towards PKS~1413+135, we find, on combining results from the two telescopes, that $\\left[ \\Delta G/G \\right] = (-1.18 \\pm 0.46) \\times 10^{-5}$ (weighted mean), where $G = g_p \\left[ \\mu \\alpha^2 \\right]^{1.85}$; this is tentative evidence (with $2.6 \\sigma$ significance, or at 99.1\\% confidence) for a smaller value of $\\alpha$, $\\mu$, and/or $g_p$ at $z \\sim 0.247$, i.e. at a lookback time of $\\sim 2.9$~Gyrs. If we assume that the dominant change is in $\\alpha$, this implies $\\left[ \\Delta \\alpha / \\alpha \\right] = (-3.1 \\pm 1.2) \\times 10^{-6}$. We find no evidence that the observed offset might be produced by systematic effects, either due to observational or analysis procedures, or local conditions in the molecular cloud. ",
        "introduction": "\\label{sec:intro}  Most modern higher-dimensional theories venturing beyond the standard model of particle physics contain the fairly generic feature that the low-energy fundamental constants should vary with time. The detection of such changes opens up an avenue to probe basic physics, especially important because other predictions of these models tend to lie at unattainably high energies ($\\gtrsim 10^{19}$~GeV).  Tests of such low-energy predictions of these models, such as changes in the constants, violation of the equivalence principles, etc, may provide the only means of distinguishing between different unified theories and are hence of much importance in physics.  Astrophysical techniques can be used to probe changes in fundamental constants like the fine structure constant $\\alpha$, the proton-electron mass ratio $\\mu \\equiv m_p/m_e$ and the proton gyromagnetic ratio $g_p$, over a large fraction of the age of the Universe (e.g. \\citealp{savedoff56,wolfe76,thompson75,dzuba99,chengalur03}). Indeed, \\citet{murphy04} applied one such technique (the ``many-multiplet method''; \\citealp{dzuba99}) to High Resolution Echelle Spectrograph (HIRES) data from the Keck telescope to obtain $\\dal = (-5.7 \\pm 1.1) \\times 10^{-6}$ from 143~absorbers at $0.2 < z < 4.2$, i.e. suggesting a smaller value of $\\alpha$ at earlier times (see also \\citealp{murphy03}). This result has not so far been confirmed with independent data on smaller samples from the Ultraviolet Echelle Spectrograph (UVES) on the Very Large Telescope (e.g. \\citealp{molaro08,srianand07b,murphy08b}. More recently, \\citet{king08} used redshifted molecular hydrogen (H$_2$) lines to place strong constraints on changes in $\\mu$, obtaining $\\dmu = (-2.6 \\pm 3.0) \\times 10^{-6}$ from three absorbers at $z \\sim 2.8$. Note that these results are all based on optical spectra, where wavelength calibration, line blending, intrinsic velocity offsets, isotopic abundances, interloping absorbers, etc, are all possible sources of systematic error, (e.g. \\citealp{murphy03,griest10}). Further, the quoted errors in the result of \\citet{murphy04} do not include effects from systematic distortions in the HIRES wavelength calibration \\citep{griest10}; similar, but smaller, distortions have also been found in the wavelength scale of the VLT-UVES spectrograph \\citep{centurion09,whitmore10}.  Given the possibility that underestimated or unknown systematics might dominate the errors from a given technique, it is important that independent techniques, with entirely different systematics, be used to test for changes in the constants. It is also crucial to probe evolution at all timescales, as the timescales of the putative changes are entirely unknown. Redshifted radio OH lines provide an independent approach to study changes in $\\alpha$, $\\mu$ and $g_p$ \\citep{chengalur03,darling03,kanekar04a}. In rare cases, the satellite OH~18cm lines (at rest frequencies of 1612.230825 (15)~MHz and 1720.529887 (10)~MHz; \\citealp{lev06}) are ``conjugate'' to each other; i.e. the lines have the same shapes, but one line is in absorption and the other in emission. This masing effect arises due to the quantum mechanical selection rules for decay routes to the {$2\\Pi_{3/2}$ (J = 3/2)}~OH ground state, after the molecules have been pumped to higher excited rotational states \\citep{elitzur92}. Crucially, the conjugate behavior {\\it ensures that the satellite lines arise from the same gas}. These lines are thus well-suited for measuring changes in $\\alpha$, $\\mu$ and $g_p$ between the source redshift and today, as systematic velocity offsets between the lines are ruled out by the maser mechanism. Any observed difference between the line redshifts must then arise due to a change in one or more of the above constants \\citep{kanekar08b}.  Only two conjugate OH systems are known at cosmological distances, at $z \\sim 0.247$ towards \\pks\\ \\citep{kanekar04b,darling04} and at $z \\sim 0.765$ towards \\pmn\\ \\citep{kanekar05}. We report here on deep Westerbork Synthesis Radio Telescope (WSRT) and Arecibo Telescope (AO) OH observations of \\pks, that yield tentative evidence for changes in $\\alpha$, $\\mu$ and/or $g_p$ over a lookback time of $\\sim 2.9$~Gyr.    \\begin{figure*} \\epsfig{file=fig1a.eps,width=3.5in,height=3.5in} \\epsfig{file=fig1b.eps,width=3.5in,height=3.5in} \\caption{Top panels: WSRT and AO satellite OH spectra towards \\pks, with optical depth ($1000 \\times \\tau$) plotted against heliocentric velocity, in km/s, relative to $z = 0.24671$. The 1612 and 1720~MHz spectra are plotted using open circles and triangles, respectively, with the 1720~MHz spectrum flipped in sign. Bottom panels: The sum of the 1612 and 1720~MHz optical depth spectra from each telescope. } \\label{fig:fig1} \\end{figure*} ",
        "conclusions": "\\label{sec:which}  If the observed ($\\sim 2.6 \\sigma$ significance) offset between the satellite lines is not caused by a ``hidden'' spectral component (or RFI) below our sensitivity threshold, it could be evidence for changes in one or  more of $\\alpha$, $\\mu$, or $g_p$. Note that the conjugate satellite technique is directly sensitive to changes in $G \\equiv g_p \\left[ \\mu \\alpha^2 \\right]^{1.85}$, and cannot provide additional information on which constant is changing. Our present result $\\Delta G/G = (-1.18 \\pm 0.46) \\times 10^{-5}$ implies the relation $3.7\\dal + 1.85\\dmu + [\\Delta g_p/g_p] = (-1.18 \\pm 0.46) \\times 10^{-5}$, suggesting that one or more of $\\alpha$, $\\mu$ and $g_p$ had smaller values at $z = 0.247$ than at the present epoch. The limiting cases, assuming that only one of $\\alpha$, $\\mu$, and $g_p$ changes, are $\\dal = (-3.1 \\pm 1.2) \\times 10^{-6}$ (for $\\dal >> \\dmu$, $[\\Delta g_p/g_p]$), $\\dmu = (-6.2 \\pm 2.4) \\times 10^{-6}$  (for $\\dmu >> \\dal$, $[\\Delta g_p/g_p]$), and $[\\Delta g_p/g_p] = (-1.18 \\pm 0.46) \\times 10^{-5}$ (for $[\\Delta g_p/g_p] >> \\dal$, $\\dmu$).  Changes in $g_p$ are expected to be far smaller than those in $\\mu$ and $\\alpha$ \\citep{langacker02}.  Further, a stringent constraint on changes in $\\mu$ from a higher redshift absorber has recently been obtained from comparisons between NH$_3$ inversion and HCN/HCO$^+$ rotational lines: \\citet{murphy08} find $\\dmu < 1.6 \\times 10^{-6}$ ($2\\sigma$) from a gravitational lens at $z \\sim 0.685$. This suggests that changes in $G \\equiv g_p [\\mu \\alpha^2]^{1.85}$ at $z \\sim 0.25$ may be dominated by changes in $\\alpha$, i.e. the first limiting case above, yielding $\\dal = (-3.1 \\pm 1.2) \\times 10^{-6}$. It is interesting that this result for $\\dal$ is similar in both amplitude and sign to the value of $\\dal$ obtained by \\citet{murphy04} using the many-multiplet method at a higher redshift, $\\dal = (-5.7 \\pm 1.1) \\times 10^{-6}$ at $\\langle z \\rangle = 1.75$.  Other sensitive astronomical techniques to probe fundamental constant evolution include those based on the many-multiplet method \\citep{murphy04}, H$_2$ lines \\citep{king08}, \\hi~21cm and OH lines \\citep{kanekar05}, inversion and rotational lines \\citep{flambaum07b}, \\hi~21cm and ultraviolet carbon lines \\citep{kanekar10}, etc. All of these are affected by systematic effects of different types (e.g. \\citealp{kanekar08b}). For example, \\citet{griest10} have recently shown that Keck-HIRES spectra are affected by wavelength calibration errors of unknown origin; the effect of these systematic errors on the many-multiplet result of \\citet{murphy04} is not known, although it seems clear that the error budget will increase. Similar problems have been found with VLT-UVES spectra (e.g. \\citealp{centurion09,whitmore10}), which would affect the H$_2$ results of \\citet{king08}.  The importance of the conjugate satellite OH technique stems from the fact that it has fewer systematic effects than other approaches and allows a test of its own applicability, in the prediction that the shapes of the satellite OH lines must be the same. We have also demonstrated that the expected null result is obtained on applying the technique to a local conjugate satellite system, Cen.A. Equally important, and again unlike most other techniques, the conjugate satellite lines allow an estimate of changes in the fundamental constants from a {\\it single space-time location}, without any averaging over multiple absorbers (which, for example, is essential in the many-multiplet method to average out local systematics). The only apparent drawback to the technique is the possibility of RFI at the line frequencies; this can be ameliorated by using multiple telescopes, as done here, or multiple observing epochs, taking advantage of the shift in the line frequency in the terrestrial frame due to the motion of the Earth around the Sun.  In summary, we have obtained deep WSRT and Arecibo Telescope spectra in the redshifted satellite OH~18cm lines from \\pks, at $z \\sim 0.247$. The lines are conjugate at our sensitivity, with the sum of the optical depths consistent with noise. We find that the peak of the cross-correlation of the satellite OH~18cm lines is offset from zero velocity (at $\\sim 2.6 \\sigma$ significance), tentative evidence for a redshift offset between them. We have examined the data for systematic effects and find no evidence that such effects might give rise to the observed velocity offset. We have also verified that the cross-correlation of the satellite OH~18cm lines in the local conjugate satellite system in Cen.A peaks at zero velocity, within the errors. The observed offset in satellite line redshifts in \\pks\\ implies $\\Delta G/G = (-1.18 \\pm 0.46) \\times 10^{-5}$, where $G \\equiv g_p [\\mu \\alpha^2]^{1.85}$, suggesting that one or more of $\\alpha$, $\\mu$, and $g_p$ had smaller values at $z \\sim 0.247$ than at the present epoch."
    },
    "1004/1004.2871_arXiv.txt": {
        "abstract": "We use a general circulation model to study the three-dimensional (3-D) flow and temperature distributions of atmospheres on tidally synchronized extrasolar planets.  In this work, we focus on the sensitivity of the evolution to the initial flow state, which has not received much attention in 3-D modeling studies.  We find that different initial states lead to markedly different distributions---even under the application of strong forcing (large day-night temperature difference with a short ``thermal drag time'') that may be representative of close-in planets.  This is in contrast with the results or assumptions of many published studies.  In general, coherent jets and vortices (and their associated temperature distributions) characterize the flow, and they evolve differently in time, depending on the initial condition.  If the coherent structures reach a quasi-stationary state, their spatial locations still vary.  The result underlines the fact that circulation models are currently unsuitable for making quantitative predictions (e.g., location and size of a ``hot spot'') without better constrained, and well posed, initial conditions. ",
        "introduction": "Understanding the flow dynamics of atmospheres is crucial for characterizing extrasolar planets. Dynamics strongly influence the temperature distribution as well as the spectral behavior.  An essential tool for studying dynamics on the large-scale is a global hydrodynamics model.  Many studies have used such a model \\citep[e.g.,][]{Showman02,Choetal03,Cooper05, Langton07, Choetal08,Dobbs-Dixon08,Showmanetal08a,Menou09}.  The models in these studies numerically solve a set of non-linear partial differential equations for the evolution of a fluid on a rotating sphere.  Hence, the initial condition (as well as the boundary conditions) needs to be specified.  Presently, physically accurate and mathematically well-posed initial conditions for the models are not known for extrasolar planets. Unlike for the solar system planets, dynamically ``balanced'' initial data\\footnote{ self-consistent set of fields which does not lead to excessive noise and deviations from accurate prediction} are not available and dominant dynamical processes, such as baroclinic instability and geostrophic turbulence, are not yet understood for the extrasolar planets \\citep{Choetal03,Choetal08, Showmanetal08b,Cho08}. Concerning initialization, there is a long history of research in geophysical fluid dynamics and numerical weather prediction, and it is still a subject of active research---even for the Earth \\citep{Holton04}.  In most simulations of close-in planets performed so far, the initial state is either at rest or with a small, randomly perturbed wind field to break the flow symmetry \\citep{Cooper05,Langton07, Dobbs-Dixon08,Showmanetal08a, Menou09}. \\citet{Choetal08} initialize their two-dimensional simulations with random eddies, and variations of the initial velocity distributions are studied. They find significant differences in the flow evolution, depending on the vigor of the eddies. On the other hand, \\citet{Cooper05} report on a three-dimensional (3-D) simulation, set up with an initial retrograde equatorial jet, and find no qualitative difference, compared with one starting from a rest state. \\citet{Showmanetal08b} and \\citet{Cho08} give summaries of the various results.  In this work, we present runs from an advanced 3-D general circulation model. As in \\citet{Cooper05}, as well as in \\citet{Showmanetal08a} and \\citet{Menou09}, the model used in this work solves the full primitive equations \\citep{Pedlosky87}. However, there are some important assets in the model used in this work (see section \\ref{sec:method}), compared with most models used so far. For example, it uses a parallel pseudospectral algorithm \\citep{Orszag70,Eliasenetal70, Canutoetal88} with better-controlled, less invasive numerical viscosity. In this regard, our model is similar to the one used by \\citet{Menou09}.  With our model, we focus on the sensitivity of the flow evolution to the initial state. The sensitivity has not been much emphasized in previous studies, particularly in those using 3-D circulation models. In order to unambiguously delineate the sensitivity effect, we set up the simulations in a manner similar to previous studies, apply idealized forcing (in many cases unencumbered by a vertical variation), and compare runs with all parameters identical---except for the initial condition.  The basic plan of the paper is as follows. We describe the model and its setup for our simulations in section \\ref{sec:method}, where we endeavor to provide enough details to facilitate reproduction of the results. In section \\ref{sec:results} we present the results of simulations initialized with different organized large-scale flow patterns, including the rest state. In this section, we also show how sensitive the flow is to small perturbations in the initial wind field. We conclude in section \\ref{sec:concl}, summarizing this work and discussing its implications for close-in extrasolar planet circulation modeling work. ",
        "conclusions": "\\label{sec:concl}  In this paper, we have shown that in a generic general circulation simulation of a tidally synchronized planet, the flow and temperature distributions depend strongly on the initial state. In all simulations initialized with a different jet configuration, large scale coherent vortices are formed; but, their location, size, and number varies, depending on the initial wind. The temperature distribution is relatively homogenized by the flow, compared to the large temperature contrast in the forcing equilibrium temperature profile. But the degree of mixing, as well as locations of temperature extremes vary between differently initialized simulations. The time variability of the atmosphere---i.e. how vortices and associated temperature patterns move around the planet or whether they stay at fixed positions---varies depending on the initial wind.  The Newtonian drag scheme used in this study is idealized and the ``correct'' parameters to use in the set-up are unknown, with many choices possible. Explorations of different parts of the parameter space will be presented elsewhere. Here, we have chosen the set-up to capture, as cleanly as possible, the effects of the initial flow in a regime of parameter space plausibly relevant for hot Jupiters with strong zonally-asymmetric forcing. We have found that the strong dependence on the initial wind is valid for a wide range of thermal drag times ($\\tau_{\\rm th}$ = 0.5--10 planet days) and with T21, T42, and T85 resolutions. Some reduction in the sensitivity is sometimes observed for the very small $\\tau_{\\rm th}$ and long time integration simulations, which require large numerical dissipation to prevent the fields from being dominated by small scale noise. This situation is, however, not physically realistic and numerically suspect.  The strong dependence on initial wind has implications for the use of general circulation models for interpretation of observations of extrasolar planet atmospheres. These results underline that, while numerical circulation models of the kind employed here are useful for studying plausible mechanisms and flow regimes, they are currently unsuitable for making ``hard'' predictions---such as exact locations of temperature extremes on a given planet."
    },
    "1004/1004.2089_arXiv.txt": {
        "abstract": "{} {The properties of the broad-band emission from the high-frequency peaked BL\\,Lac H\\,2356$-$309  ($z=0.165$) are investigated.} {Very High Energy (VHE; E $>$ 100 GeV) observations of H\\,2356$-$309 were performed with the High Energy Stereoscopic System (HESS) from 2004 through 2007. Simultaneous optical/UV and X-ray observations were made with the XMM-Newton satellite on June 12/13 and June 14/15, 2005. NRT radio observations were also contemporaneously performed in 2005.  ATOM optical monitoring observations were also made in 2007.} {A strong VHE signal, $\\sim$13$\\sigma$ total, was detected by HESS after the four years HESS observations (116.8 hrs live time). The integral flux above 240 GeV is I($>$240 GeV) = $(3.06 \\pm 0.26_{\\rm stat} \\pm 0.61_{\\rm syst}) \\times 10^{-12}$ cm$^{-2}$\\,s$^{-1}$, corresponding to $\\sim$1.6\\% of the flux observed from the Crab Nebula. A time-averaged energy spectrum is measured from 200 GeV to 2 TeV and is characterized by a power law (photon index of $\\Gamma = 3.06 \\pm 0.15_{\\rm stat} \\pm 0.10_{\\rm syst}$). Significant small-amplitude variations in the VHE flux from H\\,2356$-$309 are seen on time scales of months and years, but not on shorter time scales.  No evidence for any variations in the VHE spectral slope are found within these data. The XMM-Newton X-ray measurements show a historically low X-ray state, characterized by a hard, broken-power-law spectrum on both nights. } {The broad-band spectral energy distribution (SED) of the blazar can be adequately fit using a simple one-zone synchrotron self-Compton (SSC) model. In the SSC scenario, higher VHE fluxes could be expected in the future since the observed X-ray flux is at a historically low level.} ",
        "introduction": "H\\,2356$-$309 was suggested by \\cite{Costamante_AGN} as a strong candidate for the detection of VHE emission, based on simple one-zone synchrotron self-Compton (SSC) modeling of its spectral energy distribution (SED). Motivated by these predictions, the HESS collaboration observed this blazar in 2004, resulting in the discovery of VHE $\\gamma$-ray emission (\\cite{HESS_discovery}). Although several blazars are known to have variable VHE $\\gamma$-ray fluxes (see, e.g., \\cite{2155_flare}) and photon spectra (see, e.g., \\cite{HESS_421}), no strong evidence was found for such variability in the HESS data on this source from June through December 2004.  Hosted by an elliptical galaxy at a redshift of $z=0.165$ (\\cite{redshift}), H\\,2356$-$309 is among the more distant blazars detected at VHE energies. As a result, its observed spectrum is expected to be strongly affected (softened) by the absorption of VHE photons on the extragalactic background light (EBL). The unexpectedly hard ($\\Gamma \\sim 3.1$) VHE spectrum measured by HESS in 2004 was used to confirm the strong constraints on the density of the EBL, near lower limits based on the integrated light of resolved galaxies (\\cite{HESS_discovery}), given by the HESS spectrum from another VHE blazar (1ES\\,1101$-$232).  At X-ray energies, H\\,2356$-$309 was initially detected by the UHURU satellite \\cite{UHURU_det} and later by the HEAO-I satellite \\cite{HEAO_det}. It is among the X-ray brightest of the 150 high-frequency peaked BL Lac objects (HBL) in the Sedentary Survey \\cite{Sedentary}, which contains HBLs characterized by their exceptionally high X-ray-to-radio flux ratio. BeppoSAX observations \\cite{SAX_data} of H\\,2356$-$309 in June 1998 yielded an X-ray spectrum extending up to 50 keV, characterized by a broken power-law shape ($\\Gamma_1 = 0.78^{+0.06}_{-0.07}$; $\\Gamma_2 = 1.10 \\pm 0.04$) with a break at $1.8 \\pm 0.4$ keV.  A softer power-law spectrum between 2 and 10 keV with an X-ray photon index of ($\\Gamma_X$ = 2.43 $\\pm$ 0.11) was measured (\\cite{HESS_MWLI}) during observations by the RXTE satellite on November 11, 2004.  The RXTE spectrum is softer above 2 keV and the flux ($9.7^{+0.3}_{-1.3} \\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$) is $\\sim$3 times lower than that measured by BeppoSAX ($2.5 \\times 10^{-11}$ erg cm$^{-2}$ s$^{-1}$) in the same energy band.  The 2004 RXTE data were taken simultaneously with HESS VHE, ROTSE-IIIc optical, and Nancay radio-telescope observations. The double-peaked SED derived from this campaign (\\cite{HESS_MWLI}) was the first to include a contemporaneous VHE detection, and hence the first to sample the higher-energy SED peak for this object since H\\,2356$-$309 was not detected by the EGRET satellite.  Modeling of these multi-wavelength observations demonstrated that the SED of H\\,2356$-$309 could be adequately fit by a simple one-zone stationary SSC scenario.  Results from 76.9 hrs of new HESS observations of H\\,2356$-$309 in 2005, 2006 and 2007 are reported here.  In addition, a re-analysis of the previously published 2004 HESS data (39.9 hrs), with an improved calibration of the absolute energy scale of the detector, is presented. This improvement in the HESS calibration was not included in the previously mentioned modeling of the 2004 multi-wavelength observations, therefore this modeling is repeated.  In addition, the XMM-Newton satellite simultaneously observed H\\,2356$-$309 on two nights ($\\sim$5 hours each night) at X-ray energies, as well as at optical through ultraviolet wavelengths, during the HESS observations in June 2005. The results from this multi-wavelength observation campaign, which has much improved X-ray and optical measurements, are also discussed. ",
        "conclusions": "\\begin{figure} \\centering \\includegraphics[width=0.33\\textwidth, angle = 270]{figures/14321_f9.eps} \\caption{SED of H\\,2356-309 zoomed in on the VHE band.  The data and model curves are identical to those in Figure~\\ref{H2356_SEDs}.} \\label{VHE_zoom} \\end{figure}  Long term HESS observations of H\\,2356$-$309 have clearly confirmed the initial VHE detection from 2004.  The source has a relatively low VHE flux (1.6\\% Crab) that is weakly variable (factor of $\\sim$2) on time scales of at least months.  Although the flux is variable, no changes in the VHE spectral slope ($\\Gamma\\approx3.1$) are found. After accounting for the effects of the absorption of VHE photons on the extragalactic background light, the intrinsic VHE spectrum of H\\,2356$-$309 is found to be hard ($\\Gamma_{\\rm int}\\approx2.0$). Interestingly, the X-ray spectrum is also very hard, with the synchrotron peak located at $\\sim$1 keV, despite the source being in a historically low state in X-rays. While the synchrotron peak of H\\,2356$-$309 is located at higher than typical energies and the X-ray and VHE spectra are comparatively hard, the SED can be reasonably modeled using a simple one-zone SSC scenario. This could be expected since the luminosity of both SED peaks are similar, and their separation in frequency is not extreme."
    },
    "1004/1004.1829_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\setcounter{equation}{0} \\def\\theequation{8.\\arabic{equation}}  The attempt to explain some of the CMB ``anomalous'' features as the indication of a break of statistical isotropy is the main reason behind ours and many of the existing inflationary models populated by vector fields, but not the only one. The first one of these models \\cite{Ford:1989me} was formulated with the goal of producing inflation by the action of vector fields, without having to invoke the existence of a scalar field. The same motivations inspired the works that followed \\cite{Golovnev:2008cf,Golovnev:2008hv,Golovnev:2009ks}. Lately, models where primordial vector fields can leave an imprint on the CMB have been formulated as an alternative to the basic inflationary scenario, in the search for interesting non-Gaussianity predictions \\cite{Yokoyama:2008xw,ValenzuelaToledo:2009af,Dimopoulos:2008yv,ValenzuelaToledo:2009nq,Karciauskas:2008bc,Dimopoulos:2009am,Dimopoulos:2009vu,Bartolo:2009pa,Bartolo:2009kg}. Finally, vector fields models of dark energy have been proposed \\cite{ArmendarizPicon:2004pm,Boehmer:2007qa,Koivisto:2007bp,Koivisto:2008xf,Jimenez:2009py,Jimenez:2009zza}. All this appears to us as a rich bag of motivations for investigating these scenarios.\\\\ Before we quickly sketch some of them and list the results so far achieved in this direction, it is important to briefly indicate and explain the main issues and difficulties that these models have been facing. We will also shortly discuss the mechanisms of production of the curvature fluctuations in these models. \\\\  \\noindent Building a model where primordial vector fields can drive inflation and/or produce the observed spectrum of large scale fluctuations requires a more complex Lagrangian than the basic gauge invariant $\\mathcal{L}_{vector}=-(\\sqrt{-g}/4) F^{\\mu\\nu}F_{\\mu\\nu}$. In fact, for a conformally invariant theory as the one described by $\\mathcal{L}_{vector}$, vector fields fluctuations are not excited on superhorizon scales. It is then necessary to modify the Lagrangian. For some of the existing models, these modifications have been done to the expense of destabilizing the theory, by ``switching on'' unphysical degrees of freedom. This was pointed out in \\cite{Himmetoglu:2008zp,Himmetoglu:2008hx,Himmetoglu:2009qi}, where a large variety of vector field models was analyzed in which longitudinal polarization modes exist that are endowed with negative squared masses (the ``wrong'' signs of the masses are imposed for the theory to satisfy the constraints that allow a suitable background evolution). It turnes out that, in a range of interest of the theory, these fields acquire negative total energy, i.e. behave like ``ghosts'', the presence of which is known to be responsible for an unstable vacuum. A related problem for some of these theories is represented by the existence of instabilities affecting the equations of motion of the ghost fields \\cite{Himmetoglu:2008zp,Himmetoglu:2008hx,Himmetoglu:2009qi}.\\\\ In the remaing part of this section, we are going to present some of these models together with some recent attempts to overcome their limits. \\\\  \\noindent In all of the models we will consider, primordial vector fields fluctuations end up either being entirely responsible for or only partially contributing to the curvature fluctuations at late times. This can happen through different mechanisms. If the vector fields affects the universe expansion during inflation, its contribution $\\zeta_{A}$ to the total $\\zeta$ can be derived from combining the definition of the number of e-foldings ($N=\\int H dt$) with the Einstein equation ($H^2=(8 \\pi G/3)(\\rho_{\\phi}+\\rho_{A})$, $\\rho_{A}$ being the energy density of the vector field and $\\rho_{\\phi}$ the inflaton energy density) and using the $\\delta$N expansion of the curvature fluctuation in terms of both the inflaton and the vector fields fluctuations (see Sec.~\\ref{cf}). To lowest order we have \\cite{Dimopoulos:2008yv} \\bea\\label{direct} \\zeta_{A}=\\frac{A_{i}}{2 m_{P}^2}\\delta A_{i}, \\eea where a single vector field has been taken into account for simplicity ($m_{P}$ is the reduced Planck mass, $A$ is the background value of the field and $\\delta A$ its perturbation). When calculating the amplitude of non-Gaussianity in Sec.~\\ref{amplitude}, we will refer to this case as ``vector inflation'' for simplicity.\\\\ A different fluctuation production process is the curvaton mechanism which was initially formulated for scalar theories but it is also applicable to vectors \\cite{Dimopoulos:2006ms,Dimopoulos:2008rf}. Specifically, inflation is driven by a scalar field, whereas the curvaton field(s) (now played by the vectors), has a very small (compared to the Hubble rate) mass during inflation. Towards the end of the inflationary epoch, the Hubble rate value starts decreasing until it equates the vector mass; when this eventually happens, the curvaton begins to oscillate and it will then dissipate its energy into radiation. The curvaton becomes responsible for a fraction of the total curvature fluctuation that is proportional to a parameter, $r$, related to the ratio between the curvaton energy density and the total energy density of the universe at the epoch of the curvaton decay~\\cite{Dimopoulos:2008yv} \\bea\\label{indirect} \\zeta_{A}=\\frac{r}{3}\\frac{\\delta \\rho_{A}}{\\rho_{A}}, \\eea where $r\\equiv 3\\rho_{A}/(3\\rho_{A}+4\\rho_{\\phi})$. Anisotropy bounds on the power spectrum favour small values of $r$.\\\\ From Eqs.~(\\ref{direct}) and (\\ref{indirect}) we can see that, dependending on which one of these two mechanisms of production of the curvature fluctuations is considered, different coefficients will result in the $\\delta$N expansion (see Eq.~(\\ref{expansion})).\\\\  \\noindent In this section we will describe both models where inflation is intended to be vector-field driven and those models in which, instead, the role of the inflaton is played by a scalar field, whereas the energy of the vector is a subdominant contribution to the total energy density of the universe during the entire inflationary phase.   \\subsection{Self-coupled vector field models}  A pioneer work on vector field driven inflation was formulated by L.~H.~Ford \\cite{Ford:1989me}, who considered a single self-coupled field $A_{\\mu}$ with a Lagrangian \\bea L_{vector}=-\\frac{1}{4}F_{\\mu\\nu}F^{\\mu\\nu}+V(\\psi) \\eea where $F_{\\mu\\nu}\\equiv \\p_{\\mu}B_{\\nu}-\\p_{\\nu}B_{\\mu}$ and the potential $V$ is a function of $\\psi\\equiv B_{\\alpha}B^{\\alpha}$. Different scenarios of expansion are analyzed by the author for different functions $V$. The universe expands anisotropically at the end of the inflationary era and this anisotropy either survives until late times or is damped out depending on the shape and the location of the minima of the potential.\\\\  \\noindent The study of perturbations in a similar model was proposed by Dimopoulos in \\cite{Dimopoulos:2006ms} where he showed that for a Lagrangian \\bea\\label{dimopoulos} L_{vector}=-\\frac{1}{4}F_{\\mu\\nu}F^{\\mu\\nu}+\\frac{1}{2}m^2 B_{\\mu}B^{\\mu} \\eea and for $m^{2}\\simeq -2H^2$, the transverse mode of the vector field is governed by the same equation of motion as a light scalar field in a de Sitter stage. A suitable superhorizon power spectrum of fluctuations could therefore arise. In order to prevent production of large scale anisotropy, in this model the vector field plays the role of the curvaton while inflation is driven by a scalar field.  \\subsection{Vector-field coupled to gravity}  The Lagrangian in Eq.~(\\ref{dimopoulos}) may be also intended, at least during inflation, as including a non-minimal coupling of the vector field to gravity; indeed the mass term can be rewritten as \\bea\\label{Lag} L_{vector}\\supset\\frac{1}{2}\\left(m^2_{0}+\\xi R\\right)B_{\\mu}B^{\\mu} \\eea where, for the whole duration of the inflationary era, the bare mass $m_{0}$ is assumed to be much smaller than the Hubble rate and the Ricci scalar $R=-6\\left[\\frac{\\ddot{a}}{a}+\\left(\\frac{\\dot{a}}{a}\\right)^2\\right]$ can be approximated as $R\\simeq-12 H^2$. For the specific value $\\xi=1/6$, Eq.~(\\ref{dimopoulos}) is retrieved.\\\\  \\noindent For the Lagrangian just presented, Golovnev et al \\cite{Golovnev:2008cf} proved that the problem of excessive anisotropy production in the case where inflation is driven by vector fields can be avoided if either a triplet of mutually orthogonal or a large number $N$ of randomly oriented vector fields is considered.\\\\  \\noindent The Lagrangian (\\ref{Lag}) with $\\xi=1/6$ was also employed in \\cite{Dimopoulos:2008rf}, where inflation is scalar-field-driven and a primordial vector field affects large-scale curvature fluctuations and, similarly, in \\cite{Kanno:2008gn}, which includes a study of the backreaction of the vector field on the dynamics of expansion, by introducing a Bianchi type-I metric.    \\subsection{Ackerman-Carroll-Wise (ACW) model}  A model was proposed in \\cite{Ackerman:2007nb} where Lagrange multipliers ($\\lambda$) are employed to determine a fixed norm primordial vector field $B_{\\mu}B^{\\mu}=m^2$ \\bea\\label{Acw} L_{vector}\\supset \\lambda\\left(B^{\\mu}B_{\\mu}-m^2\\right)-\\rho_{\\Lambda} \\eea where $\\rho_{\\Lambda}$ is a vacuum energy. The expansion rate in this scenario is anisotropic: if we orient the $x$-axis of the spatial frame along the direction determined by the vector field, we find two different Hubble rates: along the $x$-direction it is equal to \\bea H_{b}^{2}=\\frac{\\rho_{\\Lambda}}{m_{P}^{2}}\\frac{1}{P(\\mu)}, \\eea and it is given by $H_{a}=(1+c\\mu^2)H_{b}$ along the orthogonal directions; $\\mu\\equiv m/m_{P}$, $P$ is a polynomial function of $\\mu$ and $c$ is a parameter appearing in the kinetic part of the Lagrangian that we omitted in (\\ref{Acw}) (see \\cite{Ackerman:2007nb} for its complete expression). As expected, an isotropic expansion is recovered if the vev of the vector field is set to zero.    \\subsection{Models with varying gauge coupling}  Most of the models mentioned so far successfully solve the problem of attaining a slow-roll regime for the vector-fields without imposing too many restrictions on the parameters of the theory and of avoiding excessive production of anisotropy at late times. None of them though escapes those instabilities related to the negative energy of the longitudinal modes (although a study of the instabilities for fixed-norm field models was done in \\cite{Carroll:2009em} where some stable cases with non-canonical kinetic terms were found). As discussed in \\cite{Himmetoglu:2008zp,Himmetoglu:2008hx,Himmetoglu:2009qi}, in the self-coupled model a ghost appears at small (compared to the horizon) wavelengths; in the non-minimally coupled and in the fixed-norm cases instead the instability concerns the region around horizon crossing.\\\\  \\noindent Models with varying gauge coupling can overcome the problem of instabilities and have recently attracted quite some attention. In \\cite{Yokoyama:2008xw}, the authors consider a model of hybrid inflation \\cite{Lyth:2005qk,Alabidi:2006wa,Salem:2005nd,Bernardeau:2004zz} with the introduction of a massless vector field \\bea L\\supset \\frac{1}{2}\\left(\\p_{\\mu}\\phi\\p^{\\mu}\\phi+\\p_{\\mu}\\chi\\p^{\\mu}\\chi\\right)-\\frac{1}{4}f^{2}(\\phi)F_{\\mu\\nu}F_{\\mu\\nu}+V(\\phi,\\chi,B_{\\mu}) \\eea where $\\phi$ is the inflaton and $\\chi$ is the so-called ``waterfall'' field. The potential $V$ is chosen in such a way as to preserve gauge invariance; this way the longitudinal mode disappears and instabilities are avoided.\\\\ Similarly, Kanno et al \\cite{Watanabe:2009ct} consider a vector field Lagrangian of the type \\bea\\label{} L_{vector}=-\\frac{1}{4}f^{2}(\\phi)F^{\\mu\\nu}F_{\\mu\\nu}, \\eea but in a basic scalar field driven inflation model. Very recently, in \\cite{Himmetoglu:2009mk,DG} the linear perturbations in these kind of models have been investigated.\\\\ Finally, in \\cite{Dimopoulos:2008yv,Dimopoulos:2009vu} varying mass vector field models have been introduced \\bea\\label{var} L_{vector}=-\\frac{1}{4}f^{2}(\\phi)F^{\\mu\\nu}F_{\\mu\\nu}+\\frac{1}{2}m^{2}B_{\\mu}B^{\\mu}, \\eea where $f\\simeq a^{\\alpha}$ and $m\\simeq a$ ($a$ is the scale factor and $\\alpha$ is a numerical coefficient). The special cases $\\alpha=1$ and $\\alpha=-2$ are of special interest. In fact, introducing the fields $\\tilde{A}_{\\mu}$ and $A_{\\mu}$, related to one another by $\\tilde{A}_{\\mu}\\equiv f B_{\\mu}=a A_{\\mu}$ ($\\tilde{A}_{\\mu}$ and $A_{\\mu}$ are respectively the comoving and the physical vectors), it is possible to verify that the physical gauge fields are governed by the same equations of motion as a light scalar field in a de Sitter background. Vector fields in this theory can then generate the observed (almost) scale invariant primordial power spectrum.    \\subsection{$SU(2)$ vector model}\\label{ourmodel}  Let us consider some models where inflation is driven by a scalar field in the presence of an $SU(2)$ vector multiplet \\cite{Bartolo:2009pa,Bartolo:2009kg}. A fairly general Lagrangian can be the following \\bea\\label{ac}\\fl S=\\int d^{4}x \\sqrt{-g}\\left[\\frac{m_{P}^{2}R}{2}-\\frac{f^{2}(\\phi)}{4}g^{\\mu\\al}g^{\\nu\\b}\\sum_{a=1,2,3}F_{\\mu\\nu}^{a}F_{\\al\\b}^{a}-\\frac{M^2}{2}g^{\\mu\\nu}\\sum_{a=1,2,3}B_{\\mu}^{a}B_{\\nu}^{a}+L_{\\phi}\\right], \\eea where $L_{\\phi}$ is the Lagrangian of the scalar field and $F_{\\mu\\nu}^{a}\\equiv\\p_{\\mu}B^{a}_{\\nu}-\\p_{\\nu}B^{a}_{\\mu}+g_{c}\\ep^{abc}B^{b}_{\\mu}B^{c}_{\\nu}$ ($g_{c}$ is the $SU(2)$ gauge coupling). Both $f$ and the effective mass $M$ can be viewed as generic functions of time. The fields $B_{\\mu}^{a}$ are comoving and related to the physical fields by $A_{\\mu}^{a}=\\left(B_{0}^{a},B_{i}^{a}/a\\right)$. The free field operators can be Fourier expanded in their creation and annihilation operators \\bea\\label{ip} \\fl \\delta A_{i}^{a}(\\vec{x},\\eta)=\\int \\frac{d^{3}q}{(2\\pi)^{3}}e^{i\\vec{q}\\cdot\\vec{x}}\\sum_{\\lambda=L,R,long}\\Big[e^{\\lambda}_{i}(\\hat{q})a_{\\vec{q}}^{a,\\lambda}\\delta A_{\\lambda}^{a}(q,\\eta)+e^{*\\lambda}_{i}(-\\hat{q})\\left(a_{-\\vec{q}}^{a,\\lambda}\\right)^{\\dagger}\\delta A_{\\lambda}^{*a}(q,\\eta)\\Big], \\eea where the polarization index $\\lambda$ runs over left ($L$), right ($R$) and longitudinal ($long$) modes and \\bea \\left[a^{a,\\lambda}_{\\vec{k}},(a_{\\vec{k}^{'}}^{a^{'},\\lambda^{'}})^{\\dagger}\\right]=(2 \\pi)^3\\delta_{a,a^{'}}\\delta_{\\lambda,\\lambda^{'}}\\delta^{(3)}(\\vec{k}-\\vec{k}^{'}). \\eea Here $\\eta$ the conformal time ($d\\eta=dt/a(t)$). Once the functional forms of $f$ and $M$ have been specified, the equations of motion for the vector bosons can be written. We compute cosmological correlation functions up to fourth order considering an action as in (\\ref{ac}). The expression of the correlators that we derive, prior to explicitating the wavefunction for the gauge bosons, apply to any $SU(2)$ theory with an action as in (\\ref{ac}), both for what we will call the ``Abelian'' and for the ``non-Abelian'' contributions. In particular, the structure of the interaction Hamiltonian is independent of the functional dependence of $f$ and $M$ and determines the general form of and the anisotropy coefficients appearing in the final ``non-Abelian'' expressions (see Sec.~\\ref{cf}). When it comes to explicitate the wavefunctions, a choice that can help keeping the result as easy to generalize as possible is the following \\bea\\label{T} \\delta B^{T}=-\\frac{\\sqrt{\\pi x}}{2\\sqrt{k}}\\left[J_{3/2}(x)+iJ_{-3/2}(x)\\right], \\eea for the transverse mode and \\bea\\label{L} \\delta B^{||}=n(x)\\delta B^{T}\\, , \\eea for the longitudinal mode ($n$ is a unknown function of $x\\equiv -k \\eta$) \\cite{Bartolo:2009pa,Bartolo:2009kg}. Let us see why. As previously stated, for $f \\simeq a^{\\alpha}$ and with $\\alpha=0,1,-2$, it is possible to verify that the (physical) transverse mode behaves exactly like a light scalar field in a de Sitter background (see also Appendix~\\ref{eomotion}). Considering the solution (\\ref{T}) then takes into account at least these special cases. As to the longitudinal mode, a parametrization was adopted as in (\\ref{L}) in order to keep the analysis more general and given that, because of the instability issues, introducing this degree of freedom into the theory requires special attention. We are going to keep the longitudinal mode ``alive'' in the calculations we present, by considering a nonzero function $n(x)$, and focus on the simplest case of $f=1$. This case is known to be affected by quantum instabilities in the longitudinal mode, anyway we choose $f=1$ for the sake of simplicity in our presentation. The results can be easily generalized to gauge invariant models (please refer to Sec.~\\ref{gen} for a sample generalization of some of the calculations to massless $f\\simeq a^{(1,-2)}$ models).  \\newpage ",
        "conclusions": "Cosmology has entered what it is called its ``precision era'': over the past few years, observations of the CMB and of other cosmological probes have been performed which have greatly improved the previous bounds on some of the fundamental parameters characterizing the early Universe cosmology. These bounds are expected to become a lot stricter with the advent of new experiments, such as the ongoing Planck satellite mission. The theoretical search for models of the early Universe has consequently been concentrated on, among other things, trying to produce more and more accurate predictions. \\\\ In this optics, alternative models to the basic single-field slow-roll inflationary scenario have been and are being investigated, with a strong focus on the computation of higher order (three and four-point) correlation functions and on higher order (beyond tree-level) corrections. All of these ``higher order'' predictions are expected to be particularly revealing of the early Universe physics, given that they are an indication of the type of fields and interactions populating the specific inflationary scenarios.   \\\\  \\noindent In this thesis I have presented an overview of the existing results concerning some of these ``higher order'' predictions; in particular, I focused on our works on loop corrections in inflationary scalar field models and on the primordial non-Gaussianity and statistical anisotropy predictions in some inflationary models populated with non-Abelian vector fields.\\\\  \\noindent Loop corrections to the power spectrum of the comoving curvature fluctuation $\\zeta$ in single-field inflation arise both from the inflaton self-interactions and from the coupling of the scalar field with gravitons. We have calculated the corrections from tensor loops, previously neglected in the literature for simplicity reasons. It turnes out that one-loop corrections from tensor-scalar interactions are of the same order of magnitude as those arising from scalar self-interactions, therefore they cannot be neglected in a self-consistent calculation.\\\\ \\noindent One loop corrections have been found to be suppressed by an $(H/m_{P})^2$ factor compared to the tree-level result; they exhibit a slightly different dependence from the external momentum, because of the presence of a logarithmic factor (which does not spoil scale invariance) and they turn out vary with time at most as fast as the logarithm of the scale factor. One-loop diagrams generally present both infrared logarithmic and ultraviolet power-law (in addition to logarithmic) divergences. The ultraviolet divergences can be treated using ordinary regularization and renormalization techniques as in flat-space quantum field theory. The infrared divergences are ``cured'' by introducing an infrared cutoff represented by the smallest observable physical mode, i.e. by considering a finite space, a sort of ``box'' of observation, with a size equal to the current horizon length. As expected, tensor modes provide, as well as the scalar modes, observably small corrections at one loop level, as long as the basic inflationary scenario is concerned. Their computation is nevertheless instructive and provides a path that can be followed in the investigation of ``non-standard'' models of inflation. More appealing results are for instance expected in inflationary scenarios described by non-canonical Lagrangians, some of which, $P(X,\\phi)$ models, have been described and analysed in this thesis. \\\\  \\noindent Motivated by an interest in models that combine non-Gaussianity and statistical anisotropy predictions for the CMB fluctuations, we have worked on theories of inflation where primordial vector fields effectively participate in the production of the $\\zeta$ perturbations, eventually focusing on some $SU(2)$ vector fields models. The two, three and four point correlation functions in these models result as the sum of scalar and vector contributions. The latter are of two kinds, ``Abelian'' (i.e. arising from the zeroth order terms in the Schwinger-Keldysh expansion) and ``non-Abelian'' (i.e. originating from the self-interactions of the vector fields). The bispectrum and the trispectrum final results are presented as a sum of products of isotropic functions of the momenta ($F_{n}$ and $G_{n}$ in the text) multiplied by anisotropy coefficients (which we indicated by $I_{n}$ and $L_{n}$) that depend on the angles among all gauge and wave vectors. \\\\ The amplitude of non-Gaussianity has been evaluated through the parameters $f_{NL}$ and $\\tau_{NL}$; in particular we have discussed the dependence of these functions from the non-angular parameters of the theory. We have provided the comparisons among the different (scalar versus vector, Abelian versus non-Abelian) contributions to $f_{NL}$ and $\\tau_{NL}$, noticing that any one of them can be the dominant contribution depending on the selected region of parameter space. In particular, we have stressed how the anisotropic contributions to the bispectrum and the trispectrum can overcome the isotropic parts. An interesting feature of these models is that the bispectrum and the trispectrum depend on the same set of parameters and their amplitudes are therefore strictly related to one another. \\\\ We have presented the shapes of both the bispectrum and the trispectrum. The isotropic functions appearing in their final expressions have been analyzed separately from their anisotropy coefficients. The bispectrum isotropic functions have been found to preferably show a local shape. The trispectrum ones have been plotted selecting equilateral and specialized planar configurations. Finally, the full expressions (i.e. completed by the anisotropy coefficients) of bispectrum and trispectrum have been presented in specific momenta configuration, in order to provide a hint of the modulation of shapes and amplitudes operated by anisotropy.\\\\ \\noindent  In our view, the most promising features of these models consists in the possibility of providing non-Gaussianity and statistical anisotropy predictions that are related to one another because of the fact that they share the same underlying theory. Models that combine both types of predictions could be more easily testable and, from non-Gaussianity measurement, more stringent statistical anisotropy predictions could be produced or viceversa.\\\\  \\noindent Inflationary models that do not spoil the current agreement with experimental data constitute a huge variety but many of them have very distinctive features that might be confirmed or ruled out by observations. Both the the non-Gaussianity and statistical anisotropy predictions and the nature and the amount of higher corrections to the cosmological correlation functions can be ranked among these distinctive features and certainly deserve further investigations, looking forward for a confront with new and promising experimental data. \\\\  \\noindent While completing this thesis, the Wilkinson Microwave Anisotropy Probe (WMAP) team published its seven-year data analyses \\cite{Komatsu:2010fb,Bennett:2010jb}. The new bounds on non-Gaussianity from this study are given by $-10<f_{NL}^{equil}<74$ and $-214<f_{NL}^{loc}<266$, at $95\\%$ CL.     \\newpage"
    },
    "1004/1004.4910_arXiv.txt": {
        "abstract": "We have investigated the oxygen and nitrogen chemical abundances in extremely compact star-forming galaxies (SFGs) with redshifts between $\\sim$0.11 and 0.35, popularly referred to as ``green peas''.  Direct and strong-line methods sensitive to the N/O ratio applied to their Sloan Digital Sky Survey (SDSS) spectra reveal that these systems are genuine metal-poor galaxies, with mean oxygen abundances $\\sim 20$\\% solar. At a given metallicity these galaxies display systematically large N/O ratios compared to normal galaxies, which can explain the strong difference between our metallicities measurements and previous ones.  While their N/O ratios follow the relation with stellar mass of local SFGs in the SDSS, we find that the mass--metallicity relation of the ``green peas'' is offset $\\ga$0.3 dex to lower metallicities.  We argue that recent interaction-induced inflow of gas, possibly coupled with a selective metal-rich gas loss, driven by supernova winds, may explain our findings and the known galaxy properties, namely high specific star formation rates, extreme compactness, and disturbed optical morphologies.  The ``green pea'' galaxy properties seem to be not common in the nearby universe, suggesting a short and extreme stage of their evolution. Therefore, these galaxies may allow us to study in great detail many processes, such as starburst activity and chemical enrichment, under physical conditions approaching those in galaxies at higher redshifts. ",
        "introduction": "The cosmological relevance of local starbursts is due to their similarities with their analogs at high redshift. They may therefore provide extremely valuable laboratories to study with high resolution and sensitivity relevant physical processes associated with the starburst activity in much better detail.  One of the most suitable environments for these purposes can be the intriguing class of very compact, extremely star-forming galaxy (SFG) at low redshift ($0.11 \\la z \\la 0.35$) recently discovered by volunteers in the ``Galaxy Zoo'' project \\citep{Lintott}.  These galaxies, popularly referred to as ``green peas'' (hereafter GPs), were first reported and studied at some length by \\citet[][hereafter C09]{C09}, who classified more than one hundred of them.  The GPs are rare objects, located in lower-density environments and spectroscopically characterized by very faint continuum emission and strong optical emission lines, namely [O{\\sc iii}]$\\lambda$5007, with an unusually large equivalent width of up to $\\sim$1000\\AA.  These properties along with their extreme compactness (unresolved in Sloan Digital Sky Survey (SDSS) imaging) are the reason for their green colors and point-like appearance in the SDSS plates. According to C09, GPs are low-mass galaxies (stellar masses $M_{\\star}$$\\sim$10$^{8.5}$ to 10$^{10}$$M_{\\odot}$) with extremely high star formation rates (SFRs up to 30 $M_{\\odot}$ yr$^{-1}$), as derived from their emission lines and UV luminosity.  Therefore, their specific SFRs (SSFRs, i.e., SFR per unit of mass) are among the highest inferred in the nearby universe.  The GP properties are then consistent with the picture of galaxy downsizing \\citep{Cowie96}, where the sites of active star formation shift from high-mass galaxies at early times to lower mass systems at later epoch \\citep[see also][]{Bundy06,Cowie08,Noeske07}.  The motivation for this Letter was one of the most striking results obtained by C09: the nearly constant solar metallicity of the GPs (log[O/H]+12$\\sim$8.7 in median) over more than two orders of magnitudes in stellar mass.  These results would imply strong constrains to the nature of GPs and their evolutionary status.  In this Letter we examine in more detail the GP chemical abundances, namely oxygen and nitrogen, and their relation with the stellar mass, using the same SDSS spectra but a different methodology than in C09. This analysis indicates instead, that GPs can be considered as a genuine population of metal-poor galaxies.  The results are compared with a larger sample of local SFGs from the SDSS and with a smaller sample of super compact UV-luminous starburst galaxies \\citep{Heckman05} placed at the same redshift range of the GPs.  In particular, the position of GPs in the fundamental relation between mass and metallicity is discussed in the context of their starburst activity and feedback processes. ",
        "conclusions": "The shape of the MZR in galaxies has been found to depend on several key processes in galaxy evolution, including the efficiency of star formation \\citep[e.g.,][]{Lequeux79,MollaDiaz05,Brooks07,Calura09}, the action of metal-rich selective outflows and metal-poor gas inflows \\citep[e.g.,][]{Larson74,Garnett02,Tremonti,Finlator08}, and the possible variations in the initial mass function \\citep{Koppen07}. Moreover, different amounts of dark matter can also assist in some of these mechanisms \\citep{Dekel86}.  The GPs follow a relation between mass and metallicity that parallels the MZR defined by the SDSS SFGs, but is offset $\\ga$0.3~dex to lower metallicities.  Interestingly, we find some remarkable similarities between the GPs and the population of [\\oiii]-selected galaxies of similar luminosity in the $z$ range 0.29--0.42 recently found by \\citet{Salzer09}.  These galaxies follow a luminosity$-$metallicity relation that parallels the one defined by SFGs, but is offset by a factor of more than 10 to lower abundances.  On the other hand, since GPs and low-mass LBAs lie in a similar offset position in the MZR relative to normal SFGs, their mentioned similarities appear even greater. One possibility to explain their offset position in the MZR is that these galaxies could be still converting a large amount of their cold gas reservoirs into stars.  In that case, their low abundances could be due to their relatively young ages compared to normal SFGs. In the same range of masses, both GPs and LBAs have much higher SSFR (typically $>$10$^{-9}$yr$^{-1}$) compared to other SFGs of similar mass (C09).  Recent studies show that galaxies with higher SSFRs or larger half-light radii for their stellar mass have systematically lower metallicities \\citep[e.g.,][]{Tremonti, Ellison08}.  However, we have found even greater under-abundances in the GPs, which have high SSFRs but are extremely compact.  Some models show that in highly concentrated (typical sizes $<$3~kpc) low-mass galaxies such as the GPs, galactic winds induced by their large SSFR are strong enough to escape from their weak potential wells, diminishing the observed global abundances \\citep[e.g.,][]{Finlator08}. In contrast, analytical models by \\citet{Dalcanton07} show that any subsequent star formation to the outflow will remove their effects on metallicity, unless galaxies have an inefficient star formation. Smoothed particle hydrodynamics (SPH) plus $N$-body simulations have shown that low star formation efficiencies, regulated by supernova (SN) feedback, could be primarily responsible for the lower metallicities of low-mass galaxies and their overall trend in the MZR \\citep{Brooks07}.  As shown by \\citet{Erb06} for SFGs at $z \\sim 2$, the constancy in the offset of the MZR suggests the presence of selective metal-rich gas loss driven by SN winds.  Inflow of metal-poor gas, either from the outskirts of the galaxy or beyond, can dilute metals in the galaxy centers, explaining an offset to lower abundances in both the MZR \\citep{Mihos96,Barnes96,Finlator08} and the N/O -- O/H diagram. In starburst galaxies, a recent cold gas accretion can be due to interactions, which eventually increases the gas surface density and consequently the star formation.  As explained by \\citet{Ellison08}, the dilution of metals due to an inflow can be restored by the effects of star formation depending on the dilution-to-dynamical timescale ratio.  Since this ratio depends inversely with galaxy radius, galaxies with smaller radius, such as the GPs, may be expected to take longer time to enhance their oxygen abundances to the values expected from the MZR.  In this line, the position of GPs in the $M_{\\star}-$N/O relation and the offset observed in the N/O -- O/H plane may favor this scenario.  Models by \\citet{Koppen05} have shown that the rapid decrease of the oxygen abundance during an episode of massive and rapid accretion of metal-poor gas is followed by a slower evolution which leads to the closed-box relation, thus forming a loop in the N/O -- O/H diagram.  The inflow hypothesis is also strongly suggested by the disturbed morphologies and close companions observed in spatially resolved {\\em HST} images for three GPs and most LBAs \\citep[C09;][]{Overzier08}.  Recent results revealed that galaxies involved in galaxy interactions fall $\\ga$0.2 dex below the MZR of normal galaxies due to tidally induced large-scale gas inflow to the galaxies' central regions \\citep[e.g.,][]{Kewley06,Michel-Dansac08,Peeples09}.  Several $N$-body/SPH simulations have shown that major interactions drive starbursts and gas inflow from the outskirts of the H {\\sc i} progenitor disks \\citep[e.g.,][]{Mihos96,Rupke10}, also supporting this scenario.  We conclude arguing that recent interaction-induced inflow of gas, possibly coupled with a selective metal-rich gas loss driven by SN winds, may explain our findings and the known galaxy properties. Nevertheless, further work is needed to constrain this possible scenario.  In particular, future assessment of the H {\\sc i} gas properties and the star formation efficiency of GPs, as well as the behavior of effective yields with mass compared with models of chemical evolution, will shed new light on the relative importance of the above processes.  Our results allow us to further constrain the evolutionary status of the GPs. These galaxies, as well as the low-mass LBAs and the [\\oiii]-selected galaxies by \\citet{Salzer09} should be analyzed and compared in more detail to elucidate whether these rare objects are sharing similar evolutionary pathways.  Even if this is not the case, their properties suggest that these galaxies are snapshots of an extreme and short phase in their evolution.  They therefore offer the opportunity of studying in great detail the physical processes involved in the triggering and evolution of massive star formation and the chemical enrichment processes, under physical conditions approaching those in galaxies at higher redshifts.  Forthcoming analysis, based on high S/N intermediate/high-resolution spectroscopy and deep NIR imaging with the Gran Telescopio Canarias (GTC), will be used to better constrain the evolutionary status of the GPs."
    },
    "1004/1004.4298_arXiv.txt": {
        "abstract": "Nonsingular cosmologies are investigated in the framework of f(R) gravity within the first order formalism. General conditions for bounces in isotropic and homogeneous cosmology are presented. It is shown that only a quadratic curvature correction is needed to predict a bounce in a flat or to describe cyclic evolution in a curved dust-filled universe. Formalism for perturbations in these models is set up. In the simplest cases, the perturbations diverge at the turnover. Conditions to obtain smooth evolution are derived. ",
        "introduction": "Singularities occuring in General Relativity (GR) may be avoided in more fundamental frameworks where GR predictions are recovered as the low energy limit. Such frameworks could be provided by String or M-theory, or loop quantum gravity. In particular, the evolution of the Big Bang cosmology might be extended to a preceding contracting phase which, due to new physics relevant at high curvature or energy scales, turns into the expanding phase our universe is experiencing now \\cite{Veneziano:1991ek,Gasperini:1992em}, thus avoiding the Big Bang singularity (which in the inflationary picture \\cite{Starobinsky:1980te} is manifest rather as geodesic incompleteness \\cite{Borde:2001nh} than divergence of curvature invariants). These scenarios are called bouncing cosmologies \\cite{Novello:2008ra}.  Bouncing cosmologies can solve the horizon problem, but to replace inflation they should, among other things, also predict a viable, nearly scale invariant spectrum of perturbations. This tricky issue can be circumvented if a curvaton field is responsible for the generation of fluctuations \\cite{Enqvist:2001zp}. Otherwise, matching conditions are often required to track the evolution of the perturbations across the bounce \\cite{Martin:2004pm}. It has been noted that the curvature perturbation may become singular at the bounce, while the gravitational potential, whose growing mode usually persists in the post-bounce era, may have regular behavior \\cite{Durrer:2002jn}. A general solution for the perturbations supports these conclusions \\cite{Bozza:2005qg}. It is also well known that the features of the spectrum can depend sensitively upon the details of the dynamics of the bounce and the physics behind it.  Hence it is useful to consider explicit examples which allow one to scrutinize the possible behaviors of fluctuations at the bounce. However, it is necessary to violate energy conditions (EC). At least the strong EC must be broken to change the sign of the expansion rate, and the null EC cannot be respected if there is no curvature. This rather generically introduces pathologies that, though one may interpret them as only a shortcoming of the effective theory, hinder from reaching definite conclusions of the evolution of the spectra \\cite{Tsujikawa:2002qc,Martin:2001ue,Kallosh:2007ad}. An example is the perturbation divergence in pre-Big Bang cosmology \\cite{Kawai:1998ab} which can be shown to be an indication of an appearance of a ghost \\cite{Koivisto:2006ai}. To avoid EC violating matter fields, one can contemplate modifications of gravity that introduce no new degrees of freedom. This can be achieved with an action involving an infinite series of d'Alembertians acting on the curvature invariants in such a way that the propagator has no poles; these string-inspired nonlocal models\\footnote{A biscalar-tensor model \\cite{Nojiri:2007uq,Koivisto:2008xfa} motivated by the nonlocal cosmology based on inverse d'Alembertian gravity \\cite{Wetterich:1997bz,Deser:2007jk} may also accommodate bounces \\cite{Farajollahi:2010pn}, though their viability remains to be shown \\cite{Koivisto:2008dh}. Bounces in modified gravity were considered also in e.g. \\cite{Aref'eva:2007uk, Saidov:2010wx,Cai:2009in,Abramo:2009qk}.} have been shown to be ghost-free and asymptotically free at their Newtonian limit \\cite{Biswas:2005qr,Biswas:2006bs}.  In the present paper we consider the simpler case of second order $f(R)$ gravity, which has similar desiderable features. The fourth order metric $f(R)$ models correspond to scalar tensor theories of the form $\\mathcal{L}_{4th} = \\phi R + V(\\phi)$. By erasing the kinetic term implicit in the nonminimal coupling, one obtains the second order theory $\\mathcal{L}_{2nd} = \\phi R - \\frac{3}{4\\phi}(\\partial\\phi)^2 + V(\\phi)$. Though this can be problematic in view of the well-posedness of the Cauchy problem \\cite{LanahanTremblay:2007sg}, these simple models may avoid generic instabilities present in higher order theories \\cite{Koivisto:2009jn} and thus provide an effective description of low energy effects of quantum gravity. In particular, loop quantum gravity is expected to modify the cosmological dynamics at high curvatures without introducing new degrees of freedom. To obtain the quadratic density correction appearing in the particular scalar loop quantum cosmology, one may need to consider an infinite number of terms in the potential $V(\\phi)$, which may be interpreted to reflect the nonlocal nature of the underlying theory as discussed in Ref.\\cite{Olmo:2008nf}. Note also the recent extension of the framework \\cite{Fatibene:2010yc}.   Low curvature corrections in these so called Palatini-$f(R)$ theories have also been considered as alternatives to dark energy \\cite{Vollick:2003aw,Allemandi:2005qs}, but though they may produce viable background expansion, they generically fail to produce the observed matter power spectrum, at least for pressureless dust cosmology \\cite{Koivisto:2006ie,Li:2006ag,Uddin:2007gj}. Problems may appear also at microscopic level, as discussions of electron-electron scattering and Hydrogen atoms, seem to imply \\cite{Flanagan:2003rb,Olmo:2008ye}, see also \\cite{Iglesias:2007nv,Barausse:2007pn}. This may be due to need to reconsider the averaging problem in these models \\cite{Li:2008fa}, or the coupling of gravity to matter taking torsion and nonmetricity into account \\cite{Sotiriou:2006qn,Capozziello:2009mq,Sotiriou:2009xt}. One may adopt the approach of considering the formalism as an effective macroscopic description, and then new phenomenology can emerge from the potentially viable high curvature corrections. Studies of spherically symmetric systems show that the classic Solar system tests are passed by these models \\cite{Kainulainen:2006wz}, while the high curvature effects have interesting predictions for white dwarfs and neutron stars \\cite{Reijonen:2009hi}. Bounces have been suggested too \\cite{Barragan:2009sq,Barragan:2009sq}.  We briefly review the first order formalism approach to nonlinear curvature gravity in section \\ref{Palatini} where we also write and solve the cosmological background equations taking into account spatial curvature. We derive the conditions for bounces to occur and confirm them numerically. In section \\ref{Perturbations} we present the equations governing the evolution of perturbations in convenient forms, based on derivations in Refs.\\cite{Koivisto:2005yc,Koivisto:2007sq}. Corrections to and generalizations of previous literature are pointed for both the background and the fluctuation equations. In section \\ref{Conclusions} we discuss these results and their implications. ",
        "conclusions": "\\label{Conclusions}  Within the Palatini framework one may consider extensions of GR without introducing new degrees of freedom. Hence they may serve as useful toy models describing more completely the cosmological evolution, with motivations from e.g. loop quantum cosmology. We showed that there are nonsingular bouncing backgrounds in simple examples of such quantum corrected gravity models, and set up the formalism for the perturbations in these models in order to monitor the evolution of their spectra across a bounce. %  The models of the type (\\ref{type1}), characterized by $F\\rightarrow 0$ at the bounce, were found to feature singular behavior of perturbations in a flat, dust-filled universe. This may be cured in the curvature-dominated case, reflecting the fact that in the $K=0$ case the effective matter sources necessarily violate the null EC, whereas this is not the case if curvature is present. Since $F$ gives the sign of the graviton action, pathologies were to be expected at $F \\rightarrow 0$. At this point the conformal relation between Einstein and Jordan frames (\\ref{conformal}) becomes ill-defined. Let us note though that as the perturbations explode, their backreaction will render the perturbative system invalid, and as these nonlinear effects are very difficult to tackle in practice, we cannot say if there is a true singularity or not and whether the bounce occurs or not. In this light the problem is rather unpredictivity.  Obviously, it would be interesting to explore possible ways to obtain bounces of the type (\\ref{type2}). Such might be constructed by considering just more general functions $f(R)$ than the one involving solely a monomial correction.  Another way to obtain smooth evolution could be to include more general sources than completely pressureless fluids. Apart from allowing $w\\neq 0$, one may also consider stabilizing the system with entropic or anisotropic stresses. Indeed, this has previously proved successful in eliminating instabilities in matter perturbations in some dark energy models \\cite{Koivisto:2007sq,Koivisto:2005mm} (however, those models based on infrared gravity modifications may be otherwise problematical as discussed in the introduction). More realistically, also radiation would be included as a source and this changes the dynamics and possibly the conclusions. Our results can be directly applied to such more general models with possibly regular evolution. This is left for future studies.   As a concluding remark we note it is possible that a complete evolution of the background and structures of the universe is not amenable to classical description by second order differential equations, and one may have to take into account in a more nontrivial way the presently unknown physics at the very high curvature scales in order to provide a fully consistent coarse-grained picture of the cosmology that emerges. Meanwhile, the quest for the effective field equations for gravitational interactions, at both high and low curvature regimes, is ongoing.  \\appendix"
    },
    "1004/1004.0191_arXiv.txt": {
        "abstract": "A {\\em Chandra} X-ray imaging observation of the jet in Pictor A showed a feature that appears to be a flare that faded between 2000 and 2002. The feature was not detected in a follow-up observation in 2009. The jet itself is over 150 kpc long and a kpc wide, so finding year-long variability is surprising. Assuming a synchrotron origin of the observed high-energy photons and a minimum energy condition for the outflow, the synchrotron loss time of the X-ray emitting electrons is of order 1200 yr, which is much longer than the observed variability timescale. This leads to the possibility that the variable X-ray emission arises from a very small sub-volume of the jet, characterized by magnetic field that is substantially larger than the average over the jet. ",
        "introduction": "Surveys using the {\\em Chandra} X-ray Observatory have been very successful at detecting knots in quasar jets \\citep[e.g.][]{2005ApJS..156...13M,2004ApJ...608..698S}. Two mechanisms are generally cited when explaining the origin of the X-rays: synchrotron emission from high energy electrons and inverse Compton scattering of Cosmic Microwave Background photons by low energy electrons in knots in relativistic bulk motion along the line of sight \\cite[see][for reviews]{2006ARA&A..44..463H,2009A&ARv..17....1W}. In the synchrotron case, variability might be expected on a time scale of years due to the short loss times for electrons of sufficient energy to produce X-rays.  In the inverse Compton model, however, variability time scales could be much larger, with rise times longer than the knot light crossing time and decay times dictated by radiative lifetimes larger than $10^4$ yr.  In both cases, adiabatic loss time scales should be of order the light crossing time.  Variability studies have generally been limited to nearby objects such as Cen A \\citep{2010ApJ...708..675G} and M 87 \\citep{2009ApJ...699..305H}. In the case of M 87, the HST-1 knot was found to flare by a factor of 50 over a period of 5 years, becoming much brighter than the nucleus. While the M 87 jet is about 7 pc across, the upstream end of the HST-1 knot is only 1.9 pc across, so a secular increase over 5 yr is reasonable. However, the increase is not smooth and there is a factor of 2 drop in less than half a year, so the variation could be up to a factor of 10 faster than the light travel time across the emission region. \\citet{2003ApJ...586L..41H} explain this by Doppler boosting by a factor of $\\delta \\approx 5$. Both Cen A and M 87 are FR I radio galaxies, which are more common -- and thus more nearby -- than FR II radio galaxies and quasars with higher power jets. Knots in high power jets are larger but less well resolved and variability has not been reported. For example, one attempt to find variability in the X-ray emission from a jet in a quasar, 3C 273, yielded a null result \\citep{2006ApJ...648..900J}.  Pictor A is a FR II radio galaxy at a redshift of 0.035. For $H_0 = 70.5$ km/s/Mpc, 1\\arcsec\\ corresponds to about 700 pc. Its pencil-like X-ray jet was found in {\\em Chandra} observations by \\citet[][hereafter, WYS01]{2001ApJ...547..740W}, extending 1.9\\arcmin\\ from the core, oriented toward the partially resolved western hot spot 4.2\\arcmin\\ from the core. They estimated that the jet's typical width was about 2.0\\arcsec, or 1.4 kpc. Based on the assumption that the jet X-ray emission was produced by the inverse Compton process, and using the radio galaxy arm-length asymmetry to infer an angle to the line of sight, they suggested that a plausible scenario was that the magnetic field in the jet was 2 $\\mu$G, about a factor 6 below the equipartition value, and that the Doppler factor in the jet was $\\sim 2.6$. However, subsequent tighter constraints on the X-ray photon index \\citep{2005MNRAS.363..649H} mean that the jet would be required to have a significantly steeper low-energy electron energy index than the lobes for an inverse Compton model to be viable. \\citet{2005MNRAS.363..649H} argue that this is implausible, and that the X-ray emission of the jet must be dominated by the synchrotron process; if so, the X-ray emission gives us no direct information on the Doppler factor or jet speed, other than ruling out highly beamed and/or sub-equipartition models. While the WYS01 estimate of the angle to the line of sight remains plausible, there are very large uncertainties associated with the use of the arm-length asymmetry to estimate $\\theta$. As Pictor A is a broad-line radio galaxy, low-luminosity unified models imply $\\theta \\la 45$\\arcdeg; the fact that the source is observed as a lobe-dominated object and the projected linear size of the source probably requires $\\theta \\ga 10$\\arcdeg.  Examining the archival {\\em Chandra} data through 2002 (observation IDs 345, 3090, and 4369), we found evidence for flares in the jet at $3\\sigma$ significance. The observations were not homogeneous, with one taken 1\\arcmin\\ from the core and the other centered at the hotspot about 4\\arcmin\\ away, so the point spread function (PSF) at the core was degraded. So, to improve the variability test and to check for new flares, we obtained new images.  Here, we report on the previous evidence for variability and the results from the new observations. A detailed analysis of the jet will be the subject of a later paper. ",
        "conclusions": "Using the formalism of \\cite{2009A&ARv..17....1W}, we estimate the average equipartition magnetic field in the jet to be $B_{\\rm eq} = 17 \\mu$G, similar to the value obtained by WYS01 in the absence of beaming. Further below we assume that the Pic A jet is indeed close to the minimum energy condition, as justified by observations of terminal hotspots and lobes in powerful radio sources \\citep[e.g.][]{2005ApJ...622..797K}. We measured the jet flux density to be 128 mJy at 1.4 GHz and assume $\\gamma_{\\rm min} = 100$, an electron energy index of 2 breaking to 3 before the X-ray band and that the radius of the jet is about 1\\arcsec, as estimated by WYS01. For these assumptions, Lorentz factors of the electrons that produce 1 keV X-rays is $\\gamma \\approx 7 \\times 10^7$. The corresponding synchrotron loss time, $\\tau \\approx 1200$ yr, is similar to the dynamical time scale $t_{\\rm dyn} = \\ell/c = 2000 $ yr, setting the knot size to the FWHM of an unresolved source (0.9\\arcsec), but both are much larger than the variability time scale of 2 yr.  Doppler beaming would reduce the estimated intrinsic value of $B_{\\rm eq}$ by a factor of $\\delta = 1/[\\Gamma (1-\\beta \\mu)]$, where $\\beta c$ is the speed of the jet moving at angle $\\theta = \\cos^{-1} \\mu$ to the line of sight, and $\\Gamma = (1-\\beta^2)^{-1/2}$. As $\\delta$ increases, the efficiency of inverse Compton scattering of microwave background photons increases rapidly, affecting the observed X-ray flux. In particular, WYS01 estimated that $\\delta \\approx 7$ would allow all the observed X-rays to be produced by the inverse Compton process with $B=B_{\\rm eq}$, although this should be taken as an upper limit if the jet X-rays are dominated by synchrotron emission as argued by \\citet{2005MNRAS.363..649H}. Increasing the jet beaming would also affect any X-ray emission resulting from the synchrotron process. For a fixed observed photon energy, $E$ (e.g. 1 keV), and the assumed minimum energy condition, the Lorentz factor of the electrons in the jet rest frame $\\gamma \\propto \\sqrt{E/B}$ is the same as that inferred in the observer frame, since the photon energy transforms by the same factor $\\delta$ as $B_{\\rm eq}$ scales. The radiative lifetimes, however, scale as $1/B^2$ and would increase as $\\delta^2$ in the rest frame or $\\delta$ in the observed frame, exacerbating the lifetime discrepancy.  Ignoring beaming for now, setting $t_{\\rm dyn}$ to 2 yr leads to an angular size of the emission region of $\\sim 1$ mas. For a knot that radiates a fraction $f$ of the total jet flux from a volume $V \\propto \\ell^3$, then $B_{\\rm eq} \\propto (f/\\ell^3)^{1/(3+\\alpha)}$. Again fixing the observed photon energy, then $\\gamma \\propto B^{-1/2}$ and $\\tau \\propto B^{-3/2} \\propto \\ell^{3\\xi}/f^\\xi$, where $\\xi = \\frac{3}{6+2\\alpha}$. Setting $\\alpha = 0.5$ for our assumed spectral shape gives $\\xi = 3/7$. So, if $\\ell$ is $10^3$ smaller, then $\\tau$ is a factor of $\\sim 7200$ smaller. For the 48\\arcsec\\ flare, $f \\approx 0.03$, so $\\tau \\sim 0.7$ yr, which is consistent with the observed flare. The local $B$ field would be $\\approx 2$ mG, $> 100$ times larger than the average in the jet. Thus, in order to explain flares on a timescale of years within a jet that is about 700 pc in diameter would require 1) X-ray emission from a small knot substantially smaller than the jet cross section and 2) local magnetic fields substantially larger than the average over the jet.  The energy loss timescale of individual relativistic electrons is unlikely to be the factor controlling the rate of fading of a flare since the energy release that is manifest as a flare is also likely to cause expansion of the emitting region. If we take this emitting region to be spherical, with radius $R$, and the magnetic field within it to be tangled, then the optically-thin synchrotron emission from that region scales as $R^{-4\\alpha -2}$ because expansion decreases both the magnetic field strength and the energies of the embedded relativistic particles.\\footnote{Different brightness decrease functions arise under different magnetic field configurations or expansion geometries, but most cases show rapid loss of power under adiabatic losses, at sufficiently high observing frequency.}  If the flaring region had $\\alpha = 0.5$, then an expansion by 50\\% would cause it to fade by a factor 5 and become undetectable against the larger-scale jet emission. Such an expansion is feasible if the flaring region is only 2~ly in diameter.  As VLBI observers expand their maps to include knots at one arcsec scales, some knots are showing mas-scale structure. For example, \\cite{2009ApJ...695..707G} found a very compact hot spot 40 kpc from the core of a quasar and \\cite{2010arXiv1002.2588C} detected the HST-1 knot in VLBI observations. Examining more cases might well prove to be a fruitful endeavor. Imaging at even 0.1\\arcsec\\ scales with the {\\em Hubble} Space Telescope and future generations of X-ray telescopes could well show that jet knots have substantial substructure."
    },
    "1004/1004.0158_arXiv.txt": {
        "abstract": "{Young stars are formed within dusty discs. The grains in the disc are originally of the same size as interstellar dust, around 0.1~$\\mu$m. Models predict that these grains will grow in size through coagulation. Observations of the silicate features at $\\mu$m wavelengths are consistent with growth from submicron to micron sizes in selected sources whereas the slope of the SED at longer wavelengths traces growth up to mm sizes and larger.} {We here look for a correlation between these two grain growth indicators.} {A large sample of T-Tauri and Herbig-Ae/Be stars, spread over five different star-forming regions, was observed with the Spitzer Space Telescope at 5--13~$\\mu$m; a subsample was observed at mm wavelengths. We complement this subsample with data from the literature to maximise the overlap between $\\mu$m and mm observations and search for correlations. Synthetic spectra are produced to determine which processes may produce the dust evolution.} {Dust disc masses in the range $<1$ to $7 \\times 10^{-4}$~M$_\\odot$ are obtained. Most sources have a mm spectral slope consistent with grain growth. There is a tentative correlation between the 10-$\\mu$m silicate feature and the slope of the SED between 1 and 3~mm. The observed sources seem to be grouped per star-forming region in the $\\mu$m-vs-mm diagram. The modelling results show that first the 10-$\\mu$m feature becomes flatter and subsequently the mm slope becomes shallower. Grain size distributions shallower than that of the ISM and/or bright central stars are required to explain specific features. Settling of larger grains towards the disc midplane affects 10-$\\mu$m feature, but hardly the mm slope.} {The tentative correlation between the strength of the 10-$\\mu$m feature and the mm slope suggests that the inner and outer disc evolve simultaneously. Dust with a mass dominated by $\\sim$mm-sized grains is required to explain the shallowest mm slopes. Other processes besides grain growth may also be responsible for the removal of small grains. Observations with future, more sensitive telescopes are required to provide the necessary statistics to study these processes of disc and dust evolution.} ",
        "introduction": "\\label{sect: introduction}  A long-standing problem in planet formation is how tiny interstellar dust particles of less than a micron in size coagulate and grow to eventually form planets, thousands of kilometres in size. It is in the very nature of this field that it has to be studied at various levels, since different physical processes dominate during the various phases. The first steps, which lead to dust grains of about a decimetre in size, are studied both in the laboratory and with computer simulations \\citep[see][for detailed reviews]{dominik:2007, blum:2008}. Local concentrations of boulders and subsequent gravitational collapse may then lead to the formation of planetesimals several hundreds of kilometres in size \\citep[e.g.,][]{johansen:2007}. This paper focuses on the observational signatures of (sub)micron-sized grains up to centimetre-sized pebbles.  The InfraRed Spectrograph (IRS) on-board the Spitzer Space Telescope has provided a wealth of mid-infrared (5--40 $\\mu$m) spectra from discs around pre-main-sequence stars \\citep[e.g.][]{kessler-silacci:2006,furlan:2006}. The spectra of these objects are often dominated by silicate emission features at 10 and 20~$\\mu$m. In young stellar objects, these features are formed in the upper atmosphere of the hot inner disc. The varying strength and the shape of these features can be naturally explained by different grain sizes in the upper layers of the inner disc, with strong, pointed features being representative of $\\sim$0.1~$\\mu$m-sized grains and flatter features coming from dust grains of several $\\mu$m in size \\citep{kessler-silacci:2006}. These results confirm earlier results from the Infrared Space Observatory \\citep[][]{bouwman:2001, meeus:2001} and from ground-based observations \\citep[e.g.,][]{przygodda:2003}. It has been suggested that crystallisation has a similar effect on the 10-$\\mu$m feature as grain growth \\citep[e.g.,][]{honda:2003, meeus:2003}. However, this effect is minimal and the dominating factor for the strength and shape of the 10-$\\mu$m feature is the characteristic grain size \\citep{olofsson:2009}.  Because the 10-$\\mu$m feature only probes the surface layers of the inner disc, a stronger, more peaked feature could also be due to the settling of larger, micron-sized grains towards the mid-plane. As the larger grains settle and the small ones remain suspended in the upper layers, the surface becomes dominated by small grains, creating a strong silicate band. \\citet{dullemond:2008} investigated this idea through theoretical models. They find that settling can in principle explain the different shapes of the\t10-$\\mu$m feature, but only in quite specific cases, so that overall grain growth is still the most likely explanation for the flattening of these features. Recent interferometric observations of the 10-$\\mu$m spectral region in discs around $\\sim$1 and 2-3-M$_\\odot$ objects show that the grains closer to the central star are both larger and more crystalline than those further out in the disc \\citep[see, e.g., the recent review by][]{vanboekel:2008}. Hence, the evolution of the 10-$\\mu$m feature may be caused by a combination of grain growth and crystallisation and appears to progress from the inner disc outwards. On the other hand, analysis of the longer wavelenth mid-infrared crystalline features indicates significant growth and crystallisation in the outer disc as well \\citep{olofsson:2009}.  Whereas the mid-infrared region potentially provides information on the growth of grains from interstellar, submicron sizes to sizes of several microns, the growth to larger sizes can only be probed by submillimetre (submm), millimetre (mm), and occasionally centimetre (cm) observations. Ground-breaking work was done by \\citet{beckwith:1990} and \\citet{beckwith:1991}, both analytically studying the emission of dust grains and obtaining the first submm slopes by observing a large sample of young stellar objects at mm wavelengths. More recently, \\citet{andrews:2005} performed a sensitive single-dish submm continuum survey of 153 young stellar objects in the Taurus-Auriga star-formation region, including a large amount of archival and literature data. They found that the submm slope between 350~$\\mu$m and 1.3~mm could be well described by $\\alpha = 2.0 \\pm 0.5$, where $F_\\nu \\propto \\nu^{\\alpha}$, while the value for the interstellar medium is $\\alpha \\approx 3.7$ \\citep[cf.][]{draine:2006}. \\citet{andrews:2005} interpreted this shallow slope as a combined effect of a contribution from optically thick regions in the disc and grain growth. It should be noted, however, that the sources in this study were spatially unresolved, and the (sub)mm emission may have a significant contribution from surrounding (envelope) material. More recently, interferometric studies of several dozen T-Tauri stars gave values of $\\alpha \\lesssim 3.0$ \\citep[][]{rodmann:2006,andrews:2007,lommen:2007}. Similar results were found for a number of more massive Herbig-Ae/Be stars \\citep[e.g.,][]{natta:2004}. From this mm slope one can estimate the opacity index $\\beta \\approx (1 + \\Delta) \\times (\\alpha - 2)$, where $\\Delta$ is found to be $\\sim$0.20 \\citep{rodmann:2006, lommen:2007}, and values of $\\beta \\approx 1.0$ for $\\lambda \\gtrsim 1$~mm\twere found. Such a slope can be naturally explained by a significant fraction of grains at least several mm in size present in the discs \\citep{draine:2006}.  A subsample of the sources observed by \\citet{lommen:2007} overlapped with the {\\em Spitzer} Infrared Spectrograph (IRS) observations published by \\citet{kessler-silacci:2006} and \\citet{lommen:2007} found a tentative correlation between the mm slope of the spectral energy distribution (SED) and the strength and shape of the 10-$\\mu$m silicate feature for these sources. Note that the 10-$\\mu$m feature primarily probes the hot surface layers of the inner disc, whereas the (sub)mm observations provide information of the cold mid-plane of the outer disc. A correlation between the two is therefore not obvious at all and a confirmation of this correlation would give very valuable information on the processes of dust growth in protoplanetary discs, as it would imply that grain growth from submicron to mm sizes is both fast and occurs simultaneously throughout the whole disc.  \\citet{acke:2004a} calculated the (sub)mm spectral indices of 26 Herbig-Ae/Be stars, for which the infrared SED could also be determined. They found a correlation between the strength of the ratio of the near- to mid-infrared excess and the slope of the (sub)mm energy distribution for these sources, which they attributed to a correlation between the disc geometry (flared versus self-shadowed) and the size of the grains in the disc. However, the authors did not find a correlation between the strength and the shape of the 10-$\\mu$m silicate feature and the (sub)mm spectral index \\citep[see also][]{acke:2004b}.  The aim of this paper is to investigate the tentative correlation between the strength and shape of the 10-$\\mu$m silicate feature and the spectral slope in the (sub)mm regime, found by \\citet{lommen:2007}, for a larger sample. A subsample of sources studied with the Spitzer IRS were observed with mm and cm interferometers (Section 2). Interferometers were used to ascertain that the emission is dominated by disc emission, since extended emission from surrounding material will be filtered out. Also, spatially resolving the disc ensures that the emission is not optically thick \\citep[e.g.,][]{natta:2004}. The results of the observations, including dust disc masses and mm slopes, are shown in Section 3, and in Section 4 we present model results for discs. The observations and models are compared and discussed in Section 5; conclusions are formulated in Section 6. ",
        "conclusions": "\\label{sect: conclusions}  We observed five binary systems and 35 single T-Tauri stars (of which one turned out to harbour two sources as well) at infrared and mm wavelengths. New {\\em Spitzer} IRS spectra of the wavelength region containing the 10-$\\mu$m silicate feature are presented for 13 sources and the slopes in the mm regime of the SED are obtained, also for 13 sources. Furthermore, dust disc masses are obtained for 21 new sources, including strict 3$\\sigma$ upper limits for nine of the sources. Combining our new observations with data from the literature, a tentative correlation between the strength of the 10-$\\mu$m feature and the mm slope is found for a sample of in total 31 T-Tauri stars. This seems to suggest that the inner and outer discs do not evolve independently.  Modelling protoplanetary discs with varying geometries and grain size distributions indicates that grain growth alone cannot explain the observed evolution in the strength of the 10-$\\mu$m feature and the mm slope of the SED. It would rather seem that as the maximum grain size increases, also the power-law slope of the grain size distribution becomes shallower. \\citet{ricci:2010} also found that for the majority of the sources in their sample a grain size distribution flatter than that of the interstellar medium is required \\citep[see also][]{natta:2007}. This is an indication that a power-law grain size distribution with a fixed minimum grain size is too simple a representation of the dust in protoplanetary discs.  The sample contains only one cold or transitional disc for which the mm slope in the SED and the strength of the 10-$\\mu$m silicate feature could be obtained properly. This source, CS~Cha, ends up rather in the centre of the 10-$\\mu$m vs mm-slope diagram, indicating that grain growth is probably not the source of the removal of dust from the inner disc. Transitional discs are often relatively weak in the (sub)mm regime and a new generation of telescopes is required to determine the mm slope for more of these sources.  Although all 10-$\\mu$m features used in this work were obtained using the {\\em Spitzer} Space Telescope, the mm fluxes and hence the mm slopes were determined with a large number of different telescopes. As this potentially leads to systematic effects, it will be necessary to observe a statistically large enough sample with only one or maybe two (sub)mm telescopes. Ideally, about a dozen sources per star-forming regions for at least four regions should be observed with the same telescope. However, due to the relative weakness of the sources under consideration, the scarcity of available telescope time, and the spread of young star-forming regions over the sky, this is currently hard to achieve. The Atacama Large Millimeter/submillimeter Array (ALMA) will, upon completion, be able to study an order of magnitude more sources than presented here, providing us with the statistics to study interesting relations such as the grouping of YSOs by star-forming region in the $\\mu$m-vs-mm diagram. ALMA will also be able to resolve most of the sources studied here.  This work has mainly focused on observations at 1 and 3~mm, at which grains with sizes of the order of 1~cm can be studied \\citep{draine:2006}. The detection of larger grains, or pebbles, in protoplanetary discs, requires observations at cm wavelengths. However, there may be other sources of emission at cm wavelengths, such as magnetic fields or stellar winds, and it is necessary to monitor sources over extended periods of time to disentangle the different emission mechanisms \\citep{wilner:2005, lommen:2009}. The maximum bandwidth of the ATCA was recently improved by a factor of 16 with the implementation of the Compact Array Broadband Backend (CABB). This will for the first time allow the survey of large samples of young stellar objects in the southern hemisphere at cm wavelengths.  Thus, in five to ten years from now we will not only be studying how grain growth occurs in protoplanetary discs, but we will actually be able to pinpoint the locations in the discs where the formation of pebbles and boulders, the precursors to planetesimals and planets, is taking place. Furthermore, on-going observations at cm wavelengths will give us a stronger handle on the processes taking place in and around the discs, affecting the growth of grains and the formation of planets."
    },
    "1004/1004.0644_arXiv.txt": {
        "abstract": "{} {Studies of the debris disc phenomenon have shown that most systems are analogous to the Edgeworth-Kuiper Belt (EKB). In this study we aim to determine how many of the IRAS 25$\\mu$m excesses towards A stars, which may be indicative of asteroid belt analogues, are real, and investigate where the dust must lie and so build up a picture of what these systems are like.} {We observe using ground-based mid-infrared imaging with TIMMI2, VISIR, Michelle and TReCS a sample of A and B-type main sequence stars previously reported as having mid-infrared excess. We combine modelling of the emission spectrum from multi-wavelength photometry with a modelling technique designed to constrain the radial extent of emission in mid-infrared imaging to constrain the possible location of the debris.} { We independently confirm the presence of warm dust around three of the candidates: HD 3003, HD 80950 and $\\eta$ Tel.  For the binary HD3003 a stability analysis indicates the dust is either circumstellar and lying at $\\sim$ 4 AU with the binary orbiting at $>$ 14AU, or the dust lies in an unstable location; there is tentative evidence for temporal evolution of its excess emission on a $\\sim$20 year timescale.  For 7 of the targets we present quantitative limits on the location of dust around the star based on the unresolved imaging. We demonstrate that the disc around HD71155 must have multiple spatially distinct components at 2 and 60AU. We model the limits of current instrumentation to resolve debris disc emission and show that most of the known A star debris discs which could be readily resolved at 18$\\mu$m on 8m instruments have been resolved, but identify several that could be resolved with deep ($>$ 8 hours total) integrations (such as HD19356, HD139006 and HD102647). } {Limits from unresolved imaging can help distinguish between competing models of the disc emission, but resolved imaging is key to an unambiguous determination of the disc location. Modelling of the detection limits for extended emission can be useful for targeting future observational campaigns towards sources most likely to be resolved. MIRI on the JWST will be able to resolve the majority of the known A star debris disc population.  METIS on the E-ELT will provide the opportunity to explore the hot disc population more thoroughly by detecting extended emission down to where calibration accuracy limits disc detection through photometry alone, reaching levels below 1 zodi for stars within 10pc. } ",
        "introduction": "\\label{s:intro}  Analysis of the IRAS database over the last 20 years has shown that over 300 main sequence stars have dust discs around them.  This material is thought to be the debris left over at the end of the planet formation process \\citep[e.g.][]{mannings}. The spectral energy distribution (SED) of this excess in the best studied cases (e.g., Vega, $\\beta$ Pictoris, Fomalhaut, $\\epsilon$ Eridani) peaks longward of 60$\\mu$m implying that this dust is cool ($<$80K), and so resides in Edgeworth-Kuiper belt (EKB)-like regions in the systems. The EKB-like location and analogy is confirmed in the majority of cases where these discs have been resolved (in thermal emission, e.g. \\citealt{holland, greaves}; in scattered light imaging, e.g. \\citealt{kalas07, schneider, schneider09}), since the dust is shown to lie $>$40AU from the stars, and its short lifetime means that it must be continually replenished by the collisional destruction of km-sized planetesimals \\citep{wyattdent}. The inner 40AU radius hole is thus thought to arise from clearing by an unseen planetary system, the existence of which is supported by the presence of clumps and asymmetries seen in the structure of the dust rings \\citep[e.g.][]{wyatt99, greaves98, augereau}.  On closer examination the inner regions of known cool debris disc systems are much more complex than simply 'dust-free inner holes'. $\\beta$ Pictoris has (a relatively small amount of) resolved dust in this inner region \\citep{lagage, telesco, boccaletti}, thought to be there because this is a young (12Myr, \\citealt{zuckermanEA}) transitional system in which these regions have yet to be fully cleared by planet formation processes. \\citet{absil} have recently presented interferometric data showing Vega (thought to be around 380-500Myr old, \\citealt{peterson}) is likely to possess extended dust emission within 8AU, with evidence for a similar warm dust propulation around Fomalhaut \\citep{absil09}, and \\citet{difolco} have also recently presented evidence for hot dust around the 10Gyr old $\\tau$ Ceti. Multiple-component discs mirror our own Solar System, with our debris disc concentrated in the asteroid belt and EKB.  It is possible that more sources with known cold discs have dust in the inner system, but the difficulty of separating hot dust emission from the stellar photosphere often limits detection. \\citet{chen09} presented a sample of 11 debris systems believed to have multiple dust belts, although confirmation of the multiple components via resolved imaging is required for several of these stars.  The presence of excess in the mid-infrared range is of particular interest.  The temperature of dust emitting at 24$\\mu$m suggests that it should lie close to the star, in regions of a few to a few tens of AU.  These are regions in which we might expect planets to reside (see above), and so the origin of the emission and possible links with any cold debris in the system must be explored.  Surveys suggest that 24$\\mu$m excess may be common, with around half of main sequence stars that exhibit excess mid-infrared emission in the IRAS database \\citep{mannings} having  an excess at 25$\\mu$m only \\citep{zuckerman}. However, excesses taken from the IRAS database cannot be used at face value.  \\citet{song02}, who searched the IRAS database for excess emission towards M-type stars, noted that when searching a large number of stars for excesses close to the detection threshold, a number of false positives must be expected due to noise.  There have also been a few instances in which the IRAS excess has been shown to be attributed to background objects that fall within the relatively large IRAS beams ($>$30''; such objects range from highly reddened carbon stars or Class II YSO's \\citealt{lisse}, to distant galaxies \\citealt{sheret}). Another possible source of mid-infrared excess emission is reflection nebulosity \\citep{kalas}.  Indeed it is now routine for papers discussing the excess sources found by IRAS to address the possibility that some of these are bogus debris discs \\citep{moor, rhee}.  Recent statistical studies using Spitzer data have revealed a large sample of A stars with excess emission at 24 and/or 70$\\mu$m \\citep{rieke,su}.  Spitzer observations benefit from smaller beamsizes and higher resolution compared to IRAS, improving the reliability of excess measurements.  The 24 and 70$\\mu$m excesses around A stars have a wide variation in levels amongst systems of similar ages, but overall there is a decrease in the upper envelope of excess inversely proportional to time \\citep{rieke,su}.  These features can be interpreted in terms of a steady-state evolution of belt-like planetesimal discs in a collisional cascade, where the fall-off with time is due to the collisional grinding away of material and the variation in excess levels between systems of similar ages can be reproduced by variations in initial disc mass and planetesimal belt radii \\citep{wyattsmith07}. Although recent work by \\citet{Currie} has shown that excess emission at 24$\\mu$m around A stars increases from 5--10 Myr and peaks around 10--15 Myr before declining with age, which is not predicted by the steady-state model, this can be explained by the delayed formation of Pluto-sized bodies in the disc \\citep{kenyon04II}. It is only when Pluto-sized bodies are formed that the orbits of planetesimals are stirred to high relative velocities and the steady-state collisional cascade can begin.  However this model does not account for the possibility of multiple belts (see above).  There have also been other models proposed to explain the variance in excess levels between similarly aged stars and the variety in the SED slopes of those stars with excess at both 24 and 70$\\mu$m in \\citet{rieke}, which can be interpreted as evidence for stochastic evolution \\citep{rhee}.  A collision similar to the Earth-Moon forming massive collision has been proposed as the explanation for the large excess and spectral features observed around HD172555 \\citep{lisse09}. Similarly short-lived dust production origins have been proposed for dust found close to the central star around several Sun-like stars. For these systems a steady-state collisional cascade production  from a spatially coincident planetesimal belt cannot explain the levels of excess \\citep{wyattsmith06, lohne}.  For debris discs with a cold dust population in addition to hot dust emitting at 24$\\mu$m the outer planetesimal belt could be feeding the hot dust population (as has been proposed for $\\eta$ Corvi, \\citealt{wyattsmith06, smithhot, smithMIDI}).  However, the mechanism that might be transporting the dust from cold outer regions to the hot inner locations observed at 24$\\mu$m is as yet unclear (possibilities include dynamical scattering by a migrating planet, see e.g. \\citealt{gomes, booth}).  To tackle these issues regarding the origin of the 24$\\mu$m dust emission, and in particular explore what its presence may reveal about planet formation and as yet undetected planetary populations, this paper looks at the constraints on the true dust distribution around a sample of A stars with 24$\\mu$m excess. This can be assessed indirectly from SED fitting to multi-wavelength infrared photometry, but uncertainties arising from degeneracies in dust model fitting and the possibility of multiple temperatures of dust mean that determining radial location from SED fitting alone is challenging (see discussions of individual sources in section 4). Resolved imaging provides more direct constraints on the dust location.   This paper is structured as follows; In \\S \\ref{s:sample} the sample selection is described. In \\S \\ref{s:obsext} we describe the various observational and analysis techniques employed for the observations, with the results and discussion of individual sources presented in \\S \\ref{s:res}.  An extension modelling technique is used to explore which of the A star discs in the literature may be fruitful subjects of future imaging in \\S \\ref{s:pred}. The implications of these and the observational results are discussed in \\S \\ref{s:dis}. Conclusions are in \\S \\ref{s:conc}. ",
        "conclusions": "\\label{s:dis}  This sample contains 7 sources with excess infrared emission confirmed either in this paper or with Spitzer data \\citep{rieke, chen06}, and a further 3 sources with excess emission in IRAS requiring confirmation.  The SED fitting indicates that these objects are surrounded by dust at a distance of between 2-60 AU (or alternatively two temperatures of dust at 2 and 61AU for HD71155 and 4 and 24 AU for $\\eta$ Tel).  These regions are those in which we might expect the formation of giant planets, and so it is important consider how the existence of this dust emitting in the mid-infrared adds to our current understanding of dust distributions in circumstellar regions.  \\citet{rieke} looked at a sample of 266 A-type stars between 5-626 Myr old with MIPS at 24 and 70$\\mu$m and examined the relationship between fractional excess and the age of the central star.  They found that the upper limit of excess emission generally fell off as $time^{-1}$ for the stars with detected excess. Assuming the fits to the SED profiles presented in this paper and plotting the predicted 24$\\mu$m excess emission compared to stellar flux versus age it is clear that the results presented here are in-line with this relationship (Figure \\ref{riekecomp}).  Combining this with the age spread in sources suggests we have a representative sample of A star debris discs.  \\begin{figure} \\includegraphics[width=8cm]{rieke_comp.ps} \\caption{\\label{riekecomp} Plot of age versus fractional excess from SED fitting at 24 $\\mu$m.  The line shows a $t^{-1}$ dependence, taken from \\citet{rieke}, and this plot can be directly compared to their Figure 2.} \\end{figure}  HD71155 and $\\eta$ Tel, the two systems with resolved dust populations, have both been found to have multiple disc components. \\citet{chen09} summarised the debris systems which have, through SED fitting or resolved imaging, been identified as multiple component discs.  Adding HD71155 to this sample brings the total number to 12 known ``Solar System analogues'' (as defined in \\citealt{chen09}, systems with multiple component discs).  The FEPS survey on Spitzer surveyed 328 solar-like stars finding $\\sim$10\\% have 70$\\mu$m emission indicating the presence of cold debris, and of these 1/3 have SEDs that are best fit by multiple temperature excesses \\citep{hillenbrand}.    From our sample of 10 debris disc sources (3 of which, $\\lambda$ Gem, HD23281 and HD141795,  still require photometric confirmation), 2 are multi-component discs (20\\%; consistent with the 1/3 rate given our small sample size).  These results are further evidence that multiple component or extended discs are common.  The traditional view of debris discs is that of a ring of planetesimals residing outside any planetary system producing cold dust, analogous to the Edgeworth-Kuiper belt (EKB) in the Solar System. Multiple component discs could be seen as analogous to the Solar System, although of course the extrasolar discs detected to date are much brighter than the asteroid belt and EKB ($L_{\\rm{IR}}/L_\\star \\ge 10^{-4}$ for detected extrasolar discs;  asteroid and Edgeworth-Kuiper belts $L_{\\rm{IR}}/L_\\star = 10^{-8}-10^{-7}$ and $10^{-7}-10^{-6}$ respectively, \\citealt{dermott02,stern96aa}).  For sources with large amounts of emission from close to the star, the origin of the dust is uncertain.  \\citet{wyattsmith06} presented a model based on the collisional evolution models of \\citet{dd03} that predicted a maximum dust luminosity dependent on age for a disc at a given radius.  This maximum brightness occurs when discs just reach collisional equilibrium, in which the size distribution of the planetesimals in a disc is fixed as mass is transferred down through a cascade of collisions to smaller and smaller sizes until they are removed by radiation pressure.  Very massive discs process their mass very quickly and are therefore short-lived, whereas sparse discs take a long time to reach collisional equilibrium but are not very luminous.  This model has been shown to accurately recreate the A star debris disc population observed with Spitzer \\citep{rieke} under assumptions of a distribution of initial disc mass, radius, and current age \\citep{wyattsmith07}.  The predicted maximum luminosity for each disc (or disc component) in this study is given in Table \\ref{results} ($f_{\\rm{max}}= 1.2\\times10^{-6}r^{7/3}t_{\\rm{age}}^{-1}$ from equation 20 in \\citealt{wyattsmith07} where $r$ is the radius of the dust belt from our fits and $t_{\\rm{age}}$ is the age of the system as given in Table \\ref{sample}) . Uncertainties in the model parameters mean that only if $f_{\\rm{IR}}/f_{\\rm{max}} > 1000 $ do we take this as evidence of transient emission (level of excess cannot be produced by a collisional cascade).  Most of our targets have $f_{\\rm{IR}}/f_{\\rm{max}} < 1000$, with $\\lambda$ Gem and the inner disc of HD71155 exceeding this value, and HD141795 having $f_{\\rm{IR}}/f_{\\rm{max}} \\sim 550$.  Of these possibly transient sources only HD71155 has confirmed excess emission (see Section 4.2), and confirmation of the other discs would be required before speculating on their origin.  At an offset of 2AU around HD71155 the collisional lifetime of bodies is short and thus we interpret the emission as evidence for a transient dust producing event.  A recent massive collision in an otherwise quiescent disc could produce a short lived increase in excess, although as such events are likely to be rare (and the resulting dust grains have short lifetimes) the probability of witnessing such an event is low (see discussion in \\citealt{wyattsmith06}).  At an age of 169Myr (Table \\ref{sample}), ongoing terrestrial planet formation could be responsible for the emission \\citep{kenyon04}, with collisions between planetary embryos resulting in large amounts of dust production.  The cooler belt at 61AU could represent a parent population of the hot dust emission as well as producing spatially coincident dust dominating the excess at longer wavelengths (a similar possibility exists for the older Sun-like star  $\\eta$ Corvi, \\citealt{smithhot, smithMIDI}).  However, the transport mechanism to get the dust to 2AU is unknown.  A dynamical instability like that thought to have caused the Late Heavy Bombardment (LHB) in the Solar System \\citep[triggered by the migration of Jupiter/Saturn, see e.g.][]{gomes, levison}, during which large amounts of debris from the EKB was thrown into the inner Solar System, could be responsible. \\citet{booth} concluded that such an event occured around at most 12\\% of Sun-like stars.  Around A stars, where the 70$\\mu$m excess has been shown to exhibit a fall-off proportional to time like the 24$\\mu$m emission \\citep[albeit with a longer decay time, ][]{su}, the statistics are less clear.  Resolving the outer disc component would allow further examination of a possible link between the two populations (for example, emission spread inwards towards the hot dust population rather than confined to a narrow could be evidence for a link).  A further possibility is that the emission arises from dust grains not produced in collisions but in the sublimation of a population of comets or one Super-comet.  These possibilities were explored as the origin of the hot dust population around HD69830 by \\citet{beichman05}, who concluded that the continuous generation of small grains by a population of comets would require too large a mass reservoir to be the likely origin of the dust. A single massive comet (Sedna-sized in the case of HD69830) could release small dust grains over a few Myr if captured into a close orbit \\citet{beichman05}.  This mechanism could be responsible for the 24$\\mu$m dust population in more systems.   Dust emission that is from a transient event will necessarily show temporal variation.  The difference between the IRAS photometry on HD3003 and the measurements taken with MIPS and those presented in this paper (see section 4.1) could be evidence of such evolution. If real, this variance could be a reflection of the binary nature of the system, with the orientation of the secondary as it proceeds on its orbit changing the overall illumination of the system.  A determination of the orbit of the binary will allow this possibility to be checked.  Alternatively, if the temporal variance reflects changing levels of dust or changing dust distributions then a transient origin is more likely.  Taking the assumed circumstellar radius of 4AU for the dust (see section 4.1) $f_{\\rm{IR}}/f_{\\rm{max}} = 379$, a high level but one at which we would not state the emission must be transient conclusively (see above and detailed discussion of uncertainties in \\citealt{wyattsmith06}).   However, if the dust is of transient origin, the dust location needs not be stable, and the constraint of $>$14.4AU for the binary semi-major axis need not hold.  In this situation the determination of the orbit of the binary could again greatly improve our understanding of this system.  Alternatively resolving the dust distribution (as would be possible with METIS on the E-ELT, see Figure \\ref{futurepred}) could also inform our models of the stability of the system, and thus the likely origin of the dust.  Of the sources considered in this paper only $\\eta$ Tel and HD71155 have had the location of the dust population confirmed by resolved imaging.  For the remaining discs constraints have been placed on the dust location, but degeneracies in the SED fitting in particular (summarised in Section 5) mean that resolved imaging is required to determine the true dust distributions in these systems, and so constrain models for the dust origin particularly where transience is inferred.  The predictions in section 5 can be used to target sources most likely to be resolved with currently available instruments, but as shown in Figure \\ref{futurepred} most currently known disc targets will require the use of MIRI to detect faint levels of extended emission or the high resolution of METIS on the E-ELT to resolve emission on small scales. High resolution will also be important for the detection of sub-structure in the discs which could indicate the presence of planets which will be important for ascertaining the nature of these systems and distinguishing between models for the origin of the dust.  Evidence for planets in the dust distribution of debris disc systems include: sharp disc edges (as seen around Fomalhaut, \\citealt{kalas05}); clumps (Vega, \\citealt{wyatt03}; similar structure may be observable in the EKB because of Neptune's resonant Plutino population if the disc were brighter); warps ($\\beta$ Pic, \\citealt{augereau}); and asymmetries (HR4796A, \\citealt{wyatt99})."
    },
    "1004/1004.0702_arXiv.txt": {
        "abstract": "We propose a two-temperature radial inflow-outflow model near Sgr A* with self-consistent feeding and conduction. Stellar winds from individual stars are considered to find the rates of mass injection and energy injection. These source terms help to partially eliminate the boundary conditions on the inflow. Electron thermal conduction is crucial for inhibiting the accretion. Energy diffuses out from several gravitational radii, unbinding more gas at several arcseconds and limiting the accretion rate to $<1\\%$ of Bondi rate. We successfully fit the X-Ray surface brightness profile found from the extensive Chandra observations and reveal the X-Ray point source in the center. The super-resolution technique allows us to infer the presence and estimate the unabsorbed luminosity $L\\approx4\\cdot10^{32}{\\rm erg~s^{-1}}$ of the point source. The employed relativistic heat capacity and direct heating of electrons naturally lead to low electron temperature $T_e\\approx 4\\cdot10^{10}$~K near the black hole. Within the same model we fit $86$~GHz optically thick emission and obtain the order of magnitude agreement of Faraday rotation measure, thus achieving a single accretion model suitable at all radii. ",
        "introduction": "Our Galaxy hosts a supermassive black hole (BH) with a mass $M=4.5\\cdot10^6 M_\\odot$ \\citep{ghez,reid} at a distance $R=8.4$~kpc. The BH exhibits low luminosity state probably due to inefficient feeding and cooling. Almost all available matter outflows from the region, whereas only the small fraction accretes \\citep{quataert_wind}. This feeding region within several arcseconds contains X-Ray emitting gas, but some X-Rays are expected from a synchrotron self-Compton (SSC) or synchrotron source from accretion at several Schwarzschild radii $r_{\\rm g}$. The study of X-Rays offers a unique opportunity to test the full range of accretion scales from several $''$ to $r_{\\rm g}=10^{-5}$$''$ and construct a single model.  Modeling the accretion flow with such a huge range of scales is a challenge.  3D SPH simulations are performed in the outer region between $1''$ and $10''$ \\citep{rockefeller,cuadra}. Latest MHD simulations \\citep{sharma_spherical} are limited to $3$ orders of magnitude in radius and axial symmetry. Only the one-dimensional calculation \\citep{quataert_wind} can in principle resolve the flow everywhere. Thus, 1D modeling is the approach we adopt extending it down to the BH horizon.  We analyze the quiescent observations \\citep{muno_diff} of X-Ray emission from central several arcseconds around Sgr A* in \\S \\ref{s_obs}. The total exposure is 25 times longer compared to previously analyzed data \\citep{baganoff}. The super-resolution processing based on spacecraft dithering helps resolving sub-pixel scales. The up-to-date data on stellar wind emitters are summarized in \\S \\ref{s_winds}. We smooth matter ejection rates of individual stars over radius and sum them into a single feeding rate, also properly averaging the wind velocity. This presents a significant improvement over an ad-hoc feeding in \\citet{quataert_wind}. The dynamical two-temperature equations are derived in \\S \\ref{s_dynam}. We consider the electron conduction the main energy transport mechanism, approximating the unsaturated heat flux by a simple formula. The Bondi flow \\citep{bondi52} without heat transport overestimates the X-Ray luminosity by a factor of $10^3.$ The other important effects considered are the relativistic heat capacity of electrons and superadiabatic heating equivalent to entropy production. The ways to solve the resulting system of equations and corresponding results are presented in \\S \\ref{s_sol}. We employ the shooting method and find the minimum $\\chi^2$ fit for X-Ray surface brightness profile, simultaneously fitting $86$~GHz flux. The best fit model requires X-Ray point source. The viability of a non-cooling radial flow is examined. ",
        "conclusions": "\\label{s_sol} We solve the derived system of equations from the outer boundary of the feeding region at $14''=1.3\\cdot10^6r_{\\rm g}$ to the inner boundary at about $1.3 r_{\\rm g},$ thus covering $6$ orders of magnitude in radius. Such a huge dynamic range requires the special solution technique, the solution of a time-dependent system of equations \\citep{quataert_wind} not being an option. We employ the shooting method and find the smooth transonic solution through the inner sonic point at $\\sim3r_{\\rm g}.$ In the presence of conduction the point, where sound speed equals inflow velocity is not special anymore, and instead the point, where isothermal speed equals the inflow velocity, plays the role of transonic surface \\citep{johnson}. The system of equations is reduced to one temperature in the outflow by setting $T_e=T_p$ and adding the equations (\\ref{energy_electr}) and (\\ref{energy_prot}). The inner boundary is set at a point $r_{\\rm in}$, where $dT_e/dr=0$ in a non-conductive solution. Then for any non-zero conductivity the zero heat flux condition $dT_e/dr=0$ is enforced at $r_{\\rm in}.$ The outer boundary condition at $r_{\\rm out}$ is uncertain. It is natural to think the outflow would be transonic \\citep{lamers}, however, significant outer pressure may hold the gas in the subsonic regime near $r_{\\rm out}.$ The position of zero velocity stagnation point $r_{\\rm st}$ determines the accretion rate $\\dot{M}.$ Instead of setting the pressure at the outer boundary we regulate that pressure by setting temperature $T_{\\rm st}$ at the stagnation point. Thus, we have 4 independent variables in the fit: accretion rate $\\dot{M},$ temperature at stagnation point $T_{\\rm st},$ the ion heating rate $f_p$ and the normalization $N$ of the point source contribution. They are all found iteratively to minimize $\\chi^2.$ We also iteratively find the positions of sonic point and stagnation point. The positions of inner boundary and outer boundary are unchanged while solving the 4-point boundary value problem.  The observed surface brightness radial profile is the data we fit. We generate a surface brightness profile corresponding to the dynamical model by performing the optically thin ray tracing of X-Rays at a set of photon energies and projected distances from the BH. We employ the up-to-date bremsstrahlung emissivities (\\citet{gould} and errata) and account for the emission by heavy elements, excluding iron.  Solar metallicity interstellar absorption \\citep{morrison} is assumed with hydrogen column $N_{\\rm H}=10^{23}$cm${}^{-2}.$ The fluxes are convolved with the response of Chandra to find counts, then blurred with the energy-independent PSF (see Figure~\\ref{fig_model}) and integrated over the radial extent of each bin.  The model with $\\dot{M}=6\\cdot10^{-8}M_\\odot{\\rm year}^{-1},$ $f_p=0.46,$ $T_{\\rm st}=3.2\\cdot10^7$~K and $550{\\rm counts~pixel}^{-2}$ produced at $r=0$ by a point source gives an excellent fit with the minimum reduced $\\chi^2=1.45$ and weighed $\\chi^2_{\\rm wei}=0.68$ with $1/r$ weights. The stagnation point is at $r_{\\rm st}=1.01''.$ The correspondent unabsorbed point source luminosity $L=4\\cdot10^{32}{\\rm erg~s}^{-1}$ is estimated for monoenergetic photons at $4$~keV and agrees with the estimates of SSC luminosity in \\citet{moscibr_sim}. Energy $4$~keV is chosen as the energy Chandra is most sensitive to for assumed $N_{\\rm H}.$ The minimum reduced $\\chi^2=15$ is achieved for the model without the point source. The models with the outer sonic point instead of finite bounding pressure underpredict the X-ray surface brightness at several arcseconds, assuming fixed $N_{\\rm H}=10^{23}$cm${}^{-2}.$ The reliable fitting for $N_{\\rm H}$ is possible only with the use of spectral data and is left for future research. The assumption $T_p=T_e$ represents the additional point of concern. Temperature equilibrium might not hold at the stagnation point at $1''$ \\citep{quataert_wind}, however the thermalization rate exceeds the outflow rate at $5''$ in our subsonic dense outflow, thus $T_p=T_e$ holds there. The reliable modeling of non-equilibrium flows requires the modeling of the whole spatial structure of the stellar winds and is left for the future research as well.  \\begin{figure}[h] \\includegraphics [scale=0.7, bb= 20 -10 175 495]{arcplot.jpg} \\caption{Radial profiles of electron density $n=n_e$ in cm~$^{-3}$ (upper panel) and electron temperature $T_e$ in keV (lower panel) in the feeding region.} \\label{fig_arcprof} \\end{figure} The profiles of electron density $n_e$ and temperature $T_e$ within several arcseconds from the BH are shown on Figure~\\ref{fig_arcprof} and compare well with the simple earlier estimates \\citep{baganoff,quataert_wind}. The difference is that our best fit is a subsonic flow supported by the outer medium with the density bounce at $5''.$ Though the achieved outflow velocity $v_{\\rm out}=300{\\rm km~s}^{-1}$ is almost independent of radius for $r>2''.$ The line cooling \\citep{sutherland} reduces the heat contents only by several percent for gas reaching $5'',$ bremsstrahlung cooling being less important.  \\begin{figure}[h] \\includegraphics [scale=0.7, bb= 20 -10 175 495]{rgplot.jpg} \\caption{Radial profiles of dimensionless electron temperature normalized to electron mass $k_B T_e/(m_e c^2)$ (upper panel) and ratio of ion to electron temperatures $T_p/T_e$ (lower panel) close to the BH. The inner sonic point is at $3r_{\\rm g}.$} \\label{fig_rgprof} \\end{figure} The profiles of dimensionless electron temperature $k_B T_e/(m_e c^2)$ and ratio $T_p/T_e$ within several Schwarzschild radii from the BH are shown on Figure~\\ref{fig_rgprof}. The electron temperature $T_e=4\\cdot10^{10}$~K and density $n_e=2\\cdot10^6{\\rm cm}^{-3}$ are found close to the BH. This dynamical model gives an excellent fit to the optically thick luminosity $L=1.73$~Jy at $86$~GHz \\citep{krichbaum} for assumed equipartition of thermal energy with the magnetic field. The model overpredicts by a factor of $20$ the observed Faraday rotation measure $RM\\sim 50{\\rm cm}^{-2}$ at $230$~GHz \\citep{marrone}, but this may well be a geometric factor. The accretion rate, temperature and density near the BH are in good agreement with more complicated models specifically focusing on sub-mm emission \\citep{sharma_spherical,moscibr_sim}. We notice that the ratio of ion and electron temperatures $T_p/T_e$ is significantly larger than predicted by \\citet{moscibr_sim}, but probably because of the significantly lower $T_p$ in their numerical simulations of the limited domain."
    },
    "1004/1004.0228_arXiv.txt": {
        "abstract": "We calculate the equation of state (EOS) of dense matter, using a relativistic mean field (RMF) model with a density dependent coupling that is a slightly modified form of the original NL3 interaction.  For nonuniform nuclear matter we approximate the unit lattice as a spherical Wigner-Seitz cell, wherein the meson mean fields and nucleon Dirac wave functions are solved fully self-consistently. We also calculate uniform nuclear matter for a wide range of temperatures, densities, and proton fractions, and match them to non-uniform matter as the density decreases. The calculations took over 6,000 CPU days in Indiana University's supercomputer clusters. We tabulate the resulting EOS at over 107,000 grid points in the proton fraction range $Y_P$ = 0 to 0.56.  For the temperature range $T$ = 0.16 to 15.8 MeV we cover the density range $n_B$ = 10$^{-4}$ to 1.6 fm$^{-3}$; and for the higher temperature range $T$ = 15.8 to 80 MeV we cover the larger density range $n_B$ = 10$^{-8}$ to 1.6 fm$^{-3}$. In the future we plan to study low density, low temperature (T$<$15.8 MeV), nuclear matter using a Virial expansion, and we will match the low density and high density results to generate a complete EOS table for use in astrophysical simulations of supernova and neutron star mergers. ",
        "introduction": "The equation of state (EOS) for hot, dense matter in massive stars relates energy and pressure to temperature, density, and composition.  It has been a long-standing problem to understand the EOS at both subnuclear and supranuclear density, to which great efforts have been devoted, from laboratory heavy ion collision experiments \\cite{heavyions}, computer simulations of supernova \\cite{simulation1,simulation2}, and theoretical many-body calculations \\cite{theory}.  The EOS of hot dense matter in supernovae (SN) and neutron star (NS) mergers encompass multi-scale physics.  Temperature can vary from 0 to as high as 100 MeV, exciting nuclei, nucleon and possibly pion and other degrees of freedom.  The density can vary from $\\approx 10^{4}$ to 10$^{15}$ g$\\cdot$ cm$^{-3}$, where matter can be in gas, liquid or solid phases.  The proton fraction can vary from 0 to 0.6, from extremely neutron rich matter to proton rich matter.  These very large parameter ranges make construction of a full EOS table difficult. It is necessary to employ different approximations for different parameter ranges. As a result, there exist only two realistic EOS tables that are in widespread use for astrophysical simulations, the Lattimer-Swesty (L-S) equation of state \\cite{LS}, that uses a compressible liquid drop model with a Skyrme force, and the H. Shen, Toki, Oyamatsu and Sumiyoshi (S-S) equation of state \\cite{Shen98a,Shen98}, that uses the Thomas Fermi and variational approximations with a relativistic mean field (RMF) model.  We plan to generate a complete equation of state, employing  relativistic mean field calculations for matter at intermediate and high density as described in this paper.  In the future we plan to use the Virial expansion of a nonideal gas to describe matter at low density.   These two parts will be matched together and we will generate a thermodynamically consistent EOS over the full range of parameters.  Finally we will generate additional EOSs from RMF models with different high density symmetry energies.  This will allow one to correlate features of astrophysical simulations with properties of the symmetry energy assumed for the EOS.    There are still large uncertainties in the EOS at supranuclear densities.  The density dependence of the symmetry energy $dS/dn_B$ is poorly known and strongly influences the stiffness of the EOS.  It can be constrained from measurements of NS radii and masses \\cite{Lattimer04}, precision determination of the neutron rms radius in $^{208}$Pb \\cite{prex}, and also heavy ion collision experiments \\cite{heavyions}. A stiff EOS (high pressure) at high density gives larger NS radii, while a stiff EOS at normal and low density favors a larger neutron radius in $^{208}$Pb \\cite{HP01}. The elliptic and transverse flow observables in heavy ion collisions are sensitive to the isospin dependence of mean fields and to the EOS at densities up to a few times nuclear saturation density.  Many nuclear many-body models fall into two categories, the non-relativistic Skyrme models (See for example, Ref.~\\cite{Talmi} for a review) and relativistic mean field models \\cite{RMFreviews}. The parameters in these models are usually fitted to nuclear properties at normal nuclear densities, afterwards they are extrapolated to study supranuclear matter.  The L-S EOS uses a Skyrme model featuring a relatively soft EOS and the S-S EOS uses the RMF interaction TM1 that features a stiffer EOS.  Since the symmetry energy is not well constrained, it is important to explore the effects of different symmetry energies on the EOS and SN simulations.  In this paper, we use a RMF model for non-uniform matter at intermediate density and uniform matter at high density.  Low density pure neutron matter is analogous to a unitary gas \\cite{HS05b}, where the neutron-neutron scattering length is much larger than both the effective range and the average inter-particle spacing. To better describe neutron-rich matter at low density, we use a density dependent scalar meson-nucleon coupling. At high density, the model reduces to the normal RMF parameter set NL3.  The unit lattice of non-uniform nuclear matter is conveniently approximated by a spherical Wigner-Seitz (W-S) cell. The meson mean fields and nucleon Dirac wave functions inside the Wigner-Seitz cell are solved fully self-consistently.  This is unlike the S-S EOS that used Thomas-Fermi and variational approximations and the L-S EOS that used a simple liquid drop model. The size of the W-S cell is found by minimization of the free energy per nucleon.  The W-S approximation provides a framework to incorporate the best known microscopic nuclear physics \\cite{Negele73}. The nuclear shell structure effects are included automatically and it is already possible for some effects of complex nuclear pasta states to be included in spherical calculations in the form of shell states \\cite{HS08}.  Full three-dimensional W-S calculations in principle could incorporate various pasta shapes \\cite{Gogelein07,Newton09}, which would make the transition to uniform matter more smooth. However this will demand much larger computational resources. In this work we use the spherical W-S approximation.  Our relativistic mean field calculations can accurately describe the radial shape of large neutron rich nuclei including the expected neutron rich skin.  In contrast the original L-S EOS is based on a very simple liquid drop model of nuclear structure that may incorrectly describe the neutron skin.  Alternatively the S-S EOS is based on a Thomas Fermi approximation that neglects shell effects.  These are included in our Hartree calculations.  Furthermore, the variational forms for the densities assumed by S-S may be a poor approximation for large proton numbers where the Coulomb repulsion is large.  Instead our exact solutions of the radial mean field equations allow richer density distributions including shell states with central depressions \\cite{HS08}.  These differences in densities may be important for neutrino interactions in Supernovae.  Finally our calculations correctly reproduce the Unitary gas limit for a low density neutron gas, see below, while both the L-S and S-S EOS reduce incorrectly to the energy of free neutrons.  One can demand that any EOS be consistent with, possibly model dependent interpretations of, observations of neutron stars.  For example, Klahn et. al. propose a series of tests that an EOS should satisfy to be consistent with observations \\cite{klahn}.  They demand that any reliable nuclear EOS be able to reproduce the recently reported high pulsar mass of $2.1\\pm$ 0.2 $M_\\odot$ for PSR J0751+1807 \\cite{nice}.  However, this observation may have been retracted \\cite{notnice}.  Furthermore, Klahn et al. require the EOS to reproduce a large binding energy for Pulsar B in J0737-3039.  However, this conclusion could be sensitive to assumptions about the system such as the amount of mass loss.  Klahn et. al. go on to demand that the EOS not allow direct URCA cooling of neutron stars of mass 1 to 1.5 $M_\\odot$.  We consider a more conservative approach.  While many stars cool slowly, observations do suggest that at least some stars have enhanced cooling.  Unfortunately observations do not directly constrain the mass that may separate enhanced from normal cooling.  Indeed, there is little direct observational evidence that more massive stars cool more quickly, although this is a theoretical prejudice.  One can also use laboratory data to constrain the EOS.  The neutron skin thickness of a heavy nucleus constrains the density dependence of the symmetry energy.  Furthermore, there are many measurements of the skin thickness with a variety of strongly interacting probes.  However, there may be important model dependence from strong interaction uncertainties.  For example $(^3He,t)$ measurements of spin dipole strength have been used to extract neutron skin thicknesses in Sn isotopes \\cite{spindipole}.  For these measurements, the spin dipole strength was assumed to be proportional to the measured cross section, and the proportionality constant was arbitrarily fixed in order to reproduce the skin thickness of $^{120}$Sn as predicted by an old Hartree-Fock calculation \\cite{oldHF}.  Presumably, if a different skin thickness in $^{120}$Sn is fit, one would also get different answers for the skin thickness in other isotopes.  This situation may soon change.  The Lead Radius Experiment (PREX) at Jefferson Laboratory is using parity violating electron scattering to measure the neutron skin thickness in $^{208}$Pb \\cite{prex}.  Parity violation is a sensitive probe of neutrons because the weak charge of a neutron is much larger than that of a proton.  Furthermore, this electro-weak reaction may have much smaller strong interaction uncertainties.  Data taking for PREX should be completed by June 2010.  Instead of trying to determine, ahead of time, the best EOS to satisfy existing observational constraints, we adopt what we hope will be a more robust approach.  We are calculating a number of EOSs based on different effective interactions.  In this paper we present first results for the NL3 interaction with a symmetry energy that is large at high densities.  In later work we will present EOSs with softer high density symmetry energies.  These different EOS will allow one to correlate features of astrophysical simulations with properties of the EOS.  Then one can draw conclusions based on combined information from laboratory experiments and astronomical observations.  In this paper we focus on nucleon degrees of freedom.  Hyperons could play a role at high densities, see for example \\cite{shen2}.  However, the contribution of hyperons could depend on uncertain hyperon interactions.  In addition, there could be pion or kaon condensates or a variety of quark matter phases.  See for example the review by Page and Reddy \\cite{theory}.  Chiral symmetry restoration and the softening of pionic or kaonic modes could be important.  Finally, thermal pions and pion interactions should be very important at high temperatures.   All of these effects may increase the uncertainties in the EOS.   We tabulate the equation of state for intermediate and high density nuclear matter over the range of temperatures $T$, densities $n_B$, and proton fractions $Y_P$ given in Table  \\ref{tab:phasespace} and described in Sec. \\ref{methodology}.  We calculate the free energy of nonuniform matter for 17021 points, and the free energy of uniform matter for 90478 points in $T$, $n_B$, and $Y_P$ space.  This took 6000 CPU days on Indiana University's supercomputer clusters.   The paper is organized as follows: in Section \\ref{formalism} the density dependent RMF model is explained in detail.  In Section \\ref{numerics} we describe the RMF parameters that we use including a density dependent coupling.  We describe the computational methodology for our large computer runs in Section \\ref{methodology}.  Section \\ref{result} shows results for RMF calculations, including the free energy and the nucleon density distributions in the non-uniform W-S cells.  Finally, Section \\ref{summary} presents a summary of our results and gives an outlook for future work.  \\begin{table}[h] \\centering \\caption{Range of temperatures $T$, densities $n_B$, and proton fractions $Y_P$ in the EOS table.} \\label{tab:phasespace}\\btab{cccc} \\hline \\hline Parameter & Low T & High T & Total \\# \\\\ \\hline log$_{10}$(T) [MeV] & -0.8 to 1.2 & 1.2 to 1.9 & 32 \\\\  log$_{10}$($n_B$) [fm$^{-3}$] &-4.0 to 0.2 & -8.0 to 0.2 & 43, 83 \\\\  $Y_P$  & 0,0.05 to 0.56 & \\ 0,0.05 to 0.56  & 53 \\\\  \\hline \\etab \\end{table} ",
        "conclusions": ""
    },
    "1004/1004.1638_arXiv.txt": {
        "abstract": "We employ the \\Chandra\\ Multiwavelength Project (ChaMP) and the Sloan Digital Sky Survey (SDSS) to study the fraction of X-ray-active galaxies in the field to $z = 0.7$. We utilize spectroscopic redshifts from SDSS and ChaMP, as well as photometric redshifts from several SDSS catalogs, to compile a parent sample of more than 100,000 SDSS galaxies and nearly 1,600 \\Chandra\\ X-ray detections. Detailed ChaMP volume completeness maps allow us to investigate the local fraction of active galactic nuclei (AGN), defined as those objects having broad-band X-ray luminosities $L_X$(0.5-8~keV)~$ \\ge 10^{42}$~erg\\,s$^{-1}$, as a function of absolute optical magnitude, X-ray luminosity, redshift, mass, and host color/morphological type. In five independent samples complete in redshift and $i$-band absolute magnitude, we determine the field AGN fraction to be between $0.16\\pm{0.06}$\\% (for $z \\le 0.125$ and $-18 > M_i > -20$) and $3.80\\pm{0.92}$\\% (for $z \\le 0.7$ and $M_i < -23$). We find excellent agreement between our ChaMP/SDSS field AGN fraction and the \\Chandra\\ cluster AGN fraction, for samples restricted to similar redshift and absolute magnitude ranges: $1.19\\pm{0.11}$\\% of ChaMP/SDSS field galaxies with $0.05 < z < 0.31$ and absolute $R$-band magnitude more luminous than $M_R < -20$ are AGN. Our results are also broadly consistent with measures of the field AGN fraction in narrow, deep fields, though differences in the optical selection criteria, redshift coverage, and possible cosmic variance between fields introduce larger uncertainties in these comparisons. ",
        "introduction": "\\label{intro}  A fundamental constraint on all theories modeling the interplay of supermassive black hole (SMBH) accretion and galaxy evolution should be the fraction of galaxies in the local universe that host actively accreting nuclei. Accretion onto a SMBH is the predominant source of energy produced by active galactic nuclei (AGN) and QSOs, so a measure of accretion activity (at a variety of redshifts and luminosities) provides a valuable constraint on the black hole mass function at the present day \\citep{Soltan82,Rees84,Marconi04}.  In both active and inactive galaxies, observations show a remarkably tight correlation between the properties of galactic spheroids and the masses of their central black holes, the $M_{BH}-\\sigma$ relation \\citep{Ferrarese00,Gebhardt00,Gultekin09}, suggesting that SMBHs play an important role in the formation of galactic bulges. Simulations indicate that the black hole and the bulge may co-evolve, each growing as a result of repeated merger events \\citep[\\eg,][]{Hernquist89,DiMatteo05,Hopkins05,Hopkins08}. In this scenario, the merger causes cold gas to fall to the center, triggering starbursts, and fueling rapid black hole growth. Even while accretion inflows increase the black hole's mass, outflows transport energy from the active nucleus to the surrounding host galaxy \\citep[\\eg,][]{Granato04}, creating a mechanism by which the galaxy and it's central black hole can ``feedback'' as they evolve. In extreme cases, these outflows can carry enough energy to alter the galaxy's evolution by ejecting gas and truncating star formation.  Several well-studied phenomena support this so-called feedback paradigm, which links star formation and SMBH growth to galaxy mergers. First, the space density of low-luminosity active galactic nuclei (LLAGN) peaks at lower host mass than that of luminous quasars \\citep{Barger05a,Hasinger05,Scannapieco05}, indicating that AGN may undergo ``cosmic downsizing'' as do galaxy spheroids \\citep{Cowie96}. Next, the local mass density of SMBHs corresponds closely to the luminosity density produced by quasars at high redshift \\citep{Yu02,Hopkins06}. Hence, the present-day population of AGN and quiescent black holes may represent the fossil record of their brighter, high-redshift counterparts (the QSO luminosity function peaks at redshifts $z \\sim 2$), formed at a time when galaxy mergers were common. This suggests that all massive galaxies may have undergone an AGN phase \\citep[\\eg, ][]{Marconi04,Shankar04}.  Though they clearly influence the evolution of galaxies and the conditions in the early universe, fundamental questions about the formation and evolution of accreting black holes remain. To be robust, the merger and feedback paradigm must reproduce the fraction of galaxies hosting AGN, the ratio of obscured to unobscured AGN necessary to match the cosmic X-ray background, the space density of SMBHs, and AGN clustering.  X-ray emission is the most reliable primary signature of AGN activity because X-rays offer the most complete and efficient marker of active accretion close to the black hole (10--100 gravitational radii). Combined with other wavelengths, X-ray observations sample a wide variety of AGNs and improve our census of AGN demographics \\citep[][and references therein]{Brandt05}.  Galaxies that do not harbor a powerful accreting black hole may also produce X-rays due to hot interstellar gas, supernova remnants, and low- and high-mass X-ray binary populations \\citep{Fabbiano06,Kim06}; and even in ``normal'' galaxies a LLAGN may persist \\citep[see][for a review]{Ho08}. In this work, we consider the X-ray-active fraction of galaxies out to $z = 0.7$, in a field covering more than $20~deg^2$, at a wide range of broad-band X-ray luminosities ($10^{40}$ \\simlt\\ $L_X$[0.5-8~keV;~erg\\,s$^{-1}$] \\simlt\\ $10^{46}$), including those characteristic of normal and starforming galaxies. Since the many X-ray-producing mechanisms described above can muddy the waters in lower-X-ray-luminosity systems, we focus the bulk of our analysis on sources with X-ray luminosities $L_X$(0.5-8~keV)~$ \\ge 10^{42}$~erg\\,s$^{-1}$, characteristic of accreting SMBHs.  Measurements of the AGN fraction as a function of environment have been made using both optical spectroscopic samples and optical/radio-selected samples. \\cite{Kauffmann04} use emission-line galaxies in the Sloan Digital Sky Survey \\citep[SDSS;][]{York00} to measure the fraction as a function of environment and find that the fraction of very luminous AGN decreases with increasing environmental density \\cite[see also][]{Popesso06}; while less luminous SDSS AGN show no such trend \\citep{Miller03,Kauffmann04}. Studies of the radio-loud SDSS AGN fraction demonstrate that the fraction of galaxies hosting radio-loud AGN increases towards richer environments \\citep[except at the highest emission-line luminosities,][]{Best05}.  As in field studies, X-ray identification of AGN activity in clusters has proved to be even more efficient than optical or radio techniques, particularly for low-luminosity AGN. A significant body of work on the cluster AGN fraction has been enabled by a combination of \\Chandra\\ X-ray observations and spectroscopic follow-up for a series of 32 clusters with $0.05 < z < 1.3$ \\citep{Martini02,Martini06,Martini07,Martini09}. These authors find larger AGN fractions in clusters than those found in optical emission line studies. One of the aims of the present work is to make a detailed comparison between the field and cluster X-ray AGN fractions based on a combination of \\Chandra\\ imaging and ground-based optical photometry.  NASA's {\\it Chandra X-ray Observatory} \\citep{Weisskopf02} has generated a rich archive of X-ray sources, detected with unprecedented resolution and sensitivity. By cross-correlating the \\Chandra\\ archive with the SDSS --- an extensive optical survey in five filters --- it is possible to identify X-ray counterparts to individual optical objects, confirming their classification as quasars or active galactic nuclei. The \\Chandra\\ Multiwavelength Project \\citep[ChaMP;][]{Green04,Green09} has carefully analyzed 323 \\Chandra\\ fields (about 30 square degrees) that overlap the SDSS and characterized all optical/X-ray matches.  We employ these X-ray and optical surveys to determine the fraction of actively accreting SMBHs (or AGN) out to redshift of $\\sim 0.7$. We also explore the AGN fraction as a function of galaxy luminosity, redshift, mass, and host galaxy morphology. Such measures of the AGN fraction will prove useful for constraining the fueling, lifetimes, and growth of central SMBHs.  In \\S \\ref{catalogs} we describe the Extended ChaMP (\\Chandra\\ Cycles 1--6) and SDSS surveys; in \\S \\ref{sample_select} we detail the ChaMP/SDSS sample developed for this study, including specifics on the spectroscopic and photometric redshift catalogs we employ. Section \\ref{properties} outlines our methods for determining k-corrections, absolute magnitudes, masses, X-ray and optical luminosities, and our definition of five independent samples complete in redshift and $i$-band absolute magnitude. In \\S \\ref{the_frac} we describe our determination of the X-ray-active and AGN fractions in these volume-limited optical samples, investigate possible evolution of the fraction with redshift, and trends with absolute magnitude and mass, in addition to providing an assessment of how photometric redshift errors and other selection effects might impact our measure of the AGN fraction. We discuss the AGN fraction as a function of galaxy color and compare our results to others in the field, groups, and clusters in \\S \\ref{discussion}; and wrap up with our conclusions in \\S \\ref{conclusions}.  We adopt an $H_0 = 70$~km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\Lambda} = 0.7$, and $\\Omega_M = 0.3$ cosmology throughout. ",
        "conclusions": ""
    },
    "1004/1004.3121_arXiv.txt": {
        "abstract": "We calculate NLTE models of stellar winds of hot compact stars (central stars of planetary nebulae and subdwarf stars). The studied range of subdwarf parameters is selected to cover a large part of these stars. The models predict the wind hydrodynamical structure and provide mass-loss rates for different abundances. Our models show that CNO elements are important drivers of subdwarf winds, especially for low-luminosity stars. We study the effect of X-rays and instabilities on these winds. Due to the line-driven wind instability, a significant part of the wind could be very hot. ",
        "introduction": "Despite being low-luminous, the horizontal branch stars may still have stellar winds due to their large effective temperatures and low stellar radii. The winds of these stars is thus similar to those of luminous OB stars \\citep{vinca,unglaub}, which are driven mainly by photon absorption in lines of heavy elements, e.g., C, N, O, or Fe (see \\citealt{owopo}; \\citealt{kkpreh}; \\citealt{pulvina} for a review; or \\citealt{wrh}).  The main wind parameters are the mass-loss rate $\\dot M$ and the terminal velocity $v_\\infty$. The mass-loss rate depends on the stellar luminosity $L$ as $\\dot M \\sim L^{1/\\alpha'}$ \\citep{pulvina}, where $\\alpha'\\approx0.6$ for O stars \\citep{nlteiii}. This scaling implies that the mass-loss rate increases with luminosity. Because the wind is driven by metal lines, the mass-loss rate depends also on the mass fraction of heavier elements, $Z$. For O stars the scaling $\\dot M \\sim Z^{0.67}$ is roughly valid \\citep{nlteii}. In contrast, the wind terminal velocity is proportional to the surface escape speed, $v_\\infty\\propto v_\\text{esc}$.  Here we complement the studies of hot subdwarf winds provided by \\citet{vinca} and \\citet{unglaub}. ",
        "conclusions": "We have shown that stellar winds from central stars of planetary nebulae resemble stellar winds of luminous O stars. In contrast, stellar winds of hot subdwarf stars are much weaker and resemble stellar winds of B~stars in several important aspects. They \\begin{itemize} \\item are driven mainly by lighter elements, \\item show multicomponent effects, \\item may be very hot in outer parts due to inefficient shock cooling. \\end{itemize}"
    },
    "1004/1004.3298_arXiv.txt": {
        "abstract": "We searched for evidence of the reddening of background QSO spectra from the Sloan Digital Sky Survey (SDSS) due to dust in intervening damped Lyman $\\alpha$ (DLA) systems. We utilise the two Data Releases DR5 and DR7 to arrive at sample sizes of 475 (DR5) and 676 (DR7) absorbers, based on two different published lists of SDSS DLAs. Both samples span roughly the redshift range of $2.2<z_{abs}<5.2$, with a mean of z$\\sim$3.0, and the majority of the DLAs (75\\%) below z=3.3. We construct geometric mean spectra in the absorber restframes ranging from 1240 to $\\sim$ 2800 \\AA, and composite spectra of samples matching the 'DLA' QSOs in $i${} band magnitude and emission redshift z$_{em}$, but without absorption lines. By comparing the slopes of these composite spectra with their matched counterparts, we find no sign of reddening in the ensemble of the absorbers from these samples. Owing to both the unprecedently large sizes of the DLA samples themselves and the non-DLA SDSS QSO sample, from which we can draw our matching spectra, we can place very tight limits for this non-detection ($<$E(B-V)$>${} =-0.0013$\\pm$ 0.0025 (DR5) and $<$E(B-V)$>${} =-0.0017$\\pm$0.0022 (DR7). Interestingly, when applying our technique to the samples of York et. al. (2006), vandenBerk et al. (2008) (intervening and intrinsic MgII absorbers) and the smaller DLA-subsample and pool of comparison QSOs of Vladilo et al. (2008), we do recover their results, i.e. detect the same amount of reddening as these authors do. Furthermore, we have tested whether subsamples of our large sample in categories involving the absorbers (HI column densities, presence or absence of accompanying metal absorption, absorber redshift) or the background quasars (emission redshift, brightness) do reveal dust extinction, but found no trends. These results are at odds with both detections of dust reddening  from previous studies, and also with expectations from observations of high-redshift galaxies. ",
        "introduction": "Detecting galaxies at high-redshift provides a powerful observational way to probe structure formation in the early Universe. There have now been several hundreds of galaxies found using the Lyman-break technique (e.g. \\citet[]{steidel1999}. Although very successful, this method is limited to luminous objects. A complementary technique to select high-redshift galaxies is the detection of systems in absorption in the spectra of background quasars. Since their selection is independent of luminosity and morphology, the study of galaxy evolution through these systems circumvents the biases inherent to traditional emission studies.\\\\  Damped Lyman-$\\alpha$ systems (hereafter DLAs) have hydrogen column densities, log N(HI) $>$ 20.3 and are thought to be associated with parts of high-redshift galaxies, just as they are at lowest redshift \\citep[]{rao2003, chen2005}. From our knowledge of the local Universe we expect that high-redshift galaxies, especially those forming stars, should contain some amount of dust. Therefore, the possibility that DLAs are dusty objects is currently hotly debated. Indeed, smoothed-particle hydrodynamic simulations \\citep[]{nagamine2004}{} predict that DLAs should be 1/3$^{\\rm rd}$ solar at $z=2.5$ and even more metal-rich toward lower redshifts whilst current observations measure typically a metallicity around 1/10$^{\\rm th}$ solar. Dust is therefore invoked to reconcile results from the simulations with current observations. In addition, one can estimate the dust content of individual DLA from the ratio of depleted onto undepleted elements it contains. It is assumed that a fraction of the elements is not detected in the gas phase because it is locked into dust grains (e.g. \\citet[]{pettini2000}. Such studies have shown that the dust depletion levels observed in DLAs are small, lower than those measured in the Galactic interstellar-medium clouds with similar N(HI) column densities. In fact, the observations resemble the dust depletion of the Galactic halo warm gas \\citep[]{kulkarni1997, savaglio2000}.  Another more direct way to probe the dust content of DLAs is to look for a 2175 \\AA\\ feature (the so-called ``dust bump'') in the quasar spectrum. This feature is known to be more predominant in the Milky Way than in the Small and Large Magellanic Clouds. Many searches lead to no-detections \\citep[]{jura1977, smith1979, boisse1988, lanzetta1989, fall1989, pei1991}. Using Sloan Digital Sky Survey (SDSS) spectra, \\citet[]{wang2004} claimed the detections of bumps in three individual MgII absorbers at z=1.5, while \\citet[]{junkkarinen2004}{} have detected a clear bump in a DLA at z=0.524 toward a BL Lac at z=0.94. Other searches remained unsuccessful, indicating that most DLAs do not show a predominant 2175 \\AA\\ feature. A promising tool for finding the signatures of dust are searches for the strong silicate absorption features in the infrared. \\citet[]{kulkarni2007}, using Spitzer data, report on the first indication of 9.7 $\\mu$m silicate absorption in a DLA at z$_{abs}=0.52$, and find it to be stronger than expected from extrapolation of the relation from Galactic diffuse ISM. % \\\\  A consequence of DLAs being dusty is that it might prevent the inclusion of their background quasar in colour-colour selected samples. In order to overcome such biases, radio-selected quasar samples have been built recently: while \\citet[]{ellison2001}{} did not discover a missing population, a more recent survey finds 4 fields (out of $\\sim$60) where the radio-flux selected quasars have no optical counterparts down to very faint limits \\citep[]{jorgenson2006}. These objects are the potential candidates of dusty DLAs obscuring background quasars.\\\\  %  In this paper, we investigate yet another possible signature of the dust content of DLAs. Indeed, dust in a quasar absorber will redden the colour of its background quasar. \\citet[]{pei1991}{} were the first to look for reddening of the quasar continua in a sample containing DLAs and detect reddening at 99.999\\% level. More recently, \\citet[]{murphy2004}{} have used 79 DLAs from SDSS Data Release 2 to compare the slope of power laws fitting quasar continua with and without DLAs. These authors found a colour excess due to Small Magellanic Cloud type of dust of reddening $<$E(B-V)$>${} $<$0.01 mag. Using colour excess measurements for SDSS sightlines containing DLAs and matching sightlines without them, \\citet[]{vladilo2008} report a detection of reddening in a sample of DR5 SDSS QSOs towards sightlines containing DLAs at the mean level of $<$E(B-V)$>${} $\\sim$6.3x10$^{-3}$. They derive their estimate by comparing colours of QSOs with intervening DLAs to the mean colour of a set of QSO of similar emission redshift and brightness, and bootstrapping the resulting broad distribution of colour excesses to ascertain the slightly redder colours of the DLA QSOs. A byproduct of the \\citet[]{pontzen2009}{} work, who performed a Bayesian analysis on a variety of optical and radio datasets to determine the loss fraction of QSOs due to dust obscuration, is an estimate of the probability density distribution for dust reddening in optical spectra, that also peaks around the same value for $<$E(B-V)$>${} . Here, we build composite spectra of quasar spectra with DLAs from SDSS Data Releases 5 and 7, and compare them with reference composite spectra built to match the original spectra in magnitudes and redshifts.  The organisation of the paper is as follows : after the introduction, we describe the selection of the DLA samples in section 2, followed by the construction of the matching non-absorber samples, and the procedure for obtaining composite spectra in section 3, where we also presents the method for deriving extinction as well as the results for the full samples. After discussing subsample analyses (section 4), we summarise and conclude in section 5.\\\\ ",
        "conclusions": "\\label{summary} While our analysis technique reproduces well past results for other classes of absorbers at different redshifts \\citep[]{vandenberk2008, york2006}, and DLAs using other methods \\citep[]{murphy2004} and \\citep[]{vladilo2008} where some extinction could be detected, we could no detect any significant extinction in samples of DLAs drawn from the Data Releases 5 and 7 of the SDSS. These findings are based on the largest DLA sample ever studied, but also on the largest sample of quasars from which to draw a reference spectrum. The implication is that most of the DLAs observed in SDSS have a very small dust content.\\\\   \\citet[]{york2006}{} have analysed SDSS spectra with intervening strong MgII absorbers, and detect clear dust reddening. More specifically, their subsamples with the highest MgII equivalent width (subsamples 8 and 25 in their notation, Table A4) show the highest degrees of reddening ($<$E(B-V)$>${} =0.032 and 0.085). These strong MgII absorbers are then classified, based on the column densities of the accompanying metal lines, as akin to low redshift DLAs. We note however, that a direct measurement of the HI absorption in these absorbers was not possible due to the low wavelength cutoff of SDSS, and the mean HI column densities for the subsamples are derived by using an SMC relationship between $<$E(B-V)$>${}  and HI. \\citet[]{rao2006}{} also caution that not all high EW MgII absorbers are indeed DLAs, although \\citet[]{york2006}{} point out that the specific requirement of the QSOs probed by \\citet[]{rao2006}{} with HST being UV-bright introduces a bias against red QSOs, that may have lead them to an underestimate of the DLA fraction. Hence, it is a priori not clear whether our results contradict these findings. Also note that there may be a redshift evolution effect. In fact, \\citet[]{menard2008}{} have noted that the rest-frame colour excess E(B-V) for their MgII samples shows a strong decline with increasing redshift (their figure 10, lower right panel). Thus, it is possible that the difference in the mean redshifts of our samples (z$_{abs} \\sim$3.1) compared to the lower z$_{abs} \\sim$1.4 of \\citet[]{york2006}{} may partly explain the seemingly discrepant results. \\\\  Yet another strategy involving  spectral analysis is fitting the QSO continuum spectra by a power-law, and checking for deviations in the slope of these for spectra with and without DLAs. Employing this technique, \\citet[]{murphy2004} can place an upper limit for a sample of 70 DLAs from SDSS DR2 of $<$E(B-V)$>${} =0.01. Note that we can rule out such a high value for our samples with a very high degree of confidence. \\\\  In addition, there have been claims of firm detections of dust reddening towards sightlines afflicted by intervening DLAs. Specifically, \\citet[]{vladilo2008}{} report a mean E(r-z)=27($\\pm$9)x10$^{-3}${} for a subsample of the same DR5 DLA sample that we use. When applying our method to these subsamples of DLAs {\\it and}{} the subsample of DR5 'pool' QSOs for the constructing of the comparison spectrum, we recover their results. Note that \\citet[]{vladilo2008}{} have removed all quasars from their samples that may have additional absorption unrelated to the DLA, reducing the number of objects in each sample by a large fraction. We conclude that the differences do not reside in the methods, but the sample selection. We note that the reddening distribution of undesired interlopers probably has a tail with a small number of objects exhibiting much larger than the typical reddening expected from DLA systems. It may hence be that for such low frequency events, that cases of e.g. BAL QSOs, despite efforts to filter them out, enter the DLA samples more often than the comparison pool, and hence could wash out the signal. Increasing the sample size reduces the possible impact of such 'special', very red QSOs. %  Our result is also at odds with expectations from the evolution of high-redshift galaxies (but see \\citet[]{pontzen2008}. It should be emphasised however that there are two separate issues: i) what is the dust content of the quasar absorbers we know of today and ii) which fraction of the quasar absorbers are missed because of the extinction they induce on their background quasar such that it is not selected in the first place. Indeed, models of galactic chemical evolution (e.g. \\citet[]{prantzos2000} predict the existence of regions of high metallicity and column density which are mostly not detected in current DLA surveys \\citep[]{boisse1998}. One way to reconcile these two findings is to invoke a dust-bias in the observations. These apparent discrepancies could indeed be explained by the dust content of the absorbers themselves: they could absorb enough light from the quasar to make the latter too faint to be included in current quasar samples, which barely go beyond magnitude g$\\sim$20.2. In fact, an extension of the \\citet[]{fall1993} calculations taking into account recent observations addresses both these issues: \\citet[]{vladilo2005} propose that whilst the dust content of the DLAs in current samples is not high, the missing fraction is possibly quite important ranging from 30\\% to 50\\% at $z=1.8-3.0$. One way to to test such a hypothesis is to look at fainter objects than what can be probed with SDSS. Using a Bayesian analysis of these issues in light of different techniques and datasets \\citet[]{pontzen2009} posit that probably less than 10\\%{} of DLAs have been missed in current optical samples, and also derive an $<$E(B-V)$>${}  posterior probability distribution peaking strongly around log $<$E(B-V)$>${} $\\sim -2.0$.\\\\  On the other hand, if the absence of the detection of extinction in DLAs is not due to an observational bias, then selecting in absorption implies that we are picking a population of high-redshift galaxies which have probably undergone little star formation, e.g. \\citet[]{weatherley2005} and \\citet[]{wolfe2006}.  This is difficult to reconcile with the fact that these system are gas-rich by construction."
    },
    "1004/1004.1853_arXiv.txt": {
        "abstract": "% Stars between about 4 and 25 solar masses spend a significant fraction of their post-main sequence lifetime as red supergiants (RSGs) and lose material via stellar winds during this period. For RSGs more massive than 10 solar masses, this mass loss becomes of evolutionary significance, and probably determines the upper mass limit of RSGs in the Hertzsprung-Russell diagram. Despite decades of observations, the driving mechanism responsible for mass loss in RSGs remains unknown. Mainly this is because the optical spectrum accessible from the ground provides almost no useful wind diagnostics, and what information is obtained is spatially averaged over the entire wind volume. However, within the last decade, {\\em Hubble Space Telescope (HST)} observations of many useful ultraviolet wind diagnostics have been obtained at a high signal-to-noise ratio and spectral resolution. In particular, RSGs in eclipsing binaries can provide spatially resolved observations of stellar chromospheres and winds. I review possible RSG wind acceleration mechanisms, discuss some observational constraints, and present some empirical models of RSG chromospheres and winds. ",
        "introduction": "Stellar winds are a ubiquitous phenomenon across much of the Hertzsprung-Russell diagram (HRD). These winds can be broadly grouped into the three main categories of hot star winds, coronal (solar-type) winds, and the cool winds of evolved, late-type stars (Figure~\\ref{fig:hr_diag}).  \\begin{figure}[htb] \\plotfiddle{bennett_fig1.eps}{3.9in}{90}{62}{64}{240}{-17} \\caption{Stellar Winds in the Hertzsprung-Russell Diagram.} \\label{fig:hr_diag} \\end{figure}  \\subsection{Hot Star Winds}  Stars earlier than spectral class B3 typically have high velocity winds, and significant mass-loss rates. Radiation pressure is substantial in these stars since the (large) opacity of iron-group metals peaks in the mid-ultraviolet where the radiation field is strong, and so $P_{\\rm rad}$ is the dominant term in the momentum equation. \\citet*{PDB_cas75} showed that the winds of hot stars are driven by the force of radiation on optically thick atomic lines in regions above the photosphere, where the continuum is optically thin.  \\subsection{Coronal (Solar-type) Winds}  Main sequence stars later than about F0 have substantial convective zones, which interact with stellar rotation to produce magnetically active photospheres. Acoustic waves and turbulence generated by convective ``noise'' propagate outward into the chromosphere, where the acoustic waves steepen into shocks and dissipate, producing localized, transient heating of the chromospheric gas (and the non-magnetic chromospheric line emission). Gas in the corona, lying above the chromosphere, is heated to million degree temperatures by dissipation of magnetic flux and reconnection of magnetic flux tubes tied by their photospheric footprint to convective motions. Coronal winds are tenuous with a mass-loss rate too small to be of evolutionary importance. These are driven mainly by the thermal pressure gradient of the hot corona, the \\citet{PDB_par58} mechanism, although dynamic events in the corona (magnetic reconnection events, flares, and coronal mass ejections) also contribute to the solar wind mass flux (and presumably to that of similar stars also).   \\subsection{Cool Winds of Late-Type Stars}  The red supergiant binary $\\alpha$~Her was first shown, by \\citet{PDB_deu56}, to be losing mass via a stellar wind. The intervening decades of research have failed to produce a viable physical process that explains the mass-loss process in red giants and supergiants \\citep{PDB_jud91}. In these stars, mass loss occurs via a low velocity (typically $\\sim$20--50 km s$^{-1}$), massive ($10^{-9}$ to $10^{-5}$ M$_\\odot$ yr$^{-1}$) wind. This cool wind regime is separated in the HRD from that of coronal winds by the \\citet{PDB_lin79} dividing line (Figure~\\ref{fig:hr_diag}), which lies near spectral class K1 for red giants and around F0 for the most luminous supergiants. Noncoronal stars show generally weaker chromospheric emission line fluxes and a lower degree of ionization than the coronal stars, and have little or no detectable X-ray flux. There is little evidence of emission from gas warmer than about 20,000 K, and most of the gas must be much cooler than this. ",
        "conclusions": ""
    },
    "1004/1004.2061_arXiv.txt": {
        "abstract": "We develop and analyze a simple cellular automaton (CA) model that reproduces the main properties of the evolution of soft X-ray coronal loops. We are motivated by the observation that these loops evolve in three distinguishable phases that suggest the development, maintainance, and decay of a self-organized system. The model is based on the idea that loops are made of elemental strands that are heated by the relaxation of magnetic stress in the form of nanoflares. In this vision, usually called ``the Parker conjecture\" (Parker 1988), the origin of stress is the displacement of the strand footpoints due to photospheric convective motions. Modeling the response and evolution of the plasma we obtain synthetic light curves that have the same characteristic properties (intensity, fluctuations, and timescales) as the observed cases. We study the dependence of these properties on the model parameters and find scaling laws that can be used as observational predictions of the model. We discuss the implications of our results for the interpretation of recent loop observations in different wavelengths. ",
        "introduction": "\\label{intro}   EUV and soft X-ray observations of the solar corona reveal a high degree of structuring and intermittency, especially in and around active regions (ARs). Due to the frozen-in condition of the coronal plasma, the emitting material is ordered by the magnetic field in elongated features commonly known as coronal loops. The study of the structure and evolution of loops is therefore fundamental for understanding coronal dynamics, particularly in relation to the origin of coronal heating (Mandrini \\etal 2000). Historically, the first attempts to physically describe coronal loops began with what is known as the Rosner, Tucker \\& Vaianna (RTV, 1978) approach, which considers loops as monolithic structures containing plasma in static equilibrium. Although soft X-ray loop observations seem to be consistent with the RTV assumptions (Porter \\& Klimchuk 1995), more recent observations from the Transition Region and Coronal Explorer (TRACE) suggest that EUV loops are actually too dense to be in static or quasi-static equilibrium (Aschwanden \\etal 2001, Winebarger \\etal 2003). The properties of these loops are more consistent with the predictions of models based on impulsive heating (Winebarger \\etal 2003, Klimchuk 2006).   TRACE observations also show a high degree of filamentation of the coronal structure, suggesting that what appear to be monolithic loops can actually be collections of unresolved individual strands with more or less independent evolutions. Although the question of whether coronal loops are isothermal or multithermal has been a matter of a heated debate during several years (Aschwanden \\& Nightingale 2005, Schmelz \\& Martens 2006), recent observations seem to indicate that both kinds of loops actually occur (Schmelz \\etal 2009). Combining all the available information, there are strong arguments in favor of many loops being made of unresolved strands that evolve more dynamically than what quasi-static models predict (for a recent review on the subject see Klimchuk 2009).   There are two main families of mechanisms proposed to explain the origin of coronal heating: the so called AC and DC heatings. Generally speaking, AC refers to the dissipation of MHD waves, while DC relates to the dissipation of highly stressed magnetic fields. The latter involve the release of free magnetic energy at locations of strong gradients in the magnetic field, i.e., current sheets. Ultimately, the source of energy is the distortion of the magnetic structures by convective motions. This is precisely the mechanism envisioned by Parker (1988). He proposed that, since the footpoints of magnetic strands are being displaced by photospheric granular motions, there must be a continuous increase of the magnetic stress between neighbor strands. These stresses accumulate until some threshold is reached and reconnection occurs. In the process, stored magnetic energy is released and the plasma is heated. Because of similarities with the usual flaring process and the energy scale involved, these reconnection events are usually called nanoflares (Cargill 1994, Cargill \\& Klimchuk 2004). One key point of the process described above is the existence of a threshold or critical value for the reconnection to occur. A theoretical support for the existence of this threshold is given in Dahlburg \\etal (2005, 2009). They found that when certain field misalignment is reached the secondary instability provides the mechanism for a rapid energy release.   The intermittent and filamentary behavior, and the possible presence of power law distributions (see e.g., Aschwanden \\& Parnell 2002), led some authors to propose that the solar corona might be a good example of a self-organized critical (SOC) system (Bak 1990). Beginning with the seminal work of Lu \\& Hamilton (1991), there have been an important number of studies involving the use of SOC models in relation to the coronal heating problem (for a review on SOC models see Charbonneau \\etal 2001).   In our recent paper, L\\'opez Fuentes \\etal (2007), we studied the evolution of coronal loops observed with the Solar X-ray Imager (SXI) on board the Geosynchronous Operational Environmental Satellite 12 (GOES-12). Since this instrument observes the sun almost continuously, it allowed us to follow the evolution of a set of coronal loops from birth to disappearance. We found that the loop evolution can be separated into three main phases, which we called: \\textit{rise}, \\textit{main} and \\textit{decay}. As we indicated there and began to explore in Klimchuk, L\\'opez Fuentes \\& DeVore (2006), the observed evolution can be related to the development, maintenance and destruction of a critical system. Here, we develop a simple model based on a cellular automaton (CA) that reproduces the observed evolution. We model the plasma response to the heating using the EBTEL hydrodynamics code (Klimchuk \\etal 2008). We then create synthetic light curves taking into account the temperature-dependent response of the SXI instrument.  We compare these with the observed light curves and analyze how the different parameters of the model influence the properties of the obtained light curves, such as the characteristic intensities and timescales, which we studied in detail in L\\'opez Fuentes \\etal (2007) for the observed cases.   In Section~\\ref{obs_loops} we sumarize the results from L\\'opez Fuentes \\etal (2007) and show examples of observed light curves that will be later compared with the model. In Section~\\ref{model} we describe the CA model and in Section~\\ref{results} we analyze our results. We discuss and conclude in Section~\\ref{discussion}. ",
        "conclusions": "\\label{discussion}   We develop a pseudo-1D model based on a cellular automaton (CA) that reproduces the main properties of the evolution of coronal loops observed in soft X-rays. As we describe in Section~\\ref{model}, the model is based on the idea that coronal heating is due to the sudden reconnection of tangled magnetic strands (Parker 1988). Using typical solar ranges for the model parameters we qualitatively reproduce the observed intensity amplitude, fluctuations, and evolution timescales of soft X-ray loops studied in L\\'opez Fuentes \\etal (2007). The synthetic loops obtained with the CA+EBTEL model combination also reproduce the measured physical properties of the observed loop plasma such as density and temperature.   Although our model is related to usual models of self-organized criticality (SOC, see e.g. Charbonneau \\etal 2001), it differs in some important aspects. In most SOC models the rules for the critical condition and the redistribution of the field among neighboring strands are based on mean values of the field strength. The strength of each strand is compared with the mean value of its neighbors, and if the critical condition is fulfilled the field is redistributed equally among those neighbors. One may classify this type of model \\textit{isotropic}. As we describe in Section~\\ref{model}, our CA model is \\textit{anisotropic} in this respect. The other important difference is that SOC models use null boundary conditions. The mesh points (strands) at the boundary are set to have zero magnetic field (what we call $B_{h}$ in Section~\\ref{model}) at all times. Here, we use instead periodic boundary conditions. It may be partly due to these differences that the nanoflares produced with our model do not follow power law distributions, which are characteristic of SOC models (for an anisotropic SOC model associated with power law distributions, see e.g. Morales \\& Charbonneau 2008). We must stress that our main goal here is not to obtain such distributions, but instead to explain the evolution of loops in the frame of the nanoflare model. It is worth pointing out, though, that our CA model reaches a steady critical state in much the same way as SOC models do (see Figure~\\ref{example}).   As we mention in Section~\\ref{intro}, in recent years there has been a debate over the issue of whether coronal loops are isothermal or multithermal. The argument goes that if loops are isothermal they are likely to be monolithic structures, while multithermality implies that loops are composed of different unresolved strands evolving independently. There seems to be evidence in favor of both kinds of observations (see e.g., Schmelz \\etal 2009), so a consensus is forming that both scenarios actually occur. Let us analyze this possibility in terms of the model presented here. For instance, if the reconnection mechanism is not very efficient (e.g., a small $f$ factor, see Section~\\ref{model}) and all nanoflares are approximately the same size, then the heating injection will be difficult to distinguish from a single source of constant heating. In that case, the strand will keep more or less the same temperature along the main phase of the evolution. If all the strands that comprise the loop have similar physical conditions (it is expected that they all have approximately the same field strength, for example), then the loop will be nearly isothermal. These loops can be easily taken as examples of quasi-static evolution. On the other hand, if nanoflares occur less frequently, so the waiting time between consecutive nanoflares is not negligible, the strands have time to cool to lower temperatures between one nanoflare and the next. Since nanoflares occur randomly, at any given time during the evolution of the loop there will be strands simultaneously having different temperatures. That would correspond to a multithermal loop observation.  There is yet a third possibility.  If the nanoflares occur only once in each strand, without repeating, and if they all occur at approximately the same time, then the loop will appear roughly isothermal at any given moment.  The temperature will decrease on a cooling timescale, however, and the loop will be short lived.  We call this scenario a short duration nanoflare ``storm\" (Klimchuk 2009). Many warm EUV loops can be explained in this way.  Hot soft X-ray loops tend to live much longer than a cooling time, on the other hand (L\\'opez Fuentes \\etal 2007). If they are indeed isothermal, then they can only be explained by rapidly repeating nanoflares or by truly steady heating.   The three alternative scenarios described above relate to the question weather the same loop is observable by instruments that are sensitive to different temperatures.  In the case of nanoflares that repeat rapidly in each strand of the loop bundle, the loop will be visible only over a narrow range of temperatures.  It will be seen either in hot soft X-ray emission or in warm EUV emission, but not both. In the case of nanoflares that repeat with a time delay longer than the cooling time, the loop will be simultaneously visible in both emissions.  In the final case of a short duration nanoflare storm, the loop will be visible first in soft X-rays and then in the EUV.  Observational evidence has been presented for all three possibilities.  The question of which is most prevalent is not yet answered.  Loops produced by short duration storms have been clearly identified and can be found at many locations within active regions (e.g., Ugarte-Urra \\etal 2006, 2009).  On the other hand, soft X-ray loops are most pronounced in the central cores of active regions, where EUV loops are less common (Schmieder \\etal 2004).  There are also reports of soft X-ray and EUV loops that are almost but not exactly co-spatial (Nitta 2000, Nagata \\etal 2003). We attempted to study the warm emission associated with the SXI loops studied in L\\'opez Fuentes \\etal (2007).  Unfortunately, TRACE data are not available, so we had to rely on lower cadence observations from SOHO/EIT. Since only four images per day were taken in each channel, we could only compare the hot and warm emission at a few times during the loop evolution. We found that some of the SXI loops had an EIT counterpart but others did not. Clearly, more systematic observations of loops in different wavelengths are needed.   It is possible that there are two well differentiated populations of AR loops. Hotter and longer-lived loops are preferentially located at the core of ARs (Antiochos \\etal 2003, Ugarte-Urra \\etal 2009), while the periphery is dominated by longer, shorter-lived, cooler loops (Del Zanna \\& Mason 2003). If this is indeed the case in general, then the different locations and evolutions of these loop classes may be related to different physical properties at their origin. The model we present in this paper applies specifically to hot loops. Whether it can explain warm loops at the periphery of ARs is unclear. We can nonetheless speculate on the difference in these two regions. The strength of the coronal magnetic field is generally greater in the core of an AR and this will tend to produce stronger heating. Our model also depends fundamentally on the properties of footpoint shuffling. It may be that these properties vary systematically across the active region. This is an important question that we hope will be answered with the latest magnetogram observations from SOT/Hinode. These observations will hopefully also confirm our proposal that the slow decay phase of hot loops is associated with the dispersal of magnetic footpoints to the point where magnetic field tangling is no longer efficient.   In Section~\\ref{results} we found a series of scaling laws relating observed properties of the loop evolution to different physical parameters of the model. Some of these parameters, such as the loop length, can actually be obtained directly from the observations. The strength and tilt of the magnetic field at the coronal base can be inferred by combining photospheric magnetogram observations with models of the coronal structure (see e.g., L\\'opez Fuentes \\etal 2006, Mandrini \\etal 2000). Therefore, these scaling laws are predictions of the model that can be tested observationally.   From the discussion in the previous paragraphs we conclude that among future important investigations is the systematic study of loop observations in different wavelengths and the analysis of possible scalings of the evolution parameters (timescales, mean intensities, fluctuations) with the loop physical properties such as length and magnetic field. From the theoretical side, the model presented here is still very simple, and more sophisticated versions should be developed, such as 2D models that better account for the vector nature of the magnetic field and the fact that each magnetic strand is wrapped around multiple other strands. We have already started work in this direction."
    },
    "1004/1004.2257_arXiv.txt": {
        "abstract": "Physical processes influencing the properties of galaxies can be traced by the dependence and evolution of galaxy properties on their environment. A detailed understanding of this dependence can only be gained through comparison of observations with models, with an appropriate quantification of the rich parameter space describing the environment of the galaxy. We present a new, multiscale parameterization of galaxy environment which retains an observationally motivated simplicity whilst utilizing the information present on different scales. We examine how the distribution of galaxy (u-r) colours in the Sloan Digital Sky Survey (SDSS), parameterized using a double gaussian (red plus blue peak) fit, depends upon multiscale density. This allows us to probe the detailed dependence of galaxy properties on environment in a way which is independent of the halo model. Nonetheless, cross-correlation with the group catalogue constructed by Yang et al, 2007 shows that galaxy properties trace environment on different scales in a way which mimics that expected within the halo model. This provides independent support for the existence of virialized haloes, and important additional clues to the role played by environment in the evolution of the galaxy population. This work is described in full by Wilman~et~al, 2010, MNRAS, in press. ",
        "introduction": "The evolution of galaxies is heavily intertwined with the growth of the dark matter dominated structure in which they live, as baryons react to their local gravitational potential. This leads to a strong and evolving dependence of galaxy properties on their environment \\citep[e.g.][]{Dressler80,Balogh04}. Environment is often defined in a observationally practical fashion, by counting the number of nearby neighbours. A different strategy is to assign galaxies to groups within a theoretical halo framework. Motivated by the need for a fair comparison of data with models, it would be preferable to define a statistical framework which retains the information contained by environment measured on different scales, without resorting to the inevitable assumptions required by a halo (group) finder during membership assignment. Simple two point galaxy statistics (such as correlation functions) trace the mean overdensity on different scales. However these are themselves closely correlated with one another. A better approach would be to examine how galaxy properties correlate with overdensity on one scale for a fixed overdensity on another, independently computed scale. Now that such a large, homogeneous sample of galaxies is available with the SDSS, this is now statistically possible \\citep{Kauffmann04,Balogh04Ha,Blanton06,Blanton07}. In Wilman et al, 2010, we take the logical next step: defining a framework for quantifiable comparisons with models. The basic principles and initial results are described here. ",
        "conclusions": "Our sample is divided into bins, spaced evenly in log-density within a grid of two-scale density. These scales are always chosen so that the smaller scale measures the density of neighbours up to a fixed distance (dividing radius $\\rdiv$), and the larger scale then measures the density in some annulus beyond $\\rdiv$. Figure~\\ref{figure:results} shows the fraction of red galaxies resulting from bimodal fitting within each bin for $\\rdiv=0.5$,$1.0$ and $2.0\\Mpc$, where the outer radius for the outer annulus is $1.0$,$2.0$ and $3.0\\Mpc$ respectively. There is also a (weak) dependence of $\\fr$ on luminosity within this sample, and so we repeat this procedure for each $0.5$mag bin of luminosity in the range $-21.5\\leq\\Mr\\leq-20.0$, with $\\geq25$ galaxies per luminosity and density/density bin required to produce reasonable fits (reduced $\\chi^2\\sim1$). For each bin we compute a residual density dependent $\\dfr=\\fr(\\Mr,\\Sigma_{0,\\rdiv},\\Sigma_{\\rdiv,r_o})-\\fr(\\Mr)$. The three luminosity bins (which have similar appearance) are then coadded ($\\dfr$ is weighted by the number of contributing galaxies) to produce the final maps (figure~\\ref{figure:results}). Colour gradients and contours illustrate the direction of change in $\\dfr$ (vertical contours illustrate no large scale dependence).  \\begin{figure} \\includegraphics[height=.45\\textheight]{dwilman_fig2.ps} \\caption{Dependence of the relative fraction of red galaxies, $\\dfr$, on multiscale density. The bimodal fitting process is repeated independently for each bin in each two dimensional grid of interior and annulus density. This is illustrated for three choices of scales, moving from smaller to larger scales going from the top left to the bottom right. Grid colour changes and overplotted contours illustrate the direction of change within this parameter space. Interestingly our results illustrate a change from positive correlation with density (more red galaxies in denser environments) on small ($<1\\Mpc$) scales, to an anti-correlation with density (less red galaxies in denser environments) on large ($>2\\Mpc$) scales! This mimics the behaviour expected within a simple halo model framework (see text).} \\label{figure:results} \\end{figure}  We find a changing dependence of $\\dfr$ on density moving from small to large scales. On $<1\\Mpc$ scales, there are more red galaxies in denser environments (a strong $\\dsca$ dependence is supplemented with a residual positive correlation with $\\dscc$). On larger scales there is no residual positive correlation, supporting previous results \\citep{Kauffmann04,Blanton06,Blanton07}. However on still larger scales $>2\\Mpc$, there is a highly significant {\\it anti-correlation} with density at fixed $<2\\Mpc$ scale density (i.e. more blue galaxies at higher $\\dsce$)!  The SDSS group (halo) catalogue of \\citet{Yang07} is projected onto our multiscale density space. For the fraction of galaxies in haloes more massive than some fixed limit (e.g. $10^{12.5}\\Msol$), we find a dependence on multiscale density remarkably similar to that seen in figure~\\ref{figure:results}. i.e. positive correlations on small scales (at larger density, more galaxies in massive haloes) are replaced by anti-correlations on large scales (at larger density, fewer galaxies in massive haloes)! More detailed analysis reveals the reason: For galaxies in the centre of haloes, density will always decrease with scale $r$, on average. Galaxies located a distance $\\rh$ from the nearest halo centre will instead show increasing density with $r$ up to $r=\\rh$ (where the annulus encompasses the halo core), and then decrease to larger $r$. i.e. The galaxies at highest density at scale $r$, at fixed smaller scale density, have $r_i< r\\sim\\rh <r_o$. The virial radius of a $10^{12.5}\\Msol$ halo is $\\sim0.2\\Mpc$, so larger scales can trace the cores of haloes to which our galaxy does not belong.  On larger scales this effect only becomes more important, leading to a lower fraction of galaxies in massive haloes at high densities on large scales - i.e. the anti-correlations. Therefore the higher fraction of blue galaxies at higher large scale density can be explained within the context of the halo model: These galaxies are in the accretion regions, and remain blue until after they have been accreted onto more massive haloes.  This work is explained in more detail by Wilman et al., 2010, MNRAS, in press. Our multiscale density parameterization of the environment will provide an observationally motivated parameter space to examine the environment and the associated dependence of galaxy properties, and to make quantitative comparisons with expectations coming from different physical models of galaxy evolution."
    },
    "1004/1004.0227.txt": {
        "abstract": "Concentrations of matter, such as galaxies and galactic clusters, originated as very small density fluctuations in the early universe. The existence of galaxy clusters and super-clusters suggests that a natural scale for the matter distribution may not exist. A point of controversy is whether the distribution is fractal and, if so, over what range of scales. One-dimensional models demonstrate that the important dynamics for cluster formation occur in the position-velocity plane. Here the development of scaling behavior and multifractal geometry is investigated for a family of one-dimensional models for three different, scale-free, initial conditions. The methodology employed includes: 1) The derivation of explicit solutions for the gravitational potential and field for a one-dimensional system with periodic boundary conditions (Ewald sums for one dimension); 2) The development of a procedure for obtaining scale-free initial conditions for the growing mode in phase space for an arbitrary power-law index; 3) The evaluation of power spectra, correlation functions, and generalized fractal dimensions at different stages of the system evolution. It is shown that a simple analytic representation of the power spectra captures the main features of the evolution, including the correct time dependence of the crossover from the linear to nonlinear regime and the transition from regular to fractal geometry. A possible physical mechanism for understanding the self-similar evolution is introduced. It is shown that hierarchical cluster formation depends both on the model and the initial power spectrum. Under special circumstances a simple relation between the power spectrum, correlation function, and correlation dimension in the highly nonlinear regime is confirmed. ",
        "introduction": "The observation that visible matter in the universe shows structure on a huge range of scales, from galaxies, to clusters, to super-clusters, to voids \\cite{peebles2}, has led Mandelbrot, Pietronero and others to conjecture that the structure of the universe may be fractal \\cite{cosfrac}. Support for this controversial conjecture is provided by the fact that the inter-galactic, two-body correlation function decays as a power-law \\cite{Martinez}. If the geometry of the universe was truly fractal on all scales, then the basic principle of cosmology, that the universe is homogeneous and isotropic on large scales, would be violated. While current observations support the existence of an upper bound for the size of the largest structures, the issue is not completely closed \\cite{bigvoid,Labini09}. Typically, in nature, observable fractal behavior is usually restricted to a finite scale range \\cite{Fed}. Regardless of whether such a bound exists, it is still important to understand how the appearance of fractal geometry develops from the underlying dynamics.  The remarkable uniformity of the temperature distribution of the cosmic background radiation CMB is a consequence of the nearly uniform energy distribution at recombination \\cite{WMAP5yr}. Over time, the small (order $10^{-4}$) density fluctuations developed into the structures we see today. Measurements and theory suggest that, following inflation, the density fluctuations were Gaussian random variables with an approximate power-law spectral density. They were subsequently modulated by baryonic acoustic oscillation before recombination \\cite{Martinez}. The commencement of structure formation in the dark matter component actually preceded recombination. These circumstances provide a natural scenario for the initial conditions of any dynamical study. With current technology astronomers can observe bright objects in the distant past. However, they cannot directly observe dynamical processes that take giga-years to unfold. As a consequence, computer simulation plays an especially prominent role in astrophysics, and has been employed to investigate complex behavior ranging from galaxy collisions and mergers to structure formation\\cite{Virgo}. Although progress has been substantial, the ability of three-dimensional simulations with on the order of $10^{9}$ particles to resolve fractal scaling laws is still hampered by limited resolution and approximations in the underlying dynamics. In contrast, simulations with one-dimensional models can incorporate up to $10^{6}$ particles without compromising the two-body gravitational interaction. Thus, although they are less realistic {}``toy'' models, they have the potential to yield insights into clustering with current computers. In addition, since theory suggests that the first collapsed objects were highly flattened {}``pancakes'', they may yet provide contact with the real universe \\cite{peebles2,Zel}.  Rouet and Feix were the first to recognize the potential for using one-dimensional models to investigate clustering in a matter-dominated cosmological setting \\cite{Rouet1,Rouet2}. Since there is no curvature in one dimension, general relativity does not provide a unique path for obtaining the correct laws of motion. Nonetheless, they showed how the transformation to co-moving coordinates could be accomplished in a completely self-consistent formulation which we refer to as the RF (Rouet-Feix) model. Starting with a spatially uniform initial distribution, they demonstrated that, as time evolved, hierarchical clustering occurs. Thus the one-dimensional evolution is similar to what is believed to have occurred in the universe following recombination. Rouet, Feix and Jamin computed the box counting dimension for the distribution of matter in the one-dimensional configuration and $\\mu$ (position, velocity) space. They obtained a result of about 0.6 for the configuration space, strongly suggesting a fractal geometry. Since their seminal work, another similar model was proposed by Fanelli and Aurell, referred to simply as the quintic or Q model \\cite{quintic}. As we will see, the Q model sacrifices internal consistency in order to maintain the correct coefficient of the average cosmological density.  A different approach employing a one-dimensional model to explore matter-dominated cosmology was taken by Gouda and co-workers. Yano and Gouda employed the Zeldovich approximation \\cite{Zel} to investigate the evolution of the one-dimensional system\\cite{Gouda_powsp}. When there is no cutoff in the initial spectra they showed that the evolution of the power spectra is self-similar, and established that three different scaling regimes occur, each with its own characteristic spectral index. Later Yano \\etal \\cite{Gouda_caus} studied the development of a single wave in the phase plane to investigate the effect of caustics on the evolution of the power spectra. In addition, Tatekawa and Meida used the Zeldovich approximation to study the self-similar evolution of a one-dimensional system with initial conditions selected from a Cantor set\\cite{Tat}.  More recently Miller and Rouet extended the original work of Rouet and Feix to include an investigation of multi-fractal properties \\cite{MRexp,MRGexp}. They considered the fact that the clusters are actually forming in $\\mu$ space. An analysis of the generalized fractal dimensions was applied to both configuration and $\\mu$ space and it was shown that multifractal geometry develops in each. Miller and Rouet carried out a multifractal analysis of both the RF and Q models, as well as a Hamiltonian version. Valageas has investigated the Hamiltonian version with identical boundary conditions, and has shown the existence of a sequence of stationary and equilibrium states\\cite{Val1,Val2}. The use of one-dimensional models has recently been extended by Sutter and Ricker to include a dark energy field in addition to dark matter \\cite{Ricker1d}. They investigated how the introduction of dark energy influences the formation of Zeldovich {}``pancakes'', presumably the first large-scale structures to emerge in the cosmos. In a recent series of articles Gabrielli \\etal have studied the one-dimensional system in the infinite particle limit \\cite{gab_pp,Gab,Joyce_1d,Gab_Joy_1DPD,Joyce_range}. A central focus of the series is the statistical properties of the force distribution obtained from sampling static distributions of particle positions from infinite perturbed lattices. In \\cite{Joyce_1d} a number of the issues mentioned here are discussed within this framework.  Theoretical cosmology suggests that during the period that density fluctuations remain small, evolution is linear and is dominated by a {}``growing mode'' and the spectral distribution of the fluctuations remains a power-law \\cite{peebles2}. However, in the simulations of Miller and Rouet described above, initial conditions were created by independently sampling the velocity of each particle from either a uniform or Gaussian distribution. In the work presented here, the initial conditions follow the cosmological picture for cold dark matter. We have performed simulations and a multi-fractal analysis of each model for a variety of initial, power-law, spectral densities. Here we describe the circumstances under which the evolution follows the accepted pattern of self-similar, hierarchical clustering. We show how the evolution of the spectral density and correlation function depend on both the model and the initial power-law index. We find that, for sufficiently large samples, the analytic relation between correlation dimension, correlation function, and spectral density is obeyed.  In the following we first explain the mathematical formulation of the different models in section \\ref{1D}. In particular, we show how to construct the exact gravitational potential and field for a one-dimensional, periodic system, \\textit{i.e.}, Ewald sums for one dimension. In Section \\ref{Simulation-IC} we describe how simulations are performed and show how to create scale-free initial conditions for the growing mode in phase space for an arbitrary initial scaling index. We also present typical results for the Q model in the highly nonlinear region, including the calculation of power spectra and the two-body correlation function. In addition a simple analytic representation of the power spectra is introduced that captures the main features of the evolution, including the correct time dependence of the crossover from the linear to nonlinear regime and the transition from regular to fractal geometry. In Section \\ref{fractal-measure} we introduce techniques for measuring generalized fractal dimensions and apply them to the Q model. In addition we give a comparison of results for related models and different initial power spectra and demonstrate that the power spectra, correlation function, and correlation dimension are interrelated in the nonlinear regime. To provide a physical basis for the observed behavior we introduce the mechanism of hierarchical virialization. Finally, in the last section, we present general conclusions that can be drawn from this investigation and provide a discussion of some open questions that still need to be addressed. ",
        "conclusions": "\\label{summary}  Here and in our earlier work we have seen that hierarchical clustering is an intricate, highly choreographed dance in phase space that requires an initially cold and uniform system. The three models that were studied here differ only in the value of the friction constant, $\\alpha$: $\\alpha=$$\\frac{1}{\\surd2,},$$\\frac{1}{\\surd6,},0$ , respectively, for the RF, Q, and Hamiltonian models. The Q (Zeldovich) model has a hybrid geometry that may actually represent the local behavior of a 3+1 dimensional system for a short time, but cannot reproduce its real clustering properties in the highly nonlinear regime. Alternatively, the RF model is completely self-consistent with the transformation to comoving coordinates in 1+1 dimensions. It is useful for unambiguously studying nonlinear behavior, but has a weaker connection to the 3+1 dimensional universe. In fact, simulations of both the Q and RF models show remarkably similar behavior. The Hamiltonian model provides a contrasting dynamical system with no dissipation in the comoving coordinate system. Like the Q and RF models, it undergoes rapid hierarchical clustering for a short time span. However, unlike them, on long time scales the system becomes space filling and appears to revert to the stable stationary state investigated by Valageas \\cite{Val1,Val2}.  To mimic cosmological investigations, the initial positions were selected by sampling a power-law spectrum with random phases. The velocities were chosen to insure that the system is in the growing mode in phase space. The three models were studied with a range of initial power-law indices $n$. Qualitative behavior was observed for $n=3,4$ while extensive studies were carried out for $n=0,\\:1,\\:2$. Both the RF and Q models exhibited very similar behavior. For the case $n=2$, initially the systems remain fairly uniform. After some time, around $\\tau=8$ for the Q model, the systems break into small clusters. As time progresses further the clusters coalesce into larger ones, supporting the well known bottom-up scenario of hierarchical cluster formation predicted for the real universe. Concurrently the variance of the velocities grows approximately exponentially. The power spectrum was monitored at various times and continued to show scale-free, power-law, behavior. In the strongly clustered epoch it turns around at small scales to a negative value supporting the observed structure of clusters of increasing size. It was shown that a simple analytic representation of the power spectra captured the main features of the evolution, including the correct time dependence of the crossover from the linear to nonlinear regime and the transition from regular to fractal geometry. The underlying mechanism for the observed behavior may be a process of {}``hierarchical virialization'' whereby virialized clusters that have {}``frozen out'' from the cosmic expansion become the new {}``particles'' for the next virialization sequence. The measurements were consistent with the self-similar evolution found in the simulations of Yano and Gouda using the Zeldovich model \\cite{Gouda_powsp}.  In the well clustered regime, a multifractal analysis of the configuration space was performed. For the RF and Q models, in contrast with our earlier work with isothermal and waterbag initial conditions \\cite{MRexp,MRGexp}, here a single dominant scaling range of the partition function $C_{q}$ was found for all values of $q$ considered. Generalized dimensions $D_{q}$ and scaling index $\\tau_{q}$ were evaluated for $-4<q<4$. Robust fractal behavior was inferred for $q>0$ from the values of $D_{q}$. For the Q model the gradual decrease in $D_{q}$ for $q>0$ suggests that the system is also weakly multifractal. The RF model exhibited the greatest inhomogeneity. It appeared nearly monofractal with $D_{2}=0.5$. In addition, the two-body correlation function was constructed and also exhibited a power-law dependence on the displacement. For large samples, on the order of $10^{5}$ or more particles, the mathematical relation between the correlation dimension $D_{2}$, the spectral index $n$ and the correlation function decay exponent $\\gamma$, $D_{2}=-n=1-\\gamma,$ was nicely satisfied. For the Q model at $\\tau=13$, $D_{2}=0.63$ over a finite scale range. For $\\mbox{\\ensuremath{n>2}}$ there was evidence of hierarchical clustering whereas for $n<2$ there was a tendency to form a few large clusters, but in the absence of self-similar, bottom-up, evolution. Like the Q and RF models, the Hamiltonian version undergoes rapid hierarchical clustering for a short time span. However, the scaling behavior was less robust, and the agreement among the exponents was weaker. In conclusion, one-dimensional models can exhibit the type of evolution characteristic of the real universe. Their evolution does reveal fractal behavior, but on a limited length scale that changes with time. It also shows that hierarchical structure formation depends sensitively on the initial conditions. Their three dimensional counterparts are useful for constraining models of the early universe.  A number of interesting questions about these models still remain. In forthcoming work we will explore the connection between symmetry and the various one-dimensional models and show how additional models analogous to the Q type can be constructed. We will also establish a connection between the increase in the variance of the velocity and virialization. While the fractal properties of the over-dense clusters appear to be well represented in the current work, the geometry of the under-dense {}``voids'' has not been revealed. This is under investigation with a greater number of particles per cluster to improve the resolution. A complementary approach that may retain more information in the under-dense regions is provided by the numerical integration of the Vlasov equation. To finally address the conjecture of Balian and Schaeffer concerning the bifractality of the universe for this class of models, this must be completed. The approximate analytic representation of the power spectra introduced in Section (III) suggests that the low density inter-cluster regions have fractal dimension unity, but these results must be considered preliminary. A compelling issue is the connection between cosmologies in different dimensions. In an earlier work we showed that the RF model was able to crudely reproduce the time of galaxy formation in a matter-dominated universe \\cite{MRexp}. It would be intriguing to be able to establish a more precise correspondence.  \\ack  The authors benefitted from interactions with Athanasios Batakis, George Gilbert, Michael Joyce, Michael Kiessling, Igor Prokhorenkov, Paul Ricker, Patrick Valageas, and Martin White. The authors appreciate support provided by the Center for Technologies Resources at Texas Christian University. B. Miller benefitted from the support of the Research Foundation of TCU, and from the hospitality of MAPMO at Universit\u00e9 d'Orl\u00e9ans during the Spring of 2009.\\\\"
    },
    "1004/1004.0252_arXiv.txt": {
        "abstract": "To investigate the black hole mass vs. stellar velocity dispersion (\\msigma) relation of active galaxies, we measured the velocity dispersions of a sample of local Seyfert 1 galaxies, for which we have recently determined black hole masses using reverberation mapping. For most objects, stellar velocity dispersions were measured from high signal-to-noise ratio optical spectra centered on the \\ion{Ca}{2} triplet region ($\\sim 8500$ \\AA), obtained at the Keck, Palomar, and Lick Observatories. For two objects, in which the \\ion{Ca}{2} triplet region was contaminated by nuclear emission, the measurement was based on high-quality $H$-band spectra obtained with the OH-Suppressing Infrared Imaging Spectrograph at the Keck-II Telescope. Combining our new measurements with data from the literature, we assemble a sample of 24 active galaxies with stellar velocity dispersions {\\it and} reverberation-based black hole mass measurements in the range of black hole mass $10^{6}< \\mbh/\\msun < 10^{9}$.  We use this sample to obtain reverberation mapping constraints on the slope and intrinsic scatter of the \\msigma\\ relation of active galaxies. Assuming a constant virial coefficient $f$ for the reverberation mapping black hole masses, we find a slope $\\beta=3.55\\pm0.60$ and the intrinsic scatter $\\sigma_{\\rm int}=0.43\\pm0.08$ dex in the relation $\\log (M_{\\rm BH} / M_{\\odot}) = \\alpha + \\beta \\log (\\sigma_{\\ast} / {200}$ \\kms), which are consistent with those found for quiescent galaxies.  We derive an updated value of the virial coefficient $f$ by finding the value which places the reverberation masses in best agreement with the \\msigma\\ relation of quiescent galaxies; using the quiescent \\msigma\\ relation determined by G\\\"ultekin et al.\\ we find $\\log f=0.72^{+0.09}_{-0.10}$ with an intrinsic scatter of $0.44\\pm0.07$ dex. No strong correlations between $f$ and parameters connected to the physics of accretion (such as the Eddington ratio or line-shape measurements) are found.  The uncertainty of the virial coefficient remains one of the main sources of the uncertainty in black hole mass determinations using reverberation mapping, and therefore also in single-epoch spectroscopic estimates of black hole masses in active galaxies. ",
        "introduction": "In the past decade, supermassive black holes have been found to be virtually ubiquitous at the centers of local quiescent galaxies (for reviews, see Kormendy 2004; Ferrarese \\& Ford 2005). Remarkably, the mass of the central black hole (\\mbh) correlates with the global properties of their host spheroids, primarily the stellar velocity dispersion ($\\sigma_*$; Ferrarese et al. 2000; Gebhardt et al. 2000a). This \\msigma\\, relation is believed to be of fundamental importance, providing one of the strongest empirical links between galaxy formation/evolution and nuclear activity.  Measuring the slope and scatter of the relation, and any possible dependence on additional parameters, is essential to make progress in our understanding of the co-evolution of galaxies and black holes.  One key open question is whether active galactic nuclei (AGNs) share the same \\msigma\\ relation as quiescent galaxies, as would be expected if the relation were universal and nuclear activity were just a random transient phase.  Initial studies using several reverberation-mapped Seyfert 1 galaxies show that the \\msigma\\ relation of active galaxies is consistent with that of quiescent galaxies (Gebhardt et al. 2000b; Ferrarese et al. 2001).  Other studies using larger samples with single-epoch black hole masses suggest that the slope of the relation may be different than that of quiescent galaxies, albeit at low significance (e.g., Greene \\& Ho 2006; Shen et al.\\ 2008). However, these later studies rely on an indirect estimator of black hole mass, based on the kinematics and size of the broad-line region (BLR) as inferred from single-epoch spectra. The method works as follows: the broad-line profile gives a characteristic velocity scale $\\Delta V$ (measured from either line dispersion, $\\sigma_{\\rm line}$, or the full-width at half-maximum intensity, $V_{\\rm FWHM}$); the average size of the BLR ($R_{\\rm BLR}$) is determined from the empirical correlation with optical luminosity (Wandel et al. 1999; Kaspi et al. 2005; Bentz et al.\\ 2009a); and the black hole mass is obtained as \\begin{equation} M_{\\rm BH} = f \\frac{(\\Delta V)^2 R_\\mathrm{BLR}}{G}, \\label{eqn:mbh} \\end{equation} where $G$ is the gravitational constant and $f$ is a virial coefficient that depends on the kinematics and the geometry of the BLR.  Although single-epoch mass estimates are believed to be accurate to $\\sim 0.5$ dex (e.g., Vestergaard \\& Peterson 2006), it is hard to quantify precisely their uncertainty and any possible biases (Marconi et al. 2008; Onken et al. 2009). Thus, the slope and intrinsic scatter of the \\msigma\\, relation of active galaxies remain highly uncertain.  Multi-epoch data provide direct measurements of the BLR size via reverberation mapping time lags ($R_{\\rm BLR}= c\\tau$, where $c$ is the speed of light and $\\tau$ is the measured reverberation time scale), and more secure measurements of the broad-line kinematics as determined from the variable component of the broad lines (Peterson et al.\\ 2004).  Onken et al.\\ (2004) combined reverberation black hole masses with measurements of stellar velocity dispersion for 14 objects to measure the slope of the \\msigma\\ relation, showing that AGNs and quiescent galaxies lie on the same \\msigma\\, relation for an appropriate choice of the virial coefficient ($ f =5.5\\pm1.8$).\\footnote {Note that the virial coefficient, $\\langle f \\rangle$, is $5.5$ when the line dispersion ($\\sigma_{\\rm line}$) is used. If $V_{\\rm FWHM}$ is used, $\\langle f \\rangle$ has to be properly scaled depending on the $V_{\\rm FWHM}/\\sigma_{\\rm line}$ ratio.  See Onken et al. (2004) and Collin et al. (2006) for details.} Remarkably, the scatter in \\mbh\\ relative to the \\msigma\\ relation is found to be less than a factor of 3. Unfortunately, their sample is very small (14 objects) and covers a limited dynamic range ($6.2 < \\rm log \\mbh\\ < 8.4$), making it difficult to simultaneously measure the intrinsic scatter and the slope and investigate any trends with black hole mass, or with other properties of the nucleus or the host galaxy (cf. Collin et al.\\ 2006).  Under the assumption that the relations should be the same for active and quiescent galaxies, the Onken et al.\\ (2004) study provides some information on the geometry of the BLR, an absolute normalization of reverberation black hole mass, and an upper limit to the intrinsic scatter in the virial coefficient. Since the reverberation-mapped AGN sample is the ``gold standard'' used to calibrate all single-epoch mass estimates (e.g., Woo \\& Urry 2002; Vestergaard et al. 2002; McLure \\& Jarvis 2002; Vestergaard \\& Peterson 2006; McGill et al.\\ 2008; Shen et al. 2008), this comparison is effectively a crucial link in establishing black hole masses for all AGNs across the universe and for all evolutionary studies (e.g., McLure \\& Dunlop 2004; Woo et al.\\ 2006, 2008; Peng \\etal\\ 2006a,b; Netzer \\etal\\ 2007; Bennert et al.\\ 2010; Jahnke et al. 2009; Merloni et al. 2010).  In this paper, we present new measurements of host-galaxy velocity dispersion obtained from deep, high-resolution spectroscopy at the Keck, Palomar, and Lick Observatories. The new measurements are combined with existing data (e.g., Onken et al. 2004; Nelson et al. 2004; Watson et al. 2008) and with recently determined black hole masses from reverberation mapping (Bentz et al.\\ 2009b) to construct the \\msigma\\ relation of broad-lined AGNs for an enlarged sample of 24 objects. For the first time, we are able to determine {\\it simultaneously} the intrinsic scatter and the slope of the \\msigma\\ relation of active galaxies. Also, by forcing the slope of the \\msigma\\ relation to match that of quiescent galaxies, we determine the average virial coefficient and its nonzero intrinsic scatter. Finally, we study the residuals from the determined \\msigma\\ relation to investigate possible trends in virial coefficient with properties of the active nucleus.  The paper is organized as follows. In \\S 2, we describe the observations and data reduction of the optical and the near-infrared (IR) data. New velocity dispersion measurements are presented in \\S 3, along with a summary of previous measurements in the literature.  In \\S 4, we investigate the \\msigma\\ relation of the present-day AGNs, and determine the virial coefficient in measuring \\mbh.  Discussion and summary follow in \\S 5 and \\S 6, respectively. ",
        "conclusions": "\\begin{figure*} \\epsscale{1.0} \\plotone{f7.eps} \\caption{The \\msigma\\ relation of active galaxies with reverberation black hole masses (blue), compared with non-active galaxies with dynamical black hole masses. The reverberation masses were determined assuming the virial coefficient, $\\log f=0.72\\pm0.10$. The solid line is the best-fit slope of the active galaxies while the dashed line is the best fit to the inactive galaxy samples from G\\\"ultekin et al. (2009).} \\label{fig_msigma} \\end{figure*}  We present the \\msigma\\ relation of the reverberation sample in Figure 7, using the slope, $\\beta=3.55\\pm0.60$, determined in \\S 4.1 and the average virial coefficient, $\\langle\\log f\\rangle=0.72\\pm0.10$, determined in \\S 4.2. Compared to the local quiescent galaxies, active galaxies follow a consistent \\msigma\\ relation with a similar slope and scatter.  Note that the mean black hole mass of the reverberation sample ($\\langle\\log \\mbh/\\msun \\rangle=7.3\\pm0.74$) is an order of magnitude smaller than that of quiescent galaxies ($\\langle\\log \\mbh/\\msun \\rangle=8.2\\pm0.79$).  The slightly shallower slope of the reverberation sample is consistent with the trend in the quiescent galaxies that the slope is shallower for galaxies with lower velocity dispersion ($\\sigma_* < 200$ \\kms); see G\\\"ultekin et al. (2009).  In contrast, the slope of late-type quiescent galaxies ($4.58\\pm1.58$) seems higher than the slope of our active galaxies, which are mainly late-type galaxies.  However, given the uncertainty in the slope ($\\beta=3.55\\pm0.60$) of the active galaxy \\msigma\\ relation, the difference between quiescent and active galaxies is only marginal. We did not attempt to divide our sample into various morphology groups or a few mass bins to test the dependence of the slope, since the sample size is still small and biased toward lower mass objects.  The smaller mean \\mbh\\ of the reverberation sample is due to the relative difficulty of measuring the stellar velocity dispersions for the more luminous and higher redshift quasars that have the largest black hole masses. Although there are 17 quasars with measured reverberation black hole masses, the stellar velocity dispersions have not been measured for most of those high-luminosity objects (see Watson et al. 2008; Dasyra et al. 2007).  In our current sample, there are only three objects above $3\\times 10^{8}$ \\msun.  In order to better constrain the \\msigma\\ relation of active galaxies, it will be necessary to pursue further measurements of velocity dispersions for reverberation-mapped AGNs at high masses.  Additional stellar velocity dispersion measurements will be available based on Keck OSIRIS spectra in the near future (Woo \\& Malkan, in preparation).  The key assumption in determining the slope of the \\msigma\\ relation is that the virial coefficient does not systematically vary as a function of black hole mass. Currently, we do not have sufficient information to test whether this assumption is valid. In an empirical study of the systematic effects of emission-line profiles, Collin et al. (2006) showed that it was necessary to use a smaller value of the virial coefficient for AGNs with larger $V_{\\rm FWHM}$.  In contrast, the change of the virial coefficient was not required if the line dispersion ($\\sigma_{\\rm line}$) was used instead.  These results reflect the fact that there is a large range of $V_{\\rm FWHM}/\\sigma_{\\rm line}$ ratio, and that systematic uncertainties are higher if $V_{\\rm FWHM}$ is used.  With the enlarged sample of 24 AGNs, we investigated the dependence of the virial coefficient on the line properties and the Eddington ratio. However, we did not clearly detect any increasing trend of the intrinsic scatter with the line profile, line width, or Eddington ratio, implying that the \\mbh\\ based on the mean virial coefficient is not systematically over- or under-estimated as a function of line width or the Eddington ratio.  We note that since our sample is still small and biased to narrower-line and lower Eddington ratio objects, further investigation with a larger sample, particularly at higher mass scales, is required to better constrain the systematic errors.  For single-epoch black hole mass determination, we suggest using $\\log f= 0.72\\pm0.10$ when the gas velocity is determined from the line dispersion ($\\sigma_{\\rm line}$). If the FWHM of an emission line ($V_{\\rm FWHM}$) is measured instead of the line dispersion, then additional information on the $V_{\\rm FWHM}/\\sigma_{\\rm line}$ ratio is needed in the black hole mass calculation. The value of $V_{\\rm FWHM}/\\sigma_{\\rm line}$ is typically assumed to be 2 (Netzer 1990), although there is a range of values. For example, the mean of $V_{\\rm FWHM}/\\sigma_{\\rm line}$ of our reverberation sample is $2.09$ with a standard deviation of $0.45$ (see also Peterson et al. 2004 and McGill et al.\\ 2008). We did not attempt to determine an average virial coefficient using the virial product based on $V_{\\rm FWHM}$. A correction factor depending on $V_{\\rm FWHM}/\\sigma_{\\rm line}$ should be applied in addition to the virial coefficient, $\\log f= 0.72$, when gas velocity is determined from $V_{\\rm FWHM}$. Our empirical measurement of the mean virial coefficient is not very different from that of Onken et al. (2004) although the sample size and the dynamical range of the reverberation sample is significantly improved.  We note that our results on the mean virial coefficient are based on broad-line dispersions measured from the rms spectra, which are not generally available for single-epoch black hole mass determinations. Detailed comparison of broad-line profiles and widths between rms and single-epoch spectra is necessary to  understand the additional uncertainty of single-epoch black hole mass estimates (see Denney et al. 2009), and we will pursue further investigations of this issue using LAMP data in the future."
    },
    "1004/1004.5525_arXiv.txt": {
        "abstract": "We present the first Bayesian constraints on the single field inflationary reheating era obtained from Cosmic Microwave Background (CMB) data. After demonstrating that this epoch can be fully characterized by the so-called reheating parameter, we show that it is constrained by the seven years Wilkinson Microwave Anisotropies Probe (WMAP7) data for all large and small field models. An interesting feature of our approach is that it yields lower bounds on the reheating temperature which can be combined with the upper bounds associated with gravitinos production. For large field models, we find the energy scale of reheating to be higher than those probed at the Large Hadron Collider, $\\rhoreh^{1/4}>17.3$ TeV at $95\\%$ of confidence. For small field models, we obtain the two-sigma lower limits $\\rho_\\ureh^{1/4}>890\\,\\TeV$ for a mean equation of state during reheating $\\bar{w}_\\ureh=-0.3$ and $\\rho_\\ureh^{1/4}>390\\,\\GeV$ for $\\bar{w}_\\ureh=-0.2$.  The physical origin of these constraints is pedagogically explained by means of the slow-roll approximation. Finally, when marginalizing over all possible reheating history, the WMAP7 data push massive inflation under pressure ($p<2.2$ at $95\\%$ of confidence where $p$ is the power index of the large field potentials) while they slightly favor super-Planckian field expectation values in the small field models. ",
        "introduction": "\\label{sec:intro}  The current ongoing flow of high accuracy astrophysical observations has important consequences for our understanding of the very early Universe. In particular, the widely accepted inflationary paradigm~\\cite{Guth:1980zm, Linde:1981mu, Albrecht:1982wi,Linde:1983gd} (for a review, see \\eg Refs.~\\cite{Linde:2007fr,Martin:2003bt,Martin:2004um,Martin:2007bw}) is now under close scrutiny. According to this scenario, the Cosmic Microwave Background (CMB) anisotropies and the large scale structures originate from the unavoidable quantum fluctuations of the inflaton and gravitational fields in the very early Universe subsequently amplified during inflation~\\cite{Mukhanov:1981xt,Hawking:1982cz,Starobinsky:1982ee, Guth:1982ec,Bardeen:1983qw}. One can show that the corresponding power spectrum of the cosmological fluctuations naturally acquires an almost scale invariant form which is fully consistent with all observations. Another crucial property of the inflationary power spectrum is that the slight deviations from scale invariance are linked to the microphysics of inflation~\\cite{Stewart:1993bc,Mukhanov:1990me,Liddle:1994dx}. Therefore, by measuring these deviations, one can probe the shape of the inflaton potential and, therefore, learn about the physical origin of the inflaton field.  \\par  It is often claimed from the above properties that observations give access to a limited part of the potential only, namely the one which is slow-rolled over by the inflaton when scales of astrophysical interest today left the Hubble radius. This observational window represents a range of approximately $7$ e-folds or three decades in wavenumbers. However, inflation does not consist of the slow-roll phase only and the pre and/or reheating period is also of fundamental importance since it allows us to understand how inflation is connected to the hot big-bang phase~\\cite{Turner:1983he,Kofman:1997yn,Bassett:2005xm, Mazumdar:2010sa}. This physical phenomenon is related to a different part of the inflationary potential, usually the one located close to its true minimum, \\ie a few decades in e-folds away from the observable window.  \\par  Observation of pre-/reheating effects can be achieved in two ways. First, the power spectrum can evolve on large scales when the inflaton field oscillates around the minimum of its potential. However, this happens only in quite complicated models, typically those containing more than one field~\\cite{Finelli:1998bu, Bassett:1998wg, Finelli:2000ya}. In fact, it was recently shown that this type of effect can also happen in single field inflation but on much smaller scales~\\cite{Jedamzik:2010dq,Jedamzik:2010hq,Easther:2010mr}. Second, the duration of the pre-/reheating phase can significantly modify the position of the observational window mentioned above. Put differently, at fixed astrophysical scales today, changing the pre-/reheating duration is equivalent to moving the window along the potential, hence probing different values of the power spectrum spectral index, amplitude of the anisotropies and tensor-to-scalar ratio. Obviously, this cannot be done arbitrarily because CMB data impose accurate bounds on their value. Conversely, this opens up the possibility to constrain the pre-/reheating duration and/or its equation of state from CMB data~\\cite{Martin:2006rs}. Notice that a direct detection of primordial gravitational waves would also allow us to probe the reheating temperature, as shown in Refs.~\\cite{Nakayama:2008wy, Kuroyanagi:2009br}.  \\par  The goal of this article is to address this question for the standard scenarios of inflation. It is traditional to study three categories of models usually considered as representative of the full inflationary space. These models are large field~\\cite{Linde:1984st}, small field~\\cite{Linde:1981mu,Albrecht:1982wi}, and hybrid inflation~\\cite{Linde:1993cn}. Hybrid scenarios involve multiple fields and, therefore, the power spectrum can change during the preheating phase. This makes this class of scenarios deserving of a separate investigation. For this reason, in this article, we limit ourselves to the class of large and small field models.  \\par  In the following, we will use the term ``reheating'' to refer to the pre-/reheating phases of the Universe defined to have occurred just after the end of inflation and just before the radiation dominated era. So far, the constraints on the reheating energy scale are not so numerous. Obviously, it should be less than the energy scale of inflation which implies that $T_\\ureh\\lesssim 10^{16}$ GeV. In addition, if one assumes that supersymmetry is the correct extension of the standard model of particle physics, then constraints from Big-Bang Nucleosynthesis (BBN) on unstable gravitinos lead to a reheating temperature $T_\\ureh\\lesssim 10^7$ GeV~\\cite{Khlopov:1984pf, Kallosh:1999jj, Giudice:1999yt, Lemoine:1999sc, Maroto:1999ch, Giudice:1999am, Buonanno:2000cp, Copeland:2005qe,Jedamzik:2006xz, Kawasaki:2008qe, Bailly:2009pe}. Notice that this constraint can nevertheless be avoided if one considers the scenario of Ref.~\\cite{Mardon:2009gw}. Reheating itself should also proceed before BBN and this implies that $T_\\ureh \\gtrsim 10$ MeV. We see that the reheating temperature is poorly constrained, in particular, its lower limit. As a matter of fact, the work presented here precisely yields a lower limit on the reheating energy scale from the current seven years Wilkinson Microwave Anisotropies Probe (WMAP7) data~\\cite{Jarosik:2010iu, Komatsu:2010fb, Larson:2010gs}.  \\par  In order to derive constraints on the reheating phase, we make use of Bayesian techniques and utilize a full numerical approach~\\cite{Ringeval:2007am}. This has several advantages. First, it is exact and rests only on the linear theory of cosmological perturbations: the method remains accurate when the slow-roll approximation breaks down, as one expects near the end of inflation. Second, and of particular importance for the present work, it permits a new treatment of reheating. Indeed, instead of viewing the reheating parameters as nuisance parameters, they can easily be included in the Bayesian data analysis process. Third, the evolution of cosmological perturbations in the hot big-bang eras already relies on numerical codes. Treating perturbations during inflation in the same way allows the whole procedure to be automatized and to be easily extended to other scenarios. Fourth, the numerical approach allows us to address the question of the priors choice in a particularly well-defined way. Indeed, from a physical point of view, our prior knowledge is on the inflationary theory and not on the shape of the primordial power spectra which is actually a model prediction. Therefore, it is better, and easier, to choose prior probability distributions directly on the model parameters, such as the power index of the large field potentials. This reflects the fact that a model of inflation is not a disembodied mathematical structure that one only needs to ``fit'' but a physical scenario rooted in high energy physics that one needs to understand.  \\par  This paper is organized as follows. In Sec.~\\ref{sec:physicalorigin}, we extend the above discussion and explain in detail why the reheating epoch can be constrained with CMB data. In particular, we introduce the so-called reheating parameter which depends on the reheating duration and on the mean equation of state of the fluid dominating the Universe during this epoch. Then, using the slow-roll approximation, we analytically demonstrate that the accuracy of the WMAP7 data is now sufficient to obtain some constraints on the reheating era. In Sec.~\\ref{sec:cmbtoreh}, using a full numerical integration of the tensor and scalar power spectra coupled to Bayesian methods, we derive the constraints that any reheating model has to satisfy. Then, assuming specific values for the mean equation of state, we translate these constraints into new lower limits for the reheating energy density and/or reheating temperature. These results significantly improve the bounds coming from the Big-Bang Nucleosynthesis. In Sec.~\\ref{sec:conclusion}, we recap our main findings and discuss how our results are modified by the inclusion of others CMB data sets. In Appendix~\\ref{appendix:lfreh}, we work out a typical example which illustrates the robustness of our assumptions: a noninstantaneous transition between reheating and the radiation dominated era when one considers the finite decay width of the inflaton field. Finally, as a by-product of our data analysis, Appendix~\\ref{appendix:srpost} presents the updated WMAP7 constraints on the spectral index, tensor-to-scalar ratio and first order slow-roll parameters marginalized over second order effects. ",
        "conclusions": "\\label{sec:conclusion}  We now conclude our investigation by revisiting our main results. The most important conclusion is that, both for large and small field scenarios, the reheating parameter $\\ln R$ is now constrained by CMB data. The physical origin of this result is clear. For fixed physical length scales today, a change in $\\ln R$ modifies the location of the CMB observable window along the inflationary potential, which is possible only for a limited range of $\\ln R$ given the data accuracy. This conclusion is general and does not depend on the details of the reheating epoch. However, if one assumes a model for reheating, typically if one chooses a specific value of the mean equation of state, then it becomes possible to express the constraints mentioned above as limits on the energy density at the end of reheating or, equivalently, as constraints on the reheating temperatures. This leads to Eqs.~(\\ref{eq:lfrhorehmin}). and~(\\ref{eq:conssf}). These results are of particular interest for the supersymmetric extension of either large or small field inflationary models. Indeed, our result limits the reheating temperature from below whereas gravitinos production gives an upper bound~\\cite{Khlopov:1984pf, Kallosh:1999jj, Giudice:1999yt, Lemoine:1999sc, Maroto:1999ch, Giudice:1999am, Buonanno:2000cp, Copeland:2005qe,Jedamzik:2006xz, Kawasaki:2008qe, Bailly:2009pe}: $T_\\ureh < 10^{4}\\,\\TeV$, where $T_\\ureh$ is given in Eq.~(\\ref{eq:Trehdef}). Therefore, assuming $g_*\\simeq 200$ in the large field case, we now have an allowed range of variation for the reheating temperature given by \\begin{equation} 6 \\, \\mbox{TeV}\\lesssim T_\\ureh\\lesssim 10^4 \\, \\mbox{TeV}. \\end{equation}  \\par  By including the reheating parameter in our analysis, we can marginalize over all reheating history to infer the inflationary parameter values in a robust way. For large field scenarios, we find that the power index is upper limited by $p<2.2$, at $95\\%$ confidence limit. This means that the prototypical model of inflation, namely, massive chaotic inflation, is now under pressure. Similarly, the small field models with sub-Planckian vacuum expectation values $\\mu$ are slightly disfavored. In fact, $\\mu$ is correlated with the potential power $p$, as represented in Fig.~\\ref{fig:sf2D}. Without marginalization over $p$, large values of $p \\gtrsim 6$ are actually necessary to allow the sub-Planckian values of $\\mu$ to be inside the $95\\%$ contour. If one has a theoretical prejudice for $\\mu< \\Mp$, and for reasonable values of $p < 6$, then small field models can also be considered under pressure.  \\par  Since the constraints on the reheating parameters are directly related to the ability of the data to determine the observable parameters, one could consider more data sets than the WMAP7 data. In fact, since solely the accuracy on the primordial parameters matters, data sets improving the constraints on the standard cosmological parameters do not change the reheating bounds. On the other hand, small scale CMB experiments may be decisive but, as mentioned in Ref.~\\cite{Komatsu:2010fb}, they do not give a significant improvement on the determination of $\\nS$ and $r$ due to their low accuracy at large multipoles. We have indeed tested that our limits do not change by including the Baryonic Acoustic Oscillation~\\cite{Percival:2009xn} or the Arcminute Cosmology Bolometer Array Receiver data~\\cite{Reichardt:2008ay} in our analysis. On the other hand, since the future Planck data are expected to improve the bounds on $\\nS$, $r$ and even on new primordial observables, we should get unprecedented information on the inflationary reheating era."
    },
    "1004/1004.5180_arXiv.txt": {
        "abstract": "Several independent cosmological tests have shown evidences that the energy density of the Universe is dominated by a dark energy component, which cause the present accelerated expansion. The large scale structure formation can be used to probe dark energy models, and the mass function of dark matter haloes is one of the best statistical tools to perform this study. We present here a statistical analysis of mass functions of galaxies under a homogeneous dark energy model, proposed in the work of Percival (2005), using an observational flux-limited X-ray cluster survey, and CMB data from WMAP. We compare, in our analysis, the standard Press-Schechter (PS) approach (where a Gaussian distribution is used to describe the primordial density fluctuation field of the mass function), and the PL (Power Law) mass function (where we apply a nonextensive q-statistical distribution to the primordial density field). We conclude that the PS mass function cannot explain at the same time the X-ray and the CMB data (even at 99\\% confidence level), and the PS best fit dark energy equation of state parameter is $\\omega=-0.58$, which is distant from the cosmological constant case. The PL mass function provides better fits to the HIFLUGCS X-ray galaxy data and the CMB data; we also note that the $\\omega$ parameter is very sensible to modifications in the PL free parameter, $q$, suggesting that the PL mass function could be a powerful tool to constrain dark energy models. ",
        "introduction": "In the last few years, several experiments and observations have shown strong evidences that the expansion of our Universe is accelerating, under the influence of a mysterious dark energy component\\cite{Observ,Sper07}. The dark energy presents an equation of state parameter $\\omega=p/\\rho$ and represents about 74\\% of the Universe. When $\\omega=-1$ we have the special case of the cosmological constant - or vacuum energy - as responsible for the Universe acceleration (it is the standard $LCDM$ cosmological model).  We also know that around $\\sim$ 26\\% of the Universe are composed by matter (baryonic + dark matter); analyzing the growth of these matter density fluctuation, we can describe the formation of the large structures, such as clusters and super-clusters of galaxies.  The mass function of galaxies is a powerful tool to study the large scale structure formation in the Universe. The standard analytical approach was developed in 1974 by Press and Schechter\\cite{PS74} (hereafter PS), and it is still used today, for its success and simplicity. The PS approach uses a Gaussian distribution to describe the primordial matter density fluctuation field.  Inspired by Tsallis $q$-nonextensive statistics\\cite{Tsal88} and kinetic theory\\cite{RS00,ZER}, we have proposed a new mass function\\cite{L04,L07,L08}, replacing the Gaussian by a non-Gaussian distribution - a Power Law (PL) distribution. The PL mass function has a physically motivated free $q$ parameter (related with the long range gravitational correlations between particles), which provides malleability to fit observational data. Also, if $q=1$ we recover the original Gaussian distribution. The PL mass function is an extension of the PS one, and presents the same analytical simplicity.  We have compared the PS mass function with the PL one, in the special case of the $LCDM$ cosmological model\\cite{L08}. Using the X-ray flux-limited sample of galaxy clusters from Reiprich and Boehringer\\cite{reiprich02} (HIFLUGCS) in a $\\chi^{2}$ statistical analysis, and applying independent cosmological tests (BAO and Shift Parameter) to better constrain the results, we concluded that the PS approach presents incompatibilities with the independent CMB data. On the other hand, using the PL mass function, we have better cosmological parameters, and we note an overlap with the CMB data for a large range of values of its $q$ free parameter. These results, although encouraging, are limited to the $LCDM$ model ($\\omega=-1$).  The mass function could be used to probe the dark energy behavior in the Universe, and works using homogeneous and non-homogeneous dark energy in the large scale structure formation are beginning to flourish in the literature. In this work we use a model of homogeneous dark energy from Percival\\cite{Perci1}. By allowing other values for the $\\omega$ dark energy parameter, we can compare the applicability of the PS and the PL mass functions to probe dark energy models. ",
        "conclusions": "\\label{Sec3}  In this work we perform a statistical analysis, using the X-ray galaxy cluster data named HIFLUGCS\\cite{reiprich02}, and the independent Cosmic Microwave Background radiation (CMB) data from WMAP\\cite{Sper07}, to study the influence of galaxy mass functions under a homogeneous dark energy background, from the work of Percival\\cite{Perci1}. Performing a $\\chi^{2}$ analysis and marginalizing over the $\\omega=p/\\rho$ dark energy equation of state parameter, we compare the standard Gaussian Press-Schechter mass function\\cite{PS74} to the PL mass function (based on the non-extensive thermodynamics and kinetic theory), which presents a physically motivated free parameter, $q$\\cite{L04,L07,L08}.  We conclude that the PS mass function cannot explain at the same time the X-ray and the CMB data, even at $3\\sigma$ confidence level, which imposes a problem to the standard Gaussian scenario. Also, the PS approach presents a best fit dark energy equation of state parameter of $\\omega=-0.58$, which is clearly distant from independent estimates of $\\omega\\sim-1$ (a cosmological constant).  Observing the contours of Figure \\ref{Fig2}, in the range $1.10<q<1.20$, the PL mass function provides better fits to the HIFLUGCS and the CMB independent data, with almost the same best-fit for $\\Omega_{m}$ and $\\sigma_{8}$, but with a huge modification in the best-fit dark energy $\\omega$ parameter, as we change $q$. This strongly suggests that the PL mass function could be a powerful tool to constrain dark energy models. Future analysis, based on the next generation of cluster surveys in medium and high redshifts, will certainly provide better constraints to the dark energy model that accelerates de Universe."
    },
    "1004/1004.5463_arXiv.txt": {
        "abstract": "Recently, Buta et al. (2009, AJ, 137, 4487) examined the question \"Do Bars Drive Spiral Density Waves?\", an idea supported by theoretical studies and also from a preliminary observational analysis \\citep{block2004}. They estimated maximum bar strengths $Q_b$, maximum spiral strengths $Q_s$, and maximum $m$=2 arm contrasts $A_{2s}$ for 23 galaxies with deep AAT $K_s$-band images. These were combined with previously published $Q_b$ and $Q_s$ values for 147 galaxies from the OSUBSGS sample and with the 12 galaxies from \\cite{block2004}.  Weak correlation between $Q_b$ and $Q_s$ was confirmed for the combined sample, whereas the AAT subset alone showed no significant correlations between $Q_b$ and $Q_s$, nor between $Q_b$ and $A_{2s}$. A similar negative result was obtained in \\cite{durbala2009} for 46 galaxies.  Based on these studies, the answer to the above question remains uncertain.  Here we use a novel approach, and show that although the correlation between the {\\it maximum} bar and spiral parameters is weak, these parameters do correlate when compared {\\it locally}.  For the OSUBSGS sample a statistically significant correlation is found between the local spiral amplitude, and the forcing due to the bar's potential at the same distance, out to $\\approx$1.6 bar radii (the typical bar perturbation is then of the order of a few percent).  Also for the sample of 23 AAT galaxies of \\cite{buta2009} we find a significant correlation between local parameters out to $\\approx$1.4 bar radii. Our new results confirm that, at least in a statistical sense, bars do indeed drive spiral density waves. ",
        "introduction": "The response of an outer disk to a rotating bar is intimately related to the maintenance of long lasting spiral arms in galaxies.  Optical photometry \\citep{schweizer1976} established that besides the gas and young stars, the spirals are present also in the old population. This was confirmed by near-infrared surveys \\citep{esk2002}, and is particularly true for grand-design spirals \\citep{knapen1996}.  Short-lived stellar density waves can be induced via disk instabilities \\citep{bertin1977,tremaine1978}, or by galaxy interactions \\citep{toomre1972}, but such transient patterns fade after $\\sim$10 galaxy rotations \\citep{sellwood1984}, unless they are maintained by some feedback cycle, e.g. due to the swing amplification \\citep{toomre1981}. On the other hand, spiral arms are excited by a growing bar, as demonstrated by the very first N-body simulations \\citep{hohl1971} and by analytical calculations \\citep{atha1980}.   Support for the bar/spiral connection is provided by the examples where prominent spirals extend from the ends of the bar (see e.g. NGC 1300, p. 525 in Binney \\& Tremaine 2008, or NGC 986 in Buta et al. 2010).  Also, grand-design spirals are more frequent in barred than in non-barred galaxies \\citep{elm1982}. Nevertheless, a direct connection between bars and spirals has been difficult to prove observationally.    The possible driving of spirals by bars was addressed by \\citet{block2004}, who compared the maximum of bar-related torque strength $Q_b$ (the maximum of tangential force amplitude normalized by mean radial force) with the maximum associated with the spirals ,$Q_s$, after separating the bar and spiral components based on their Fourier density amplitude profiles. Near-IR observations were used, to show the effect of the bar on the surrounding mass, rather than on the gas or the formation of young stars.  Based on 12 galaxies observed in the $K_s$-band, \\citet{block2004} found a strikingly clear correlation between $Q_b$ and $Q_s$. However, the correlation is less clear in later studies using larger samples and deeper images. In \\cite{buta2005} we analyzed the $H$-band images from the Ohio State University Bright Spiral Galaxy Survey (OSUBSGS, Eskridge et al. 2002) for 147 galaxies, and were able to ``weakly verify a possible correlation between $Q_s$ and $Q_b$''. More recently, in \\citet{buta2009} we analyzed deep $K_s$ Anglo-Australian Telescope (AAT) observations for 23 galaxies, with no statistically significant correlation between $Q_b$ and $Q_s$, nor between $Q_b$ and $A_{2s}$ (the maximum of m=2 Fourier density amplitude of spirals).  Similarly, \\cite{durbala2009} found no correlations when analyzing Sloan $i$-band data for 46 isolated barred galaxies; a lack of correlation between bar and spiral arm strengths was seen also by \\cite{seiger2003} who analysed 41 galaxies.  Nevertheless, \\citet{buta2009} showed that the correlation is present when combining the AAT data with the previous data sets of \\citet{block2004} and \\citet{buta2005}.  In this paper the bar/spiral connection is re-investigated using the same samples which were used by \\cite{buta2005,buta2009}: the OSUBSGS sample containing $\\sim$100 barred galaxies, and our AAT sample of 23 barred galaxies. A novel approach is used: instead of comparing the {\\em maximum} bar strength with the {\\em maximum} spiral density amplitude, we compare the locally measured bar forcing and spiral amplitude as a function of distance. The locally felt forcing due to bar is a more important parameter than the maximum forcing, since the $Q_b$ is typically attained well inside the spiral structure. Also, we examine the spiral density rather than force amplitude, since the former measures more directly the possible {\\em response} to bar forcing.  Using our approach, a statistically significant correlation is demonstrated to exist between the bar forcing and the spiral amplitude, up to a considerable distance beyond the end of the bar. ",
        "conclusions": "A connection between bar forcing and spiral density amplitudes was investigated for two near-infrared galaxy samples: 103 barred galaxies from the magnitude-limited OSUBSGS survey, and 23 barred galaxies in our AAT Survey. The main results are:  (1) {\\em For both samples a statistically significant correlation is found between the local tangential bar forcing, $Q_{bar}(r)$, and the local spiral amplitude, $A_2(r)$, up to a radial distance of $r\\sim$1.5$R_{bar}$. }  \\noindent The correlation suggests that, at least in a statistical sense, the stellar spirals of the disk are not transient features but rather represent a continuation of the bar mode itself, or are driven by the bar through some mechanism.  Further out, the spirals may be either independent modes or transient wave packets.   (2) {\\em The obtained range of the correlation is similar for early and late-type spirals ($0\\le T\\le3$ and $4\\le T\\le9$) , and also for small and large bars ($R_{bar}/h_r$ smaller/larger than unity).}  \\noindent This does not favor the idea that only certain types of bars could drive spiral structure, or that the forcing on the stellar component requires the presence of significant gas component. Nevertheless, the current samples are small, and a larger number of galaxies is needed to draw reliable conclusions about morphological type dependencies, or to probe whether the statistical correlation extends to even larger distances beyond the bar.  In this respect the forthcoming S$^4$G survey (The Spitzer Survey of Stellar Structure in Galaxies; Sheth 2009, Sheth et al. 2010) will be extremely useful, providing unprecedentedly deep 3.6 and 4.5 micron observations for nearly 2300 nearby galaxies."
    },
    "1004/1004.3559_arXiv.txt": {
        "abstract": "In this paper I review the theory and numerical simulations of non-linear dynamics of preheating, a stage of dynamical instability at the end of inflation during which the homogeneous inflaton explosively decays and deposits its energy into excitation of other matter fields. I focus on preheating in chaotic inflation models, which proceeds via broad parametric resonance. I describe a simple method to evaluate Floquet exponents, calculating stability diagrams of Mathieu and Lame equations describing development of instability in $m^2\\phi^2$ and $\\lambda\\phi^4$ preheating models. I discuss basic numerical methods and issues, and present simulation results highlighting non-equilibrium transitions, topological defect formation, late-time universality, turbulent scaling and approach to thermalization. I explain how preheating can generate large-scale primordial (non-Gaussian) curvature fluctuations manifest in cosmic microwave background anisotropy and large scale structure, and discuss potentially observable signatures of preheating.  \\vspace{2em} \\noindent\\sl Dedicated to the memory of Lev Kofman, a friend and a mentor, without whom this article would never have been written. ",
        "introduction": "The idea of inflation (a period of rapid quasi-exponential expansion of the Universe) neatly solves several long-standing issues in cosmology \\cite{Linde:2005ht, Linde:2005dd}, and has been spectacularly confirmed by observations of the Cosmic Microwave Background (CMB) anisotropies \\cite{Komatsu:2010fb, Bennett:2010jb}. While the Universe is inflating, its contents is cold. But eventually, inflation has to end and the field driving the inflation must decay, depositing energy into high-energy particles. This process, known as reheating, ``boils'' the vacuum and starts the thermal history of the universe with the hot big bang. As universe continues to cool down, it could undergo more phase transitions, which would happen at symmetry breaking points of the theory. Very little is known about the fundamental physics at these energy scales, and cosmological observations could be our only source of information for the foreseeable future. No photons reach us directly from this epoch as the universe is filled with hot plasma and is opaque until recombination. Nevertheless, the expansion history of the early universe is imprinted on the sky in the form of primordial curvature fluctuations. With success of WMAP and the data from Planck soon to come, CMB observations are reaching precision required to disentangle other subdominant effects from Gaussian fluctuations due to simple inflation \\cite{Komatsu:2009kd}.  The most basic models of reheating involve inflaton decaying into one or more other scalar fields. Among the most interesting are the ones where decay is non-perturbative, for example proceeding through parametric resonance naturally happening in chaotic inflation models \\cite{Dolgov:1989us, Traschen:1990sw, Kofman:1994rk, Shtanov:1994ce, Kofman:1995fi, Kofman:1997yn, Greene:1997fu}, or tachyonic instability in hybrid inflation models \\cite{Linde:1993cn, GarciaBellido:1997wm, Felder:1998vq}. For all their simplicity, these models have surprisingly rich physics involving non-equilibrium phase transitions. While initial stages of preheating are linear and instability development can be understood analytically \\cite{Kofman:1997yn, Greene:1997fu}, dynamics could be chaotic \\cite{Podolsky:2002qv}, and the field evolution quickly becomes inhomogeneous and non-linear, so non-perturbative decay of the inflaton has to be studied numerically \\cite{Khlebnikov:1996mc, Prokopec:1996rr, Kasuya:1998td, Tkachev:1998dc, Felder:2000hj, Felder:2001kt, Copeland:2002ku, GarciaBellido:2002aj, Podolsky:2005bw, Dufaux:2006ee, Felder:2006cc}. In this paper I briefly review the theory of parametric resonance, go over general results of numerical simulations of non-linear field evolution during preheating, and discuss signatures of preheating that could potentially be observed. ",
        "conclusions": "Very little is known about how inflation actually ended, and what was the high energy physics like at those energy scales (or even what is the inflaton itself). Traces of reheating are hidden from us by opaque plasma in nearly thermal state, and are unobservable directly. One must seek signatures of preheating that survive thermalization and could be detected. These could include stable relics like topological defects \\cite{Kasuya:1998td, Tkachev:1998dc, Battye:1998xe} or primordial black holes \\cite{Rubin:2000dq, Suyama:2004mz, Suyama:2006sr}, stochastic gravitational wave background produced by inhomogeneities during reheating \\cite{Easther:2006gt, Easther:2006vd, GarciaBellido:2007dg, GarciaBellido:2007af, Dufaux:2007pt, Caprini:2007xq}, or anomalies in the expansion history of the universe imprinted in the primordial curvature fluctuations \\cite{Dvali:2003em, Kofman:2003nx, Chambers:2007se, Chambers:2008gu, Bond:2009xx, Bond:preview, Kohri:2009ac, Chambers:2009ki}. Of these, the last effect appears to be the most promising observationally, as stable relic formation is difficult without spoiling cosmology, and stochastic gravitational waves are very hard to detect.  Curiously enough, the simple model of preheating (\\ref{eq:V:L4G22}) could generate non-Gaussian curvature fluctuations of observable amplitude which are intermittent, producing primordial ``cold spots'' \\cite{Bond:2009xx}. Although arguable \\cite{Bennett:2010jb}, the CMB temperature map appears to have slight statistical anomalies in the form of a cold spot in the southern hemisphere \\cite{Cruz:2009nd}, and slight discrepancy in north-south temperature anisotropy spectra \\cite{Eriksen:2007pc}. Exciting possibility is that primordial ``cold spot'', whether from preheating or some other early universe source, could potentially explain both. A tell-tale signature of \\textit{primordial} non-Gaussian cold spot is the associated $E$-mode polarization pattern around it, which might be possible to test with Planck data \\cite{Vielva:2010vn}. Primordial potential ``dips'' would also manifest in formation of large scale structure."
    },
    "1004/1004.3838_arXiv.txt": {
        "abstract": "Electron acceleration in collisionless shocks with arbitrary magnetic field orientations is discussed. It is shown that the injection of thermal electrons into diffusive shock acceleration process is achieved by an electron beam with a loss-cone in velocity space that is reflected back upstream from the shock through shock drift acceleration mechanism. The electron beam is able to excite whistler waves which can scatter the energetic electrons themselves when the \\Alfven Mach number of the shock is sufficiently high. A critical Mach number for the electron injection is obtained as a function of upstream parameters. The application to supernova remnant shocks is discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.4904_arXiv.txt": {
        "abstract": "V1408 Aql (= 4U 1957+115) is a low-mass X-ray binary with an orbital period near 9.3 hr, whose compact star is a black hole candidate. The system shows a large-amplitude orbital photometric modulation at optical wavelengths. We have obtained new optical photometry of V1408 Aql from which we derive the orbital light curve and an improved orbital ephemeris. We show that the orbital light curve can be reproduced by a model in which the accretion disk around the compact star is thin, axisymmetric, and uneclipsed. The secondary star is heated by X-rays from the compact star and the accretion disk. The orbital modulation is produced entirely by the changing aspect of the irradiated secondary star with orbital phase. Because the system does not eclipse, the fits of the model light curves are insensitive to the orbital parameters, allowing a wide range of orbital inclinations and mass ratios. ",
        "introduction": "Low-mass X-ray binaries (LMXBs) are close binary stars containing a black hole or neutron star surrounded by an accretion disk fueled by mass transferred from a low-mass companion. V1408~Aql is an LMXB with an orbital period of $P =9.329 \\pm 0.011$~h, derived by \\citet{tho87} from a roughly-sinusoidal modulation he found in its optical light curve. \\citet{tho87} attributed the modulation to the varying aspect of the secondary star, which is heated by flux from the primary star and its accretion disk. \\citet{hak99} obtained two nights of $UBVRI$ photometry, covering one orbital period in total, and also saw an orbital variation, but the light curve was more asymmetric than sinusoidal. Noting that V1408~Aql became bluer when fainter, \\citet{hak99} interpreted the minimum in the light curve as a grazing eclipse of a cool, red rim on the accretion disk by the secondary star. In their model the ratio of the rim height to disk radius is large, $H/R\\sim0.2-0.5$, and the orbital modulation outside eclipse is caused by variations in the height of the rim, which obscures more or less of the brighter inner disk. Because it invokes an eclipse, the model implies a high orbital inclination, the best fitting model having an inclination near $75^{\\circ}$.  4U~1957+115, the X-ray counterpart of V1408~Aql, is a persistent X-ray source, showing neither large transient increases in brightness nor faint quiescent states. Its X-ray flux does, however, vary on all time scales from 10~Hz (1-2 percent rms amplitude) to years (by a factor of four) \\citep{wij02, rus10}. The orbital period has not been detected at X-ray wavelengths but \\citet{lev06} reported a ``marginal detection'' of a periodicity at $9.3175\\pm 0.0005$ h, perhaps related to the orbital period; and \\citet{now99} reported a 117-day period in the X-ray light curve, although this period appears to be evanescent \\citep{wij02}. The distance to 4U~1957+115 is unknown and its mean X-ray luminosity consequently uncertain; but for a distance of 7~kpc the mean unabsorbed luminosity is $\\log(\\textrm{L}_X) \\approx 36.5$ (erg~sec$^{-1}$, 1.5-12 keV), while for a distance of 40 kpc it is $\\log(\\textrm{L}_X) \\approx 38.1$ \\citep{rus10}.  The X-ray spectrum of 4U~1957+115 has been measured numerous times. It is unusually soft, placing 4U~1957+115 in the region of the X-ray color-color diagram occupied by black hole binaries, including black hole soft X-ray transients (SXTs) in their outburst states, leading to the suspicion that 4U~1957+115 harbors a black hole \\citep{whi84, sch89, mcc03}. If it does contain a black hole, it would be the only known black hole binary in the Galaxy that has not been seen to cycle between low and outburst states. The nearly featureless spectrum can be successfully fitted with a variety of models that include some combination of three components: a soft component from a single- or multi-temperature black body, presumably coming from the accretion disk;  a hard component that is either a power law, with or without a cutoff, or a thermal Comptonization spectrum; and a broad emission feature near 6.5~keV from Fe K-line emission \\citep{yaq93,singh94,ric95,now99,wij02}. The spectrum is softer when 4U~1957+115 is fainter, harder when it is brighter.  Most recently \\citet{now08} fit high-quality Chandra, XMM-Newton, and RXTE observations of 4U~1957+115 with {\\tt diskbb} \\citep{mit84} and {\\tt kerrbb} models \\citep[see][]{li05,dav06} in the XSPEC package, fixing the inclination at $75^{\\circ}$, and found that these black hole models fit the observed X-ray spectrum well. The fits required high inner-disk temperatures, $\\textrm{kT}_{in} = 1.2 - 1.8\\ \\textrm{keV}$, and a low normalization factor, $(R_{in}/D_{10})^2 \\cos i = 5 - 25$, where $R_{in}$ is the inner radius of the disk in km, $D_{10}$ is the distance in units of 10 kpc, and $i$ is the orbital inclination. There are degeneracies among the derived values for the mass of the compact object, the mass accretion rate, and the distance, but the best fitting models ranged from a $3\\ \\textrm{M}_\\odot$ black hole with a spin of $a^{*}>0.82$ at a distance of 10~kpc to a $16\\ \\textrm{M}_\\odot$ black hole with $a^{*}\\approx 1$ at a distance of 22~kpc. Nevertheless, the question remains: Is the compact object in 4U~1957+115 really a black hole? The evidence for a black hole is hardly conclusive. A soft spectrum does not necessarily require a black hole -- many SXTs contain neutron stars \\citep{tan96}, and \\citet{now08} did not explore models for disks around neutron stars. Furthermore, \\citet{yaq93} argued that the high values of $\\textrm{kT}_{in}$ and low values of $R^2_{in}\\cos i$ deduced from fits to the X-ray spectral energy distribution are more typical of neutron star systems than black hole systems.  In this paper we present new optical photometry of V1408~Aql. We refine the orbital period and derive the mean orbital light curve. There are no eclipses in the light curve. We achieve good fits to the orbital light curve with a model similar to that proposed by \\citet{tho87} in which the orbital modulation is due to the varying aspect of the irradiated face of the secondary star. We believe this model is more physically realistic than models that invoke a tall disk rim to produce the orbital modulation, but it increases the range of possible orbital inclinations and mass ratios. ",
        "conclusions": "We obtained high-speed photometry of V1408~Aql on nine nights in June and August of 2008. The optical light curve flickered by at least 5\\% on time scales of minutes to hours and the orbit-averaged flux varied by 30\\% on time scales of a few days. The mean orbital light curve was roughly sinusoidal with a peak-to-peak amplitude near 23\\% and was not measurably different from the light curve observed by \\citet{tho87} in 1985. We attribute the lack of an orbital modulation reported by \\citet{rus10} and the different shape of the orbital light curve found by \\citet{hak99} to the sparseness of their data sets and to the distortions of the light curve induced by flickering and night-to-night variations. The refined orbital period is $0.38877\\pm0.00013$~d (equation~\\ref{ephem}). There are no significant periodicities in the RXTE ASM X-ray light curve at or near the orbital period.  The orbital light curve is consistent with a model in which the accretion disk around the compact star is thin, axisymmetric, and uneclipsed. The secondary star is irradiated and heated by X-rays from the compact star and the accretion disk. The orbital modulation is caused entirely by the varying aspect of the heated face of the secondary star with orbital phase. This simple model yields consistent fits for all mean light levels and it avoids an unphysically tall disk rim.  At inclinations greater than $70^{\\circ}$ the synthetic orbital light curves flatten at minimum light. The lack of a flat minimum in the observed light curve places an upper limit of $\\sim 70^{\\circ}$ on the orbital inclination. We cannot place an unimpeachable lower limit on the orbital inclination but argue that both the amplitude of the flickering and the lack of color variations over the orbit imply that the orbital inclination is greater than $20^{\\circ}$. Thus, the orbital inclination lies between $20^{\\circ}$ and $70^{\\circ}$. We achieved good fits to the light curves with mass ratios from 0.025 to 0.3. If the radius of the accretion disk is close to the tidal truncation radius, the smaller mass ratios, which correspond to systems in which the primary is a black hole, require X-ray luminosities nearer $10^{38}$~erg~s$^{-1}$ and orbital inclinations in the upper half of the permitted range.  Figure 8 in \\citet{rus10} plots optical luminosity against X-ray luminosity for V1408~Aql and other LMXBs. The distance to V1408~Aql is unknown, but \\citet{rus10} adopt three distances for comparison. If the distance is 7~kpc, the X-ray luminosity is $\\sim 10^{36.5}$~erg~s$^{-1}$, and at at 20~kpc, it is $\\sim 10^{37.5}$~erg~s$^{-1}$, placing V1408~Aql in a regions of the figure occupied by neutron star binaries but not black hole binaries. At 40~kpc, the X-ray luminosity is $\\sim 10^{38.5}$~erg~s$^{-1}$ and V1408~Aql falls in a region occupied by some soft-state black hole binaries and the neutron star Z-sources. All these possibilities remain viable with our new model.  If, however, V1408~Aql contains a black hole it would be unusual in two ways. First, as mentioned earlier it would be the only known black hole binary in the Galaxy that has not been seen to cycle between low and outburst states. Second, irradiative heating is more typical of neutron star X-ray binaries than black hole binaries. This is shown most clearly by SXTs during their outbursts. Large amplitude modulations from irradiative heating have been observed during outbursts of Aql~X-1 and XTE~2123-058 and possibly X1608-52, all of which are classical SXTs containing neutron stars; and have also been observed in the related neutron star X-ray transient XTE~J2129+47 \\citep{tho79,wel00,zur00,wac02}. In contrast, orbital modulation from irradiative heating of the secondary star is rarely observed in black hole SXTs and, if it is observed, its amplitude is low \\citep{nei07}. This is due at least partly to the lower mass ratios of black hole binaries. At lower mass ratios the secondary star is comparatively smaller. As shown in section 5.3, the smaller secondary can produce a large-amplitude orbital modulation only if its irradiated face is heated to high temperatures.  Finally, the light curves of the black hole SXTs often show superhumps \\citep{has01,zur06,kat09}. Such superhumps are ubiquitous in mass transfer binaries with disks whose radii cross the 3:1 orbital resonance, which requires a mass ratio less than $\\sim 0.25$ \\citep{lub91,kat09}. The failure of V1408~Aql to show superhumps in its light curve argues that its mass ratio is larger than $\\sim 0.25$ and the mass of its primary star is less than $\\sim 1.6\\ M_\\odot$. All the known black holes in binary stars have masses greater than about $4\\ M_\\odot$ \\citep{cas06}. While none of these arguments are conclusive, together they weigh heavily in favor of a neutron star and against a black hole in V1408~Aql.  We thank R. I. Hynes for helpful discussions.  \\newpage"
    },
    "1004/1004.3119_arXiv.txt": {
        "abstract": "{} {We assess the ability of a solid ring to model a global perturbation induced by several thousands of main-belt asteroids.} {The ring is first studied in an analytical framework that provides an estimate of all the ring's parameters excepting mass. In the second part, numerically estimated perturbations on the Earth-Mars, Earth-Venus, and Earth-Mercury distances induced by various subsets of the main-belt population are compared with perturbations induced by a ring. To account for large uncertainties in the asteroid masses, we obtain results from Monte Carlo experiments based on asteroid masses randomly generated according to available data and the statistical asteroid model.} {The radius of the ring is analytically estimated at 2.8 AU. A systematic comparison of the ring with subsets of the main belt shows that, after removing the 300 most perturbing asteroids, the total main-belt perturbation of the Earth-Mars distance reaches on average 246 m on the 1969 - 2010 time interval. A ring with appropriate mass is able to reduce this effect to 38 m. We show that, by removing from the main belt $\\sim$240 asteroids that are not necessarily the most perturbing ones, the corresponding total perturbation reaches on average 472 m, but the ring is able to reduce it down to a few meters, thus accounting for more than 99$\\%$ of the total effect.} {} ",
        "introduction": "Asteroid perturbations are considered the major obstacle for achieving a satisfactory prediction accuracy in planetary ephemerides. The asteroid problem persists despite the ephemerides being fitted to a variety of highly accurate observations available over almost 40 years. The most constraining observations are tracking data from spacecraft orbiting Mars, which over the past 8 years are available with metric precision. Despite this accuracy, today's ephemerides are unable to extrapolate the Mars ranging data one year into the future with a precision better than 15 m \\citep[see][]{folkner2008,fienga2009}. The obstacle to improving the situation is to correctly account for a large number of similar asteroid effects without any accurate knowledge of the asteroid masses. In modern ephemerides, usually about 300 asteroids are modeled individually and the rest of the main-belt is represented by a circular ring. Our objective in this study is to estimate the ability of a ring to model large numbers of asteroids.  In section \\ref{sec_ring} we set a simple analytical framework and show that a ring is in fact a first-order representation of the main-belt perturbation. Section \\ref{sec_num} describes a new implementation of the ring in the INPOP ephemeris \\citep[see][]{fienga2009}. To evaluate the ring model, we chose to compare the perturbation induced by the ring on the Earth-Mars distance with the perturbations induced by a test model containing thousands of asteroids. Section \\ref{sec_mb_model} describes the chosen test model that consists of 24635 asteroids assigned each according to the available data with reasonable random mass values. In section \\ref{sec_comp} the effect of the ring is compared with the perturbation induced by the test model after removing various asteroid subsets. The comparison is repeated for different asteroid masses in Monte-Carlo-like experiments. We are thus able to estimate the number of asteroids that need to be individualized in the asteroid model so the ring is able to satisfactorily represent the remaining global perturbation. We also derive estimates for the amplitude of the global perturbation, together with the corresponding mass of the ring and the amplitude of the residual perturbation unmodeled by the ring. The perturbations on the Earth-Mercury and Earth-Venus distances are considered at the end of the section. Some applications of the obtained results are discussed in section \\ref{sec_appli} . ",
        "conclusions": "A ring is an implementation of an averaged orbit, which is a very good model of the perturbation induced by the main-belt asteroids. After removing less than 300 objects from the main-belt, the ring is able to account for more than 99$\\%$ of the remaining perturbation on all the inner planets. Since the amplitude of the global effect can reach several hundreds of meters in terms of the Earth-Mars distance, it is advisable to keep the ring in an ephemeris model of the Solar System.  \\begin{appendix}"
    },
    "1004/1004.1436_arXiv.txt": {
        "abstract": "We have identified a sample of cool field brown dwarf candidates using IRAC data from the {\\it Spitzer} Deep, Wide-Field Survey (SDWFS).  The candidates were selected from 400,000 SDWFS sources with $[4.5] \\leq 18.5$ mag and required to have $[3.6]-[4.5] \\geq 1.5$ and $[4.5] - [8.0] \\leq 2.0$ on the Vega system.  The first color requirement selects objects redder than all but a handful of presently known brown dwarfs with spectral classes later than T7, while the second eliminates 14 probable reddened AGN. Optical detection of 4 of the remaining 18 sources implies they are likely also AGN, leaving 14 brown dwarf candidates. For two of the brightest candidates (\\teight\\ and \\teightnew), the spectral energy distributions including near-infrared detections suggest a spectral class of $\\sim$ T8. The proper motion is $< 0\\farcs25$ yr$^{-1}$, consistent with expectations for a luminosity inferred distance of $>70$~pc.  The reddest brown dwarf candidate (\\ydwarf) has $[3.6] - [4.5]=2.24$ and $H - [4.5] > 5.7$, redder than any published brown dwarf in these colors, and may be the first example of the elusive Y-dwarf spectral class. Models from \\citet{Burrows:03} predict larger numbers of cool brown dwarfs should be found for a \\citet{Chabrier:03} mass function. Suppressing the model [4.5] flux by a factor of two, as indicated by previous work, brings the Burrows models and observations into reasonable agreement. The recently launched Wide-field Infrared Survey Explorer ({\\it WISE}) will probe a volume $\\sim40\\times$ larger and should find hundreds of brown dwarfs cooler than T7. ",
        "introduction": "}  Although first predicted to exist in 1963 \\citep{Kumar:63, Hayashi:63}, brown dwarfs were not discovered until decades later.  The first viable brown dwarf candidate was GD 165B \\citep{Becklin:88}, an L dwarf whose exact nature as hydrogen-burning star or brown dwarf has yet to be ascertained \\citep{Kirkpatrick:99}.  The first undisputed brown dwarf, and the first T dwarf, was Gl 229B \\citep{Nakajima:95}, whose telltale methane absorption implied an effective temperature too low for a normal star.  In the late 1990's the advent of large-area surveys with near-infrared capability -- the Two Micron All-Sky Survey \\citep[2MASS;][]{Skrutskie:06}, the Sloan Digital Sky Survey \\citep[SDSS;][]{York:00}, and the Deep Near-Infrared Survey of the Southern Sky \\citep[DENIS;][]{Epchtein:97} -- uncovered hundreds more examples and enabled the study of brown dwarfs as a population in their own right \\citep{Kirkpatrick:05}\\footnote{see {\\tt DwarfArchives.org} for a full list}.  The latest spectral type brown dwarfs currently known are T8 and T9 dwarfs found by 2MASS, the United Kingdom Infrared Deep Sky Survey \\citep[UKIDSS;][]{Lawrence:07}, and the Canada France Brown Dwarf Survey \\citep{Delorme:08b}. The coolest of these have effective temperatures of $\\sim$550 K and implied masses of around 15-35 $M_{\\rm Jupiter}$ for assumed ages of 1-5 Gyr \\citep{Warren:07, Burgasser:08, Burningham:08, Delorme:08a, Leggett:09}. Cooler field brown dwarfs must exist, however, as objects of much lower implied mass have been identified in young clusters such as the Orion Nebula Cluster \\citep{Zapatero:02, Weights:08}, Upper Scorpius \\citep{Lodieu:07} and Chamaeleon I \\citep{Luhman:05}, or as companions to other low-mass cluster members \\citep[e.g.,][]{Luhman:06}\\footnote{see also {\\tt http://vlmbinaries.org}}. Finding and characterizing such colder field objects will set important boundary conditions on star formation processes and determine the total amount of mass in stars, a key ingredient in modeling galaxy formation. Identifying examples of such objects is also important to the study of very cold, planet-like atmospheres. A leading question is whether a new spectral class beyond T, dubbed ``Y'', will be needed \\citep{Kirkpatrick:08}.  The two shortest wavelength bands in the {\\it Spitzer} Infrared Array Camera \\citep[IRAC;][]{Fazio:04a} were designed to identify cool brown dwarfs from the signature due to strong methane absorption at $3.6 \\mu$m coupled with a relative lack of absorption at $4.5 \\mu$m \\citep{Fazio:98}.  Finding the coolest and nearest brown dwarfs is a key objective for the {\\it Wide-field Infrared Survey Explorer} \\citep[WISE,][]{Liu:08}, which launched 2009 Dec. 14, and hence two of its four imaging bands are at similar wavelengths (3.4 and $4.6 \\mu$m).  While a number of brown dwarf companions have been found using IRAC, prior to the {\\it Spitzer} Deep, Wide-Field Survey \\citep[SDWFS;][]{Ashby:09} only a single isolated field brown dwarf, of spectral class T4.5, has been identified to date on the basis of IRAC data \\citep{Stern:07}.  This object, IRAC~J1429050.8+333011, was found in the IRAC Shallow Survey \\citep{Eisenhardt:04} and was required to be unresolved in complementary NOAO Deep Wide-Field Survey \\citep[NDWFS;][]{Jannuzi:99} $I$-band data, which necessitated $I < 23$.  Here we remove the limitation of requiring optical detection and use the deeper and more reliable SDWFS IRAC data to search for cooler brown dwarf candidates. ",
        "conclusions": "The SDWFS search confirms the impression from the IRAC Shallow Survey \\citep[][Fig. 4b]{Eisenhardt:04} that at high Galactic latitude, objects with IRAC colors as red as the coolest known brown dwarfs are rare.  Of 367,176 SDWFS sources meeting criteria (i) and (iii) in \\S\\ref{sec:select}, less than one in ten thousand is a real source meeting criterion (ii), i.e. ($[3.6] - [4.5] \\geq 1.5$, equivalent to $F_{\\nu}(4.5)/F_{\\nu}(3.6) > 2.5$). Only two real objects have $[3.6] - [4.5] > 2$ --- presumably the realm inhabited by the elusive Y-dwarfs --- making them an order of magnitude rarer still. Although we were careful {\\it not} to require [3.6] detection, all brown dwarf candidates, and all but one DOG (SDWFS~J143819.58+340957.3 --- Figure \\ref{fig.stamps}) are in fact clearly detected in [3.6].  A blackbody with $[3.6] - [4.5] \\geq 1.5$ would have $T_{\\rm BB} \\simlt 500$K, while a power-law spectrum would need $\\alpha > 4$ where $F_{\\nu} \\propto \\nu^{-\\alpha}$. Such spectra might arise from cool brown dwarfs or warm dust, or from obscuration of hotter spectra by dust. For brevity, we often substitute $C$ for $[3.6] - [4.5]$ in the remainder of the discussion.    \\subsection{Dusty Sources}  Dust enshrouded carbon stars and AGB stars \\citep[e.g.][]{Cutri:89}, or class 0 or I protostars \\citep[e.g.][]{Enoch:09} can have very red IRAC colors due to warm dust emission (Figure \\ref{fig:colcol}). However for the $T_{\\rm BB} \\simlt 500$K needed to produce $[3.6] - [4.5] \\geq 1.5$, the corresponding  $[4.5] - [8]$ is  $> 2.3$, violating criterion (iv). Significant emission from cooler dust is typical for such sources, which is even less consistent with the longer wavelength photometry for the brown dwarf candidates. Hence it is possible that some of the sources classified as DOGs may in fact be stellar.  Nevertheless Figure~1 shows that the locus of likely AGB stars from \\citet{Robitaille:07} lies near the brown dwarf selection region, with three out of 23 with $C \\equiv [3.6] - [4.5] \\geq 1.5$ falling inside it. AGB stars associated with our Galaxy are unlikely contaminants, both due to their faintness and to the high Galactic latitude of the field (67 degrees). The absolute [4.5] magnitude of a typical AGB star is $\\sim -10$ \\citep{Cutri:89}, putting AGBs with the characteristic $[4.5] \\sim 18$ values found in Table~1 at a distance of order 3 Mpc. \\citet{Mauron:08} finds three AGB stars more than 100 kpc from the Sun, presumably from the disruption of tidally captured dwarf galaxies. If this space density is typical, there could be of order 10 AGB stars in the SDWFS volume at a distance near 3 Mpc. However, using the NASA Extragalactic Database, we find no galaxies brighter than 15th mag (optical, i.e $\\sim 1000\\times$ less luminous than the Milky Way at 3 Mpc) and with redshifts $< 1000$ km/s within 5 degrees ($\\sim300$ kpc at 3 Mpc) of the field.  Note that AGB stars are often large amplitude variables. \\citet{Rejkuba:03} and \\citet{Davidge:04} find that most AGB stars in NGC 5128 and M32 respectively are variable. \\citet{Rejkuba:03} give an average K band amplitude of 0.77 mag and a period of 395 days, similar to values found by \\citet{Glass:95} for Galactic AGB stars. The threshold for the \\citet{Robitaille:07} sources is a factor of two (0.75 mag). The peak to peak variation in the [4.5] mags between the four SDWFS epochs (which span four years) exceeds the factor of two level for two of the sources in Table~1. For one of these,  SDWFS~J143712.48+334516.5, inspection of the data shows this is because its [4.5] brightness is spuriously high in one epoch due to a cosmic ray in the aperture, which was rejected in the combined 4-epoch measurement. For the other source, SDWFS~J143821.36+353523.3, although the peak to peak variation is 1.5 mag, the rms is 0.66 mag, which is only slightly more than a one sigma excess above the median variability at [4.5] for this magnitude \\citep{Kozlowski:10}. Because there is little evidence for excess variability in the Table~1 sources, and no obvious source for intergalactic AGB stars, we consider such stars unlikely to be a significant contaminant for our brown dwarf candidate sample.  \\citet{Dey:08} describe a sample of 2603 objects in the Bo\\\"otes field selected to have $R - [24] \\ge 14$, which they interpret as dust obscured galaxies (DOGs). Of the 2491 DOGs from \\citet{Dey:08} with SDWFS IRAC photometry, 12 have $C  \\ge 1.5$, but all of these have $[4.5] - [8] > 2$ so they do not appear in Table~1. DOGs with $C \\ge 1.5$ (either from the \\citet{Dey:08} $R - [24]$ selection, or classified as DOGs here from due to $[4.5] - [8] > 2$) are plotted as open circles in Figure \\ref{fig:colcol}. No attempt is made to distinguish AGN with $C > 1.5$ from DOGs, since DOGs include AGN. Objects plotted as dots in Figure 1 satisfy the \\citet{Stern:05b} AGN criteria.  Of these four have $C > 1.5$, and with $[4.5] - [8] \\sim 3$ they fall in the midst of the DOG colors.  \\citet{Dey:08} suggest the $24\\mu$m emission arises from warm dust heated by an AGN for the brighter sources, or from redshifted PAH emission at $z\\sim2$ for the fainter sources. The red IRAC colors for these fainter DOGs are likely due to obscuration of stellar light by dust. We estimate $A_V \\simgt 6$ to produce $C \\ge 1.5$ at $z \\sim 2$. For the reddest DOG, SDWFS~J143819.58+340957.3 (Figures \\ref{fig.stamps} and \\ref{fig:BD_sed}), $A_V$ well above 10 is indicated. As noted in \\S\\ref{sec:OtherData}, {\\em none} of the objects in Table~1 are detected at $24\\mu$m, while of the 14 sources classified here as DOGs from IRAC photometry (i.e $C \\ge 1.5$ and $[4.5] - [8] > 2$), 10 of the 13 observed at $24\\mu$m were detected.  The 14 IRAC-classified DOGs were also detected in $B_W$ in 9 cases, likely due to redshifted Ly-$\\alpha$ emission, with 8 of the $B_W$ detected objects having detections in $R$ and $I$ as well. The bottom portion of Table~1 lists three brown dwarf candidates (based on their IRAC photometry) with faint optical detections, and a fourth with a marginal optical detection, under the heading ``Likely AGN.\"  These have been marked with shaded gray circles in Figure \\ref{fig:colcol} to indicate that they are likely DOG variants. The potential for additional DOG/AGN contamination of the sample is discussed in \\S \\ref{sec:complete}.  \\subsection{Brown Dwarfs\\label{sec:BDs}}  Even for $T_{\\rm eff} < 500$K, brown dwarf spectra are expected to have $[4.5] - [8] < 2$ due to molecular absorption features \\citep[Figures \\ref{fig:colcol} and \\ref{fig:BD_sed},][]{Burrows:03}. If the sources in Table~1 are brown dwarfs, they have spectral classes later than T6, based on their $[3.6] - [4.5]$ colors \\citep[Figure 10,][]{Patten:06}. Only two of the 86 sources in \\citet{Patten:06} have $C \\equiv [3.6] - [4.5] \\geq 1.5$: GJ 570D (T7.5) with a color of 1.68, and 2MASS 0415-0935 (T8.0) with a color of 1.82.  At the time this paper was submitted, only three other brown dwarfs had colors this red: ULAS0034 with a color of 1.81 and spectral class T8.5 \\citep{Warren:07}; ULAS1335 \\citep[$C=2.05$, T9;][]{Burningham:08}; and 2MASS 0939-24 \\citep[$C=2.10$, a possible T8 binary;][]{Burgasser:08}\\footnote{ After submission, five others appeared in \\citet{Leggett:10}: 2MA 0348-60, $C=1.53$, T7; ULAS 2321+13, $C=1.64$, T7.5; ULAS 1238+09, $C=1.75$, T8.5; 2MA 1114-26, $C=1.78$, T7.5; and Wolf 940B, $C=2.01$, T8.5. }. If it is confirmed, with $C = 2.24$ \\yd\\ would have the reddest IRAC color of any brown dwarf yet found (Figure \\ref{fig:colcol}).  While the optical photometry does not reach the depths expected for T  and late L dwarfs with $[4.5] \\simgt 18$, the typical $I - [4.5] \\simgt 6$ limits argue against the sources in Table~1 being significantly hotter than brown dwarfs. From Figure 8 of \\citet{Hawley:02} and Fig. 10 of \\citet{Patten:06}, the onset of L dwarfs is at $i-[4.5] \\sim 6.2$, equivalent to $I-[4.5] \\sim 5.8$.  The $J - [4.5]$ color ranges from 2.49 to 3.99 for the previously published brown dwarfs with $C \\geq 1.5$. For \\teightshort\\ we find $J - [4.5] = 3.56$, which together with $C = 1.84$ is very similar to the observed values for ULAS0034. Detailed modeling of ULAS0034 lead \\citet{Warren:07} and \\citet{Leggett:09} to conclude that it has $T_{\\rm eff} \\approx 600$K. For \\teightshortnew\\ the $J - [4.5] = 3.11$ and $C = 1.91$ are each about 0.1 mag redder than the corresponding 2MA 0415-0935 (T8.0) values. Figure \\ref{fig:BD_sed} shows the photometry for \\teightshort\\ overplotted with a 600K model from \\citet{Burrows:03}. The $J - [4.5]$ limits for other brown dwarf candidates range from $>2.9$ to $>4.4$, consistent with spectral classes beyond T6.  \\citet{Burningham:08}, \\citet{Warren:07}, and \\citet{Leggett:10} suggest that $H - [4.5]$ is tightly correlated with $T_{\\rm eff}$. Figure~\\ref{fig:hm4p5} shows $H - [4.5]$ as a function of spectral type. The published brown dwarfs with $C \\geq 1.5$ have $H - [4.5]$ colors ranging from 2.98 to 4.34, while the $2\\sigma$ limits here range from $>2.1$ to $>3.7$ ($>5.7$ for \\yd\\, see below), consistent with spectral classes beyond T5. For \\teightshort\\ $H - [4.5] = 3.49$ is similar to the T8 brown dwarf 2MA 0415-0935. We include \\teightshortnew\\ in figure~\\ref{fig:hm4p5} as a grey-shaded point with $H-[4.5]=3.46\\pm0.7$. This assumes $J-H \\sim -0.35$ based on photometry for late-T objects from \\citet{Patten:06} and \\citet{Leggett:10} together with our $J=21.17$ detection to estimate $H=21.52$, with the error bar consistent with the observed $2\\sigma$ limit of $H>20.8$.  The available $K_s$ data offer fewer constraints than $J$ and $H$. The published brown dwarfs with $C \\geq 1.5$ have $K - [4.5]$ colors ranging from 2.98 to 5.17, while the $2\\sigma$ limits here range from $K_s>1.1$ to $>2.4$, consistent with spectral classes beyond L5.  From their $J - [4.5]$, $H - [4.5]$, and $[3.6] - [4.5]$ colors and limits, we associate a spectral class of T8 for \\teightshort, and (more tentatively) T8.5 for \\teightshortnew. Assuming $M_{4.5} = 13.5$ for \\teightshort\\ \\citep[based on][]{Patten:06}, the luminosity distance is $\\sim 70$ pc. As noted in \\S \\ref{sec:prop_mot}, neither source shows significant proper motion, with an estimated $0\\farcs25\\, {\\rm yr}^{-1}$ limit for sources with $[4.5] < 18$. This is unsurprising, since typical proper motions should be $\\sim 0\\farcs1\\, {\\rm yr}^{-1}$ at this distance, using an average tangential velocity from a volume-limited sample of stars of 37 km~s$^{-1}$ \\citep[or $\\sim 4$ AU in 6 months]{Reid:97}.   With $C=2.24$, $[4.5] - [8] < 1.6$, and $H - [4.5] > 5.7$, the SED for \\yd\\ falls in previously unpopulated regions of color space (Figures \\ref{fig:colcol} and \\ref{fig:hm4p5}), and hence this object may be a member of the long-sought Y-dwarf class. Figure \\ref{fig:BD_sed} shows a 400K \\citet{Burrows:03} model which roughly agrees with the observations for \\yd. If the 1.3 mag drop in [4.5] luminosity between the 400K and 600K \\citet{Burrows:03} models is applicable, $M_{4.5} \\sim 15$ for \\yd\\, implying a distance $\\sim 50$ pc.    Again this is consistent with the $0\\farcs4\\, {\\rm yr}^{-1}$ proper motion limit noted in \\S \\ref{sec:prop_mot} for the $[4.5] = 18.47$ mag of \\yd, vs. the typical expected proper motion of $\\sim 0\\farcs15\\, {\\rm yr}^{-1}$ at 50 pc. Alternative explanations for \\yd\\ must account for the non-detections in [5.8], [8], and [24].  \\subsection{Reliability and  Completeness\\label{sec:complete}}  The presence of three (possibly four) optical detections amongst the 18 sources which satisfy the IRAC color selection criteria shows that those criteria do not produce a pure brown dwarf sample, and suggests that additional AGN /DOGs may have scattered into the brown dwarf selection region. Monte-Carlo simulations using the existing SDWFS catalog and error distribution were carried out to evaluate the expected level of such contamination.  The pool of 11,907 SDWFS sources was identified which satisfied criterion (iii) in \\S \\ref{sec:select} and relaxed versions of criteria (i) and (ii), i.e.  $[4.5] \\leq 19.0$ and $C \\equiv [3.6] - [4.5] \\geq 1.0$. For each source in this pool, the associated flux errors was used to generate 10,000 realizations of the IRAC photometry and to find the likelihood that each source would meet the full color selection criteria given in \\S \\ref{sec:select}.  The summed likelihood of was 67.3, with 34 sources having likelihoods greater than $50\\%$ (and a summed likelihood of 26.3). The summed likelihood of selection for the 5202 sources with $1.0 < C < 1.1$ was less than 0.1, indicating contamination from bluer sources is not important. Note that $C > 1$ corresponds to brown dwarf spectral types later than T5 \\citep{Patten:06, Leggett:10}. Visual inspection was carried out in the same manner described in \\S \\ref{sec:select} for the 86 sources with individual likelihood $\\geq 20\\%$, and for 40 representative sources with lower likelhooods.  Based on this inspection, one third of the sources meeting the color criteria would be classified as artifacts, with the remainder equally divided between objects in or very near the cold brown dwarf color selection space, objects with AGN colors scattering into the selection criteria (e.g., with faint detections at $5.8$ and $8.0\\, \\mu$m), and objects for which the distinction between AGN and cold brown dwarf colors was uncertain. We infer from this that 1/3 to 2/3 of the 18 objects in Table~1 are likely to be objects whose true colors are consistent with cold brown dwarfs, (i.e. that 6 to 12 of the 14 sources in the upper part of Table~1 are likely real brown dwarfs).  A complementary Monte-Carlo calculation was made to assess the completeness of the color-selected sample given typical photometric errors as a function of magnitude for the SDWFS data.  The probability that a source would meet the color selection criteria was evaluated using 10,000 realizations of sources as a function of magnitude and color over the range $15 \\leq [4.5] \\leq 19.5$ and $1.0 \\leq [3.6] - [4.5] \\leq 2.6$.  As expected, the probability is $\\sim50\\%$ for bright sources with $[3.6] - [4.5] = 1.5$. For objects $[3.6] - [4.5] = 1.6$, 90\\% will be selected at $[4.5] = 17.0$, and 90\\% of objects with $[3.6] - [4.5] = 1.9$ are selected at $[4.5] = 18.0$. Applying the appropriate percentages as a function of magnitude and color to the sources in the upper part of Table~1, we find an average completeness of $\\sim60\\%$.  However it is also the case that warmer brown dwarfs whose true color is bluer than $C=1.5$ can scatter into the sample. Using models for the true distribution of brown dwarf magnitudes and colors (\\S \\ref{sec:counts}) in conjunction with the completeness calculations, we find that this effect closely compensates for losses due to incompleteness.  With the finding that sources hotter than T6 do not contribute appreciably, and the reliability estimate, this implies that the true population of cool brown dwarfs meeting the selection criteria is between 6 and 12. Some of these are likely to be unresolved binary brown dwarf systems \\citep[see e.g.][]{Burgasser:07}, but this has a relatively small effect on the number density because the increase in numbers due to binaries is compensated for by the larger volume over which they are detectable in a flux limited sample.   If a fraction $B$ of the sample is equal mass binaries, the net effect is a reduction of $B(1 - \\sqrt(2)/2) \\approx 0.3B$ in a volume limited sample. Hence we do not correct for binarity, and in the following section we take 9 brown dwarfs with $C > 1.5$ as representative, of which 8 have $1.5 < C < 2$ and one has $C > 2$.  \\subsection{Brown Dwarf Counts\\label{sec:counts}}  We compare our source counts to the models of \\citet{Burrows:03}, who give a grid of 32 cool brown dwarf models with cooling curves parameterized by mass $\\mu$ and age $t$. From the tabulated effective temperature $T_{\\rm eff}$, gravity $g$ and mass we have computed the luminosity $L$, and set up a linear interpolation in $\\log L$ and $\\log T_{\\rm eff}$ {\\it vs.} $\\log \\mu$ and $\\log t$ to give the luminosity and effective temperature for any mass and age.  The expected number $N$ of detectable brown dwarfs can be computed for any mass function and age distribution using the distance $r$ as a function of brown dwarf magnitude $m$ \\begin{eqnarray} r & = & 10^{1+(m-M(\\mu,t))/5} \\nonumber \\\\ N & = & \\Omega \\int \\int \\int p[m,C(\\mu,t)]\\ n(\\mu,t)\\ r^2 \\frac{dr}{dm}\\ dm\\ d\\mu\\ dt \\end{eqnarray} where $\\Omega$ is the survey area, $M(\\mu,t)$ is the absolute magnitude as a function of mass and age, $C(\\mu,t)$ is the color as a function of mass and age, $n(\\mu,t)$ is the number density of brown dwarfs per unit mass and age, and $p(m,C)$ is the probability from the Monte-Carlo completeness calculation in \\S \\ref{sec:complete} of a brown dwarf with magnitude $m$ and color $C$ being selected. In general the number of sources scattered into the acceptance region was quite similar to the number scattered out of the region, so the Monte-Carlo completeness corrections were small.  Assuming a uniform distribution in age between 100 million and 10 billion years, a \\citet{Chabrier:03} log normal mass function peaking at $0.079 M_\\odot$, $\\Omega = 10\\, {\\rm deg}^2$, and a magnitude limit of $[4.5] < 18.5$ or a flux $> 7.15 \\mu$Jy, the predicted numbers are 55 sources with $1.5 < C < 2$ (where $C \\equiv [3.6]-[4.5]$) and 63 with $C > 2$.  Since we estimate only 8 sources are brown dwarfs with $1.5 < C < 2$ and one with $C > 2$, the hypothesis that \\citet{Chabrier:03} and \\citet{Burrows:03} are both correct can be rejected.  There are two problems: the predicted ratio $N(C > 2)/N(1.5 < C < 2) \\approx 1.1$ is much higher than the observed $1/8$, and the predicted $N(1.5 < C < 2)$ is too high.  While power law mass functions ($n \\propto M^{-\\alpha}$) with $\\alpha$ near 0 predict lower counts, the ratio of counts in the color bins is still too high, and has only a weak dependence on $\\alpha$.  However \\citet[]{Patten:06} and \\citet{Golimowski:04} show that the observed [4.5] or $M$-band flux is substantially lower than predicted by these models.  This is also apparent in Figure \\ref{fig:colcol}. Detailed modeling for the known $C \\geq 1.5$ brown dwarfs \\citep[e.g. Figure 7 of ][]{Leggett:09} finds $T_{\\rm eff} = 550 - 800$K for these objects. The Burrows models for these temperatures predict $[3.6] - [4.5]$ colors which are significantly redder than observed, and $[4.5] - [8]$ colors which are significantly bluer than observed, and suppressing the model [4.5] flux corrects this.  Suppressing the model [4.5] flux has a strong effect on both the ratio problem and the number problem.  This flux suppression, presumably due to the CO fundamental at $4.7\\mu$m, has been attributed to non-equilibrium chemistry altering the expected CO absorption depths \\citep{Hubeny:07}. Suppression of flux in the spectrum causes some backwarming, so the effective temperature increases to \\begin{equation} T_{\\rm eff}^\\prime = T_{\\rm eff}/(1-Sf_2/R_2)^{1/4} \\end{equation} where $f_2$ is the fraction of the bolometric luminosity in the IRAC channel 2 ([4.5]) detection band from the model, $S$ is the [4.5] suppression, and $R_2$ is the resolution of the IRAC $4.5 \\mu$m filter.  We use $R_2 = 3$ which is somewhat lower than the actual $R_2 = 4.5$, to allow for suppression of flux outside of the [4.5] passband.  The $4.5 \\mu$m flux fraction decreases to $f_2^\\prime = (1-S)f_2(T_{\\rm eff}^\\prime)$. Increased suppression decreases the ratio $N(C > 2)/N(1.5 < C < 2)$ and also reduces the expected number counts, as shown by the heavy solid curve in Figure \\ref{fig:number-vs-supression}.  The horizontal band shows the estimated range of 5 to 11 with $1.5 < C <2$.  The lighter solid curve peaking near $S = 0.52$ shows the Poisson likelihood of a given flux suppression based on the numbers seen in the two color bins, assuming that the \\citet{Chabrier:03} single object mass function and uniform age distribution are correct.  The likelihood is maximized by a suppression of $S = 0.52$. Similar results are obtained for a power law mass function with $\\alpha = 1.3$. This suppression of $S = 0.52$ agrees well with the estimate by \\citet{Golimowski:04} that the [4.5] flux is suppressed by a factor between 1.5 and 2.5, which corresponds to $S = 0.33$ to 0.60 in our terminology. In other words, suppressing the \\citet{Burrows:03} model by a factor of two brings both the models into agreement with both the observed mid-infrared colors and number counts.  Another potential solution is to adopt a different mass function, which, like the predicted luminosities and colors, is not well known for these very low mass objects. Since the \\citet{Burrows:03} models give the luminosity vs. mass and age for single brown dwarfs, it is not strictly correct to use system mass functions in these calculations, but to provide a range of examples we have included  in Figure 6 both the \\citet{Chabrier:03} single object and system mass function, as well as the \\citet{Bochanski:09} mass function. Based on SDSS observations of late M dwarfs, \\citet{Bochanski:09} find a log normal mass function peaking at $0.27 M_\\odot$. This results in a much smaller predicted number density of brown dwarfs than for the \\citet{Chabrier:03} single object mass function, and requires a smaller 4.5 $\\mu$m flux suppression ($S =0.26$) to agree with the observed SDWFS counts, as shown by the dotted curves in Figure \\ref{fig:number-vs-supression}. However, this would not account as well for the observed $4.5\\mu$m flux discrepancy. The \\citet{Chabrier:03} system mass function, which peaks at $0.2 M_\\odot$, shows intermediate results, matching the SDWFS counts for a flux suppression of $S = 0.31$, as shown by the dashed curves in Figure \\ref{fig:number-vs-supression}. The number density data shown in \\citet{Bochanski:09} flatten significantly at the low mass end, so this log normal mass function may not be reliable in the brown dwarf regime. For both the \\citet{Bochanski:09} mass function and significant flux suppression of the Burrows models to be correct would suggest that nearly all of the SDWFS brown dwarf candidates are dusty galaxies with no evidence for star formation in the rest UV or mid-IR.  Scaling from the SDWFS counts to the all-sky {\\it WISE} survey, which launched  on 2009 Dec. 14, is simpler than comparing to models.  The {\\it WISE} sensitivity requirement at $4.6\\mu$m is $160\\mu$Jy, resulting in a surveyed volume which is $\\sim40 \\times$ greater than SDWFS, and hence $\\sim 250$ to 500 similarly cool brown dwarfs for a Euclidean distribution. These {\\it WISE} brown dwarfs will be the nearest examples, with correspondingly brighter fluxes and larger parallaxes and proper motions, making followup observations much easier. This should enable a definitive determination of the properties of the ultracool brown dwarf population."
    },
    "1004/1004.4213_arXiv.txt": {
        "abstract": "Light gravitinos, with mass in the eV to MeV range, are well-motivated in particle physics, but their status as dark-matter candidates is muddled by early-Universe uncertainties.  We investigate how upcoming data from colliders may clarify this picture.  Light gravitinos are produced primarily in the decays of the next-to-lightest supersymmetric particle, resulting in spectacular signals, including di-photons, delayed and non-pointing photons, kinked charged tracks, and heavy metastable charged particles.  We find that the Tevatron with $20~\\ifb$ and the 7 TeV LHC with $1~\\ifb$ may both see evidence for hundreds of light-gravitino events.  Remarkably, this collider data is also well suited to distinguish between currently viable light-gravitino scenarios, with striking implications for structure formation, inflation, and other early-Universe cosmology. ",
        "introduction": "\\label{sec:intro}  Supersymmetry is one of the most promising ideas for new physics beyond the Standard Model.  Supersymmetric theories that incorporate local supersymmetry (or supergravity) predict the existence of the gravitino, the spin-3/2 superpartner of the graviton.  When supersymmetry is broken, the gravitino acquires a mass through the super-Higgs mechanism, ``eating'' the spin-1/2 goldstino, the Goldstone fermion associated with spontaneously-broken local supersymmetry \\cite{Volkov:1973jd,Fayet:1974jb,de Wit:1975th,% Deser:1977uq}.  In contrast to other superpartners, the gravitino can have a mass $\\mG$ that is not at the weak scale ${\\mweak \\sim 100~\\GeV - 1~\\TeV}$, and viable models exist for gravitino masses as low as the eV scale and as high as 100 TeV.  In this work, we consider light gravitinos, with mass in the eV to MeV range.  Such gravitinos are highly motivated in particle physics, as they emerge in models with gauge-mediated supersymmetry breaking (GMSB), in which constraints on flavor violation are naturally satisfied~\\cite{Dine:1981za,% Dimopoulos:1981au,Nappi:1982hm,AlvarezGaume:1981wy,Dine:1994vc,% Dine:1995ag}.  Light gravitinos also have cosmological motivations.  In particular, they are the original supersymmetric dark-matter candidate~\\cite{Pagels:1981ke}.  Assuming a high reheating temperature, gravitinos are initially in thermal equilibrium and then freeze out while still relativistic.  As we discuss in detail below, their resulting relic density is \\begin{equation} \\Omega_{\\tilde{G}} h^2 \\simeq \\left[ \\frac{\\mG}{1~\\keV} \\right] \\left[ \\frac{106.75}{g_{*S,f}} \\right] \\ , \\label{omega} \\end{equation} where $g_{*S,f}$ is the number of relativistic degrees of freedom at freeze out, and has been normalized to the total number of degrees of freedom in the Standard Model.  When originally proposed in the 1980's, uncertainties in $h$ and the total matter relic density allowed $\\mG \\sim \\keV$.  This led to a simple and attractive gravitino--dark-matter scenario, consistent with standard Big Bang cosmology, in which the Universe cooled from some high temperature, and keV gravitinos froze out and now form all of the dark matter.  In the intervening years, however, a variety of astrophysical constraints have greatly complicated this picture.  First, the dark-matter relic density is now known to be ${\\Omega_{\\textrm{DM}} h^2 \\simeq 0.11}$.  Second, constraints on structure formation, as probed by galaxy surveys and Lyman-$\\alpha$ forest observations, require that the bulk of dark matter be cold or warm~\\cite{Viel:2005qj}.  As we will discuss more fully below, this leads to three scenarios of interest: \\begin{enumerate} \\setlength{\\itemsep}{1pt}\\setlength{\\parskip}{0pt}\\setlength{\\parsep}{0pt} \\item $\\mG \\alt 15-30~\\eV$: Gravitinos are produced by the standard cosmology leading to \\eqref{omega}; they are hot dark matter, but their contribution is small enough to be consistent with the observed small-scale structure.  Some other dark-matter particle is required. \\item $15-30~\\eV \\alt \\mG \\alt \\textrm{few}~\\keV$: Non-standard cosmology and a non-standard gravitino production mechanism are required, both to avoid overclosure and to cool the gravitinos to satisfy small-scale-structure constraints.  Some other dark-matter particle may be required. \\item $\\mG \\agt \\textrm{few}~\\keV$: Non-standard cosmology is required to dilute the thermal relic density of \\eqref{omega}.  Gravitinos produced by thermal freeze out are cold enough to be all of the dark matter. \\end{enumerate} Note that the original ``keV gravitino'' scenario, previously favored, is now the most disfavored, in the sense that it is excluded by both overclosure and small-scale-structure constraints.  All of the possibilities are rather complicated, however, as in each case, some additional physics is required, either to provide the rest of the dark matter or to modify the history of the early Universe to allow gravitinos to be all of the dark matter.  In this paper, we discuss how collider data may help clarify this picture.  Light gravitinos are primarily produced at colliders in the decays of the next-to-lightest supersymmetric particle (NLSP).  It is a remarkable coincidence that modern particle detectors, with components placed between 1 cm to 10 m from the beamline, are beautifully suited to distinguish between the NLSP decay lengths predicted in scenarios 1, 2, and 3.  For example, the decay length of a Bino NLSP decaying to a gravitino is~\\cite{Cabibbo:1981er} \\begin{equation} c \\tau \\simeq 23~\\cm \\left[ \\frac{\\mG}{100~\\eV} \\right]^2 \\left[ \\frac{100~\\GeV}{m_{\\tilde{B}}} \\right]^5 \\ . \\end{equation} This implies that scenarios 1, 2, and 3 make distinct predictions for collider phenomenology, and the identification of the gravitino collider signatures realized in nature may have far-reaching implications for the early Universe.  Of course, this requires that gravitinos can be produced in sufficient numbers and distinguished from Standard Model backgrounds.  In this work, we determine event rates for a variety of signatures, including prompt di-photons and delayed and non-pointing photons (relevant for neutralino-NLSP scenarios), as well as kinked charged tracks and heavy metastable charged particles (relevant for stau-NLSP scenarios).  We present results for an assumed final Tevatron dataset ($20~\\ifb$ of $2~\\TeV$ $p\\bar p$ collisions), an early LHC dataset ($1~\\ifb$ of $7~\\TeV$ $pp$ collisions), and a future LHC dataset ($10~\\ifb$ of $14~\\TeV$ $pp$ collisions).  We find that the final Tevatron and early LHC data have roughly equivalent sensitivity to these events, with both capable of seeing hundreds of distinctive light gravitino events. The full LHC data greatly extends the reach in parameter space, and may also allow precision measurements of NLSP lifetimes and gravitino masses.  We begin in \\secref{cosmology} by reviewing the cosmological bounds on light gravitinos and discussing how these bounds are relaxed in early-Universe scenarios that differ from the canonical one.  In \\secref{colliders} we then discuss NLSP decays to gravitinos, GMSB models, and current collider constraints.  In \\secref{prospects} we present our results for the number of light-gravitino events at colliders, based on collider simulations, and discuss the cosmological implications.  We summarize our conclusions in \\secref{conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  Light gravitinos in the mass range $\\eV$ to $\\MeV$ appear in GMSB models that naturally avoid flavor violation.  We have examined the decay of supersymmetric particles to light gravitinos at colliders such as the Tevatron and the LHC.  These decays will give rise to dramatic signatures, such as prompt di-photons or non-prompt photons, if the NLSP is a neutralino, or kinked charged tracks or heavy metastable charged particles, if the NLSP is a stau (or some other charged particle).  We find large regions of the gravitino-mass--NLSP-mass parameter space in which the rate for such events may be appreciable at the Tevatron and LHC and which are consistent with current null supersymmetry searches.  Given that ${\\mG\\ll \\mNLSP}$ for these events, the decay kinematics of individual events cannot be used to determine the gravitino mass.  However, the event rate and the distribution of decay locations may be used to narrow the range of NLSP and gravitino masses.  Information about the nature of the NLSP may also be gleaned from the Standard Model decay products.  One of the attractions of supersymmetry has been its ability to provide a natural candidate for the cold dark matter required by a wealth of cosmological observations.  Unfortunately, despite being well-motivated in GMSB models, the canonical light-gravitino scenario does not provide a natural cold-dark-matter candidate.  Nevertheless, this canonical scenario does allow gravitinos with masses ${\\mG \\alt 30~\\eV}$ that compose a fraction of the total dark matter, as determined by current astrophysical constraints on the relic abundance and small-scale structure.  Given that upcoming structure formation observations are expected to probe hot-dark-matter masses as low as ${\\mG \\sim \\eV}$, detection of a gravitino in the mass range ${\\eV \\alt \\mG \\alt 30~\\eV}$ via prompt signals at colliders would have implications for future small-scale-structure measurements.  And although masses ${\\mG \\agt 30~\\eV}$ are disfavored, they may still be possible if the pre-BBN history was different than in the canonical scenario.  Detection of gravitinos in this mass range via non-prompt and metastable signals at colliders would thus have serious implications for early-Universe cosmology, and may provide some insight into the reheating and inflationary eras.  And who knows? There may indeed be new early-Universe physics that results in a gravitino that has the right cosmological abundance and temperature to be the dark matter. \\\\"
    },
    "1004/1004.1093_arXiv.txt": {
        "abstract": "{ Part of the very high energy $\\gamma$-ray radiation coming from extragalactic sources is absorbed through the pair production process on the extragalactic background light photons. Extragalactic magnetic fields alter the trajectories of these cascade pairs and, in turn, convert cosmic background photons to gamma-ray energies by inverse Compton scattering. These secondary photons can form an extended halo around bright VHE sources. } {We searched for an extended emission around the bright blazars Mrk 421 and Mrk 501  using the MAGIC telescope data.} {If extended emission is present, the angular distribution of reconstructed gamma-ray arrival directions around the source is broader than for a point-like source. In the analysis of a few tens of hours of observational data taken from Mrk~421 and Mrk~501 we used a newly developed method that provides better angular resolution. This method is based on the usage of multidimensional decision trees. Comparing the measured shapes of angular distributions with those expected from a point-like source one can detect or constrain possible extended emission around the source. We also studied the influence of different types of systematic errors on the shape of the distribution of reconstructed gamma-ray arrival directions for a point source.} {We present upper limits for an extended emission calculated for both sources for various source extensions and emission profiles. We obtain upper limits on the extended emission around the Mrk~421 (Mrk~501) on the level of $<$ 5\\% ($<$ 4\\% ) of the Crab Nebula flux above the energy threshold of 300 GeV. Using these results we discuss possible constraints on the extragalactic magnetic fields strength around a few times $10^{-15}$~G. } {} ",
        "introduction": "Blazars are well-known extragalactic sources of Very High Energy (VHE) $\\gamma$-rays. Their radiation traverses over large distances through the extragalactic space filled with CMB (Cosmic Microwave Background) and EBL (Extragalactic Background Light) photons and could be absorbed via the pair production process. This effect can lead to a change in the spectral shape of observed radiation and it can be used to constrain the EBL density (see e.g. \\cite{stecker1992}, \\cite{eblpaper}).  A possible extended emission around extragalactic sources of VHE $\\gamma$-rays  was first discussed by \\citet{aharonian_halo}. $\\gamma$-rays with energies larger than 10 TeV are strongly absorbed via a pair production process on EBL and CMB photons relatively close to the source. Secondary $\\gamma$-rays can be produced in a cascade initiated by those primary photons. Note that the redshift dependent energy for which the gamma-ray opacity is equal 1 vary within a factor of $\\sim 3$ in different EBL models \\citep{Kneiske:2003tx, Stecker:2005qs, Franceschini:2008tp, gilmore08, Primack:2008nw, finke10}.  Magnetic fields are very non-uniform and their strength can vary within many orders of magnitude depending on the location in the large scale structure. So far they have been measured only in the galaxies with strength $\\sim\\mu$G \\citep{kz08, beck08}, the cores of galaxy clusters (within the inner 100 kpc) $\\sim$ several $\\mu$G \\citep{ct02}, and near borders of few clusters with strength $10^{-8}$ -- $10^{-7}$~G (on Mpc scales) \\citep{xkhd06,kksp07}. On larger distance scales the magnetic field strengths are not known but they have to be much weaker (see e.g. \\cite{angelis}). Moreover, the theoretical models predict very weak ($B  \\ll 10^{-12}$~G) extragalactic magnetic fields (EGMF) in voids in large scale structures, outside galaxies and galaxy clusters \\citep{kronberg94,grasso00,widrow2002,neronov09}. It is believed that voids make up a significant part of the space volume.  Development of a secondary gamma-ray cascade depends on its original energy and the strength of surrounding magnetic field. Two cases can be considered. First, if the TeV source is located in an intense magnetic field region and the maximal gamma-ray energy is large ($E_{max}\\gg 50$ TeV), gamma-rays would produce $e^+e^-$ pairs directly near the source (within a few Mpc). These will be isotropized in the strong magnetic field which might exist around the source host galaxy or galaxy cluster \\citep{aharonian_halo}.  Second, the mean free path for gamma-rays of moderate energies ($E<50$~TeV) is longer, therefore they can traverse larger distances from the source. Therefore $e^+e^-$ pairs will be produced in regions with much weaker magnetic fields \\citep{plaga95, neronov07, neronov2, kachelriess}. In this case the deflections of the pair trajectories by EGMF are not large enough to make the secondary cascade emission isotropic. Instead the inverse Compton scattering of CMB photons by $e^+e^-$ pairs will produce secondary $\\gamma$-rays with a slightly different direction with respect to the primary photons. The re-direction of the cascade photons into the field of view of the telescope can lead to the appearance of an extended emission around the point-like source, even in the case $B\\le 10^{-12}$~G \\citep{neronov07,neronov2,kachelriess}. In fact the extended emission would be produced between the source and the observer. For a distant observer that extended emission will appear as superimposed onto the point source thus mimicking a halo.  The secondary $\\gamma$-rays also can initiate cascades, provided that the optical depths in the EBL radiation field are still large enough. The energies of further generation of $\\gamma$-rays may be below the energy threshold of VHE $\\gamma$-ray instruments.  Blazars are known to be strongly variable in particular at the $\\gamma$-ray energies (e.g. \\cite{aharonian_pks}, \\cite{quantum_gravity}). The extended emission component cannot follow the original time profile of the emission. This is because the secondary cascade photons do not propagate along the same path as the direct $\\gamma$-rays from the primary source. Instead it will be delayed and stretched to much longer time scales (up to $\\sim10^6$ years at $\\sim 100$~GeV energy for $B\\sim 10^{-12}$~G, \\cite{neronov07,neronov09}, see also \\citet{plaga95,dai02,murase08}). Thus the direct emission from the point source will be overlayed on the extended emission component. The latter can constitute a part of the ``quasi'' constant, quiescent emission.  The first attempt to detect extended emission around extragalactic TeV $\\gamma$-ray sources was performed by the HEGRA instrument for Mrk~501. It yielded only an upper limit of 5-10\\% of the Crab Nebula flux (at energy $\\geq 1$ TeV) on angular scales of 0.5 to 1$^\\circ$, \\citep{hegra_halo}.  In this paper we report on our search for extended emission of VHE $\\gamma$-rays from the bright blazars Mrk~421 and Mrk~501 using the 17m diameter MAGIC Imaging Atmospheric Cherenkov Telescope (IACT). The details of the telescope, its performance and the standard analysis chain are described in \\cite{crab_paper}. Due to the large size of the mirror dish and improved light sensors, the MAGIC trigger threshold ($\\sim 50$~GeV) is $\\sim 2-3$ times lower compared to other operating IACTs.   This paper is structured in the following way: In section 2. we describe the analysis method, which was used for searching for an extended emission. Also, we present a new method that improves the angular resolution. In  section 3. we perform analyses of possible systematic effects which can mimic the existence of an extended emission. Then, in section 4. we describe the data sample used for the analysis. In section 5. we present results of the analysis: the upper limits on the extended emission from Mrk~421 and Mrk~501. Finally, in section 6. we discuss how those upper limits constrain the strength of the EGMF. This is followed by a short discussion on future prospects and conclusions. ",
        "conclusions": "We have developed a novel method based on the Random Forest algorithm for the estimation of the source position. This method improves the angular resolution by $\\sim 20-30\\%$.  In the used data samples consisting of 26h observations of Mrk~501 and 38h of Mrk~421 no extended emission has been detected around these sources. Our study showed that if there is an extended emission around Mrk 501, then its flux is $<$ 4\\% C.U. (see fig.~\\ref{fig_uplim} left), for the analysis energy threshold of 300 GeV. For Mrk 421 the upper limits are less stringent, because the source was in a high emission state, which causes the tail of the point-like emission to extend into the halo region thus creating additional background. The constraint on the extended emission flux from Mrk~421 is $< 5\\%$ C.U. (see fig.~\\ref{fig_uplim} right).  We analyzed different types of systematic errors connected with the observations of extended sources. We found them to be negligible compared to the upper limits obtained on the flux of the extended emission component of Mrk 421 and Mrk 501.  With the second telescope, the MAGIC system can be operated in the stereo mode. By combining simultaneous information from both telescopes, the angular resolution will be improved. This will allow us to perform a search for the AGN halo with a better sensitivity. The stereo observations improve the sensitivity 2.5 times at lower energies ($\\sim 100$  GeV). Moreover, the better signal-to-noise ratio reduces the possible bias due to systematic errors. It will allow us to perform a sensitive search for halos at lower energies, thus extending the range of the EGMF strengths investigated with this method.  A future $\\gamma$-ray telescope project - the Cherenkov Telescope Array is aiming for a much better angular resolution. In this case lower halo extensions can be investigated with an improved sensitivity."
    },
    "1004/1004.3575_arXiv.txt": {
        "abstract": "We present the results of a study of the late-type spiral galaxy NGC\\,0959, \\emph{before} and \\emph{after} application of the pixel-based dust extinction correction described in Tamura \\etal\\ 2009 (Paper~I).  \\emph{Galaxy Evolution Explorer} (\\emph{GALEX}) far-UV (FUV) and near-UV (NUV), ground-based Vatican Advanced Technology Telescope (VATT) \\emph{UBVR}, and \\emph{Spitzer}/Infrared Array Camera (IRAC) 3.6, 4.5, 5.8, and 8.0\\,\\mum\\ images are studied through pixel Color-Magnitude Diagrams (pCMDs) and pixel Color-Color Diagrams (pCCDs).  We define groups of pixels based on their distribution in a pCCD of \\bmir\\ versus \\fuvu\\ colors \\emph{after} extinction correction.  In the same pCCD, we trace their locations \\emph{before} the extinction correction was applied. This shows that selecting pixel groups is not meaningful when using colors uncorrected for dust. We also trace the distribution of the pixel groups on a pixel coordinate map of the galaxy.  We find that the pixel-based (two-dimensional) extinction correction is crucial to reveal the spatial variations in the dominant stellar population, averaged over each resolution element.  Different types and mixtures of stellar populations, and galaxy structures such as a \\emph{previously unrecognized} bar, become readily discernible in the extinction-corrected pCCD and as coherent spatial structures in the pixel coordinate map. ",
        "introduction": "Studying the distribution of different stellar populations within galaxies is crucial to understand the formation and evolutionary history of galaxies. One of the obstacles that prevents us from directly observing these stellar populations is the attenuation and reddening by dust that is randomly distributed among, or projected along the line-of-sight toward, these stellar populations.  This has remained a major problem for nearly a century \\citep[e.g.,][]{trumpler30, mathis77, viallefond82, caplan85, witt92, roussel05, driver08}.  While the main effect of dust is to attenuate the light emitted by stars, dust also scatters light out of the line-of-sight, and may re-direct light from nearby regions into our line-of-sight \\citep[e.g.,][]{witt92, witt96, witt00}.  The impact of extinction---the wavelength-dependent net effect of absorption and scattering---correlates with the spatial distribution of star-forming regions \\citep[e.g.,][]{waller92, deo06}, as it depends on both the geometry of the dust distribution \\citep[e.g.,][]{elmegreen80, walterbos88, calzetti94, witt96, witt00} and on the physical and chemical properties of the dust \\citep[e.g.,][]{vanhouten61, witt92, whittet01, whittet04}.  On typical observational scales, the dust will be intermixed with the stellar populations as thin layers, filaments, and dense clumps.  Even though some light may be scattered into the line-of-sight from the rear and from regions near the line-of-sight, most of the extinction by far will be due to dust that is distributed \\emph{in front} of the stellar populations of interest \\citep{byun92}.  Popular methods to measure dust extinction, such as the use of Hydrogen recombination line ratios \\citep[e.g.,][]{rudy84, scoville01, calzetti05, calzetti07, kennicutt09}, the ultraviolet (UV) spectral slope \\citep[e.g.,][]{calzetti94, kong04}, or ratios of the UV and total infrared fluxes \\citep[e.g.,][]{buat96, calzetti00, boissier04, boissier05}, have various constraints and limitations \\citep[e.g.,][]{xu96, petersen97, regan00, price02, boissier04, rieke04}.  In \\citet[][hereafter Paper~I]{tamura09a}, we introduced a new method to measure the dust extinction in galaxies, and demonstrated that the spatial distribution of dust extinction in a late-type spiral galaxy, NGC\\,0959, can be mapped using only the \\emph{flux ratio of optical \\vv-band and mid-infrared (mid-IR) 3.6\\,\\mum\\ images}.  In the present paper, we present the results of a multi-wavelength (UV--optical--mid-IR) study of the effect of the pixel-based two-dimensional extinction correction of Paper~I on color composite images, a pixel Color--Color Diagram (pCCD), and a pixel coordinate map of NGC\\,0959.  These results show that a detailed pixel-based extinction correction is essential to reveal the nature and distribution of stellar populations in galaxies. ",
        "conclusions": "We have presented the results of a study of NGC\\,0959 using color composite images, a pCCD, and a pixel coordinate map, and demonstrated the importance and potential of a pixel-based extinction correction.  Our study combined ground- and space-based surface photometry, ranging in wavelength from the far-UV (\\emph{GALEX}) through mid-IR (\\emph{Spitzer}/IRAC).  Among the possible combinations of color-magnitude and color-color diagrams examined, we found that the \\bmir\\ versus \\fuvu\\ pCCD proved the most powerful diagnostic of differences in the stellar populations contributing to the flux in a pixel.  After applying the pixel-based, two-dimensional extinction correction described in Paper~I, we defined six different pixel groups in this diagnostic diagram, ranging from pixels for which both \\fuvu\\ and \\bmir\\ colors indicate that their fluxes are dominated by very young, massive stellar populations, to pixels that appear to sample only light from evolved stellar populations.  We demonstrated that it was not possible to meaningfully define such pixel groups \\emph{before} extinction correction.  We then showed that pixels that belong to a given pixel group form well-defined, contiguous regions in a pixel coordinate map, revealing systematic spatial variations in the dominant stellar populations that would not be readily discernible without a two-dimensional correction for extinction.  We were able to report the presence of a previously unrecognized stellar bar, for example.  Our pixel-based two-dimensional method to correct for extinction, based on only 3.6\\,\\mum\\ and $V$-band images (adding $U$ improves fidelity), has the advantage that it is applicable to any galaxy that is significantly resolved at rest-frame $\\sim$3.6\\,\\mum\\ and that it is not restricted to a few individual sightlines within a galaxy, nor to the very nearest galaxies. Although seemingly simple and crude, we have demonstrated that application of this method allowed us to uncover relatively detailed spatial information on the nature of stellar populations and on the star formation history within NGC\\,0959, despite significant spatial variations in the attenuation by intervening dust.  While the spatial resolution in the present study was limited by the \\emph{GALEX} NUV PSF of 5\\farcs3 (FWHM) to linear scales of $\\sim$250\\,pc, in the near future, combination of \\emph{HST}/WFC3 and \\emph{JWST} will provide much higher resolution (FWHM\\,$\\simeq$\\,0\\farcs06--0\\farcs11) at rest-frame UV through mid-IR wavelengths.  These instruments will thus allow us to study the distribution of dust extinction and the underlying stellar populations at high fidelity and at much finer spatial resolution within nearby galaxies, as well as at similar scales of a few hundred pc within more distant galaxies.  Since galaxies at $z$\\,$\\gtrsim$\\,1 are important building blocks of the galaxies that we see today at $z$\\,$\\simeq$\\,0, studying these higher redshift objects with proper two-dimensional extinction estimates will be of significant importance to reveal the nature of galaxy assembly and evolution.    \\vspace{0.5cm} {\\small  This work is funded by NASA/ADP grant NNX07AH50G and NASA/JWST grant NAG\\,5-12460.  R.A.W.\\ was supported in part by NASA/JWST grant NAG\\,5-12460.  We thank Violet Mager for providing the data observed at the Vatican Advanced Technology Telescope (VATT): the Alice P.\\ Lennon Telescope and the Thomas J.\\ Bannan Astrophysics Facility.  We also thank Daniela Calzetti for useful discussions that helped improve our paper.  This study has made use of the NASA/IPAC Extragalactic Database (NED), which is operated by the Jet Propulsion Laboratory (JPL), California Institute of Technology, under contract with the National Aeronautics and Space Administration (NASA).  This study has also made use of NASA's Astrophysics Data System (ADS).  We thank the anonymous referee for constructive comments that helped improve this paper.  }    \\renewcommand{\\baselinestretch}{0.90} {\\footnotesize"
    },
    "1004/1004.3743_arXiv.txt": {
        "abstract": "We announce the identification of a new cataclysmic variable star in the field of the Kepler Mission, KIC J192410.81$+$445934.9.  This system was identified during a search for compact pulsators in the Kepler field.  High-speed photometry reveals coherent large-amplitude variability with a period of 2.94 h.  Rapid, large-amplitude quasi-periodic variations are also detected on time scales of $\\approx 1200$ s and $\\approx 650$ s.  Time-resolved spectroscopy covering one half photometric period shows shallow, broad Balmer and \\ion{He}{1} absorption lines with bright emission cores as well as strong \\ion{He}{2} and Bowen blend emission. Radial velocity variations are also observed in the Balmer and \\ion{He}{1} emission lines that are consistent with the photometric period.  We therefore conclude that KIC J192410.81$+$445934.9 is a nova-like variable of the UX UMa class in or near the period gap, and it may belong to the rapidly growing subclass of SW Sex systems. Based on 2MASS photometry and companion star models, we place a lower limit on the distance to the system of $\\sim 500$ pc.  Due to limitations of our discovery data, additional observations including spectroscopy and polarimetry are needed to confirm the nature of this object. Such data will help to further understanding of the behavior of nova-like variables in the critical period range of $3-4$ h, where standard cataclysmic variable evolutionary theory finds major problems.  The presence of this system in the Kepler mission field-of-view also presents a unique opportunity to obtain a continuous photometric data stream of unparalleled length and precision on a cataclysmic variable system. ",
        "introduction": "} Cataclysmic variables (CVs) are close binary systems in which a white dwarf (WD) accretes material from a companion star.  The secondary star is a late-type dwarf filling its Roche lobe, and, in most cases, the overflowing material forms an accretion disk around the WD primary.  One class of CVs is the nova-like (NL) system.  Historically, a NL system not showing magnetic characteristics is classified as a \\object{UX UMa} star \\citep{1996ASSL..208....3D}.  Non-magnetic NLs exhibit the same variety of optical spectra as the dwarf novae in outburst, ranging from pure emission-line spectra (high inclination) to pure absorption-line spectra (low inclination).  Therefore, NLs are characterized by relatively high mass accretion rates.  \\begin{deluxetable}{cccl} \\tablecolumns{4} \\tablewidth{0pt} \\tablecaption{Astrometric and Photometric Properties of KIC J1924$+$4459\\label{tab.kic_data}} \\tablehead{\\colhead{Parameter} & \\colhead{Value} & \\colhead{Date} & \\colhead{Source}} \\startdata RA(J2000)      & $19^{\\rm h}24^{\\rm m} 10\\fs 82$ & \\nodata & KIC10\\tablenotemark{a}\\\\ Dec(J2000)     & $+44\\degr 59\\arcmin 34\\farcs 9$ & \\nodata & KIC10\\tablenotemark{a} \\\\ $\\mu_{\\rm RA}$  & $24\\pm 8$ mas yr$^{-1}$  & 1980.90 & USNO-B1.0 \\\\ $\\mu_{\\rm Dec}$ & $34\\pm 17$ mas yr$^{-1}$  & 1980.90 & USNO-B1.0 \\\\ $B_p$          & 15.22 & 1951.5195 & USNO-B1.0 \\\\ $R_p$          & \\nodata & 1951.5195 & USNO-B1.0 \\\\ $B_p$          & 15.59 & 1976.56 & YB6\\tablenotemark{b} \\\\ $V_p$          & 15.79 & 1976.56 & YB6\\tablenotemark{b} \\\\ $B_p$          & 15.30 & 1988.4695 & USNO-B1.0 \\\\ $R_p$          & 15.36 & 1991.5795 & USNO-B1.0 \\\\ $I_p$          & 15.72 & 1990.4463 & USNO-B1.0 \\\\ $J$            & $16.839\\pm 0.137$ & 1998.436 & 2MASS \\\\ $H$            & $\\geq 16.736$ & 1998.436 & 2MASS \\\\ $K$            & $\\geq 16.536$ & 1998.436 & 2MASS \\\\ $g$            & 16.182  & 2007.1\\tablenotemark{c} & KIC10\\tablenotemark{a} \\\\ $g-r$          & $-0.116$  & 2007.1\\tablenotemark{c} & KIC10\\tablenotemark{a} \\\\ $r-i$          & $-0.091$  & 2007.1\\tablenotemark{c} & KIC10\\tablenotemark{a} \\\\ $r-z$          & $-0.214$  & 2007.1\\tablenotemark{c} & KIC10\\tablenotemark{a} \\\\ \\enddata \\tablenotetext{a}{Kepler Input Catalog v.~10, \\citealt{2008PhST..130a4034L}} \\tablenotetext{b}{USNO Yellow-Blue Catalog v.~6, D.~Monet (private communications)} \\tablenotetext{c}{KIC10 photometry is a weighted average of data obtained 2006.281, 2007.412, and 2007.486} \\end{deluxetable}  The classification scheme may improve with consideration of more precise photometric and spectroscopic characteristics.  Systems in the 3--4 h orbital period range can be identified as \\object{VY Scl} stars when they are found to undergo low states or can be identified as \\object{SW Sex} stars when their spectra display single-peaked emission features irrespective of inclination angle with absorption components variable along the orbital period \\citep{Thorstensen91,1995cvs..book.....W,1996ASSL..208....3D}. Recently, the SW Sex class has been claimed to be the dominant population of the 3--4 h orbital period CV systems \\citep[e.g.,][]{2007MNRAS.377.1747R,2007MNRAS.374.1359R}.  NLs, generally not (yet) known to display the above characteristics are classified as UX UMa stars from the presence of strong absorption features in Balmer lines with superposed narrow emission cores as well as emission lines of \\ion{He}{1} and \\ion{He}{2}.  The detection of emission in these lines are clear signs of accretion and therefore cannot be mistaken with non-accreting WD$+$M binaries.  Although the historical definition of NL systems is of a non-magnetic CV, there is considerable evidence that magnetism is present in at least some NLs, and that magnetism may be the force behind many of the defining characteristics of SW Sex systems.  Based on profiles of emission lines during eclipses, \\citet{Williams89} proposed that magnetic accretion columns could be present in NLs.  Magnetism was also invoked to explain the complex behavior in the SW Sex star \\object{V795 Her} \\citep[e.g.,][]{Thorstensen86,Casares96} and confirmed with the detection of variable circular polarization in the system \\citep{Rodriguez02}.  Variable circular polarization has also been detected in the SW Sex stars \\object{LS Peg} \\citep{Rodriguez01} and \\object{RX J1643.7+3402} \\citep{Rodriguez09}.  Other evidence of magnetism has also been observed in multiple SW Sex systems, including emission line flaring in systems such as \\object{BT Mon} \\citep{Smith98} and \\object{V533 Her} \\citep{RodMar02} and far UV / X-ray signatures of magnetism in objects like \\object{DW UMa} \\citep{Hoard03} and LS Peg \\citep{Baskill05}.  Clearly magnetism plays an important role in many, if not all, SW Sex stars.  We report on the discovery of a NL, UX UMa-type system, KIC J192410.81$+$445934.9 (Kepler ID 8751494, hereafter KIC J1924$+$4459), in the field-of-view of the Kepler Mission.  This discovery was made as part of an as-yet unsuccessful search for compact pulsators (WDs and subdwarf B stars) in the Kepler Input Catalog.  Time-series photometric observations of compact pulsator candidates were obtained at the TNG, McDonald Observatory, and the BOAO Observatory (Seung Lee-Kim, private communication) without any clear detection of pulsations in any of the observed targets.  Some preliminary TNG results are presented by \\citet{2009CoAst.159...97S}.  The detection of a photometric periodicity at 2.94 h in KIC J1924$+$4459, potentially also detected in the radial velocities of the emission lines, and the detection of rapid non-periodic variability hints that this star may belong to the growing class of SW Sex systems. ",
        "conclusions": "} \\subsubsection{Short-term variability} The light curves display large amplitude variations of 0.12 mag that appear to be periodic.  The light curve shape is not sinusoidal, but shows sharp rises and structured decays.  We then performed a Scargle analysis of the Loiano data that revealed a strong peak at 94.4 $\\mu$Hz and a secondary peak at 188 $\\mu$Hz.  The latter peak is close to twice the frequency of the strong peak.  We therefore fitted a composite sinusoidal function with a fundamental period of $2.9422\\pm 0.0016$ h (frequency $\\nu =8.1560\\pm 0.0051\\, {\\mathrm d}^{-1}$ ) and the first harmonic.  \\begin{figure} \\begin{center} \\includegraphics[angle=270, width=0.45\\textwidth]{f3.eps} \\end{center} \\caption{Discrete Fourier transforms of the combined McDonald and Loiano data sets (top) and of the synthetic light curve generated using the sinusoidal fits including the fundamental and first harmonic to the combined McDonald and Loiano photometry (bottom). Vertical lines indicate the frequencies of the fundamental and first harmonic.  Also indicated in the DFT are two broad peaks of power likely related to the quasi-coherent variations observed in the light curves.  \\label{fig.dft}} \\end{figure}  As the amplitudes of the variability were similar in both the unfiltered Loiano photometry and the BG40 McDonald photometry, and as the two data sets were interleaved, we combined the two data sets to more precisely constrain the periodicities in the data.  Due to the large gap between the discovery photometry and our subsequent follow-up and the short duration of observation, we did not include the discovery TNG data in the combined analysis.  A discrete Fourier transform (DFT) of the combined data (Fig.~\\ref{fig.dft}) shows a strong peak very near the frequency of the Loiano data peak.  Again, we simultaneously fit a fundamental frequency and its first harmonic \\begin{equation} f=A_0\\sin(2\\pi \\frac{t-t_{0,\\rm phot}}{P_{\\rm phot}}) + A_1\\sin(4\\pi \\frac{t-t_1}{P_{\\rm phot}})\\, , \\label{eqn.phot} \\end{equation} where $f$ is the fractional amplitude. For the fundamental frequency, we derive a photometric period $P_{\\rm phot}=2.9358\\pm 0.0017$ h ($\\nu= 8.1749\\pm 0.0070 \\,{\\mathrm d}^{-1}$) with an amplitude $A_0=0.0556\\pm 0.0014$ and $t_{0,\\rm phot} ({\\rm BJD}) = 2454715.52547(74)$.  For the first harmonic, $A_1=0.0133\\pm 0.0014$ and $t_1 ({\\rm BJD}) = 2454715.57688(148)$. Figure \\ref{fig.ts_phot} shows this fit overplotted on the photometric data.  Although the formal errors indicate a significant difference between this period and that derived from the Loiano photometry alone, we believe the true error in the period is larger than the formal errors. In Figure \\ref{fig.ts_phot}, the model light curves for both periods are plotted over the observed photometry.  Even when considering the TNG data, neither period is obviously preferred.  In addition, light curves folded at both periods look qualitatively similar.  We therefore consider a realistic estimate on the systematic error in the period to be 0.0064 h, the difference between the two photometric periods.  Additional observations will be needed to determine the photometric period more precisely. If the photometric period of 2.94 h is the orbital period of the system, then KIC J1924$+$4459 lies within the $2-3$ h CV period gap.  No deep eclipses are observed.  Small dips which could be grazing eclipses are visible just after minimum light in the Loiano data. These have been indicated in Figure \\ref{fig.ts_phot} with vertical tick marks.  The location of the tick mark at T=24.175 h after BJD 2454714.5 was set by visual inspection; all other tick marks are spaced by integer multiples of the photometric period of 2.9358 h. Thse dips are consistent with the photometric period throughout both Loiano data sets, but are not seen in the McDonald data.  A longer data set, such as that being obtained by the Kepler mission, therefore will be required to address this issue.  After removing this long-period variability, the residual light curves from all four nights of data also show significant amplitude ($\\sim 5\\%$) variability on time scales of $\\sim 650s$ and $\\sim 1200s$. Significant power is visible at these frequencies in the DFTs of the combined data (Fig.~\\ref{fig.dft}) and of each individual run. However, this variability is not fully coherent, even within a single observation, indicating that these are not non-radial pulsations on the surface of the WD primary.  \\subsubsection{Long-Term Variability} We have been able to identify KIC J1924$+$4459 in several archival catalogs, including both epochs of the Palomar Observatory Sky Survey (POSS), the Lick Northern Proper Motion study, and 2MASS.  Historic photometric data culled from the the Naval Observatory Merged Astrometric Dataset (NOMAD) are presented in Table \\ref{tab.kic_data}. We note that the USNO-A2 catalog does not resolve KIC J1924$+$4459 and a similarly-bright optical companion star; we therefore do not include those data in this summary.  Visual inspection of both epochs of the POSS reveals significant proper motion of KIC J1924$+$4459 relative to the optical companion, so these two sources are likely unrelated.   Even allowing for uncertainty in the photographic plate magnitudes, residual contamination from the neighbor star, and the uncertain phase of the recent digital photometry, it appears that KIC J1924$+$4459 may have faded by $\\sim 0.5-1$ magnitude between the second-epoch POSS observations of the early 1990s and the 2MASS/KIC observations of the late 1990s and mid-2000s.  Such fading episodes are observed in NLs of the \\objectname{VY Scl} subclass.  Given the uncertainties, however, we cannot conclusively state whether KIC J1924$+$4459 entered a VY Scl-like low state in the mid- to late 1990s."
    },
    "1004/1004.2868_arXiv.txt": {
        "abstract": "Non-linear and non-Gaussian signal inference problems are difficult to tackle. Renormalization techniques permit us to construct good estimators for the posterior signal mean within information field theory (IFT), but the approximations  and assumptions made are not very obvious. Here we introduce the simple concept of minimal Gibbs free energy to IFT, and show that previous renormalization results emerge naturally. They can be understood as being the Gaussian approximation to the full posterior probability, which has maximal cross information  with it. We derive optimized estimators for three applications, to illustrate the usage of the framework: (i) reconstruction of a log-normal signal from Poissonian data with background counts and point spread function, as it is needed for gamma ray astronomy and for cosmography using photometric galaxy redshifts, (ii) inference of a Gaussian signal with unknown spectrum and (iii) inference of a Poissonian log-normal signal with unknown spectrum, the combination of (i) and (ii). Finally we explain how Gaussian knowledge states constructed by the minimal Gibbs free energy principle at different temperatures can be combined into a more accurate surrogate of the non-Gaussian posterior. ",
        "introduction": "\\label{sec:intro} \\subsection{Abstract inference problem} \\label{sec:generic inference problem}  Measurements provide information on the signals we are interested in, encoded in the delivered data. How can this information be best retrieved? Is there a generic and simple principle from which optimal data analysis strategies derive? Can an information energy be constructed which -- if minimized -- provides us with the correct knowledge state given the data and prior information? And if this exists, how can this information ground state be found at least approximatively?  An information energy, to be minimized, would be very useful to have, since many of the existing minimization techniques, analytical and numerical, can then be applied to it. A number of such functions to be extremized to solve inference problems were proposed in the literature, like the likelihood, the posterior, or the entropy. The likelihood is the probability that the data has resulted from some signal. The posterior is the reverse, it is the  probability that  given the data some signal was the origin of it. Extremizing either of them certainly makes sense, but often ignores the presence of slightly less probable, but much more numerous possibilities in the signal phase space. Those have a much larger entropy and are therefore favored by maximum entropy methods. However, maximum entropy alone can not be the inference determining criterion, since it favors states of complete lack of knowledge, irrespective of the data. Thus some counteracting energy is required which provides the right amount of force to the inference solution. Here, we argue that the ideal information energy is provided by the Gibbs free energy, which combines both maximum entropy and maximum a posteriori (MAP) principles.  The Gibbs free energy has to be regarded as a functional over the space of possible probability density functions (PDF) of the signal given the data. The result of the minimization is therefore a PDF itself, and not a single signal estimate. Minimizing the Gibbs free energy maximizes the entropy within the constraints given by the internal energy. The latter is understood as the average of the negative logarithm of the joint probability function of signal and data weighted with the PDF.  The usage of thermodynamical concepts for inference problems is not new, see e.g. \\cite{1957PhRv..106..620J, 1957PhRv..108..171J}. What is new here, is that we develop this for signals which are fields, spatially distributed quantities with an infinite number of degrees of freedom, while using an approximate Gaussian ansatz for the PDF to be inferred. We thereby connect information field theory (IFT) \\cite{1985ApJ...289...10F, 1987ApJ...323L.103B, 1987PhRvL..58..741B, 1988PhRvL..61.1512B, 1996PhRvL..77.4693B, 1999physics..12005L, Lemm2003, 2009PhRvD..80j5005E}, as a statistical field theory dealing with a huge number of microscopic degrees of freedom, to thermodynamics, as a means to generate simplified, but macroscopic descriptions of our knowledge. Thereby we find that former IFT results obtained with complex renormalization schemes in \\cite{2010arXiv1002.2928E,2009PhRvD..80j5005E} can easily be reproduced, and even be extended to more complicated measurement situations.  In the remainder of Sect. \\ref{sec:intro} we briefly introduce to IFT, MAP, and Maximum Entropy. This motivates the minimal Gibbs free energy principle, which we formally derive in Sect. \\ref{sec:thermoinf}, and show its equivalence to maximal cross information. The application of this principle to optimize approximations of the posterior  of concrete inference problems is provided in Sect. \\ref{sec:application}. There, the log-normal Poisson problem (Sect. \\ref{sec:plndata}) and the problem to reconstruct without known signal power spectrum (Sect. \\ref{sec:rec. without spec. knowledge}), as well as their combination (Sect. \\ref{sec:LNPandspecrec}) are addressed. Finally, we show how approximate posteriors obtained at different temperatures can be combined into a better posterior surrogate in Sect. \\ref{sec:infosynth} before we conclude Sect. \\ref{sec:conclusions}.  \\subsection{Information field theory}  Information theory describes knowledge states with probabilities. If $\\Omega$ is the complete set of possibilities, and $A \\subset \\Omega$ is a subset, then $P(A)\\in [0,1]$ describes the plausibility of $A$ being the case, with $P(A)=1$ denoting $A$ being assumed to be sure, $P(A)=0$  denoting $A$ being (assumed to be) impossible, and $0<P(A)<1$ describing uncertainty about the truth of $A$. Obviously $P(\\Omega)=1$ and $P(\\emptyset)=0$. The usual rules of probability theory apply, and generalize the binary logic of Aristotle to different degrees of certainty or uncertainty \\cite{Cox1946,Cox1963}. In case the set of possibilities is a continuum, it makes sense to introduce a PDF $\\p(\\psi)$ over $\\Omega$, so that  $P(A) = \\int_A d\\psi \\, \\p(\\psi)$. Each possible state $\\psi$ can be a multi-component vector, containing all aspects of reality which are in the focus of our inference problem.  We might be interested in a sub-aspect of $\\psi$ which we call our signal  $s=s(\\psi)$. The induced signal PDF is retrieved from a functional or path integral over all the phase spaces of the possibilities of $\\psi$ via $P(s) = \\int \\D\\psi \\, \\p(\\psi)\\, \\delta(s-s(\\psi))$. If $s$ is a field, a function over a physical space $V$, then $s = (s_x)_{x\\in V}$ might be a vector in the Hilbert space  $\\Omega$ of all $L^2$-integrable functions over $V$ and $\\p(s)$ is then a probability density functional. Information theory for $s$ becomes IFT, which is a statistical field theory.  Inference on the signal $s$ from data $d$ is done from the posterior probability $\\p(s|d)$, which can be constructed from the joint PDF of signal and data $\\p(d,s)$ via \\begin{equation}\\label{eq:psd} \\p(s|d) = \\frac{\\p(d,s)}{\\p(d)} = \\left. \\frac{e^{-\\beta\\,H(d,s)}}{Z_\\beta}\\right|_{\\beta = 1}, \\end{equation} where $\\p(d,s)=\\int_\\Omega \\D\\psi\\, \\p(d|\\psi)\\,\\delta(s-s(\\psi))\\, \\p(\\psi) = \\p(d|s)\\, \\p(s)$ and $\\p(d)=\\int \\D s\\, \\p(d,s)$. The second equality in \\eq{eq:psd} is just a renaming of the numerator and denominator of the first fraction, which highlights the connection to statistical mechanics. Thus we define the information Hamiltonian \\begin{equation} H(d,s) = - \\log \\p(d,s), \\end{equation} the partition function including a moment generating source term $J$ \\begin{equation}\\label{eq:Zbeta} Z_\\beta(d,J) = \\int \\D s\\, e^{-\\beta \\, (H(d,s) + J^\\dagger s)}, \\end{equation} and the inverse temperature $\\beta = 1/T$ as usual in statistical mechanics. Here $s^\\dagger$ is the transposed and complex conjugated signal vector $s$, leading to a scalar product $j^\\dagger s = \\int_V dx \\,\\bar{j}_x s_x$. The ad-hoc notion of temperature is as in standard simulated annealing practice. It permits to narrow (for $T<1$) or widen (for $T>1$) the explored phase space region with respect to the one of the joint PDF and therefore is a useful auxiliary parameter. We show  in Sect.~\\ref{sec:tempered-posterior} that the well known thermodynamical equipartition theorem holds: \\begin{equation} \\label{eq:equi-partition} \\expectWRT{H(s,\\,d)}{(s|d)} - H(m,\\,d) \\approx  \\frac{1}{2}\\, N_\\mathrm{dgf}\\,T. \\end{equation} where $N_\\mathrm{dgf}$ is the number of degrees of freedom and $m$ is the mean signal field as defined below in \\eq{eq:mean-field}, which defines the ground state energy. This relation can e.g.~be used to check the correctness of an implementation of a signal phase-space sampling algorithm.   \\subsection{Maximum a posteriori} \\label{sec:MAP}  The first guess for a suitable energy to be minimized to obtain the information state might be the Hamiltonian. Minimizing the Hamiltonian with respect to $s$, while keeping $d$ at their observed values, is equivalent to maximizing the joint probability $\\p(d,s)$ and also the posterior $\\p(s|d)$. The classical field emerging from this is called the MAP signal reconstruction in signal processing. For a detailed discussion of the usage of the MAP principle in IFT see \\cite{Lemm2003}. The MAP field is often a very good approximation of the mean field \\begin{equation}\\label{eq:mean-field} m = \\expectWRT{s}{(s|d)} \\equiv \\int \\D s \\, s\\, \\p(s|d), \\end{equation} which is the optimal estimator of the signal in a statistical $L^2$ error norm sense: \\begin{equation} m = \\mathrm{argmin}_{\\tilde{s}} \\expectWRT{\\int_V dx \\,(s_x-\\tilde{s}_x)^2}{(s|d)}. \\end{equation} The MAP estimator on the other hand can be shown to optimize the statistical $L^0$ norm\\footnote{The $L^0$ norm measures the amount of exact agreement via $\\| f \\|_{0} = \\lim_{\\varepsilon\\rightarrow 0} \\frac{1}{\\varepsilon} \\int\\! dx \\, \\theta(f^2(x)-\\varepsilon^2)$, with $\\theta$ denoting the Heaviside function.}, the result of which may strongly deviate from the mean $m$, if the posterior is highly asymmetric around its maximum. Thus we can regard the MAP estimator as a good reference point, but not as the solution we are seeking in general. It is, however, accurate (in the $L^2$ error norm sense) in case the posterior around its maximum is close to a Gaussian. In this case, the MAP field can easily be augmented with some uncertainty information from the Hessian of the Hamiltonian \\begin{equation} \\mathcal{H} = \\left. \\frac{\\delta^2 H(d,s)}{\\delta s\\, \\delta s^\\dagger}\\right|_{s=m}, \\end{equation} as an approximation of the two point function of the signal uncertainty \\begin{equation} D \\equiv  \\expectWRT{(s-m)\\,(s-m)^\\dagger}{(s|d)}. \\end{equation} Thus we set $D\\approx \\mathcal{H}^{-1} $ in \\begin{equation}\\label{eq:psdg} \\p(s|d)  \\approx \\tilde{\\p}(s|d) = \\G(s-m, D), \\end{equation} where we introduced the Gaussian \\begin{equation}\\label{eq:Gaussian} \\G(\\phi, D) \\equiv \\frac{1}{|2\\pi\\, D|^\\frac{1}{2}} e^{-\\frac{1}{2}\\, \\phi^\\dagger D^{-1} \\phi}. \\end{equation} Unfortunately, the MAP estimator can perform suboptimally in cases where the Gaussian approximation does not hold, see e.g.  \\cite{2010arXiv1002.2928E}.  \\subsection{Maximum Entropy} \\subsubsection{Image entropy} \\label{sec:MaxEnt}  Another quantity often extremized in image reconstruction problems is the so-called \\textit{image entropy} (iE) \\cite{% 1978Natur.272..686G, % 1979MNRAS.187..145S, % 1980MNRAS.191...69B, % 1983CVGIP..23..113B, % 1983iimp.conf..267G, % 1984Natur.311..446S, % 1984Natur.312..381T, % 1984MNRAS.211..111S, % 1986JMOp...33..287B, % gull1989, % gullskilling,  % 1998mebm.conf....1S}. % In classical maximum (image) entropy (MiE here, usually ME) methods the iE is defined for a strictly positive signal via \\begin{equation}\\label{eq:MaxEntEnt} S_\\mathrm{iE}(s) = - \\int_V dx\\, s_x \\,\\log(s_x/\\tilde{s}_x) \\equiv - s^\\dagger \\log(s/\\tilde{s}), \\end{equation} where $\\tilde{s}_x$ is the reference image, which is used to model some prior information. In the second equality we have defined the component-wise application of functions on fields, e.g. $(f(s))_x = f(s_x)$, which we use throughout this work.  We note, that the iE is actually not a physical entropy. Usually its usage is argued for by ad hoc assumptions on the distribution entropy of photon packages in the image plane, rather than being a well motivated description of the signal prior knowledge (or lack thereof). In the following we will reveal the implicitly assumed prior of MiE methods.  The data enter the MiE method in form of an image energy, which is ideally chosen to be the negative log-likelihood, \\begin{equation} E(d|s) = - \\log\\bigl(\\p(d|s)\\bigr), \\end{equation} in order to ensure the best imprint of the data on the reconstruction. The entropy is then maximized with the energy constraint given by minimizing \\begin{equation}\\label{eq:EME} E_\\mathrm{iE}(d,s) = E(d|s) - T\\, S_\\mathrm{iE}(s) \\end{equation} with respect to $s$. Here $T$ is some adjustable temperature-like parameter, permitting us to choose the relative weight of image entropy and image energy. Low temperature means that the MiE map follows the data closely, high temperature that the map space wants to be more uniformly occupied by the signal reconstruction.  The prior information on the signal, $\\p(s)$, does not enter the MiE formalism explicitly. Actually, an implicit prior can be identified, assuming that MiE is actually a MAP principle. In that case the implicitly assumed Hamiltonian is $H_\\mathrm{iE}(d,s) \\cong E_\\mathrm{iE}(d,s)$, where $\\cong$ denotes equality up to an irrelevant, since $s$-independent, additive  constant, and we find \\begin{equation} \\p_\\mathrm{iE}(s) \\propto e^{T\\, S_\\mathrm{iE}(s)} \\propto  \\prod_x  \\,\\left(\\frac{s_x}{\\tilde{s}_x}\\right)^{-T\\, s_x}. \\end{equation} This is not a general prior, but a very specific PDF. Although there is some flexibility to adopt its functional form by choosing $\\tilde{s}$, $T$, and the image space (pixel space, Fourier space, wavelet space, etc.) in which \\eq{eq:MaxEntEnt} holds, $\\p_\\mathrm{iE}(s) $ can not be regarded as being generic. The MiE prior strongly suppresses large values in the MiE map. If a data feature can be either explained by a single map pixel exhibiting a peak value or by several pixels dividing that value among themselves, MiE will usually prefer the second option, leading to blurred reconstructed images.  We conclude, that the term \\textit{maximum entropy} commonly used in image reconstruction is very misleading. A more accurate term would be \\textit{minimal dynamical range}, since the implicitly assumed prior states that pixels carrying larger than average signal $s_x$ are extremely unlikely.  \\subsubsection{Physical entropy}  A physical entropy should measure the distribution spread of a PDF using a phase space integral over its phase space. In fact, the latter is given by the Boltzmann entropy as given by the negative Shannon information, \\begin{equation}\\label{eq:SB} S_\\mathrm{B} = - \\int \\D s \\, \\p(s|d) \\, \\log \\p(s|d), \\end{equation} which is a functional of the signal posterior, $S_\\mathrm{B}= S_\\mathrm{B}[\\p(s|d)] $, and not of the signal map. Inserting \\eq{eq:psd} yields \\begin{equation}\\label{eq:E-F} S_\\mathrm{B} = \\expectWRT{H(d,s)}{(s|d)} +  \\log\\, Z_1(d,0) = U - F, \\end{equation} where we introduced the internal energy $U=  \\expectWRT{H(d,s)}{(s|d)}$ and the Helmholtz free energy $F= F_1(d,0)$ with \\begin{equation}\\label{eq:defF} F_\\beta(d,J) = - \\frac{1}{\\beta} \\,\\log\\, Z_\\beta(d,J). \\end{equation} The fully $J$-dependent Helmholtz free energy provides the field expectation value via \\begin{equation}\\label{eq:m=dFdJ} m= \\expectWRT{s}{(s|d)}= \\left. \\frac{\\partial F_\\beta(d,J)}{\\partial J}\\right|_{\\beta=1, J=0}. \\end{equation} The entropy is also given in terms of the free energy via \\begin{equation} S_\\mathrm{B} = \\left. \\frac{\\partial F_\\beta(d,J)}{\\partial \\beta}\\right|_{\\beta=1, J=0}. \\end{equation} The entropy as well as the free energy are functionals of the posterior and not of the signal. Maximizing or minimizing them does not provide a signal estimator, but singles out a PDF. If we restrict the space of PDFs to the ones we can handle analytically, namely Gaussians as given in \\eq{eq:psdg} and \\eq{eq:Gaussian}, we might obtain a suitable approximation scheme to the full field theoretical inference problem.  Maximizing the entropy alone does not lead to a suitable algorithm, since the maximal entropy state is that of complete lack of knowledge, with a uniform probability for every signal possibility. The internal energy, however, favors knowledge states close to the posterior maximum and would return the MAP solution if extremized alone. Thus the right combination of entropy and internal energy is to be extremized. We would expect a free energy of the form $U- T\\, S_\\mathrm{B}$ to be this function, in analogy to the energy \\eq{eq:EME} used in MiE methods. Thermodynamics teaches us that the Gibbs free energy is the quantity to be minimized (which is identical to the Helmholtz free energy in case $J=0$). Since we are going to calculate this for an approximation of the real PDF, it is necessary to go through the derivation in order to make sure we do this in the right fashion and understand all implications. ",
        "conclusions": "\\label{sec:conclusions}  We have shown that the minimal free Gibbs energy principle in information field theory can be used to obtain approximate knowledge states with maximal cross-information to the exact posterior. The construction of such knowledge states with Gaussian PDF is relatively straightforward: \\begin{enumerate} \\item The joint PDF of signal and data $\\p(d,s)$ has to be specified, e.g. by specifying a data likelihood $\\p(d|s)$ and signal prior $\\p(s)$, and using $\\p(d,s)= \\p(d|s)\\,\\p(s)$. \\item The information Hamiltonian $H$ is the negative logarithm of this, $H(d,s)= -\\log(\\p(d,s))$. \\item A suitably parametrized PDF as a surrogate for the posterior has to be specified, e.g. a Gaussian with its mean and dispersion as degrees of freedom. \\item The internal energy $U$ and entropy $S_\\mathrm{B}$ of this PDF have to be calculated as the PDF-average of the Hamiltonian and the negative log-PDF, respectively. \\item The Gibbs free energy, $G=U-T\\,S_\\mathrm{B}$, has then to be minimized with respect to all degrees of freedom of the surrogate PDF. \\item Any statistical summary like mean and variance can now be extracted from the surrogate PDF. \\end{enumerate} The minimal free energy principle is therefore well suited to tackle statistical inference problems. We have demonstrated this with two different problems and their combination: reconstructing a log-normal field from Poisson data subject to a point spread function and reconstruction without prior knowledge on the signal power spectrum. Earlier results from renormalization calculations in \\cite{2009PhRvD..80j5005E,2010arXiv1002.2928E} have been reproduced. The there used renormalization schemes can therefore be understood as aiming for a surrogate Gaussian PDF which has maximal cross information to the correct posterior. Since these results were previously shown to reconstruct well, also the here proposed method for the more complicated combined case  can be expected to work. However, a detailed implementation and verification of this was left for future work.  Finally we have sketched how Gaussian knowledge states obtained at different thermodynamical temperatures can be combined into a more accurate representation of the posterior, from which moments of the signal uncertainty distributions can easily be extracted.  The minimal Gibbs energy and maximal cross information principle introduced here to IFT should allow the construction of novel reconstruction schemes for statistical inference problems on spatially distributed signals. The thermodynamical language may help to clarify concepts and to simplify applications of IFT, since it permits us to tackle non-linear inverse problems without the need to use diagrammatic perturbation theory and renormalization schemes."
    },
    "1004/1004.2059_arXiv.txt": {
        "abstract": "{Supergiant fast X-ray transients are a subclass of high mass X-ray binaries that host a neutron star accreting mass from the wind of its OB supergiant companion. They are characterized by an extremely pronounced and rapid variability in X-rays, which still lacks an unambiguous interpretation. A number of deep pointed observations with \\XMM\\ have been carried out to study the quiescent emission of these sources and gain insight into the mechanism that causes their X-ray variability.} {We continued this study by using three \\XMM\\ observations of the two supergiant fast X-ray transient prototypes \\xte\\ and \\igrB\\ in quiescence.} {An in-depth timing and spectral analysis of these data have been carried out.} {We found that the quiescent emission of these sources is characterized by both complex timing and spectral variability, with multiple small flares occurring sporadically after periods of lower X-ray emission. Some evidence is found in the \\XMM\\ spectra of a soft component below $\\sim$2~keV, similar to that observed in the two supergiant fast X-ray transients AX\\,J1845.0-0433 and IGR\\,J16207-5129 and in many other high mass X-ray binaries.} {We suggest some possible interpretations of the timing and spectral properties of the quiescent emission of \\xte\\ and \\igrB\\ in the context of the different theoretical models proposed to interpret the behavior of the supergiant fast X-ray transients.} ",
        "introduction": "\\label{sec:intro}  Supergiant fast X-ray transients are a subclass of supergiant X-ray binaries (SGXBs) that host a neutron star (NS) accreting from the wind of its OB supergiant companion \\citep[SFXT,][]{sguera05}. In contrast to the previously known supergiant X-ray binaries, SGXBs, \\citep[i.e., the so-called ``classical SGXBs'', see e.g.,][]{white95}, SFXTs are characterized by a pronounced transient-like activity. These sources undergo few hours long (as opposed to weeks-months long) outbursts with peak X-ray luminosities of $\\gtrsim$10$^{36}$~erg/s, and exhibit a large dynamic range in X-ray luminosity \\citep[$\\gtrsim$10$^4$ between outburst and quiescence;][]{walter07}. The origin of this extreme variability is still debated, and different models have been developed to interpret it. One of these models involves a NS accreting matter from the extremely clumpy wind of its supergiant companion \\citep{zand05, walter07, negueruela08}. According to this interpretation, the sporadic capture and accretion of these clumps by the compact object can produce the observed fast X-ray flares. To reach the required X-ray luminosity swing, very high density clumps are required \\citep[a factor 10$^4$-10$^5$ denser than the homogeneous stellar wind,][]{walter07}. Numerical simulations of supergiant star winds indicate that these high density clumps might be produced by instabilities in the wind \\citep[][and references therein]{runacres02, oskinova07}. \\citet{bozzo08} proposed that the X-ray variability of the SFXT sources might be driven by centrifugal and/or magnetic ``gating'' mechanisms that can halt most of the accretion flow during quiescence, and only occasionally permit direct accretion onto the NS \\citep[see also][]{grebenev07}. The properties of these gating mechanisms depend mainly on the NS magnetic field and spin period. At odds with the extremely clumpy wind model, in the gating scenario a transition from the regime in which the accretion is (mostly) inhibited to that in which virtually all the captured wind material accretes onto the NS requires comparatively small variations in the stellar wind velocity and/or density, and can easily give rise to a very large dynamic range in X-ray luminosity. Yet another model was proposed for IGR\\,J11215-5952, so far the only SFXT to display regular periodic outbursts, which envisages that the outbursts take place when the NS in its orbit crosses a high density equatorial region in the supergiant's wind \\citep{sidoli07}.  Except for their peculiar X-ray variability, SFXT sources share many properties with the previously known SGXBs. Measured orbital periods in SFXTs range from 3 to 30~days \\citep[IGR\\,J16479-4514: 3.32~days, IGR\\,J17544-2619: 4.92~days; IGR\\,J18483-0311: 18.5~days; SAX\\,J1818.6-1703: 30.0~days;][]{jain09,clark09,bird09,zurita09,sguera07}, and are thus similar to those of other SGXBs. The only exception is the SFXT IGR\\,J11215-5952, which has an orbital period of 165~days \\citep{sidoli07}. The relatively high eccentricities inferred for two SFXTs \\citep[$\\sim$0.3-0.7;][]{zurita09,rahoui08} suggest that these systems might be characterized by somewhat more elongated orbits than classical SGXBs. The spin period of the NS hosted in these sources has been measured only in four cases, the periods ranging from 4.7 to 228~s \\citep[IGR\\,J1841.0-0536: 4.7~s; IGR\\,J1843-0311: 21 s; IGR\\,J16465-4507: 228~s; IGR\\,J11215-5952: 186.78~s;][]{bamba01, sguera07, lutovinov05, swank07}, and thus being similar to the spin periods measured in the classical SGXBs. However, owing to the limited duration of most observations, it is not possible to exclude that a number of SFXTs have much longer spin periods \\citep[see e.g.,][]{smith98,bozzo08}.  A number of pointed \\XMM\\ observations of several SFXTs were carried out to study the quiescent emission of these sources and gain insight into the mechanism that drives their peculiar X-ray activity. During the observation of the SFXT IGR\\,J16479-4514, \\XMM\\ captured the source undergoing a very steep luminosity decay from the end of an outburst to a much lower state. The latter was interpreted in terms of an eclipse of the source by the companion star \\citep{bozzo08,bozzo09,jain09}. This observation revealed that in at least one case the X-ray variability of a SFXT was due to the obscuration by the companion star. In the case of IGR\\,J18483-0311, \\XMM\\ helped identify pulsations in the quiescent X-ray flux of this source \\citep{giunta09}, and thus provided strong support for the idea that SFXTs also accrete matter during their quiescent states \\citep[see e.g.,][]{sidoli07,bozzo08}.  To study the low level emission of SFXT sources, we present in this paper quiescent state \\XMM\\ observations of the prototypical SFXTs XTE\\,J1739-302 and IGR\\,J08408-4503. In Sect.~\\ref{sec:sources}, we summarize previous observations of these sources, and in Sect.~\\ref{sec:data} and \\ref{sec:results} we present our data analysis and results. In particular, we find that \\begin{figure*}[t!] \\centering \\includegraphics[scale=0.26,angle=-90]{background_obs1.ps} \\includegraphics[scale=0.26,angle=-90]{background_obs2.ps} \\includegraphics[scale=0.26,angle=-90]{background_igrj08.ps} \\label{fig:backselection} \\caption{Selection of the high background intervals during the two \\XMM\\ observations (Epic-PN camera) of \\xte\\ (OBS1 on the left and OBS2 in the middle) and the observation of \\igrB\\ (on the right). In each case, we reported the source lightcurve not corrected for the selection of the good time intervals and not subtracted for the background in the 0.2-12~keV ({\\it upper panel}) and 10-12~keV energy band ({\\it middle panel}), and the count rate of the total FOV in the 10-12~keV energy band ({\\it lower panel}). In all cases, the time bin is 100~s. Only the observational intervals in which the total FOV count rate in the 10-12 keV energy band was below the threshold indicated with a dashed line were considered for the timing and spectral analysis of the sources.} \\end{figure*} the quiescent spectra of these sources contain a soft component below $\\sim$2~keV. We discuss some possible interpretations of this component in Sect.~\\ref{sec:discussion}. A comparison is also carried out between the quiescent and outburst spectral properties of \\xte\\ and \\igrB.\\ Our conclusions are summarized in Sect.~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  The three \\XMM\\ observations that we have analyzed in the present work, indicate that the quiescent spectra of the two  prototypical SFXT \\xte\\ and \\igrB\\ are characterized by two different spectral components, one dominating the spectrum at the softer energies ($\\lesssim$2~keV) and the other one being more prominent above 2~keV.  The properties of the soft component ($\\lesssim$2~keV) could be reasonably well constrained in the case \\igrB,\\ where the absorption column density was relatively low ($<$10$^{22}$~cm$^{-2}$), whereas in the case of \\xte\\ the detection of this component is less significant. However, the similarity in the timing and spectral behavior observed in the quiescent state of the two sources argues in favor of adopting the same spectral model for both of them. We suggested that the model comprising a CUTOFFPL component at the higher energies plus a MKL component would provide a reasonable description of the data and a plausible physical explanation of the properties observed in the two sources (Sect.~\\ref{sec:discussion}). According to this interpretation, the MKL component would represent the contribution to the total X-ray emission of the shocks in the wind of the supergiant companion. The results of the fits with this model to the data of the three observations inferred a temperature of the MKL component and an emitting region comparable with the values found also in the case of the SFXT AX\\,J1845.0-0433.  Similar soft spectral components have been detected in many other HMXBs and SGXBs. In a few cases, the detection of a number of prominent emission lines in the high resolution X-ray spectra of these sources carried out with the gratings onboard \\chan\\ and the RGS onboard \\XMM\\ \\citep[see e.g.,][]{watanabe06} have convincingly demonstrated that these components are produced by the stellar wind around the NS, and proved to be a powerful diagnostic to probe the structure and composition of the stellar wind in these systems. The statistics of the present \\XMM\\ observations is far too low to permit a similar in-depth study of the stellar wind in the case of \\xte\\ and \\igrB.\\ Furthermore, because of the relatively low luminosity and the high absorption that characterize the emission of these sources in quiescence, observations at higher spectroscopic resolution are probably too challenging for the present generation of X-ray satellites, and the improved sensitivity of the X-ray spectrometers planned for future X-ray missions (e.g. IXO) is probably required to firmly establish the presence of a soft spectral component in the quiescent emission of the SFXT sources and shed light on its nature.  If the harder X-ray emission (2-10~keV) detected from the \\XMM\\ observations of \\xte\\ and \\igrB\\ was really produced by residual accretion as we argued in the previous section, then the accretion process in these sources would take place over more than 4 orders of magnitude of X-ray luminosity\\footnote{That we did not find any evidence for X-ray eclipses in the three \\XMM\\ observations that we analyzed (this is unlike the case of IGR\\,J16479-4514, see Sect.~\\ref{sec:intro}) is also consistent with this interpretation.}. This is similar to the results reported for the SFXT IGR\\,J17544-2619 \\citep{rampy09} and, possibly, for the SFXT SAX\\,J1818.6-1703 \\citep[in the latter case the origin of the lowest quiescent emission remains unclear,][]{bozzo09}. In the case of IGR\\,J17544-2619, \\citet{rampy09} ascribed the high dynamic range in the X-ray luminosity to the accretion of clumps from the wind of the supergiant star with a high density contrast with respect to the surrounding homogeneous wind. However, it was also suggested that a similar variability might result from the transition across different accretion regimes onto the NS \\citep{bozzo08}. We note that a similar scenario can be envisaged for interpreting the variations in the X-ray flux observed during the multiple small flares detected in the present observations. Even though they took place at a much lower luminosity level than the brightest outbursts (a factor of $\\sim$10$^3$-10$^4$), our analysis showed that all these events shared a number of similar timing and spectral properties. In particular, the timescales on which the smaller flares develop is comparable with the decay timescale of the source luminosity during the outbursts (see Sect.~\\ref{sec:intro}), and the spectral photon indices and absorption column densities measured from the \\XMM\\ observations are also in qualitative agreement with those reported previously when the sources were observed at a much higher X-ray luminosity level (see Sect.~\\ref{sec:intro}). It is thus most likely that the transitions between the lower quiescent states and the small flares detected by \\XMM\\ from \\xte\\ and \\igrB\\ might have been triggered by the same mechanism that sometimes gives rise to the brightest outbursts (i.e., the accretion of clumps from the stellar wind and/or the transition between different accretion regimes of the NS, see Sect.~\\ref{sec:intro}).  In contrast to the case for the SFXT IGR\\,J18483-0311, we did not detect any pulsation in the quiescent emission of either \\xte\\ or \\igrB,\\ and provided in Sects.~\\ref{sec:xteresults} and \\ref{sec:igrresults} the corresponding upper limits to the spin periods and pulsed fractions we were able to infer from the present data.  Deep pointed observations of SFXTs in quiescence are still required in order to understand the origin of the peculiar X-ray variability of these sources and distinguish between different models proposed to interpret their behavior."
    },
    "1004/1004.5509_arXiv.txt": {
        "abstract": "Assuming that the universe contains a dark energy fluid with a constant linear equation of state and a constant sound speed, we study the prospects of detecting dark energy perturbations using CMB data from Planck, cross-correlated with galaxy distribution maps from a survey like LSST.  We update previous estimates by carrying a full exploration of the mock data likelihood for key fiducial models. We find that it will only be possible to exclude values of the sound speed very close to zero, while Planck data alone is not powerful enough for achieving any detection, even with lensing extraction.  We also discuss the issue of initial conditions for dark energy perturbations in the radiation and matter epochs, generalizing the usual adiabatic conditions to include the sound speed effect. However, for most purposes, the existence of attractor solutions renders the perturbation evolution nearly independent of these initial conditions. ",
        "introduction": "The biggest problem in cosmology today is the understanding of the accelerated expansion of the universe. Although one could try to attack this question leaving aside the cosmological principle or modifying Einstein's gravity, the most classical approach consists of assuming a perturbed FLRW universe with a negative pressure component. The minimal model (in terms of number of free parameters) compatible with the current data is a cosmological constant, which should be perfectly homogeneous by definition.  Other candidates (which may or may not alleviate the fine--tuning and coincidence problems of the cosmological constant) include, for instance, scalar field models, or effective descriptions in terms of a fluid with free parameters yet to be measured. A canonically kinetic normalized scalar field would fluctuate, but since in that case the sound speed $c_s^2 \\equiv \\delta p / \\delta \\rho$ (computed in the rest frame of the scalar field) is equal to one, local pressure would prevent density contrasts to grow significantly. In an effective fluid description, the sound speed is a free parameter, and dark energy clustering can be more efficient in the limit in which overdensities are not balanced by local pressure perturbations ($c_s \\rightarrow 0$)\\,.  Generally speaking, the study of small perturbations could be used as a tool for discriminating between various models with a negative pressure component (cosmological constant, dark energy fluid, quintessence or k--essence fields, coupled dark energy, etc.) or a modified theory of gravity. One of the major difficulties comes from the fact that the expansion history predicted by a given Lagrangian theory of gravity can be reproduced in General Relativity by a dark fluid having an appropriate (time varying) equation of state, or by a scalar field with an adequate kinetic term and potential. Fortunately, a very precise measurement of clustering properties in our universe could at least help to discard some models in favor of others at the level of perturbations \\cite{Bertschinger:2008zb,Kunz:2006ca}. However, the spatial fluctuations of typical dark energy models are very much suppressed with respect to those of dark matter, and detecting their effect is a real challenge.  The effects of quintessence perturbations (for which $c_s=1$) on the CMB and LSS power spectra were discussed in \\cite{DeDeo:2003te}. For an (uncoupled) dark energy fluid, there have been several studies on the possibility to measure $c_s\\,$, but its value remains unconstrained with present data (see e.g. \\cite{Bean:2003fb,Weller:2003hw,Hannestad:2005ak,Corasaniti:2005pq,Xia:2007km,Mota:2007sz,dePutter:2010vy,Li:2010ac}). In a recent attempt, the authors of~\\cite{dePutter:2010vy} used present Cosmic Microwave Background (CMB), Large Scale Structure (LSS) and supernovae data (including CMB$\\times$LSS cross--correlation), and showed that it is possible to see some weak preference for $c_s\\neq 1\\,$, but only for a certain kind of early dark energy model in which the equation of state is not constant.  We may still hope to discriminate between different values of $c_s$ using combinations of future CMB data with 3--dimensional galaxy clustering data \\cite{Takada:2006xs}, with CMB$\\times$LSS cross--correlation data \\cite{Hu:2004yd}, or with results from a large neutral hydrogen survey such as that of the SKA project \\cite{TorresRodriguez:2008et} (see also \\cite{TorresRodriguez:2007mk}).  In this work, we focus on an effective description which has already been studied by several authors: namely, a dark energy fluid with a linear equation of state $p=w\\rho\\,$, a constant equation of state parameter $w$ close to $-1$ and a constant sound speed defined in the range $0 \\leq c_s^2 \\leq1\\,$.  In Section \\ref{sspeed} we review the concept of sound speed for a cosmological fluid. The purpose of Section \\ref{initcond} is to clarify the non-trivial issue of initial conditions for dark energy perturbations in the radiation era, which are a priori non--adiabatic since $c_s^2 > w\\,$. We present for the first time the initial conditions fulfilled by the dark energy fluid in the synchronous gauge (i.e., in the gauge used by most Boltzmann codes), when all other fluids have adiabatic primordial perturbations. In Section \\ref{attractors} we study analytically the evolution of these perturbations during the matter epoch. We derive approximations for the attractor solutions followed by dark energy perturbations (both in the Newtonian and synchronous gauges). These new results can be used in the future for analytical estimates of the impact of dark energy on structure formation. In Section \\ref{detect}, in order to update the analysis of~\\cite{Hu:2004yd}, we carry a full Monte-Carlo exploration of the likelihood of future mock CMB and LSS data, in order to infer the sensitivity of these data to the dark energy sound speed, and to investigate possible parameter degeneracies. Finally, in Section \\ref{conclude}, we summarize our findings. ",
        "conclusions": "\\label{conclude}  In sections \\ref{initcond} and \\ref{attractors}, we have studied analytically the behavior of dark energy perturbations in the radiation and matter epochs in the synchronous and conformal Newtonian gauges. We have written down the formulas that generalize the usual adiabatic initial conditions for small inhomogeneities when a dark energy fluid with constant equation of state and speed of sound is considered. Our equations~(\\ref{initx}, \\ref{initxtheta}) can be readily implemented in the initial conditions of e.g. {\\sc camb}, which applies to early radiation domination. Instead, equations~(\\ref{suphor}, \\ref{supsound}) -- or their counterpart (\\ref{newt_md_super}, \\ref{supsound}) in the Newtonian gauge -- provide correct initial conditions above the sound horizon for codes simulating only the growth of matter fluctuations during radiation and dark energy domination. The fact that there is an attractor behavior for the evolution of the perturbations guarantees that dark energy fluctuations are driven to solutions dictated by the dark matter ones. This implies that the final state of dark energy today is independent on its possible initial conditions, provided that the attractor solution is reached way before dark energy comes to dominate the background expansion; using the above set of initial conditions guarantees that such a condition is fullfilled.  After clarifying these issues, we have studied in section \\ref{detect} the prospects for detecting dark energy perturbations using CMB data from the Planck satellite, cross-correlated with galaxy distribution maps obtained with an LSST-like instrument.  We have chosen to focus on models of dark energy consisting in a fluid with constant equation of state $w$ and constant sound speed $\\hat{c}_s$. We have performed a full exploration of the likelihood of our mock data using a Monte Carlo approach, in order to cross-check previous Fisher matrix estimates in a more robust way. We find that the confidence interval for $\\hat{c}_s$ inferred from the data will potentially allow us to put a lower bound on $\\hat{c}_s$, and to constrain at least its order of magnitude (although with a slightly smaller significance than in earlier analytic estimates).  Our results build upon previous works on the same topic \\cite{Bean:2003fb,Hu:2004yd,Takada:2006xs,dePutter:2010vy} and complement other researches were different models have been studied (see for instance \\cite{dePutter:2010vy} for early dark energy and \\cite{Martinelli:2010rt} for the case in which there is a coupling to dark matter)."
    },
    "1004/1004.1073.txt": {
        "abstract": "We report optical long-slit spectra and direct imaging (ground-based and with $HST$) of the pre-planetary nebula (pPN) M2-56 obtained at different epochs. The optical nebula is composed by shock-excited material distributed in two pairs of nested lobes with different sizes and surface brightness. The compact, bright inner lobes (ILs) have an angular size of $\\sim$1\\farcs5$\\times$1\\arcsec\\ each and display closed, bow-shaped ends.  The extended, faint outer lobes (OLs), which enclose the inner ones, have an agular size of $\\sim$13\\arcsec$\\times$10\\arcsec. %Our long-slit spectra reveal a %complex expansive kinematics in the lobes. Within the ILs and the OLs the velocity increases with the distance to the center, however, the ILs show expansion velocities larger than the OLs. %The expansion %velocity reaches a maximum value of $\\sim$350\\,\\kms\\ at the tips of %the ILs. Consistent with the large speeds reached by the ILs (of up to $\\sim$350\\,\\kms\\ at the tips), we have measured the expansive proper motions of the knots ($\\Delta \\theta_{\\rm t}$$\\sim$0\\farcs03 yr$^{-1}$) by comparing $HST$ images taken in 1998 and 2002. Moreover, we have discovered remarkable changes with time in the continuum and line emission spectrum of M\\,2-56. In 1998, we detected a burst of \\hal\\ emission from the nebula nucleus that is interpreted as an indication of a dense, fast ($\\sim$350-500\\,\\kms) bipolar wind from the nebula's core (referred to as ``F1-wind''). Such a wind has been recently ejected (after 1989) probably as short-duration mass-loss event. Our data also reveal an optically thick compact cocoon (or shell-like structure) and a \\ion{H}{2} region around the central star that result from further post-AGB mass-loss after the F1-wind. Recent brightening of the scattered stellar continuum as well as an increase of scattered \\hal\\ emission along the lobes is reported, both results pointing to a decrease of the optical depth of the circumstellar material enshrouding the star. %The data presented here have been %analized extensively to derive the spatio-kinematic structure of the %different nebular components of M\\,2-56 and their physical properties %(density, mass, mass-loss rate, and momentum) with unprecedented %detail. The data presented here unveil the complex post-AGB mass-loss history of this object, whose rapid evolution is driven by multiple episodes of mass outflow, not regularly spaced in time, leading to: ($i$) acceleration of the molecular envelope that surrounds the optical nebula (kinematical age $t_{\\rm k}$$\\sim$1400\\,yr -- Castro-Carrizo et al., 2002), ($ii$) the OLs ($t_{\\rm k}$$\\sim$350-400\\,yr), ($iii$) the ILs ($t_{\\rm k}$$\\sim$40\\,yr), ($iv$) the F1-wind ($t_{\\rm k}$$<$10\\,yr), and ($v$) the nuclear cocoon and \\ion{H}{2} region ($t_{\\rm k}$$\\la$2\\,yr?). %The presence of the cocoon and the HII region %suggest current post-AGB wind %activity at the nebula's heart. %%not regularly spaced in %time, some of which are continuously active. The sucessive multiple post-AGB winds in M\\,2-56 are characterized by ejection speeds increasing with time. % from 200 to $\\ga$\\,500\\,\\kms. %and mass-loss rates decreasing from $\\approx$10$^{-4}$ to $\\approx$10$^{-7}$\\,\\my. %\\,\\my\\ in the molecular envelope, %to $\\approx$10$^{-6}$\\,\\my\\ (in the OLs and ILs), and to %$\\approx$10$^{-7}$\\,\\my\\ (in the younger F1-wind). In contrast, the mass-loss rate and linear momentum show a time decreasing trend. %\\mloss$_{\\rm OLs}$$\\sim$\\mloss$_{\\rm ILs}$$\\approx$10$^{-6}$\\,\\my\\  $>$  \\mloss$_{\\rm F1}$$\\approx$10$^{-7}$\\,\\my\\ and %$P_{\\rm mol}$$\\sim$10\\,\\msun\\kms\\ $>$ %$P_{\\rm OLs}$$\\sim$0.17\\,\\msun\\kms\\ $>$ $P_{\\rm %ILs}$$\\sim$0.05\\,\\msun\\kms\\ $>$ $P_{\\rm %F1}$$\\sim$2$\\times$10$^{-3}$\\,\\msun\\kms. %The relatively low values of the momentum found lead us to conclude that %none of the post-AGB ejections probed by our optical data could have transferred its large %linear momentum ($P_{\\rm mol}$$\\sim$10\\,\\msun\\kms) to the molecular envelope. %The ``lobe inside %lobe'' structure suggests multiple episodes of mass outflows in %M\\,2-56. %% tkin=1200yr for the slowest, outermost regions (at 10arcsec) %% tkin=300-500 yr for the OLs %% tkin=40 yr for the ILs %% tkin=8.5 yr %% Our data provide evidence of simultaneous wind ejection in the axial direction and in a perpendicualar plane.  %The flux ratio of the different lines observed is consistent %with shock-excitation. As it is generally believed, shocks in pPNs %likely result from the impact by fast post-AGB ejections on the slow %envelope ejected during the AGB phase. By comparing our flux %measurements with the predictions by shock-excitation models we have %constrained the shock velocity (50-80\\kms) and pre-shock density %(RANGE\\cm3). The electron density, atomic and ionized mass in the %shocked lobes have been estimated... (GIVE VALUES). {\\bf %Interpretation post-AGB evolution; compare with similar objects: CRL618, %OH231, He3-1475, M1-92???}   %We believe that %these differences indicate a change in the on-going stellar wind %activity at the nucleus, most likely, an episode of intense and fast %(probably jet-like) mass ejection by the central post-AGB star. The %duration of the hypothetic post-AGB wind that caused the burst is ~<10 %years. Rapid recombination due to high densities at the core could %explain why the burst became unobservable (in Halpha) in only 2 years %after the burst. The complex nebular morphology of M2-56 (with %multiple collimated and bow-shaped structures with different %orientations) indicate that multiple jet-like ejections and Halpha %bursts like the one discovered by us have happened several times in %the past, leading to the acceleration of the nebula through shock wind %interactions. ",
        "introduction": "\\label{intro} Planetary nebulae (PNe) evolve from Asymptotic Giant Branch (AGB) stars after a brief ($\\approx$10$^3$\\,yr), intermediate stage known as the pre-Planetary Nebula (pPN) or post-AGB phase. During the AGB to PN transition dramatic transformations occur: the spherical, slowly expanding ($V_{\\rm exp}$\\,$\\sim$\\,15\\,km/s) circumstellar envelope (CSE) expelled during the AGB becomes a nebula with, usually, clear departures from sphericity and fast ($\\sim$\\,100\\,km/s) outflows directed along one or more axes. %%As the central star progressively gets hotter the circumstellar material %% gets ionized. The varied PN and pPN morphologies include not only axisymmetric (elliptical and bipolar) shells but also multipolar (multiaxial) structures, multiple co-axial shells, highly collimated jet-like ejections sometimes arranged in a point-symmetric structure, etc \\citep[e.g.][]{sah98,sahs07,ue07,sio08}. Most of the observed PN and pPN morphologies are difficult to explain without invoking the action of precessing, perhaps episodic, collimated outflows \\citep[e.g.][]{mir99,gue01,mir06}. In fact, collimated outflows are considered as the main shaping agent of PNe through their interaction with the spherical AGB CSE \\citep{st98}.  The existence of very young, high velocity bipolar jets in some PPNe and evolved AGB stars is evidenced, for example, by water maser observations \\citep[e.g.][and references therein]{ima02,ima07,sua09}. In some pPNs/PNs, fast, post-AGB winds are also revealed by P-Cygni profiles in the \\hal\\ (and other recombination lines) emission close to the central star \\citep{san01,san08}. %One of the most exciting results recently provided by %long-slit spectra with $HST$/STIS is the first direct observation of %the spatio-kinematic structure of a fast collimated pAGB wind (jet) in %a `pristine' stage, i.e$.$, not strongly altered by the interaction %with the AGB shell \\citep{san01}. In spite of the growing evidence of jet-like ejections in PPNs, their origin, nature (episodic or continuous?), typical mass-loss rates, life-times, etc, are still very poorly known. %It is believed that post-AGB wind interaction takes place quickly, and then the %shocked gas expands freely - this is deduced from the ``Hubble-law'' %of the pPN kinematics. In many objects, interation times are %$\\la$100\\,yr as inferred from the analysis of the CO envelope dynamics %\\citep{buj01}. This means that the life-time of the post-AGB winds is also %$\\la$100\\,yr, after that time the wind is switched off. We suspect %there may be multiple episodes of post-AGB mass-outflows in some %objects (based on their multipolar morphologies), don't know the %frequency of such process or the duration of each post-AGB mass loss %episode. Detailed studies of PPNs are crucial to obtain information about the properties and dynamics of post-AGB winds and their evolution. Optical spectroscopic observations of pPNs and PNs are particularly useful for probing the post-AGB winds and their interaction with the CSE formed in the previous AGB phase.  %** General about post-AGB evolution. Focus on what we know about %post-AGB mass-loss, which is very little: we believe post-AGB winds %are fast and collimated (but continuous?, episodic?... what is the %duration of the post-AGB ejection?? -- %of course we don't know their origin either...    This work focuses on the pPN M\\,2-56 (= IRAS\\,23541+7031= PK 118+08), which so far remains relatively poorly characterized, especially in the visible. A ground-based optical image obtained by Goodrich (1991, hereafter G91) showed a $\\sim$10\\arcsec$\\times$2\\arcsec\\ bipolar nebula with its long axis oriented in the East-West direction (P.A.=90\\degr). The limited quality of this image (PSF$\\sim$1\\farcs6 and low S/N) impeded a correct identification of the different nebular components in M\\,2-56 revealed by data presented in this work. High-angular resolution optical images obtained with the $HST$ have been used as complementary data by \\cite{sio08} and \\cite{cc02} although a detailed study based on these images has not been done so far.  The $HST$ images show that the bright nebulosity referred to as the ``West lobe'' by G91 is resolved into two compact lobes ($\\sim$1\\farcs5$\\times$1\\arcsec) emerging from the center.  The emission line optical spectrum of M\\,2-56, with strong forbidden lines of low-excitation ions, indicates major excitation by shocks %% with velocities in the range $\\sim$40-100\\,\\kms\\ \\citep{goo91,tra93,rie06}. The lack of substantial polarization in the observed line spectrum implies that the latter is locally produced in the shock-excited lobes of M\\,2-56 and that there is not a significant contribution by line emission from the central regions reflected by the nebular dust \\citep[''scattered spectrum'',][]{tra93}.  (In this work we show that this situation has recently changed). The optical spectrum of M\\,2-56 also shows a weak red continuum that is the light from a B-type central star attenuated and/or partially reflected by circumstellar dust \\citep{coh77}.  A detailed study of the dynamics and morphology of the molecular envelope of M\\,2-56 with high-angular resolution was carried out by \\cite{cc02}. Interferometric CO mapping shows an incomplete $\\sim$30\\arcsec$\\times$15\\arcsec\\ hourglass-like envelope surrounding the optical lobes and similarly oriented to the latter in the plane of the sky. The inclination of the symmetry axis of the CO envelope with respect to the plane of the sky is estimated to be $i$=17\\degr$\\pm$2\\degr, with the East side of the nebula tilted towards us. The densest molecular material lies at the center, where the intersection of the two bipolar lobes shapes an equatorial ring/torus expanding at 7-8\\,\\kms. Two compact CO clumps are found at the tips of each lobe expanding axially with velocities of $\\sim$100-200\\,\\kms. The total mass and linear momentum of the molecular envelope of M\\,2-56 is $M_{\\rm mol}$=0.05\\,\\msun\\ and $P_{\\rm mol}$=2\\,\\msun\\kms, adopting a distance to the source of $d$=2.1\\,kpc and a standard value of the CO-to-H$_2$ relative abundance of X(CO)=2$\\times10^{-4}$. The morphology and dynamics of the bipolar nebula is interpreted as the result of the interaction between fast, collimated post-AGB ejections and the slow CSE formed during the AGB phase. The Hubble-type expansion %% (outflow velocity is proportional to distance from the source) %% (\\vexp $\\propto$ r) %% Hubble-like velocity law/field of the molecular outflow suggests that such an interaction took place approximately 1400\\,yr ago in a relatively short time, $\\la$300\\,yr. %Most of the material in the bipolar nebula is therefore shocked-AGB matter.  In this paper we will adopt a distance to M\\,2-56 of $d$=2.1\\,kpc based on the detailed discussion and final estimate by \\cite{cc02}. As estimated by these authors, for such a distance, the luminosity of M\\,2-56 obtained by integrating its spectral energy distribution (SED) from the optical to the mm-wavelength range is $\\sim$5500\\,\\ls. %The relatively low luminosity %and long post-AGB lifetime deduced from the CO data as compared to the  In this work, we report multi-epoch optical long-slit spectra and images of M\\,2-56 that %provide an unprecedented detailed picture of the %post-AGB wind interaction process in this object have allowed us to describe the structure, kinematics, physical conditions, and mass-loss history of this object with unprecedented detail. %Observations and main results %are shown in Sections \\ref{obs} and \\ref{res}, respectively. In %Section \\ref{components}, we describe the spatio-kinematic structure of %the main nebular components. In Section \\ref{extinction}, the %extinction produced by distinct components and its spatial %variation is estimated and discussed. The distribution of the electron %density and the total atomic and ionized mass are computed in Section %\\ref{nemass}. In Section \\ref{dis}, our results are discussed in terms of the formation and %evolution of M\\,2-56, which is currently in rapid transition from the %AGB to the PN stage and presents clear signs of on-going, variable %wind activity at the nucleus. A summary of our findings %is given in Section \\ref{summ}. Our results are discussed in terms of the formation and evolution of M\\,2-56, which is currently in rapid transition from the AGB to the PN stage and presents clear signs of on-going, variable wind activity at the nucleus. A summary of our findings is given in Section \\ref{summ}.   %% fast post-AGB wind from a central star sweeps up a shell in a %% surrounding slow, dense AGB wind. ",
        "conclusions": "\\label{dis}  We interpret the nebular morphology and kinematics of M\\,2-56 in the context of the AGB-to-post-AGB wind interaction scenario, i.e., resulting from the hydrodynamical interaction between fast, collimated (jet-like) post-AGB ejections and the slowly expanding envelope expelled during the previous AGB phase (see \\S\\,\\ref{intro}). %The passage of shock %waves produced by such a wind interation through the nebula does not %only affect its dynamics but it is reponsible for the %heating/ionization and of the gas (\\S\\,\\ref{excitation}). %% ----- a\u00f1adir una frase de este tipo -------------- %Our data have provided unprecedented details on how %the poorly known processes %has taken place in this object The remarkable ``lobe inside a lobe'' appearance of M\\,2-56 together with the discovery of the recent \\hal\\ emission burst, and the revelation of other very small-scale structures at the nebula's core points to multiple episodes of mass outflow during the post-AGB phase. %In the following, we describe the complex mass-loss %history of this object with unprecedented detail.    The earliest post-AGB mass-loss event in M\\,2-56 (that we have knowledge of from existing data) is probably the one that shaped and accelerated the extended hourglass-like molecular outflow enclosing the optical nebula (\\S\\,\\ref{intro}). As discussed by \\cite{cc02}, such an interaction started approximately $\\sim$1400\\,yr ago and took place quickly, in less than 300\\,yr \\citep[see also][]{buj01}. %% Molecular envelope is thought to be made mainly of shocked material %% from the AGB CSE, so-called ``shocked-AGB gas''. Mass and Momentum carried by %% this component, and expansion velocities. %% Vmax~100 km/s, tkin~ 1400 yr, Mass=0.05x6\\,msun, P=2x6\\,\\msun\\kms. - factor 6 due to XCO The smaller kinematical ages of the optical OLs and ILs derived from our long-slit spectra, $t_{\\rm kin}$$\\sim$380 and $\\sim$40\\,yr, respectively, suggest that the latter resulted from two consecutive and more recent post-AGB ejections. The duration of such ejections cannot be asserted based on our data but it may be a small fraction of their kinematical ages as in the case of the fast molecular outflow. We find that the caps of the optical OLs, which are located inmediately behind of two diametrally oposed dense axial clumps traced by CO, show a different kinematics than the main body of the OLs. In particular, the OLs-caps are characterized by expansion velocities close to those of the aforementioned CO condensations. The kinematical age derived for the OLs-caps, $t_{\\rm kin}$$\\sim$1200\\,yr, is thus similar to that of the CO axial clumps, which may suggest that both were shaped and accelerated in the same wind interaction process between the primer post-AGB bipolar flow(s) and the AGB envelope. In this scenario, the massive molecular envelope would be mainly composed of shocked-AGB wind whereas the OLs-caps could be a remnant of the primer post-AGB wind. %excited by the passage of backward-shocks %(i.e.\\ mainly shocked-'post-AGB' material). Alternatively, the material in the OLs-caps could have been ejected at the same time as the rest of the OLs, i.e.\\ $\\sim$380\\,yr ago, and subsequently decelerated (from expansion speeds of $\\sim$400\\kms\\ down to the observed velocities in these regions, \\vexp$\\sim$110\\kms) as a consequence of its interaction with the pre-existing CO clumps. %% Vmax~250km/s; tkin~380 yr; M=Mion+Matom=7e-4\\msun; P < 7e-4x250km/s~ 0.17 msun km/s %% OLs-caps: V~110\\kms. %% ILs, vmax~400\\kms, tkin~40yr, M=1.2e-4\\msun, P<1.2e-4 x 350\\kms~0.04 msun km/s %%{\\bf ** Valentin's comment: hablar mas de los ILs.**} The fast ILs probably trace a more recent shock interaction between an ensuing, fast post-AGB wind with the material at the base of the pre-existing and slower OLs.  Our data also uncover a dense equatorial structure at the nebula center that is expanding in the direction perpendicular to the lobes with radial velocities of up to $\\sim$100\\,\\kms\\ (\\S\\,\\ref{ne} and Fig.\\,\\ref{f_ne}). %This structure may form part of a more extended %hourglass-shaped component sourrounding the optical lobes. The velocity gradient observed, which is approximately linear, indicates a kinematical age of $\\sim$300-400\\,yr (adopting $i$=90\\degr$-$18\\degr=72\\degr). The similar age of the equatorial flow and the OLs is consistent with both components resulting from a single, sudden mass-ejection event that took place simultaneously along the nebula axis and in the perpendicular plane. This scenario has been proposed previously by \\cite{alc07} to explain the radial acceleration along the equator and bipolar flows of the molecular envelope of the pPN M\\,1-92. %%happen at the same time the OLs were shaped given In M\\,2-56, the compact equatorial structure unveiled by the [\\ion{S}{2}] doublet may represent the inner (warmer) regions of the dense, molecular ring/torus probed by CO emission (\\S\\,\\ref{intro}). The lower expansion velocity of the latter (\\vexp=7-8\\,\\kms) may indicate substantial deceleration of the equatorial flow after interaction with the dense, slow AGB envelope. %Although our spatial resolution %does not enable detecting a central cavity, The density is found to increase outwards along the optical equatorial flow, in support of the presence of compressed (swept up) material at the boundary with the dense molecular torus. %_____________________________________________________________________________ % %Finally, the small blue-shift (with respect to \\vsys) of the %center of expansion of the optical equatorial flow may %indicate the presence of a binary system at the core of M\\,2-56. %% with an orbital period comparable to the life-time of the sudden ejection % that led to the OLs+disk; a small fraction of tkin~300-400,yr. %_____________________________________________________________________________   %We found that the equatorial structure dividing the two lobes is a %thin flat disk, which expands radially with a velocity proportional to %the distance to the central stellar system. The kinetic age of this %equatorial flow is very similar to that measured in the two lobes, %suggesting that the whole structure was formed as a result of a single %event some 1200 yr ago,  The discovery of the \\hal\\ emission burst from the nebular core in 1998 (feature F1) brings to light a latter stellar wind. As discussed in \\S\\,\\ref{f1wind}, the F1-emitting wind is probably bipolar.  The linear size and mean expansion velocity of this fast, compact flow allow us to derive an upper limit to its kinematical age of only $<$14\\,yr. The upper limit arises because of possible deceleration of the F1-wind by interation with the ILs. The absence of feature F1 in the \\hal\\ spectrum obtained in 1989 by G91, in fact, suggests that the F1 wind could have been ejected only $<$10\\,yr before it was first identified in our run\\,\\#1. In this case, the initial expansion speed of the F1-wind (before substantial deceleration took place) should have been \\vexp$\\sim$500\\kms\\ to explain the axial extent of this component. This value of \\vexp\\ is indeed in agreement with that inferred from the broad wings and the P-cygni absorption component of feature F1, which indicate that a fraction of the F1-flow is actually moving at this relatively large speeds (\\S\\,\\ref{f1wind}). %% strings of knots... % multiple episodes of mass-loss are believed to % be required for stars proceeding from a LBV to a WR phase and % have been observed around ~ Car and P Cygni (e.g., Walborn % 1976; Smith & Hartigan 2006).   The absence of feature F1 in 2000 is probably due to fast recombination of the gas in the dense ($>$2.6$\\times$10$^4$\\,\\cm3) F1-wind (\\S\\,\\ref{ne}). In order for the F1-wind to be completely or considerably neutral already in 2000, most of it should have been far away (detached) from the central star at that time, otherwise one would expect to observe emission from the dense photoionized material at the base of the wind close to the stellar surface as long as the wind is on-going. We conclude, therefore, that the rapid recombination observed, in less than $\\sim$3\\,yr, is most consistent with the F1-wind being a short-time duration event, e.g.\\ a pulsed-jet or bullet-type ejection, leading to detached or clumpy structures, rather than being blown in a continuous way. %The latter would lead to a %Stromgrem sphere uniformly filled with wind material that would remain %photoionized close to the stellar surface and, therefore, continuously %emitting in \\hal\\ as long as the wind is on-going. {\\bf Se entiende %este parrafo anterior??} %In the latter case, %the wind would be expected to fill uniformly the stromgrem sphere and beyond. %it would be impossible for the whole wind to become unobservable in %2.9 yrs, because the wind at the stellar surface would not be able to %recombine!! it has to be a clump that has moved away from the star, %going beyond the very small region that is photoionized by the star! %% calcular el radio de stromgrem !! We believe that the clumpy, jet-like structures observed in the direct $HST$ images of M\\,2-56 at the base of the eIL could represent the footprints of such a fast F1-wind. This is based on the similar size and location of both structures in the images and the spectra. If this is correct, the emission that we see from the prominent three-finger like structure in the F606W $HST$ image must be mainly stellar light $scattered$ by the F1-wind, after recombination, since that image was obtained in 2002 once the \\hal\\ emission feature F1 had dissapeared.  %% Mass loss rate of the F1-wind %% Method A) %% Mloss=PI rj^2 rho Vexp; rho>3e4cm-3, vexp=500\\kms, rj<=0.05arcsec ---> Mloss=3e-7\\,\\my!! %% Method B) %% M(F1)/tkin, M(f1)<4.1e-6 Msun, tkin<=10yr --> Mj~ 4e-7\\,\\my... %% aunque ambas estimaciones son muy inciertas, coinciden EXTRAORDINARIAMENTE BIEN!!   %This is not the wind that shaped the inner and outer lobes, since it is %too young compared with the kinematical age of the lobes. So, there %must have been at least two other ejections that led to the extended %and inner lobes (the sharp bright contrast, remarkable differning %shape, size, etc, suggest that both the inner and outer resulted in %two different wind-interaction events).  We have estimated a lower limit to the mass-loss rate that led to the main nebular components identified in our optical imaging and spectroscopic data, namely, the OLs, the ILs and the F1-wind, from the ratio between the total mass contained in each component and their kinematical ages (Table\\,\\ref{t_masses}). The lower limit arises because the duration of each of the multiple post-AGB mass-loss episodes undergone by M\\,2-56 could be shorter than their kinematical ages. % derived from %our long-slit spectra. %, as shown by %\\cite{buj01}, the dynamics of most pPNs %indicates that the strong mass acceleration happens %quickly, in a time (referred to as $acceleration time$, $t_{\\rm acc}$) shorter than %total post-AGB lifetime, that is, $t_{\\rm acc}$ $<$\\,$t_{\\rm kin}$. %% que determina la duracion de la interacion? por un lado, %Free motion after a sudden acceleration is in fact the easiest explanation %to the linear dependence of the velocity on the distance to the %central star observed in M\\,2-56 (\\S\\,\\ref{}) %and in most post-AGB flows (referred to as the ''Hubble-law'' kinematics of PNs). We obtain very similar values of the mass-loss rate for the OLs and ILs, namely, \\mloss$_{\\rm OLs}$$\\sim$2$\\times$10$^{-6}$ and \\mloss$_{\\rm ILs}$$\\sim$3$\\times$10$^{-6}$\\,\\my. For the F1-wind, we compute a value of \\mloss$_{\\rm F1}$$\\sim$4$\\times$10$^{-7}$\\,\\my, however, we note the larger uncertainties in this case mainly due to our vague estimate of the mass of the F1-wind. We have attempted a different estimate of the mass-loss rate of the F1-wind as given by %\\begin{equation} $\\mloss=\\pi r_j^2 \\vexp \\rho$, %% \\\\ %\\end{equation} %\\noindent where the geometry of the F1-wind is approximated by a cylindrical structure with radius $r_j$, density $\\rho$, and expanding at \\vexp. Adopting $r_j$$\\sim$0\\farcs05, $\\rho$$>$3$\\times$10$^4$\\,\\cm3 and \\vexp$\\sim$500\\,\\kms, we derive \\mloss$_{\\rm F1}$$\\sim$4$\\times$10$^{-7}$\\,\\my. The radius adopted for the F1-wind is that of the bright, jet/finger-like features in the F606W $HST$ images that may be footprints of the F1-wind. The agreement between our two estimates of \\mloss$_{\\rm F1}$ suggests that the order of magnitude obtained may be correct.  In general, very little is known about mass-loss rates of post-AGB winds, their continuous or episodic nature, and about the structure of the resulting flows after the AGB-to-'post-AGB' wind interaction.  Our data indicate that the various post-AGB mass-loss episodes experienced by M\\,2-56 did not happen at regular time intervals, in particular, the time span between two consecutive post-AGB ejections has shortened with time. Also, according to the different expansion velocities of the distinct nebular components of M\\,2-56, including the molecular envelope, the succesive post-AGB winds seem to be characterized by ejection speeds exponentially increasing with time. %, with rates that %had gone from $\\la$0.13\\,\\kms\\,yr$^{-1}$, between the primer post-AGB %%wind that shaped the CO bipolar flow, up to 2.5\\,\\kms\\,yr$^{-1}$, %between the post-AGB ejection that led to the optical ILs and the %F1-wind. Recently, the expansion velocity of the stellar wind could %have raised up to $\\sim$1200\\,\\kms\\ in only $\\la$11\\,yr, if the broad %\\hal\\ profile from the compact \\ion{H}{2} region of M\\,2-56 obseved in %2009 is primarily due to fast expansive motions %(\\S\\,\\ref{res_hii}). It is also worth mentioning that, unlike most pPNs and PNs with multiple optical lobes, the OLs and the ILs of M\\,2-56 are oriented almost identically (with their main symmetry axis inclined with respect to the line-of-sight by $i$=18\\degr; \\S\\,\\ref{inc}). Since the molecular envelope has also a similar orientation, we conclude that the different post-AGB bipolar ejections in this object (except maybe for the F1-wind) have taken place along the same (or nearly the same) direction/axis over the last $\\sim$1400\\,yr. We note that, in contrast, signatures of significant jet directional changes are common amongst pPNs and PNs \\cite[e.g.][and references therein]{mir99,sahs07}. Another example of directionally stable jets is found in the PN He\\,2-90 \\citep{gue01}. %___________________________________________________________________________ %If the three-finger like structure at the %base of the bright eIL is the footprint of the F1-wind, then the %different orientation in the plane of the sky of the individual %finger/jet-like components may suggest that the directional stability %of the stellar wind could have been recently altered, although %an open-angle bipolar F1 wind oriented along the same axis i=18\\degr\\ %but with a non-uniform (clumpy) density structure is also possible. %___________________________________________________________________________ % o que el viento F1- es multidireccional! %% ____________________ INTERESANTE?? ______________________________________ %% podria haber habido precesion de manera continuada con un periodo <40yr (para que no se aprecien sus efectos en los ILs NI EN LOS ols) %% y mayor o similar a ~10yr, que es la edad cinematica del F1-wind. %% si Pprecesion = Porbital --> ?? %%  %We show the previous results in Fig.\\,\\ref{f_vt}. Note that only an upper limit %to expansion velocity of the most recent post-AGB wind can be obtained from %the FWZI of the \\hal\\ profile that arises from the compact %\\ion{H}{2} region visible in our 2009 spectroscopic data. %: $i$) from \\vexp(CO)$\\sim$110\\,\\kms\\ to %\\vexp(OLs)$\\sim$250\\,\\kms\\ (in $\\sim$1400-380=1120\\,yr); $ii$) from %\\vexp(OLs)$\\sim$250\\,\\kms to %\\vexp(ILs)$\\sim$400\\,\\kms\\ (in $\\sim$380-40=340\\,yr); $iii$) and from %\\vexp(ILs)$\\sim$400\\,\\kms\\ to \\vexp(F1)$\\sim$500\\,\\kms (in %$\\sim$40-<10\\,yr). %{\\bf Que os parece esta figurilla? la dejamos?} %____ EL PROBLEMA ES QUE LO QUE MEDIMOS NO SIEMPRE ES LA VELOCIDAD DEL VIENTO PAGB %____ SINO LA DEL VIENTO CHOCADO. ES UN LIMITE INFERIOR...  We have obtained a rough estimate of the scalar linear momentum carried by the bipolar OLs, the ILs, and the F1-wind from the product of their masses (Table\\,\\ref{t_masses}) by their maximum expansion velocities, which have been taken to be $\\sim$250, 350, and 500\\,\\kms, respectively. In this calculation we assume that the nebular components are elongated structures with a velocity field that is mainly axial (\\S\\,\\ref{model}). The values found ($P_{\\rm OLs}$$\\sim$0.17\\,\\msun\\kms\\ $>$ $P_{\\rm ILs}$$\\sim$0.05\\,\\msun\\kms\\ $>$ $P_{\\rm F1}$$\\sim$2$\\times$10$^{-3}$\\,\\msun\\kms) are orders of magnitude smaller than the momentum carried by the molecular outflow, $P_{\\rm mol}$$\\approx$10\\,\\msun\\kms\\ (in deriving this value, the linear momentum computed by Castro-Carrizo et al., 2002, $P_{\\rm mol}$$\\sim$2\\,\\msun\\kms, has been corrected for our lower estimate of the CO-to-H$_2$ molecular abundance -- see \\S\\,\\ref{extinction}). %%, i.e.\\ could be $P_{\\rm mol}$$\\approx$12\\,\\msun\\kms. According to this, we must conclude that none of the post-AGB ejections probed by our optical data could have transferred its large linear momentum to the CO outflow and, therefore, a primer more energetic post-AGB ejection is necessary to explain the dynamics of the molecular envelope. This result is in good accordance with the time decreasing trend of the linear momentum of the optical nebular components inferred by us ($P_{\\rm OLs}$$>$$P_{\\rm ILs}$$>$$P_{\\rm F1}$), which independently suggest (by extrapolation) that earlier post-AGB flows could have carried a larger linear momentum.  The presence of a compact dusty structure and a nuclear \\ion{H}{2} region unveiled by our data (\\S\\,\\ref{cocoon} and \\ref{res_hii}) show evidence of further post-AGB mass-loss after the F1-wind in M\\,2-56. The decrease of the optical depth along the line-of-sight and along the nebular symmetry axis observed (\\S\\,\\ref{cocoon_09}) is consistent with a detached/shell-like (cocoon?) structure moving away from the central star. Assuming that the cocoon expands at constant velocity, $v_{\\rm c}$, and that the density varies with the radial distance to the star as $\\rho \\propto r^{-2}$ (i.e.\\ the mass-loss rate is constant), it can be easily demonstrated that for a thick and a thin shell-like geometry:  \\begin{eqnarray} v_{\\rm c}=r_{98}\\times \\left(\\sqrt{\\frac{N_{98}}{N_{09}}}-1\\right) / \\Delta t \\hspace{1cm} {\\rm (thick\\ shell)} \\\\ v_{\\rm c}=r_{98}\\times \\left(\\frac{N_{98}}{N_{09}}-1\\right) / \\Delta t  \\hspace{1cm} {\\rm (thin\\ shell)}, \\end{eqnarray}  \\noindent where $r_{98}$ is the inner and mean radius of the thick and thin cocoon, respectively, in 1998, $N_{98}$ ($N_{09}$) is the column density in 1998 (2009), and $\\Delta t$ is the time scale of the variation of the optical depth. In 1998 and 2009, the total extinction towards the star was A$^{\\star}_{V}$=12\\,mag and 8.8\\,mag, respectively (\\S\\,\\ref{cocoon_09}). Subtracting the component of the extinction produced by the CO envelope and the ISM at the center, $\\sim$3\\,mag, we estimate that the $V$-band extinction produced by the cocoon was A$^{\\rm cocoon}_{V}$=12-3=9\\,mag in 1998 and %%, which  is equivalent to a column density of $N_{98}$=2$\\times$10$^{22}$\\,\\cm2. A$^{\\rm cocoon}_{V}$=8.8-3=5.8\\,mag in 2009. %In 2009, the extiction produced by the cocoon is %8.8-3=5.8\\,mag, i.e. Nh2=1.33e22. Since the extinction is proportional to the column density, we find $N_{98}/N_{09}$=9/5.8=1.55.  Adopting $r_{98}$$\\sim$0\\farcs1 ($\\sim$3.14$\\times$10$^{15}$\\,cm, \\S\\,\\ref{cocoon}) and $\\Delta t$$\\sim$10\\,yr, the variation of the optical depth observed can be explained for moderate expansion velocities of the cocoon of $v_{\\rm c}$$\\sim$30 and 60\\,\\kms\\ for the thin and thick shell approximation, respectively. %% ============================================================================= %If the expansion %velocity of the cocoon is larger, which is not unreasonable given the %increasing trend of the expansion velocity of the multiple post-AGB winds of M\\,2-56, %e.g.\\ if $v_{\\rm c}$=500\\,\\kms\\ as for the F1-wind, then the observed decrease of the optical depth could %have taken place in only $\\sim$0.5\\,yr (or less if $r_{98}$$<$0\\farcs1). %% =============================================================================  In the previous expanding cocoon scenario, the detached geometry (needed to explain the decrease of the optical depth) implies that the mass-loss process that led to this structure had already ended in 2009. The compact size of the cocoon, which is well inside the ILs and plausibly closer to the star than the F1-wind, suggest it is at least as young as the latter (i.e., $t_{\\rm kin}$$<$10\\,yr). Adopting an expansion velocity of $v_{\\rm c}$=500\\,\\kms\\ as for the F1-wind, which is not unreasonable given the increasing trend of the expansion velocity of the multiple post-AGB winds of M\\,2-56, we obtain $t_{\\rm kin}$$\\la$\\,2\\,yr.  %% ============================================================================= %Using the cocoon column density derived from its optical %depth applying the standard conversion factor between these two magnitudes (see %\\S\\,\\ref{extinction}), $N_{\\rm 98}$$\\sim$2$\\times$10$^{22}$\\,\\cm2 and %$N_{\\rm 09}$$\\sim$1.3$\\times$10$^{22}$\\,\\cm2, we deduce a lower limit %to the average number density in the cocoon of $\\rho$\\,$\\approx$10$^6$\\,\\cm3\\ taking %into account a generous upper limit to the angular size/thickness of this structure of 0\\farcs1. %Assuming that the mass-loss rate that led to the cocoon is given by %\\mloss$_{\\rm c}$=4$\\pi$$r^2$$\\rho$$v_{\\rm c}$, and adopting a thick shell-like geometry %(in which case $N$$\\sim$$\\rho$$\\times$$r$, with $r$ being the inner radius of the %shell), then:%  %\\begin{eqnarray*} %\\mloss_{\\rm c} \\sim N_{\\rm 98} 4 \\pi r_{\\rm 98} v_{\\rm c} \\sim %2\\times10^{-5} \\left[\\frac{r_{\\rm 98}}{\\arcsec}\\right]\\left[\\frac{v_{\\rm c}}{\\kms}\\right]\\,\\my. \\\\ %\\end{eqnarray*} %% ============================================================================= %% the mass of the cocoon is M=4$\\pi$N r_{\\rm in} r_{\\rm out}; if rout and rin < 0.1arcsec, then %% the mass is approx. <1.6e-3\\msun!!  % since the INNER radius of the cocoon is unknown, and could be as small as the stellar radius, 3e-6\\arcsec, and % the \\vexp\\ is also undetermined, 30-$>$500\\kms??, we cannot compare % the mass-loss rate of the cocoon with that of the earlier F1-wind.  %Adopting a density of $\\rho$\\,$\\approx$10$^6$\\,\\cm3 at a radius of $\\sim$0\\farcs1, we obtain %a mass-loss rate for the cocoon of \\mloss$\\sim$4$\\times$10$^{-4}$$\\times$$v_{\\rm c}$[\\kms]\\,\\my.  %Finally, we note that although expansion of the cocoon is a simple, plausible explanation %of the decrease of the optical depth towards the star, the cocoon could have been pierced %by  The presence of compact cocoons enshrouding the central source has been suggested in several pPNs. For example, the so-called {\\sl searchlight beam} features that are observed emerging from the nebular core in an increasing number of pPNs and PNs (of which CRL\\,2688 remains as the best known example) have been hypothesized to result from starlight escaping through holes/cavities in an inner dust cocoon \\citep{sah98,san07}. Since these cocoons have not been {\\sl directly} detected, their existence remains speculative so far. %How these cocoons form in the post-AGB phase is a %mystery, especiallyi f they happen to have a spherical (or nearly spherical) %symmetry in which case it would be particularly intrincate %understanding the recovery of isotropic mass-loss $\\sim$1400\\,yr after %the breaking up of the spherical symmetry (evidenced by the bipolar %molecular envelope). %A reasonable alternative to the renaissance of a spherical wind in the %post-AGB phase could be the simultaneous ejection of multiple bipolar %flows along the nebula axis AND along the equator, which would also %explain the large obscuration of the central source observed along %both directions.  Distinct, high-velocity outflows oriented in almost %perpendicular directions are also detected in some pPNs, e.g.\\ %CRL\\,2688 \\citep{cox00}. %% podria el cocoon ser alargado, e.g. como un ellipsoide (en lugar de una shell esferica)?? %% si pero tambien chocaria con la pinched waist of the optical nebula. %Maybe the simplest explanation is obscuration by 1) bipolar flow + 2) expanding disk!! %% As mentioned in Section \\ref{cocoon}, We cannot rule out that the large obscuration of the central source of M\\,2-56 both along the lobes and the equatorial plane could also be produced by {\\sl two} distinct nebular components, namely, a dusty bipolar flow and an equatorial disk/torus, rather than by a unique cocoon-like structure. Direct observations of the obscured, innermost regions of M\\,2-56 using, for example, high-angular resolution techniques in the infrarred and mm-wavelength range are crucial to characterize the circumstellar geometry at the nebula's heart. %and to attempt understanding the complex %mass-loss history of the latest post-AGB ejections in this object. %%%% ______________________ NO DEBERIA EL COCOON EMITIR EN EL NIR?? _______________ %%%% PODRIAMOS CALCULAR LA TEMP. DEL POVO ESPERADA PARA ESE COCOON, Y SI SERIA O NO DETECTABLE EN EL NIR. %%%% ESTARIA BIEN USAR AO O INTERFEROMETRIA PARA VER SI EXISTE ESE COCOON!! %%%% mm-wave interf. could also help. %%%% LA FOTOMETRIA NIR AND MIR **INDICA** UNA COMPONENTE DE POLVO CALIENTE Y POCO MASIVA DE ACUERDO A LOS CALCULOS DE VALENTIN!!  Finally, our spectra taken in 2009 have revealed a compact \\ion{H}{2} region around the central star of M\\,2-56 that may represent the very latest (current?) mass-loss episode of this object.  It is possible that such a compact \\ion{H}{2} region was already formed in 1998 but its light, like the stellar radiation, was heavily blocked from our view by the dusty cocoon. Only due to the recent decrease of the optical depth of the latter, the emission from the \\ion{H}{2} region has been able to reach the lobes where it is scattered by the nebular dust, enabling its indirect detection. Alternatively, it is also possible that the \\ion{H}{2} region has formed recently (after 2000). In the following, we discuss this possibility in some detail. The recent formation of the \\ion{H}{2} region would imply an equally recent and rapid evolution (in less than $\\sim$10\\,yr) of the central star towards higher effective temperatures, needed to explain the raise of photonionizing UV radiation and detectable \\hal\\ emission from the ionized regions. Although a large increase of \\teff\\ is not deduced from our multi-epoch spectroscopic data, which do not show significant changes in the stellar spectrum\\footnote{except for its brightening as a result of the decrease of the cocoon optical depth -- see \\S\\,\\ref{extinction}.}, a low/moderate increase of the stellar temperature from a B\\,0.5 to a B\\,0 spectral type ($\\Delta$\\teff$\\sim$3000\\,K) is simultaneously consistent with a stable B-type classification and with a large increase of the \\hal\\ equivalent width from $W_\\lambda$=50 to 300\\AA\\ \\citep[for details, see][]{san08}. %In fact, the featureless stellar spectrum of M\\,2-56 %has remained consistent with a reddened B-type central star at least %over the last $\\sim$30\\,yr \\cite[see %e.g.][]{coh77,goo91,tra93}. %%% ---> Our measurement of $W_\\lambda$=145\\AA\\ in 2009 <---- The \\teff\\ evolutionary rate value deduced in this case, $\\Delta$\\teff/$\\Delta$t$\\ga$3000/10$\\sim$300\\,K\\,yr$^{-1}$, is larger than that expected for the central star of M\\,2-56 based on theoretical evolutionary models for post-AGB objects taking into account the low luminosity and, thus, initial mass of this object \\citep[e.g.][]{blo95,hoof97}. However, %According to these models, the low luminosity of M\\,2-56, $\\sim$5300\\,\\ls, %implies a low initial (final) mass of 3\\msun\\ (0.6\\,\\msun), for which %$\\Delta$\\teff/$\\Delta$t$\\la$30\\,K\\,yr$^{-1}$ is predicted \\citep[see %Fig.\\,6 in][]{blo95}. these models assume constant values for the post-AGB mass-loss rate of 10$^{-7}$-10$^{-8}$\\my, single-star evolution scenarios, and other simplifying hypothesis that may not be appropriate for M\\,2-56. In particular, we do not rule out that the lattest sudden mass ejections leading to the F1-wind and the compact cocoon, could have sped up/boost %% skyrocketed the evolution of the central star towards slightly higher \\teff\\ and the recent emergence of a nuclear \\ion{H}{2} region in just a few years.  In addition to the central \\ion{H}{2} region, there are probably shocks currently happening at the innermost circumstellar regions of this fascinating object as evidenced by the recent emergence of \\ion{Fe}{2} lines from the nucleus (\\S\\,\\ref{res_09}). These \\ion{Fe}{2} lines are known to be good tracers of astrophysical shocks, e.g., in supernovae, Herbig Haro objects, PNe, etc \\citep[e.g.][]{wel99,rei00}. %: the [\\ion{Fe}{2}] lines are very sensitive to the high %densities and temperatures in the shocked gas; moreover, fast %shocks are able to extract substantial amounts of Fe from dust grains The presence of shocks at the nebula's core suggests current, fast stellar ejections (in agreement with the broad nuclear \\hal\\ profile) interacting hydrodynamically with the material in the close environment of the star ejected in earlier mass-loss episodes with smaller velocities. %(Note that it would only take $\\sim$10\\,yr for a $\\sim$1000\\,\\kms\\ wind to reach %a distance from the star of $\\sim$1\\arcsec, for d=2100\\,pc) %% Backward shocks are also expected in the current post-AGB wind itself. High-angular resolution optical imaging and spectroscopy are needed to properly study the spatio-kinematic structure of the innermost nebular regions and to attempt understanding the complex mass-loss history of the latest post-AGB ejections in M\\,2-56.  %{\\bf Ve pensando en esto... tienes las bibliografia de OH231 y CRL618 %(articulos parecidos), puedes ir tambien pensando en como se compara %este objeto con estos otros dos... en terminos de la estrella central, %masa nebular deducida a partir de Halpha y OI, masa molecular (de CO), %velocidades de expansion, morfologia nebular, edad dinamica...} % %% el objeto mas similar es M2-9 y tambien algunas PNs como Hb 12, Hen2-104 y MyCN 18 (que se %supone que son binarias simbioticas. En todos estos casos las nebulosas presentan %``nested bipolar lobes'' or ``nested hourglass structures''. Los lobulos estan orientados %de la misma manera."
    },
    "1004/1004.4096_arXiv.txt": {
        "abstract": "We have created synthetic neutral hydrogen ({\\sciHI}) Galactic Plane Survey data cubes covering $90^\\circ \\le \\ell \\le 180^\\circ$, using a model spiral galaxy from SPH simulations and the radiative transfer code {\\Tor}.  The density, temperature and other physical parameters are fed from the SPH simulation into {\\Tor}, where the {\\sciHI} emissivity and opacity are calculated before the 21-cm line emission profile is determined. Our main focus is the observation of Outer Galaxy `Perseus Arm' {\\sciHI}, with a view to tracing atomic gas as it encounters shock motions as it enters a spiral arm interface, an early step in the formation of molecular clouds.  The observation of {\\sciHI} self-absorption features at these shock sites (in both real observations and our synthetic data) allows us to investigate further the connection between cold atomic gas and the onset of molecular cloud formation. ",
        "introduction": "The formation of molecular clouds out of cold, dense atomic gas is a key process in the cyclic evolution of the interstellar medium. Above column densities of $N_H \\approx 10^{20}$ cm$^{-2}$, atomic hydrogen gas transitions to molecular H$_2$, a critical step towards the formation of stellar systems \\citep{dix98}.  For the molecular fraction of the gas to remain stable depends largely on the ability of the molecular gas to self-shield against dissociating photons from the local radiation field.  Moreover, the importance of dust grains in H$_2$ formation requires their presence in these transition interfaces.  The 21-cm neutral hydrogen line is the main tracer of interstellar structure. The ubiquitous nature of atomic hydrogen in the Galaxy was demonstrated in early surveys \\citep{Oort58}, and more recent surveys imaging the {\\sciHI} distribution at moderately high resolution (parsec scales at distances of kpc) have steadily improved upon that view.  Almost all of the {\\sciHI} in the Galaxy's disc has been sampled by the International Galactic Plane Survey (IGPS; \\citet{vgps06,sgps05,tayetal03}).  Beyond the Galactic Plane, all-sky mapping efforts of the Southern (GASS; \\citet{gass09}) and Northern skies (EBHIS; \\citet{ebhis09}) have improved vastly on the previous generation's all-sky {\\sciHI} maps.  Galactic 21 cm surveys at Arecibo with its ALFA instrument \\citep{galfapaper,peek09,vanL09,putmanetal09} are revealing {\\sciHI} structures in our Galactic neighbourhood with unprecedented sensitivity.  Observational evidence of the interface between atomic and molecular gas relies on a good H$_2$ tracer. The $J=1-0$ rotational transition of CO is the most common surrogate of molecular hydrogen, yet there are conditions under which its efficacy as a reliable H$_2$ tracer is undermined \\citep{dou07,kla05,dob07,pel09}. For example, in photon-dominated regions, the CO molecule is largely dissociated, while H$_2$ self-shielding allows for a significant amount of molecular gas to remain.  In very cold, dense molecular cores, CO will freeze onto dust grains along with other larger molecules, and so the H$_2$ content in such cores would be underestimated by traditional CO observations.  In addition to the above scenarios, which assume a molecular cloud has already formed, a newly forming molecular cloud must have an epoch wherein a significant quantity of H$_2$ has formed before an appreciable (detectable) amount of CO has been reached \\citep{Ber04}, and so the evolution of protostellar objects may be already underway in molecular clouds before the onset of critical CO densities. OH is another molecule which has been used to probe the H-H$_2$ transition region, with some success \\citep{Mag90,Wan93}. Neutral and ionised carbon atoms also have potential to be good tracers of diffuse, pre-CO molecular gas \\citep{oka05}.  In order for appreciable H$_2$ to form from a largely atomic hydrogen reservoir, the gas must be cooled to temperatures on the order of 20 K, among other favourable conditions (sufficient number densities, H$_2$ formation sites). Models of gas flowing into the Perseus spiral arm \\citep{Roberts1972} predict compression of the {\\sciHI} gas at the interface of a ``shock ridge.\"  Though his model was isothermal, such conditions are favourable for strong cooling of {\\sciHI}, leading to self-absorption of 21-cm emission, and possibly enhanced H$_2$ production.  Hence {\\sciHI} self-absorption (HISA) observations may help to trace the earliest stages of molecular cloud formation, in regions where CO is not a workable tracer for molecular hydrogen.  Widespread HISA has been observed in the Outer Galaxy \\citep{Gibsonetal2005}, and its association with CO is anything but straightforward.  HISA and CO emission are often found to be coincident spatially and at the same radial velocities, but the correspondence is complicated.  There are regions where only one is present, which may indicate differences in the evolutionary state of these clouds.   Studies of HISA and OH emission have found the OH emission to be more localised to the HISA \\citep{Kav06,Li03}, compared to CO.  As an extra complication, the detection of HISA requires background {\\sciHI} emission at the same velocity as the absorbing gas, which may not be available along many lines of sight.  In addition to analytical models, computational efforts have investigated this scenario in detail. \\citet{DobbsGloverClarkKlessen2008} performed galaxy-scale SPH simulations that trace the dynamics of interstellar clouds as they encounter shocks in spiral arms.  To enhance resolution in the spiral arms, their model traces only those SPH particles in a galactocentric annulus between $R_g = 5$ kpc out to $R_g = 10$ kpc. In addition to strong cooling of the atomic gas, molecular hydrogen forms in the vicinity of the shocks.  Smaller clouds aggregrate into larger ones within the spiral shock, where the molecular fraction is $\\ge 0.5$.  As the clumps leave the spiral arm they are sheared into smaller spurs and lose most of their molecular gas.  The link between these models and observations can be strengthened by deriving predicted observables directly from the simulations.  We do not observe densities, temperatures and pressures of clouds directly---rather, we rely on spectral line and continuum processes to deliver photons containing the valuable information, from which we infer the state of the object under study.  As radiative transfer is a guiding principle for this propagation of information, we endeavour to simulate the creation of observable photons via radiative transfer, within the modelled physical representation of the ISM in our Galaxy.  \\citet{gib06} demonstrated the production of HISA by spiral shocks with a two-dimensional, purely atomic (cold and warm {\\sciHI}) radiative transfer model.  In this paper we present results from generating synthetic {\\sciHI} spectral line datacubes from a multiphase numerical galaxy calculation.  The aforementioned SPH calculation by \\citet{DobbsGloverClarkKlessen2008} is used as a basis to create a synthetic Galactic Plane survey, fashioned after the {\\sciHI} data products created for the IGPS.  Simulating such surveys requires the ``observer\" to be placed at an appropriate position inside the simulated Galaxy.  Such challenges are accomplished using {\\Tor}, an Adaptive Mesh Refinement (AMR) code which is capable of performing radiative transfer calculations to produce synthetic spectra \\citep{Harries2000}. Moreover, the added flexibility of operating on a numerical model enables us to use the same code to observe the Galaxy from multiple positions, as well as to emulate extragalactic {\\sciHI} observations \\citep{Acr10}.  This paper is organised in the following way.  Section 2 reviews the method used to create our synthetic Galactic Plane Survey, and the data cubes produced are described in section 3.  In section 4 we discuss the use of HISA profiles as a diagnostic of early molecular cloud formation.  In section 5 we compare briefly our synthetic data to observational {\\sciHI} profiles of the Milky Way, setting the basis for a discussion of how future refinements to our investigation, such as the incorporation of feedback mechanisms, will bring closer agreement between the synthetic and observed {\\sciHI} distributions.  \\section[]{Numerical Work: SPH Simulations and TORUS}  The method for creating synthetic {\\sciHI} cubes has been outlined previously in \\citet{Acr10}, hereafter called Paper I.  We briefly summarise the main points here.  An SPH spiral galaxy calculation is fed into {\\Tor}, wherein the particle representation of the gas is transformed into a gridded representation, aided by AMR to provide the best resolution in regions of highest density.  Next, radiative transfer calculations and ray tracing algorithms produce synthetic spectra for lines of sight chosen from an arbitrary observing position.  The intensities determined by the radiative transfer and ray tracing are converted to a brightness temperature $T_b$, using the Rayleigh-Jeans approximation. In Paper I we demonstrated this method by simulating 21-cm {\\sciHI} observations of nearby spiral galaxies, assigning the SPH galaxies orientations and systemic velocities matching M31 and M33.  Moreover, the observer can be placed within the galaxy, allowing for the simulation of {\\it Galactic} {\\sciHI} observations to be performed.  In such cases, some adaptations of {\\Tor} algorithms must be made. The AMR grid is generated at a much finer resolution than that used in Paper I, in order to provide better resolution for structure close to the observer.  The maximum {\\sciHI} mass per cell for the AMR grids used in this paper is $2.5\\times 10^{34}$ g ($12.6$ $M_\\odot$), which is a factor of 100 smaller than that used in paper I.  In order to cope with the increased number of cells in the AMR grid, the domain of the grid is restricted to exclude regions of the galaxy which are not relevant to the cube being calculated.  A different AMR grid is used for generating each data cube but they typically comprise $6\\times10^6$ unique voxels.   Also, while the same ray-tracing method is used for both the external and internal views, its execution must be handled differently if we are to consider multiple lines of sight.  For the external case the rays are essentially parallel, whereas in the internal case there is an obvious angular dependence to the observations. Additionally, by reversing the direction of ray trace, relative to that employed in Paper I (i.e.~back to front rather than front to back), we are able to calculate the contribution of each cell of the AMR grid to either the emission or absorption in a given data cube pixel.  This feature is used to decompose spectral features into separate emission and absorption components, as part of our investigation of HISA in spiral arms.  These differences are critical when we extend the creation of single spectra to making fully three-dimensional datacubes of simulated {\\sciHI} emission.  The cubes are discussed in the next section.    \\begin{figure} \\includegraphics[angle=-90,width=84mm]{fig1.eps} \\caption{A synthetic {\\sciHI} spectrum measured from within the model galaxy, toward a position near Galactic coordinates $\\ell = 135.5^\\circ, b = -0.3^\\circ$.} \\label{figsp1} \\end{figure}  Figure \\ref{figsp1} shows an {\\sciHI} spectrum measured from a position within the SPH galaxy, corresponding roughly to the Earth's position in our Galaxy, toward a direction corresponding to Galactic longitude $\\ell \\approx 135.5^\\circ$ and latitude $b \\approx -0.3^\\circ$.  We see ``local\" spiral arm gas in a radial velocity range $-20 < v_r < 0$ km s$^{-1}$, and an outer arm beyond $v_r \\approx -38$ km s$^{-1}$.  The outer arm shows some evidence of {\\sciHI} absorption, with a narrow dip in $T_b$ appearing near $v_r = -41$ km s$^{-1}$. This HISA signature is common in synthetic spectra all across our region of interest. Since HISA is a main focus of our investigation into molecular cloud formation, the {\\Tor} code was modified so that in addition to total brightness temperature cubes, we create separate datacubes where only positive or negative intensity contributions ($dI$) for a given grid cell are mapped.  For this paper, the simulated galaxy has reached a time of 250 Myr. At this age the distribution of gas in different phases, and the fraction of molecular gas, have reached a roughly steady state. While tracing the evolution of specific ISM structures using many timesteps from an SPH simulation is the subject of future investigations, the present work concentrates on deriving {\\sciHI} properties from a single epoch, as is the case with most Galactic radio observations. As mentioned previously, the models are restricted to Galactocentric radii of 5 to 10 kpc.  For this reason it is reasonable to aim to reproduce observations of the Outer Galaxy, toward its Anticenter ($\\ell = 180^\\circ$). In this paper, we survey the entire second quadrant of our SPH galaxy, covering longitudes between $\\ell = 90^\\circ$ and $\\ell = 180^\\circ$. ",
        "conclusions": "\\subsection{Comparison with Observational Data}  The synthetic datacubes we have produced correspond to the Second Galactic Quadrant, for which the best comparison set is the Canadian Galactic Plane Survey, part of the International Galactic Plane Survey \\citep{tayetal03}.  We compare below the general characteristics and features observed in both the synthetic and real data.  By the very fact that our SPH galaxy is {\\it not} the Milky Way, no strong agreement is expected {\\it a priori}.  The main purpose of the synthetic models is to investigate whether the geometry, kinematics, chemistry and thermodynamics of a spiral shock can produce HISA, as has been observed in the Perseus Arm of our Galaxy. In comparing the two datasets, we can discern what characteristics this relatively simple galaxy model has in common with Galactic {\\sciHI} observations.   \\begin{figure} \\centering \\includegraphics[angle=0,scale=0.7]{fig11.eps} \\caption{CGPS Channel map of Perseus Arm gas centered on $\\ell = 135^\\circ$, $b = 0^\\circ$.} \\label{figCgGHI} \\end{figure}     In figure \\ref{figCgGHI} we show the Galactic analogue of the synthetic {\\sciHI} channel map for region GHI that was shown in figure \\ref{figgalxy}. The main difference observed across the Second Quadrant is the apparent confinement of {\\sciHI} gas to the plane in our synthetic data cubes, compared to the widespread and more diffuse distribution of atomic gas seen in the CGPS data. This is not surprising, given the lack of feedback mechanisms in the simulations that would serve to stir up the interstellar medium more vigourously---we will discuss this point further below.  For now, we consider the consequences of this scale-height mismatch.  For one, with a greater column density of {\\sciHI} toward the overdense midplane, the brightness temperatures for the synthetic line profiles are larger than what is observed in the Milky Way.  Within the CGPS data brightness temperatures above 150 K are not common, yet in the synthetic cubes we encounter many regions with $T_b > 200$ K.    \\begin{figure*} \\centering \\includegraphics[angle=-90,scale=0.64]{fig12.eps} \\caption{Longitude-velocity diagram of Milky Way {\\sciHI} emission near $\\ell=135^\\circ, b=0^\\circ$ as observed with CGPS data.} \\label{figCgGHILV} \\end{figure*}  Unlike the scale height differences observed in the channel maps, the spiral structure in the SPH galaxy and the Milky Way show good agreement, in terms of the Perseus Arm, the main feature in the Second Quadrant of the Outer Galaxy.  Figure \\ref{figCgGHILV} shows the Milky Way's mid-Second Quadrant midplane longitude-velocity distribution, which portrays a similar separation, on the order of $\\Delta v_r \\sim 40$ km s$^{-1}$, between local and Perseus Arm gas to that observed in figure \\ref{figgalxz}.  The presence of gas beyond the Perseus Arm is a main difference exhibited by the real data, as well as distinct {\\sciHI} features in the interarm gas. Additionally, the overall velocity width of the Galaxy's spiral arms is in general greater than that of the simulated galaxy. In the SPH galaxy, particles were given circular velocities with a velocity dispersion assigned to be $\\sigma = 6 $ km s$^{-1}$, while values of $\\sigma \\sim 8$ km s$^{-1}$ or greater typically describe the {\\sciHI} of our Outer Galaxy \\citep{Saha09}. Many factors contribute to such a dispersion value, including feedback mechanisms and a warped disk; \\citet{DobbsPrice2008} also found that magnetic fields broaden spiral arms in SPH galaxy simulations.        \\subsection{The Role of Feedback}  The interstellar medium of our Galaxy is constantly disrupted, by both external and internal agents.  The inflow of gas from satellite galaxies and high-velocity clouds provides fuel for future generations of molecular clouds and stellar systems.  Numerous feedback processes from stellar evolution also add energy and momentum to the ISM from within: outflows, supernova shocks and winds from individual stars or stellar associations.  With these agents absent from the simulated galaxy, our synthetic observations correspond to a rather quiescent version of a Milky Way-type galaxy.  The consequences are abundantly clear in our synthetic datacubes.  {\\sciHI} clouds collect toward the midplane and cool, with nothing to destroy them, resulting in larger than expected column and number densities.  Thus the lower scale height seen in our simulated galaxy is well explained.  If our synthetic data had shown agreement with the {\\sciHI} scale height suggested by the CGPS data, then the importance of feedback in contributing to the overall structure of the ISM would be severely diminished.  Much of the {\\sciHI} distribution in our Galaxy can be described with words such as filaments, shells, and bubbles. Interstellar bubbles are created by outward forces originating from a central nidus.  Shells and filaments form as a result of bubbles encountering each other, or other ISM structures, and their evolution is also aided by radiative and magnetic field effects.  Without the injection of supernova-like events or winds from OB associations, the simulated galaxy cannot be expected to recreate the filamentary structure observed with {\\sciHI} observations of the Milky Way. Hence feedback is a possible solution to both the overconfinement of {\\sciHI} in the midplane, as well as the non-filamentary nature of the {\\sciHI} channel maps. Feedback should also increase the velocity dispersion of the simulated galaxy's spiral arms, since the mixing of the {\\sciHI} will certainly result in gas with a wider range of radial velocities.   \\subsection{Future Work}  The simulation used in this paper provides a basis with which to compare future simulations which include more physical processes.  As highlighted above, feedback appears to be instrumental in reproducing the correct structure (holes, filaments) and vertical distribution of the ISM. Moreover, magnetic fields are an important source of pressure, supporting the gas in the vertical direction \\citep{cox05}.  In addition, magnetic fields are likely to reduce thermal instability in the gas \\citep{fie65,hen00,sto09}, which could reduce the amount of HISA seen in our simulated observations. On the other hand, self gravity will promote the formation of dense {\\sciHI} or molecular structures in the midplane, particularly at higher surface densities.   Our synthetic observations of the {\\sciHI} distribution and the resulting HISA properties in these future simulations will demonstrate the cumulative effect of the added physics.   As the molecular fraction of the SPH galaxy evolves, formation of H$_2$ and CO take place using prescribed rate equations \\citep{DobbsGloverClarkKlessen2008}.  The simplicity of the two-level hyperfine 21-cm transition makes the production of the synthetic {\\sciHI} spectra rather straightforward, compared to the detailed balance calculations required for other ISM tracers.  Still, in principle the {\\Tor} code can be adapted to produce synthetic observations of any species present in the original simulation.  \\citet{run10} have achieved synthetic molecular line observations for several species, including CO, for far-field view simulated observations of molecular clouds with {\\Tor}.  We will integrate their work to produce a synthetic CO galactic plane survey to complement our {\\sciHI} data. Besides spectral lines, we will also aim to produce realistic maps of continuum emission from dust.  The use of numerical simulations also offers the unique opportunity to trace individual features in the time domain.  The simulation we used here from \\citet{DobbsGloverClarkKlessen2008} and others like it show the evolution of atomic hydrogen clouds as they encounter shocks in spiral arms, form molecular clouds, and eventually disperse as they exit the spiral arm on the other side.  Radio observations can only piece together this scenario by observing multiple clouds at different stages of this evolution.  In a future study we will recreate our synthetic Galactic Plane survey for several timesteps of a numerical solution, and track the evolution of the {\\sciHI} profiles of gas clouds which pass through spiral shocks.  In all of the above studies, we can also extend our synthetic Galactic Plane survey to other parts of the simulated Galaxy.  From our chosen observer position, only lines of sight within $|\\ell| < 46^\\circ$ are affected by the absence in the model domain of the inner 5 kpc of the galaxy.  Thus we can extend our investigation to include {\\sciHI} observations from other components of the IGPS for comparison."
    },
    "1004/1004.1957_arXiv.txt": {
        "abstract": "Given the importance of shear flows for astrophysical gas dynamics, we study the evolution of the Kelvin-Helmholtz instability (KHI) analytically and numerically. We derive the dispersion relation for the two-dimensional KHI including viscous dissipation. The resulting expression for the growth rate is then used to estimate the intrinsic viscosity of four numerical schemes depending on code-specific as well as on physical parameters. Our set of numerical schemes includes the Tree-SPH code VINE, an alternative SPH formulation developed by Price (2008), and the finite-volume grid codes FLASH and PLUTO. In the first part, we explicitly demonstrate the effect of dissipation-inhibiting mechanisms such as the Balsara viscosity on the evolution of the KHI. With VINE, increasing density contrasts lead to a continuously increasing suppression of the KHI (with complete suppression from a contrast of 6:1 or higher). The alternative SPH formulation including an artificial thermal conductivity reproduces the analytically expected growth rates up to a density contrast of 10:1. The second part addresses the shear flow evolution with FLASH and PLUTO. Both codes result in a consistent non-viscous evolution (in the equal as well as in the different density case) in agreement with the analytical prediction. The viscous evolution studied with FLASH shows minor deviations from the analytical prediction. ",
        "introduction": "Shear flows are an integral part of many astrophysical processes, from jets, the formation of cold streams, to outflows of protostars \\citep{Dekel_Birnboim_2009, Agertz_2009, Diemand_2008, Walch_Burkert_Naab_Gritschneder_2009}, and cold gas clouds falling through the diffuse hot gas in dark matter halos \\citep{Bland-Hawthorn_2007,Burkert_2008}. Jets and outflows of young stars can entrain ambient material, leading to mixing and possibly the generation of turbulence in e.g. molecular clouds \\citep{Burkert_2006, Banerjee_Klessen_Fendt_2007, Gritschneder_a, Carroll_Frank_Blackman_Cunningham_Quillen_2009}, while the dynamical interaction of cold gas clouds with the background galactic halo medium can lead to gas stripping of e.g. dwarf spheroidals (e.g. Grcevich \\& Putman 2009), and the disruption of high-velocity clouds \\citep{Quilis_Moore_2001,Heitsch_Putman_2009}. The KHI is believed to significantly influence the gas dynamics in all of these different scenarios. \\\\ Moreover, viscous flows play a crucial role in e.g. gas accretion onto galactic discs (\\citealp{Das_2008,Park_2009,Heinzeller_Duschl_2009}), as well as in dissipative processes like the turbulent cascade. Typically, the gas viscosity seems to be rather low in the interstellar medium, with typical flow Reynolds numbers of $10^5$. \\\\ To describe these complex processes in detail, numerical schemes are applied to follow the hydrodynamical evolution. Numerous simulations use smoothed-particle hydrodynamics (SPH), (\\citealp{Gingold_Monaghan_1977,Lucy_1977,Review_Benz_1990,Review_Monaghan_1992,Review_Monaghan_2005}), because its Lagrangian approach allows us to follow the evolution to high densities and small spatial scales. In combination with N-body codes, it is a perfect tool for cosmological simulations (e.g. \\citealp{Hernquist_Katz_1989, Couchmann_1995, Springel_2002, Marri_2003, Serna_2003}) and galaxy formation and evolution \\citep{Katz_1992, Evrard_1994, Navarro_1995, Steinmetz_1999, Thacker_2000, Steinmetz_2002, Naab_Jesseit_Burkert_2006}. SPH describes the physical properties of a fluid by smoothing over a representative set of particles. However, this can lead to several problems. It can fail to correctly model sharp density gradients such as contact discontinuities, or velocity gradients occurring in e.g. shear flows (see \\citealp{Agertz_2007}), thus suppressing shear instabilities such as the KHI. \\\\ An interesting problem to test the limitations of SPH as well as grid codes is the passage of a cold dense gas cloud moving through a hot and less dense ambient medium \\citep{Murray_1993, Vietri_Ferrara_Miniati_1997, Agertz_2007}. Such a configuration would be typical for gas clouds raining onto galactic protodisks, for High-Velocity Clouds in the Milky Way and for cold HI clouds in the Galactic disk. \\citet{Murray_1993} demonstrated using a grid code that in the absence of thermal instabilities and/or gravity clouds moving through a diffuse gas should be disrupted by hydrodynamical shear flow instabilities within the time they need to travel through their own mass. Agertz et al. (2007) have shown that the KHI, and therefore the disintegration of such clouds is suppressed in SPH simulations. This problem, in particular the suppression of the KHI, has been subject to recent discussion in the literature. Several solutions have been proposed, e.g. \\cite{Price_2007} discusses a mechanism, which involves a special diffusion term (see also \\citealp{Wadsley_2008}). \\\\ Furthermore, \\cite{Read_Hayfield_Agertz_2009} identify two effects occurring in the SPH formalism, each one separately contributing to the instability suppression. The first problem is related to the leading order error in the momentum equation, which should decrease with increasing neighbor number. However, numerical instabilities prevent its decline. By introducing appropriate kernels, \\cite{Read_Hayfield_Agertz_2009} showed that this problem can be cured. The second problem arises due to the entropy conservation. Entropy conservation inhibits particle mixing and leads to a pressure discontinuity. This can be avoided by using a temperature weighted density following \\cite{Ritchie_2001}. Recently, \\cite{Abel_2010} has shown to solve this problem by evaluating the pressure force with respect to the local pressure. In contrast to standard SPH schemes this applies forces to particles only if there is a net force acting upon them. \\\\ Another characteristic of SPH is the implementation of an artificial viscosity (AV) term \\citep{Monaghan_Gingold_1983}, which is necessary in order to treat shock phenomena and to prevent particle interpenetration. AV can produce an artificial viscous dissipation in a flow corresponding to a decrease of the Reynolds-number and therefore a suppression of the KHI \\citep{Review_Monaghan_2005}. To confine this effect, a reduction of viscous dissipation was proposed by \\citet{Balsara_1995} and improved by \\citet{Colagrossi_2004}. \\citet{Thacker_2_2000} studied different AV-implementations in SPH and pointed out that the actual choice of the AV-implementation is the primary factor in determining code performance. An extension of SPH which includes physical fluid viscosities was discussed by e.g. \\citet{Takeda_1994}, \\citet{Flebbe_1994}, \\citet{er03}, \\citet{Sijacki_Springel_2006} and \\citet{Lanzafame_2006}.\\\\ An alternative to conventional numerical schemes may arise from a new class of hybrid schemes based on unstructured grids and combining the strengths of SPH and grid codes \\citep{Springel_2009}. Some of the problems listed above might be solved with this type of implementation.\\\\\\\\ In this paper we determine how accurate shear flows and the corresponding incompressible KHI are described in common numerical schemes. Therefore, in \\S~\\ref{analytics}, we analytically derive the growth rates of the KHI including viscosity. In \\S~\\ref{numdesc} we briefly describe the numerical schemes and outline how the simulations have been analyzed. We then discuss our results.  At first, we concentrate on the standard SPH implementation, which does not contain a physical viscosity but instead uses AV. However, as mentioned above, AV does influence the evolution of the flow. In \\S~\\ref{SPH_simulations}, we discuss the ability of two numerical SPH-schemes to model the incompressible KHI, namely the Tree-SPH method VINE \\citep{Wetzstein_Nelson_Naab_Burkert_2008, Nelson_Wetzstein_Naab_2008}, and the SPH code of \\cite{Price_2007}.\\\\ By comparing to the derived analytical solution, we asses the effects of AV in VINE and estimate the intrinsic physical viscosity caused by AV (\\ref{SPH_simulations_equal_density}). We then study the development of the KHI for different density contrasts (\\ref{SPH_diff_dens_layers}). We show that the instability is suppressed for density contrasts equal to or larger than $6:1$. We also discuss the remedy suggested by \\cite{Price_2007}, hereafter P08.\\\\ In \\S\\ref{Grid_simulations} we then study the same problem with two grid codes, FLASH \\citep{Fryxell_2000} and PLUTO \\citep{Mignone_2007}.  As the intrinsic artificial viscosity is negligible in these schemes, we study the non-viscous as well as the viscous evolution of the KHI for equal (\\ref{FLASH_equal}) as well as non-equal (\\ref{FLASH_diff}) density layers. We summarize our findings in \\S\\ref{conclusions}. \\section[analytics]{KHI -- analytical description} \\label{analytics} \\begin{figure*} \\begin{center} \\includegraphics[width=0.5\\textwidth]{Shear_Flow_new_black_white} \\caption[Shear flow initial configuration] {Sketch of the initial conditions considered: Two fluid layers with constant densities $\\rho_1$ and $\\rho_2$ flowing in opposite directions with uniform velocities $U_1$ and $U_2$.} \\label{Shear_Flow} \\end{center} \\end{figure*} To derive the growth rate of the KHI in two dimensions including viscosity, we follow the analysis of \\citet{Chandrasekhar} (for a related analysis see also \\citealp{Funada_Joseph_2001} and \\citealp{Kaiser_2005}). The fluid system is assumed to be viscous and incompressible. We use Cartesian coordinates in $x$ and $y$, with two fluids at densities $\\rho_1$, $\\rho_2$, and velocities $U_1$, $U_2$ moving anti-parallel along the $x$-axis, separated by an interface layer at $y=y_s$ (see Fig.~\\ref{Shear_Flow}). We neglect the effect of self-gravity. The hydrodynamical equations for such a system are then given by the continuity equation \\begin{equation} \\label{continuity} \\frac{\\partial}{\\partial t}\\rho + {\\bf{\\nabla}} \\cdot \\left( \\rho {\\bf{v}} \\right)=0, \\end{equation} and the momentum equation \\begin{equation} \\label{navstokes} \\rho \\cdot \\left[ \\frac{\\partial {\\bf{v}}}{\\partial t}+ \\left( {\\bf{v}} \\cdot {\\bf{\\nabla}} \\right) {\\bf{v}} \\right]= -{\\bf{\\nabla}} p +\\rho \\nu \\bigtriangleup {\\bf{v}}, \\end{equation} with the flow density $\\rho$, velocity ${\\bf{v}}$, the thermal pressure $p$ and the kinematic viscosity $\\nu$. \\subsection{Linear Perturbations} \\label{perturb} We linearize equations~\\ref{continuity}, and \\ref{navstokes} with the perturbations \\begin{eqnarray} {\\bf{v}} &\\rightarrow& {\\bf{v}} + \\delta {\\bf{v}} = (U(y)+u;w)\\\\ \\rho &\\rightarrow& \\rho + \\delta \\rho,\\\\ p &\\rightarrow& p +\\delta p. \\end{eqnarray} $u$, $w$ express the perturbation in the velocity, $\\delta \\rho$ and $\\delta p$ in the density and pressure, respectively. This yields the system of linearized equations as \\begin{eqnarray} \\rho\\partial_{t} u + \\rho U \\partial_{x} u + \\rho w \\partial_{y} U &=&-\\partial_{x} \\delta p+ \\nu(\\rho+\\delta \\rho) \\partial^2_{y}U+\\nonumber \\\\ & &\\rho\\nu (\\partial^{2}_{x}+\\partial^{2}_{y}) u,\\label{vx} \\\\ \\rho\\partial_{t} w+ \\rho U \\partial_{x} w &=&-\\partial_{y}\\delta p+ \\rho\\nu(\\partial^{2}_{x}+ \\partial^{2}_{y}) w,\\label{vz}\\\\ \\partial_{t} \\delta \\rho + U \\partial_{x}\\delta \\rho &=& - w \\partial_{y} \\rho, \\label{lincont}\\\\ \\partial_{t}\\delta y_{s} + U_{s} \\partial_{x} \\delta y_{s} &=& -w (y_{s}),\\label{intz}\\\\ \\delta_{x} u + \\delta_{y} w &=& 0.\\label{divv} \\end{eqnarray} Eqs.~\\ref{vx} and \\ref{vz} represent the linearized Navier-Stokes equations, where the density may change discontinuously at the interface positions denoted by $y_s$. Eq.~\\ref{lincont} is the linearized continuity equation. In Eq.~\\ref{intz} the subscript $s$ distinguishes the value of the quantity at $y=y_s$ (the interface layer). The last equation, Eq.~\\ref{divv} expresses the incompressibility of the fluid. With perturbations of the form \\begin{equation} \\label{perturb_quantities} u,w,\\delta \\rho, \\delta p, \\delta y_s \\sim exp[i(k_x x +n t)], \\end{equation} and assuming that the flow is aligned with the perturbation vector, i.e. $k=k_x$, we arrive at \\begin{eqnarray} \\label{dispersionsrelation} && D \\left \\{ \\rho (n + k U) (D w) - k \\rho (D U)w \\right \\} - \\rho k^2 (n + k U)w = \\nonumber \\\\ & & i D \\left \\{ \\rho \\nu k^2 (D w) \\right \\} - i D \\left \\{ \\rho \\nu (D^3 w) \\right \\} - \\nonumber \\\\ & & D \\left\\{ k \\nu (\\rho + \\delta \\rho) (D^2 U) \\right \\} + i \\rho \\nu k^2 (D^2 w) - i \\rho \\nu k^4 w, \\end{eqnarray} where $D\\equiv d/dy$. The term, $i \\rho \\nu k^2 (D^2 w)$ in Eq.~\\ref{dispersionsrelation} can be replaced with \\begin{equation} i \\rho \\nu k^2 (D^2 w) = i k^2 D(\\rho \\nu (Dw)) - i k^2 (D w)(D(\\rho \\nu)). \\end{equation} The boundary condition at $y=y_s$ is determined by an integration over an infinitesimal element ($y_s-\\epsilon \\mbox{ to } y_s + \\epsilon$), for the limit $\\epsilon \\rightarrow 0$. Please note, that with Eq.~\\ref{lincont} it follows for $\\delta \\rho$, \\begin{equation} \\delta \\rho = i\\frac{w}{(n+k_x U)} (D \\rho). \\end{equation} After integration, the boundary condition becomes, \\begin{eqnarray} \\label{boundary} && \\Delta_s \\left\\{ \\rho (n+k U)(D w) -\\rho k (D U)w \\right\\} = \\nonumber \\\\ & & i k^2 \\Delta_s \\left\\{ \\nu \\rho (D w) \\right\\} - i \\Delta_s \\left\\{ \\nu \\rho (D^3 w) \\right\\} - \\nonumber \\\\ & & k \\Delta_s \\left\\{ \\nu \\rho (D^2 U) \\right\\}- i k \\Delta_s \\left\\{ \\nu \\frac{w}{(n+k U)} (D \\rho) (D^2 U) \\right\\} + \\nonumber \\\\ & & i k^2 \\Delta_s \\left\\{ \\nu \\rho (D w) \\right\\} - i k^2 \\lim_{\\epsilon \\rightarrow 0}{ \\int_{y_s-\\epsilon}^{y_s+\\epsilon}{(D w) D(\\nu \\rho)\\,dy}} \\end{eqnarray} where $\\Delta_s$ is specifying the jump of any continuous quantity $f$ at $y=y_s$, \\begin{equation} \\Delta_s(f)=f_{(y=y_s+0)}-f_{(y=y_s-0)}. \\end{equation} For $\\nu\\equiv 0$ we retrieve the corresponding expression as given by \\citet{Chandrasekhar}. \\subsection{Special case: constant velocities and densities} To simplify the derivation of the growth rate $n$ further, we consider the case of two fluid layers with constant densities $\\rho_1$ and $\\rho_2$, and constant flow velocities $U_1$ and $U_2=-U_1$. In each region of constant $\\rho_{1,2}$ and $U_{1,2}$, Eq.~\\ref{dispersionsrelation} reduces to, \\begin{eqnarray} \\label{dispersion_constantregion} & &\\left[ (n+kU_{1,2})\\rho_{1,2}-2i \\nu k^2 \\right] (D^2 w) + i \\nu (D^4w) - \\nonumber \\\\ & & k^2 \\left[ (n+kU_{1,2}) - i \\nu k^2 \\right]w=0 \\end{eqnarray} The layers are separated at $y=y_s=0$, and $w/(n+k U)$ must be continuous at the interface. Also, $w$ must be finite for $y\\rightarrow \\infty$, so that the solution of Eq.~\\ref{dispersion_constantregion} has the following form, \\begin{eqnarray} w &=& A(n+k U_1)e^{+k y} \\mbox{         } (y<0)\\\\ w &=& A(n+k U_2)e^{-k y} \\mbox{         } (y>0). \\end{eqnarray} We assume that $\\nu_1=\\nu_2=\\nu$ (which is the case if we consider two media with the same viscous properties). Inserting this in Eq. ~\\ref{boundary}, the characteristic equation yields, \\begin{eqnarray} && n^2 + 2 \\left[ k (\\alpha_2 U_2 + \\alpha_1 U_1) -\\frac{i k^2 \\nu}{2} \\right] n + \\\\ & & k^2 \\left( \\alpha_2 U^2_2 + \\alpha_1 U^2_1 \\right) - i k^3 \\nu \\left( \\alpha_2 U_2 + \\alpha_1 U_1 \\right) =0. \\end{eqnarray} The parameters $\\alpha_1$, $\\alpha_2$ are defined by, \\begin{eqnarray} \\label{alpha} \\alpha_1 = \\frac{\\rho_1}{\\rho_1+\\rho_2}, \\mbox{      } \\alpha_2 = \\frac{\\rho_2}{\\rho_1+\\rho_2}. \\end{eqnarray} \\begin{table*} \\label{table_conversion} \\begin{center} \\begin{tabular}{|c|c|c|}\\hline physical parameters & dimensionless &in cgs units \\\\\\hline\\hline Box size & 2 & 2 cm     \\\\ \\hline Mass & 4 & 2780.81 g \\\\ \\hline velocity &  0.387  & 0.40 km/s\\\\ \\hline time     &   1 & 9.8 $\\cdot 10^{-6}$s\\\\ \\hline \\end{tabular} \\caption[Conversion of code units to physical units] { Initial conditions in dimensionless units (first column) and in cgs units (second column). In the text we always refer to dimensionless units. } \\end{center} \\end{table*} Solving for n, we get the expression for the mode of the linear KHI: \\begin{eqnarray} \\label{KHI_mode} && n = - \\left[ k(\\alpha_2 U_2 + \\alpha_1 U_1) - \\frac{i k^2 \\nu}{2} \\right] \\pm \\nonumber \\\\ & & \\sqrt{ - k^2 \\alpha_1 \\alpha_2 (U_1-U_2)^2 - \\frac{k^4 \\nu^2}{4}}, \\end{eqnarray} applying $U_2=-U_1=U$ leads to \\begin{eqnarray} \\label{KHI_analytical} n &=& \\left[ k^2 U^2 (\\alpha_2-\\alpha_1) + \\frac{i k^2 \\nu}{2} \\right] \\pm \\nonumber \\\\ & & \\sqrt{ - 4 k^2 \\alpha_1 \\alpha_2 U^2 - \\frac{k^4 \\nu^2}{4}}. \\end{eqnarray} The mode is exponentially growing/decaying with time, if the square root of $n$ becomes imaginary, \\begin{eqnarray} \\label{mode_diff_layers} n &=& \\left[ k^2 U^2 (\\alpha_2-\\alpha_1) \\right] + \\nonumber \\\\ & & i \\left[ \\frac{\\nu k^2}{2} \\pm \\sqrt{\\frac{\\nu^2 k^4}{4}+4k^2U^2 \\alpha_1 \\alpha_2 } \\right]. \\end{eqnarray} The first term describes oscillations (which is not of interest for the growth), the second term the growth/decay, with a damping due to the viscosity. We use this formula for the comparison with our numerical studies for different density shearing layers. For equal density shearing layers $\\rho_1=\\rho_2=\\rho$, Eq.~\\ref{mode_diff_layers} leads to \\begin{equation} \\label{slope_equal_densities} n = i \\left[ \\frac{\\nu k^2}{2} \\pm \\left( \\frac{\\nu^2 k^4}{4}+ k^2 U^2 % \\right)^\\frac{1}{2} \\right]. \\end{equation} In \\S~\\ref{SPH_simulations} and \\S~\\ref{Grid_simulations} we use the velocity in direction of the perturbation, which in the above analysis refers to the $y$-direction and therefore, to the $v_y$-velocity component ($w$) when comparing with simulations. The exponential term in Eq.~\\ref{perturb_quantities} ($\\sim \\exp{(i \\cdot n \\cdot t)}$) describes the time evolution of the KHI. In the following, we therefore compare $\\ln{\\left(v_y\\right)}$ with the analytical expectation $\\ln{(w)} \\sim i \\cdot n \\cdot t$. ",
        "conclusions": "\\label{conclusions} We have studied the Kelvin-Helmholtz instability applying different numerical schemes. We use two methods for our SPH models, namely the Tree-SPH code VINE \\citep{Wetzstein_Nelson_Naab_Burkert_2008, Nelson_Wetzstein_Naab_2008}, and the code developed by \\cite{Price_2007}. The grid based simulations of the KHI rely on FLASH \\citep{Fryxell_2000}, while as a test for the non-viscous evolution we also apply PLUTO \\citep{Mignone_2007}. \\\\ We first extended the analytical prescription of the KHI by \\cite{Chandrasekhar} to include a constant viscosity. With this improvement we were able to measure the intrinsic viscosity of our subsequently performed numerical simulations. We test both SPH as well as grid codes with this method. \\\\ We then concentrated on the KHI-evolution with SPH. We performed a resolution study to measure the dependence of the KHI growth on the mean number of SPH neighbors ($\\bar{n}_{\\mathrm SPH}$) and the total number of particles, respectively. We found that our simulations were well resolved and that a different number of $\\bar{n}_{\\mathrm SPH}$ did not significantly influence the KHI growth rate. \\\\ In case of equal density shearing layers we then measured the intrinsic viscosity in VINE by evaluating our simulations against the analytical prediction in the linear regime. Without using the Balsara viscosity the AV parameters $\\alpha$ and $\\beta$ effectively lead to a damping of the KHI. The commonly suggested and used settings of $\\alpha=1$, and $\\beta=2$ result in a strong suppression of the KHI. More quantitatively, we derive values of $0.065 < \\nu_{\\mathrm{SPH}} < 0.1$ for $0< \\alpha < 1$. Different values of $\\beta$ do not have a strong impact on $\\nu_{\\mathrm{SPH}}$. By introducing the Balsara-viscosity the dissipative effects of the AV can be reduced significantly, effectively rendering the results to be independent of $\\alpha$ and $\\beta$. However, the constant floor viscosity of $\\nu_{\\mathrm{SPH}}=0.065$ prevails. Furthermore for a given $\\alpha$, we estimated the effective Reynolds-number ($Re$) of the flow. For the minimum SPH viscosity of $nu_{\\mathrm{SPH}}=0.065$ we derive a maximum Reynolds number of $12$. This is very small compared to typical Reynolds numbers of real turbulent flows ($Re > 10^5$). For different density shearing layers we confirmed the results discussed in \\cite{Agertz_2007}, i.e. the KHI is completely suppressed for shear flows with different densities (in the case of VINE for $DC\\ge 6$). Here, using the Balsara switch does not solve the problem. This indicates that other changes to the SPH formalism are required in order to correctly model shearing layers of different densities. To demonstrate this we applied the solution of \\cite{Price_2007} to our initial conditions for DC=10. In this case the KHI was suppressed in VINE. However, we found good agreement between the analytically predicted amplitude evolution and the simulation of \\cite{Price_2007} for $DC=10$. \\\\ The second part of this paper addresses the non-viscous- and viscous KHI evolution using grid codes. In the case of equal density shearing layers, we found the non-viscous growth rates for shear flows with FLASH and PLUTO to be in good agreement with the analytical prediction. In the viscous case, the FLASH-amplitudes show only a minor dependency on the viscosity if $\\nu < 0.03$. Increasing the viscosity leads to a damped evolution, with the simulated growth coinciding with the analytical prediction. \\\\ For non-viscous shear flows (with a density contrast of $DC=10$) the KHI does develop for FLASH and PLUTO in agreement with the analytical prediction. In the viscous case FLASH (also analyzed with $DC=10$) slightly overpredicts the corresponding growth rates for $\\nu < 0.03$ by a constant factor of $\\sim 0.12$. \\\\ The comparison between VINE, FLASH and PLUTO in the equal density case, where $AV=0$ and $\\nu=0$, demonstrated that VINE does have an intrinsic viscosity (which we estimated to $\\nu_{\\mathrm{int}} \\sim 0.065$)."
    },
    "1004/1004.1166_arXiv.txt": {
        "abstract": "Recent work on the engines of active galactic nuclei jets suggests their power depends strongly and perhaps counter-intuitively on black hole spin.  We explore the consequences of this on the radio-loud population of active galactic nuclei and find that the time evolution of the most powerful radio galaxies and radio-loud quasars fits into a picture in which black hole spin varies from retrograde to prograde with respect to the accreting material.  Unlike the current view, according to which jet powers decrease in tandem with a global downsizing effect, we argue for a drop in jet power resulting directly from the paucity of retrograde accretion systems at lower redshift $z$ caused by a continuous history of accretion dating back to higher $z$. In addition, the model provides simple interpretations for the basic spectral features differentiating radio-loud and radio-quiet objects, such as the presence or absence of disk reflection, broadened iron lines and signatures of disk winds.  We also briefly describe our models' interpretation of microquasar state transitions.  We highlight our result that the most radio-loud and most radio-quiet objects both harbor highly spinning black holes but in retrograde and prograde configurations, respectively. ",
        "introduction": "Over the past decade and a half, our understanding of the dynamics of active galaxies (AGN) has made significant strides.  Once of marginal interest astrophysically, black holes have taken center stage, becoming an integral part of galactic dynamics and evolution.  The picture that has emerged involves the presence of supermassive black holes at the center of most if not all galaxies, with active galaxies interacting with these black holes via accretion, producing winds and jets.  There is overwhelming evidence that galaxies are interconnected with black holes to the extent that the cosmic evolution of both is coupled (Kormendy \\& Richstone 1995; Magorrian et al 1998; Gebhardt et al 2000; Ferrarese \\& Merritt 2000; Tremaine et al 2002; Marconi \\& Hunt 2003).  Despite this, there are uncertainties in our understanding of how black hole engines produce jets, how such structures are collimated over large scales, and how galaxies that harbor these powerful engines evolve over cosmic time.  The theoretical framework in which we have attempted to address such questions over the past two decades involves the ``spin paradigm'' (Blandford 1990; Wilson \\& Colbert, 1995; Moderski, Sikora \\& Lasota, 1998), whereby high black hole spin can lead to powerful, radio-loud, jetted AGN, while low black hole spin to radio-quiet, weak or non-jetted AGN where jet power is related to black hole spin via the Blandford-Znajek mechanism (Blandford \\& Znajek 1977; henceforth BZ). The literature on the BZ effect and black hole spin-related processes for jet formation in relation to radio-loud AGN is extensive.  In its most recent guise, the spin paradigm enters the literature in various places and with differing results.  Using the BZ mechanism, Mangalam, Gopal-Krishna \\& Wiita (2009) propose a scenario leading to low black hole spins in advection dominated accretion flows under the assumption that the magnetic flux threading the black hole is insensitive to spin.  Similar ideas grounded in general relativistic magnetohydrodynamic simulations suggest that high spin for higher order spin power dependence in thick accretion geometries vs thin accretion flows, produce the conditions for the $3$ order of magnitude difference in power between radio-loud and radio-quiet AGN (Tchekhovskoy et al 2010).  In work that closely follows the results of numerical simulations that supports the work of Nemmen et al (2007) on high black hole spin in geometrically thick flows, Benson \\& Babul (2009), also using advection dominated thick accretion flow models, arrive at a maximum spin of $0.92$ for black holes, suggesting that a spin-equilibrium scenario is compatible with the linear accretion rate dependence of jet power found in recent observations of AGN jets (Allen et al 2006).  In addition to the question of the radio-loud/radio-quiet division in active galaxies, a division within the group of jetted AGN is observed and a classification is produced according to differences in power and jet collimation (Fanaroff \\& Riley 1974).  Attempts at understanding this Fanaroff-Riley (FR) division have also been framed in the context of the spin-paradigm in an attempt to explore the possibility of differences in spin between FRII and FRI objects.  Meier (1999) suggested a magnetic switch operating at high spin triggers the production of relativistic jets.  Using the BZ effect and hybrid Meier model, Daly (2009b) calculated the cosmological evolution of black hole spins, showing that they are largest (about unity) for the FRII sources at high redshift and decrease (to about 0.7) for radio-loud sources at lower redshift.  The idea that the engine of jets determines their FR morphology spans the works from Rees (1982) to Hardcastle et al (2007).  A theoretical approach to determining black hole spin that is model-independent, but assumes that spin changes only by extraction of the reducible black hole mass, applied to a small subset of powerful radio galaxies, finds that they harbor low spinning black holes (Daly 2009a).  This work suggests also that the FRI/FRII division is due to environment.  Such environmental effects on jet morphology, unrelated at least directly to the properties of the central engine of the jet, have been addressed starting with De Young (1993) and more recently with Gopal-Krishna and Wiita (2000). Black hole mass dependence (Ghisellini \\& Celotti 2001) and accretion dependence of the FRI/FRII division have also been explored (Marchesini et al 2004).  Under the assumption that jets are produced by the combined effort of the BZ and Blandford-Payne (Blandford \\& Payne 1982; henceforth BP) effects, recent numerical studies of general relativistic magnetohydrodynamics of black hole accretion flows suggest a tight link between jet power and black hole spin (Garofalo, 2009a, 2009b). Whereas BZ involves jets produced via spin-energy extraction from the black hole, BP is a jet mechanism that originates in the accretion disk via mass-loading of disk gas onto large scale magnetic fields. The black hole spin dependencies of these processes have direct implications for the jets produced in AGN.  According to these studies, largest BZ power occurs for highly retrograde accretion systems with respect to the black hole and less so for prograde ones, transitioning to zero power at zero spin while the BP power is also maximized for highly retrograde accretion but monotonically decreases toward high prograde ones.  In the context of this framework, we attempt in this paper to construct a physical foundation for the morphology and evolution of jet-producing AGN.  At the heart of our picture lies the notion that accretion onto spinning supermassive black holes tends, in time, to produce increasingly prograde accretion systems, in which lower jet output is hosted by more stable accretion configurations.  In terms of the radio-loud AGN population, ``high excitation radio galaxies'', evolve over cosmic timescales toward ``low-excitation radio galaxies'', as black holes conspire with their host galaxy in evolving toward prograde accretion states from an earlier phase of retrograde accretion flow.  In addition, we suggest a retrograde vs prograde division for the engines of radio-loud objects whose interaction with the external environment produces the distribution of the Fanaroff \\& Riley classification, according to which the powerful, highly collimated jets of the FRII class, are retrograde spin systems while the mostly less powerful, less collimated FRI sources embedded in gas-rich environments are the late-state evolution toward prograde accretion systems.  Our interpretation of the radio-loud/radio-quiet division will be that the most radio-loud and the most radio-quiet AGN both harbor rapidly rotating black holes with their accretion angular momentum vectors determining the difference, while the FRI/FRII division in our model is based on a combination of nuclear properties and environment. In addition to this, our model unifies radio-loud, radio-quiet, and FRI, FRII objects, to the extent, as we will illustrate, that radio-loud objects of the FRII class evolve either into radio-loud FRI objects or into radio-quiet AGN. We point out, furthermore, that the theoretical frameworks grounded in the spin-paradigm discussed above, all suffer from a ``spin paradox'' (David L. Meier) to be discussed in a follow-up paper in which we illustrate its resolution within our framework (Meier \\& Garofalo, in preparation).  In section \\ref{observation} we highlight the relevant observed properties of radio-loud AGN which constitute the pieces of the theoretical puzzle we construct in Section \\ref{Theory}.  Section \\ref{Conclusions} briefly addresses the radio-loud/radio-quiet dichotomy in this model and its extension to microquasars and then concludes. ",
        "conclusions": "We presented a framework for the cosmological evolution of radio-loud galaxies in which the gap region between accretion disks and black holes is key.  In this gap paradigm, jet production is most effective when black hole magnetospheres conspire with their large gap regions to produce strong black hole-threading magnetic fields that are geometrically favorable to both the BZ and BP effects.  As time evolution through accretion spins the black hole up in the prograde direction and the gap region size decreases, the interaction of rotating black holes with their magnetospheres becomes inefficient and jets weaken, become less collimated and if accretion spins the black hole up to high prograde values, likely fails entirely to produce jets.  Maximally spinning prograde black hole magnetospheres that are the result of cosmological evolution via accretion, thus, may be systems that have now become inactive.  The radio-quiet quasars formed in ellipticals may, instead, involve the population of post-merger systems characterized by high accretion onto spinning black holes in a prograde configuration.  According to our scenario, the origin of the radio-loud/radio-quiet dichotomy may lie in the existence of two types of engine efficiencies.  On the one hand, retrograde black holes conspiring with their magnetized flows are highly efficient in producing non-thermal, jet outflows, but shorter-lived configurations. These engines produce the observed radio-loud population.  On the other hand, prograde, high accretion systems with small gap regions, while incapable of producing powerful jets, are capable of tapping into the gravitational potential energy of their inner accretion flow, making them highly efficient in producing strong thermal emission. The reprocessed gravitational potential energy of the inner regions near the rapidly spinning black hole is fundamental in producing the disk winds from larger disk radii.  It has recently been shown that GRS 1915+105, a galactic microquasar, exhibits a soft state in which the radiation field drives a hot wind off the accretion disk carrying enough mass to halt the flow of matter into the radio jet (Neilsen \\& Lee 2009), a picture that is underscored by theoretical work (Meier 1996) and for which observation is mounting in other systems as well (Blum et al. 2010).  This scenario suggests that most FRI HERGs are systems in which the disk winds eventually inhibit jet production making such objects part of the radio-quiet population, thereby explaining the rarity of FRI HERGs.  The handful of observed FRI HERGs are then systems where the black hole spin is small, producing relatively larger gap regions, which decreases the production of winds further out from the disk and allows the presence of an FRI jet.  In other words, in the cosmological evolution of radio-loud objects, FRI HERGs are objects that are transitioning away from strong jet production and weak winds, to a scenario involving weaker jets and stronger winds but that have not reached the high prograde spin states that would produce strongest winds and weakest or no jets.  We highlight the compatibility of this scenario with the existence of FRI type radio quasars (Heywood et al 2007) as pointed out to us by David L. Meier.  The fact that accretion states transition from HERGs to LERGs, however, allows such systems to remain radio-loud as they increase their prograde spins as a result of the failure of ADAFs to produce the strong jet-inhibiting disk winds.  Thus, spinning black holes produce both FRII HERGs and radio-quiet quasars depending on retrograde vs prograde accretion.  If mergers are responsible for both the high accretion rates and the occurrence of retrograde flows, it does not surprise that spiral galaxies are not competitive in their radio-loudness, and less accretion powered than the post-merger ellipticals.  If retrograde accretion is formed in mergers, it also does not suprise that spirals are predominantly radio-quiet.  In fact, only the low prograde spin configurations would lead to radio-loudness.  Given that higher prograde accretion systems surrounded by radiatively efficient thin disks have smallest gap regions and the strongest jet-inhibiting winds, higher prograde spin systems in spirals should be radio-quiet; and, therefore, maximally spinning prograde configurations should be the most radio-quiet sources.  In short, prograde vs retrograde, radiatively efficient accretion flows, produce the radio-quiet/radio-loud dichotomy.  And finally, we conclude our discussion with a glance at the implications for microquasars.  Because such objects are prograde accretion systems, their accretion state changes produce the microquasar counterpart to AGN HERGs and LERGs, or more appropriately, a transition between high prograde spin, high excitation accretion states (or radiatively efficient accretion states) to high prograde spin, low excitation accretion states (or radiatively inefficient accretion states).  If the parallel with the FRII feedback phase in AGN is valid, these high prograde spin, high excitation accretion states in microquasars produce strong disk winds that expel the gas (Neilsen \\& Lee 2009), thereby producing an ADAF phase during which disk winds drop and jets are no longer hindered while the system enters the radio dominated hard state.  Eventually, the accretion flow from the companion allows the system to re-enter the soft or radiatively efficient accretion state.  The extension of our model to microquasars suggests that the approach to the high excitation accretion (soft) state from the low excitation accretion (hard) state, involves the brief combined presence of a BZ component and a BP-like component produced by the onset of inner disk winds.  In other words, there is a short time during which the microquasar mimicks the conditions in FRII HERGs, thereby producing a powerful, collimated jet.  Note that no spin change occurs.  This suggests that microquasars undergo state transitions that take them from high prograde spinning high excitation accretion states (soft states) to high prograde spinning low excitation accretion states (hard states) and back, and that during the low excitation to high excitation transition, the physical conditions of an FRII HERG are briefly produced.  We summarize our main conclusions as follows:  \\begin{enumerate} \\item FRIIs have retrograde spin while FRIs have prograde spin \\item FRII HERGs evolve toward FRI LERGs \\item FRIIs have larger jet efficiencies while FRIs have lower jet efficiencies \\item High prograde spin systems in HERG states have large disk efficiencies (and thus strong disk winds) while LERG states and retrograde accretion states have lower disk efficiencies (and thus lower disk winds) \\item Highly efficient jet engines are highly inefficient disk engines and viceversa \\item FRI HERGs have low prograde spin \\item Radio quiet AGN are high spinning prograde, radiatively efficient systems \\item Maximally spinning prograde, radiatively efficient systems, are the most radio quiet AGN \\item Spiral galaxies involve prograde accretion \\item Spiral galaxies should be more radio loud when the spin is lower and more radio quiet, quasar-like, when the spin is high \\item A radiatively inefficient to radiatively efficient transition, in high prograde systems, produces a short-lived, powerful, collimated jet as observed in microquasars. \\end{enumerate}  We conclude by highlighting the fundamental role or impact of general relativity in galaxy evolution in this model.  The validity of our paradigm suggests that aspects of the spacetime metric that are assumed to govern physics in regions provincially relegated to the near black hole, are in fact dominant on much larger scales, to the extent that galaxy morphology, energetics, and evolution, are tightly linked to the details of strong-field general relativistic effects. Whereas a Newtonian treatment of spacetime is successful just a few gravitational radii from the black hole in the sense that relativistic corrections are negligible there, the scale of influence of black hole spin that our paradigm forces us to grapple with is daunting, with up to more than eight orders of magnitude beyond its local sphere of influence.  It is the details of tiny regions of highly curved spacetime that have the greatest effect on the large-scale properties of galaxies.  \\label{Conclusions}"
    },
    "1004/1004.4137.txt": {
        "abstract": "\\begin{list}{ } {\\rightmargin 1in} %{\\leftmargin -1in} \\baselineskip = 11pt %rule{4.75in}{0.5pt} %\\vskip 1pt \\parindent=1pc {\\small Planet migration is the process by which a planet's orbital radius changes in time. The main agent for causing gas giant planet migration is the gravitational interaction of the young planet with the gaseous disk from which it forms. We describe the migration rates resulting from these interactions based on a simple model for disk properties. These migration rates are higher than is reasonable for planet survival. We discuss some proposed models for which the migration rates are lower. There are major uncertainties in migration rates due to a lack of knowledge about the detailed physical properties of disks. We also describe some additional forms of migration. %Write your concise abstract here. \\\\~\\\\~\\\\~}%leave this in to get the correct vertical space after the abstract  \\end{list} ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.2215_arXiv.txt": {
        "abstract": "{% The dominant part of the difference between the observed and model frequencies of the Sun can be approximated by a power law. We show that when this empirical law is employed to correct the model frequencies and then the small frequency separations, $\\delta\\nu_{02}(n)$, are used for solar age determination, the results are consistent with the meteoritic age (4.563 Gyr $<$ t $<$ 4.576 Gyr). We present the results and compare with those obtained by using the ratios, $r_{02}(n)$, of small to large frequency separations.} ",
        "introduction": "\\sloppy  It is known that there is a systematic offset between the observed and model frequencies of the Sun. This offset increases with increasing frequency and is shown (Kjeldsen et al.~2008) to be fitted well with a power law as \\begin{equation} \\nu_{\\rm{obs}}(n)-\\nu_{\\rm{best}}(n)=a\\left[\\frac{\\nu_{\\rm{obs}}(n)}{\\nu_{0}}\\right]^b, \\end{equation} where the power $b$ is determined to be 4.90 for the Sun. Here $\\nu_{\\rm{obs}}(n)$, and $\\nu_{\\rm{best}}(n)$, represent the observed, and the best model, frequencies with spherical degree $l$=0, and radial order $n$; $\\nu_{0}$ is a constant frequency (chosen to be 3100 $\\mu$Hz for the Sun), which corresponds to the frequency for peak power in the spectrum. This difference between the observed and calculated frequencies exists due to improper modelling of the outer turbulent convective layers of the Sun. The outer layers affect the high frequencies most, as the upper turning point of the high-frequency waves are closer to the surface. Since all stellar models are calibrated with respect to the Sun, and since, thanks to the recent developments, we are at a stage to have observations of individual frequencies of stars other than the Sun, it is important to understand the effects of near-surface stellar layers on the oscillation frequencies. These effects also influence the small frequency separations, $\\delta\\nu_{l \\hspace{1pt} l+2}(n)=\\nu_{nl}-\\nu_{n-1\\hspace{2pt}l+2}$, and hence, for instance, inferences of stellar ages. Here $\\nu_{nl}$ is the frequency of a mode with spherical degree $l$ and radial order $n$. The use of $\\delta\\nu_{l \\hspace{1pt} l+2}(n)$ for the purpose of solar-age determination has been presented in earlier works (e.g. Dziembowski et al.~1999; Bonanno et al.~2002). It has also been shown by several authors (e.g. Roxburgh \\& Vorontsov 2003; Ot\\'{i} Floranes et al.~2005; Christensen-Dalsgaard 2009) that the frequency separation ratios, $r_{l \\hspace{1pt} l+2 }(n)=(\\nu_{nl}-\\nu_{n-1\\hspace{1pt}l+2})/(\\nu_{nl}-\\nu_{n-1\\hspace{1pt}l})$, are not as sensitive to the near-surface layers. Christensen-Dalsgaard (2009) provides a comparison between the use of $\\delta\\nu_{l \\hspace{1pt} l+2}(n)$ and $r_{l \\hspace{1pt} l+2 }(n)$. Here, we present our preliminary results on comparing different ways of using the seismic data for determination of the solar age, including the application of near-surface correction.  \\fussy  \\vspace{-8pt} ",
        "conclusions": "The results of the age determinations using different input, as described in Section 2, are summarized in Table 1, and one of the cases is illustrated in Fig. 1.  The ages of our best models found using the two different evolution codes are consistent within the error limits, differing by 0.7\\% at most, and they are also compatible with the values in the literature obtained by employing small frequency separations: 4.57 $\\pm$ 0.11 Gyr (Bonanno et al.~2002), and 4.66 $\\pm$ 0.11 Gyr (Dziembowski et al.~1999). Moreover, our two sets of results agree with the meteoritic age, 4.563 Gyr $<$ t $<$ 4.576 Gyr (Wasserburg, in Bahcall \\& Pinsonneault 1995), within the uncertainty limits.  In principle we can determine the duration of the PMS evolution, up to the ZAMS, by calibrating the models using PMS evolution to get the meteoritic age of the Sun (see, e.g. Morel et al.~2000, for a discussion on the definition of ZAMS). Our preliminary analysis suggests that the PMS evolution for the Sun is 0.05$\\pm$0.11 Gyr in order to be consistent with the meteoritic age. A more precise and detailed analysis is obviously needed for defining ZAMS.  The results obtained using separation ratios are consistent with those obtained using the small frequency separation of the corrected frequencies. However, since the application of the surface correction has no significant effect on the age determined by the separation ratios, using these ratios yields a more robust age determination."
    },
    "1004/1004.1107.txt": {
        "abstract": "We have examined the luminosity-size relationship as a function of environment for 12150 SDSS galaxies with precise visual classifications from the catalog of \\cite{Nair:2010p22790}. Our analysis is subdivided into investigations of early-type galaxies and late-type galaxies. Early-type galaxies reveal a surprisingly tight luminosity-size relation. The dispersion in luminosity about the fiducial relation is only $\\sim 0.14$ dex (0.35 mag), even though the sample contains galaxies which differ by a factor of almost $100$ in luminosity. The dispersion about the luminosity-size relation is comparable to the dispersion about the fundamental plane, even though the luminosity-size relation is fundamentally simpler and computed using purely photometric parameters. The key contributors to the dispersion about the luminosity-size relation are found to be color and central concentration. Expanding our analysis to the full range of morphological types, we show that the slope, zero point, and scatter about the luminosity-size relation is independent of environmental density. Our study thus indicates that whatever process is building galaxies is doing so in a way that preserves fundamental scaling laws even as the typical luminosity of galaxies changes with environment. However, the distribution of galaxies {\\em along} the luminosity-size relation is found to be strongly dependent on galaxy environment. This variation is in the sense that, at a given morphology, larger  and more luminous galaxies are rarer in sparser environments. Our analysis of late-type galaxy morphologies reveals that scatter increases towards later Hubble types. Taken together, these results place strong constraints on conventional hierarchical models in which galaxies are built up in an essentially stochastic way. ",
        "introduction": "\\label{sec:introduction}  In hierarchical models the formation of galaxies is driven by multiple mergers and complicated feed-back effects. As a result the star-formation and mass-building history of individual galaxies can differ in a multitude of ways. However, a common feature of these models is that the end-product of this process is generally a system which lies on (or at least near) a fundamental scaling relation e.g. the fundamental plane for elliptical galaxies \\citep{BoylanKolchin:2006p21539,Robertson:2006p15098,Ciotti:2007p22639,Covington:2008p1439}, or the Tully-Fisher relation for spiral galaxies \\citep{Dalcanton:1997p21815}). In some models this agreement is somewhat artificial, because the form of the scaling relation is taken as an input parameter, but in other cases it emerges because some fundamental relations are closely linked to simple underlying physics (e.g. the Faber-Jackson relation emerges from the virial theorem coupled with assumptions about mass-to-light ratio being a weak function of mass).  How can measurements of the fundamental scaling relations of galaxies best be used to constrain ideas for galaxy formation? It is arguable that the most direct way forward is simply to measure how the form of these relations changes with environment.  Environment is the central parameter in hierarchical models, since in this picture the ultimate fate of a galaxy is determined mainly by its merger history, and the rate of merging is accelerated in richer environments, which are expected to host the oldest and most massive galaxies at any time \\citep{Mo:1996p22126}. It is perhaps even more interesting to measure the {\\em scatter} about the fundamental relations as a function of environment, because even if a relation's mathematical form is set by simple physics, the process by which a galaxy winds up on that relation is governed by random processes. A galaxy's stochastic merger history sets the timing over which it builds, but its final state is the product of a complex interplay between a large number of parameters, such as its initial mass, angular momentum and dark halo concentration, not to mention feedback to the intergalactic medium, so it would be surprising if the scatter about the fundamental relations were very small.  In a recent paper \\cite{vandenBergh:2008p11891} found that the luminosity-diameter relation provides the tightest of all the purely photometric correlations used to characterize galaxies. This fact was first noted by \\cite{Giuricin:1988p12909}, who also claim to have detected ``appreciable differences in the galaxy-luminosity relationships for different clusters''. On the other hand,  \\cite{Girardi:1991p12914} found no significant differences between the luminosity-diameter relations in a variety of environments. \\cite{Gavazzi:1996p21749} and \\cite{vandenBergh:2008p11891} have reached similar conclusions, albeit with small samples, and both suggest that galaxies are best viewed as complex systems linked by a single fundamental property, namely stellar mass. This viewpoint is supported by  \\cite{Disney:2008p12920}, who concluded the same thing on the basis of a principal component analysis of data from a sample of galaxies selected from a large blind survey of neutral hydrogen gas emission.  On the other hand, recent work \\citep{Shen:2003p70,Kauffmann:2003p7199,Bernardi:2003p1624,Desroches:2007p19582,Hyde:2008p12893} using much larger samples drawn from the SDSS \\citep{York:2000p3192} have shown the size-luminosity and size-mass relations exhibit curvature for both early and late type galaxies. In addition, \\cite{Bernardi:2007p1618,Bernardi:2009p21529,Lauer:2007p21894,vonderLinden:2007p21830,Liu:2008p22026,Coenda:2009p18836} have found that the central or brightest cluster galaxies (BCGs) exhibit a steeper size-luminosity relation than field early type galaxies, though \\cite{Guo:2009p19543} and \\cite{Weinmann:2009p19068} do not find the trend.  The works just described are based on either small ($<$1000) samples of visually classified objects or much larger samples ($>$50,000) with very broad morphological segregations based on measurements of central concentration. The recent publication of the detailed morphological catalog of \\cite{Nair:2010p22790} offers us an opportunity to revisit the question of the environmental dependence of the galactic scaling relations as a function of morphology and environment using samples with high-precision classifications that are an order of magnitude larger than those available to previous authors. Our aim in the present paper is to determine whether the form and scatter of the luminosity-size relation for galaxies remains constant as a function of environment. We will explore a range of Hubble types from E to Sc, but our main emphasis will be on early-type systems. This is because in hierarchical models elliptical galaxies are the end-product of major mergers, so they are a crucible for testing the central idea of merger-driven evolution. Early-type galaxies also provide us with the closest connection with existing work on this subject, because they dominate the populations of rich clusters, and most previous work on the environmental dependence of the luminosity-size relationship has focused on simple comparisons between rich clusters and the field. Our own sample differs from this earlier work in that it does not focus on rich clusters, and instead spans a range of environments from sparse regions in the field up to systems that could best be characterized as poor-intermediate clusters. In our study there are only 700 galaxies in clusters with $\\log(M_{halo}/ M_\\odot)>14.0$ and our data includes galaxies from only 47 clusters with more than 50 members (described below).  The plan for this paper follows: Our sample is described in Section~\\ref{sec:sample}, followed by our measurements of the fiducial luminosity-size relationships in Section~\\ref{sec:LR}. The dependence of these relationships on environment is explored in Section~\\ref{sec:LRenv}. Section~\\ref{sec:LRcc} looks into the importance of color and central concentration on scaling relations. Section~\\ref{sec:BCG} investigates the luminosity-size relationship for the brightest cluster galaxies (BCGs) in comparison to field and cluster satellite galaxies. Our results are discussed in Section~\\ref{sec:discussion}, and conclusions are presented in Section~\\ref{sec:conclusion}. We compare the two SDSS size measures, R(90) and R(dev), in Appendix A and compare the publicly available environment estimates used in this paper in the Appendix B. Throughout this paper we assume a flat dark energy-dominated cosmology with $h=0.7$ and $\\Omega_\\Lambda$=0.7. All magnitudes quoted are in the AB system. ",
        "conclusions": "\\label{sec:conclusion}  The present investigation is based on a sample of 12150 galaxies that are contained in the Sloan Digital Sky Survey \\citep{York:2000p3192}. Visual morphological classifications by \\cite{Nair:2010p22790}, and environmental information from \\cite{Blanton:2005p79}, \\cite{Baldry:2006p103} and \\cite{Yang:2007p19054}  were employed to study the dependence of the luminosity-radius relation of these objects on both their morphological type and environmental density. The main conclusions of this work are: \\begin{enumerate} \\item The luminosity-size relation of elliptical galaxies is well represented by a simple power law. The scatter about this relationship is very small: for elliptical galaxies, the dispersion is only 0.14 dex (0.35 mag), making this relationship as tight as that defined by the fundamental plane, even though it is based on purely photometric parameters. \\item The scatter about this luminosity-size relation is due to variations in both galaxy central concentration and star-formation history. Increasing central concentration has the effect of making modest changes in the slopes and intercepts of the fiducial relations. Although for all subsets of early-type galaxies the scatter remains small, the scatter in the luminosity-size relationship is $\\sim50$\\% larger for a small fraction ($<10\\%$) of early-type galaxies in the blue cloud than on the red sequence. \\item The slope and dispersion about the luminosity-size relationship increases as one proceeds along the Hubble sequence $E \\Rightarrow {\\rm Sa} \\Rightarrow {\\rm Sb}  \\Rightarrow {\\rm Sc}$. \\item Elliptical galaxies are presumed to have formed by the chaotic mergers of either smaller ellipticals or of disk-like ancestral objects, but it is difficult to see how such a process could result in a very tight luminosity-size relationship. \\item For galaxies of a given Hubble type, the slope of the luminosity-size relationship seems independent of environment. However the range spanned by the luminosities (or sizes) of that class of galaxy seems to be a strong function of environment. \\item The de Vaucouleurs radius, Rdev, and Petrosian half light radius, R(50), are more sensitive to changes in central concentration and hence merger history than R(90). Thus it appears mergers only influence the internal structure of galaxies and not their global luminosity-size relation.  \\end{enumerate} It is not clear why objects that might have assembled in such very different ways, from differing ancestral objects, should have had evolutionary tracks that converged to  show small dispersions around simple power-law forms for their luminosity-size relationships. Perhaps elliptical galaxies  started out very differently from each other at $z\\sim3$, but their morphologies converged to similar configurations as they expanded between redshifts $z = 3$ and $z = 0$ \\citep{Johansson:2009p18681}."
    },
    "1004/1004.2492_arXiv.txt": {
        "abstract": "% Gamma Ray Bursts (GRBs) show evidence of different light curves, duration, afterglows, host galaxies and they explode within a wide redshift range. However, their spectral energy distributions (SEDs) appear to be very similar showing a curved shape. Band et al. (1993) proposed a phenomenological description of the integrated spectral shape for the GRB prompt emission, the so called {\\it Band function}. In this letter we suggest an alternative scenario to explain the curved shape of GRB SEDs: the {\\it log-parabolic model}. In comparison with the Band spectral shape our model is statistically favored because it fits the GRB spectra with one parameter less than the Band function and it is motivated by a theoretical acceleration scenario. The new \\fer~ observations of GRBs will be crucial to disentangle between these two models. ",
        "introduction": "The physical mechanisms behind the GRBs prompt emission are still under debate. Band et al. (1993), investigating the BATSE GRBs sample proposed a phenomenological description of the integrated spectral shape for the GRB prompt emission, the so called {\\it Band function}. The introduction of this function was strongly suggested by the observational evidence that the shape of the Spectral energy Distribution (SED) of the GRB prompt emission is convex and broadly peaked. It is remarkable that there has not been physical explanation in terms of accelerations processes and non-thermal radiative losses that can led to the Band spectral shape.  In this letter we propose to describe the curved shape of GRB prompt emission using the {\\it log-parabolic} function (Massaro et al. 2004), successfully used to describe the SEDs of BL Lac objects over several decades. First, we consider the differences between this model in comparison with the Band function investigating the different $\\gamma$-ray flux predictions in the \\fer~LAT energy range. Second, we point toward the physical interpretation of the log-parabolic shape in terms of Fermi acceleration mechanisms. Finally, we present a simple synchrotron emission model to explain the GRB prompt emission, that appears to be the most reasonable scenario.  For our numerical results, we use cgs units unless stated otherwise and we assume a flat cosmology with $H_0=72$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M}=0.27$ and $\\Omega_{\\Lambda}=0.73$ (Spergel et al. 2007). ",
        "conclusions": "In the present letter we propose to interpret the SEDs of GRBs prompt emission using the log-parabolic shape. In comparison with the Band function (Band et al. 1993) the log-parabolic shape is favored for two main reasons.  First, it is statistically better, because it requires only 3 parameters, namely the curvature $b$, the peak energy $E_p$ and the height of the SED evaluated at the peak energy $S_p$ (see Eq. 6 and Eq. 15), while the usual Band model needs 4 spectral parameters, namely: $\\alpha$, $\\beta$, $E_c$ and $F_0$ or 5 if another high energy exponential cut off is introduced. Second, the proposed function has a strong physical motivation. This shape is directly related to the solution of kinetic equation for the particles accelerated by Fermi mechanisms when the random acceleration is also taken into account with all the other terms (e.g. first order Fermi mechanisms) (Kardashev 1962).  In the recent \\fer~LAT observations only few GRBs have been detected in contrast with the predictions of the Band function. The high energy curvature of the log-parabolic shape has a natural explanation for the \\fer~observations without the introduction of any new parameter, as the exponential cut-off in the Band function.  As shown in Massaro et al. (2006) or more recenly in Tramacere et al. (2009) and Paggi et al. (2009), the synchrotron emission of a log-parabolic electron distribution yields a curved SED near its peak, well described in terms of the same spectral shape. In this letter, we also presented how the synchrotron scenario from a log-parabolic PED can describe the spectrum of the GRBs prompt emission, successfully.  Finally, we remark that a different scenario, including other synchrotron or inverse Compton components and their spectral evolution with time, can make the spectral shape of the GRB prompt emission more complex in the \\fer~LAT energy range as for example recently observed in the case of GRB 090902B (Abdo et al. 2009)."
    },
    "1004/1004.0945_arXiv.txt": {
        "abstract": "We discuss how simultaneous observations by multiple heliospheric imagers can provide some important information about the azimuthal properties of Coronal Mass Ejections (CMEs) in the heliosphere. We propose two simple models of CME geometry that can be used to derive information about the azimuthal deflection and the azimuthal expansion of CMEs from SECCHI/HI observations. We apply these two models to four CMEs well-observed by both {\\it STEREO} spacecraft  during the year 2008. We find that in three cases, the joint {\\it STEREO-A} and {\\it B} observations are consistent with CMEs moving radially outward. In some cases, we are able to derive the azimuthal cross-section of the CME fronts, and we are able to measure the deviation from self-similar evolution. The results from this analysis show the importance of having multiple satellites dedicated to space weather forecasting, for example in orbits at the Lagrangian L4 and L5 points. ",
        "introduction": "\\label{intro}  With the launch of the Solar Mass Ejection Imager (SMEI) and the {\\it Solar Terrestrial Relations Observatory (STEREO)} in 2003 and 2006, respectively, Coronal Mass Ejections (CMEs) can now be observed remotely all the way to the Earth \\citep[]{Harrison:2009}, and their properties compared with in-situ observations at 1~AU \\citep[]{Davis:2009, Rouillard:2009c}. Although the Heliospheric Imagers \\citep[HI-1 and HI-2, see][]{Eyles:2009}, part of the Sun Earth Connection Coronal and Heliospheric Investigation (SECCHI) suite onboard {\\it STEREO} \\citep[]{Howard:2002}, enable the tracking of CMEs, they cannot provide directly the CME position because what is observed is the Thomson scattered signal integrated over a line-of-sight. Therefore, the CME position and kinematics have to be derived from a variety of models, which make assumptions regarding the CME shape and its direction of propagation \\citep[]{Vourlidas:2006, Sheeley:1999, Tappin:2009}. In \\citet{Lugaz:2009b}, using real and simulated data from \\citet{Lugaz:2009a}, we compared some of the assumptions typically used: i) Fixed-$\\phi$, where it is assumed that the HIs track a single plasma element moving radially outward, ii) Point-P, where the CME is treated as an expanding sphere centered on the Sun, and iii) a newly proposed harmonic mean (HM) model, where the CME cross-section is a circle anchored at the Sun and its diameter corresponds to the harmonic mean of the positions derived by the Fixed-$\\phi$ and Point-P approximations. We found that the HM model gives the best results for a wide CME observed by one spacecraft on the limb, while the Fixed-$\\phi$ approximation, which is the most widely used, gives comparable results up to approximatively 0.5~AU.  Most of the work so far on SECCHI/HI data has focused on observations by a single spacecraft, whereas stereoscopic CME observations by the SECCHI coronagraphs (COR-1 and COR-2) have resulted in a number of studies to improve our understanding of CME physical parameters. For example, \\citet{Thernisien:2009} used multi-spacecraft observations to determine the CME direction, speed and orientation, and \\citet{Colaninno:2009} used multi-spacecraft observations to determine the CME direction and mass. In \\citet{Lugaz:2005}, we showed how CME images made by wide-angle white-light imagers from different viewpoints can help in determining the CME direction of propagation in the heliosphere. Recently, \\citet{Wood:2009} proposed a technique similar to that of \\citet{Thernisien:2009} but which is expanded to the HI fields-of-view and with a full treatment of the Thomson scattering. The authors use a geometrical model, which is fitted visually to joint observations by the two {\\it STEREO} spacecraft to derive the CME aspect, density structure and orientation. \\citet{Liu:2010}, in a recent study, calculated the CME central longitude by triangulation in the COR and HI field-of-views for one CME event and obtained good results up to approximatively 0.5~AU.  \\begin{figure*}[ht*] \\begin{center} {\\includegraphics*[width=8.cm]{f1a.eps}} \\hspace{0.1cm} {\\includegraphics*[width=8.cm]{f1b.eps}} \\caption{Sketches of the the two models used and described in this work, illustrating two ways to explain asymmetric observations. On the left, the CME is modeled as a sphere connected to the Sun with a varying direction of propagation $\\phi$. On the right, the CME is modeled as a sphere on a fixed direction of propagation $\\phi = -21^\\circ$ but with a radius, $R_1$ varying with time. The sketches are to scale except for the size of the Sun, the Earth and the spacecraft and correspond approximatively to the geometry on April 26, 2008. The angles are $\\epsilon_A = 39^\\circ$ and $\\epsilon_B = 18^\\circ $. The varying $\\phi$ model is with $\\phi = -39^\\circ$ and $R = 135~R_\\odot$ and the varying radius model with $R_1 = 35~R_\\odot$ and $R_2 = 90~R_\\odot$.} \\end{center} \\end{figure*}  In this article, we discuss how simultaneous observations of a CME from the two {\\it STEREO} spacecraft can be used to derive the CME central longitude or its radius of curvature, in addition to its radial distance. To do so, we propose two simple models of the CME described in section~\\ref{model}. Contrary to triangulation techniques, we do not assume that the HIs are able to track the exact same feature for both spacecraft, but that they observe different parts of the same structure. In the first model, the CME cross-section is treated as an expanding-propagating circle attached to the Sun as proposed separately in \\citet{Webb:2009}, \\citet{Tappin:2009} and \\citet{Lugaz:2009b}, but with a varying central position. In the second model, the CME has a fixed direction but it is not attached to the Sun and its diameter can vary freely. We apply these two models to the analysis of 4 CMEs observed in 2008 in section~\\ref{data}. The conclusions of this investigation are drawn in section \\ref{conclusions}. ",
        "conclusions": "In this article, we propose two models to derive information about the azimuthal properties of CMEs from multi-spacecraft observations in the heliosphere using simple geometrical considerations. These models can be used to derive the CME radius or central position from SECCHI observations without having to rely on any extra human judgement or fitting procedure, after obtaining the time-elongation measurements. The main hypothesis of these two models is that the HI instruments onboard {\\it STEREO-A} and {\\it B} do not observe the same plasma element, as is likely to be the case in the COR fields-of-view and as is assumed in \\citet{Liu:2010}. The assumption here is that the CME front is locally circular and that it is viewed in the HIs at the elongation angle corresponding to the tangent to this front. We applied these two new models to six features belonging to four CMEs observed in 2008 by both STEREO spacecraft.  For five of the six studied features, the April 26 (center and front), June 2 and December 12, 2008 (front 1 and front 2) CMEs, we find that the measurements can be explained as being from CMEs propagating close to radially outward. However, for two of the CMEs (the June 2 CME and both features from the December 12 CME), we do not assume a radial propagation and we derive the CME central longitude directly with model 1. We find that the CME central longitude remains within 15$^\\circ$ of radial, with the June 2 CME, in particular, being deflected monotonically towards the east by about 5$^\\circ$ in 0.5~AU. For these CMEs, model 2, which assumes radial propagation but does not assume a self-similar expansion, does not provide additional information. Also, the two features associated with the December 12 CME are found to be separated by about 10$^\\circ$ at all time, which was also found using direct triangulation by \\citet{Liu:2010}. %   For the April 26, 2008 CME, the model assuming self-similar expansion (model 1) appears to fail, because it predicts a large deflection of the CME which is not physically expected nor confirmed by in-situ measurements. Self-similar expansion of CMEs in the heliosphere has long been assumed \\citep[]{Xue:2005, Krall:2006, Wood:2009}, but in general, it has not been tested with heliospheric observations. We analyze this CME with model 2 , where self-similar expansion is not assumed and using the direction of propagation derived by \\citet{Thernisien:2009}.  We derive the change over time of the CME cross-section, which is found to decrease with a rate of about $4^\\circ$/day. The cause of this decreasing cross-section has to be further studied, but it shows that the CME expands slower than self-similarly. This may, for example, reflect a variation of the CME radius of curvature in the ecliptic plane. This result is found for solar minimum conditions when the background solar wind is more simply structured.  For the other CME (August 30, 2008) the measurements can only be explained, with model 1, by a a deflection towards the East of $30^\\circ$ in 0.5~AU of the CME. It is also possible that an instance of CME-CME interaction locally deformed this CME front away from circular which would render our models less accurate.     In this study, we ignore the effect of the Thomson sphere and the angular dependence of the Thomson scattering intensity. To apply the new analysis techniques presented here, the CME has to propagate between the {\\it STEREO-A} and {\\it STEREO-B} spacecraft. For this study focusing on CMEs observed in 2008, this means that the CMEs propagate within 40$^\\circ$ from the Sun-Earth line, and we believe that ignoring the difference in Thomson scattering between the two spacecraft is justified as a first approximation. However, as the two {\\it STEREO} spacecraft continue to separate, CMEs propagating farther away from the Sun-Earth line will be observed simultaneously by both spacecraft and a more thorough analysis of the effect of the two different Thomson spheres should be made. Additionally, effects associated with the interaction of a CME with the Thomson sphere have been found to happen at elongation angles greater than $45^\\circ$ (e.g. \\citet {Manchester:2008}, \\citet{Lugaz:2008}), which is farther than that analyzed here.  We believe that, by providing a value of the CME radius in the ecliptic plane or the time-dependent variation of the central longitude which can be simply computed from time-elongation measurements,  these models could improve space weather forecasting of CMEs. However, to use such techniques on a real-world basis, it might be useful to have dedicated spacecraft making heliospheric observations at fixed locations, for example from the Lagrangian points L4 and L5. In general, we believe our model 1 can be used to derive the CME central longitude and its temporal variation but is not applicable for these CMEs whose expansion is not self-similar. When this is the case, for example because of interaction with the structured solar wind or with previous CMEs, we propose a second model, which can quantify the deviation from self-similar, assuming radial propagation. Future work should focus on the radius and central longitude of a CME at different position angles and on further testing and validation of the models."
    },
    "1004/1004.4757_arXiv.txt": {
        "abstract": "The feedback between massive stars and the interstellar medium is one of the most important processes in the evolution of dwarf galaxies. This interaction results in numerous neutral and ionised gas structures that have been found both in the disc and in the halo of these galaxies. However, their origin and fate are still poorly understood.\\\\ We here present new \\HI\\ and optical data of two Magellanic irregular dwarf galaxies in the Local Volume: IC\\,4662 and NGC\\,5408. The \\HI\\ line data were obtained with the Australia Telescope Compact Array and are part of the ``Local Volume \\HI\\ Survey''. They are complemented by optical images and spectroscopic data obtained with the ESO New Technology Telescope and the ESO 3.6m telescope. Our main aim is to study the kinematics of the neutral and ionised gas components in order to search for outflowing gas structures and to make predictions about their fate. Therefore, we perform a Gaussian decomposition of the \\HI\\ and \\Ha\\ line profiles.\\\\ We find the \\HI\\ gas envelopes of IC\\,4662 and NGC\\,5408 to extend well beyond the optical discs, with \\HI\\ to optical diameter ratios above four. The optical disc is embedded into the central \\HI\\ maximum in both galaxies. However, higher resolution \\HI\\ maps show that the \\HI\\ intensity peaks are typically offset from the prominent \\HII\\ regions.\\\\ While NGC\\,5408 shows a fairly regular \\HI\\ velocity field, which allows us to derive a rotation curve, IC\\,4662 reveals a rather twisted \\HI\\ velocity field, possibly caused by a recent merger event. We detect outflows with velocities between 20 and 60\\skms\\ in our \\Ha\\ spectra of both galaxies, sometimes with \\HI\\ counterparts of similar velocity. We suggest the existence of expanding superbubbles, especially in NGC\\,5408. This is also supported by the detection of \\emph{FWHMs} as high as 70\\skms\\ in \\Ha, which cannot be explained by thermal broadening alone. In case of NGC\\,5408, we compare our results with the escape velocity of the galaxy, which shows that the measured expansion velocities are in all cases too low to allow the gas to escape from the gravitational potential of NGC\\,5408. This result is consistent with studies of other dwarf galaxies. ",
        "introduction": "The interplay between massive stars and the interstellar medium (ISM) has a large effect on the formation and the evolution of galaxies, especially of dwarf galaxies.  Photoionisation is the most likely interaction process. However, stellar winds and supernovae (SNe) explosions of massive stars also contribute significantly to the energy input into the ISM. The radiative and mechanical feedback heat the gas and drive it outwards, sweeping up the ambient gas into a thin shell. A superbubble filled with hot gas evolves and begins to expand into the ISM. Due to Rayleigh-Taylor instabilities, the outer shell can rupture and the hot gas can vent out through so-called \\emph{chimneys} into the halo \\citep{Norman1989}.  Numerous ionised gas structures such as supergiant shells or filaments close to the galactic disc of dwarf galaxies, but also at kpc-distances from any place of current star formation, have been detected on deep \\Ha\\ images \\citep[e.g.,][]{Hunter1993,Bomans1997, Hunter1997, Lopez-Sanchez2008a}. Spectroscopic observations of the \\Ha\\ line revealed that most of the ionised gas structures expand from the disc into the halo of their host galaxies \\citep[e.g.,][]{Marlowe1995,Martin1998,Bomans2001,vanEymeren2007}. This leads to the question of what the fate of the gas is. One scenario is that the gas cools down in the halo and eventually falls back onto the galactic disc \\citep[outflow, \\emph{galactic fountain} scenario,][]{Shapiro1976}. However, it might also be possible that the gas escapes from the gravitational potential by becoming a freely flowing wind (galactic wind). The detection of large amounts of hot gas in the intergalactic medium (IGM) and the generally low metal content of dwarf galaxies support this scenario.  Hydrodynamic simulations modelling superbubble blow-out in dwarf galaxies \\citep{MacLow1999} predict that at least a part of the hot gas has enough kinetic energy to leave the gravitational potential of its host galaxy and to enrich the IGM. The relatively low escape velocities of dwarf galaxies should facilitate the removal of substantial amounts of gas \\citep{Larson1974}. According to \\citet{Ferrara2000} galactic winds occur in dwarf galaxies with gas masses up to $10^9$\\,\\Msun. However, the fate of the gas does not only depend on the mass of the host galaxy and therefore its gravitational potential, but it also strongly depends on the morphology of the ISM distribution. For a spherical distribution, the ISM seems to be more resistant to ejection \\citep{Silich2001}.  A recent detailed kinematic study of the neutral and ionised gas in the two irregular dwarf galaxies NGC\\,2366 and NGC\\,4861 revealed several outflows in each galaxy \\citep{vanEymeren2009a,vanEymeren2009c}. The measured expansion velocities were of the order of 20 to 50\\skms, which is, in comparison with the escape velocities of the host galaxies, too low to allow the gas to leave the gravitational potential. This result confirms earlier observations: no convincing case of a galactic wind in local dwarf galaxies has been reported so far \\citep{Bomans2005}. Note that galaxies like M\\,82 are not typical dwarf galaxies as they show strong starbursts, are more luminous and have a higher mass.  Altogether, this shows that despite all the detailed observational studies and simulations, we do not fully understand what is going on in these galaxies. Galactic winds are still thought to be necessary ingredients for their formation and evolution, but direct evidence seems to be difficult to get. In order to improve our knowledge about the processes happening in dwarf galaxies, we performed a multi-wavelength study of altogether four irregular dwarf galaxies \\citep{vanEymeren2008PhD}. The results for NGC\\,2366 and NGC\\,4861 have already been published \\citep[][see also above]{vanEymeren2009a,vanEymeren2009c}. We here concentrate on the two remaining galaxies IC\\,4662 and NGC\\,5408. We obtained \\HI\\ line data as well as optical images and spectroscopic data. Some basic properties of both galaxies are given in Table~\\ref{basics}.  IC\\,4662 (HIPASS J1747--64) is classified as a barred irregular galaxy of Magellanic type (IBm). Its distance of $D_{\\rm TRGB}$\\,=\\,2.44\\,Mpc was obtained by \\citet{Karachentsev2006} and hence makes it the nearest known representative of blue compact dwarfs \\citep{Karachentsev2006}. It seems to be a rather isolated galaxy, belonging to no known groups. \\citet{deVaucouleurs1975} describes IC\\,4662 as a foreground galaxy in the direction of the NGC\\,6300 group.  NGC\\,5408 (HIPASS J1403--41) is classified as an IB(s)m galaxy. It was first studied by \\citet{Bohuski1972} who found that its nucleus consists of several bright \\HII\\ regions and appears to be undergoing a violent burst of star formation. As the galaxy reveals an ultra-luminous X-ray source, NGC\\,5408\\,X-1, very close to the main \\HII\\ regions, it was a popular object to study over the last decade \\citep[e.g.,][]{Soria2004,Soria2006,Strohmayer2007,Lang2007,Kaaret2009}. This X-ray source has recently been argued to harbour an intermediate-mass black hole \\citep{Strohmayer2009}. The distance of NGC\\,5408 of $D_{\\rm TRGB}$\\,=\\,4.81\\,Mpc was obtained by \\citet{Karachentsev2002}. Its position on the sky puts NGC\\,5408 in the Centaurus A group. The closest known neighbour appears to be ESO\\,325-G?001 at a projected distance of 208\\arcmin.\\\\  This paper is organised as follows: in Sect.~2 we give an overview over the observations and the data reduction. In Sect.~3 we describe and compare the morphology of the neutral and ionised gas. Section~4 contains the kinematic analysis of both gas components including a search for expanding gas. The results are discussed in Sect.~5, which is followed by a summary in Sect.~6. \\begin{table*} \\centering \\caption{The basic properties of IC\\,4662 and NGC\\,5408.} \\label{basics} $$ \\begin{tabular}{lccccc} \\hline \\hline \\noalign{\\smallskip} Optical name         & IC\\,4662  & NGC\\,5408  & Ref. \\\\ HIPASS name          & HIPASS\\,J1747--64 & HIPASS\\,J1403--41 &              \\\\ \\hline \\noalign{\\smallskip} opt. centre:      & & & (1)\\\\ ~~$\\alpha$ (J2000)& 17$\\rm ^h$ 47$\\rm ^m$ 08.8$\\rm ^s$ & 14$\\rm ^h$ 03$\\rm ^m$ 20.9$\\rm ^s$ &\\\\ ~~$\\delta$ (J2000) & --64\\degr\\ 38\\arcmin\\ 30\\arcsec & --41\\degr\\ 22\\arcmin\\ 40\\arcsec & \\\\ distance [Mpc]       & 2.44          & 4.81        & (2),(3) \\\\ $v_{\\rm opt}$ [\\kms]& $336\\pm29$    & $537\\pm33$  & (1) \\\\ type                 & IBm           & IB(s)m         & (1) \\\\ opt. diameter [\\arcmin]    & $3.0 \\times 1.6$ & $2.6 \\times 1.6$ & (4) \\\\ opt. diameter [kpc]             & 2.1 $\\times$ 1.1 & 3.6 $\\times$ 2.2 \\\\ inclination [\\degr]         &  58 & 52 & (4) \\\\ position angle [\\degr]      & 104 & 62 & (4) \\\\ \\AB\\ [mag]           & 0.303 & 0.298 & (5) \\\\ $m_{\\rm B}$ [mag]    & $12.33\\pm0.09$ & $12.59\\pm0.09$ & (4) \\\\ B-V & 0.41 & 0.56\\\\ $M_B$ [mag]    & $-14.91\\pm0.1$  & $-16.12\\pm0.09$ & \\\\ \\LB\\ [10$^9$\\Lsun]   & $0.14\\pm0.01$    & $0.44\\pm0.03$ & \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} $v_{\\rm HI}$ [\\kms] & $302\\pm3$  & $506\\pm3$  & (6) \\\\ $v_{\\rm LG}$ [\\kms] & 153        & 314        & (6) \\\\ $w_{\\rm 50}$ [\\kms] &  86        &  62        & (6) \\\\ $w_{\\rm 20}$ [\\kms] & 133        & 112        & (6) \\\\ $F_{\\rm HI}$ [Jy\\skms]       & $130.0\\pm12.0$ & $61.5\\pm6.7$ & (6) \\\\ $M_{\\rm HI}$ [10$^8$\\Msun]    & $1.83\\pm0.17$  & $3.36\\pm0.37$ \\\\ \\noalign{\\smallskip} \\hline \\hline \\end{tabular} $$ \\flushleft\\footnotesize{ Note: The blue luminosity is calculated using a solar $B$-band magnitude of 5.48 mag. -- References: (1) \\citet{deVaucouleurs1991}, (2) \\citet{Karachentsev2006}, (3) \\citet{Karachentsev2002}, (4) \\citet{Lauberts1989} [ESO Uppsala], (5) \\citet{Schlegel1998}, (6) \\citet{Koribalski2004} [HIPASS BGC].} \\end{table*} ",
        "conclusions": "Our analysis has shown that both galaxies show differences in the kinematic behaviour of the neutral and ionised gas components. Often, this can be interpreted as evidence for outflows of the ionised gas. Sometimes, these outflows are detected in both \\Ha\\ and \\HI. In this section, we discuss our results and make some predictions about the fate of the gas. We will also assess our results with respect to other studies of dwarf galaxies. Finally, we will come back to the peculiar \\HII\\ region in the south of IC\\,4662. \\subsection{The outflows and the special case of IC\\,4662} \\label{Sectouflows} The outflow velocities vary from 20 up to 60\\skms\\ in both galaxies. In some cases we found an \\HI\\ counterpart for an expanding ionised gas structure, moving at about the same velocity. NGC\\,5408 shows little evidence for line-splitting in \\HI. The neutral gas distribution seems to be warped in the outer parts, which might indicate gas accretion, interaction with gas-rich companions or infall of \\HI\\ clouds \\citep[see ][and references therein]{Sancisi2008}. However, NGC\\,5408 appears to be quite isolated and no \\HI\\ clouds could be found in its near vicinity on our deep \\HI\\ maps. Furthermore, gas infall would result in a lopsided gas distribution, i.e., more gas on one side than on the other, which is definitely not seen in NGC\\,5408. Therefore, the outflow scenario is much more likely.  From the morphology and the kinematics, we suggest that NGC\\,5408 harbours at least one expanding superbubble. In comparison to IC\\,4662, this galaxy shows outflows almost everywhere across the disc.  For IC\\,4662 no XMM Newton observations are available and the ROSAT X-ray All-Sky Survey does not show a source that could be associated with this galaxy. Nevertheless, IC\\,4662 is a very complex system: the comparison of the neutral and ionised gas kinematics revealed that it is very difficult to take the \\HI\\ velocities as a reference value for the \\Ha\\ velocities as the \\HI\\ velocity field is quite distorted. Especially the discrepancy in slit\\,3 (see Fig.~\\ref{ic4662pv}, lower right panel) is probably due to the large-scale \\HI\\ velocity field and not due to an outflow of the ionised gas. The same might be true for slit\\,2 (between 30\\arcsec\\ and 70\\arcsec) where the \\HI\\ velocities again drop significantly. However, here, the ionised gas seems to follow that movement. If we just look at the velocities of the ionised gas and try to define a velocity gradient across each slit, a velocity offset of about 20\\skms\\ blue-shifted can be seen, which coincides with the north-eastern filamentary area. From the morphology alone expanding gas is expected, but because of the distorted \\HI\\ velocity field, and the lack of spatial coverage and spectral resolution of the long-slit spectra, it is difficult to quantify the outflow velocities.  As can be seen very nicely in NGC\\,5408 (Fig.~\\ref{n5408pv}, lower row), the kinematics of the ionised gas are dominated by stellar feedback causing offsets from the rotation velocity of the galaxy. This implies that the \\Ha\\ velocities can not be assumed to reflect the rotation of a galaxy; instead, we rely on the \\HI\\ velocities. However, in the case of IC\\,4662, the \\HI\\ velocity field is very distorted as well, which is why we are very careful about our detections.  Looking at the large-scale structure, IC\\,4662 shows line-splitting of the \\HI\\ line at many positions across the galaxy (Fig.~\\ref{IC4662profiles}). As already mentioned above (Sect.~\\ref{I4662HIkinematics}), the two components are often separated by about 70\\skms. Given the size of the optical galaxy in comparison to the distribution of the neutral gas, this is probably not caused by an interaction of stars and the ISM. As we see global line-splitting, the infall of single \\HI\\ clouds is also highly unlikely. Instead, the line-splitting is probably a consequence of a large scale distortion like, e.g., a merger process. As IC\\,4662 is a very isolated galaxy, we may rule out any kind of current tidal interaction with another object. We think that it is more likely that two dwarf galaxies collided in the past and are now in the process of merging. This hypothesis explains the distorted kinematics we observe in IC\\,4662, the two perpendicular systems as observed in the \\HI\\ distribution, and also the \\HI\\ tail we found in the eastern part of the galaxy (see Fig.~\\ref{IC4662mom}). There is increasing evidence that interactions and mergers of dwarf galaxies trigger star-formation activity in dwarf objects, but such features can only be detected when optical and radio observations are carefully compared \\citep[e.g.,][]{Lopez-Sanchez2009}. \\subsection{Large \\emph{FWHMs} in IC\\,4662 and NGC\\,5408} The \\Ha\\ line profiles of both galaxies show very large \\emph{FWHMs} (see Fig.~\\ref{fwhm}). In the case of IC\\,4662, the \\emph{FWHMs} of the low intensity profiles are at about 70\\skms. Even for high intensity profiles, this number does not decrease, which means that the high \\emph{FWHMs} are not an effect of low S/N. 70\\skms\\ are far too high to be explained by thermal broadening alone as this would require an electron temperature of 125000\\degr! Taking a reasonable value of the electron temperature of 10000\\,K gives a contribution of the thermal component of about 23\\skms. This still leaves us with 66\\skms\\ that have to be explained by other mechanisms. As shown in Sect.~\\ref{I4662longslit}, the residuals after subtracting a single Gaussian from the observed line profile usually indicate that the \\Ha\\ line is a superposition of at least two components. As the line has a symmetric shape, this may indicate expanding bubbles.  NGC\\,5408 also shows quite high \\emph{FWHMs}. Here, it becomes obvious that this result is also not an effect of the spectral resolution as the line widths remain equally broad when measured at higher spectral resolution (from the archival spectrum). In comparison to IC\\,4662, the \\emph{FWHMs} are more scattered, although they are preferably located in a band centred at 70\\skms, which is similar to the average value found for IC\\,4662. Additionally, the ionised gas observed in slit~1b reaches \\emph{FWHMs} of up to 100\\skms. These high values have mainly been found in the western \\HII\\ region, which is where the X-ray source is located.  Detailed studies of the velocity dispersion in \\HII\\ regions have been carried out by \\citet{Munoz-Tunon1996} and later by \\citet{Martinez-Delgado2007}. Their diagrams show very similar patterns compared to ours: a narrow horizontal band and inclined bands at higher velocity dispersion, which they interprete as coming from faint expanding bubbles and shells. Indeed, we found the largest \\emph{FWHMs} to coincide with areas of filamentary structure (IC\\,4662) or potential superbubbles (NGC\\,5408). \\subsection{Outflow or galactic wind?} As we could not derive a reasonable rotation curve from the distorted velocity field of IC\\,4662, we focus this analysis on NGC\\,5408.  Following the method described in \\citet{vanEymeren2009a}, we compared the expansion velocities with the escape velocity of NGC\\,5408 in order to make predictions about the fate of the outflowing gas structures. The escape velocity was estimated by using the cored pseudo-isothermal halo model, which represents the shape of the rotation curves best, especially in the inner few kpcs \\citep[e.g.,][]{vanEymeren2009b}, which is where we detected the outflows. Only the circular velocity and the maximum halo radius $r_{\\rm max}$ have to be known in order to calculate the escape velocity. The circular velocity was estimated from the plateau of the \\HI\\ rotation curve. We derived the escape velocity for two different radii. First, we took the outer radius of the \\HI\\ intensity distribution (9.5\\,kpc) as a lower limit for the halo size. As a second estimate we calculated the escape velocity using twice the \\HI\\ radius (19\\,kpc), which we consider to be a more physical assumption. The resulting curves are shown in Fig.~\\ref{ICNGCesc} (dotted and solid lines respectively). The observed rotation curve is indicated by small grey triangles; the outflows with their distances from the dynamic centre and their velocities compared to the \\HI\\ velocities are marked by large black triangles. A higher $r_{\\rm max}$ results in an increase of the escape velocity. As the radius of the dark matter halo $r_{\\rm max}$ is probably much higher than assumed, the escape velocities will be even higher than the ones we derived. The curve with $r_{\\rm max}=r_{\\rm HI,max}$ gives therefore a lower limit for the escape velocity values. As can be seen, the expansion velocities of the outflows, although partially quite significant, stay far below the escape velocity, even for the lower limit case. This is mainly due to the fact that they were found close to the dynamic centre where the escape velocities are high. \\begin{figure} \\centering \\includegraphics[width=.48\\textwidth]{vaneymeren16.pdf} \\caption[Outflow or galactic wind?]{Outflow or galactic wind? The expansion velocities of the outflowing gas structures in NGC\\,5408 are compared with the escape velocities calculated by using the pseudo-isothermal halo model. The escape velocity is calculated for an isothermal halo of $r_{\\rm max}\\rm=9.5\\,kpc$ (dotted line), which fits the size of the \\HI\\ distribution, and $r_{\\rm max}\\rm =19\\,kpc$ (solid line). The observed rotation curve is indicated by small grey triangles. The error bars represent receding and approaching side. The expanding gas structures are marked by black large triangles.} \\label{ICNGCesc} \\end{figure} \\subsection{Comparison with previous work} We now compare our results for IC\\,4662 and NGC\\,5408 with similar studies from the literature. Many dwarf galaxies reveal filaments and shell-like structures up to kpc-size on deep \\Ha\\ images \\citep[see, e.g., ][]{Hunter1993,Bomans1997,vanEymeren2007}. Spectral line data show \\Ha\\ line splitting, indicating outflowing gas, which is often associated with these filaments and shells \\citep[e.g.,][]{Martin1998}. Partly, the neutral gas is visibly affected by the expanding ionised gas as can, e.g., be seen in the starburst dwarf galaxy NGC\\,5253 \\citep[][]{Kobulnicky2008}. The neutral gas is pushed by the ionised gas, which leads to the observed expansion velocities. On the other hand, the surrounding neutral gas decelerates the expanding bubble so that it might eventually be stalled by the neutral gas envelope \\citep[][]{Kobulnicky2008}. Figure~\\ref{momhres} shows very nicely how the expanding gas can affect the neutral gas, as, e.g., NGC\\,5408 has a local \\HI\\ minimum in the centre of the potential superbubble. A very similar morphology has been found in NGC\\,5253 \\citep{Lopez-Sanchez2010}. NGC\\,2366 \\citep{vanEymeren2009a} and NGC\\,4861 \\citep{vanEymeren2009c} also show \\HI\\ minima coinciding with expanding ionised gas structures.  The expansion velocities of the galaxies in our sample (including NGC\\,2366 and NGC\\,4861) vary between 20\\skms\\ and 60\\skms. These values are in good agreement with previous studies: \\citet{Heckman1995} and \\citet{Westmoquette2008} found expansion velocities between 50 and 100\\skms\\ in NGC\\,1569; \\citet{Schwartz2004}, using absorption line studies, found the expansion velocities of a sample of dwarf galaxies to be about 30\\skms. An exception seems to be the Wolf-Rayet blue compact dwarf galaxy He\\,2-10, where expansion velocities up to 250\\skms\\ have been detected \\citep[][]{Mendez1999}. Note that more massive galaxies like M\\,82 show outflows with expansion velocities higher than 100\\skms\\ \\citep[][]{Schwartz2004,Strickland2007,Westmoquette2009}. In our sample, the highest velocities have been found in NGC\\,5408.  How high do the expansion velocities have to be to turn an outflow into a galactic wind? Coming back to our sample (including NGC\\,2366 and NGC\\,4861, but excluding IC\\,4662 whose velocity field is too distorted to derive a rotation curve and the associated escape velocity), it could be shown that the expansion velocities of the outflows only account for 30 to 50\\%\\ of the escape velocities. Note that these velocities were calculated under the assumption that the radius of the dark matter halo $r_{\\rm max}$ equals the \\HI\\ radius. As $r_{\\rm max}$ is probably much higher, the resulting escape velocities will be higher, which therefore further decreases the possibility of a galactic wind. This result is consistent with studies of other dwarf galaxies \\citep[e.g.,][]{Martin1998,Schwartz2004}. An exception is NGC\\,1569 which rotates very slowly with only 30\\skms \\citep{Heckman1995}, but harbours outflows with expansion velocities of up to 100\\skms\\ \\citep[][]{Heckman1995,Westmoquette2008}. Still, a galactic wind could not be confirmed, although the chances are much higher than in other dwarf galaxies. Another exception is He\\,2-10 \\citep{Mendez1999}, which however, as mentioned before, shows very fast expanding gas with velocities up to 250\\skms.  At first glance, all of these non-detections of a galactic wind look like a contradiction with theoretical models like, e.g., the superbubble model by \\citet{MacLow1999}. Note, however, that outflowing gas is mostly detected close to the dynamic centre, where the escape velocities are very high. Only a few kpcs away from the dynamic centre, the escape velocity drops significantly (see, e.g., Fig.~\\ref{ICNGCesc}), increasing the chance of a galactic wind. As the surface brightness goes with $n_e^2$ with $n_e$ being the electron density, a less dense gas will be much more difficult to detect. This implies that the outer regions of an outflow will easily escape the detection in emission line studies \\citep[e.g.,][]{Bomans2007}, which means that we consistently miss the fastest moving gas. A study of the interstellar NaI absorption line (which is as an absorption line not affected by the $n_e^2$ problem) does not help since NaI traces much cooler, denser gas, which is probably not as much accelerated as the diffuse HII.   Additionally, we have to take into account the dynamic masses of our sample galaxies, which are of the order of a few times 10$^9$\\Msun, i.e., close to the upper limit of the mass range of dwarf galaxies. As Fig.~2 in \\citet{MacLow1999} shows, at total masses of 10$^9$\\Msun, most of the gas is still bound to the disk when the energy input stops. This strong dependence of blow-out on the mass of the galaxy has already been noted by \\citet{DeYoung1994}. Furthermore, the central star clusters in, e.g., NGC\\,2366 \\citep{Drissen2001} or NGC\\,4861 \\citep[e.g.,][]{Barth1994} are only a few Myrs old. If we assume that the current star formation episode is the driver of the outflows, it makes a superbubble blowout even more unlikely.  It is also possible that the expansion velocity of the outflows will never get close to or beyond the escape velocity due to the dark matter halo that slows down the gas that was expelled from the disc \\citep[][]{Silich1998}. Their 2d calculations of multi-supernova remnants evolving in dwarf galaxies show that the total retention of the outflowing gas is very likely for galaxies with an ISM mass of about 10$^9$\\Msun.  Altogether, this means that our observations (including NGC\\,2366 and NGC\\,4861) are in good agreement with other observations and with the simulations. All sample galaxies harbour filaments and shells that could be shown to expand, but no evidence for a galactic wind could be found.  As the expansion velocities in IC\\,4662 are comparable to those detected in the other three galaxies, we suppose that we would find similar results for this galaxy. However, if this galaxy is a merging system, processes with much more impact than star formation activity probably dominate the kinematic behaviour. \\subsection{The southern \\HII\\ region in IC\\,4662} \\label{southHII} The southern \\HII\\ region in IC\\,4662 has been and still is a subject of great interest. It appears to be detached from the main body. Furthermore, \\citet{Hidalgo-Games2001} pointed out that it differs significantly from the main body in oxygen abundance. Therefore, the question arises whether it is an own system interacting with IC\\,4662 (e.g., H. Lee, private communication) or just an unusually placed OB association. Our deep \\Ha\\ image shows that this bright \\HII\\ region is at least connected to the main complex by a chain of small \\HII\\ regions and diffuse filamentary gas structures \\citep[see Sect.~\\ref{LVHISgenmorph} and ][]{Noeske2003}. Moreover, together with the main body this region is embedded into the highest column density \\HI. Additionally, its kinematics do not differ from the main body, but follow the rotation of the whole optical system. All these detections suggest that the southern \\HII\\ region is part of IC\\,4662.  Similar features have been found in the blue compact dwarf galaxy NGC\\,1140: a chain of \\HII\\ regions that emanates south of the main body, and a core which is kinematically decoupled from the rest of the galaxy \\citep[][]{Hunter1994,Westmoquette2010}. Therefore, this galaxy was suggested to be a merger remnant \\citep{Hunter1994}. A study of 20 starburst galaxies by \\citet{Lopez-Sanchez2010b} has shown that in fact almost all of them show signs of interaction or of merging with small companions. They conclude that interactions play an important role in the triggering mechanism of the strong star-formation activity observed in dwarf galaxies.  Taking into account all the observational peculiarities and this similarity to NGC\\,1140, we definitely need to discuss the detached \\HII\\ region in IC\\,4662 in the context of a possible merger. If IC\\,4662 indeed consists of two merging or merged systems, the difference in the chemical composition might indicate that the main \\HII\\ region complex belongs to one galaxy and the southern \\HII\\ region to the second galaxy.  Detailed multi-wavelength observations of IC\\,4662 (and other blue compact dwarf galaxies) are on the way \\citep[see also ][]{Lopez-Sanchez2009}."
    },
    "1004/1004.3341_arXiv.txt": {
        "abstract": "The integrated Sachs-Wolfe (ISW) effect is an important implication for dark energy. In this paper, we have calculated the power spectrum of the ISW effect in the time varying vacuum cosmological model, where the model parameter $\\beta=4.407$ is obtained by the observational constraint of the growth rate. It's found that the source of the ISW effect is not only affected by the different evolutions of the Hubble function $H(a)$ and the dimensionless matter density $\\Omega_m(a)$, but also by the different growth function $D_+(a)$, all of which are changed due to the presence of matter production term in the time varying vacuum model. However, the difference of the ISW effect in $\\Lambda(t)\\textmd{CDM}$ model and $\\Lambda \\textmd{CDM}$ model is lessened to a certain extent due to the integration from the time of last scattering to the present. It's implied that the observations of the galaxies with high redshift are required to distinguish the two models. ",
        "introduction": "Since 1998, the type Ia supernova (SNe Ia) observations \\cite{ref:Riess98,ref:Perlmuter99} have showed that the expansion of our universe is speeding up rather than slowing down. During these years from that time, many additional observational results, including current Cosmic Microwave Background (CMB) anisotropy measurement from Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{ref:Spergel03,ref:Spergel06}, and the data of the Large Scale Structure (LSS) \\cite{ref:Tegmark1,ref:Tegmark2}, also strongly support this suggestion. In order to understand the mechanism of the accelerating expansion of the universe, the common recognition has been received that there exists an exotic energy component with negative pressure, called dark energy (DE), whose density accounts for two-thirds of the total energy density in the universe according to the current observations. It has been the most important thing to make an effort to probe into what the nature of DE is in the modern cosmology. The simplest and naturalest candidate is the cosmological constant ${\\Lambda}$ (vacuum energy) \\cite{ref:LCDM1,ref:LCDM2}, with the equation of state (EOS) $w=\\frac{P}{\\rho}=-1$. However, it's pity that it confronts with two difficulties: the fine-tuning problem and the cosmic coincidence problem. An alternative offer is the dynamical DE to alleviate the difficulties above, such as quintessence \\cite{ref:quintessence}, phantom \\cite{ref:phantom}, quintom \\cite{ref:quintom}, generalized Chaplygin gas (GCG) model \\cite{ref:GCG}, Holographic DE \\cite{ref:holo}, and so on.  At present, although in the presence of numerous dark energy models, the concordance model still fits best well with the current combined observations \\cite{ref:WMAP5}. So it's necessary to review the cosmic constant problems: The fine-tuning problem is the vacuum energy density predicted by quantum field theory (QFC) is $10^{74}GeV^4$, which is about $10^{121}$ orders of magnitude larger than the observed value $10^{-47}GeV^4$, so there must be so-called \"fine-tuning\" while looking for an balance between two numbers which exits so much discrepancy. The cosmic coincidence problem says: Since the energy densities of vacuum energy and dark matter turn on different evolutions during the expansion history of the universe, why are they nearly equal today and why it happens just now? To get this coincidence, it appears that their ratio must be set to a specific, infinitesimal value in the very early universe. Therefore, in nature the second difficulty is also \"fine-tuning\" problem. In the final analysis, both of the two difficulties are related to the energy density of vacuum energy. On the basis of this consideration, a phenomenological and natural attempt at alleviating the above problems is allowing vacuum energy $\\Lambda$ to vary \\cite{ref:LCDM2, ref:LtCDM}. In $\\Lambda(t)\\textmd{CDM}$ model, we keep unvaried for the EOS of the vacuum energy, equaling to $-1$, and have an time-evolving energy density. Although the weak point in this ideology is the unknown functional form of the $\\Lambda(t)$ parameter, we can deal with the $\\Lambda(t)$ parameter by different phenomenologically parameterized form, referring to \\cite{ref:LCDM2,ref:LtCDMparam1}. Meanwhile, the underly explanation for $\\Lambda(t)$ parameter has also been explored in the framework of QFT using the renormalization group \\cite{ref:RG}.  The ISW effect \\cite{ref:ISW1} arises from a time-dependent gravitational potential at late time and leads to CMB temperature anisotropies. From the physical picture point of view, the frequency of the photon released by CMB has been redshifted or blueshifted, which is caused by the decaying gravitational potential, while propagating along the line of light-like from the last scattering surface to now. The ISW effect is a significant implication for dark energy. Since the gravitational potential was an constant during matter-dominated. At the late time, the gravitational potential is time-dependent when dark energy begin to dominate in the universe. So, the detection of ISW effect will provide great support and more information for dark energy. Now, the ISW effect is detected by the cross-correlation between CMB and LSS observations \\cite{ref:LSS1,ref:densitycontrast,ref:LSS3,ref:LSS4,ref:LSS5}. In this way, we can separate CMB temperature anisotropies caused by ISW effect from the primary temperature anisotropies.  In this paper, we consider the time varying vacuum model, where a parameterized form of $\\Lambda(t)$ is used. The densities of dark matter and vacuum energy don't conserve independently owing to the time varying vacuum term, which plays the role of the interaction between dark matter and dark energy. The matter production term leads to the different evolutions of the Hubble parameter $H(a)$ and the dimensionless matter density $\\Omega_m(a)$, and the different structure formation, comparing with the concordance model. Meanwhile, there being a relation between these three quantities and the source term of the ISW effect, we explore the impact from the change of these three quantities on ISW effect in the $\\Lambda(t)\\textmd{CDM}$ model in theory.  The paper is organized as follows. In next section, we briefly review of the cosmological evolution and structure growth in time varying vacuum model. In section III, we focus on the theory and detection of ISW effect and its power spectrum. The last section is the conclusion. ",
        "conclusions": "In summary, we have calculated the power spectrum of the ISW effect in the time varying vacuum energy model. It's shown that the source of the ISW effect is not only affected by the different evolutions of the Hubble function $H(a)$ and the dimensionless matter density $\\Omega_m(a)$, but also by the different growth function $D_+(a)$, all of which are changed due to the presence of matter production term in the time varying vacuum model. In FIG. \\ref{fig:growthrate} and FIG. \\ref{fig:D}, it's seen that the difference between the growth function in $\\Lambda(t)\\textmd{CDM}$ model and that in $\\Lambda \\textmd{CDM}$ model is not minor. What's more, the left panel in FIG. \\ref{fig:D} shows there is a crosspoint at $a=0.4$ for the evolutions of the growth function in two models. The growth function in $\\Lambda \\textmd{CDM}$ model is greater than that in $\\Lambda(t) \\textmd{CDM}$ model before $a=0.4$. When $a>0.4$, the growth function in $\\Lambda \\textmd{CDM}$ model is smaller than that in $\\Lambda(t) \\textmd{CDM}$ model. However, since the ISW effect is calculated by integrating time-dependent gravitational potential from the time of last scattering to the present, the difference of the ISW effect caused by the difference of the growth function in the two models within the scale factor ranges $[0.4,1]$ can partly counteract that caused by the discrepancy in their growth functions at $a<0.4$ though the integration. Therefore, the difference of the ISW effect in $\\Lambda(t)\\textmd{CDM}$ model and $\\Lambda \\textmd{CDM}$ model is lessened to a certain extent due to the integrated effect. As shown in FIG. \\ref{fig:Tg}, there is not much discrepancy between the availably observed quantities $C_{Tg}$ in the two models for the SDSS galaxies and the sample from NVSS. For the SDSS quasars with high redshift, the amplitudes of the ISW-cross power spectrum in $\\Lambda(t)\\textmd{CDM}$ model and $\\Lambda \\textmd{CDM}$ model obviously decrease. However, the relative discrepancy between ISW-cross power spectrums in $\\Lambda(t)\\textmd{CDM}$ model and $\\Lambda \\textmd{CDM}$ model for the SDSS quasars is larger than those for the SDSS galaxies and the sample from NVSS. Therefore, at present the ISW effect is weaker than the growth history in distinguishing $\\Lambda \\textmd{CDM}$ model and $\\Lambda(t)\\textmd{CDM}$ model. However, it's worth expecting the observational data of the ISW-cross power spectrum for the observed galaxies with high redshift in the future."
    },
    "1004/1004.0612_arXiv.txt": {
        "abstract": "{In the last years we have witnessed a dramatic change in the research infrastructures: Advances in communication networks, computational resources and data storage devices are fostering new and more efficient science. In this new scenario, the Virtual Observatory (VO) is the framework where a new methodology for astronomical research is being built. This poster shows the natural connection between Asteroseismology and VO. We describe the current status of a project developed by the Spanish Virtual Observatory in which, for the first time, asteroseismic models together with visualization tools for asteroseismology are managed within VO.} ",
        "introduction": "}  The twentieth century has provided an important qualitative and quantitative leap in the development of science and technology. The perfect symbiosis between the two has o\\-pen\\-ed new horizons in the knowledge of nature:  from the study of the fundamental components of matter, through life sciences, to better understand of the origin and evolution of the Universe and its components, like galaxies, stars, planets, to name a few.  In the last decades we have witnessed a significant pro\\-gress of Stellar Physics, largely due to the development of its great laboratory: the stellar seismology (also known as asteroseismology). Today, several space missions have been designed to study the interior of stars and search for extra-solar pla\\-nets like MOST (Matthews et al.~1998), CoRoT (Baglin et al.~2003), or Kepler (Gilliland et al.~2010) and some others are in preparation like PLATO (Catala 2009). Besides this, se\\-ve\\-ral precision instruments are being used for the first time for asteroseismology (CRIRES, UVES, FOCES, FEROS, HARPS) and others are currently in phase of development: NAHUAL (Mart\\'{\\i}n et al.~2005), GIANO (Oliva et al.~2006), CARMENES (Quirrenbach et al.~2009), etc.  Moreover, one of the great revolutions in data trans\\-mi\\-ssion is undoubtedly the emergence of the Internet and its boundless possibilities. Day after day, new Internet-based technologies such as Web, Web 2.0, social networks, etc. come up with new useful tools. In astrophysics, this breakthrough has its own name: the Virtual Observatory\\footnote{http://ivoa.net/} (VO). Today, this tool allows us to handle hundreds of terabytes in a single click, and access to thousands of astronomy da\\-ta\\-ba\\-ses interlinked through a single web portal.  We present a new generation tool for stellar physics: the fusion of Virtual Observatory technologies with as\\-te\\-ro\\-seis\\-mo\\-lo\\-gy. The result is a tool conceived to easily handle observations and models in the framework of the large research programs like the above mentioned space missions. We seek to offer to the community a tool from which users have access to very different models without building interfaces, easily find models representative of the studied stars, and even compare online global and shell physical variables  of a selection of models. ",
        "conclusions": ""
    },
    "1004/1004.2617_arXiv.txt": {
        "abstract": "{  We present the results of H$\\alpha$ monitoring of the BL Lac object OJ~287 with the VLT during seven epochs in 2005-08. We were able to detect five previously undetected narrow emission lines, $\\lambda\\lambda$6548, 6583[NII], $\\lambda6563$H$\\alpha$ and $\\lambda\\lambda$6716, 6731 [SII] during at least one of the epochs and a broad H$\\alpha$ feature during two epochs.  The broad H$\\alpha$ luminosity was a factor $\\sim$ 10 lower in 2005-08 than in 1984 when the line was previously detected and a factor $\\sim$ 10 lower than what is observed in quasars and Seyfert galaxies at the same redshift. The data are consistent with no change in the position or luminostity of the H$\\alpha$ line in 2005-08. The width of the H$\\alpha$ line was $4200 \\pm 500$ km s$^{-1}$, consistent with the width in 1984. } ",
        "introduction": "The BL Lac object \\object{OJ 287} (z = 0.306) has received some attention after the discovery of the recurring outbursts every $\\sim$ 12 years in the optical data over the past century \\citep{1988ApJ...325..628S}.  This observation has led many authors to consider a binary black hole model for OJ~287 \\citep{1988ApJ...325..628S,1996ApJ...460..207L,1997ApJ...478..527K,1997ApJ...484..180S,1998MNRAS.293L..13V,2000ApJ...531..744V,2006ApJ...646...36V,2007ApJ...659.1074V}. The exact mechanism of the outburst differs from one model to another. Some models attribute the flux increase to changes of the direction of the relativistic jet \\citep[e.g.][]{1997ApJ...478..527K}, other models to the impact of the secondary black hole onto the accretion disk of the primary followed by tidal effects and increased luminosity of the jet \\citep[e.g.][]{1996ApJ...460..207L}. The most detailed model so far \\citep{2007ApJ...659.1074V} suggests a very high mass for the primary black hole $M = 1.8 \\times 10^{10}\\ M_{\\odot}$ and a very eccentric orbit for the secondary ($e = 0.66$). To explain the aperiodicity of the outbursts the precession of the orbit has to be very strong, $\\Delta \\phi = 40^{\\circ}$ per orbit, in this model.  Testing these models was so far concentrated mainly on the timing of the outbursts with broadband photometric observations \\citep[e.g.][]{2009ApJ...698..781V} with other aspects receiving less attention. One interesting constraint to the models could be obtained by monitoring the emission line variability during the outbursts. The luminosity of the emission lines could change in response to varying continuum radiation or, if the size of the secondary black hole orbit and the characteristic broad line region (BLR) size are of the same order, the secondary black hole could affect the velocity field of the BLR.  Furthermore, if the BLR velocity dispersion was measured, it could be used to constrain the mass of the primary black hole in OJ~287.  However, there are no published results of spectroscopic monitoring of OJ~287 in the literature.  Thus we started a project in 2005 at the VLT to monitor the broad line region of OJ~287 during the expected outbursts in 2005-08 when the two black holes were expected to be at closest approach by various models.  The main aim of this project was to see if any changes could be observed in the BLR of OJ~287 during the expected closest approach of the binary.  This kind of monitoring project is quite challenging because the emission lines are very weak in BL Lacs by definition.  There are only two reported detections of broad lines in OJ~287 in the literature. \\cite{1989A&AS...80..103S} detected the H$\\beta$ line with a FWHM of 56 \\AA\\ with an equivalent width (EW) of 1.1 \\AA. \\cite{1985PASP...97.1158S} also detected a weak H$\\beta$, but the most notable feature in their spectrum is the prominent broad (FWHM $\\sim$ 120 \\AA) H$\\alpha$ line. It is impossible to know if these detections represent typical levels or very high states of the lines. Given the weakness of emission lines in BL Lacs in general, the latter alternative seems more probable as non-detections are not always reported. Based on this past information the H$\\alpha$ line was chosen for the monitoring because it can be easily detected with modern spectrographs at sufficiently large telescopes if the luminosity of the line is close to what was observed by \\cite{1985PASP...97.1158S}. The results of this monitoring are reported in this paper.  Throughout this paper we use the cosmology $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M}$ = 0.3 and $\\Omega_{\\Lambda}$ = 0.7. ",
        "conclusions": "The most notable result of our monitoring is the dramatic decrease in broad H$\\alpha$ luminosity between 1984 and 2005-08. In Fig. \\ref{kontiviiva} we show the continuum and broad H$\\alpha$ line flux of OJ~287 in 1980-2010. In the upper panel three outburst seasons of OJ~287 can be seen: 1982-84, 1994-95 and 2005-08. The high H$\\alpha$ flux appeared just after the strong outbursts in 1982-83, 290 days after the second peak.  Unfortunately, the sampling is too poor to draw any conclusions on the possible connection between continuum and BLR luminosity in OJ~287. Such a connection would not even be necessarily expected because in BL Lac objects the vast majority of the continuum arises from the jet and is highly beamed by relativistic effects due to the small angle between the line of sight and the jet axis. Thus the observed continuum variations are not necessarily connected to the changes of the continuum source illuminating the H$\\alpha$ emitting clouds, most likely the inner parts of the accretion disk. It is nevertheless an intriguing observation that the high line luminosity in 1984 was observed right after a luminous continuum outburst. Given the weakness of the broad H$\\alpha$ line we are unable to study a possible connection between BLR velocity field and the suggested periastron of the secondary black hole in 2005-07 \\citep{2007ApJ...659.1074V}.  \\begin{figure} \\centering \\epsfig{file=14198fg5.eps,width=9cm} \\caption{\\label{kontiviiva} Continuum (upper panel) and H$\\alpha$ (lower panel) flux of OJ~287 in 1980-2010. In the lower panel crosses with arrows mark 3$\\sigma$ upper limits and solid circles detections. The error bar of the leftmost point in the lower panel has been arbitrarily set to 10\\%.} \\end{figure}  The average line luminosities in 2005-08 were $\\log(L/{\\rm erg\\ s}^{-1})$ = 41.8, 41.1, 41.2 and 40.8 for broad H$\\alpha$, narrow H$\\alpha$, $\\lambda$6583[NII] and $\\lambda$6548[NII], respectively. The H$\\alpha$ luminosities of Seyfert galaxies and quasars are generally much higher as shown by he comparisons in Table \\ref{llumtable} and Fig. \\ref{lumvertailu}. The quasar data were obtained from \\cite{2007AJ....134..102S}, who give data of 77~429 quasars from the Fifth Data Release of the Sloan Digital Sky Survey \\citep{2007ApJS..172..634A}. We selected quasars at the redshift interval z = $0.30 - 0.31$ (147 quasars) and obtained the H$\\alpha$ line fluxes and widths from the SpecLine tables in the SDSS archive. Figure \\ref{lumvertailu} shows only quasars with FWHM $>$ 1000 km s$^{-1}$ (101 quasars). In addition we plot data for 17 z = 0.36 Seyfert galaxies in \\cite{2008ApJ...673..703M}.  As can be seen from Fig. \\ref{lumvertailu} the H$\\alpha$ luminosity of OJ~287 was lower than in typical quasars and Seyfert galaxies by a factor of $\\sim$ 10 in 2005-08. In December 1984, however, the H$\\alpha$ luminosity was comparable to that of quasars and Seyferts.  We finally note that the luminosities of the two narrow lines $\\lambda$6548 and $\\lambda$6583 [NII], however, are comparable to the narrow-line luminosities of the quasars in \\cite{2007AJ....134..102S}.  \\begin{table} \\caption{\\label{llumtable} Comparison of line luminosities of OJ~287 to those of quasars and Seyfert galaxies.} \\centering \\begin{tabular}{llll} \\hline & \\multicolumn{3}{c}{Log L}\\\\ Line & OJ~287 & Quasars & Seyferts\\\\ \\hline H$\\alpha$ broad & 41.8 & 42.9 $\\pm$ 0.4 & 43.1 $\\pm$ 0.3\\\\ $\\lambda$6548[NII] & 40.8 & 41.3 $\\pm$ 0.8 & \\ldots \\\\ $\\lambda$6583[NII] & 41.2 & 41.4 $\\pm$ 0.7 & \\ldots \\\\ \\hline \\end{tabular} \\end{table}  \\begin{figure} \\centering \\epsfig{file=14198fg6.eps,width=9cm} \\caption{\\label{lumvertailu} Comparison of broad H$\\alpha$ luminosities of OJ~287 (open circles connected by a line), z = 0.30-0.31 quasars from \\citet[][filled circles]{2007AJ....134..102S} and z = 0.36 Seyfert galaxies from \\citet[][open triangles]{2008ApJ...673..703M}.} \\end{figure}"
    },
    "1004/1004.5171_arXiv.txt": {
        "abstract": "{ We analyze large scale mapping observations of the molecular lines in the \\coTwo, \\coThr, \\TcoTwo, and  \\TcoThr transition emissions toward the Cepheus B molecular cloud with the KOSMA 3m-telescope. The integrated intensity map of the \\coTwo transition has shown a structure with a compact core and a compact ridge extended in the north-west of the core. The cloud is surrounded by an optically bright rim, where the radiation-driven implosion (RDI) may greatly change the gas properties. The intensities of the \\NcoThr transition are higher than those of the \\NcoTwo transition along the rim area. We find characteristic RDI structure in position-velocity diagrams. Non-LTE Large velocity gradient (LVG) model analysis shows that the density and temperature at the edge are higher than that in the center. Our results provide evidences that RDI is taking place in Cepheus B molecular cloud. ",
        "introduction": "% \\label{sect:intro}  The association Cepheus OB3 is one of the youngest groups of early type stars associated with the HII region S155 \\citep{1959Sharpless} at a distance of 0.7 kpc \\citep{1964Blaauw}. The observations in the \\coOne line have shown a few bright components named Cepheus A, B, C, D, E, F (Sargent 1977, 1979). Of the six components, the hottest CO feature and abundant dust emission were detected in the Cepheus B, which was surrounded by an arc-like optical HII region S155. The molecular cloud/HII region complex with such a morphology is called bright rimmed cloud (BRC), where star formation activity can be triggered by the radiation-driven implosion (RDI) mechanism.  The radio continuum observations \\citep{1978Felli} by Westerbork Synthesis Radio Telescope (WSRT) confirmed that the UV radiation from a bright O7 star created an arc-shaped ionization front which surrounded the Cepheus B molecular cloud. The observations of the H$_{2}$CO and recombination lines \\citep{1981Panagia} have suggested that the ionized material was flowing away at about 11 \\kms, and the ionization front was moving into the molecular cloud at a velocity of about 2 \\kms. \\coThr and \\TcoTwo observations together with the far-infrared analysis showed that the cloud was externally heated and the edge of the cloud was compressed by the expansion of S155 \\citep{1992Minchin}.   \\citet{2000Beuther} observed the Cepheus B area at a large scale in the $J$=3-2 and 2-1 transitions of $^{12}$CO, $^{13}$CO and C$^{18}$O with the KOSMA 3m-telescope. Based on the volume densities derived from PDR model, they found that Cepheus B was highly clumped and the volume of these clumps fills only 2-4\\% of the whole cloud. \\citet{2006Mookerjea} confirmed this with further observations of [C\\,I] $^3$P$_2$ - $^3$P$_1$ and \\coFou. Based on the studies of protoplanetary disks, \\citet{2009Getman} found a spatial-temporal gradient of young stars from the molecular core toward HD217086, which identified HD217086 as the primary ionizing source of the cloud.   Although a lot of efforts have been made to characterize the physical conditions of the Cepheus B, the mechanism of the star formation in this region is still not quite clear. We present analysis of a larger scale observation of multiple CO lines, and discuss the RDI as the possible trigger of the star formation in this region. ",
        "conclusions": "\\label{sect:discussion}    As seen in Figure \\ref{Fig2}, the observed molecular cloud is surrounded by an ionized bright rim facing the exciting source with a hot core close to the rim and a cometary shaped tail on the side opposite to the rim. This scenario matches well with BRCs which is modeled as externally illuminated, photo-evaporated and ablated into elongated head-tail morphologies by UV radiation of OB stars \\citep{1983Reipurth}. The isolation and simple morphology of BRCs make them ideal environments to investigate the RDI \\citep{1989Bertoldi} mode of star formation, where pressure from the ionization shock front at the surface propagates through a cloud and overcomes the magnetic, turbulent and thermal pressure trigger cloud collapse, thereby triggering local star formation.    In our case, the spatial shift of v$_{\\rm LSR}$ shows strong connections with the RDI mechanism. The PV diagrams in Figure \\ref{Fig4} shows a velocity gradient from the head (north-west) of the cloud to the tail (south-east). The blueward elongation toward the head by about 3 \\kms from the center indicates that the head part is expanding towards us, while the redward elongation toward the tail by about 4 \\kms indicates the tail part is flowing away to the opposite direction. The velocity gradients of the head and tail matches with the RDI model prediction of \\citet{1994Lefloch}, where the red-shift of the tail is caused by the acceleration of the gas across the tail, and the blue-shift of the head is because the compressed gas in the head moves faster than the gas immediately behind it. Such PV diagrams are reported by various BRC and cometary globule studies \\citep{1995Lefloch, 1997Sugitani, 1997White, 2002Bachiller} which are explained by the RDI mechanism, and so are taken as a characteristic structure of the RDI.    The reversed line ratios across the bright rim are another interesting feature. The spatial distribution of the reversed line ratios suggest that they are physically connected with the RDI along the rim area. Higher transitions are more readily excited by higher temperatures or higher densities (Petitpas, 1998), either of which can be a result of RDI. The highly luminous ionizing shock front from the RDI interacts with the molecular cloud, pumping more particles from a lower level to a higher one, which makes the $J$=3-2 emission stronger along the rim area.   In order to investigate RDI effects and its physical characteristics quantitatively, we made further analysis with a large velocity gradient \\citep[LVG,][]{1974Goldreich, 2008Qin} radiative transfer model in the MIRIAD software package. In this model, the molecular cloud is considered spherically uniform and the non-LTE excitation properties are calculated based on the kinetic temperature (T$_{\\rm kin}$), the column density per unit velocity interval (N$_{\\rm CO}$/$\\Delta$V), and the H$_{2}$ density (n$_{\\rm H_{2}}$). In this case, we use the line intensity of \\coTwo  and the line ratio of \\coThr/\\coTwo to determine the volume density n$_{\\rm H_{2}}$ and kinetic temperature T$_{\\rm kin}$ at the two selected positions A and B, assuming the column density in a range of $1\\times10^{14}$ to $1\\times10^{18} cm^{-2}$, and $\\Delta$V=1 \\kms. Representative solutions that can be fitted to the observations are listed in Table \\ref{Table 2}. A typical case is shown in Figure \\ref{Fig5}.      \\begin{figure} \\includegraphics[width=75mm, angle=0]{ms468fig5a.eps} \\includegraphics[width=75mm, angle=0]{ms468fig5b.eps} \\caption{LVG model analysis at position A (left panel, $N_{CO}=4.6\\times10^{16} cm^{-2}$) and B (right panel, $N_{CO}=1.2\\times10^{16} cm^{-2}$) using the \\coTwo and \\coThr line. The horizontal and vertical axes stand for volume density n(H$_{2}$) and kinetic temperature T$_{\\rm kin}$ in log scale, respectively. Red curves are line intensity of \\coTwo and blue curves are the line ratios of \\coThr/\\coTwo. The results of T$_{\\rm kin}$ and n(H$_{2}$) are taken from the cross points of certain red and blue curves. The error bars indicate the range of the possible solutions listed in Table \\ref{Table 2}. } \\label{Fig5} \\end{figure}  \\begin{table}[h!!!]  \\small \\centering   \\begin{minipage}[]{120mm} \\caption[]{ Solutions of LVG analysis } \\label{Table 2}\\end{minipage}  \\tabcolsep 6mm \\begin{tabular}{clcl} \\hline\\noalign{\\smallskip} Position    &N$_{\\rm CO}$ (10$^{16}$ cm$^{-2}$)     &n(H$_{2}$) (10$^{4}$ cm$^{-3}$)    &T$_{\\rm kin}$ (K)                 \\\\ \\hline\\noalign{\\smallskip} A  &3.2  &1.6    &40   \\\\ &6.8  &3.2    &31   \\\\ &10  &4.0    &31   \\\\ &22  &2.0    &30   \\\\ &46  &1.0    &30   \\\\ &68  &0.63    &30   \\\\ &100   &0.20    &32   \\\\ B  &1.1  &4.0    &71   \\\\ &1.2  &8.9    &50   \\\\ &1.3  &50    &45   \\\\  \\noalign{\\smallskip}\\hline \\end{tabular} \\end{table}     For position A, various column densities are possible to explain the observation results self-consistently covering a range from $3\\times10^{16}$ to $1\\times10^{18} cm^{-2}$, with corresponding volume densities ranging from $0.2\\times10^4$ to $4.0\\times10^4 cm^{-3}$. The average value of the possible volume densities is $2.0\\times10^{4} cm^{-3}$. For position B, the few self-consistent solutions require the column densities in a tiny range of $1.1\\times10^{16}$  to $1.3\\times10^{16} cm^{-2}$. The average density derived is $1.5\\times10^{5} cm^{-3}$, nearly one order of magnitude higher than that at position A. The average kinematic temperature of B (61 K) is also twice as high as that of A (32 K). Even in the cases that both positions have similar densities (e.g. $n(H_2)=4.0\\times10^4 cm^{-3}$), the kinematic temperature required at position B is still 2.3 times than that at position A. Therefore we can conclude that the reversed line ratios along the bright rim are produced by either high densities, or high temperatures, and likely both. Both possibilities can be explained by radiation and pressure of the ongoing RDI, where ionizing shock wave of the RDI propagates into the molecular cloud, compressing it at the edge, and thus makes the density and temperature of the edge even higher than the center.  In comparison with the LVG analysis results, we have also calculated the column densities assuming LTE conditions. Taking the upper energy level temperature of 16.6 K \\citep{20081Qin}, the column density of CO is given by \\citep{1991Garden}: \\begin{equation} N_{CO}=1.08\\times10^{13}\\frac{T_{ex}}{exp(-16.6/T_{ex})} \\int T_{mb}dv \\,(cm^{-2}) \\label{Eq1} \\end{equation} and \\begin{equation} T_{ex}=\\frac{hv}{k}(ln\\{1+\\frac{hv}{k}[\\frac{T_{mb}}{f}+\\frac{hv}{k}(exp(\\frac{hv}{kT_{bg}})-1)^{-1}]^{-1}\\})^{-1} \\,(K) \\label{Eq2} \\end{equation} where the beam filling factor $f$ is assumed to be 1, and the cosmic background temperature $T_{bg}$ is taken as 2.732 K. With the \\coTwo line properties listed in Table \\ref{Table 1}, we derive a CO column density of $4.0\\times10^{16} cm^{-2}$ for position A, and $1.9\\times10^{16} cm^{-2}$ for position B. Because \\coTwo is always considered optically thick, the column density derived here should be a lower limit. Compared with the column densities from LVG analysis, we find that LVG results at position A roughly agree with the LTE calculation, but for position B, all the LVG solutions are lower than the LTE lower limit. This result agrees with the scenario that at position B the LVG calculation requires less column density than LTE to explain the observing results because it takes into account enhanced photon escape probability due to greater velocity gradient at the bright rim, where the nearby HII region agitates the molecular cloud with strong FUV radiation pressure and heating, exhibiting a case of RDI.      We present analysis of large scale mapping observations of the molecular lines \\coTwo \\coThr \\TcoTwo and  \\TcoThr transition emissions toward the Cepheus B molecular cloud. The target lies at the edge of HII region S155, and is surrounded by an optically bright rim. The CO $J$=3-2 emission is stronger than $J$=2-1 emission along the bright rim. PV diagrams of this region show a velocity gradient from the head of the cloud to the tail, which is characteristic of RDI. According to the LVG analysis, the density and temperature at the edge of the cloud are higher than that at the center. All of our results provide evidences that RDI is taking place in Cepheus B molecular cloud.   \\normalem"
    },
    "1004/1004.5347_arXiv.txt": {
        "abstract": "Primordial black hole (PBH) abundance limits constrain the primordial power spectrum, and hence models of inflation, on scales far smaller than those probed by cosmological observations.  Single field inflation models which are compatible with all cosmological data can have large enough perturbations on small scales to overproduce PBHs, and hence be excluded. The standard formulae for the amplitude of perturbations do not hold for modes that exit the horizon close to the end of inflation however.  We use a modified flow analysis to identify models of inflation where the amplitude of perturbations on small scales is large. For these models we then carry out a numerical evolution of the perturbations and use the PBH constraints on the power spectrum to eliminate models which overproduce PBHs. Significant PBH formation can occur in models in which inflation can continue indefinitely and is ended via a secondary mechanism. We demonstrate that PBHs constrain these types of inflation models and show that a numerical evaluation of the power spectrum decreases  the number of otherwise viable models of inflation. ",
        "introduction": "Primordial black holes (PBHs) can form in the early Universe via the collapse of large density perturbations~\\cite{Carr:1974nx,Carr:1975qj}. There are tight constraints on the abundance of PBHs formed from their present day gravitational effects and the consequences of their evaporation. These limits can be used to constrain the power spectrum of the primordial density, or curvature, perturbations.  The PBH constraints on the curvature power spectrum are fairly weak, being many orders of magnitude larger than the measurements on cosmological scales. They do, however, apply over a very wide range of scales  and therefore provide a useful constraint on models of inflation.  Peiris and Easther~\\cite{Peiris:2008be} have shown that there are single field inflation models, which are compatible with all cosmological observations, for which the perturbation amplitude on small scales is large enough to produce a significant density of PBHs.  For scales which exit the horizon close to the end of inflation the standard (Stewart-Lyth~\\cite{Stewart:1993bc}) formulae for the amplitude of perturbations do not hold. Leach and Liddle~\\cite{Leach:2000yw} carried out a numerical calculation of the evolution of perturbations for a quadratic inflationary potential. They found that the perturbations on scales which exit the horizon close to the end of inflation were roughly an order of magnitude larger than predicted by the Stewart-Lyth formula (see also Ref.~\\cite{Bugaev:2008bi}). Therefore to fully exploit the power of PBH constraints on inflation models, a numerical calculation of the amplitude of perturbations on small scales is required. It has recently been shown~\\cite{Lyth:2005ze,Zaballa:2009xb} that PBHs can also form on scales which never leave the horizon.  We do not consider this possibility here.  In this paper we use a modified flow analysis to identify inflation models where the perturbations at the end of inflation may be large enough for primordial black holes to be overproduced. For these models we carry out a numerical evolution of the primordial perturbations and use the PBH constraints on the power spectrum to eliminate models which overproduce PBHs.  We describe the modified flow analysis in Sec.~\\ref{flow} and the evolution of perturbations and the calculation of the power spectrum in Sec.~\\ref{pert}. We apply the primordial black hole abundance constraints and present our results in Sec.~\\ref{results} and conclude with discussion in Sec.~\\ref{discuss}. ",
        "conclusions": "\\label{discuss}  We have applied constraints on the primordial power spectrum from the overproduction of primordial black holes to inflation models generated by a modified flow algorithm. The amplitude of inflationary perturbations is usually calculated using the Stewart-Lyth~\\cite{Stewart:1993bc} expression, however for scales which exit the horizon close to the end of inflation the assumptions underlying this expression are violated. A numerical calculation is therefore required, and the amplitude of the perturbations on small scales can be significantly enhanced~\\cite{Leach:2000yw,Bugaev:2008bi}. The models generated by the modified flow algorithm which end naturally (roughly $7\\%$ of the total) generally have a red spectrum of perturbations on all scales and so are unaffected by PBH constraints. The remaining $93\\%$ of models equations have a late time attractor with $\\epsilon_{H} \\rightarrow 0$ and it is assumed that an auxiliary mechanism terminates inflation. The majority of these models are however excluded due to PBH over-production. The number of viable models decreases if the power spectrum is calculated numerically. Of the models generated using our modified flow analysis which have cosmological observables within the 3$\\sigma$ ranges found by WMAP7~\\cite{Komatsu:2010fb} $19\\%$ are excluded by PBH over-production.  We conclude that PBH constraints provide a significant constraint on models of inflation. Furthermore to exploit their full power an accurate numerical calculation of the amplitude of primordial perturbations on small scales, which exit the horizon close to the end of inflation, is required."
    },
    "1004/1004.0754_arXiv.txt": {
        "abstract": "{ We carefully analyze how the abundance of Nitrogen over Oxygen evolves when dependent on metallicity stellar yields with a primary component of N proceeding from AGBs stars are used. We show the results obtained with a chemical evolution models grid, calculated with variable star formation efficiencies, which produce different star formation histories. Finally we see how the N/O abundance is related on the evolutionary history. ",
        "introduction": "Nitrogen abundances when compared with the Oxygen one, N/O, may, in principle, inform about the time in which low and intermediate mass (LIM) stars formed, and, consequently, when the widest episode of star formation took place. This idea is based on the stellar mean-lifetimes for LIM stars, main producers of N, which are longer than the ones of massive stars, which eject O to the interstellar medium (ISM).  The data, however, must be carefully interpreted before to reach misleading conclusions based in this simple scheme. If a mass of gas form a bulk of stars simultaneously, in a Single Stellar Populations (SSP), it is true that N will appear after O in the ISM, with a certain time delay. However, when data proceeding from different galaxies or different regions within a galaxy, are compared, the previous scenario is not longer valid, since a) most of data refers to spiral and irregular galaxies, where the star formation is a continuous process and, due to that, there are a mix of SSPs; and b) the observations correspond to the final stage after an evolutionary path object, and galaxies follow different tracks. ",
        "conclusions": "The present day data as well as the high-redshift trends in the plane O/H-N/O may be reproduced with chemical evolution models using stellar yields where LIM stars create a certain quantity of NP  Differences in the star formation histories of galaxies are essential to reproduce the data and the observed dispersion, the present position of a galaxy in the diagram N/O vs O/H being determined by this evolutionary history  It is possible to have NP ejected by LIM stars even at a low O abundance since O abundance is not a time scale."
    },
    "1004/1004.4341_arXiv.txt": {
        "abstract": "s{ Even 100 years after the discovery of cosmic rays their origin remains a mystery. In recent years, TeV gamma-ray detectors have discovered and investigated many Galactic sources where particles are accelerated up to energies of 100 TeV.  However, it has not been possible up to now to identify these sites unambiguously as sources of hadronic acceleration. The observation of cosmic high-energy neutrinos from these or other sources will be a smoking-gun evidence for the sites of the acceleration of cosmic rays.} ",
        "introduction": "Up to now, information on objects in our galaxy and beyond has nearly exclusively been obtained using electromagnetic waves as cosmic messengers. Modern astronomy started in 1610 with optical photons, when Galileo Galilei took one of the first telescopes and pointed into the sky. Only in the last century, people started to extended the observations to lower and higher energies of the electromagnetic spectrum, where today's instruments span an enormous range of 20 orders of magnitude from radio waves to TeV gamma-rays.  In addition, we know from measurements of the cosmic-ray spectrum that there exist sources in the universe which accelerate protons or heavier nuclei up to energies of $\\sim 10^{20}$\\,eV, $10^7$ times higher than the most energetic man-made accelerator, the LHC at CERN. These highest energies are believed to be reached in extra-Galactic sources like gamma-ray bursts or active galactic nuclei whereas Galactic sources like supernova remnants (SNRs) or micro-quasars are thought to accelerate particles at least up to energies of $3\\times10^{15}$\\,eV, also called the ``knee'' region of the cosmic-ray spectrum. Though these phenomena are quite different, the basic acceleration mechanism is believed to be very similar. Particles are injected into shock fronts which develop when fast moving matter collides with other matter. In these shock fronts, the injected particles are accelerated in a repeating process (Fermi acceleration) where the energy gain per cycle is only small.  Other phenomena where a ``one-shot'' acceleration of cosmic rays to high energies might take place are objects with strong magnetic fields (up to $10^{15}$\\,G) like pulsars and magnetars. However, despite the detailed measurements of the cosmic-ray spectrum and 100 years after their discovery by Victor Hess, we still do not know what the sources of the cosmic rays are as they are deflected in the Galactic and extra-Galactic magnetic fields and hence have lost all information about their origin when reaching Earth. Only at the highest energies beyond $\\sim 10^{19.6}$\\,GeV cosmic rays might retain enough directional information to locate their sources.  Alternative messengers for locating the sources of the cosmic rays must have two distinct properties: they have to be electrically neutral and essentially stable. Only two of the known elementary particles meet these requirements: photons and neutrinos. Both particles are inevitably produced when the accelerated protons or nuclei collide with matter or photons inside or near the source. In these reactions neutral and charged pions are produced which then decay into high-energy photons and neutrinos, respectively. The well known ratio between the production of neutral and charged pions yields a direct link between the photon and neutrino flux from proton-proton and proton-gamma interactions, respectively.  The number of sources detected in TeV gamma-rays has increased dramatically during the last decade which was made possible by significant advances in the technique of air Cherenkov telescopes. Today, over 100 sources are known, both Galactic (e.g., SNRs, micro-quasars, pulsars) and extra-Galactic (e.g., active galactic nuclei, starburst galaxies). However, it has not been possible up to now to unambiguously identify these sources as sites of cosmic-ray acceleration as the observed TeV photons could also have been generated in the up-scattering of photons in reactions with accelerated electrons (inverse-Compton scattering). High-energy neutrinos, on the other hand, are only produced in reactions of accelerated hadrons and hence are a smoking-gun evidence for the sources of cosmic rays. ",
        "conclusions": "The observation of cosmic high-energy neutrinos will open a new window to the universe and help to solve long-standing mysteries like the question of the origin of the cosmic rays. With the recent advances in gamma-ray astronomy we were able to identify many good Galactic candidate sources of neutrino emission. However, due to the expected low fluxes and the early high-energy cut-offs of the predicted neutrino spectra the detection of these sources is challenging.  IceCube will be the first detector to advance into the ``discovery region'' but will likely only scratch it. Therefore, a factor five to ten more sensitive detector is needed in the Northern hemisphere which will also cover the inner Galactic Plane and Center. This detector, KM3NeT, is supposed to start data taking around 2014. With IceCube and KM3NeT combined, neutrino astronomy will hopefully soon become reality."
    },
    "1004/1004.1087_arXiv.txt": {
        "abstract": "We study the excitation of fine structure levels of \\Cboth\\ and \\OI\\ by ultraviolet (UV) photons around strong UV sources which are also ionizing sources of the cosmological reionization at redshift of $\\sim 10$.  The evolutions of ionized regions around a point source are calculated by solving rate equations for non-equilibrium chemistry. Signals of UV photons through the fine structure lines are considered to be stronger at locations of more abundant chemical species of \\Cboth\\ and \\OI.  Such environments would be realized where strong fluxes of non-ionizing UV line photons available for the pumping up of fine structure levels exist, and simultaneously ionizing UV photons are effectively shielded by dense \\HI\\ regions.  Signals from \\HI\\ regions of moderately large densities induced by redshifted UV photons emitted at the point sources are found to be dominantly large over those of others.  We discuss the detectability of the signals, and show that signals from idealized environments will be possibly detected by radio observations with next-generation arrays to come after the Atacama Large Millimeter/submillimeter Array (ALMA). ",
        "introduction": "\\label{sec1}  At the epoch of Big Bang Nucleosynthesis of cosmic temperature $T \\ga 0.1$~MeV, light elements of $^2$H,$^{3,4}$He and $^{7}$Li are produced in abundances larger than $10^{-10}$ times that of hydrogen.  As the temperature decreases, light elements recombine with electrons, and trace amounts of chemical molecules form during the Dark Ages between the recombination and the reionization of the universe~\\citep{von2009}.  The universe is then considered to experience the reionization at temperature of $\\sim 30$~K, i.e., cosmic redshift of $z\\sim 10$~\\citep{lar2010}. Sources of the reionization are perhaps ultraviolet (UV) photons from quasi-stellar objects (QSOs), young galaxies and hypothetical low metallicity objects in the early universe, i.e., Population III stars, or some exotic energy injection processes triggered, for example, by the particle decay.  If stars including Population III have contributed to the reionization, the metal enrichment in the universe would occur simultaneously.  The reionization history as well as the metal enrichment of the universe is, however, not yet determined precisely.  Many suggestions have been made for investigations of the cosmic reionization using atomic physics with future observations in radio frequency.  The first method is observations of intergalactic medium (IGM) through the redshifted 21~cm line of neutral hydrogen~\\citep*{mad1997}.  The neutral gas regions could be quickly preheated to temperature larger than that of cosmic background radiation (CBR) in the presence of a strong flux of UV photons through the photoionization heating~\\citep[e.g.,][]{che2004}.  In such environments, the spin temperature determined by the population fractions of the upper and lower states of hyperfine levels is forced to deviate from the CBR temperature by level-mixing via Ly$\\alpha$ photon scattering, i.e., Wouthuysen-Field effect~\\citep{wou1952,fie1958}.  There would be brief chances of observing them in absorption, and long chances in emission against the CBR.  Future radio observations at meter wavelengths of IGM would produce plenty of information on the epoch, picture and sources of the reionization.  Although the 21~cm signal of the reionization of the universe is expected to be relatively strong, the strong contamination from Galactic and extragalactic foregrounds makes it difficult to observe the absolute signal over the whole sky except for fluctuations around it. Measurements of angular fluctuations in the \\HI\\ signal are very difficult, while fluctuations in frequency would be easily separated and be detectable with future large arrays~\\citep{gne2004}.  As for angular fluctuation, \\citet*{dim2004} have shown that angular fluctuations in 21~cm emission at $\\theta \\ga 1\\arcmin$ can be detected if efficient subtraction of foreground sources is performed successfully.  The 21~cm signal not by Wouthuysen-Field effect but by an effect of collision in earlier universe of higher densities has also been suggested to provide a valuable informations in the epoch when the spin temperature traced the gas kinetic temperature~\\citep{loe2004}.  Long-lived metastable 2$s$ state of hydrogen has been studied, and signals of the reionization through 3~cm fine structure line (2$s_{1/2} \\rightarrow 2p_{3/2}$) have been found to be undetectable with existing radio telescopes~\\citep{dij2008}.  The reason is that a Ly$\\beta$ photon which is the pumping source of the ground state can excite only one hydrogen atom into the 2$s$ state, while a Ly$\\alpha$ photon can excite many atoms into the excited spin state.  The signals through hyperfine transitions of \\DI, \\HeIIthree\\ \\citep{deg1985} also exist.  \\citet{sig2006} have suggested a challenging possibility of determination of D/H abundance ratio by looking at hyperfine structure line at high redshift of both the H and D.  Signals from IGMs and dense objects at high redshifts through the $^3$He hyperfine transition has been studied.  Detailed physics and suggested observational plans were presented~\\citep{mcq2009}, and the anticorrelation of \\HI\\ 21 cm maps and \\HeIIthree\\ 3 cm maps was suggested~\\citep{bag2009}.  The infrared fine structure lines of \\Cboth\\ and \\OI~\\citep{bah1968} play an important role in cooling of neutral regions~\\citep{hol1999}. Astronomical observations of those lines provide rich information about physical conditions in observed objects~\\citep{kau1999}.  The signatures of existence of heavy chemical species like C and O through resonant scatterings with CBR imprinted on the CBR spectrum have been investigated~\\citep*{bas2004,her2006}.  With the Planck~\\footnote{http://www.rssd.esa.int/index.php?project=planck.} spacecraft, the signatures can be searched for in the temperature anisotropy of CBR, and constraints on abundances of heavy species will be derived from the observations in reverse.  Idea of UV pumping through fine structure lines of metals during reionization as a way to distort the CMB during a reionization epoch has been suggested~\\citep{her2007}.  They focused on the \\OI\\ fine structure, and provided estimations for the order of magnitude of the distortion for different levels of UV fluxes.  The angular power spectrum of the distortion introduced by the \\OI\\ hosting regions at high redshift has subsequently been studied~\\citep{her2008}.  Observations of bright objects in radio frequency have been performed in search of emission lines of fine structures of heavy chemical species.  The \\CII\\ 157$\\mu$m line emissions of high redshift QSOs [J1148+5251 at $z=6.42$~\\citep{mai2005}, BR~1202-0725 at $z=4.7$~\\citep{ion2006}] and a galaxy [BRI~0952-0115 at $z=4.43$~\\citep{mai2009}] have been detected.  \\citet{wal2009} observed the high redshift quasars (J1148+5251 at $z=6.42$) and derived an upper limit on a \\NII\\ 205$\\mu$m line luminosity.  The targets of these observations were, in fact, star-forming interstellar medium which emits through atomic fine structure lines.  There are dense systems detected which absorb photon spectra of QSOs at fine structure transitions of \\Cboth\\ and their physical conditions have been constrained~\\citep*{sil2002,qua2002,rei2005}.  In this paper, we focus on cosmological signals of only UV sources without star forming activities through fine structure lines and their detectability.  The structure of this paper is as follows.  In Section\\ \\ref{sec2}, the pumping mechanism of fine structure levels and the dependence of the pumping efficiency as a redshift are shown.  In Section\\ \\ref{sec3}, the model for calculation of reionization around a point source is introduced.  In Section\\ \\ref{sec4}, the result of the calculation is presented.  Intensities of signals of UV radiation from the source through fine structure lines are estimated, and the dependence of signals on the redshift is shown.  We discuss the detectability of signals, and claim that signals may be detected by future radio measurements in $O(100-1000~{\\rm GHz})$ range.  In Section\\ \\ref{sec5}, we summarize our conclusions. ",
        "conclusions": "\\label{sec5}  The reionization history of the universe at early epoch of redshift $z \\sim 10$ is not know precisely yet.  We study signals of the cosmological reionization which would have been left on the cosmic background radiation (CBR) through a series of excitations of fine structure levels of \\Cboth\\ and \\OI\\ and following emissions (absorption) of line photons.  Since the reionization of the universe would have proceeded inhomogeneously in space, regions of neutral or low ionization states irradiated by non-ionizing ultraviolet (UV) photons naturally exist during the reionization epoch.  Non-ionizing UV photons interact with C and O to excite them to unstable excited states.  This excitation followed by decays of the excited states leads to an excitation of fine structure levels.  Since the UV photons as a source of the reionization produce signals through fine structure lines in \\CI, \\CII\\ or \\OI\\ regions, some information on the reionization may be obtained by observations for these lines.  Essentially, one ionizing photon can produce one line photon for a fine structure transition by scattering with C and O species.  Strong signals are, therefore, emitted at locations near strong sources of UV emission with large flux of UV photon.  We then assume a strong point source such as quasi-stellar object (QSO) which starts lighting in the early universe, and calculate the evolution of ionized region utilizing a non-equilibrium chemical reaction network.  A rough picture of ionization is shown by this calculation that \\Cboth\\ and \\OI\\ regions irradiated by nonionizing UV line photons can exist at locations where ionizing UV photons are effectively shielded by dense \\HI\\ regions.  There are two classes of UV photons available for the UV pumping of fine structure levels:  1. photons of frequencies just around the fine structure lines at the boundary between ionized region and non-ionized region for \\Cboth\\ and \\OI,  2. photons which are emitted as more energetic photons at point source and redshifted toward the transition energies of fine structures in \\Cboth\\ and \\OI\\ regions.  At ionization boundaries, UV photons which can excite the fine structure levels are immediately lost since they are used by the UV pumping.  It is then predicted that very small regions emitting line photons energized by UV sources would exist at ionization boundaries.  However, the signals of the latter class of UV photons is expected to be stronger than those of the first.  Outside the ionization boundaries, redshifted UV photons can leave their signatures possibly over wide region.  The ratios between UV intensities of respective fine structure lines are changed in the relaxation process through scattering with the C and O species, and are affected by the spin states of C and O.  Neutral \\HI\\ regions near strong UV sources emit signals of UV photons available until the relaxation is completed.  Those signals of UV photons could be contaminated by signals of collisional excitations if densities in observed regions were large. Regions of intermediate densities might be good sites to investigate for signals of the reionization since an effective UV pumping is possible in denser region where an effect of collisional excitation is also larger.  The dependence of magnitudes of signals on the source redshift is shown, and the detectability of such signals is discussed taking the Atacama Large Millimeter/submillimeter Array (ALMA) as an example.  Although the magnitudes of signals depend on physical environments of observed points, which are roughly described by many fixed parameters in this study, detections of fine structure line of \\Cboth\\ and \\OI\\ might be possible with the ALMA or future projects."
    },
    "1004/1004.4207_arXiv.txt": {
        "abstract": "\\noindent The design, optimisation and construction of an anti-coincidence veto detector to complement the ZEPLIN--III direct dark matter search instrument is described. One tonne of plastic scintillator is arranged into 52 bars individually read out by photomultipliers and coupled to a gadolinium-loaded passive polypropylene shield. Particular attention has been paid to radiological content. The overall aim has been to achieve a veto detector of low threshold and high efficiency without the creation of additional background in ZEPLIN--III, all at a reasonable cost. Extensive experimental measurements of the components have been made, including radioactivity levels and performance characteristics. These have been used to inform a complete end-to-end Monte Carlo simulation that has then been used to calculate the expected performance of the new instrument, both operating alone and as an anti-coincidence detector for ZEPLIN--III. The veto device will be capable of rejecting over 65\\% of coincident nuclear recoil events from neutron background in the energy range of interest in ZEPLIN--III. This will reduce the background in ZEPLIN--III from $\\simeq$0.4 to $\\simeq$0.14 events per year in the WIMP acceptance region, a significant factor in the event of a non-zero observation. Furthermore, in addition to providing valuable diagnostic capabilities, the veto is capable of tagging over 15\\% for $\\gamma$-ray rejection, all whilst contributing no significant additional background. In conjunction with the replacement of the internal ZEPLIN--III photomultiplier array, the new veto is expected to improve significantly the sensitivity of the ZEPLIN--III instrument to dark matter, allowing spin independent WIMP-nucleon cross sections below 10$^{-8}$~pb to be probed. ",
        "introduction": "\\noindent ZEPLIN--III~\\cite{1,z3sim} is a two-phase (liquid/gas) xenon detector developed to observe low energy nuclear recoils resulting from the elastic scattering of galactic weakly interacting massive particles (WIMPs)~\\cite{3}. A key feature of ZEPLIN--III is the ability to discriminate between incident particle species for each event by recording both vacuum ultraviolet scintillation light and electroluminescence from ionisation of the xenon target~\\cite{4,5}. The ratio in signal strength from these channels differs for electron and for nuclear recoil interactions, allowing the efficient rejection of most background events. However, in the absence of an electromagnetic interaction, weak and strong elastic scattering events appear identical, making the rejection of single scatter background neutron events difficult~\\cite{6,7,8}. Additionally, even with good discrimination, for large exposures a small fraction of electron-recoil events can be misidentified as nuclear recoils. With predicted WIMP event rates of less than 0.1 events/kg/day, and with energy depositions of the order of 10 keV, dark matter detectors such as ZEPLIN--III must reduce the number of background events, especially due to neutrons, by a very large factor. Typically this is achieved through the use of high purity components of low radionuclide content, hydrocarbon and lead shielding external to the detector to attenuate and moderate neutrons and $\\gamma$-rays originating in the local environment, as well as the operation of the detector at a site with large rock overburden to reduce the cosmic muon flux. Having been manufactured from low background components such as oxygen free copper, shielded with 30~cm of high density polypropylene and lead, and situated at the Boulby Underground Laboratory, UK, the background event rate of ZEPLIN--III is already greatly reduced. Sensitivity to spin-independent WIMP-nucleon cross sections as low as 8.1$\\times$10$^{-8}$~pb for a WIMP mass of 55 GeV/c$^2$ has been demonstrated~\\cite{z3fsr}, as well as a spin-dependent WIMP-neutron cross section as low as 1.8$\\times$10$^{-2}$~pb~\\cite{z3fsrsd}. However, to improve the sensitivity of the device further, ZEPLIN--III is being retrofitted with an active veto detector. Since WIMPs are highly unlikely to interact within a detector system twice, any event doing so may be discarded.  Following this upgrade, the predicted combined background event rate in ZEPLIN--III translates into much less than one count within the WIMP acceptance region for a one year dataset. The implementation of this veto system, in conjunction with new custom-built ultra-low background photomultiplier tubes (PMTs), will allow ZEPLIN--III to probe spin independent WIMP-nucleon cross sections below the 10$^{-8}$~pb level.  Experimental measurements have been made of the individual performance and radiological content of all components. These results have been fed into an end-to-end Monte Carlo simulation of the ZEPLIN--III instrument~\\cite{z3sim}, coupled with the entire veto detector array located in the Boulby mine, to evaluate the overall performance. An iterative process has been followed in which, with the geometry of the scintillator defined, alternative additional components and light collection efficiency were evaluated, informed by further experimental data where necessary.  In the following sections we first describe the general design of the veto.  We then provide details of the measurements of the radioactivity content of the components, the performance of components, and the optimisation process.  Finally we provide results from the evaluation of the predicted overall performance of the veto. ",
        "conclusions": "\\noindent The second science run of the ZEPLIN--III project will feature two main upgrades: lower background photomultiplier tubes and the use of an active veto. Here, the design of the veto has been presented, together with details of the radiological content and performance of the components to be used. Detailed GEANT4 Monte Carlo simulations have been used to aid in the characterisation of the veto, and to estimate its overall performance. In addition to providing valuable diagnostic information, it is expected that the veto will be able to reject over 65\\% of neutrons, and over 15\\% of $\\gamma$-rays from background radioactivities, whilst contributing negligibly to the ZEPLIN--III acceptance region background, expected to be $\\simeq$0.4 events per year.  The veto will reduce this background to less than $\\simeq$0.14\\ events per year.  This is a significant factor in the event of a non-zero observation. The veto has been fully assembled and systems integrated, and will be installed around ZEPLIN--III shortly for the commencement of the second science run as this instrument probes yet greater sensitivity."
    },
    "1004/1004.3611_arXiv.txt": {
        "abstract": "If the orientations of galaxies are correlated with large-scale structure, then anisotropic selection effects such as preferential selection of face-on disc galaxies can contaminate large scale structure observables.  Here we consider the effect on the galaxy bispectrum, which has attracted interest as a way to break the degeneracy between galaxy bias and the amplitude of matter fluctuations $\\sigma_8$.  We consider two models of intrinsic galaxy alignments: one where the probability distribution for the galaxy's orientation contains a term linear in the local tidal field, appropriate for elliptical galaxies; and one with a term quadratic in the local tidal field, which may be applicable to disc galaxies.  We compute the correction to the redshift-space bispectrum in the quasilinear regime, and then focus on its effects on parameter constraints from the transverse bispectrum, i.e. using triangles $(\\bmath k_1,\\bmath k_2,\\bmath k_3)$ in the plane of the sky. We show that in the linear alignment model, intrinsic alignments result in an error in the galaxy bias parameters, but do not affect the inferred value of $\\sigma_8$. In contrast, the quadratic alignment model results in a systematic error in both the bias parameters and $\\sigma_8$.  However, the quadratic alignment effect has a unique configuration dependence that should enable it to be removed in upcoming surveys. ",
        "introduction": "While the evolution of dark matter perturbations in the current $\\Lambda$CDM model is well understood theoretically, the relation between the galaxy distribution and the large scale (dark) matter distribution is complicated by the detailed physics of galaxy formation and different models may lead to different clustering properties of galaxies.  In particular, while local theories of galaxy formation predict that the galaxy density fluctuations trace the matter fluctuations on large scales, they also predict that the two are related by the {\\em bias parameter} $b$, which is in general not known {\\em a priori} \\citep{K84}.  The unknown bias parameter represents a key problem for attempts to measure the growth of cosmological perturbations using galaxies.  In combination with the galaxy power spectrum, third order galaxy clustering measures such as the bispectrum or (equivalently) the 3-point correlation function can be used to measure non-linear galaxy bias and break the degeneracy between the normalization of the matter power spectrum, $\\sigma_8$, and the linear galaxy bias. This enables one remove the effects of galaxy biasing and measure the cosmological growth of structure from the galaxy distribution \\citep{F94, VHM98, SCF99, V00}, and thus constrain dark energy \\citep[e.g.][]{DTJ06}. Recently third order galaxy clustering has been analyzed by several authors using the bispectrum \\citep{SFF01,FFS01,VHP02,KNS07} and the three point correlation function \\citep{JB04, KSN04,NSS06}. Using mock catalogs from numerical simulations, \\citet{SCPS06} show that a combined analysis of the galaxy power spectrum and bispectrum including their cross-correlation contains significant information on galaxy bias and fundamental cosmological parameters and helps break parameter degeneracies of other cosmological probes.  The most important systematic errors in interpreting the observed galaxy clustering arise in the non-linear regime,where the behavior of galaxy biasing and models of the (redshift space) galaxy power spectrum and bispectrum are difficult to model \\citep[see][for the complications of a current model of the redshift space Bispectrum]{S+08}. Recently \\citet{Hirata09} showed that the alignment of galaxies by large scale tidal fields can cause a systematic error in the determination of the linear redshift space distortion parameter $\\beta$ \\citep{K87}: the alignment of galaxies with the tidal field (along the stretching axis of the field for large elliptical galaxies) in combination with a viewing direction dependent galaxy selection effect, e.g. preferential selection of galaxies which are observed along their long axis, will lead to a selection probability for galaxies which is modulated by the tidal field along the line of sight.  This results in an anisotropy in redshift-space clustering with the same scale and angular dependence as the linear redshift-space effect. In this paper we will explore the implications of such a tidal alignment contamination for the observed galaxy bispectrum and how it affects the measurement of galaxy bias parameters.  Throughout this paper we assume a standard $\\Lambda$CDM cosmology with the best-fit WMAP 7 \\citep{WMAP7} parameters, and assume Gaussian initial density perturbations.  We begin in Section~\\ref{sec:theory} with a derivation of the standard redshift space galaxy bispectrum and discuss toy models for physical processes that cause alignments of galaxy orientations with large scale structure. In Section \\ref{sec:TA} we explain how tidal alignments of galaxies in combination with an orientation dependent galaxy selection modify the observed galaxy distribution and calculate the corresponding corrections to the galaxy bispectrum. Using a Fisher matrix technique we then estimate the systematic error induced by tidal alignments to measurements of galaxy bias parameters from angular clustering in Section~\\ref{sec:fisher}. We conclude and discuss mitigation strategies in Section~\\ref{sec:theend}. ",
        "conclusions": "\\label{sec:theend}  Using simple toy models for intrinsic alignment and the local bias approximation we have analyzed the effect of tidal alignment on the galaxy bispectrum. If the orientation of galaxies depends on the surrounding tidal field, and if the detection probability for galaxies is orientation dependent, the observed clustering of galaxies is modified by tidal alignments. This astrophysical contaminant can introduce systematic errors to parameters derived from the bispectrum.  \\begin{figure} \\includegraphics[width = 3.2in]{test.eps} \\caption{Removal of quadratic alignment bias. Dotted ellipses show the biased parameter estimates and their $95\\%$ contour regions in the presence of quadratic alignment contamination with $A_2 = 1$ which is unaccounted for in the analysis. The solid ellipses illustrate the $95\\%$ contour regions of the unbiased parameter estimates in an analysis which includes a quadratic alignment contamination and marginalizes over $A_2$.} \\label{fig:A2} \\end{figure}  A toy model for linear alignments \\citep{CKB}, which is based on the assumption that tidal fields elongate/compress haloes and thus determine galaxy shapes, results in a rescaling of linear and non-linear galaxy bias parameters that is proportional to the strength of the halo shape distortion. The presence of this systematic error in the observed galaxy bias measurements cannot be detected from projected clustering data as the strength of the alignment contamination is completely degenerate with the unobservable true bias parameters and outside information will be necessary to remove it. Normalizing the strength of the linear tidal alignment toy model to measurements of intrinsic alignments in weak lensing observations, we find that linear alignment may introduce systematic errors to galaxy bias measurements at the percent level (again using only the real-space observables), and thus will likely not be significant.  Using a simple model for quadratic alignment based on galaxy spin correlations in linear tidal torque theory we calculate a systematic contamination which modifies the shape of the galaxy bispectrum. Depending on survey characteristics, we find that quadratic alignment may introduce significant systematic errors to the galaxy bias parameters and the normalization of the power spectrum derived from the angular galaxy bispectrum. As the quadratic alignment contamination has different shape than the galaxy bispectrum, one can include a model for the contamination in the analysis and marginalize over its normalization. Figure~\\ref{fig:A2} illustrates how such a marginalization may remove the systematic bias at the cost of larger statistical errors. The biased data points and contour levels (dashed lines) are taken from Fig.~\\ref{fig:bias} for a fiducial model with $b_1 = 1$ and $b_2 = 0$. The new statistical errors including marginalization over $A_2$ are calculated by adding $A_2$ as a nuisance parameter and including the contamination signal in the fiducial model of the Fisher matrix analysis ($\\vec{\\mathcal B} \\rightarrow \\vec{\\mathcal B}+\\vec{\\Delta \\mathcal B}$ in Eq.~\\ref{eq:fisher}).  This analysis lives in the weakly nonlinear regime to enable the use of simple models for linear and quadratic alignment. As the information content of the bispectrum increases dramatically with the maximal spatial frequency that is included in an analysis, any realistic analysis will have extend well into the quasilinear regime. While models from the redshift space bispectrum on these scales \\citep{S+08} approach the required accuracy for such analyses, the treatment of tidal alignments including the non-Gaussian nature of the angular moment distribution and non-linear stages of galaxy formation requires further work."
    },
    "1004/1004.0539_arXiv.txt": {
        "abstract": "We have performed a survey for substellar companions to young brown dwarfs in the Taurus star-forming region using the Wide Field Planetary Camera 2 on board the {\\it Hubble Space Telescope}. In these data, we have discovered a candidate companion at a projected separation of $0\\farcs105$ from one of the brown dwarfs, corresponding to 15~AU at the distance of Taurus. To determine if this object is a companion, we have obtained images of the pair at a second epoch with the adaptive optics system at Gemini Observatory. The astrometry from the {\\it Hubble} and Gemini data indicates that the two objects share similar proper motions and thus are likely companions. We estimate a mass of 5--10~$M_{\\rm Jup}$ for the secondary based on a comparison of its bolometric luminosity to the predictions of theoretical evolutionary models. This object demonstrates that planetary-mass companions to brown dwarfs can form on a timescale of $\\tau\\lesssim1$~Myr. Companion formation on such a rapid timescale is more likely to occur via gravitational instability in a disk or fragmentation of a cloud core than through core accretion. The Gemini images also reveal a possible substellar companion ($\\rho=0\\farcs23$) to a young low-mass star that is $12\\farcs4$ from the brown dwarf targeted by {\\it Hubble}. If these four objects comprise a quadruple system, then its hierarchical configuration would suggest that the fragmentation of molecular cloud cores can produce companions below 10~$M_{\\rm Jup}$. ",
        "introduction": "\\label{sec:intro}   Most of the known planets outside of the solar system have been detected through radial velocity measurements, and thus reside close to their parent stars \\citep[$a\\lesssim5$~AU,][]{mar05,udr07}. Meanwhile, a small number of companions with masses below 10~$M_{\\rm Jup}$ have been detected at large separations through high-resolution imaging \\citep[$a\\sim30$--300~AU,][]{cha04,laf08,kal08,mar08}. The available data tend to favor core accretion as the mechanism responsible for the close-in giant planets \\citep{san04,fis05}, but gravitational instability in disks and fragmentation of molecular cloud cores are plausible (perhaps likely) alternatives for producing the wide companions \\citep{lod05,dob09}. The formation models for wide planetary-mass companions are particularly in need of observational tests.    Companions are expected to form much more quickly through disk instability and cloud core fragmentation than core accretion \\citep[$\\tau\\sim0.1$~Myr vs. $\\tau\\sim5$--10~Myr,][]{boss98,boss06,bur06,pol96}. Thus, measurements of the frequency of planetary companions at very young ages can help identify the dominant mode for the formation of these objects. To provide data of this kind, we have performed a search for substellar companions to 32 young brown dwarfs in the Taurus star-forming region ($\\tau\\sim1$~Myr, $d\\sim140$~pc) using high-resolution images obtained with the {\\it Hubble Space Telescope}.  In this Letter, we present the discovery of a companion that has a mass of 5--10~$M_{\\rm Jup}$ and a projected separation of 15~AU from one of the brown dwarfs in this survey. ",
        "conclusions": "The first planetary-mass companion to be detected through direct imaging was 2MASS J12073346-3932539~B \\cite[$a=40$~AU, $M\\sim4$~$M_{\\rm Jup}$,][]{cha04,bil07,giz07,duc08}, which is bound to a young brown dwarf in the TW Hya association \\citep[$\\tau\\sim8$~Myr, $M\\sim25$~$M_{\\rm Jup}$,][]{giz02}. The new companion that we have presented, 2M~J044144~B, is a younger analog of 2MASS J12073346-3932539~B, thus demonstrating that wide planetary-mass companions to brown dwarfs can form on a timescale of $\\tau\\lesssim1$~Myr. This result is consistent with the rapid formation expected from both gravitational instability in disks and fragmentation of cloud cores. The latter is more likely to have produced these two secondaries given their relatively large masses compared to the primaries \\citep{lod05}. Indeed, if 2M~J044144~A and B are members of a quadruple system, then its hierarchical configuration further suggests that 2M~J044144~B formed by cloud core fragmentation. Additional data are needed to determine more definitively whether 2M~J044144~A/B and 2M~J044145~A/B comprise a quadruple system."
    },
    "1004/1004.0649_arXiv.txt": {
        "abstract": " ",
        "introduction": "}} Recently several particle physics models have {been} constructed that connect the {\\smnew{(SM)}} to  hidden sectors and lead to massive narrow vector boson resonances {as well as} other signatures which can be detected at colliders { \\cite{kn12,FLNPRL,kctc}}. The connection to the hidden sector arises via mass mixings and kinetic mixings\\cite{kn12,FLNPRL,kctc,Holdom,Dienes:1996zr,FLN07} and via higher dimensional operators. {Models with the above forms of communication between the  sectors} also have important implications for dark matter {\\cite{FKN,kctc,FLN07}} {{(for a review see \\cite{white,Nath})}.} In this work we show that multicomponent dark matter can arise {from}  $U(1)^n$ extensions of the {minimal supersymmetric standard model } (MSSM)  with {Abelian} hidden {sectors} which include hidden sector matter. {Our} motivation stems in part from {the} results {of several} dark matter experiments that have {recently} appeared. Thus the PAMELA Collaboration  \\cite{PAM} has observed a positron excess improving previous results from HEAT and AMS experiments { \\cite{heatams}}. One possible explanation of such an excess is via the annihilation of dark matter in the galaxy\\cite{Feldman:2008xs}. Additionally, recent data from {CDMS-II} hints  at  the possibility of dark matter events above the background, and this will be explored  further by the upgraded XENON experiment  \\cite{CDMSXENON,Aprile:2010um}.  For a thermal relic, the PAMELA data and  {CDMS-II} data  taken together at face value do raise a theoretical puzzle if indeed both signals arise  from the annihilation of cold dark matter. Thus most models {which aim to explain the PAMELA  positron excess} do not give a significant number of dark matter events in  the direct  detection experiments currently operating. {Conversely,} models which {can} give {a} detectable signal in direct detection experiments {typically} do not explain the PAMELA data without the use of enormous so-called boost factors. As we will show here, this can  {be} circumvented in models where the dark matter has several components. Thus, motivated in part by the recent cosmic anomalies  we develop supersymmetric models which contain minimally a hidden Abelian sector broken at the sub-TeV scale where the mass generation of the hidden states involves nontrivial mixings with the field content of the electroweak sector of the minimal supersymmetric extension of the \\smnew leading  to dark matter which can have several components  which can be both bosonic and fermionic.     More specifically, in this work we  go beyond the simple theoretical construction that thermal dark matter compatible with  WMAP  observations is composed of a single fundamental particle. There is no overriding principle that requires such a restriction, and nonbaryonic dark matter (DM) may indeed be  constituted of several components, so in general one has $(\\Omega h^2)_{DM} = \\sum_{i} (\\Omega h^2)_{DMi}$, where $i$ {refers} to the various species of dark particles  that can contribute to the total nonbaryonic $(\\Omega h^2)_{DM}$. In fact we already know that neutrinos do contribute to dark matter although their contribution is relatively small.  Thus   we propose here a new class of multicomponent  cold dark matter models in Abelian $U(1)$ extensions of MSSM which can simultaneously provide an explanation of the PAMELA and WMAP data  { through a Breit-Wigner enhancement \\cite{Feldman:2008xs}}, while producing detectable  signals for the direct searches for dark matter with CDMS/XENON and other dark matter experiments.  A simultaneous satisfaction of  the PAMELA positron excess and the satisfaction of  WMAP relic density constraints can also occur if there is  a nonthermal mechanism for the annihilation of dark matter with a wino lightest (R parity odd) supersymmetric particle (LSP) \\cite{MoroiRandall,Kane,Acharya,Hisano,white,Nath}. However, a detectable spin-independent cross  section in such a nonthermal  framework does require that a pure wino is supplemented by a suitable admixture of Higgsino content as in the analysis of  \\cite{FKRN} and in  \\cite{Feldman:2009wv}{,} the later for a thermal relic. We remark that multiple  $U(1)$ factors and its influence on dark matter have very recently been studied \\cite{Feldman:2009wv,Arvanitaki}. We  also remark, some other works have recently looked at dark matter with more than 1 component \\cite{Hur:2007ur}.  The models proposed {and analyzed} here are very different from these.     The outline of the rest of the paper is as follows: In Sec.(\\ref{hidden}) we give a detailed description of the two models one of  which is based on a $U(1)_X$ extension of the MSSM where $U(1)_X$ is a hidden sector gauge group with  Dirac fermions in the hidden sector. This model allows for dark matter consisting of Dirac, Majorana, and spin zero {particles}. The second model is based on a $U(1)_X \\times U(1)_C$ extension of MSSM, where $U(1)_C$ is a {gauged} leptophilic symmetry and $U(1)_X$, as before, is the hidden sector gauge group   which also contains Dirac particles  in the hidden sector. This model too has Dirac, Majorana, and spin zero particles  as possible dark matter. In both {cases} we will {primarily} focus on the possibility that dark matter  consists  of Dirac  and Majorana particles, {and  we will not discuss in detail  the possibility of  dark matter with bosonic degrees of freedom.} In Sec.(\\ref{relic}) we discuss the relic densities  in the two component models. In Sec.(\\ref{pamela}) we give  an analysis of the positron, antiproton, and photon fluxes in the two models. In Sec.(\\ref{events}) we  give an analysis of event rates  for the proposed models for {CDMS-II} and for {XENON-100.} We give the analysis within the {framework of} supergravity grand {unified models}  \\cite{Chamseddine:1982jx,Arnowitt:1992aq}  defined by the parameters $m_0, m_{1/2}, A_0, \\tan\\beta$, and sign$(\\mu)$  with nonuniversalities (NUSUGRA) defined by $\\delta_{1,2,3}$  {in the gaugino sector} {so that} $U(1)_{{Y}}\\times SU(2)_L\\times SU(3)_C$ gaugino masses at the grand unified theory (GUT) scale are given by $\\tilde m_{i} = m_{1/2} (1+\\delta_i)$ {($i=1,2,3$)} (see, e.g., \\cite{Corsetti:2000yq} and references {therein}). {We} also discuss the possible  new physics one might {observe} at the LHC (for a recent review see also  \\cite{Nath}) and elsewhere for these models. Conclusions  are given in Sec.(\\ref{conclusions}). ",
        "conclusions": "{\\label{conclusions}} {In this work} we  have proposed a new class of models with dark matter consisting of two, three, or even four components.  We considered the two component model consisting of Dirac and  Majorana  particles in detail. We showed that this {two component} model can fit the positron excess seen in the PAMELA experiment as well as can produce detectable signals in the current direct detection experiments while satisfying WMAP relic density constraints.   Thus. the   Dirac  component of  the two component dark matter {model} allows  a fit to the PAMELA data via annihilation of the Dirac particles close to a Breit-Wigner pole. On the other hand,  the Majorana component of dark matter plays the dominant role in the generation of events in dark matter detectors. Specifically, we  showed that in the two component picture it is possible to generate events of size 1-2 in 612 Kg-d of data in the {CDMS-II} detector as well as event rates that can be tested by the results of XENON-100, and an observable number of events in other ongoing direct detection experiments. Further, it was shown that models which lead to detectable signals in direct detection experiments are typically associated  with a relatively light spectrum which is discoverable at the LHC with modest luminosity. Further, one class of models discussed in this work produces  a $Z^{\\prime \\prime}$ vector boson which has visible decays only to $e^+e^-$ and $\\mu^+\\mu^-$. The proposed models contains massive  scalar fields which are also possible candidates for dark matter. {Thus these spin zero fields in combinations with Dirac and Majorana particles  present the possibility of a multicomponent dark matter.} Finally, we note that it would be very interesting to investigate phenomena where both components play a significant role, i.e., regions of the parameter space of the model which allow both dark matter candidates to appear in the data analysis.  Such a possibility may appear in certain decay fragments at the LHC where the missing energy signals from the two dark matter particles would be different because of their different masses and interactions. \\\\    \\noindent {\\em {Acknowledgements}}:  DF thanks Kenji Kadota  for a communication. This research is  supported in part by DOE {Grant No.}  DE-FG02-95ER40899 (Michigan - MCTP), NSF {Grant No.} PHY-0757959 {(Northeastern)}, and NSF {Grant No.} PHY-0653342 (Stony Brook - YITP)."
    },
    "1004/1004.3156_arXiv.txt": {
        "abstract": "{} {$J$, $H$, and $K^{'}$ images obtained from the near-infrared imager CFHTIR on the Canada-France-Hawaii Telescope are used to derive the morphological parameters of the red giant branch (RGB) in the near-infrared color-magnitude diagrams for 12 metal-poor globular clusters in the Galactic bulge direction. Using the compiled data set of the RGB parameters for the observed 12 clusters, in addition to the previously studied 5 clusters, we discuss the properties of the RGB morphology for the clusters and compare them with the calibration relations for the metal-rich bulge clusters and the metal-poor halo clusters. } {The photometric RGB shape indices such as colors at fixed magnitudes of $M_K=M_H=(-5.5, -5, -4,$ and $-3)$, magnitudes at fixed colors of $(J-K)_o = (J-H)_o = 0.7$, and the RGB slope are measured from the fiducial normal points defined in the near-infrared color-magnitude diagrams for each cluster. The magnitudes of RGB bump and tip are also estimated from the differential and cumulative luminosity functions of the selected RGB stars. The derived RGB parameters have been used to examine the overall behaviors of the RGB morphology as a function of cluster metallicity.} {The correlations between the near-infrared photometric RGB shape indices and the cluster metallicity for the programme clusters compare favorably with the previous observational calibration relations for metal-rich clusters in the Galactic bulge and the metal-poor halo clusters. The observed near-infrared magnitudes of the RGB bump and tip for the investigated clusters are also in accordance with the previous calibration relations for the Galactic bulge clusters. } {} ",
        "introduction": "\\begin{table} \\caption{Observational log of the target clusters}      % \\label{tbl1} \\centering \\begin{tabular}{c c c c c c} \\hline\\hline Target & Filter & Exp. time (sec) & FWHM$(^{''})$ & Year \\\\    % \\hline NGC 6333 & $J$     & 4$\\times$1, 8$\\times$30 & 0.60, 0.67 & 2002\\\\      % & $H$     & 4$\\times$1, 8$\\times$30 & 0.61, 0.58 &\\\\ & $K^{'}$ & 4$\\times$1, 8$\\times$30 & 0.58, 0.61 &\\\\ NGC 6626 & $J$     & 4$\\times$1, 8$\\times$30 & 0.67, 0.62 & 2002\\\\      % & $H$     & 4$\\times$1, 8$\\times$30 & 0.57, 0.65 &\\\\ & $K^{'}$ & 4$\\times$1, 8$\\times$30 & 0.62, 0.65 &\\\\ NGC 6235 & $J$     & 4$\\times$2, 8$\\times$30 & 0.73, 0.74 & 2003\\\\      % & $H$     & 4$\\times$2, 8$\\times$30 & 0.77, 0.81 &\\\\ & $K^{'}$ & 4$\\times$2, 8$\\times$30 & 0.72, 0.77 &\\\\ NGC 6266 & $J$     & 4$\\times$2, 8$\\times$30 & 0.82, 0.93 & 2003\\\\      % & $H$     & 4$\\times$2, 8$\\times$30 & 0.73, 0.88 &\\\\ & $K^{'}$ & 4$\\times$2, 8$\\times$30 & 0.77, 0.92 &\\\\ NGC 6273 & $J$     & 4$\\times$2, 8$\\times$30 & 0.86, 0.87 & 2003\\\\      % & $H$     & 4$\\times$2, 8$\\times$30 & 0.83, 0.74 &\\\\ & $K^{'}$ & 4$\\times$2, 8$\\times$30 & 0.83, 0.74 &\\\\ NGC 6287 & $J$     & 4$\\times$2, 8$\\times$30 & 0.61, 0.76 & 2003\\\\      % & $H$     & 4$\\times$2, 8$\\times$30 & 0.62, 0.73 &\\\\ & $K^{'}$ & 4$\\times$2, 8$\\times$30 & 0.65, 0.73 &\\\\ NGC 6293 & $J$     & 4$\\times$2, 8$\\times$30 & 0.89, 1.18 & 2003\\\\      % & $H$     & 4$\\times$2, 8$\\times$30 & 0.95, 0.87 &\\\\ & $K^{'}$ & 4$\\times$2, 8$\\times$30 & 0.66, 0.73 &\\\\ NGC 6325 & $J$     & 4$\\times$2, 8$\\times$30 & 0.73, 0.77 & 2003\\\\      % & $H$     & 4$\\times$2, 8$\\times$30 & 0.65, 0.78 &\\\\ & $K^{'}$ & 4$\\times$2, 8$\\times$30 & 0.68, 0.73 &\\\\ NGC 6355 & $J$     & 4$\\times$2, 8$\\times$30 & 0.74, 0.80 & 2003\\\\      % & $H$     & 4$\\times$2, 8$\\times$30 & 0.73, 0.87 &\\\\ & $K^{'}$ & 4$\\times$2, 8$\\times$30 & 0.73, 0.83 &\\\\ NGC 6401 & $J$     & 4$\\times$2, 8$\\times$30 & 0.66, 0.68 & 2003\\\\      % & $H$     & 4$\\times$2, 8$\\times$30 & 0.61, 0.72 &\\\\ & $K^{'}$ & 4$\\times$2, 8$\\times$30 & 0.60, 0.72 &\\\\ NGC 6558 & $J$     & 4$\\times$2, 8$\\times$30 & 0.95, 0.99 & 2003\\\\      % & $H$     & 4$\\times$2, 8$\\times$30 & 0.80, 0.99 &\\\\ & $K^{'}$ & 4$\\times$2, 8$\\times$30 & 0.73, 0.89 &\\\\ Terzan 4 & $J$     & 4$\\times$1, 8$\\times$30 & 0.68, 0.75 & 2003\\\\      % & $H$     & 4$\\times$1, 8$\\times$30 & 0.65, 0.73 &\\\\ & $K^{'}$ & 4$\\times$1, 8$\\times$30 & 0.80, 0.72 &\\\\ \\hline \\end{tabular} \\end{table}  The current view of the Galaxy formation is mainly focused on the hierarchical merging paradigm in the cold dark matter cosmology. Globular clusters, as tracers of the early formation and the current structure of the Galaxy, play a key role in studies of the paradigm, because they are present from the central bulge to the outer halo with various metallicities. Particularly, the Galactic bulge harbors a globular cluster population with a broad metallicity distribution that extends from about twice solar to less than one-tenth solar abundance~\\citep{Ort99}, while most field stars in the bulge have near-solar metallicity~\\citep{Mcw94,Zoc03}. The metal-rich globular clusters in the the Galactic bulge share the kinematics, spatial distribution, and composition of the bulge field stars~\\citep[e.g.,][and references therin]{Min08}. This indicates that metal-rich globular clusters are associated with the Galactic bulge recognized as the dominant proto-Galactic building block~\\citep[e.g.,][]{Cot00}. On the other hand, the origin of the metal-poor globular clusters in the Galactic bulge direction is still a subject of debate since accurate measurements of kinematics and high resolution chemical abundances are lacking. In the hierarchical model of the Galaxy formation, however, old metal-poor field stars in the bulge form via merging and accretion events in the early Universe~\\citep{Nak03}. In this sense, the metal-poor clusters currently located in the central region of the Galaxy might be the oldest objects which did not form originally in the Galactic bulge. Thus, the metal-poor clusters in the bulge region can play a key role to understand the early epoch of the formation of the Galactic bulge.   In the Milky Way, about 150 globular clusters are listed in the database of ~\\citet{Har96}, which was revised in 2003. Recently, new faint clusters and cluster candidates have also been found ~\\citep[e.g.,][]{Car05,Kob05,Wil05,Fro07}. Out of 43 globular clusters located within 3 kpc of the Galactic center~\\citep{Har96}, 22 are metal-poor ([Fe/H] $<-1.0$) and 21 are metal-rich ([Fe/H] $>-1.0$). Recently, ~\\citet{Val07} presented near-infrared color-magnitude diagrams (CMDs) and physical parameters for a sample of 24 globular clusters toward the Galactic bulge and located within $|b|\\leqq10^{\\circ}$ and $|l|\\leqq20^{\\circ}$. They discussed the near-infrared properties of the red giant branch (RGB) for 12 observed clusters, in addition to those previously studied by their group ~\\citep[e.g.,][]{Fer00,Val04a,Val04b,Val04c,Val05,Ori05}. We note, however, that their sample of the clusters have high priority to the metal-rich population, i.e., 17 out of the 24 are relatively metal-rich with [Fe/H] $>-1.0$, taking into account of a bulge origin for the metal-rich globular clusters.  In our research, we have focused on obtaining a moderately deep homogeneous photometric data set in the near-infrared regime for the metal-poor clusters in the bulge direction. Near-infrared photometry offers advantages for a study of the cool population of the RGB stars in the Galactic globular clusters, because of its high sensitivity to low temperature. In addition, high extinctions toward the bulge can be reduced by observing the near-infrared wavelengths, as the extinction in the $K$ band is only $\\sim10$ percent of that in the $V$ band~\\citep{Rie85}. Using these bases, ~\\citet{Kim06} presented the morphological properties of the RGB in the near-infrared CMDs for five metal-poor clusters of the Galactic bulge (NGC 6541, NGC 6642, NGC 6681, NGC 6717, and NGC 6723) and also for three halo clusters.  In this paper, we report new results of the near-infrared photometry for 12 metal-poor clusters and present a homogeneous photometric data set of the RGB morphology for 17 globular clusters, covering $\\sim75\\%$ of the total 22 metal-poor globular clusters in the Galactic bulge direction. The results of the RGB morphology for the programme clusters are compared with the previously published calibrations of ~\\citet{Val04a,Val04b,Val07}. The observations, procedures for data reduction, and phtometric measurements are described in Sect. 2. In Sect. 3, we describe the near-infrared CMDs and the fiducial normal points of target clusters. In Sect. 4, the morphological properties of CMDs such as RGB shape feature, RGB bump, and RGB tip are presented. Finally, the results are discussed and summarized in Sect. 5. ",
        "conclusions": "Detailed analyses of the RGB morphology for 12 metal-poor ([Fe/H] $\\leq-1.0$) globular clusters in the Galactic bulge direction have been performed using the high-quality near-infrared $JHK^{'}$ photometry. From the study of the RGB shapes in the near-infrared CMDs for each cluster, we measured photometric parameters, such as, the colors at different magnitude levels, the magnitudes at different colors, and the RGB slopes. The magnitudes of the RGB bump and tip, as major RGB evolutionary features, have also been determined from the LFs of the selected RGB stars in each cluster. The determined indices of the RGB morphology for the 12 observed clusters have been combined with the results for 5 bulge clusters in~\\citet{Kim06}, thus the entire dataset comprises $\\sim75\\%$ of the total 22 metal-poor ([Fe/H] $\\leq-1.0$) globular clusters within 3 kpc from the Galactic center. The behavior of the RGB morphology for the programme clusters has been compared with the previous empirical calibration relations as a function of cluster metallicity for the Galactic bulge globular clusters by~\\citet{Val04a,Val07} and theoretical predictions of the Yonsei-Yale isochrones~\\citep{Kim02,Yi03}. The results are summarized as follows:  \\begin{enumerate} \\item Photometric indices for the RGB color at fixed magnitudes, $M_K=M_H=(-5.5, -5, -4,$ and $-3)$, and the RGB magnitudes at fixed colors, $(J-K)_o = (J-H)_o = 0.7$ have been measured from the fiducial normal points of the near-infrared $(J-K,K)$ and $(J-H,H)$ CMDs. Our results indicate that the correlations between the derived RGB indices and the cluster metallicity for the metal-poor globular clusters in the Galactic bulge direction are consistent with previous observational calibration relations for a sample of the metal-rich bulge clusters and the halo clusters~\\citep{Val04a}. The trends of the theoretical models reliably represent the observed RGB color and magnitude indices, although there appears to be systematic shifts in color and magnitude, as a result of the uncertainties in the theoretical calculations and observational measurements. \\item The RGB slopes have been estimated from the determined fiducial normal points at the magnitude range between 0 and 5 magnitude fainter than the brightest point of the RGB. The distribution of the RGB slopes for the observed clusters show an expected evolutionary feature, i.e., the lower metallicity of the cluster, the steeper the RGB slope, while the RGB slopes for the programme clusters tend to be slightly flatter than those in the previous calibrations of~\\citet{Val04a}. \\item The absolute $M_K$ and bolometric $M_{bol}$ magnitudes of the RGB bump and tip for the observed clusters have been determined from the differential and cumulative LFs of the selected RGB stars. The correlations between the cluster metallicity and the derived magnitudes of the RGB bump and tip for the metal-poor clusters in the Galactic bulge direction are consistent with the recent calibration relations for the Galactic bulge clusters~\\citep{Val07}. \\end{enumerate}  Of a total of 17 metal-poor clusters presented in this paper, only two clusters NGC 6266 and NGC 6723 have the cluster's orbital data~\\citep{Din99,Din03}, indicating that NGC 6723 is a halo member passing the Galactic bulge at this moment and NGC 6266 is associated with the motion of the Galactic thick disk. Together with the derived RGB morphological properties, further information about detailed orbital data will provide more robust constraints on the role of the metal-poor globular clusters in the formation of the Galactic bulge."
    },
    "1004/1004.3683_arXiv.txt": {
        "abstract": "We study the spectroscopic properties of a large sample of Low Surface Brightness galaxies (LSBGs) (with B-band central surface brightness $\\mu_0(B)>$22 mag arcsec$^{-2}$) selected from the Sloan Digital Sky Survey Data Release 4 (SDSS-DR4) main galaxy sample. A large sample of disk-dominated High Surface Brightness galaxies (HSBGs, with $\\mu_0(B)<$22 mag arcsec$^{-2}$) are also selected for comparison simultaneously. To study them in more details, these sample galaxies are further divided into four subgroups according to $\\mu_0(B)$ (in units of mag arcsec$^{-2}$): vLSBGs (24.5-22.75), iLSBGs (22.75-22.0), iHSBGs (22.0-21.25), and vHSBGs ($<$21.25). The diagnostic diagram from spectral emission-line ratios shows that the AGN fractions of all the four subgroups are small ($<$9\\%). The 21,032 star-forming galaxies with good quality spectroscopic observations are further selected for studying their dust extinction, strong-line ratios, metallicities and stellar mass-metallicities relations. The vLSBGs have lower extinction values and have less metal-rich and massive galaxies than the other subgroups. The oxygen abundances of our LSBGs are not as low as those of the H~{\\sc ii} regions in LSBGs studied in literature, which could be because our samples are more luminous, and because of the different metallicity calibrations used. We find a correlation between 12+log(O/H) and $\\mu_0(B)$ for vLSBGs, iLSBGs and iHSBGs but show that this could be a result of correlation between $\\mu_0(B)$ and stellar mass and the well-known mass-metallicity relation. This large sample shows that LSBGs span a wide range in metallicity and stellar mass, and they lie nearly on the stellar mass vs. metallicity and N/O vs. O/H relations of normal galaxies. This suggests that LSBGs and HSBGs have not had dramatically different star formation and chemical enrichment histories. ",
        "introduction": "Low Surface Brightness Galaxies (LSBGs) are important populations in the field environments. However, their properties were seldom studied and their contributions to the galaxy population were underestimated since they are hard to find owing to their faintness compared with the night sky. An initial quantitative study was done by Freeman (1970), who noticed that the central surface brightness of their 28 out of 36 disc  galaxies fell within a rather narrow range, $\\mu_0$(B)=21.65$\\pm$0.3 mag arcsec$^{-2}$. This could be caused by selection effects as pointed out by Disney (1976), and had been previously recognized by Zwicky (1957).  Since then, many efforts were made to search for LSBGs from surveys, and a sample of LSBGs were gathered after then. Specially, Impey and his colleagues adopted the Automated Plate Measuring (APM) mechanism to scan UK Schmidt plates and discovered 693 LSB field galaxies which forms the most extensive catalog of LSBGs to that date (Impey et al. 1996). Then, their group studied the properties of this sample of LSBGs in a series of works. O'Neil et al. (1997a,b in ``Texas survey\") firstly found the red LSB galaxy populations. Monnier Ragaigne et al. (2003a,b,c) selected a sample of about 3800 LSBGs from the all-sky near-infrared 2MASS survey, and then obtained the 21 cm H~{\\sc i} observations and estimated the H~{\\sc i} masses for two subsamples of 367 and 334 LSBGs. More related references about LSBG searching and catalogs could be found in Bothun, Impey, McGaugh (1997), Impey \\& Bothun (1997) and Zhong et al. (2008).  The modern Sloan digital sky survey (SDSS, York et al. 2000) has its great advantages to extend the sample of LSBGs and study their physical properties in details. It provides to the public the homogeneous data of photometric and spectroscopic observations for several hundred thousands galaxies. Therefore, it will be very powerful and appropriate to search for a large sample of LSBGs from the SDSS.  Kniazev et al. (2004) developed a method to search for LSBGs from the SDSS early data release (EDR, Stoughton et al. 2002) field images and used the sample of Impey et. al. (1996) to test their method. They recovered 87 same objects as in Impey et al. and 42 new LSBGs. Bergvall et al. (2009) studied the red halos of 1510 nearly edge-on LSBGs from SDSS. Mattsson et al. (2008) discussed the N, O abundances of a subsample of them. Rosenbaum \\& Bomans (2004) studied the large-scale environment of LSBGs based on SDSS-EDR data, and Rosenbaum et al. (2009) extend the sample to Data Release 4 (DR4, Adelman-McCarthy et al. 2006) to get more detailed results. Our group has successfully selected a large sample of nearly face-on disc LSBGs from SDSS-DR4 main galaxy sample (Strauss et al. 2002). In Zhong et al. (2008), we present their basic photometric properties including the clear correlations of disk scalelength versus B-band absolute magnitude and distance (also see the similar results of Graham \\& Worley 2008 for high surface brightness galaxies, HSBGs), and the stellar populations from colors. This large sample includes 12,282 LSBGs with $\\mu_0$(B)$<$22 mag arcsec$^{-2}$, and another 18,051 HSBGs with $\\mu_0$(B)$>$22 mag arcsec$^{-2}$ are also selected for comparisons.  Chemical abundance is an fundamental property of galaxy. The chemical properties of stars and gas within a galaxy provide both a fossil record of its star formation history and information on its present evolutionary status. Some researchers have obtained the metallicities of some H~{\\sc ii} regions in a small sample of LSBGs. For example, McGaugh (1994) obtained the oxygen abundances of 41 H~{\\sc ii} regions in 22 LSBGs, Roennback \\& Bergvall (1995) obtained oxygen abundances for 24 H~{\\sc ii} regions in 16 blue LSBGs. de Blok \\& van der Hulst (1998) present measurements of the oxygen abundances of 64 H~{\\sc ii} regions in 12 LSBGs, Kuzio de Naray et al. (2004) reported the oxygen abundances of 16 H~{\\sc ii} regions in 6 LSBGs. Their results show that these H~{\\sc ii} regions in the sample of LSBGs have low metallicities, the 12+log(O/H) is about 8.06 to 8.20 (in median in the samples), which could mean that the LSBGs are metal-poor and unevolved systems.  In this second paper of our series work about the large sample of nearly face-on LSBGs selected from the SDSS-DR4 main galaxy sample (Zhong et al. 2008, Paper I), we will study their spectroscopic properties, in particular, quantities that can be derived from nebular emission line ratios. We will measure nuclear activity, dust extinction, oxygen abundances, nitrogen to oxygen ratios, and examine the stellar mass-metallicity relations, the relations of metallicities vs. $\\mu_0$(B) and stellar masses vs. $\\mu_0$(B). The same property parameters are also obtained for the HSBGs for comparisons. To be in more details, the LSBGs and HSBGs are taken as a whole sample to be divided into four subgroups according to $\\mu_0$(B) for the studies mentioned above.  This paper is organized as follows. In Sect.2, the sample selection criteria are given and the Active Galactic Nucleus (AGN) fractions are estimated from BPT (Baldwin, Phillips, Terlevich 1981) diagram. In Sect.3, we present the dust extinction, diagnostic diagram of emission-line ratios, oxygen abundances and log(N/O) abundance ratios of the sample galaxies. In Sect.4, we present the stellar mass-metallicity relations, and the relations of 12+log(O/H) vs. $\\mu_0$(B), and stellar mass log$M_*$ vs. $\\mu_0$(B) for the sample galaxies. The discussions are given in Sect.\\ 5. In Sect.\\,6, we summarize and conclude the work. Throughout this paper, a cosmological model with  $H_0$=70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega _M$=0.3 and $\\Omega _\\Lambda =0.7$ has been adopted.   \\section[]{The sample}  The parent sample of this work is the 30,333 nearly face-on disc galaxies selected from SDSS-DR4 main galaxy sample by Zhong et al. (2008). Here we match this sample with the emission line catalog of SDSS-DR4 to get a subsample of star-forming galaxies to study their spectral properties, especially metallicities. The fluxes of the emission-lines of the sample galaxies are taken from publication of the MPA/JHU collaboration (the MPA SDSS website\\footnote{http://www.mpa-garching.mpg.de/SDSS/}, Kauffmann et al. 2003a,b; Brinchmann et al. 2004; Tremonti et al. 2004 etc.). These measurements were obtained from the stellar-feature subtracted spectra with the spectral population synthesis code of Bruzual \\& Charlot (2003). The sample selection criteria of this work are given as follows. To be clear, the selection criteria in Zhong et al. (2008) are also simply mentioned.  \\begin{enumerate}  \\item  $fracDev_r$ $<$ 0.25  The parameter $fracDev_r$ indicates the fraction of luminosity contributed by the de Vaucouleurs profile relative to exponential profile in the $r$-band. This means the selected sample almost having an exponential light profile rather than a de Vaucouleurs profile. This can also minimize the effect of bulge light on the disk galaxies.   \\item $b/a$ $>$ 0.75  This is for selecting the nearly face-on disc galaxies, and this is corresponding to the inclination $i<$41.41 degree, where $a$ and $b$ are the semi-major and semi-minor axes of the fitted exponential disk, respectively.  \\item $M_B$ $<$ -18  This  B-band absolute magnitude cut excludes the few dwarf galaxies ($\\sim$6\\%) contained in the sample by the previous two steps of selection.  Up to now, 30,333 nearly face-on disc galaxies are selected. Taking $\\mu_0$(B)=22 mag arcsec$^{-2}$ as the criterion, 12,282 objects are classified as LSBGs and other 18,051 objects as HSBGs. More details can be found in Zhong et al. (2008).  \\item 0.04$<z<$0.25  This redshift range may ensure a covering fraction $>$20\\% of the galaxy light in the SDSS spectral observations, and allow us to get reliable metallicity estimates for the galaxies (Kewley et al. 2005; Liang et al. 2006a). Then 28,148 (92.8\\% of 30,333) of our sample galaxies are within this redshift range, which includes 11,086 (90.3\\% of 12,282) LSBGs and 17,062  (94.5\\% of 18,051) HSBGs. The aperture effects will be specially discussed in Sect.\\,5.  \\item  Emission line selection  We require the fluxes of the emission lines have been measured for [O~{\\sc ii}]3727, H$\\beta$, [O~{\\sc iii}]5007, H$\\alpha$, [N~{\\sc ii}]6583, [S~{\\sc ii}]6717,6731, and they should have higher S/N ratios by requiring the lines of H$\\beta$, H$\\alpha$, and [N~{\\sc ii}]6583 detected at greater than 5$\\sigma$ (following Tremonti et al. 2004 and Liang et al. 2006a). Then, it results in 22,757 (81\\% of 28,148) nearly face-on disc galaxies, which include 7,419 (67\\% of 11,086) LSBGs and 15,338 (90\\% of 17,062) HSBGs. We have checked the effects of the S/N ratios of these three emission lines for the sample selection. If only H$\\alpha$ is detected at greater than 5$\\sigma$, the obtained LSBGs and HSBGs are 8,005 and  15,527, respectively; if only [N~{\\sc ii}]6583 has S/N$>$ 5$\\sigma$, the corresponding numbers are 7,916 and  15,540; and if only H$\\beta$ has S/N$>$ 5$\\sigma$, the corresponding numbers are 7,538 and 15,383. This means that H$\\beta$ plays larger role than H$\\alpha$ and [N~{\\sc ii}] in limiting the sample. The effect of [O~{\\sc iii}]5007 will be further discussed in the later part of this section, and the effect of emission line selection for LSBGs will be further discussed in Sect.\\,5.  \\item Four subgroups according to $\\mu_0$(B)  By following McGaugh (1996) and the discussion in Zhong et al. (2008), we try to further divide our total sample of nearly face-on disc galaxies (22,757) into four subgroups according to their B-band central surface brightness $\\mu_0$(B) (the method for calculating $\\mu_0$(B) can be found in Sect.2.2 of Zhong et al. 2008). We divide LSBGs into ``vLSBGs\" (very Low Surface Brightness Galaxies) and ``iLSBGs\" (intermediate Low Surface Brightness Galaxies), HSBGs into ``iHSBGs\" (intermediate High Surface Brightness Galaxies) and ``vHSBGs\" (very High Surface Brightness Galaxies). These acronyms differ slightly from those adopted by McGaugh et al. (1996). The numbers of galaxies in the four subgroups are:  1,364 (18\\% of 7,419) vLSBGs with 22.75$<$$\\mu_0(B)$$<$24.5 mag arcsec$^{-2}$,  6,055 (82\\% of 7,419) iLSBGs with 22$<$$\\mu_0(B)$$<$22.75 mag arcsec$^{-2}$),  9,107 (59\\% of 15,338) iHSBGs with 21.25$<$$\\mu_0(B)$$<$22 mag arcsec$^{-2}$, and  6,231 (41\\% of 15,338) vHSBGs with $\\mu_0(B)$$<$21.25 mag arcsec$^{-2}$.  Indeed, McGaugh (1996) named these four subgroups of $\\mu_0$(B) bins as LSBGs, ISBGs, HSBGs and VHSBGs consequently.   \\item Star-forming galaxies  For metallicity estimates of galaxies, only star-forming galaxies are selected. The traditional line diagnostic diagram [N~{\\sc ii}]/H$\\alpha$ vs. [O~{\\sc iii}]/H$\\beta$ (BPT: Baldwin, Phillips, Terlevich 1981; Veilleux \\& Osterbrock 1987) is used to separate the star-forming galaxies from AGNs.  We adopt the formula given by Kauffmann et al. (2003a, the solid line in Fig.~\\ref{fig.diag}): $\\rm log([O~III]5007/H\\beta)<0.61/(log([N~II]6583/H\\alpha)-0.05)+1.3$, to select the star-forming galaxies. Kewley et al. (2001) also gave a diagnostic line, which often is assumed as the upper limit for star-forming galaxies (see the dashed line in Fig.~\\ref{fig.diag}).   Finally we select 21,032 (92.4\\%) star-forming galaxies from the 22,757 galaxies having good quality spectral observations. In each of the four subgroups of surface brightness, there are 1,299 (95.2\\% of 1,324) vLSBGs, 5,551 (91.7\\% of 6,055) iLSBGs, 8,310 (91.2\\% of 9,107) iHSBGs, and 5,872 (94.2\\% of 6,231) vHSBGs. These mean that the AGN fractions of these sample galaxies are quite small, less than 9\\%, and vLSBGs has the least fraction of AGNs, $<$5\\%. Fig.~\\ref{fig.diag} presents these results. The AGN fraction here is much lower then that (33\\%) found by Tremonti et al. (2004) for the whole sample of star-forming galaxies. The reason could be that our $fracDev_r$  cut has selected against galaxies with bulges, and we only select the nearly face-on disc galaxies. The bulge mass is very tightly correlated with black hole mass (Kormendy \\& Richstone 1995; Magorrian et al. 1998).   \\end{enumerate}   The median values of the B band central surface brightnesses $\\mu_0(B)$ (in units of mag arcsec$^{-2}$) in the four sub-groups are: 22.96 for vLSBGs, 22.28 for iLSBGs, 21.64 for iHSBGs, and 20.86 for vHSBGs. The corresponding mean values of them are 23.02, 22.31, 21.64 and 20.77, respectively. We will study these star-forming galaxies in the following parts of this work.  The MPA/JHU group has provided the 12+log(O/H) abundances for the SDSS star-forming galaxies as given in Tremonti et al. (2004). For our sample galaxies, we will adopt the oxygen abundances measured by them. The reason will be discussed in Sect.~\\ref{secn2o2r23} in details. Due to the number limits of the sample with reliable oxygen abundance measurements, the numbers of our sample galaxies in each of the four subgroups are then reduced to 601 (46\\% of 1,299) of vLSBGs, 2,630 (47\\% of 5,551) of iLSBGs, 5,517 (66\\% of 8,310) of iHSBGs, and 5,317 (90\\% of 5,872) of vHSBGs.  The reason for the large fraction of reduced numbers here could be mainly due to the S/N cut of [O~{\\sc iii}]5007$>$ 5$\\sigma$, which probably was used in producing the MPA/JHU metallicity catalog, although was not used by Tremonti et al. (2004). We have tried to check this guess. In criterion (v), when we only consider [O~{\\sc iii}]5007$>$ 5$\\sigma$ as the limiting for selecting the sample from S/N ratio, we will obtain 3,440 LSBGs and  11,463 HSBGs, which are much smaller than only considering other three emission lines. Furthermore, if we consider all the four strong emission lines H$\\beta$, H$\\alpha$, [N~{\\sc ii}]6583 and [O~{\\sc iii}]5007 detected at greater than 5$\\sigma$, the resulted sample will be 3,332 LSBGs and 11,386 HSBGs, which are similar to what obtained by only considering [O~{\\sc iii}]5007$>$ 5$\\sigma$ as given above. This means, indeed, [O~{\\sc iii}]5007 plays largest role in limiting the sample. These samples (3,332 and 11,386) are further classified as 3,014 (90.4\\%) star-forming LSBGs and 10,439 (91.6\\%) star-forming HSBGs, which still suggest the small AGN fraction ($<$10\\%) of the sample galaxies. Then, we divided these star-forming galaxies into four subgroups following their surface brightness. There are 567 vLSBGs, 2,447 iLSBGs, 5,181 iHSBGs and 5,258 vHSBGs. These four values are quite similar to the numbers of 601, 2,630, 5,517 and 5,317 given in last paragraph, which are obtained by matching the samples from criteria (i)-(vii) with the metallicity catalog of MPA/JHU, although very slight differences still exist.  The stellar masses of the sample galaxies are taken from the MPA/JHU database as well (Kauffmann et al. 2003b, Gallazzi et al. 2005). To estimate stellar masses of the galaxies, they firstly estimated the dust-correction to the observed $z$-band magnitude of the galaxy, and then the stellar mass was computed by multiplying the dust-corrected luminosity of the galaxy by the stellar mass-to-ligh ratio predicted by model. In their work, the M/L ratios are estimated by also considering the spectral features. But nearly all other methods use colors to estimate M/L. Due to the limits of reliable stellar mass estimates, our sample galaxies having both metallicities and stellar masses are then reduced to 597 of vLSBGs, 2,609 of iLSBGs, 5,486 of iHSBGs, and 5,291 of vHSBGs.  All the corresponding numbers of sample galaxies and their median values of the property parameters are presented in Table~\\ref{tab1}.   \\begin{figure} \\centering \\includegraphics [width=7.5cm, height=6.5cm]{f1a.eps} \\\\ \\vspace{-1.8cm} \\includegraphics [width=7.5cm, height=6.5cm]{f1b.eps} \\\\ \\vspace{-1.8cm} \\includegraphics [width=7.5cm, height=6.5cm]{f1c.eps} \\\\ \\vspace{-1.8cm} \\includegraphics [width=7.5cm, height=6.5cm]{f1d.eps} \\vspace{-1.8cm} \\caption {The diagnostic diagrams for star-forming galaxies and AGNs for the sample galaxies in four bins of surface brightness $\\mu_0(B)$: (a) vLSBGs, (b) iLSBGs, (c) iHSBGs, (d) vHSBGs. The diagnostic lines come from Kauffmann et al. (2003a, the solid line) and  Kewley et al. (2001, the dashed line). The galaxies classified as AGNs fall to the upper right of the lines, and the star-formings fall to the lower left of the lines. We adopt the solid line as the diagnostic line in this work.} \\label{fig.diag} \\end{figure}    \\begin{table} \\begin{center} \\caption{Basic information of the sample galaxies in the four subgroups with $\\mu_0(B)$. All the estimates of parameters are in median. } \\label{tab1} \\scriptsize \\begin{tabular}{lcccc} \\hline &  vLSBG      &    iLSBG     &   iHSBG      &    vHSBG     \\\\  \\hline $\\mu_0(B)$    & 24.5-22.75  & 22.75-22.0   & 22.0-21.25  &  $<$21.25   \\\\ All galaxies  &  1,364      & 6,055       & 9,107       &   6,231    \\\\ Star forming &  1,299      & 5,551       & 8,310       &  5,872     \\\\ AGN fraction &  4.8\\%      &  8.3\\%       & 8.8\\%       & 5.8\\%      \\\\ A$_V$  & 0.46      &  0.63  &     0.76  & 0.83   \\\\    \\hline with O/H    &   &    &   &     \\\\ Star forming &  601      & 2,630       & 5,517       &  5,317     \\\\ 12+log(O/H)  &   8.77   &  8.94    &   9.03 &  9.06 \\\\ log(N/O)      &  -1.05   & -0.94    &   -0.85  & -0.83   \\\\  \\hline with stellar mass    &   &    &   &     \\\\ Star forming &  597      & 2,609       & 5,486       &  5,291     \\\\ log(M$_*$)   &    9.55   &  9.91    &    10.21   & 10.29    \\\\ light\\_fraction   &   0.12   &  0.15    &    0.21   & 0.30    \\\\  \\hline \\end{tabular} \\end{center} $Note: $ {\\footnotesize Raw~(1)-(5) consequently refer to the ranges of $\\mu_0(B)$ in units of mag arcsec$^{-2}$, the numbers of all the sample galaxies; the numbers of star-forming galaxies; AGN fractions and median values of dust extinction. Raw~(7)-(9) refer to the numbers of the galaxies having oxygen abundance estimates, the median values of oxygen abundances and log(N/O) consequently. Raw~(11)-(13) refer to the numbers of the galaxies having both oxygen abundance and stellar mass estimates, median values of stellar masses and light fractions in the fiber observations consequently.} \\end{table} ",
        "conclusions": "In this second paper of our series work about a large sample of LSBGs ($\\mu_0(B)>$22 mag arcsec$^{-2}$) selected from the SDSS-DR4 main galaxy sample (Zhong et al. 2008, the Paper I), which are low-inclination, disk-dominated galaxies, we study the spectroscopic properties of the star-forming sample galaxies including dust extinction, strong emission-line ratios, oxygen abundances, log(N/O) abundance ratios, and the relations of 12+log(O/H) vs. stellar mass,  12+log(O/H) vs. $\\mu_0(B)$, and stellar mass vs. $\\mu_0(B)$. For comparison, a large sample of HSBGs ($\\mu_0(B)<$22 mag arcsec$^{-2}$) is also selected simultaneously and done similar analyses. To be in more details, the entire sample galaxies of LSBGs and HSBGs are further divided into four subgroups according to their  $\\mu_0(B)$ (in units of mag arcsec$^{-2}$): i.e., vLSBGs with 24.5-22.75, iLSBGs with 22.75-22.0, iHSBGs with 22.0-21.25, and vHSBGs with $<$21.25. Their resulted properties are summarized as follows.  \\begin{enumerate}  \\item The AGN fractions of the sample galaxies are small, less than 9\\%, verified by the BPT diagrams from emission-line ratios. The reason could be that our $fracDev_r$  cut has selected against galaxies with bulges. We select the star-forming galaxies for further studies.  \\item LSBGs span a wide range in dust attenuation, metallicity, N/O and stellar mass. The median values of these property parameters all increase with surface brightness as can be seen in Table\\,1. However, these trends can, for the most part be accounted for by the differences in stellar mass among the samples. Most of our sample galaxies have log([N~{\\sc ii}]/ [O~{\\sc ii}])$>-$1.2, which means they belong to the upper branch of $R_{23}$ vs. metallicity relation. Thus our vLSBGs and iLSBGs are not as metal-poor as the H~{\\sc ii} regions in a small sample of LSBGs studied by McGaugh (1994) and some following works (8.06-8.20 of 12+log(O/H) in median). One reason of this discrepancy could be that our sample galaxies are more luminous than theirs intrinsicly. The second reason could be the different calibrations of $R_{23}$ for metallicity, previous studies, such as McGaugh et al. (1994), adopted log([N~{\\sc ii}]/ [O~{\\sc ii}])$>-$1.0 (higher than $-1.2$, what we used) to judge and found most of their samples were in the lower branch of $R_{23}$ vs. metallicity relation, while we adopt the oxygen abundances of the galaxies obtained by using the Bayesian approach by the MPA/JHU group as Tremonti et al. (2004). This could be well explained by Fig.~\\ref{fig.com.oh.lower} for those our sample galaxies having $-$1.2$<$log([N~{\\sc ii}]/[O~{\\sc ii}])$<-$1.0.  \\item The log(N/O) abundances of our sample galaxies are more consistent with the combination of the ``primary\" and ``secondary\" components, but the ``secondary\" component dominates. In the log(N/O) vs. 12+log(O/H) relations of the four subgroup galaxies, the medain values of them at given O/H show slight differences ($<$0.074\\,dex) following surface brightness as shown in Fig.~\\ref{fig.no.oh} and Fig.~\\ref{fig.dis.on.referee}a. These may mean that the sample galaxies may have some but not quite much difference in star formation history, stellar populations, and enrichment history of chemical elements, such as nitrogen and oxygen, as well as the ratio of intermediate-mass to massive stars.  \\item The relations between stellar masses and 12+ log(O/H) for our sample galaxies show that the vLSBGs have less metal-rich and massive ones than other three subgroups, And the iLSBGs have also less such objects than the iHSBGs and vHSBGs. We would like to highlight that the LSBGs in our sample span a wide range of stellar mass and metallicity although having less massive and metal-rich ones; LSBGs are not much more metal poor than HSBGs of the same mass (less than 0.11\\,dex shown by the median-value points in Fig.~\\ref{fig.dis.on.referee}b).  \\item For the vLSBGs and iLSBGs, the galaxies with lower surface brightnesses have lower metallicities, and have smaller stellar masses generally. The iHSBGs also follow this trend slightly, but the vHSBGs do not. These trends could be intrinsicly linked to stellar masses. If the effects of stellar masses are removed by calculating the residuals between the measured oxygen abundances and those calculated from stellar masses using formula, then the four subgroups do not show much difference (from +0.035 to -0.039\\,dex as given in Fig.~\\ref{fig.mu.oh.mass}c). In a word, the metallicity of LSBGs is correlated with $\\mu_0(B)$, but more tightly correlated with stellar mass; the apparent correlation of $\\mu_0(B)$ and O/H is a consequence of the correlation with stellar mass; the residuals of the MZR show very slightly correlation with $\\mu_0(B)$, which means that these LSBGs and HSBGs may have some variation but not quite much in star formation history and chemical enrichment history.  \\end{enumerate}"
    },
    "1004/1004.4776_arXiv.txt": {
        "abstract": "Two new luminosity functions of galaxies can be built starting from three and four parameter generalized gamma distributions. In the astrophysical conversion, the number of parameters increases by one, due to the addition of the overall density of galaxies. A third new galaxy luminosity function is built starting from a three parameter generalized gamma distribution for the mass of galaxies once a simple nonlinear relationship between mass and luminosity is assumed; in this case the number of parameters is five because the overall density of galaxies and a parameter that regulates mass and luminosity are added. The three new galaxy luminosity functions were tested on the Sloan Digital Sky Survey (SDSS) in five different bands; the results always produce a \"better fit\" than the Schechter function. The formalism that has been developed allows to analyze the Schechter function with a transformation of location. A test between theoretical and observed number of galaxies as a function of redshift was done on data extracted from a two-degree field galaxy redshift survey. ",
        "introduction": "The luminosity function of galaxies $ \\Phi (L) dL$ (LF) is the number of galaxies per unit volume, $Mpc^3$, whose luminosity is comprised between $L$ and $L+dl$, see Section 3.7 in \\cite{Longair2008} or Section 1.12 in \\cite{Padmanabhan_III_2002}. The luminosity, $L$, has physical units of $Watt~Hz^{-1}$ and  therefore the band being considered should always be specified. In our case, we selected the data of the  Sloan Digital Sky Survey (SDSS) which has  five bands $u^*$ ($\\lambda = 3550\\AA{})$, $g^*$ ($\\lambda = 4770\\AA{})$, $r^*$ ($\\lambda = 6230\\AA{})$, $i^*$ ($\\lambda = 7620\\AA{})$ and $z^*$ ($\\lambda = 9130\\AA{})$ with  $\\lambda$ denoting the wavelength of the CCD camera, see \\cite{Gunn1998}.  The Schechter LF of galaxies, see \\cite{schechter}, is the more widely used function \\begin{equation} \\Phi (L) dL = (\\frac {\\Phi^*}{L^*}) (\\frac {L}{L^*})^{\\alpha} \\exp \\bigl ( {- \\frac {L}{L^*}} \\bigr ) dL \\quad, \\label{equation_schechter} \\end {equation} where $\\alpha$ sets the slope for low values of $L$, $L^*$ is the characteristic luminosity and $\\Phi^*$ is the normalization. In the formula above, $ L^* $ characterizes the break of  the LF. As an example, $L \\ge 0.1 L^*$ defines a giant  galaxy , see the text after formula (1.18) in \\cite{Gallagher2000}. An astronomical form of equation~(\\ref{equation_schechter}) can be deduced by introducing the distribution in absolute magnitude \\begin{eqnarray} \\Phi (M)dM=&&(0.4  ln 10) \\Phi^* 10^{0.4(\\alpha +1 ) (M^*-M)}\\nonumber\\\\ && \\times \\exp \\bigl ({- 10^{0.4(M^*-M)}} \\bigr)  dM \\quad, \\label{equation_schechter_M} \\end {eqnarray} where $M^*$ is the characteristic magnitude as derived from the data.  This function has three parameters that can be found by fitting the data. Over the years, many modifications have been made to the Schechter LF in order to improve its fit: we report two of them. When the fit of the rich clusters LF is not satisfactory a two-component Schechter-like function is introduced, see~\\cite{Driver1996} \\begin{eqnarray} L_{max} > L > L_{Dwarf} : \\quad \\Phi (L) dL = (\\frac {\\Phi^*}{L^*}) (\\frac {L}{L^*})^{\\alpha} \\exp \\bigl ( {- \\frac {L}{L^*}} \\bigr ) dL \\quad, \\nonumber \\\\ \\\\ L_{Dwarf} > L > L_{min} : \\quad \\Phi (L) dL = (\\frac {\\Phi_{Dwarf}}{L^*}) (\\frac {L}{L_{Dwarf}})^{\\alpha_{ Dwarf}} dL \\quad, \\nonumber \\end {eqnarray} where \\begin{eqnarray} \\Phi_{Dwarf} = \\Phi^* (\\frac {L_{Dwarf} }{L^*})^{\\alpha} \\exp \\bigl ( {- \\frac {L_{Dwarf} }{L^*}} \\bigr ) \\quad. \\nonumber \\end{eqnarray} This two-component function defined between the maximum luminosity, $L_{max}$, and the minimum luminosity, $L_{min}$, has five parameters because two additional parameters have been added: $L_{Dwarf}$ which represents the magnitude where dwarfs first dominate over giants and ${\\alpha_{Dwarf}}$ which regulates the faint slope parameter for the dwarf population.   Another LF introduced in order to fit the case of extremely low luminosity galaxies is the double Schechter function with five parameters, see~\\cite{Blanton_2005}: \\begin{equation} \\Phi(L) dL = \\frac{dL}{L_\\ast} \\exp(-L/L_{\\ast}) \\left[ \\phi_{\\ast,1} \\left( \\frac{L}{L_{\\ast}} \\right)^{\\alpha_1} + \\phi_{\\ast,2} \\left( \\frac{L}{L_{\\ast}} \\right)^{\\alpha_2} \\right] \\quad, \\end{equation} where the parameters $\\Phi^*$ and $\\alpha$ which characterize the Schechter function have been doubled in $\\phi_{\\ast,1}$ and $\\phi_{\\ast,2}$. The strong dependence of LF on different environments such as voids, superclusters and supercluster cores was analyzed by \\cite{Einasto_2009} with \\begin{equation} F (L) \\mathrm{d}L \\propto (L/L^{*})^\\alpha (1 + (L/L^{*})^\\gamma)^{(\\delta-\\alpha)/\\gamma} \\mathrm{d}(L/L^{*}), \\label{eq:abell} \\end{equation} where $\\alpha$ is the exponent at low luminosities $(L/L^{*}) \\ll 1$, $\\delta$ is the exponent at high luminosities $(L/L^{*}) \\gg 1$, $\\gamma$ is a parameter of transition between the two power laws, and $L^{*}$ is the characteristic luminosity. The previous LFs leave a series of questions unanswered or partially answered: \\begin {itemize} \\item What is the function of introducing a transformation of location in the LF? \\item Is it possible to model the data with just a single LF? \\item Is it possible to deduce a LF for galaxies starting from the mass distribution of galaxies? \\item Is it possible to improve the Schechter function by introducing a transformation of location? \\item Does the new LF match the observed behavior of the number of galaxies for a given solid angle and flux as a function of the redshift? \\end {itemize} In Section~\\ref{gammagene}, this paper explores how a generalized gamma distribution can model two galaxy LFs. The method which allows the LF to be deduced for galaxies starting from a mass distribution as given by a generalized gamma is presented in Section~\\ref{masses}. Section~\\ref{location} reports the analytical and numerical results on the Schechter function with transformation of location. In Section~\\ref{secz}, the redshift dependence of the Schechter function and one of the four new LFs are compared with data from the two-degree Field Galaxy Redshift Survey in the 2dFGRS catalogue. ",
        "conclusions": " \\begin{itemize} \\item The generalized gamma galaxy LF with five parameters, formula (\\ref{pdf5magni}), produces a \"better fit\" than the Schechter LF in all five bands considered, see Table~\\ref{dataLF5}. As an example, in the $u^*$ band, $\\chi^2$ and  $\\chi_{red}^2$  are 282 and 0.59 against 330 and 0.68 for the Schechter LF.  \\item The generalized gamma galaxy LF with four parameters, formula (\\ref{pdf4magni}), produces a \"better fit\" than the Schechter LF in all five bands considered, see Table~\\ref{dataLF4}. As an example, in the $g^*$ band, $\\chi^2$ and  $\\chi_{red}^2$  are 747 and 1.263 against  753 and 1.263 for the Schechter LF. Because the number of parameters is less, $\\chi^2$, $AIC$ and $BIC$ are higher than in the case of the generalized gamma galaxy LF with five parameters, see Table~\\ref{dataLF5}. \\item The LF of galaxies obtained from the generalized gamma PDF for the mass of galaxies assuming $L \\propto M^d$, formula~(\\ref{pdfml5}), produces values of $d$ that are slightly greater than one and a  \"better fit\" than the Schechter LF in all five bands considered, see Table~\\ref{dataMLLF5}. As an example, in the $r^*$ band, $\\chi^2$ and  $\\chi_{red}^2$  are 2217 and 3.31 against  2260 and 3.36 for the Schechter LF.   \\end{itemize}  The Schechter LF of galaxies can also be improved by incorporating a transformation of location, formula~(\\ref{schechter_loc4}), and the numerical analysis produces a \"better fit\" than the standard Schechter LF in three out of the five bands analyzed, see Table~\\ref{schlocation}. As an example, in the $z^*$ band, $\\chi^2$ and  $\\chi_{red}^2$  are 3133 and 4.25 against  3245 and 4.4 for the Schechter LF.   One of the four new LFs, more precisely the generalized gamma with four parameters, (equation~(\\ref{lf4})), was tested on the 2dFGRS database and the analysis of $\\chi^2$ on the histogram of observed frequencies versus redshift produces $\\chi^2$=6654 for the new LF against $\\chi^2$=8078 for the Schechter function."
    },
    "1004/1004.1824_arXiv.txt": {
        "abstract": "In this paper we set out to measure time dilation in quasar light curves. In order to detect the effects of time dilation, sets of light curves from two monitoring programmes are used to construct Fourier power spectra covering timescales from 50 days to 28 years.  Data from high and low redshift samples are compared to look for the changes expected from time dilation.  The main result of the paper is that quasar light curves do not show the effects of time dilation.  Several explanations are discussed, including the possibility that time dilation effects are exactly offset by an increase in timescale of variation associated with black hole growth, or that the variations are caused by microlensing in which case time dilation would not be expected. ",
        "introduction": "\\label{sec1}  Time dilation (the stretching of time by a factor of $(1+z)$) is a fundamental property of an expanding universe.  Given the success of the the currently accepted cosmological model, which certainly implies expansion, it is perhaps surprising that more attention has not been paid to making direct measures of time dilation.  This must surely be due in part to the fact that measures of time dilation can tell little or nothing about cosmological parameters within the framework of a Big Bang universe, but only whether or not the Universe is expanding.  Also, it turns out to be surprisingly hard to formulate a conclusive test for time dilation.  What is needed is an event or fluctuation of known rest frame duration which can be observed at sufficiently high redshift with an accuracy which enables the predicted stretching by a factor of $(1+z)$ to be observed.  The study of gamma-ray bursts has generated considerable interest in time dilation.  The uncertainty in the intrinsic timescales of the bursts has has made it difficult to measure time dilation directly, so earlier papers (for example \\cite{d98}) concentrated on demonstrating consistency when time dilation was allowed for.  When redshifts of individual bursts became available more definitive tests were possible \\cite{c01}, but correcting the raw data for selection effects involving an inverse correlation between luminosiy and timescale has made the results hard to interpret with confidence.  Rather than seeing gamma-ray bursts as providing a convincing test of time dilation, it is probably safer to say that if time dilation is a property of the Universe, then observations of gamma-ray bursts are consistent with this.  The light curves of distant supernovae provide a much more promising test for time dilation.  Early work by \\cite{g01} which involved measuring the light curve widths of a sample of distant supernovae covering a wide range of redshifts, provided convincing evidence for the presence of cosmological time dilation.  More recently, \\cite{f05} have examined in detail the spectral evolution of a high redshift supernova, and shown that it is not consistent with no time dilation at a very high confidence level.  In this paper we address the question of whether time dilation is seen in quasar light curves.  We use the light curves of over 800 quasars monitored on timescales from 50 days to 28 years to construct Fourier power spectra for high and low redshift samples, and compare their Spectral Energy Distributions (SEDs) to look for the effects of time dilation. ",
        "conclusions": "\\label{sec7}  In this paper we have used Fourier power spectrum analysis of over 800 quasar light curves to measure timescales of variation at different redshifts.  The expected effects of time dilation are absent, the SEDs at high and low redshift being essentially identical.  There seems to be no explanation for this within the conventional cosmological framework, and so various other possibilities are considered.  These include the idea that the effects of time dilation are exactly offset by an increase in timescale of variation associated with black hole growth. Alternatively, the observed variations could be caused by microlensing, in which case time dilation would not be expected."
    },
    "1004/1004.0319_arXiv.txt": {
        "abstract": "We provide a physical interpretation and explanation of the morphology-density relation for galaxies, drawing on stellar masses, star formation rates, axis ratios and group halo masses from the Sloan Digital Sky Survey.  We first re-cast the classical morphology-density relation in more quantitative terms, using low star formation rate (quiescence) as a proxy for early-type morphology and dark matter halo mass from a group catalog as a proxy for environmental density: for galaxies of a given stellar mass the quiescent fraction is found to increase with increasing dark matter halo mass.  Our novel result is that - at a given stellar mass - quiescent galaxies are significantly flatter in dense environments, implying a higher fraction of disk galaxies.  Supposing that the denser environments differ simply by a higher incidence of quiescent disk galaxies that are structurally similar to star-forming disk galaxies of similar mass, explains simultaneously and quantitatively these quiescence -nvironment and shape-environment relations.  Our findings add considerable weight to the slow removal of gas as the main physical driver of the morphology-density relation, at the expense of other explanations. ",
        "introduction": "It has been known for many decades that galaxy morphology and environment are correlated \\citep{smith35, zwicky42, sandage61}.  In dense environments the fractions of elliptical and S0 galaxies are higher than in low-density environments, at the expense of a decreased fraction of spiral and irregular galaxies \\citep{dressler80, vogt04b}. This trend is universal and persists over a large dynamic range in density \\citep{postman84}.  S0 and spiral galaxies have in common that they have rotationally supported stellar disks.  This suggests that there may be a direct evolutionary link, in the sense that S0 galaxies could be the descendants of spiral galaxies that had their star-formation activity truncated.  The existence of the morphology-density relation (MDR) suggests that this process may be related to the environment \\citep[see, e.g.,][for an overview of the field]{boselli06}.  Such a process generally involves the partial or entire removal of the gaseous interstellar medium from galaxies that become satellites in larger dark matter halos, through interaction with the intergalactic medium.  Fast, rather violent stripping of a galaxies' entire interstellar medium is often considered a viable option \\citep{gunn72, quilis00}, but also more gentle stripping of hot gas outside the cold disk \\citep{larson80, bekki02, mccarthy08}, often called 'starvation' or 'strangulation', and the gradual stripping of neutral gas from the outer parts of disks \\citep[e.g.,][]{chung07} are commonly invoked. In any case, gas-deficient spiral galaxies in clusters do exist \\citep[e.g.,][]{giovanelli85}, which indicates that gas stripping occurs.  Eventually, in all these scenarios, gas stripping removes the fuel for star formation, presumably producing a quiescent galaxy which is similar to a normal spiral galaxy in terms of its structural properties such as its bulge-to-disk ratio.  Early on, \\citet{vandenbergh76} described such gas-free, 'amenic' spirals, and \\citet{balogh98} observed how cluster galaxies have lower star formation rates than field galaxies with the same bulge-to-disk ratio.    However, if S0s occur when a spiral galaxy is stripped from its gas by interaction with the intragalactic medium, how do we explain the ubiquitous presence of S0s outside dense environment?  Moreover, over the years, it has become clear that S0 galaxies differ from the spiral galaxies in several important ways, implying that the S0 population as a whole is not simply a population of disk galaxies with little or no star formation.  S0 galaxies have more pronounced thick disks than spiral galaxies, as was first shown by \\citet{burstein79d}. Furthermore, the bar fraction among S0s is significantly smaller than that among spiral galaxies \\citep{aguerri09, laurikainen09}.  Perhaps the most important difference is that S0s have larger bulge-to-disk ratios than spiral galaxies \\citep{dressler80}.  This is supported by increasingly convincing evidence that the \\citet{tully77} relation for S0 galaxies differs from that of spiral galaxies \\citep[][, M. Williams et al., in prep.]{neistein99, hinz03, bedregal06}. Although some of these differences may be the result of secular evolution and evolving stellar populations after the truncation of star formation, difficulties remain \\citep[see, e.g.,][for a discussion]{dressler84}.  In particular, the large bulges of S0 galaxies cannot be explained.  Besides the disconnect between S0s and spirals, the increased fraction of elliptical galaxies in dense environments \\citep{dressler80} also challenges the idea that gas stripping explains the MDR.  This concern, as well as the different properties of spirals and S0s, could be accommodated by a invoking a process that affects both the structure and gas content of a galaxy.  Tidal interactions with other galaxies or the potential of a large halo can strip an in-falling galaxy from gas, and reduce the stellar disk as well \\citep{moore96}.  There are two reasons why it is surprisingly difficult to interpret the MDR, constrain its origin, and understand its implications for galaxy formation and evolution in general.  First, 'morphology' is a phenomenological parameter that is the combination of several physical quantities (structure, that is, bulge-to-disk ratio or concentration, and star formation activity). Second, many galaxy properties depend on one another, and some of these dependencies are much stronger than the dependency of morphology on environment.  More specifically, although structure and star formation activity depend strongly on one another \\citep[e.g.,][]{kauffmann03b, hogg04, kauffmann06}, these two parameters behave distinctly different as a function of environment. \\citet{kauffmann04} and \\citet{blanton05b} showed that structure depends only weakly on environment, whereas star formation activity, usually as traced by color, decreases significantly from low- to high-density environments \\citep[e.g.,][]{lewis02, hogg03, gomez03, hogg04, balogh04, blanton05b, baldry06, skibba09a}.  Moreover, \\citet{balogh98} showed that galaxies that are similar in structure have lower star formation rates if they are situated in a cluster. The net effect of these correlations is that the morphological mix changes with environment \\citep{vanderwel08a, bamford09, skibba09c}.  To summarize the above, we have an apparent contradiction: on the one hand, the structural properties of galaxies do not show strong environmental dependencies; on the other hand, the relative elliptical, S0 and spiral galaxy fractions clearly do vary with environmental density, and these different types of galaxies have different structural properties.  In this paper, we address this issue.  We explicitly address the question whether gas stripping can explain the MDR, even in the face of the evidence that S0s and spiral galaxies are structurally different, that S0s occur in both low- and high-density environments, and that elliptical galaxies also prefer dense environments.  Perhaps contrary to expectation, the listed evidence does not rule out gas stripping as an explanation for the MDR.  One should distinguish between the S0 population as a whole, and those galaxies that have been affected by environmental processes, giving rise to the MDR.  In other words, the origin of S0 galaxies and the origin of the MDR are not necessarily the same.  This distinction, which was already noted by \\citet{postman84}, will prove to be critical.   In this paper, we examine the shapes of a large sample of galaxies with low star formation activity (quiescent galaxies), as inferred from their spectroscopic properties, in different environments. The sample selection from the Sloan Digital Sky Survey (SDSS) is described in Section \\ref{sample}.  The shape parameter we use is the projected axis ratio. For an individual galaxy this may not contain much information, but for large samples it becomes a powerful diagnostic. We take advantage of the uniform and well-calibrated data set provided by the SDSS, which is ideally suited to disentangle projection effects and the intrinsic shape distribution of galaxies, as has previously been demonstrated by \\citet{vincent05} and \\citet{padilla08}.  Such an approach is complementary to studies of morphology, defined either visually or otherwise, because it makes no assumptions about the connection with galaxy structure.  Because, in addition, we select our sample spectroscopically, our analysis is completely independent of the structural properties of different types of galaxies.  The first goal of this paper is to confirm that the decreased star formation rate in dense environments is reflected as an increased fraction of quiescent galaxies in massive dark matter halos (Section \\ref{secmdr}).  A halo-based description of environment is more physical than estimates of the local surface number density of galaxies, and, moreover, the correlation between galaxy color and halo mass, and not galaxy density, has been demonstrated to be the principal driving factor behind the observed correlation between galaxy color and environment in general \\citep{blanton07b}.  The second goal is to determine whether or not gas stripping can account for the dependence of star formation activity on group mass. We analyze the axis ratio distribution of quiescent galaxies and its dependence on halo mass, in Sections \\ref{disks} and \\ref{mdrdisks}. Gas stripping of spiral galaxies, which presumably leaves their structural properties mostly unchanged, and alternative scenarios invoked to explain the MDR, such as harassment, which does change galaxy structure, will have a different effects on the axis ratio distribution of quiescent galaxies in high-mass halos.  Finally, we discuss our findings in the context of previous work and identify slow gas stripping as the process that is likely responsible for shaping the MDR (Section \\ref{disc}).  This claim is then explicitly shown to be consistent with the different properties of spiral and S0 galaxies, and the increased fraction of elliptical galaxies in dense environments.  \\begin{figure} \\epsscale{1.2} \\plotone{f1.eps} \\caption{Fraction of spectroscopically selected, quiescent (\\textit{black lines and error bars}) and red (\\textit{red lines and error bars}) galaxies with mass $2.5\\times 10^{10} < M_*/\\msol < 10^{11}$ as a function dark matter halo mass from the \\citet{yang07} group catalog.  The error bars indicate standard binomial errors in the fraction (these are henceforth used in all figures which show fractions).  Both the quiescent and red fraction increase with halo mass, and it is apparent that 'quiescence' is a more stringent criterion than 'redness'.} \\label{mdr} \\end{figure} ",
        "conclusions": "We use stellar masses, star formation activity, axis ratios, and group halo masses of galaxies in the SDSS to provide a physical interpretation of the morphology-density relation and its origin.  Our findings are as follows. \\begin{itemize} \\item{The fraction of galaxies with low specific star formation rates (quiescent galaxies) increases with halo mass. This also holds at fixed galaxy mass and for satellite galaxies (Section \\ref{secmdr})}. \\item{Quiescent galaxies in high-mass dark matter halos ($M_{\\rm{halo}} > 2.5\\times 10^{14}\\msol$) are significantly more often disk dominated than quiescent galaxies in lower-mass halos (Section \\ref{disks}).} \\item{This additional population of disk-dominated quiescent galaxies quantitatively matches the increased fraction of quiescent galaxies in high-mass halos (Section \\ref{mdrdisks}).} \\end{itemize} Hence, our findings show that the morphology-density relation arises as a result of the increased fraction of disk-dominated, quiescent galaxies in high-mass halos, at the expense of disk-dominated, star forming galaxies, which are more frequently found in low-mass halos. Other studies (see Section 4.2) provide evidence that slow stripping of the interstellar medium is the most likely explanation for the decreased star formation activity of spiral galaxies in massive groups.  We conclude that the slow stripping of gas from spiral galaxies, which does not strongly alter its structural properties of a galaxy, likely explains the MDR.  These findings are discussed in the context of the rich history of studies on galaxy morphologies and their environmental dependence (Section 4).  In particular, we demonstrate that our conclusions are not incompatible with the suite of evidence that S0 galaxies and spiral galaxies have systematically different properties.  Such evidence has often, and correctly, been invoked to argue that S0 galaxies cannot, generally speaking, be stripped spiral galaxies. Rather, the stripped, quiescent galaxies that drive the morphology-density relation are only a small subset of the entire population quiescent galaxies with disks."
    },
    "1004/1004.0069_arXiv.txt": {
        "abstract": " ",
        "introduction": "Anisotropic turbulence is a fundamental feature in many astro- or geophysical systems. Well known realizations are the turbulent motions in the solar convection zone or the flow of liquid iron in the outer core of the Earth. Both kinds of convectively driven flows are assumed to be responsible for dynamo action. Flow and field are attuned in a complex non-linear system where the flow produces the magnetic field which in turn backreacts on the field producing flow. Direct observations of the sun offer the possibility to examine the details of the complicated interactions between convection and magnetic fields. On the top of the solar convection zone, in the photosphere, where rotational effects are rather unimportant, occasionally strong localized and radial oriented magnetic fields inhibit the convective transport of energy which results in a cooler (and therefore darker) area, a sunspot (see e.g. Weiss, 1990). Numerical simulations of non-linear magnetoconvection confirm that a vertical magnetic field suppresses turbulent motions as well as it considerably reduces the size of the convection cells, a result that has been known from the linear stability analysis of Chandrasekhar (1961). The smaller cells come along with reduced variations of the temperature so that the correlations of velocity and temperature fluctuations, the turbulent heat flux, decrease with increasing field (Cattaneo et al., 2003). Towards the edge of the sunspot the field changes its orientation to a more horizontal direction. As a result the flow pattern changes to small brighter and darker elongated filaments that surround the sunspot (Weiss et al., 2004). Thermodynamic properties like the heat transport and the temperature distribution are undoubtedly linked to the anisotropy induced by the dominant magnetic field component. Closely related to the interaction of convection, magnetic field and turbulent heat transport are questions about the existence of warmer (and therefore brighter) rings around sunspots (Eschrich \\& Krause, 1977; R\\\"udiger \\& Kitchatinov, 2000) or the behavior of (star-) spots in fast rotating stars. In contrast to the sun, observations of fast rotating stars using Doppler-Imaging techniques reveal an accumulation of star spots close to the poles (Strassmeier, 2002, 2006). This phenomenon clearly indicates the different behavior of convectively driven turbulence in fast rotating objects and points out that a different dynamo mechanism is running where differential rotation is of minor importance (Bushby, 2003). Comparable considerations might also be of interest for the Earth because observations of the radial field component at the core mantle boundary as well as highly resolved simulations of the geodynamo show strong localized flux patches whose pattern and pair-like occurrence resemble the behavior of sunspots (Roberts \\& Glatzmaier, 2000; Jackson, 2003).    However, the particular properties of a fast rotating planetary body like the Earth lead to physical conditions that differ significantly from the sun or stellar objects. A small density contrast between the top and the bottom of the fluid outer core and a nearly incompressible fluid result in a small Mach number flow which is essentially influenced by the non-linear back-reaction of a strong magnetic field. Convection in the fluid outer core is mainly affected by two dominant forces, the Lorentz force due to the strong magnetic field and the Coriolis force due to the fast rotation of the Earth. In such a rapidly rotating convection driven dynamo the turbulence is subject to three preferred directions defined by magnetic field, rotation axis and gravity. The arising anisotropies result in a plate like form of convection cells which are aligned along the rotation axis and elongated in the direction of the dominant magnetic field component (Braginsky \\& Meytlis, 1990; St. Pierre, 1996; Matsushima et al., 1999).  Despite the progress in computational and numerical techniques that have been made during the recent years, it is still impossible to simulate turbulent convection in rotating spherical shells with a sufficient resolution so that the resolved scale range of the turbulence is rather restricted (Hollerbach, 2003). Furthermore, anisotropy usually is neglected by an implicit assumption of a large eddy simulation where scalar parameters resemble the turbulent values of the diffusivities. Beside the fact that mostly the diffusivities -- for the reason of numerical stability -- have to be chosen much larger than even the (poorly known) ``real'' turbulent values, this oversimplifying assumption describes an isotropic transport of flux and an isotropic dissipation of energy. This is only justified if either the resolution is high enough so that all dynamical important modes are resolved or if no preferred direction originated by external forces, stratification or boundary conditions exists. Both, in general, are not the case. In order to include anisotropic effects in a more sophisticated way, turbulence models have to be considered where tensorial expressions for the diffusivities take care of the directional dependence of the turbulence and the influence of the unresolved scales is parameterized in terms of the resolved large-scale fields.  Tensor coefficients for the viscous and thermal diffusivity of a fast rotating and a strong field model of the Earth's core have been derived by Phillips \\& Ivers (2001, 2003). They present expressions that describe enhanced diffusion along a dominant  azimuthal magnetic field. Non-diagonal elements are neglected so that diffusion in the horizontal directions due to a coupling between vertical and horizontal components of the turbulent flow is not included. The results are intended for the use in pseudo-spectral codes but at this stage no applications in numerical investigations are available so that the consequences for global simulations remain unknown.  A different approach is examined by Buffett (2003) and Matsushima (2004, 2005). They present a subgridscale (SGS) model which is essentially based  on the self-similarity of the turbulence. By comparison of a highly resolved direct numerical simulation (DNS) with a large eddy simulation based on the SGS model, they show that the SGS model is able to predict appropriate anisotropic heat- and momentum fluxes. An advanced version of the SGS model that includes the Lorentz- and induction terms is introduced by Matsui \\& Buffett (2005). They apply the so called non-linear gradient model, an adaption of the similarity model that additionally is based on the local character of the turbulence. In general, the SGS models coincide rather well with the direct numerical simulations, however, both models are not able to reproduce all details of the spatiotemporal behavior. Furthermore, systematic deviations are obtained in mean quantities, like kinetic and magnetic energy, which are situated between the values from the resolved DNS and the values of a truncated, unresolved DNS.   In order to evaluate the functionality of SGS models it is useful to  further investigate the development of a convectively driven turbulence in view of the dependence on field strength and latitude. Subject matter of the present paper is the connection between turbulence characteristics, that are caused by the average action of the small-scale fluctuations and the turbulent transport of heat which may essentially contribute to the conditions that determine the temperature distribution within the fluid outer core. A simplified Cartesian system is examined where a convectively driven turbulence in a conducting fluid is subject to fast rotation and an externally imposed magnetic field. As we do not intend to reproduce geophysical realities in all details, the presented investigation is based on a highly idealized system where the emphasis lies on the anisotropy inducing effects of rotation and magnetic field. Properties like the exact behavior of the temperature gradient, the equation of state or the influence of curvature are assumed to be of secondary importance. ",
        "conclusions": "Rotating magnetoconvection has been examined for a wide range of parameters and due to the varieties of the results it is nearly impossible to describe the outcome uniformly. The convectively driven flow changes its characteristic properties considerably in dependence of the applied magnetic field and/or rotation rate. The behavior of the cell pattern indicates that in case of a dominant toroidal field convection in a fast rotating system should occur in extremely thin sheet like structures. In the presented simulations the vertical extension of the cells is fixed by the vertical boundary conditions so that the average size of the convection cells in $z$-direction is extended over the whole box domain. It is hardly believable that such structures are stable over the largest possible scales given by the size of the fluid outer core  so that the ``real'' vertical structure of the convection remains dubious and might only be obtained from, at present unrealistic, global high-resolution simulations.    For faster rotation the effects of the magnetic field on the fluid flow via the Lorentz force become more significant. This has been  shown by the behavior of the anisotropy functions $A_{\\mathrm{H}}$ and $A_{\\mathrm{V}}$. Particular attention deserves the vertical anisotropy. The isotropic behavior on the small scales caused by rotation vanishes with the introduction of a magnetic field. As a consequence, the vertical component of the turbulence intensity dominates the flow which leads to an enhanced vertical mixing of the convective layer. Turbulent transport in a rotating system is facilitated through the introduction of a magnetic field.   The turbulent heat flux has been calculated as an example for a diffusive process that occurs on scales that are not resolved in global simulations. Most remarkably, the vertical heat flux at the pole slightly increases with increasing magnetic field strength (if the rotation is fast enough) and remains at a high level even for very strong magnetic fields. Thus, cooling of the Earth's interior is enhanced by the presence of a dominant toroidal field. Comparison of the results from non-rotating magnetoconvection with simulations that involve rotation and a horizontal (Fig.~\\ref{hor_intensity__b}) or a vertically oriented magnetic field (Fig.~\\ref{vert_heatflux_vert_field}), respectively, shows that an enhancement of the vertical heat flux at high field strength only occurs in combination of a horizontally oriented magnetic field, a sufficient fast rotation rate and close to the north-pole. However, the effects of $B_y$ on the vertical heat transport are relatively small compared to the strong quenching for $\\Lambda\\ga 2 (B\\ga 0.3B_{\\mathrm{eq}})$ in case of an imposed $B_z$.  Likewise of interest is the behavior of the meridional heat flux which is always negative, indicating that heat is transported towards the poles. From comparison with slower rotating magnetoconvection we know that the rotational quenching of $\\langle u'_xT'\\rangle$ is weaker than for $\\langle u'_yT'\\rangle$ and $\\langle u'_zT'\\rangle$, so that it might be possible that in a fast rotating system a significant amount of heat is transported towards the poles. Therfore, the poles should be warmer than the equator. This results in a non-radial oriented large-scale temperature gradient which acts as a source for a meridional large-scale flow. Such a meridional flow -- if strong enough -- interacts with the column like large-scale structure of the convective motions and might be important for the understanding of the dynamo mechanism in the Earth's core (see e.g. Sarson \\& Jones (1999), who identified a fluctuating meridional flow as an explanation for the occurrence of dipole reversals). Indeed, global simulations of Olson \\& Christensen (2002) yield a warmer northpole (see their Fig.~5) but it is not clear if this effect is a result of turbulent heat transport. A large scale meridional flow is also found in the solar convection zone, giving rise to important consequences for the solar rotation law (R\\\"udiger et al., 2005).    The distribution of the temperature on the top of the fluid core depends on the heat flow at the core mantle boundary that is usually prescribed as a boundary condition. Seismological measurements indicate a non-uniform heat flow at the core mantle boundary, whose influence has been examined in numerical simulations of Olson \\& Christensen (2002). The lateral and longitudinal variations of the boundary heat flux affect the local behavior of the fluid flow: convection is enhanced (inhibited) where the boundary heat flow is high (low). In case of non-axisymmetric  boundary conditions at the top of the fluid core, the azimuthal heat flux becomes important. The azimuthal heat flux is always negative and increases towards the equator. This leads to a latitude dependent westward transport of heat which should interact with the constraints given by the inhomogeneous thermal conditions at the core mantle boundary.  Most considerations within the presented work rely on the assumption of a dominant toroidal field within the fluid core. This does not have to be necessarily the case. Due to the insulating mantle the toroidal component of the Earth's magnetic field is hidden from any observations. Indeed, most of the authors favor a geodynamo-model of $\\alpha\\Om$-type (Fearn 1998) which is characterized by a dominant azimuthal field. However, estimations of the toroidal field strength from the extrapolation of the dc electric potential near the top of the Earth's mantle to the core mantle boundary result in comparable strength of the toroidal and the poloidal field components (Levy \\& Pearce, 1991). Furthermore, recent simulations of Christensen \\& Aubert (2006) that cover a very large parameter space, yield solutions without significant differential rotation where poloidal and toroidal field strength are comparable ($\\alpha^2$-Dynamo).  Independent from the orientation of the dominant field component, the direction of the horizontal components of the heat flux (towards the poles and westwards, i.e. against the direction of the rotation) is a consequence of the rotation and not caused by the magnetic field which only enhances or reduces the magnitude of the heat flux. Therefore the horizontal components of the heat flux should show a similar sign even in case of comparable  toroidal and poloidal field strength.  The circumstances within a rotating spherical shell are more complicated as they appear in the presented Cartesian model where relevant characteristics have not been concerned for simplicity. In particular, the  combined effects of all three components of the turbulent heat flux cannot be inferred from a simple local model so that the consequences of the common behavior could only be examined by a global simulation. However, the results should at least qualitatively resemble the combined effects of the small scale fluctuations on large scale fields and comparable effects are expected to be present in spherical geometries. To what extend the reported results for the heat flux will endure in simulations in spherical geometry, where also possible larger scale flows will occur, again can only be clarified in highly resolved global simulations."
    },
    "1004/1004.3540_arXiv.txt": {
        "abstract": "We present a consistent three-dimensional model for the head-tail radio galaxy NGC 1265 that explains the complex radio morphology and spectrum by a past passage of the galaxy and radio bubble through a shock wave. Using analytical solutions to the full Riemann problem and hydrodynamical simulations, we study how this passage transformed the plasma bubble into a toroidal vortex ring. Adiabatic compression of the aged electron population causes it to be energized and to emit low surface brightness and steep-spectrum radio emission. The large infall velocity of NGC 1265---which is barely gravitationally bound to the Perseus cluster at its current position---and the low Faraday rotation measure values and variance of the jet strongly argue that this transformation was due to the accretion shock onto Perseus situated roughly at $R_{200}$. Estimating the volume change of the radio bubble enables inferring a shock Mach number of $\\M\\simeq 4.2_{-1.2}^{+0.8}$, a density jump of $3.4_{-0.4}^{+0.2}$, a temperature jump of $6.3_{-2.7}^{+2.5}$, and a pressure jump of $21.5\\pm10.5$ while allowing for uncertainties in the equation of state of the radio plasma and volume of the torus. Extrapolating X-ray profiles, we obtain upper limits on the gas temperature and density in the infalling warm-hot intergalactic medium of $kT\\lesssim0.4$~keV and $n\\lesssim 5\\times 10^{-5}\\,\\rmn{cm}^{-3}$. The orientation of the ellipsoidally shaped radio torus in combination with the direction of the galaxy's head and tail in the plane of the sky is impossible to reconcile with projection effects. Instead, this argues for post-shock shear flows that have been caused by curvature in the shock surface with a characteristic radius of 850~kpc. The energy density of the shear flow corresponds to a turbulent-to-thermal energy density of 14\\%---consistent with cosmological simulations. The shock-injected vorticity might be important in generating and amplifying magnetic fields in galaxy clusters.  We suggest that future polarized radio observations by, e.g., LOFAR of head-tail galaxies can be complementary probes of accretion shocks onto galaxy clusters and are unique in determining their detailed flow properties. ",
        "introduction": "Head-tail radio galaxies show spectacular asymmetric radio morphologies and occur in clusters of galaxies.  The favored interpretation of these head-tail sources is radio jets ejected by an active core of a galaxy. The jets are bent at some angle toward one direction giving rise to a ``head'' structure and fan out at larger distances in a characteristic tail that extends over many tens to hundreds of kpc. If the flow impacting these galaxies is supersonic, ram pressure of the intracluster medium (ICM) causes the jets to bend \\citep{1979Natur.279..770B}; if the flow is trans-sonic with a Mach number $\\M\\sim1$, the thermal pressure gradient of the interstellar medium of these galaxies due to their motion through the ICM causes the bending \\citep{1979ApJ...234..818J}.  In this model, the supersonic inflow is decelerated and heated by a bow shock in front of the galaxy which also generates a turbulent wake that re-accelerates the relativistic particle population in the tail and illuminates the large-scale tail structure of these sources. The Perseus cluster is an outstanding example, as it hosts many of these head-tail sources: most prominently NGC 1265 and IC 310 \\citep{1968MNRAS.138....1R, 1975A&A....38..381M}. High-resolution observations of the head structure of NGC 1265 reveal two radio jets emerging from the galaxy which are at an angle of $90\\degr$ to the tail where the two jets are first parallel to each other and then merge \\citep{1986ApJ...301..841O}. \\citet{1979A&A....76..109G} found an extension of this radio tail to the north-east. Surprisingly, at lower surface brightness, \\citet{1998A&A...331..901S} find an additional large scale structure that arches toward the east around the steep spectrum tail of NGC 1265 and forms a closed ring with a very steep, but constant spectral index across. Using rotation measure (RM) synthesis, \\citet{2005A&A...441..931D} find a very weak diffuse polarized structure with a Faraday depth of approximately $50~\\rmn{rad\\,m}^{-2}$ at the angular position coincident with the steep spectrum tail as well as the ring structure of NGC 1265. Follow-up mosaic observations of the area around the Perseus cluster show that most if not all of this emission is of Galactic origin \\citep{2011A&A...526A...9B}.  \\citet{1998A&A...331..901S} discuss a total of four models to explain this puzzling observation. They immediately discard two models---one model of chance superposition of several independent head-tail galaxies due to the lack of strong radio sources in this field (Model 1) as well as another model that hypothesizes reacceleration of mildly relativistic electrons in the turbulent wake of a galaxy due to contrived projection probabilities and implausible energetics (Model 2).  The previously favored pair of models discussed in \\citet{1998A&A...331..901S}---apparently simple---do in fact appear to have considerable amount of fine-tuning and have severe physical problems associated with them as we will point out here.  Model 3: in this model, they postulate a helical cluster wind that has to be aligned with the line-of-sight (LOS) to produce the observed ring structure in projection: such a configuration is highly unlikely, and this model has problems in explaining the constancy of the spectrum and the surface brightness along the radio ring. To balance the synchrotron and inverse Compton cooling of an electron population, this model would need to postulate a fine-tuned acceleration process that however must not fan out the well-confined radio emission along the arc. Model 4: in this model, the ``radio tail'' would outline the ballistic orbit of NGC 1265. To sustain such a helical orbit of NGC 1265 over $360\\degr$, this model would require an undetected dark object of mass $M\\gtrsim M_\\rmn{NGC~1265}\\simeq 3\\times10^{12}M_\\odot$ orbiting the galaxy.\\footnote{Since the velocity dispersion of NGC 1265 is unknown, we use the Faber-Jackson relation in combination with the $M_\\rmn{bh}$--$\\sigma$ relation \\citep{2002ApJ...574..740T} to obtain the stellar mass of NGC 1265 of $M_\\ast\\simeq 2.5\\times10^{11}\\,M_\\odot$. We estimate the total mass by using the universal baryon fraction and a stellar-to-baryon mass yield of 0.5 so that the resulting halos mass should be considered a lower limit to the true one.} The change of the direction in the brighter part of the tail remains unexplainable and there is also the problem in explaining the constancy of the spectrum and the surface brightness along the radio ring.  Given these difficulties, it is attractive to consider alternative possibilities that may explore the interaction between a radio galaxy with the outskirts of the Perseus ICM. This has been foreshadowed in a remark by \\citet{2005A&A...441..931D} who speculate whether the large scale polarized structure that arches around the steep spectrum tail of NGC 1265 is indeed the remains of an earlier phase of feedback. Our work will demonstrate that this picture is not only the simplest consistent explanation for the radio morphology and spectrum, but we also use it to indirectly infer the presence of a cluster shock wave and measure its properties.  Previously, the morphology of a giant radio galaxy has already been used to indirectly detect a large scale shock at an intersecting filament of galaxies \\citep{2001ApJ...549L..39E} by using the radio galaxy as a giant cluster weather station \\citep{1998Sci...280..400B}. Gravitationally driven, supersonic flows of intergalactic gas follow these filaments toward clusters of galaxies that represent the knots of the cosmic web \\citep{1996Natur.380..603B}.  The flows will inevitably collide and form large-scale shock waves \\citep{1998ApJ...502..518Q, 2000ApJ...542..608M, 2003ApJ...593..599R, 2006MNRAS.367..113P}. Of great interest is the subclass of accretion shocks that are thought to heat the baryons of the warm-hot intergalactic medium (IGM) when they are accreted onto a galaxy cluster. Formation shocks have also been proposed as possible generation sites of intergalactic magnetic fields \\citep{1997ApJ...480..481K, 1998A&A...335...19R, 2008Sci...320..909R}. Prior to this work, only discrete merger shock waves have been detected in the X-rays \\citep[e.g.,][]{2002ApJ...567L..27M} and there was no characterization possible of the detailed flow properties in the post-shock regime as to test whether shear flows---the necessary condition for generating magnetic fields---are present.  The structure of the paper is as follows. In Section~\\ref{sec:idea}, we present the basic picture of our model in a nutshell while we derive the detailed three-dimensional (3D) geometry of NGC 1265 within our model in Section~\\ref{sec:geometry}. In Section~\\ref{sec:shock}, we work out the properties of the accretion shock onto the Perseus cluster including those of the post-shock flow and present the implications for the infalling warm-hot IGM.  In Section~\\ref{sec:model}, we carefully examine the underlying physics as well as the hydrodynamic stability of our model and discuss our findings in Section~\\ref{sec:conclusions}.  Throughout this work, we use a Hubble constant of $H_{0} = 70\\,\\rmn{km~s}^{-1}\\rmn{Mpc}^{-1}$. For the currently favored $\\Lambda$CDM cosmology with the present day density of total matter, $\\Omega_m=0.28$, and the cosmological constant, $\\Omega_\\Lambda=0.72$, we obtain an angular diameter distance to Perseus ($z=0.0179$) of $D_\\rmn{ang}=75\\,\\rmn{Mpc}$; at this distance, $1\\arcmin$ corresponds to $21.8\\,\\rmn{kpc}$. X-ray data estimates a virial radius and mass for Perseus\\footnote{We define the virial mass $M_\\Delta$ and the virial radius $R_\\Delta$ as the mass and radius of a sphere enclosing a mean density that is $\\Delta=200$ times the critical density of the Universe.}  of $R_{200}=1.9\\,\\rmn{Mpc}$ and $M_{200}=7.7\\times10^{14}M_\\odot$ \\citep{2002ApJ...567..716R}. ",
        "conclusions": "\\label{sec:conclusions}  In providing a quantitative 3D model for NGC 1265, we provide conclusive observational evidence for an accretion shock onto a galaxy cluster.  The accretion shock is characterized by an intermediate-strength shock with a Mach number of $\\M\\simeq 4.2_{-1.2}^{+0.8}$. In the context of the IGM, it appears that even weak to intermediate strength shocks of $\\M \\sim 2.3$ are able to accelerate electrons as the radio relic in A521 suggests \\citep{2008A&A...486..347G}.  This supports the view that these intermediate strength shocks are very important in understanding non-thermal processes in clusters; in particular as they are much more abundant than high-Mach number shocks \\citep[e.g.,][]{2003ApJ...593..599R, 2006MNRAS.367..113P}. The presented technique offers a novel way to find large-scale formation shock waves that would likely be missed if we only saw giant radio relics. These are thought to be due to shock-accelerated relativistic electrons at merging or accretion shock waves \\citep{1998A&A...332..395E, 2006Sci...314..791B}, whereas the discrimination of the two types of shock waves is difficult on theoretical as well as observational grounds \\citep{2006Sci...314.772E}.  The rich morphology of the large scale radio structure that arches around steep-spectrum tail of NGC 1265 enables a qualitative argument for the need of shear to explain the orientation of its ellipsoidally shaped torus.  The energy density of the shear flow corresponds to a turbulent-to-thermal energy density of 14\\%---consistent with estimates in cosmological simulations.  However, simulations are needed to quantify the effect of magnetic helicity internal to the bubble on the torus morphology. In any case, the presence of post-shock shear implies the amplification of weaker seed magnetic fields through shearing motions as well as the generation of magnetic fields through the Biermann battery mechanism if the vorticity has been generated at the curved shock. Even more interesting, the shock passage of multiphase gas with a large density contrast necessarily yields a multiply curved shock surface that causes vorticity injection into later accreted gas on the corresponding curvature scales. Hence, shearing motions are superposed on various scales which implies amplification of the magnetic fields on these scales.  More sensitive polarized and total intensity observations of NGC 1265 are needed to confirm our model predictions with radial polarization vectors of the radio torus and a polarization degree of around 5\\%. In addition, future observations of different systems similar to NGC 1265 would be beneficial to get a statistical measurement of properties of the accretions shocks and its pre-shock conditions. This provides indirect evidence for the existence of the warm-hot IGM.  Our work shows the potential of these kind of serendipitous events as ways to explore the outer fringes of clusters and dynamical features of accretion shocks that are complementary to X-ray observations."
    },
    "1004/1004.4683_arXiv.txt": {
        "abstract": "We present a statistically complete sample of very X-ray luminous galaxy clusters detected in the MAssive Cluster Survey (MACS). This second MACS release comprises all 34 MACS clusters with nominal X-ray fluxes in excess of $2\\times 10^{-12}$ erg s$^{-1}$ cm$^{-2}$ (0.1--2.4 keV) in the ROSAT Bright Source Catalogue; two thirds of them are new discoveries. Extending over the redshift range from 0.3 to 0.5, this subset complements the complete sample of the 12 most distant MACS clusters ($z>0.5$) published in 2007 and further exemplifies the efficacy of X-ray selection for the compilation of samples of intrinsically massive galaxy clusters. Extensive follow-up observations with Chandra/ACIS led to three additional MACS cluster candidates being eliminated as (predominantly) X-ray point sources. For another four clusters --- which, however, remain in our sample of 34 --- the point-source contamination was found to be about 50\\%. The median X-ray luminosity of $1.3\\times 10^{45}$ erg s$^{-1}$ (0.1--2.4 keV, Chandra, within $r_{\\rm 500}$) of the clusters in this subsample demonstrates the power of the MACS survey strategy to find the most extreme and rarest clusters out to significant redshift. A comparison of the optical and X-ray data for all clusters in this release finds a wide range of morphologies with no obvious bias in favour of either relaxed or merging systems. ",
        "introduction": "Clusters of galaxies have long been recognised to offer exceptional opportunities for cosmological and astrophysical studies of remarkable diversity.  As the largest gravitationally bound entities in the universe, they are rare objects, originating from extreme overdensities in the primoridal density field, and growing through continuous accretion as well as serial mergers into mass concentrations of $10^{14-15}$ $M_{\\odot}$ at the present epoch.  While in-depth studies of individual clusters are invaluable for our understanding of the physical processes governing the interaction between the three principal cluster constituents (dark matter, intracluster gas, and galaxies), it is only through observations of well defined, large samples of clusters that we can (a) obtain statistically meaningful information about the properties of clusters as a class of objects evolving on cosmological timescales, and (b) hope to find exceptional systems, such as, e.g., the Bullet Cluster (Markevitch et al.\\ 2004; Brada\\v{c} 2006; Clowe et al,\\ 2006), that enable us to conduct quantitative measurements of fundamental astrophysical parameters. For many decades, such studies had to be based on optically selected cluster samples which, while large, have the distinct disadvantage of being inherently affected and biased by projection effects (van Haarlem, Frenk \\& White 1997; Hicks et al.\\ 2008).   A nearly unbiased way of selecting statistical cluster samples is through X-ray surveys, as the X-ray emission, which originates from the diffuse intra-cluster gas trapped in the clusters' gravitational potential well and heated to virial temperatures of typically $10^{7-8}$ K, represents direct proof of the existence of a three-dimensionally bound system. Also, the X-ray emission is much more peaked at the cluster centre than is the projected galaxy distribution, making projection effects in X-ray selected cluster samples highly improbable.   \\begin{figure*} \\parbox{0.49\\textwidth}{ \\includegraphics[width=0.49\\textwidth]{pdcs-z}} \\parbox{0.5\\textwidth}{ \\includegraphics[width=0.51\\textwidth]{macs-z}}\\hfill \\caption{Histograms of galaxy redshifts in the fields of the nine optically selected systems of the PDCS (left; Oke et al.\\ 1998) and the twelve most distant X-ray selected MACS clusters (right; Ebeling et al.\\ 2007). For ease of comparison the MACS data are shown over the same redshift range and with the same binning as used in the published PDCS figure. Note that both surveys used similar criteria to select galaxies for spectroscopic follow-up observation.\\label{fig:optvsx}} \\end{figure*}  The advantage of X-ray cluster surveys over optical surveys is illustrated in Fig.~\\ref{fig:optvsx} which compares the redshift histograms of the nine optically selected systems of the Palomar Distant Cluster Survey (PDCS; Oke, Postman \\& Lubin, 1998) with those of the twelve most distant X-ray selected MACS clusters (Ebeling et al.\\ 2007). The severe contamination by fore- and background structures seen in projection in the PDCS is endemic in optically selected cluster samples. Pure projection effects like, e.g., CL0231+0048 (Fig.~\\ref{fig:optvsx}, left) can be largely eliminated by including information on galaxy colors or redshifts (photometric or spectroscopic) in the original cluster detection phase. However, even the latest, state-of-the-art optical cluster samples remain biased, as they are prone to select intrinsically poor systems whose apparently compact cluster core, high optical richness, and high velocity dispersion are inflated by line-of-sight alignment and infall (Hicks et al.\\ 2008; Horesh et al.\\ 2009). By contrast, X-ray selected cluster samples are almost entirely free of projection effects since they, by virtue of the X-ray selection criteria, comprise exclusively intrinsically massive, gravitationally collapsed systems. ",
        "conclusions": "We present the second statistically complete MACS subsample, comprising the 34 clusters with X-ray detect fluxes in excess of $2\\times 10^{-12}$ erg s$^{-1}$ cm$^{-2}$ (0.1--2.4 keV) in the RASS Bright Source Catalogue. All clusters feature redshifts of $z\\ge 0.3$, and 22 of the 34 are new discoveries. Chandra observations of the entire sample confirmed the cluster origin of the emission and allowed the elimination of three additional candidates whose X-ray emission was found to be dominated by point sources.  A comparison of the appearance of MACS clusters in the RASS and in Chandra observations confirms that all but the most disturbed clusters at $z>0.3$ appear point-like at the angular resolution of the RASS. We find the original RASS count rates as listed in the BSC to be less accurate than manual measurements within the BSC aperture and using a local annulus for background subtraction. RASS fluxes based on recomputed count rates are in good agreement with the respective Chandra values, except for four clusters for which the RASS count rate is significantly contaminated by point sources within the BSC detect cell, and one system whose very extended emission is not fully captured by the Chandra measurement. For the remainder of the sample, X-ray point sources contribute, on average, only about 3\\% to the flux within $r_{\\rm 500}$. All clusters of this second MACS subsample feature X-ray luminosities (within $r_{\\rm 500}$) in excess of $4\\times 10^{44}$ erg s$^{-1}$ (0.1--2.4 keV) after correction for X-ray point sources, and are thus considerably more X-ray luminous than the Coma cluster. The sample's median X-ray luminosity of $1.3\\times 10^{45}$ erg s$^{-1}$ confirms the efficiency of our survey technique to identify massive clusters well beyond the redshift limits of previous RASS-based cluster surveys. A first assessment of the optical and X-ray morphology of the clusters in this sample finds both fully virialized and heavily disturbed systems to be well represented, arguing against a strong bias in favour of either cool-core clusters or extreme mergers. A more quantitative analysis and discussion of the morphology and relaxation state of MACS clusters will be presented in a forthcoming paper.  When combined with the most X-ray luminous clusters in the local Universe ($z<0.3$) from the eBCS and REFLEX surveys (Ebeling et al.\\ 1998, 2000; B\\\"ohringer et al.\\ 2004) and the complete set of the 12 most distant MACS clusters ($z>0.5$) released earlier (Ebeling et al.\\ 2007), the sample presented here allows cosmological and astrophysical studies of the properties and evolution of the most massive galaxy clusters over a contiguous redshift range from $z\\sim 0$ to $z=0.7$ (Allen et al.\\ 2003, 2004, 2008; Ebeling et al.\\ 2009; Mantz et al.\\ 2008, 2010a,b, Rapetti et al.\\ 2009). For convenience, we have supplied a tabulated version of the appropriate selection function."
    },
    "1004/1004.1284_arXiv.txt": {
        "abstract": "{ We study quasinormal modes and scattering properties via calculation of the $S$-matrix for scalar and electromagnetic fields propagating in the background of spherically and axially symmetric, traversable Lorentzian wormholes of a generic shape. Such wormholes are described by the Morris-Thorne ansatz and its axially symmetric generalization. The properties of quasinormal ringing and scattering are shown to be determined by the behavior of the wormhole's shape function $b(r)$ and shift factor $\\Phi(r)$ \\emph{near the throat}. In particular, wormholes with the shape function $b(r)$, such that $b'(r) \\approx 1$, have very long-lived quasinormal modes in the spectrum. We have proved that the axially symmetric traversable Lorentzian wormholes, unlike black holes and other compact rotating objects, do not allow for superradiance. As a by product we have shown that the 6th order WKB formula used for scattering problems of black or wormholes provides high accuracy and thus can be used for quite accurate calculations of the Hawking radiation processes around various black holes.} }  \\begin{document}  \\PRDonly{ \\title{Passage of radiation through wormholes of arbitrary shape} \\author{R. A. Konoplya}\\email{konoplya@tap.scphys.kyoto-u.ac.jp} \\affiliation{Department of Physics, Kyoto University, Kyoto 606-8501, Japan  \\\\ \\& \\\\ Theoretical Astrophysics, Eberhard-Karls University of T\\\"{u}bingen, T\\\"{u}bingen 72076, Germany} \\author{A. Zhidenko}\\email{zhidenko@fma.if.usp.br} \\affiliation{Instituto de F\\'{\\i}sica, Universidade de S\\~{a}o Paulo \\\\ C.P. 66318, 05315-970, S\\~{a}o Paulo-SP, Brazil}  \\begin{abstract} We study quasinormal modes and scattering properties via calculation of $S$-matrix for scalar and electromagnetic fields propagating in the background of spherically symmetric and axially symmetric traversable Lorentzian wormholes of a generic shape. Such wormholes are described by the general Morris-Thorne ansatz. The properties of quasinormal ringing and scattering are shown to be determined by the behavior of the wormhole's shape function $b(r)$ and shift-factor $\\Phi(r)$ \\emph{near the throat}. In particular, wormholes with the shape function $b(r)$ such that $b'(r) \\approx 1$, have very long-lived quasinormal modes in the spectrum. We have proved that the axially symmetric traversable Lorentzian wormholes, unlike black holes and other compact rotating objects, does not allows for super-radiance. As a by product we have shown that the 6-th order WKB formula used for scattering problems of black or wormholes gives quite high accuracy and thus can be used for quite accurate calculations of the Hawking radiation processes around various black holes. ",
        "introduction": "A wormhole is usually a kind of ``shortcut'' through space-time connecting distant regions of the same universe or two universes. General Relativity though does not forbid wormholes, requires presence of some exotic matter for their existence. A wide class of wormholes respecting the strong equivalence principle is called \\emph{Lorentzian wormholes}, whose space-time, unlike that of \\emph{Euclidean wormholes}, is locally a Lorentzian manifold. A physically interesting subclass of Lorentzian wormholes consists of \\emph{traversable wormholes} through which a body can pass for a finite time. This last class of wormholes will be probed here by the passage of test fields in the wormhole vicinity and observing their response in the form of the waves scattering and quasinormal ringing. However, before discussing the scattering processes, let us briefly review some basic features and observational opportunities of wormholes.  The static spherically symmetric Lorentzian traversable wormholes can be modeled by a generic Morris-Thorne metric \\cite{Morris-Thorne}, which is fully determined by the two functions: the shape function $b(r)$ and the redshift function $\\Phi(r)$. The axially symmetric traversable wormholes can be described by generalization of the Morris-Thorne ansatz, suggested by Teo \\cite{Teo:1998dp}. In order to support such a wormhole gravitationally stable against small perturbations of the space-time one needs an exotic kind of matter which would violate the weak energy condition. The latter requires that for every future-pointing timelike vector field, the matter density observed by the corresponding observers is always non-negative. Thus for supporting a traversable wormhole one needs a source of matter with negative density such as, for example, \\emph{dark energy} \\cite{Ida:1999an}. Wormholes might appear as a result of inflation of primordial microscopic wormholes to a macroscopic size during the inflationary phase of the Universe. If wormholes exist, they could appear as a bubble through which one could see new stars. Massive, rotating wormholes surrounded by some luminous matter could have a visible ``accretion disk'' near them, similar to that surrounding supermassive black holes. Thus, gravitational lensing should give some opportunities for detecting wormholes \\cite{Nandi:2006ds,Safonova:2001vz}. General features of various wormholes were investigated in a great number of papers and we would like to refer the reader to only a recent few, where further references can be found \\cite{Balakin:2010ar,Richarte:2010bd,Bandyopadhyay:2009ek,Bronnikov:2010hu,Usmani:2010cd,Sarbach:2010we,Mimoso:2010yp,Hayward:2002sz,Lobo:2009du}.   As we mentioned before, stability of wormholes is the crucial criterium for their existence \\cite{Gonzalez:2008wd,Gonzalez:2008xk,Doroshkevich:2008xm,Amendo}. In order to judge about a wormholes stability, one has to treat the complete set of the Einstein equations with some non vanishing energy momentum tensor representing the matter source. Thereby, stability depends on particular geometry of the wormhole metric as well as on the equation of state for matter. In this paper, we shall set up a more generic problem: implying the existence of some matter that provides stability of generic traversable wormholes, whose metric is given by some generic redshift $\\Phi(r)$ and shape $b(r)$ functions; we are interested in scattering processes of test fields in the vicinity of such wormholes. In other words, we shall consider the passage of test radiation, which does not influence the wormhole geometry, through wormholes of various shapes. Thus, our aim is to learn how the wormhole shape correlates with the scattering properties of various fields in its background. This is a kind of probing of a wormhole's geometry by test fields.   The scattering phenomenon around wormholes can be descried by a number of characteristics, two of which, the \\emph{quasinormal modes} and the \\emph{\\emph{S-matrix}}, we shall consider here. The quasinormal modes are proper oscillation frequencies, which appear in a black hole's response to initial perturbation. The quasinormal modes of black holes do not depend upon the way of the mode's excitation but only on the black hole parameters and are, thereby, ``fingerprints'' of a black hole. It can be shown \\cite{Konoplya:2005et} that the response of a wormhole to initial perturbations also is dominated by the quasinormal modes. The quasinormal modes are defined mathematically in the same way for wormholes as they are for the black holes \\cite{Konoplya:2005et}. The S-matrix for such problems can be fully described by the wave's reflection/transmission coefficients. We shall show that both the quasinormal modes and reflection coefficients of traversable wormholes are determined by the geometry of a wormhole near its throat. The geometry far from the throat is insignificant for scattering in the same manner as it happens for black holes \\cite{Konoplya:2006ar}.  Rotating black holes are known to be spending their rotational energy on amplification of incident waves of perturbation \\cite{Starobinsky1}. This phenomenon occurs also for various rotating compact bodies, such as, for example, conducting cylinders, and is called superradiance \\cite{Zeldovich,Starobinsky1,Starobinsky2}. When considering rotating traversable wormholes, one could probably expect that the same superradiance should take place. However, we shall show here that rotating axially symmetric traversable wormholes do not allow for the superradiance.  The paper is organized as follows: Sec. II gives some basic information about traversable wormholes. Sec. III is devoted to quasinormal modes of various traversable wormholes. Sec. IV is about correlation of the scattering coefficients with a wormhole's shape. Sec. V considers scattering around rotating traversable axially symmetric wormholes. Finally, in the Appendix we demonstrate the utility of the 6th order WKB formula for calculations of the graybody factors of black holes, which gives quite accurate ``automatized'' formulas for calculations of the intensity of the Hawking radiation. ",
        "conclusions": "We have considered the propagation of test scalar and Maxwell fields in the background of the traversable spherically symmetric and axially symmetric wormholes of a generic form. The majority of our conclusions are independent of the concrete shape of a wormhole. We have found that when the shape function $b(r)$ approaches the regime $b'(r) =1$ at the wormhole's throat $b=b_0$, the quasinormal modes asymptotically approache pure real nondecaying modes, meaning the appearance of long-lived modes. The latter might give much better observational opportunities. The reflection coefficient approaches the step function which depends on the values of the shape and redshift function at the throat. We have proved that independently on their particular forms, traversable axially symmetric wormholes do not allow for superradiance. We were limited here by wormholes whose shape is symmetric with respect to their center (throat). Spherically (but not axially) symmetric wormholes are also necessarily symmetric with respect to the throat."
    },
    "1004/1004.4468_arXiv.txt": {
        "abstract": "{There is observational evidence of propagating kink waves driven by photospheric motions. These disturbances, interpreted as kink magnetohydrodynamic (MHD) waves are attenuated as they propagate upwards in the solar corona.}{To show that resonant absorption provides a simple explanation to the spatial damping of these waves.}{Kink MHD waves are studied using a cylindrical model of solar magnetic flux tubes which includes a non-uniform layer at the tube boundary. Assuming that the frequency is real and the longitudinal wavenumber complex, the damping length and damping per wavelength produced by resonant absorption are analytically calculated in the thin tube (TT) approximation, valid for coronal waves. This assumption is relaxed in the case of chromospheric tube waves and filament thread waves.}{The damping length of propagating kink waves due resonant absorption is a monotonically decreasing function of frequency.  For kink waves with low frequencies the damping length is exactly inversely proportional to frequency and we denote this as the TGV relation. When moving to high frequencies the TGV relation continues to be an exceptionally good approximation of the actual dependency of the damping length on frequency. This dependency means that resonant absorption is selective as it favours low frequency waves and can efficiently remove high frequency waves from a broad band spectrum of kink waves. The efficiency of the damping due to resonant absorption depends on the properties of the equilibrium model, in particular on the width of the non-uniform layer and the steepness of the variation of the local Alfv\\'{e}n speed.}{Resonant absorption is an effective mechanism for the spatial damping of propagating kink waves. It is selective as the damping length is inversely proportional to frequency so that the damping becomes more severe with increasing frequency. This means that radial inhomogeneity can cause solar waveguides to be a natural low-pass filter for broadband disturbances. Hence kink wave trains travelling along, e.g., coronal loops, will have a greater proportion of the high frequency components dissipated lower down in the atmosphere. This could have important consequences with respect to the spatial distribution of wave heating in the solar atmosphere.} ",
        "introduction": "The first observations of post-flare transversal coronal loop oscillations by the Transition Region and Coronal Explorer (TRACE) \\citep[e.g.,][]{aschetal99,nakaetal99,aschetal02}, inspired much development in magnetohydrodynamic (MHD) wave theory. This observational breakthrough was important since estimated wave parameters, such as frequency and amplitude allowed us to implement magnetoseismological techniques to probe the plasma fine structure of the Sun's atmosphere, an idea initially proposed by e.g., \\citet{uchida1970} and \\citet{robetal84}. It is now commonly accepted that these transversal waves are the kink mode from MHD wave theory \\citep[see e.g.,][]{edrob83}, a highly magnetically dominated, i.e., Alfv\\'{e}nic wave \\citep[see][for a discussion on the nature of kink waves]{goossensetal09}. The observed post-flare kink waves in coronal loops have two main defining characteristics, firstly they are standing modes and secondly, they are strongly damped oscillations \\citep[in about 1-4 periods, see e.g.,][]{aschetal03}. Initially there were several physical mechanisms proposed to explain the observed damping, e.g., footpoint leakage \\citep{bergbruyne95,depontmart01}, phase mixing \\citep{heypri83,roberts02,ofasch02} and resonant absorption \\citep{rudrob02,goossetal02} and more recently loop cooling \\citep{mortonerd09}. Thus far, resonant absorption, caused by plasma inhomogeneity in the direction transverse to the magnetic field, has proved the most likely candidate in explaining the observed short damping times in coronal loops \\citep[see][for review]{goossens08}. Consistent seismological studies based on resonant absorption using the observed values of periods and damping times of standing kink waves were carried out by \\citet{arreguietal07} and \\citet{goossetal08}. These two studies show that, at least for a collection of 11 loops, resonant absorption provides an explanation of the observed periods and damping times. The observational signatures of the alternative cooling loop damping mechanism proposed by \\citet{mortonerd09}, differ from resonant absorption by the fact that frequency changes as a function of time (the resonant damping theory developed so far is restricted to static equilibria and therefore frequency is expected to be constant in time). Extensive MHD modelling has also shown that resonant absorption is the most likely explanation for the damping of transverse oscillations in prominence fine structures \\citep[see][]{arretal08,soleretal09b,soleretal09c,arrball10}.      If indeed, coronal loop kink oscillations are being attenuated by the process of resonant absorption, one can exploit this to estimate the transverse plasma inhomogeneity length scales using observed frequencies and damping times \\citep{goossensetal06}. The original equilibrium models by e.g., \\citet{rudrob02}, to study resonant absorption consisted of monolithic loop structures. However, there has been some observational evidence by e.g., \\citet{aschw05} that coronal loops are composed of many different strands, possibly at different temperatures and densities. To model this loop multi-thread structure, there was an initial numerical study by \\citet{terrarretal08} and it was found that the process of resonant absorption was still an efficient damping mechanism in more complex and realistically structured loop models. Further study into the properties of standing kink waves was also undertaken relating plasma inhomogeneity in the direction of the magnetic field caused by e.g., density \\citep{diazetal04,andriesetal05,arretal05,dymova06,erdverth07,verthetal07} or magnetic \\citep{vertherd08,ruderetal08} stratification. It was shown that eigenfrequencies and eigenfunctions of coronal loops with longitudinal stratification, were altered in such a way that one could determine, e.g., the coronal density scale height by estimating the ratio of the fundamental mode to that of higher overtones, a technique first proposed by \\citet{andriesetal05a} and later developed by \\citet{verthetal08} to correct for magnetic stratification.  All this theoretical development to describe kink waves in coronal loops with realistic plasma stratification in the transverse and longitudinal direction was restricted to the study of standing waves, since this was what was detected in TRACE data. However, recently it has come to light that there are also ubiquitous small amplitude propagating transversal MHD waves in the solar atmosphere. These were first observed by the novel Coronal Multi-Channel Polarimeter (CoMP) instrument \\citep{tomczetal07}. Moreover, \\citet{tomcmac09} have been able to separate outward and inward propagating wave power. It was found that the outward power was greater than the inward power by about a factor of two and this can only be explained if the waves are damped in situ \\citep[see also][]{pascoeetal10}. The reason we could not detect these propagating waves previously with TRACE is that the amplitudes are of the order 50 km, while the TRACE resolution is only about 800 km. \\citet{tomczetal07} originally interpreted these wave as Alfv\\'{e}n waves, however, \\citet{vandetal08b} subsequently argued that the observed waves were actually more consistent with the propagating kink mode. Although both modes are dominated by the restoring force of magnetic tension, in the geometry of a solar flux tube, e.g., a magnetic cylinder, a pure Alfv\\'{e}n wave is strictly torsional with no transverse component and therefore completely incompressible. On the other hand, a kink wave propagating in a flux tube has a transverse perturbation component and is weakly compressible, at least in the linear regime \\citep[see e.g.,][]{goossensetal09}.  In the present paper we shall restrict the study to investigating the effect of transverse plasma inhomogeneity on the propagating kink mode. Although the two problems of standing and propagating kink waves are closely related, since a standing wave is a superposition of two propagating waves, there are some differences that need to be considered. The standing transverse oscillation is the result of an initial value problem, i.e., an initial disturbance in the solar corona, e.g., a CME or a flare, induces the oscillations of the loops at their natural frequencies or eigenmodes. Alternately, the transverse travelling waves have a forced nature since the photosphere acts as a driver. The frequencies of the kink waves observed by CoMP show a peak around 5 min, indicating a $p$-mode driven photospheric origin. The spatial scale of the driver at the base of the tube is also important in order to excite kink oscillations. From the properties of MHD waves in a flux tube we know that a purely incompressible excitation excites purely incompressible Alfv\\'en waves if the driver is strictly localised inside or outside the tube (assuming a homogeneous loop model). On the other hand, an excitation located both inside and outside the tube invariably excites transverse oscillations since an almost incompressible surface wave between the two media will be established, i.e., the kink mode. ",
        "conclusions": "The spatial damping due to resonant absorption of driven kink waves has been investigated. The main conclusion of the work is that the damping length of propagating kink waves due resonant absorption is a monotonically decreasing function of frequency. The TGV relation for kink waves was derived, demonstrating that for low frequencies the damping length is exactly inversely proportional to frequency. In the high frequency range the TGV relation continues to be an excellent approximation of the actual dependency of the damping length on frequency. Certainly, for all physically relevant frequencies the dependency of damping length on frequency is accurately described by the TGV relation. This dependency means that resonant absorption is selective as it favours low frequency waves and can efficiently remove high frequency waves from a broad band spectrum of kink waves. This has great significance for solar atmospheric kink waves, since high frequency waves will tend to lose more power than their low frequency counterparts before reaching high altitudes in the solar corona, with the exact percentage power loss depending on the properties of the equilibrium, in particular the width of the non-uniform layer and steepness in variation of the local Alfv\\'{e}n speed. With respect to mode conversion, the process of resonant absorption will cause the higher frequency waves to be attenuated more because the global kink mode will be converted into localised Alfv\\'enic modes at lower heights. If the energy of these Alfv\\'enic motions is eventually dissipated, then resonant absorption should  produce a characteristic distribution of the energy as a function of height in the solar atmosphere. This could have important consequences with respect to the spatial distribution of wave heating in the solar atmosphere.  It was also shown that spatial and temporal damping are basically equivalent. In the TT approximation the damping per period and the damping per wavelength are exactly the same. The differences in these two quantities arise in the regime where the TT is not valid, but even in this situation it is easy to relate the spatial and the temporal damping rates through the group and phase speeds of the kink MHD waves. This allows us to translate the results from the temporally damped waves ($\\omega$ complex, $k$ real) to spatially attenuated waves ($\\omega$ real, $k$ complex) due to resonant absorption. This mechanism requires the frequency of the driver to be between the internal and the external Alfv\\'en frequency of the tube. This might seem a very restrictive condition, but in fact it is just the opposite. In the driven problem the frequency is fixed but the system chooses the proper wavelength (along the waveguide) to accommodate the kink mode in the tube. This kink mode generated at the base of the loop propagates upwards along the tube and at the same time is attenuated due to the inhomogeneity at the tube boundary.  An interesting result is that both the damping length in the spatial problem and the damping time in the temporal problem are always smaller than in the TT approximation (when $f\\rightarrow 0$), meaning that waves with short wavelengths or large frequencies are always more efficiently damped. Note that the observations of standing kink waves observed with TRACE and for the propagating kink waves detected with the CoMP instrument are precisely in the regime where the TT is applicable, i.e., where the waves are less affected by resonant absorption. The fact that overtones of standing kink waves and high frequency propagating waves have proved difficult to detect may be a direct consequence of the filtering due to resonant absorption. From a different perspective, we have also shown that the damping per wavelength (and the damping per period) has a weak dependence on the frequency.  It is necessary to remark that our results are based on a simple magnetic flux tube model, i.e., a straight cylinder, with no gravity and pressure. The only inhomogeneity in this model is plasma density in the radial direction, however this a characteristic property of solar waveguides that is observed at all atmospheric heights, e.g., chromospheric magnetic bright points or coronal loops. Thus the theory of resonant damping of propagating kink waves due to radial plasma density inhomogeneity offers a natural explanation for the dissipation observed by e.g., \\citet{tomcmac09}. However, a detailed comparison between the observations and the expected frequency dependent response due to resonant absorption is needed to quantify the precise spatial distribution of wave heating due to this mechanism in the solar plasma, e.g., kink wave dissipation as a function of both frequency and height. This problem will be the subject of a future work (Verth et al. 2010, in preparation)."
    },
    "1004/1004.1551_arXiv.txt": {
        "abstract": "We present samples of starburst galaxies that represent the extremes discovered with infrared and ultraviolet observations, including 25 Markarian galaxies, 23 ultraviolet luminous galaxies discovered with GALEX, and the 50 starburst galaxies having the largest infrared/ultraviolet ratios.  These sources have z $<$ 0.5 and cover a luminosity range of $\\sim$ 10$^{4}$. Comparisons between infrared luminosities determined with the 7.7 \\um PAH feature and ultraviolet luminosities from the stellar continuum at 153 nm are used to determine obscuration in starbursts and dependence of this obscuration on infrared or ultraviolet luminosity.  A strong selection effect arises for the ultraviolet-selected samples: the brightest sources appear bright because they have the least obscuration.  Obscuration correction for the ultraviolet-selected Markarian+GALEX sample has the form log[UV(intrinsic)/UV(observed)] = 0.07($\\pm$0.04)M(UV)+2.09$\\pm$0.69 but for the full infrared-selected $Spitzer$ sample is log[UV(intrinsic)/UV(observed)] = 0.17($\\pm$0.02)M(UV)+4.55($\\pm$0.4). The relation of total bolometric luminosity $L_{ir}$ to M(UV) is also determined for infrared-selected and ultraviolet-selected samples.  For ultraviolet-selected galaxies, log $L_{ir}$  = -(0.33$\\pm$0.04)M(UV)+4.52$\\pm$0.69.  For the full infrared-selected sample, log $L_{ir}$ = -(0.23$\\pm$0.02)M(UV)+6.99$\\pm$0.41, all for $L_{ir}$ in \\ldot~ and M(UV) the AB magnitude at rest frame 153 nm.  These results imply that obscuration corrections by factors of two to three determined from reddening of the ultraviolet continuum for Lyman Break Galaxies with z $>$ 2 are insufficient, and should be at least a factor of 10 for M(UV) $\\sim$ -17, with decreasing correction for more luminous sources. ",
        "introduction": "The formation and evolution of galaxies in the Universe is now being observed to redshifts z $\\sim$ 8 \\citep[e.g.][]{bou09,ouc09} using photometric imaging to wavelengths $\\sim$ 1 \\um in search of \"Lyman Break Galaxies\" (LBGs; Steidel et al. 1999).  These observations reveal the rest frame ultraviolet luminosity of star-forming galaxies and allow the measurement of changes in the star formation rate density of the Universe as a function of redshift \\citep{mad98}.  (Because such star-forming galaxies are observationally characterised by indicators of young, massive, short-lived stars, we refer to them in the remainder of this paper as \"starburst\" galaxies.)  Observations of LBGs indicate an increase in the characteristic luminosity of starburst galaxies for 0 $<$ z $\\la$ 2, and then a decrease in luminosity for z $\\ga$ 2 until the observational limit at z $\\sim$ 8 \\citep{red09}.  Is this evolution of starburst galaxies determined correctly based on LBGs?  The most uncertain factor in quantifying the luminosities is determining an accurate correction for obscuration by dust, because this obscuration can be extreme in the rest frame ultraviolet.   The importance of dust in the early universe is emphasized by the result that the most luminous galaxies known have most of their luminosity emerging as reradiation by dust of intrinsic luminosity initially generated at much shorter wavelengths and absorbed by the dust.  Initially termed \"Ultraluminous Infrared Galaxies\" (ULIRGS; Soifer et al. 1987) when discovered with the Infrared Astronomical Satellite (IRAS), such galaxies are found in large numbers in surveys at mid-infrared wavelengths at all redshifts to z $\\sim$ 3 \\citep[e.g.][]{lef05,dey08} and millimeter wavelengths \\citep{cha05,pop08,men09}. The evolution of starbursts among ULIRGs has been tracked to z $\\sim$ 3 \\citep{wee08} with the $Spitzer$ Space Telescope using redshifts measured with The Infrared Spectrograph on $Spitzer$ (IRS; Houck et al. 2004).  Because of the inability to measure redshifts of dusty starbursts with z $>$ 3 using the IRS, these have not yet been found at redshifts as high as LBGs.  Combining our understanding of dusty galaxies dominated by infrared luminosity with galaxies discovered by ultraviolet observations is essential to a full understanding of the formation and evolution of galaxies (as well as the formation and evolution of dust).  The availability of IRS spectroscopy for the same galaxies which are detected in the rest frame ultraviolet by the Galaxy Evolution Explorer (GALEX, Martin et al. 2005) allows a direct comparison between the infrared luminosity from dust and the ultraviolet luminosity from hot stars.  Our initial comparison in this way for 287 starburst galaxies having both $Spitzer$ spectra and GALEX ultraviolet measures found large differences in the star formation rates (SFRs) deduced from infrared compared to ultraviolet measures of SFR \\citep{sar09}.  These differences are attributed to extreme obscuration of the ultraviolet luminosity.  Our previous sample arose primarily from selection in the infrared, chosen as starbursts having IRS spectra available.  In the present paper, we use IRS spectra for samples of ultraviolet-selected starburst galaxies to study obscuration effects in sources with similar selection criteria as LBGs, in order to understand how conclusions depend on selection criteria.  Our objective is to use observations in the infrared and ultraviolet to make an independent test of dust corrections for starbursts detected as LBGs at high redshift, such as in \\citet{red09}, and to determine reliable corrections for dust obscuration which are independent of selection technique.  We make measures only of the obscuration in order to avoid uncertain assumptions regarding transformation of luminosities to SFRs.  In section 2, below, we describe the sample selection which includes 25 Markarian galaxies, 23 ultraviolet luminous galaxies discovered with GALEX, and the 50 starburst galaxies having the largest infrared/ultraviolet ratios.  In section 3, we present the infrared spectra of all 98 sources and compare the characteristics of infrared and ultraviolet spectra in section 4.  Results are used in section 5 to develop an empirical method for determining obscuration corrections as a function of M(UV), and in section 6 to compare bolometric luminosities $L_{ir}$ with M(UV).  Our most important result (section 7) indicates that the bolometric luminosities deduced for starburst galaxies using LBGs at z $>$ 2 have been underestimated, because obscuration corrections   have been significantly underestimated for starbursts.  This means that SFRs have also been underestimated at high redshifts, and we propose a new observational test to verify these conclusions. ",
        "conclusions": "1. We have selected starburst galaxies discovered using both infrared and ultraviolet selection for comparison of spectral parameters measured with the $Spitzer$ Infrared Spectrograph to ultraviolet photometry from GALEX.  Our primary objective was to determine how the luminosity and dust obscuration of starbursts depend on selection technique.  Sources have z $<$ 0.5 and cover a luminosity range of $\\sim$ 10$^{4}$; the highest luminosities are comparable to starburst luminosities found at z $\\sim$ 3.  2. We measure and illustrate spectra of 26 Markarian galaxies, 23 ultraviolet luminous galaxies discovered with GALEX, and 50 galaxies with the most extreme infrared/ultraviolet ratios discovered with $Spitzer$.  We find (section 4.3) that these samples differ in the nature of the ionizing stars, with hotter stars in the ultraviolet-selected sources as evidenced by stronger [NeIII] emission and stronger mid-infrared dust continuum (section 4, Figures 6 - 8.)  3. It is found that a strong selection effect arises for the ultraviolet-selected samples: the most luminous UV sources appear luminous because they have the least obscuration, not because they have the largest bolometric luminosities (section 5.3).  Comparisons between rest-frame infrared fluxes f$_{\\nu}$(7.7 \\ums) arising from PAH emission and ultraviolet fluxes f$_{\\nu}$(153 nm) arising from the stellar continuum allow an empirical method for determining obscuration in starbursts and dependence of this obscuration on infrared or ultraviolet luminosity.    Even sources selected in the ultraviolet and which resemble Lyman Break Galaxies require obscuration corrections of about a factor of 10.  These results are contrary to those adopted in studies of LBGs, which find that the most luminous galaxies in M(UV) are also the most obscured. Obscuration correction for the ultraviolet-selected Markarian+GALEX sample has the form log[UV(intrinsic)/UV(observed)] = 0.07($\\pm$0.04)M(UV)+2.09$\\pm$0.69; the full infrared-selected $Spitzer$ sample is log[UV(intrinsic)/UV(observed)] = 0.17($\\pm$0.02)M(UV)+4.55($\\pm$0.4); and is log [UV(intrinsic)/UV(observed)] = 0.14($\\pm$0.03)M(UV)+4.68$\\pm$0.50 for the extreme infrared-selected $Spitzer$ sample (Figure 13).  4. Analysis of obscuration determined from the H$\\alpha$/H$\\beta$ ratio in SDSS spectra confirms the trends observed by comparing infrared/ultraviolet ratios, although obscuration measured by H$\\alpha$/H$\\beta$ is much less than measured by infrared/ultraviolet (section 5.4 and Figure 15).  5. Stellar luminosities measured at rest frame 1.6 \\um are exceptionally uniform compared to starburst luminosity over a redshift range 0 $<$ z $<$ 2 and luminosity range of 10$^{4}$, but the ratio from different samples indicates that the 1.6 \\um stellar luminosity is also obscured by dust in proportion to the ultraviolet obscuration.  This leads to our suggesting that stellar luminosities measured at 1.6 \\um reveal the recently evolved stars from the starburst rather than an underlying old population (section 5.5 and Figure 16).  6. We examine the relation of total bolometric luminosity $L_{ir}$ to M(UV) for infrared-selected and ultraviolet-selected samples.  Spectroscopy of Markarian galaxies confirms the transformation assumed for $\\nu$L$_{\\nu}$(7.7 $\\mu$m) to $L_{ir}$ and implies that the hotter ionizing stars in the uv-selected samples do not affect this calibration.  For ultraviolet-selected galaxies, log $L_{ir}$  = -(0.33$\\pm$0.04)M(UV)+4.52$\\pm$0.69.  For the full infrared-selected sample, log $L_{ir}$ = -(0.23$\\pm$0.02)M(UV)+6.99$\\pm$0.41, and log $L_{ir}$ = -(0.25$\\pm$0.03)M(UV)+7.2$\\pm$0.48 for the extreme infrared-selected $Spitzer$ sample, all for $L_{ir}$ in \\ldot and M(UV) the AB magnitude at rest frame 153 nm (section 6 and Figure 18).  7. Our most important summary result (section 7 and Figure 20) indicates that the bolometric luminosities deduced for starburst galaxies using LBGs at z $>$ 2 are insufficient, by a factor of at least 10 for M(UV) $\\sim$ -17.  This also implies that obscuration corrections by factors of two to three applied for 3 $<$ z $<$ 7 to track the formation and evolution of the earliest galaxies are insufficient, if dusty galaxies exist at such epochs comparable to those with z $<$ 3.  We propose an observational test which can determine more accurate dust corrections and $L_{ir}$ for large samples of infrared-selected starbursts and LBGs at z $\\la$ 2.  This test is possible because existing $Spitzer$ surveys allow reliable discovery of dusty starbursts with z = 1.8 $\\pm$ 0.2 based only on photometric properties."
    },
    "1004/1004.5234_arXiv.txt": {
        "abstract": "Nanometer- and micrometer-sized solid particles play an important role in the evolutionary cycle of stars and interstellar matter. The optical properties of cosmic grains determine the interaction of the radiation field with the solids, thereby regulating the temperature structure and spectral appearance of dusty regions. Radiation pressure on dust grains and their collisions with the gas atoms and molecules can drive powerful winds. The analysis of observed spectral features, especially in the infrared wavelength range, provides important information on grain size, composition and structure as well as temperature and spatial distribution of the material.  The relevant optical data for interstellar, circumstellar, and protoplanetary grains can be obtained by measurements on cosmic dust analogs in the laboratory or can be calculated from grain models based on optical constants. Both approaches have made progress in the last years, triggered by the need to interpret increasingly detailed high-quality astronomical observations. The statistical theoretical approach, spectroscopic experiments at variable temperature and absorption spectroscopy of aerosol particulates play an important role for the successful application of the data in dust astrophysics. ",
        "introduction": "\\label{sect:intro}  % Dust grains exist in a wide variety of cosmic environments ranging from the diffuse interstellar medium and molecular clouds in our galaxy to circumstellar shells and disks around young stellar objects, dusty outflows from evolved stars, atmospheres of brown dwarfs, the dust tori of active galactic nuclei and starburst galaxies, and high-redshift quasars. The particles are especially important during the birth process of stars and as precursor material for the formation of planets. In the late stages of stellar evolution they condense in molecular outflows and are important for the mass loss process and the spectral appearance of AGB and post-AGB stars.  The field of cosmic-dust research has been the subject of intense research over the last decades. The international conferences ``Astrophysics of Dust'' in Estes Park, Colorado (2003) and ``Cosmic Dust - Near and Far'' in Heidelberg, Germany (2008) both reflected the rapid development of the field. Astronomical dust research is of course triggered by the progress in astronomical observations, mainly by those in the infrared range, because of the opportunities to (1) observe dusty objects optically thick in the visible and (2) to analyse the thermal emission of dust by infrared spectroscopy. The ISO and Spitzer space telescopes have played a major role in this respect, and current and future projects such as the Herschel mission and JWST will continue to ensure progress in this field.  The description of the interaction of radiation fields with a system of solid particles always requires knowledge of their absorption and scattering cross sections which depend on their chemical composition, solid-state structure and morphology (see Table~\\ref{tbl-1}). Any continuum  radiative transfer code for the calculation of intensity/polarization maps and spectral energy distributions needs the specification of these properties. In addition, the transfer of momentum from a radiation field to a grain is determined by the extinction and scattering  efficiency and the average cosine of the scattering angle.  \\begin{table} \\caption{Chemical and structural properties of solid particles and the description of optical properties on different length scales } \\label{tbl-1} \\begin{center} \\begin{tabular}{lll} Length scale &  Structural properties &  Description\\\\ &&\\\\ \\hline &&\\\\ ``Atomic''          & Chemical composition and bonding  &  Dielectric function $\\epsilon(\\lambda)$\\\\ & Crystal structure                 &  Complex refractive index $m(\\lambda$)\\\\ & Defects and impurities            &        \\\\ &&\\\\ ``Mesoscopic''     & Inhomogeneities                   &  Theory of effective media\\\\ & Porosity                          &  $\\epsilon_{\\rm eff}(\\lambda)$,\\\\ & Mantles                           &  Core-mantle description\\\\ & Surface states                    &  \\\\ &&\\\\ ``Macroscopic'',    & Size distribution                 &  Mie theory, T-matrix method,\\\\ Morphology         & Grain shapes                      &  Separation of variables method,\\\\ & Agglomeration                     &  Discrete dipole approximation\\\\ & Coalescence        &  $C_{\\rm ext}(\\lambda), C_{\\rm abs}(\\lambda), C_{\\rm sca}(\\lambda)$\\\\ \\end{tabular} \\end{center} \\end{table}   In this review we will deal with the progress in theoretical and experimental work aimed to providing data and tools for the interpretation of spectroscopic dust-related observations, and for modeling of dusty astrophysical objects. Current models go significantly beyond the assumption of spherical grains characterized by artificially constructed optical constants, as has been the prefered approach for many years. Effects of grain morphology, temperature, and composition have been studied experimentally, and complex light-scattering models have been used to explore the dependence of scattering and absorption cross sections on the structure of grains on all length scales (see Table~\\ref{tbl-1}), including porosity, inhomogeneity, and agglomeration. Such grain models have been used in sophisticated multi-dimensional radiative transfer codes. Among the experimental data, which can be included as inputs into the computations, are also measurements of the scattering matrices of cosmic dust analogs which have been performed in Amsterdam for many years and have been made available through the Amsterdam Light Scattering Database (http://www.astro.uva.nl/scatter). Although this review unfortunately cannot cover this and some other experimental fields devoted to interpret observational data of cosmic dust, all general considerations presented in the next sections are valid for them as well, especially those on the relation between calculations and measurements.  It is beyond the scope of this paper to discuss extensively the astronomical dust models and optical data derived from observations. For a description of the various ``classical'' dust models, we refer to the contributions in Ref. \\nocite{alti89}1. More recent models have been developed by Mathis \\cite{mathis96}, Li \\& Greenberg \\cite{ligb97}, Zubko \\textit{et al.} \\cite{zubko04}, Draine \\& Li \\cite{drli07} and can also be found in the proceedings of the two dust conferences \\cite{AoD03,DNF08}. ",
        "conclusions": "The measurements and the theoretical calculations of optical properties of cosmic grains play an important role in a wide range of astrophysical applications. The collection of relevant data for well-characterized materials over a broad wavelength range has made great progress over the last two decades. Here we want to stress that the characterization of the physical and chemical structure of the materials is an important part of the study of cosmic dust analogs. An interesting direction of future research will be the comparison of the properties of cosmic dust grains collected by interplanetary space missions such as the Stardust mission (e.g., Ref. \\nocite{bradley10}76) with synthetic materials produced in the laboratory.  The combination of condensation experiments carried out under conditions similar to cosmic environments and advanced spectroscopic methods is certainly the future of solid-state astrophysics. Mass spectroscopy has to be used to characterize the transition from atoms and molecules to small clusters and solids. In addition, the change of material properties under the bombardment of cosmic rays and irradiation with intense X-ray and UV photons has to be better understood. Another important factor can be thermal annealing cycles of the grains.  The application of laboratory data for the interpretation of astronomically measured fluxes relies on the knowledge of the density and temperature structure of the considered dusty regions. Radiative transfer calculations have to be carried out for deeply embedded objects such as protostars and active galactic nuclei. In the case of very small particles the effects of quantum heating have to be taken into account. Even for the interpretation of spectra of optically thin configurations, the dust temperature and density distributions have to be known."
    },
    "1004/1004.1976_arXiv.txt": {
        "abstract": "{The present work performs an asteroseismic study of a $\\gamma$\\,Doradus star member of the open cluster NGC~2506. This works conjugates the observational information of the star and the strong constraints imposed by the cluster memebership, with the latest developed theoretical asteroseismic tools for the modelling of this type of objects, in particular the Frequency Ratio Method (FRM) and Time-Dependent Convection (TDC), to study the $\\gamma$\\,Dor stars found in the NGC 2506 stellar cluster. The use of both techniques gives us the opportunity of constructing a self-consistent procedure, allowing the mode identification and improving the modelling of the $\\gamma$\\,Dor members of the cluster. In this work we present the result of the analysis of the first target of our project, showing the advantadge of modelling cluster member's $\\gamma$\\,Dor stars. This is the first step in a more ambitious project. In the near future, we plan to expand the study to other variable members of the NGC 2506 cluster. This will improve our knowledge of $\\gamma$\\,Dor pulsatorsand probe the self-consistency of our procedures.} ",
        "introduction": "The $\\gamma$\\,Doradus (Dor) stars are suitable objects for astroseismic study. They are high-order gravity ($g$)~mode pulsators, which permits the probing of their stellar coreand the use of the asymptotic approximation, for which analytical solutions of the oscillation equations exist (\\citealt{Smey07}). The Frequency Ratio Method (FRM; \\citealt{Moya05, Suarez05}), based on this property, is particularly adapted for obtaining asteroseismic information on $\\gamma$\\,Dor pulsating stars that show, at least, three pulsation frequencies. The method provides an estimate of the identification of the radial order $n$ and the spherical degree $\\ell$ of observed frequenciesand of the integral of the buoyancy frequency (Brunt-V\\\"ais\\\"al\\\"a) weighted over the stellar radius along the radiative zone ($I_{\\rm{obs}}$).  The excitation mechanism of $\\gamma$\\,Dor stars was an outstanding problem until the works of \\cite{Guzik00} and \\cite{Dupt05} who used respectively, frozen and Time-Dependent Convection (TDC) theories showed that the position of the base of the convective envelope is the key for driving $\\gamma$\\,Dor $g$~modes. Moreover, TDC is formulated in the framework of mixing length theory and therefore the $\\gamma$\\,Dor instability strip is predicted to be sensitive to the mixing-length parameter $\\alpha_{\\rm MLT}$.  The asteroseismic modelling of $\\gamma$\\,Dor, even armed with these techniques, still needs the global parameters of the observed stars. One of the ways to obtain a significant enhancement is to use the strong constraints imposed by the cluster membership for example, the age, metallicity and distance (\\citealt{2002A&A...390..523S, 2007MNRAS.379..201S,2009A&A...507..901C}).  At the beginning of the recognition of the $\\gamma$\\,Dor stars as a new class of variable stars, a great effort has been devoted to their detection in clusters. At that time the main aim of the community was to check whether this type of pulsations was linked or not to age and obviously, stellar clusters were the most suitable targets for this purpose (\\citealt{2002ASPC..259..152M}).  One of the richest cluster in $\\gamma$\\,Dor members is the old open cluster NGC 2506, for which photometric observations provide a considerable number of variable stars (\\citealt{2007A&A...465..965A}). Among them, 15 have been identified as $\\gamma$\\,Dor stars candidates (labeled as V$_{11}$ -- V$_{25}$ in that paper).   This work attempts to investigate the impact of using a procedure which includes FRM and TDC on the modelling of $\\gamma$\\,Dor, stars taking advantage of their membership in clusters. In particular, we here provide preliminary results for the $\\gamma$\\,Dor star V$_{11}$, belonging to the open cluster NGC 2506. ",
        "conclusions": "With this prospective study we showed that the application of FRM and TDC to the cluster member $\\gamma$\\,Dor candidate NGC 2506 V$_{11}$ enhances significantly our determination of the mass, radius and age. Specifically, its metallicity is confirmed to be sub-solar. The energy transport efficiency of the outer convective zone is high ($\\alpha_{\\rm MLT}=1.5$). These results must be confirmed with the extension of this study to other $\\gamma$\\,Dor stars of this cluster.  This is a first step of a more ambitious project. In the near future, we plan to expand the study to other variable members of the NGC 2506 cluster. A better knowledge of variable membership (\\citealt{Ottosen08}) and their global parameters will improve our knowledge on $\\gamma$\\,Dor pulsators and probe the self-consistency of our procedures.  \\begin{table*} \\centering{ \\caption{Minimum and maximum values of the physical parameters of NGC 2506 V$_{11}$ fulfilling all the constraints} \\label{tab3} \\begin{tabular}{lccccccc}\\hline Mass  & ${\\rm T}_{\\rm eff}$(K)  &   [Fe/H]  &  L/L$_{\\sun}$  &  R/R$_{\\sun}$   &  Age (Myr) & $\\alpha_{\\rm MLT}$ & $\\alpha_{\\rm ov}$ \\\\ \\hline \\hline 1.25 -- 1.29 & 6997 & -0.52 & 6.07 & 1.6   & 1832 & 1.5 & 0.1\\\\ 1.33 -- 1.38 & 7223 & -0.32 & 6.83 & 1.774 & 2585 & 1.5&  0.3\\\\ \\hline \\end{tabular}} \\end{table*}"
    },
    "1004/1004.0694_arXiv.txt": {
        "abstract": "Scaling relations of observed galaxy cluster properties are useful tools for constraining cosmological parameters as well as cluster formation histories.  One of the key cosmological parameters, $\\sigma_8$, is constrained using observed clusters of galaxies, although current estimates of $\\sigma_8$ from the scaling relations of dynamically relaxed galaxy clusters are limited by the large scatter in the observed cluster mass-temperature ($M-T$) relation.  With a sample of eight strong lensing clusters at $0.3 < z <0.8$, we find that the observed cluster concentration-mass relation can be used to reduce the $M-T$ scatter by a factor of 6.  Typically only relaxed clusters are used to estimate $\\sigma_8$, but combining the cluster concentration-mass relation with the $M-T$ relation enables the inclusion of unrelaxed clusters as well.  Thus, the resultant gains in the accuracy of $\\sigma_8$ measurements from clusters are twofold: the errors on $\\sigma_8$ are reduced and the cluster sample size is increased. Therefore, the statistics on $\\sigma_8$ determination from clusters are greatly improved by the inclusion of unrelaxed clusters. Exploring cluster scaling relations further, we find that the correlation between brightest cluster galaxy (BCG) luminosity and cluster mass offers insight into the assembly histories of clusters.  We find preliminary evidence for a steeper BCG luminosity - cluster mass relation for strong lensing clusters than the general cluster population, hinting that strong lensing clusters may have had more active merging histories. ",
        "introduction": "\\label{intro}  As the most massive bound systems known, galaxy clusters provide an important link in understanding the composition and growth of structure in the Universe.  Clusters follow a variety of observational scalings of mass with temperature, luminosity, or cluster counts, and these scalings are sensitive to cosmological parameters including the matter density parameter $\\Omega_{\\rm m}$, the cosmological constant density parameter $\\Omega_\\Lambda$, the dark energy equation-of-state parameter $w$, and the normalization of the matter power spectrum $\\sigma_8$ (e.g., \\citealt{HA01.2,BA02.2,LE02.1,SC03.2,AL04.2,VI09.1}).  Such constraints from galaxy clusters complement the constraints on cosmological parameters from Type Ia supernovae and cosmic microwave background observations.  However, useful galaxy cluster constraints on cosmological parameters depend primarily on accurate determinations of cluster masses.  Observationally, cluster masses are typically measured in one of three ways.  A long-established method for determining cluster masses employs the virial theorem and the measurement of velocities of the galaxies that constitute the cluster.  Based on the three assumptions that the cluster is in virial equilibrium, the galaxy distribution efficiently traces the cluster mass distribution, and the velocity dispersions $\\sigma$ of the galaxies are isotropic, the cluster mass contained within a radius $r$ is estimated $M \\sim \\sigma^2 r/G$.  However, these mass estimates may be biased as a result of galaxy velocity anisotropies or if the galaxy distribution does not follow the total mass distribution (e.g., \\citealt{BA82.1}).  A second method uses cluster X-ray emission as a tracer of cluster masses. The hot intracluster gas, which is the dominant baryonic component of a cluster and is typically twice the mass of the total mass of the galaxies in a cluster, emits X-rays via bremsstrahlung radiation and atomic line emission.  With the temperature $T$ and radial density $\\rho(r)$ profiles determined from X-ray spectra and surface brightness distributions, the cluster mass is given by $M \\sim r^2/\\rho(r) \\; d(-\\rho T)/ dr$.  This method assumes that the intracluster gas is spherically distributed and is in hydrostatic equilibrium \\citep{EV96.1}.  However, these assumptions may be incorrect. If the gas distribution is not spherical, X-ray mass estimates will be biased by projection effects.  Many galaxy clusters are also not in hydrostatic equilibrium, in particular dynamically unrelaxed clusters that are undergoing mergers.  There is evidence that the bias of hydrostatic equilibrium mass is linked to the dynamical state of the galaxy cluster (e.g., \\citealt{AN09.1,ZH09.1}). In addition, the hot gas of galaxy clusters with buoyant bubbles near their cores might indicate a departure from hydrostatic equilibrium (e.g., \\citealt{CH01.1}).  The most direct estimates of cluster masses employ gravitational lensing distortions of background galaxies.  This technique is free of assumptions about the dynamical state of the cluster, which enables it in principle to yield more consistent mass estimates, though it is also sensitive to projection effects.  More accurate cluster mass estimates can in turn provide tighter constraints on cosmological parameters, and therefore it is of key importance to reduce the errors in cluster mass estimates.  For example, the primary source of error in cluster-based determinations of $\\sigma_8$ is the error in the mass-temperature relation for relaxed clusters (e.g., \\citealt{PI03.1,HE04.2,VO05.1}).  Recent studies show that an X-ray independent mass approach such as gravitational lensing provides a unique tool to calibrate the mass - temperature relation (e.g., \\citealt{SM05.1, MA08.1, ZH08.1}). Here, we use strong gravitational lensing mass measurements of a sample of eight strong lensing clusters at $0.3 < z < 0.8$ to accurately measure the galaxy cluster mass-temperature relation.  We also include the effects of cluster concentrations in an effort to further reduce the scatter in the cluster mass-temperature relation, which would ultimately enable tighter constraints on $\\sigma_8$.  In addition to the correlations that exist between cluster properties, some observational properties of brightest cluster galaxies (BCGs) also scale with properties of the host clusters.  Whereas scalings between cluster properties are sensitive to cosmological parameters, scalings between BCGs and their host clusters provide constraints on BCG formation and the evolution of clusters.  BCGs are a unique population: they are the most massive and luminous galaxies in the Universe.  They are typically located near the centers of clusters, which suggests that a BCG's formation history is intricately linked to the formation of the cluster itself. However, the formation of BCGs is still poorly understood.  BCGs may form after their host clusters assemble in one of two ways.  First, a BCG may be the first galaxy to be dragged in by dynamical friction to the center of the dark matter halo destined to become a cluster, where it then grows through galactic cannibalism by merging with subsequent galaxies that fall to the center (e.g., \\citealt{OS75.1,HA78.1}).  However, this scenario typically requires more than a Hubble time to form a BCG because much of the mass of the infalling galaxy is tidally stripped, which reduces the dynamical friction effect and slows the infall \\citep{ME85.1}.  BCG formation may also occur after cluster formation if the host cluster's central cooling flow forms stars at the cluster center and those stars build the BCG \\citep{CO77.1}. There are several instances of ongoing or recent star formation in BCGs that occupy cooling-flow clusters (e.g., \\citealt{CA98.1,CR99.1, HI05.1,MC06.1}), but it is unclear whether the star formation is fueled by the cooling flows or by cold gas brought in through recent galaxy mergers \\citep{BI08.2}.  In another scenario, BCGs might form in concert with their host clusters. A BCG may begin with several galaxies merging together in a group to form a large galaxy, and then when groups merge as hierarchical structure formation continues, this large galaxy eventually becomes a BCG in a massive cluster (e.g., \\citealt{ME85.1, DU98.1, BO06.1}).  Here, we examine the correlation between BCG luminosity and cluster mass in eight strong lensing clusters at $0.3 < z <0.8$.  This will enable constraints not only on BCG and cluster formation in general, but also on how the BCGs in strong lensing clusters may have formed and evolved differently than BCGs in the general cluster population.  The rest of this paper is organized as follows. In Section~\\ref{sample} we describe the selection of our cluster sample, and Section~\\ref{properties} gives the masses, dynamical states, and X-ray temperatures for these clusters.  In Section~\\ref{mtrelation} we find the $M-T$ relation for the relaxed clusters in our sample and show how the inclusion of cluster concentrations both significantly reduces the scatter in the $M-T$ relation and lifts the restriction on cluster dynamical state.  In Section~\\ref{bcgprop} we identify the BCGs in our sample and measure their luminosities. We use these luminosities in Section~\\ref{lm} to measure the correlation between BCG luminosity and cluster mass, and we find preliminary evidence that strong lensing clusters may have more active merging histories than the general cluster population.  Section~\\ref{conclude} presents our conclusions. Throughout this paper, we adopt a spatially flat cosmological model dominated by cold dark matter and a cosmological constant ($\\Omega_{\\rm m}=0.3$, $\\Omega_\\Lambda=0.7$, $h=0.7$). ",
        "conclusions": "\\label{conclude}  We have determined the scaling of cluster mass with cluster temperature and the scaling of BCG luminosity with cluster mass for eight observed strong lensing galaxy clusters imaged with {\\it HST} and at redshifts $0.3 < z < 0.8$.  We explored cluster concentrations as a means of reducing the scatter in the $M-T$ relation and enabling more precise constraints on cosmological parameters, and we used the $L-M$ relation as an indicator of the formation histories of strong lensing BCGs and clusters. Our main results are:  \\begin{enumerate}  \\item  The best-fit cluster mass-temperature relation for our four dynamically relaxed strong lensing clusters is \\begin{equation} \\label{eq:mt} E(z) \\left( \\frac{M_{2500}} {10^{14} \\, h_{70}^{-1} \\, M_\\odot} \\right) = 1.60 \\pm 3.42 \\left( \\frac{k T} {5 \\, \\mathrm{keV}} \\right)^{1.43 \\pm 1.28} \\, , \\end{equation} which is consistent with the theoretical expectation of the $M-T$ relation for relaxed clusters, as well as the $M-T$ relations determined by other observations and simulations.  We find that relaxed strong lensing clusters do not deviate from the $M-T$ relation for the general population of relaxed clusters.  Significantly, we find an inverse correlation between cluster temperature and cluster concentration that, when incorporated into the $M-T$ relation, reduces the $M-T$ scatter by a factor of 6, from 360$\\%$ to 60$\\%$.  By correcting cluster temperatures according to the temperature-concentration relation, we find that the $M-T$ relation given in Equation~\\ref{eq:mt} describes not only the relaxed strong lensing clusters, but the entire cluster population regardless of dynamical state.  Specifically, the scatter in unrelaxed clusters decreases by a factor of 30, from 500$\\%$ in the uncorrected $M-T$ relation to 15$\\%$ in the temperature-corrected $M-T$ relation.  Incorporating concentration effects into the $M-T$ relation tightens the $M-T$ relation for all clusters, which in turn reduces the error in the determination of $\\sigma_8$ from cluster counts.  Whereas accurate cluster determinations of $\\sigma_8$ were previously made only with relaxed clusters, concentrations enable the inclusion of unrelaxed clusters.  The larger cluster samples possible with the inclusion of unrelaxed clusters offer yet more precise $\\sigma_8$ estimates from cluster observations.  \\item The best-fit relation between BCG luminosity and cluster mass for our sample of strong lensing clusters is \\begin{equation} \\frac{L_\\mathrm{K, BCG}} {10^{11} \\; h_{70}^{-2} \\; L_\\odot} = 0.97 \\pm 0.17 \\left( \\frac{M_{200}} {10^{14} \\; h_{70}^{-1} \\; M_\\odot} \\right)^{0.48 \\pm 0.09} \\; , \\end{equation} which is $\\sim 85\\%$ steeper than the correlations predicted for non-strong-lensing clusters by other observations, theory, and cosmological simulations.  This result supports the current evidence that BCGs are built up through mergers with massive galaxies in other groups and clusters, and also hints that strong lensing clusters may have more active merging histories than typical clusters.  A larger sample of strong lensing clusters might enable more definite conclusions about the formation histories of strong lensing clusters.  \\end{enumerate}  Accurate cluster mass measurements and full use of the range of cluster property interdependencies are key components in the calibration of clusters as tracers of cosmological parameters.  As we have shown, gravitational lensing enables the most direct measurements of cluster mass, without assumptions about the cluster's dynamical state that are inherent in other methods.  We have also shown that the correlation between cluster temperature and concentration can significantly reduce the scatter in the cluster $M-T$ relation, enabling more precise estimates of $\\sigma_8$.  It may be that other cluster scalings can be effectively combined to reduce the error on additional cosmological parameter estimates."
    },
    "1004/1004.5414_arXiv.txt": {
        "abstract": "In the recent years, the ``Nice\" model of solar system formation has attained an unprecedented level of success in reproducing much of the observed orbital architecture of the solar system by evolving the planets to their current locations from a more compact configuration. Within the context of this model, the formation of the classical Kuiper belt requires a phase during which the ice giants have a high eccentricity. An outstanding question of this model is the initial configuration from which the Solar System started out. Recent work has shown that multi-resonant initial conditions can serve as good candidates, as they naturally prevent vigorous type-II migration. In this paper, we use analytical arguments, as well as self-consistent numerical N-body simulations to identify fully-resonant initial conditions, whose dynamical evolution is characterized by an eccentric phase of the ice-giants, as well as planetary scattering. We find a total of eight such initial conditions. Four of these primordial states are compatible with the canonical \"Nice\" model, while the others imply slightly different evolutions. The results presented here should prove useful in further development of a comprehensive model for solar system formation. ",
        "introduction": "The question of how the solar system formed dates back centuries. The last decade, however, has seen a considerable amount of progress made on this issue. Notably, the development of the Nice model (Tsiaganis et al. 2005) has proven to be a milestone. The scenario foretold by the Nice model is as follows: the giant planets form in a compact configuration, and driven by planetesimal scattering (Fernandez \\& Ip 1984, Malhotra 1995) begin migrating divergently. Eventually, Jupiter and Saturn cross their mutual 2:1 mean-motion resonance (MMR), which results in an acquisition of eccentricities for both planets. Subsequently, the whole outer solar system undergoes a brief period of dynamical instability, during which Uranus and Neptune are scattered to their current orbits.  There are a few aspects to the success of the Nice model. First and foremost, it has been able to replicate the architecture of the secular dynamics of the outer solar system (Morbidelli et al. 2009a). Second, it provides a semi-quantitative description of the formation of the Kuiper belt (Levison et al. 2008). Third, the inward flux of planetecimals during the phase of dynamical instability allows for chaotic capture of Jupiter's and Neptune's Trojan populations (Morbidelli et al. 2005, Nesvorni{\\'y} et al. 2007). Finally, if the resonance crossing between Jupiter and Saturn is timed appropriately, the global mayhem provides a natural trigger for Late Heavy Bombardement (LHB) (Gomes et al. 2005). There are other observational constraints that should be reproduced in a model for the solar system's formation, such as the dynamical structure of the inner solar system, and considerable progress has already been made in this direction (Brasser et al. 2009, Morbidelli et al. 2009b). At the same time, it is also crucial to explore the unobservable aspect of the model, namely the initial conditions. This is the purpose of our study.  In this paper, we consider various multi-resonant configurations as possible initial conditions for the Nice model. We investigate the early stages of dynamical evolution for a large number of candidate systems and show that only eight configurations appear to be consistent with the formation of the Kuiper belt in the framework of the Nice model. It is noteworthy that running simulations of the Nice model to completion is computationally expensive. Consequently, as we seek to examine a large array of initial conditions, we are forced to utilize early dynamical events, namely planetary scattering, as proxies for successful formation of the solar system. In this manner, we limit the duration of each simulation to only a few tens of millions of years. Within the context of our integrations, this is long enough for the system to pass through the epoch of dynamical instability, but not long enough to scatter away all of the planetesimals which end up on long-term unstable orbits.  The plan of the paper is as follows: in section 2, we explain how multi-resonant configurations prevent type-II migration, and our approach to their assembly. In section 3, we discuss the evolution scenarios of the considered systems, estimate the amplitudes of eccentricity jumps in relevant cases, and present the results of N-body simulations. We discuss our results and conclude in section 4. ",
        "conclusions": "The calculations presented in this work aim to place constraints on the early dynamical state of the solar system. We began by constructing a large array of multi-resonant systems. We then considered each one of these systems individually and tested them against two inter-related constraints: scattering of an ice-giant from a gas-giant and the outer ice giant undergoing a transient phase of high eccentricity. Both events are needed to reproduce the current dynamical architecture of the solar system.  We showed numerically that three configurations in which Jupiter \\& Saturn start out in a 3:2 MMR follow evolutionary tracks which are compatible with our constraints. The three systems are those in which (1) Saturn \\& Uranus are in a 3:2 MMR while Uranus \\& Neptune are in a 5:4 MMR, (2) Saturn \\& Uranus are in a 4:3 MMR while Uranus \\& Neptune are in a 3:2 MMR and (3) Saturn \\& Uranus are in a 4:3 MMR while Uranus \\& Neptune are also in a 4:3 MMR. Two of these configurations were previously thought to be unstable (Morbidelli et al. 2007).  We showed that the classic Nice model can be reproduced from multi-resonant initial conditions where Juipter \\& Saturn start out in a 5:3 MMR. The four particular primordial states that we found were those where initially, (1) Saturn \\& Uranus are in a 3:2 MMR while Uranus \\& Neptune are in a 5:4 MMR, (2) Saturn \\& Uranus are in a 3:2 MMR while Uranus \\& Neptune are in a 6:5 MMR, (3) Saturn \\& Uranus are in a 4:3 MMR while Uranus \\& Neptune are in a 3:2 MMR, and (4) Saturn \\& Uranus are initially in a 4:3 MMR while Uranus \\& Neptune are also in a 4:3 MMR.  Finally, we used an analytical technique to rule out a large portion of the generated multi-resonant systems, in which Jupiter and Saturn are initially in the 2:1 MMR, based on an argument that the eccentricities generated by any Uranus/Neptune resonant encounters are too small. Simultaneously, we showed that the considered constraints can be satisfied by the dynamical evolution whose initial condition has Jupiter \\& Saturn locked in a 2:1 MMR, and the other pairs of the planets in 4:3 MMR's.     The calculations presented here are intended in part as a point of departure for future research. There is indeed a large array of unexplored issues. For instance, in the case of the system where Jupiter \\& Saturn start out in a 2:1 MMR, it is not clear if a single scattering event alone is enough to correctly reproduce the secular dynamics of Jupiter and Saturn, or if a resonant encounter is also required. In the initial condition where Jupiter and Saturn start out in a 5:3 MMR, dynamical stability may pose an issue (Morbidelli \\& Crida 2007). A further criterion of interest is LHB. While we do not attempt to time the onset of instability with LHB in these simulations, it is certainly fair to ask if the dynamical evolutions presented here are compatible with a long quiescent period preceding any resonant encounters. Two related issues immediately follow. First, will changing the placement of the inner boundary of the planetesimal swarm qualitatively change the process responsible for the onset of the instability of the dynamical evolutions? Second, how will the effective viscosity that arises from the self-gravity of the disk affect our results? A large-scale numerical modeling effort will be instrumental in providing these answers.  There is certainly room for broader study of the current setup of the problem as well. In this work, we have restricted ourselves to four-planet multi-resonant configurations. Certainly, the idea of initially forming more than two ice giants is not unreasonable (Ford \\& Chiang 2007, see however Levison \\& Morbidelli 2007). Although the results of Morbidelli et al. (2007) suggests that additional planets in a compact multi-resonant system compromise dynamical stability, more work is needed to obtain a good handle on this part of the problem.  In conclusion, the determination of a small subset of initial conditions allows for a much more efficient survey of the parameter space. In this work, we have taken a step in this direction. We must keep in mind that the system at hand is highly chaotic, and must in the end be studied numerically. The resulting determinations are often probabilistic rather than conclusive, however the results are certainly bound to gain statistical weight as the number of completed simulations increases. Thus, while much progress is yet to be made, additional research carries great value since a solid understanding of initial conditions plays an unavoidably important role in further development of a comprehensive model for solar system's formation."
    },
    "1004/1004.2833_arXiv.txt": {
        "abstract": "{} { The elemental abundance structure of the Galactic disc has been extensively studied in the solar neighbourhood using long-lived stars such as F and G dwarfs or K and M giants. These are stars whose atmospheres preserve the chemical composition of their natal gas clouds, and are hence excellent tracers of the chemical evolution of the Galaxy. As far as we are aware, there are no such studies of the inner Galactic disc, which hampers our ability to constrain and trace the origin and evolution of the Milky Way. Therefore, we aim in this study to establish the elemental abundance trend(s) of the disc(s) in the inner regions of the Galaxy. } { Based on equivalent width measurements in high-resolution spectra obtained with the MIKE spectrograph on the Magellan II telescope on Las Campanas in Chile, we determine elemental abundances for 44 K-type red giant stars in the inner Galactic disc, located at Galactocentric distances of 4-7\\,kpc. The analysis method is identical to the one recently used on red giant stars in the Galactic bulge and in the nearby thin and thick discs, enabling us to perform a truly differential comparison of the different stellar populations. } { We present the first detailed elemental abundance study of a significant number of red giant stars in the inner Galactic disc.  We find that these inner disc stars show the same type of chemical and kinematical dichotomy as the thin and thick discs show in the solar neighbourhood.  The abundance trends of the inner disc agree very well with those of the nearby thick disc, and also to those of the Bulge. The chemical similarities between the Bulge and the Galactic thick disc stellar populations indicate that they have similar chemical histories, and any model trying to understand the formation and evolution of either of the two should preferably incorporate both of them. } {} ",
        "introduction": "The inner Galactic disc is one of the least studied regions of the Milky Way because of high interstellar extinction and contamination by background Bulge stars. Apart from a few studies of bright hot OB stars \\citep[e.g.,][]{daflon2004} and Cepheids \\citep[e.g.,][]{luck2006}, which both trace the most recent young disc stellar population, almost no information is available about the detailed abundance structure of the inner Galactic disc. Open questions are for instance, whether the inner Galactic disc show the same clear kinematic and chemical dichotomy as the Galactic disc in the solar neighbourhood, where the thin and thick discs stand out as two distinct stellar populations?  Recent studies have revealed that the Galactic bulge and the Galactic thick disc have very similar abundance trends which reflect similar, and possibly even a shared, chemical histories \\citep{melendez2008,bensby2009,bensby2010,alvesbrito2010}. A restriction of these studies is that their thick disc samples have been observed in the solar neighbourhood, and if the Bulge has a secular origin \\citep[e.g.,][]{kormendy2004,howard2009}, models show that it likely has to be from gas and stars in the inner parts of the Galactic disc \\citep[e.g.,][]{rahimi2010}. Both the inner and the local disc will help us put constraints on how these Galactic components formed, if we can verify the existence of an inner Galactic thick disc and differentially compare it with the Bulge.  Here we will present the first results regarding detailed elemental abundances in 44 red giant stars that are located at 4-7\\,kpc from the Galactic centre. They have been analysed using the same method as in the recent study of red giants in the Bulge and nearby thin and thick discs by \\cite{alvesbrito2010}. We will focus on four $\\alpha$-elements (Mg, Si, Ca and Ti) and omit most of the analysis details and results for other iron-peak elements for a coming publication. ",
        "conclusions": "\\subsection{Distinct populations in the inner disc?}   In the Toomre diagram in Fig.~\\ref{fig:toomre} the stars have been coded according to the simple assumption that those with $v_{\\rm tot}>85\\,\\kms$ are thick disc stars, and those with lower velocities  are thin disc stars \\citep[e.g.,][]{fuhrmann2004}. Because we do not know the properties of the inner thick disc, the coding should not be taken literally. It is also obvious that the errors in the calculated space velocities make this classification uncertain. Hence, we just coded those stars that move on more circular orbits and those that have more kinematically hot orbits. Below we will call them kinematically hot stars (black circles) and kinematically cold stars (empty circles).  In Fig.~\\ref{fig:alfazmax}a we see that stars with distances greater than 2.5\\,kpc from the Sun have consistently high $\\rm [\\alpha/Fe]$ values ($\\rm \\alpha\\equiv (Mg + Si + Ti)/3$)). These distant stars are all located around or more than 1\\,kpc from the plane (Fig.~\\ref{fig:astrometry}), and they are essentially all kinematically hot stars. Then we see in Fig.~\\ref{fig:alfazmax}b that all stars located more than 1\\,kpc from the plane have high $\\rm [\\alpha/Fe]$ values, and that most of these are kinematically hot stars. We also see a few kinematically hot stars that are located close to the plane and also have low $\\rm [\\alpha/Fe]$ values. However, Fig.~\\ref{fig:alfazmax}c shows that these stars have kinematic properties that allow them to reach as far as $\\sim$2\\,kpc from the plane. At the same time, stars that are close to the plane, have low $\\rm [\\alpha/Fe]$ values, and are kinematically cold, stay within 1\\,kpc from the plane. We also note that we have a few stars with high $\\rm [\\alpha/Fe]$ values, which are kinematically hot, but which remain close to the the plane. These are stars that have highly eccentric orbits. Figure~\\ref{fig:alfazmax}d then shows that stars that move on highly eccentric orbits all have high $\\rm [\\alpha/Fe]$ values, and are all classified as kinematically hot. For stars with less eccentric orbits there is a gradual decrease in $\\rm [\\alpha/Fe]$ as the orbits become more circular. With a few exceptions, the stars with the least eccentric orbits have the lowest $\\rm [\\alpha/Fe]$ values. Figure~\\ref{fig:alfazmax}e  shows that stars with low [Fe/H] have high $\\rm [\\alpha/Fe]$, with a flat trend that eventually starts to decrease for metallicities higher than $\\rm [Fe/H]\\approx-0.3$.  Also, the stars with cold kinematics generally have higher [Fe/H] and lower $\\rm [\\alpha/Fe]$.  These connections and correlations between kinematics and chemistry that we see for the inner disc sample is what we see for disc stars in the solar neighbourhood. Stars with orbits that are highly eccentric and/or reach far from the plane generally have high $\\rm [\\alpha/Fe]$ values, and those on more circular orbits, which stay closer to the plane, have low $\\rm [\\alpha/Fe]$ values. Stars with these properties are generally classified as thick disc and thin disc stars, respectively \\citep[e.g.,][]{bensby2005,fuhrmann2004}. That we see the same correlations in the inner Galactic disc strongly suggests that we have two distinct disc populations also in the inner disc, an inner thin disc and an inner thick disc, similar to those in the solar neighbourhood.     \\subsection{The Galactic bulge -- thick disc connection}  In Fig.~\\ref{fig:haltplottar} we show the detailed abundance trends for four $\\alpha$-elements, comparing the 44 inner disc K giants to the Bulge giants and nearby thin and thick disc giants from \\cite{alvesbrito2010}. We emphasise that all stars in these plots have been analysed with the exact same methods, allowing truly differential comparisons between the different populations.  We note that especially the Mg abundance trend shows very little scatter, and that the inner disc giants have high [Mg/Fe] ratios for $\\rm [Fe/H]<-0.3$ and lower enhancements for higher [Fe/H]. This is a signature of enrichment by massive stars at low metallicities, and a delayed contribution from low mass stars at higher metallicities, consistent with the same signature seen in the nearby thick disc \\citep[e.g.,][]{feltzing2003}. This points to the existence of an inner thick disc and moreover that this thick disc does not differ much in terms of abundance trends, from the thick disc we see in the solar neighbourhhod. The same trend that is seen for Mg can also be seen in the Si and Ti plots, but with larger scatters. No clear trend can be seen in the Ca plot.  Furthermore, the abundance trends of the inner disc appear to be very similar to those of the Bulge. This inevitably points to a possible connection between the thick disc and the Bulge, implying they both might have formed at the same time \\cite[e.g.,][]{genzel2008}, sharing a similar star-formation rate and initial mass function. A possible scenario could be that the sub-solar part of the Bulge has a secular origin, and has formed from inner disc material \\citep[e.g.,][]{shen2010}.  \\begin{figure} \\resizebox{\\hsize}{!}{ \\includegraphics[bb=18 155 592 718,clip]{14809f3.eps}} \\resizebox{\\hsize}{!}{ \\includegraphics[bb=18 155 592 515,clip]{14809f4.eps}} \\caption{[$\\alpha$/Fe] versus $X$, $|Z|$, $Z_{\\rm max}$, $e$, and [Fe/H]. Symbols as in Fig.~\\ref{fig:toomre}. \\label{fig:alfazmax} } \\end{figure}  \\begin{figure*} \\resizebox{\\hsize}{!}{ \\includegraphics[bb=10 204 570 420,clip]{14809f5.eps} \\includegraphics[bb=18 204 592 420,clip]{14809f6.eps}} \\resizebox{\\hsize}{!}{ \\includegraphics[bb=10 155 570 405,clip]{14809f7.eps} \\includegraphics[bb=18 155 592 405,clip]{14809f8.eps}} \\caption{Abundance trends for our inner disc giants (red filled circles) together with Bulge giants (asterisks), nearby thin disc giants (green empty triangles); and nearby thick disc giants (blue empty squares), all from \\cite{alvesbrito2010}. Typical error bars are shown in each plot. \\label{fig:haltplottar} } \\end{figure*}  The agreement between the Bulge and the thick disc has recently also been seen in studies that compare Bulge stars with nearby thick disc stars. For instance, \\cite{bensby2010} presented a detailed abundance analysis of 15 microlensed dwarf stars in the Galactic bulge. These stars were found to share the same abundance trends as were traced by kinematically selected thick disc dwarf stars in the solar neighbourhood \\citep[][and 2010, in prep.]{bensby2003,bensby2005}. Similarly, \\cite{melendez2008}  and \\cite{alvesbrito2010} found very good agreement between the abundance trends of red giants in the Bulge and thick disc red giants in the solar neighbourhood (see also \\citealt{ryde2010}). Similar to this study, the analysis methods of these studies are internally fully consistent (same methods, model stellar atmospheres, atomic data, etc.). They compare dwarfs with dwarfs, and giants with giants. Other studies of red giants in the Bulge \\citep[e.g.][]{fulbright2007,zoccali2006,lecureur2007} have found that the Bulge is significantly more $\\alpha$-enhanced at higher metallicities than thin and thick disc stars. As discussed in \\cite{bensby2010} and \\cite{alvesbrito2010}, it is likely that those studies suffer from problems with the analysis (especially line blending). They also compare their Bulge giant samples with disc dwarf samples. The combined effect is that their Bulge stars seem spuriously more enhanced in the $\\alpha$-elements than the thick disc stars.  We further note that none of the inner disc giants are as metal-rich as some of the most metal-rich Bulge giants. As the metallicity distribution of the thick disc peaks at $\\rm [Fe/H]\\approx -0.6$ \\citep{carollo2010}, it is not surprising that our sample does not contain many metal-rich (thick disc) stars. Instead the upper metallicity limit appears to be close to, or slightly above, solar values (similar to what is seen for nearby thick disc dwarfs stars, \\citealt{bensby2007letter2}). A possible connection between the thick disc and the metal-rich part of the Bulge is therefore dubious. In that case, the metal-rich parts of the Bulge must have another origin, which possibly could be from accreted (extra-galactic?) material \\citep[see, e.g., the models by][]{rahimi2010}. Evidence for two co-existing formation scenarios within the Bulge was recently shown by \\cite{bensby2010} and \\cite{babusiaux2010}. With our result for the inner Galactic disc, the bonds between the the metal-poor part of the Bulge and the Galactic thick disc have grown even stronger."
    },
    "1004/1004.4318.txt": {
        "abstract": "We present rest-frame optical images and spectra of the gravitationally lensed, star-forming galaxy J0900+2234 ($z=2.03$). The observations were performed with the newly commissioned LUCIFER1 near-infrared (NIR) instrument mounted on the Large Binocular Telescope (LBT). We fitted lens models to the rest-frame optical images and found the galaxy has an intrinsic effective radius of $7.4\\pm0.8$~kpc with a lens magnification factor of about 5 for the A and B components. %We fit lens models to the rest-frame optical images and find the galaxy has an intrinsic %effective radius of 6~kpc with a $\\sim 5$ lens magnification of the A and B components. %We found that the C component is from a new galaxy at a higher redshift ($z\\sim3.5$) We also discovered a new arc belonging to another lensed high-z source galaxy, which makes this lens system a potential double Einstein ring system. Using the high S/N rest-frame optical spectra covering $H$+$K$ band, we detected H$\\beta$, [\\ion{O}{3}], H$\\alpha$, [\\ion{N}{2}], and [\\ion{S}{2}] emission lines. Detailed physical properties of this high-z galaxy were derived. The extinction towards the ionized \\ion{H}{2} regions ($E_g(B-V)$) was computed from the flux ratio of H$\\alpha$ and H$\\beta$ and appears to be much higher than that towards stellar continuum ($E_s(B-V)$), derived from the optical and NIR broad band photometry fitting. The metallicity was estimated using N2 and O3N2 indices. It is in the range of $\\frac{1}{5}-\\frac{1}{3}$ solar abundance, which is much lower than the typical $z\\sim2$ star-forming galaxies. From the flux ratio of [\\ion{S}{2}]$\\lambda6717$ and [\\ion{S}{2}]$\\lambda6732$, we found that the electron number density of the \\ion{H}{2} regions in the high-z galaxy were $\\simeq1000$~cm$^{-3}$, consistent with other $z\\sim2$ galaxies but much higher than that in local \\ion{H}{2} regions. The star-formation rate was estimated via the $H\\alpha$ luminosity, after correction for the lens magnification, to be about $365\\pm69~M_{\\sun}~yr^{-1}$. Combining the FWHM of H$\\alpha$ emission lines and the half-light radius, we found the dynamical mass of the lensed galaxy is $(5.8\\pm0.9)\\times10^{10}~M_{\\sun}$. The gas mass is $(5.1\\pm1.1)\\times10^{10}~M_{\\sun}$ from the H$\\alpha$ flux surface density by using global Kennicutt-Schmidt Law, indicating a very high gas fraction of $0.79\\pm0.19$ in J0900+2234. ",
        "introduction": "The redshifts between $1<z<3$ mark the time of peak cosmic star formation \\citep{madau98} and quasar activity \\citep{fan01}. The current-day Hubble sequence was being built up, and a large fraction of present-day stars were formed \\citep{dickinson03}. Observations of galaxies in this redshift range provide a direct picture of the formation and evolution of galaxies and the assembly of their central supermassive black holes.  By obtaining rest-frame optical spectra, we can perform detailed studies of the star formation history, as well as properties of the ISM and stellar populations in galaxies. However, well-known diagnostic optical emission lines are shifted into the near infrared (NIR) for $z=1-3$. Since high-z galaxies (e.g., Lyman Break Galaxies or LBGs) are faint, it is difficult to obtain high signal to noise ratio (S/N) NIR spectra with current facilities, especially for weak emission and absorption lines. Therefore, many previous NIR spectroscopic studies have been focused on the strong [\\ion{O}{3}] or H$\\alpha$ emission lines \\citep[e.g.,][]{pett01, erb03, erb06b, erb06c}, from which only limited inferences can be drawn.  One method of addressing this problem is to observe high-z star-forming galaxies which have been gravitationally lensed by foreground massive galaxies or clusters. Lensing can boost the observed flux of high redshift galaxies by several tens or more, so that high S/N spectroscopy in the NIR becomes feasible. One of the first lensed galaxies to be studied in this way was the serendipitously discovered MS1513-cB58 at $z=2.73$ \\citep{yee96, bechtold97, teplitz00}. Subsequently, several groups began systematic searches for strongly lensed high redshift galaxies towards clusters of galaxies \\citep[]{sand05, richard08} and red galaxies in the Sloan Digital Sky Survey \\citep[SDSS,][]{york00} images \\citep[e.g.,][]{belo09,kubo09}. These searches are finding lensed star-forming galaxies at redshifts $z>2$ \\citep[e.g.,][]{alla07,lin09, dieh09,koes10} which can then be followed-up in the NIR, or rest-frame optical and UV \\citep[e.g.,][]{smai07, hain09, fink09,quid09,quid10,pett10}.  The object of this paper is the lensed system J0900+2234 \\citep[z=2.03;][]{dieh09} which was discovered in SDSS images, and then confirmed with follow-up deep optical $g$, $r$, $i$ imaging and spectroscopy. The lens galaxy is at $z = 0.49$ based on SDSS spectroscopy. In this paper, we report $J$, $H$, $Ks$ imaging and $H$+$K$ spectroscopy of J0900+2234 with the newly commissioned LUCIFER instrument \\citep{mand08,Ager10} at the Large Binocular Telescope \\citep[LBT;][]{hill08}. %Large Binocular Telescope \\citep[LBT;][]{hill08} % Near-Infrared %Spectroscopic Utility with Camera and Integral Field Unit \\citep[LUCIFER;][]{Ager10}). We detected the nebular emission lines H$\\beta$, [\\ion{O}{3}]$\\lambda\\lambda$4959,5007, H$\\alpha$, [\\ion{N}{2}]$\\lambda$6583, and [\\ion{S}{2}]$\\lambda\\lambda$6717,6732 with simultaneous $H$+$K$ band spectral coverage from 1.40-2.20$\\mu m$. The wide simultaneous wavelength range coverage not only improves observing efficiency, but also helps to reduce the measurement uncertainties of the line ratios (e.g., H$\\alpha$/H$\\beta$) resulting from the flux calibration and slit losses variation between different exposures.  This paper is organized as follows: In \\S~\\ref{observation}, we describe the observations and reduction of the LUCIFER data. In \\S~\\ref{measurements}, we describe the photometry, reconstruction of the source with lens modeling and line flux measurements. In \\S~\\ref{lensmodel}, we construct a simple lensing model using the J0900+2234 NIR imaging. In \\S~\\ref{physics}, we report the detailed physical properties of J0900+2234 and discuss the systematic uncertainties of the physical property measurements. In \\S~\\ref{conclusion}, our main results are summarized. Throughout this paper, we use a cosmology with $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{\\rm M}=0.3$, and $\\Omega_{\\Lambda}=0.7$. The magnitudes are all AB magnitudes. ",
        "conclusions": "We present LBT/LUCIFER1 NIR (rest-frame optical) imaging and spectroscopy of the lensed galaxy J0900+2234 ($z=2.03$). The lensed components A and B were placed in the slit to obtain NIR spectra covered from 1.40 to 2.20$~\\mu$m. The detailed physical properties of the lensed star-forming galaxy were studied using the rest-frame optical spectra (Table~\\ref{phys}). We summarize the main results as follows.  \\begin{itemize}  \\item The new imaging was used to construct a lensing model. The magnification factor for the fluxes of the lensed galaxies (A plus B) is estimated to be 4.8.   \\item A new lensed arc A1 was discovered by the deep NIR $J$, $H$, and $Ks$-band images. The colors and position imply that this arc does not have the same source galaxy as the other four components discovered before, suggesting the presence of a rare double Einstein ring. Follow-up spectroscopy shows that the C component does not have the same source galaxy as A and B.   \\item We fit the optical and NIR broad band photometry to theoretical stellar population spectral templates (BC03), and found the galaxy age to be 180~Myr, $E_s(B-V) = 0.07$ ($E_s(B-V)=0.20$) for the A (B) component. The stellar mass is 1.9$\\times10^{10}~M_\\sun$, and the intrinsic luminosity of the galaxy is L/L* $\\approx$ $5-10$.  \\item Using the flux ratio of H$\\alpha$ and H$\\beta$, we found that the emission line gas suffers an extinction of $E_g(B-V)$ is $0.84\\pm0.31$ ($0.59\\pm0.08$) for the A (B) component, which is much higher than $E_s(B-V)$. This result implies that there is more extinction towards the ionized gas than the stellar continuum.  \\item The oxygen abundance was estimated using the N2 and O3N2 indices. We found that the metallicity in this star forming galaxy is only $21-35\\%$ solar, which is somewhat lower than typical star-forming galaxies at $z\\sim2$. The electron number density, $n_e$, was measured using the flux ratio of [\\ion{S}{2}]$\\lambda6717$ and [\\ion{S}{2}]$\\lambda6732$. The $n_e$ is in the range of $1029^{+3333}_{-669}$~cm$^{-3}$ ($1166^{+7020}_{-855}$ ~cm$^{-3}$) for the A (B) component. These values are consistent with other studies of high-z galaxies \\citep[e.g.,][]{hain09,quid09}, but much larger than that in local \\ion{H}{2} regions.  \\item The extinction-corrected luminosity of H$\\alpha$ emission line was used to estimate the SFR, and the intrinsic SFR of J0900+2234 is $365\\pm69~M_\\sun~\\rm yr^{-1}$ which is one magnitude higher than a typical $z\\sim2.0$ UV-selected star-forming galaxy. We estimate the Virial mass of J0900+2234 from the FWHM of the H$\\alpha$ emission line and find a dynamical mass of $(5.8\\pm0.9)\\times10^{10}~M_\\sun$. The gas mass of the galaxy is $(5.1\\pm1.1)\\times10^{10}~M_\\sun$ estimated from the H$\\alpha$ flux by assuming the global Kennicutt-Schmidt law.    \\end{itemize}  Comparison to the local star-forming galaxies, the J0900+2234 (z=2.03) features as much higher electron density. The high electron density and high ionization state \\citep[e.g.,][]{hain09} probably cause the offset between the \\ion{H}{2} regions in the high-z galaxies and those in the local galaxies in the BPT diagram. This star-forming galaxy also has a much higher SFR, and a relatively lower age, stellar mass, and metallicity than those in the typical high-z star-forming galaxies.     The commissioning of LUCIFER1 on the 8.4m LBT left primary mirror (SX) has allowed a detailed study of the $z=2$ galaxy J0900+2234 in relatively short exposure times. We look forward to future NIR observations of other lensed galaxies, as well as field galaxies which are not lensed, which will give a similarly richly detailed picture of galaxies in this important stage of galaxy evolution."
    },
    "1004/1004.3311_arXiv.txt": {
        "abstract": "The Diffuse Supernova Neutrino Background (DSNB) is the weak glow of MeV neutrinos and antineutrinos from distant core-collapse supernovae. The DSNB has not been detected yet, but the Super-Kamiokande (SK) 2003 upper limit on the $\\bar{\\nu}_e$ flux is close to predictions, now quite precise, based on astrophysical data. If SK is modified with dissolved gadolinium to reduce detector backgrounds and increase the energy range for analysis, then it should detect the DSNB at a rate of a few events per year, providing a new probe of supernova neutrino emission and the cosmic core-collapse rate. If the DSNB is not detected, then new physics will be required. Neutrino astronomy, while uniquely powerful, has proven extremely difficult -- only the Sun and the nearby Supernova 1987A have been detected to date -- so the promise of detecting new sources soon is exciting indeed. \\newpage ",
        "introduction": "The Diffuse Supernova Neutrino Background (DSNB) is the flux of neutrinos and antineutrinos emitted by all core-collapse supernovae in the causally-reachable universe; it will appear isotropic and time-independent in feasible observations. (Hereafter, neutrinos means neutrinos and antineutrinos, and supernovae means core-collapse supernovae, unless specified.) It results from the quasi-thermal MeV neutrino spectrum emitted by newly-formed neutron stars convolved with the rapid redshift evolution of the supernova rate. The cosmic energy density in neutrinos from core-collapse supernovae is comparable to that in photons from stars, $\\sim 0.01$ eV cm$^{-3}$, and is $\\sim 10$ times less than that of the cosmic microwave background.  Some earlier literature referred to ``supernova relic neutrinos,\" but ``relic\" often caused confusion with the $10^{-4}$ eV neutrinos from the Big Bang.  \\medskip  Why is detection of the DSNB an important and relevant experimental goal?  \\smallskip $\\bullet$ {\\bf Understanding supernovae is crucial to astrophysics and physics.} Supernovae are central to the cosmic history of stellar birth and death; the production of chemical elements, neutron stars and black holes, cosmic rays, and gravitational waves; and to exploiting the extreme physical conditions of core collapse to probe the properties of neutrinos and hypothetical particles.  New observations and syntheses are needed.  \\smallskip $\\bullet$ {\\bf We cannot understand supernovae without detecting neutrinos.} The optical supernova reveals little about the collapsing core itself, which is embedded in the stellar envelope; the total energy in photons is much less than that in neutrinos, leaves over a vastly longer time, and depends on the properties of the progenitor star. Neutrinos are emitted from the core, revealing its properties.  \\smallskip $\\bullet$ {\\bf Detecting bursts of neutrinos from nearby supernovae is difficult.} Prodigious emission and large detectors partially compensate the tiny detection cross section.  For Milky Way supernovae (distance $D \\sim 10$ kpc; $1 {\\rm\\ pc} = 3.1 \\times 10^{18} {\\rm\\ cm}$), a neutrino burst would be detected easily, but the burst rate is only a few per century. For supernovae in nearby galaxies ($D \\sim$ 1--10 Mpc), where the total burst rate is above one per year, much larger detectors will be required.  \\smallskip $\\bullet$ {\\bf The DSNB is a guaranteed steady source of supernova neutrinos.} While the probability of detecting even a single neutrino from a distant supernova is infinitesimally small, the number of supernovae per year is astronomically large, yielding a nonzero rate. The total rate and spectral shape of the DSNB are new probes of supernova neutrino emission and the cosmic core-collapse rate.  \\medskip  {\\it The DSNB has not been detected yet, but the discovery prospects are excellent.} Super-Kamiokande (SK) is large enough to detect the DSNB at the level of a few events per year.  These events are hidden by detector backgrounds, which are already low and could be substantially reduced with added gadolinium to detect neutrons, a proposed upgrade for which there is active research and development.  Detection of the DSNB will yield science that is difficult to obtain otherwise. It will measure the average supernova neutrino emission spectrum, and comparison to burst detections will probe the star-to-star variation.  It will also be sensitive to the optically-dark failed supernovae expected from high-mass stars.  This review is written for experimentalists, observers, and theorists in nuclear physics, particle physics, and astrophysics, and covers the many aspects of the DSNB.  It is intended to be an accessible overview of the main themes needed to predict, detect, and interpret the DSNB, with details given in references. ",
        "conclusions": "Neutrino astronomy can revolutionize astrophysics and physics, due to its unique abilities to see unseen processes and to probe subtle new particle interactions. Stellar fusion reactions and the core-collapse mechanism were revealed by the detection of neutrinos from the Sun and the nearby SN 1987A. Solar neutrinos were crucial to the discovery of neutrino mixing, and supernova neutrinos probed exotic physics far beyond the reach of laboratory experiments.  However, the great promise of neutrino astronomy remains largely unfulfilled, as no other neutrino sources have yet been detected. Recall the courage of the pioneers of this field, who, more than forty years ago, had neither adequate detectors nor accurate astrophysics to fortify their hopes.  We have advantages of many orders of magnitude on each, and can see the end of the desert.  To make progress on our understanding of supernovae and their role in astrophysics, as well as of the neutrinos themselves, we must detect more supernova neutrinos.  \\medskip  Fortunately, first detection of the DSNB is within reach. Why are we so sure?  \\smallskip $\\bullet$ {\\bf The framework for calculating the DSNB is solid.} \\\\ It is built on simple principles of cosmology, well tested with photons.  \\smallskip $\\bullet$ {\\bf Supernova neutrino emission has been measured from SN 1987A.} \\\\ Theory and indirect evidence suggest it should typically be larger.  \\smallskip $\\bullet$ {\\bf The uncertainty on the cosmic supernova rate is modest.} \\\\ Observations of massive stars and supernovae are precise and quickly improving.  \\smallskip $\\bullet$ {\\bf All aspects of DSNB signal detection are understood.} \\\\ Neutrino interactions and detector properties have each been well measured.  \\medskip Therefore, even including uncertainties, there is a robust prediction that {\\it there must be a few DSNB neutrino events per year in SK}. These events are hidden by backgrounds -- what about the experimental prospects for revealing them?  \\smallskip $\\bullet$ {\\bf SK has already shown sensitivity quite close to these predictions.} \\\\ Importantly, SK has controlled detector backgrounds nearly to the required level.  \\smallskip $\\bullet$ {\\bf There are good prospects for improving the SK DSNB sensitivity.} \\\\ With gadolinium, SK would reduce its backgrounds and energy threshold.  \\medskip  Therefore, the DSNB could be discovered in SK.  With sufficient detector background reduction, even a few signal events will be statistically significant and scientifically important. If the DSNB is not found, there must be surprising new physics, requiring an upheaval in astrophysics or exotic new particle physics.  The enduring value of a DSNB measurement will come from its comparison to other observations and theory, which include the SN 1987A data and numerical simulations of core collapse. The SK 2003 limit is already an important new result, as it excludes neutrino emission as large as often assumed, and is only a factor 2--4 away from more conservative cases. Significant developments in the theoretical predictions are also eagerly anticipated, as, after more than forty years of work, recent results have increasingly realistic physics and accurate calculations, show explosions in some cases, and some are able to integrate to several seconds. These models will provide a rich context for interpreting DSNB results and, eventually, to novel tests of neutrino properties and other new physics.  In the next several years, we will gain a newly deep and comprehensive understanding of supernovae and their neutrinos, following from the above improvements converging with unprecedented electromagnetic data, results from gravitational wave observatories, and, with a little luck, a Milky Way neutrino burst.   \\newpage  {\\sc Disclosure Statement} \\\\ The author is not aware of any affiliations, memberships, funding, or financial holdings that might be perceived as affecting the objectivity of this review.  \\bigskip  {\\sc Acknowledgments} \\\\ The author's research was supported by NSF CAREER Grant PHY-0547102.  I am grateful to my collaborators on the DSNB -- Shin'ichiro Ando, Eli Dwek, Brian Fields, Andrew Hopkins, Shunsaku Horiuchi, Amy Lien, Louie Strigari, Mark Vagins, Terry Walker, Hasan Y\\\"uksel, and Pengjie Zhang -- as well as to my additional collaborators on closely-related topics -- Dick Boyd, Will Farr, Joe Formaggio, Matt Kistler, Chris Kochanek, Tony Mezzacappa, Jos\\'{e} Prieto, Matthew Sharp, Kris Stanek, Yoichiro Suzuki, Todd Thompson, and Petr Vogel -- for sharing their knowledge, enthusiasm, and wisdom.  I especially thank Shin'ichiro Ando, Andrew Hopkins, Shunsaku Horiuchi, Todd Thompson, Mark Vagins, and Hasan Y\\\"uksel for helpful comments on the draft manuscript."
    },
    "1004/1004.2796.txt": {
        "abstract": "Microlensing light curves are typically computed either by ray-shooting maps or by contour integration via Green's theorem. We present an improved version of the second method that includes a parabolic correction in Green's line integral. In addition, we present an accurate analytical estimate of the residual errors, which allows the implementation of an optimal strategy for the contour sampling. Finally, we give a prescription for dealing with limb-darkened sources reaching arbitrary accuracy. These optimizations lead to a substantial speed-up of contour integration codes along with a full mastery of the errors. ",
        "introduction": "Microlensing is one of the most promising methods for finding the first Earth-like extrasolar planet \\citep{Gaudi2010,DomRev}. When a compact object transits very close to the line of sight of a background source star, the flux coming from the source is amplified by gravitational lensing and follows a typical bell-shape light curve, analytically described by \\citet{Pacz1986}. If the lens is a star accompanied by a secondary body like a planet, an additional bump or dip appears on the light curve \\citep{MaoPacz}. The timescale of such features ranges from a few hours to a few days, depending on the square root of the mass of the planet \\citep{GouLoe}. Such short timescales require intensive monitoring by telescopes situated all over the Earth.  Nowadays, more than 30 telescopes are involved in microlensing searches towards the Galactic bulge. More than 600 events are discovered every year. Out of these, roughly from 10 to 20 events show anomalies that can be interpreted as due to binary lenses. Some of these are finally accepted as showing evidence of an extrasolar planetary system. Since the start of microlensing searches, 26 events have been reported as containing planetary candidates with stronger or weaker evidence \\cite{DomRev}. Nine of these events have also been included in the most updated exoplanet list available on the web (http://exoplanet.eu) \\citep{Beaulieu2006,Bennett2008,Bond2004,Dong2009,Gaudi2008,Gou2006,Janczak2009,Sumi2010,Udalski2005}. Finally, one event brings the spectacular signature of two planets in the same system \\citep{Gaudi2008}.  Unfortunately, interpreting binary microlensing events is a very long process that may even take from one to four years for a single event. This is due to several reasons: one is related to the difficulty of getting rid of all systematic errors in the photometry of each dataset. Some sets of images need to be reduced with different methods in order to compare the effects of systematics, and subtract them from the final result. After the reduction process is complete, the final datasets may contain hundreds or thousands of data points, which are then ready for the modelling process.  The modelling process is typically driven by the criterium of $\\chi^2$ minimization, which can be achieved either by downhill algorithms or by Markov Chain Monte Carlo (MCMC) methods. In order to evaluate the $\\chi^2$ for a single tentative model, it is necessary to compute the microlensing magnification at each data point. However, as is well known, microlensing magnification cannot be calculated analytically if the lens is a binary system. As the angular extension of the source plays a major role in the observed magnification, the computation of even one single model point is relatively time-consuming. Multiplying this basic time unit for the number of data points, the number of models within a MCMC simulation, the number of different hypotheses to be checked (parallax, xallarap, orbital motion, limb darkening, binary source, $\\ldots$), we can easily imagine why the study of a single event takes so long. It is then critical to reduce the computational time of a single model point as much as possible, so that the whole modelling process is cut down to a more reasonable duration.  There are basically two classes of methods for the computation of the microlensed flux of a source. The first is based on the construction of ray-shooting maps \\citep{Kayser}. In practice, rays are shot back from the observer to the lens plane and then deflected to the source plane. If they intercept the source disk they are counted as contributing to the total magnification. This method has three main advantages: it is conceptually simple, it can naturally take into account the limb darkening profile of the source, maps at fixed lens configurations can be re-used for different source positions. Numerous optimized versions have appeared in the literature, designed for re-use of maps \\citep{Wambsganss,Rattenbury,Dong2006}. An alternative strategy is to speed-up the computation for a single source position giving up the map re-use \\citep{BenRhi}. In this case, starting from the positions of the centers of the images, one shoots rays only where really needed. A further improvement of this method employing a polar coordinate-grid with an optimized prescription to handle limb darkening has been recently presented by \\citet{Bennett2009}.  The second method is based on an application of Green's theorem (which can be viewed as a two-dimensional version of Stokes' theorem). In practice, one can find the area of an image by calculating a Riemann integral along the image contour. The use of contours in gravitational lensing dates back to \\citet{SchKay}. Green's theorem was then used for calculating areas by \\citet{Dom93}. The method has been refined by \\citet{Dom95} and then applied to microlensing by \\citet{GouGau} and \\citet{Dom98}. The appeal of this method is that a two-dimensional calculation is turned into a one-dimensional calculation, which is much faster, in principle. Related to this approach are the algorithms presented by \\citet{Dong2006} and the adaptive grid search by \\citet{Dominik2007}. The main advantages come from the potentially high computational velocity and the high flexibility for models in which the lens configuration changes (e.g. in the treatment of planet orbital motion). However, the contour integration approach requires an images reconstruction procedure (which can be sometimes complicated); in addition, limb darkening cannot be naturally incorporated in the algorithm. In particular, the latter limitation has oriented the community to give a general preference to ray-shooting methods.  Nevertheless, apart from its undisputed elegance, the contour integration approach is still competitive for obtaining preliminary microlensing models very quickly, which is particularly interesting in view of the realization of real-time modelling of binary microlensing events. Furthermore, when orbital motion is relevant, traditional ray-shooting methods typically become definitely too heavy. In this case, only adaptive methods \\citep{BenRhi,Bennett2009} can compete with Green's theorem algorithms.  Finally, besides ray-shooting and contour integration methods, it is worth mentioning that when the source size is only marginally relevant (for sources not too close to caustics), one can approximate its effects by quadrupole or hexadecapole approximations \\citep{Gould2008,PejHey}. These methods allow to obtain a substantial speed-up of the code avoiding useless heavy computations when the source size correction is small. They can be used in combination with other methods that may intervene when the source gets closer to a caustic.  In this paper we present four new ideas for boosting codes based on contour integration approach. In Section 2 we show how Green's line integral can be approximated to third order introducing a parabolic correction, with a substantial improvement in accuracy. In Section 3 we present accurate estimates of the residual errors in Green's integral. These are used to implement an optimal sampling strategy that allows to minimize the calculations for a given required accuracy, as explained in Section 4. In Section 5 we suggest an easy prescription for the treatment of limb darkening that achieves a fixed accuracy avoiding lengthy calculations. The benefits achieved by all these innovations are documented by several numerical examples in Section 6. ",
        "conclusions": "In this paper we have presented four innovations for codes attempting microlensing calculations based on contour integration. In this class of codes the contours of the images are obtained by inversion of the lens map at the source boundary; the area of the images is then obtained by a simple contour integration rather than by a surface integration.  We have introduced a parabolic third order correction in the evaluation of Green's line integral, which leaves a residual of the fifth order in the interval size $\\Delta \\theta$. The speed-up with respect to the classical trapezium approximation is around 4 for a target accuracy in the magnification of $\\delta \\mu=10^{-2}$.  We have introduced accurate error estimators, whose reliability has been shown by comparing several magnification maps obtained at different levels of target accuracy.  Thanks to these error estimators, we have proposed a new optimal sampling strategy driven by the error estimates in each sampling interval. With respect to a uniform sampling we get a speed-up ranging from 3 to 20 at $\\delta \\mu=10^{-2}$ depending on how large is the magnification experienced by the source in its microlensing trajectory.  Finally, we have faced the problem of limb darkening, which is the hardest obstacle for codes based on contour integration. Also in this case we have introduced error estimators and an optimal sampling strategy with the aim of minimizing calculation keeping full control of the accuracy. As a result, we have a very reliable limb darkening approximation to any desired accuracy. At $\\delta \\mu=10^{-2}$ the slow-down with respect to a uniform brightness disk ranges from 2 to 11.  Summing up, the speed-up gained by the optimal sampling is sufficient to compensate the slow-down due to the migration from uniform brightness disks to realistic limb darkened sources. In addition, the parabolic correction guarantees a net speed-up with respect to traditional linear Green's theorem codes. Finally, the full control of the errors is an invaluable help in order to avoid redundant calculations and concentrate the efforts where it is really needed. Its impact with respect to traditional codes with fixed sampling is definitely huge and difficult to quantify.  All these innovations should put codes based on contour integration in the front line for binary and planetary microlensing events modelling."
    },
    "1004/1004.2856_arXiv.txt": {
        "abstract": "We have investigated the variation of magnetic helicity over a span of several days around the times of 11 X-class flares which occurred in seven active regions (NOAA 9672, 10030, 10314, 10486, 10564, 10696, and 10720) using the magnetograms taken by the Michelson Doppler Imager (MDI) on board the $Solar$ $and$ $Heliospheric$ $Observatory$ ($SOHO$). As a major result we found that each of these major flares was preceded by a significant helicity accumulation, (1.8--16)$\\times$10$^{42}$ Mx$^2$ over a long period (0.5 to a few days). Another finding is that the helicity accumulates at a nearly constant rate, (4.5--48)$\\times$10$^{40}$ Mx$^2$ hr$^{-1}$, and then becomes nearly constant before the flares. This led us to distinguish the helicity variation into two phases: a phase of monotonically increasing helicity and the following phase of relatively constant helicity. As expected, the amount of helicity accumulated shows a modest correlation with time-integrated soft X-ray flux during flares. However, the average helicity change rate in the first phase shows even stronger correlation with the time-integrated soft X-ray flux. We discuss the physical implications of this result and the possibility that this characteristic helicity variation pattern can be used as an early warning sign for solar eruptions. ",
        "introduction": "Magnetic helicity is a measure of twists, kinks, and inter-linkages of magnetic field lines (Berger \\& Field 1984) and has been an important parameter in solar dynamo theories (Parker 1955). While the source of the magnetic helicity lies below the surface of the Sun, it was recently recognized as a useful parameter in describing solar features observed above surface such as spiral patterns of sunspot fibrils, helical patterns in filaments and coronal mass ejections (CMEs; for a review, see Brown et al. 1999). Naturally magnetic helicity studies have been directed to the energy buildup and instability leading to eruptions and CMEs (e.g., Rust 2001; Kusano et al. 2004; Phillips et al. 2005). More recently, several studies were carried out to relate the change of magnetic helicity to the problem of impending or triggering solar flares. Moon et al. (2002a) studied the magnetic helicity change around major flares to find its rapid helicity change before flares, and concluded that a sudden helicity injection may trigger flares. Moon et al. (2002b) applied the same approach to seven homologous flares in the active region, NOAA 8100, over a period of 6.5 hr to find a good correlation between the amount of incremental helicity and the soft X-ray flux during each homologous flare. The results from both studies thus point to the idea that the helicity change occurring over short timescales (around a half hour) can be a significant factor in triggering flares. Kusano et al. (2003) proposed annihilation of magnetic helicity as a triggering mechanism for solar flares. Numerical simulations were carried, which show that, if the helicity is sharply reversed within a magnetic arcade, reconnection quickly grows in the helicity inversion layer, driving explosive dynamics. Yokoyama et al. (2003) studied flare activities in active region NOAA 8100 to find that most of the flare events occurred about half a day after the helicity injection rate changed its sign, and the positions of H$\\alpha$ emission in flares well correspond to the helicity inversion lines in space.  Sakurai \\& Hagino (2003) studied two active regions appeared in 2001 (NOAA 9415 and 9661), both of which have produced X-class flares. Their finding was, on the contrary, that the magnetic helicity integrated over the regions evolved slowly and did not show abrupt changes at the time of the flares, although the distributions of magnetic helicity changed significantly over a few days in the regions.  In this paper, we study long term (a few days) variations of the magnetic helicity around major X-class flares. While some of the above studies suggest short-term helicity change as an important topic for flare triggering, Hartkorn \\& Wang (2004) found that the rapid helicity change at the time of a flare can occur as an artifact under the influence of flare emission on the spectral line adopted in MDI measurements. This means that a short term variation during strong flares can hardly be measured with enough accuracy. However we can study a long-term variation of the magnetic helicity with active region magnetograms gathered over many days excluding the times of flares. Our research is therefore focused on a possible characteristic helicity evolution pattern that is associated with flare impending mechanisms. ",
        "conclusions": "We have investigated the variation of magnetic helicity over a time span of several days around the times of 11 X-class flares which occurred in seven active regions using MDI magnetograms to find the following results.  First, a substantial amount of helicity accumulation is found before the flare in all the events. The helicity increases at a nearly constant rate, (4.5--48)$\\times$10$^{40}$ Mx$^2$ hr$^{-1}$, over a period of 0.6 to a few days, resulting in total amount of helicity accumulation in the range of (1.8--16)$\\times$10$^{42}$ Mx$^2$. Such a wide range of helicity accumulation indicates that each active region has its own limit of helicity storage to keep a stable magnetic structure in the corona. The finding of a monotonically increasing phase is similar to the earlier one by Sakurai \\& Hagino (2003) that the magnetic helicity integrated over the regions evolved slowly and did not show abrupt changes at the time of the flares. The helicity increase over days before the flares reconfirms the conventional idea that helicity accumulation by a certain amount is necessary for a large flare to occur (Kusano et al. 1995; Choe \\& Lee 1996).  Second, there is a strong positive correlation between the average helicity change rate of phase I and the corresponding GOES X-ray flux integrated over the flaring time. The amount of helicity accumulation during phase I also correlates with the soft X-ray flux, as expected, but the correlation is stronger with the helicity change rate. This result probably implies that the helicity change rate is more accurately determined than the amount of helicity change itself as the initial time of helicity build-up is poorly determined.  If the above correlations hold for a large number of events, we may predict the flare strength (the integrated X-ray flux) based on the helicity change rate. Monitoring of helicity variation in target active regions may also aid the forecasting of flares. A warning sign of flares can be given by the presence of a phase of monotonically increasing helicity, as we found that all the major flares occur after significant helicity accumulation. As a reference we have checked helicity variation of the six active regions in non-flaring times to find much lower helicity change rates compared with those around the major flares. We thus conclude that the relative magnetic helicity can be a powerful tool for predicting major flares."
    },
    "1004/1004.3688_arXiv.txt": {
        "abstract": "We explain the motivation and main results of our work in Ref.~\\refcite{Lehners:2009ja}. Using the covariant formalism, we derive the equations of motion for adiabatic and entropy perturbations at third order in perturbation theory for cosmological models involving two scalar fields, and use these equations to calculate the trispectrum of ekpyrotic and cyclic models. The non-linearity parameters $f_{NL}$ and $g_{NL}$ are found to combine to leave a very distinct observational imprint. ",
        "introduction": "The analysis of non-Gaussianity \\cite{Komatsu:2009kd}, {\\it i.e.} the deviation from perfect Gaussian statistics, is becoming a strong discriminator of competing models of the early universe. Whilst most efforts until now have been devoted to the study of the bispectrum -- the Fourier transform of the 3-point function of the curvature perturbation $\\zeta$ -- recent progress in observational cosmology is such that it may be possible in future to detect the trispectrum (4-point correlation function of $\\zeta$), if it is large enough. However, the calculational technology required to make the corresponding predictions is not yet in place. In particular, analytic formulae can be obtained from the $\\delta N$ formalism, the most widely used technique, only in very special cases and furthermore its numerical implementation is demanding. Here we report on a recent development: the derivation of the equations of motion to third order in perturbation theory, which are needed to calculate the non-Gaussianity in two-scalar-field cosmological models. The resulting equations are generally applicable and well suited to numerical analyses, as demonstrated by our calculation of the trispectrum generated by ekpyrotic and cyclic cosmological models --- scenarios in which the primordial fluctuations are created during a collapsing phase before the Big-Bang (see Ref.~\\refcite{Lehners:2008vx} for a review). ",
        "conclusions": "Motivated by huge observational perspectives in coming years concerning non-Gaussianity measurements, we have developed in Ref.~\\refcite{Lehners:2009ja} the necessary tools to compute the trispectrum generated in  two-field cosmological models. Our results are generally applicable and well suited for numerical analyses and we used them to calculate the trispectrum generated in ekpyrotic and cyclic models. The combined consideration of the bi- and trispectrum in that case results in a distinct observational imprint that should enable one to select or rule these models on observational grounds in the foreseeable future."
    },
    "1004/1004.1644_arXiv.txt": {
        "abstract": "The Milky Way's dark matter halo is thought to contain large numbers of smaller subhalos. These objects can contain very high densities of dark matter, and produce potentially observable fluxes of gamma rays. In this article, we study the gamma ray sources in the Fermi Gamma Ray Space Telescope's recently published First Source Catalog, and attempt to determine whether this catalog might contain a population of dark matter subhalos. We find that, while approximately 20-60 of the catalog's unidentified sources could plausibly be dark matter subhalos, such a population cannot be clearly identified as such at this time. From the properties of the sources in the First Source Catalog, we derive limits on the dark matter's annihilation cross section that are comparably stringent to those derived from recent observations of dwarf spheroidal galaxies. ",
        "introduction": "In the standard theory of hierarchical structure formation, cold dark matter particles gather to form small halos, which later merge to form ever more massive halos. As a result of this process, dark matter halos (such that that hosting the Milky Way) are expected to contain many subhalos, ranging from dwarf spheroidal galaxies ($M \\gsim 10^7 M_{\\odot}$) down to objects with masses as small as $10^{-3}-10^{-8} M_{\\odot}$. The minimum mass of dark matter subhalos is determined by the temperature at which the dark matter particles became kinetically decoupled from the cosmic neutrino background, which in turn depends on the particle physics of the dark matter candidate.  With the exception of dwarf spheroidal galaxies, efforts to directly observe dark matter subhalos have failed to detect any such objects. A promising method to search for nearby dark matter substructures is to use gamma ray telescopes, such as the Fermi Gamma Ray Space Telescope (FGST), to detect the products of dark matter annihilations. To such a telescope, relatively large ($\\sim 10^3-10^7 M_{\\odot}$) and nearby ($\\sim 0.01-10$ kpc) subhalos could appear as as bright and point-like gamma ray sources, without counterparts in other wavelengths. For reasonable assumptions (based on the results of numerical simulations) regarding the number of subhalos in the Milky Way and the dark matter distribution within those subhalos, one is led to expect a handful of subhalos to be observable by the FGST. For example, a dark matter candidate in the form of a 50 GeV particle with a annihilation cross section of $\\sigma v \\sim 3 \\times 10^{-26}$ cm$^3$/s is predicted to provide on the order of a few subhalos that would be observable at or above the $5\\sigma$ level~\\cite{Pieri:2007ir}.  The FGST collaboration has recently published a catalog of 1451 point sources~\\cite{catalog}, including 630 that are not associated with sources in other astronomical catalogs. Of these unidentified sources, 368 have been detected with greater than 5$\\sigma$ significance, and are more than 10 degrees away from the Galactic Plane. In this article, we study this collection of gamma ray sources and consider whether any significant fraction might be the result of dark matter annihilations taking place in nearby subhalos.  The remainder of this article is structured as follows. In Sec.~\\ref{theory}, we calculate the number of dark matter subhalos predicted to be observable by FGST, for various astrophysical assumptions and particle dark matter candidates. In Sec.~\\ref{catalog}, we determine which of the objects contained within the FGST's First Source Catalog can be fit by dark matter annihilations, for various combinations of mass and annihilation channel. We use this information in Sec.~\\ref{limits} to derive upper limits on the dark matter's annihilation cross section. In Sec.~\\ref{hints}, we discuss some of the more interesting features of the First Source Catalog, and consider whether they might result from a population of relatively large and nearby subhalos. We summarize our results in Sec.~\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  In this article, we have studied the recently published Fermi Gamma Ray Space Telescope's First Source Catalog, and considered whether a significant number of the sources in this catalog might be dark matter subhalos. For a typical thermal dark matter candidate with a mass of 50 GeV, and reasonable astrophysical assumptions, one predicts that a few relatively large ($\\sim 10^3-10^7 M_{\\odot}$) and nearby ($\\sim 0.01-10$ kpc) subhalos would be bright enough to appear within such a catalog. It is unlikely, however, that such a small number of sources could be identified from among the hundreds of unidentified sources in the Fermi First Source Catalog. If the dark matter's annihilation cross section were larger than that typically predicted for a thermal relic, however, it is possible that a larger number of subhalos could be detectable by gamma ray telescopes.  From among Fermi's First Source Catalog, we studied 368 sources as dark matter subhalo candidates, each of which has been detected with high significance ($> 5 \\sigma$), is located away from the Galactic Plane ($|b| > 10^{\\circ}$), and is not associated with any source observed at other wavelengths. Although we find that the spectrum of many of these sources could be well fit by that predicted from annihilating dark matter, this is most likely the result of the relatively large error bars involved, and on the wide range of dark matter masses and annihilation channels we have considered.  By studying the distribution of source spectra in the First Source Catalog, we derived an upper limit on the annihilation cross section of dark matter, as a function of its mass and dominant annihilation channel. For dark matter particles with relatively low masses (less than a few hundred GeV), we find constraints that are comparably stringent to those derived from observations of dwarf spheroidal galaxies. For example, in the case of a 50 GeV dark matter particle that annihilates to $b\\bar{b}$, we conclude that $\\sigma v \\lsim 7 \\times 10^{-26}$ cm$^3$/s.  Lastly, we note the appearance of bump-like features in the distribution of spectra observed in the Fermi First Source Catalog. Such features could be explained ({\\it ie.} flattened) if the dark matter takes the form of a 500 GeV particle with an annihilation cross section to $\\tau^+ \\tau^-$ of $\\sigma v \\sim 6 \\times 10^{-23}$ cm$^3$/s, or a 180 GeV particle with an annihilation cross section to $W^+ W^-$ of $\\sigma v \\sim (1-2) \\times 10^{-24}$ cm$^2$/s (such as is predicted for a wino-like neutralino). These scenarios could also potentially account for the cosmic positron excess observed by PAMELA.   \\bigskip   {\\it Acknowledgements:} MB is supported by the US Department of Energy, under grant DE-FG03-92-ER40701. DH is supported by the US Department of Energy, including grant DE-FG02-95ER40896, and by NASA grant NAG5-10842."
    },
    "1004/1004.4899_arXiv.txt": {
        "abstract": "We conduct numerical experiments by evolving gaseous/stellar disks embedded in live dark matter halos aiming at quantifying the effect of gas spatial resolution and gas content on the bar evolution. Three model sequences have been constructed using different resolution, and gas fraction has been varied along each sequence within the range of \\fg$= 0\\%-50\\%$, but keeping the disk and halo properties unchanged. We find that the spatial resolution becomes important with an increase in the gas content.  For the higher resolution model sequences, we observe a bimodal behavior in the bar evolution with respect to the gas fraction, especially during the secular phase of this evolution. The switch from the gas-poor to gas-rich behavior is abrupt and depends on the resolution used, being reasonably confined to \\fg$\\sim 5\\%-12\\%$. The diverging evolution has been observed in nearly all basic parameters characterizing bars, such as the bar strength, central mass concentration, bar vertical buckling amplitude, bar size, etc. We find that the presence of the gas component severely limits the bar growth and affects its pattern speed evolution. Gas-poor models display rapidly decelerating bars, while gas-rich models exhibit bars with constant or even slowly accelerating tumbling. We also find that the gas-rich models have bar corotation (CR) radii within the disk at all times, in contrast with gas-poor and purely stellar disks. In addition, the CR-to-bar size ratio is less than 2 for gas rich-models. Next, we have confirmed that the disk angular momentum within the CR remains unchanged in the gas-poor models, as long as the CR stays within the disk, but experiences a sharp drop before leveling off in the gas-rich models. Finally, we discuss a number of observed correlations between various parameters of simulated bars, such as between the bar sizes and the gas fractions, between the bar strength and the buckling amplitude, the bar strength and its size, etc. ",
        "introduction": "Many issues related to the formation and evolution of galactic bars remain unsettled (e.g., Villa-Vargas, Shlosman \\& Heller 2009, hereafter Paper~I). In Paper~I, we have revisited the properties of barred disks embedded in the dark matter (DM) halos, with an emphasis on the angular momentum redistribution in {\\it collisionless} systems. Here we attempt to understand some effects of the gas component on the bar evolution. In a follow-up work, we investigate the effect of varying the halo parameters in the gaseous/stellar disk models.  Due to its viscous and dissipative nature, the gas can influence the stellar component in the disk, well beyond its observed mass fraction (e.g., Shlosman \\& Noguchi 1993; Friedli \\& Benz 1993; Heller \\& Shlosman 1994; Berentzen et al. 1998; Bournaud \\& Combes 2002; Berentzen et al. 2007; Curir et al. 2007). The gas dissipation, especially in the bar's presence, triggers the mass redistribution in the disk. A substantial fraction of barred orbits are self-intersecting, which while irrelevant for stars has a strong effect on the gas. The gas populating such orbits is short-lived, resulting in shocks and loss of rotational support. Different viscosities in gas and stellar `fluid' lead to a delayed response to the gravitational torques and to an exchange of angular momentum between these components. The loss of angular momentum in the gas induces a flow down to the inner disk, forming a central mass concentration (CMC), and possibly forming and fueling the central supermassive black hole (SBH) via nested bars mechanism (e.g., Shlosman et al. 1989, 1990; Pfenniger \\& Norman 1990; Friedli \\& Martinet 1993; Knapen et al. 1995; Begelman \\& Shlosman 2009; Hopkins \\& Quataert 2010), and additional processes closer to the SBH (e.g., Shlosman 1999; Hopkins et al. 2010). The gas is capable of modifying the disk orbital structure (e.g., Berentzen et al. 1998). Moreover, the cold gas is inherently clumpy which effectively heats up the stellar `fluid' and populates the disk with an increased fraction of chaotic orbits, resulting in a decay of the bar strength (e.g., Shlosman \\& Noguchi 1993; Berentzen et al. 2007). The vertical buckling instability in the bar is progressively damped by the two-fluid gas-star interaction (e.g., Berentzen et al. 2007).  Formation and evolution of stellar bars is intricately related to the redistribution of angular momentum, $J$, in the disk-halo system. This means the presence of sources and sinks of angular momentum. Sellwood (1981) has shown that the former reside in the disk and the latter in the DM halo (see also Debattista \\& Sellwood 1998; Athanassoula 2002). Interactions of sources and sinks of $J$ are dominated by resonances (Lynden-Bell \\& Kalnajs 1972; Athanassoula 2002; Martinez-Valpuesta, Shlosman \\& Heller 2006; Paper~I). Details of these resonance interactions are not fully understood and quantified. This is especially true in the presence of the gas component, as the above cited works are based on a purely collisionless modeling.  In this paper we analyze the effect of the gaseous component on the evolution of galactic bars and take a numerical approach. In particular, we vary the gas content in the disk and advance models with various numerical resolution, but keep the stellar disk and DM halo paramaters unchanged, using the Standard Model of Paper~I. We are especially interested in how the gas presence and the associated numerical resolution affect the basic parameters of a stellar bar, as well as the angular momentum transfer process in a barred disk and in a disk-halo system. For this purpose, we perform a detailed comparison of models with gas to the standard collisionless model published in Paper~I. ",
        "conclusions": "\\label{s_disc}  \\begin{figure*}[!t] \\centering \\includegraphics[angle=0,scale=1.0]{fig9.ps} \\caption{Bar strengthening during $(a)$ dynamical and $(b)$ secular phases of its evolution as a function of the gas fraction, \\fg, for different spatial resolution in the gas, \\egrav, namely \\egrav$=0.016$ (crosses, solid lane), 0.05 ($\\times $, short dashed lane), and 0.1 (circles, long dashed lane). The filled square represents the SD model (Paper~I) with no gas. The model SD\\_G8S1 has been omitted (see the text). } \\label{figmtot1} \\end{figure*}  We have studied the effects of gas fraction, \\fg, and spatial resolution, \\egrav, on some aspects of bar evolution in galactic disks embedded in DM halos. Specifically, we aimed at understanding their effect on the basic parameters characterizing stellar bars, e.g., bar size, corotation radius, strength, etc. We have also followed the angular momentum redistribution in the disk-halo systems as a function of \\fg{} and \\egrav. The SD pure stellar model (Paper~I) has been used as a template --- the disk and halo parameters in current models have been fixed at those of the SD. As before, the bar evolution has been separated into two main phases --- the dynamical phase, which represents the bar instability itself de facto, and the secular phase. These phases are separated by the vertical buckling instability in the bar. We investigate a number of new correlations which are corollaries to evolution discussed in section~3. In the follow-up work (Villa-Vargas et al., in preparation), we analyze the angular momentum transfer and bar evolution as a function of disk gas fraction and by varying the basic parameters of the DM halo.  Overall, there appears to be a substantial difference in the evolution of the gas-poor and rich models, as has been shown in the previous section. This is discussed below in more detail. The boundary between the gas-poor and gas-rich models depends on the spatial resolution for the gas component --- it lies around 5\\%--7\\% for \\egrav=0.016, and shifts to $\\sim 10\\%-12\\%$ for \\egrav=0.1. The loss of spatial resolution, in a way, changes the nature of the gas component --- it becomes less dissipative.  We define the bar strengthening during its dynamical phase as $\\Delta $\\Atworb~$\\equiv $\\Atworb ($t=t_{\\rm peak}$)\\ --\\Atworb($t=0$), where $t_{\\rm peak}$ is the time when the bar has reached its maximal amplitude prior to the onset of buckling; the second term \\Atworb($t=0$)$=0$. Similarly, the bar growth during the secular phase is defined as $\\Delta $\\Atworb~$\\equiv $\\Atworb($t=270$) -- \\Atworb ($t=t_{\\rm min}$), where $t_{\\rm min}$ is the time of the minimal bar amplitude immediately following the buckling.  In the {\\it dynamical} stage of the bar evolution, the largest differences have been observed in the gas-rich models, where $\\Delta $\\Atworb{} drops dramatically with \\fg{} for \\egrav$=0.016$ and 0.05, while it exhibits no dependence on \\fg{} for \\egrav$=0.1$ (Figs.~1, 9a). In the gas-poor models, the maximal bar amplitude achieved in the dynamical phase is nearly independent of the gas fraction, nor does it depends on the spatial resolution used in the modeling of the gas component (Figs.~1, 9a). Fig.~9a supports the Berentzen et al. (1998, 2007) analysis that the increasing presence of the gas component makes the bar instability milder. In the {\\it secular} stage, this sharp decrease in $\\Delta $\\Atworb{} with \\fg{} persists, including the lowest resolution sequence as well. The bar growth is severely restrained in the gas-rich models in this phase. It either had vanished or even becomes negative. The underlying physical process has been quantified by Berentzen et al. (1998): the {\\it gas affects the orbital structure} of the collisionless components, which in turn modifies the bar properties. However detailed understanding of how this influences particular parameters of the bars requires more work.  We have attempted to clarify this dependence of the bar strength on \\fg{} and \\egrav. As a working hypothesis, we have tested whether the final mass of the CMC, \\mcmc, in each model can serve as an underlying hidden parameter replacing \\fg{} and \\egrav. The dependency of the final bar strength on the CMC mass in the pure stellar disks has been analyzed by Athanassoula, Lambert \\& Dehnen (2005). Extensive study of the bar evolution in the presence of an analytical CMC has been performed by Shen \\& Sellwood (2004). Bornaud et al. (2005) have attributed the bar weakening to both the CMC and the angular momentum transfer from the gas to the stellar bar, using analytical CMC and DM halo. However, the latter process has been found not important by Berentzen et al. (2007) following a detailed analysis of the angular momentum transfer in live potentials of the CMC and DM halo. Moreover, Berentzen et al. (1998, 2007) have used self-consistently growing CMCs in the gaseous/stellar disks to demonstrate that the secular bar growth is strongly affected. The set of numerical models analyzed here is in fact the most controlled experiment performed so far to test the influence of gas fraction and its resolution on the bar evolution.  \\begin{figure*}[!t] \\centering \\includegraphics[angle=0,scale=1.0]{fig10.ps} \\caption{$(a).$ The secular strengthening of the bar, $\\Delta$\\Atworb, after the first buckling and until the end of the simulations, as a function of the final central mass concentration CMC within $R=0.1$, \\mcmc{}, for different spatial resolution in the gas, abbreviated as in Fig.~9. The filled square represents the SD model (Paper~I) with no gas. The model SD\\_G8S1 has been omitted (see the text). $(b).$ Final \\mcmc{} normalized by the total mass of the disk at $t=0$ as a function of \\fg. } \\label{figmtot1} \\end{figure*}  \\begin{figure*}[!t] \\centering \\includegraphics[angle=0,scale=1.0]{fig11.ps} \\caption{First vertical buckling amplitude of the stellar bar, $A_{\\rm 1,z}$, as a function of the bar strengthening during dynamical and secular phases of its evolution, $\\Delta $\\Atworb, for different spatial resolution in the gas, abbreviated as in Fig.~9. The filled square represents the SD model with no gas. } \\label{figmtot1} \\end{figure*}  Fig.~10a shows the dependence of the bar strengthening after the first buckling, $\\Delta$\\Atworb{} on final \\mcmc{}, with all other parameters characterizing the stellar disk and the DM halo being fixed. One model, SD\\_G8S1, has been omitted, as its CMC secular evolution is unusually flat and falls out of the sequence with \\egrav$=0.016$. If \\mcmc{} is the single defining variable which controls $\\Delta$\\Atworb, all three \\egrav-sequences are expected to merge into a single sequence in Fig.~10a. What is actually observed is that the three sequences exhibit a very similar behavior --- each curve is flat for the gas-poor models, experiences an abrupt drop, and flattens out. However, the curves do not coincide completely as would be expected if \\mcmc{} were the {\\it unique} underlying parameter. In fact, in some cases the points corresponding to the same \\mcmc{} but different \\egrav{} have a substantial vertical dispersion in $\\Delta$\\Atworb{}. Hence, the issue remains inconclusive. In contrast, we show the final \\mcmc{} value normalized by the disk mass at $t=0$ as a function of \\fg{} (Fig.~10b). This fractional \\mcmc{} value is indeed unique for the three sequences, as all three curves have nearly merged.  We now turn to a plausible correlation between the buckling amplitude and the bar strengthening. The vertical buckling (e.g., Toomre 1966; Combes et al. 1990; Raha et al. 1991) is a recurrent instability (Martinez-Valpuesta et al. 2006) that was recently analyzed by Berentzen et al. (2007) in the presence of gas (see also Berentzen et al. 1998). The gas component, it has been concluded, leads to a milder instability. Here we have attempted to relate the buckling amplitude, $A_{\\rm 1,z}$, to the change in the bar amplitude, $\\Delta $\\Atworb, in the dynamical and secular phases of evolution (Fig.~11). In both phases we observe a clear correlation between \\Aone{} and $\\Delta $\\Atworb{}. Stronger bar instability leads to a stronger buckling, and indeed, increasing \\fg{} makes the bar and buckling instabilities milder (Figs.~11a; see also Fig.~9a). On the other hand, stronger buckling goes in tandem with the bar secular growth (Fig.~11b). This trend shows saturation for the strongest bars. Higher resolution sequences with \\egrav$= 0.016$ and 0.5 behave in a very similar fashion, while the lowest resolution sequence stands out of this correlation. We have also checked the value of the drop in \\Atworb{} during the buckling and find a clear match between the higher resolution sequences, while the lower resolution models behave differently.  The physical extent of the bar depends on its ability to capture additional orbits. While in principle this capture can proceed at all radii, the fertile region lies between the bar's end and its CR radius, where various families of orbits can be easily destabilized. Therefere, the  bar growth due to the orbit capture should go in tandem with the angular momentum {\\it influx} because the near-CR orbits will have a larger momentum-to-energy, J/E, ratio than the bar orbits. In Paper~I, we have shown, and this is confirmed here, that the influx of angular momentum across the CR and into the bar is able to maintain \\Jdin$\\sim $const. in time, as long the CR radius lies within the disk, but in the presence of the gas {\\it this statement is limited to the gas-poor models only}. This happens despite that the CR remains within the disk at all times and for all gas-rich models (Fig.~4). In other words, in gas-rich models the $J$ influx across the CR cannot compensate for its loss within the CR.  The amplitude \\Atworb{} is expected to be related to the ability of the bar to capture additional orbits in the bar CR region, leading to its geometrical growth. In both the dynamical and secular phases of the bar evolution we find that the final bar amplitude, \\Atworb, correlates with the final bar size\\footnote{The only exception appears to be the low resolution sequence in the dynamical stage whose points cluster in the same region (Fig.~12a)} (Fig.~12). So stronger bars appear to be longer as well. Taken at face value, this property of bar evolution appears to be supported by recent observations in that the $K$-band images of barred galaxies exhibit a correlation between the bar amplitude, measured by the $m=2$ Fourier component, and its size (Elmegreen et al. 2007). However, there are caveats.  \\begin{figure*}[!t] \\centering \\includegraphics[angle=0,scale=0.9]{fig12.ps} \\caption{The bar amplitude, \\Atworb, at its pre-buckling peak (left) and at the end of the simulation (right) as a function of the bar size (radius), taken at the same time, during dynamical and secular phases of the bar evolution.  Different spatial resolution in the gas has been abbreviated as in Fig.~9. The full square represents the SD model with no gas. We have subtracted the {\\it ansae} from the bar sizes given in Fig.~3 (more details in the text). } \\label{figmtot1} \\end{figure*}  \\begin{figure}[!t] \\centering \\includegraphics[angle=0,scale=1.0]{fig13.ps} \\caption{Bimodal distribution of the final bar sizes, \\Rbar($t=270$), as a function of the gas fraction, \\fg.  We have subtracted the {\\it ansae} from the bar sizes given in Fig.~3 (more details in the text). } \\label{figmtot1} \\end{figure}   We find that the bar sizes are substantially smaller in the gas-rich models compared to the gas-poor ones (Figs.~3, 13). While the median bar size at the end of the simulation in the latter models appears to be $\\sim 1.2$, it is around 0.5 in the former ones, i.e., 12~kpc vs 5~kpc, embedded in the identical disks and halos. This difference arises almost entirely during the secular phase of the bar evolution. More precisely, by the end of the dynamical stage, one can notice that the bar size slightly anti-correlates with \\fg{} for the gas-rich models and only for \\egrav$=0.016$ and 0.05 sequences. This trend is sufficiently weak and bars differ by not more than 10\\%--15\\% of their length. So essentially, the bar growth in the dynamical phase is independent of \\fg.  What is rather striking, is the abrupt change in the bar's growth `habit' in the secular phase --- bars grow in the gas-poor models and stagnate in the gas-rich ones. This leads to a {\\it bimodal} evolution of the final bar sizes with respect to the gas mass fraction (Fig.~13). We note that this diverging evolution should be observable only in the long-lived bars which naturally reside in the older disks. The relevant evolutionary timescale is that of a few Gyrs. The Milky Way disk is a good candidate, as it is probably about 10~Gyr old and was not affected by major mergers over this time period (e.g., Gilmore, Wyse \\& Norris 2002).  The anti-correlation between the final bar sizes and disk gas fractions, found here, complicates the simplistic picture of bar evolution. The reason for this is that our gas-rich models lead to more massive CMCs which increase bulge-to-disk mass ratios and, therefore, should be associated with earlier type disks. This means that smaller bars in our models lie in early-type disks (Fig.~13), while they are expected to be found in late-type disks, as noted by a number of surveys (e.g., Erwin 2005; see also Laine et al. 2002). The simplest resolution of this discrepancy can lie in the omission of star formation processes and especially the feedback from stellar evolution and from the central supermassive black holes in this work. This leads to a gross over-estimate of the amount of gas which reaches the central kpc and hence contributes to the growth of the CMC by dragging stars and even DM inwards. We note, however, that the purpose of this numerical exercise was to understand the effect of gas fraction and gas spatial resolution on the disk-bar-halo interactions in the system. Therefore, we have simplified the long list of processes known to affect the galaxy evolution at large.  In summary, we find that the spatial resolution in the gas component becomes increasingly important for the bar evolution in the gas-rich disks. This is true for the dynamical but especially for the secular phase of evolution. In most cases, model sequences with \\egrav$= 0.016$ and 0.05 show a similar behavior, while differing substantially from the sequence with \\egrav$=0.1$.  A bimodal behavior has been found for models based on their gas fractions. The border line between the gas-poor and gas-rich systems appears to lie around $5\\%-7\\%$ for higher resolution models. It shifts to $\\sim 10\\%-12\\%$ for the lowest resolution models. The switch from a gas-poor to a gas-rich behavior appears to be sufficiently abrupt. It is clearly visible in all basic characteristics of bar evolution, such as the bar strength, the CMC mass, the bar buckling amplitude, the bar size, etc. The largest differences in the evolution have been found in the secular phase.  We find that the presence of the gas component severely limits the bar growth and affects its pattern speed evolution. While pure stellar models (Paper~I) exhibit a rapid slowdown of the bar tumbling, as known for a long time (e.g., Debattista \\& Sellwood 1998; Athanassoula 2003), the addition of a substantial amount of gas reverses this trend completely. Furthermore, the CR-to-disk size ratio, \\Rcr/\\Rd, was determined to be an important dynamic discriminator between various phases of barred disk evolution. Here we find that the gas-rich models are characterized by \\Rcr/\\Rd~$ < 1$ and by $\\Omega_{\\rm b}\\sim$~const. In these models, $\\Omega_{\\rm b}$ can even slightly increase with time. In addition, the gas-rich models maintain \\Rcr/\\Rbar~$ < 2$ which is much more in agreement with observations of stellar bars being fast rotators.  The angular momentum evolution displays the same degree of bimodality with the gas fraction. For low \\fg, the total disk $J$ drops steeply and monotonically with time after the buckling, while it decreases weakly for the gas-rich models. The reason for this behavior is of course that the bar amplitude is substantially lower in the gas-rich models. Our attempt to explain this difference between the gas-poor and rich models in terms of the more massive CMCs in the latter ones has been rather inconclusive. We shall return to this issue in the forthcoming work.  Next, we have confirmed our previous claim (Paper~I) that the angular momentum, \\Jdin, within the CR radius is maintained at a constant level due to the influx of angular momentum across the CR as it expands. We have extended this statement to the gas-poor models. For the gas-rich models, \\Jdin{} drops abruptly to a lower level and stays constant thereafter.  A number of corollaries follow from the above results. We find that the bar strength inversely correlates with the gas fraction, both in the dynamical and secular phases of bar evolution. The only exception seems to be the lowest resolution sequence in the dynamical phase which fails to capture this trend. We also find that the buckling amplitude becomes larger for stronger bars prior to the onset of buckling. On the other hand, the secular growth is most prominent in bars which show a large buckling amplitude. Finally, we show that stronger bars are also the longest ones throughout both evolutionary phases, and that bar sizes anti-correlate with the gas fraction."
    },
    "1004/1004.4916_arXiv.txt": {
        "abstract": "Although they are but a small fraction of the mass ejected in core-collapse supernovae, neutrino-driven winds (NDWs) from nascent proto-neutron stars (PNSs) have the potential to contribute significantly to supernova nucleosynthesis.  In previous works, the NDW has been implicated as a possible source of r-process and light p-process isotopes.  In this paper we present time-dependent hydrodynamic calculations of nucleosynthesis in the NDW which include accurate weak interaction physics coupled to a full nuclear reaction network.  Using two published models of PNS neutrino luminosities, we predict the contribution of the NDW to the integrated nucleosynthetic yield of the entire supernova. For the neutrino luminosity histories considered, no true r-process occurs in the most basic scenario.  The wind driven from an older $1.4 M_\\odot$ model for a PNS is moderately neutron-rich at late times however, and produces $^{87}$Rb, $^{88}$Sr, $^{89}$Y, and $^{90}$Zr in near solar proportions relative to oxygen.  The wind from a more recently studied $1.27 M_\\odot$ PNS is proton-rich throughout its entire evolution and does not contribute significantly to the abundance of any element.  It thus seems very unlikely that the simplest model of the NDW can produce the r-process.  At most, it contributes to the production of the N = 50 closed shell elements and some light p-nuclei. In doing so, it may have left a distinctive signature on the abundances in metal poor stars, but the results are sensitive to both uncertain models for the explosion and the masses of the neutron stars involved. ",
        "introduction": "The site where r-process nuclei above A=90 have been synthesized remains a major unsolved problem in nucleosynthesis theory \\citep[e.g.,][]{Arnould07}.  Historically, many possibilities have been proposed \\citep[see][]{Meyer94}, but today, there are two principal contenders - neutron star mergers \\citep{Lattimer77,Freiburghaus99} and the NDW \\citep{Woosley94,Qian96,Hoffman97,Otsuki00,Thompson01}.  Observations of ultra-metal-poor stars suggest that many r-process isotopes were already quite abundant at early times in the galaxy \\citep{Cowan95,Sneden96,Frebel07}, suggesting both a primary origin for the r-process and an association with massive stars. NDWs would have accompanied the first supernovae that made neutron stars and, depending upon what is assumed about their birth rate and orbital parameters, the first merging neutron stars could also have occurred quite early.  Both the merging neutron star model and the NDW have problems though. In the simplest version of galactic chemical evolution, merging neutron stars might be capable of providing the necessary integrated yield of the r-process in the sun, but they make it too rarely in large doses and possibly too late to be consistent with observations \\citep{Argast04}.  On the other hand, making the r-process in NDWs requires higher entropies, shorter time-scales, or lower electron mole numbers, $Y_e$, than have been demonstrated in any realistic, modern model for a supernova explosion \\citep[though see][]{Burrows06}.  Previous papers and models for nucleosynthesis in the NDW have focused on the production of nuclei heavier than iron using either greatly simplified dynamics \\citep{Beun08,Farouqi09} or nuclear physics \\citep{Qian96,Otsuki00,Arcones07,Fischer09,Huedepohl09}. Post processing nuclear network calculations have been performed using thermal histories from accurate models of the dynamics \\citep{Hoffman97,Thompson01,Wanajo06}, but the calculations sampled only a limited set of trajectories in the ejecta. No one has yet combined the complete synthesis of a realistic NDW with that of the rest of the supernova.  To address this situation, and to develop a framework for testing the nucleosynthesis of future explosion models, we have calculated nucleosynthesis using neutrino luminosity histories taken from two PNS calculations found in the literature \\citep{Woosley94,Huedepohl09}. This was done using a modified version of the implicit one-dimensional hydrodynamics code Kepler, which includes an adaptive nuclear network of arbitrary size.  This network allows for the production of both r-process nuclei during neutron-rich phases of the wind and production of light p-elements during proton-rich phases.  Since the results of wind nucleosynthesis depend sensitively on the neutrino luminosities and interaction rates \\citep{Qian96,Horowitz02}, we have included accurate neutrino interaction rates that contain both general relativistic and weak magnetism corrections.  The synthesis of all nuclei from carbon through lead is integrated over the history of the NDW and combined with the yield from the rest of the supernova, and the result is compared with a solar distribution. If a nucleus produced in the NDW is greatly overproduced relative to the yields of abundant elements in the rest of the supernova, there is a problem. If it is greatly underproduced, its synthesis in the NDW is unimportant.  If it is co-produced, the NDW may be responsible for the galactic inventory of this element.  An important outcome of this study are the yields expected from a ``plain vanilla'' model for the NDW. Are there any elements that are robustly produced and thus might be used as diagnostics of the wind in an early generation of stars?  In \\S\\ref{wind_physics}, we discuss the general physics of neutrino driven winds and analytically delineate the regions in neutrino temperature space were different modes of nucleosynthesis occur.  We then discuss our numerical model in \\S\\ref{computational_method}.  In \\S\\ref{results}, the results of the time dependent models are presented.  We conclude with a discussion of how the NDW might affect galactic chemical evolution and consider if this allows the strontium abundance in low metallicity halo stars to be used as a tracer of supernova fallback at low metallicity.  Additionally, we investigate if observed abundances in SN 1987A can put constraints on late time neutrino luminosities from PNSs.  Finally, we discuss some possible modifications of the basic model that might improve the r-process production. These ideas will be explored more thoroughly in a subsequent paper. ",
        "conclusions": "We have performed calculations of the dynamics and nucleosynthesis in time dependent neutrino driven winds.  This was done for two sets of neutrino spectra calculated in one-dimensional supernova models taken from the literature.  The nucleosynthesis in these models was compared with supernova yields to determine if these models were consistent with observations.  Additionally, we compared the results of these numerical models to analytic models of the neutrino driven wind and found good agreement.  Similar to most of the work on the NDW after \\cite{Woosley94}, we find that it is unlikely that the r-process occurs in the neutrino driven wind unless there is something that causes significant deviation from a purely neutrino driven wind.  Additionally, in the simplest case, there is little production of p-process elements at early times in the wind.  In our calculation that used spectra from a more massive neutron star, the wind only produces the N=50 closed shell elements $^{87}$Rb, $^{88}$Sr, $^{89}$Y, and $^{90}$Zr.  This result is sensitive to small changes in the neutrino interaction rates (i.e. the inclusion of weak magnetism) and changes to the neutrino temperature of order 10\\%.  Comparing our models with the over abundance of strontium seen in SN 1987A suggests that the difference between the electron and anti-electron neutrino temperatures in the model of \\cite{Woosley94} may have been to large.  We also find that the effect of a wind termination shock on the wind nucleosynthesis is small.  Using neutrino spectra from an $8.8 M_\\odot$ supernova that drives a wind which is proton rich throughout its duration \\citep{Huedepohl09}, we find that no significant $\\nu$p-process occurs and the wind does not contribute to the yields of the supernova. The neutrino spectra from this model are probably more accurate than the spectra from the model of \\cite{Woosley94}.  We also investigated the effect of an outer boundary pressure which resulted in a wind termination shock.  This had a negligible effect on the nucleosynthesis.  However, one also expects that the nucleosynthesis in the NDW will vary considerably from event to event, especially with the mass and possibly the rotation rate of the PNS. The winds from more massive PNS have greater entropy and might, in general, be expected to produce heavier elements and more of them. The neutrino spectral histories of PNS as a function of mass have yet to be determined over a wide range of parameter space.  Currently, the neutrino luminosities and temperatures are the largest uncertainties in models of the NDW."
    },
    "1004/1004.5121_arXiv.txt": {
        "abstract": "We present a detailed analysis of the dayside atmosphere of the hot-Neptune GJ~436b, based on recent {\\it Spitzer} observations. We report statistical constraints on the thermal and chemical properties of the planetary atmosphere, study correlations between the various molecular species, and discuss scenarios of equilibrium and non-equilibrium chemistry in GJ~436b. We model the atmosphere with a one-dimensional line-by-line radiative transfer code with parameterized molecular abundances and temperature structure. We explore the model parameter space with $10^6$ models, using a Markov chain Monte Carlo scheme. Our results encompass previous findings, indicating a paucity of methane, an overabundance of CO and CO$_2$, and a slight underabundance of H$_2$O, as compared to equilibrium chemistry with solar metallicity. The concentrations of the species are highly correlated. Our best-fit, and most plausible, constraints require a CH$_4$ mixing ratio of $10^{-7} to 10^{-6}$, with CO $\\geq 10^{-3}$, CO$_2$ $\\sim 10^{-6} to 10^{-4}$, and H$_2$O $\\leq 10^{-4}$; higher CH$_4$ would require much higher CO and CO$_2$. Based on calculations of equilibrium and non-equilibrium chemistry, we find that the observed abundances can potentially be explained by a combination of high metallicity ($\\sim 10~\\times$ solar) and vertical mixing with $K_{zz} \\sim 10^6 - 10^7$ cm$^2$/s. The inferred metallicity is enhanced over that of the host star which is known to be consistent with solar metallicity. Our constraints rule out a dayside thermal inversion in GJ~436b. We emphasize that the constraints reported in this work depend crucially on the observations in the two {\\it Spitzer} channels at 3.6 $\\micron$ and 4.5 $\\micron$. Future observations with warm {\\it Spitzer} and with the {\\it James Webb Space Telescope} will be extremely important to improve upon the present constraints on the abundances of carbon species in the dayside atmosphere of GJ~436b. ",
        "introduction": "\\label{sec:intro}  The last decade in exoplanetary science has demonstrated our capability in detecting and characterizing atmospheres of transiting extrasolar giant planets. Several observations have been reported using the {\\it Hubble Space Telescope (HST)}, the {\\it Spitzer Space Telescope (Spitzer)}, and from ground. Beginning with the first detection of sodium in the atmosphere of HD~209458b in transit (Charbonneau et al. 2002), and the first detections of dayside thermal emission from hot Jupiters TrES-1 and HD~209458b (Charbonneau et al. 2005; Deming et al. 2005), atmospheric observations of giant exoplanets today are a norm, albeit still very challenging (e.g. Knutson et al. 2008; Charbonneau et al. 2008; Swain et al. 2008; Grillmair et al. 2008; Desert et al. 2009; Swain et al. 2009a). The intensity of observational efforts have been matched with equally challenging accomplishments in theoretical modeling of exoplanet atmospheres, and data interpretation (Seager \\& Sasselov, 2000; Seager et al. 2005; Burrows et al. 2006 \\& 2008; Fortney et al. 2006; Barman et al. 2005; Tinetti et al. 2007; Showman et al. 2009; Madhusudhan \\& Seager, 2009). Several inferences of gaseous H$_2$O, CH$_4$, CO and CO$_2$, and thermal inversions have subsequently been made in hot Jupiter atmospheres (Burrows et al. 2007; Tinetti et al. 2007; Barman, 2007; Grillmair et al. 2008; Swain et al. 2008; Madhusudhan \\& Seager, 2009; Swain et al. 2009a \\& 2009b).  A new era in exoplanetary science has now dawned. Latest observations are leading to discovery and characterization of transiting exoplanets much less massive than the archetypal hot Jupiters, namely extrasolar Neptunes and super-Earths. Several low-mass transiting planets are presently known, e.g. hot Neptunes GJ~436b (Butler et al. 2004; Maness et al. 2007), HAT-P-11b (Bakos et al. 2010), HAT-P-26b (Hartman et al. 2010), and Kepler-4b (Borucki et al. 2010), and super-Earths CoRoT-7b (Leger et al. 2009) and GJ~1214b (Charbonneau et al.  2009). These planets have opened a new regime in atmospheric modeling and data interpretation. However, observations of thermal emission have been reported for only one of these planets so far, GJ~436b (Deming et al. 2007, Stevenson et al. 2010).  The first transiting hot Neptune known, GJ~436b, with a mass of 22.6~M$_\\oplus$ and radius of 4.2~R$_\\oplus$, orbits an M Dwarf at an orbital separation of 0.03 AU (Butler et al. 2004; Gillon et al. 2007; Maness et al. 2007). The host star has an effective temperature of about 3500 K and a metallicity consistent with solar (Torres et al. 2008). The average density of the planet is 1.5 g/cc (Torres et al. 2008), i.e., similar to the bulk density of Neptune, an ice giant. At an equilibrium temperature of $\\sim$ 700 K, assuming zero albedo and efficient energy circulation, the density originally hinted of a hot-ice interior (Gillon et al. 2007). More detailed studies of the possible bulk composition of GJ~436b indicate that an additional layer of H/He atmosphere would be needed to account for the observed radius (Figueira et al.~2009; Rogers \\& Seager.~2009).  The atmosphere of GJ~436b has been a subject of substantial interest in the recent past. Deming et al. (2007) and Demory et al. (2007) reported independent detections of thermal emission from the dayside of GJ~436b in the {\\it Spitzer} 8 $\\micron$ IRAC channel. Although no meaningful inferences about the molecular compositions can be drawn from a single data-point, model fits to the 8 $\\micron$ flux contrast have favored the interpretation of inefficient day-night energy redistribution for GJ~436b (Demory et al. 2007; Spiegel et al. 2010, unless there are additional unknown optical absorbers in the atmosphere). The models used in Demory et al. 2007 and Spiegel et al. 2010 had assumed equilibrium chemistry. Given the low temperatures of GJ~436b compared to hot Jupiters, equilibrium chemistry suggests that the planet's atmosphere must be abundant in methane and water vapor, and be scarce in carbon monoxide (e.g Burrows and Sharp, 1999).  Recent observations have suggested distinct departures from predictions of equilibrium chemistry models. Pont et al. (2009) reported a transmission spectrum of GJ~436b obtained in the 1.1-1.9 $\\micron$ bandpass using the {\\it HST} NICMOS instrument, but found no significant feature in the 1.4 $\\micron$ water band; the spectrum was flat at the 2-$\\sigma$ uncertainties. More recently, Stevenson et al. (2010) reported planet-star flux contrasts of the dayside atmosphere of GJ~436b in six channels of {\\it Spitzer} broadband photometry, and inferred a deficiency of methane in the atmosphere of GJ~436b, using models based on Madhusudhan \\& Seager (2009) and the present work. The high planet-star flux contrast observed in the 3.6 $\\micron$ IRAC channel was central to the low methane requirement. Even though our models do not assume equilibrium chemistry, the inferred methane mixing ratio of $\\sim 10^{-7}$ in a hydrogen dominated atmosphere at $\\sim 700 K$ signals a surprising new regime in atmospheric chemistry of extrasolar planets.  In this work, we report detailed statistical constraints on the atmospheric properties of  GJ~436b, and explore channels of equilibrium and non-equilibrium chemistry that might explain the observed chemical abundances. We first estimate the atmospheric chemical composition and temperature structure at different levels of fit to the data, using a 1-D line-by-line radiative transfer model for exoplanet atmospheres (Madhusudhan \\& Seager, 2009). Our constraints result from exploring the model parameter space with $\\sim 10^6$ models, optimized using a Markov chain Monte Carlo optimization scheme. We then use detailed calculations of equilibrium and non-equilibrium chemistry, and with different metallicities, to explain the observed constraints on the various molecular species.  Our results indicate that a high metallicity and non-equilibrium chemistry are required to explain the molecular abundances constrained by the observations. The requirement of non-equilibrium chemistry is consistent with our findings in Stevenson et al. (2010). The observations require the presence of CO and CO$_2$ and a substantial depletion of CH$_4$ in the dayside atmosphere of GJ~436b. In this work, we find that the constraints on the CO and CO$_2$ mixing ratios can be explained by vertical mixing in the atmosphere (with $K_{zz} = 10^6 - 10^7$ cm$^2$/s) and a high metallicity ($10 \\times $ solar). And, we suggest that the low CH$_4$ required can potentially be explained by non-equilibrium chemistry. We also find that the observations require inefficient day-night energy redistribution in GJ~436b, which can be confirmed by future observations of phase curves.  In what follows, we first give a brief overview of model independent interpretation of {\\it Spitzer} observations in section~\\ref{sec:data_interpretation}. We then describe our atmosphere model, the parameter space optimization scheme, and the models for equilibrium and non-equilibrium chemistry we use in this work, in section~\\ref{sec:model}. We present the results of our analysis in section~\\ref{sec:results}. Finally, in section~\\ref{sec:discussion}, we present a summary of our work and discuss consequences for future observations and theoretical models.  \\begin{figure}% \\centering \\includegraphics[width = 0.5 \\textwidth]{fig1.eps} \\caption{ Molecular absorption features in {\\it Spitzer} photometric bandpasses. The red dotted lines at the top show the six {\\it Spitzer} bandpasses. The black lines show the extent of absorption features due to the corresponding molecules. The gray curve shows a hypothetical model spectrum of GJ~436b based on equilibrium chemistry, and the black filled circles show the corresponding integrated points in the {\\it Spitzer} channels. The red filled circles with error bars show the observations of GJ~436b reported by Stevenson et al. (2010).} \\label{fig:spitzer} \\end{figure} ",
        "conclusions": "\\label{sec:discussion}  We presented a detailed analysis of the dayside atmosphere of GJ~436b. Our results show that a high metallicity along with non-equilibrium chemistry are required to explain the observations. We also studied the correlations between the various molecular species, and reported detailed constraints on the metallicity, chemical processes, and day-night energy circulation. Although our results come from observations in six channels of {\\it Spitzer} photometry, some channels are more important than others. Here, we discuss the relative importance of the different {\\it Sptizer} channels to our conclusions. We also discuss some potential alternate interpretations of the data, followed by a summary of our results.  \\subsection{Sensitivity of Results to {\\it Spitzer} Observations}  The constraints reported in this work depend critically on the two {\\it Spitzer} IRAC observations at 3.6 $\\micron$ and 4.5 $\\micron$. The high planet-star flux contrast in the 3.6 $\\micron$ channel is responsible for the constraints of low methane abundance and low energy circulation.  The low flux contrast in the 4.5 $\\micron$ channel is responsible for the requirement of high CO and/or CO$_2$ in GJ~436b. While the observation in the 3.6 $\\micron$ channel was reported to be of the higher S/N of all channels, the 4.5 $\\micron$ channel was a non-detection (Stevenson et al. 2010). Nevertheless, future observations in both these channels would be extremely important to confirm that these fluxes actually represent the steady state atmosphere in GJ~436b. The observations in the remaining four channels (5.8 $\\micron$  - 24 $\\micron$) are much less constraining, although still very useful. For instance, the moderate flux observed in the 8 $\\micron$  channel, where methane absorbs strongly, is important to the conclusion that the very high flux in the 3.6 $\\micron$ channel could not have been due to a thermal inversion causing methane emission.  Our constraints on the molecular abundances are fairly robust with respect to the observational uncertainties. The high flux in the 3.6 $\\micron$ channel cannot be explained by equilibrium chemistry. A model with equilibrium chemistry and solar or 30 $\\times$ solar metallicities predicts planet-star flux contrasts that are lower than the observed value by over 4 -$\\sigma$ (also in agreement with models of Demory et al. 2007; Spiegel et al. 2010). On the other hand, the contrasts predicted in the 4.5 $\\micron$ channel based on equilibrium chemistry alone would be higher than the observed non-detection by over 3-$\\sigma$ (also see Spiegel et al. 20010). We have been conservative in our analysis by allowing our best fits to the 4.5 $\\micron$ point to lie within the 3-$\\sigma$ $\\pm$ 1-$\\sigma$ upper-limits. Had we considered this point to be a strict non-detection at 1-$\\sigma$, our results would predict even higher CO and/or CO$_2$. Finally, the large uncertainties in the observed fluxes in the 16 $\\micron$ and 24 $\\micron$ channels provide only fiducial constraints on the temperature structure and the H$_2$O and CO$_2$ abundances.  \\subsection{Alternate Interpretations}  The high planet-star flux contrast observed in the 3.6 $\\micron$ IRAC channel is central to most of the constraints reported in this work. Our inferences could partly be restricted by the specific choices that are inherent to our models. For example, our models do not include clouds or hazes. Although scattering from hazes has been suggested to be potentially relevant in the optical and near-IR (e.g. Sing et al. 2009), a high contribution at the longer wavelengths of the 3.6 $\\micron$ channel is unlikely. We have also assumed the planet atmosphere to be in local thermodynamic equilibrium (LTE). Swain et al. (2010) reported a potential signature due to non-LTE methane fluorescent emission in HD~189733b observed in the range of the 3.6 $\\micron$ channel (but c.f. Mandell et al. 2010, who did not find such a feature in follow-up observations). The dayside atmosphere of GJ~436b might also be variable, as has been reported previously for hot Jupiters (Grillmair et al. 2008; Madhusudhan \\& Seager, 2009). However, for variability to explain the observed flux in the 3.6 $\\micron$ channel, the temperature at the $\\sim$ 1 bar level in the atmosphere of GJ~436b has to exhibit fluctuations greater than 400 K, between subsequent observations of Stevenson et al. (2010). These and other alternate explanations are worth exploring with more data at different epochs.  \\subsection{Summary}  We have presented constraints on the chemical composition and temperature structure of the dayside atmosphere of hot Neptune GJ~436b, based on recent {\\it Spitzer} observations. One of our key findings is the strict upper limit on the mixing ratio of methane. We find that models fitting the observations require a methane mixing ratio below $10^{-6}$. Slightly higher methane mixing ratios require CO$_2$ $\\sim 10^{-3} - 10^{-2}$, which is implausible in the hydrogen rich atmosphere with the temperature structure of GJ~436b. The abundances of all the molecules are highly correlated. Applying nominal conditions of physical plausibility, we find the constraints on the molecular mixing ratios to be CH$_4$ $ \\sim 10^{-7} -  10^{-6}$, CO $\\geq 10^{-3}$, CO$_2$ $\\sim 10^{-6} - 10^{-4}$, and H$_2$O $\\leq 10^{-4}$. These constraints on the molecular abundances cannot all be explained based on equilibrium chemistry, for any metallicity, as reported in Stevenson et al. (2010). At the temperatures of GJ~436b, equilibrium chemistry with solar abundances predicts CH$_4$, CO and CO$_2$ mixing ratios to be $\\sim 5 \\times 10^{-4}$, $\\sim 10^{-5}$ and $\\sim 10^{-7}$, respectively, contrary to the observed abundances.  The observed constraints on the molecular abundances can be explained by a combination of high metallicity and non-equilibrium thermochemistry. A high metallicity is required for a high CO$_2$ abundance. Vertical mixing along with a high metallicity is required to dredge up the high CO abundance from the lower layers of the atmosphere to observable layers. Finally, vertical mixing and photochemistry can cause substantial depletion of CH$_4$, as reported by Zahnle et al. (2009b). At the temperatures (T $>$ 1100 K) we obtain for the lower atmosphere of GJ~436b, the results of Zahnle et al. (2009b) indicate the depletion of equilibrium CH$_4$ via oxidation to CO, caused by overabundance of the H radical. Our joint analysis of the parameters of non-equilibrium chemistry required to explain the abundances of all the species, suggests that the dayside atmosphere of GJ~436b has a high metallicity of $\\sim 10 \\times$ solar and a diffusion coefficient of $K_{zz} = 10^6 - 10^7$ cm$^2$/s. The metallicity is substantially enhanced over that of the host star which is consistent with solar metallicity (Torres et al. 2008).  Our results also constrain the dayside temperature structure and the day-night energy redistribution in the atmosphere of GJ~436b. A temperature inversion is ruled out by the current observations; although, small inversions which are not observable at the resolution of the current photometric data cannot be conclusively ruled out. The observations also suggest inefficient day-night energy redistribution ($f_r$) in GJ~436b, requiring $(1-A_B)(1-f_r) = \\eta \\geq 0.7$, at the 1-$\\sigma$ fit. Thus, the maximum $f_r$ allowed by the data at the 1-$\\sigma$ fits is 0.3 for zero bond albedo ($A_B$), and 0.23 for $A_B$ = 0.1. Future observations of thermal phase curves in the available warm {\\it Spitzer} channels will be instrumental in validating the low redistribution requirement. A finding on the contrary might indicate a substantial interior energy source. We emphasize that the constraints reported in this work depend primarily on the two {\\it Spitzer} channels (3.6 $\\micron$ and 4.5 $\\micron$). Thus, future observations in these channels will be extremely important in confirming the present and previous results on the atmosphere of GJ~436b.  The atmosphere of the hot-Neptune GJ~436b presents new challenges and opportunities for detailed modeling of exoplanet atmospheres. As low-mass transiting planets continue to be discovered, unexpected findings are likely to continue. The next generation of models and observations will help unravel those mysteries, and help put our solar system in perspective."
    },
    "1004/1004.5358.txt": {
        "abstract": "We present near-infrared (NIR) broadband (0.80--2.42 $\\mu$m) spectroscopy of two low mass X-ray binaries: V404 Cyg and Cen X-4. One important parameter required in the determination of the mass of the compact objects in these systems is the binary inclination. We can determine the inclination by modeling the ellipsoidal modulations of the Roche-lobe filling donor star, but the contamination of the donor star light from other components of the binary, particularly the accretion disk, must be taken into account. To this end, we determined the donor star contribution to the infrared flux by comparing the spectra of V404 Cyg and Cen X-4 to those of various field K-stars of known spectral type. For V404 Cyg, we determined that the donor star has a spectral type of K3 III. We determined the fractional donor contribution to the NIR flux  in the H- and K-bands as $0.98 \\pm .05$ and $0.97 \\pm .09$, respectively. We remodeled the H-band light curve from \\citet{sanwal1996} after correcting for the donor star contribution to obtain a new value for the binary inclination. From this, we determined the mass of the black hole in V404 Cyg to be $M_{BH}= 9.0^{+.2}_{-.6}M_{\\odot}$. We performed the same spectral analysis for Cen X-4 and found the spectral type of the donor star to be in the range K5 -- M1V. The donor star contribution in Cen X-4 is $0.94\\pm.14$ in the H-band while in the K-band, the accretion disk can contribute up to 10\\% of the infrared flux. We remodeled the H-band light curve from \\citet{shahbaz1993}, again correcting for the fractional contribution of the donor star to obtain the inclination. From this, we determined the mass of the neutron star as $M_{NS}= 1.5^{+.1}_{-.4}M_{\\odot}$. However, the masses obtained for both systems should be viewed with some caution since contemporaneous light curve and spectral data are required to obtain definitive masses. ",
        "introduction": "V404 Cyg and Cen X-4 belong to the class of low-mass X-ray binaries (LMXBs) in which mass transfer from a late type star to a black hole (BH) or a neutron star (NS) takes place via an accretion disk. By determining the mass of the compact object, we can constrain models for the formation and evolution of black holes and neutron stars in close binary systems \\citep{brown98, fryer2001, nelemans2001}. Of the 16 BH systems with BH masses determined to accuracies of 5--30\\%,  most of the masses range from 4 -- 16 M$_{\\odot}$ with only one system, J0422+32, with a BH mass $<$5 M$_\\odot$ \\citep{charlescoe06, casares2006, gelino2003}.  This result is not in agreement with the predictions of BH mass distributions from theoretical modeling which predict a continuous decreasing mass distribution from neutron star masses up to 10 -- 20$M_{\\odot}$ \\citep{fryer2001, belczynski08}. However, there currently are not enough systems with precise mass determinations to allow us to distinguish between observational biases and/or theoretical limitations as the source of the disparity.  To increase the number of LMXBs with precise compact object mass determinations, we have undertaken NIR spectroscopy of several systems to determine the relative contributions of the donor star and the accretion disk to the NIR light curves, from which accurate binary inclinations and, in turn, stellar masses are derived.  In this paper, we present results for two LMXBs, the BH binary V404 Cyg and the NS binary Cen X-4.  V404 Cyg went into outburst on May 22, 1989 and was subsequently discovered with the all-sky monitor aboard the Ginga satellite. A high X-ray luminosity of $\\approx 10^{38}$ ergs/sec suggested that the system was a close binary system with a black hole or a neutron star as the accretor \\citep{marsden89}. During the outburst, the visual magnitude of V404 Cyg increased from V=18.3 to V=11.6. One year after its discovery, it had fallen back to V=18.0 \\citep{wagner92}. Since then, V404 Cyg has been studied in quiescence at optical and infrared wavelengths, establishing most of the system parameters. \\citet{casares92} established the mass function, the minimum mass of the compact object, as $f(M) = 6.26 \\pm 0.31 M_{\\odot}$ which was later refined by \\citet{casares94} to be $f(M) = 6.08 \\pm 0.06 M_{\\odot}$. They determined the donor star to be of spectral type K0IV and obtained the binary mass ratio from the rotational broadening of the absorption lines in the donor star spectrum.  The ellipsoidal modulation of the donor star has been modeled in the NIR to determine the binary inclination and the black hole mass: in the K-band by \\citet{shahbaz1994} and in H by \\citet{sanwal1996}.  The resulting black hole mass determinations were 12 M$_{\\odot}$ and $<$12.5 M$_{\\odot}$, respectively, where one of the key differences between their results was that \\citet{shahbaz1994}  assumed that the donor star was the only source of flux in the NIR, while \\citet{sanwal1996} found evidence for accretion disk contamination.  \\citet{shahbaz1996} tested this assumption by performing NIR spectroscopy of V404 Cyg.  By comparing the NIR spectrum to scaled template spectra from spectral type standard stars, they determined that the donor star contributes 100$\\pm$11 \\% of the K-band flux.  However, for this analysis they relied on fits to the $^{12}$CO bands in K, which in other LMXBs (and cataclysmic variables) have been found to be affected by metallicity variations in the carbon abundance \\citep{froning2007}.  As a result, a broader spectral analysis that examines multiple atomic and molecular species is required to firmly tie down the donor star contribution in the NIR.  Cen X-4 was discovered on May 23, 1969 by the ${\\it Vela}$ ${\\it 5B}$ satellite \\citep{Conner69} when it went into X-ray outburst. During its second outburst in 1979, it increased in visual magnitude by $\\sim$6 magnitudes from its pre-outburst phase \\citep{canizares80}. \\citet{matsuoka80} identified the mass accreting star as a neutron star due to the fact that it displayed a Type I burst. \\citet{shahbaz1993} modeled the H-band light curve to obtain a binary inclination of  31$^\\circ$--54$^\\circ$ and the mass of the neutron star as 0.5--2.1M$_{\\odot}$. In quiescence, \\citet{torres2002} performed high-resolution optical spectroscopy on Cen X-4 and determined the mass of the compact object as $0.49M_{\\odot} < M < 2.49 M_{\\odot}$. They determined the mass function to be $f(M) = 0.220 \\pm 0.005M_{\\odot}$, which was later refined by \\citet{davanzo05} to $ f(M) = 0.201 \\pm 0.004M_{\\odot}$. The system parameters were again refined by \\citet{casares2007} and much tighter constraints were imposed on the radial velocity semi-amplitude, $K_2=144.6\\pm.3$ km s$^{-1}$, the orbital period, $P=.6290522\\pm.0000004$ d and the binary mass ratio, $q=.20\\pm.03$. The mass function adopted for this paper is calculated from the $P$ and $K_2$ values of \\citet{casares2007}, which is $ f(M) = 0.197 \\pm 0.001M_{\\odot}$.  However, as with V404 Cyg, in all the above cases, it was assumed that the donor star was the only NIR flux source, an assumption that has not been tested in Cen X-4.  In this paper, we present broadband (0.8--2.42 $\\mu$m) NIR spectroscopy of V404 Cyg and Cen X-4 obtained with SpeX on the IRTF.  By comparing the spectra to field stars of known spectral type,  we determine the donor star contribution to the infrared flux for both systems. We use these results to remodel previously obtained light curve data to determine the binary inclinations and to derive precise compact star masses for both targets. ",
        "conclusions": "We have obtained broadband near infrared spectroscopy of V404 Cyg and Cen X-4 in order to directly evaluate the relative contributions in the infrared light due to the donor star and the accretion disk and determine precise compact star masses. We determined the spectral type of the donor star in V404 Cyg to be a K3 III. We performed dilution analysis based on individual prominent metal lines and groups of lines and determined the donor star fraction to be $0.98 \\pm .05$ in the H-band and $0.97\\pm.09$ in the K-band. With previous light curve data in the H-band, we obtained the inclination of the binary as $i= (67^{+3}_{-1})^\\circ$ and a mass for the BH of $9.0^{+.2}_{-.6}$ M$_\\odot$ . In Cen X-4 we performed a similar analysis but restricted our fits to the strongest absorption lines due to the lower S/N of the spectrum. We determined the spectral type of the donor star to lie within the range K5 -- M1 V but could not constrain it further. The dilution fraction was determined to be $0.94\\pm.14$ in the H-band and  $1.09\\pm.20$ in the K-band. From past H-band data we modeled the light curve with a donor star plus a constant flux from the accretion disk and determined the mass of the neutron star in Cen X-4 as $1.5^{+.1}_{-.4} M_{\\odot}$ which agrees well with past estimates. These mass measurements should be viewed with some caution however, since contemporaneous light curve and spectral data are required to obtain definitive masses.    %% In this section, we use  the \\subsection command to set off %% a subsection.  \\footnote is used to insert a footnote to the text.  %% Observe the use of the LaTeX \\label %% command after the \\subsection to give a symbolic KEY to the %% subsection for cross-referencing in a \\ref command. %% You can use LaTeX's \\ref and \\label commands to keep track of %% cross-references to sections, equations, tables, and figures. %% That way, if you change the order of any elements, LaTeX will %% automatically renumber them.  %% This section also includes several of the displayed math environments %% mentioned in the Author Guide."
    },
    "1004/1004.2715_arXiv.txt": {
        "abstract": "We present high resolution echelle spectra of 7 proximate damped Lyman alpha (PDLA) systems whose relative velocity separation from the background quasar is $\\Delta V < 3000 $ \\kms.  Combining our sample with a further 9 PDLAs from the literature we compare the chemical properties of the proximate systems with a control sample of intervening DLAs.  The PDLAs are usually excluded from statistical studies of absorption selected galaxies and this sample constitutes the first systematic study of their chemical and ionization properties.  Taken at face value, the sample of 16 PDLAs exhibits a wide range of metallicities, ranging from $Z \\sim 1/3 Z_{\\odot}$ down to $Z \\sim 1/1000 Z_{\\odot}$, including the DLA with the lowest N(SiII)/N(HI) yet reported in the literature.  However, some of these abundances may require ionization corrections. We find several pieces of evidence that indicate enhanced ionization and the presence of a hard ionizing spectrum in PDLAs which lead to properties that contrast with the intervening DLAs, particularly when the \\nhi\\ is low.  The abundances of Zn, Si and S in PDLAs with log \\nhi\\ $>$ 21, where ionization corrections are minimized, are systematically higher than the intervening population by a factor of around 3.  We also find possible evidence for a higher fraction of NV absorbers amongst the PDLAs, although the statistics are still modest.  6/7 of our echelle sample show high ionization species (SiIV, CIV, OVI or NV) offset by $>$100 \\kms\\ from the main low ion absorption.  We analyse fine-structure transitions of CII$^{\\star}$ and SiII$^{\\star}$ to constrain the PDLA distance from the QSO.  Lower limits range from tens of kpc up to $>$160 kpc for the most stringent limit.  We conclude that (at least some) PDLAs do exhibit different characteristics relative to the intervening population out to 3000 \\kms\\ (and possibly beyond).  Nonetheless, the PDLAs appear distinct from lower column density associated systems and the inferred QSO-absorber separations mean they are unlikely to be associated with the QSO host.  No trends with $\\Delta V$ are found, although this requires a larger sample with better emission redshifts to confirm. We speculate that the PDLAs preferentially sample more massive galaxies in more highly clustered regions of the high redshift universe. ",
        "introduction": "Absorption systems with $13 <$ log \\nhi\\ $<$ 17 and within $\\sim$ 5000 \\kms\\ of a background QSO are often considered to be `associated' with the quasar host galaxy, its immediate environment or its ejecta. Evidence for this conclusion comes from time variability of the absorption (Hamann et al. 1997a; Wise et al. 2004; Lundgren et al. 2007), partial coverage of the emission source (Hamann et al. 1997b,c; Fox et al. 2008), elevated incidence of highly ionized species (e.g. Ganguly et al. 2001; Richards 2001; Fechner \\& Richter 2009), possible dependence of absorber incidence on QSO properties such as radio loudness or orientation (M\\o ller \\& Jakobsen 1987; Richards et al. 2001; Baker et al. 2002; Vestergaard 2003; Wild et al. 2008) and solar or super-solar metallicities (e.g. Petitjean, Rauch \\& Carswell 1994; Hamann et al. 1997a; Srianand \\& Petitjean 2000; D'Odorico et al. 2004; Fechner \\& Richter 2009).  The associated population of absorbers appears to extend out to at least 10,000 \\kms\\ (Petitjean et al. 1994; Richards et al. 1999; Wild et al. 2008; Tytler et al. 2009).  Although the $z_{\\rm abs} \\sim z_{\\rm em}$ \\lya\\ forest, CIV, MgII and Lyman limit systems have been well-researched in the last 20 years, the same is not true of the proximate damped Lyman alpha systems (DLAs).  This is partly a historical bias.  DLAs are excellent tools for characterising the high redshift galaxy population.  Knowing that associated systems may be connected with the QSO or its host galaxy, statistical samples of DLAs have therefore excluded these potentially special systems (e.g. Wolfe et al. 1995; Ellison et al. 2001; Prochaska, Herbert-Fort \\& Wolfe 2005; Jorgenson et al. 2006; Noterdaeme et al. 2009).  Furthermore, the chemical abundances of intervening DLAs can be determined fairly easily because their gas is predominantly neutral.  Proximity to a QSO could invalidate that assumption and lead to erroneous abundance determinations, unless photoionization modelling is used.  However, it is for exactly these reasons that the study of proximate DLAs (PDLAs) is interesting. PDLAs provide a potential probe of galaxies clustered around (or associated with) QSOs, and even provide insight into the interplay between ionizing radiation and the ISM.  One of the difficulties in studying the proximate DLAs is the small redshift path for each line of sight, which means that PDLAs are very rare. For example, assuming the same number density as intervening systems at $z = 3$, $\\sim$ 100 QSOs would need to be surveyed to find just one PDLA.  A second challenge is constructing a sample of PDLAs selected with a homogeneous QSO-absorber velocity offset ($\\Delta V$) criterion.  Since the derived emission redshift of a QSO depends critically on the emission lines used (e.g. Gaskell 1982; Tytler \\& Fan 1992), the value can be uncertain by over 1000 \\kms\\ ($\\sim$ 3 Mpc proper at $z=3$).  The first systematic study of PDLAs was conducted with the Complete Optical and Radio Absorption Line System (CORALS) Survey data by Ellison et al. (2002) who found a factor of 4 excess in the number density of PDLAs in their radio-selected sample.  Russell, Ellison \\& Benn (2006) extended this work to identify 33 PDLAs in the Sloan Digital Sky Survey (SDSS) Data Release 3 (DR3) and confirmed the PDLA excess, independent of radio loudness, at 3.5$\\sigma$ significance.  In the largest PDLA survey conducted so far, Prochaska, Hennawi \\& Herbert-Fort (2008) use the SDSS DR5 to identify 108 PDLAs with $\\Delta V <$ 3000 \\kms.  In addition to its larger size, this work re-calculated the QSO redshifts in order to obtain improved estimates of $\\Delta V$.  Prochaska et al. (2008b) found that PDLAs outnumber the intervening DLAs at $z \\sim 3$ by a factor of 2, but found no statistically significant excess at $z<2.5$ or $z>3.5$.  In addition to the PDLA number density, Prochaska et al. (2008b) investigated whether their \\hi\\ column density distribution function differed from the intervening DLAs.  They found marginal evidence for an excess of high \\nhi\\ PDLAs.  Previously, Tytler (1982) had found no excess for the slightly lower \\nhi\\ Lyman limit systems but with a much larger sample Prochaska, O'Meara \\& Worseck (2010) find a deficit.  Apart from their incidence and \\nhi\\ distribution, few studies of PDLAs exist to assess whether or not they are consistent with the intervening population.  In particular, PDLAs are rarely targetted for echelle spectroscopy, so little is known about their chemical enrichment. Nonetheless, a small number of PDLAs have appeared in chemical abundance studies (see Table \\ref{lit_z_table} for those that meet our criteria) and they do not appear to have strikingly different properties to the intervening DLAs (Lu et al. 1996; Akerman et al. 2005; Rix et al. 2007).  The range of metallicities, $\\alpha$ element enhancements and dust depletions are all consistent with the range seen amongst intervening DLAs.  Rix et al. (2007) also found negligible ionization corrections were necessary for the PDLA in their study, as is usually the case for other DLAs.  It has therefore usually been concluded that the PDLAs do not constitute a distinct population from the intervening systems (e.g. M\\o ller et al. 1998).  In this work, we present echelle spectra for seven PDLAs (Section \\ref{obs_sec}), determine improved emission redshifts (Section \\ref{z_sec}) and measure column densities for a range of metal line species (Section \\ref{coldens_sec}).  We supplement our sample with other PDLAs critically selected from the literature to form a sample of 16 proximate absorbers (Section \\ref{lit_sec}).  Ionization properties and metal abundances of PDLAs are compared to the intervening systems in Sections \\ref{ion_sec} and \\ref{abund_sec}. Finally, we use measurements of CII$^{\\star}$ and upper limits to SiII$^{\\star}$ to constrain the distance between the PDLA and the background QSO (Section \\ref{dist_sec}).  Unless otherwise stated, we assume $\\Omega_M = 0.3$, $\\Omega_{\\Lambda}=0.7$ and H$_0$ = 70 \\kms\\ Mpc$^{-1}$. ",
        "conclusions": "We have presented new high resolution echelle spectra for seven proximate damped Lyman alpha systems with $\\Delta V < 3000$ \\kms.  The metal column densities derived from Voigt profile fits and the apparent optical depth method are complemented with abundances for a further 9 PDLAs taken from the literature.  Our PDLA sample is compared to the most complete sample of intervening DLAs currently available.  Our principal conclusions are:  \\begin{enumerate} \\item PDLAs exhibit a range of metallicities at a given redshift, ranging from $\\sim$ 1/3 to 1/1000 of the solar value (Section \\ref{metal_sec} and Figure \\ref{Fe_Si_z}).  One of the PDLAs in our sample exhibits the lowest N(SiII)/N(HI) of any known DLA and has a value similiar to the intergalactic medium at this redshift.  \\item Based on this modest-sized sample, there is a general trend (with the exception of one fairly high velocity PDLA) for low metallicities ($Z \\sim 1/50 Z_{\\odot}$) in PDLAs with log \\nhi $<$ 21 and higher metallicities ($Z \\sim 1/10 Z_{\\odot}$) when the HI column density is higher (Figure \\ref{Fe_Si_S_Zn_HI}).  At these high \\hi\\ column densities, the metallicities of PDLAs are systematically higher than the intervening sample by a factor of around three.  \\item At least half of the PDLAs with \\nhi\\ $<$ 20.8 have sub-solar ratios of [S/Si] which can not be easily explained by known dust or nucleosynthetic trends (Section \\ref{ion_alpha} and Figure \\ref{S_Si_HI}).  We interpret the low values as resulting from ionization by a hard spectrum.  Sub-solar values of [S/Si] can be present even in PDLAs with $\\Delta V$ $>$ 2000 \\kms, and with no obvious trend with velocity.  \\item In addition to the dependence of metallicity on \\nhi\\ and the sub-solar [S/Si] ratios, other indications of enhanced/hard ionization in the PDLAs which distinguish them from the intervening DLAs are: 1) a possibly higher fraction of NV absorbers (tentative based on the small number statistics of our sample), 2) higher fractions of SiIV/SiII at low \\nhi, 3) similar velocity structure in SiIV as SiII in 3/5 QSOs where the comparison can be made, 4) low ratios of [Ar/Si] (Section \\ref{ion_sec}).  \\item Most of the PDLAs (6/7) in our sample of seven have additional high ionization gas at large negative velocities of $-100$ to $-825$ \\kms\\ (Section \\ref{disc_high} and Table \\ref{high_ion_tab}).  The most extreme examples both have either a confirmed (Q0151+048) or tentative (J1131+6044) close companion QSO.  \\item One of the PDLAs in our sample has very sub-solar [Fe/Zn] consistent with a large dust depletion fraction.  However, in general, the [Fe/Zn] ratios of the PDLAs are consistent with the intervening population (Section \\ref{dust_sec} and Figure \\ref{Zn_Fe}).  The PDLA towards J1604+3951 has very different depletion in its two main components.  \\item The range of alpha-to-iron ratios in the PDLAs is also consistent with the intervening DLAs, although there are a few notably high values (Section \\ref{alpha_sec} and Figure \\ref{Si_Fe_HI}).  \\item From analyses of the fine structure lines of SiII and CII, QSO--absorber distances at least a few tens of kpc are determined, with one case being constrained to have a separation of $>$ 160 kpc (Section \\ref{dist_sec}).  \\end{enumerate}  We conclude that PDLA properties are generally consistent with an origin external to the QSO host (in contrast with, e.g. the narrow line associated systems, D'Odorico et al. 2004).  However, the larger abundances (at least of high \\nhi\\ PDLAs) imply that they may not be representative of the intervening sample.  We suggest that the PDLAs may preferentially sample overdense environments where biased galaxy formation has assembled more massive galaxies with higher metallicities.  Indeed, at low redshift, some associated absorbers have been identified with small clusters of galaxies (e.g. Bergeron \\& Boisse 1986; Hamann et al. 1997a).  If confirmed with larger larger samples, this means that PDLAs could be used as probes of massive galaxies at high redshift.  PDLAs have been widely excluded from most DLA surveys.  The results presented here indicate that, depending on the science objective, this is a valid approach.  We have argued that PDLAs may sample a rather special population of DLAs, possibly those in proto-clusters. Although this renders the PDLAs an interesting probe of high redshift galaxies, our results reveal some of the biases that could be introduced into statistical surveys that do not impose a $\\Delta V$ limit in their DLA selection.  A high priority for future work will be the improvement of emission redshift determinations. The results presented here imply that even at fairly large relative velocities ionization may affect abundance determinations. Determining more accurate $\\Delta V$ values in larger PDLA samples will provide an empirically motivated cut-off for studies of intervening DLAs."
    },
    "1004/1004.5135_arXiv.txt": {
        "abstract": "We present HST/WFPC2 observations across the disk of the nearby isolated dwarf S0 galaxy NGC~404, which hosts an extended gas disk. The locations of our fields contain a roughly equal mixture of bulge and disk stars.  All of our resolved stellar photometry reaches $m_{F814W}=26$ ($M_{F814W}=-1.4$), which covers 2.5 magnitudes of the red giant branch and main sequence stars with ages $<$300 Myr.  Our deepest field reaches $m_{F814W}=27.2$ ($M_{F814W}=-0.2$), sufficient to resolve the red clump and main-sequence stars with ages $<$500~Myr. Although we detect trace amounts of star formation at times more recent than 10 Gyr for all fields, the proportion of red giant stars to asymptotic giants and main sequence stars suggests that the disk is dominated by an ancient ($>$10~Gyr) population. Detailed modeling of the color-magnitude diagram suggests that $\\sim$70\\% of the stellar mass in the NGC~404 disk formed by $z\\sim2$ (10 Gyr ago) and at least $\\sim$90\\% formed prior to $z\\sim1$ (8 Gyr ago).  These results indicate that the stellar populations of the NGC~404 disk are on average significantly older than those of other nearby disk galaxies, suggesting that early and late type disks may have different long-term evolutionary histories, not simply differences in their recent star formation rates.  Comparisons of the spatial distribution of the young stellar mass and FUV emission in GALEX images show that the brightest FUV regions contain the youngest stars, but that some young stars ($<$160~Myr) lie outside of these regions.  FUV luminosity appears to be strongly affected by both age and stellar mass within individual regions.  Finally, we use our measurements to infer the relationship between the star formation rate and the gas density of the disk at previous epochs. We find that most of the history of the NGC~404 disk is consistent with star formation that has decreased with the gas density according to the Schmidt law.  However, $\\sim$0.5--1 Gyr ago, the star formation rate was unusually low for the inferred gas density, consistent with the possibility that there was a gas accretion event that reignited star formation $\\sim$0.5 Gyr ago.  Such an event could explain why this S0 galaxy hosts an extended gas disk. ",
        "introduction": "While both redshift surveys and stellar population studies agree that late-type disks form the majority of their stars by $z\\sim1$ \\citep[e.g.,][]{ravindranath2004,barden2005,melbourne2007,williams2008}, the ages of early-type disks are not as well constrained.  The theoretical and high redshift observational data suggest that the old populations of early type disks may be systematically older than those of later type disks.  On the other hand, the morphological differences between early and late type disks may be of recent origin, due only to differences in star formation over the past few hundred Myr.  Such a scenario has has been suggested for the S0 galaxy NGC~5102 \\citep{davidge2008}, as well as for the differences between dEs and dIrrs.  Unfortunately, little is known about the evolutionary history of S0 galaxies, especially in the field.  Although typical S0 galaxies appear to differ significantly from E galaxies, it is not yet clear where S0 galaxies fit into the Hubble tuning fork \\citep{vandenbergh2009} or what amount of star formation they typically contain \\citep{temi2008}.  While non-interacting S0 galaxies tend to have disks similar in thickness to those of late-type galaxies \\citep{pohlen2004,li2009}, other properties of S0 disks differ significantly from those of late-type galaxies.  For example, their color gradients are less pronounced than those of late-type disks \\citep{macarthur2004}, and S0 galaxies are typically fainter than later type disks \\citep{vandenbergh2009}. S0 galaxies are predominantly found in galaxy clusters and are mostly stripped of {\\sc H i} \\citep{poggianti2009,vandenbergh2009}, whereas field S0 galaxies like NGC~404 are less common and could be a type of post-starburst dwarf \\citep{davidge2008}.  As the nearest face-on example of an S0 galaxy \\citep[3.05~Mpc, m-M$_0$ = $27.42\\pm0.02$; M$_K$ = -18.9; $A_V$=0.08; $i$=11$^{\\circ}$; $V_{c}=190$~km~s$^{-1}$][]{dalcanton2009,schlegel1998,delrio2004}, NGC~404 has been previously studied in some detail. Although its optical morphology suggests that NGC~404 is a classic ``red and dead'' early-type galaxy, the center appears to have a significant young component.  With long-slit spectroscopy of the galaxy core, \\citet{bouchard2010} found evidence for intermediate-aged, and old stellar populations, along with a much younger component that fades very quickly with radius.  The nucleus also has a dominant 1 Gyr old population \\citep[$\\sim$50\\% of the mass,][]{seth2010}.  UV spectroscopy also indicates some very young ($\\leq$ 10 Myr) stars in the galaxy center \\citep{maoz1998}. Farther out, NGC~404 has a very large ($>$20 kpc) stellar disk that is known to be dominated by red giants \\citep{tikhonov2003}.  This disk also contains a large amount of {\\sc H i} gas in a well-defined inner disk ($<$5 kpc) and a warped outer disk \\citep{delrio2004}, although the density of gas is low ($<$2$\\times$10$^{20}$ M$_{\\odot}$ pc$^{-2}$ throughout) and only a faint ring of star formation is detected in the far ultra-violet (FUV)\\footnote{http://www.galex.caltech.edu/newsroom/glx2008-02r.html} \\citep{wiklind1990}.  The origin of the gas disk of NGC~404 is somewhat mysterious. \\citet{delrio2004} suggest that the disk is likely due to a recent (0.5--1 Gyr ago) merger event with a dIrr galaxy, motivated by their detection of a possible warp in the disk.  However, NGC~404 is very isolated for a galaxy of its mass \\citep[$\\Theta$=-1.0; see][for details]{karachentsev2004}. With no other galaxies of any kind detected within 1~Mpc of NGC~404, such a merger is not obviously consistent with the galaxy's environment.  The gas disk could therefore simply be due to late time gas infall from filaments rather than accretion of bound objects.  Furthermore, a gas disk of such low density could have a very long lifetime given the low star formation rate, suggesting that it could have been accreted far earlier. By comparing the stellar populations and gas densities, we can constrain the origin of this unusual disk.  We examine the stellar populations of NGC~404 in detail using deep observations from HST.  We determine the star formation history (SFH) of several portions of the galaxy by fitting the distribution of stars in color-magnitude diagrams (CMDs) with model distributions determined from stellar evolution isochrones \\citep[as described in][]{dolphin2002}.  Section 2 describes our data set and analysis procedures.  Section 3 presents the results of our measurements. Section 4 interprets the measurements in the context of the gas properties and our understanding of S0 galaxies, and Section 5 summarizes our conclusions.  We assume a distance of 3.05~Mpc \\citep{dalcanton2009} for conversions of angular measurements to physical distances and adopt an inclination angle $i$=11$^{\\circ}$ \\citep{delrio2004} for surface density measurements.  We adopt a five-year WMAP \\citep{dunkley2009} cosmology for all conversions between time and redshift. ",
        "conclusions": "We have measured resolved stellar photometry for $\\sim$15 arcmin$^2$ of the inner disk and $\\sim$10 arcmin$^2$ of the outer disk of the S0 galaxy NGC~404.  Detailed fitting of the resulting CMDs shows that the disk is dominated by stars older than $\\sim$10~Gyr, with all of our inner disk data constraining the age of $\\sim$90\\% of the stellar mass to $>$8~Gyr and our best data constraining the age of 75\\% of the stellar mass to $>$10~Gyr.  These results show that the difference between this disk and most later-type disks that have been studied in detail is seen in the old stellar population as well as the young.  We found no trends between the age distributions of the young or old populations and distance from the galaxy center within the observed regions.  However, as seen by \\citet{tikhonov2003} the mean RGB color and AGB/RGB ratio of the innermost regions are consistent with a bulge population in NGC~404 that is younger than the disk.  Despite its low gas density, NGC~404 appears to have been forming stars throughout its history on timescales of Gyr, even in the inner disk ($\\sim$1--1.5$'$).  The SFH and the current density of the gas disk are consistent with NGC~404 having evolved roughly in a closed box over much of its history, starting as a large disk of high surface density and approximately following the Schmidt law as its overall star formation rate and gas density decreased.  Such long-lived disks have been observed in isolated galaxies before, as in the compact dwarf ADBS 113845+2008 \\citep{cannon2009}, whose gas disk is also of low density.  Perhaps these low density disks have very long lifetimes due to very low star formation efficiencies at these densities.  Our SFHs indicate that star formation in the low density relic disk may be episodic, going through minor episodes of star formation that last a few hundred Myr every Gyr or so.  On the other hand, in NGC~404 there is some evidence that this passive evolution was slightly disrupted by an event 0.6-0.9 Gyr ago, possibly a small merger.  Such a merger event would alleviate the necessity for very low star formation efficiency $\\sim$1~Gyr ago by allowing that the currently-observed gas disk was not fully in place at that time.  Support for this work was provided by NASA through grants GO-10915 and GO-11719 from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Incorporated, under NASA contract NAS5-26555."
    },
    "1004/1004.4419_arXiv.txt": {
        "abstract": "The OH molecule, found abundantly in the Milky Way, has four transitions at the ground state rotational level($J = 3/2$) at cm wavelengths. These are E1 transitions between the $F^+$ and $F^-$ hyperfine levels of the $\\Lambda$ doublet of the $^{2}\\Pi_{\\frac{3}{2}}, J=3/2$ state. There are also forbidden M1 transitions between the hyperfine levels within each of the doublet states occuring at frequencies 53.171~MHz and 55.128~MHz. These are extremely weak and hence difficult to detect. However there is a possibility that the level populations giving rise to these lines are inverted under special conditions, in which case it may be possible to detect them through their maser emission. We describe the observational diagnostics for determining when the hyperfine levels are inverted, and identify a region near W44 where these conditions are satisfied. A high-velocity-resolution search for these hyperfine OH lines using the low frequency feeds on four antennas of the GMRT and the new GMRT Software Backend(GSB) was performed on this target near W44. We place a 3$\\sigma$ upper limit of $\\sim$17.3 Jy (at 1 $km~s^{-1}$ velocity resolution) for the 55~MHz line from this region. This corresponds to an upper limit of 3  $\\times 10^8$ for the amplification of the Galactic synchrotron emission providing the background. ",
        "introduction": "The ground state rotational level (J=3/2) of the $^{2}\\Pi_{\\frac{3}{2}}$ ladder of the OH molecule is split by $\\Lambda$-doubling into two states, each of which is further split into two states because of the hyperfine interaction between the spin of the unpaired $2p$ electron of the O atom and the nuclear magnetic moment of the H atom. The four well-known 18~cm OH lines, namely the ``main lines'' at 1665~MHz \\& 1667~MHz and the ``satellite lines'' at 1612 MHz \\& 1720 MHz arise from transitions between these levels.  The main lines, associated with star-forming regions, are sometimes seen strongly mased. The 1612 MHz and 1720 MHz satellite line masers are usually associated respectively with late-type stars and shocked regions, and sometimes seen as a mased-thermal pair\\citep{masers}. Transitions between the hyperfine states of the same $\\Lambda$ doublet state (M1) occur at 53.171 MHz($J=3/2,\\ F:2^- \\rightleftarrows 1^-$) and 55.128 MHz($J=3/2,\\ F:2^+ \\rightleftarrows 1^+$). The thermal emission or absorption from these lines is extremely weak given that they are highly forbidden M1 transitions with $A \\leq 10^{-18} s^{-1}$\\citep{destombes}. The expected thermal line brightness temperature for the 54 MHz lines is 1~K for an OH column density of 10$^{17}$ cm$^{-2}$\\citep*{roshi}. It is not possible to detect such weak thermal lines at 54 MHz, given the high Galactic background. Hence, the lines have to be necessarily mased inorder to be detectable in reasonable integration times. We discuss below the conditions for inversion of the populations in the levels from which the 53~MHz and 55~MHz lines arise.  \\begin{figure*} \\begin{center} \\includegraphics[scale=0.6]{OH_ladder.eps} \\caption{Ground state energy levels(J=3/2) of the $^2\\Pi_\\frac{3}{2}$ ladder of the OH molecule. n$_1$, n$_2$, n$_3$ and n$_4$ are the hyperfine level populations. The frequencies, in MHz, of the transitions between these levels are also indicated.} \\label{ladder} \\end{center} \\end{figure*}  The J=3/2 state of the $^2\\Pi_{\\frac{3}{2}}$ ladder is given in Figure~\\ref{ladder}.  For the 53 MHz line to be mased we require the corresponding level populations be inverted:  \\begin{equation} \\frac{n_{2}/g_{2}}{n_{1}/g_{1}} > 1 \\label{oh53} \\end{equation}  which implies  \\[ \\frac{n_{2}/g_{2}}{n_{3}/g_{3}}  \\times \\frac{n_{3}/g_{3}}{n_{1}/g_{1}} > 1 \\ \\ \\ and \\ \\ \\ \\frac{n_{2}/g_{2}}{n_{4}/g_{4}}  \\times \\frac{n_{4}/g_{4}}{n_{1}/g_{1}} > 1 \\] \\\\ Condition \\eqref{oh53} is satisfied if the populations in the levels corresponding to the 1665 MHz and 1720 MHz lines are inverted while those in the levels corresponding to the  1612 MHz and 1667 MHz lines are not. Observationally if the 1665 MHz and 1720 MHz lines are mased, while the 1612 MHz and 1667 MHz lines are not, we could expect the 53 MHz line levels to be inverted. While the inequalities are necessary and sufficient conditions, the requirement on the individual lines are merely sufficient, eg.  \\[ \\frac{n_{2}/g_{2}}{n_{3}/g_{3}} \\leq 1 \\ \\ \\ and \\ \\ \\ \\frac{n_{3}/g_{3}}{n_{1}/g_{1}} \\gg 1 \\] \\\\ could still give \\[ \\frac{n_{2}/g_{2}}{n_{1}/g_{1}} > 1 \\] \\\\ Similarly, for the 55 MHz line to be mased we require the corresponding level populations be inverted:  \\begin{equation} \\frac{n_{4}/g_{4}}{n_{3}/g_{3}} > 1 \\label{oh55} \\end{equation}  which implies  \\[ \\frac{n_{4}/g_{4}}{n_{2}/g_{2}}  \\times \\frac{n_{2}/g_{2}}{n_{3}/g_{3}} > 1 \\ \\ \\ and \\ \\ \\ \\frac{n_{4}/g_{4}}{n_{1}/g_{1}}  \\times \\frac{n_{1}/g_{1}}{n_{3}/g_{3}} > 1 \\] \\\\ Condition \\eqref{oh55} is satisfied if the populations in the levels corresponding to the 1667 MHz and 1720 MHz lines are inverted while those in the levels corresponding to the  1612 MHz and 1665 MHz lines are not. Observationaly if the 1667 MHz and 1720 MHz lines are mased, while the 1612 MHz and 1665 MHz lines are not, we could expect the 55 MHz line levels to be inverted. \\begin{table} \\begin{center} \\begin{tabular}{|c|c|c|c|c|c|} \\hline 1612 & 1665 & 1667 & 1720 & 53  & 55 \\\\ MHz &  MHz &  MHz &  MHz & MHz & MHz \\\\ \\hline Not        & Inverted       & Not       & Inverted &$\\checkmark$& -- \\\\ inverted &                  & inverted                   & & & \\\\ \\\\ Not      & Not       & Inverted       & Inverted      & -- & $\\checkmark$ \\\\ inverted & inverted                 &                    & & & \\\\ \\\\ \\hline \\end{tabular} \\end{center} \\caption{The conditions on the level populations corresponding to the 18~cm OH lines, required for {\\bf maser emission} of the 53 and/or 55 MHz hyperfine OH lines} \\end{table} In summary, a region where the 1720 MHz line is mased, but the 1612 MHz line is not, would be a promising region to look for mased emission from the 53/55 MHz lines. Which of these is likely to be mased depends on whether the 1665 MHz or 1667 MHz line is mased. It is interesting to note in this context that there are regions in our Galaxy\\citep{turner} and in external galaxies\\citep*{langevelde95,chengkan} where the satellite lines are conjugate, viz. where their profiles are mirror images of one another, i.e. when one is in emission the other is in absorption, and the sum of the two profiles is consistent with noise. Regions such as these are promising ones to search for the 53 MHz and 55 MHz OH lines. If the hyperfine OH 53 MHz or 55 MHz lines are strongly amplified, then the lines may be detectable. However, the amplification factor is a matter of speculation. \\citet{roshi} attempted to detect the 53~MHz line, but could only place an upper limit to the amplification factor. ",
        "conclusions": "Figure \\ref{psfit} shows the final cumulative averaged spectrum obtained from seven out of the eight available polarizations of the four antennas. The velocity resolution is $0.25\\ km~s^{-1}$ per channel. The ordinate is the baseline-subtracted line-to-continuum flux ratio, after fitting for the baseline with a second order polynomial for the passband shown in the plot. At around the expected LSR velocity of 46 $km~s^{-1}$, where the 1720 MHz line is seen inverted towards W44 in as many as 25 hotspots\\citep{claussen}, there is a weak $\\sim$4$\\sigma$ spectral emission feature, whose peak is 0.006 in units of line-to-continuum temperature ratio, (T$_l$-T$_c$)/T$_c$. The RMS optical depth is 0.0014 units over a 0.25 $km~s^{-1}$ channel. We conclude the 55 MHz OH line is not detected to the 4$\\sigma$ limit. However, this feature becomes more prominent(see Figure \\ref{psfit_smoothed}) when the spectrum is smoothed to $1\\ km~s^{-1}$,  commensurate with the velocity widths of the hotspots listed by \\citet{claussen}. We assume that the system temperature is dominated by the sky temperature at these frequencies. The sky temperature, 32100 K, is obtained by scaling the temperature of the target region, using a power law index of -2.5, from the 34.5 MHz survey of \\citet{geetee}. For the measured aperture efficiency of $\\eta_{A}$=0.7\\citep{amiri} and $A = 1590\\ m^{2}$ per polarisation, we obtain, using the equation $S A_{e\\!f\\!f} = kT_{l}$,  a 3$\\sigma$ flux limit of 26.5 Jy for the observed line-to-continuum ratio. On smoothing to 1 $km~s^{-1}$ we obtain a flux limit of 17.3 Jy per channel. \\begin{figure} \\begin{center} \\includegraphics[width=3.35in]{postfit.eps} \\caption{Baseline-subtracted line-to-continuum flux ratio around G34.3+0.1 for the 55 MHz OH line. The velocity resolution is 0.25 $  km~s^{-1}$ per channel.} \\label{psfit} \\end{center} \\end{figure} \\begin{figure} \\begin{center} \\includegraphics[width=3.35in]{postfit_smoothed.eps} \\caption{Baseline-subtracted line-to-continuum flux ratio around G34.3+0.1, smoothed to a velocity resolution of 1.0 $km~s^{-1}$ per channel.} \\label{psfit_smoothed} \\end{center} \\end{figure}  The equation of radiative transfer gives \\[ T_l = T_{bg}e^{-\\tau_{\\nu} } + T_x(1-e^{-{\\tau}_{\\nu}}) \\]  where $T_{bg}$ is the brightness temperature of the background source providing the radiation and $T_x$ is the excitation temperature.  Both the background temperature $T_{bg}$ and the excitation temperature $T_x$ contribute to the maser line brightness, but when $T_{bg}$ far exceeds $T_x$, say $T_{bg} > 10|T_x|$, we can omit the contribution of $T_x$ to $T_l$. The 3$\\sigma$ limit to the line brightness temperature from our observation is $T_l$ = 74 K for a 0.25 $km~s^{-1}$ channel which translates into 26.5 Jy. At a resolution of 1 $km~s^{-1}$, the limiting flux is 17.3 Jy. For comparison for the 53~MHz line, \\citet{roshi} place a limit of 49~Jy at a velocity resolution of 4.6 $km~s^{-1}$.  For a maser that is $100\\ mas$\\citep{hoffman} in size, the beam dilution factor for the $10 \\degr$ GMRT beam is $1.3 \\times 10^{1\\!1}$. Given $T_{bg}$ = 32100 K, we obtain the maser amplification, omitting the contribution from $T_x$, as \\[ G = \\frac{74 \\times 1.3 \\times 10^{1\\!1}}{32100} \\sim 3 \\times 10^8 \\] providing an upper limit to the optical depth as $log(G)$ = 19.52. At a resolution of 1.0 $km~s^{-1}$, we get an upper limit of $\\tau = log(G) - log(74/16.49) = 18.58$ given that \\mbox{$T_l = 16.49$ K.} The calculation above assumes that (some reasonable fraction of) the Galactic synchrotron emission provides the background radiation that is amplified by the maser."
    },
    "1004/1004.3239_arXiv.txt": {
        "abstract": "We propose a closure model for the transport of entropy and momentum in astrophysical turbulence, intended for application to rotating stellar convective regions. Our closure model is first presented in the Boussinesq formalism, and compared with laboratory and numerical experimental results on Rayleigh-B\\'enard convection and Homogeneous Rayleigh-B\\'enard convection. The predicted angular momentum transport properties of the turbulence in the slowly rotating case recover the well-known $\\Lambda-$effect, with an amplitude uniquely related to the convective heat flux. The model is then extended to the anelastic case as well as the fully compressible case. In the special case of spherical symmetry, the predicted radial heat flux is equivalent to that of mixing-length theory. For rotating stars, our model describes the coupled transport of heat and angular momentum, and provides a unified formalism in which to study both differential rotation and thermal inhomogeneities in stellar convection zones. ",
        "introduction": "Turbulent convection occurs frequently in stellar interiors and other astrophysical fluid flows.  While convective motion naturally transports heat and chemical elements, the transport of angular momentum by convection in rotating bodies is a more subtle issue.  It is of particular interest in the case of the Sun, where the internal pattern of rotation has been measured but remains incompletely understood.  It may also play an significant role in accretion flows.  Numerical simulations of astrophysical convection are becoming increasingly powerful and capable of resolving a widening range of length and time scales.  Nevertheless, a simpler, statistical description of turbulent transport is desirable in order to treat the effects of convection on the structure and evolution of stars.  It almost goes without saying that such a description cannot be derived strictly from the equations of fluid dynamics but must involve some modelling or parametrization.  The mixing-length theory of turbulent transport was developed by Prandtl (1925) and applied to stellar convection by Biermann (1932). It is still the basic model used in most calculations of stellar structure and evolution, usually in the form devised by B\\\"ohm-Vitense (1958).  The main purpose of mixing-length theory is to relate the convective heat flux to the superadiabatic gradient; in this context it does not usually deal with the transport of (angular) momentum that arises in the presence of shear or rotation.  A standard theoretical approach to convection in differentially rotating stars is set out in the monograph by R\\\"udiger (1989). Angular momentum transport is described by a Reynolds stress tensor whose components can be related to the large-scale mean flows and thermodynamical gradients. A first contribution to the Reynolds stress is typically proportional to the angular velocity gradient through a turbulent viscosity coefficient. An important additional contribution comes from the $\\Lambda$-effect (named after Lebedinsky), whereby even uniformly rotating convection transports angular momentum by virtue of its anisotropy. Attempts to constrain or parameterize these quantities have been made through local numerical simulations (e.g.\\ K\\\"apyl\\\"a, Korpi \\& Tuominen 2004) or theoretical models (e.g.\\ Kitchatinov \\& R\\\"udiger 1993). Mean-field models of stellar rotation (e.g. Kitchatinov \\& R\\\"udiger 1999, Rempel 2005) have been developed which use such parameterized expressions for the Reynolds stress and heat flux.  Reynolds-stress models of turbulent flows have been developed in the engineering community over several decades (e.g.\\ Pope 2000). The exact equation governing the Reynolds stress in a turbulent fluid cannot be solved because of the well known closure problem whereby an infinite hierarchy of correlations is involved.  Nevertheless, by parametrizing the difficult terms in this equation, models can be constructed that bear some fidelity to the turbulent dynamics.  From a more physical point of view, what is obtained is a time-dependent constitutive equation for the turbulent fluid, which relates the turbulent stress to the local history of deformation. There is a close similarity with models of non-Newtonian fluids (Ogilvie \\& Proctor 2003). The advection and deformation of the turbulent stress are accurately represented since they derive from linear terms in the Reynolds-stress equation, while the nonlinear `relaxation' effects are only modelled (as is also true for non-Newtonian fluids).  A similar approach can be applied to turbulent convection in which buoyancy forces play an essential role.  The additional correlations that must be considered are the flux and the variance of entropy (or temperature, in the Boussinesq approximation).  This approach offers some benefits over the conventional description in terms of a turbulent viscosity and a $\\Lambda$-effect.  It can be formulated in a covariant manner and is not tied to the spherical geometry of a slowly rotating star.  It starts from a more fundamental description and allows phenomena such as the $\\Lambda-$effect to emerge in a natural way from more elementary considerations.  It may also allow a more unified approach to be taken towards problems involving astrophysical turbulence.  In this paper we explore some of the consequences of a simple dynamical model of astrophysical convection of this type.  The model derives from one originally conceived for magnetohydrodynamic turbulence in accretion discs (Ogilvie 2003) and later applied to rotating shear flows without magnetic fields (Garaud \\& Ogilvie 2005, GO05 herafter). Our motivation is to develop and test a model that can be applied to the convective zone of the Sun, to other stars or to accretion discs. We emphasize, however, that our model is chosen to be as simple as possible for the purposes of this investigation.  In contrast with some of the engineering literature, we restrict the algebraic complexity in order to retain a physical understanding of the terms in the equations.  Further refinements are likely to be required in order to provide an accurate match to a wide range of data.  In comparing a closure model of astrophysical convection with experimental and numerical results, we face certain difficulties. Astrophysical convection usually takes place at very high Rayleigh number, in a highly turbulent regime.  Experiments have been conducted at very high Rayleigh number but mainly for the Rayleigh--B\\'enard problem in which the flow is dominated by boundary layers, which may not be relevant in an astrophysical context, or by mean flows not represented in the closure model.  An alternative system is provided by the homogeneous Rayleigh--B\\'enard problem, which has periodic boundary conditions in all directions.  This model, however, has certain peculiarities of its own.  These issues will be addressed in the sections that follow.  In the remainder of the paper, we develop the closure model first in the Boussinesq approximation (Section~2) and apply it to the standard Rayleigh--B\\'enard problem (Section~3).  We then consider the homogeneous Rayleigh--B\\'enard system with triply periodic boundary conditions (Section~4); in this section we also introduce rotation and discuss the $\\Lambda-$effect.  We then adapt the model to the anelastic approximation for use in stars and other astrophysical flows (Section~5) and finally draw conclusions (Section~6).  A number of technical details are covered in the appendices. ",
        "conclusions": "We have laid out the foundations of a new second-order closure model for the dynamics of turbulent convection, with future applications to stellar convective regions in mind. This model is a direct extension of the work of Ogilvie (2003) and GO05, and has similar properties. The proposed closure has a straightforward physical interpretation, and well-understood limitations.  Comparison with laboratory and numerical experiments reveals good overall agreement of the model predictions with known properties of rotating shear flows (GO05) and high Rayleigh-number rotating convection (this work). In particular, our model naturally reproduces the standard scaling relationships between the Rayleigh and Nusselt numbers for Rayleigh-B\\'enard convection and for Homogeneous Rayleigh-B\\'enard convection, and contains the well-known $\\Lambda$-effect describing angular momentum transport in a rotating turbulent fluid.  When extended to the anelastic (or fully compressible) case, our formalism can be applied to study convection in stellar interiors. Note that the effects of Maxwell stresses can also straightforwardly be included following Ogilvie (2003) if needed. We show that the model naturally reduces to a version of mixing-length theory when applied in a one-dimensional framework. In the presence of rotation it becomes a powerful tool to study within a single framework the multi-dimensional balance involving large-scale mean quantities such as the entropy profile, the meridional circulation and the differential rotation. Future work applying our closure model in a spherical shell geometry will help understand some of the trends seen in the increasingly large number of available observations of stellar differential rotation."
    },
    "1004/1004.1185_arXiv.txt": {
        "abstract": "{Vortex-type motions have been {measured} by tracking bright points in high-resolution observations of the solar photosphere. These small-scale motions are thought to be determinant in the evolution of magnetic footpoints and their interaction with plasma and therefore likely to play a role in heating the upper solar atmosphere by twisting magnetic flux tubes. We report the observation of magnetic concentrations being dragged towards the center of a convective vortex motion in the solar photosphere from high-resolution ground-based and space-borne data.  We describe this event by analyzing a series of images at different solar atmospheric layers. By computing horizontal proper motions, we detect a vortex whose center appears to be the draining point for the magnetic concentrations detected in magnetograms and well-correlated with the locations of bright points seen in G-band and CN images.} ",
        "introduction": "Convection is an important mechanism of energy exchange responsible for the formation of the extensively studied solar granulation. The interaction between the convective plasma flows and solar magnetic fields causes the appearance of  solar structures on many spatial scales, sunspots and pores being the most conspicuous ones. Magnetism is often understood to be the main process controlling the whole  Sun's configuration and behavior, even in the so-called quiet regions where small magnetic elements such as bright points (hereafter BPs) are detected mainly along dark lanes in-between granules. Small-scale ($\\lesssim$0.5~Mm) convectively driven vortex-type motions were discovered by \\cite{bonet2008}.  {These motions were suggested by \\citet{sturrock81} to account for about half of the photospheric kinetic energy, but this could not be confirmed due to the limitations of early observations. On larger spatial scales, i.e., granular sizes, vortical motions were observed by \\citet{brandt1988}, \\citet{tarbell91}, and \\citet{title92}. Signatures of photospheric vortex flows have also been observed on even larger scales such as supergranular junctions \\citep{attie2009}. Vortical motions in the solar photosphere were studied by simulations \\citep[][and references herein]{zirker93} and are understood to play an important role in the evolution of magnetic footpoints and heating in the upper solar atmosphere \\citep{brandt1988,wang1995}}. Here we describe a vortex motion in the solar photosphere that persists for about 20 minutes. It seems to affect the small-scale magnetic concentrations, which appear to be dragged towards the vortex center.   \\begin{figure*} \\centering \\includegraphics[angle=-90,width=1.\\linewidth]{balmacedafig1.ps} \\caption{Quasi-simultaneous observations showing the evolution of a quiet sun region on 29 September 2007.  \\emph{From top to bottom}: Mg I magnetogram, G-band, CN, Ca~II~H. Every panel displays the average image over 4-min intervals and the time stamps correspond to the initial time for each interval, respectively.  BPs are detected in all filtergrams and are co-spatial with the magnetic concentrations observed in magnetograms. Spatial scales in both axis are in arcsec and correspond to solar coordinates (FOV $\\sim$6$\\arcsec$$\\times$5$\\arcsec$). The black box at 09:04 UT shows the FOV covered by the SP in Fig.~\\ref{lilia}. Encircled in white in G-band frames at 08:48, 08:52, and 08:56 UT is the location of the detected vortex (Sect.~\\ref{S:3.3}) for reference.} \\label{secuencia} \\end{figure*} ",
        "conclusions": "\\label{S:4}  In Fig.~\\ref{barycenter}, we plot the trajectories followed by the \\emph{centroid} of the distribution of {magnetic flux density} for lobes L1 and L2. The \\emph{centroids} denoted with plus-signs and asterisks for the respective lobes are now calculated for the individual maps obtained from SP data with a cadence of $\\sim$36~s using the magnetic flux density $M$ for weighting. They are defined by $x_{c}=\\sum_{i=1}^{N} x_{i} M(x_{i},y_{i}) / \\sum M(x_{i},y_{i})$ and $y_{c}=\\sum_{i=1}^{N} y_{i} M(x_{i},y_{i}) / \\sum M(x_{i},y_{i})$, where ($x{i}$, $y_{i}$) are the coordinates of the $i$th pixel for each lobe. In spite of its lower resolution, we use the SP dataset because its time series extends until 09:44 UT, allowing us to track the dynamics of L1 and L2 over a longer period. The trajectories are superimposed on both the respective flow maps computed for the SST time series (\\emph{s1} and \\emph{s2}) and the divergence field for each series. {The comparison with the centroid positions calculated using the MgI magnetograms (top row in Fig.~\\ref{secuencia}) allowed us to estimate the errors to be $\\epsilon_{x1} = \\pm0.35'' $, $\\epsilon_{y1} = \\pm0.38''$ for L1 and $\\epsilon_{x2} = \\pm0.32''$, $\\epsilon_{y2} = \\pm0.28''$ for L2.} {From the centroid positions, it was possible to estimate the horizontal velocities for each lobe}. While L1 motion is confined to a radius of approximately 1\\arcsec, L2 travels a much larger distance. In the time interval prior to 09:14, its motion seems to be strongly influenced by the vortex formed in the coordinates [3.5, 1.5] in the left panel of Fig.~\\ref{barycenter}, increasing its velocity {from $\\sim$1.5 to 3~km~s$^{-1}$} as it approaches this location. After the vortex has disappeared, L2 continues its movement {with decreasing velocity}  probably influenced by the continuous evolution of the surrounding granules towards a new sink denoted by the region of negative divergence whose center is located at coordinates [1.0, 2.5] in Fig.~\\ref{barycenter} (right). {Because of the lack of space here, additional analysis of centroids velocities will be deferred to Paper~I. }  {A possible explanation of the observed magnetic concentrations being dragged towards the vortex center could be that this is the result of the stochastic evolution of granules that allow the BPs to approach the vortex influence, increase their velocities, and eventually \\emph{fall} into the vortex  (\\emph{surface scenario}).  On the other hand, long-lasting ($>1$ h) vertical magnetic flux tubes anchored beneath the surface might be directly influenced by the vortex action due to some interaction mechanism taking place underneath. Although presenting additional evidence to support this scenario is beyond the scope of this paper, this possibility should not be ruled out.}"
    },
    "1004/1004.3149_arXiv.txt": {
        "abstract": "Within 5-10 years, very long baseline interferometry facilities will be able to observe the ``shadow'' of super-massive black hole candidates. This will allow, for the first time, to test gravity in the strong field regime. In this paper, we study numerically the photon orbits in the $\\delta = 2$ Tomimatsu-Sato space-time. The $\\delta = 2$ Tomimatsu-Sato space-time is a stationary, axisymmetric, and asymptotically flat exact solution of the vacuum Einstein equations. We compare the associated shadow with the one of Kerr black holes. The shape of the shadow in the $\\delta = 2$ Tomimatsu-Sato space-time is oblate and the difference between the two axes can be as high as 6\\% when viewed on the equatorial plane. We argue that future space sub-mm interferometers (e.g. VSOP-3) may distinguish the two cases, and thus are able to test the Cosmic Censorship Conjecture. ",
        "introduction": "Today gravity is relatively well tested in the weak field limit, while little or nothing is known when it becomes strong~\\cite{will}. One of the most outstanding issues is the actual nature of the final product of the gravitational collapse. Currently there are many known ``black hole candidates'', but we do not know if these objects are really the black holes predicted by general relativity and, more in general, if they have an event horizon. At the observational level, the situation may however change soon: the capability of very long baseline interferometry (VLBI) has improved significantly at short wavelength and it is now widely believed that within 5-10 years it will be possible to observe the direct image of the accretion flow around a black hole with a resolution comparable to its event horizon~\\cite{doe1, doe2}. Such observations will allow to test gravity in the strong field regime and investigate the nature of our black hole candidates.   One of the main targets of these VLBI experiments is the observation of the black hole ``shadow''~\\cite{bardeen, melia, rohta}, i.e. a dark area over a bright background. The exact shape and position of the shadow directly depend on the metric of the space-time very close to the massive object, being determined by the innermost photon orbits. In light of this, it is important to understand how to use future observations to test general relativity and what questions can be addressed. For example, in~\\cite{sha1, sha2}, it is shown that the shadow can be used to test the Kerr bound $|a_*| \\le 1$, where $a_*$ is the dimensionless spin parameter. In~\\cite{psaltis}, the authors suggest the test of the no-hair theorem. In this paper we study the shadow in the $\\delta = 2$ Tomimatsu-Sato (TS2) space-time~\\cite{tom1,tom2}, which is {\\it not} a black hole solution. We argue that future sub-mm interferometers will be able to measure the shape of the shadow with an accuracy at the level of 1\\% and thus may distinguish the Kerr from the TS2 space-time.   TS2 is a stationary, axisymmetric, and asymptotically flat exact solution of the vacuum Einstein equations. The Kerr metric describes the gravitational field produced by a spinning mass, while the TS2 one describes the gravitational field of a spinning and deformed mass. Despite such appealing features, the solution violates the Cosmic Censorship Conjecture~\\cite{penrose}, which precludes the formation of naked singularities in the Cauchy development starting from regular initial data. In 4D general relativity, the Kerr space-time is the only stationary and asymptotically flat solution with a regular event horizon~\\cite{carter, robinson}. In presence of naked singularities, there is no uniqueness theorem and the Tomimatsu-Sato space-times are just the simplest extension of the Kerr solution. In general, one can indeed expect that the massive object is described by an infinite number of free parameters~\\cite{m-n}. However, it turns out that the introduction of a larger number of parameters do not change significantly the orbits of the space-time: basically they represent corrections to higher order multipole moments. So, as a first approximation, deviations from the Kerr metric can be studied by considering the Tomimatsu-Sato space-times and here, for the sake of simplicity, we discuss the TS2 solution, even if this does not mean that TS2 is the only alternative to the Kerr space-time when the Cosmic Censorship Conjecture is relaxed. It is however remarkable that TS2 can be obtained from the Neugebauer-Kramer solution representing a superposition of two Kerr black holes, in the limit in which the centers of the two black holes coincide. This fact has been interpreted in~\\cite{kodama} as an indication that TS2 can be a good candidate for the final stage of the gravitational collapse.   The Cosmic Censorship Conjecture is motivated by the fact that space-times with naked singularities present several kinds of pathologies. In the case of the Tomimatsu-Sato space-times, the most suspicious property is the existence of closed time-like curves close to the massive object. Nevertheless, one can even argue that the restriction imposed by such conjecture is more motivated by the limitations of our current knowledge of gravity rather than by true physical reasons. For example, the physical interpretation of space-times with causality violating regions has been recently investigated by a couple of authors in the framework of string theories~\\cite{horava, israel, drukker, gimon}. In these works, it was found that such pathological regions go to a new phase and a domain wall forms. Across the domain wall, the metric is non-differentiable and the expected region with closed time-like curves arises from the naive continuation of the metric ignoring the domain wall. A similar mechanism could work for the Tomimatsu-Sato space-times. A domain wall as a surface of the astrophysical black hole candidates would be also consistent with the non-observation of X-ray bursts from X-ray binaries with a black hole candidate: the accreting matter could be converted to exotic stuff as soon as it hits the surface of the compact object and no thermonuclear reaction would occur~\\cite{abramo}.    The paper is organized as follows. In Sec.~\\ref{s-ts2}, we review the basic properties of the TS2 space-time. In Sec.~\\ref{s-motion}, we present our approach and, in Sec.~\\ref{s-shadow}, we determine the shape and the position of the shadow in TS2 and we compare with the one of Kerr black holes. In Sec.~\\ref{s-disc}, we discuss our results and the possibility of detecting deviations from the Kerr metric with future experiments. Summary and conclusions are reported in Sec.~\\ref{s-con}. Throughout the paper we use natural units $G_N = c = 1$ and metric with signature $(-+++)$. ",
        "conclusions": "}  Within the next decade, very long baseline interferometry techniques will likely be able to image the surrounding environment of some super-massive black hole candidates, with resolution at the level of the black hole event horizon. These experiments will open the possibility of studying strong gravity and the actual nature of the final product of the gravitational collapse. One of the main goal is the observation of the shadow, whose shape and position directly depend on the metric of the space-time around the massive object.   In this paper, we have investigated the shape and the position of the shadow which would be observed in a $\\delta = 2$ Tomimatsu-Sato space-time, and compared with the one expected in the Kerr case. The $\\delta = 2$ Tomimatsu-Sato space-time is a stationary, axisymmetric, and asymptotically flat exact solution of the vacuum Einstein equations which violates the Cosmic Censorship Conjecture. It is not a black hole solution. Despite several significant differences between the two space-times, the associated shadows are quite similar. In the $\\delta = 2$ Tomimatsu-Sato space-time, the shadow is more flattened and an accurate measurement of its shape, at the level of $\\sim 1\\%$, may observe deviations from the Kerr metric, at least for certain values of the spin parameter $q$ and the viewing angle $i$. So, near future ground based experiments or space missions like VSOP-2 will not be able to detect any difference. With more advanced space missions in the future, it is still challenging, but not out of reach.   However, our conclusion that we have to wait for 15-20 years may be too pessimistic. Since around naked singularities the gravitational force can be repulsive~\\cite{sim1}, outflows of gas can be produced~\\cite{sim2}. Similar phenomena are among the science goals of experiments like VSOP-2.   \\begin{figure} \\par \\begin{center} \\includegraphics[height=7cm,angle=0]{a00i90.ps} \\hspace{-2.5cm}\\includegraphics[height=7cm,angle=0]{a07i90.ps} \\\\ \\vspace{.3cm} \\includegraphics[height=7cm,angle=0]{a00i45.ps} \\hspace{-2.5cm} \\includegraphics[height=7cm,angle=0]{a07i45.ps} \\\\ \\vspace{.3cm} \\includegraphics[height=7cm,angle=0]{a00i00.ps} \\hspace{-2.5cm} \\includegraphics[height=7cm,angle=0]{a07i00.ps} \\end{center} \\par \\vspace{-5mm} \\caption{Contour of the shadow in $\\delta = 2$ Tomimatsu-Sato space-time (solid lines) and in Kerr space-time (dotted lines) in the $XY$-plane of the distant observer, for spin parameter $q = 0$ (left column) and $q = 0.7$ (right column). The viewing angle is respectively $i = 90^\\circ$ (top panels), $45^\\circ$ (central panels), and $0^\\circ$ (bottom panels). Horizontal and vertical axis in units $M = 1$.} \\label{f-main} \\end{figure}    \\begin{figure} \\par \\begin{center} \\includegraphics[height=6cm,angle=0]{ration.ps} \\hspace{-0.5cm} \\includegraphics[height=6cm,angle=0]{ratior.ps} \\end{center} \\par \\vspace{-5mm} \\caption{$R_{TS2}/R_{Kerr}$ as a function of the polar coordinate $\\Theta$ for the shadows shown in Fig.~\\ref{f-main}. Left panel: spin parameter $q = 0$. Right panel: spin parameter q = 0.7.} \\label{f-ratio} \\end{figure}"
    },
    "1004/1004.1650_arXiv.txt": {
        "abstract": "Between 1996 July and 2002 April, one or more spacecraft of the interplanetary network detected 787 cosmic gamma-ray bursts that were also detected by the Gamma-Ray Burst Monitor and/or Wide-Field X-Ray Camera experiments aboard the \\it BeppoSAX \\rm spacecraft.  During this period, the network consisted of up to six spacecraft, and using triangulation, the localizations of 475 bursts were obtained. We present the localization data for these events. ",
        "introduction": "Between 1996 July and 2002 April, the Wide Field X-Ray Camera (WFC) and Gamma-Ray Burst Monitor (GRBM) aboard the \\it BeppoSAX \\rm mission detected 62 and 1092 cosmic gamma-ray bursts, respectively, and localized many of them  to accuracies which ranged from arcminutes to tens of degrees \\citep{vetere07,frontera09}; instrument descriptions may be found in \\citet{feroci97}, \\citet{frontera97}, and \\citet{jager97}.  These detections were used to initiate searches through the data of the spacecraft comprising the interplanetary network (IPN). In 475 cases localizations could be obtained by triangulation, and successful multiwavelength counterpart searches were initiated for some of them. The IPN contained between 4 and 6 spacecraft during this period.  They were, in addition to \\it BeppoSAX \\rm: \\it Ulysses, \\rm in heliocentric orbit at distances between 670 and 3180 light-seconds from Earth \\citep{hurley92}; \\it Konus-Wind \\rm, in various orbits up to around 4 light-seconds from Earth \\citep{aptekar95}; \\it HETE-II - FREGATE \\rm , in low Earth orbit \\citep{ricker03,atteia03}; the \\it Near-Earth Asteroid Rendezvous \\rm mission (NEAR), at distances up to 1300 light-seconds from Earth \\citep{trombka99}; \\it Mars Odyssey\\rm, launched in 2001 April and in orbit around Mars starting in 2001 October, up to 1250 light-seconds from Earth \\citep{hurley06a}; the \\it Compton Gamma-Ray Observatory \\rm (the Burst and Transient Source Experiment, BATSE - \\citet{fishman92}; and \\rm the \\it Ramaty High Energy Solar Spectroscopic Imager \\rm (RHESSI) both in low Earth orbit \\citep{smith02}. Their timelines are presented in figure 1.  In this paper, we present the localization data obtained by the IPN for these bursts.  At least three other spacecraft recorded GRB detections during this period, although they were not used for triangulation and therefore were not, strictly speaking, part of the IPN.  The \\it Rossi X-Ray Timing Explorer \\rm (RXTE)  All Sky Monitor detected and localized some \\it BeppoSAX \\rm bursts \\citep{smith99}.  It operated in the low energy X-ray range, where the light curves of gamma-ray bursts differ significantly from the high energy range where the other IPN instruments operate.  The \\it Defense Meteorological Satellite Program \\rm (DMSP) \\citep{terrell96, terrell98, terrell04} and the \\it Stretched Rohini Satellite Series \\rm (SROSS) \\citep{marar94} spacecraft detected, but did not localize bursts. ",
        "conclusions": "This is the eleventh in a continuing series of IPN catalogs \\citep{hurley99a, hurley99b, hurley00a, hurley00b, hurley00c, hurley05, hurley06b, laros97, laros98, hurley09}; the localization data for all of them can be found in electronic form at the IPN website. The IPN is, in effect, a full-time, all-sky monitor, when the duty cycles and viewing constraints of all its instruments are considered.  Its fluence and flux thresholds for 50\\% detection efficiency are about $\\rm 6 \\times 10^{-7} erg \\, cm^{-2}\\, and \\, 1 \\, photon \\, cm^{-2}\\, s^{-1}$, \\,respectively.  Over the \\it BeppoSAX \\rm mission, 787 bursts were detected by the GRBM and/or the WFC and at least one other IPN instrument and 475 of them could be localized to some extent by triangulation.  The more precise and/or rapid localizations were announced in over 50 IAU and GCN Circulars (in 1997, and in 1998 -- 2002, respectively), resulting in multiwavelength counterpart searches. Regardless of precision and speed of the localizations, however, burst data such as this are useful for numerous studies, such as searching for indications of activity from previously unknown soft gamma repeaters, associating supernovae with bursts, or searching for neutrino and gravitational radiation associated with bursts."
    },
    "1004/1004.5379_arXiv.txt": {
        "abstract": "\\noindent In this Paper we report the discovery of \\src; a peculiar X-ray source with a position 3.6$\\pm$0.2\\arcsec\\,off-nuclear from an SDSS DR7 z=0.0447 galaxy. The 3.6\\arcsec\\, offset corresponds to 3.2 kpc at the distance of the galaxy.  The 0.3--8 keV X-ray flux of this source is 5$\\times 10^{-14}$ \\flx\\, and its 0.3--8 keV luminosity is 2.2$\\times 10^{41}$ \\lum\\, (2.7$\\times 10^{41}$ \\lum; 0.5--10 keV) assuming the source belongs to the associated galaxy. We find a candidate optical counterpart in archival HST/ACS g$^\\prime$--band observations of the field containing the galaxy obtained on June 16, 2003. The observed magnitude of g$^\\prime=26.4 \\pm 0.1$ corresponds to an absolute magnitude of $-10.1$. We discuss the possible nature of the X-ray source and its associated candidate optical counterpart and conclude that the source is either a very blue type IIn supernova, a ULX with a very bright optical counterpart or a recoiling super-massive black hole. ",
        "introduction": "Off-nuclear ultra--luminous X--ray sources (ULXs) have been found in different galaxies. It has been found that many of the ULXs are associated with active star--forming regions (e.g.~in the Antennae, \\citealt{2003ApJ...591..843F}), optical emission line nebulae (e.g.~\\citealt{2006IAUS..230..293P}) and radio halos (e.g.~\\citealt{2005ApJ...623L.109M}; \\citealt{2007ApJ...666...79L}).  In addition, six off-nuclear 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}; in NGC~2276, \\citealt{2004ApJ...604..653D}; ESO~243-49, \\citealt{2009Natur.460...73F}; in NGC~5775, \\citealt{2008MNRAS.390...59L} and M101~ULX-1, \\citealt{2004ApJ...617L..49K}).  They are called hyper-luminous X-ray sources. The optical counterpart has been found and studied in only a sub-sample of these (M101~ULX-1, \\citealt{2004AJ....128.2783K}, \\citealt{2009ApJ...704.1628L}; ESO~243-49, \\citealt{2009arXiv0910.1356S}; and possibly M~82~X-1, \\citealt{2004MNRAS.348L..28K}).  An important question in the work on the ULXs is whether the emission is isotropic or not. If beaming is significant (\\citealt{2001ApJ...552L.109K}) an intermediate--mass black hole (IMBH) or super-massive black hole (SMBH) is not needed to explain sources with luminosities up to $\\approx 10^{40}$\\lum. For several sources, ionizing luminosities derived from the optical emission lines and the resolved radio bubbles indicate that beaming is not important (\\citealt{2003RMxAC..15..197P}; \\citealt{2005ApJ...623L.109M}; \\citealt{2007ApJ...666...79L}) and hence these sources are good candidates for IMBHs. For the sample of hyper-luminous ULXs, their X-ray luminosity seems too high for a stellar mass black hole even in the presence of some beaming. For their classification there are currently three possible scenarios: very bright type IIn supernovae, IMBHs and recoiling SMBHs. We will briefly introduce these options.  The X-ray brightest supernovae are of type IIn (\\citealt{2003LNP...598...91I}). Therefore, in principle, they could be responsible for a sub-sample of the very bright ULXs. Time variability in the X-ray as well as in the optical band can be used to constrain this possibility.  In the cold dark matter (${\\mathrm \\Lambda CDM}$) cosmological scenario current (z=0) galaxies are the product of hierarchical mergers of smaller galaxies. These smaller building blocks also host black holes in their centers (\\citealt{1995ARA&A..33..581K}). Evidence for this comes from the observed M--$\\sigma$ relation (\\citealt{1998AJ....115.2285M}; \\citealt{2000ApJ...539L...9F}). Furthermore, several so called dual Active Galactic Nuclei (AGNs) have been found (cf.~\\citealt{2009ApJ...702L..82C} and references therein). After a galaxy merger, the two black holes will eventually merge as well (see \\citealt{1980Natur.287..307B}).  Recent fully relativistic numerical simulations allow for the calculation of the linear angular momentum that is transported by gravitational wave radiation during the final plunge in the black hole -- black hole merger (\\citealt{2005PhRvL..95l1101P}). The transport of linear angular momentum acts as a kick on the newly formed black hole merger product: gravitational wave recoil.  This effect had been estimated before by for instance \\citet{1973ApJ...183..657B} and \\citet{1989ComAp..14..165R}. When a black hole merger takes place in the presence of an accretion disc, the recoiling black hole will take along the part of the nuclear star cluster and accretion disc that falls within its gravitational influence area (${\\rm R_{infl}\\approx 0.3 M_8 / v_{1000}^2\\,pc}$; where $M_8$ is the mass in units of $10^8$\\msun\\,and $v_{1000}$ is the velocity in units of 1000 km s$^{-1}$; Bonning, Shields \\& Salviander 2007). Binary SMBHs surrounded by an accretion disc will have emptied the inner portion leaving a gap in the disc. Upon a kick this gap will refill on a short timescale (\\citealt{2007PhRvL..99d1103L}) and AGN activity will resume.  Subsequently, such recoiling black holes may become visible as off--nuclear AGNs (\\citealt{2007ApJ...666L..13B}; \\citealt{2008ApJ...687L..57V}; \\citealt{2009ApJ...691.1050F}). The disc mass will allow for the AGN activity to last tens of millions of years (cf.~\\citealt{2007ApJ...666L..13B}).   The source SDSS~J092712.65+294344.0 has been proposed as a recoiling massive black hole (\\citealt{2008ApJ...678L..81K}). However, different interpretations as a binary black hole (\\citealt{2009ApJ...697..288B}; \\citealt{2009MNRAS.398L..73D}) or a chance alignment with a more distant source in the same cluster have also been proposed (\\citealt{2009ApJ...696.1367S}). Similarly, SDSS~J105041.35+345631.3 and SDSS~J153636.22+044127.0 have been suggested as recoiling SMBH candidates (\\citealt{2009ApJ...707..936S}; \\citealt{2009Natur.458...53B}, respectively). The main difference between the IMBH and the SMBH scenario for the hyper-luminous sources lies in their optical properties. The recoiling SMBHs should carry with it the broad line region. Therefore, they should contain broad lines in their optical spectra whereas this is not the case for the IMBH scenario.  In this Manuscript we discuss the nature of another hyper-luminous off-nuclear X-ray source: \\src. ",
        "conclusions": "We have located a peculiar X-ray source with a position 3.6\\arcsec\\, off-nuclear from an SDSS DR7 z=0.0447 galaxy (see Fig.~\\ref{ima} and \\ref{hstima}). The 3.6\\arcsec\\, corresponds to 3.2 kpc at the distance of the galaxy. The presence of another X-ray source in the same observation coinciding with a bright optical point source, allows us to register the \\chan\\, observation, reducing the astrometric uncertainty in the position of the off-nuclear AGN-candidate to less than 0.2\\arcsec, making the 3.6\\arcsec\\, offset highly significant. The 0.3--8 keV X-ray flux of this source is 5.4$\\times 10^{-14}$ \\flx\\, and its 0.3--8 keV (0.5--10 keV) luminosity is 2.2$\\times 10^{41}$ \\lum\\, (2.7$\\times 10^{41}$ \\lum\\,) given the redshift of the galaxy (z=0.0447, d=182.6 Mpc). We find a candidate optical counterpart in archival HST/ACS g$^\\prime$--band observations of the field containing the galaxy obtained on June 16, 2003. The observed magnitude of g$^\\prime=26.4 \\pm 0.1$ corresponds to an absolute magnitude of -10.1 taking a foreground extinction from our Galaxy of ${\\rm A_{g^\\prime}=0.18}$ magnitudes into account. These findings make this source an unusually bright ULX, a very bright supernova, a recoiling black hole or a background AGN (with higher luminosity). We discuss these possibilities below.  If this X-ray source is due to a supernova, the 0.5--10 keV X-ray luminosity is among the highest measured for supernovae. The X-ray brightest supernovae are of type IIn, this implies that the maximum optical g$^\\prime$-band magnitude was $\\approx$17.3 (with mean absolute blue magnitude of -19.15 type IIn are also optically among the brightest supernovae; ~\\citealt{2002AJ....123..745R}). Now, more than 6 years after the initial HST/ACS observations, the source magnitude will have changed considerably. The exact way and amount are difficult to predict as many scenarios are possible since the optical lightcurves of type~IIn supernovae are heterogeneous and we do not know the explosion date.  The supernova could have exploded in between the serendipitous HST/ACS and the X-ray observation or before the HST/ACS observation. The X-ray luminosity of the X-ray bright type~IIn supernovae peaks around 400-1000 days after the explosion making it impossible to discriminate between the different scenarios on the basis of the single epoch X-ray and optical observations. The g$^\\prime - z^\\prime \\approxlt 0.7$ colour of the optical counterpart candidate is blue for a type IIn supernova origin of the counterpart (cf.~\\citealt{2008PZ.....28....6T} for typical type IIn colours). If the supernova occurred in between the HST/ACS and the \\chan\\, observation the HST/ACS images are of the supernova progenitor star (cluster).  The SDSS imaging data of this galaxy was taken a few weeks before the HST data and thus cannot help us deciding between these scenarios.  In the ULX case the majority of known optical counterparts have blue colours and they are often embedded in an ionized optical nebula (\\citealt{2008A&A...486..151G} and references therein). In addition, the counterparts are usually part of small young star clusters or OB associations (e.g.~\\citealt{2002ApJ...577..726Z}). The FWHM of the ACS image at the distance of the galaxy corresponds to more than 100 pc hence star clusters would appear as unresolved point sources. The ULX counterparts currently known would be below the detection threshold of the HST/ACS images as their absolute optical magnitude ${\\rm M_V}$ is between $-4$ and $-9$ (cf.~\\citealt{2008MNRAS.387...73R}) while the source present in the HST/ACS observation has an absolute $g^\\prime$ magnitude of $-10.1$.  This absolute magnitude is similar to the absolute magnitude of the counterpart of ESO~243-49 proposed by \\citet{2009arXiv0910.1356S}. The blue colours of the candidate counterpart to \\src\\, make it hard to reconcile \\src\\,with the scenario of an IMBH in a Globular Cluster (cf.~the colours of Globular Clusters in \\citealt{1996yCat.7195....0H}). Comparing the limit on the $g^\\prime - z^\\prime <0.7$ colour with that calculated using the Padua evolution tracks (\\citealt{2008A&A...482..883M}; \\citealt{2008PASP..120..583G}) we find that the blue colour is consistent with an IMBH in a young, massive, star cluster.  In the recoiling black hole case the optical counterpart is probably related to the accretion disc, broad line region and the nuclear cluster that are retained. The absolute magnitude of the nuclear region in AGNs varies by more than one order of magnitude, but the faintest absolute magnitudes are ${\\rm M_B}=-9.8$ (NGC~4395; \\citealt{1989ApJ...342L..11F}) and $-11.6$ for NGC~3031/M81 reported by \\citet{1997ApJS..112..391H} in line with our finding. At an off-nuclear distance of 3.2 kpc, a recoil velocity of $\\sim$300 km s$^{-1}$ is implied if the lifetime of the AGN activity is limited to $\\sim10^7$ years. The hard X-ray spectral power-law index of 0.9$\\pm0.3$ that we find implies that there is additional extinction beyond the Galactic value if the intrinsic source spectrum is compatible with that of an AGN.  One has to worry about chance alignments on the sky. At an 0.5--10 keV flux of $5.4\\times 10^{-14}$ \\flx\\,the chance to find a background AGN in a circle with radius of 3.6\\arcsec\\, is very small. From the 0.5--10 keV $\\log$N -- $\\log$S relation found by \\citet{2008A&A...492...51M} there are approximately 76 AGN per square degree.  Using this we derive that in a circle with radius of 3.6\\arcsec\\, on average 2$\\times10^{-4}$ AGN is found. However, we searched 17 X-ray sources in galaxies with fluxes equal to or above 5$\\times 10^{-14}$\\flx, so considering this number of trials, leads to a probability of 4$\\times 10^{-3}$ that the source is a background AGN.  The X-ray spectrum of the source argues against a background AGN. The hard X-ray spectrum of the source together with an 0.5-2 keV flux of $9\\pm3\\times 10^{-15}$\\flx~make that the source falls well below the trend between power-law index and source 0.5-2 keV flux found by \\citet{2005A&A...444...79M} in their study of the Lockman Hole IV (see their figure 4). The X-ray to optical flux ratio in the $g^\\prime$-band is 87 and in the $z^\\prime$-band this is larger than 84, whereas most \\chan\\, selected AGNs have values in the $R$-band below 10 (e.g.~\\citealt{2003AJ....126..632B}; \\citealt{2009ApJS..180..102L}). If the optical counterpart is unrelated to the X-ray source the X-ray to optical flux ratio is larger still. However, if there is additional absorption local to the AGN, the X-ray to optical flux ratio has been observed to be larger than 90 (\\citealt{2005MNRAS.358..693C}), although there are very few if any of these sources known at the X-ray flux level of \\src~(cf.~\\citealt{2005A&A...444...79M}). If the optical source is related to the X-ray source but due to a background AGN shining through the disc of the foreground galaxy that adds a column of 5$\\times 10^{21} {\\rm cm^{-2}}$, the $g^\\prime-z^\\prime$ colour limit would become $<-1.3$ which is uncharacteristically blue for an AGN (cf.~the $g^\\prime-z^\\prime$ colours of ROSAT - SDSS AGN presented in \\citealt{2007AJ....133..313A} have a mean $g^\\prime-z^\\prime$ colour of $\\approx 0.7\\pm 0.5$).  We conclude that \\src\\, is a strong candidate for a recoiling SMBH, a bright ULX with a bright optical counterpart or a very blue type IIn supernova."
    },
    "1004/1004.2559_arXiv.txt": {
        "abstract": "We present deep, accurate and homogeneous multiband optical (\\U,\\B,\\V,\\I) photometry of the Carina dwarf spheroidal galaxy, based on more than 4,000 individual CCD images from three different ground-based telescopes.  Special attention was given to the photometric calibration, and the precision for the \\B, \\V, and \\I bands is generally better than 0.01 mag.  We have performed detailed comparisons in the \\V, \\bmv\\ and \\V, \\bmi\\ color-magnitude diagrams (CMDs) between Carina and three old, metal-poor Galactic Globular Clusters (GGCs, M53, M55, M79).  We find that only the more metal-poor GCs (M55, [Fe/H]=--1.85; M53, [Fe/H]=--2.02 dex) provide a good match with the Carina giant branch.  We have performed a similar comparison in the \\V, \\vmi\\ CMD with three intermediate-age clusters (IACs) of the Small Magellanic Cloud (Kron~3, NGC~339, Lindsay~38). We find that the color extent of the SGB of the two more metal-rich IACs (Kron~3, [Fe/H]=--1.08; NGC339, [Fe/H]=--1.36 dex) is smaller than the range among Carina's intermediate-age stars. Moreover, the slope of the RGB of these two IACs is shallower than the slope of the Carina RGB. However, the ridge line of the more metal-poor IAC (Lindsay~38, [Fe/H]=--1.59 dex) agrees quite well with the Carina intermediate-age stars. These findings indicate that Carina's old stellar population is metal-poor and appears to have a limited spread in metallicity ($\\Delta$[Fe/H]=0.2--0.3 dex). The Carina's intermediate-age stellar population can hardly be more metal-rich than Lindsay~38 and its spread in metallicity also appears modest.  We also find that the synthetic CMD constructed assuming a metallicity spread of 0.5~dex for the intermediate-age stellar component predicts evolutionary features not supported by observations.  In particular, red clump stars should attain colors that are redder than red giant stars, but this is not seen.  The above results are at odds with recent spectroscopic investigations suggesting that Carina stars cover a broad range in metallicity ($\\Delta$[Fe/H]$\\sim$1--2 dex).  We also present a new method to estimate the metallicity of complex stellar systems using the difference in color between the red clump and the middle of the RR Lyrae instability strip.  The observed colors of Carina's evolved stars indicate a metallicity of [Fe/H]=--1.70$\\pm$0.19 dex, which agrees quite well with spectroscopic measurements. ",
        "introduction": "Dwarf galaxies play a fundamental role in several astrophysical problems. Current cosmological simulations predict dwarf satellite populations significantly larger than the number of dwarfs observed near giant spirals like the Milky Way and M31. This discrepancy is called the ``missing satellite problem'' and is a challenge to the currently most popular cosmological model: the $\\Lambda$ Cold Dark Matter paradigm \\citep{kly99,moore06,mandau08}. However, it has not yet been established whether this discrepancy is due to limitations in the theoretical modeling or to observational bias at the faint end of the dwarf galaxy luminosity function \\citep{kli09,kop09,krav10}.  Moreover, detailed analyses of the distribution of dwarf galaxies in multi-dimensional parameter space indicate that they follow very tight relations. In particular, \\citet{prada02} suggested that Local Group (LG) dwarf galaxies show strong correlations (Fundamental Line - FL) among mass-to-light (M/L) ratio, surface brightness, and metallicity. They explained the correlation between M/L ratio and metallicity using a simple chemical enrichment model in which the hot metal-enriched gas, at the end of the star-formation epoch, is lost via continuous galactic winds \\citep{maclow99}. In a recent investigation \\citet{woo09} reached similar conclusion, i.e., LG dwarf galaxies follow a one-parameter relation driven by the total stellar mass. They also found that dwarf spheroidals (dSphs) appear to be, at fixed stellar mass, systematically more metal-rich than dwarf irregulars (dIs, see also \\citet{mateo98}). Futhermore, the FL of dIs in a five-dimensional parameter (total space mass, surface brightness, rotation velocity, metallicity, star formation rate) is linear and tight, and extends the scaling relations of giant late-type galaxies into the low-mass regime. On the other hand, the FL of dSphs as defined in a four-dimensional parameter space (total mass, surface brightness, rotation velocity, metallicity) is also linear and very tight, but it does {\\it not\\/} lie on an extrapolation of  the scaling relations of giant early-type galaxies. Iron and heavy-element abundances, the spread in these quantities, and the star-formation history (SFH) of LG dwarf galaxies, are crucial observational constraints on the physical mechanisms responsible for the scaling relations described above, and on their implications for different cosmological models \\citep{orb08}. Spectroscopic investigations based on high-resolution spectra of bright red giants (RGs) in several LG dSphs indicate a large spread in iron abundance \\citep{shet01,shet03}. This is also supported by metallicity estimates from the CaII triplet method \\citep{bat08}. In this context, the Carina dSph is particularly relevant: ({\\em i}) It is relatively close and its central density is modest ($\\rho$=0.17$M_\\odot/pc^3$, \\citet{mateo98}) so it is easily resolved with ground-based telescopes \\citep{mou83}. ({\\em ii}) It is {\\it the\\/} LG dSph most clearly showing multiple, widely separated star formation episodes(\\citet{mig90,mig97}; \\citet{smec94,smec96}; \\citet{hur98}; \\citet{her00}; \\citet{har01}; \\citet{dol02}; \\citet{riz03}) with a tail that might extend to less than one Gyr ago (\\citet{mon03}). ({\\em iii}) It is also a benchmark to constrain the pulsation properties of old (RR Lyrae, \\citet{saha86}) and intermediate-age (dwarf Cepheids, \\citet{matal98}; Anomalous Cepheids, \\citet{ora03}) variable stars. ({\\em iv}) High-resolution spectra are available for a sample of ten bright RGs. The mean metallicity  is [Fe/H]=--1.69 with $\\sigma$=0.51 dex \\citep{koch08}. Independent measurements by \\citet{shet03}, using high resolution spectra for five bright RGs, provided a mean metallicity of [Fe/H]=--1.64 and $\\sigma$=0.2 dex. ({\\em iv}) Medium-resolution calcium triplet measurements are also available for a large sample (437) of RG stars \\citep{koch06}; their metallicity distribution shows a peak at \\hbox{[Fe/H]=--1.72$\\pm$0.01} (metallicity scale by Carretta \\& Gratton 1997) and the metallicity ranges from $\\sim$--2.5 to $\\sim$--0.5 dex. Using the same spectra, but different selection criteria (364 candidate Carina stars), Helmi et al.\\ (2006) found the same peak in the metallicity distribution ([Fe/H]=--1.7$\\pm$0.1 dex) and metallicities ranging from $\\sim$--2.3 to $\\sim$--1.3 dex (see their Fig.~2).  Deep, accurate photometric investigations indicate an old (13 Gyr) stellar population and several distinct intermediate-age populations (2--6 Gyr) together with a blue plume of younger stars. The occurrence of both old and intermediate-age populations is also indicated by two distinct samples of core helium burning stars:  an old, wide HB and a separate red clump (RC). The shape and thickness of the RGB ($\\Delta$\\bmv$\\sim$0.02 for 20.5$\\lesssim$$V$$\\lesssim$20.75 mag) suggest that the old and intermediate-age stellar populations have a minimal spread in chemical composition. This seems inconsistent with the spectroscopic observations that indicate a metallicity distribution possibly reaching extreme values near --3.1 and +0.1~dex \\citep[scale]{zinn84}, or near --2.8 and --0.2~dex \\citep[scale]{car97}.  Despite the criticisms raised by \\citet{koch06}, conclusions based on purely photometric indicators are hampered by only two major factors: ({\\em i}) A decrease in photometric precision when moving from bright to faint RG stars might produce a spread in color mimicking a spread in metal content. In particular, Carina's stellar density is very low, requiring that photometry be obtained with multi-chip CCD cameras or with many different pointings of a single CCD. A high degree of consistency in the photometric zero-points across the field is essential to avoid spurious broadening of the RG and Main Sequence (MS) loci. ({\\em ii}) Due to its low surface density, large extent, and low Galactic latitude (--22\\deg), the Carina sample is contaminated by foreground field stars. Their colors, unfortunately, are similar to the Carina RGB, once again mimicking a broadening in color due to a spread in age/metallicity.  Our group is involved on a long-term investigation of the stellar populations in the Carina dSph. In particular, we have assembled our own and archival {\\it UBVI\\/} images covering the entire body of the galaxy. In this investigation we focus on comparing Carina with old and intermediate-age template clusters. We also present a new method for estimating the metal content of complex stellar systems  from the difference in color between the middle of the RR Lyrae instability strip and the peak of RC stars. ",
        "conclusions": "For a robust discrimination of candidate Carina members from foreground field stars and background galaxies we adopted the \\umv,\\bmi\\ color-color diagram (C-CD). The top panel of Fig.~1 shows that this is effective for distinguishing candidate Carina and field stars. Carina RGs (\\umv$\\ge$0.6, \\bmi$\\ge$1.4) attain, at fixed \\bmi, bluer \\umv\\ colors than field stars. The difference is mainly caused by a difference in the mean metallicity between Carina and the more metal-rich Galactic disk stars, but the difference in surface gravity between galaxy giants and foreground dwarfs may also play a role.  In contrast, hot horizontal-branch (HB) stars and relatively young MS stars in Carina (\\bmi$\\lesssim$0.4) are relatively free from field star contamination, since they are bluer than the thick disk and halo turnoffs; compact background galaxies represent the principal source of contamination among these bluer objects.  Fortunately, this C-CD is also a good diagnostic to separate candidate Carina stars from blue background galaxies since the latter show, at fixed \\bmi, systematically bluer \\umv\\ colors than stars of any luminosity class or metallicity (apart from white dwarfs, which are not an issue here).  Based on the above empirical evidence, we performed a series of tests to identify candidate Carina stars using various color-magnitude diagrams (CMDs) together with the C-CD.  The solid-outlined box in the top panel of Fig.~1 shows the final selection, while the bottom panels show the $V$,\\bmv\\ CMD of the total sample (left), the candidate Carina stars (middle), and the candidate non-member objects (right). Note that after the selection there remain some field stars with colors similar to Carina's RGB ( $V$$\\le$21.5, \\bmv$\\sim$0.4-0.6; middle panel). However, RG stars are distributed along a very narrow sequence over the entire magnitude range (18$\\le$$V$$\\le$23), and the number of field stars lying near that RGB in both color and magnitude is appreciably reduced. The same statement applies to the intermediate-age red clump (RC: $V$$\\sim$20.5, \\bmv$\\sim$0.65) and to the old HB ($V$$\\sim$20.75, \\bmv$\\sim$0.65). These evolutionary phases are characterized by very narrow distributions in either magnitude or color or both. The improvement compared to data already available in the literature is due to smaller photometric uncertainties ($\\sigma_{\\bmv}$$\\le$0.02 for $V$$\\le$21, see error bars in Fig.~1). Plain evolutionary arguments \\citep{mon03,zocc03} suggest that small dispersions in magnitude and color in these evolutionary phases imply a negligible spread in chemical composition. This finding supports the results based on B,V,I photometry of Carina RG stars obtained by \\citet{riz03}. The presence of a few stars in the right panel with magnitudes and colors typical of red HB ($V$$\\sim$20.75, \\bmv$\\sim$0.4) and faint subgiant branch (SGB) stars ($V$$\\sim$23, \\bmv$\\sim$0.5) demonstrates that our C-CD selection is imperfect---especially at intermediate colors---and some real Carina members have probably been erroneously rejected. However, in these photometric selections it is more important to keep probable non-members out than to keep possible members in \\citep{cal09}.  To provide robust constraints on Carina's stellar content, we compared it with stellar systems characterized by different ages and metal abundances. In particular, to constrain the old stellar population we selected three globular clusters (GCs)---M79 = \\ngc{1904}, M55 = \\ngc{6809}, and M53 = \\ngc{5024}---with metal abundances ranging from \\hbox{--1.64}$\\pm$0.15 to --2.02$\\pm$0.15, and low foreground reddenings (E(\\bmv)$\\le$0.09, see Table~1). The optical {\\it BVI\\/} photometry of these GCs is also on the \\citet{stet05} system. The accuracy of the photometry is certainly better than our knowledge of the foreground reddening toward these clusters and Carina.  The $V$, \\bmv\\ (left panels) and $V$, \\bmi\\ (middle panels) CMDs of the selected GCs are shown in Fig.~2, together with their individual ridge lines (red curves). We also selected three intermediate-age clusters (IACs) belonging to the Small Magellanic Cloud (SMC)---Kron~3, \\ngc{339}, Lindsay~38---with similar ages and iron contents ranging from --1.08$\\pm$0.12 to --1.59$\\pm$0.10 dex \\citep[see Table~1]{glatt09}. The photometry for these clusters is from ACS@HST data transformed into the $V$, $I$ Johnson-Cousins bands using relations by \\citet{sir05}. These clusters are {\\it not\\/} included in the homogenous photometry project (the original sources of the data are listed in Table~1), so we are unable to independently confirm the precision of the photometry or the accuracy of the calibrations.  Data plotted in the right panels of Fig.~2 display the $V$,\\vmi\\ CMDs for these clusters together with their ridge lines and the outlines of their RCs.  To compare the GCs with the old stellar populations in Carina, we arbitrarially shifted their ridge lines in magnitude and color to provide a good fit with the SGB (23.25$\\lesssim$$V$$\\lesssim$23.75, 0.45$\\le$\\bmv$\\le$0.60) in Carina. Then we inferred the difference in distance and in reddening between Carina ($\\mu$=20.15, E(\\bmv)=0.03; \\citet{ora03,piet09}) and the individual GCs. The distances and the cluster reddenings we found following this approach agree, within the errors, with similar estimates available in the literature (see columns 4 and 5 in Table~1). The top and the middle left panels of Fig.~3 show that the GCs M79 and M55 provide a poor match to Carina, since hot HB stars are systematically brighter than Carina HB stars. On the other hand, both the GB and the HB of the more metal-poor GC (M53) match Carina quite well. The same outcomes apply if the comparison is performed using the $V$, \\bmi\\ CMDs (see the middle panels in Fig.~3). The extinction in the {\\it BVI\\/} bands was estimated using the empirical relations provided by \\citet{mccall04}.  The same approach was followed to compare the IACs with Carina stellar populations. The ridge lines were shifted in magnitude and color to provide a good fit to the lower envelope of intermediate-age SGB (22.5$\\lesssim$$V$$\\lesssim$23.0, 0.45$\\lesssim$\\vmi$\\lesssim$0.80) stars in Carina. The distances that we estimate are larger than those estimated by \\citet{glatt09} (see Table~1) using isochrones. However, the current distances agree, within the errors, with SMC distances available in the literature \\citep{roman09}. The same outcome applies for the cluster reddenings (see the right panels in Fig.~3). This comparison between Carina and the IACs revealed the following: ({\\em i}) the color extents of the SGBs of the the two more metal-rich clusters (Kron~3 and \\ngc{339}; i.e., the length of the horizontal sequence between the upper MS and the lower GB) are smaller than the color range covered by Carina's intermediate-age SGB stars; ({\\em ii}) the slope of the RGB in the two more metal-rich clusters (Kron~3, \\ngc{339}) is shallower than the observed slope of Carina's RGB; ({\\em iii}) the ridge line of Lindsay~38 is a fairly good match to Carina's intermediate-age population. ({\\em iv}) the RC contour of Kron~3 agrees with Carina's RC, while for the other two clusters the clumps are either slightly redder and brighter (\\ngc{339}) or marginally fainter (Lindsay~38).  These comparisons indicate that Carina's old population is relatively metal-poor, and its spread in metallicity is modest: at most 0.2--0.3 dex. Carina's intermediate-age population seems to have the same metallicity as Lindsay~38, or perhaps slightly lower. The spread in metal abundance in this stellar component also seems very limited, since the RGBs of the two more metal-rich IACs are considerably redder than the red envelope of Carina's RGB. The lack of accurate photometry for metal-poor ([Fe/H]$<$--1.6 dex) IACs with well populated RGBs prevents us from providing more firm constraints on Carina's intermediate-age population. These findings do not depend on the adopted metallicity scale, and if we adopt instead iron abundances on the \\citet{car09} metallicity scale (see column 3 in Table~1), the metallicity range inferred for Carina is very similar.  For a new estimate of the metal abundance of Carina we decided to use a new indicator. Evolutionary prescriptions for old and intermediate-age stellar structures indicate that the difference in color between the center of the RR Lyrae instability strip and the RC stars is strongly correlated with iron abundance. The correlation is caused by the mild dependence of the color of the instability strip on iron abundance and by the significant dependence of the color of RC stars when moving from metal-poor to metal-rich compositions (see Fig.~4). We estimated the predicted difference in color between RC stars (at the beginning of central helium burning) and the center of RR Lyrae strip ($\\log T_e$=3.85) from the large set of scaled-solar evolutionary models \\citep{pie04} available in the BaSTI database\\footnote{Evolutionary models can be download from http://www.oa-teramo.inaf.it/BASTI}. We adopted evolutionary prescriptions that cover a wide range of iron abundances (--2.27$\\le$[Fe/H]$\\le$+0.06 dex), and an age range of 1--6 Gyr for the intermediate-age population and 13 Gyr for the old population. Note that the instability strip is minimally affected by cluster age when moving from 7--8 to 12--13 Gyr. In this context two robust evolutionary predictions need to be underlined. ({\\em i}) When moving from metal-poor to metal-rich compositions, the intermediate-age RC becomes  redder than the old RGB for [Fe/H]=--1.79 dex (see the top panel of Fig.~4). Current evolutionary models predict that a change in metallicity from [Fe/H]=--1.96 to [Fe/H]=--1.49 produces a change in $V$ magnitude of RC stars of 0.03 mag, but a change of 0.12 mag in the \\bmv\\ color. ({\\em ii}) The spread in magnitude of HB stars at the mean color of the RR Lyrae strip increases, as expected, with the spread in metallicity. Note that a change in metallicity from [Fe/H]=--1.96 to [Fe/H]=--1.49 implies a change in $V$ magnitude of HB stars of 0.10 mag, but a change smaller than 0.01 mag in the \\bmv\\ color of the instability strip (see the bottom panel of Fig.~4).  Finally, we performed a linear regression between metallicity and the RC--HB difference in three different colors:  \\begin{tabular}{r l} $[Fe/H]$ =&$ (-3.02\\pm 0.03)+(3.89\\pm 0.05)\\Delta(\\bmv)_{HB}^{RC}$ \\\\ $[Fe/H]$ =&$ (-3.47\\pm 0.04)+(2.37\\pm 0.04)\\Delta(\\bmi)_{HB}^{RC}$ \\\\ $[Fe/H]$ =&$ (-4.13\\pm 0.08)+(6.02\\pm 0.15)\\Delta(\\vmi)_{HB}^{RC}$ \\end{tabular}  The above relations (see also right panels in Fig.~5) indicate that the difference in the \\bmi\\ color is significantly more sensitive to [Fe/H] than the \\bmv\\ color. The \\vmi\\ is the least sensitive of all and shows a stronger nonlinear trend in the metal-poor ([Fe/H]$\\le$--1.5) regime. We estimated the difference in color between the RC and the middle of the RR Lyrae instability strip (see the left panels in Fig.~5). We found $\\Delta$\\bmv$^{RC}_{HB}$=0.37$\\pm$0.05, $\\Delta$\\bmi$^{RC}_{HB}$=0.75$\\pm$0.05 and $\\Delta$\\vmi$^{RC}_{HB}$=0.40$\\pm$0.05 mag, where the errors account for uncertainties in the photometry and in the estimate of the mean color of the RR Lyrae strip. We applied the above relations and we found the following mean metallicities: [Fe/H]=--1.60$\\pm$0.24($\\Delta$\\bmv$^{RC}_{HB}$), [Fe/H]=--1.70$\\pm$0.19($\\Delta$\\bmi$^{RC}_{HB}$) and [Fe/H]=--1.76$\\pm$0.44($\\Delta$\\vmi$^{RC}_{HB}$). We adopted the estimate based on \\bmi, since it is the most precise.  As a further test to constrain the spread in metallicity of the Carina sub-populations, we computed a series of synthetic CMDs changing the spread both in age and in metal content. The set of evolutionary models used for these numerical simulations are the same used to derive the new metallicity indicator (Pietrinferni et al.\\ 2004). The code adopted to compute the synthetic CMDs has already been discussed by Pietrinferni et al.\\ (2004) and by Cordier et al.\\ (2007). We only briefly mention the initial parameters adopted to compute the synthetic CMDs. The Initial Mass Function was modeled using a power law with a Salpeter exponent ($\\alpha$=-2.35), and we assumed a fraction of unresolved binaries of the order of 10\\% with a minimum mass ratio for the binary systems equal to 0.7. The synthetic CMDs were constructed assuming two star formation events, with a mean age of 2.5--4 Gyr for the intermediate-mass and of 11--13 Gyr for the low-mass sub-population (Monelli et al.\\ 2003). We also assumed that the two star formation episodes include the same fraction of stars. Moreover, we assumed mean metallicities of [Fe/H]=--1.5 and of --2.0 dex for the intermediate-age and the old burst, respectively. In accordance with the above findings, we assumed the same spread in metallicity---$\\pm$0.1 dex---for the intermediate and the old star formation episode. Finally, to mimic actual observations, synthetic photometric errors were added using a Gaussian in the \\B\\ and \\V\\ bands with a dispersion equal to 0.02 mag. The above assumptions adopted to simulate the Carina observations are crude, but a detailed analysis is beyond the aims of this investigation. We plan to study the Carina star formation history in a forthcoming paper (Monelli et al.\\ 2010). Data plotted in the top panel of Fig.~6 show that the aforementioned assumptions provide evolutionary features that are in reasonable agreement with the observations (see Fig.~2). In particular, both RC and HB stars are clearly separated in magnitude and in color from each other and from RG stars. Moreover, we selected two bins along the RGB for 20.75$\\le$$m_v$$\\le$21.25 and 21.50$\\le$$m_v$$\\le$22.00 and we found that the spread in \\bmv\\ color is $\\sim$0.06 mag. This value is quite similar, within the errors, to the spread in \\bmv\\ color of the observed CMD at the same magnitude intervals, i.e., 0.06 mag. To constrain the impact of a spread in metal content, we used the same evolutionary ingredients adopted to construct the above synthetic CMD, but we assumed for each of the two stellar components a spread in metallicity of 0.2 dex.  The spread in \\bmv\\ color at the same magnitude intervals is slightly larger than observed, namely 0.08 mag.  Finally, we adopted for the intermediate-age stellar component alone a spread in metallicity of 0.5 dex. Data plotted in the bottom panel of Fig.~6 clearly show that such a spread in metallicity causes the RC stars to attain colors that are redder than the RG stars. The spread in \\bmv\\ color at the same magnitude intervals is a factor of two larger than observed, i.e., 0.12 mag. Moreover, the synthetic CMD suggests the occurrence of subgiant and giant branch stars with colors that are redder than the subgiant of the old sub-population. The observed CMD does not show these features.  The comparison with empirical calibrators further supports the contention that Carina hosts predominantly metal-poor stellar populations, in agreement with the spectroscopic investigations. However, the limited range in color covered by the RC stars appears incompatible with the spectroscopic results: the spread in iron abundance of the intermediate-age population is either small, or is counterbalanced by some other---as yet unrecognized---variable. The same conclusion applies for the old population, based on the limited magnitude spread among the old HB stars. These conclusions are also supported by synthetic CMDs constructed assuming a different spread in metallicity between the intermediate and the old sub-populations. The simulations indicate a spread in metallicity of the two sub-populations smaller than 0.2 dex."
    },
    "1004/1004.0076_arXiv.txt": {
        "abstract": "We present a model for high-energy emission sources generated by a standing magnetohydrodynamical (MHD) shock in a black hole magnetosphere.  The black hole magnetosphere would be constructed around a black hole with an accretion disk, where a global magnetic field could be originated by currents in the accretion disk and its corona.  Such a black hole magnetosphere may be considered as a model for the central engine of active galactic nuclei, some compact X-ray sources and gamma-ray bursts. The energy sources of the emission  from the magnetosphere are the gravitational and electromagnetic energies of magnetized accreting matters and the rotational energy of a rotating black hole.  When the MHD shock generates in MHD accretion flows onto the black hole, the plasma's kinetic energy and hole's rotational energy can convert to radiative energy. In this letter, we demonstrate the huge energy output at the shock front by showing {\\it negative energy}\\/ postshock accreting MHD flows for a rapidly rotating black hole.  This means that the extracted energy from the black hole can convert to the radiative energy    % at the MHD shock front. When axisymmetric shock front is formed, we expect a ring-shaped region with very hot plasma near the black hole; the look would be like an ``aurora''. The high energy radiation generated from there would carry to us the information for the curved spacetime due to the strong gravity. ",
        "introduction": "\\label{sec:intro}   A number of observations provides evidence for super-massive black holes in galactic centers or stellar-mass black hole in galaxies \\citep[e.g.][]{Miyoshi+95,Ghez+98,Ghez+05,Genzel+00,Schodel+03, Bender+05,Eisenhauer+05}. Several observations show high-energy phenomena around black hole candidates where some amount of energy in accretion flows is released. In this paper, to understand the high-energy phenomena near the black hole and the nature of the curved spacetime due to the hole's strong gravity, we consider ingoing magnetohydrodynamic (MHD) flows in the magnetosphere of a black hole, which is considered as the transit region from an accretion disk to the event horizon. The global magnetic field in the magnetosphere could be generated by currents streaming in the disk and its corona around the black hole.   The structure of a steady and axisymmetric black hole magnetosphere with a thin disk has been discussed by \\citet{TT01} as a vacuum magnetosphere, \\citet{Uzdensky05} as a force-free magnetosphere, and \\cite{Camenzind87,NTT91} as a MHD magnetosphere. These models show the disk--black hole connecting magnetic field lines, as shown in Figure~\\ref{fig:BH-Mag}a. The loop-like shaped magnetic field is generated around the disk's inner edge region. Some magnetic field lines connect to the event horizon. Due to the plasma inertia effect, the magnetic field lines are forced to direct toward the event horizon, and then almost radial flows would be reasonable at least near the event horizon. Such a configuration of the loop-like magnetic field lines indicates the interaction between the black hole and the disk surface (or corona) located at several times the inner-edge radius of the thin disk. Along this loop-like magnetic field lines, the MHD fluid ejected from the disk surface streams inward due to the dominant gravitational force by the black hole.   Although in many previous works equatorial accretion flows have been discussed, by considering the black hole magnetosphere, we expect naturally off-equatorial inflows streaming along the loop-like magnetic field lines. So, we also expect the shock formation near the polar region of the black hole \\cite[see][for adiabatic shocks]{FTT07}. Of course, the shock formation by equatorial accretion flows would be possible within the required physical conditions.  However, the emission from the off-equatorial MHD shock would include the information for the ``magnetosphere'' (i.e., the configuration of the magnetic fields) around a black hole, and would enable us to find the information for the black hole spacetime related to the radiation.   We are interested in the energy conversion from the accreting plasma's energy, which is composed of the kinetic part and magnetic part, to the radiative energy at the inner-most region of the black hole magnetosphere.  This is because such energy conversion by the shock formation would be related to the generation of huge radiative energy in the hole's deep gravitational potential. Furthermore, in addition to the plasma's kinetic energy, the rotational energy of the black hole may also be converted to the radiative energy (see Fig.~\\ref{fig:BH-Mag}b). For a fast MHD shock on accretion onto a black hole, about 10\\% of the upstream flow's energy (including the rest mass energy of a particle) can convert to the thermal energy of the downstream flow \\citep{TGFRT06}. This means that the plasma temperature becomes so high across the shock front. Such a hot plasma region can be considered as a source of high energy radiation, which gives us both the information for the strong gravitational field and the state of magnetized plasma around the black hole. Of course, some part of the radiation emitted from the hot plasma will fall into the black hole because of the strong gravitational lens effects. \\footnote{ The escape rate of photons from the ergosphere is almost 10\\% and this rate becomes larger for a larger spin value \\cite[see][]{RT-MT10}. } Nevertheless we can expect that huge radiative energy can be released at the very hot shocked plasma region, and considerable radiation flux will be obtained for us.   \\begin{figure}[t]% \\begin{center} \\includegraphics[scale=0.6,angle=0]{fig-1a.eps} \\includegraphics[scale=0.3,angle=0]{fig-1b.eps} \\end{center} \\caption{ (a) A model of the black hole magnetosphere. Ingoing MHD flows stream along the disk--black hole connecting magnetic field lines. Because of the loop-shaped magnetic field configuration, off-equatorial shock would be expected. (b) MHD shock in black hole accretion.  After passing through {\\it first}\\/ fast magnetosonic point, the MHD inflow  can make the MHD shock. The postshocked sub-fast magnetosonic flows must pass through the {\\it second}\\/ fast magnetosonic point again before reaching the event horizon. } \\label{fig:BH-Mag} \\end{figure} %   In this paper, we discuss energy release by shocked MHD accretion onto a rapidly rotating black hole. In \\S 2, we introduce trans-fast magnetosonic accretion flows and apply the shock condition to the flows. Here, we adopt the method by \\cite{TT08} for solving trans-fast magnetosonic flow solutions. Then, we can easily obtain the solutions without critical condition analysis at the magnetosonic points. In \\S 3, we show {\\it negative}\\/ energy postshock MHD inflows that means the energy extraction from a rotating black hole and discuss the released energy at the shocked region. In \\S 4, we summarize our MHD shock models. ",
        "conclusions": "We have discussed global shocked accretion solutions in a black hole magnetosphere, which are composed of two trans-magnetosonic solutions with a MHD shock where the general relativistic MHD has been applied to the accretion flows.  Then, we have shown the possibility of a very hot shocked plasma region very close to the event horizon. At this stage, in the study of the black hole magnetosphere, the expected energy spectrum from the MHD shock is not clear. Nevertheless, we newly point out the possibility of the energy extraction from the black hole through the MHD shock with the negative energy in the postshock flow. Although, in this paper, we show the possibility of shocked accretion solutions with negative energy ($E<0$), the dependence of flow parameters (i.e., $E_{\\rm up}$, $\\Omega_F$, $a$, $L_{\\rm up}$) and flow's latitude will be presented in the next paper (Takahashi \\& Takahashi, in preparation).    To complete the black hole's rotational energy release at the MHD shock, we must investigate the structure of the shocked hot plasma region.  Although the value of the released energy depends on the field-aligned flow parameters, we can estimate the plasma density and temperature there by the help of the accretion disk models that give the boundary condition for the magnetosphere.  We need a realistic model of the black hole magnetosphere with a magnetized accretion disk. Although we assume a given magnetic field line, the configuration may be unstable. Then, the dynamical behaviors of the field configuration should be also investigated by helping numerical simulations \\citep[e.g.,][]{Koide06,MG04,McKinney06}. Then, a significant amount of X-rays and gamma-rays would be estimated. In order to take into account such radiation effects in the general relativistic plasma, we should formulate the local radiation effects of MHD fluid by using the energy-momentum tensor with the local heat flux, in addition to the fluid part and the electromagnetic part. Then, we will find the evidences for {\\it real}\\/ black holes in the observational data % (e.g., the spectrum in high-energy range, the short time variability, the radio image of black hole shadow, X-ray polarization, etc). That is, we expect that the high energy emission from this hot plasma brings to us {\\it additional}\\/ information about the black hole spacetime; especially, the polar region emission including this information can be distinct from the emissions from the equatorial plasma source, which have been investigated by many authors in the models of accreting gas disk."
    },
    "1004/1004.2842_arXiv.txt": {
        "abstract": "We report the discovery by B. G. Harris and S. Dvorak on JD 2455224.9385 (2010 Jan 28.4385 UT) of the predicted eruption of the recurrent nova U Scorpii (U Sco).  We also report on 815 magnitudes (and 16 useful limits) on the pre-eruption light curve in the UBVRI and Sloan r' and i' bands from 2000.4 up to 9 hours before the peak of the January 2010 eruption.  We found no significant long-term variations, though we did find frequent fast variations (flickering) with amplitudes up to 0.4 mag.  We show that U Sco did not have any rises or dips with amplitude greater than 0.2 mag on timescales from one day to one year before the eruption.  We find that the peak of this eruption occurred at JD 2455224.69$\\pm$0.07 and the start of the rise was at JD 2455224.32$\\pm$0.12.  From our analysis of the average B-band flux between eruptions, we find that the total mass accreted between eruptions is consistent with being a constant, in agreement with a strong prediction of nova trigger theory.  The date of the next eruption can be anticipated with an accuracy of $\\pm$5 months by following the average B-band magnitudes for the next $\\sim$10 years, although at this time we can only predict that the next eruption will be in the year 2020$\\pm$2. ",
        "introduction": "Recurrent novae (RNe) are ordinary novae (binary systems with mass accreting onto a white dwarf until thermonuclear runaway is triggered) for which the recurrence time scale is between a decade and a century, such that more than one eruption has been observed (Payne-Gaposchkin 1964; Bode \\& Evans 2008; Evans et al. 2008).  To have the fast recurrence time scale, the novae must have the white dwarf near the Chandrasekhar mass and have a high accretion rate.  These properties, at face value, imply that the white dwarf will soon exceed the Chandrasekhar mass and become a Type Ia supernova, and thus RNe are one of the premier candidates for the progenitor class of these supernovae.  RNe typically have relatively fast eruptions, high ejection velocities, and small eruption amplitudes when compared to ordinary novae.  Only ten RNe are known with certainty in our Milky Way (Schaefer 2010).  U Scorpii (U Sco) previously erupted in March 1999 with a peak at V=7.5 mag (Schaefer 2010).  In quiescence, it has V$\\approx$17.6 and has deep {\\it total} eclipses taking it down to V=18.9 mag (Schaefer 2010) with an orbital period of 1.23 days (Schaefer 1990; Schaefer \\& Ringwald 1995).  U Sco is the fastest of all known novae, fading by three magnitudes from peak in just 2.6 days, while its rise from minimum to peak is 6-12 hours (Schaefer 2010).  No light echo was detected to deep limits after the 1987 eruption (Schaefer 1988).  U Sco has now had ten known eruptions, in the years 1863, 1906, 1917, 1936, 1945, 1969, 1979, 1987, 1999 (Schaefer 2010), and now 2010 as we report in this paper.  With the discovery of the 1917, 1945, and 1969 eruptions (Schaefer 2001; 2004), it has become apparent that U Sco has outbursts at intervals of $10\\pm2$ years since 1900.  The exceptions to this are the two intervals of 19 and 24 years, which are easily interpreted as being double intervals, with eruptions around 1927 and 1957 having been missed.  (U Sco is $3.6\\degr$ from the Sun every 28 November, so a significant fraction of its very fast eruptions must be missed.)  With this, it became apparent that the next eruption of U Sco should occur in the year $2009\\pm2$.  Schaefer (2005) made a better prediction, on the physical basis that the time between eruptions scales as the inverse of the average mass accretion rate between eruptions (as measured from the B-band flux), with the scaling determined by the inter-eruption light curves from prior eruptions.  The predicted eruption date was $2009.3\\pm1.0$.  This is the first time that a specific star has been predicted to have an eruption on a specific date.  With this advance notice, a large international collaboration was formed to provide detailed photometry and spectroscopy in the X-ray, ultraviolet, optical, and infrared bands.  With U Sco going from quiescence to peak to one magnitude below peak in 24 hours, we realized that we must have frequent monitoring of U Sco to get a fast alert of an eruption.  To this end, we mobilized daily and hourly photometry with the SMARTS 1.3-m telescope in Chile, the fully-robotic 2.0-m Liverpool telescope (Steele et al. 2004) in the Canary Islands, and the four ROTSE 0.45-m telescopes in Australia, Texas, Namibia, and Turkey.  In addition, we mobilized a large number of observers through the American Association of Variable Star Observers (AAVSO).  For the seven months each year centered on the opposition of U Sco, we got hourly data.  The headquarters of the AAVSO served as the international clearinghouse for discovery reports and delivery of alerts to the world.  In addition, U Sco was heavily monitored from 2001 to 2009 with long time series photometry, where the main goal was to precisely measure the timing of the eclipses.  The result of all this activity from 2000-2010 is the all-time best pre-eruption light curve for any nova.  This paper presents all the magnitudes and an analysis of this large data set. ",
        "conclusions": ""
    },
    "1004/1004.4525_arXiv.txt": {
        "abstract": "The viability of a possible cosmological scenario is investigated. The theoretical framework is the constrained next-to-minimal supersymmetric standard model (cNMSSM), with a gravitino playing the role of the lightest supersymmetric particle (LSP) and a neutralino acting as the next-to-lightest supersymmetric particle (NLSP). All the necessary constraints from colliders and cosmology have been taken into account. For gravitino we have considered the two usual production mechanisms, namely out-of equillibrium decay from the NLSP, and scattering processes from the thermal bath. The maximum allowed reheating temperature after inflation, as well as the maximum allowed gravitino mass are determined. ",
        "introduction": "There is accumulated evidence both from astrophysics and cosmology that about 1/4 of the energy budget of the universe consists of so called dark matter, namely a component which is non-relativistic and neither feels the electromagnetic nor the strong interaction. For a review on dark matter see e.g.~\\cite{Munoz:2003gx}. Although the list of possible dark matter candidates is long, it is fair to say that the most popular dark matter candidate is the lightest supersymmetric particle (LSP) in supersymmetric models with R-parity conservation~\\cite{Feng:2003zu}. The superpartners that have the right properties for playing the role of cold dark matter in the universe are the axino, the gravitino and the lightest neutralino. By far the most discussed case in the literature is the case of the neutralino (see the classic review ~\\cite{Jungman:1995df}), probably because of the prospects of possible detection.  However, the gravitino is another very interesting candidate for cold dark matter, since its interactions are completely determined by the supergravity lagrangian~\\cite{Cremmer:1982en}, in contrast to what happens to the neutralino or the axino case, where the interactions depend on the chosen model. Unfortunately, gravitino belongs to the class of new exotic particles that can be potentially dangerous for cosmology, and it is therefore escorted by the so-called gravitino problem~\\cite{Ellis:1982yb}. The mass of the gravitino strongly depends on the SUSY-breaking scheme, and can range from eV scale to scales beyond the TeV region~\\cite{Giudice:1998bp, Nilles:1983ge, Randall:1998uk}. In particular, in gauge-mediated SUSY-breaking schemes~\\cite{Giudice:1998bp} the gravitino mass is typically less than 100~MeV, while in gravity-mediated schemes~\\cite{Nilles:1983ge} it is expected to be in the GeV to TeV range. Finally, it must be noted that there are hybrid models of gauge- and gravity-mediation, in which gravity provides sub-dominant and yet non-negligible contributions~\\cite{Hiller:2008sv}. Therefore, according to the precise mechanism for supersymmetry breaking, the gravitino can be either stable or unstable, with the corresponding gravitino cosmology. In general, the gravitino problem requires that the reheating temperature after inflation should be lower than $10^6-10^7$~GeV~\\cite{Kawasaki:2004yh, Pradler:2006hh}, which poses serious difficulties to the thermal leptogenesis scenario~\\cite{Buchmuller:2002rq}.  In the present work we want to consider gravitino dark matter in the constrained next-to-minimal supersymmetric standard model (cNMSSM), assuming that the lightest neutralino is the next-to-lightest supersymmetric particle (NLSP), and taking into account the two usual gravitino production mechanisms to be discussed later on, namely the out-of-equillibrium decays of the NLSP, as well as scattering processes from the thermal bath.  This article is organized as follows. In the next section we present the theoretical framework. In section 3 we discuss all the relevant constraints from colliders and from cosmology, and we show our results. Finally, we conclude. ",
        "conclusions": "In the framework of the cNMSSM, we have considered a possible cosmological scenario with the gravitino LSP and the neutralino NLSP. The gravitino is stable and plays the role of cold dark matter in the universe, while the neutralino is unstable and decays to gravitino. We have taken into account the relevant gravitino production mechanisms, which are i) the NLSP decay, and ii) scattering processes from the thermal bath. Our results can be seen in the figures. We have found that i) the gravitino is necessarily very light, and ii) the reheating temperature after inflation is two orders of magnitude lower than the temperature required for thermal leptogenesis. The singlino scenario must be excluded, while in the bino case it is possible to have a gravitino in the gravity-mediated SUSY-breaking scheme."
    },
    "1004/1004.4149_arXiv.txt": {
        "abstract": "The nuclear physics input used to compute the Monte Carlo reaction rates and probability density functions that are tabulated in the second paper of this series (Paper II) is presented. Specifically, we publish the input files to the Monte Carlo reaction rate code \\texttt{RatesMC}, which is based on the formalism presented in the first paper of this series  (Paper I). This data base contains overwhelmingly experimental nuclear physics information. The survey of literature for this review was concluded in November 2009. ",
        "introduction": "For a number of reasons, the present evaluation of charged-particle thermonuclear reaction rates represents a significant step forward compared to previous work. First, we developed a new method of computing reaction rates, which is based on Monte Carlo techniques and assigns to each nuclear input quantity a physically motivated probability density function. The method is described in the first paper of this series (Paper I) and the numerical results for reaction rates and rate probability density functions are presented in the second paper of this series (Paper II). Second, a number of years have passed since the last two evaluations \\cite{Ang99,Ili01} of similar scope as the present work have been published. The rapid progress seen in the field of nuclear astrophysics over the past few years clearly warrants a new charged-particle reaction rate evaluation. Third, thermonuclear reaction rates are not directly measured quantities, but are derived from a multitude of different nuclear physics input quantities. Consequently, the quality and reliability of a reaction rate evaluation hinges directly on the transparency and reproducibility of the input data. The last two aspects will be addressed in the present paper.  We present here the input files to the Monte Carlo code \\texttt{RatesMC} containing the nuclear physics data used to compute our reaction rates. For each reaction we list nuclear properties; nonresonant S-factors; recommended resonance energies, strengths and partial widths; upper limits of resonance contributions; numerical rate integrations; and interferences between resonances of same spin and parity. Section 2 contains a brief discussion of general procedures, literature sources and the status of data. In Sec. 3 we focus attention on a number of important issues. A brief summary is given in Sec. 4. The meaning of each input row in a sample input file to \\texttt{RatesMC} is explained in App. \\ref{sec:inptfiles}. The actual nuclear physics input data for each of the reactions evaluated in the present work are listed in App. \\ref{sec:nuclinput}. ",
        "conclusions": "In the present paper of the series (Paper III) we publish the nuclear physics input data used to compute the Monte Carlo reaction rates presented in  Paper II. The reaction rates are calculated using the new method discussed in Paper I. Our input data, listed in the appendix, are based on the most reliable information that we are able to extract from the published literature. By making our recommended input data available to the community we intend to improve the reproducibility and transparency of the evaluated reaction rates. The reaction rate uncertainties given in Paper II are statistical in nature and do not account for unknown systematic errors in the nuclear input data listed here. The survey of literature for this review was concluded in November 2009."
    },
    "1004/1004.4180_arXiv.txt": {
        "abstract": "\\noindent   The Eurybates family is a compact core inside the Menelaus clan, located in the L$_4$ swarm of Jupiter Trojans. Fornasier et al. (2007) found that this family exhibits a peculiar abundance of spectrally flat objects, similar to Chiron-like Centaurs and C-type main belt asteroids. On the basis of the visible spectra available in literature, Eurybates family's members seemed to be good candidates  for having on their surfaces water/water ice or aqueous altered materials. \\\\ To improve our knowledge of the surface composition of this peculiar family, we carried out an observational campaign at the Telescopio Nazionale Galileo (TNG), obtaining near-infrared spectra of 7 members. Our data show a surprisingly absence of any spectral feature referable to the presence of water, ices or aqueous altered materials on the surface of the observed objects. Models of the surface composition are attempted, evidencing that amorphous carbon seems to dominate the surface composition of the observed bodies and some amount of silicates (olivine) could be present. ",
        "introduction": "Jupiter Trojans (JTs) are small bodies of the Solar System located in the Jupiter's Lagrangian points L$_4$ and L$_5$. Their origin is not yet well understood and it is still matter of debate. Several mechanisms were proposed to model their origin (Marzari \\& Scholl, 1998a,b, 2000, 2007; Marzari et al., 2002; Morbidelli et al., 2005), and it is widely accepted that they formed in the outer Solar System, in regions rich in frozen volatiles. The JT population is supposed to have undergone a significant collisional evolution, and to be at least as collisionally evolved as main belt asteroids. The discovery of dynamical families in both L$_4$ and L$_5$ clouds supports this hypothesis (e.g. Milani, 1993; Milani \\& Kne$\\breve{z}$evi\\'c, 1994; Beaug\\'e \\& Roig, 2001; Dell'Oro et al., 1998).  Physical properties of JTs are poorly known. The presently available  spectroscopic data set is largely unsatisfactory, covering only about 10\\% of the entire JT population. To improve our knowledge of the nature of these bodies, in the last years several surveys have been carried out, both in visible and infrared wavelengths (Dotto et al., 2008, and reference therein). Although JTs formed at large heliocentric distances, the data so far acquired have shown a lack of any evidence of ices on their surfaces. Emery \\& Brown (2003, 2004) analysed the content of water ice and hydrated materials on the surface of 17 JTs, obtaining upper limits of a few\\% for water ice and of 30\\% for hydrated materials. More recently, Yang \\& Jewitt (2007) suggested that water ice can occupy no more than 10\\% of the total surface of (4709) Ennomos. JTs belonging to dynamical families do not exibit any spectral feature related to the presence of ices on their surface (Dotto et al., 2006).  The data so far available in the literature put in evidence a great homogeneity in the whole population: all of the known JTs are low albedo bodies belonging to the primitive $C$, $P$ or $D$ classes. The same uniformity  is also found in JTs belonging to dynamical families (Fornasier et al., 2004, 2007; Dotto et al., 2006). However, some differences between the L$_4$ and L$_5$ swarms are evident: as discussed by Fornasier et al. (2007), the majority of L$_5$ JTs are D-types, while an higher presence of C and P types is observed among the L$_4$ objects.  A peculiar case is given by the Eurybates dynamical family, in the L$_4$ swarm. This family is a strong cluster inside the Menelaus clan (Roig et al., 2008), which survives also at a very low relative velocity cut-off, as defined by Beaug\\'e and Roig (2001). The family population, up to date, is composed by 28 members at a cut-off of 70 m/s, 22 of them surviving also at 40 m/s. On the basis of the dynamical properties, it is still not possible to understand if the Eurybates members constitute a distinct family that lies in the same space of proper elements of Menelaus or, as suggested by Roig et al. (2008), they formed by a secondary break-up of a former Menelaus member.  Although the Eurybates family is clustered, in the space of proper elements, in a small portion of the region occupied by the Menelaus clan, its members show spectral properties quite different from those of Menelaus: as shown by Roig et al. (2008) the Menelaus clan evidences a larger diversity of taxonomic classes including C, P, and D-type objects in agreement with the whole JT population, while the Eurybates members are characterized by almost flat visible spectra (see e.g. Fig. 12 in Roig et al., 2008), with spectral slopes strongly clustered around 2 \\%/10$^3$ \\AA, and  spectral behaviors similar to those of C--type main belt asteroids and/or Chiron-like Centaurs (Fornasier et al., 2007).  The Eurybates family assumes a great importance in the study of JTs because such a peculiar clustering of spectrally flat objects strongly affects the color-size-orbital parameter distributions of the whole JT population investigated up to now. Fornasier et al. (2007) noted how this family fills the distribution of spectrally neutral JTs at low inclination and appears to be the major responsible of a color-inclination trend (bluer bodies concentrated at lower inclination) of the whole JT population. In the same paper, the Eurybates family appears also to be the major cause of the abundance of C- and P-types among the L$_4$ objects, which would imply a more heterogeneous composition of this swarm than the L$_5$ one. Moreover, the Eurybates family strongly contributes to the population of L$_4$ small JTs (with a D$<$40~km) having low spectral slopes.  The observations made by Fornasier et al. (2007) showed the presence of a drop-off of reflectance shortward of 0.52~$\\mu$m in the visible spectra of four Eurybates members (18060, 24380, 24420 and 39285). This behavior is detected on the spectra of many main belt C-type asteroids (Vilas, 1994; 1995), and it is often associated to other spectral features due to aqueous alteration products. Since no other absorption features were found on the visible spectra of Eurybates members, we still do not have a final proof that aqueous alteration processes occurred on the surface of these bodies. Nevertheless, the presence of the ultraviolet drop-off could suggest that subsurface water or water ice could have been present on Eurybates members at a certain moment of their life, in order to cause aqueous alteration on their surfaces. They are therefore good candidates to preserve still detectable spectral signatures of water/water ice or aqueous altered materials. ",
        "conclusions": ""
    },
    "1004/1004.5052_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{chap:Introduction}  \\label{sect:fR:Introduction}  As was commented in Chapter 1, when the modified Einstein equations were rewritten {\\it \\`{a} la Einstein}, the presence of a function $f(R)$ in the gravitational sector modifying the usual EH Lagrangian may be understood as the introduction of an effective fluid which is not restricted to hold the usual energy conditions. Therefore $f(R)$ functions may be used to explain the present cosmological acceleration. Historically, some $f(R)$ models were proposed to modify GR at short scales, i.e., high energies trying to explain inflation, as for instance $f(R)\\propto R^{2}$, but no interest was paid in those models to provide a mechanism to cause late time acceleration. First attempts to induce cosmological acceleration considered $f(R)\\propto 1/R$ but those models turned out to be in conflict with solar system tests \\cite{Chiba} and even to be unstable when matter is introduced \\cite{Dolgov&Kawasaki}.   Before studying these issues, let us mention that $f(R)$ models, apart from satisfying those gravitational and cosmological conditions given in Section \\ref{sec:Int:fR:Constraints}, should verify some extra conditions of cosmological viability. For instance, they have to include a background evolution with Big Bang nucleosynthesis (BBN) and both radiation and matter dominated cosmological eras. This fact will be explicitly studied in this chapter in Section \\ref{sect:fR:Reconstruction}. On the other hand, they must provide cosmological perturbations compatible with cosmological constraints from CMB and large scale structures (LSS). This fact will be studied thoroughly in Chapter 3.   The present chapter is organized as follows: in Section \\ref{sect:fR:Standard Einstein's eqns for FLRW} we shall revise the standard approach to describe the cosmological evolution in the $\\Lambda\\text{CDM}$ model in a homogeneous, isotropic and spatially flat metric. In the following Section \\ref{sect:fR:Modified Einstein's equations for FLRW universe} we shall generalize the usual Einstein equations when $f(R)$ gravity theories are present. Then we shall study in Section \\ref{sect:fR:Cosmological viability} the cosmological viability conditions for $f(R)$ theories to hold a dust matter dominated era followed by a late time acceleration and some interesting models which have been proposed to be viable will be provided in Section \\ref{sect:fR:Some viable f(R) models}. Then, Section \\ref{sect:fR:Reconstruction} will be the core of the chapter and it will be devoted to study if $f(R)$ theories are able to mimic standard $\\Lambda\\text{CDM}$ evolution. These $f(R)$ models will possess vacuum solutions with null scalar curvature what allows to recover some GR solutions usually considered.  To finish this chapter, we shall study in Section \\ref{sect:fR:Arbitrary fluid} how the modification of the gravitational sector by $f(R)$ models may mimic the influence of perfect fluids (parameterized by a constant equation of state) in the cosmological evolution without any presence of such a fluid in the fluid content. The chapter will finish with Section \\ref{sect:fR:Conclusions} by drawing some attention over the main obtained conclusions.   The results presented in this chapter were originally published in \\cite{Dobado&Dombriz}.  This chapter will be devoted to a study of the evolution of scalar cosmological perturbations in $f(R)$ theories. To do so, a completely general procedure will be implemented and several consequences will be analyzed. The importance for addressing the present problem lies in the necessity to discriminate among different DE models, including $f(R)$ modified gravities, by using observations. It is well-known that by choosing adequate $f(R)$ functions, one can mimic any expansion history, and in particular that of the $\\Lambda\\text{CDM}$ model. Accordingly, the exclusive use of % observations from SNIa \\cite{SN}, baryon acoustic oscillations \\cite{BAO} or CMB shift factor \\cite{WMAP}, based upon different distance measurements which are sensitive only to the cosmological expansion history, cannot settle the question of the DE nature \\cite{Linder2005}.  However, there exists a different type of observations which are sensitive, not only to the cosmological expansion history, but also to the evolution of matter density perturbations. Pioneering work % on density perturbations in FLRW cosmological models was presented in \\cite{Lifshitz}. The fact that the evolution of perturbations depends on the specific gravity model, i.e., it differs in general from that of Einstein's gravity even though the background evolution is the same, means that this kind of observations will help distinguishing between different models capable to explain cosmic acceleration. This is therefore the aim of the present chapter: to provide a completely exact method to determine how cosmological perturbations grow in $f(R)$ theories and to settle the extracted consequences from this result. Such a problem has been exhaustively considered in the literature \\cite{perturbations}. In this chapter we shall show that for $f(R)$ theories the differential equation for the matter density contrast is a fourth order differential equation in the longitudinal (also called conformal or Newtonian) gauge but it reduces to a second order differential equation for sub-Hubble modes.  This general equation will be compared with the standard simplification procedure (so-called quasi-static approximation) widely used in the previous literature. This approximation considers simplifications in the density equation determination from the very beginning. We shall show that for general $f(R)$ functions, the quasi-static approximation is not justified. However, for those $f(R)$ adequately describing the present phase of accelerated expansion and satisfying local gravity tests, it does give a correct description for the evolution of perturbations.  Once the general results are presented, some immediate applications may be implemented. For instance, our analysis may also be used to settle the validity of some proposed $f(R)$ models, by comparing the predicted matter spectra with recent observations of LSS \\cite{Tegmark}.   The present chapter is organized as follows: in Section \\ref{sec:Perturbations:Theory perturbations} we briefly review the theory of cosmological perturbations for the standard $\\Lambda\\text{CDM}$ model, introducing the gauge-invariant variables and revising the well-known results for the EH theory in order to establish a comparison with $f(R)$ gravities. Next, Section \\ref{sec:Perturbations:Perturbations in f(R)} will be devoted to thoroughly study cosmological perturbations in $f(R)$ theories. The perturbed modified Einstein equations are obtained through a completely general procedure for those theories in the Subsection \\ref{subsec:Perturbations:Einstein's eqns f(R) perturbations}. In Subsection \\ref{subsec:Perturbations:Density equation in f(R)} the density perturbations equation is obtained whereas the Subsection \\ref{subsec:Perturbations:Quasi-static approximation} compares the obtained exact results with the ones given by using the quasi-static approximation. Finally in this section, in Subsection \\ref{subsec:Perturbations:Some proposed models} we shall study the growth of perturbations for some particular $f(R)$ models. Section \\ref{sec:Perturbations:Techniques to rule out} will then show how our previous results may be used to constrain or rule some $f(R)$ models out and finally some general conclusions for the presented results will be given in Section \\ref{sec:Perturbations:Conclusions perturbations}.   This chapter is based upon the results presented in references \\cite{Cruz, Perturbations_proceedings, PRLcomment09}. \\par \\label{sec:BH:Introduction} We finish the exposition of our research on $f(R)$ modified gravities by considering some aspects derived from the study of BHs in these theories.   Following the motivation already explained in Section \\ref{sec:Int:Motivation}, $f(R)$ models may present BH solutions as GR and other alternative gravity theories do. Therefore it is quite natural to ask about BHs features in those gravitational theories since, on the one hand, some BHs signatures may be peculiar to Einstein's gravity and others may be robust features of all generally covariant theories of gravity. On the other hand, the results obtained may lead to rule out some models which will be in disagreement with expected physical results. For those purposes, research on thermodynamical quantities of BHs is of particular interest.  These attempts to detect particular signatures from these objects could be experimentally implemented, as was explained in Section \\ref{sec:Int:BW:BH}, at the LHC in the coming years. Therefore the generation of mini BHs could provide important information about the correct underlying gravity theory.   Previous literature on $f(R)$ theories \\cite{Whitt}  proved, by previously performing a conformal transformation in the gravitational action, that  Schwarzschild solution is the only  static spherically symmetric solution for an action  of the form $R+aR^2$ in $D=4$. Also by using this conformal transformation, uniqueness theorems of spherically symmetric solutions for general polynomial actions in arbitrary dimensions % were proposed in \\cite{Mignemi} (see also \\cite{Multamaki} for additional results and \\cite{olmo} for spherical solutions with sources).  Using the Euclidean action method \\cite{Witten, Hawking&Page} in order to determine different thermodynamical quantities, anti-de Sitter ($AdS$) BHs in $f(R)$ models have been studied \\cite{Cognola}. In  \\cite{Briscese} the entropy of Schwarzschild-de Sitter BHs was calculated for some particular cosmologically viable models in vacuum and their cosmological stability was discussed.   BH properties have also been widely studied in other modified gravity theories: for instance, \\cite{cvetic,Cai_GaussBonet_AdS} studied BHs in Einstein's theory with a Gauss-Bonnet term plus a cosmological constant. Different results were found depending on the dimension $D$ and the sign of the constant horizon curvature $k$. For $k=0,-1$, the Gauss-Bonnet term does not modify $AdS$ BHs thermodynamics at all (only the horizon position is modified with respect to the EH theory) and BHs are not only locally thermodynamically stable, but also globally preferred. Nevertheless, for $k=+1$ and $D=5$ (for $D\\geq6$ thermodynamics is again essentially that for $AdS$ BH) there exist some features not present in the absence of Gauss-Bonnet term. Gauss-Bonnet and/or Riemann squared interaction terms were studied in \\cite{Cho}, where the authors concluded that in this case phase transitions may occur with $k=-1$ .  Another approach is given by Lovelock gravities, which are free of ghosts and where the field equations contain no more than second derivatives of the metric. These theories were for instance studied in \\cite{Matyjasek} and the corresponding entropy was evaluated.   The layout of this chapter is as follows: In Section \\ref{sec:BH:Constant curvature black hole solutions} some generalities about BHs in $f(R)$ theories such as several aspects of constant curvature solutions for static spherically symmetric cases with and without electric charge are presented. Then in Section \\ref{sec:BH:Perturbative results} a perturbative approach around the EH action, with no previous imposition of constant curvature, is performed. There we shall find that up to second order in perturbations only BH solutions of the Schwarzschild-$AdS$ type are present. Explicit expressions for the effective cosmological constant are obtained in terms of the $f(R)$ function. To deal with such differential equations we have used the algebraic manipulation package $\\tt Mathematica$~\\cite{mathematica}.   Finally, we shall consider in Section \\ref{sec:BH:Black hole thermodynamics} the BHs thermodynamics in $AdS$ space-time. There it will be found that this kind of solutions can only exist provided the theory satisfies $R_0+f(R_0)<0$ where $R_0$ holds for the constant curvature solution. Interestingly, this expression is proved to be related to the condition which guarantees the positivity of the effective Newton's constant in this type of theories. In addition, it also ensures that the thermodynamical properties in $f(R)$ gravities are qualitatively similar to those of standard GR. Then, some consequences for local and global stability for some particular $f(R)$ models will be provided in Section 4.5 and figures of thermodynamical regions will be shown in Section 4.6. The present chapter is finished by Section \\ref{sec:BH:Conclusions} remarking some conclusions of the presented results.   This chapter is mainly based upon the results that have been presented in references \\cite{BH_Dombriz} and \\cite{BH_Dombriz_Proceeding}.  \\label{sec:CS:Introduction} Although the properties of BH solutions in BW models have been exhaustively studied in the literature, this is not the case of a different type of (topological) spherically symmetric solutions, the so-called brane-skyrmions. These configurations may appear in a natural way in a broad class of BW theories.  On the one hand, this type of textures can be understood as holes in the brane which make possible to pass through them along the extra-dimensional space. For such objects, cosmological involved scales and consequences have not been already studied in the literature. Therefore, the {\\it leitmotiv} of this chapter will be precisely to study some astrophysical and cosmological effects that those configurations may have in the CMB data. %  On the other hand, some striking features have recently been discovered \\cite{Vielva} in the CMB data such as the presence of an anomalous cold spot (CS) in the WMAP temperature maps. This CS has been interpreted in different manners by making use of different physical mechanisms. Among them, one of the most exciting from the physical point of view could justify the appearance of cold spots as the result of a collapsing texture, perhaps coming from some early universe Grand Unified Theory (GUT) phase transition.  In this chapter we shall propose an alternative explanation to those in the existing literature: it will be shown that the brane-skyrmions provide a natural scenario to reproduce the CS features and that typical involved scales, needed for the proposed brane-skyrmions to describe correctly the observed CS, can be as low as the electroweak scale.   The present chapter is thus organized as follows: brane-skyrmions in brane-worlds are presented in Section \\ref{sec:CS:CS as skyrmions} and then we shall determine the possible distortion in the fractional profile of temperatures generated by the presence of these objects. Next, in Section \\ref{sec:CS:Cold spot in WMAP} we shall present the most recent results that claim the existence of a CS in the CMB data. Different possible explanations that have been previously proposed in the literature will be summarized in Section \\ref{sec:CS:Cosmic texture}.  The subsequent physical interpretation of the calculations performed in Section \\ref{sec:CS:CS as skyrmions} will be presented in Section \\ref{sec:CS:Physical interpretation}. In this section, we will also show that those results are in agreement with the CMB data and with other theoretical proposals. To conclude the chapter, possible future detection of these brane-skyrmions and some conclusions will be studied in Section \\ref{sec:CS:Future prospects and conclusions}.   The results of this chapter were originally published in \\cite{Dombriz&Cembranos:CS}. ",
        "conclusions": "\\label{sect:fR:Conclusions} In this chapter we have found the $f(R)$ gravity which exactly reproduces the same evolution of the universe, from BNN up to the present time, as standard $\\Lambda\\text{CDM}$ model does, but without the introduction of any form of DE or cosmological constant. The gravitational Lagrangian $R+f(R)$ is analytical at the origin and consequently $R=0$ is a vacuum solution for the field equations. Therefore Minkowski, Schwarzschild and other important $R=0$ GR solutions, with $\\Lambda_{}=0$, are also solutions for this $f(R)$ gravity. This result was originally presented in \\cite{Dobado&Dombriz}.\\\\   The price that we have to pay for all those good properties is that our Lagrangian, considered as a function of $R$, has a very complicated analytical structure with cuts along the real axis from infinity to $R=3\\Lambda_{}$ and from $R=4\\Lambda_{}$ to infinity. Obviously the only reasonable interpretation of our action is as some kind of effective action. In classical physics one typically starts from some action principle, obtains the corresponding field equations and finally solves them for some initial or boundary conditions. In this work we have proceeded in the opposite way: we started from solutions obtained in the standard cosmological model and then we have searched for an action that, possessing certain properties, gives rise to field equations having the same solutions.   Classical actions are of course real but effective quantum actions usually have a complex structure coming from loops and related to unitarity. The presence of an imaginary part in the action, evaluated on some classical configuration, indicates quantum lost of stability by particle emission of this configuration \\cite{MD}. Therefore it is tempting to think that our action could have some interpretation in terms of an effective quantum action. However, our action determination procedure does not allow to make such a kind of statement. Complications in the action could be just an artifact of our construction. In any case it was shown that such action exists and it reproduces the present universe evolution without DE having $R=0$ as a vacuum solution. As a drawback of this result, we have shown that this $f(R)$ function cannot be considered as a fully consistent gravitational theory since it does not obey the viability conditions revised in Chapter 1 and that therefore it should be regarded as an effective model to mimic $\\Lambda\\text{CDM}$ cosmological background evolution.   To conclude the chapter a completely general procedure to reproduce EH gravity with an arbitrary perfect fluid by using $f(R)$ theories has been implemented. It has been explicitly shown that any perfect fluid when is described by a constant equation of state can be mimicked by an appropriate $f(R)$ model. Standard cases for perfect fluids such as dust matter, radiation and cosmological constant have been presented in this analysis.  \\chapter[Cosmological perturbations in $f(R)$ theories]{Cosmological perturbations in $f(R)$ theories} \\label{chap:Perturbations}  \\label{sec:Perturbations:Conclusions perturbations} In this chapter we have studied the evolution of matter density perturbations in $f(R)$ theories of gravity. Thus we have presented a completely general procedure to obtain the exact fourth order differential equation for the evolution of scalar perturbations in the longitudinal gauge. This expression is valid for any general $f(R)$ theory and applicable at any scale $k$. If the EH gravitational action -- both with and without cosmological constant -- is considered in this general expression, well-known standard results are recovered.   We have also shown that for sub-Hubble modes, the obtained expression reduces to a second order differential equation. Hence, we have been able to compare this result with that obtained within the quasi-static approximation, widely used in the literature. Our research has explicitly shown that for arbitrary $f(R)$ functions such an approximation is not justified.   However, if we limit ourselves to $f(R)$ theories for which $\\vert f_R\\vert \\ll 1$ today, then the perturbative calculation for sub-Hubble modes requires to take into account, not only the first terms,  but also higher-order terms in the $\\epsilon\\equiv{\\cal H}/k$ parameter. In that case, the resummation of such terms modifies the equation. Thus, this equation can be seen to be equivalent to the quasi-static case but only if the universe expands approximately as in a matter dominated phase or in a $\\Lambda\\text{CDM}$ model. Finally, the fact that for $f(R)$ models with $\\vert f_R\\vert \\ll 1$ the background behaves today precisely as that of $\\Lambda\\text{CDM}$ makes the quasi-static approximation correct in those cases.   We have finally applied our results to prove that no significant departure of $f(R)$ theories from $\\Lambda\\text{CDM}$ is allowed for those models that intend to be cosmologically viable according to recent data. In fact, the strong $k$-dependence appearing in the evolution of perturbations has allowed to rule out $f(R)$ gravities which have recently been claimed to be cosmologically viable. \\chapter[Black holes in $f(R)$ theories]{Black holes in $f(R)$ theories} \\label{chap:Black holes}  \\label{sec:BH:Conclusions} In this chapter we have considered static spherically symmetric solutions in $f(R)$ theories of gravity in arbitrary dimensions. Firstly, we have discussed the constant curvature case (including charged BH in 4 dimensions). Then, the general case, without imposing a priori the condition of constant curvature, has also been studied.  Another important result of this chapter was obtained by performing a perturbative analysis around the EH case, assuming regular $f(R)$ functions. We have found explicit expressions up to second order for the metric coefficients. These coefficients only gave rise to constant curvature (Schwarzschild-$AdS$) solutions as in the EH case.  On the other hand, we have also calculated thermodynamical quantities for the $AdS$ BHs and considered the issue of the stability of this kind of solutions. We have found that the condition for a $f(R)$ theory of gravity to support this kind of BHs is given by $R_0+f(R_0)<0$ where $R_0$ is the constant curvature of the $AdS$ space-time. This condition also implies that the effective Newton's constant is positive and that the graviton does not become a ghost. Consequently thermodynamical and cosmological viabilities of $f(R)$ theories turned out to be related as we have shown.   Finally we have considered several explicit examples of $f(R)$ functions and studied the parameter regions in which BHs in such theories are locally stable and globally preferred. It was found that the qualitative behaviour is the same as in standard EH gravity but the thermodynamical regions are modified depending on the parameters values in each case. \\newpage \\chapter[Brane-skyrmions and the CMB cold spot]{Brane-skyrmions and the CMB cold spot} \\label{chap:Cold spot}  \\label{sec:CS:Future prospects and conclusions} In this section we dedicate some lines to mention how both present and future experiments may confirm the validity of the BW models. Then we shall sketch the most important conclusions and consequences of the brane-skyrmion model presented in this chapter as a viable explanation for the observed CS.  On the one hand, the fact that the fundamental scales of the theory are of the order of the TeV opens the possibility to test this explanation with collider experiments through the production of real or virtual branons and KK-gravitons. The expected signatures of the model \\cite{Mirabelli:1998rt} from the production of KK-gravitons come fundamentally from the single photon channel studied by LEP, which restrict ${M}_D> 1.2$ TeV at 95\\% of confidence level.  On the other hand, the LHC accelerator will be able to test the model up to ${M}_D=3.7$ TeV, analysing single photon and monojet production to establish the number of extra dimensions as was explained in Section \\ref{sec:Int:BW:BH}.  Finally, there also exists the possibility to find signatures of the model at low energies associated with branon phenomenology. This case is more interesting from the cosmological point of view, since branons can constitute the nonbaryonic dark matter abundance as typical WIMPs \\cite{DM2,Cembranos:2004jp}. Present constraints coming from the single photon analysis realized by L3 (LEP) imply $f > 122$ GeV (at 95 \\% C.L.) \\cite{Alcaraz:2002iu,Achard:2004uu} and the LHC will be able to check this model up to $f=1080$ GeV through monojet production \\cite{Cembranos:2004jp}. The idea to test the physics associated with the CS with the next generation of colliders at the TeV scale is a very intriguing and distinctive property of this texture.   \\chapter{Conclusions and prospects} \\label{chap:Conclusions} In this thesis we have studied some cosmological and astrophysical consequences in two different types of modified gravity theories: $f(R)$ theories and brane-world extra dimensions theories.   First, we have shown some relevant results for $f(R)$ modified gravities. Thus, we have studied how some features of general relativity can be mimicked by $f(R)$ models if certain conditions are imposed on those models. We have also proved that there exists a class of $f(R)$ models which can indeed reproduce the present cosmological evolution --from matter-radiation equality till today -- as described by general relativity within the standard cosmological concordance model. These $f(R)$ models allow to remove the cosmological constant term and therefore the observed acceleration is understood to have a purely geometrical interpretation with no dark energy origin. Furthermore, initial conditions may be imposed on such a class of theories in order to recover a null scalar curvature solution for a vacuum scenario. Unfortunately such functions are not cosmologically viable and should be considered as effective models to reproduce $\\Lambda\\text{CDM}$ evolution.  Also, it has been proved that any perfect fluid parameterized by a constant equation of state can be reproduced by the presence of $f(R)$ terms. Thus, these functions could provide a mechanism to reproduce the cosmological behaviour of dark energy type fluids. It therefore seems worthwhile to try to search for a unique phenomenological $f(R)$ theory able to explain the main features in the cosmological evolution, from inflation till late time observed acceleration.  Next, we have considered the modification introduced by the new $f(R)$ terms in the evolution of cosmological perturbations. Here, an analysis concerning first order cosmological scalar perturbations has been performed. We have presented a general method in order to obtain the evolution equation of the matter density contrast for arbitrary $f(R)$ theories. That procedure was proven to be valid regardless of the chosen $f(R)$ model or the size of the involved scales. An important by-product of the calculations is the conclusion that the usual approximations made in the literature for sub-Hubble modes are not always valid. In fact, the required hypotheses to get the so-called quasi-static approximation have been established here. % For those sub-Hubble modes, it was found that only for $f(R)$ models satisfying the local gravity constraints, the evolution of perturbations is indistinguishable from that obtained using the quasi-static approximation but perfectly distinguishable from the evolution obtained by the $\\Lambda\\text{CDM}$ model.   Now that those theoretical calculations were made, the correct evolution for matter perturbations in sub-Hubble modes for $f(R)$ models is obtained. Therefore some robustness tests may be used to compare experimental data with expected theoretical results. Thus both the validity and certain constraints on those $f(R)$ functions may be established. This procedure shows how our theoretical results may be used in order to confront the predictions made by a given $f(R)$ theory with large scale structure observations.    As a final subject within our $f(R)$ models research, we dealt with some features of black holes in arbitrary dimensions. First, the scalar constant curvature case for static and spherically symmetric solutions in vacuum was studied. It was determined that the only possible solution for the modified Einstein equations in this case was the Schwarzschild-(anti-)de Sitter solution. As a complementary result to the previous research, we studied the static and spherically symmetric case without imposing the constant curvature condition. To do so, a perturbative approach was performed around the standard Schwarzschild-anti-de Sitter solution of general relativity. It was found that, up to second order in perturbations, the only solution was the generalized Schwarzschild-anti-de Sitter solution with modified coefficients in terms of $f(R)$ and its derivatives evaluated at the background curvature.  This part of the thesis was finished by studying the thermodynamic properties of Schwarzschild-anti-de Sitter black holes in $f(R)$ theories. Thus accessible thermodynamical regions for several $f(R)$ models were studied in detail. In that realm, a very interesting property, which encourages further investigation, was found: it turned out that thermodynamical viability of $f(R)$ theories in constant curvature solutions was embodied by the condition $1+\\text{d}f(R)/\\text{d}R > 0$. This condition was also required to ensure gravitational viability for $f(R)$ models. Consequently two different aspects of the viability of $f(R)$ theories have been brought together.   The last part of the thesis was devoted to a study of some consequences of the existence of topologically nontrivial brane configurations, the so-called brane-skyrmions, in brane-world theories. In particular we have considered the effects of brane-skyrmions on the temperature fluctuations of the CMB and their potential connection with the so-called CMB cold spot. We have shown that brane-world theories can naturally accommodate this striking cosmological feature by invoking these nontrivial topological configurations. A different type of textures related to Grand Unified Theories has also been considered in the literature, but the model presented here shows that this type of defect could be explained from the electroweak scale physics.   To conclude, the results included in this thesis have shown that modified gravity theories remain compelling candidates to describe the properties of the gravitational interaction on very large scales. Both the validity and viability of these theories have still to be subjected to many theoretical and experimental tests. To do so, the possibility of reproducing standard cosmological results, first order perturbations, black holes and features on the CMB temperature maps were considered as interesting aspects which could shed some light on the fundamental properties of the gravitational interaction.   \\newpage $ $ \\newpage \\makeatletter \\renewcommand{\\chapter}{ \\thispagestyle{plain}% \\global\\@topnum\\z@ \\@afterindentfalse \\secdef\\@chapter\\@schapter} \\def\\@chapter[#1]#2{ \\refstepcounter{chapter}% \\typeout{\\@chapapp\\space\\thechapter.}% \\chaptermark{#1} \\@makechapterhead{#2}% \\@afterheading \\addcontentsline{toc}{chapter}{\\protect\\numberline{\\thechapter}#1} } \\renewcommand{\\@makechapterhead}[1]{% {\\ifnum \\c@secnumdepth >\\m@ne \\Large\\bfseries \\@chapapp\\space \\thechapter \\par\\nobreak \\fi \\interlinepenalty\\@M \\Large\\bfseries #1\\par\\nobreak \\vskip 0\\p@ }} \\renewcommand{"
    },
    "1004/1004.0677_arXiv.txt": {
        "abstract": "{\\textit{Spitzer Space Telescope} observations revealed powerful mid-infrared (mid-IR) H$_{2}$ rotational line emission from the Stephan's Quintet (SQ) X-ray~emitting large~scale~shock ($\\sim15\\times35\\,$kpc$^2$) associated~with~a~collision~between~two galaxies. Because  H$_2$ forms on dust grains, the presence of H$_2$ is physically linked to the survival of dust, and we expect some dust emission to come from the molecular gas. } { To test this interpretation, IR~observations and dust modeling are used to identify and characterize the thermal dust emission from the shocked molecular gas. } {The spatial distribution of the IR emission allows us to isolate the faint PAH and dust continuum emission associated with the molecular gas in the SQ shock. We model the spectral energy distribution (SED) of this emission, and fit it to \\textit{Spitzer} observations. The radiation field is determined with \\textit{GALEX} UV, \\textit{HST}~$V$-band, and ground-based near-IR observations. We  consider two limiting cases for the structure of the H$_2$ gas. It is either diffuse, penetrated by UV radiation, or fragmented into clouds optically thick to UV. } { Faint PAH~and~dust continuum~emission~are~detected in the SQ shock, outside star-forming regions. The $12/24\\,\\mu$m flux ratio in the shock is remarkably close to that of the diffuse Galactic interstellar medium, leading to a Galactic PAH/VSG abundance~ratio. However, the properties of the PAH emission spectrum in the shock differ from that of the Galaxy, which may suggest an enhanced fraction of large and neutrals PAHs. In both models (diffuse or clumpy H$_2$ gas), the IR~SED is consistent with the expected emission from dust associated with the warm ($>150\\,$K) H$_2$ gas, heated by a UV radiation field of intensity comparable to that of the solar neighborhood. This is in agreement with \\textit{GALEX UV} observations that show that the intensity  of the radiation field in the shock is $G_{\\rm UV} = 1.4 \\pm 0.2$ [Habing units]. } { The presence of PAHs and dust grains in the high-speed ($\\sim\\!1000\\,$km~s$^{-1}$) galaxy collision suggests that dust survives. We propose~that~the dust that survived destruction was in pre-shock gas at densites larger than a few $0.1\\,$cm$^{-3}$, which was not shocked at velocities larger than $\\sim \\!200\\,$km~s$^{-1}$. Our model assumes a Galactic dust-to-gas mass~ratio and size~distribution, and present data do not allow us~to identify any~significant deviations of the abundances and size distribution of dust grains from that of the Galaxy. Our model calculations show that far-IR \\textit{Herschel} observations will help constraining the structure of the molecular gas, and the dust size distribution, and thereby to look for signatures of dust processing in the SQ shock. } ",
        "introduction": "Stephan's Quintet (Hickson Compact Group \\object{HCG92}, Arp 319, hereafter SQ) is an extensively studied compact group of four interacting galaxies that have a complex dynamical history \\citep[e.g.][]{Allen1980, Moles1997}. A remarkable feature of SQ is that a giant ($\\approx 15 \\times 35$~kpc) shock is created by an intruding galaxy, \\object{NGC 7318b} (Sbc pec), colliding into \\object{NGC 7319}'s tidal tail, at a relative velocity of $\\sim 1\\,000$~km~s$^{-1}$. Evidence for a group-wide shock comes from observations of an extended X-ray ridge containing shock-heated ($\\sim 5\\times 10^6$~K) gas \\citep{Pietsch1997, Trinchieri2003, Trinchieri2005, O'Sullivan2009}, strong radio synchrotron emission from the radio emitting plasma \\citep{Allen1972, Sulentic2001, Williams2002, Xu2003} and shocked-gas excitation diagnostics from optical emission lines \\citep{Xu2003}. This extended region is denoted \\textit{``ridge''} or simply \\textit{``SQ shock''} in this paper.  Observations with the infrared spectrograph \\citep[\\textit{IRS},][]{Houck2004} onboard the \\textit{Spitzer Space Telescope} have revealed a powerful mid-infrared (mid-IR) $\\rm H_2$ rotational line emission from warm ($\\sim 10^{2} - 10^{3}\\,$K) molecular gas in the SQ shock \\citep{Appleton2006, Cluver2010}. The H$_2$ emission is extended not only over the whole ridge, but in several other structures, including an extension towards the Seyfert  galaxy NGC~7319,  and the  intergalactic starburst SQ-A \\citep{Xu1999}. The latter structure, beyond the northern tip of the ridge, has been shown to contain significant CO-emitting gas \\citep{Gao2000, Smith2001a, Lisenfeld2002}. To explain the H$_2$ emission from the SQ ridge, \\citet{Guillard2009} considered the collision of two flows of multiphase dusty gas and proposed a model that quantifies the gas cooling, dust destruction, H$_{2}$ formation and excitation in the postshock medium. H$_2$ gas can form out of gas that is shocked to velocities sufficiently low ($V_{\\rm s} < 200$~km~s$^{-1}$) for dust to survive. Because H$_2$ molecules form on dust grains \\citep[e.g.][]{Cazaux2004}, dust is a key element in this scenario.  \\citet{Xu1999, Xu2003} reported detection of diffuse far-infrared (hereafter FIR) emission from the intergalatic medium (hereafter IGM) with the \\textit{Infrared Space Observatory (ISO)}, and proposed that the dust emission in the shock region would arise from dust grains that efficiently cool the X-ray emitting plasma via collisions with hot electrons. The FIR emission would then trace the structure of the shock, as suggested by \\citet{Popescu2000a} for the case of shocks driven into dusty gas that is accreting onto clusters of galaxies. However, the poor spatial resolution of these observations makes it difficult to separate the dust emission associated with star formation (in the neighborhood galaxies, or SQ-A) from that really associated with the shock. In addition, \\citet{Guillard2009} show that the dust contribution to the cooling of the hot ($\\sim 5 \\times 10^{6}$~K) plasma  is expected to be low, because of efficient thermal sputtering of the grains. If we assume that the age of the galaxy collision is $\\sim 5 \\times 10^{6}$~yrs, grains  smaller than $0.1\\,\\mu$m in radius must have been destroyed. However, the plasma could still be dusty if, before the shock, a significant fraction of the dust mass was  in larger grains and/or if dust destruction is balanced by mass exchange between the cold and the hot gas phases \\citep{Guillard2009}.  The discovery of bright H$_2$ emission in the ridge set new perspectives on the origin of the dust emission in the SQ ridge. In the context of our model for the H$_2$ formation in the SQ shock, we expect some dust emission from the molecular gas. We use \\textit{Spitzer} observations and a model of the dust emission to test this expectation. \\textit{Spitzer} observations show that the bright Polycyclic Aromatic Hydrocarbons (henceforth PAHs) and mid-IR continuum emitting regions are spatially correlated with UV emission and associated with star-forming regions mainly associated with the individual sources in the group (in particular the spiral arm of the intruder galaxy NGC~7318b) and with SQ-A \\citep{Cluver2010, Natale2010}. These IR-bright regions do not correlate with the radio, X-ray, or H$_2$ line emission, which trace the shock structure. In this paper we focus on the fainter dust emission from the the SQ ridge itself, outside these star-forming regions. The spatial distribution of the mid-IR emission allows us to isolate the dust emission from the shock. UV, visible and near-IR observations determine the spectral energy distribution (hereafter SED) of the radiation \ufb01eld used as an input to the dust model. We consider two limiting cases for the structure of the molecular gas, either diffuse, or fragmented into  clouds that are optically thick to UV light. An updated version of the \\citet{Desert1990} model is used to compute the dust emission from molecular gas in these two cases, and fit the model results to the observed IR SED in the SQ shock. The Galactic dust size distribution is taken as a reference.      This paper is organized as follows. Section~\\ref{sec:UVIRobservations} presents the new IR \\textit{Spitzer} observations of SQ, and the UV, optical and near-IR ancillary data we use in this paper. The method used to perform photometry within the SQ shock region and the results are described in sect.~\\ref{sec:photometry}. Section~\\ref{sec:dustemission} presents the \\textit{Spitzer} imaging and spectroscopy results pertaining to the dust emission in the shock structure, emphasizing the PAH properties in the shock. The physical framework and inputs of the dust modeling are outlined in  sect.~\\ref{sec:models}, and the results are discussed and compared to the mid-IR \\textit{Spitzer} observations in sect.~\\ref{sec:results}. In sect.~\\ref{sec:discussion} we discuss the dust processing  in the shock. Then, we  present our conclusions in sect.~\\ref{sec:conclusion} and propose new observations to constrain the physical structure of the molecular gas in the shock.  In this paper we assume the distance to the SQ group to be 94~Mpc (with a Hubble constant of 70 km~s$^{-1}$~Mpc$^{-2}$) and a systemic velocity for the group as a whole of $6\\,600$~km~s$^{-1}$. At this distance, $10\\ \\rm arcsec=4.56\\,$kpc. ",
        "conclusions": "\\label{sec:conclusion}  In this paper, we present new \\textit{Spitzer} imaging and spectroscopic observations that reveal PAH and VSG emission from the galaxy-scale shock in Stephan's Quintet (SQ). Here are our main observational results:  \\begin{itemize} \\item Faint dust emission is detected in the center of the H$_2$-bright SQ shock structure, outside star-forming regions lying in the SQ halo. The $12 / 24 \\,\\mu$m flux density ratio in the SQ ridge is remarkably similar to that of the diffuse Galactic ISM. This suggests that the PAH to VSG abundance ratio is similar to that of the diffuse ISM of the Galaxy. \\item The global  mid-IR SED is consistent with the expected dust emission from the amount of warm H$_2$ detected by \\textit{Spitzer} ($N_{\\rm H} \\simeq 2 \\times 10^{20}$~cm$^{-2}$) for a UV radiation field intensity of $G_{\\rm UV} \\sim 1$ [Habing unit], which is consistent with UV observations of the shock.  \\item The PAH emission spectrum in the SQ shock is significantly different from that of the diffuse Galactic ISM. The $7.7$, $11.3$ and $17\\,\\mu$m aromatic bands are detected, but the $6.2\\,\\mu$m band is absent. The $17\\,\\mu$m complex is prominent, but the $16.4\\,\\mu$m is not detected. Interestingly,  the $7.7 / 11.3\\,\\mu$m flux ratio in the SQ shock  is a factor $\\sim 2$ lower than that of the diffuse Galactic PAHs. These characteristics  may suggest an enhanced fraction of neutral and large PAHs. \\end{itemize}   \\textit{Spitzer} imaging and spectroscopy reveal powerful H$_2$ emission in the Stephan's Quintet X-ray giant shock that extends over the full area ($\\approx 35 \\times 15$~kpc$^2$) of the ridge \\citep{Cluver2010}. We expect some dust emission to come from molecular gas because H$_2$ forms on dust grains. In this paper we test this interpretation by modeling the IR emission from dust associated with the H$_2$  gas present in the SQ shock structure, and by comparing the model calculations with  \\textit{Spitzer} observations.  \\begin{itemize} \\item We model the emission from dust associated with diffuse or clumpy molecular gas, embedded within H{\\sc ii} gas and X-ray emitting plasma. The model SED is consistent with mid-IR \\textit{Spitzer} observations for both cases, for a Galactic dust-to-gas mass ratio and a Galactic dust size distribution. For diffuse gas, the best-fit column density is $N_{\\rm H} = 1.8 \\pm 0.5 \\times 10^{20}$~cm$^{-2}$, which is close to the value derived from warm H$_2$ observations. For clumpy molecular clouds that are optically thick to UV radiation, we find that the H$_2$ surface filling factor is $f \\sim 0.03 $. So far, the present data and the degeneracy between the dust size distribution and the cloud structure do not allow us to decide whether the molecular gas is diffuse or fragmented into clouds that are optically thick to UV photons.  \\item The presence of dust in the SQ shock shows that dust is able to survive in such a violent environment. We propose that at the time of the high-speed galaxy encounter, the dust that survived destruction was in pre-shock gas at densites larger than a few $0.1\\,$cm$^{-3}$, which was not shocked at velocities larger than $\\sim \\!200\\,$km~s$^{-1}$ \\citep[see ][]{Guillard2009}. Present data do not allow us to indentify a plausible impact of the shock on the dust size distribution, but the peculiar properties of the PAH emission in the SQ ridge (summarized above) may be the result of PAH processing in shocks.   \\end{itemize}   Future Far-IR observations are needed to constrain both the structure of the H$_2$ gas and the dust properties in the SQ shock. The PACS and SPIRE instruments on \\textit{Herschel} would provide SED of the Far-IR dust emission resolved on scales of $5-10$~arcsecs, crucial  for helping to distinguish how clumpy the dust distribution is in the shock structure. The flux in the FIR is different for the two models (diffuse or clumpy, see Fig.~\\ref{Fig_dust_spec_comp}). The spatial resolution of the PACS [$60-210 \\, \\mu$m] instrument on-board the \\textit{Herschel Space Observatory} at $70\\,\\mu$m is comparable to that of \\textit{Spitzer} at 24~$\\mu$m. It will thus be possible to extract the FIR emission from the shock with much more accuracy than with  \\textit{Spitzer}. This is needed to possibly break some degeneracies of the model and better constrain the key-parameters (e.g. dust size distribution).   The Stephan's Quintet galaxy-collision is a unique environment to study dust survival in shocks,  on galactic scales. The presence of dust in the SQ shock stresses the need to revisit  the standard calculations of dust survival timescales in the ISM, by taking into account its multiphase structure."
    },
    "1004/1004.1671_arXiv.txt": {
        "abstract": "We describe deep, new, wide-field radio continuum observations of the Great Observatories Origins Deep Survey -- North (GOODS-N) field.  The resulting map has a synthesized beamsize of $\\sim$1.7\\arcsec\\ and an r.m.s.\\ noise level of $\\sim$3.9\\,$\\mu$Jy\\,beam$^{-1}$ near its center and $\\sim$8\\,$\\mu$Jy\\,beam$^{-1}$ at 15\\arcmin\\, from phase center. We have cataloged 1,230 discrete radio emitters, within a 40\\arcmin\\,$\\times$\\,40\\arcmin region, above a 5-$\\sigma$ detection threshold of $\\sim$20\\,$\\mu$Jy at the field center. New techniques, pioneered by \\cite{owe08}, have enabled us to achieve a dynamic range of 6800:1 in a field that has significantly strong confusing sources. We compare the 1.4-GHz (20-cm) source counts with those from other published radio surveys. Our differential counts are nearly Euclidean below 100\\,$\\mu$Jy with a median source diameter of $\\sim$1.2\\arcsec. This adds to the evidence presented by \\cite{owe08} that the natural confusion limit may lie near 1\\,$\\mu$Jy.  If the Euclidean slope of the counts continues down to the natural confusion limit as an extrapolation of our log N - log S, this indicates that the cutoff must be fairly sharp below 1 $\\mu$Jy else the cosmic microwave background temperature would increase above 2.7K at 1.4 GHz. ",
        "introduction": "The GOODS-N field \\citep{dic03,gia04} covers $\\approx160$\\,arcmin$^{2}$ centered on the Hubble Deep Field North \\citep{wil96} and is unrivaled in terms of its ancillary data. These include extremely deep {\\em Chandra}, {\\em Hubble Space Telescope} and {\\em Spitzer} observations, deep $UBVRIJHK$ ground-based imaging and $\\sim$3,500 spectroscopic redshifts from 8--10-m telescopes. Previous radio observations of this region, however, fell  short of complementing this unique dataset.  \\begin{deluxetable}{cccc} \\tablecolumns{4} \\tablewidth{0pt} \\tablecaption{GOODS-North VLA observing log\\label{TABLE1}} \\tablenum{1} \\pagestyle{empty} \\tablehead{ \\colhead{Observing dates}& \\colhead{Integration time (h)}& \\colhead{NRAO project code}& \\colhead{Configuration}} \\startdata  Nov-Dec 1996 & 42  & AR368& A\\\\ Feb 2005     & 28  & AM825& B\\\\ Aug 2005     & 7   & AM825& C\\\\ Dec 2005     & 2   & AM825& D\\\\ Feb-Apr 2006 & 86  & AM857& A\\\\ \\enddata \\end{deluxetable}  Radio emission is a relatively unbiased tracer of star formation and can probe heavily obscured active galactic nuclei (AGN) -- objects that are missed by even the deepest X-ray surveys. Radio observations thus allow us to fully exploit the wealth of data taken at X-ray--through--millimeter wavelengths, providing a unique extinction-free probe of galaxy growth and evolution through the detection of starbursts and AGN.  The recent imaging of \\cite{owe08} (OM08) ($\\sigma = 2.7\\,\\mu$Jy\\,beam$^{-1}$ at 1.4\\,GHz) have shown that the techniques exist to make radio images that approach the theoretical noise limit. To this end, we have obtained new, deep radio imaging of the GOODS-N field.  \\begin{figure*}[!ht] \\begin{center} \\epsscale{2.0}\\plotone{figure1.eps} \\caption{Contours of constant r.m.s.\\ noise after correction for the shape of the primary beam. The GOODS-N area covered by {\\em HST}/ACS is outlined near the field center, with an average sensitivity of $\\sim$3.9\\,$\\mu$Jy\\,beam$^{-1}$.  At a radius of $\\sim$15\\arcmin\\ (roughly the half-power point of the primary beam) the r.m.s.\\ sensitivity drops to $\\sim$7\\,$\\mu$Jy\\,beam$^{-1}$. The strong inward deflection of the 7-$\\mu$Jy and 10-$\\mu$Jy contours is due to the 265-mJy radio source discussed in \\S\\ref{DynRange}.}\\label{fig1} \\end{center} \\end{figure*}  While GOODS-N was selected to be free from bright sources at optical wavelengths, the field contains several very bright radio sources which place severe limitations on the dynamic range that can be obtained and hence the ultimate sensitivity of the radio map. Before new techniques were developed to deal with this issue, only moderately-deep radio imaging was possible \\citep{ric00}.  The earliest VLA data in the GOODS-N field were reprocessed using new techniques by \\cite{big06} and \\cite{mor08}, achieving a noise level of 5.7--5.8\\,$\\mu$Jy\\,beam$^{-1}$ -- a 23-25\\% improvement on the original map by \\citet{ric00}, and close to the theoretical noise limit when one considers the increase in system temperature at the low elevations permitted during the original observations.  While the reduction of \\cite{big06} provided improved access to the $\\mu$Jy radio population, even deeper radio imaging is required to properly complement the extremely deep GOODS {\\em Spitzer} mid-infrared data, as well as forthcoming deep observations at far-infrared and submillimeter wavelengths from {\\em Herschel}, SCUBA-2, the Large Millimeter Telescope, and other facilities. To this end, we have added 123\\,hr to the existing data. The reduced full resolution image (beam=1.7\\arcsec) and its rms map are available online\\footnote{http:/www.ifa.hawaii.edu/$\\sim$morrison/GOODSN and at the NASA/IPAC Infrared Science Archive (IRSA) http://irsa.ipac.caltech.edu/ as an ancillary data product associated with the GOODS Spitzer Legacy survey}.   The paper is laid out as follows: in \\S2 we describe the observations and the reduction of the data. In \\S3 we discuss the cataloging of radio emitters. \\S4 contains the results and some discussion of the catalog. We present our conclusions in \\S5. ",
        "conclusions": "We have presented a deep, new, high-resolution radio image of the GOODS-N field and a catalog containing 1,230 secure radio sources in the central 40\\arcmin\\,$\\times$\\,40\\arcmin\\ region.  The image reaches an r.m.s.\\ of $\\sim$3.9\\,$\\mu$Jy\\,beam$^{-1}$ in the best central region. It is thus one of the deepest 1.4-GHz radio surveys ever undertaken, well matched to the excellent {\\em Spitzer}, {\\em Hubble Space Telescope} and {\\em Chandra} data in GOODS-N, and to the upcoming {\\em Herschel}, SCUBA-2, and the Large Millimeter Telescope imaging at 100--850\\,$\\mu$m and  $\\sim$1.1-4mm, respectively, providing an unbiased probe of star formation out to $z\\sim3$ and a means to test the calibration of other star-formation rate indicators.  Our analysis of the source size distribution and the counts provide further evidence that the faint 1.4-GHz emitters have a median angular size of $\\sim$1.2\\arcsec, that log $N$--log $S$ remains flat below 100\\,$\\mu$Jy, and that the natural confusion limit at 1.4\\,GHz must be near 1\\,$\\mu$Jy -- a prediction that can be tested in the coming years using the Expanded VLA (EVLA) and e-MERLIN."
    },
    "1004/1004.3438_arXiv.txt": {
        "abstract": "The influence of the convective structure of the solar photosphere on the shifts of spectral lines of iron was studied. Line profiles in the visible and infrared spectrum were synthesized with the use of 2-D time-dependent hydrodynamic solar model atmospheres. The dependence of line shifts on excitation potential, wavelength, and line strength was analyzed, along with the depression contribution functions. The line shifts were found to depend on the location of the line formation region in convective cells and the difference between the line depression contributions from  granules and intergranular lanes. In visible spectrum the weak and moderate lines are formed deep in the photosphere. Their effective line formation region is located in the central parts of granules, which  make the major contribution to the absorption of spatially unresolved lines. The cores of strong lines are formed in upper photospheric layers where is  formed  reversed granulation due to convection reversal and physical conditions change drastically there. As a consequence the depression contributions in the strong line from intergranular lanes with downflows substantially increase. This accounts for smaller blue shifts of strong lines. In infrared spectrum the observed decrease in the blue line shifts is explained by the fact that their effective line formation regions lie higher in the photosphere and extend much further into the reversed granulation region due to the opacity rise with the increase of line wavelength.  Additionally the effective line formation depths of the synthesized visible and infrared Fe~I lines and their dependence on line parameters is discussed. ",
        "introduction": "If the motion of the Sun-Earth system is taken into account as well as  the correction for the gravitational red shift (636 m/s) caused by the difference between the gravitational potential of the solar atmosphere and the Earth is introduced, the  most absorption lines observed at the center of the solar disk are shifted  towards shorter wavelengths  with respect to laboratory lines. The observed blueshifts of the solar iron lines are 300--400 m/s on the average \\cite{9,1,11,12,13,26,28,32}. They were measured to an accuracy of 100 m/s in the range from 0 to 1 km/s. It is found that the blueshifts decrease with growing height of line formation as suggested by its dependence on excitation potential.  The largest blueshifts are observed in weak lines with high excitation potentials, while the smallest blueshifts are in strong lines. The redshifts were found to be small in some strong lines in the visible spectrum, but they can be as large as 1 km/s in the ultraviolet lines in the range 195--200 nm \\cite{32}. Some ``reddening'' of blueshifts was also noted in active photospheric regions  \\cite{1,11}. The blueshifts observed in weak and moderate lines can be explained by a decrease with height in the velocity of the convective motions along the line of sight towards the observer  \\cite{12,13}. The formation region of strong photospheric lines extends to the layers of  reversed granulation, a pattern comparable  to the granulation but with reversed brightness modulation. The matter at the center of convection cells becomes cooler than in descending intergranular lanes, but it still moves upwards. An inverse correlation between brightness and shifts was found in the cores of very strong lines in the visible and infrared spectrum  \\cite{22}. The effects of convection overshoot in the solar photosphere  \\cite{27} is believed to be the cause of the observed line redshifts.  The objective of this study is the effect of reversed granulation and the role of granules and intergranular lanes in the formation of spatially unresolved spectral lines and the cause of the ``reddening'' of blueshifts of strong lines. Such a study is possible today due to the remarkable progress in the simulation  of solar granulation and its evolution \\cite{10,14,27,30,33} and in the development of the idea of depression contribution functions  \\cite{19,20,23,25}. Based on 2-D hydrodynamic model atmospheres  \\cite{18}, we synthesized Fe~I lines and calculated their shifts. We used contribution functions to investigate the region of line formation in convection cells and to analyze the influence of physical conditions in the granular photosphere on the shifts of the synthesized lines. ",
        "conclusions": "We synthesized the Fe I lines in the framework of the 2-D HD models of solar granulation and determined the wavelength shifts of these lines and their dependence on the  line parameters. The derived relations fit quite well the observation data. Additionally we calculated the depression contribution functions to demonstrate that the line shifts are determined by the location of effective line formation region in convection cells. Obtained  results  confirmed the convective character of the line shifts.  Analysis of the line profiles and the corresponding contribution functions showed that the contribution from granules and intergranular lanes to the depression of spatially unresolved spectral lines is essentially different for weak and strong lines as well as for lines in the visible and infrared ranges. In the visible range weak and moderate lines originate deep in the photosphere, where the classical convection structure is not disrupted, i.e., hot masses move upwards and cool masses descend. Cores of strong iron lines are formed in the upper photosphere, in the convection reversal zone, where  occurs a temperature inversion as well as granulation inversion. The temperature in the ascending central regions of a convection cell is lower than in descending regions. The temperature inversion increases the role of intergranular lanes in the formation of the depression of the spatially averaged line. As a result, the depression contribution from the regions with downflows grows, and the line blueshifts diminish to zero or even become red. In the infrared spectrum the effect of reserved granulation on the line formation is stronger.  Due to growth of line opacity with wavelength the level of effective line formation is shifted higher into the photosphere  where are manifested the reserved granulation zone. Therefore the line blueshifts in the infrared spectrum are smaller than in the visible spectrum.  Convective motions cause the asymmetry  of spectral line profiles. The blue wings and cores of spectral lines are mainly formed in the central parts of granules, where upflows prevail. The  red wings are mainly formed in intergranular lanes with downflows. Greater  contributions from  intergranular lanes in the depression of  strong lines results in the strengthening of red wings of absorption lines and consequently a greater asymmetry, as compared to weak and moderate lines. Therefore the red wing of a spectral line can be regarded as an indicator of the additional absorption  in intergranular lanes.  We have found that the weak infrared Fe~I lines with $\\lambda \\approx$~1600 nm and very high excitation potentials (about 6.4 eV) are formed at the deepest photospheric layers. Such lines may be of prime importance in diagnostics of stellar atmospheres."
    },
    "1004/1004.1117_arXiv.txt": {
        "abstract": " ",
        "introduction": "In an high-mass X-ray binary system (HMXB), the compact object is embedded in its companion's stellar wind and accretes from it. This environment can modulate the observed \\mbox{X-ray} emission if our line of sight probes different parts of the wind -- differing in, e.g., density and ionization state -- along the binary orbit. As the accretion flow may depend on these properties of the wind -- \\emph{both} are likely to differ, e.g., between the low/hard and high/soft states (e.g., Smith et al.~2002; Gies et al.~2003) -- a thorough understanding of the wind is required for a complete picture of HMXBs and their state(-transition)s.  Due to its persistent brightness, the HMXB Cyg\\,X-1 allows for a detailed study of its wind. Based on a \\Chandra-HETGS observation in 2003, Hanke et al.~(2009) report a highly photoionized wind seen at orbital phase $\\phi$\\,=\\,\\,0.93--0.03 in the low/hard state; the high-resolution spectrum shows a multitude of absorption lines of H- and He-like ions. Inhomogeneities in the wind -- dense clumps at lower ionization stage -- lead to absorption dips in the soft X-ray light curve (Hanke et al. in prep.).  While \\Chandra-HETGS is suitable for narrow absorption lines, other instruments are required to measure the broad-band spectrum. Due to the violent variability during absorption dips, \\emph{simultaneous} observations are indispensable. We have therefore scheduled a joint multi-satellite observation (comprising \\emph{all} X-/$\\gamma$-ray missions) on 2009 April 18/19 (MJD\\,54\\,574/5) at $\\phi$\\,=\\,\\,0.97--0.21. At this time, Cyg\\,X-1 was in the low/hard state as well (see also Hanke et al.~2008).  \\begin{figure}\\centering \\psbox[xsize=8cm]{Hanke-Manfred_fig_1.ps} \\caption{Scattering of hard X-rays in the wind: The hard X-ray light curves of \\Suzaku-PIN (blue) and \\INTEGRAL-IBIS (black) show a scattering trough around $\\phi$=0.03, indicated by a linear fit (red). The lowest count rate of this fit is 69\\,\\% of its highest count rate. } \\label{fig:scatteringTrough} \\vskip-.5\\baselineskip \\end{figure} ",
        "conclusions": ""
    },
    "1004/1004.3923_arXiv.txt": {
        "abstract": "Rate coefficients for state-to-state rotational transitions in CO induced by both para- and ortho-H$_2$ collisions are presented. The results were obtained using the close-coupling method and the coupled-states approximation, with the CO-H$_2$ interaction potential of Jankowski \\& Szalewicz (2005). Rate coefficients are presented for temperatures between 1 and 3000 K, and for CO($v=0,j$)  quenching from $j=1-40$ to all lower $j^\\prime$ levels. Comparisons with previous calculations using an earlier potential show some discrepancies, especially at low temperatures and for rotational transitions involving large $|\\Delta j|$. The differences in the well depths of the van der Waals interactions and the anisotropy of the two potential surfaces lead to different resonance structures in the energy dependence of the cross sections which influence the low temperature rate coefficients. Applications to far infrared observations of astrophysical environments are briefly discussed. ",
        "introduction": "A number of PESs \\citep{pris78, flower79, poul82, schin84, bacic85, jankow98, jankow05} have been developed for the CO-H$_2$ complex. One of the most accurate surfaces was given by \\citet{jankow98} who calculated a four-dimensional, rigid-rotor PES referred to as V$_{98}$. This surface was constructed using the symmetry-adapted perturbation theory (SAPT) method  with high-level electron correlation effects. An analytical fit of the {\\it ab initio}  PES has a global minimum of 109.3 cm$^{-1}$ at the intermolecular separation of 7.76~a$_0$  for the linear geometry with the C atom pointing toward H$_2$. \\citet{anto00} adopted V$_{98}$ in their calculations of  state-to-state cross sections for CO rotational excitation at collision energies  between 795 and 991 cm$^{-1}$  which were found to be in good agreement with their measurements. The V$_{98}$ PES was also utilized in second virial coefficient calculations for mixtures of hydrogen and CO \\citep{gott00}. A comparison with experimental data suggested that the van der Waals well of the V$_{98}$ PES is too deep by 4\\%-9\\%, though it represents an improvement over earlier surfaces. The V$_{98}$ surface has also been used in full coupled-channel cross section and rate coefficient calculations for rotationally inelastic scattering of CO by  para- and ortho-H$_2$ \\citep{flower01,meng01}. However, because the attractive well of V$_{98}$ was presumed to be too  deep, \\citet{meng01} modified the PES by merely multiplying the surface by a constant factor of 0.93 which was subsequently used in their scattering calculations.  A newer CO-H$_2$ PES was reported by \\citet{jankow05}, referred to as V$_{04}$. To improve the surface accuracy, they adopted the coupled-cluster method with single, double, and noniterative triple excitations [CCSD(T)] and the supermolecular approach. The V$_{04}$ surface was obtained on a five-dimensional grid including the dependence on the H-H separation, but with the CO molecule taken to be rigid with the C-O separation set to the value of the C-O distance averaged over the CO ground state vibrational wave function. The PES was then obtained by averaging over the intramolecular vibration of H$_2$ to yield a four-dimensional rigid-rotor surface. Similar to V$_{98}$, an analytical fit of the V$_{04}$ was presented which has a global minimum of 93.05 cm$^{-1}$, or 16.22 cm$^{-1}$ shallower than $V_{98}$. \\citet{jankow05} used V$_{04}$ to calculate the rovibrational energy levels for the para- and ortho-H$_2$-CO complex as well as the second virial coefficient. The rovibrational energies were found to agree with the experimental values of \\citet{mcke98} for para-H$_2$ to within 0.1 cm$^{-1}$ while a scaling of V$_{04}$, resulting in a deepening of the well depth by 4 cm$^{-1}$, was needed to match the measured second virial coefficients. V$_{04}$ may be too shallow though it is possible that the experimental second virial coefficients are systematically too low \\citep{jankow05}. Comparative studies of collisional cross sections using the two surfaces have been presented in our earlier work \\citep{yang061} and by \\citet{wern06}. In \\citet{yang061}, calculations with the V$_{04}$ surface gave cross sections in excellent agreement with the state-to-state measurements of \\citet{anto00}. ",
        "conclusions": "Rotational quenching of CO due to para- and ortho-H$_2$ collisions has been studied using an explicit  quantum-mechanical close-coupling approach and the coupled-states approximation on the potential surface, V$_{04}$, of \\citet{jankow05}.  State-to-state quenching rate coefficients for initial rotational levels $j_2=$1, 2, $\\cdots$, 40 of CO are obtained over a wide temperature range and available in tables formatted for astrophysical applications.  Resonances result in undulations in the temperature dependence of the rate coefficients with the amplitudes of the undulations typically decreasing with $j_2$. For temperatures less than $\\sim$50~K, the current state-to-state rotational quenching rate coefficients obtained with V$_{04}$ are found to depart from the results of \\citet{flower01} obtained with the earlier V$_{98}$ PES. The discrepanices are likely related to the differences in the well depth and anisotropy of the two potentials."
    },
    "1004/1004.3865_arXiv.txt": {
        "abstract": "We present an updated analysis of the M31 pixel lensing candidate event OAB-N2 previously reported in \\cite{novati09}. Here we take advantage of new data both astrometric and photometric. Astrometry: using archival 4m-KPNO and HST/WFPC2 data we perform a detailed analysis on the event source whose result, although not fully conclusive on the source magnitude determination, is confirmed by the following light curve photometry analysis. Photometry: first, unpublished WeCAPP data allows us to confirm OAB-N2, previously reported only as a viable candidate, as a well constrained pixel lensing event. Second, this photometry enables a detailed analysis in the event parameter space including the effects due to finite source size. The combined results of these analyses allow us to put a strong lower limit on the lens proper motion. This outcome favors the MACHO lensing hypothesis over self lensing for this  individual event and points the way toward distinguishing between the MACHO and self-lensing hypotheses from larger data sets. ",
        "introduction": "Gravitational microlensing is an established tool of research. The original motivation \\citep{pacz86} has been the search for dark matter in the form of massive compact halo objects (MACHOs) (for a recent discussion from a larger point of view on dark matter and gravitational lensing we refer to \\citealt{massey10}). The first lines of sight explored to this purpose have been those toward the Magellanic Clouds (LMC and SMC). Microlensing campaigns toward the Galactic center have then been used to constrain Galactic structure (\\citealt{moniez10} and references therein) and, more recently, for the detection of extra-solar planets (\\citealt{gaudi10,dominik10} and references therein).  Results on MACHOs toward the LMC and the SMC have been reported by the MACHO \\citep{macho00}, the EROS \\citep{eros07} and the OGLE collaboration \\citep{ogle09}. There is agreement that MACHOs are excluded as a major component of dark matter for a wide range of possible mass values (down to about $10^{-6}~\\mathrm{M}_\\odot$). There remains however a range, $(0.1-1)~\\mathrm{M}_\\odot$, in which the results of the MACHO collaboration, as confirmed by \\cite{bennett05}, indicate that MACHOs might compose about a $f=20\\%$ fraction of the Milky Way (hereafter MW) halo mass at least up to the distance of the LMC. This outcome is challenged by the results of the EROS and the OGLE collaborations. A possible contamination of the microlensing signal due to MACHOs, once the background of intrinsic variable stars that can mimic microlensing is excluded, is \\emph{self lensing}, wherein both the source and lens belong to some luminous population\\footnote{Hereafter we will refer to ``self lensing'', broadly speaking, for any microlensing event with a \\emph{stellar} lens, not necessarily belonging to the \\emph{same} stellar population of the source.}. (This possibility has been first considered in \\citealt{sahu94,gould95} and further studied by several authors, e.~g. \\citealt{gyuk00,mancini04}; we refer also to \\citealt{novati09b} for a specific analysis of the OGLE-II results, to \\citealt{novati06} for a study of the possible contribution of the LMC own dark matter halo to the observed MACHO events and to \\citealt{moniez10} for an updated review on this issue). Indeed, the fact that the only remaining allowed mass range for MACHOs, from the LMC analyses, corresponds to that of the stars that might also act as lenses may be indicative of some observational bias (we recall that the characteristics of the microlensing events, in particular their duration, depend on the lens mass). On the other hand, if the reported events are really to be attributed to MACHOs, a fortiori for such a sizeable fraction as that implied by the MACHO collaboration results, this might have some deeper astrophysical meaning.  Beyond the Magellanic Clouds, the next suitable target for a microlensing search is our nearby and similar galaxy of Andromeda, M31 \\citep{crotts92,agape93,jetzer94}. It allows one to explore a different line of sight through the MW halo; we can fully map M31's own dark matter halo (which is not possible for the MW, this being perhaps the most severe limitation of LMC studies); finally, the inclination of M31 is expected to give rise to a characteristic signature in the spatial distribution of M31 halo events, so facilitating their identification against the contamination of self-lensing events. Because of the distance of M31, the sources for microlensing events are no longer resolved objects, at least for ground-based telescopes, so that we enter the \\emph{pixel lensing} regime \\citep{gould96}. Several observational campaigns have already been undertaken along this line of sight (for a review we refer to \\citealt{novati10}). The results obtained up to now on the MACHO issue are controversial. The POINT-AGAPE collaboration claimed evidence of a MACHO signal in the same mass range indicated by the MACHO LMC analysis \\citep{novati05}. This outcome, however, has been challenged by the MEGA collaboration who concluded that the detected signal is compatible with the expected self-lensing rate \\citep{mega06}. Finally, the WeCAPP collaboration is expected to present the final analysis of their 11-year campaign\\footnote{A.~Riffeser and S.~Seitz, private communication.}.  The issue of self lensing for M31 microlensing is even more difficult than for the LMC for two principal reasons. First, the expected self-lensing rate is quite large with respect to MACHO lensing, at least in the inner part of M31 (in which both signals have the larger expected rate because of the huge number of available sources). The exact ratio depends not only on the unknown halo fraction in the form of MACHOs and on the MACHO mass but also on the not fully understood contribution of M31 \\emph{stellar} lensing. (In fact, a possible approach to deal with this problem is that followed in the already cited analyses of \\citealt{novati05} and \\citealt{mega06}, where statistical arguments were used for a full set of events. Interestingly, a main point of disagreement between these two analyses is the role played by self lensing with respect to MACHO lensing, as discussed also in \\citealt{ingrosso06,ingrosso07}). Second, the lack of knowledge of the source flux, specific for the ``pixel'' lensing regime, adds a further degeneracy in the event parameter space that makes  the disentangling of the two signals more difficult. In particular, this degeneracy implies that, rather than determining the Einstein time, $t_\\mathrm{E}$, the impact parameter, $u_0$, and the source flux $\\phi^*$, from the light curve  one can usually reliably estimate only the full-width-half-maximum duration, $t_{\\mathrm{FWHM}}=t_\\mathrm{E}\\,f(u_0)$ \\citep{gondolo99}, and the flux deviation at maximum amplification, $\\Delta\\phi=\\phi^* \\, A_\\mathrm{max}(u_0)$, usually expressed in term of magnitude, e.~g. $\\Delta R$ \\citep{gould96,wozniak_pacz97}.  This problem of interpretation of the lens nature, together with the overall small number of \\emph{reliable} microlensing candidate events reported toward M31 so far, make the issue of a correct astrophysical understanding of \\emph{single} events that may contain precious additional information particularly relevant. For M31 lensing this issue has been first emphasized by the WeCAPP collaboration in \\cite{arno06}. The WeCAPP collaboration, furthermore, presented an extremely detailed analysis of the microlensing event PA-S3/GL1 (first reported by POINT-AGAPE, \\citealt{paulin03,belokurov05}, and subsequently, \\citealt{wecapp03}, detected in their own data set by the WeCAPP collaboration itself). Also thanks to the excellent sampling along the bump assured by the two independent data sets, and on the basis of a study of the \\emph{source finite size effect} \\citep{witt_mao94}, \\cite{arno08} showed that this extremely bright microlensing event is more likely to be attributed to MACHO lensing than to self lensing. Interestingly, this would hold even though the event is located very near to the M31 center, $d\\sim 4'$, where the expected self-lensing rate is larger.  In fact, the first detailed analysis of a M31 pixel lensing event aimed at the characterization of the lens nature has been presented by the POINT-AGAPE collaboration for the event PA-N1 \\citep{point01}. With the source magnitude and color fixed thanks to an analysis of archival \\emph{Hubble Space Telescope} (HST) data, and based on a specific Monte Carlo simulation of the experiment, where the observed characteristics of the event had been used to delimit the possible parameter space (using in particular a lower limit for the lens proper motion), \\cite{point01} conclude that the PA-N1 lens is more likely a MACHO if the mass halo fraction in the form of compact halo object is above $20\\%$. The results of this analysis have been however challenged by the MEGA collaboration (who also reported this event as MEGA-ML16, \\citealt{mega06}), in particular by \\cite{mega05_hst} who argued, on the basis of new HST observations, against the source identification reported in \\cite{point01}.  Following the strong interest in the characterization of single events, in the present work we will present an updated analysis of the microlensing candidate event OAB-N2 first reported by the PLAN (Pixel Lensing ANdromeda)  collaboration in \\cite{novati09}. First, we present a new analysis related to the issue of the source identification and characterization using KPNO and HST archival images, Sect.~\\ref{sec:source}. The driving motivation for a new analysis comes from the significantly improved sampling along the flux deviation due to the access to unpublished data of the WeCAPP collaboration \\citep{wecapp01}. As we show, these new data allow us to firmly establish  the microlensing nature of this flux variation and to carry out an improved and more detailed analysis of the lensing parameter space with the inclusion of the finite source size effect. This allows us to put some constraints on the relative lens-source \\emph{proper motion} \\citep{gould94}, Sect.~\\ref{sec:light}. In Sect.~\\ref{sec:res} we discuss the outcomes of the analysis within the framework of the Monte Carlo simulation of our M31 pixel lensing experiment. The limits on the lens proper motion are here relevant as they can be used as a tool to distinguish bewteen the self lensing and MACHO lensing nature of this microlensing event. ",
        "conclusions": "In the present paper we have presented an updated analysis of the M31 microlensing candidate event OAB-N2 first reported by the PLAN collaboration in \\cite{novati09}. In this previous work, in particular, the insufficient sampling along the flux variation did not allow us to probe unambiguously the microlensing nature of this flux variation. Here we have taken advantage of newly available WeCAPP data to complete the light curve sampling. From a methodological perspective, we observe that the importance of merging different data sets for microlensing observation might look even trivial when considering that Galactic Bulge events are routinely observed by multiple telescopes when searching for planetary anomalies. The present analysis, however, shows the extent to which this is suitable, feasible and, as this case nicely shows, possibly essential also for M31 pixel lensing events. (Such a strategy, besides, is unavoidable in order to achieve the challenging purpose of a systematic research for exoplanets in M31, \\citealt{covone00,baltz01,chung06,ingrosso09}. A first attempt in this sense is being carried out by the ANGRSTOM collaboration \\citealt{kerins06,darnley07,kim07}. The usually short expected timescale, $t_\\mathrm{FWHM}$, makes however potentially important such a strategy also for ``normal'' single-lens events.).  A preliminary analysis on archival images of the 4m Mayall-KPNO telescope and of WFPC2/HST allowed us to reject the tentative identification of the OAB-N2 source made in \\cite{novati09} with a rather bright \\cite{massey06} M31 catalog source. The event source, in fact, appears to be a nearby fainter blended object on the KPNO image, as confirmed also by the HST analysis. Although we can not conclusively determine the source magnitude, because of blending and the huge background level, even this partial result is a first step toward a better characterization of this event. Indeed, the new light curve photometry analysis confirms that the source must be a fainter object than previously concluded.  The improved sampling along the OAB-N2 light curve then allows us to reach a few important conclusions. First, we find a strong confirmation of the microlensing nature of this flux variation. The flux variation characteristics, the large signal to noise ratio, the short duration and the position, within the inner $3'$ of M31 where the expected rate is larger, together with the further evidence provided by the joint analysis of two completely independent data sets, put this flux variation among the most convincing and interesting M31 microlensing events reported so far. Second, we have carried out a more detailed analysis in the event microlensing parameter space with the specific purpose to constrain the lens nature, whether due to self lensing or to MACHO lensing. The main outcome we get to is a \\emph{lower} limit for the relative source-lens proper motion, $\\mu$. As we have shown, this parameter, independent of  the lens mass and therefore less subject to model issues, not only may allow one to distinguish M31 lensing from MW lensing (where the main expected signal is MACHO lensing and for which $\\mu$ is significantly larger than for M31 lensing) but also among different M31 lens populations. Specifically, at 95\\% level, we find $\\mu  > 0.64~\\mathrm{km}\\,\\mathrm{s}^{-1}\\,\\mathrm{kpc}^{-1}$. Comparing to the output of the Monte Carlo simulation of the experiment we conclude that, if not to the MW halo population, this value is more compatible with M31 MACHO lensing rather than with self lensing, for which only about 2\\% of the overall population is expected to fulfill this constraint. Although the self lensing hypothesis can not be definitively ruled out for this single specific event so that we are not yet in a position to strongly discriminate self lensing from MACHO lensing, this analysis shows the extent to which the lens proper motion, with no other prior assumptions on MACHOs but their spatial distribution and kinematics, is a powerful tool toward this purpose.  As an intermediate step of the light curve fit $\\chi^2$-based analysis we have looked for possible signatures of finite size source effect (through a study of the normalized angular radius, $\\rho$) for which however we find only marginal indications. We recall that an evidence in this sense would imply that the lens belongs to some M31 populations. The robust outcome we get to is rather an \\emph{upper} limit for $\\rho$ and on the related quantity, the source radius crossing time, $t_*=\\rho\\,t_\\mathrm{E}$ which, together with an estimate, from the source magnitude and color, of the physical angular radius, finally gives us the aforementioned \\emph{lower} limit for the lens proper motion.  The outcome of this analysis is  affected by the lack of an accurate and precise estimate of the unresolved source magnitude and color, independent of the light curve fit, for which a focused proposal on HST data would be extremely welcomed.  In the framework of M31 pixel lensing, the present analysis of OAB-N2 is the third example, following those for the events PA-N1 \\citep{point01} and PA-S3/GL1 \\citep{arno08}, where a detailed investigation of the lensing parameter space allowed to draw conclusions on the lens nature. Interestingly, although in the present work with a worse confidence level with respect to that reported in \\cite{arno08} for PA-S3/GL1 and with the caveats discussed in the Introduction as for PA-N1, in all cases the analyses point toward the same outcome, namely that MACHO lensing should be preferred over self lensing."
    },
    "1004/1004.4442_arXiv.txt": {
        "abstract": "We observed the new X-ray transient and black-hole candidate XTE J1652$-$453 simultaneously with \\textit{XMM-Newton} and the \\textit{Rossi X-ray Timing Explorer} (\\textit{RXTE}). The observation was done during the decay of the 2009 outburst, when XTE J1652$-$453 was in the hard-intermediate state. The spectrum shows a strong and broad Fe emission line with an equivalent width of $\\sim 450$ eV. The profile is consistent with that of a line being produced by reflection off the accretion disc, broadened by relativistic effects close to the black hole. The best-fitting inner radius of the accretion disc is $\\sim 4$ gravitational radii. Assuming that the accretion disc is truncated at the radius of the innermost stable circular orbit, the black hole in XTE J1652$-$453 has a spin parameter of $\\sim 0.5$. The power spectrum of the \\textit{RXTE} observation has an additional variability component above 50 Hz, which is typical for the hard-intermediate state. No coherent quasi-periodic oscillations at low frequency are apparent in the power spectrum, which may imply that we view the system at a rather low inclination angle. ",
        "introduction": "\\label{sec:intro} Most binaries with a stellar-mass black-hole primary are X-ray transients in which matter is accreted from the companion star onto the black hole (BH). These systems are characterised by a sudden and rapid rise in intensity (reaching a maximum in about one to several weeks) followed by a gradual decline. The so-called outbursts typically last several months; eventually the source becomes too weak to be detected with current X-ray instruments \\citep[for typical light curves, see Fig. 4--9 in][]{Remillard06}. The large excursions in the intensity seen in X-ray transients are likely due to changes in the accretion rate onto the BH, triggered by disc instabilities due to hydrogen ionisation \\citep[see Section 2.2.1 in][]{Done07,Lasota01}.  During an outburst, not only the intensity, but also the X-ray spectrum and X-ray variability changes \\citep[and references therein]{Mendez97,McClintock2006,Belloni10}. In the bright phase of an outburst, at X-ray luminosities $L_{\\rm X} \\sim10^{37}$--$10^{39}$ erg s$^{-1}$, the X-ray spectrum is dominated by thermal emission below $\\sim 5$ keV originating from an optically thick and geometrically thin accretion disc, with a typical temperature of $\\sim 1$ keV, that extends down to the marginally stable orbit \\citep{Shakura73}. This phase is generally referred to as the soft or high/soft state, and is further characterised by a low level of variability, with typical fractional rms values of a few per cent. When the luminosity decreases to about 1--4 per cent of the Eddington luminosity \\citep[][]{Maccarone03}, the source makes a transition to the so-called hard or low/hard state, in which the X-ray spectrum is dominated by hard power-law like emission with a typical power-law photon index of 1.5--1.7. The high-energy emission is produced by inverse Compton scattering of photons by a population of hot electrons in a region referred to as the corona \\citep{Thorne75}. The hard state is further characterised by a high level of variability, with rms amplitudes of $> 20$ per cent, with strong quasi-periodic oscillations (QPO) at low frequencies \\citep[$\\sim 0.1$--10 Hz; e.g.,][]{Homan05}.  X-ray spectra of BH transients sometimes show broad Fe emission lines at around 6--7 keV \\citep[for a review see][]{Miller07}. These lines, from fluorescent-K$\\alpha$ emission of neutral or ionised Fe, are thought to be produced through irradiation of the disc by the source of hard photons. The Fe line is broadened by relativistic effects and becomes asymmetric \\citep{Fabian89,Laor91}, and its profile provides information on the disc properties, such as the inclination angle, the emissivity, and the inner radius. The latter can be used to measure the BH spin \\citep[e.g.,][]{Reynolds08}, assuming that the disc extends down to the innermost stable circular orbit (ISCO) in the soft states.  In this paper we report on the timing and spectral analysis of a contemporaneous \\textit{XMM-Newton} and \\textit{RXTE} observation of the new X-ray transient XTE J1652$-$453. This is the first (and so far the only) time that this source has been observed with an instrument with a good spectral resolution to study the Fe line which was previously detected in \\textit{RXTE} Proportianal Counter Array (PCA) data \\citepalias{ATel2107}. Furthermore, the broadband coverage of the simultaneous \\textit{XMM-Newton} and \\textit{RXTE} data allows to properly model the underlying continuum spectrum. In Section~\\ref{sec:obs&red} we describe the outburst evolution from \\textit{RXTE} observations, we give details about the \\textit{XMM-Newton} and \\textit{RXTE} observations, and we describe the steps that we followed to generate the power and energy spectra. We continue with the data analysis in Section~\\ref{sec:data analysis} where we test different line and continuum models, and where we consider several broadening mechanisms. In this section we also present our results from the spectral and timing analysis separately. In Section~\\ref{sec:discussion} we discuss our results in the context of the spectral state, disc geometry, and inferred black hole spin. In the appendix we discuss how the choice of background spectrum and abundances influences the continuum and line parameters inferred from the spectrum. \\subsection{The 2009 outburst of XTE J1652$-$453}\\label{sec:2009outburst} XTE J1652$-$453 was discovered in the \\textit{RXTE}/PCA Galactic bulge scan on June 28 2009 \\citepalias{ATel2107}. Follow-up observations with the \\textit{Swift} X-ray Telescope (XRT) and \\textit{RXTE}/PCA measured a spectrum consistent with a stellar-mass black hole in the canonical high/soft state \\citepalias{ATel2108,ATel2120}, with the spectrum below 10 keV dominated by emission from an accretion disc with a temperature of $\\sim 0.6$ keV. The \\textit{RXTE}/PCA spectrum also showed evidence of a relatively strong emission line at 6.4 keV, most likely originating from an Fe-K$\\alpha$ line. However, the line could not be studied in detail because of the limited spectral resolution of \\textit{RXTE}.  A near infrared (NIR) imaging observation of the field of XTE J1652$-$453 on July 11 2009, provided a possible NIR counterpart of XTE J1652$-$453 \\citep{ATel2125}. However, from NIR observations on July 15 and August 28 2009, the suggested NIR counterpart showed no significant variability, suggesting that the previous detection might have been an unrelated interloper star \\citep{ATel2190}. Radio observations with the Australia Telescope Compact Array on July 6, 7, and 13 2009, showed emission that may be indicative of the decay of an optically thin synchrotron event associated with the activation of XTE J1652$-$453 \\citep{ATel2135}.  On August 22 2009, we triggered an \\textit{XMM-Newton} observation of XTE J1652$-$453, simultaneously with \\textit{RXTE}, to study the Fe emission line and the disc properties in the decline of the outburst. In this paper we present the results of these observations.  Around September 10 2009, the Burst Alert Telescope (BAT) on-board \\textit{Swift} showed an increase (by a factor of 1.5--2) in the hard X-ray count rate (15--50 keV) of XTE J1652$-$453, indicating a transition to the hard state \\citep{ATel2219}. On September 14 2009, the \\textit{RXTE} spectrum was still consistent with a disc blackbody ($kT \\sim 0.6$ keV), a power law (index of $\\sim 1.8$), and a line feature at $\\sim 6$ keV. Later \\textit{RXTE} observations did not require a disc component, and no line feature was apparent \\citep{ATel2219}. ",
        "conclusions": "\\label{sec:discussion} The \\textit{XMM-Newton}/\\textit{RXTE} energy spectrum of XTE J1652$-$453 shows a broad and strong Fe emission line with an equivalent width of $\\sim 450$ eV, which is among the strongest ever seen from a black-hole candidate \\citep[see e.g.,][]{Rossi05,MillerChandra04}. The line is consistent with being produced by reflection off the accretion disc, broadened by relativistic effects. The power spectrum from \\textit{RXTE} data can be fitted with two Lorentzians below 10 Hz, and shows an additional variability component above 50 Hz. The simultaneous \\textit{XMM-Newton}/\\textit{RXTE} observations were performed when XTE J1652$-$453 was in the hard-intermediate state \\citep[HIMS;][]{Mendez97,Homan05}, as the outburst decayed from the disc dominated high/soft state towards the low/hard state. Our high-resolution, broadband observations substantially increase the number of good datasets from this state \\citep[the other is GX 339$-$4]{MillerChandra04}. The strong line as well as the additional component at high frequencies are both characteristic of the hard-intermediate state. \\citet{MillerChandra04} show a \\textit{Chandra} grating spectrum of the black-hole candidate GX 339$-$4, with a very similar continuum spectrum as XTE J1652$-$453, where the line has an equivalent width of $\\sim 600$ eV. \\citet{KleinWolt08} show that there is an additional noise component at high frequencies in the HIMS, although it is somewhat weaker than the component seen in our data. Power spectra taken in the HIMS also typically show the strongest, most coherent, low-frequency QPO \\citep[e.g.,][]{KleinWolt08}. The lack of any such coherent low-frequency QPO in our data may imply that we view the system at a rather low inclination angle. This is consistent with the inclination inferred from fitting the Fe line profile with relativistic models.%  Understanding the intermediate state is important as it marks the onset of the spectral transition, and the appearance of the radio jet. While there is still debate on the nature of the transition, most models currently explain it by the inner disc evaporating into a hot coronal flow \\citep[see e.g.,][]{Liu06,Meyer00}. The intermediate state is therefore the one in which the disc is only slightly recessed from the last stable orbit, so there is strong disc emission, but there is also strong hard X-ray flux from the re-emerging hot flow. This geometry also predicts that there should be strong reflection features, and that these should be strongly smeared by relativistic effects. As the outburst declines further, the models predict that the inner disc quickly recedes, being replaced by the hot flow. With the receding disc, the solid angle subtended by the optically thick disc decreases, the amount of reflection decreases, and the larger radius means that the line is less strongly smeared. These models then suggest that the strongest and broadest Fe lines should be seen in the intermediate states, where the disc is closest to the last stable orbit, but co-exists with enough of the hard X-ray emitting corona to produce the required illumination of the disc.%  From low-resolution (\\textit{RXTE}) spectral monitoring it is indeed seen that for the high/soft states, as there is very little hard X-ray emission illuminating the disc, the Fe lines are generally weak, while they appear strongly at the intermediate states, as coronal X-ray emission re-appear. The line width and strength then decrease as the outburst declines \\citep{Gilfanov10}. However, such low-resolution data cannot distinguish changes in the reflection broadening from changes in the ionisation state, and since the disc temperature and/or strength of the irradiation change, the ionisation state should also change. Even if there is no change in geometry, an increase in disc temperature gives an increase in ionisation state, and so the strength of the consequent reflection features increases \\citep{Ross93,Ross07}, as well as the line width increases due to increased Compton broadening \\citep[e.g.,][]{Young05}. Moderate to high resolution data show that the Fe lines are ionised in the intermediate state \\citep[e.g.,][and our data]{MillerChandra04}, and are predominantly neutral in the hardest low/hard state \\citep[e.g.,][]{Tomsick09,Hiemstra09}.%  The high-resolution, broadband spectrum of XTE J1652$-$453 allows us to separately constrain ionisation, solid angle and relativistic smearing. The ionisation is high with log$\\xi \\sim 3.5$, and the very broad Fe line requires strong relativistic effects. The fits using a self-consistent reflection model give $R_{\\rm in} \\simeq 4$ $GM/c^{2}$ corresponding to the last stable orbit for a black hole of spin $\\sim 0.5$. This is a lower limit since we expect the disc to be slightly recessed. None the less, the \\textit{RXTE} monitoring of XTE J1652$-$453 shows that the direct disc emission has similar temperature and normalisation as in the high/soft state earlier in the outburst \\citepalias{ATel2107,ATel2120}, supporting the idea that the disc is very close to the last stable orbit in our data.%  However, the fits with the self-consistent reflection model show clear negative residuals at $\\sim7.1$ keV. This could correspond to a blue-shifted absorption line, such as the ones that are often seen in high inclination systems \\citep[e.g.,][]{Kubota07}. This absorption feature is not obviously apparent in the raw data, and it is only required when the smeared \\textsc{reflionx} model is used since the intrinsic Compton broadening in this model predicts no sharp drop on the blue wing of the line \\citep{Done06}. The \\textsc{reflionx} model underestimates the Compton broadening as it calculates the temperature of the disc from irradiation only, for $kT$ of 0.1--0.3 keV \\citep{Ross93}. This is obviously lower than the observed inner disc temperature of 0.6 keV. The Compton broadened line width should be $\\sigma \\propto E_{\\rm line}\\sqrt{(2kT_{e})/(m_{e}c^2)}$ $\\sim 0.32$ keV for He-like Fe. This is negligible compared to the 1.7--2 keV broadening implied by $R_{\\rm in} \\sim 4$ $GM/c^{2}$, and should not significantly change the derived inner radius. The Compton broadening smooths out the sharp drop from the blue wing of the line even more, and increases the mismatch with the data at $\\sim 7.1$ keV. Since the model does not describe the data in detail, we should be cautious in interpreting the reflection parameters in detail.  We observed the black-hole candidate and X-ray transient XTE J1652$-$453 simultaneously with \\textit{XMM-Newton} and \\textit{RXTE}. During this observation the source was in the hard-intermediate state \\citep{Homan05}. We found a strong and asymmetric Fe emission line in the X-ray spectrum of XTE J1652$-$453, with an equivalent width of $\\sim 450$ eV, one of the strongest lines ever seen from a black-hole candidate. From fits to only \\textit{XMM-Newton}/PN spectrum, the equivalent width is slightly lower, $\\sim 380$ eV.  The line is also broad, giving an inner disc radius of $\\sim 4$ $GM/c^{2}$, from either a single \\textsc{Laor} line profile or from a self-consistent ionised reflection model. This implies a spin parameter of $\\sim 0.5$, which is a lower limit as the disc may be slightly recessed in this state. In any case, neither model fits the data well in detail. This is not at all surprising for the single \\textsc{Laor} line as the line must be accompanied by a reflected continuum. Be that as it may, the self-consistent ionised reflection model used here incorporates all the known effects such as multiple lines and edges from the different ionisation states of Fe (and other elements), together with Compton broadening \\citep{Ross05}, such that the line is intrinsically broad. The characteristic skewed profile expected from a single narrow line is smeared out, so there is no sharp drop in the model at the blue side of the line. However, the data clearly prefer such a feature, indicating that even the best current reflection models may not be able to describe the detailed shape of the Fe line seen in the spectrum of this state.  The Fourier power spectrum of the \\textit{RXTE} observation simultaneous with that of \\textit{XMM-Newton} shows additional variability above 50 Hz. No coherent QPOs at low frequency are accompanying this, which may suggest that we see the source at a fairly low inclination angle. The $3.4\\sigma$ (single trial) component at $\\sim 130$ Hz is too broad to be the high-frequency QPO, but in case it is, and assuming that it reflects the Keplerian frequency of the inner disc radius of $\\sim 4$ gravitational radii (inferred from the line profile), then the estimated upper limit on the black-hole mass is $\\sim 30$ $M_{\\odot}$. None the less, since the source is in the hard-intermediate state rather than in the soft-intermediate state where the high-frequency QPOs are occasionally seen \\citep{KleinWolt08}, this additional high-frequency component may be a related feature at lower frequencies and lower coherence, and it might be connected to the re-emergence of the jet."
    },
    "1004/1004.3330_arXiv.txt": {
        "abstract": "We model multiwavelength afterglow data from the short Gamma-Ray Burst (GRB) 090510 using a combined leptonic-hadronic model of synchrotron radiation from an adiabatic blast wave.  High energy, $\\gtrsim 100$~MeV, emission in our model is dominated by proton-synchrotron radiation, while electron-synchrotron radiation dominates in the X ray and ultraviolet wavelengths.  The collimation-corrected GRB energy, depending on the jet-break time, in this model could be as low as $3\\times 10^{51}$~erg but two orders of magnitude larger than the absolute $\\g$ ray energy.  We also calculated the opacities for electron-positron pair production by $\\g$ rays and found that TeV $\\g$ rays from proton-synchrotron radiation can escape the blast wave at early time, and their detection can provide evidence of a hadronic emission component dominating at high energies. ",
        "introduction": "Gamma-Ray Burst science has entered a new era with launch of the {\\em Fermi} $\\g$-ray space telescope.  The main instrument, Large Area Telescope (LAT), is more sensitive than any previous instrument in the 20~MeV -- 300~GeV range~\\citep{LAT}, whereas the {\\em Fermi} Gamma-ray Burst Monitor (GBM) observes the whole unocculted sky in the 8~keV -- 40~MeV range~\\citep{GBM}.  With the advent of the Burst Alert Telescope (BAT), X-Ray Telescope (XRT) and UV-Optical Telescope (UVOT) onboard the {\\em Swift} satellite~\\citep{Swift} it is now possible to obtain simultaneous multiwavelength data in the optical to multi-GeV $\\g$-ray energy range from GRBs.  GRB 090510 is the first GRB to provide data from simultaneous observtions by {\\em Fermi}~\\citep{grb090510_Fermi} and {\\em Swift}~\\citep{grb090510_Swift}, as well as by a couple of other satellites.  At a redshift $z=0.903\\pm 0.003$~\\citep{grb090510_redshift} the isotropic-equivalent $\\g$-ray energy release from this short GRB ($T_{90} \\lesssim 2$~s) is $E_{\\g,\\rm iso} = (1.08\\pm 0.06)\\times 10^{53}$~erg with a fluence of $(5.03\\pm 0.25)\\times 10^{-5}$~erg~cm$^{-2}$ in the 10~keV -- 30~GeV range~\\citep{grb090510_Fermi}.  {\\em Fermi} LAT detected long-lived emission up to $\\sim 200$~s after trigger ($T_0 =$~00:23:00~UT, 2009 May 10) in the $\\gtrsim 100$~MeV range.  While such high-energy emission, that is temporally extended beyond the keV -- MeV emission, was first detected in GRB 940217 by the {\\em Compton Gamma-Ray Observatory}~\\citep{hur94}, this feature is common to most GRBs detected with {\\em Fermi} LAT.  {\\em Swift} XRT and UVOT collected data from GRB 090510 between $\\sim T_0+97$~s and $T_0+1.9$~ks before an Earth Occultation (EO), and again after $T_0+5.1$~ks~\\citep{grb090510_Swift}. {\\em Swift} BAT collected most data within $T_{90} = 0.3\\pm 0.1$~s (15 -- 350~keV), and sparsely between $T_0+0.4$~s and $\\sim T_0+100$~s~\\citep{grb090510_Swift}.  Smooth temporal evolution of the flux, $F\\propto t^{-\\a}$, of the long-lived emission in {\\em Fermi} LAT ($\\a_\\g = 1.38\\pm 0.07$), and {\\em Swift} XRT ($\\a_{X,1} = 0.74\\pm 0.03$ before EO) and UVOT ($\\a_{O,1} = -0.50^{+0.11}_{-0.13}$ before EO) observations strongly suggest an afterglow origin~\\citep{grb090510_afterglow}.  Synchrotron radiation by shock-accelerated electrons in a decelerating GRB blast wave~\\citep{mr97,spn98} have successfully explained much of the broadband afterglow data at radio, optical, and X-ray frequencies in the pre-{\\em Fermi} era.  However fitting combined {\\em Fermi} and {\\em Swift} data from GRB 090510 with simple $e$-synchrotron model results in unusual parameter values, and most importantly it is difficult to reconcile the $F_\\nu \\propto t^{-\\alpha} \\nu^{-\\beta}$ temporal relations~\\citep{grb090510_afterglow}.  More complex scenarios have been proposed to model GRB 090510 data such as a radiative fireball in an $e^\\pm$ pair dominated environment~\\citep{ggn09,ggnc09}, adiabatic fireball in a low density medium and small magnetic field~\\citep{kb09a,kb09b,gmxf09}, and two component jet~\\citep{cgp09}.  Here we present a combined leptonic- and hadronic- afterglow model to fit multiwavelength data from GRB 090510.  Inclusion of ion acceleration and radiation in the GRB blast wave is a natural and simple extension of the $e$-synchrotron blast wave model, and has been discussed by a number of authors~\\citep{bd98,tot98a,zm01,wan09,rdf10}. We show that LAT emission in the $\\gtrsim 100$~MeV range is dominated by synchrotron radiation from protons accelerated in the external forward shock of a decelerating blast wave.  (Note that \\citet{rdf10} considered proton-synchrotron radiation from a coasting blast wave.) The XRT and UVOT light curves can be reasonably reproduced by synchrotron radiation from electrons accelerated in the same external forward shock.  We present the $e$- and ion- synchrotron afterglow model in Sec.\\ 2, compare this model with GRB 090510 afterglow data in Sec.\\ 3, and discuss our results in Sec.\\ 4. ",
        "conclusions": "We have fitted the LAT and XRT light curves from $p$- and $e$- synchrotron emissions, respectively, before the EO from an adiabatic blast wave in a constant density medium (Fig.~\\ref{fig:synchrotron}). The photon index for $\\gtrsim 100$~MeV emission is $1+\\b_\\g = 1.92\\pm 0.05$ for the proton index $k_2 = 2.84\\pm 0.09$, and is consistent within $1\\sigma$ of the measured value of the power-law component in the 0.9~s -- 1.0~s time interval~\\citep{grb090510_Fermi}.  At earlier time, 0.6~s -- 0.8~s, where the power-law component is significant and harder, the agreement is within $3\\sigma$.  However emission below $\\sim 20$~keV in that time interval, which we do not model, may contribute to the hardening of the power-law component.  Our model is compatible with photospheric emission which can dominate in the GBM and BAT range~\\citep[e.g.][]{rm06,gmxf09,twm10}.  The isotropic-equivalent $\\g$-ray energy release from GRB 090510 is 2--4 orders of magnitude larger than the typical range of $10^{49}$ - $10^{51}$ erg for short GRBs, derived from pre-{\\em Fermi} era data~\\citep{n07}.  Our model requires even larger, $10^{55}$ erg, of isotropic-equivalent energy release in the jet~\\citep[see, however,][]{fw01,ej05}.  Thus the jet of GRB 090510 must be strongly beamed for our model to be viable.  Depending on the jet opening angle derived in Sec.\\ 3 from our model fit, the absolute jet energy is in between $\\approx (3-7)\\times 10^{51}$~erg if the jet break takes place between 1.4 -- 5.1~ks.  For a jet break at $\\gtrsim 100$~ks, the absolute energy release is $\\gtrsim 6\\times 10^{52}$~erg.  Note that the inferred absolute energy of the {\\em Fermi}-LAT GRB sample is generally high, reaching $10^{53}$~erg in some cases~\\citep{c10}.  We also note that the microphysical parameter $\\epsilon_e \\approx 10^{-4}$ that we derive for GRB 090510 using the combined $p$- and $e$- synchrotron model, is much lower than the typically assumed value for short GRBs in the range $10^{-1}$ - $10^{-2}$ which is derived from long GRB data in the pre-{\\em Fermi} era using $e$-synchrotron model only~\\citep{n07}.  Finally, a confirm detection of large-amplitude and short-time variability, synchronized to the keV - MeV emission, in LAT data will argue against their afterglow origin~\\citep[see, e.g.,][]{dm99} as we modeled here.  A crucial test for the hadronic emission model in the afterglows of GRBs at redshift $\\lesssim 0.5$, thus avoiding absorption in the extragalactic background light~\\citep{rdf09,frd10}, may come from ground-based TeV telescopes such as MAGIC, VERITAS, HESS and HAWC.  An extrapolation of the LAT flux in the $\\gtrsim 200$~GeV range, where $e$-synchrotron radiation flux is negligible and internal $\\g\\g$ opacity is small (Fig.~\\ref{fig:opacity}), should produce similar spectra as in LAT and lower flux from $p$-synchrotron radiation.  On the contrary, a synchrotron-self-Compton model~\\citep[e.g.][]{dcm00,zm01} is expected to produce a spectral hardening and flux increase in the TeV range.  \\vskip 10pt  I thank {\\em Fermi} LAT and GBM, and {\\em Swift} team members, especially M.~De Pasquale, C.~D.~Dermer, J.~D.~Finke, N.~Gehrels, S.~ Guiriec, V.~Pelassa and F.~Piron, for providing data and useful comments.  This work is supported by the Office of Naval Research and NASA {\\em Fermi} Cycle II GI program."
    },
    "1004/1004.1711.txt": {
        "abstract": "\\noindent % We report on work to increase the number of well-measured Type Ia supernovae (\\sneia) at high redshifts. Light curves, including high signal-to-noise HST data, and spectra of six \\sneia that were discovered during 2001 are presented. Additionally, for the two \\sne with $z>1$, we present ground-based $J$-band photometry from Gemini and the VLT. These are among the most distant \\sneia for which ground based near-IR observations have been obtained. %, having different sensitivity and possibly other %systematic uncertainties than space measurements. We add these six \\sneia together with other data sets that have recently become available in the literature to the Union compilation \\citep{2008ApJ...686..749K}. We have made a number of refinements to the Union analysis chain, the most important ones being the refitting of all light curves with the \\saltii fitter and an improved handling of systematic errors. We call this new compilation, consisting of $\\totalsne$ supernovae, the \\compname compilation. The flat concordance $\\Lambda$CDM model remains an excellent fit to the \\compname data with the best fit constant equation of state parameter $w=\\bestfitw^{\\bestfitwustaterr}_{\\bestfitwdstaterr}\\mathrm{(stat)}^{\\bestfitwusyserr}_{\\bestfitwdsyserr} \\mathrm{(stat+sys\\ together)}$ for a flat universe, or $w=\\bestfitcurvew^{\\bestfitcurvewustaterr}_{\\bestfitcurvewdstaterr}\\mathrm{(stat)}^{\\bestfitcurvewusyserr}_{\\bestfitcurvewdsyserr} \\mathrm{(stat+sys\\ together)}$ with curvature. We also present improved constraints on $w(z)$. While no significant change in $w$ with redshift is detected, there is still considerable room for evolution in $w$. The strength of the constraints depend strongly on redshift. In particular, at $z \\gtrsim 1$, the existence and nature of dark energy are only weakly constrained by the data. %   The value of $w$ at $z \\gtrsim 1$ is well-enough constrained %    to suggest the existence of dark energy. ",
        "introduction": "Type Ia supernovae (\\sneia) are an excellent tool for probing the expansion history of the Universe. % % It was not long after it was discovered that these objects could be % used as standard candles that About a decade ago, combined observations of nearby and distant \\sneia led to the discovery of the accelerating universe \\citep{1998Natur.391...51P,1998ApJ...509...74G,1998ApJ...507...46S, 1998AJ....116.1009R,1999ApJ...517..565P}.  Following these pioneering efforts, the combined work of several different teams during the past decade has provided an impressive increase in both the total number of \\sneia and the quality of the individual measurements. At the high redshift end ($z \\gtrsim 1$), the Hubble Space Telescope (\\hst) has played a key role. It has successfully been used for high-precision optical and infrared follow-up of \\sne discovered from the ground \\citep{2003ApJ...598..102K,2003ApJ...594....1T,2004ApJ...602..571B,% 2009ApJ...700.1415N}, and, by using the Advanced Camera for Surveys (\\acs), to carry out both search and follow-up from space \\citep{2004ApJ...607..665R,2007ApJ...659...98R,2008ApJ...673..981K,dawson:2009}. At the same time, several large-scale ground-based projects have been populating the Hubble Diagram at lower redshifts. The Katzman Automatic Imaging Telescope \\citep{2001ASPC..246..121F}, the Nearby Supernova Factory \\citep{2006NewAR..50..436C}, the Center for Astrophysics \\sn group \\citep{2009ApJ...700..331H}, the Carnegie Supernova Project \\citep{2006PASP..118....2H,2010AJ....139..120F}, and the Palomar Transient Factory \\citep{2009PASP..121.1395L} are conducting searches and/or follow-up for \\sne at low redshifts ($z < 0.1$). The \\sn Legacy Survey (\\snls) \\citep{2006AA...447...31A} and \\essence \\citep{2007ApJ...666..674M,2007ApJ...666..694W} are building \\sn samples over the redshift interval $0.3 < z < 1.0$, and the \\sdss \\sn Survey \\citep{2008AJ....136.2306H,2009ApJS..185...32K} is building a \\sn sample over the redshift interval $0.1 < z < 0.3$, a redshift interval that has been relatively neglected in the past. These projects have discovered $\\sim700$~well-measured \\sne. The number of well-measured \\sne beyond $z\\sim1$ is approximately 20 and is comparatively small. % Defined as z>1 and surviving the Kowalski et al. cuts.  \\citet{2008ApJ...686..749K} (hereafter K08) provided a framework to analyze these and future datasets in a homogeneous manner and created a compilation, called the ``Union'' \\sneia compilation, of what was then the world's SN data sets. Recently, \\citet{2009ApJ...700.1097H} (hereafter H09) added a significant number of nearby \\sne to a subset of the ``Union'' set to create a new compilation, and similarly the \\sdss \\sn survey \\citep{2009ApJS..185...32K} (hereafter KS09) carried out an analysis of a compilation including their large intermediate-$z$ data set \\citep{2008AJ....136.2306H}. When combined with baryon acoustic oscillations \\citep{2005ApJ...633..560E}, the H09 compilation leads to an estimate of the equation of state parameter that is consistent with a cosmological constant while KS09 get significantly different results depending on which light curve fitter they use.  An important role for \\sneia beyond $z \\sim 1$, in addition to constraining the time evolution of $w$, is their power to constrain astrophysical effects that would systematically bias cosmological fits. Most evolutionary effects are expected to monotonically change with redshift and are not expected to mimic dark energy over the entire redshift interval over which SNe Ia can be observed. Evolutionary effects might also have additional detectable consequences, such as a shift in the average color of \\sneia or a change in the intrinsic dispersion about the best fit cosmology.  Interestingly, the most distant \\sneia in the Union compilation (defined here as \\sneia with $z\\gtrsim 1.1$) are almost all redder than the average color of \\sneia over the redshift interval $0.3 < z < 1.1$. The result is unexpected as bluer \\sneia at lower redshifts are also brighter \\citep{1998A&A...331..815T,2005A&A...443..781G} and should therefore be easier to detect at higher redshifts. Possible explanations for the redder than average colors of very distant \\sneia range from the technical, such as an incomplete understanding of the calibration of the instruments used for obtaining the high redshift data, to the more astrophysically interesting, such as a real lack of bluer \\sne at high redshifts. % We discuss why they could be redder, they still seem to land on the % the Hubble diagram, so I do not think we should mention that this % could have anything to do with evolution of the color-mag relation. % %% or evolution in the relationships used to standardize observed \\sne %% magnitudes.  The underlying assumption in using \\sneia in cosmology is that the luminosity of both near and distant events can be standardized with the same luminosity versus color and luminosity versus light curve shape relationships. While drifts in \\snia populations are expected from a combination of the preferential discovery of brighter \\sneia and changes in the mix of galaxy types with redshift \\citep{2007ApJ...667L..37H} --- effects that will affect different surveys by differing amounts --- a lack of evolution in these relationships with redshift has not been convincingly demonstrated given the precision of current data sets. This assumption needs to be continuously examined as larger and more precise \\snia data sets become available.  % %% This is not true since they are still red after these two technical %% issues have been corrected. %% % As we will discuss later in the paper, a combination of two technical % issues made these SNe redder that they really were.  %% Compute the mean and standard deviation of the offset. A common %% denominator is the NICMOS calibration. Mention this and the general %% opt-IR offset.  % Clearly, the reasons for the paucity of blue SNe at $z\\gtrsim 1.1$ is % needed if \\sne are to play an important role in future dark energy % experiments. The first step is to verify the veracity of the effect, % and this requires additional and independently measured sets of very % distant \\sneia.  In this paper, we report on work to increase the number of well-measured distant \\sneia by presenting \\sneia that were discovered in ground based searches during 2001 and then followed with WFPC2 on \\hst. Two of the new \\sneia are at $z \\sim 1.1$ and have high-quality ground-based infrared observations that were obtained with \\isaac on the \\vlt and \\niri on \\gemini. %These data can be used to shed light on the reasons for the unusual %red colors of \\sneia at $z \\sim 1.1$, because they do not rely on the %calibration of NICMOS, as all other \\sneia at $z \\sim 1.1$ do. This paper is the first paper in a series of papers that will provide a comparable sample of $z>1$ \\sneia to the \\sne now available in the literature. The \\sneia in this series of papers were discovered in 2001 (this paper), 2002 (Suzuki \\etal in preparation) and from 2005 to 2006 during the Supernova Cosmology Project (\\scp) cluster survey \\citep{dawson:2009}.  %Currently, no single \\snia survey contributes more than the whole, so %combining the results from all \\snia surveys still leads to an %improvement in the constraints that can be set on the properties of %dark energy. In addition to performing a homogeneous analysis of the %world sample of \\snia, K08 also studied the tension between different %surveys. It is a recipe that we will follow in this paper.  %In this paper we continue the work of building a high precision Hubble %Diagram at high redshift, by presenting six \\sneia that were %discovered in ground based searches and then followed with \\hst. We %also present ground based infrared observations of the \\sneia at %$z>1$. We will publish additional \\sneia at $z>1$ discovered by the %\\scp in \\citet{dawson:2009} and Suzuki et al. (in preparation).  %Not only is it important to discover \\sneia at high redshifts, but %also to be able to do obtain high precision follow-up photometry in %multiple filters. The statistical weight of each individual \\sn %depends on how accurate its maximum magnitude, light curve shape and %color are measured. Currently, the color uncertainty is the dominating %the error budget when these parameters are combined to form the %quantity that is used for probing cosmology.  The paper is organized as follows. In Section \\ref{sec:search}, we describe the \\sn search and the spectroscopic confirmation, while Sections \\ref{sec:followup}, \\ref{sec:groundphoto} and \\ref{sec:spacephoto} contain a description of the follow-up imaging and the \\sn photometry. The light curve fitting is described in Section \\ref{sec:lcfitting}. In Section \\ref{sec:union} we update the K08 analysis both by adding new data and by improving the analysis chain. The paper ends with a discussion and a summary.  %The paper is outlined as follows: In Section~\\ref{sec:search} we %describe the search and spectroscopic confirmation, while %Section~\\ref{sec:followup} covers the follow-up data. %Sections~\\ref{sec:groundphoto} and~\\ref{sec:spacephoto} present the %ground and space based photometry respectively. The light curve %fitting is described in section~\\ref{sec:lcfitting} and added to the %K08 sample in~\\ref{sec:union}. In Section~\\ref{sec:salt2} we discuss %the uncertainties from light curve fitting, and the paper is concluded %in section~\\ref{sec:conclusions}. % and~\\ref{sec:summary}. ",
        "conclusions": ""
    },
    "1004/1004.5385_arXiv.txt": {
        "abstract": "We present a new mechanism for slow-roll inflation based on higher dimensional supersymmetric gauge theory compactified to four dimensions with twisted (supersymmetry breaking) boundary conditions. These boundary conditions lead to a potential for directions in field space that would have been flat were supersymmetry preserved. For field values in these directions much larger than the supersymmetry-breaking scale, the flatness of the potential is nearly restored. Starting in this nearly flat region, inflation can occur as the theory relaxes towards the origin of field space. Near the origin, the potential becomes steep and the theory quickly descends to a confining gauge theory in which the inflaton does not exist as a particle. This confining gauge theory could be part of the  Standard Model (QCD) or a natural dark matter sector; we comment on various scenarios for reheating.  As a specific illustration of this mechanism, we discuss 4+1 dimensional maximally supersymmetric gauge theory on a circle with antiperiodic boundary conditions for fermions. When the theory is weakly coupled at the compactification scale, we calculate the inflaton potential directly in field theory by integrating out the heavy W-bosons and their superpartners. At strong coupling the model can be studied using a gravity dual, which realizes a new model of brane inflation on a non-supersymmetric throat geometry. Assuming there exists a UV completion that avoids the $\\eta$-problem, predictions from our model are consistent with present observations, and imply a small tensor-to-scalar ratio. ",
        "introduction": "Cosmic inflation \\cite{Guth:1980zm,Linde:1981mu,Albrecht:1982wi} provides an attractive explanation for the flatness and homogeneity of the observable universe and gives rise in a natural way to primordial inhomogeneities that lead to structure formation (for a recent review, see \\cite{Baumann:2009ds}). Many effective field theory models of inflation have been suggested, however, one would like to have some guiding principles that narrow the field and therefore allow more specific predictions. From the field theory point of view, an attractive condition to demand is naturalness: the parameters in the effective field theory should not require a significant amount of fine-tuning. Another possible set of constraints may arise by demanding that the theory has a consistent UV completion including gravity. One possibility for such a UV completion is string theory; if we assume that string theory gives the correct microscopic theory of our universe, then the correct model of inflation must arise in some way from string-theoretic degrees of freedom.  In this paper, we consider a new scenario for inflation that can be naturally embedded in string theory, but more generally can be understood as a mechanism to generate a naturally flat inflaton potential from supersymmetry-breaking compactification of higher-dimensional supersymmetric gauge theory. The model has a strongly-coupled version best described as inflation coming from wrapped branes on a non-supersymmetric throat geometry, and a weakly-coupled version that can be analyzed directly in field theory. For any coupling, the model has a natural embedding in string theory based on the low energy dynamics of D-branes or M5-branes. We begin by describing the basic idea in a general context. \\vskip 0.1 in \\noindent {\\bf The brane inflation scenario} \\vskip 0.1 in \\noindent Suppose that as part of some warped compactification of string theory, we have a supersymmetric throat geometry which includes a cycle upon which a brane is wrapped (it is also extended in the four non-compact directions). This is schematically depicted in \\figref{fig-throat}. Suppose further that there is a moduli space for the brane, so that it can move freely up and down the throat. Now, consider a related geometry for which we choose some different supersymmetry-breaking boundary conditions on this cycle, such that in the modified non-supersymmetric geometry, the cycle contracts to zero size at the tip of the throat. Far away from this tip, where the cycle is large, the geometry is very similar to the supersymmetric case, so the radial potential for the brane is almost flat. But eventually, the brane will roll down to the tip and disappear, giving all its energy to the other degrees of freedom in the theory when it contracts. Thus, the scalar field describing the radial position of the brane is a natural candidate for an inflaton field. \\vskip 0.1 in \\noindent {\\bf The field theory scenario} \\vskip 0.1 in \\noindent Consider a higher-dimensional supersymmetric field theory with a moduli space (i.e. flat directions in the potential) parameterized by some scalar fields. Now, compactify to 3+1 dimensions on a manifold for which supersymmetry is broken by boundary conditions (e.g. a circle with antiperiodic boundary conditions for fermions). Suppose that the compactified theory flows in the IR to a massive (confining) theory.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=0.7\\textwidth]{brane.pdf} \\end{center} \\caption{\\label{fig-throat} Geometrical picture of inflation in the strongly coupled regime.} \\end{figure}  In the uncompactified theory, if we take one of the moduli scalar fields $\\phi$ to have a large expectation value, then typically there will be some W-bosons (and their superpartners) that become very massive. With enough supersymmetry, integrating these out does not spoil the existence of a moduli space, so the contribution to the effective potential coming from the massive bosons must exactly cancel the contribution from the fermions. This remains true in a supersymmetric compactification but is generically modified by boundary conditions that break supersymmetry.  To give a concrete example, consider  a ``twisted\" circle compactification with periodic boundary conditions  for the bosons and  antiperiodic boundary conditions for the fermions on circle. The mode numbers for the Kaluza-Klein (KK) modes of the W-bosons are shifted relative to their fermionic superpartners, so the masses are \\bea M_b &=& \\sqrt{\\phi^2 + {n^2 \\over R^2}} ~, \\cr M_f &=& \\sqrt{\\phi^2 + {(n+ {1 \\over 2})^2 \\over R^2}}~, \\eea where $1/R = M_{KK}$ is the KK compactification scale.  Thus, there is no longer an exact cancelation in calculating the effective potential. However, boson and fermion modes with $n \\ll \\phi R$ will be almost degenerate. As a result, cancelations return for large $\\phi R$ and the potential becomes flat here.  Thus, scalar fields which would be moduli in a supersymmetric compactification give rise to fields that can act as inflaton fields in the twisted compactification. If these fields start with large values $\\phi \\gg 1/R$, they will roll slowly back towards the origin of moduli space. At the origin, we get to the theory which is massive in the IR, so there are no more light excitations. \\vskip 0.1 in \\noindent {\\bf A specific model} \\vskip 0.1 in \\noindent In the bulk of this paper, we will illustrate these two scenarios in a specific model. We will see that the brane inflation picture and the field theory picture represent strongly coupled and weakly coupled versions of the same mechanism.  The model we consider is maximally supersymmetric 4+1 dimensional $U(N)$ Yang-Mills theory compactified on a spatial circle of radius $R = 1/M_{KK}$ with antiperiodic boundary conditions for fermions \\cite{Witten:1998zw}.\\footnote{In the UV, this can be completed as the 5+1 dimensional $(0,2)$ CFT (the low-energy theory of M5-branes) compactified on a torus with periodic boundary conditions on one cycle and antiperiodic boundary conditions on the other cycle.}  The effective dimensionless 't Hooft coupling of the compactified theory is $\\lambda$.  We consider initial conditions in which one of the adjoint scalar fields is $X = diag(\\phi,0,\\cdots,0)$. With supersymmetry-preserving boundary conditions, the scalar  $\\phi$ would parameterize an exact flat direction in the potential. In the case of broken supersymmetry, a potential is generated in the quantum effective action, but this turns out to be quite flat for $\\phi \\gg M_{KK}$ . When the theory is weakly coupled at the Kaluza-Klein scale ($\\lambda \\ll 1$), we can compute the effective potential by integrating out the W-bosons and their superpartners at one loop. We find that the potential for large $\\phi$ (canonically normalized) takes the form \\be V(\\phi) = V_0 \\left(1- A (\\phi / \\phi_0)^2 e^{-{\\phi \\over \\phi_0}} \\right) \\; , \\ee which quickly approaches a constant for large $\\phi$.  At strong coupling ($\\lambda \\gg 1$), the field theory has a gravity dual description as type IIA string theory on a non-supersymmetric D4-brane throat geometry or M-theory on a non-supersymmetric M5-brane throat at very strong coupling.\\footnote{We will make these statements more precise in the body of the paper.} In either case, the circle direction corresponding to the antiperiodic compactified direction of the field theory contracts smoothly to a point at the IR end of the geometry. Assuming that $N$ is large, taking one scalar to have an expectation value corresponds to introducing a probe brane in this throat geometry which is extended along all the field theory directions, including the compact ones, as shown in figure \\figref{fig-throat}. This is the wrapped brane that we described in general above. The quantum effective potential for our scalar is then calculated using the effective action for this probe brane (Born-Infeld plus Chern-Simons type couplings to the form fields) in the throat geometry. In either the type IIA string theory regime or the M-theory regime, we find a potential that behaves for large $\\phi$ as \\be V(\\phi) = V_0(1-  \\phi_0^3/\\phi^3) \\; , \\ee again asymptotically approaching a constant.    Starting with the effective actions we derive and {\\it minimally coupling} to gravity,\\footnote{A matter field is minimally coupled to gravity if its action does not involve explicit curvature tensors.} it is straightforward to show that we can achieve slow-roll inflation with an acceptable spectrum of fluctuations for natural values of the parameters. For example, the predictions for the scalar spectral index $n_s$ range continuously from $n_s \\approx 0.967$ in the weakly coupled model to $n_s \\approx 0.973$ at strong coupling. The predictions and constraints for the minimally coupled model are summarized in \\tableref{table:summary}.  \\begin{table}[t] \\centering \\vskip 0.1 in \\begin{tabular}{|  c | c | } \\hline $\\lambda \\ll 1  $ & $\\lambda \\gg 1$  \\\\   \\hline   \\, & \\\\  $\\qquad n_s = 1 - {2 \\over {\\cal N}_{CMB}} \\sim 0.967 \\qquad$  &  $\\qquad n_s = 1 - {8 \\over 5 {\\cal N}_{CMB}} \\sim 0.973 \\qquad$ \\\\ & \\\\ $\\alpha_s = - {2 \\over {\\cal N}^2_{CMB}} \\sim -0.00055$  &  $\\alpha_s = - {8 \\over 5 {\\cal N}^2_{CMB}} \\sim -0.00044$ \\\\ & \\\\ $r  \\sim 8.6 \\times 10^{-12}  N^{-1}$   &  $r  \\sim 2.5 \\times 10^{-9}  (\\lambda N)^{3 \\over 7}$ \\\\ \\hline \\end{tabular} \\caption{Weak and strong coupling predictions for the scalar spectral index $n_s$, the running of the spectral index $\\alpha_s$, and the tensor/scalar ratio $r$. For $\\lambda \\gg 1$ we require $N \\gg 1$ for the analysis to be reliable, and $N \\lesssim 10^4 \\lambda^{-{3 \\over 5}}$ to obtain enough e-foldings. Numerical values are given assuming CMB perturbations at the pivot scale left the horizon at ${\\cal N}_{CMB} = 60$ e-foldings before the end of inflation. } \\label{table:summary} \\end{table}  On the other hand, it is well known that once the field theory is coupled to gravity and treated in an inflating spacetime, there will generically be $H^2 \\phi^2$ terms in the effective potential (here $H$ is the Hubble parameter during inflation) that can spoil slow-roll inflation. The coefficient of such a term will in general depend on the UV completion of our model (the full string compactification in the brane picture). Thus, before taking the predictions seriously, one should investigate whether there exists a more complete theory including gravity for which the coefficient of $H^2 \\phi^2$ is either naturally small or can be fine-tuned to an acceptable level. \\vskip 0.1 in \\noindent {\\bf The end of inflation} \\vskip 0.1 in \\noindent A particularly nice feature of our mechanism is that inflation ends naturally -- with no remaining light degrees of freedom in the inflationary sector -- without having to invoke any additional fields, antibranes, etc. In the string theory scenario, this occurs since the brane self-annihilates when it reaches the IR where the cycle on which the brane is wrapped contracts to a point. In the field theory scenario, the theory at the origin of moduli space is a confining theory, in which the confinement scale can be far below the scale of inflation.\\footnote{In this case, the theory will be heated well into its deconfined phase when inflation ends. In the string theory scenario, the scale of inflation is generally not high enough to deconfine the gauge theory dual to the throat geometry. For details see \\appref{app-reheat}.}  It is conceivable that the non-supersymmetric confining gauge theory left over at the end of inflation could be part the Standard Model (e.g. the Yang-Mills part of QCD), or perhaps a separate sector that gives rise to the observed dark matter. \\vskip 0.1 in \\noindent {\\bf Relation to previous work} \\vskip 0.1 in \\noindent There exist many previous discussions of inflation models based on supersymmetric gauge theory (see \\cite{Lyth:1998xn} for an early review) and of string theory models based on branes in warped throat geometries (for reviews of inflation in the context of string theory, see \\cite{Baumann:2009ni, HenryTye:2006uv, Burgess:2007pz, McAllister:2007bg}), but none employ the basic mechanism we have described. From a field theory point of view, the work \\cite{Inami:2009bs} is somewhat related in that it uses a higher-dimensional gauge theory with a supersymmetry-breaking compactification, but their candidate inflaton comes from the Wilson line of the higher-dimensional gauge field around the circle. In string theory, models based on wrapped branes have been discussed before \\cite{Becker:2007ui}, but here the wrapped branes already have a (relatively steep) potential in a supersymmetric throat geometry. The well-known brane-antibrane inflation model \\cite{Kachru:2003sx} is similar in that the inflaton field is the radial position of a brane in a warped throat, and the potential is generated only after supersymmetry breaking. In that case, the supersymmetry breaking occurs by the addition of an anti-brane at the tip of the throat, and inflation ends when the brane and antibrane annihilate. Our model is slightly simpler in that the only brane required is the one that gives rise to the inflaton. Another rather novel feature of our model relative to models studied previously is that we are able to carry out the analysis both in a weakly coupled limit, where we can use field theory methods, and at strong coupling, where we exploit the gravitational dual descriptions. \\vskip 0.1 in \\noindent {\\bf Outline} \\vskip 0.1 in \\noindent The outline of the paper is as follows. We begin in \\secref{sec-setup} by describing the field theory on which our model is based, and giving the gravity dual descriptions that may be used to analyze the theory at strong coupling. In \\secref{sec-actions}, we derive the effective action for our candidate inflaton field. We obtain the effective action by directly integrating out W-bosons at one loop and by a dual gravity calculation at strong coupling. The effective actions are in general not of canonical form, but we show in \\secref{sec-pot} that during any possible period of inflation, the theory is well-approximated by an effective action with canonical kinetic terms and a simple potential, which we derive in each limit. In \\secref{sec-analysis}, we write down the conditions that arise from demanding that slow-roll parameters are small, and derive the predictions for inflationary perturbations in terms of the parameters of our model, assuming a minimal coupling to gravity. In \\secref{sec-numbers}, we compare these predictions with observations, thereby constraining the model parameters. We find that our theory gives a viable model of inflation consistent with observations for a broad range of parameters, and summarize the final predictions as a function of the parameters once all observational constraints have been taken into account. In \\secref{sec-reheat}, we make some comments about the end of inflation for the two regimes of our model, and discuss a few scenarios for how this model may be integrated with the degrees of freedom of the Standard Model in a more complete theory. In \\secref{sec-eta}, we recall that a consistent  UV completion of our field theory may include terms beyond the minimal coupling to gravity that we have assumed, and comment on the potential ``$\\eta$-problem'' in which some of these terms may spoil slow-roll inflation. In \\secref{sec-related}, we discuss various ways in which our model may be generalized. In \\secref{sec-disc}, we make some concluding remarks and discuss future directions. ",
        "conclusions": "\\end{table}  On the other hand, it is well known that once the field theory is coupled to gravity and treated in an inflating spacetime, there will generically be $H^2 \\phi^2$ terms in the effective potential (here $H$ is the Hubble parameter during inflation) that can spoil slow-roll inflation. The coefficient of such a term will in general depend on the UV completion of our model (the full string compactification in the brane picture). Thus, before taking the predictions seriously, one should investigate whether there exists a more complete theory including gravity for which the coefficient of $H^2 \\phi^2$ is either naturally small or can be fine-tuned to an acceptable level. \\vskip 0.1 in \\noindent {\\bf The end of inflation} \\vskip 0.1 in \\noindent A particularly nice feature of our mechanism is that inflation ends naturally -- with no remaining light degrees of freedom in the inflationary sector -- without having to invoke any additional fields, antibranes, etc. In the string theory scenario, this occurs since the brane self-annihilates when it reaches the IR where the cycle on which the brane is wrapped contracts to a point. In the field theory scenario, the theory at the origin of moduli space is a confining theory, in which the confinement scale can be far below the scale of inflation.\\footnote{In this case, the theory will be heated well into its deconfined phase when inflation ends. In the string theory scenario, the scale of inflation is generally not high enough to deconfine the gauge theory dual to the throat geometry. For details see \\appref{app-reheat}.}  It is conceivable that the non-supersymmetric confining gauge theory left over at the end of inflation could be part the Standard Model (e.g. the Yang-Mills part of QCD), or perhaps a separate sector that gives rise to the observed dark matter. \\vskip 0.1 in \\noindent {\\bf Relation to previous work} \\vskip 0.1 in \\noindent There exist many previous discussions of inflation models based on supersymmetric gauge theory (see \\cite{Lyth:1998xn} for an early review) and of string theory models based on branes in warped throat geometries (for reviews of inflation in the context of string theory, see \\cite{Baumann:2009ni, HenryTye:2006uv, Burgess:2007pz, McAllister:2007bg}), but none employ the basic mechanism we have described. From a field theory point of view, the work \\cite{Inami:2009bs} is somewhat related in that it uses a higher-dimensional gauge theory with a supersymmetry-breaking compactification, but their candidate inflaton comes from the Wilson line of the higher-dimensional gauge field around the circle. In string theory, models based on wrapped branes have been discussed before \\cite{Becker:2007ui}, but here the wrapped branes already have a (relatively steep) potential in a supersymmetric throat geometry. The well-known brane-antibrane inflation model \\cite{Kachru:2003sx} is similar in that the inflaton field is the radial position of a brane in a warped throat, and the potential is generated only after supersymmetry breaking. In that case, the supersymmetry breaking occurs by the addition of an anti-brane at the tip of the throat, and inflation ends when the brane and antibrane annihilate. Our model is slightly simpler in that the only brane required is the one that gives rise to the inflaton. Another rather novel feature of our model relative to models studied previously is that we are able to carry out the analysis both in a weakly coupled limit, where we can use field theory methods, and at strong coupling, where we exploit the gravitational dual descriptions. \\vskip 0.1 in \\noindent {\\bf Outline} \\vskip 0.1 in \\noindent The outline of the paper is as follows. We begin in \\secref{sec-setup} by describing the field theory on which our model is based, and giving the gravity dual descriptions that may be used to analyze the theory at strong coupling. In \\secref{sec-actions}, we derive the effective action for our candidate inflaton field. We obtain the effective action by directly integrating out W-bosons at one loop and by a dual gravity calculation at strong coupling. The effective actions are in general not of canonical form, but we show in \\secref{sec-pot} that during any possible period of inflation, the theory is well-approximated by an effective action with canonical kinetic terms and a simple potential, which we derive in each limit. In \\secref{sec-analysis}, we write down the conditions that arise from demanding that slow-roll parameters are small, and derive the predictions for inflationary perturbations in terms of the parameters of our model, assuming a minimal coupling to gravity. In \\secref{sec-numbers}, we compare these predictions with observations, thereby constraining the model parameters. We find that our theory gives a viable model of inflation consistent with observations for a broad range of parameters, and summarize the final predictions as a function of the parameters once all observational constraints have been taken into account. In \\secref{sec-reheat}, we make some comments about the end of inflation for the two regimes of our model, and discuss a few scenarios for how this model may be integrated with the degrees of freedom of the Standard Model in a more complete theory. In \\secref{sec-eta}, we recall that a consistent  UV completion of our field theory may include terms beyond the minimal coupling to gravity that we have assumed, and comment on the potential ``$\\eta$-problem'' in which some of these terms may spoil slow-roll inflation. In \\secref{sec-related}, we discuss various ways in which our model may be generalized. In \\secref{sec-disc}, we make some concluding remarks and discuss future directions.  \\label{sec-disc}  In this paper, we have described a new mechanism for slow-roll inflation. To conclude, it is worth emphasizing a few of its special features: \\begin{itemize} \\item The model arises quite naturally with only two ingredients: supersymmetry and extra dimensions. In compactifying to four dimensions, supersymmetry-breaking boundary conditions are no less natural (and perhaps even more generic) than supersymmetry preserving boundary conditions. \\item These essential features are both crucial ingredients of string theory; thus the model arises quite naturally from string-theoretic degrees of freedom \\item As a brane-inflation model, it is very simple since it requires only a single brane on a natural throat geometry; no additional ingredients are required to ensure a graceful exit at the end of inflation. \\item The model can be treated analytically -- and provides a viable model of inflation -- over a very broad range of parameters. \\item There are no problematic light degrees of freedom in the inflationary sector left over at the end of inflation; the energy in the inflaton field is quickly transferred to other degrees of freedom at the end of inflation \\end{itemize}  Of course, as for most models of inflation involving supersymmetry, our scenario may require some modest fine-tuning (of order one part in one hundred) to avoid the $\\eta$-problem, as we have discussed in \\secref{sec-eta}. An interesting future direction would be to carry out an analysis of allowed perturbations in the UV part of our throat geometry (following \\cite{Baumann:2008kq, Baumann:2010sx}) to check that operators able to counteract the effects of the curvature couplings discussed in \\secref{sec-curv} can arise in the context of consistent compactifications of string theory. Other interesting future directions include fleshing out some of the reheating scenarios discussed in \\secref{sec-reheat}, and exploring some of the other models of section \\secref{sec-related}. Finally, it will be interesting to see whether our model continues to satisfy the observational constraints (particularly on $n_s$) once the more accurate data from the Planck satellite becomes available."
    },
    "1004/1004.2046_arXiv.txt": {
        "abstract": "The mounting evidence for anomalously large peculiar velocities in our Universe presents a challenge for the $\\Lambda$CDM paradigm. The recent estimates of the large scale bulk flow by Watkins {\\it et al.} are inconsistent at the nearly $3\\sigma$ level with $\\Lambda$CDM predictions. Meanwhile,  Lee and Komatsu have recently estimated that  the occurrence of high-velocity merging systems such as the Bullet Cluster (1E0657-57) is unlikely at a  $6.5 - 5.8\\sigma$ level, with an estimated probability between $3.3\\times 10^{-11}$ and $3.6\\times 10^{-9}$ in $\\Lambda$CDM cosmology. We show that these anomalies are alleviated in a broad class of infrared-modifed gravity theories, called brane-induced gravity, in which gravity becomes higher-dimensional at ultra large distances. These theories include additional scalar forces that enhance gravitational attraction and therefore speed up structure formation at late times and on sufficiently large scales. The peculiar velocities are enhanced by $24-34$\\% compared to standard gravity, with the maximal enhancement nearly consistent at the $2\\sigma$ level with bulk flow observations. The occurrence of the Bullet Cluster in these theories is $\\approx 10^4$ times more probable than in $\\Lambda$CDM cosmology. ",
        "introduction": "\\label{massgrav}  The defining feature of theories with infinite-volume extra dimensions, such as DGP and Cascading Gravity, is that 4$D$ gravity is mediated by a resonance graviton --- a continuum of massive states --- with general propagator \\be \\frac{1}{k^2 +m^2(k)}\\,. \\label{prop} \\ee In DGP, for instance, $m^2(k) = r_c^{-1}k$.  In real space, this gives a gravitational potential that interpolates from the $4D$ scaling, $1/r$, at short distances to the $5D$ scaling, $1/r^2$, at large distances, with a cross-over scale set by $r_c$. In Cascading Gravity, the soft mass term $m^2(k)$ is a more complicated function of $k$, involving multiple cross-over scales~\\cite{cascade1,cascade2}. For simplicity, we shall assume that all cross-over scales are comparable and denote this common scale by $r_c$.  Because $4D$ gravity is massive, each graviton has 5 polarizations: the usual 2 helicity-2 states of GR, 2 helicity-1 states and 1 helicity-0 degree of freedom. At distances $r\\ll r_c$, only the helicity-2 and helicity-0 degrees of freedom are relevant --- the helicity-1 states are very weakly coupled to matter and can be safely ignored. In the DGP model, the helicity-0 mode has a nice geometrical interpretation. It measures the extrinsic curvature of the brane in the extra dimension and is thus referred to as a brane-bending mode. Cascading Gravity theories have $D-5$ additional helicity-0 or scalar modes, accounting for the higher number of extra dimensions in which the brane can bend. This counting of degrees of freedom can alternatively be understood from a decomposition of the $D$-dimensional massless spin-2 representation~\\cite{cascade1,nonFP}.  These scalar modes couple to the trace of the stress-energy tensor of matter on the brane and, combined with the helicity-2 states, result in a one-graviton exchange amplitude between conserved sources having the tensor structure of $D$-dimensional massless gravity~\\cite{cascade2}: \\be {\\cal A} \\sim T^{\\mu\\nu} \\cdot \\frac{1}{k^2}\\cdot \\tilde{T}_{\\mu\\nu} - \\frac{1}{D-2}T^\\mu_{\\; \\mu} \\cdot \\frac{1}{k^2}\\cdot \\tilde{T}^\\nu_{\\; \\nu}\\,, \\label{amp} \\ee where we have neglected the graviton mass term for the scales of interest~\\cite{footnote1}. For non-relativistic sources, this corresponds to the modified Poisson equation \\be k^2\\Psi_{\\rm dyn} = -4\\pi G \\left(1 + \\frac{D-4}{D-2}\\right) \\rho \\,. \\label{poissonflat} \\ee In other words, the gravitational attraction is a factor of $1+\\frac{D-4}{D-2}$ stronger than in standard Newtonian gravity.  If~(\\ref{amp}) were valid in the solar system, the theory would already be ruled out by post-Newtonian constraints, for arbitrarily small $m$ --- this is the famous van Dam-Veltman-Zhakarov (VDVZ) discontinuity~\\cite{vDVZ}. The resolution, first proposed by Vainshtein for massive gravity~\\cite{vainshtein}, is that the weak-field/linear assumption implicit in~(\\ref{amp}) is actually a poor approximation for the scalar modes in the vicinity of astrophysical bodies or in the early universe. Instead, as shown explicitly in the DGP model~\\cite{resolve,resolveothers}, non-linearities in these modes in high-density regions result in them decoupling from matter, leading to an approximate recovery of Einstein gravity in, {\\it e.g.}, the solar system.  Let us start with the DGP case. On scales $\\ll H^{-1}$ (and thus $\\ll r_c$), we can neglect time derivatives relative to gradients of the graviton helicity-0 mode. This mode, denoted by $\\chi$, satisfies the approximate equation~\\cite{lue,Fabian,Roman} \\bea \\nonumber & & \\nabla^2 \\chi + \\frac{r_c^2}{3\\beta_{\\rm DGP}a^2} [ (\\nabla^2\\chi)^2 - (\\nabla_i\\nabla_j\\chi)(\\nabla^i\\nabla^j\\chi) ] \\\\ && \\qquad \\qquad\\qquad\\qquad= \\frac{8\\pi Ga^2\\delta\\rho}{3\\beta_{\\rm DGP}} \\,, \\label{pieqn} \\eea where the $\\nabla$'s denote spatial derivatives, and where \\be \\beta_{\\rm DGP} \\equiv 1 + 2 Hr_c\\left(1+\\frac{\\dot{H}}{3H^2}\\right)\\,. \\label{betaDGP} \\ee The overdot denotes a derivative with respect to proper time. The Vainshtein effect appears in two guises in the above. Firstly, in the early universe when $Hr_c \\gg 1$, the coupling to matter density becomes vanishingly small since $\\beta_{\\rm DGP} \\gg 1$. Secondly, even at late times when $Hr_c\\lsim 1$ and $\\beta_{\\rm DGP} \\approx1$, sufficiently large overdensities trigger non-linearities in $\\chi$ and result in its decoupling.  The analogue of~(\\ref{pieqn}) is not yet known for Cascading Gravity. There are multiple scalars in this case, each of which is expected to exhibit Vainshtein screening. There have not yet been any successful calculations that keep the non-linearities of all scalar modes. For the purpose of this paper, we shall take a phenomenological approach and assume that all cascading scalar degrees of freedom obey an equation of the form~(\\ref{pieqn}). This is consistent with our earlier assumption of a single cross-over scale $r_c$. In other words, we collectively denote the scalars by $\\chi$ and assume that they satisfy \\bea \\nonumber & & \\nabla^2 \\chi + \\frac{r_c^2}{3\\beta a^2} [ (\\nabla^2\\chi)^2 - (\\nabla_i\\nabla_j\\chi)(\\nabla^i\\nabla^j\\chi) ] \\\\ && \\qquad \\qquad\\qquad\\qquad= \\frac{8\\pi G a^2\\delta\\rho}{3\\beta}\\,. \\label{pieqn2} \\eea Using the scaling of the resulting deviation from GR in the solar system~\\cite{degrav}, we can infer that the cosmological deviation scales as $\\beta \\sim H^2r_c^2$. We then generalize~(\\ref{betaDGP}) to \\be \\beta \\equiv 1 + 2 H^2r_c^2\\left(1+\\frac{\\dot{H}}{3H^2}\\right)\\,. \\label{beta} \\ee This modification makes a sharper transition in time than~(\\ref{betaDGP}). This is related to the fact that gravity rapidly weakens on long length scales in models that can successfully degravitate the cosmological constant. This also means that whatever replaces~(\\ref{pieqn2}) in Cascading Gravity will likely exhibit even more efficient Vainshtein screening within collapsed objects. The effect of the $\\chi$ field on matter is through its contribution to the actual potential $\\Psi_{\\rm dyn}$ that moves particles: \\be \\Psi_{\\rm dyn} = \\Psi_{\\rm N} + \\frac{3}{2} \\(\\frac{D-4}{D-2} \\)\\chi\\,, \\label{dyn} \\ee where $\\Psi_{\\rm N}$ is the usual Newtonian potential. The $D$-dependence here is chosen to recover~(\\ref{poissonflat}) in the linearized limit. When performing N-body simulations, we solve~(\\ref{pieqn2}) using a multi-grid relaxation scheme similar to the one described in~\\cite{Fabian}. Our N-body code is an updated version of the one used in previous work~\\cite{markw}, revised to solve~(\\ref{pieqn2}) exactly and without resorting to the approximation of spherical symmetry. Meanwhile, we solve for $\\Psi_{\\rm N}$ in the usual way through Fourier transforms on a particle mesh. The code and further numerical results will be described in more detail elsewhere~\\cite{marcos}.  Most of our results are derived within linear theory, since this is a valid approximation for bulk flows over the scales of interest. In this regime, we can obtain a modified evolution equation for the density perturbations $\\delta\\equiv \\delta\\rho/\\bar{\\rho}$. Using the fact that the energy-momentum tensor on the brane is covariantly conserved, density and velocity perturbations evolve as usual via \\be \\dot{\\delta} = -\\frac{k}{a}v\\;;\\qquad \\dot{v} + Hv = \\frac{k}{a}\\Psi_{\\rm dyn}\\,. \\label{emomcons} \\ee These can be combined with the linearized version of~(\\ref{pieqn}) to yield \\be \\ddot{\\delta} + 2H\\dot{\\delta} = 4\\pi G \\bar{\\rho} \\left(1+\\beta^{-1}\\frac{D-4}{D-2}\\right)\\delta\\,. \\label{deltaevolve} \\ee Note that this is consistent with~(\\ref{poissonflat}) in the flat space limit $\\beta\\rightarrow 1$. As another quick check, letting $\\beta\\rightarrow\\beta_{\\rm DGP}$ and setting $D=5$, this agrees with explicit cosmological results in DGP~\\cite{lue}. In this paper, we will solve for peculiar velocities using linear theory and cross-check the results with simulations. ",
        "conclusions": "In this paper, we have explored how peculiar velocities are affected in a broad class of IR-modified gravity theories called brane-induced gravity. On the scales of interest, these theories admit a local $4D$ description in terms of weak-field gravity plus $D-4$ scalar fields coupled to the trace of the matter stress-tensor. These scalar degrees of freedom effectively strengthen the gravitational attraction at late times and speed up structure formation. As a result, peculiar velocities are systematically larger than those expected in standard gravity. Comparisons between N-body simulations and linear theory calculations show that linear theory gives an excellent description of the physics of bulk flows.  We have found that large-scale bulk flows can be enhanced up to $\\sim 40\\%$ relative to $\\Lambda$CDM cosmology. The enhancement is limited by observational constraints on $\\sigma_8$, the tightest limit coming from X-ray cluster abundance. The predicted peculiar velocities alleviate the tension with recent observations of the bulk flow on $50h^{-1}$~Mpc scales by Watkins {\\it et al.}~\\cite{hudson1,hudson2}, from a $\\approx 3\\sigma$ discrepancy in $\\Lambda$CDM gravity to a $\\approx 2\\sigma$ difference in Cascading Gravity. The agreement between theory and data is improved by $\\Delta \\chi^2 \\simeq -4.6$ in our model. Although modest, this improvement offers further motivation for more accurate bulk flow observations. Peculiar velocities are also enhanced on much larger scales ($\\sim 300h^{-1}$~Mpc) probed by kinetic Sunyaev-Zel'dovich observations, but this is insufficient to explain the enormous bulk flow inferred by~\\cite{kash1,kash2}.  Violent merging systems, such as the Bullet Cluster, are much more probable in Cascading Gravity. Drawing on the recent analysis of~\\cite{Lee:2010hj}, we have found that the occurrence of the Bullet Cluster in our theories is $\\gtrsim 10^4$ times more likely than in standard gravity.  We are pursuing various parallel tracks to improve upon the preliminary analysis presented here. As mentioned above, the tightest constraint on our models comes from cluster counts. Through N-body simulations and various semi-analytical techniques, we are currently determining the halo mass function in Cascading Gravity as a function of $D$ and $r_c$ \\cite{marcos}. This will allow a more accurate comparison with X-ray data. With regards to the Bullet Cluster, we can use similar simulations to calculate the distribution of initial velocities for subclusters, following the standard gravity analysis of~\\cite{Lee:2010hj}.  On the theoretical side, the calculations presented here rely on some phenomenological input. While we are confident that our results capture the essence of Cascading cosmology predictions, they do not derive from a rigorous cosmological analysis of the complete higher-dimensional theory. An important first step in this direction would be to obtain the decoupling limit of Cascading Gravity while keeping the non-linearities in all scalar degrees of freedom. This would allow a derivation of the modified Friedmann equation, as in~\\cite{galnic} for DGP, as well as the perturbation equations in the Newtonian regime. This is work in progress~\\cite{kurtmarkdan}.  The observations of large bulk flows and violent cluster mergers offer tantalizing evidence that structure is evolving more rapidly than predicted by $\\Lambda$CDM cosmology. As we extend our measurements of large-scale bulk flows beyond our local region and discover an increasing number of merging systems, it will become clear whether these are statistical flukes or the first indication of a new realm of gravitational physics on cosmological distances in the late universe.   \\vspace{0.5in}"
    },
    "1004/1004.4871_arXiv.txt": {
        "abstract": "Thermal emission during X-ray bursts is a powerful tool to determine neutron star masses and radii, if the Eddington flux and the apparent radius in the cooling tail  can be measured accurately, and distances to the sources are known. We propose here an improved method of determining the basic stellar parameters using the data from the cooling phase of photospheric radius expansion bursts covering a large range of luminosities. Because at that phase the blackbody apparent radius depends only on the spectral hardening factor (color-correction), we suggest to fit the theoretical dependences  of the color-correction versus flux in Eddington units to the observed variations of the inverse square root of the apparent blackbody radius with the flux. For that we use a large set of atmosphere models for burst luminosities varying by three orders of magnitude and for various chemical compositions and surface gravities. We show that   spectral variations observed during a long photospheric radius expansion burst from \\source\\ are entirely consistent with the theoretical expectations for the passively cooling neutron star atmospheres. Our method allows us to determine both the Eddington flux (which is found to be smaller than the touchdown flux by 15\\%) and the ratio of the stellar apparent radius to the distance much more reliably. We then find a lower limit on the neutron star radius of 14 km for masses below 2.2$\\msun$, independently of the chemical composition. These results suggest that the matter inside neutron stars is characterized by a stiff equation of state. We also find evidences in favour of hydrogen rich accreting matter and obtain an upper limit to the distance of 7 kpc. We  finally show that  the apparent blackbody emitting area  in the cooling tails of the short bursts from \\source\\ is two times smaller  than that for the long burst and their evolution  does not follow the theory. This makes their usage  for determination of the neutron star parameters questionable and casts serious doubts on the results of previous works that used  for the analysis similar bursts from other sources. ",
        "introduction": "Studies of the thermal emission from neutron stars (NSs) have been used extensively to determine their masses and radii \\citep[e.g.][]{Damen90,vP90,LvPT93,R02, HR06,WB07,OGP09,GO10}, which can provide constraints on the properties of the matter at supranuclear densities \\citep{HPY07,LP07}. Thermonuclear X-ray bursts at NS surfaces and, particularly, the photospheric radius expansion (PRE) bursts exceeding the Eddington limit at high-flux phases \\citep{Kuulkers03} are excellent laboratories for such detailed studies. The observed Eddington flux gives a constraint on the NS mass-radius relation. This method, however, suffers from the uncertainty in determination of the exact moment when the luminosity reaches the Eddington limit at the NS surface, which is often assumed to coincide with the moment of ``touchdown'', when  the measured color temperature is highest and the apparent radius is smallest \\citep{Damen90}. This interpretation is uncertain, as the touchdown flux typically  coincides with the  maximum flux \\citep{GOP08}, while the latter is expected to be larger  than the surface Eddington flux by the redshift factor $1+z$ \\citep{LvPT93}.   The second constraint can be obtained from the apparent radius of the NS at late stages of the burst. This method also has systematic uncertainties related to the color-correction factor $f_{\\rm c}=T_{\\rm c}/T_{\\rm eff}$ (the ratio of the color temperature to the effective temperature of the star),  which is a function of the burst luminosity.  To minimize the theoretical  uncertainties in modeling the burst atmospheres at nearly Eddington luminosities, we proposed \\citep{SPW11} to use the whole cooling track and check that the evolution of the blackbody normalization with flux is consistent with the theoretically predicted evolution for a passively cooling neutron star.  The unique determination of   mass and radius requires one more constraint. Knowing the distance to a burster breaks the degeneracy, and therefore bursters located in globular clusters would be of  interest. For the analysis in the present paper we choose \\source\\ which resides in the globular cluster Terzan 2. We analyze three PRE bursts from that source and show that they follow different cooling tracks. We also speculate on the reasons for such a discrepancy. We demonstrate that the apparent blackbody radius for the long PRE burst  evolves according to the predictions for the passively cooling neutron star with a constant apparent surface and use these data to determine the NS parameters.  \\newpage ",
        "conclusions": "We applied a recently computed detailed set of NS atmosphere models covering a large range of luminosities \\citep{SPW11} to the data of the PRE bursts of \\source.  We showed that the variation of the apparent blackbody radius during the cooling stage of the 150 s long PRE X-ray bursts in \\source\\ is entirely consistent with the theoretical  color-correction--flux  dependence at intermediate fluxes. We thus obtained the Eddington flux and the apparent NS radius (divided by the distance to the source). We find that the Eddington flux  (for Thomson opacity) is reached not at the so called ``touchdown'', but later at a 15\\% lower luminosity. We  constrained  mass and radius of the NS using the estimated distance to the source. We find a lower limit on the stellar radius of $\\sim$14 km for $M<2.3\\msun$ at 90\\% confidence independently of chemical composition. (If the NS in \\source\\ has a spin of 500 Hz, the radius of the non-rotating star would be about 1 km smaller.) Smaller radii are possible only for more massive NS. These results support a stiff equation of state of the NS matter. We showed that  hydrogen rich accreting matter is preferred  and  obtained an upper limit on the distance to \\source\\ of about 7 kpc.  We have also demonstrated that the cooling tracks of the short PRE bursts from \\source\\ that occurred during the high/soft state do not follow the evolution expected from theory and the NS apparent areas are a factor of two smaller than that for the long burst. The probable reason is the partial eclipse  of the NS surface by the optically thick accretion disk. An additional spectral hardening during the cooling tails associated with the influence of the boundary/spreading layer can be also important. Therefore, the constraints on the NS mass and radius obtained from such bursts are not  reliable.  Finally, we showed that previous analyses of the short PRE bursts from three sources EXO\\,1745--248, 4U\\,1820--30  and 4U\\,1608--52 are  questionable, because they ignore the fact that spectral evolution during the bursts is not consistent with the  theory for the passively cooling unobscured NS that is the base for the analysis. Assuming that the touchdown flux is reached when the photosphere is detached from the NS surface \\citep{SLB10} does not solve the problem because the determined blackbody normalization in the cooling tail and the Eddington flux are not reliable in these short bursts.  We suggest that only PRE bursts showing spectral evolution consistent with the theory should be used when estimating NS masses and radii.  Further improvement in accuracy of determination of the NS parameters requires new atmosphere models with the exact Compton scattering kernel, as well as accounting for the Doppler effect due to the rapid NS rotation. Rotation also introduces additional distortion to the NS spectra, because of the difference in the effective gravity at the stellar poles and the equator due to the centrifugal force. We plan to investigate the importance of these effects in future work."
    },
    "1004/1004.4106_arXiv.txt": {
        "abstract": "In the optical, the spectrum of symbiotic binaries consists of contributions from the cool giant, symbiotic nebula and the hot star. Strong emission lines are superposed on the continuum. In this paper we introduce a simple method to extract individual components of radiation from photometric $UBV$ magnitudes. We applied the method to classical symbiotic stars \\astrobj{AX Per}, \\astrobj{AG Dra}, \\astrobj{AG Peg} and \\astrobj{Z And}, the symbiotic novae \\astrobj{RR Tel} and \\astrobj{V1016 Cyg} and the classical nova \\astrobj{V1974 Cyg} during its nebular phase. We estimated the electron temperature and emission measure of the nebula in these systems and the $V$ magnitude of the giant in the symbiotic objects. Our results are in a good agreement with those obtained independently by a precious modelling the UV--IR SED. ",
        "introduction": "Symbiotic stars are long-period ($P_{orb}\\sim$ 1 -- 3 years or more) interacting binary systems, which comprise a late-type giant (usually a red giant of the spectral type M) and a hot compact star, most probably a white dwarf. If the giant of the spectral type G or K is present, we call them 'yellow symbiotics'. During quiescent phases, the white dwarf accretes a fraction of the wind from the giant. Typical mass loss rate from the giant in the symbiotic binaries is of a few times $10^{-7}$\\myr. The accretion process leads to heating up the white dwarf's surface to a very high temperature of $T_{h}\\sim 10^{5}$\\,K, and increases its luminosity to $L_{h} \\sim 10^{2}-10^{4}$\\lo. Such the hot and luminous WD is capable of ionizing neutral wind particles of both the stars giving rise to a strong nebular radiation. During quiescent phases, the symbiotic nebula represents mostly the ionized part of the stellar wind from the giant. During active phases, the mass loss rate from the hot component increases, and can temporary exceed that from the giant. As a result, properties of the symbiotic nebula significantly change during the active phases.  Accordingly, the observed flux, $F_{\\lambda}$ (corrected for the interstellar extinction) is given by the superposition of three basic components of radiation -- from the nebula, $F_{\\lambda}^{nebula}$, the cool giant, $F_{\\lambda}^{giant}$, and the hot stellar source, $F_{\\lambda}^{hot}$, i.e. \\begin{equation} F_{\\lambda} = F_{\\lambda}^{nebula} + F_{\\lambda}^{giant} + F_{\\lambda}^{hot}. \\label{eqn:totalflux} \\end{equation} Contributions of the individual components of the radiation depend on the wavelength. The hot stellar source dominates the far-$UV$ spectrum, while the nebula is usually dominant in the near-$UV$/$U$ domain, and the cool giant (especially the red giant) dominates the near-$IR$ spectral region. In the case of yellow symbiotics the contribution of the giant to the Johnson's $V$ filter is very strong relatively to that from the nebula. However, contributions from different sources in a symbiotic system depend also on the level of the activity.  During active phases of symbiotic binaries with a high orbital inclination, we observe narrow minima in their light curves. They can be interpreted as eclipses of the hot object by the cool giant. This implies that during the activity the hot stellar source produces a dominant amount of its radiation within the optical region. It is believed that an optically thick disk is formed around the hot star at the orbital plane during active phases \\citep[see][]{sk05}. Its relatively cool rim mimics a warm pseudophotosphere, while the circumstellar material above/below the disk can easily be ionized by its hot central part. Therefore the observed depths of the minima are partially filled in by the nebular radiation. Thus during the eclipses, the observed light consists of the contributions from the giant and the nebula. During quiescent phases the optical region is dominated by the extensive nebula, which thus prevents from observing eclipses. Corresponding colour indices are very different from those observed during active phases.  The aim of this paper is to propose a method of disentangling photometric $UBV$ magnitudes of symbiotic stars into their components from contributing sources as given by Eq. (\\ref {eqn:totalflux}). This means to determine physical parameters of their radiation. In particular, the electron temperature, $T_{\\rm e}$, and emission measure, $EM$, of the symbiotic nebula, and the brightness of the cool giant in the $V$ passband ($V^{giant}$). We describe our method in the following section. First, we formulated some simplifying assumptions (Sect.~2.1), and then we derived a system of equations, which relates theoretical properties of the radiative components to their observational characteristics (Sect.~2.4). In Sect.~3 we applied this approach to selected objects and compared the results with those achieved independently by another methods. Discussion and summary of the results are found in Sects.~4 and 5, respectively. ",
        "conclusions": "In this contribution we employed a method of disentangling the composite spectrum of the optical continuum on the basis of simple multicolour photometric measurements. Our method (Sect.~2) allowed us to determine the physical parameters of the main contributing components of the radiation into the optical region -- the nebula and the giant. Main results may be summarized as follows. \\begin{enumerate}[(i)] \\item On the basis of the $UBV$ photometric measurements, we determined the true optical continuum of the selected symbiotic (-like) binaries. \\item We compared the observed continuum by a model, which includes contributions from the nebula and the giant (Sect.~3). In this way we determined the electron temperature and emission measure of the nebula, and the $V$ magnitude of the giant. \\item Our model parameters are well comparable with those determined independently by another method. In particular, by a precise modelling the UV--IR SED as introduced by \\citet[][]{sk05}. \\item Our approach thus provides a good estimate of the physical parameters of contributing sources of radiation into the optical on the basis of a simple $UBV$ photometry. \\end{enumerate}"
    },
    "1004/1004.1515_arXiv.txt": {
        "abstract": "{} {In this paper we present a detailed study of the giant radio halo in the galaxy cluster Abell 697, with the aim to constrain its origin and connection with the cluster dynamics.} {We performed high sensitivity GMRT observations at 325 MHz, which showed that the radio halo is much brighter and larger at this frequency, compared to previous 610 MHz observations. In order to derive the integrated spectrum in the frequency range 325 MHz--1.4 GHz,  we re--analysed archival VLA data at 1.4 GHz and made use of proprietary GMRT data at 610 MHz.} {Our multifrequency analysis shows that the total radio spectrum of the giant radio halo in A\\,697 is very steep, with $\\alpha_{\\rm~325 MHz}^{\\rm~1.4 GHz} \\approx 1.7-1.8$. % Due to energy arguments, a hadronic origin of the halo is disfavoured by such steep spectrum. Very steep spectrum halos in merging clusters are predicted in the case that the emitting electrons are accelerated by turbulence, observations with the upcoming low frequency arrays will be able to test these expectations.} {} \\keywords {radiation mechanism: non-thermal -- galaxies: clusters: general -- galaxies: clusters: individual: A\\,697} ",
        "introduction": "\\label{sec:intro}   Radio halos are diffuse low surface brightness radio sources extended on Mpc scales, observed in the central regions of a fraction of X--ray luminous (i.e. massive) galaxy clusters. Their synchrotron emission is not associated with individual galaxies, but it origins from the non--thermal component of the intracluster medium (ICM). They provide the most important piece of evidence for the presence of relativistic ($\\sim$GeV) electrons and $\\mu$G magnetic fields in galaxy clusters (see Ferrari et al. \\cite{ferrari08} for a recent review). To date radio halos have been found from large sky surveys (such as the Northern VLA Sky Survey, NVSS, Condon et al. \\cite{condon98} and the Westerbork Northern Sky Survey, WENSS, Rengelink et al. \\cite{rengelink97}; i.e. Giovannini et al. \\cite{giovannini99}; Kempner \\& Sarazin \\cite{kempner01}), and from pointed observations of individual clusters, mostly carried out at 1.4 GHz (Govoni et al. \\cite{govoni01}; Bacchi et al. \\cite{bacchi03}; and the recent compilation by Giovannini et al. \\cite{giovannini09}) and at 610 MHz (the GMRT Radio Halo survey, Venturi et al. \\cite{venturi07} and \\cite{venturi08}, hereinafter VGB07 and VGD08). They have steep radio spectra, with $\\alpha \\sim 1.2-1.4$ ($S\\propto \\nu^{-\\alpha}$) from integrated flux density measurements. \\newline \\indent The origin of radio halos is a long--standing problem, since the radiative life--time of the relativistic electrons responsible for their synchrotron emission is much shorter than the particle diffusion time required to cover their Mpc extent (Jaffe \\cite{jaffe77}). This requires some form of particle re--acceleration. So far two main models have been proposed to explain the origin of radio halos: (1) the \\textquotedblleft{re--acceleration model}\\textquotedblright, whereby radio halos originate from the in--situ re--acceleration of pre--existing electrons by magneto--hydrodynamic (MHD) turbulence injected in the ICM during cluster mergers (Brunetti et al. \\cite{brunetti01}; Petrosian \\cite{petrosian01}; Brunetti \\& Blasi \\cite{brunetti05}); (2) \\textquotedblleft{secondary electron models}\\textquotedblright, predicting that relativistic electrons are injected in the cluster volume by hadronic collisions between relativistic cosmic rays and the thermal protons of the ICM (e.g. Dennison \\cite{dennis80}; Blasi \\& Colafrancesco \\cite{blasi99}).\\\\ Recent radio observations provide indirect evidence for turbulent acceleration in the ICM (e.g. Brunetti et al. \\cite{a521nature08}). Observational support to the ``re--acceleration scenario'' is also given by the growing evidence for the connection between cluster mergers and radio halos, based on high sensitivity radio and X--ray observations (Buote \\cite{buote01}; Govoni et al. \\cite{govoni04}; VGD08 and references therein; Giacintucci et al. \\cite{giaci09}; Cassano \\cite{cassano09}), and by constraints to the evolution of radio halos (Brunetti et al. \\cite{brunetti09}).  The analysis of the radio spectrum of halos is important to address the question of their origin. The turbulent re--acceleration model provides the unique expectation of spectra much steeper than those found to date, as a consequence of less energetic merger events (e.g. Cassano \\cite{cassano09}). The detection of radio halos with very steep spectrum ($\\alpha >$ 1.6) would be a major piece of evidence in support of this scenario, and at the same time it would disfavour a secondary origin of the electrons, which requires a very large proton energy budget (e.g. Brunetti \\cite{brunetti04}). The prototype of these sources, which we refer to as \\textit{ultra steep spectrum radio halos} (hereinafter USSRH), has been recently discovered in the merging cluster A\\,521 and has a spectrum with $\\alpha\\sim$1.9 (Brunetti et al. \\cite{a521nature08}; Dallacasa et al. \\cite{dallacasa09}). \\\\ In this paper we present a detailed study of the merging cluster A\\,697, which hosts a candidate steep spectrum radio halo. Our work is based on proprietary GMRT radio observations at 610 and 325 MHz, and data from the VLA public archive at 1.4 GHz.  A\\,697 is a rich and massive galaxy cluster located at z=0.282, belonging to the Abell, Corwin and Olowin catalogue (ACO catalogue, Abell et al. \\cite{abell89}). The cluster is hot and luminous in the X--ray band, and is part of the ROSAT Brightest Cluster Sample (BCS; Ebeling et al. \\cite{ebeling98}). Its general properties are summarized in Table \\ref{tab:cluster_prop} \\footnote{We adopt the $\\Lambda$CDM cosmology with H$_0$=70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m=0.3$ and $\\Omega_{\\Lambda}=0.7$. At the redshift of A\\,697 (z=0.282), this cosmology leads to a linear scale of $1^{\\prime \\prime}=4.26$ kpc.}, which provides: coordinates; morphological classification of the cluster and richness; redshift z; cluster velocity dispersion $\\sigma_v$ (from optical spectroscopy); X--ray luminosity L$_{\\rm X}$ (taken from the BCS catalogue); virial mass M$_V$ and the corresponding virial radius R$_V$. The presence of diffuse cluster--scale radio emission in A\\,697 was first suggested in Kempner \\& Sarazin (\\cite{kempner01}) by inspection of the NVSS and the WENSS, and further confirmed by observations at 610 MHz with the Giant Metrewave Radio Telescope (GMRT) as part of the GMRT Radio Halo Survey (VGB07 and VGD08). From those observations the extended radio emission at the cluster centre has been unambiguously classified as a giant\\footnote{Linear size \\gtsim~ 1 Mpc h$^{-1}_{\\rm 50}$.} radio halo. \\begin{table}% \\label{tab:cluster_prop} \\caption[]{General properties of the galaxy cluster A\\,697.} \\begin{center} \\footnotesize \\begin{tabular}{ll} \\hline\\noalign{\\smallskip} RA$_{J2000}$  &  08h 42m 53.3s \\\\ DEC$_{J2000}$ & $+$36$^{\\circ}$ 20$^{\\prime}$ 12$^{\\prime \\prime}$ \\\\ Bautz--Morgan Class & II--III \\\\ Richness & 1 \\\\ z & 0.282  \\\\ $\\sigma_v$  & 1334 km s$^{-1}$ (a) \\\\ L$_{\\rm X[0.1-2.4 keV]}$  & 10.57 $\\times$ 10$^{44}$ erg s$^{-1}$\\\\ M$_{\\rm V}$ &  2.25 $\\times$ 10$^{15}$ M$_{\\odot}$ (b)\\\\ R$_{\\rm V}$ &  2.90 Mpc (b)\\\\ \\hline\\smallskip \\end{tabular} \\end{center} Notes to Table \\ref{tab:cluster_prop}: (a): Girardi et al. \\cite{girardi06}; (b): Estimated from the L$_X$--M$_V$ relation, see Eq. 6 and 7 in Cassano et al. (\\cite{cassano06}). \\end{table}  The paper is organized as follows. In Sect. \\ref{sec:GMRT325} we present the new 325 MHz GMRT data; in Sect. \\ref{sec:spectrum} and \\ref{sec:injections} we derive the integrated spectrum and discuss the sources of uncertainty in the flux density measurements; in Sect. \\ref{sec:discussion} we discuss the origin of the A\\,697 radio halo in the framework of the turbulent re--acceleration model. Summary and conclusions are given in \\ref{sec:results}. ",
        "conclusions": "\\label{sec:results}  In this paper we presented 325 GMRT high sensitivity observations of the giant radio halo in A\\,697, and performed an accurate study of its integrated spectrum.  The largest extent of the radio halo at 325 MHz is $\\sim 5^{\\prime}$ (corresponding to 1.3 Mpc h$_{\\rm 70}^{-1}$), and it is considerably larger than at higher frequencies. The radio spectrum, determined with three data points in the frequency range 325 MHz--1.4 GHz, has a spectral index $\\alpha_{\\rm 325~MHz}^{\\rm 1.4~GHz} = 1.8\\pm 0.1$. An accurate analysis of the fraction of flux density that might be missed in our observations due to the u--v coverage at short spacings led to an estimated spectral index  $\\alpha_{\\rm rev}$=1.7. We thus conclude that A\\,697 hosts a very steep spectrum radio halo.  A qualitative analysis of the X--ray emission from the ICM suggests that this cluster is unrelaxed. Comparison with the 325 MHz emission from the radio halo shows that the non--thermal and thermal emissions have a similar morphology. Though at a qualitative level, the observations are in agreement with the idea that A\\,697 is the result of a multiple merger.  Similarly to the case of A\\,521, we showed that the very steep spectrum of the halo disfavours a hadronic origin, which would require an unplausibly large energy budget. On the other hand, models based on turbulent acceleration for the origin of radio halos predict a large number of halos in the Universe with very steep spectrum. This can be tested as soon as LOFAR and LWA will become operational."
    },
    "1004/1004.1209_arXiv.txt": {
        "abstract": "SN 2007od exhibits characteristics that have rarely been seen in a Type IIP supernova (SN).  Optical $V$-band photometry reveals a very steep brightness decline between the plateau and nebular phases of $\\sim$4.5 mag, likely due to SN 2007od containing a low mass of $^{56}$Ni. The optical spectra show an evolution from normal Type IIP with broad H$\\alpha$ emission, to a complex, four component H$\\alpha$ emission profile exhibiting asymmetries caused by dust extinction after day 232. This is similar to the spectral evolution of the Type IIn SN 1998S, although no early-time narrow ($\\sim$200 km s$^{-1}$) H$\\alpha$ component was present in SN 2007od. In both SNe, the intermediate-width H$\\alpha$ emission components are thought to arise in the interaction between the ejecta and its circumstellar medium (CSM). SN 2007od also shows a mid-infrared excess due to new dust. The evolution of the H$\\alpha$ profile and the presence of the mid-IR excess provide strong evidence that SN 2007od formed new dust before day 232. Late-time observations reveal a flattening of the visible lightcurve. This flattening is a strong indication of the presence of a light echo, which likely accounts for much of the broad, underlying H$\\alpha$ component seen at late-times. We believe the multi-peaked H$\\alpha$ emission is consistent with the interaction of the ejecta with a circumstellar ring or torus (for the inner components at $\\pm$1500 km s$^{-1}$) , and a single blob or cloud of circumstellar material out of the plane of the CSM ring (for the outer component at -5000 km s$^{-1}$). The most probable location for the formation of new dust is in the cool dense shell created by the interaction between the expanding ejecta and its CSM. Monte Carlo radiative transfer modeling of the dust emission from SN 2007od implies  that up to $\\sim$4 x 10$^{-4}M_{\\sun}$ of new dust has formed. This is similar to the amounts of dust formed in other core-collapse supernovae such as SNe 1999em, 2004et, and 2006jc. ",
        "introduction": "The discovery of massive amounts of dust in luminous infrared (IR) high-z galaxies, seen less than 1 Gyr after the Big Bang \\citep{2003A&A...406L..55B}, challenges our understanding of major dust contributors in the Universe. After only 1 Gyr, there has not been time for low-mass stars to form and evolve to the asymptotic giant branch, which accounts for about half of the dust production in the Galaxy at the current epoch \\citep{1998ApJ...501..643D}. Therefore, it has been suggested that core-collapse supernovae (CCSNe) may be significant contributors to the dust production in high-z galaxies. Surprisingly, dust mass estimates derived from the observed SEDs (spectral energy distributions) of recent SNe have been much lower than necessary to account for  the dust observed in early-time galaxies.  The sample of measured SNe is small but it appears that a typical Type II SN is producing only 10$^{-4}-10^{-5}M_{\\sun}$ of dust, well below the 1 M$_{\\sun}$ needed to account for the large amounts of dust seen in the early Universe \\citep{1989ApJ...344..325K, 2001MNRAS.325..726T}. SN 2004et has been found to have formed only a few 10$^{-4}M_{\\sun}$ of dust  \\citep{2009ApJ...704..306K}. \\citet{2006Sci...313..196S} showed that larger masses (10$^{-2}M_{\\sun}$) of dust can be derived in the case of SN 2003gd if the dust is considered to be clumpy and mostly silicate, but \\citet{2007ApJ...665..608M} questioned these findings and derive smaller dust masses (10$^{-4}M_{\\sun}$) for SN 2003gd. Finally, \\citet{2007MNRAS.375..753E}, using similar methods to \\citet{2006Sci...313..196S}, did not find that large amounts of dust had condensed in the Si-poor envelope of SN 1987A.  There are three distinct signatures of dust formation in CCSNe. These include a drop in visual brightness due to increased dust extinction accompanied by an infrared (IR) excess, and the appearance of asymmetric emission line profiles caused by the greater attenuation by dust in the ejecta of redshifted emission than blueshifted emission.  These signatures usually appear a few hundred days after the explosion and only SN 1987A \\citep{1989LNP...350..164L,1993ApJS...88..477W}, SN 2003gd \\citep{2006Sci...313..196S, 2007ApJ...665..608M}, and SN 2004et \\citep{2006MNRAS.372.1315S,2009ApJ...704..306K} have shown all three.  A few CCSNe have shown evidence for early dust formation. SN 1998S formed dust sometime between days 140-268 \\citep{2000ApJ...536..239L}, while SN 2006jc showed dust formation signatures between days 50-75 \\citep{2008MNRAS.389..141M,2008ApJ...680..568S}.  The SNe which show signs of early dust formation are all of type IIn, where the ejecta interacts with the circumstellar medium (CSM) of the SN. The dust may form in a cool dense shell (CDS) between the forward and reverse shocks \\citep{1994ApJ...420..268C,2004MNRAS.352..457P}. The formation of dust in the CDS allows for an additional avenue of dust formation not available in normal Type II SNe. The capability to form dust both within the SN ejecta and in the CDS suggests that these interacting systems have potential for forming dust in greater quantities than normal Type II SNe, although to date they have produced dust amounts comparable to normal Type II SNe.    In this paper, we present new spectroscopic and photometric data covering 800 days of the evolution of SN 2007od, an unusual Type IIP SN that shows evidence for CSM interaction and early-time dust formation. ",
        "conclusions": "\\subsection{Lightcurve Evolution} A typical Type IIP SN lightcurve shows a plateau for the first $\\sim$60-100 days, followed by a sharp drop in luminosity of $\\sim$1.5-3 mag as the SN transitions into the nebular phase, after which the lightcurve exhibits an exponential decline powered by $^{56}$Co decay \\citep{1986ARA&A..24..205W}. The lightcurve of SN 2007od (Figure 2) shows a normal plateau phase followed by an exceptionally large drop in brightness as it is entering the nebular phase. The last photometry before the gap was obtained on day 82 (Chornock 2009, personal communication). If we assume the end of the plateau phase took place shortly after that time and we extrapolate the slope of the $^{56}$Co decay  backwards into the gap (days 82-232) when no photometry is available, we find that the $V$ brightness fades a minimum of 4.5 mag over the typical twenty day drop period  \\citep{2003A&A...404.1077E}. \\citet{1986ARA&A..24..205W}, and more recently \\citep{2009ApJ...703.2205K}, have suggested that if only a small amount of  $^{56}$Ni is synthesized in the SN explosion the lightcurve will show a steep drop between the plateau and nebular phases due to a lack of sufficient thermal energy. Generally, such a steep drop is only seen in low-luminosity (M$_{v}\\sim$-14) SNe such as 1999br, 1999eu and 2005cs, that have very low inferred $^{56}$Ni masses (2 -- 8 x 10$^{-3}$M$_{\\sun}$)  \\citep{2004MNRAS.347...74P}. SN 1994W, which dropped $\\sim$3.5 mag from days 110 to 122 \\citep{1998ApJ...493..933S} is the only other luminous (M$_{v}\\sim$-18) Type IIP besides SN 2007od to show such a steep drop in brightness. Its $^{56}$Ni mass was estimated to be $\\sim$2.6 x 10$^{-3}$M$_{\\sun}$. \\citet{1998ApJ...493..933S} suggested that after its 3.5 mag drop from the plateau, SN 1994W declined much faster than the $^{56}$Co decay once in the nebular phase, possibly due to dust formation.  After day 458 the lightcurve of SN 2007od shows a deviation from  $^{56}$Co decay in the form of a flattening of the lightcurve between days 659 and 811. By day 811, assuming a constant brightness has been maintained from day 672, the difference in an extrapolated $^{56}$Co decay is then $\\sim$2.2 magnitudes in $V$. At late times, the Gemini and HST $V$ magnitudes differ significantly due to differences in the filter bandpasses as discussed in Section 2. However, the constancy of the three epochs of HST/ACS $V$ and $I$ photometry from day 659 to 811 show that the flattening of the lightcurve is real. It is likely being caused by a light pulse from the SN reflecting off of surrounding material and creating a light echo. Given the distance of the SN (24.5 Mpc), platescale of the WFC detector (0\\farcs05~pix$^{-1}$) and FWHM of the stellar PSF (1.9 pix), a source could have a spatial extent on the plane of the sky of up to 0.12\\arcsec or 46 lt-yr and still be unresolved on the detector.  In particular, we find through experimentation that a source with the flux of this SN and with a FWHM up to 1.2 times that of the PSF will not leave detectable residuals when subtracted away, therefore we allow the physical size of the emitting region to be up to 0.12\\arcsec.  Using the light echo equation \\citep{1939AnAp....2..271C,2003AJ....126.1939S}, dust could be up to 475 lt-yr in front of the SN on day 811 and still remain unresolved in the HST imaging.  Thus, the echoes could be produced from either circumstellar or interstellar dust.  Given that the foreground extinction to the SN is low ($A_V \\sim 0.4$), any dust producing a light echo should be optically thin. Typically, the contribution from a light echo becomes significant when a SN has faded $\\sim$8 mag from maximum \\citep{2006MNRAS.369.1949P}. SN 2007od was more than 10 mag below maximum light by day 672  when the contribution of the echo to the lightcurve became significant. A qualitative comparison of the lightcurve and evolution at late times of SN 2007od to models by \\citet{2005MNRAS.357.1161P} suggest that the echo is more likely due to an interstellar sheet than circumstellar dust.  We discuss possible contributions of an echo to the spectral evolution of SN 2007od  in Section 3.2.  As seen in Figure 2, the evolution of the $V-I$ colors is normal for a dust forming Type IIP SN. We have only included the $V-I$ colors from the Gemini observations due to the bandpass considerations of HST described in Section 2. At the time of dust formation, previous photometry of SNe 1999em and 2004et showed only small ($<$0.1 mag) changes in their $V-I$ colors.  The $V-I$ color evolution of SN 2007od before and after the gap in observations (days 82-232) when we think dust formation began, is consistent with these other SNe. The lack of change in the $V-I$ color evolution indicates that, in general, the color evolution is not sensitive to dust formation. This is likely caused by either a very small amount of dust being formed, or larger amounts forming in clumps. If the dust is clumpy, then the color evolution would not reflect the optical depth of dust present and the colors would not be as reddened as they would be if the dust were uniformly distributed e.g., \\citet{1996ApJ...463..681W}.    \\subsection{Spectral Evolution}  The spectral evolution of SN 2007od shows behavior not previously seen in other SNe. Figure 3 shows the full optical spectrum of SN 2007od at 10, 232, and 458 days. Early-time spectra showed SN 2007od to be a normal Type II SN with a blue continuum, broad H$\\alpha$ emission with a P-Cygni profile, and an expansion velocity of $\\sim$10000 km s$^{-1}$  \\citep[Chornock 2009, personal communication]{2007CBET.1119....1B}. Other notable spectral features at early times include broad H$\\beta$ emission, weak Na I D $\\lambda\\lambda$5890,5896 absorption, and [O I] $\\lambda\\lambda$6300,6363 emission that was weaker than that normally seen in Type II SNe. There was no sign of either a narrow (FWHM  $\\sim$ 200 km s$^{-1}$) H$\\alpha$ component nor X-ray emission \\citep{2007CBET.1119....1B,2007ATel.1259....1I}, typically seen in a Type IIn SN.   If present, the X-ray luminosity of SN 2007od was $<$ 1.0 $\\times$ 10$^{40}$ erg s$^{-1}$, significantly smaller than luminosities measured in Type IIn SNe such as 2006jd and 2007pk \\citep{2007ATel.1290....1I,2007ATel.1284....1I}.  By the time of the first Gemini/GMOS epoch (day 232) the spectrum of SN 2007od showed evidence for strong interaction with its CSM, an event usually associated with Type IIn SNe.  The H$\\alpha$ emission profile had changed dramatically from the early-time spectra, with the single broad component evolving into at least three intermediate width (FWHM $\\sim$ 1700 km s$^{-1}$) emission components on top of a fainter, broader emission component (Figure 4).  Two of these intermediate width components are arranged symmetrically about zero velocity at  $\\pm$1500  km s$^{-1}$ and the third, which becomes more prominent at later epochs, lies at -5000 km s$^{-1}$. We should note that the narrow emission at zero velocity is HII emission from the host galaxy not associated with the SN.  Underlying the three intermediate width components, the broad H$\\alpha$ emission present in the early-time spectra can still be seen, and it is likely that it is responsible for the broad red wing of the H$\\alpha$ emission profile (see Figure 4). Broad, blended Fe II emission between 5100-5400 \\AA\\, similar to that seen in SN 1998S \\citep{2004MNRAS.352..457P} appeared by day 232, and [Ca II] $\\lambda\\lambda$7291,7324 emission became prominent by day 458 (Figure 3).  A similar multi-component structure was seen in the H$\\alpha$ profile of SN 1998S.  Unlike SN 2007od, SN 1998S was classified as a Type IIn, and showed both broad and narrow H$\\alpha$ emission components at early-times due to the expanding ejecta and photoionized CSM material, respectively. The narrow component disappeared before day 108 \\citep{2000ApJ...536..239L}, and sometime before day 250, a triple-peaked H$\\alpha$ emission profile had developed \\citep{2000AJ....119.2968G} (Figure 4).  The three intermediate width components persisted in SN 1998S from approximately day 250 to day 660, exhibiting a progressive fading of the red and central components with respect to the blue component; a phenomenon that was attributed to dust formation \\citep{2000AJ....119.2968G}.  The timescale for the evolution of the SN 1998S spectrum is similar to that of SN 2007od, which showed a single broad H$\\alpha$ emission component from days 10-114, followed by a period (114-232 days) when the SN was unobservable, and finally the emergence of  multiple-component emission by day 232.  In SN 2007od, the blueshifted component  at -1500 km s$^{-1}$ is much stronger than its corresponding component at +1500 km s$^{-1}$.  Like SN 1998S, the simplest explanation for this asymmetry is that newly formed dust is extinguishing the redshifted, far-side emission more than the blueshifted, near-side emission.  The ratios of the inner peak heights would require an optical depth of approximately 1.0 at H$\\alpha$ for the asymmetry to be attributable to dust attenuation.  If the dust existed in the CSM before the SN explosion, then it would lie outside of the ejecta and all of the emission components would be equally reddened.  It is thought that the outer peaks in the H$\\alpha$ profile of SN 1998S originated in the interaction of the ejecta with a circumstellar ring or torus \\citep{2000AJ....119.2968G,2000ApJ...536..239L,2002ApJ...572..932P,2004MNRAS.352..457P,2005ApJ...622..991F}.  If we assume the same is true of the asymmetric inner peaks of SN 2007od, then there must be dust between the source of the near-side (blueshifted) and the far-side (redshifted) emission components in the torus (Figure 6). Moreover, if this dust is forming in the CDS, as was suggested by \\citet{2004MNRAS.352..457P} for SN 1998S, then the emission would likely be coming from a radiative forward shock \\citep{2004MNRAS.352..457P,2005ApJ...622..991F} as one would expect equal attenuation of both redshifted and blueshifted emission originating in the reverse shock (see Figure 6). It is also worthwhile to note that \\citet{2009ApJ...704..306K} has recently alluded to CSM interaction producing a similar multi-component H$\\alpha$ profile in SN 2004et, but occuring at least 400 days after it was seen in SNe 1998S and 2007od.  In SN 2007od,  the $\\pm$1500  km s$^{-1}$ peaks appear sometime between day 114 and 232.  Given the maximum velocity of the ejecta of 10000 km s$^{-1}$, this constrains the distance of the inner radius of the CSM responsible for these profile peaks to $\\sim$700-1300 AU.  Assuming a progenitor wind of 10 km s$^{-1}$, this CSM was formed during a mass loss event ending 300-600 years prior to the SN explosion.  With the emergence of the -5000 km s$^{-1}$ components between days 114 and 309, a similar analysis places this CSM at 700-1700 AU and the formation period between 300-800 years prior to explosion.  The appearance of the asymmetric components between days 114 and 232 implies that dust formation began much earlier than typically seen for non-interacting CCSNe.  These SNe (SN 1987A, 2003gd, 2004et) showed developing asymmetries between days 300 and 600 post explosion  due to dust condensation in the ejecta \\citep{1989LNP...350..164L,2006Sci...313..196S,2006MNRAS.372.1315S}.  Furthermore, the [Ca II] $\\lambda$7291 emission profile in the day 458 spectrum shows only the one broad component, with no intermediate-width components present (Figure 3). This implies that the complex emission seen in H$\\alpha$ is arising in the hydrogen-rich CSM lost from the star before the explosion and not in the metal-enriched ejecta \\citep{2000AJ....119.2968G}. These observations are more consistent with dust formation in the CDS as has been suggested in SN 1998S \\citep{2004MNRAS.352..457P,2005ApJ...622..991F}.   The H$\\alpha$ profile of SN 1998S also showed a third peak near zero velocity that has been attributed to the interaction of the ejecta with a spherical distribution of shocked CSM clouds \\citep{2000AJ....119.2968G,2000ApJ...536..239L,2002ApJ...572..932P,2004MNRAS.352..457P,2005ApJ...622..991F}.  In the case of SN 2007od, Figure 4 also shows the emergence of a third peak at -5000  km s$^{-1}$ that was hidden by the broad ejecta emission on day 232.  This blue peak becomes increasingly more prominent as the broad underlying ejecta emission begins to fade over time.   By day 699, when we believe the light echo is contributing significantly to the flux, this feature is still prominent. Because a light echo depicts only the early-time emission, it is not possible to explain the -5000 km s$^{-1}$ peak with a light echo scenario. The nature of this third peak represents the starkest difference between SNe 2007od and 1998S.  As previously mentioned, in SN 1998S the peak not associated with the CSM torus was observed at zero velocity and attributed to a spherical distribution of shocked CSM clouds.  Although a similar solution can not be invoked to explain a peak centered at -5000 km s$^{-1}$ as seen in SN 2007od, it is possible that this peak comes from the emission of a single cloud or blob of shocked CSM.  If this scenario is correct, because the torus and the blob are seen at different velocities, the blob cannot be located in the plane of the torus. At the same time, the blob would be required to have a velocity component along our line of sight larger than that of the torus.  Therefore, if the torus is nearly edge on as viewed from Earth (Figure 6), as hypothesized in SN 1998S, the blob cannot be moving orthogonal to the plane of the torus as this would severely limit its velocity component along our line of sight.  On the other hand, if the inclination$^{2}${\\footnotetext[2]{\\textit{i} = 90$^{\\circ}$ is edge-on} of the ring is less than 90$^{\\circ}$, the blob could be moving orthogonally, but it would require a sufficiently thick torus to retain the dust extinction necessary to explain the asymmetric inner peaks. The scenario presented here to account for the observations is closely based on the scenario used to explain the observations of SN 1998S. This torus + blob scenario is introduced in order to illuminate the manner in which the multi-component H$\\alpha$ emission may arise, but the existing data do not allow the detailed geometry of SN 2007od to be ascertained.    From days 232 to 699, the relative strength of the emission component at -5000 km s$^{-1}$ is increasing as the inner pair of emission components ($\\pm$ 1500 km s$^{-1}$) slowly disappear. Again if we assume the SN 1998S scenario, then this is explained by the ejecta moving through, and completely destroying, the torus but not the blob. The underlying broad component is also present at all epochs. The broad component is likely a combination of the residual ejecta emission and, given the flattening of our lightcurve $\\sim$10 mag below maximum, a light echo.  The contribution to the total broad component by the light echo is expected to grow with time as the ejecta emission fades. Based on the lightcurve, on day 238 the light echo flux would be $\\sim$ 20x fainter than the directly transmitted ejecta emission, and on day 348 the echo would still be $\\sim$8x fainter. However, by day 672/699, the echo flux will likely be greater than the directly transmitted ejecta emission. Therefore the echo should not be a major contributor to the H$\\alpha$ profile at early times, but would be a major contributor of the late-time profile. Figure 4 shows the fits to the day 232 and 348 H$\\alpha$ profiles.  We have used Gaussians to approximate the intermediate width peaks and a scaled day 50 spectrum to approximate the residual ejecta emission.  We should note that the scaled day 50 spectrum is only an approximation to the day 232 and 358 ejecta emission.  Realistically, in addition to fading over time, the profile would likely lose its P Cygni absorption, and its overall width would likely decline. The flat-topped profile of the -5000 km s$^{-1}$ component is shown to be well fit by two intermediate width Gaussians possibly due to density enhancements within the CSM cloud or blob.  At later times the profile becomes too complex to fit.  On days 458 and 672/698, the -5000 km s$^{-1}$ peak begins to dominate and its two component nature becomes more clear.  Moreover, there is a continuing fading of the inner peaks to the point where they  disappear into the broad emission by day 672/698.  As previously mentioned, by this last epoch, the broad emission is a combination of the light echo spectrum and the weak residual ejecta emission.  However, the residual flux beyond 5000 km s$^{-1}$ is unlikely to originate from either of these two sources and remains unexplained. There is also an indication that the radial velocities of both sets of components get smaller with time. This is similar to the behavior of the blue component of SN 1998S which moved to smaller velocities from day 249 to day 658, and was explained as an interaction of the shock with more massive sections of CSM \\citep{2004MNRAS.352..457P}.   \\subsection{Radiative Transfer Modeling}  The SEDs of SN 2007od shown in Figure 5 all show strong IR excesses attributable to warm dust. The SED at 300 days can be fitted well by the sum of a blackbody with a temperature of 5100 K representing hot, optically-thick gas in the ejecta and a modified blackbody with a temperature of 580 K subject to a $\\lambda$$^{-1}$ emissivity law representing the dust.  A similar fit to the day 455 SED uses a 5100 K blackbody and a 490 K modified blackbody, and for day 667 a 5100 K blackbody and a 600 K modified blackbody.  The fits are shown in Figure 5 and the relevant parameters are compiled in Table 4.  The IR emission could arise from pre-existing dust in the CSM, and/or new dust condensing in either or both of the CDS and the ejecta itself.  IR echoes from pre-existing CSM dust heated by the early-time luminosity are often seen in other SNe. The position of the dust in SN 2007od cannot be inferred from IR observations alone, but when combined with the evidence provided by the asymmetry in the H$\\alpha$ emission components, it suggests that the IR emission is coming from dust forming in the ejecta or the CDS, and not from an IR echo.  The very early onset of dust formation (between 114-232 days) further suggests that the dust formation is occurring in the CDS since the ejecta at this time would still be too warm for dust condensation \\citep{1991A&A...249..474K}.  In order to better understand the amount and composition of the dust forming within SN 2007od we used our 3D Monte Carlo radiative transfer (RT) code MOCASSIN \\citep[and references therein]{2005MNRAS.362.1038E}.  For the purposes of our RT models we will assume the dust and source luminosity is spread throughout a spherical, expanding shell with an inner radius, R$_{in}$, and an outer radius radius, R$_{out}$.  The luminosity at each location is proportional to the local density. With this simple model, we can quantify the total amount of dust in the system as a whole (ejecta + CDS) and monitor its variation from day 300 to 667.  We have modeled two different dust distributions.  Following radiative transfer models of other Type II SNe \\citep{2006Sci...313..196S,2007MNRAS.375..753E}, the ``smooth\" model for SN 2007od has the dust uniformly distributed throughout the shell according to a r$^{-2}$ density profile, while the ``clumpy\"  model reflects an inhomogeneous distribution of spherical clumps embedded within an interclump medium of density $\\rho$.  The irradiating photons are only produced in the interclump medium, while the clumps are assumed to be dark (see \\citet{2007MNRAS.375..753E} for further discussion of this issue). We have adopted a standard grain size distribution of $a^{-3.5}$ between 0.005 and 0.05 $\\mu$m  \\citep{2006Sci...313..196S, 2007ApJ...665..608M,2009ApJ...704..306K} and have tested two different concentrations, i.e. a dust model dominated by silicates (75\\% Si; 25\\% AC) and one dominated by amorphous carbon (75\\% AC; 25\\% Si).  For each of these models, we used the optical grain constants of \\citet{1988ioch.rept...22H} and \\citet{1984ApJ...285...89D}, respectively.  The main input parameters for the model are the ejecta temperature and luminosity, the inner and outer radii of the shell, and the mass of dust present.  Initial values of the luminosity, temperature, and approximate size of the dust shell are based on the results of our blackbody analysis above.  Tuning the dust mass, along with slight variations in the other parameters, allowed us to obtain a good fit to the optical and Spitzer/IRAC points for the first two epochs. We should point out that without the 5.8 and 8.0 $\\mu$m data on day 667, we were unable to make reliable constraints on the dust mass at this epoch. The best fit to the data was obtained with a smooth dust distribution and is shown in Figure 5 and the corresponding input parameters are listed in Table 5. This table also lists the dust mass calculated by our clumpy model.  Although no attempts were made to fine tune the models to match the 16 and 24 $\\mu$m flux upper limits on day $\\sim$455, care was taken to ensure that our preferred model did not predict emitted fluxes that would have been detected given our experimental setup.  We found that we were unable to account for this relatively low flux at 16 and 24 $\\mu$m with compositions dominated by silicate dust (see Figure 5).  All other parameters being equal, transitioning from amorphous carbon dust to silicate dust has the effect of broadening the IR peak as the SED will then be increased by strong silicate emission features at 10 and 18 $\\mu$m.  These two emission bumps are clearly seen in the Si model in Figure 5.  The conclusion is that the silicate dominated model predicts too much flux beyond 10 $\\mu$m at day 455 leading us to favor the amorphous-carbon dominated model with 75\\% AC and 25\\% Si.  Given this concentration, our best smooth model predicts a total of 1.7 x 10$^{-4}$ M$_{\\sun}$ and 1.9 x 10$^{-4}$ M$_{\\sun}$ of dust on days 300 and 455, respectively, with an increase in $\\tau$$_{v}$ from 1.5 to 2.0 over the two epochs. This implies a total of 2 x 10$^{-5}$ M$_{\\sun}$ of new dust forming between day 300 and 455 after explosion.  The values of $\\tau$ at 6600 \\AA\\ predicted by our MOCASSIN fits were 1.3 and 1.8 for days 300 and 460, respectively. These opacities are adequate to explain the inner peak height ratios discussed in Section 3.2.  Although the constraints were less reliable on the dust mass at day 667, we were able to show that the available data is consistent with our previous epochs with a predicted dust mass of 1.8 x 10$^{-4}$ M$_{\\sun}$ and $\\tau$$_{v}$ = 2.75.  The best clumpy model gave a total dust mass of 2.5 x 10$^{-4}$ M$_{\\sun}$ and 4.2 x 10$^{-4}$ M$_{\\sun}$ on days 300 and 455/667, respectively, implying the formation of 1.7 x 10$^{-4}$ M$_{\\sun}$ of new dust formed. In each case, the quantity of dust formed is several orders of magnitude lower than that required to contribute significantly to the dust budget of the early universe \\citep{2003MNRAS.343..427M}."
    },
    "1004/1004.2720_arXiv.txt": {
        "abstract": "We use high-resolution relativistic MHD simulations coupled with a radiative transfer code to compute multiwavelength afterglow light curves of magnetized ejecta of gamma-ray bursts interacting with a uniform circumburst medium. The aim of our study is to determine how the magnetization of the ejecta at large distance from the central engine influences the afterglow emission, and to assess whether observations can be reliably used to infer the strength of the magnetic field. We find that, for typical parameters of the ejecta, the emission from the reverse shock peaks for magnetization $\\sigma_0\\sim 0.01-0.1$ of the flow, and that it is greatly suppressed for higher $\\sigma_0$.  The emission from the forward shock shows an achromatic break shortly after the end of the burst marking the onset of the self-similar evolution of the blast wave. Fitting the early afterglow of GRB~990123 and 090102 with our numerical models we infer respective magnetizations of $\\sigma_0\\sim 0.01$ and $\\sigma_0\\sim 0.1$ for these bursts. We argue that the lack of observed reverse shock emission from the majority of the bursts can be understood if $\\sigma_0\\simmore 0.1$, since we obtain that the luminosity of the reverse shock decreases significantly for $\\sigma_0\\sim 1$. For ejecta with $\\sigma_0\\simmore 0.1$ our models predict that there is sufficient energy left in the magnetic field, at least during an interval of $\\sim 10$ times the burst duration, to produce a substantial emission if the magnetic energy can be dissipated (for instance, due to resistive effects) and radiated away. ",
        "introduction": "Gamma ray bursts (GRBs) are thought to be caused by the energy dissipated by an ultra-relativistic outflow. Two alternative scenarios for powering the flow have been extensively considered, resulting either in the thermal-energy dominated fireball \\citep{Goodman:1986eq,Paczynski:1986qd} or in the Poynting-flux dominated flow \\citep[PDF:][]{Usov:1992hp,Thompson:1994ft,Meszaros:1997bx}. In the fireball model, magnetic fields are not dynamically important, and at large distances the flow is weakly magnetized $\\sigma_0 \\ll 1$. Radiation pressure drives the acceleration of the flow until the bulk Lorentz factor becomes $\\Gamma_0 \\simeq h_{\\rm in}$ \\citep[where $h_{\\rm in}$ is the initial specific enthalphy of the flow; see e.g.,][]{Aloy:2005zp,Aloy:2007p7001}, and internal collisions power the prompt emission \\citep{Rees:1994ca}. MHD models for GRB jets assume that the flow is launched Poynting-flux dominated ($\\sigma_{\\rm in}\\gg 1$; just as in the widely accepted theoretical model for relativistic jets from active galactic nuclei). The acceleration of the flow is model dependent and therefore uncertain. PDF flows have in common the fact that the acceleration to ultra-relativistic speed happens by converting part of the magnetic energy into kinetic energy \\citep[e.g.,][]{Li:1992ez}. However, there is no full consensus on which is the magnetization of the flow ($\\sigma_0$) when the acceleration is completed. \\citet{Lyutikov:2003lk} consider a sub-fast magnetosonic flow that maintains $\\sigma_0\\gg 1$ up to distances of $r> 10^{16}\\,$cm, where it interacts with the surrounding medium. In contrast, a large number of ideal and non-ideal MHD models for jet acceleration are super-fast magnetosonic. Still, they support a picture where the acceleration/dissipation processes are not 100\\% efficient in consuming the magnetic energy and, therefore, the flow magnetization $\\sigma_0\\sim 1$ and the bulk Lorentz factor is $\\Gamma_0\\simeq \\sigma_{\\rm in}/(1+\\sigma_0)\\sim \\sigma_{\\rm in}$ at large distances \\citep{Drenkhahn:2002zm, Komissarov:2009p2165, Tchekhovskoy:2009p7068, Lyubarsky:2010p3273}.  During or after the acceleration process magnetic dissipation \\citep{Thompson:1994ft,Spruit:2001hi,Giannios:2008p7077} or internal shocks \\citep{Fan:2004p1007,Mimica:2010p6852} may power the GRB.  Here we focus on the dynamics of the deceleration of $\\sigma_0 \\simless 1$ flows with focus on the role of magnetic field strength to the resulting emission.   If the magnetic fields of the ejecta are dynamically unimportant (e.g., $\\sigma_0\\ll 1$) at large distance where the interaction with the circumburst medium becomes substantial, a forward shock (FS) and a reverse shock (RS) are expected to form at the interface between the GRB ejecta and circumburst (external) medium \\citep{Sari:1995oq}. If the ejecta is strongly magnetized, we expect a weak or absent RS \\citep{Giannios:2008zl}. Observationally, there is a wealth of recent evidence that magnetic fields shall be dynamically important. Most bursts show no trace of the RS emission, in contrast to predictions of the fireball model \\citep[see, however, ][]{Nakar:2005p4038,McMahon:2006p4021}.  When the bump in the early optical light curve is observed (and resembles that expected by the RS), modeling of the light curves indicates that the magnetic field is larger in the shocked ejecta than in the shocked external medium \\citep{Fan:2002p7207,Zhang:2003wc,Panaitescu:2004p7344}.  More conclusive is the detection of $\\sim 10\\%$ polarization in the early optical afterglow emission of GRB 090102, consistent with the ejecta containing quasi-coherent, large-scale magnetic fields \\citep{Steele:2009p5925}. Since dynamically dominant magnetic fields would suppress the RS emission altogether, these authors give a rough estimate of $\\sigma_0 \\sim 1$ for GRB 090102, supporting super-fast magnetosonic MHD models for GRBs.  To make quantified inferences for the degree of the magnetization of the ejecta, a detailed understanding of the dynamical interaction between magnetized ejecta and external medium is needed. This has been done at an increasing level of sophistication. As a first approach to the problem, the dynamical effects of magnetic fields have been ignored, assuming a purely hydrodynamic behavior of the ejecta \\citep{Rees:1992kx,Sari:1995oq,Kobayashi:1999ly,Granot:2001fq,Meliani:2007qm}. Such an approximation has lead to infer that the magnetic field strength at the RS and at the FS is different (i.e., the parameter $\\epsilon_{B}$ is larger at the RS than at the FS \\citep{Fan:2002p7207,Zhang:2003wc}).  A more accurate approach takes into account the modification in the shock jump conditions because of the presence of magnetic fields \\citep{Fan:2004if,Zhang:2005ts}.  \\citet{Lyutikov:2006wj,Genet:2006fe} studied the afterglow from sub-fast ($\\sigma_0\\gg 1$) MHD models. \\citet{Giannios:2008zl} found that rarefaction waves, instead of shocks, can also form in magnetized ejecta \\citep[see also][]{Mizuno:2009p210,Aloy:2008kx}. These studies have left several issues open (e.g., that of the timescale of transfer of the magnetic energy into the shocked external medium) or debated \\citep[see][]{Lyutikov:2005tx}.  The complete dynamical evolution of homogeneous, magnetized ejecta colliding with the ISM has been simulated by \\citep[][MGA09 in the following]{Mimica:2009qa} for a few characteristic cases, showing how the magnetic energy is transferred in the shocked ISM until a self-similar blast wave forms at later time.  Although bolometric light curves of the expected emission of magnetized ejecta were computed by MGA09, a detailed calculation of the synthetic light curves at different wavelengths is missing. In this work we have coupled the SPEV (\\emph{SP}ectral \\emph{EV}olution) code \\citep{Mimica:2009p1445}, used to compute the non-thermal emission from relativistic outflows, to the MRGENESIS \\citep{Mimica:2005sp,Mimica:2007db} RMHD code in order to study in accurately the signatures of magnetic fields and the (non)existence of shocks in afterglow ejecta.  The paper is organized as follows: in \\secref{sec:numaftsim} we give a brief overview of the results in MGA09 and present our extended set of models. \\secref{sec:non-thermal} gives details of the model we use to compute the non-thermal synchrotron emission from relativistic shocks. In \\secref{sec:dynamics} we discuss the energy content of GRB ejecta and of the external medium as a function of the distance and of the initial flow magnetization. Light curves and spectra are presented in \\secref{sec:lightcurves}. We apply our models to the GRBs 990123 and 090102 in \\secref{sec:application}. Discussion and conclusions are given in \\secref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  The theoretical modeling of GRB central engines is still far from predicting whether the GRB flow is thermally or magnetically dominated in the launching region. Furthermore, prompt GRB emission mechanisms are not understood well enough to enable probing of the magnetization of the emitting plasma.  A more promising way to infer the magnetization of the GRB flow may come through studying interactions with the circumburst medium. Quantitative inferences, however, critically depend upon the detailed modeling of the dynamics of such interaction, as well as the emission mechanisms from shocked plasma. In this work we presented relativistic MHD models of ejecta-medium interactions coupled to a radiative transfer code to calculate multi-wavelength synchrotron lightcurves. The most salient results of our parametric study are the following:  \\begin{enumerate}  \\item For very weakly magnetized ejecta, the standard forward/reverse shock structure \\citep{Sari:1995oq} is reproduced by our simulations. The reverse shock can dominate the early optical/IR emission while the forward shock dominates the subsequent emission at all wavelengths.  Increasing the magnetization of the ejecta to $\\sigma_0\\sim 0.01-0.1$, the reverse shock remains strong while the powerful fields result in a brighter synchrotron emission, typically, in the optical band. At still stronger magnetization $\\sigma_0\\simmore 1$, the dissipation at the reverse shock and, therefore its emission are greatly reduced.  \\item Observationally, the majority of GRBs do not show bright optical emission at early times. \\citet{Klotz:2009p4221} present a systematic study of early optical observations and tight upper limits tens of seconds after GRB emission. We interpret this lack of early optical detections as evidence for weak/absent reverse shock as expected from GRB flows characterized by $\\sigma_0\\simmore 1$ at large distance from the central engine. Nevertheless, a few GRBs such as 990123 and 090102 do show bright optical emission that is likely coming from the reverse shock \\citep{Akerlof:1999yf,Meszaros:1999p4433,Nakar:2005p4038,Steele:2009p5925}.  \\item Our detailed modeling of early afterglow emission from GRB~990123 and 090102 indicates that the magnetization of the ejecta is $\\sigma_0\\sim 0.01-0.1$ for these bursts, i.e., close to the value that maximizes the RS brightness (see \\figref{fig:R-band-app}). For GRB~090102, the presence of a large scale field is strongly supported by the detection of $\\sim 10$\\% optical polarization \\citep{Steele:2009p5925} during a $60\\,$s exposure shortly after the burst. It, therefore, appears that there is a range of magnetizations of $\\sigma_0 > 0.01-0.1$ for the GRB ejecta at the radius of deceleration. This values are compatible with super-fast magnetosonic MHD models for GRBs.  \\item The reverse shock slows down and compresses the shell.  As a result, during the RS crossing the magnetic energy of the ejecta is actually increased \\citep[][MGA09]{Zhang:2005ts}.  Subsequently, the bulk of the magnetic energy is diluted because of a rarefaction that crosses the ejecta (produced when the RS reaches the innermost radial edge of the ejected shell). The remaining ($\\sim$10\\%) of the magnetic energy decays slowly with time (see Fig.~1). If a substantial fraction of this energy is radiated away because of dissipative MHD processes \\citep{Giannios:2006hb}, it may account for the X-ray flares observed minutes to hours after many GRBs \\citep{Burrows:2005kk}.  \\item The bulk of the magnetic energy is transferred into the shocked ISM shortly after the end of the burst. Only at this point, the blast wave establishes the well known BM self-similar behavior. A flat or rising light curve until well after the GRB is over at high-energies is the signature imprinted on the FS emission by the aforementioned energy transfer. An initially flat high-energy light curve may be the signature of $\\sigma_0\\sim 1$ ejecta while a rising one comes from smaller magnetization (assuming always a constant density ISM).  When the energy transfer is completed, an achromatic break takes place. After the break, the light curve declines as expected from the BM solution.  \\item It has been proposed \\citep{Kumar:11p7654,Ghisellini:2010p7532} that the GeV emission observed by the LAT detector on FERMI in several bursts is result of the FS. The model can account for GeV observations shortly after the end of the burst and and for the late optical and X-ray data. In its simplest form, however, the FS model cannot explain the peak of the GeV emission at times $t_{peak}<t_{\\rm GRB}$ seen in GRB~080916C, 090510 and likely in 090902B (since, as we have shown here, the energy transfer to the FS results in $t_{peak}\\simmore t_{\\rm GRB}$). On the other hand, from the \\citet{Ghisellini:2010p7532} sample of several LAT bursts, at least three GRBs (090323, 090626, and 091003) show a rather flat GeV light curve that steepens only after the MeV emission (i.e. the traditional GRB) is over. Though more detailed modeling is needed, the emission from these bursts appears compatible with coming from the early FS driven by ejecta with $\\sigma_0\\ge 1$ in agreement with inferences we have made from the lack of observed of RS emission.  \\end{enumerate}  The results outlined so far suggest that to explain the paucity of observed optical flashes one needs to argue that most of the observed bursts posses substantial magnetizations (point (ii) above). However, our models are sensitive to a number of parameters, among which, $\\epsilon_B$ is probably the most crucial one. The value that we use as representative in our models ($\\epsilon_B \\sim 5\\times 10^{-3}$; Sect.~\\ref{sec:lightcurves}) results in powerful optical flashes even for the $\\sigma_0=0$ model.  Values of $\\epsilon_B\\ll 5\\times10^{-3}$ reduce the strength of the RS emission. We find that for $\\epsilon_B\\simless 10^{-6}-10^{-5}$ the RS emission falls below that of the FS, for all our purely hydrodynamic models ($\\sigma_0=0$). Therefore, even in fireball models one may argue that the optical emission of the RS is suppressed if mildly relativistic shocks are not able to efficiently generate stochastic magnetic fields.  Another issue that has to be taken into consideration is the fact that we are modeling the ejecta as initially being cold. Moderately hot or {\\it warm} ejecta (i.e., in which $P/\\rho c^2\\simless 1$) may result, e.g., if the conversion of internal to kinetic energy is slower than predicted by the standard fireball model (which might not be unlikely by looking at the numerical models of, e.g., \\citealt{Aloy:2000ad,Aloy:2005zp}), or if internal shocks heat up the ejecta shortly before the reverse shock crossing. In such scenario, a weaker RS may be envisioned, since it is more difficult to shock warm than cold matter. Thus, we may not exclude the possibility that the absence of a clear RS signature in many GRBs arises (assuming that the ejecta is not magnetized) from the fact that the ejecta is warm by the time the afterglow emission takes place."
    },
    "1004/1004.0725_arXiv.txt": {
        "abstract": "Linear acceleration emission occurs when a charged particle is accelerated parallel to its velocity. We evaluate the spectral and angular distribution of this radiation for several special cases, including constant acceleration (hyperbolic motion) of finite duration. Based on these results, we find the following general properties of the emission from an electron in a linear accelerator that can be characterized by an electric field $E$ acting over a distance $L$: (i) the spectrum extends to a cut-off frequency $\\hbar\\omega_{\\rm c}/mc^2\\approx L \\left(E/E_{\\rm Schw}\\right)^2/\\lambdaC$, where $E_{\\rm Schw} =1.3\\times10^{18}\\,\\textrm{V\\,m}^{-1}$ is the Schwinger critical field and $\\lambdaC=\\hbar/mc=3.86\\times10^{-13}\\,\\textrm{m}$ is the Compton wavelength of the electron. (ii) the total energy emitted by a particle traversing the accelerator is $\\frac{4}{3}\\alpha_{\\rm f}\\hbar\\omega_{\\rm c}$, in accordance with the standard Larmor formula, where $\\alpha_{\\rm f}$ is the fine-structure constant. (iii) the low frequency spectrum is flat for hyperbolic trajectories, but in general depends on the details of the accelerator. We also show that linear acceleration emission complements curvature radiation in the strongly magnetized pair formation regions in pulsar magnetospheres.  It dominates when the length $L$ of the accelerator is less than the formation length $\\rho/\\gamma$ of curvature photons, where $\\rho$ is the radius of curvature of the magnetic field lines, and $\\gamma$ the Lorentz factor of the emitting particle. In standard static models of pair creating regions linear acceleration emission is negligible, but it is important in more realistic dynamical models in which the accelerating field fluctuates on a short length-scale. ",
        "introduction": "Linear acceleration emission is the radiation produced when a charged particle is accelerated in a direction parallel to its velocity. The resulting rectilinear trajectory is a special case. In astrophysics, it is expected to occur in the inner parts of a pulsar magnetosphere, where the strong magnetic field enforces one-dimensional motion on the electrons, and there is a non-zero electric field aligned with the magnetic field.  The calculation of linear acceleration emission is complicated by the fact that relatively large sections of the trajectory radiate coherently (in a sense that will be made more precise in Section~\\ref{equations}). In the celebrated case of \\lq hyperbolic\\rq\\ motion, in which the acceleration measured in the instantaneous rest frame of the particle is constant, the entire trajectory remains coherent and it is impossible to define a local photon emission rate. Linear acceleration emission is a potentially important component of the electromagnetic cascades that occur in pulsars. These cascades are responsible for the materialization of the electrons and positrons that radiate in pulsar wind nebul\\ae, and, ultimately, make their way to Earth as a high-energy cosmic-ray component \\citep[e.g.,][]{profumo09}.  Early work on this problem in the context of pulsars treated the {\\em relativistic oscillator}, in which the electric field is a linear function of position \\citep{wagoner69}. Linear solutions were considered by \\citet{melrose78} and \\citet{kuijpersvolwerk96} who emphasized linear acceleration emission masers. Recently, the focus has shifted to superluminal, large-amplitude electrostatic waves. The particle trajectory and resulting emission in these waves was calculated using a semi-classical formalism by \\citet{rowe95} and has been re-analyzed by \\citet{melroseetal09} and \\citet{melroseluo09} who found and discussed some significant differences with Rowe's results.  In this paper we compute the emissivity for linear acceleration in several special cases. The results enable us to estimate the emission in the general case, and we use this estimate to discuss the importance of the process in both static and dynamic models of the pair producing regions in pulsar magnetospheres, which we call \\lq\\lq gaps\\rq\\rq. Our method differs from that of \\citet{melroseetal09} and \\cite{melroseluo09} in several respects. In particular, by avoiding the use of the \\lq Airy integral approximation\\rq\\ we resolve a puzzle they encountered concerning the applicability of Larmor's formula to linear acceleration. We also drop the restriction to superluminally propagating disturbances, which enables us to treat static gaps, as well as those containing spatially isolated electric field structures such as double layers. Although the method and the details of the emissivity we find are substantially different, the main conclusions concerning the importance of linear acceleration emission in pulsars are consistent with those of \\citet{melroseluo09}: it is unimportant for the parameters currently under discussion in static gap models of pulsars, but could be the dominant radiation mechanism in dynamic gap models.  The paper is organized as follows: in section~\\ref{equations} we briefly review the formalism of classical electrodynamics concerning the radiation of test particles in prescribed vacuum fields, emphasizing the role of the photon formation or coherence length. In section~\\ref{hyperbolic} we apply this formalism to the well-known problem of hyperbolic orbits (motion with constant acceleration in the particle's rest frame) and derive a useful approximate representation of the angular and spectral distribution of radiated photons. In section~\\ref{special} we use the method of steepest descents to compute an approximate emissivity for a simple wave form that can describe a stationary or moving double layer.  We also present an exact analytic solution for the spectrum emitted by a particle undergoing simple harmonic motion.  In each case we demonstrate that the results agree with simple estimates based on the spectrum found for hyperbolic orbits.  Finally, we discuss in section~\\ref{pulsars} the role of linear acceleration and curvature emission in pulsar gap models, comparing the relative importance of each process. Our conclusions are presented in section \\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  The radiation produced by a particle whose acceleration is parallel to its velocity, linear acceleration emission, has been investigated for a number of special cases. Particle motion of this type is expected to occur in the intense magnetic fields close to the surface of a pulsar. The emission produced by these particles is conceptually difficult due to the macroscopic size of the photon formation length. For the particular case of hyperbolic motion, corresponding to acceleration in a uniform electric field, the formation length is comparable to the length of the entire accelerating region. Linear acceleration emission is similar to both synchrotron and curvature radiation in the sense that a photon is emitted with a probability $\\alpha_{\\rm f}$ when a particle traverses a formation length. The essential properties of the radiation produced by a particle on a hyperbolic trajectory can be understood in terms of two key parameters: the length $L$ of the hyperbolic part of the trajectory, and the length $\\accel$ over which the particle increases its energy by $mc^2$. Some generic features of the emission mechanism in different electric field configurations can also be described using these quantities. In particular, provided  $L$ and $\\accel$ are interpreted appropriately (see section \\ref{special} for some examples) the critical frequency can be expressed as $\\omega_{\\rm crit}\\approx Lc/\\accel^2$.   Current models of pair production in pulsar magnetospheres typically assume that high energy photons that induce the secondary pair cascade, are produced via either curvature radiation or inverse-Compton scattering.  An alternative to these processes is linear acceleration emission, as investigated by \\citet{melroseluo09}. We have demonstrated that this process is in fact complementary to curvature radiation, the transition between the two regimes depending only on the ratio of the system size to the formation length of curvature photons.  For existing static polar gap models, linear acceleration is unlikely to play an important role, although this relies on the environmental parameters being close to the canonical values. The limits for both vacuum gap models and space-charge limited models are presented in section \\ref{pulsars}.   Recent numerical studies provide evidence that the above stationary models are unstable to perturbations \\cite[e.g.][]{levinsonetal05, timokhin09}. Several analytic models have been developed in which particles are accelerated in large amplitude electrostatic waves that propagate parallel to the magnetic field \\cite[see e.g.][and references therein]{luomelrose08}. Since the characteristic length scale of the electric fields in these models is on the order of the electron inertial length, linear acceleration emission is likely to be a more appropriate description of the resulting radiation."
    },
    "1004/1004.0380_arXiv.txt": {
        "abstract": "To understand the correlation and the radiation mechanism of flare emission in different wavelength bands, we have coordinated a number of telescopes to observe Sgr~A* simultaneously. We focus only one aspect of the preliminary results of our multi-wavelength observing campaigns, namely, the short time scale variability of emission from Sgr~A* in near-IR, X-ray and radio wavelengths.  The structure function analysis indicate most of the power spectral density  is detected on hourly time scales in all wavelength bands. We also report minute time scale variability at 7 and 13mm placing a strong constraint on the nature of the variable emission.  The hourly time scale variability can be explained in the context of a model in which the peak frequency of emission shifts toward lower frequencies as a self-absorbed synchrotron source expands adiabatically near  the acceleration site.  The short time scale variability, on the other hand, places a strong constraint on the size of the emitting region.  Assuming that rapid minute time scale fluctuations of the emission is optically thick in radio wavelength, light travel arguments requires relativistic particle energy, thus suggesting the presence of outflow from Sgr~A*. ",
        "introduction": "%  Observations of stellar orbits in the proximity of the enigmatic radio source Sgr~A* located at the very dynamical center of our galaxy (Reid and Brunthaler 2004) have established that it is a 4 $\\times 10^6$\\msol\\ black hole (Ghez et al. 2008; Gillessen et al. 2009). The relative proximity of Sgr~A* compared to AGNs allows us to observe the spectacular activities of the central engine at remarkably small spatial scales. The variable continuum flux of AGNs is known to signal activities of the central engine and is detected throughout the entire electromagnetic spectrum with time scales  ranging from days to years (e.g., Arshakian et al. 2010).  The bulk of the continuum flux is believed to be generated in the accretion disk, where the localization of the source of variable continuum emission becomes essential for our understanding of the launching and transport of energy in active galaxies.  Sgr~A* provides the best laboratory to study the properties of low-luminosity accreting black holes.  The emission from Sgr~A* is therefore the subject of intense scrutiny, which will have a long-lasting impact on our understanding of the radiatively inefficient accretion flow into, as well as outflows from, massive black holes.  The time scale for variability is proportional to the mass of the black hole, thus  Sgr~A*, which is a hundred times closer to us than the next nearest example, presents an unparalleled opportunity to closely study the processes by which gas is captured, accreted or ejected, by characterizing the time variability over scales of minutes to years.  The luminosity of Sgr~A* is thought to be due to partial capture of thermal winds from a neighboring cluster of massive stars.  However, the bolometric luminosity of Sgr~A* ($\\sim 100$ \\lsol) is several orders of magnitude lower than expected for the estimated accretion rate, prompting a number of theoretical models to explain its very low radiative efficiency and matching the spectral energy distribution (SED) of its quiescent emission. The broad-band spectrum of Sgr~A* peaks at sub-millimeter wavelengths, which is the dividing line between optically thin infrared and optically thick millimeter and radio emission.  The existence of flare emission from Sgr~A* is now well established in both optically thick and thin regimes.  What is really exciting is that we are beginning to peer into the closest supermassive black hole and are discovering rapid, time variability that likely stems from gas dynamical flow and radiation emitted very near the event horizon. ",
        "conclusions": "Our multi-wavelength  monitoring of Sgr A* characterizes  the intrinsic time variability of Sgr~A* by studying structure function of the observed light curves. Unlike near-IR structure function plots that can be fit by single power laws, radio and X-ray structure function plots are fit by multiple power laws in different time lag domains. Structure function analysis of the data presented in near-IR, X-ray and radio data suggest that most of the power falls in the long time scale fluctuation of the emission from Sgr~A* and no obvious QPO activity. We also note short time scale variability at these wavelengths when the data are sampled properly. In particular, the minute-time scale variability is best seen at NIR wavelengths in Fig. 1a. The  structure function also  goes down well below one minute at radio wavelengths showing a slope which is not completely flat at the shortest lags. The best case for such minute time scale variability is seen in Figure 3a at 7mm  where the slope is not consistent with being flat for lag times less than 0.3 minute.  The square of the mean of the measurement errors for 3.3sec sample data at 7mm is $\\sim0.0007$ Jy$^2$ which is more than a factor of two  less than the   amplitude of the structure function 2$\\times10^{-3}$ Jy$^2$ at the shortest lag time.    Assuming that the sub-minute time scale variability in NIR and radio wavelengths are  significant, they place a strong constraint on the size of the region from which variable emission arises. Using light travel time arguments, the minute-time scale variability gives an upper limit to the scale length of radio variation $\\sim$1/8 AU. For comparison, the Schwarzschild radius of Sgr~A* is 10 $\\mu$as or 0.1 AU at the distance of 8 kpc (Doeleman et al. 2008). Assuming that the emission is optically thick at 7mm, then the inferred brightness temperature is $> 5\\times10^{10}$K corresponding to particle energy $> 4.5$MeV.  These high energy particles have relativistic speeds and can in principle leave the gravitational potential of Sgr~A* in the form of outflow. The inferred scale length corresponding to one-minute light travel time is comparable to  to the time averaged spatially resolved 0.1AU  scale  observed at 1.3mm by Doeleman et al. (2008). The quiescent variable emission from Sgr~A* could then be interpreted mainly as an ensemble average of numerous flares that are detected on minute-time scale.  Alternatively, the short time scale emission or quiescent variability could be due to fluctuations in the accretion flow of Sgr~A* due to magnetic field fluctuations resulting from MRI, as recent MHD simulations in a number of studies indicate.  A more detailed account of these results will be given elsewhere."
    },
    "1004/1004.2776_arXiv.txt": {
        "abstract": "{We present high quality long slit spectra along the major and minor axes out to 1.5-2 $R_e$ (14-22 kpc) of three bright elliptical galaxies (\\object{NGC1600}, \\object{NGC4125}, \\object{NGC7619}) obtained at the Hobby-Eberly Telescope (HET). We derive stellar kinematic profiles and Lick/IDS indices ($H\\beta,Mgb,Fe_{5015},Fe_{5270},Fe_{5335}, Fe_{5406}$). Moreover, for \\object{NGC4125} we derive gas kinematics and emission line strengths. We model the absorption line strengths using Simple Stellar Populations models that take into account the variation of $[\\alpha/Fe]$ and derive ages, total metallicity and element abundances. Overall, we find that the three galaxies have old and $[\\alpha/Fe]$ overabundant stellar populations with no significant gradients. The metallicity is supersolar at the center with a strong negative radial gradient. For \\object{NGC4125}, several pieces of evidence point to a recent dissipational merger event.  We calculate the broad band color profiles with the help of SSP models.  All of the colors show sharp peaks at the center of the galaxies, mainly caused by the metallicity gradients, and agree well with the measured colors.  Using the Schwarzschild's axisymmetric orbit superposition technique, we model the stellar kinematics to constrain the dark halos of the galaxies. We use the tight correlation between the Mgb strength and local escape velocity to set limits on the extent of the halos by testing different halo sizes. Logarithmic halos -- cut at 60~kpc -- minimize the overall scatter of the Mgb-$V_{esc}$ relation. Larger cutoff radii are found if the dark matter density profile is decreasing more steeply at large radii. } ",
        "introduction": "\\label{sec_intro} One of the most important remaining questions in contemporary astrophysics concerns the formation and evolution of stellar populations in elliptical galaxies.\\footnote{Tables 3, 4 and 6 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/} One can think of two globally different formation scenarios: the classic \\emph{monolithic-collapse} model for the formation and evolution of early-type galaxies \\citep{1972ApJ...178..319T,1975MNRAS.173..671L, 1996A&A...311..361T} and the \\emph{hierarchical merging} scenario \\citep{1978MNRAS.183..341W,1993MNRAS.264..201K}. In the former, early-type galaxies build most of their stars during a single event in the early universe while in the latter a substantial fraction of the stellar populations is formed from multiple mergers and the accretion of smaller objects over an extended period.  Statistically, the extremely small scatter of the observed color magnitude relation of elliptical galaxies or bright cluster members suggests that the stellar populations are surprisingly homogeneous \\citep{1992MNRAS.254..601B,1997A&A...320...41K}.  This uniformity of stellar populations in ellipticals is also supported by the internal tightness of the so called Fundamental Plane \\citep{1987ApJ...313...42D,1987ApJ...313...59D, 1992ApJ...399..462B, 1993A&A...279...75S}. These findings are consistent with a \\emph{monolithic-collapse}. Moreover, the tightness of the \\begin{math}Mg-\\sigma\\end{math} relation for massive elliptical galaxies observed in the local universe \\citep{1993ApJ...411..153B,2007MNRAS.377..759S} but holding up to the intermediate redshift \\begin{math}z\\thickapprox 1\\end{math} \\citep{1997MNRAS.291..527Z,1998ApJ...493..529B} requires a picture of a short, and highly efficient star formation process at high redshift and passive evolution since then.  In contrast, a large number of elliptical galaxies are observed to have kinematically decoupled cores and were explained as the result of dissipational major merging events by star dominated systems \\citep{1992ApJ...393..484B,1989A&A...217...35B,1998A&A...332...33M}. Furthermore, boxy galaxies are mostly more massive, more radio-loud, more strong X-ray emitters, more frequently disturbed, and less flattened than disky early-type galaxies, suggesting that these two subclasses reflect different formation mechanisms and evolution history \\citep{1992A&A...258..250B,1989A&A...215..266N}. Recently, \\citet{2000MNRAS.315..184K,2005ApJ...621..673T, 2006MNRAS.370.1213C, 2006AJ....131.1288B,2006MNRAS.370..702C,2010MNRAS.tmp..387R} showed that early-type galaxies in low density and in high density environments might exhibit different formation ages, and \\citet{2009A&A...499...47S,2009ApJ...691.1862M} found evidence that lower mass galaxies have more extended star formation histories.  In standard closed-box models of chemical enrichment, the metallicity is a function of the yield and of how much gas has been locked after star formation has ceased \\citep{1980FCPh....5..287T}.  Therefore, the metallicity strongly depends on the dynamical parameters -- the deeper a potential is, the easier it is to retain gas in the galaxy. Then, galaxies formed in a \\emph{monolithic-collapse} should have steep radial metallicity gradients \\citep{1976MNRAS.176...31L,1999MNRAS.306..655T}, correlated with the local escape velocity -- since the latter is a measure of the local depth of the potential well. On the other hand, major mergers, which come along with \\emph{hierarchical merging} will dilute stellar population gradients and their connection with the escape velocity as a result of violent relaxation \\citep{1980MNRAS.191P...1W}. In \\emph{hierarchical merging}, however, particularly for high-mass galaxies a significant amount of mass accretion can be due to minor mergers. If this deposites stellar populations that originally formed in small-mass halos preferentially in the outer parts of the final massive ellipticals, then this might enhance stellar-population gradients and one would again expect deviations from correlations with the local escape velocity, but in the opposite direction as induced by major mergers. Accordingly, the observation of stellar population gradients and their connection to dynamical parameters can give crucial insight into the formation paths of individual galaxies.  The comparison of measured line strengths to simple stellar population models \\citep[][hereafter TMB03]{1994ApJS...95..107W,1999ApJ...513..224V,2003MNRAS.339..897T} yields an estimate of the age, metallicity and element abundances and the galaxy formation scenarios, formation time scales and star formation histories can be investigated in detail from the study of the kinematic properties and line strength features \\citep{1992A&A...258..250B,1992ApJ...399..462B,2004MNRAS.354..753M}. \\citet{2007MNRAS.374..769S} show that when two bursts are considered, the metallicity and element abundances derived with such a procedure are roughly luminosity weighted, while the stellar ages can be biased towards the ones of the younger component.  There is a considerable number of works which focus on the study of kinematic profiles, line strength indices and stellar population parameters in early-type galaxies \\citep{1992A&A...258..250B,1994MNRAS.269..785B,1995ApJ...448..119F, 1998A&AS..130..267L,1998A&A...333..419T,2000AJ....119.1645T,2000AJ....120..165T, 2000A&A...353..917L,2000MNRAS.315..184K,2002MNRAS.330..547T,2005MNRAS.358..813D, 2007A&A...464..853L, 2008MNRAS.386..715T}.  However, the previous measurements are concentrated within $R_{e}$/2 of galaxies.  From a comparison of the stellar parameters within $R_{e}$/8 with these within $R_{e}$/2 of some early-type galaxies sample, \\citet{1993MNRAS.262..650D,2000AJ....119.1645T,2005MNRAS.358..813D} found that the outer regions of the elliptical galaxies are slightly more metal poor than the centers, and that ages are likely to increase slightly outwards.  The same trends were detected by \\citet{1995ApJ...448..119F}. But still only 1/3 of the star mass is contained within $R_e$/2.  So far the exact behaviour of the stellar kinematics and the abundance ratios at larger radii is still poorly known.  Extended stellar kinematic profiles are also welcome to probe the dynamics of the outer regions of elliptical galaxies, where dark matter halos start to dominate the gravitational potential.  Together with other kinematic tracers, like the radial velocities of planetary nebulae \\citep{2009MNRAS.393..329N} and globular clusters \\citep{2006MNRAS.373..157B}, the density and temperature profiles of X-ray halos  \\citep{arXiv:0911.0678} or gravitational lensing effects \\citep{2009ApJ...703L..51K}, extended stellar kinematics profiles have been used to measure the dark matter halos densities and constrain their assembly epoch \\citep{Gerhard01, Tho07, Tho09}.  A question of particular interest is the DM halo sizes. It has been probed statistically using gravitational lensing, and it can be addressed, if accurate X-ray observations sample the very outer regions of ellipticals.  Through gravitational lensing a statistical value for the cutoff radius of $64^{+15}_{-14}$ kpc for the average cluster galaxy with a velocity dispersion of 220 km/s has been estimated \\citep{2007ApJ...656..739H}. X-ray measurements, on the other hand, do not seem to detect the ends of halos \\citep{arXiv:0911.0678}.  Here we attempt to improve on the leverage of stellar dynamical models, that can not constrain the mass profiles accurately beyond the last kinematic data point \\citep{2005MNRAS.360.1355T,2007MNRAS.381.1672T}, by adding the further constraint coming from the measured absorption line strengths.   \\citet{1990ApJ...359L..41F} found that the local color is a function of the local escape velocity. \\citet{1993MNRAS.262..650D} discovered a tight correction between $Mg_2$ and the local escape velocity $V_{esc}$ in 8 galaxies. More recently, \\citet{2009arXiv0906.0018W} also reported a tight relation between Mg\\emph{b} and $V_{esc}$ in NGC3379 while \\cite{2009MNRAS.398.1835S} explored correlations between stellar population parameters and $V_{esc}$ in the SAURON sample.  Dynamical measurements can only constrain the density profile of the dark matter halo out to radii for which data are available.  In contrast, as the Mgb line strength is a function of the escape velocity at a given radius, it probes the gravitational potential directly rather than its gradient, as it is the case with dynamical measurements. Therefore, since the gravitational potential depends on the whole density distribution (in principle out to infinity), the Mgb-$V_{esc}$ provides a novel tool to constrain the density distribution of dark matter in the outer regions of galaxies.  For this work we obtained deep long slit spectra of three giant local galaxies, \\object{NGC1600}, \\object{NGC4125} and \\object{NGC7619} (see Table \\ref{tab_prop}). They were chosen as representative of the class massive ellipticals, that could be easily fit into the service mode schedule \\citep{2007PASP..119..556S} of the Hobby-Eberly Telescope (HET) \\citep{1998AAS...193.2602B} at McDonald Observatory.  These galaxies have been well studied in several works \\citep{1994MNRAS.269..785B,1995ApJ...448..119F}.  The galaxy \\object{NGC1600}, has boxy isophotes \\citep{2002A&A...391..531R} and does not show rotation both along major and minor axes. The galaxy \\object{NGC4125}, has dust and ionized gas aligned along the stellar major axis \\citep{1989ApJ...346..653K}; \\citet{1994A&AS..105..341G} also found that the dust lane is inclined by 10\\degr with respect to the apparent major axis; in addition, this galaxy shows strong soft (0.1-2.4~keV) X-ray and far-infrared ($\\lambda$ = 60 $\\mu$m) emission \\citep{1998AJ....116.2682C}.  For \\object{NGC7619} \\citet{2002A&A...391..531R} found evidence of tidal interactions with the galaxy NGC7626. Recently, \\citet{2006ApJ...636..698F} analyzed \\emph{Chandra} data and obtained that \\object{NGC1600} and \\object{NGC7619} show a positive temperature gradient towards the outer parts. On the other hand \\object{NGC4125} has a declining temperature profile.  Our spectra extend out to 1.5 $R_e$ for \\object{NGC1600} and more than 2$R_e$ for \\object{NGC4125} and \\object{NGC7619}.  This dataset allows us to constrain the general characteristics of the stellar kinematics and populations extending to large radii. By combining stellar dynamical modeling with the measured line strength indices we derive constraints on the extent of the dark matter halos in these galaxies.   The paper is organized as follows. In Sect. \\ref{sec_obsdat} we describe the observations (Sect.  \\ref{sec_obs}) and the data reduction (Sect. \\ref{sec_dat}). In Sect.  \\ref{sec_kinlin} we present the kinematics (Sect. \\ref{sec_kin}) and the line strength measurements (Sect. \\ref{sec_lin}). Comments on individual galaxies are given in Sect. \\ref{sec_comments}.  We analyze the Lick indices and derive ages, metallicities, and $\\alpha/Fe$ ratios in Sect. \\ref{sec_stelpop}. Section \\ref{sec_SPPmod} describes the models and the method used; Sect. \\ref{sec_SPPresults} presents the related results; Sect. \\ref{sec_colors} derives colors and mass-to-light ratios.  In Sect. \\ref{sec_dynamics} we produce dynamical models of the galaxies. We discuss the adopted method in Sect. \\ref{sec_dynmet}, the resulting models (Sect. \\ref{sec_dynres}), the constraints coming from the correlation between Mg\\emph{b} and local escape velocity (Sect. \\ref{sec_vesc}) and the resulting degeneracy between dark halo cut-off radii and outer density slope (Sect. \\ref{sec_deg}). A discussion of the results and a summary of this work are presented in Sect.  \\ref{sec_summary}  \\begin{table*} \\caption{List of the morphological classifications, B band luminosity, luminosity distance, receding velocity and effective radius of \\object{NGC1600}, \\object{NGC4125} and \\object{NGC7619}. They are taken from \\citet{1989ApJS...69..763F,2001MNRAS.328..461O} and NED(NASA/IPAC Extragalactic Database). \\label{tab_prop}} \\begin{center} \\begin{tabular}{lccccccr} \\hline Name   &Ra  & Dec & Type & log($L_B$/$L_\\odot$) & Dis & Vel(km/s)  &$R_e$( arcsec)\\\\ \\hline \\object{NGC1600} &  04h31m39s  &-05d11m10s &E3 &11.03 &64Mpc  &4681 &45   \\\\ \\object{NGC4125}  & 12h08m06s  & 65d10m26s &E6 &10.80 &20Mpc  &1356 &58  \\\\ \\object{NGC7619}  & 23h20m14s  &8d12m22s   &E2 &10.58 &46Mpc  &3760 &32  \\\\ \\hline \\end{tabular} \\end{center} \\end{table*} ",
        "conclusions": "\\label{sec_summary}  Accurate kinematic profiles extending out to 1.5 to 2 $R_e$ along the major and minor axis for the three giant elliptical galaxies \\object{NGC1600}, \\object{NGC4125} and NGC7169 have been measured, together with 6 lines strength Lick indices. For \\object{NGC4125} we also detected gas emission along the major axis and measured its kinematics and equivalent width strength. From the comparison of the [NI]$\\lambda\\lambda$5197,5200/H$\\beta$ vs [OIII]$\\lambda$5007/H$\\beta$ diagnostic diagram to the previous work, we concluded that the emission region is probably caused by LINER-like emission.  With the help of the SSP models, we derived the stellar population parameters, the M/L ratios and the broad band colors of our galaxies. We found the galaxies \\object{NGC1600} and \\object{NGC7619} have high metallicities, in contrast, the galaxy \\object{NGC4125} is just slightly above solar metallicity. The three objects have significant metallicity gradients.  All of the galaxies are overabundant, \\object{NGC4125} just by $0.1$ dex, and do not have significant gradients of element abundances along the axes.  These galaxies have sharp red peaks at the center, reflecting the steep change of metallicity and in agreement with the models predicted color profiles.  According to the simple element enrichment scenario, the $\\alpha$ elements are mainly delivered by Type II supernovae explosions of massive progenitor stars, and a substantial fraction of Fe peak elements come from the delayed exploding Type Ia supernovae \\citep{1984ApJ...286..644N,1996ApJ...460..408T}. Thus the $\\alpha/Fe$ can be used as an indicator to constrain the formation timescale of stars. Hence the absence of radial variations of $\\alpha$/Fe ratio in our galaxies likely suggests that there is no radial variation of star formation time scales.  The radial metallicity and lines strength gradients give one of the most stringent constraints on the galaxy formation. The galaxies that form monolithically have steeper gradients and the galaxies that undergo major mergers have shallower gradients.  The mean metallicity gradients for non-merger and merger galaxies derived by theoretical simulation in \\citet{2004MNRAS.347..740K} are $\\triangle$[Z/H]/$\\triangle$ log(r) $\\sim -0.30 \\pm 0.2$ and $-0.22 \\pm 0.2$, respectively. The author found that the galaxies with gradients steeper than -0.35 are all non-major merger galaxies.  As it can be seen from Table \\ref{tab_grad} the gradients $\\triangle [Z/H]/\\triangle \\log(r)$ of our galaxies are compatible with numerical simulations.  At face value, there is a weak indication that \\object{NGC1600} and \\object{NGC7619} were formed through a monolithic collapse process, while \\object{NGC4125} was shaped via a (recent) major merger as also indicated by minor-axis rotation, the presence of gas, young central stellar ages, just mild $\\alpha$/Fe-overabundance and the unrelaxed appearance of the outer stellar envelope.  Using the axisymmetric Schwarzschild's orbits superposition technique \\citep{1979ApJ...232..236S} with and without dark halos, we derived the local escape velocity of these galaxies.  From the correlation between Mg\\emph{b}, and the local escape velocity, we further confirmed the suggestion of \\citet{1990ApJ...359L..41F} that metallicity and lines strength are a function of local escape velocity. Moreover, we considered models with logarithmic dark matter halos of different sizes. The best fitting Mg\\emph{b}-$V_{esc}$ relation comes from models which cut off the logarithmic halo at 60 kpc. Similar cut-off radii are obtained when using the correlation between $[Z/H]$ and $V_{esc}$, but the scatter in this relation is larger, because it also depends on the stellar population age (see also \\citealt{2009MNRAS.398.1835S}).  We find that -- at a given escape velocity -- the youngest of our galaxies, \\object{NGC4125}, is also the most metal-rich one, which could be explained within a merger scenario as star formation from enriched material.  Our cut-off radii agree with the lensing analysis of cluster galaxies by \\citet{2007ApJ...656..739H} but not with the results from X-rays and strong-lensing \\citep{arXiv:0911.0678,2009ApJ...703L..51K}. Larger cut-off radii are compatible with the data if dark matter halos with steeper outer density profiles are considered.  We discover deviations from a linear correlation between log Mg\\emph{b} and log $V_{esc}$ in the outer parts of our galaxies, where the measured Mg\\emph{b} are lower than predicted by the extrapolation from the correlation further inside. The physical link between $V_{esc}$ and Mg\\emph{b} is established at the main time of star formation activity of the galaxy. Dry mergers happening after this episode will increase $V_{esc}$ without increasing Mg\\emph{b}. Hydrodynamical cosmological simulations by \\citet{2009ApJ...699L.178N} indicate that the outer parts of massive ellipticals (beyond $\\ga 5 \\, \\mathrm{kpc}$) might be dominated by stars which were born in low-mass halos and accreted by frequent minor mergers. If this would be the dominant process for the formation of the outer parts of giant ellipticals one would expect such a break-down of the linear correlation between log Mg\\emph{b} and log $V_{esc}$ as observed, because the outer Mg\\emph{b} would come from stars which were formed in small halos with low $V_{esc}$ (and accordingly low Mg\\emph{b}).  We plan to further investigate the connection between line strengths and dark matter halos with a more extended galaxy sample."
    },
    "1004/1004.1405_arXiv.txt": {
        "abstract": "{We present an analytical derivation of the Sachs Wolfe effect sourced by a primordial magnetic field. In order to consistently specify the initial conditions, we assume that the magnetic field is generated by a causal process, namely a first order phase transition in the early universe. As for the topological defects case, we apply the general relativistic junction conditions to match the perturbation variables before and after the phase transition which generates the magnetic field, in such a way that the total energy momentum tensor is conserved across the transition and Einstein's equations are satisfied. We further solve the evolution equations for the metric and fluid perturbations at large scales analytically including neutrinos, and derive the magnetic Sachs Wolfe effect. We find that the relevant contribution to the magnetic Sachs Wolfe effect comes from the metric perturbations at next-to-leading order in the large scale limit. The leading order term is in fact strongly suppressed due to the presence of free-streaming neutrinos. We derive the neutrino compensation effect dynamically and confirm that the magnetic Sachs Wolfe spectrum from a causal magnetic field behaves as $\\ell(\\ell+1)\\,C^\\B_\\ell \\propto \\ell^2$ as found in the latest numerical analyses.}    \\begin{document} ",
        "introduction": "The origin of the large scale magnetic fields observed in galaxies and clusters is still unknown: one of the possible explanations is that they have been generated in the primordial universe. A primordial magnetic field of the order of the nanoGauss could leave a detectable imprint in the cosmic microwave background (CMB) anisotropies. This has been analysed in several works: for the scalar mode, see for example \\cite{Kahniashvili:2006hy,Yamazaki:2008gr,Giovannini:2007qn,Finelli:2008xh,PFP,Shaw:2009nf,yamafin}. Here we concentrate on the effect that a primordial magnetic field can have on the temperature CMB spectrum at large scales: in particular, we focus on the Sachs Wolfe effect. The motivation is, that conflicting results are present in the literature regarding the $\\ell$-dependence of the Sachs Wolfe effect induced at large scales by a primordial magnetic field generated by a phase transition: the analytical analysis of \\cite{Kahniashvili:2006hy} found $\\ell(\\ell+1)\\,C^\\B_\\ell$ scaling as $ \\ell^{-1}$ or more negative, and the same result was found in the numerical calculation of \\cite{Yamazaki:2008gr}; on the other hand, \\cite{Finelli:2008xh,PFP,Shaw:2009nf,yamafin,Caprini:2009vk} found $\\ell(\\ell+1)\\,C^\\B_\\ell$ scaling as $\\ell^2$.  The aim of this paper is to solve this discrepancy analytically, and we find that the relevant contribution to the Sachs Wolfe effect is the one found in \\cite{Finelli:2008xh,PFP,Shaw:2009nf,yamafin,Caprini:2009vk}. Naively, the magnetic field anisotropic stress would induce a CMB spectrum $\\ell(\\ell+1)\\,C^\\B_\\ell\\propto \\ell^{-1}$ at large scales, basically due to the fact that the metric perturbation $\\Phi$ is proportional to $\\Phi \\propto \\Pi_\\B (\\HH/k)^2$ at leading order in the large scale expansion $k/\\mathcal{H}\\ll 1$, where $\\Pi_\\B$ is the magnetic field anisotropic stress. However, as soon as neutrinos decouple and start to free-stream, they develop a non-zero anisotropic stress which adjusts to compensate the one coming from the magnetic field, see \\cite{Kojima:2009gw}. We demonstrate here that this compensation cancels the leading order contribution to the CMB spectrum, and the dominant contribution becomes the one from the next-to-leading order in the $k/\\mathcal{H}\\ll 1$ expansion, namely $(k/\\mathcal{H})^0$, which induces then $\\ell(\\ell+1)\\,C^\\B_\\ell \\propto \\ell^2$. A residual $\\ell^{-1}$ contribution to the Sachs Wolfe remains: this is the relic effect of the period of time after the magnetic field generation but before the decoupling of neutrinos. However, this contribution is completely unobservable in the CMB today.  In this paper we concentrate on causal magnetic fields, generated at a phase transition in the radiation dominated universe before neutrino decoupling (see for example \\cite{phase}). We make this choice in order to be able to consistently set the initial conditions (in a following work we will analyse the inflationary generation case \\cite{inprep}). We assume that the magnetic energy momentum tensor is first order in perturbation theory, and that the magnetic field has no background contribution: it is a stochastic primordial magnetic field with spectral index $n\\geq 2$ because of its causal generation \\cite{Durrer:2003ja} \\be \\vev{B_i(\\bk)B_j^*(\\bq)}=(2\\pi)^3\\de(\\bk-\\bq) (\\de_{ij}-\\hat{k}_i\\hat{k}_j) \\, A \\, k^n~, \\ee where $A$ is the amplitude of the spectrum. The magnetic energy momentum tensor is therefore automatically gauge invariant. Moreover, we work under the one-fluid MHD approximation, meaning that the conductivity of the universe is high, so that we can neglect the electric field and charge separation phenomena, occurring at very small scales. Since we are ultimately interested in the CMB spectra at large scales, we also neglect the presence of baryons and work under the tight coupling approximation. The universe is therefore composed only of a radiation component (photons and massless neutrinos - we also neglect neutrino masses), a pressure-less matter component (cold dark matter), and the magnetic field.  In the first part of the paper, we solve Einstein's equations in the long wavelength limit, \\ie consistently neglecting all terms proportional to $(k/\\mathcal{H})^2$. By doing so, we can find analytic solutions for the gauge invariant variables describing the scalar metric perturbations and those of the total (radiation plus matter) fluid, which hold at scales larger than the horizon. Indeed, to calculate the Sachs Wolfe effect, we need the metric perturbations at scales larger than the horizon at recombination.  The solutions are completely determined once the initial conditions have been specified. To set the initial conditions, we follow what has been done for the topological defect case \\cite{Deruelle:1997py,Uzan:1998rt}. First of all, we assume a sudden phase transition, as for example the electroweak phase transition. The magnetic field is then instantaneously generated, if one considers perturbations with length scale relevant for the CMB today: $k / \\mathcal{H}_{\\rm B} \\ll 1$, where $ \\mathcal{H}_{\\rm B}$ is the Hubble time during the phase transition, corresponding approximatively to its duration. Before the phase transition, the universe is the usual perturbed FLRW universe with only the standard inflationary adiabatic perturbations; after the phase transition, the magnetic field energy momentum tensor contributes at first order to the metric perturbations. In order to connect these two stages, we match the geometry and the fluid variables on the surface of constant density, so that the induced three metric and the extrinsic curvature are continuous \\cite{Deruelle:1997py}. This implies the conservation of the total energy momentum tensor at the magnetic field generation time.  With this procedure, we find the solutions for the metric and fluid perturbation variables. Each variable is given by the sum of the adiabatic inflationary mode plus the magnetic field contribution, which is null before the transition. We have no freedom on how to add the magnetic field contribution: this is entirely specified by the matching. The solutions are such that the total fluid energy density and momentum are conserved at large scales, and the magnetic field does not alter the total curvature perturbation at large scales: this is given only by the inflationary contribution. Therefore, the matching procedure selects an isocurvature magnetic mode which leaves the curvature unchanged.  After neutrino decoupling, the solutions change again due to the non-zero neutrino anisotropic stress. In order to include analytically the free-streaming neutrinos in our analysis, we solve the Bardeen equation combined with the neutrino conservation equations, and find an analytical fit for the time evolution of the neutrino anisotropic stress. Then we introduce this fit back into Einstein's equation and solve for the metric and fluid perturbations including the contribution from the free-streaming neutrinos. With these complete, analytical solutions we then evaluate the Sachs Wolfe effect induced by the presence of the causal magnetic field.  We confirm that the free-streaming neutrinos have a fundamental impact on the magnetic Sachs Wolfe \\cite{Finelli:2008xh,PFP,Shaw:2009nf,yamafin}. Neglecting their presence one would conclude that the temperature anisotropy due to the magnetic field is proportional to $\\Pi_\\B/(k\\eta_1)^2$, where $\\eta_1$ denotes approximatively the conformal time at recombination. Therefore, the magnetic anisotropic stress would completely dominate the large scale Sachs Wolfe effect with respect to the magnetic energy density, and would induce a CMB spectrum of the type $\\ell(\\ell+1)\\,C^\\B_\\ell\\propto \\ell^{-1}$. However, as already pointed out in \\cite{Shaw:2009nf,Kojima:2009gw}, the neutrino anisotropic stress acts to compensate and reduce the magnetic field one. As a result, the Sachs Wolfe contribution correctly accounting for neutrinos becomes of the form $f(\\eta_\\B,\\eta_\\nu,\\eta_{\\rm rec})\\,\\Pi_\\B/(k\\eta_1)^2$, where $f(\\eta_\\B,\\eta_\\nu,\\eta_{\\rm rec})$ is a function of conformal time at recombination $\\eta_{\\rm rec}$, of the magnetic field generation time $\\eta_\\B$ and of the neutrino decoupling time $\\eta_\\nu$, which strongly suppresses the magnetic Sachs Wolfe contribution. Therefore, the contribution to the magnetic Sachs Wolfe behaving as $1/(k\\eta_1)^2$ is only due to the evolution of the metric perturbations between the magnetic generation time and neutrino decoupling.  This contribution has been neglected in the literature \\cite{Finelli:2008xh,PFP,Shaw:2009nf,yamafin}, because it has been interpreted as connected to a decaying mode. However, with our analytical approach, we can show that the magnetic field solution does not behave as the standard inflationary one: one cannot really distinguish a growing and decaying mode, since all modes have a comparable amplitude at horizon crossing, and would therefore leave the same imprint on the CMB. The reason why this contribution is negligible only resides in the action of free-streaming neutrinos, the anisotropic stress of which counteracts the magnetic field one and effectively cancels the large scale temperature anisotropy. The new contribution to the Sachs Wolfe effect can only act as a source of anisotropy after the magnetic field generation time but before neutrino decoupling, and is negligible. Therefore, we finally confirm the result present in the literature \\cite{Finelli:2008xh,PFP,Shaw:2009nf,yamafin}: the dominant contribution to the magnetic Sachs Wolfe effect becomes the one coming from the metric perturbations at the following order in the long wavelength expansion, \\ie at order $(k\\eta_1)^0$. This is what has been calculated in all the numerical analyses of the CMB temperature anisotropy. In the second part of the paper, we therefore proceed to evaluate analytically this contribution.  In order to solve Einstein's equation at next-to-leading order, \\ie including terms of the order $(k/\\mathcal{H})^2$, the easiest way is to solve for the curvature perturbation equation. This allows us to find the solution for the Bardeen potentials at order $(k\\eta_1)^0$, and consequently all the other fluid variables. We can then evaluate the Sachs Wolfe contribution at this order.  The structure of the paper is the following: in section \\ref{sec:Einstein} we derive Einstein's and conservation equations in the presence of a magnetic field. In section \\ref{init_cond}, we explain the matching procedure, used then in section \\ref{long} to calculate the metric and fluid variables at leading order in the $k\\eta_1\\ll 1$ expansion, both before and after neutrino decoupling. In section \\ref{sec:SW} we compute the Sachs Wolfe effect from the leading order solution, and we conclude that the relevant contribution comes from the next-to-leading order. Therefore, in section \\ref{nexttoleading} we calculate the next-to-leading order solutions, and we conclude in section \\ref{sec:conc}. We consider scalar metric perturbations on a spatially flat Friedmann background, and we work with gauge invariant variables using the notations of \\cite{KS} (with respect to \\cite{Ma:1995ey}, we have $\\Phi$ with the opposite sign while $\\Psi$ is the same). We normalise the scale factor to one at equality, so that in a matter plus radiation universe $a=y^2+2y$ in terms of the dimensionless variable $y=\\eta/\\eta_1$, where $\\eta$ denotes conformal time, $\\eta_1=\\eta_{\\rm eq}/(\\sqrt{2}-1)\\simeq \\eta_{\\rm rec}$ and $\\eta_{\\rm eq}$ represents conformal time at equality. Scales larger than the horizon at recombination satisfy $x_1 \\equiv k\\eta_1\\ll 1$. A dot denotes derivative with respect to conformal time, while a prime denotes derivative with respect to $y$. Greek indexes go from 0 to 3, while latin ones from 1 to 3. The neutrino background energy density fraction is denoted $R_\\nu=\\bar\\rho_\\nu/\\bar\\rho_{\\rm rad}\\simeq 0.4$. ",
        "conclusions": "\\label{sec:conc}  In this work we present an analytical computation of the Sachs Wolfe effect induced by a primordial magnetic field. We have restricted our analysis to a magnetic field generated by a causal process, which acts on a time scale shorter than the Hubble time at generation, such as a first order phase transition in the early universe. The reason is that, under this hypothesis, the initial conditions for the metric and fluid variables are determined unambiguously by imposing conservation of the total energy momentum tensor across the transition, just as in the topological defects case \\cite{Deruelle:1997py,Uzan:1998rt}. After setting the initial conditions with this matching, we have solved Einstein's and conservation equations analytically and computed the magnetic effect on the gauge invariant metric and fluid perturbations in the large scales limit $k\\eta_1\\ll 1$, \\ie for scales larger than the horizon at recombination. Using these solutions, we have then determined the Sachs Wolfe effect. This is sourced only by the magnetic anisotropic stress: the temperature perturbation behaves as $\\Delta T/T\\propto \\Pi_\\B/(k\\eta_1)^2$ at leading order in the $k\\eta_1\\ll 1$ expansion. A scaling of the CMB spectrum as $\\ell(\\ell+1)C^\\B_\\ell\\propto \\ell^{-1}$ trivially follows.  However, as already pointed out in \\cite{Shaw:2009nf,Kojima:2009gw}, the magnetic anisotropic stress is quickly compensated by the neutrino one, once neutrinos decouple from the primordial fluid. We have verified this compensation analytically, by solving the neutrino evolution equations in the limit $k\\eta_1\\ll 1$, where we can consistently neglect multipoles higher than the anisotropic stress. Consequently, the contribution to the Sachs Wolfe at leading order in $k\\eta_1\\ll 1$ gets strongly suppressed after neutrino decoupling, and the next-to-leading order in $k\\eta_1\\ll 1$ becomes the relevant one: the one  which could in principle provide an observable effect in the CMB. In order to derive the magnetic Sachs Wolfe effect at next-to-leading order in $k\\eta_1\\ll 1$ we had to compute the metric and fluid variables consistently at this order. This can be done analytically by solving the evolution equation for the curvature perturbation, and leads to a temperature anisotropy which is sourced both by the magnetic anisotropic stress and the magnetic energy density: we found the approximate result $\\Delta T/T\\simeq -0.2 \\, \\Om_\\B -2.7 \\,\\Pi_\\B$. The new scaling of the CMB spectrum,  arising from this contribution, takes the form $\\ell(\\ell+1)C^\\B_\\ell\\propto \\ell^{2}$.  The matching conditions which we used to determine the initial conditions for the magnetic field perturbations imply that the curvature perturbation is continuous across magnetic field generation time. The subsequent evolution of the curvature perturbation is then completely determined. We found that, at leading order in $k\\eta_1\\ll 1$, the magnetic field does not contribute to the curvature, which is therefore conserved at large scales and given only by the inflationary mode (for which we have imposed adiabatic initial conditions). At next-to-leading order in $k\\eta_1\\ll 1$, the curvature perturbation is instead affected by the presence of the magnetic field. As already found in \\cite{Shaw:2009nf}, we confirm that, between the magnetic field generation time and neutrino decoupling, the curvature is sourced by the magnetic anisotropic stress and therefore grows (starting from zero at magnetic field generation time). After neutrino decoupling, when the magnetic anisotropic stress gets compensated, the part sourced by it remains constant, and the curvature perturbation gets a contribution from the magnetic energy density which is no longer compensated.  Our analysis provides a clarification of the results present in the literature on the Sachs Wolfe effect arising from a primordial magnetic field. The magnetically induced CMB anisotropies have been studied in details in the literature by numerical integration of Boltzmann codes \\cite{Finelli:2008xh,PFP,Shaw:2009nf,yamafin}, and in most of the cases the large scale CMB spectrum behaves as  $\\ell(\\ell+1)C^\\B_\\ell\\propto \\ell^{2}$, corresponding to a temperature anisotropy which does not depend on wavenumber. This is a consequence of the choice of the initial conditions. A Boltzmann code needs the initial conditions to be specified after neutrino decoupling, so the contribution to the Sachs Wolfe effect arising from the period of time between the magnetic field generation and the neutrino decoupling has been neglected in \\cite{Finelli:2008xh,PFP,Shaw:2009nf,yamafin}, and the magnetic field anisotropic stress is exactly compensated by the free-streaming neutrino one since the beginning. This sets to zero the temperature anisotropy coming from the metric and fluid perturbations at leading order in $k\\eta_1\\ll 1$ to start with. This choice of the initial conditions has been motivated by an analogy with the inflationary case, for which one can identify a growing and a decaying mode, and consistently neglect the decaying one. In fact, the inflationary solution for the Bardeen potential is, for $y\\ll 1$ \\be \\Phi^- (y) \\simeq  \\frac{3}{4\\,y^3}\\,\\frac{b_1}{x_1^2}+\\frac{20}{3}\\,\\frac{a_1}{x_1^2}= \\frac{3}{4}\\, \\frac{\\eta_1}{\\eta}\\,\\frac{b_1}{(k\\eta)^2}+\\frac{20}{3}\\left(\\frac{\\eta}{\\eta_1}\\right)^2 \\frac{a_1}{(k\\eta)^2}~, \\ee where the mode proportional to $b_1$ is decaying with time, while the one proportional to $a_1$ is constant. One expects that from inflation the two modes are generated with the same amplitude: consequently, $b_1$ must be much smaller than $a_1$ since, at very early times (\\eg reheating), $\\eta_1/\\eta_{\\rm reh}\\gg 1$ \\cite{libroruth}. The Sachs Wolfe effect from both modes at leading order in $k\\eta_1\\ll 1$ is given by (\\cf also \\eqref{dToverTinf}) \\be \\left. \\frac{\\De T}{T}\\right|_{\\rm inf} \\simeq -\\frac{12(1+y_{\\rm rec})}{y_{\\rm rec}^3(2+y_{\\rm rec})^3} \\, \\frac{b_1}{x_1^2}-\\frac{2(40+40y_{\\rm rec}+18y_{\\rm rec}^2+3y_{\\rm rec}^3)}{3(2+y_{\\rm rec})^3}\\, \\frac{a_1}{x_1^2}~. \\ee The two modes have the same dependence on $k$, and about the same amplitude for $y_{\\rm rec}\\simeq 1$: therefore, knowing that $b_1\\ll a_1$, one can neglect the contribution of the decaying mode (for an analysis of the effect of the inflationary decaying mode on the CMB, see \\cite{Amendola:2004rt}).  The situation for the metric perturbation induced by the magnetic field is a bit different. In this case, the solution at leading and next-to-leading order in $k\\eta_1\\ll 1$ reads, for $y\\ll 1$ (\\cf also \\eqref{Phiplus} and \\eqref{PhiBnext}) \\bea \\Phi^B (y) & \\simeq & -\\frac{3}{y^2}\\left(1-\\frac{y_\\B}{y}\\right) \\frac{\\Pi_\\B}{x_1^2}+\\left[ -\\frac{5}{9}+\\frac{2}{3}\\log\\left(\\frac{y}{y_\\B}\\right)+\\frac{y_\\B}{2y}\\right]\\Pi_\\B \\\\ & = & -3\\left(1-\\frac{\\eta_\\B}{\\eta}\\right) \\frac{\\Pi_\\B}{(k\\eta)^2}+\\left[ -\\frac{5}{9}+\\frac{2}{3}\\log\\left(\\frac{\\eta}{\\eta_\\B}\\right)+\\frac{\\eta_\\B}{2\\eta}\\right]\\Pi_\\B~. \\nonumber \\eea The term at leading order in $k\\eta_1\\ll 1$ is decaying with time as $y^{-2}$, whereas the next-to-leading order is growing logarithmically. One could therefore be tempted to neglect the first term on the right hand-side of the above equation. However, the key point here is that, contrary to the inflationary case, the two modes not only have no independent amplitudes (they are both proportional to $\\Pi_\\B$), but also do not share the same $k$-dependence. Consequently, we see that the `decaying' mode always has the same amplitude of the `growing' one at horizon crossing $k\\eta\\simeq 1$, proportional to a few $\\Pi_\\B$. Since the CMB in general selects the contribution of modes which cross the horizon at recombination time $\\eta_1$ (the first peak), this means that the two modes can have the same impact on the CMB for $k\\eta_1\\simeq 1$. Moreover, for the Sachs Wolfe effect, which selects modes still outside the horizon at recombination, we see that the `decaying' mode could even dominate.  However, the analyses of \\cite{Finelli:2008xh,PFP,Shaw:2009nf,yamafin} do get the correct result: as we have seen, this is due to the compensating effect of the neutrinos, which exactly cancels $\\Pi_\\B$ at leading order in $k\\eta_1\\ll 1$, strongly suppressing the amplitude of the `decaying' mode. The relevant contribution to the CMB anisotropies is therefore caused by the `growing' mode\\footnote{Note that, at next-to-leading order in $k\\eta_1\\ll 1$ the cancellation from the neutrinos is not active: \\cf the different amplitudes of \\eqref{pinuapp} and \\eqref{pinunext}}. The correct CMB spectrum at large scales from the Sachs Wolfe effect from a causally generated magnetic field with spectral index $n\\geq 2$ is therefore $\\ell(\\ell+1)C^\\B_\\ell\\propto \\ell^{2}$. Note however, that in order for this effect to be observable, the magnetic field should have an amplitude of at least a few nanoGauss: causal magnetic fields generated before Nucleosynthesis are unfortunately strongly constrained to amplitudes far too small to leave any observable effect in the CMB \\cite{Caprini:2001nb}."
    },
    "1004/1004.4016_arXiv.txt": {
        "abstract": "We define geometric axioms for the metric and the connection of a spacetime where the gravitational influence of the connection may be interpreted as dark matter.  We show how these axioms lead to the Einstein-Klein-Gordon equations with a cosmological constant, where the scalar field of the Klein-Gordon equation represents the deviation of the connection from the standard Levi-Civita connection on the tangent bundle and is interpreted as dark matter.  This form of dark matter, while not quantum mechanical, gives virtually identical predictions to some other scalar field dark matter models, including boson stars, which others have shown to be compatible with the $\\Lambda CDM$ model on the cosmological scale. In addition, we quantify the already known fact that this scalar field dark matter, unlike the WIMP model, is automatically cold in a homogeneous, isotropic universe, sufficiently long after the Big Bang.  With these motivations in mind, we show how this scalar field dark matter, which naturally forms dark matter density waves due to its wave nature, may cause the observed barred spiral pattern density waves in many disk galaxies and triaxial shapes with plausible brightness profiles in many elliptical galaxies.  If correct, this would provide a unified explanation for spirals and bars in spiral galaxies and for the brightness profiles of elliptical galaxies. We compare the results of preliminary computer simulations with photos of actual galaxies. ",
        "introduction": "There are three ideas, each of independent interest, which, if correct, would create a new connection between differential geometry and astronomy, very much in the tradition of general relativity's previous successes at describing the large-scale structure of the universe.  We begin by discussing these ideas in general terms.  \\underline{Idea 1}:  Natural geometric axioms motivate studying the Einstein-Klein-Gordon equations with a cosmological constant.  The Klein-Gordon equation is a wave type of equation for a scalar field which we propose as a model for dark matter.  While this geometric motivation is new, modeling dark matter with a scalar field satisfying the Klein-Gordon equation is not (see \\cite{JWLee} for a survey of scalar field dark matter and boson stars).  Hence, ideas 2 and 3 apply to these other works as well.  \\underline{Idea 2}:  Wave types of equations for matter fields, such as the Klein-Gordon equation, naturally form density waves in their matter densities because of constructive and destructive interference, like waves on a pond, or the Maxwell equations for electromagnetic radiation. Unlike waves on a pond or the Maxwell equations, however, the group velocities of wave solutions to the Klein-Gordon equation (with positive ``mass'' term) can be anything less than the speed of light, and can be arbitrarily slow for wavelengths which are long enough. This allows for the possibility of gravitationally bound ``blobs'' of dark matter to form.  In this paper we make some conjectures about these scalar field dark matter density waves.  \\underline{Idea 3}:  Density waves in dark matter, through gravity, naturally form density waves in the regular baryonic matter.  In the case of disk galaxies, where friction in the interstellar medium of gas and dust is important \\cite{BT}, we exhibit examples where barred spiral density wave patterns form in the regular matter, as seen in figures 1, 2, 3, and 4. In the case of elliptical galaxies, where the interstellar medium is believed to be mostly irrelevant \\cite{BT}, we show how these dark matter density waves tend to produce triaxial ellipsoidal shapes for the regular matter with plausible brightness profiles, as seen in figures 5, 6, and 7.  \\begin{figure}\\label{NGC1300} \\begin{center} \\includegraphics[height=75mm]{NGC1300.jpg} \\includegraphics[height=75mm]{NGC1300_1_75000_1_-1_2000_1990_25000000_8p7e-013_7500_5000_25000000_10000P10p1147_1000_5Pinv.jpg} \\end{center} \\caption{NGC1300 on the left, simulation on the right.  The simulated image on the right results from running the Matlab function spiralgalaxy(1, 75000, 1, -1, 2000, 1990, 25000000, 8.7e-13, 7500, 5000, 25000000, 10000) described in section \\ref{SGS1}. Left photo credit: Hillary Mathis/NOAO/AURA/NSF.  Date:  December 24, 2000. Telescope:  Kitt Peak National Observatory's 2.1-meter telescope. Image created from fifteen images taken in the BVR pass-bands. } \\end{figure}  This paper is an attempt to put the above three ideas together in as precise a way as possible.  The rules of the game we choose to play here are strict:  define a concise set of geometric axioms, and then try to understand the implications of those axioms.  The axioms we choose, stated in the next section, are more fundamental than defining an action. Instead, we declare the spacetime metric and connection on the tangent bundle as the fundamental objects of our universe, and then define the properties that the action for these two objects must have.  In this paper we make the case that the simulated images in figures 1-6 as well as the computed brightness profile described by figure 7 are all consequences of these axioms.  The resulting theory is a generalization of the vacuum Einstein equations with a cosmological constant, which already famously explains gravity, 73\\% of the mass of the universe as dark energy \\cite{WMAPobservations}, the accelerating expansion of the universe, black holes, and all other vacuum general relativity effects. If successful, this generalized theory would then, by describing dark matter, account for 95\\% of the mass of the universe \\cite{WMAPobservations} and explain some portion of the structures of galaxies.  \\begin{figure}\\label{NGC4314} \\begin{center} \\includegraphics[height=75mm]{NGC4314.jpg} \\includegraphics[height=75mm]{NGC4314_1_75000_1_-0p15_2000_1990_25000000_8p7e-013_7500_5000_30000000_50000P10p1147_1000_5P.jpg} \\end{center} \\caption{NGC4314 on the left, simulation on the right.  The simulated image on the right results from running the Matlab function spiralgalaxy(1, 75000, 1, -0.15, 2000, 1990, 25000000, 8.7e-13, 7500, 5000, 30000000, 50000) described in section \\ref{SGS2}. Left photo credit:  G. Fritz Benedict, Andrew Howell, Inger Jorgensen, David Chapell (University of Texas), Jeffery Kenney (Yale University), and Beverly J. Smith (CASA, University of Colorado), and NASA.  Date:  February 1996. Telescope:  30 inch telescope Prime Focus Camera, McDonald Observatory.} \\end{figure}  However, strictly speaking, quantum mechanics is not part of the theory we propose here, as should be expected since general relativity and quantum mechanics have yet to be unambiguously unified.  Hence, our theory is clearly incomplete.  This is not so bad considering that every theory known today is, in the strictest sense, incomplete. However, it is still a reasonable question to wonder if our theory does a good job of describing dark matter, even if it does not describe regular particulate matter. Hence, as a way of testing our dark matter model, we treat the remaining 4\\% of the mass of the universe, composed of particles of various kinds, in the traditional manner similar to test particles, but with mass, which arguably makes this theory compatible with $\\Lambda CDM$ models, as will be explained in section \\ref{cosmologicalpredictions}.  The simulated images in figures 1-6 are pictures of the effect of the dark matter on the regular matter that we have sprinkled into the theory. Finding a less contrived way of getting regular matter into the next theory, while still respecting the idea of keeping our axioms as simple as possible, is an important open problem  So, does the Klein-Gordon equation accurately describe dark matter and predict some observed properties and structures of galaxies? In this paper we present evidence of this possibility by trying to understand the effect that this model of dark matter would have on the structure of galaxies, which is a reasonable idea since galaxies have large components of dark matter.  In doing so we have had to make approximations and educated guesses, so the comparisons in figures 1-6, while encouraging, should be taken in this context. Also, our ``simulations'' of galaxies only simulate the effect of the dark matter on the regular matter and hence are very primitive.  Perhaps a better name would be ``numerical experiments.'' However, one has to start somewhere, and it is already interesting that compelling patterns very much resembling actual galaxies have emerged.  We will describe the models we have used in sections \\ref{GalacticScale}, \\ref{SpiralGalaxies}, and \\ref{EllipticalGalaxies} and the assumptions we have made. We will do our best to clearly label where we have had to make approximations and educated guesses, as well as rigorous arguments, so that readers may make their own judgments about what is presented here.  \\begin{figure}\\label{NGC3310} \\begin{center} \\includegraphics[height=75mm]{NGC3310.jpg} \\includegraphics[height=75mm]{NGC3310_1_75000_1_-0p15_2000_1990_100000000_8p7e-013_7500_5000_45000000_50000P2p5275_1000_5P.jpg} \\end{center} \\caption{NGC3310 on the left, simulation on the right.  The simulated image on the right results from running the Matlab function spiralgalaxy(1, 75000, 1, -0.15, 2000, 1990, 100000000, 8.7e-13, 7500, 5000, 45000000, 50000) described in section \\ref{SGS3}.  Left photo credit: NASA and The Hubble Heritage Team (STScI/AURA).  Acknowledgment: G.R. Meurer and T.M. Heckman (JHU), C. Leitherer, J. Harris and D. Calzetti (STScI), and M. Sirianni (JHU).  Dates:  March 1997 and September 2000.  Telescope:  Hubble Wide Field Planetary Camera 2.} \\end{figure}      In this appendix we provide a brief introduction to general relativity relevant to the next section where we show how the Einstein-Klein-Gordon equations with a cosmological constant result from Axiom \\ref{A1}.  The differential geometry discussed in these two appendices is detailed but is also very standard.  Our discussion will not require much knowledge beyond the first three chapters of \\cite{ONeill}, which we recommend as an excellent text book for those who want to use differential geometry to study the large scale structure of the universe.  Special relativity, Einstein's theory which unified space and time, is really all about the study of the geometry of the flat Minkowski spacetime \\cite{ONeill}.  The defining idea leading to general relativity is the idea that spacetime is not necessarily flat. Philosophically, the question ``Why should spacetime be curved?'' has a great response:  ``Why should spacetime be flat?''  What is remarkable is that removing the assumption that spacetime is flat leads to general relativity as arguably the simplest way of defining what the spacetime metric should be, as we will explain.  Moreover, experiments have shown that general relativity describes gravity as accurately as can be measured.  The central equation of general relativity is the Einstein equation \\[ G = 8\\pi T, \\] where $G = Ric - \\frac12 R g$ is the Einstein curvature tensor and $T$, called the stress energy tensor, describes the local energy and momentum density.  However, it is important to realize that this equation by itself is only half a theory in that this equation describes how the matter curves the metric but does not explain how the matter evolves forward in time.  From this point on we will only discuss complete theories where matter evolution equations are specified, fully realizing that one can always add in an unspecified stress energy tensor later if desired.  For example, \\[ G = 0, \\] the vacuum Einstein equation, is a complete theory in our terminology because there is no matter term to evolve.  It is worth pointing out that vacuum general relativity is already an intricate and very interesting theory.  For example, black holes are solutions to the vacuum Einstein equation when all of the matter has already fallen into the black holes.  This demonstrates an interesting characteristic of general relativity - zero local energy and momentum density does not imply that the system has zero total mass. In fact, by the positive mass theorem \\cite{SY3}, \\cite{SY5}, an isolated asymptotically flat vacuum solution must have strictly positive total mass, unless the spacetime is flat everywhere (and hence is the Minkowski spacetime).  Also, matter can be approximated by black holes of various sizes. For example, one could replace the Sun and the planets by black holes of equal mass, and the corresponding n-body problem with black holes is a solution to $G=0$.  For that matter, one could replace every star, planet, meteor, and piece of dust in a galaxy with a black hole of equal mass, and the resulting system would be a solution to $G=0$.  The dynamics would be a good approximation to the original system except when the objects got close to one another.  All of the above statements are true assuming that an existence theory for the vacuum Einstein equation $G = 0$ generically exists, which is currently an open problem.  The difficulty in even proving an existence theory for $G = 0$ is another indication of the highly nontrivial nature of the equation.  Einstein chose $G = Ric - \\frac12 R g$ once he learned that the divergence of $G$ equaled zero by the second Bianchi identity \\cite{ONeill}. This zero divergence property can be interpreted as a local conservation property for energy and momentum density, a understandably desirable property. Hilbert realized that the zero divergence property of $G$ was a direct result of the fact that metrics which are critical points of the functional \\begin{equation}\\label{eqn:vacuumGR} H_U(g) = \\int_{U} R \\; dV, \\end{equation} satisfy the corresponding Euler-Lagrange equation $G = 0$ on the spacetime $N$.  The role of the open set $U$ in the above equation, whose closure is required to be compact, is to keep the integral finite. We define a metric $g$ to be a critical point of $H_U(g)$ if this quantity does not change to first order for all smooth variations of the metric with compact support in the interior of $U$.  Note that we are not minimizing or maximizing anything but are simply looking for critical points.  Moreover, Hilbert realized that the zero divergence property is true for any Euler-Lagrange equation resulting from an action which is invariant under reparameterizations of the metric, as the Einstein-Hilbert action in equation \\ref{eqn:vacuumGR} is.  Today we think of this fact as a special case of Noether's theorem.  Hence, the zero divergence property of $G$, while nice, is nothing particularly special and is not the fundamental point. Instead, vacuum general relativity is best understood as defining the metric to be a critical point of the Einstein-Hilbert action in equation \\ref{eqn:vacuumGR}. This raises the truly relevant question:  What is so special about the Einstein-Hilbert action?  We begin by noting that in some ways the Einstein-Hilbert action is the simplest nontrivial scalar valued functional of a metric. The scalar curvature of a spacetime metric is a scalar at each point, so it is perfectly natural to integrate it.  However, there is another characterizing property of this functional which is seen by expressing it in a coordinate chart.  Given a coordinate chart mapping $\\Phi: \\Omega \\rightarrow R^4 = \\{(x^0, x^1, x^2, x^3) \\}$, define \\begin{equation} g_{ij} = g(\\partial_i, \\partial_j),\\;\\;\\; \\mbox{ where } \\;\\;\\; \\partial_i = D\\Phi^{-1} \\left( \\frac{\\partial}{\\partial x^i} \\right) \\end{equation} to be the components of the metric in this coordinate chart.  Then standard calculations show that the formula for the scalar curvature in terms of the metric in a coordinate chart is \\begin{eqnarray}\\nonumber R  &=& (g^{ik}g^{jl} - g^{ij}g^{kl}) g_{ij,kl} + g_{ij,k}g_{ab,c} \\cdot \\\\&& \\left(\\frac34 g^{ia}g^{jb}g^{kc} - \\frac12 g^{ia}g^{jc}g^{kb} -g^{ia}g^{jk}g^{bc} - \\frac14 g^{ij}g^{ab}g^{kc} + g^{ij}g^{ac}g^{kb} \\right) \\end{eqnarray} where commas denote coordinate chart derivatives and the Einstein summation convention of summing over indices which are both raised and lowered is in effect.  As usual, we define $[g^{ij}] = [g_{ij}]^{-1}$ as matrices. Also, the formula for the metric volume form in terms of the metric in a coordinate chart is \\begin{equation} dV = |g|^{1/2}  \\; dV_{R^4} \\end{equation} where $|g| = |det([g_{ij}])|$.  Using $|g|_{,l} = |g|\\,g^{mn}g_{mn,l}$ and $g^{ij}_{\\hspace{.1in},k} = - g^{ia}g^{jb}g_{ab,k}$ we can integrate by parts over an open set $U$ (whose closure is compact) in the coordinate chart to get \\begin{eqnarray}\\nonumber \\int_U R \\; dV  &=& \\int_{\\partial(\\Phi(U))} (g^{ik}g^{jl} - g^{ij}g^{kl}) g_{ij,k} \\nu_l  |g|^{1/2} dA_{R^4} +\\int_{\\Phi(U) } g_{ij,k}g_{ab,c}\\cdot \\\\&& \\left(-\\frac14 g^{ia}g^{jb}g^{kc} + \\frac12 g^{ia}g^{jc}g^{kb} + \\frac14 g^{ij}g^{ab}g^{kc} -\\frac12 g^{ij}g^{ac}g^{kb} \\right) |g|^{1/2}  \\; dV_{R^4} \\end{eqnarray} where $\\nu_l = \\langle \\frac{\\partial}{\\partial x^l}, \\nu \\rangle_{R^4}$ and $\\nu$ is the outward unit normal to $\\Phi(U)$ in the $R^4$ coordinate chart.  Note that our formula for the Einstein-Hilbert action is of the form \\begin{equation}\\label{eqn:scalarcurvquad} \\int_{U} R \\; dV = \\int_{\\Phi(U)} \\mbox{Quad}_M(M') \\; dV_{R^4} \\;\\;\\;+\\;\\;\\; \\mbox{boundary term}, \\end{equation} where \\[ M = \\{g_{ij}\\} \\;\\;\\;\\mbox{ and }\\;\\;\\;  M' = \\{g_{ij,k}\\} \\] are the components of the metric in the coordinate chart and all of the first derivatives of these components in the coordinate chart, and where \\begin{equation} \\mbox{Quad}_{Y}(\\{x_\\alpha\\}) = \\sum_{\\alpha,\\beta} F^{\\alpha\\beta}(Y) x_\\alpha x_\\beta \\end{equation} for some functions $\\{F^{\\alpha\\beta}\\}$ is a quadratic expression of the $\\{x_\\alpha\\}$ with coefficients in $Y$. What is interesting about this particular quadratic expression is that while its value at each point depends on the choice of coordinate chart, by equation \\ref{eqn:scalarcurvquad} its integral over $\\Phi(U)$ is invariant with respect to smooth variations of $\\Phi$ compactly supported in the interior of $U$ (which hence do not affect the boundary term). Even more intriguing, this quadratic expression is the unique one with this property, up to multiplication by a constant, which follows from the works of Cartan \\cite{Car22}, Weyl \\cite{Wey22}, and Vermeil \\cite{Ver17} (and pursued further by Lovelock \\cite{Lov72}).  Hence, if we define \\begin{equation} F_{\\Phi,U}(g) = \\int_{\\Phi(U)} \\mbox{Quad}_M(M') \\; dV_{R^4} \\end{equation} for that same quadratic expression, then \\begin{equation} {H}_U(g)= \\,F_{\\Phi,U}(g) \\;+\\; \\mbox{boundary term} \\end{equation} by equation \\ref{eqn:scalarcurvquad}. Hence, the Euler-Lagrange equation for $F_{\\Phi,U}(g)$, for variations away from the boundary, must be the same as the one for ${H}_U(g)$, which is $G=0$.  More generally, let \\begin{equation} F_{\\Phi,U}(g) = \\int_{\\Phi(U)} \\mbox{Quad}_M(M') \\; dV_{R^4} \\end{equation} for \\emph{any} nontrivial quadratic expression, but require that $g$ is at a critical point of $F_{\\Phi,U}(g)$ for \\emph{all} coordinate charts $\\Phi$ and all open sets $U$ (whose closure is compact and in the interior of $\\Omega$).  Then the resulting Euler-Lagrange equation can not depend on any one particular choice of coordinate chart, and so by uniqueness as proved by Cartan \\cite{Car22}, Weyl \\cite{Wey22}, and Vermeil \\cite{Ver17} up to a multiplicative constant, we must have \\begin{equation} c \\cdot {H}_U(g)= \\,F_{\\Phi,U}(g) \\;+\\; \\mbox{boundary term} \\end{equation} for some nonzero $c$. Hence, again the Euler-Lagrange equation for $F_{\\Phi,U}(g)$, for variations away from the boundary, must be the same as the one for ${H}_U(g)$, which again is $G=0$.  Now suppose we let \\begin{equation} F_{\\Phi,U}(g) = \\int_{\\Phi(U)} \\mbox{Quad}_M(M' \\cup M) \\; dV_{R^4} \\end{equation} for any quadratic expression (which is nontrivial on the $M'-M'$ part), and continue to require that $g$ is at a critical point of $F_{\\Phi,U}(g)$ for all coordinate charts $\\Phi$ and all open sets $U$ (whose closure is compact and in the interior of $\\Omega$). Then the resulting Euler-Lagrange equation still can not depend on any one particular choice of coordinate chart. Furthermore, the uniqueness theorems of Cartan \\cite{Car22}, Weyl \\cite{Wey22}, and Vermeil \\cite{Ver17} imply the even more general result that \\begin{equation} c \\int_{U} (R - 2 \\Lambda) \\; dV= \\,F_{\\Phi,U}(g) \\;+\\; \\mbox{boundary term} \\end{equation} for two constants $c \\ne 0$ and $\\Lambda$.  Hence, the Euler-Lagrange equation for $F_{\\Phi,U}(g)$, for variations away from the boundary, will be the same as the expression on the left. The expression on the left is called the Einstein-Hilbert action with cosmological constant and is well known to have Euler-Lagrange equation \\begin{equation}\\label{eqn:cosmoeqn} G + \\Lambda g = 0, \\end{equation} which is the equation for vacuum general relativity with a cosmological constant $\\Lambda$.  When the cosmological constant term is moved to the other side of the equation and interpreted as matter as in $G = -\\Lambda g = 8\\pi T$, this matter is given the name ``dark energy'' (not to be confused with dark matter) and has the same stress energy tensor as a fluid with positive energy density and negative pressure (for $\\Lambda > 0$). The idea that a fluid can have negative pressure is confusing to many, but the point is that dark energy is not a fluid. Dark energy is simply the phenomenon resulting from the constant in equation \\ref{eqn:cosmoeqn} above. Standard theory today assigns a very small but positive value for $\\Lambda$ to explain the accelerating expansion of the universe. While its density is very small, it is everywhere, and so ends up accounting for roughly $73\\%$ of the mass of the universe in the standard $\\Lambda CDM$ model \\cite{WMAPobservations}.  The above discussion provides the motivation for Axiom \\ref{A1} from section \\ref{GeometricMotivation}.  It is also helpful for understanding the implications of Axiom \\ref{A1}, which we discuss in the next appendix. ",
        "conclusions": ""
    },
    "1004/1004.4150_arXiv.txt": {
        "abstract": "We compare our Monte Carlo reaction rates (see Paper II of this series) to previous results that were obtained by using the classical method of computing thermonuclear reaction rates. For each reaction, the comparison is presented using two types of graphs: the first shows the change in reaction rate uncertainties, while the second displays our new results normalized to the previously recommended reaction rate. We find that the rates have changed significantly for almost all reactions considered here. The changes are caused by (i) our new Monte Carlo method of computing reaction rates (see Paper I of this series), and (ii) newly available nuclear physics information (see Paper III of this series). ",
        "introduction": "In the final paper of this series, referred to as Paper IV, we compare our Monte Carlo reaction rates to previously derived results. In Paper I we point out that previously reported (or ``classical\") reaction rates have a different meaning from our Monte Carlo rates. The latter are presented in Paper II and are subject to a straightforward statistical interpretation, while the former are statistically meaningless. Consequently, one has to be very careful when comparing such qualitatively and quantitatively different results. Since we did not want to complicate the discussion in Paper II by comparing ``apples with oranges\", we avoided so far a comparison between present and previous reaction rates.  On the other hand, we realize that for practical purposes many readers are precisely interested in such a comparison: by how much did a particular reaction rate change? How accurate was a previous reaction rate? Does a given rate used in a particular stellar model need updating or not? How significant is this new Monte Carlo method compared to classical techniques of computing thermonuclear reaction rates? In order to answer these questions, we present a comparison of our new rates with previous results.  Section 2 provides a few useful comments. A brief summary is presented in Sec. 3. Results are presented in graphical form in the Appendix, together with information on how to interpret the figures. ",
        "conclusions": "The comparison of Monte Carlo with classical reaction rates is presented in graphical form in the Appendix. For each reaction considered, two types of graphs are presented. The first shows the change in reaction rate uncertainties, while the second graph displays the new results normalized to the previously recommended reaction rate. Detailed information on how to interpret the graphs is also given in the Appendix.  Generally, the previous results used for the comparison are adopted from the latest reference that lists recommended reaction rates and, if available, associated uncertainties in {\\it tabular form}. This reference is provided after each rate comparison figure. Most of the previous reaction rates are adopted from Angulo et al. \\cite{Ang99} or Iliadis et al. \\cite{Ili01}. For some reactions, newer rates have been presented in the literature, but in graphical form only. In such cases we requested (and obtained) from the original authors the numerical reaction rates used to produce the literature plots.  Visual inspection of the graphs reveals that for almost all reactions our new thermonuclear rates differ substantially from previous results. The changes are caused both by our new Monte Carlo method of computing reaction rates (see Paper I of this series), and by newly available nuclear physics information (see Paper III of this series). We did not analyze in detail for each reaction the sources of a rate change at a given temperature. Such a comprehensive analysis would require considerably more efforts in computing time and manpower. Nevertheless, we feel that many readers may find the graphs presented here useful for a variety of purposes."
    },
    "1004/1004.3546_arXiv.txt": {
        "abstract": "We confirm an eighth gravitational lens system in the CAmbridge Sloan Survey Of Wide ARcs in the skY (CASSOWARY) catalogue.  Exploratory observations with the X-shooter spectrograph on the Very Large Telescope (VLT) of the European Southern Observatory (ESO) show the system CSWA\\,5 to consist of at least three images of a blue star-forming galaxy at $z = 1.0686$, lensed by an apparent foreground group of red galaxies one of which is at $z = 0.3877$.  The lensed galaxy exhibits a rich spectrum with broad interstellar absorption lines and a wealth of nebular emission lines. Preliminary analysis of these features shows the galaxy to be young, with an age of $\\sim 25$--50\\,Myr. With a star-formation rate of $\\sim 20$\\,M$_\\odot$~yr$^{-1}$, the galaxy has already assembled a stellar mass $M_\\ast \\sim 3 \\times 10^9$\\,M$_\\odot$ and reached half-solar metallicity.  Its blue spectral energy distribution and Balmer line ratios suggest negligible internal dust extinction.  A more in-depth analysis of the properties of this system is currently hampered by the lack of a viable lensing model. However, it is already clear that CSWA\\,5 shares many of its physical characteristics with the general population of UV-selected galaxies at redshifts $z = 1$--3, motivating further study of both the source and the foreground mass concentration responsible for the gravitational lensing. ",
        "introduction": "\\begin{figure*} \\begin{center} \\vspace{-0.5cm} \\includegraphics[width=12.5cm]{lc_fig1.eps} \\end{center} \\vspace{-0.25cm} \\caption{\\textit{Left: } $20\\times20$\\,arcsec SDSS $g$, $r$, $i$ colour composite image of the CSWA\\,5 system.~ \\textit{Right: } $r^\\prime$-band image of CSWA\\,5 recorded with a 20\\,s long exposure of the X-shooter acquisition and guiding camera, at position angle P.A.\\,=\\,8\\,degrees.  The scale of the right hand image is extended by a factor of two relative to the left one.  In the right-hand figure, the lensed images of the source are labelled im1--im4, reflecting the sequence of light arrival times (on the assumption that im4 is one of the lensed images).  Two foreground massive red galaxies are labelled L1 and L2. Note that the blue image im4 is not clearly visible in the SDSS composite on the left.  The X-shooter slit is overlaid on the right-hand image of the field; we have also indicated the apertures used in the extraction of the one-dimensional spectra of im1, im2, and im3, as described in Section~\\ref{sect:data}.  Coordinates of all six labelled objects are given in Table~\\ref{tab:maggies}.  } \\label{fig:SDSS_im} \\end{figure*}   \\begin{table*} \\begin{minipage}{150mm} \\caption{SDSS coordinates and magnitudes for the different components of the CSWA\\,5 lens system.} \\begin{tabular}{lllccccc} \\hline Object & RA (J2000) & Dec (J2000) & $~~~u~~~$ & $g$ & $r$ & $i$ & $z$ \\\\ \\hline \\hline L1  & ~~12:44:51.35~~ & ~~01:06:43.0~~ & ~~22.34$\\pm$0.50~~ & ~~20.61$\\pm$0.04~~ & ~~19.12$\\pm$0.02~~ & ~~18.44$\\pm$0.02~~ & ~~18.02$\\pm$0.04~~ \\\\ L2  & ~~12:44:51.00~~ & ~~01:06:44.0~~ & 22.33$\\pm$0.79 & 20.43$\\pm$0.06 & 18.44$\\pm$0.02 & 17.87$\\pm$0.01 & 17.67$\\pm$0.04 \\\\ im1 & ~~12:44:51.22~~ & ~~01:06:40.1~~ & 22.07$\\pm$0.20 & 21.74$\\pm$0.06 & 22.48$\\pm$0.17 & 23.85$\\pm$0.67 & 21.75$\\pm$0.45 \\\\ im2 & ~~12:44:51.27~~ & ~~01:06:45.2~~ & 20.79$\\pm$0.08 & 21.09$\\pm$0.04 & 20.91$\\pm$0.05 & 20.67$\\pm$0.07 & 20.42$\\pm$0.19 \\\\ im3 & ~~12:44:51.23~~ & ~~01:06:44.2~~ & \\\\ im4 & ~~12:44:51.44~~ & ~~01:06:43.5~~ &  \\\\ \\hline \\end{tabular} \\label{tab:maggies}  \\textit{Notes.}~ The magnitudes quoted are SDSS model magnitudes and are considerably more uncertain than the formal errors quoted because of blending. No SDSS magnitudes are available for images im3 and im4.\\\\ ~~~~The coordinates of the objects were determined from the X-shooter acquisition image (see Figure~\\ref{fig:SDSS_im}), relative to the SDSS photometric source catalogue. Positional uncertainties are $\\sim 0.1$\\,arcsec.  \\end{minipage} \\end{table*}  Galaxies at redshifts $z \\gtrsim 1$ are generally too faint for their rest-frame ultraviolet (UV) spectra to be recorded with the high resolution and signal-to-noise (S/N) ratios required for an in-depth study of their most important physical properties.  Consequently, our knowledge of the high redshift galaxy population is still largely based on analyses of stacked spectra which average together observations of many galaxies in order to improve the S/N ratio \\citep[e.g.][]{Shapley03, Vanzella09, weiner09}.  A few exceptions are sources which, thanks to their fortuitous location behind foreground mass concentrations such as massive galaxies, or groups and clusters of galaxies, are strongly lensed so that their fluxes are significantly boosted.  Observations of what has become the archetypal strongly lensed galaxy, MS\\,1512-cB58 (cB58 for short) at $z=2.7276$ \\citep{pettini00,teplitz00,pettini02,baker04,siana08}, have provided clear demonstrations of the power of this technique in giving insights into the interstellar medium and stellar populations of distant galaxies with a level of detail which would otherwise be unobtainable until the next generation of 30+\\,m optical-infrared telescopes.  The large area of sky sampled by the Sloan Digital Sky Survey \\citep[SDSS,][]{york00} has allowed systematic searches for strongly lensed galaxies \\citep[e.g.][]{allam07,lin09}, with the dual interest of identifying and studying the most massive galaxies in the lensing objects, and distant galaxies in the lensed sources. The Sloan Lens ACS Survey, or SLACS,\\footnote{\\tt http://www.slacs.org/} targets galaxies with emission lines at more than one redshift, or redshifted nebular emission lines superimposed on the spectra of early type galaxies.  These candidate gravitationally lensed systems are then followed up with high spatial resolution images from the Advanced Camera for Surveys on the \\textit{Hubble Space Telescope} \\citep[\\textit{HST;} e.g.][]{bolton06,willis05}.  In a complementary approach, The CAmbridge Sloan Survey Of Wide ARcs in the skY (CASSOWARY) targets multiple, blue companions around massive ellipticals in the SDSS photometric catalogue as likely candidates for wide-separation gravitational lens systems \\citep[see][]{belokurov09}.  These and other search strategies have led to numerous recent discoveries of strongly lensed galaxies at $z \\gtrsim 1$ which have been followed up in detail over a range of wavelengths \\citep[see][and references therein]{lemoine03,cabanac05,belokurov07,coppin07,swinbank07,cabanac08,stark08,finkelstein09,hainline09, quider09,siana09,dessauges09,quider10}.  However, all of these studies have of necessity focused only on limited wavelength intervals, mostly in the rest-frame UV (redshifted to optical wavelengths) or the rest frame optical (redshifted into the near-infrared).  It is only with the recent advent of the X-shooter spectrograph on the Very Large Telescope (VLT) facility of the European Southern Observatory (ESO), that both wavelength ranges can now be recorded at once with good sensitivity. The obvious advantages have been demonstrated in our recent study of CASSOWARY\\,20 \\citep[CSWA\\,20,][]{pettini10}, a blue star-forming galaxy at $z_{\\rm em} = 1.433$ lensed into a wide separation ($\\sim 6$\\,arcsec) Einstein Cross by a massive luminous red galaxy at $z_{\\rm abs} = 0.741$ with velocity dispersion $\\sigma_{\\rm lens} \\simeq 500$\\,km~s$^{-1}$.  \\begin{figure*} \\begin{center} \\includegraphics[width=17cm, bb=100 375 1020 700, clip]{lc_fig2.ps} \\end{center} \\caption{SDSS spectrum of galaxy L2 at $z_{\\mathrm{abs}}=0.3877$ (shifted to the rest frame). The best fit stellar template with a velocity dispersion of $\\sigma = 278$\\,km~s$^{-1}$ is shown in red. } \\label{fig:SDSS_spec} \\end{figure*}  The current CASSOWARY catalogue\\footnote {\\tt http://www.ast.cam.ac.uk/research/cassowary/} includes seven spectroscopically confirmed lensed galaxies at $z > 1$.  In this paper we confirm an eighth case with X-shooter observations of CSWA\\,5.  As can be seen from Figure~\\ref{fig:SDSS_im}, this multiple system consists of two foreground red galaxies one of which (L2) has an SDSS spectrum which shows it to be at $z_{\\rm abs} = 0.3877$, and four fainter sources, separated by $\\sim 1$--5\\,arcsec.  Three of these images, im1, im2 and im3, are shown by the observations reported here to be gravitationally lensed images of an emission line galaxy at $z_{\\rm em} = 1.0686$, while the fourth image, not covered by our observations, remains to be confirmed spectroscopically.  Coordinates and magnitudes of the different components of CSWA\\,5 are given in Table~\\ref{tab:maggies}.  However, we caution that the SDSS pipeline magnitudes, which are based on exponential profile fits, are likely to be significantly more uncertain than formal errors quoted, because of blending between the different components of this complex system.  In particular, while image im1 may appear to be bluer than im2 from the magnitudes listed in Table~\\ref{tab:maggies}, we suspect that this is just due to blending with the foreground red galaxy L1, because our X-shooter spectra of all three images, im1, im2, and im3, exhibit the same continuum slope.   This paper is organised as follows. We describe the X-shooter observations in Section~\\ref{sect:data} and consider the lensing galaxies in Section~\\ref{sect:lens}.  The main focus of the paper is the rich absorption and emission line spectrum of the source, described in Sections~\\ref{sect:source} and \\ref{sect:phys_con} respectively.  We discuss our main results in Section~\\ref{sec:discuss} and conclude in Section~\\ref{sec:conclude}. Unless otherwise stated, we assume a flat cosmological model with $H_0=73$~km~s$^{-1}$~Mpc$^{-1}$ and $\\Omega_{\\Lambda}$~=~0.72 \\citep{spergel07}. ",
        "conclusions": "\\label{sec:conclude}  It is encouraging and satisfying that through dedicated searches of large area surveys the number of gravitationally lensed galaxies at high redshifts continues to increase.  The work presented here has not only identified a new such example in the CASSOWARY catalogue but, more importantly, has demonstrated the wealth of information on the physical conditions of these objects which can be gleaned from even relatively short exposures with the X-shooter spectrograph, thanks to its high efficiency and wide wavelength coverage.  The main impediment we have found to the full characterization of CSWA\\,5 is the lack of a lensing model which can reproduce the deflections of the images. Higher spatial resolution imaging of the field (achievable with the \\textit{Hubble Space Telescope}) and more redshift determinations of galaxies in the field are needed for a comprehensive census of the mass distribution along the line of sight and an accurate estimate of the magnification of CSWA\\,5.  Such information, combined with higher S/N ratio spectra of the four lensed images (the work presented in this paper is based on only 40 minutes of integration on target), would lead to a full mapping of the properties of this star-forming galaxy.  As a final note, we should like to stress that the use of gravitational lensing to access the properties of galaxies at high redshifts is in some way still in its infancy.  While the few cases studied so far have been very successful in providing detailed views of these systems which would otherwise have simply been unobtainable, it is also the case that nearly all of them appear to be drawn from the bright end of the galaxy luminosity function. Fainter galaxies \\textit{are} accessible with 8--10\\,m telescopes equipped with modern spectrographs, but are often overlooked in searches for strong gravitational lenses.  By systematically extending such searches to fainter limits with current and future (e.g. with the Large Synoptic Survey Telescope) surveys, it will be possible to identify intrinsically fainter galaxies, particularly at epochs when the Universe was forming most of its stars, and thereby assemble a full picture of the galaxy formation process in action.  \\begin{figure} \\begin{center} \\includegraphics[width=7.5cm, bb=69 365 410 658, clip]{lc_fig9.ps} \\end{center} \\caption{Emission line flux ratios of SDSS galaxies shown in grey scale \\citep{tremonti04}. The solid and dotted lines represent the theoretical \\citep{kewley01} and empirical \\citep{kauffmann03} demarcations between star forming galaxies and galaxies harbouring AGN.  Also overplotted are the values for CSWA\\,5 (large square), and those measured from spectra of star-forming galaxies at $z \\simeq 2$ \\citep[][circles]{erb06c}, and $z \\simeq 1$--1.4 \\citep[][diamonds]{liu08} .} \\label{fig:bpt} \\end{figure}"
    },
    "1004/1004.1069_arXiv.txt": {
        "abstract": "{} {The PASTEL catalogue is an update of the \\FeH catalogue, published in 1997 and 2001. It is a bibliographical compilation of stellar atmospheric parameters providing \\tgm determinations obtained from the analysis of high resolution, high signal-to-noise spectra, carried out with model atmospheres. PASTEL also provides determinations of the one parameter \\Teff based on various methods. It is aimed in the future to provide also homogenized atmospheric parameters and  elemental abundances, radial and rotational velocities. A web interface has been created to query the catalogue on elaborated criteria. PASTEL is also distributed through the CDS database and VizieR.} {To make it as complete as possible, the main journals have been surveyed, as well as the CDS database, to find relevant publications. The catalogue is regularly updated with new determinations found in the literature.} {As of Febuary 2010, PASTEL includes 30151 determinations of either  \\Teff or \\tgm for 16649 different stars corresponding to 865 bibliographical references. Nearly 6000 stars have a determination of the three parameters \\tgm with a high quality spectroscopic metallicity. } {} ",
        "introduction": "The knowledge of stellar atmospheric parameters \\tgm is mandatory to estimate the abundance of any chemical element in a stellar atmosphere, which then traces either the composition of the interstellar medium from which the star formed or the various physical processes which may alter the initial composition.  Atmospheric parameters are thus essential in many reasearch areas related to the physics of stars and galaxies.  The detailed analysis of high resolution, high signal-to-noise spectra is the only primary method to estimate the chemical composition of stellar photospheres, from which other indirect methods of metallicity determinations can be calibrated. The general approach of a spectroscopic detailed analysis is to derive the iron abundance by matching equivalent widths of weak lines or spectral intervals of an observed spectrum to those computed from a grid of model atmospheres of various effective temperatures, gravities and metallicities. The higher the resolution and signal-to-noise ratio, the better are supposedly the results. However, each individual study has its own observational characteristics, preferred stellar atmosphere models, atomic line data and analysis methods. The consequence is a lack of homogeneity in the results from one study to the next which makes it very cumbersome to combine the determinations. Defining a common scale for atmospheric parameters is a mandatory but difficult task, which starts with the compilation of the best studies available in the literature. This is the aim of the PASTEL catalogue, which follows the previous \\FeH catalogue.  The publication of the \\FeH catalogue started with a first version in 1980 \\citep{cay80}, continued with three other versions \\citep{cay81,cay85, cay92} superseded by the 1996 edition \\citep{cay97}. Then the content of the 2001 edition \\citep{cay01} was modified by limiting the temperature range to stars cooler than 7000K and by removing references older than 1980. It was time to update the \\FeH catalogue. Since 2001, spectroscopic observations of individual stars at high resolution and high signal-to-noise (S/N) have been intensive. Several research topics are particularly productive regarding the number of analysed stars. Studies on the chemical composition of stars with planets, on the metallicity distribution and gradients in the Galaxy and on the chemical composition of metal-poor stars have produced extensive lists of stars, based on high-quality data, with a good fraction of stars which have been observed at high spectral resolution for the first time.  The methods have evolved. Due to the growing number of spectra to be processed, it becomes more frequent that spectroscopic analyses are automatised. It can be either the measurement of equivalent widths, or the fitting of individual lines or spectral regions to computed ones. As a consequence, recent papers tend to present larger number of stars than before. A typical example of such an extensive paper is the one by \\citet{Val05} presenting a uniform catalogue of stellar properties for 1040 nearby F, G, and K stars.  Effective temperature is a critical parameter in spectroscopic analyses, because errors in \\Teff lead to significant errors in measured abundances. Several recent studies have contributed to a significant increase of the number of stars with precise determinations of \\Teff. We have included these studies in PASTEL, even if they are not based on high-resolution spectra.  In this new catalogue we come back to the content of the 1996 edition of the \\FeH catalogue including hot stars, and references older than 1980. We complete it with the 2001 edition and with new references gathered from the literature. Thus, PASTEL supersedes the two previous versions of the \\FeH catalogue. In the two latest \\FeH catalogues, field and cluster stars were presented in separate tables. Here we have them in a single catalogue.   A webserver has been created for a convenient query of the catalogue. It can be queried by a list of identifiers resolved by Simbad, or by contraints on the atmospheric parameters, or on equatorial coordinates, B and V or 2MASS J, H, K magnitudes, as well as by name of authors, year of publication or bibcode.  The numbers given in this paper are those corresponding to the version of PASTEL in Febuary 2010. Because the catalogue is updated regularly, these numbers are changing accordingly. ",
        "conclusions": "We have presented the PASTEL catalogue, which is to date a bibliographical compilation with 30151 entries for 16649 different stars. \\Teff determinations  based on various methods are available for all stars. Some 5954 different stars have determinations of the full set of atmospheric parameters \\tgm with metallicity based on high resolution, high signal-to-noise spectra. The users are encouraged  to cite the original analysis papers when using the determinations compiled in PASTEL.  PASTEL offers a useful database for mining stars with known atmospheric parameters, in particular with a high-quality spectroscopic metallicity. The users have to keep in mind that the content of the catalogue is not homogeneous. The stellar content of catalogue, mainly bright stars, is biaised towards stars which are massively studied in peculiar spectroscopic programmes, like solar type stars in planet searches for instance. The \\tgm determinations are also not homogeneous and should not be simply averaged for most applications which look into detailed chemical composition of stars.  We have compared metallicity determinations available for the same stars from recent, homogeneous and high quality studies. We have found that in general these studies do not agree well. In some cases the offset can reach 0.07 dex and the dispersion 0.1 dex, larger than the individual errors quoted in the considered studies. The lack of a common \\tgm scale results from the variety of observational characteristics, model atmospheres, line data and methods of analysis which are used in spectroscopy. A serious effort should be undertaken to build extensive and homogeneous catalogues of \\tgm covering the whole HR diagram and metallicity range."
    },
    "1004/1004.3293_arXiv.txt": {
        "abstract": "The nebula J222557+601148, tentatively identified by Morris et al.\\ (2006) as a young Galactic supernova remnant (SNR) from {\\sl Spitzer} Galactic First Look Survey images and a follow-up mid-infrared spectrum, is unlikely to be a SNR remnant based on H$\\alpha$, [\\ion{O}{3}], [\\ion{S}{2}] images and low dispersion optical spectra.  The object is seen in H$\\alpha$ and [\\ion{O}{3}] $\\lambda$5007 images as a faint, roughly circular ring nebula with dimensions matching that seen in 24 $\\mu$m {\\sl Spitzer} images.  Low-dispersion optical spectra show it to have narrow H$\\alpha$ and [\\ion{N}{2}] $\\lambda\\lambda$6548,6583 line emissions with no evidence of broad or high-velocity (v $\\geq 300$ km s$^{-1}$) line emissions.  The absence of any high-velocity optical features, the presence of relatively strong [\\ion{N}{2}] emissions, a lack of detected [\\ion{S}{2}] emission which would indicate the presence of shock-heated gas, plus no coincident X-ray or nonthermal radio emissions indicate the nebula is unlikely to be a SNR, young or old.  Instead, it is likely a faint, high-excitation planetary nebula (PN) as its elliptical morphology would suggest, lying at a distance $\\sim 2 - 3$ kpc with unusual but not extraordinary mid-IR colors and spectrum.  We have identified a $m_{r'} = 22.4 \\pm 0.2$ star as a PN central star candidate. ",
        "introduction": "Of the 274 Galactic supernova remnants (SNRs) currently known \\citep{Green09}, few are believed to be relatively young with ages of 3000 years or less.  The list of young Galactic SNRs includes the historic guest stars of SN 1604 (Kepler's SNR), SN 1572 (Tycho's SNR), SN 1054 (Crab Nebula), SN 1006 (G327.6+14.6), SN 386 (G11.2-0.3), and SN 185 (RCW 86).  There are also about a dozen remnants which do not have firmly established ages but are probably less than a few thousand years old.  These include the recently recognized very young remnant G1.9+0.3 (age $\\simeq$ 150 yr; \\citealt{Reynolds08}, \\citealt{Green08}) and Cassiopeia A (SN $\\approx$ 1680; \\citealt{Thor01}, \\citealt{Fesen06}).  With so few young SNRs known, the discovery of even one more young Galactic SNR is significant.  Young SNRs are of special interest since they offer a host of details on supernova explosions on finer spatial scales than from extragalactic SNR investigations, including expansion asymmetries, ejecta abundances, and clues regarding the nature of the progenitor star and its pre-SN environment.  Hence the discovery of a small symmetric nebula, SSTGFLS J222557+601148 (hereafter J222557), detected in 24 $\\mu$m {\\sl Spitzer} Galactic First Look Survey (GFLS) images and tentatively identified as a possible SNR with an age $\\sim 1000$ yr by \\citet{Morris06} is worth investigating. Here we present optical images and spectra which indicate this nebula is likely not a supernova remnant but instead a faint, high-excitation planetary nebula (PN).  \\begin{figure*} \\plotone{t1.eps} \\caption{H$\\alpha$ image of the J222557+601148 region showing the detection of a faint emission shell. North is up and East to the left.} \\label{fig:figure0} \\end{figure*}   \\begin{figure*} \\plotone{t2.eps} \\caption{Enlarged optical and infrared images of J222557+601148. North is up and East to the left. {\\it{Top row:}} MDM 1.3 m images showing faint H$\\alpha$ (left) and [\\ion{O}{3}] (right) emission shells. {\\it{Middle row:}} H$\\alpha$ minus 6510 \\AA \\ continuum (left), and 24 $\\mu$m {\\sl Spitzer} MIPS (right). The position where we obtained the optical spectrum using a N--S slit placed across the nebula is indicated in the H$\\alpha$ minus 6510 \\AA \\ continuum image. {\\it{Bottom row:}} Broadband Harris B and SDSS r filter images with the central star PN (CSPN) candidate indicated by the arrows. This star's position relative to the emission shell is shown above in the H$\\alpha$ minus continuum image by a `$+$' symbol. } \\label{fig:images} \\end{figure*} ",
        "conclusions": "Morris et al. (2006) tentatively identified the Galactic nebula J222557+601148 as a young supernova remnant (SNR) based {\\sl Spitzer} Galactic First Look Survey images. However, H$\\alpha$, [\\ion{O}{3}], [\\ion{S}{2}] images and low dispersion optical spectra reveal it to only exhibit narrow H$\\alpha$ and [\\ion{N}{2}] $\\lambda\\lambda$6548,6583 line emissions with no evidence of broad or high-velocity line emissions.  The absence of any high-velocity optical or infrared features, the presence of relatively strong [\\ion{N}{2}] emissions, a lack of detected [\\ion{S}{2}] emission which would indicate the presence of shock-heated gas, plus no coincident X-ray or nonthermal radio emissions suggest the nebula is unlikely to be a SNR, young or old.  Instead, it is likely a faint, high-excitation PN as its elliptical morphology would suggest, lying at a distance $\\sim 2.5$ kpc with strong [\\ion{O}{4}] emission rather than dust continuum emission dominating it detection in 24 $\\mu$m {\\sl Spitzer} images.  We have identified a possible central star candidate with $m_{r'} \\simeq 22.4$."
    },
    "1004/1004.1775_arXiv.txt": {
        "abstract": "{} {Report Low-resolution optical spectroscopic observations for 78 very low-mass star and brown dwarf candidates that have been photometrically selected using the DENIS survey point source catalogue.} {Spectral types are derived for them using measurements of the PC3 index. They range from M6 to L4. H$\\alpha$ emission and NaI subordinate doublet (818.3 nm and 819.9 nm) equivalent widths are measured in the spectra. Spectroscopic indices of TiO, VO, CrH and FeH molecular features are also reported. } {A rule-of-thumb criterion to select young very low-mass objects using the NaI doublet equivalent width is given. It is used to confirm seven new members of the Upper Sco OB association and two new members of the R Cr-A star-forming region. Four of our field objects are also classified as very young, but are not members of any known nearby young association. The frequency of lower-gravity young objects in our field ultracool sample is 8.5\\% . Our results provide the first spectroscopic classification for 38 ultracool dwarfs in the solar vicinity with spectrophotometric distances in the range 17 pc to 65 pc (3 of them are new L dwarfs within 20 pc).} {} ",
        "introduction": "A complete census of the solar neighborhood is needed for many different purposes. To mention just a few examples: the understanding of the star formation history of the Milky Way; the identification of bright benchmark objects of different spectral types and evolutionary history, and the search for extrasolar planets. Henry et al. (2002) estimated that more than half of the stellar systems within 25 pc are still unknown, and most of them are thought to be ultracool dwarfs (UDs; defined as those with spectral types M6--M9, L and T; Mart\\'\\i n et al. 1996, 1997, 1999; Kirkpatrick et al. 1999, 2000), which include a mixture of very low-mass stars  and substellar-mass brown dwarfs.  Indeed, many UDs within 25 pc have been found in the last few years (Phan-Bao et al. 2008; Reid et al. 2008).  Large optical/infrared surveys have made a significant impact in the identification of our coolest neighbours. The Deep Near Infrared Survey of the Southern Sky (DENIS; Epchtein et al. 1997) has enabled the discovery of many UDs in the solar vicinity (Delfosse et al. 1997, 2001; Crifo et al. 2005; Phan-Bao et al. 2001, 2003, 2008; Kendall et al. 2004) and prompted the development of a new spectral class, the L dwarfs (Mart\\'\\i n et al. 1997, 1999; Delfosse et al. 1999), which is characterized by the condensation of dust grains in the atmospheres (Allard et al. 2001; Marley et al. 2002; Tsuji 2005).  The Sloan Digital Sky Survey (SDSS; York et al. 2000) and the Two Micron All Sky Survey (2MASS; Skrutskie et al. 1997) have also brought about the discoveries of many UDs  (Kirkpatrick et al. 1999, 2000; Fan et al. 2000; Knapp et al. 2004, Cruz et al. 2007, Reid et al. 2008) and have provided the identification of the coolest spectral class, the T dwarfs, characterized by the presence of methane bands in the near-infrared spectrum (Burgasser et al. 2006; Kirkpatrick 2005). The new generation large infrared surveys CFHTLS and UKIDSS are starting to identify even cooler dwarfs (Warren et al. 2007), the so-called Y class, for  which the distinguishing characteristic may be the emergence of ammonia molecular bands in the near infrared. Two objects with possible ammonia absorption have recently been identified (Delorme et al. 2008), but their classification as Y dwarfs remains controversial because they are very similar to late T dwarfs (Burningham et al. 2008).   The characterization of the nearby ultracool population continues to be a goal of recent papers. Besides those mentioned above, it is worthwhile to mention a few more. Kendall et al. (2007) presented twenty-one southern ultracool dwarfs (M7--L5.5) selected from 2MASS and SuperCOSMOS point source databases according to their colors and proper motions, and confirmed via low-resolution near-infrared spectrosocopy. Costa et al. (2006) and Henry et al. (2006) reported trigonometric parallaxes for several UDs, including the closest known L dwarf. Jameson et al. (2008) have provided proper motion measurements for over a hundred L and T dwarfs.  In this paper we present low-resolution optical spectra of 78 UD candidates selected from the DENIS point source catalog using photometric color criteria. For a subsample of them (50 objects with galactic latitude between 30 and 15 degrees), the Maximum Reduced Proper Motion (MRPM) method (Phan-Bao et al. 2003, 2008) is used to reject giants. We also present spectra of a few DENIS UD candidates in the general area of the Upper Sco OB associations. ",
        "conclusions": ""
    },
    "1004/1004.2360_arXiv.txt": {
        "abstract": "Under the assumption that cold dark matter and dark energy interact with each other through a small coupling term, $Q$, we constrain the parameter space of the equation of state $w$ of those dark energy fields whose variation of the field since last scattering do  not exceed Planck's mass. We use three parameterizations of $w$ and two different expressions for $Q$. Our work extends previous ones. ",
        "introduction": "The equation of state parameter of dark energy -the mysterious agent responsible for the current accelerated expansion of the Universe- ranks among the biggest unknowns in cosmology. It is usually written as $w= p_{\\phi}/\\rho_{\\phi}$ and lies in the range $(-1, -1/3)$ if the dark energy is quintessence, and $(- \\infty, -1)$ if it is of phantom type; in either case it may run with expansion. By contrast, in the very particular (but observationally favored) instance of a cosmological constant $\\, w$ stays fixed at $\\, -1$. It is then apparent that some knowledge on the nature of dark energy can be gained by setting limits on $w$.  Observational constraints on this quantity from supernovae type Ia (SN Ia), temperature anisotropies of the cosmic microwave background radiation (CMB), baryon acoustic oscillations (BAO), etc, suggest that nowadays $w$ is not far from the cosmological constant value. However, all these studies rest on some or other world model and on inevitable priors imposed upon the parameters space to ease the data analysis.  Recently, bounds on $w$ have been derived  from the very reasonable requirement that the variation experienced by the dark energy scalar field $\\phi$ (regardless it may be phantom or quintessence) from any redshift, $z$, within the classical expansion era, till now should not exceed Planck's mass (see Huang \\cite{huang}, and Saridakis \\cite{saridakis}). This is to say \\begin{equation} |\\Delta \\phi(z)|/M_{P} < 1 \\, , \\label{condition} \\end{equation} \\noindent where -as can be readily deduced- \\begin{equation} \\frac{|\\Delta \\phi (z)|}{M_{P}} = \\frac{1}{M_{P}} \\, \\int_{t}^{t_{0}} {\\dot{\\phi} \\, dt} = \\int_{0}^{z}{\\frac{\\sqrt{3 \\, |1+w(x)| \\, \\Omega_{\\phi}(x)}}{1+x}}\\, dx \\, . \\label{Deltaphi} \\end{equation}  Huang's analysis confines itself to quintessence fields with various $w$ parameterizations (viz: $w$ constant, a linear function of redshift, and the Chevallier-Polarski-Linder parametrization \\cite{ch-p,linder}) \\cite{huang}. Saridakis considers phantom fields only and uses identical parameterizations than Huang plus a further one in which $w$ depends linearly  on the logarithm of the scale factor \\cite{saridakis}.   The bound (\\ref{condition}) looks a rather natural condition and persuasive arguments  have been advanced in its favor \\cite{banks,huang1,yun-gui}. Moreover, we have found a further, solid, motivation for it. Namely: if condition (\\ref{condition}) is violated, then dark energy dominates at very early times which cannot be reconciled with observation (see Fig. \\ref{fig4} and ``forbidden\" regions on the top right corner of the panels of figures \\ref{fig6}, \\ref{fig9}, and \\ref{fig11}, below). In particular, as noted by Bean {\\it et al.} \\cite{rachelbean}, the fractional density of dark energy cannot exceed $5 \\%$ and $39 \\%$ at the primeval nucleosynthesis and recombination epochs, respectively, -we shall elaborate on this in a future publication. Nevertheless, we refrain ourselves from imposing (\\ref{condition}) on every dark energy field. Its seems safer to demand that $|\\Delta \\phi|/M_{P}$  be not much larger than order unity. However, the latter requirement cannot be easily implemented and, at any rate, it would translate on rather loose constraints on $w$. Therefore, we choose to circumscribe our analysis to dark energy fields that fulfill condition (\\ref{condition}).  In Huang's as well as in Saridakis' work matter and dark energy interact with each other only gravitationally. Certainly this is a reasonable assumption but nonetheless minimalist since, while dark energy interactions with baryons are severely restricted by solar system experiments \\cite{solar}, there is nothing against a possible coupling (interaction) with dark matter. On the contrary, a transfer from dark energy to dark matter alleviates the cosmic coincidence problem (the observational fact that both energy densities are comparable today \\cite{paul}) and may  solve it in full \\cite{rapidc}. Besides, as revealed by optical, x-ray and weak-lensing data, the internal dynamics of 33 relaxed galaxy clusters seems to favor such interaction \\cite{binw}. Likewise, a recent analysis using the 397 SN Ia of the ``Constitution\" set \\cite{constitution}, BAO and CMB data adds significant weight to the existence of the coupling \\cite{shafieloo}. In fact, the subject has evolved into a field of active research -see e.g. \\cite{active,gabriela1} and references therein; a recent review can be found in \\cite{review}.  In the presence of a coupling, say $Q$, the continuity equations for the three main contributors to the present cosmic budget adopt the form, \\begin{eqnarray} \\dot{\\rho}_{b}\\, &+&\\, 3H \\rho_{b} = 0 \\, , \\nonumber \\\\ \\dot{\\rho}_{c}\\, &+& \\, 3 H \\rho_{c} = Q \\, , \\nonumber \\\\ \\dot{\\rho}_{\\phi}\\, &+& \\, 3 H (1+w) \\rho_{\\phi} = - Q \\, , \\label{continuity} \\end{eqnarray} where the subscripts, $b$, $c$, and $\\phi$ stands for baryons, cold dark matter, and dark energy, respectively. This paper aims to constrain  different parameterizations of $w$ under diverse expressions for $Q$. Regrettably, guidelines about the latter are rather loose: it must be positive-semidefinite and small. If $Q$ were negative, the energy densities could become negative, and the second law of thermodynamics get violated \\cite{db}. If it were large, dark energy would have dominated the expansion at early times and possibly not today. Thus, for the time being, we must content ourselves with guessing plausible expressions for $Q$ just on phenomenological bases.  In view of the above equations the coupling must be a function of $H \\rho_{c}$ and $H \\rho_{\\phi}$. By power law expanding it up to first order on these quantities we write $Q(H \\rho_{c}, H \\rho_{\\phi}) \\simeq \\epsilon_{c}\\, H \\rho_{c} \\, + \\, \\epsilon_{\\phi}\\, H \\rho_{\\phi}$, where both $\\epsilon$ coefficients ought to be non-negative and small (i.e., not larger than, say, $10^{-1}$). To simplify the analysis we will consider in turn that one of the two coefficients vanishes. So, we will take \\be Q = 3 \\, \\epsilon H \\rho_{\\phi}\\, , \\qquad {\\rm and} \\qquad Q = 3 \\, \\epsilon H \\rho_{c} \\, , \\label{int} \\ee -the factor $3$ being introduced for mathematical convenience.  At this point one may object that if $Q$ is small, no significant difference with the results of Huang \\cite{huang} and Saridakis \\cite{saridakis} should  be expected. However, as we shall see, this is not the case. {\\em A priori} we can argue in favor of some non small departure from the findings of \\cite{huang} and \\cite{saridakis} by considering the ratio between $Q$ and the second term in last equation of (\\ref{continuity}). Indeed, dividing the right hand side of (\\ref{int}.1) by the second term of the said equation yields $\\epsilon/(1+w)$. In general, this ratio cannot be neglected and it can be large in absolute value when $w$ is close to $-1$. Likewise, using  instead (\\ref{int}.2) we have $\\epsilon \\, \\rho_{c}/(1+w)\\rho_{\\phi}$. Its absolute value can be of order unity, or even larger, for extended periods of the cosmic history.  Although  expressions (\\ref{int}) were proposed on phenomenological grounds they can be obtained in scalar-tensor gravity from the action (in the Einstein frame) \\cite{kaloper} \\begin{equation} S = \\int{d^{4}x\\, \\sqrt{-g}}\\, \\left\\{ \\frac{1}{16 \\pi G} R - \\frac{1}{2} \\,\\partial_{a}\\phi \\,\\partial^{a} \\phi \\, + \\, \\frac{1}{\\chi^{2}(\\phi)}\\, L_{c}(\\zeta, \\partial \\zeta, \\chi^{-1} g_{ab})\\, + \\, L_{b}(\\xi, \\partial \\xi, \\chi^{-1} g_{ab}) \\right\\} \\, . \\label{action} \\end{equation} Here $R$ denotes the Ricci scalar, $L_{c}$ and $\\zeta$ stand for the dark matter Lagrangian and the dark matter degrees of freedom, respectively (corresponding meanings have $L_{b}$ and $\\xi$); on the other hand, $\\chi(\\phi)$ couples dark matter with the dark energy field $\\phi$. From (\\ref{action}) the interaction term between dark matter and dark energy can be expressed as \\begin{equation} Q = H \\rho_{c} \\left[\\frac{d(\\ln \\bar{\\chi}(a))}{d \\ln a}\\right] \\, , \\label{Q-action} \\end{equation} where $\\bar{\\chi}(a) = \\chi(a)^{(3w_{c} - 1)/2}$ with $w_{c}$ the equation of state parameter of dark matter (zero in our case). By choosing \\begin{equation} \\bar{\\chi}(a) = \\bar{\\chi}_{0} \\, \\exp \\left[ 3 \\, \\int \\frac{\\epsilon_{\\zeta}\\,  \\rho_{\\zeta}\\, + \\,\\epsilon_{c} \\, \\rho_{c}} {\\rho_{c}} \\; d \\ln a \\right] \\, , \\label{bychoosing} \\end{equation} the first (second) expressions for $Q$ in (\\ref{int}) follows after using (\\ref{Q-action}) and setting $\\epsilon_{c} = 0$ ($\\epsilon_{\\zeta} = 0$). Similar expressions for $Q$ have been obtained from the above action in Refs. \\cite{curbelo} and \\cite{h-zhang}.  In this paper we shall assume a spatially flat Friedmann-Robertson-Walker universe with present fractional densities of baryons and cold dark mater $\\Omega_{b0} = 0.04$, and $\\Omega_{c0} = 0.24$, respectively, throughout. As usual, a zero subindex means the current value of the corresponding quantity; likewise we normalize the scale factor of the metric by setting $a_{0} = 1$. ",
        "conclusions": "Motivated by the reasonable assumption that the variation $\\Delta \\phi$ of the field driving the present phase of cosmic accelerated expansion should not exceed Planck's mass we have numerically calculated the said evolution since last scattering till now for different expressions of the equation of state parameter $w(a)$ and two couplings, $Q$, between dark matter and dark energy. This constrains the parameter space $(w_{1}, w_{0})$ as shown in the six cases studied. Figures \\ref{fig1}, \\ref{fig3}, \\ref{fig6}, \\ref{fig7}, \\ref{fig9}, and \\ref{fig11} explicitly exhibit the regions of the parameter space $(w_{1}, w_{0})$ that fulfill the bound $|\\Delta \\phi|/M_{P} < 1$. Our analysis extends those of Huang \\cite{huang} and Saridakis \\cite{saridakis} who assumed that dark matter and dark energy evolved separately and did not consider Barboza and Alcaniz's parametrization \\cite{b-a}.  The main results of this work can be summarized as follows:\\\\ \\noindent $(i)$ When the interaction obeys (\\ref{int}.1), models that always stay phantom respect condition (\\ref{condition}) at any redshift and for any $\\epsilon$ value. $(ii)$ Likewise, depending on the values taken by $w_{0}$ and $w_{1}$ models that always stay as quintessence as well as models that evolve from quintessence to phantom and models that evolve in the opposite sense can violate the said condition. $(iii)$ When $Q$ obeys (\\ref{int}.2) and $w(a)$ is given by the Chevallier-Polarski-Linder parametrization (\\ref{chevallier}), for small $\\epsilon$, all models that ever stay phantom and all models in the mixed regions I and II respect bound (\\ref{condition}); -top panels of Figs. \\ref{fig9}. However, for not so small $\\epsilon$ values, models in these regions may also violate (\\ref{condition}) -bottom panels of Fig. \\ref{fig9}. $(iv)$ When $Q$ is given by (\\ref{int}.2) and $w(a)$ obeys the Barboza-Alcaniz parametrization (\\ref{alcaniz}), for small $\\epsilon$, all models that ever stay phantom and models in the mixed region I satisfy bound (\\ref{condition}) -top panels of Fig. \\ref{fig11}. By contrast, for not so small $\\epsilon$ values, models in these regions may also violate (\\ref{condition}) -bottom panels of Fig. \\ref{fig11}. $(v)$ The dark energy field $\\phi$ appears to evolve faster when: $(a)$ the coupling $Q$ depends explicitly on $\\rho_{c}$ than when it does on $\\rho_{\\phi}$, and $(b)$ when $w(a)$ obeys the parametrization of Barboza-Alcaniz than when it does that of Chevallier-Polarski-Linder. Point (a) is in line with Caldera-Cabral {\\it et al.}  result that interacting models with $Q = 3 \\, \\epsilon \\,  H \\rho_{\\phi}$ ``work better\" than models with $Q = 3 \\, \\epsilon \\, H \\rho_{c}$ \\cite{gabriela1}. Further, these models may help explain the non-vanishing temperature (about $0.6$ Kelvin \\cite{sterile}) of sterile neutrinos \\cite{zhou}.  Our finding that $|\\Delta \\phi|$ increases with $\\epsilon$ strengthens the view expressed in the Introduction  that $Q$ must be small. Otherwise (except for (\\ref{int}.1) and $w = {\\rm constant} < -1$, right panel of Fig. \\ref{fig1}) condition (\\ref{condition}) would be violated for whatever $w$ not far from $-1$.  We have confined ourselves to models in which the interaction term is proportional to the Hubble factor (Eqs. (\\ref{int})) whence our results apply to them only. Other interaction terms can be found in the literature, among others, $Q \\propto \\dot{\\phi} \\, \\rho_{c}$ \\cite{among-others}, and $Q \\propto (\\rho_{c} \\, + \\, \\rho_{\\phi})$ -see e.g. \\cite{gabriela1, gabriela2}. However, the former class of models present the drawback of being unable to simultaneous lead to a correct sequence of cosmic eras (radiation, matter, and dark energy) and solve the coincidence problem \\cite{challenges}, while the models we have considered do not suffer from that \\cite{drawback-free}. As for the latter class of models, they generally lead to negative energy densities, either of dark matter or dark energy, at early or late times \\cite{gabriela2}. Still, for some specific values of the parameters entering the interaction these models are free of that problem.  As is well known, due to quantum instabilities phantom fields may find no place in Nature -see e.g. \\cite{cline}. However, some phantom models based on low-energy effective string theory may not suffer from such pathologies \\cite{federico} whence this issue is not settled as yet. Moreover, higher derivative terms in the Lagrangian may render phantom models stable \\cite{creminelli}.  Before closing, it is sobering to recall that variable dark energy fields are afflicted by the problem of their small mass ($m_{\\phi} \\sim 10^{-33}$ eV). Due to this unpleasant feature they are looked upon as not more natural than the cosmological constant, despite the enormous fine tuning scale of the latter. Admittedly, this is an unsolved problem that affects most (if not all) candidates of evolving dark energy. However, recently, a supergravity based approach that might point to the solution has been proposed \\cite{brax}. It provides a very small mass for the field and a seemingly natural mechanism for a weak coupling between the field and matter."
    },
    "1004/1004.2967_arXiv.txt": {
        "abstract": "We present a core-collapse supernova model for the extremely luminous Type Ic supernova 2007bi.  By performing numerical calculations of hydrodynamics, nucleosynthesis, and radiation transport, we find that SN 2007bi is consistent with the core-collapse supernova explosion of a 43 \\Msun~ carbon and oxygen core obtained from the evolution of a progenitor star with a main sequence mass of 100 \\Msun~and metallicity of $Z=Z_\\odot/200$, from which its hydrogen and helium envelopes are artificially stripped. The ejecta mass and the ejecta kinetic energy of the models are 40~\\Msun~and $3.6\\times 10^{52}$ erg.  The ejected \\Ni~mass is as large as ~6.1~\\Msun, which results from the explosive nucleosynthesis with large explosion energy. We also confirm that SN 2007bi is consistent with a pair-instability supernova model as has recently been claimed.  We show that the earlier light curve data can discriminate between the models for such luminous supernovae. ",
        "introduction": "A massive star with the main sequence mass (\\Mms) in the range of 10 $-$ 140 \\Msun~ forms an Fe core in its center and eventually collapses. This collapse is thought to end up with the core-collapse supernova (SN) of Type II, Ib, or Ic (Filippenko 1997 for a review). If a star is as massive as \\Mms $= 140 - 300$ \\Msun, the oxygen-rich core becomes dynamically unstable owing to the electron-positron pair creations (Rakavy \\& Shaviv 1967; Barkat, Rakavy, \\& Sack 1967). As the internal energy is spent by the pair creations, the core loses the stability and starts to collapse. When the central temperature exceeds $\\sim 5\\times 10^{9}$ K, the core becomes stable but the temperature is so high that oxygen burning becomes explosive to produce enough energy to unbind the star. A large amount of \\Ni~is synthesized by the explosive burning ({\\it e.g.}, Umeda \\& Nomoto 2002, hereafter UN02; Heger \\& Woosley 2002) and the subsequent radioactive decays power the light curve. Thus this event is theoretically predicted to be observed as a pair-instability supernova (PISN). Some luminous SNe like SN 2006gy (Ofek et al. 2007; Smith et al. 2007) have been suggested to be PISNe (see Section 4.2), but no clear consensus has been reached ({\\it e.g.}, Kawabata et al. 2009).  Recently, Gal-Yam et al. (2009) (G09 hereafter) suggested that the extremely luminous Type Ic SN 2007bi is the first observed example of the PISN. They showed that the PISN model is consistent with the observed light curve (LC) and the nebular spectra of SN 2007bi. They estimated the masses of C, O, Na, Mg, Ca, and \\Ni~ from the observed optical spectra. Other elements with no strong emission lines in the optical range, Si and S, are assumed to be the same as the PISN model of Heger \\& Woosley (2002). Young et al. (2010) (Y10 hereafter) showed multi-color observations of SN 2007bi and the metallicity of the host galaxy. However, the above observations of SN 2007bi have not quantitatively been compared with the core-collapse SN models. In view of the importance of clarifying the final fates of very massive stars, we examine how strong the observational constraint on the theoretical models are.  The aim of this Letter is to show that a core-collapse SN model is indeed consistent with the observed properties of SN 2007bi, if the progenitor is as massive as \\Mms $\\sim 100$ \\Msun~ and the explosion energy is large. This would imply that SN 2007bi might not necessarily be a PISN. In Section 2, we summarize the progenitor models for SN 2007bi and numerical methods used for our calculations of hydrodynamics, nucleosynthesis, and the LC.  The core-collapse SN models of SN 2007bi are presented in Section 3 and the results are discussed in Section 4. ",
        "conclusions": "In this Letter, we have shown that the LC and photospheric velocity of SN 2007bi are well-reproduced by the core-collapse SN model CC100. As some gamma-ray bursts are connected to such high energy Type Ic SNe, the extremely luminous SNe like SN 2007bi could also be connected to gamma-ray bursts which result from very massive stars. If this is the case, extremely luminous SNe like SN 2007bi could be connected to gamma-ray bursts of much more massive star origin than known SNe associated with a gamma-ray burst. Even stars more massive than 300 \\Msun~could be the origin of luminous SNe ({\\it e.g.}, Ohkubo et al. 2006, 2009).  We note, however, that, although SN 2007bi may not necessarily be a PISN, the observational data available for SN 2007bi is not sufficient to single out the explosion mechanism.  In fact, Kasen \\& Bildsten (2009) suggested that the magnetar-powered light curve model (also, Maeda et al. 2007; Woosley 2009) might explain the LC of SN 2007bi.  Here we show the comparison between our PISN model and SN 2007bi, and discuss how to distinguish the models for luminous SNe.  We also apply such LC comparison to SN 2006gy.  \\begin{figure}[t] \\begin{center} \\plotone{fig3.eps} \\caption{ A PISN model for SN 2006gy (PISN270). The bolometric LC of SN 2006gy is taken from Y10 and Agnoletto et al. (2009). The first point of SN 2006gy is the $R$ band upper limit from Smith et al. (2007). The CC100 no-mixing model and observed bolometric LC of SN 2007bi in Figure \\ref{07bi-ccLC} are also shown. The horizontal axis shows the days in the rest frame. }\\label{fig3} \\end{center} \\end{figure}   \\subsection{Pair-Instability Supernova Models for SN 2007bi}  In Section 3 (Figure \\ref{07bi-ccLC}), we have shown that observations of SN 2007bi are well reproduced by the core-collapse SN model (CC100).  Here we confirm the claim made by G09 that a PISN model can also be consistent with the bolometric LC of SN 2007bi by using the approximate PISN model PISN270.  The PISN270 model is constructed by scaling the physical structure of the homologously expanding model CC100 to the ejecta model with \\Mej~=~121 \\Msun~and \\Ekin~=~$7\\times 10^{52}$ erg.  The ejecta mass \\Mej~is the same as the C+O core mass of the PISN model with \\Mms~=~270 \\Msun~(UN02), and \\Ekin~is obtained from the nuclear energy released by explosive nuclear burning of the C+O core (UN02).  Here the same amount of \\Ni~(\\Mni~=~9.8~\\Msun) as in the 270 \\Msun~ model (UN02) is assumed to be synthesized in the inner layers.  Note that the 270 \\Msun~ model of UN02 still has the H-rich and He envelopes at the time of explosion and, here, we assume that the envelopes were stripped off by some mechanism.  Figure \\ref{fig3} shows that the bolometric LC of PISN270 (the red line) is consistent with the bolometric LC of SN 2007bi (red open circles). The rise time to the LC peak for PISN270 is $\\sim 150$ days, being consistent with the PISN model in G09.  This rise time is longer than the core-collapse SN model CC100 (Figure \\ref{fig3}), because the photon diffusion takes more time in more massive PISN270. Although \\Mni~of PISN270 is $\\sim$ 1.6 times larger than that of the core-collapse SN CC100 model, the longer rise time lowers the peak brightness powered by the radioactive decay.  These two effects make the peak magnitude of PISN270 similar to CC100.  This difference in the rising part of the LC is important to discriminate between the core-collapse SN and the PISN models. Although SN 2007bi was not observed early enough, much earlier observations before the peak could constrain the SN type from the LC. In addition, as already mentioned in Section 3, the abundance of Si and S would also be a key to distinguish between the two models.   \\subsection {Models for SN 2006gy}  As mentioned in Section 1, there has been some suggestions that the luminous Type IIn SN 2006gy is a PISN ({\\it e.g.}, Smith et al. 2007).  We thus apply our LC models for comparison with SN 2006gy.  Figure \\ref{fig3} shows the bolometric LC of SN 2006gy (filled circles and triangles).  As only the $R$ band magnitude was observed in the early epochs of SN 2006gy (the filled square), we cannot construct the bolometric LC at the early epochs but we can constrain the rise time of the LC.  Our calculations show that the rise time of the PISN model is too slow to be consistent with SN 2006gy.  Although our PISN model does not have a H-rich envelope, the presence of the H-rich envelope could even slow down the brightening ({\\it e.g.}, Kawabata et al. 2009). Woosley et al. (2008) showed that the interaction between the pulsating core and the envelope can power the LC of SN 2006gy.  As a similar mechanism, the interaction of a SN ejecta with its very dense circumstellar matter could convert the kinetic energy of ejecta directly to radiation energy and could also be the origin a luminous SN like SN 2006gy."
    },
    "1004/1004.0223_arXiv.txt": {
        "abstract": "We present a close companion search around sixteen known early-L dwarfs using aperture masking interferometry with Palomar laser guide star adaptive optics.  The use of aperture masking allows the detection of close binaries, corresponding to projected physical separations of 0.6-10.0 AU for the targets of our survey.  This survey achieved median contrast limits of $\\Delta$K $\\sim$ 2.3 for separations between 1.2 - 4 $\\lambda$/D, and $\\Delta$K $\\sim$ 1.4 at $\\frac{2}{3}\\lambda /D$.  We present four candidate binaries detected with moderate to high confidence (90-98\\%).  Two have projected physical separations less than 1.5 AU.  This may indicate that tight-separation binaries contribute more significantly to the binary fraction than currently assumed, consistent with spectroscopic and photometric overluminosity studies.  Ten targets of this survey have previously been observed with the Hubble Space Telescope as part of companion searches.  We use the increased resolution of aperture masking to search for close or dim companions that would be obscured by full aperture imaging, finding two candidate binaries.  This survey is the first application of aperture masking with laser guide star adaptive optics at Palomar.  Several new techniques for the analysis of aperture masking data in the low signal to noise regime are explored. ",
        "introduction": "The mass determinations of stars through binary studies have provided numerous mass-luminosity benchmarks for the testing and calibration of stellar models.  Such studies have only recently begun for the regime of brown dwarfs.  The empirical calibration of brown dwarf models is generally made more difficult by the added dependency on age in the mass-luminosity relationship.  For example, an object spectroscopically classified as a late-M dwarf may be a young brown dwarf or an old star just above the hydrogen burning limit.  This broadens the range of potential physical properties (mass, age, composition) that generate the same observable spectrum.  Conversely, photometry can only very generally reveal the objects' physical properties.  Measurements of brown dwarf masses through the tracking of binary orbits provide the strongest constraints on stellar models, \"mass benchmarks\" that reduce the degeneracy of photometric studies even for targets with unknown ages \\citep{Liu:2008}.    Mass measurements of brown dwarfs by \\citet{Konopacky:2010} show systematic discrepancies between models and measurements; late-M through mid-L systems tended to be more massive than models predict, while one T dwarf system was less massive than its model prediction.  This collection of mass benchmarks grows even more important as brown dwarf models are extended to infer the masses of directly imaged planets, such system HR 8799 \\citep{Marois:2008}.  Binary surveys have also begun to turn up interesting statistical results that may suggest the brown dwarf binary formation mechanism is different than that for solar type binaries (see \\citet{Burgasser:2006} for a summary of results from low mass surveys, including many results presented in this section).  Few surveys, however, have produced results for very low mass binaries, especially those with very tight separations ($\\lesssim$ 3.0 AU).  This regime of short period binaries is particularly challenging to achieve with ground-based direct imaging.  While the companion fraction of brown dwarfs is proposed to be low ($\\approx$ 15\\%) and peaked within a narrow separation range, 3-10 AU \\citep{Burgasser:2008}, little conclusive results are known for separations less than 3 AU despite preliminary evidence that many additional companions are likely to reside at very close distances (\\citet{Jeffries:2005,Pinfield:2003,Chappelle:2005}).  Over 90\\% of known very low mass binaries have less than 20 AU \\citep{Burgasser:2006}.  Competing theories of stellar formation aim to explain the observed companion statistics of brown dwarfs.  As a general trend, stars appear to have a binary fraction that decreases with mass.  This can partly be explained by the decreased binding energy of lower mass primaries, and thus a maximum binary separation that decreases as a function of total mass (\\citet{Reid:2000,Close:2003}).  For very low mass stars, the companion fraction peaks near 3-10 AU, and exhibits a significant (and statistically significant) drop at separations beyond 20 AU separation.  Slightly more than half of known very low mass binaries lay within this narrow separation range \\citep{Burgasser:2006}.  On the near side of this peak, the data collected is very likely incomplete, where the necessary resolution (300 mas) stretches the limitations of HST/NICMOS and ground-telescopes with AO alone.  What data has been collected suggests direct imaging may have missed companions at very close separations.  Spectroscopic, spectral morphology, and Laser Guide Star AO surveys suggest that very tight binaries within 3 AU may be as plentiful as binaries of moderate separation.  \\citet{Burgasser:2007} has used spectral features of unresolved sources to indicate composite spectra, implying multiplicity.  This technique has suggested numerous early/mid-L dwarfs with potential mid-T dwarf secondaries and systems of equal-mass L/T transition objects.  \\citet{Jeffries:2005} used sparse radial velocity data-sets of very low mass systems to predict an additional 17-30\\% binaries at separations less than 2.6 AU.  Photometric overluminosity studies by \\citet{Pinfield:2003} and \\citet{Chappelle:2005} have also hinted at surprisingly larger binary fractions (up to 50\\%) in the Pleiades and Praesepe, though concerns over membership contamination and the influence of variability limit the conclusiveness of the results.  In each study, with the exception of the Burgasser mid-L/mid-T systems, very low mass binaries tend towards equal mass pairs ($q\\sim$ 1) at close separations, just as is the case at moderate separation.  These preliminary results contrast those of previous, observationally complete surveys that focused on moderate and wide separation binaries. Those surveys predict that fewer than 3\\% of very low mass companions sit at separations closer than 3 AU \\citep{Allen:2007}.  This discrepancy speaks to the importance of additional, observationally complete surveys searching for binaries at close separations.  Non-Redundant Aperture Masking (NRM) on 5-10m class telescopes, combined with LGS AO, allows sub-diffraction limit resolution observations at contrasts high enough to search for most binaries in this potentially fertile, unresolved region.  The detection of close brown dwarf binaries, with a typical period of 1-2 years, also allows mass measurements of late-L or T dwarfs, providing particularly valuable empirical benchmarks for the study of low mass stellar models.  To put into perspective the dearth of benchmarks, the mass measurements of fifteen very low mass systems (including six with L or T dwarf components) using LGS AO alone by \\citet{Konopacky:2010} has tripled the number of very low mass systems with mass measurements.  In Section~\\ref{Obs} we describe the sixteen field L-dwarf targets imaged at Palomar using aperture masking with laser guide star adaptive optics and outline the data analysis techniques used to determine the binarity of the targets.  In Section~\\ref{results}, we present the results of our survey, which operated in the range of 60-320 mas (1.1-8.4 AU @ 18.4 pc, the median distance of our targets).  We identify four new candidate L dwarf-brown dwarf binaries at moderate or high (90-98\\%) confidence.  This survey achieved median contrast limits of $\\Delta$K$\\sim$2.3 between 1.2 $\\lambda/D$ and 4 $\\lambda/D$, ruling out companions down to approximately .06 \\msun for old (5 Gyr) systems and .03 \\msun for young (1 Gyr) systems.  In Section~\\ref{discussion}, we discuss the aperture masking techniques employed in this paper and present recommendations for future faint target observations.  In Section~\\ref{conclusion}, we summarize the results of this survey and discuss its implications for future companion searches around brown dwarfs.  Finally, we discuss in detail our methods for calculating detection confidence and contrast limits in the Appendix. ",
        "conclusions": "\\label{conclusion}  We present the results of a close companion search around nearby L dwarfs using aperture masking interferometry and Palomar laser guide star adaptive optics.  The combination of these techniques yielded typical detection limits of $\\Delta$K$_s$ = 1.5-2.5 between 1-4$\\lambda/D$ and limits of $\\Delta$K$_s$ = 1.0-1.7 at 0.6 $\\lambda/D$.  Our survey revealed four candidate binaries with moderate to high confidence (90-99\\%) and favorable Bayes' Factors.  Ten of the targets have previously been observed with the Hubble Space Telescope as part of companion searches.  As such, we did not expect to find bright or distant companions around these targets which would have been identified in the previous surveys.  But as demonstrated in this paper, the detection profile of aperture masking is capable of revealing close or dim companions which are obscured by the point spread function of full aperture imaging.  Aperture masking demonstrates an increase in formal resolution and detectable contrast at close separations over laser guide star adaptive optics alone.  Our survey indicated two previously observed targets as candidate binaries.  Our survey indicated one companion around 2M 0355+1133 within the formal diffraction limit of the HST and one companion around 2M 0345+2540 below the background detection threshold of the previous survey.  Two other targets, 2M 0345+2540 and 2M 2238+4354, also indicated the presence of companions, both with contrast ratios greater than 15:1.  Aperture masking is most sensitive to companions between $\\lambda / 2D$ and $4\\lambda / D$, corresponding to angular separations of 60-450 mas in K$_s$ at Palomar and physical projected separations ranging from 0.6-10 AU for the targets in our survey.  Two candidate binaries presented in this paper have projected separations less than 1.5 AU.  The results suggest a favorable target set for future companion searches.  Their candidacy is consistent with the conjecture that tight binaries are underrepresented in the current tally of low mass binaries.  Spectroscopic surveys, which focus on separations within 3 AU, are necessary to conclusively answer this question.  Extending the use of aperture masking with laser guide star AO is a rewarding approach for detecting companions within this range, and facilitating the measurements of their masses."
    },
    "1004/1004.1822_arXiv.txt": {
        "abstract": "\\noindent We begin with the premise that the law of entropy could prove to be fundamental for the evolution of intelligent life and the advent of technological civilization. Building on recent theoretical results, we combine a modern approach to evolutionary theory with Monte Carlo Realization Techniques. A numerical test for a proposed significance of the law of entropy within the evolution of intelligent species is performed and results are compared with a neutral \u201ctest\u201d hypothesis. Some clarifying aspects on the emergence of intelligent species arise and are discussed in the framework of contemporary astrobiology. \\\\  \\noindent \\textbf{Keywords:} astrobiology, SETI, evolution, entropy, Monte Carlo, extraterrestrial intelligence, Fermi paradox, exoplanets ",
        "introduction": "\\noindent Although panspermia \\citep{Hoyle1977} and the Rare Earth hypothesis \\citep{rare_Earth} are two of the most popular paradigms regarding the evolution of ETIs, modern biology indicates that alternative theories on the development of intelligent life may also carry weight. Recent papers \\citep{JaakkolaSalla2008,JaakkolaSalla2009,Kaila2008,Fu2007,NamBozhilov2009} have demonstrated that up until now, one fundamental principle of nature \u2013 the second law of thermodynamics, has not yet been fully incorporated in evolutionary algorithms. \\\\  \\noindent Entropy is fundamental to thermodynamics and its importance in physics and in astrophysics is indisputable. Biology is now incorporating it into the evolutionary paradigm. Numerous analyses (ibid.) show evolution might be understood as a process of constant complexity increase, further developing the correlation between entropy and evolution on both microscopic and macroscopic scales. \\\\  \\noindent Therefore, it may prove worthwhile to study the effect of the 2nd Law on the latter stages of evolution \u2013 in particular the evolution of intelligent technological civilizations. This paper aims to explore a new hypothesis regarding the emergence of intelligent life, which views the process of natural selection and evolution as intrinsically linked with the second law of thermodynamics.  \\\\  \\noindent The paper is divided into 6 sections. The \u201cEntropy Hypothesis\u201d, which links the second law of thermodynamics with the evolution of intelligent life, is postulated in section 2.  In section 3, we outline the numerical method by which the development of life and intelligence in the Galaxy can be simulated.  In section 4, we display the results of numerical simulations describing the development of \u201centropy-driven\u201d ETIs in the Galaxy. This data is compared with a neutral \u201cbaseline\u201d simulation to study the unique features of \u201centropy-driven\u201d ETIs.  We discuss the results and provide conclusions in sections 5 and 6 respectively.  \\\\ ",
        "conclusions": "\\noindent The entropy principle is fundamental in modern evolutionary theory: natural selection and biological evolution are deeply correlated and can be understood within the second law of thermodynamics. The direction of genomic evolution, just as other evolutionary processes, toward more probable distributions can be deduced from the logarithmic probability measure known as entropy \\citep{Sharma2007123,Kaila2008,JaakkolaSalla2009,Fu2007}.   \\\\  \\noindent We investigated the extension of this principle into intelligent technological civilizations through Monte Carlo Realization techniques. By comparing this entropy hypothesis with a more neutral baseline, it can be seen that the entropy hypothesis drives the creation and development of technology (as a new resource of increasing complexity).  This enhances the number of technological civilizations that reach an advanced stage of evolution (under a sociological pressure of \u201cinnovate or die\u201d).  \\\\  \\noindent  While the emergence of intelligent species might not be inevitable, but the process of creating such \u201creplicators\u201d in evolution is likely to be favoured, because it assures the maximum amount of free energy and will generate more entropy. \\textbf{This leads to a general definition of intelligence based on entropy arguments (discussed in section 2), which provides further insight into evolutionary theory}. However, the detectability of this sociological pressure is low: in general, although the entropy hypothesis deviates slightly from the baseline data, this deviation is too small to be able to distinguish it and prove it observationally using current instruments.   \\\\  \\noindent This work highlights two important points: firstly, that the influence of entropy may go further than was previously thought, especially in understanding the evolution of intelligent species. Secondly, our results demonstrate that sociological pressure will affect (if only slightly) future development of the Galactic community of civilizations over cosmic time, and should be considered in further studies of this type."
    },
    "1004/1004.1539_arXiv.txt": {
        "abstract": "For nearby K dwarfs, the broadening of the observed Main Sequence at low metallicities is much narrower than expected from isochrones with the standard helium--to--metal enrichment ratio $\\DYDZ \\sim 2$. Though the latter value fits well the Main Sequence around solar metallicity, and agrees with independent measurements from HII regions as well as with theoretical stellar yields and chemical evolution models, a much higher $\\DYDZ\\sim10$ is necessary to reproduce the broadening observed for nearby subdwarfs. This result resembles, on a milder scale, the very high $\\DYDZ$ estimated from the multiple Main Sequences in $\\omega$Cen and NGC~2808. Although not ``inverted'' as in $\\omega$Cen, where the metal-rich Main Sequence is bluer than the metal-poor one, the broadening observed for nearby subdwarfs is much narrower than stellar models predict for a standard helium content. We use this empirical evidence to argue that a revision of lower Main Sequence stellar models, suggested from nearby stars, could significantly reduce the helium content inferred for the subpopulations of those globular clusters. A simple formula based on empirically calibrated homology relations is constructed, for an alternative estimate of $\\DYDZ$ in multiple main sequences. We find that, under the most favourable assumptions, the estimated helium content for the enriched populations could decrease from $Y \\simeq 0.4$ to as low as $Y \\simeq 0.3$. ",
        "introduction": "\\label{sec:intro}  Helium is the second most abundant element in the Universe, having been produced by Big Bang Nucleosynthesis (BBN) with a universal mass fraction of $Y_P \\sim$ 0.24 (Steigman 2007); after that, successive stellar generations synthesize metals ($Z$) and an additional amount of helium $\\Delta Y$. The characteristic helium--to--metal enrichment rate $\\DYDZ$ is estimated to be around 2, both observationally from HII regions (e.g.\\ Fukugita \\& Kawasaki 2006; Peimbert \\etal 2007; Izotov, Thuan \\& Stasinska 2007) and theoretically, from models of stellar nucleosynthesis and chemical evolution of the Galactic disc (e.g.\\ Chiosi \\& Matteucci 1982; Maeder 1992; Carigi \\& Peimbert 2008; Casagrande 2008)  In spite of its large abundance, helium is an elusive element. It can be measured directly, from spectroscopic lines, only in stars hotter than $\\sim 10,000$~K: young, massive stars (or their surrounding HII regions) or blue Horizontal Branch stars (where though, apart from a small temperature range, the surface abundance may not trace the original one, due to helium sedimentation and metal levitation: Michaud, Vauclair \\& Vauclair 1983; Michaud, Richer \\& Richard 2008; Villanova, Piotto \\& Gratton 2009 and references therein). Only indirect methods can be used for stars of lower mass and cooler temperatures, which constitute the bulk of the Galaxy's stellar population. One such method relies on the broadening of the lower Main Sequence (MS): in fact the location of low--mass, long--lived stars in the HR diagram depends on their metallicity $Z$ and their helium content $Y$  --- or equivalently, on $\\DYDZ$. Concomitant increases in $Z$ and $Y$ shift the MS in opposite directions, so for a given variation $\\Delta Z$ the two MSs are more spread apart (the metal rich being cooler at a given luminosity) if the corresponding $\\Delta Y$ is lower; conversely, if $\\Delta Y$ is large ``enough'', the MSs may even invert, with the more metal (and helium) rich being bluer (i.e.\\ hotter; e.g.\\ Fernandes \\etal 1996).  This method dates back to Faulkner (1967); early studies based on ground--based parallaxes deduced a large and quite uncertain value of $\\DYDZ\\sim 5 \\pm 3$ from the apparent overlap of MSs of all metallicities (e.g.\\ Perrin \\etal 1977; Fernandes \\etal 1996). The situation greatly improved with Hipparcos parallaxes, so that a net separation of the MS as a function of $Z$ could be detected in the HR diagram. Studies based on Hipparcos distances concluded $\\DYDZ=2-3$ (Pagel \\& Portinari 1998; Jimenez \\etal 2003). More recently, Casagrande \\etal (2006, 2007) further improved the analysis by compiling a much larger sample of K dwarfs with homogeneous multi--band photometry, for which bolometric magnitudes and effective temperatures were derived with the InfraRed Flux Method (IRFM). The interpretation of the broadening of the lower MS could thus be performed in the theoretical HR diagram ($M_{\\rm{Bol}}$ vs.\\ $\\log \\teff$) where comparison to stellar models is more straightforward and the effects of $\\DYDZ$ more prominent (Castellani, Degl'Innocenti \\& Marconi 1999). Around solar metallicities this analysis yielded $\\DYDZ \\sim2$; but at lower $Z$ Casagrande \\etal (2007) found that the observed Main Sequence is narrower than expected from stellar models computed under the standard assumption $\\DYDZ \\sim 2$, thus implying a very steep $\\DYDZ \\sim 10$ (Fig.~\\ref{fig:YZfield}). This change in slope is at odds with Galactic chemical evolution models, that predict a ratio substantially constant with $Z$ (Carigi \\& Peimbert 2008; Casagrande 2008); a constant $\\DYDZ \\sim 2$ is also found for both metal-rich (e.g.\\ Peimbert 2003; Balser 2006) and metal-poor (e.g.\\ Fukugita \\& Kawasaki 2006; Peimbert \\etal 2007; Izotoz \\etal 2007) HII regions. More importantly, $\\DYDZ \\sim 10$ at low metallicities implies a helium content for nearby subdwarfs much lower than $Y_P$, in awkward constrast with the cosmological floor set by standard BBN.  A combination of small sample size and $Z$--dependent bolometric corrections in the observational HR diagram had masked this effect in previous studies of $\\DYDZ$, although very low $Y$ values had already been noticed in a handful of low metallicity stars with available IRFM $M_{\\rm{bol}}$ and $\\teff$ (Fernandes \\etal 1998; Lebreton \\etal 1999). This result points toward inadequacies in Main Sequence stellar models of low metallicity, and we now plan to investigate the issue further, as this may lead to reconsideration of the problem of the helium enrichment in globular clusters.  In fact, at absolute magnitudes comparable to those of the local stars studied in Casagrande \\etal (2007, $M_V \\sim 5.5-7.5$), multiple Main Sequences have been discovered in some globular clusters and interpreted as evidence for huge helium enhancement in a sub-population (Bedin \\etal 2004; Norris 2004; Piotto \\etal 2005, 2007); see Fig.~\\ref{fig:WCen}. The required $\\DYDZ \\gtrsim 100$ is extremely difficult to explain with stellar nucleosynthesis and chemical evolution models (Karakas \\etal 2006; Prantzos \\& Charbonnel 2006; Maeder \\& Meynet 2006; Choi \\& Yi 2007, 2008; Romano \\etal 2007, 2009; Renzini 2008; Yi 2009; Marcolini \\etal 2009; Peng \\& Nagai 2009).  The helium content of the sub-population also has a major role in shaping other regions of the HR diagram (Catelan, Valcarce \\& Sweigart 2009a and references therein); in particular it is reflected in the morphology of the Horizontal Branch  (D'Antona \\& Caloi 2004; Caloi \\& D'Antona 2005, 2007; Lee \\etal 2005). While it is likely that a significant helium enhancement is present in those stellar populations, the purpose of this work is to show that the revision of low metallicity MS stellar models, needed to cure the problem of the high $\\DYDZ$ and sub-primordial $Y$ deduced in nearby K dwarfs, may significantly reduce current estimates of $\\DYDZ$ in globular clusters with multiple MSs, easing their theoretical interpretation.  We show this by means of an exercise based on homology relations. Theoretical homology relations, describing how the location of the low MS in the HR diagram depends on metallicity and helium content, are briefly recalled in Section~\\ref{sec:homology}. In Section~3, we apply homology relations to interpret the split of multiple MSs in globular clusters, and show that the inferred helium enrichment for the blue population(s) agrees with the results of detailed isochrone analysis. In Section~4 we confirm that, in general, homology relations render very well the behaviour of stellar models as a function of the helium content. However, both isochrones and homology relations fail the interpretation of the low MS of nearby stars ---  as we cannot accept sub--primordial helium contents for low metallicity stars. Therefore, in Section~5 we proceed to calibrate {\\it empirical} homology relations, that return the expected $\\DYDZ = 2$ for the low MS of nearby stars with $Z<Z_{\\odot}$. Such empirically calibrated homology relations also return, for the helium--rich subpopulations of globular clusters, a helium content $Y \\sim 0.3$, rather than the much higher $Y \\sim 0.35-0.4$ of standard analysis. This lower value of helium enrichment can be reconciled far more easily with present chemical evolution models (Karakas \\etal 2006; Renzini 2008; Yi 2009). Finally, we conclude in Section~6 with a discussion of the possible physical processes that could improve stellar models for low MS stars, so as to solve the ``helium problem'' both for nearby stars and for globular clusters.  \\begin{figure} \\begin{center} \\includegraphics[scale=0.43]{figure1.eps} \\caption{{\\it Left panel}: HR diagram of nearby K dwarfs (filled circles) from the sample of Casagrande \\etal (2007), overplotted on the MSs photometry of $\\omega$Cen (Sollima \\etal 2007), adopting $E(B-V)=0.11$ (Lub 2002) and dereddened distance modulus of 13.7 (Bellazzini \\etal 2004; Del Principe \\etal 2006). The maximum split between the red and blue MS of $\\omega$Cen $\\Delta (B-R)=0.1$ occurs around $M_R \\sim M_{\\rm{Bol}} \\sim 6.5$ (Sollima \\etal 2007). {\\it Right panel}: nearby K dwarfs in the theoretical plane with isochrones having $Z=0.001, Z_{\\odot}, 0.04$ and $\\Delta Y / \\Delta Z=2$ overplotted. Clear mismatch in the broadening appears for low metallicities. Below $M_{\\rm{Bol}}=5.4$ the broadening of the MS is age independent. We also show the reddening/extinction correction corresponding to $E(B-V)=0.10$ that would be needed to reconcile the metal-poor stars with theoretical isochrones.} \\label{fig:WCen} \\end{center} \\end{figure}  \\begin{figure*} \\begin{center} \\includegraphics[scale=0.43]{figure2a.eps} \\includegraphics[scale=0.43]{figure2b.eps} \\caption{Broadening of the ZAMS in log $\\teff$ as a function of metallicity, as predicted by homology relations (Eq.~\\protect{\\ref{eq:DlogTeff}}). The broadening is referred to the primordial ZAMS ($Z_P$, $Y_P$; left panel) or to the solar ZAMS ($Z_\\odot=0.017$, $Y_\\odot=0.263$; right panel) for $\\frac{\\Delta Y}{\\Delta Z}$ ranging from 0 to 10 as indicated. Dotted lines show the approximate relation proposed in Eq.~\\protect{\\ref{eq:DlogTeff_2terms}}. The thick red line highlights the case $\\frac{\\Delta Y}{\\Delta Z}=5$, which corresponds to MS inversion around $Z_\\odot$. At increasing $\\DYDZ$, MS inversion occurs at lower and lower $Z$.} \\label{fig:homology} \\end{center} \\end{figure*} ",
        "conclusions": "In this paper we draw attention on the possible connection between the HR diagram of globular clusters with multiple Main Sequences ($\\omega$Cen and NGC~2808) and the broadening of the low MS defined by nearby stars. Although not ``inverted'' as in $\\omega$Cen, where the metal-rich Main Sequence is bluer than the metal-poor one, the broadening defined by nearby subdwarfs is much narrower than expected for a standard helium content and helium--to--metal enrichment law $\\DYDZ \\sim 2$. At low metallicities, $\\DYDZ\\sim 10$ is formally necessary to reproduce the observed broadening, but this is unacceptable as it implies helium fractions much lower than the cosmological BBN floor.  It is worth remarking here that the results by Casagrande \\etal (2007) on the broadening of the low MS  were corroborated by the analysis of a number of binary stars, whose masses were well reproduced by theoretical isochrones --- again at the expense of assuming, for some low Z binaries, sub-primordial helium abundances. Also, one may argue that the problem lies in an incorrect $\\teff$ scale for low Z stars. Without entering here a detailed discussion on the robustness of the IRFM scale by Casagrande \\etal (2006, 2010), suffice here to mention that this is one of the hottest scales available (as discussed in the original papers and in Sousa et al.\\ 2008). It is hotter by 100~K than other IRFM renditions (the now superseded scales by Alonso, Arribas \\& Martinez--Roger 1996 and Ram\\'irez \\& Mel\\'endez 2005), and it is comparable within $\\pm$50~K to various spectroscopic scales in the metallicity range relevant for this work.  Therefore, any other $\\teff$ scale will just worsen the K dwarf problem, with real stars even cooler than the models. Also, it does not seem plausible that all the (independent) scales available are systematically offset by $> 200$~K, which is what is needed to bring the K dwarfs of Casagrande et al.\\ (2007) in line with stellar models.  We estimate the possible extent of the required revision of low MS stellar models on the base of homology relations. First we show that theoretical homology relations properly reproduce the response of stellar models to the helium content: in particular, analyzing the split of the multiple MS in $\\omega$Cen and NGC~2808 by homology relations yields the same conclusions as full isochrone analysis (i.e.\\ a helium content $Y \\sim$ 0.4 for the blue subpopulations). Then, since both isochrones and theoretical homology relations fail the interpretation of the nearby low--Z Main Sequence, we calibrate {\\it empirical} homology relations to yield consistently $\\DYDZ\\sim 2$ for nearby stars, and inspect the consequences for the MS splits in globular clusters.  {\\it If} we entirely impute the failure of the low-MS stellar models to a wrong response to the helium abundance (and correspondingly re-calibrate the helium--sensitive term of the homology relations), the consequences for globular clusters are highly significant: the helium content of the blue sub-populations is reduced from 40\\% to 30\\%, which is far easier to explain with chemical evolution models (Renzini 2008; Yi 2009).  Alternatively, if stellar models for the low MS are assumed to fail in their metallicity dependence (i.e.\\ we re-calibrate only the homology term expressing the metallicity dependence of opacity) the consequences for globular clusters are negligible --- as the metallicity differences between the subpopulations are minimal or vanishing.  As the solution for K dwarf models can be intermediate between these two extreme assumptions, we suggest that the helium rich populations in globular clusters are likely to have a helium content $Y$ in between 0.3--0.4; but altogether, there is room to decrease their estimated helium content from the extreme $Y=0.4$ that is the commonly quoted value.  In summary, the purpose of this exercise is to draw the attention of model--makers to the problem of the HR diagram of nearby low--Z stars: improvements in this respect are potentially important also for the riddle of the helium self-enrichment in globular clusters.  We remind indeed that, while a helium content as high as $Y \\sim 0.35-0.4$ seems to nicely account for the multiple Main Sequences and the morphology of the Horizontal Branch (HB), other observations provide some counter-evidence for such a helium rich sub-population. In $\\omega$Cen, the location of the RGB bumps of the metal--poor and metal--intermediate populations is consistent with a maximum $\\Delta Y < 0.1$, less than what is derived from MS and HB analysis (Sollima \\etal 2005). Sollima \\etal (2006) have also identified RR Lyr\\ae\\ stars with metallicity corresponding to that of the blue MS, but normal helium content --- which would require the metal--intermediate population to be further split into a helium rich and a helium normal subcomponents.  NGC 6752 is another cluster suspected to host a helium enriched population, due to the morphology of its HB and the presence of a broadened, possible multiple MS (Milone et al.\\ 2010); however, Villanova \\etal (2009) did not detect large helium abundances in the spectra of blue HB stars in the (narrow but crucial) temperature range $8500 < \\teff < 11500$, hot enough to produce helium lines but still unaffected by helium sedimentation. Other studies of clusters suspected to host helium rich sub-populations have not confirmed their presence (e.g.\\ Lee \\etal 2009 for NGC 1851; Catelan \\etal 2009b for M3). While awaiting for independent proof of high helium abundances from direct spectral measurements or other features in the HR diagram (see the recent review by Catelan \\etal 2009a), we remark that the helium abundances obtained by our proposed revision, $Y \\lsim 0.3$, are consistent with all of the above--mentioned constraints.  The need for sub-primordial helium abundances to fit a handful of nearby subdwarfs on the HR diagram was first noticed by Fernandes \\etal (1998) and Lebreton \\etal (1999) who advocated the inclusion of additional physical processes not implemented in standard stellar evolutionary calculations. In the following, we discuss a number of possible solutions to the problem. Broadly speaking, we can classify them as metallicity--dependent solutions, which shall not concern the helium rich populations of globular clusters, and helium--dependent solutions, which we expect also to impact the interpretation of globular clusters. We may add that, considering the nice agreement between the multiple MS and the HB morpology in $\\omega$ Cen and NGC~2808, the optimal helium--dependent solution should preferably preserve this relation; this can be achieved if the new stellar models will not just shift the low metallicity ZAMS to lower effective temperatures, but also affect the luminosity of individual stars and accelerate their evolution, which helps to populate the blue side of the HB. Casagrande \\etal (2007) discussed a number of possible solutions (all in the metallicity--dependent class). We briefly recall them and discuss a few more here below.  \\paragraph*{Mixing length} A mixing length parameter decreasing with metallicity would make low Z models redder; for the isochrones we used, the variation should be from the reference solar value $\\alpha$=1.68 to $\\alpha$=1.0 at low Z. While any dependence of $\\alpha$ on metallicity, mass or other physical parameters is still highly disputed (Casagrande \\etal 2007 and references therein) this solution would have no consequence for the helium content of globular clusters, as it is metallicity dependent and does not significantly affect the luminosity and lifetime of the stellar models.  \\paragraph*{Diffusion + non--LTE effects} Element diffusion affects the location of low MS stars in the HR diagram: it makes stellar models redder, and moreover it lowers the measured surface metallicity with respect to the real intrinsic one. In this scenario, low Z subdwarfs look redder than we expect, mostly because they are actually more metal rich than we measure. Casagrande \\etal (2007) discussed this possibility resorting to literature models with fully efficient (``maximal'') diffusion, and found that it does not completely solve the problem, unless it is combined with additional errors on the observed metallicity due to non--LTE effects. (This solution also would have no effect on globular clusters, since it is metallicity and age dependent, and the gap in both is small for the multiple populations of globular clusters.)  However, as discussed in that paper, there is evidence that both of these mechanisms are not fully efficient in real stars, as one would need to solve the K dwarf problem. For diffusion in particular, observations of field and cluster metal-poor dwarfs point toward inhibited efficiency and additional --- yet ad hoc --- processes are nowadays invoked to contrast diffusion (e.g.\\ Chaboyer \\etal 2001; Richard, Michaud \\& Richer 2002, 2005; Korn \\etal 2007).  \\paragraph*{Boundary conditions} also play a role in stellar evolutionary tracks. A solar scaled $T-\\tau$ relation provides a good description of the atmospheric structure also in metal-poor stars, and may contribute to improve low Z models with respect to gray atmosphere boundary conditions. However, the effect appears to be quite small for MS stars, when each set of isochrones consistently adopts a mixing length parameter calibrated on the solar model (Vandenberg \\etal 2008).  Besides, boundary conditions are expected to become more relevant at lower stellar masses and luminosities (due to the deeper surface convection), while in Casagrande \\etal (2007) we verified that there is no systematic trend of the derived $\\DYDZ$ with $M_{bol}$.  Also this solution relies on a metallicity--dependent effect and would not bring substantial revisions to the helium content in the multiple MS of GCs.  \\paragraph*{Convection} One can always consider changing the boundary of convective regions as a viable working hypothesis. (In fact, the very presence of a convective envelope renders real K dwarfs not rigorously representable by homology relations.) Considering that the convective envelope gets thinner at lower metallicities, any change in the convection scheme (e.g. extra-mixing, or undershooting) is expected to affect solar metallicity objects more, creating a differential effect in metallicity that could change the relative location of the low MS as a function of Z. Also this solution falls in the metallicity-dependent category and is not expected to impact on the GC analysis.  \\paragraph*{Opacity} As K dwarfs are largely radiative structures, in particular in the regions where nuclear energy is produced, opacity is a leading ingredient in determining their luminosity. An incorrect metallicity dependence of opacity would seriously affect the broadening of the low MS. The idea is tempting, as the ``break'' in the estimated $\\DYDZ$ occurs close to $Z_0=0.01$, which is the classic divide between free-free and bound-free dominated opacity in stellar interiors (Cox \\& Giuli 2004). The parameter $P_3$ in the original homology relations is related to this characteristic metallicity as $P_3=1/Z_0$ (Section~\\ref{sec:homology_num} and Appendix) and we have seen that modern isochrones are better described by $P_3 \\simeq 670$ rather than 100, effectively reducing $Z_0$ to 0.0015 (Section~\\ref{sec:homology_num}; Pagel \\& Portinari 1998). However, optimizing empirical homology relations on the opacity term, restores a value for $P_3$ or $Z_{ff/bf}$ closer to the theoretical one ($P_3=150$ or $Z_0\\simeq 0.007$). This suggests that the culprit might be the metallicity  dependence of opacity (which would then be irrelevant for globular clusters). It is easier to figure errors in the bound-free contribution to opacity, rather than in the free-free component, at least for the bulk of the star wher H and He are completely ionized. Lowering the bound-free contribution to opacity would imply a recalibration of solar metallicity isochrones, after which high Z (bound-free dominated) and low Z (free-free dominated) model MS may fall closer to each other, thus reducing the broadening. However, we wonder if there is much room for profound changes in opacity, given the excellent agreement between the major current, independent databases (Opacity Project, Seaton 2005 and references therein; and OPAL, Iglesias \\& Rogers 1996 and references therein) and keeping in mind that improvement in atomic data, input physics etc. most often leads to an increase in the opacity, as more and more opacity sources are taken into account. Helioseismology has shown extant OPAL and OP opacities to be fully adequate for Solar models with the ``classic'' solar composition; while an increase in opacity by 10--20\\% in a suitable temperature range has been invoked for the Solar model with new, lower metallicity --- possibly in connection with increased neon abundance (Basu \\& Antia 2008; Asplund et~al.\\ 2009; Serenelli et~al.\\ 2009; and references therein). How these or other fundamental changes of the Solar model, to recover agreement with helioseismological constraints, will impact low Z subdwarfs and the {\\it relative} location of MS of subsolar metallicity on the HR diagram, remains to be explored.  \\medskip \\noindent Finally, we notice that the ``helium problem'' may not be limited to K dwarfs: systematic temperature offsets from the theoretical Main Sequence at low Z have been highlighted for FG dwarfs in the Geneva--Copenhagen survey (N\\\"ordstrom et~al.\\ 2004); and a helium content Y=0.23--0.24 (i.e.\\ slightly sub-cosmological) has been recently suggested for a slighly metal-poor F dwarf binary ([Fe/H]=--0.25; Clausen et~al.\\ 2010). If the problem indeed extends to FG dwarfs, clearly some of the suggested solutions (for instance, those related to diffusion) are not viable.  \\medskip \\noindent All of the solutions suggested above rely, more or less indirectly, on metallicity dependence; this is the standard way we think of stellar models. However, our exercise in this paper highlights that it is worth thinking of other possible systematics in the stellar models, especially connected to the helium abundance, for their interesting impact on the multiple MS of globular clusters.  As mentioned above, opacity has a key role in the structure of lower MS stars, so we may wonder whether the issue can be the helium dependence of opacity. This can be hardly modified in regions of complete ionization; more promising are regions where He is partially recombined and contributes to the boud-free opacity.  Other helium--dependent effects may be considered, such as diffusion mechanisms that would act differently for helium as for metals. (However, preliminary tests on only--helium diffusion seem to go in the opposite direction as needed for globular clusters, with the helium rich MS getting proportionally more red than the helium poor; A.\\ Serenelli, priv.\\ comm.)  Possibly, other of the above mentioned solutions (mixing length, convection etc.) can be recast and explored in terms of helium dependence, rather than metallicity dependence. Any ideas in this direction are highly desired: hopefully, astrophysicists will be as creative in solving the problem of the HR diagram of low metallicity K dwarfs, as they have been in tackling the riddle of the extreme helium rich populations in globular clusters!"
    },
    "1004/1004.4636_arXiv.txt": {
        "abstract": "We present new results from our search for Giant H\\,{\\sc ii} Regions in galaxies visible from the southern hemisphere. In this work we study two galaxies: NGC\\,7479 and NGC\\,6070. Using high-resolution spectra, obtained with different instruments at Las Campanas Observatory, we are able to resolve the emission-line profile widths and determine the intrinsic velocity dispersion of the ionised gas. We detect profile widths corresponding to supersonic velocity dispersions in the six observed H\\,{\\sc ii} regions. We find that all of them show at least two distinct kinematical components: a relatively narrow feature (between ~11 and ~22\\kms) and a broader (between ~31 and ~77\\kms) component. Two of the regions show a complex narrow profile in all ion lines, which can be further split into two components with different radial velocities. Whereas the wing broadening of the overall profile can be fitted with a low-intensity broad component for almost all profiles, in one region it was better reproduced by two separate shell-like wings. We have analysed the impact that the presence of multiple components has on the location of the H{\\sc ii} regions in the $\\log(L) - \\log(\\sigma)$ plane. Although the overall distribution confirms the presence of a regression, the precise location of the regions in the plane is strongly dependent on the components derived from the profile fitting. ",
        "introduction": "The most extended and luminous objects which are observed on the disks of spirals, in irregulars and in starburst galaxies are called Giant Extragalactic H{\\sc ii} Regions (GH{\\sc ii}Rs). GH{\\sc ii}Rs  are places of very active star formation and they are characterised by the large emission of ultraviolet photons from a large number of young and massive stars. Owing to the large number of ionising photons most of the surrounding gas is ionised, and the recombination process provides a strong signature of recent or ongoing massive star formation.  The emission-line profile widths of the giant H{\\sc ii} regions imply the existence of supersonic motions in the gas \\citep{SW70}. Furthermore, \\cite{TM81} found a correlation between the gas velocity dispersion ($\\sigma$) and the total luminosity (L) emitted in the respective line. \\cite{H86} and \\cite{1988A&A...201..199A} confirmed the existence of such a regression, but there was no agreement on the values derived for its slope so no conclusion could be reached regarding the origins of this motion. \\cite{2000AJ....120..752F} and \\cite{B02} using different techniques on different samples agreed on a slope close to a value of 4, which seemed to favour the gravitational source of energy for the observed supersonic motion. Even if the nature of this behaviour is not fully understood, it still allows us to distinguish giant H{\\sc ii} regions from an agglomeration of classical H{\\sc ii} regions by means of high resolution spectroscopy. GH{\\sc ii}Rs  provide the link between small-scale star forming regions, such as Orion in our Galaxy and violent star formation processes taking place in starburst galaxies.  In \\cite{Firpo05} we have confirmed the giant nature of three candidates to giant H{\\sc ii} regions in the southern sky which were identified as very luminous H{\\sc ii} regions by \\cite{F97}.  Here we continue our search and detailed analysis of GH{\\sc ii}Rs  in local Universe galaxies. We have selected from the Feinstein catalogue, a sample of the brightest H{\\sc ii} regions from two spiral galaxies: NGC\\,7479 and NGC\\,6070.  NGC\\,7479 is a barred spiral galaxy, at a distance of 31.92 Mpc \\citep{1999A&AS..135..145R}, classified as SB(s)c by \\cite{dV91}. \\cite{1989ApJ...346..126D} and \\cite{1997A&A...323..363M} have classified it also as a starburst galaxy, whereas \\cite{1983ApJ...269..466K} and \\cite{1997ApJS..112..315H} have identified it as a LINER and a Seyfert 1.9 galaxy respectively. Many studies have been performed on this galaxy owing to the presence of a certain asymmetry in the spiral structure \\citep{1995ApJ...441..549Q} that leads several authors to suggest this might be owed to a recent merger event \\citep{1996PhDT........82L,1999MNRAS.308..557L}. This scenario can also account for perturbations in the velocity field \\citep{1998MNRAS.297.1041L} and the latest minor merger model indicates that the remnant may be situated within the bar \\citep{2001Ap&SS.276..667L} as no remnant of the merger could be identified in optical images \\citep{2003A&A...409..899S}.  NGC\\,6070, at a distance of 29.8 Mpc \\citep{2002A&A...389...68G}, is classified as SA(s)cd by \\cite{dV91} although \\cite{2004A&A...423..849G} state that a bar with two main arms can be identified. \\cite{2006A&A...454..759P} further studied the bar structure classifying it as Type II.o.CT according to \\cite{2008AJ....135...20E}. The galaxy shows a normal arm distribution, with two arms that break into several on each side of the central region (see Figure \\ref{fig6070chart}). The disk dynamics was studied by \\cite{2002A&A...393..389M}, and \\cite{2002A&A...389...68G} studied the distribution of \\HII\\ regions in the galaxy.  In order to verify the nature of our giant H{\\sc ii} region candidates, we obtained high resolution spectra to measure the emission-line profile widths and to estimate if the velocity dispersion is indeed supersonic. To determine the optimum number of Gaussian components to fit each line profile in order to minimise the uncertainty in the velocity dispersion of the ionised gas, we have considered information from other studies available in the literature.  Previous papers have dealt with the fact that a single Gaussian profile might not be realistic enough when fitting the observed emission lines. Several studies favour the presence of an omnipresent broad component which explains the integral profile wings. They include \\cite{M99} who identified a broad component in the recombination lines in the central region of the 30 Doradus nebula, \\cite{1987MNRAS.226...19D} and \\cite{1996MNRAS.279.1219T} in the M33 giant \\HII\\ region NGC\\,604, \\cite{1997ApJ...488..652M} in four Wolf-Rayet galaxies, and \\cite{1999ApJ...522..199H} in the starburst galaxy NGC\\,7673. More recently, \\cite{H07,2008PhDT........35H,2009MNRAS.396.2295H} and \\cite{Hagele+10} found that the best Gaussian fits in circumnuclear star-forming regions involved the presence of broad and narrow components for the emission lines of the ionised gas.  A variety of line broadening mechanisms have been proposed to interpret the existence of the broad supersonic component measured in the emission line profile of GH{\\sc ii}Rs. Dynamics of virialised systems (Tenorio-Tagle et al.~\\citeyear{TT93}), superposition of multiple gas bubbles in expansion (Chu \\& Kennicutt~\\citeyear{CHK94}), or turbulence of the same interstellar gas (Medina Tanco et al.~\\citeyear{MT97}). More recently, Westmoquette et al.~\\citeyear{2007MNRAS.381..894W,2007MNRAS.381..913W} concluded that the narrow component represents the general disturbed ionised interstellar medium (ISM), arising through a convolution of the stirring effects of the starburst and gravitational virial motions. On the other hand, the broad component results from the highly turbulent velocity field associated with the interaction of the hot phase of the ISM with cooler gas knots, setting up turbulent mixing layers.  Other studies support the existence of shell-like profiles on the blue and red wings of the main feature as found by \\cite{CHK94}, who state that the integrated line profiles of the entire 30 Dor region may be fitted with a broad Gaussian with low-intensity wings. \\cite{Relano05} claim that a large fraction of the H{\\sc ii} regions in three barred spiral galaxies shows high-velocity and low-intensity features in its integrated line profiles and \\cite{2006A&A...455..539R} find that the great majority of their sampled H{\\sc ii} regions show that the brightest principal component is accompanied by low-intensity components symmetrically shifted in velocity.  In this paper we present echelle data obtained at Las Campanas Observatory (LCO) of six GH{\\sc ii}Rs candidates in the two mentioned galaxies. We have determined the velocity dispersion and we have evaluated the possible presence of a broad component or two symmetric low-intensity components in the fit with the observed emission line profile widths. In section \\ref{sec:observations} we present the observations and the data reduction. Section \\ref{sec:analysis} details the analysis performed on the emission-line profiles and discusses each conspicuous case individually. The summary and conclusions of this work are included in Section 4. ",
        "conclusions": "From new high resolution spectra of the H{\\sc ii} regions  NGC7479\\,I, NGC7479\\,II, NGC7479\\,III, NGC6070\\,I, NGC6070\\,II and NGC6070\\,IV obtained at the 100-inch du Pont Telescope, LCO, we have confirmed the giant nature of all these regions. We have also found that all of them show evidence of wing broadening evident mainly in the H$\\alpha$ line and confirmed in other emission lines. In five of them we are able to fit a broad component which explains the integral profile wings. Only in NGC7479\\,II can we fit the integrated profile wings with two narrow components symmetrically shifted in velocity with respect to the average of the two intense components.  NGC7479\\,I and NGC7479\\,II reveal the presence of two separated kinematic components with relatively narrow profiles in all analysed emission lines. For NGC7479\\,I region we can also generate 2D velocity images in H$\\alpha$, [\\NII]$\\lambda$6584\\AA\\ and [\\SII]$\\lambda$6717\\AA\\ lines. These images show that the two narrow components are spatially resolved and the kinematical information helps to split the knot in the velocity-distance plane. There is also evidence of the presence of a broad component in all the narrower extracted sections, at least in the H$\\alpha$ emission line. For NGC7479\\,II region we can also fit two relatively narrow Gaussian components. In both regions the radial velocity average value of our components A and B derived from the H$\\alpha$ emission line matches the global velocity in the galaxy when we project our value over the isovelocity contours in H$\\alpha$ velocity field map. A similar comparison, using the H\\,I map yields a value much closer to the value measured for component B. Then, we could only suggest that component A shows an odd kinematic behaviour. New observations with better spatial and spectral resolution are needed to clarify this point.  In the rest of the studied H{\\sc ii} regions, we find one narrow Gaussian component together with an underlying broad component, with no evidence of multiple narrow contributions. The radial velocity average derived for the narrower component corresponds to the one expected from galaxy rotation curves, within the observational errors, in all of these cases.  In most cases, the velocity dispersions of the broad components for these HII regions derived by optimal Gaussian fit are in agreement with \\cite{2009MNRAS.396.2295H,Hagele+10} for circumnuclear star-forming regions with $\\sigma_{int}$ in H$\\beta$ line $\\sim$34 to 65\\kms, and those derived by \\cite{CHK94} and \\cite{M99} for 30 Doradus nebula with $\\sigma_{int}$ in H$\\alpha$ line $\\sim$45\\kms.  The estimated offsets between the narrow and broad components in NGC\\,7479 III, NGC\\,6070 II, and NGC\\,6070 IV are in complete agreement with the values found in NGC\\,2903 and NGC\\,3310 by \\cite{2009MNRAS.396.2295H,Hagele+10} respectively, which are between -25 and 35 \\kms.  The distribution of the regions in the $\\log(L) - \\log(\\sigma)$ plane are strongly dependent on the components derived from the profile fitting. Individual components have smaller fluxes and velocity dispersions than the global profile and points are therefore shifted in the diagram. In NGC\\,6070 the values derived from the global profile seem to lie, within observational errors, in the expected location if they follow the relation expected for virialised systems. When the individual components are plotted in the same diagram, only NGC\\,6070\\,II lies close to the regression, and NGC\\,6070\\,I and IV seem to be too luminous for their measured velocity dispersion. On the other hand, two out of three regions in NGC\\,7479 lie well away this regression if parameters are derived from a single Gaussian, but four out of five individual components show a very good agreement when plotted in the same diagram. Definitely, the presence of more than one Gaussian component rules the final position of the H{\\sc ii} regions in the $\\log(L) - \\log(\\sigma)$ plane."
    },
    "1004/1004.4546_arXiv.txt": {
        "abstract": "Fitting synthetic spectral energy distributions (SED) to the multi-band photometry of galaxies to derive their star formation rates (SFR), stellar masses, ages, etc. requires making a priori assumptions about their star formation histories (SFH). A widely adopted parameterization of the SFH, the so-called $\\tau$-models where SFR $\\propto {\\rm e}^{-t/\\tau}$ is shown to lead to unrealistically low ages when applied to a sample of actively star forming galaxies at $z\\sim 2$, a problem shared by other SFHs when the age is left as a free parameter in the fitting procedure. This happens because the SED of such galaxies, at all wavelengths, is dominated by their youngest stellar populations, which outshine the older ones. Thus, the SED of such galaxies conveys little information on the beginning of star formation, i.e., on the age of their oldest stellar populations.  To cope with this problem, besides $\\tau$-models (hereafter called direct-$\\tau$ models), we explore a variety of SFHs, such as constant SFR and inverted-$\\tau$ models (with SFR $\\propto {\\rm e}^{+t/\\tau}$), along with various priors on age, including assuming that star formation started at high redshift in all the galaxies in the test sample. We find that inverted-$\\tau$ models with such latter assumption give SFRs and extinctions in excellent agreement with the values derived using only the UV part of the SED, which is the one most sensitive to ongoing star formation and reddening.  These models are also shown to accurately recover the SFRs and masses of mock galaxies at $z\\sim 2$ constructed from semi-analytic models, which we use as a further test.  All other explored SFH templates do not fulfil these two test as well as inverted-$\\tau$ models do. In particular, direct-$\\tau$ models with unconstrained age in the fitting procedure overstimate SFRs and underestimate stellar mass, and would exacerbate an apparent mismatch between the cosmic evolution of the volume densities of SFR and stellar mass.  We conclude that for high-redshift star forming galaxies an exponentially increasing SFR with a high formation redshift is preferable to other forms of the SFH so far adopted in the literature. ",
        "introduction": "The evolution of the baryonic component in galaxies is the hardest part to model in galaxy formation theories. How gas is accreted onto dark matter halos, its thermal history, how it is turned into stars, and if, how and when such star formation is quenched cannot be reliably predicted from first principles. The physical processes that are involved are too complex and non linear, with hydrodynamic simulations failing by a large margin to cover the extremely wide dynamical range that would be required to describe phenomena ranging from the formation of stars and supermassive black holes to the behaviour of multiphase gas on megaparsec scales. Where  it is hard to progress with pure theory alone, observations can help and lead to further advances in our understanding of how galaxies form and evolve. Indeed, over the last two decades a wealth of multiwavelength observations from both ground and space facilities have provided us with a rich vision of the galaxy populations in the high redshift universe. Once correctly interpreted, such photometric, spectroscopic, and high spatial resolution data can provide us with estimates of stellar mass ($M_*$), star formation rate (SFR) and star formation history (SFH), structure, dynamics and nuclear activity, for a very large number of galaxies, and do so as a function of redshift and environment.  In particular, the  accurate measurement of stellar mass and SFR is critical for trying to establish the evolutionary links connecting galaxy populations at one redshift with those at another redshift. Setting constraints on the previous SFH of individual galaxies is also relevant in this context. All this is currently obtained by fitting the spectral energy distribution (SED) of synthetic stellar populations to the SED of galaxies from their multi-band photometry. In practice, a large set of template synthetic SEDs is constructed with different SFHs, extinctions and metallicities, and a best fit is sought by picking the SFH that minimises the $\\chiquadr$.  In this paper we focus on star forming galaxies at $1.4\\lapprox z\\lapprox 2.5$, hence covering the epoch of major star formation activity, and explore a wide set of SFHs showing advantages and disadvantages of different parameterizations of them. The case of passively evolving galaxies at $z\\gapprox 1.4$ was already addressed in Maraston et al. (2006).  It is well known that at $z\\sim 2$ galaxies with SFRs as high as a few $\\sim 100 M_{\\odot}/{\\rm yr}^{-1}$ are quite common (e.g., Daddi et al. 2005), and by analogy with the rare objects at $z\\simeq 0$ with similar SFRs (like Ultra Luminous Infra-Red Galaxies, ULIRGs), it was widely believed that such galaxies were caught in a merging-driven starburst. However, integral-field near-infrared spectroscopy has revealed that at least some of these galaxies have ordered, rotating velocity fields with no kinematic evidence for ongoing merging (Genzel et al.\\ 2006). Still, the disk was shown to harbour several star-forming clumps and to have high velocity dispersion and gas fraction. All this makes such disk quite different from local disk galaxies, and it is now well documented that the same properties apply to many similar objects at $z\\sim 2$ (F\\\"orster-Schreiber et al. 2009).  That high SFRs in $z\\sim 2$ galaxies do not necessarily imply {\\it starburst} activity became clear from a study of galaxies in the GOODS fields (Daddi et al. 2007a). Indeed, for starforming galaxies at $1.4\\lapprox z\\lapprox 2.5$ the SFR tightly correlates with stellar mass (with SFR $\\propto\\;\\sim M_*$), with small dispersion ($\\sim 0.3$ dex), as shown in Figure~\\ref{sfrm}. Only a few galaxies lie far away from the correlation: a relatively small number of passive galaxies (with undetectable SFR, not shown in the figure), and sub-mm galaxies (SMG) with much higher SFRs, which may indeed be the result of gas-rich major mergers.  Among starforming galaxies, the small dispersion of the SFR for given $M_*$ demonstrates that these objects cannot have been caught in a special, starburst moment of their existence.  Rather, they must sustain such high SFRs for a major fraction of the time interval between $z=2.5$ and $z=1.4$, i.e. for some $1-2$ Gyr, instead of the order of one dynamical time ($\\sim 10^8$ yr) typical of starbursts. Similar correlations have also been found at lower redshifts, notably at $z\\sim 1$ (Elbaz et al. 2007), $0.2\\sim z\\sim 1$ (Noeske et al.  2007), and $z\\sim 0$ (Brinchmann et al. 2004).  \\begin{figure} \\includegraphics[width=0.49\\textwidth]{Fig1.eps} \\caption{The SFR vs. stellar mass of star forming $BzK$-selected galaxies in the GOODS-South field (from Daddi et al. 2007a). The sub-sample with spectroscopic redshifts is indicated, and represents the set of galaxies for which various SED fits are attempted in this paper. The SFRs are in $\\msun$/yr and masses in $\\msun$ units. \\label{sfrm}} \\end{figure}  In a recent study, Pannella et al. (2009) have measured the average SFR vs. stellar mass for $1.4\\lsim z\\lsim 2.5$ galaxies using the 1.4 GHz flux from the VLA coverage of the COSMOS field (Schinnerer et al. 2007). When combined also with data at lower redshifts (see references above), Pannella et al. derive the following relation for the average SFR as a function of galaxy mass and time: \\begin{equation} <{\\rm SFR}>\\simeq 270\\times (M_*/10^{11}\\msun)\\times (t/3.4\\times 10^9 {\\rm yr})^{-2.5}, \\end{equation} where the SFR is in $\\msun$/yr units and $t$ is the cosmic time, i.e., the time since the Big Bang. Beyond $z\\sim 2.5$ ($t\\lsim 2.7$ Gyr) the specific SFR (= SFR/$M_*$) appears to flatten out and remain constant all the way to very high redshifts (Daddi et al. 2009; Stark et al. 2009; Gonzalez et al. 2009).  In parallel with these observational evidences, theorists are shifting their interests from (major) mergers as the main mechanism to grow galaxies, to continuous {\\it cold stream} accretion of baryons, that are then turned into stars in a quasi-steady fashion (e.g., Dekel et al. 2009). Clearly, a continuous, albeit fluctuating SFR such as in these cold stream models provides a far better match to the observed tight SFR$-M_*$ relation, compared to a scenario in which star formation proceeds through a series of short starbursts interleaved by long periods of reduced activity. This is not to say that major mergers do not play a role. They certainly exist, and can lead to real giant starbursts bringing galaxies to SFRs as high as $\\sim 1000\\; M_{\\odot}/{\\rm yr}$, currently identified with SMGs (e.g. Tacconi et al. 2008; Menendez-Delmestre et al. 2009).  Deriving SFRs, ages, stellar masses, etc. from SED fitting requires making assumptions on the previous SFH of galaxies. A widespread approach is to fit the SED of galaxies at low as well as high redshifts with so called ``$\\tau$-models\", i.e., synthetic SEDs in which the SFH is described by an exponentially declining SFR, starting at cosmic time $t_\\circ$, i.e.: \\begin{equation} {\\rm SFR}=A{\\rm e}^{-(t-t_\\circ)/\\tau}, \\end{equation} where $A=$SFR($t=t_\\circ$). For a galaxy at cosmic time $t$ a $\\chi^2$ fit then gives the age (i.e., $t-t_\\circ$, the time elapsed since the beginning of star formation), the SFR e-folding time $\\tau$, the reddening $E(B-V)$, the metallicity $Z$ and finally the stellar mass $M_*$ via the scale factor $A$. The SFR then follows from Equation (2), where $t$ is the cosmic time corresponding to the observed redshift of each galaxy. Of course, the reliability of the results depends on the extent to which the actual SFH is well represented by a declining exponential.  It is worth recalling that the first applications of $\\tau$-models were to figure out the ages of local elliptical galaxies, and the typical result was that the age is of the order of one Hubble time, and $\\tau$ of the order of 1 Gyr or less (e.g., Bruzual 1983). This approach confirmed that the bulk of stars in local ellipticals are very old, hence formed within a short time interval compared to the present Hubble time. Later, the usage of $\\tau$-models was widely extended also to actively star forming galaxies at virtually all redshifts, to the extent that it became the default assumption in this kind of studies (e.g., Papovich et al. 2001; Shapley et al. 2005; Lee et al. 2009; Pozzetti et al. 2009; F\\\"oster-Schreiber et al.  2009; Wuyts et al. 2009b).  Such an assumed SFH may give reasonable results for local spirals, as their SF activity has been secularly declining for an order of one Hubble time (e.g., Kennicutt 1986), but we shall argue that it may be a rather poor representation of the SFH of high-$z$ galaxies, and may lead to quite unphysical results.  Cimatti et al. (2008) noted that the age of elliptical galaxies at $z\\sim 1.6$ turns out to be $\\sim 1$ Gyr both when using only the rest-frame UV part of the SED, and when using the whole optical-to-near-IR SED in conjunction with $\\tau$-models.  However, they also noted that the former ``age'' is actually the age of the population formed in the last significant episode of star formation, while the latter ``age'' corresponds to the time elapsed since the beginning of star formation. The near equality of these two ages suggests that the SFR peaked shortly before being quenched, rather than having peaked at an earlier time and having declined ever since. Moreover, using $\\tau$ models one implicitly assumes that galaxies are all caught at their minimum SFR, which is possibly justified for local ellipticals and spirals, but not necessarily for star-forming galaxies at high redshifts that may actually be at the peak of their SF activity.  Indeed, an integration of $dM_*/dt = <{\\rm SFR}>$ where $ <{\\rm SFR}>$ is given by Equation (1) shows that the SFR can {\\it increase} quasi-exponentially with time before the effect of the declining term $t^{-2.5}$ takes over, or star formation is suddenly quenched and the galaxy turns passive (Renzini 2009). Mass and SFR formally increase exponentially when SFR $\\propto\\sim M_*$, independent of time, as appears to be the case for $z\\gsim 2.5$ (Gonzales et al. 2010). Thus, the observations of both passive and star-forming galaxies at $1.4\\lapprox z\\lapprox 2.5$ suggest that the SFRs of these galaxies may well have increased with time, rather than decreased. For these reasons, in this paper we use both {\\it direct}-$\\tau$ models, with the SFR given by Equation (2), as well as {\\it inverted}-$\\tau$ models in which the SFR increases exponentially with time, i.e.: \\begin{equation} {\\rm SFR}=Ae^{+(t-t_\\circ)/\\tau}. \\end{equation}  Thus, direct-$\\tau$ models assume that galaxies are caught at their minimum SFR and had their maximum SFR at the beginning, whereas inverted-$\\tau$ models assume that galaxies are caught at their maximum SFR and had their minimum SFR at the beginning. These two extreme assumptions may to some extent {\\it bracket} the actual SFHs of real galaxies,  or at least of the majority of them which, because of the tight SFR$-M_*$ relation, must have a relatively smooth SFH. Here we explore which of the two assumptions gives the better fit to the SED of $z\\sim 2$ galaxies, and discuss the astrophysical plausibility of the relative results. Besides these exponential SFHs we also consider the case of constant SFRs. In principle, other, differently motivated SFHs could also be explored, but in this paper we restrict the comparison to these three simple options, with SFRs increasing with time, decreasing, or constant.  The paper is organised as follows. Section 2 provides information on the galaxy data base that is used for our SED fitting experiments. Section 3 illustrates the procedure of SED fitting and describes the models. Section 4 presents the results we obtain for the different star formation history templates and their comparisons. Section 5 is devoted to a general discussion and presents our conclusions.  Finally, we adopt a cosmology with $\\Omega_{\\Lambda}$, $\\Omega_{M}$ and \\\\$h~=H_{0}/(100~{\\rm km s^{-1} Mpc^{-1}})$ equal to 0.7, 0.3 and 0.75, respectively, for consistency with most previous works. The age of the best-fit model is required to be lower than the age of the Universe at the given spectroscopic redshift. The SFR and masses are always in $\\msun$/yr and $\\msun$ units, respectively. ",
        "conclusions": "We have tried different priors for the SFH of $z\\sim 2$ galaxies, and used a $\\chi^2$ best fit approach to derive the basic properties of such galaxies: SFR, stellar mass, reddening, age, etc. We do not expect that any of the simple mathematical forms adopted for the SFH (constant SFR, exponentially decreasing, exponentially increasing) is strictly followed by real galaxies. We expect, however, that the two opposite choices for the exponential case are sufficiently extreme to encompass the real behavior of galaxies, or at least those having experienced a quasi-steady SFR over most of their lifetime. We try various tests to distinguish which of the various options is the most acceptable one (or the least unacceptable) for giving the astrophysically most plausible values for the basic properties of star forming galaxies in the redshift range between $\\sim 1.4$ and $\\sim 2.5$.  We find that by leaving age as a free parameter the best fit procedure delivers implausibly short ages, no matter which SFH template is adopted. This is so because the SED is dominated by the stars that have formed most recently, and they outshine the older stellar populations that may inhabit these galaxies. Thus, introducing a prior on the beginning of star formation appears to be necessary.  Assuming that star formation started at high redshift (i.e., $z>>\\sim 2$, the precise value being almost irrelevant) provides a more credible framework, with models with exponentially increasing SFR (that we call inverted-$\\tau$ models) giving systematically higher SFRs for the $z\\sim 2$ galaxies in the test sample, and lower stellar masses, compared to models with exponentially decreasing SFR. These two systematic differences together imply a specific SFR that is a factor of $\\sim 5$ higher in inverted-$\\tau$ models, compared to direct-$\\tau$ models, when star formation is assumed to start at high redshift (e.g., $z=5$) in both cases.   \\begin{figure*} \\includegraphics[width=0.8\\textwidth]{Fig26.eps} \\caption{Left panel: The SFR from the full optical-to-near-IR SED fit derived from inverted-$\\tau$ models in which $\\tau$ is allowed to take very small values, compared to the SFR derived only from the UV part of the SED. Right panel: The frequency histogram of the corresponding values of $\\tau$. \\label{smalltau}} \\end{figure*}  \\begin{figure*} \\includegraphics[width=0.8\\textwidth]{Fig27.eps} \\caption{A comparison of the stellar masses derived for the GOODS galaxies from inverted-$\\tau$ models, with $\\zform=5$ and $\\tau\\ge 0.3$ Gyr,  with those derived from other adopted SFHs. Left panel: vs. inverted-$\\tau$ models with age $>\\tau=0.5$ Gyr; Central panel: vs. models with SFR=const. and age$>0.1$ Gyr; Right panel: vs. direct-$\\tau$ models and unconstrained age. \\label{masse}} \\end{figure*}  On the other hand, inverted-$\\tau$ models offer various notable advantages: 1) they indicate SFRs and extinctions in excellent agreement with those derived from the rest-frame UV part of the SED, which most directly relates to the ongoing star formation and reddening, 2) they fairly accurately recover the SFRs and masses of mock galaxies constructed from semi-analytic models, in which SFR are indeed secularly increasing in most cases, and 3) have systematically better reduced $\\chi^2$ values compared to models with exponentially decreasing SFRs, though one cannot consider them fully satisfactory. These advantages are somewhat reduced if the best fit procedure is allowed to pick very small values of $\\tau$, down to 0.05 Gyr, as in such case derived SFRs tend to be somewhat overestimated compared to those derived from the UV.  We have also explored star formation histories (such as constant SFR, or exponentially decreasing) in which age is a priori constrained to avoid very small values (e.g., age $>0.1$ Gyr, or age $>\\tau$, or fixed $\\tau$).  Adopting such priors, though it may seem quite artificial, gives more plausible solutions with respect to cases in which age is left fully free. This is due to the mentioned outshining effect of very young stellar populations over the older ones. Still, models with exponentially increasing SFRs tend to give systematically better results, as judged from the consistency with the UV derived SFRs.  Besides the above arguments, there are other reasons to prefer stellar population models constructed with exponentially increasing SFRs, including: \\par\\noindent 1) The almost linear relation between SFR and stellar mass (SFR $\\propto\\;\\sim M_*$), and the small scatter about it, that are empirically established for redshift $\\sim 2$ galaxies implies that they can indeed experience a quasi-exponential growth at these cosmic epochs. Later, some process may quench their star formation entirely and turn them into passively evolving {\\it ellipticals}, or keep growing until the secular decrease of the specific star formation rate (the $t^{-2.5}$ term in Equation (1)) takes over, and galaxies (such as spirals) continue forming stars at a slowly decreasing rate all the way to the present epoch (Renzini 2009).  \\par\\noindent 2) The direct observation of the evolution of the luminosity function in rest-frame UV (which is directly related to SFR) shows that the characteristic luminosity at 1600 \\AA\\ ($M^*_{1600}$) brightens by over one magnitude between $z=6$ and $z=3$ (Bouwens et al. 2007).  Since extinction is likely to increase with time following metal enrichment, the increase of the rest-frame UV luminosity can only be due to an increased SFR in individual galaxies.  The aim of this paper is to explore which assumptions and procedures are likely to give the most robust estimates of the basic stellar population parameters, when using data that extend from the rest-frame UV to the near-IR. We argue that ongoing SFRs are best derived at wavelengths below about one micron by using the UV part of the SED of galaxies, and inverted-$\\tau$ models give SFRs in excellent agreement with them. However, the fit to the full optical to near-infrared SED is required to derive the mass in stars. We show that the SFHs, with a prior on age or $\\tau$ to reduce the outshining effect, give results that agree quite well with each other, giving some confidence on the reliability of the derived masses. However, it is still possible that the mass may be somewhat underestimated for those galaxies in which a substantial fraction of the stellar mass was produced in a strong burst, early in their evolution.  \\begin{figure} \\includegraphics[width=0.49\\textwidth]{Fig28.eps} \\caption{Stellar masses for the $z\\sim2$~GOODS star-forming galaxies as derived from inverted-$\\tau$ models, with $\\zform\\gapprox5$ and $\\tau\\ge 0.3$~Gyr. \\label{massefinali}} \\end{figure}  At first sight the higher SFRs and lower stellar masses indicated by inverted-$\\tau$ models (compared to direct-$\\tau$ models, both with fixed formation redshift) may exacerbate an apparent mismatch between the {\\it cosmic} (volume averaged) SFR history ($\\dot{\\rho_*}(z)$) and the empirical stellar mass density ($\\rho_*(z)$), such that the integration of SFR$(z)$ tends to overproduce the stellar mass density at most/all redshifts (e.g., Wilkins, Trentham \\& Hopkins 2008). However, SFRs are more often derived from direct SFR indicators, such as the rest-frame UV, radio flux, H$\\alpha$ flux, or mid/far-infrared, rather than from SED fits. Therefore the 0.5 dex effect on the SFRs relative to the direct-$\\tau$ models should not affect this discrepancy. There may remain the 0.2 dex effect on the masses, but inverted-$\\tau$ models, by construction minimising star formation at early times, may systematically underestimate the stellar mass of galaxies, in particular if part of it was formed in early bursts.  Thus, with this caveat in mind, the use of inverted-$\\tau$ models with fixed formation redshift should not appreciably exacerbate the mismatch problem mentioned above. Instead, direct-$\\tau$ models with unconstrained age most certainly worsen this problem as they underestimate the stellar mass (see Figure~\\ref{masse}, right panel), and overestimate the SFR (see Figure~\\ref{sfrcomp}).  We conclude that the use of synthetic stellar populations with an exponentially declining SFR should be avoided in the case of star forming galaxies at redshift beyond $\\sim 1$. Exponentially increasing SFRs, with e-folding times of $\\sim 0.3-1$ Gyr, and a high starting redshift appear to provide astrophysically more plausible results, and while nature may not closely follow such a simple functional behavior, they should be preferred in deriving SFRs and stellar masses of high redshift galaxies. In any event, we believe that maximum likelihood parameters are not necessarily unbiased estimators of star formation rates and masses of high-redshift galaxies, but the broadest possible astrophysical context should be considered."
    },
    "1004/1004.1680_arXiv.txt": {
        "abstract": "We present magneto-hydrodynamic simulation results for heterogeneous systems. Heterogeneous architectures combine high floating point performance many-core units hosted in conventional server nodes.  Examples include Graphics Processing Units (GPU's) and Cell. They have potentially large gains in performance, at modest power and monetary cost.  We implemented a magneto-hydrodynamic (MHD) simulation code on a variety of  heterogeneous and multi-core  architectures --- multi-core x86, Cell, Nvidia and ATI GPU --- in different languages, FORTRAN, C, Cell, CUDA and OpenCL. We present initial performance results for these systems. To our knowledge, this is the widest comparison of heterogeneous systems for MHD simulations. We review the different challenges faced in each architecture, and potential bottlenecks. We conclude that substantial gains in performance over traditional systems are possible, and in particular that is possible to extract a greater percentage of peak theoretical performance from some systems when compared to x86 architectures. ",
        "introduction": "Magneto-hydrodynamics (MHD) studies the dynamics of magnetized conducting fluids. If there is no magnetic field present, the problem reduces to traditional fluid dynamics.   However, in most astrophysical settings for instance, the fluids are highly conductive, and observed to be magnetized. Electro-motive forces generated by magnetic fields will modify the flow, which will in turn affect the field.  As a result, one has to solve both the Euler equations and Maxwell's equations simultaneously.  The MHD equations are nonlinear, and cannot in general solved analytically. Thanks to the increasing power of computers, three dimensional simulations can be used to model these dynamics numerically. Numerical simulations are crucial both for understanding the theory of such fluids, and for use in directing real world experiments.  However, in order to achieve realistic parameter regimes to solve real world problems, numerical experiments must push the limits of computational hardware resources.  Increasingly, considerations of compute power per watt or per dollar mean that new architectures are being considered to perform these calculations.   In particular, we examine here heterogeneous systems, which consist of two different kinds of processors; one or more general-purpose conventional processors which control the overall computation, and specialized, usually multi-core, processor units to which the numerically intensive computing is offloaded  \\cite{shan2006heterogeneous}.  There are several heterogeneous platforms in use currently and we will focus on Cell/B.E. \\cite{ibmcell:website} and graphics processing units (GPU) in this paper, and use a multi-core x86 system for comparison.  Complicating comparison between these systems is that typically they are programmed using different platform-specific languages.    The Open Computing Language (OpenCL) \\cite{OpenCL:website}, is cross-platform application programming interface (API), which is designed for heterogeneous systems, including GPU and Cell, has been released by Khronos Group, and which ameliorates this problem to a degree; however we will see here that at the current time this does not solve the problem, and we must in general still use the platform-specific languages for performance.  We discuss our reference implementation of a solver for the MHD equations in \\S II; in \\S III we discuss our implementation on several architectures.    We summarize our results in \\S IV and discuss future work; in \\S V we conclude. ",
        "conclusions": "We presented magneto-hydrodynamics simulations on heterogeneous systems, e.g. Cell/B.E., Nvidia and ATI GPU.  These heterogeneous systems share a similar structure that they all have a control processor for mission management and many computing intensive processors for calculations.  Correspondingly, the memory system is also complicated, which is a challenge for programmers.  We present the results on different architectures for comparison; 10 times, 105 times and 68 times speed-up for Cell, Nvidia, and ATI GPU were achieved.  The CUDA on Nvidia GPU has the best performance on both code and fractional speed-up, and the ATI GPU improves with larger size simulation. The 2.4 fractional speed-up for CUDA on Nvidia GPU shows that a greater percentage of peak theoretical performance compared to x86 architecture was achieved.  These performance numbers were obtained with an algorithm which was directly translated from a CPU code.  Designing algorithms with heterogeneous architectures in mind may also improve performance."
    },
    "1004/1004.1694_arXiv.txt": {
        "abstract": "Using the South Pole Acoustic Test Setup (SPATS) and a retrievable transmitter deployed in holes drilled for the IceCube experiment, we have measured the attenuation of acoustic signals by South Pole ice at depths between 190 m and 500 m.  Three data sets, using different acoustic sources, have been analyzed and give consistent results.  The method with the smallest systematic uncertainties yields an amplitude attenuation coefficient $\\alpha=3.20 \\pm 0.57~\\rm km^{-1}$ between 10 and 30~kHz, considerably larger than previous theoretical estimates.  Expressed as an attenuation length, the analyses give a consistent result for $\\lambda \\equiv 1/\\alpha$ of $\\sim$300~m with 20\\,\\% uncertainty.  No significant depth or frequency dependence has been found. ",
        "introduction": "Experiments to study ultra-high-energy neutrinos have been the subject of increasing interest during recent years \\cite{AR05,AR06,AR08}. As observed by HiRes \\cite{HI08} and the Pierre Auger Observatory \\cite{AU10}, the charged cosmic ray flux decreases steeply above $10^{\\,19.5}$\\,eV. This is most likely due to the interactions of charged cosmic particles with the cosmic microwave background radiation, known as the GZK effect \\cite{GE66,ZK66}. The detection of neutrinos from this interaction would confirm this explanation.  Spectral, temporal, and directional distributions of such neutrinos, enabled by detecting a significant number of them, would address important questions of cosmology, astrophysics and particle physics \\cite{RI05}.  Estimates of the small flux of GZK neutrinos vary by an order of magnitude. The results from Engel, Seckel, and Stanev~\\cite{ESS} are often used as a standard for the discussion of possible detector scenarios. In all cases, detector effective volumes of 100~km$^3$ or larger are required.  Several past and current experimental projects seek to detect GZK neutrinos. The best limits currently come from searches for radio signals from neutrino interactions. Presently the most stringent flux limit is from the ANITA project \\cite{AN10}. However, all experiments searching for weak particle fluxes must contend with challenging systematic effects and background separation. This may be overcome in the future by using a hybrid detector. The radio technology could be complemented  by adding acoustic detectors, searching for the sound produced by the interactions of neutrinos above $10^{\\,18}$\\,eV \\cite{AS79,LE79}. Simulations of a radio-acoustic hybrid detector encompassing the optical IceCube neutrino observatory at the South Pole \\cite{IC06}, assuming an attenuation length on the order of a few kilometers as theorized in \\cite{PR06}, gave a predicted detection rate of 20 neutrino events per year, half of them detected by both radio and acoustic sensors~\\cite{BE05}.  Ice seems to be a favorable medium for the application of  the optical, radio and acoustic detection techniques. Ice properties have been measured for optical signals by the AMANDA experiment \\cite{AC06} and for radio waves by the RICE experimental program \\cite{RAT,RI06}. As far as the propagation of acoustic signals is concerned, theoretical estimates~\\cite{PR06} have indicated low absorption and scattering corresponding to an attenuation length greater than 1~km.  To test the theoretical estimates, the South Pole Acoustic Test Setup (SPATS) was deployed in January and December 2007 in the shallowest 500 m of IceCube holes at the South Pole. SPATS is used to study the following four properties: \\begin{itemize} \\item the speed of sound as a function of depth in the ice, important for the expected neutrino signal strength, event localization and reconstruction; \\item\tthe acoustic noise level at the South Pole, determining the energy threshold of a future neutrino detector; \\item\tthe rate and nature of transient acoustic signals which could mimic neutrino interactions, presenting a serious background source; \\item the attenuation of acoustic signals by the ice, which determines the necessary density of acoustic sensors in the ice for a reasonable detection efficiency for neutrino interactions. \\end{itemize} The sound speed has been measured in \\cite{SPATS2}. Noise and transients are currently under study. Our measurement of attenuation by the ice is reported here. ",
        "conclusions": "\\label{secDiscussion}  The results from the various analyses described above are summarized in Table~\\ref{summaryTable}.  \\begin{table}[hbtp] \\caption{\\label{summaryTable} Summary of results from various attenuation analyses.} \\begin{center} \\begin{tabular}{l c c c} Analysis & Measurements & $\\alpha$ [km$^{-1}$]  \\\\ \\hline \\hline Pinger (time dom.)     & 48     & 3.20 $\\pm$ 0.57    \\\\ \\hline Pinger (freq. dom.)    & 39     & 3.75 $\\pm$ 0.61    \\\\ \\hline Inter-string direct    & 12     & 3.16 $\\pm$ 1.05    \\\\ \\hline Inter-string ratio       & 1    & 4.77 $\\pm$ 0.67    \\\\ \\hline Transients               & 13     & 3.64 $\\pm$ 0.29   \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  All of the analyses yield consistent results. However, we cannot simply average them, due to the quite different systematic effects (as well as correlations among them) which affect each of them as discussed in detail in previous sections. In Table~\\ref{systematicsTable} the different effects are summarized and attributed to the individual studies they affect. The pinger approach uses the same sensor and the same transmitter at the same zenith angle and nearly the same azimuthal angle at different distances, thereby mitigating all effects connected with unknown sensitivities of emitter and receiver. For this reason we believe it provides the most reliable result. Relying on the same data set, the pinger analyses in time and frequency domain are not independent but cross check one another using different signal and background averaging and subtraction methods.  Because there is a higher number of measurements passing selection criteria in the time domain analysis than in the frequency domain analysis (see Table 2), we quote this result as our best estimate of the attenuation coefficient:  \\begin{equation} \\left< \\alpha \\right> = 3.20 \\pm 0.57~{\\rm km^{-1} }. \\end{equation}  \\noindent which expressed as an attenuation length is:  \\begin{equation} \\left< \\lambda \\right> = 312^{+68}_{-47}~{\\rm m}. \\end{equation}   \\noindent The results of all analyses are consistent with an attenuation length of $\\sim$300 m $\\pm$ 20\\,\\%.  Up to 30 kHz no strong frequency dependence of $\\alpha$ has been found. There are also no indications of depth dependence of the attenuation up to 500 m depth.   \\begin{table*} \\resizebox{0.95\\textwidth}{!}{ \\begin{minipage}[h]{0.95\\linewidth} \\caption{\\label{systematicsTable}Systematic effects present in individual analyses.} \\begin{center} \\begin{tabular}{ p{.6\\textwidth}  p{.1\\textwidth}  p{.1\\textwidth} p{.1\\textwidth}  p{.1\\textwidth}} Systematic Effect \t\t\t\t& pinger\t& inter-string single level& inter-string ratio & transients\\\\ \\hline \\hline Channel-to-channel sensitivity variation\t& no \t& yes \t& no & minimal \\\\ Azimuthal sensitivity variation\t\t& minimal  \t& yes \t& yes & minimal \\\\ Polar sensitivity variation \t\t& minimal \t& no \t& yes & minimal \\\\ \\hline Channel-to-channel transmittivity variation\t& no \t& no \t& no & no \\\\ Azimuthal transmittivity variation\t& no \t& yes \t& yes & no \\\\ Polar transmittivity variation \t\t& no \t& no\t& yes & no \\\\ \\hline \\end{tabular} \\end{center} \\end{minipage} } \\end{table*}  Our measured value for the attenuation coefficient is an order of magnitude larger than expected.  In~\\cite{PR06}, it was estimated that South Pole ice grains are sufficiently small that Rayleigh scattering is negligible and attenuation is dominated by absorption due to proton reorientation.  However, new data~\\cite{ME09} from the SPRESSO site near the South Pole indicate that the ice grains are larger than previously estimated.  Because Rayleigh scattering increases with the cubic power of the grain length, it could be that the attenuation we have measured is dominated by scattering.  Another possible mechanism of absorption, not previously considered, is internal friction at linear crystallographic dislocations.  The weak frequency dependence observed in our measurements below 30 kHz disfavors the Rayleigh scattering hypothesis.  New pinger measurements have been taken in the 2009-2010 South Pole summer in order to clarify the situation. A modified pinger, emitting a sequence of acoustic pulses at three different frequencies up to 60 kHz, was deployed up to 1000 m depth in three additional bore holes. Analysis of these data is currently underway, and should allow for a more conclusive study of both the frequency and depth dependence of acoustic attenuation in South Pole ice.  Given some of the inherent advantages of the acoustic technique relative to the radio technique (such better shielding from anthropogenic surface backgrounds), this study was undertaken primarily to determine whether the acoustic method could be a basic ingredient of a 100 km$^3$ hybrid detector for ultra-high energy neutrino detection at the South Pole. The design chosen in \\cite{BE05} with horizontal string distances of order 1~km will not reach the necessary sensitivity, given the 300~m attenuation length reported here. New geometries are under study to assess whether detectors with closer string spacing but larger area at smaller depth could be an alternative solution.  Measurements are also in preparation to measure the absolute noise level of the Gaussian noise, which is the most important remaining ingredient to determine the feasibility of the acoustic method for detecting neutrinos in South Pole ice."
    },
    "1004/1004.1141_arXiv.txt": {
        "abstract": "We analyze radio continuum and line observations from the archives of the Very Large Array, as well as X-ray observations from the \\emph{Chandra} archive of the region of massive star formation W75N. Five radio continuum sources are detected: VLA 1, VLA 2, VLA 3, Bc, and VLA 4. VLA 3 appears to be a radio jet; we detect J=1-0, v=0 SiO emission towards it, probably tracing the inner parts of a molecular outflow. The radio continuum source Bc, previously believed to be tracing an independent star, is found to exhibit important changes in total flux density, morphology, and position. These results suggest that source Bc is actually a radio Herbig-Haro object, one of the brightest known, powered by the VLA~3 jet source. VLA 4 is a new radio continuum component, located a few arcsec to the south of the group of previously known radio sources. Strong and broad (1,1) and (2,2) ammonia emission is detected from the region containing the radio sources VLA~1, VLA~2, and VLA~3. Finally, the 2-10 keV emission seen in the \\emph{Chandra}/ACIS image shows two regions that could be the termination shocks of the outflows from the multiple sources observed in W75N. ",
        "introduction": "The compact, thermal (i. e. free-free) radio sources found at centimeter wavelengths in regions of star formation can have different natures. Some of them are ultracompact or hypercompact HII regions, photoionized by an embedded hot luminous star. Other sources are clumps of gas or even circumstellar disks that are being externally ionized by a nearby star, such as the Orion proplyds (e. g. O'Dell \\& Wong 1996; Zapata et al. 2004) and the bright-rimmed clouds (e.g. Carrasco-Gonzalez et al. 2006). An additional class of sources is formed by the thermal jets, collimated outflows that emanate from young stars and whose ionization is most probably maintained by the interaction of the moving gas with the surrounding medium (Eisloffel et al. 2000). Finally, there are also some examples of ``radio Herbig-Haro'' objects, obscured knots of gas that are being collisionally ionized by the shock produced by a collimated outflow (e. g. Curiel et al. 1993). There are also compact non-thermal sources, of which the most common are the young low-mass stars with active magnetospheres that produce gyrosynchrotron emission (e.g. Feigelson \\& Montmerle 1999). There are also emission regions where fast shocks may produce synchrotron emission (e.g. Henriksen et al. 1991; Garay et al. 1996).  To advance in the understanding of the nature of these compact radio sources it is necessary to have good quality data that allows the observer to establish the angular size, the morphology, the spectral index, the time variability, the polarization, and the presence of proper motions in them. Only a handful of the known star-forming regions have been studied carefully enough to clearly establish the nature of its compact radio sources.  The massive star-forming region W75N is part of the Cygnus X complex of dense molecular clouds. Its distance is estimated to be 2 kpc (Fischer et al. 1985). Haschick et al. (1981) detected at 6 cm a source that was interpreted as an ultracompact HII (UC HII) region named W75N(B), that later was resolved into three small diameter radio continuum sources: Ba, Bb and Bc in the high angular resolution ($0\\rlap.{''}5$) observations of Hunter et al. (1994). Two of these sources, Ba and Bb, were also detected at 1.3 cm (with a $0\\rlap.{''}1$ resolution) by Torrelles et al. (1997) (who named them as VLA 1 and VLA 3) along with another fainter and more compact source, located in between VLA 1 and VLA 3, and named as VLA 2. These compact sources were believed to be UC HII regions in an early phase, based on the presence of H$_2$O and/or OH masers (Baart et al. 1986, Hunter et al. 1994, Torrelles et al. 1997) in its close vicinity.  At radio wavelengths, VLA 1 has a structure elongated at a PA of $\\sim 43^\\circ$ (Torrelles et al. 1997), approximately in the direction of the observed high-velocity three pc-scale bipolar outflow (Hunter et al. 1994; Shepherd et al. 2003). However, Shepherd (2001) had pointed that VLA 1 does not appear to contribute significantly to the energetics of the large-scale outflow. VLA 2 shows a symmetric roundish shape. A detailed Very Long Baseline Interferometry (VLBI) study of the water maser cluster morphology and kinematics in the vicinity of these two radio sources by Torrelles et al. (2003) showed remarkably different outflow ejection geometries among them (VLA 1 has a collimated, jetlike outflow, while VLA 2 has a shell outflow expanding in multiple directions). This is a surprising result given that both objects are separated by only 1400 AU and share the same environment. VLA 3 was found to have deconvolved dimensions of $0\\rlap.{''}09 \\times 0\\rlap.{''}04$ $(PA = 152^\\circ)$ at 1.3 cm by Torrelles et al. (1997). This source is also detected at mid-infrared wavelengths, with a bolometric luminosity of ~ 750~L$_{\\odot}$ (Persi, Tapia, \\& Smith 2006). The physical parameters of the ionized gas in VLA 1, VLA 2, VLA 3, and Bc, as well as the more evolved W75N(A) are all consistent with each being powered by a ZAMS star of spectral type between B0 and B2 (Shepherd et al. 2004). The present evidence is consistent with an independent exciting source of this spectral type for each radio continuum component, but the results presented in this paper suggest a likely alternative interpretation.  In this paper we analyze Very Large Array (VLA) and \\emph{Chandra}  archive data of the region, mostly searching for variability and proper  motions that could help in a better understanding of the radio sources. ",
        "conclusions": "We have analyzed VLA and \\emph{Chandra} archive data of the W75N region. In the radio continuum, we detect five sources: VLA1, VLA2, VLA3, Bc, and VLA4, from north to south respectively. Our main conclusions follow:  1) We detect important changes in total flux density, morphology, and position in the source Bc, suggesting that it is not tracing an independent star but actually is a radio HH object powered by VLA~3. Its average velocity in the plane of the sky is 220$\\pm$70 km s$^{-1}$. If our interpretation is correct, this is one of the brightest radio HH objects known.  2) We detect J = 1--0; v = 0 emission of SiO centered on VLA 3, probably tracing the inner parts of an outflow associated with this radio jet. Our results strongly support the identification of VLA 3 as a radio jet.  3) The large scale ammonia emission shows that the W75N region contains two filamentary molecular clouds with different radial velocities. The star formation is taking place at the intersection of these two filaments, which is also the hottest region. This suggests that the star formation could be triggered by the collision of the two filamentary clouds.  4) Strong and broad (1,1) and (2,2) ammonia emission is detected from the region containing the radio sources VLA~1, VLA~2, and VLA~3.  5) We detected a new radio continuum component, VLA 4, a few arcsec to the south of the group of previously known sources.  6) Two sources are detected in the 2-10 keV band \\emph{Chandra}/ACIS image. These two sources could be tracing the termination shocks of outflows in the region, but better data are needed to understand the nature of the X-ray emission."
    },
    "1004/1004.3058.txt": {
        "abstract": "We present the compact radio structure of three radio-loud narrow line Seyfert 1 galaxies from VLBA archive data at 2.3, 5 and 8.4 GHz. In RXS J16290+4007, the radio structure is mostly unresolved. The combination of compact radio structure, high brightness temperature and inverted spectrum between simultaneous 2.3 and 8.4 GHz, strongly favors jet relativistic beaming. Combining with the VLBI data at 1.6 and 8.4 GHz from literatures, we argued that RXS J16333+4718 may also harbor a relativistic jet, with resolved core-jet structure in 5 GHz. B3 1702+457 is clearly resolved with well defined jet component. The overall radio steep spectrum indicates that B3 1702+457 is likely a source optically defined as NLS1 with radio definition of compact steep spectrum sources. From these three sources, we found that radio loud NLS1s can be either intrinsically radio loud (e.g. B3 1702+457), or apparently radio loud due to jet beaming effect (e.g. RXS J16290+4007 and RXS J16333+4718). %At 5 GHz radio images, the jet component of B3 %1702+457 apparently moves toward the core with $\\sim 3.7$ mas during %about two month. Several possibilities for this jet inward motion %are briefly discussed. ",
        "introduction": "While both permitted and forbidden optical emission lines are present in narrow line Seyfert 1 galaxies (NLS1), their broad Balmer lines are narrower than that of normal broad line Seyfert 1 galaxies with the FWHM of the $\\rm H\\beta$ line less than 2000 $\\rm km ~s^{-1}$ (Osterbrock \\& Pogge 1985; Goodrich 1989). Moreover, NLS1 galaxies also exhibit other extreme observational properties, such as the relatively weak forbidden-line emission, i.e., $\\rm [O III] 5007/H\\beta < 3$ (Goodrich 1989), strong permitted optical/UV Fe II emission lines (Boroson \\& Green 1992; Grupe et al. 1999; V\\'{e}ron-Cetty et al. 2001), steep soft X-ray spectra (Wang et al. 1996; Boller et al. 1996; Grupe et al. 1998), and rapid X-ray variability (Leighly 1999; Komossa \\& Meerschweinchen 2000). NLS1 galaxies are thought to be young AGNs with relatively small black hole masses and high accretion rate (indicated by Eddington ratios: the bolometric-to-Eddington luminosity ratio $L_{\\rm bol}/L_{\\rm Edd}\\sim1$) (e.g., Boroson 2002; Collin \\& Kawaguchi 2004). %They %were found to locate at one extreme end of eigenvector 1 of the %correlation matrix which is believed to be driven primarily by the %Eddington ratio (Boroson \\& Green 1992; Sulentic et al. 2000).  %\\textbf{the radio-loud NLS1 populations:} Conventionally, NLS1 galaxies were thought to be radio quiet. With the discovery of radio loud NLS1s, it is realized that NLS1s simply have low probability to be radio loud, instead of completely radio quiet \\cite[e.g.][]{kom06}. Compared to about 10\\%-15\\% for normal broad line AGNs and quasars (e.g., Ivezi\\'{c} et al. 2002), only 7\\% of all NLS1 galaxies are radio loud (Komossa et al. 2006; Zhou et al. 2006), while only $\\sim2.5\\%$ of the NLS1s are `very' radio loud (radio loudness $R>100$). As shown in \\cite{kom06}, the radio-loud NLS1 galaxies are generally compact, steep spectrum sources in the radio regime, therefore, they are likely associated with compact steep-spectrum (CSS) radio sources. In contrast, observational evidences have shown that several of the radio- loudest NLS1 galaxies display blazar characteristics and harbor relativistic jets \\citep{doi07,zho07,yua08}. Yuan et al. (2008) presented a comprehensive study of a sample of 23 genuine radio-loud NLS1 AGNs with radio loudness larger than 100. The radio sources of the sample are ubiquitously compact, unresolved on scales of several arcseconds. Some of these objects show interesting radio to X-ray properties that are unusual to most of the previously known radio-loud NLS1 AGNs, but are reminiscent of blazars, including flat radio spectra, large-amplitude flux and spectral variability, compact VLBI cores, very high variability brightness temperatures, enhanced optical continuum emission, flat X-ray spectra, and blazar-like spectral energy distributions (SEDs) \\cite[also see][]{fos09a}. Despite that the number of radio loud NLS1s is rapidly growing \\citep{zho06}, mechanisms driving the NLS1 radio properties is still not clear, such as, accretion rate, black hole spin, host galaxy properties and merger history. %We interpret them %as signatures of the postulated blazar nature of some of the %objects, which likely possess at least moderately relativistic jets. %Intrinsically, some of them have relatively low radio power and %would have been classified as radio-intermediate AGNs. The black %hole masses are estimated to be within $10^6-10^8 M_{\\odot}$, and %the Eddington ratios close to unity, as in normal NLS1 AGNs. The %results imply that radio-loud AGNs may be powered by black holes of %moderate masses ($\\sim10^6-10^7 M_{\\odot}$) accreting at high rates. %Some of the objects, despite having strong emission lines, resemble %high-energy-peaked BL Lac objects (HBLs) in their SEDs. Given the %peculiarities of blazar-like NLS1 AGNs, questions arise as to %whether they are plain downsizing extensions of normal radio-loud %AGNs, or form a previously unrecognized population. %Estimates of black hole masses %show that the radio-louds are at the upper range of NLS1 black hole %masses, and are located in a previously scarcely populated area of %MBH-R diagrams. Optical properties of these radio-loud NLS1 galaxies %are similar to the NLS1 population as a whole (Komossa et al. 2006; %the same holds for a radio-selected NLS1 sample of Whalen et al. %2006). A surprising exception is that radio-loud NLS1 galaxies %preferentially show moderate-strong Fe II emission when compared to %radio-quiet NLS1s (Komossa et al. 2006).  %\\textbf{the VLBI observations of radio-loud NLS1s and their possible %nature - Doppler boosted/intrinsically strong radio sources %\\citep{doi09}, \\cite{doi06,doi07} -- VLBA for J0948+0022:} %Using information on the radio flux %variations, very high radio brightness temperatures ($\\gtrsim %10^{13}$ K) were inferred for NLS1s PMN J0948+0022 (Zhou et al. %2003) and 0846+513 (Zhou et al. 2005), arguing for relativistic %beaming (see also Wang et al. 2006). %These enigmatic objects present a challenge to current models of %both NLS1 galaxies and blazars, and render a unique opportunity to %study jet formation in black hole systems accreting at high rates.  %Either radio loud NLS1s are found to be associated with CSS sources %\\cite[][]{kom06}, or blazars-like sources with relativistic beaming %(e.g. PMN J0948+0022 - Zhou et al. 2003; 0846+513 - Zhou et al. %2005; Yuan et al. 2008), they are mostly inferred indirectly. Although a relativistic jet has been proposed to exist in radio-loud NLS1s \\cite[e.g.][]{kom06,yua08,fos09a}, only until recently its existence can be strongly proved through the $\\gamma-$ray emission successfully detected in several radio-loud NLS1s \\citep{abd09a,abd09b,abd09c,fos09b}. Especially from these observations, radio-loud NLS1s were claimed to be a third class of $\\gamma-$ray AGNs, besides blazars and radio galaxies \\cite[e.g.][]{abd09c}. To explore the jet properties in radio-loud NLS1s, the global e-VLBI observations have been made for the first $\\gamma-$ray detected radio-loud NLS1 PMN J0948+0022 \\citep{abd09b}, and more observations are proposed\\footnote{http://www.evlbi.org/gallery/images.html.}. However in general, still not much is known for the VLBI compact radio structure of radio-loud NLS1s, from which the jet properties can be directly explored, such as jet orientation, and jet beaming effect. Recently, Doi et al. (2006, 2007) performed high-resolution VLBI observations for several radio-loud NLS1 AGNs and found that they are unresolved with milliarcsecond resolutions, setting direct lower limits on the brightness temperatures in the range of $10^{7}-10^{9}$ K. They also found that inverted radio spectra are common in the radio-loudest objects. The authors suggested that Doppler beaming, presumably resulting from highly relativistic jets, can explain the observations naturally. Most recently, \\cite{doi09} found significant pc-scale structures with high brightness temperatures for five radio-loud NLS1s from the VLBA high-resolution observations, indicating nonthermal jets in these sources. Moreover, the authors claimed that the central engines of some radio loud NLS1s can generate mildly- or highly-relativistic jets, which make them apparently radio-loud. However, the nuclei of some NLS1s have the ability to generate jets significantly strong enough to make them intrinsically radio-loud.    In this paper, we present the high resolution radio structure of three radio-loud NLS1s from the unpublished VLBA archive data. The observations and data reduction is shown in Section 2, and the results is given in Section 3. The last section is dedicated to discussions. Throughout the paper, we assume a cosmology with $H_{0}=71~\\rm km~s^{-1}~Mpc^{-1}$, $\\Omega_{\\rm M}=0.27$, and $\\Omega_{\\Lambda}=0.73$ \\citep{spe03}. The spectral indices $\\alpha$ is defined as $f_{\\nu}\\propto\\nu^{-\\alpha}$, in which $f_{\\nu}$ is the flux density at frequency $\\nu$. ",
        "conclusions": "\\subsection{Brightness temperature}  From the high-resolution VLBA images, the brightness temperature of radio core $T_{\\rm B}$ in the rest frame can be estimated with \\citep{ghi93}  \\begin{equation} T_{\\rm B}=\\frac{S_{\\nu}\\lambda^2}{2k\\Omega_{\\rm s}}=1.77\\times10^{12}(1+z)(\\frac{S_{\\nu}}{\\rm Jy})(\\frac{\\nu}{\\rm GHz})^{-2}(\\frac{\\theta_{d}}{\\rm mas})^{-2} \\end{equation} in which $z$ is source redshift, $S_{\\nu}$ is core flux density at frequency $\\nu$, and $\\theta_{\\rm d}$ is source angular diameter $\\theta_{\\rm d}=\\sqrt{\\rm ab}$ with a and b being major and minor axis, respectively. The intrinsic brightness temperature $T_{\\rm B}^{'}$ can be related with $T_{\\rm B}$ by  \\be T_{\\rm B}^{'}=T_{\\rm B}/\\delta \\label{tb} \\ee in which $\\delta$ is Doppler factor.  The brightness temperature can also be measured from variability, which is correspondingly called variability brightness temperature $T_{\\rm B,var}$ and can be estimated as \\citep{yua08} \\be T_{\\rm B,var}\\gtrsim \\frac{\\Delta P_{\\nu e}}{2\\pi^{2}k\\nu^{2}(\\Delta t)^{2}}=\\frac{2D_{\\rm L}^{2}\\Delta S_{\\nu}}{(1+z)\\pi k\\nu^{2}(\\Delta t)^{2}} \\ee where $k$ is the Boltzmann constant, $D_{\\rm L}$ the luminosity distance, $\\Delta S_{\\nu}$ the variable portion of the observed flux density, $\\Delta P_{\\nu e}$ the corresponding radio power at the emission frequency in the source rest frame, $\\nu$ the observing frequency, and $\\Delta t$ the time span in the observer's frame. The intrinsic brightness temperature $T_{\\rm B}^{'}$ can be related with $T_{\\rm B,var}$ by  \\be T_{\\rm B}^{'}=T_{\\rm B,var}/\\delta^{3} \\label{tbvar} \\ee where $\\delta$ is Doppler factor.  Normally, the upper limit of physically realistic brightness temperature of nonthermal radio emission can be taken as the equipartition brightness temperature $T_{\\rm in}=5\\times10^{10}$ K \\citep{rea94}, or the inverse Compton catastrophic brightness temperature $T_{\\rm in}\\sim10^{12}$ K (Kellermann \\& Pauliny-Toth 1969, but see Singal 2009 for other interpretations). Consequently, the Doppler factor can be constrained either by equation (\\ref{tb}) or equation (\\ref{tbvar}). However, it should be noted that the Doppler factor constrained in these ways is generally smaller than that from gamma rays \\citep{abd09c}. This is due to the fact that gamma rays are produced in very compact regions, whose compactness is required for gamma rays to escape, but this in turn means that these compact blobs are also optically thick to radio frequencies (synchrotron self-absorption). As the blob is sufficiently expanded to be optically thin for radio emission, it is likely that it has decelerated \\cite[e.g.][]{bla79}.  \\subsection{Individual objects}  %$\\it RXS~ J08066+7248$ --- The radio structure is shown in Fig. %\\ref{fig1}. The source is unresolved in 1.6 GHz in the beam size of %$\\sim 4\\times10$ (mas). However, the source has a steep spectrum %between 1.4, 1.6, 5 and 8.4 GHz. Therefore, it is resemble to %compact spectrum sources, which is thought to be young radio source. %The X-ray photon index is $\\Gamma\\sim2.3$ (Xu et al. 2003). %\\\\ %\\\\ $\\it RXS~ J16290+4007$ --- We show the radio structure of this source in Fig. \\ref{fig2}, in which the source is compact and unresolved in all images, except that it is slightly resolved with an eastern component at 5 GHz on December 2005. The simultaneous observations at 2.3 and 8.4 GHz show an inverted spectrum $\\alpha=-0.10$, which is consistent with the inverted spectrum between FIRST 1.4 and GB6 5 GHz $\\alpha=-0.28$ \\cite[see Table \\ref{tab1}, see also][]{zho02}. The brightness temperature at all wavebands are high $T_{\\rm b}\\gtrsim 10^{11}~\\rm K$, especially it reaches $T_{\\rm b}= 10^{12.4}~\\rm K$ at 8.4 GHz. The inverted spectrum, compact structure and high brightness temperature all together suggest a Doppler beaming effect in this source. This is further supported by the flux variations. Although the variation is not evident during two epoch 5 GHz observations (see Table \\ref{tab3}), large variations are found when comparing with the previous data at 5 GHz \\cite[22 mJy,][]{ver01} and 8.4 GHz \\cite[JVN 26.3 mJy,][]{doi07}. At 8.4 GHz, the variability brightness temperature is quite high $T_{\\rm b,var}=10^{12.2}$ K (see Table \\ref{tab3}). The high $T_{\\rm B}$ values are commonly explained as emission originating from relativistic jets \\citep{bla79}. Conservatively taking the inverse Compton limit $10^{12}~ \\rm K$, the minimum Doppler factor can be estimated as $\\delta_{\\rm min}=T_{\\rm b}/10^{12}$ or $\\delta_{\\rm min}=(T_{\\rm b,var}/10^{12})^{1/3}$. This results in $\\delta_{\\rm min}=1.2-2.3$, which can be larger if the equipartition brightness temperature $5\\times10^{10}$ K is used as limit \\citep{rea94}. To conclude, this source can be a blazar-like NLS1s with jet pointing towards us with small viewing angle \\cite[see also][]{doi07,yua08,doi09}, which is consistent with the results presented in \\cite{pad02}.  Tentatively putting our data at 2.3, 5 and 8.4 GHz together, we find the radio spectrum are resemble to that of High Frequency Peakers (HFPs), i.e. compact objects with a simple convex radio spectrum turning over at frequencies well above a few GHz, which likely represent the earliest stage in individual radio source evolution \\cite[e.g.][]{dal00,ori09}. While this should be checked with further simultaneous multiwaveband observations, as a matter of fact, its spectral shape above 5 GHz is inverted, as measured by simultaneous multiwavelength observations with Effelsberg 100-m telescope by \\cite{neu94}. Alternatively, it is likely that the beaming effect boosted the radio emission as well as the peak frequency of the synchrotron self-absorption spectra \\cite[e.g.][]{doi07}. As pointed out by \\cite{ori09}, HFPs are selected on the basis of both the convex shape of their spectra and the high frequency at which the spectral peak occurs. In this way, it is possible that boosted objects matching the selection criteria contaminate the HFPs sample. Indeed, from the analysis of the radio properties, \\cite{ori09} found that HFPs galaxies are not variable, are unpolarized and exhibit a Double/Triple structure, typical of young radio sources. On the other hand, the majority of HFPs quasars are strongly variable, polarized and with Core-Jet structure, as expected in beamed objects. Therefore, RXS J16290+4007 is likely a boosted object rather than a genuine HFPs young radio sources. \\\\ \\\\ $\\it RXS~ J16333+4718$ --- This source is unresolved at JVN 8.4 GHz milliarcsecond resolution (Doi et al. 2007), and at 1.6 VLBA GHz \\citep{doi09}. Its spectral shape above 5 GHz is inverted, as measured by simultaneous multiwavelength observations by Neumann et al. (1994). We present 5 GHz radio structure in Fig. \\ref{fig3}, at which it is slightly resolved into core-jet structure. Combining JVN 8.4 GHz \\citep{doi07} and VLBA 1.6 GHz data \\citep{doi09} with our 5 GHz data, we find that the core spectral index is flat between 1.6 and 5 GHz $\\alpha=0.30$, which is consistent with that of \\cite{zho02} (see Table \\ref{tab1}). The high brightness temperature at 5 GHz $T_{\\rm b}=10^{11.3}$ K in combination with the flat spectrum indicate that an at least mild relativistic jets is likely responsible for the radio emission and structure. However, the spectra is steep between 1.6 and 8.4 GHz $\\alpha=0.73$ \\citep{doi09}, and between 5 and 8.4 GHz $\\alpha=1.73$. The steep spectrum and compact radio structure make this source resemble to CSS source. This however is not conclusive, since the observations were not performed simultaneously. Indeed, the variation can be clearly seen when comparing GB6 with VLBA 5 GHz data. Therefore, further simultaneous multi-epoch VLBI observations are needed to investigate the spectral index of this source, then to explore the source nature. %J1633+4718 (RX J16333+4719) The object was identified with a NLS1 %galaxy in an optical identification of ROSAT All-Sky Survey (RASS) %sources (Moran et al. 1996; Wisotzki \\& Bade 1997). The AGN (see %Fig. 1 for its SDSS spectrum) is actually one of the two nuclei, %separated by 400, residing in an interacting/merger system. The %other nucleus has a starburst spectrum (Bade et al. 1995).  The %system was noticed to be associated with the IRAS source IRAS %16319+4725 (R.A. = 16h33m23.03s, decl. = +47 19003.500), whose %positional uncertainty\u00a1\u00aamuch larger than the separation of the two %optical nuclei\u00a1\u00aaencompasses both of the nuclei (Boller et al. 1992; %Moran et al. 1996; Condon et al. 1998a). The [O iii] doublet shows a %noticeable blue wing in its SDSS spectrum. \\\\ \\\\ $\\it B3~ 1702+457$ --- This source was classified as Compact Steep Spectrum (CSS) sources with turnover frequency $\\nu_{\\rm peak}<150$ MHz in the sample of Compact Radio sources at Low Redshift (CORALZ) \\citep{sne04}. The VLBA 5 GHz radio structure is shown in Fig. \\ref{fig4}, in which the source is resolved into core-jet structure with weak radio jet components clearly seen in two images. The spectral index between FIRST 1.4 GHz and GB6 5 GHz show a steep spectrum with $\\alpha=1.23$ (see Table \\ref{tab1}). % steep between GB6 5 and JVN 8.4 GHz $\\alpha=0.56$. Consistently, the VLBA 1.6 GHz \\citep{doi09} and JVN 8.4 GHz \\citep{doi07} show a steep index $\\alpha=0.84$, which is also present between VLBA 5 GHz and JVN 8.4 GHz $\\alpha=2.12$ (see Tables \\ref{tab1} and \\ref{tab3}). In contrast, the spectral index of the prominent radio component become flat $\\alpha\\sim0.26$ with VLBA 1.6 and 5 GHz data, from which the radio core can be identified. However, VLBA 1.6 GHz and 5 GHz were not obtained simultaneously, therefore, the steep spectrum of 1.6 GHz and 5 GHz can not be completely excluded due to the variability. Indeed, the flux variation is apparent at 5 GHz. The GB6 flux density is 24.7 mJy (see Table \\ref{tab1}), whereas the flux density of VLBA core is 56.8 mJy on June 2000 (see Table \\ref{tab3}). It thus is possible that the VLBA 5 GHz data were taken when the source was at high state, while 1.6 GHz data was at relatively low state, which results in a flat spectrum between VLBA 1.6 GHz and 5 GHz. The brightness temperature of this source does not show a severe beaming effect, which is not inconsistent with the steep spectrum nature of this source. In order to investigate the genuine spectral index, it is no doubt that the simultaneous observations are required to explore the nature of the radio structure. In addition, the flux monitoring also enable us to explore the nature of this source since genuine CSS sources usually do not show voilent variations \\citep{ode98}.  %Intriguingly, the weastern jet component obviously moves towards %core in two 5 GHz images with time separation of about two months. %This inward motion has been found in various works %\\cite[e.g.][]{kel04,pin07,lis09}. While the position angle changes %about $10^{\\circ}$, the distance of radio jet components to core %decreased by about 3.7 mas in 66 days (i.e. Jun. 16, 2000 - Aug. 21, %2000), from which the proper motion can be tentatively estimated %$\\sim-80c$ with $c$ being the speed of light. However, only two %epoch observations preclude us to investigate the kinematics of jet %component in details. The multi-epoch VLBI observations are needed %to confirm the inward motion or trace the motion of jet components.  %This high apparent speed implies a %large Lorentz factor, and then a high Doppler factor, which however %is not seen since the brightness temperature is not very high %$\\sim10^{11}$ K. Moreover, it seems that the peak apparent speeds in %blazar jets are about $30c$ (Piner et al. 2006), although the %apparent speeds in excess of $40c$ were reported in some cases (An %et al. 2005; Jorstad et al. 2005). If the components is not the same %one, however, one should explain the disappearance of previous %component and appearance of the new component. Alternatively, it can %be a projection effect.  %The X-ray spectral was investigated by Vaughan et al. (1999), and %the photon index from the power-law fit is $\\Gamma\\sim2.2$.  %% This section contains more display math examples, including unnumbered %% equations (displaymath environment). The last paragraph includes some %% examples of in-line math featuring a couple of the AASTeX symbol macros."
    },
    "1004/1004.0001_arXiv.txt": {
        "abstract": "As some of the first known objects to exist in the Universe, Lyman alpha emitting galaxies (LAEs) naturally draw a lot of interest.  First discovered over a decade ago, they have allowed us to probe the early Universe, as their strong emission line compensates for their faint continuum light.  While initially thought to be indicative of the first galaxies forming in the Universe, recent studies have shown them to be increasingly complex, as some fraction appear evolved, and many LAEs appear to be dusty, which one would not expect from primordial galaxies.  Presently, much interest resides in discovering not only the highest redshift galaxies to constrain theories of reionization, but also pushing closer to home, as previous ground-based studies have only found LAEs at $z > 3$ due to observational limitations.  In this review talk I will cover everything from the first theoretical predictions of LAEs, to their future prospects for study, including the HETDEX survey here in Texas. ",
        "introduction": "As new and more powerful instruments and telescopes are developed, we now probe deeper into the Universe than ever before.  While once redshifts of a few tenths were difficult to observe, we can now routinely probe to $z = 4$--$7$, corresponding to a time when the Universe was only $\\sim 1$\\,Gyr old.  The goal is simple in principle.  When we look at the galaxies around us, we see a well defined morphological sequence \\citep{hubble26}, extending out to $z\\geq 1$.  We would like to understand how the first galaxies formed, and through what evolutionary means they evolved into the galaxies we see today.  To study the early Universe, we need techniques to accurately determine the distances to galaxies.  Locally, various distance ladder rungs are used, but the most powerful of these, SNe Ia, taps out at $z\\sim 1$.  At high redshift, a common method is to use the expected spectral features of the galaxies we hope to find, such as with the Lyman break technique \\citep{steidel93}.  At $z \\geq 2$, the Lyman limit, an intrinsic spectral break at 912\\,\\AA, shifts into the observed optical.  Though the intergalactic medium (IGM) is largely ionized, enough remnant neutral gas is in a typical line-of-sight to efficiently absorb photons bluer than the Lyman limit, increasing the amplitude of this break.  Additionally, from $912$--$1216$\\,\\AA, a forest of absorption lines appear due to absorption of intervening H\\,{\\sc i} along the line of sight, known as the Ly$\\alpha$~forest.  At higher redshifts, the number of absorbers increases, and the opacity in the Ly$\\alpha$~forest increases dramatically, such that the spectral break begins at 1216\\,\\AA, and is called the Ly$\\alpha$~break.  Using this spectral feature, one can discover large samples of Lyman break galaxies (LBGs) by observing with three broadband filters.  For example, at $z \\sim 3$, the Lyman break is in the $u^{\\prime}$-band, thus one would observe with $u^{\\prime}$, $g^{\\prime}$ and $r^{\\prime}$ filters.  As the break passes through the $u^{\\prime}$-band, the $u^{\\prime}-g^{\\prime}$ color will increase.  By examining how the colors of stellar population model tracks evolve with redshift, one can select a $u^{\\prime}-g^{\\prime}$ color which will select galaxies at $2.5 \\leq z \\leq 3.5$.  However, this single color can also include red galaxies at lower redshifts, thus one typically also demands a blue rest-frame ultraviolet (UV) color, $g^{\\prime}-r^{\\prime}$ in our example.  With this method, one can efficiently select large numbers of high-redshift star-forming galaxies\\footnote[1]{See a visualization of this selection at: http://shepherd.physics.tamu.edu/select$\\_$movie.html}.  However, the LBG selection method relies on robust continuum detections in multiple images, as well as a very deep dropout filter.  Thus, typical LBG studies place a brightness limit on their sample, typically $R\\leq 25.5$ \\citep{steidel96}, leaving a large population of faint galaxies unexplored.  Additionally, though galaxies selected with this method should lie in a redshift slice of $\\Delta$$z \\sim 1$, one requires followup spectroscopy to confirm the selection criteria and determine the redshift selection function.  This can be costly, as even ``bright'' LBGs ($m_{\\rm AB}\\sim 23$) can require many hours on 8--10m class telescopes to confirm their redshifts via absorption lines.  If the goal is to study the earliest galaxies, then one can rely on some additional spectral features.  In particular, they should be heavily star-forming, and thus bright in numerous nebular emission lines.  \\citet{partridge67} predicted that one could select the first galaxies on the basis of their Ly$\\alpha$~emission, which at a rest wavelength of 1216\\,\\AA~shifts into the optical from $2 \\leq z \\leq 6$.  With a typical initial mass function (IMF), many stars will form with sufficient mass (and thus temperature) to emit large quantities of hydrogen-ionizing photons ($\\lambda \\leq 912$\\,\\AA).  These photons will ionize H\\,{\\sc i} out to approximately the Str$\\mathrm{\\ddot{o}}$mgren radius.  At this point, there is a balance between ionization and recombination, and when a proton and electron recombine, a Ly$\\alpha$~photon will be emitted $\\sim 67$\\% of the time \\citep[Case B, ][]{osterbrock89}.  \\citet{partridge67} predicted that $> 5$\\% of the total luminosity of these galaxies could emit in this single line.  Spurred by these predictions, a number of studies searched for these Lyman alpha emitting galaxies (LAEs), but to no avail \\citep[e.g.,][]{koo80,meier81}.  A number of causes were suggested to explain the lack of LAEs, such as the possibility that these galaxies could be dusty \\citep[e.g.,][]{meier81}.  In a uniform interstellar medium (ISM), Ly$\\alpha$~photons will resonantly scatter off H\\,{\\sc i}, resulting in longer path lengths, and thus they stand a much greater chance of being absorbed by dust.  Or, some reasoned that perhaps these objects simply didn't exist, or were far too faint to be observed.  Thankfully, fortunes turned in the late 1990's, as two groups searching for LAEs, \\citet{cowie98} and \\citet{rhoads00}, discovered large samples.  Why were they able to do this while previous groups were not?  The answer is a matter of technological advancement.  \\citet{cowie98} used the new Keck 10m telescope, while \\citet{rhoads00} used the new MOSAIC large format detector on the Kitt Peak National Observatory 4m telescope.  The deeper imaging afforded by Keck or the wider field-of-view provided by MOSAIC both resulted in a larger volume being probed, and thus gave these studies a greater chance of discovery.  These, and subsequent studies, used the method of narrowband selection to identify these galaxies on the basis of their bright Ly$\\alpha$~lines.  To do this, one observes with a narrowband filter, where the redshift of the LAEs one finds is defined by the central wavelength and width of the filter.  For example, \\citet{rhoads00} used a set of filters around 6500\\,\\AA, thus the LAEs they found were at $z \\sim 4.5$.  One also observes the field with a broadband filter encompassing the narrowband filter.  Examining sources in both images, one selects objects with a flux excess in the narrowband image as possible emission line candidates.  However, low-redshift interlopers are possible; for example, with a narrowband filter at 6500\\,\\AA, one could also select [O\\,{\\sc ii}], [O\\,{\\sc iii}] or H$\\beta$ emitters.  To screen against these, one typically enforces a dropout criterion similar to the Lyman break method, requiring an additional broadband image blueward of the line.  Using this method of narrowband selection, \\citet{cowie98} and \\citet{rhoads00} both found that one could discover large samples of LAEs fairly easily, at least at redshifts where Ly$\\alpha$~is observed in the optical.  One of the first interesting discoveries about LAEs was that their Ly$\\alpha$~equivalent widths (EWs) were strong (the EW is a measure of the line strength relative to the continuum).  In many cases, the EWs exceeded predictions from normal stellar population models \\citep[e.g.,][]{charlot93}.  Figure~\\ref{ewdist} shows the distribution of Ly$\\alpha$~EW versus age for different star-formation histories (SFHs).  Unless one is observing a LAE in the first few Myr, the Ly$\\alpha$~EW should be $<$ 200\\,\\AA.  However, this is frequently not the case \\citep[e.g.,][]{kudritzki00,malhotra02}.  Thus, something unexpected is occurring in many LAEs.  \\begin{figure}[!ht] \\centering \\includegraphics[width=5in]{finkelstein.fig1.eps} \\caption{Model variation of Ly$\\alpha$~EW with age for three different star formation histories: constant, exponentially decaying, and a simple stellar population, or burst (from Finkelstein et al.\\ 2008).  All of these models have Z = 0.02 Z$_{\\odot}$, and include the effect of IGM absorption at $z \\sim 4.5$, where we assume that half of the line flux (i.e.~the blue wing) will be subject to IGM attenuation.}\\label{ewdist} \\end{figure} ",
        "conclusions": "Selecting galaxies on the basis of their Ly$\\alpha$~emission allows the discovery of a plethora of faint, high-redshift galaxies.  Numerous studies over the past few years have shown that these galaxies primarily appear to be a unique population of star-forming galaxies, which are young ($< 100$\\,Myr) and low mass ($< 10^{9} M_{\\odot}$), although a minority can be older and more massive.  Many of the LAEs which have been discovered exhibit strong Ly$\\alpha$~equivalent widths which cannot be explained by normal stellar populations.  While AGN contamination could reconcile this problem, AGN (at least luminous ones) appear to be rare among LAEs.  Very low metallicities could also cause high EWs, but the dustiness of many LAEs likely rules out this scenario.  Finally, this dust, if in a clumpy geometry, can enhance the Ly$\\alpha$~EW by increasing the escape fraction of Ly$\\alpha$~relative to the continuum, and this scenario can explain the SEDs of a sample of high-redshift LAEs.  However, to truly determine the make-up of LAEs, including their dust properties, we require direct measurements, which is available presently with NIR spectroscopy, and in the near future with ALMA.  Pushing galaxy studies to the highest redshifts, ground-based narrowband selection of LAEs becomes extremely difficult due to the bright NIR sky background, with the deepest studies finding only a handful of candidate $z > 7$ LAEs.  However, the addition of WFC3 to {\\it HST} provides the capability to obtain large samples of $z > 7$ galaxy candidates.  Initial studies of these galaxies have shown them to be somewhat primitive in their make-up, with young ages ($< 100$\\,Myr), low masses ($10^{8}$--$10^{9}$ $M_{\\odot}$), low metallicity ($< 0.1Z_{\\odot}$), and consistent with little-to-no dust extinction.  Comparing the physical properties of these most distant galaxies to those at lower redshifts, they appear more similar to LAEs at all redshifts than LBGs at any redshift.  This implies that typical galaxies at $z > 7$ are less evolved than at lower redshifts.  At some point in the past we should begin to see the first building blocks of galaxies, and these observations show that $z=$7--8 may be this epoch, although larger samples of spectroscopically confirmed galaxies will be necessary to make a definitive determination."
    },
    "1004/1004.5138_arXiv.txt": {
        "abstract": "We give an explicit realization of the ``String Axiverse\" discussed in Arvanitaki et. al \\cite{Arvanitaki:2009fg} by extending our previous results on moduli stabilization in $M$ theory to include axions. We extend the analysis of \\cite{Arvanitaki:2009fg} to allow for high scale inflation that leads to a moduli dominated pre-BBN Universe. We demonstrate that an axion which solves the strong-CP problem naturally arises and that both the axion decay constants and GUT scale can consistently be around $2\\times 10^{16}$ GeV with a much smaller fine tuning than is usually expected. Constraints on the Axiverse from cosmological observations, namely isocurvature perturbations and tensor modes are described. Extending work of Fox et. al \\cite{Fox:2004kb}, we note that {\\it the observation of tensor modes at Planck will falsify the Axiverse completely.} Finally we note that Axiverse models whose lightest axion has mass of order $10^{-15}$ eV and with decay constants of order $5\\times 10^{14}$ GeV require no (anthropic) fine-tuning, though standard unification at $10^{16}$ GeV is difficult to accommodate. ",
        "introduction": "\\label{intro}  The dimensionless QCD $\\theta$-angle is a source of CP violation in the Standard Model, highly constrained by measurements of the electric dipole moment (EDM) of the neutron and $^{199}Hg$: $|\\theta_{qcd}| \\lesssim 10^{-10}$ \\cite{Baker:2006ts}. This is the strong CP problem, {\\it why is $\\theta_{qcd}$ so small?}.  An elegant solution to the problem might be provided by the Peccei-Quinn (PQ) mechanism \\cite{Peccei:1977hh}, in which $\\theta_{qcd}$  is promoted to a dynamical field known as the QCD axion ($a_{qcd}\\equiv f_a^{qcd}\\,\\theta_{qcd}$). The axion is charged under an anomalous global $U(1)$ symmetry, broken by non-perturbative QCD effects at the mass-scale $f_a^{qcd}$. The potential generated by the QCD instantons is: \\ba \\label{Vqcd}V \\sim \\Lambda_{qcd}^4\\,\\left[1-\\cos\\left(\\frac{a_{qcd}}{f_a^{qcd}}\\right)\\right]\\,,\\ea which is minimized at $\\theta_{qcd} = 0$, solving the strong CP problem. Various properties of the QCD axion have been considered in \\cite{oldaxion}, \\cite{Raffelt}.  However, in order to really solve the problem, this PQ symmetry must be of incredibly high quality. In particular, all contributions to the axion potential from other PQ breaking sources must be at least ten orders of magnitude suppressed compared to that from QCD instantons.  Phenomenological approaches in field theory trying to realize such a PQ symmetry have been pursued \\cite{axionpheno}; see \\cite{Cheung:2010hk} for a recent approach. However, since global symmetries are believed to be always broken by quantum gravitational effects, it is worthwhile exploring axions in string theory and under what circumstances they could solve the strong-CP problem.  Pseudoscalar fields with axion-like properties generically arise in string/$M$ theory as the zero modes of antisymmetric tensor fields along the extra dimensions \\cite{Witten:1984dg}. The number of axions is determined essentially by the topology of the extra dimensions. This number, like the gauge group and number of fermion generations, can be viewed as a {\\it discrete},  UV boundary condition and it is fairly common to have hundreds, if not thousands, of axions present. There is a separate PQ symmetry for each axion field, inherited from the gauge symmetry associated with the higher dimensional tensor field. Often these axions pair up with geometric moduli fields to form complex chiral mutliplets in ${\\cal N}=1,D=4$ supergravity, the effective theory describing  many string compactifications. Although a plethora of axions occur in string/$M$ theory, they do not automatically satisfy the criteria required by the PQ mechanism for the QCD axion. The absence of PQ breaking operators is usually guaranteed to all orders in perturbation theory because the PQ symmetries are remnants of (higher dimensional) gauge symmetry, surviving as global symmetries below the Kaluza-Klein scale. As in ordinary QCD, this global symmetry can be broken by non-perturbative effects, such as those arising from instantons of various sorts - worldsheet, brane, gauge and gravitational. One has to ensure that in addition to being valid to all orders in perturbation theory, the contributions to the QCD axion potential from these non-perturbative effects are negligible compared to that from QCD instantons. Axions in various string theoretic contexts have been studied in \\cite{axionstring}.  A qualitative picture of the mass spectrum of axions in string theory can be given. Since the potential of a given axion $a$ is generated by non-perturbative effects, its mass-squared will be of order $\\frac{\\Lambda^4}{f^2_a}$ -- where $\\Lambda^4 \\sim M^4 e^{-bS_{inst}}$, where $S_{inst}$ is the action of the instanton,  which dominates the potential of $a$ and $b$ is a number, which characterises the \"charge of the instanton\" that generated the potential \\cite{Svrcek:2006yi}. $M$ is given by the geometric mean of the supersymmetry breaking scale ($\\sqrt{F}$) and the `fundamental' scale, eg. the Planck scale. We will derive a precise formula later with the above general form. $S_{inst}$ is usually the volume of a submanifold of the extra dimensions in fundamental units. We will later see that $f_a$ is typically of order the $GUT$ scale. Therefore, the spectrum of axion masses will typically be exponentially hierarchical. Roughly speaking, with a large number of axions, one expects a spectrum, which is uniformly distributed on a log scale, rather like the Yukawa couplings in $M$ theory \\cite{Acharya:2003gb}. The axion masses are thus expected to span many orders of magnitude. These sorts of observations inspired the authors of \\cite{Arvanitaki:2009fg} to propose a variety of astrophysical tests, which could probe the large range of axion masses from $10^{-33}$eV to $10^8$ eV. These tests could provide evidence for the existence of an Axiverse. These include observations of the polarization of the CMB, suppression of the matter power spectrum at small scales and the spectrum of gravitational waves from rapidly rotating astrophysical black holes. The purpose of this paper is not to discuss these phenomena and tests in detail, but to address in more depth the nature of the mass spectrum of axions in string/$M$ theory and the solution to the strong CP problem. At the end, however, we will summarize some of the important astrophysical observables, which could either completely falsify the framework or provide strong evidence for it.   The requirement of stabilizing moduli with a sufficiently large mass in realistic string compactifications complicates the above picture for generating masses for the axions. For example, simple mechanisms of stabilizing moduli with fluxes in a supersymmetric minimum can completely break the PQ symmetry and give axions masses comparable to that of the moduli \\cite{Acharya:2002kv}. Stabilizing certain moduli by non-perturbative effects in a supersymmetric minimum, such as in the KKLT idea \\cite{Kachru:2003aw}, also gives axion masses comparable to that of moduli; hence these axions \\emph{cannot} solve the strong CP-problem. In fact, it can be shown that there does not exist any \\emph{supersymmetric} minimum within ${\\cal N}=1,D=4$ supergravity with phenomenologically allowed values of stabilized moduli but unfixed axions (down to the QCD scale), which could solve the strong CP problem \\cite{Conlon:2006tq}. In particular, for supersymmetric vacua it was shown that either all moduli appear in the superpotential in which case the axions are very massive, or some moduli do not appear in the superpotential in which case the potential is tachyonic \\cite{Conlon:2006tq}.  Therefore, in order to look for QCD axion candidates in string theory, one should consider compactifications in which moduli are stabilized in a {\\it non-supersymmetric} minimum. Such vacua have been studied in detail in the context of $M$ theory compactifications on $G_2$ manifolds without fluxes \\cite{Acharya:2006ia, Acharya:2008hi}. In these compactifications \\emph{all} the axions pair up with geometric moduli fields in such a way that all the moduli superfields enjoy a PQ symmetry, implying that the entire moduli superpotential can only arise from non-perturbative effects. In a series of papers \\cite{Acharya:2006ia, Acharya:2008hi} it has been proven that strong dynamics in the hidden sector can a) generate a potential, which stabilizes {\\it all} the moduli fields and b) generates a hierarchically small supersymmetry breaking scale. Moreover, the potential stabilizes just one out of the many axion fields. Happily, the QCD axion belongs to the set of unfixed combinations. Higher order non-perturbative effects, which are generically present, but not considered in \\cite{Acharya:2006ia, Acharya:2008hi} because they are sub-dominant, {\\it will} then stabilize the remaining axions (including the would-be QCD axion) at an exponentially smaller scale than the moduli mass scale (which is set by the gravitino mass $m_{3/2}$). In other words, one considers the moduli to be stabilized at the  supersymmetry breaking scale $(m_{3/2})$ with sub-leading non-perturbative effects being responsible for stabilizing most of the axions. Being non-perturbative, the resulting spectrum of axions will be distributed (roughly) uniformly on a logarithmic scale - as we indeed verify in a detailed model - thereby giving a detailed realization of the Axiverse. It can be shown that certain classes of Type IIB compactifications inspired by the moduli stabilization mechanism in $M$ theory \\cite{Acharya:2006ia, Acharya:2008hi} also exhibit the above properties \\cite{Bobkov:2010rf}. \\emph{Therefore, the analysis and results of this paper can be directly applied to the those Type IIB compactifications as well}. For concreteness, however, we will study the effective theory arising from $M$ theory compactifications here.   The axion decay constants $f_a$ are constrained by various observations. The decay constant for the QCD axion $f_a^{qcd}$ must be bounded from below by $\\sim 10^9$ GeV because of processes like axion emission from stars and supernovae. An upper bound on {\\it all} axions arises from the requirement of `not overclosing the Universe'. The thermal population of light axions ($\\lesssim 1$ eV for $f_a \\sim 10^9$ GeV) is typically quite negligible. However, being extremely weakly coupled to matter, axion fields begin coherent oscillations, which begin when the Hubble scale $H \\sim m_a$, and the constraint arises from not storing too much energy density in the axion fields. This energy density is quadratic in the initial vev (called mis-alignment angle) of the axion field. The value of the upper bound obtained clearly depends on the detailed cosmological evolution i.e. matter vs. radiation domination and entropy releases during the oscillations. Assuming a cosmology characterized by a radiation-dominated phase after inflation, one gets $f_a^{qcd} \\lesssim 10^{12}$ GeV (the bound for general axions is given in section IV.) Thus, in order to satisfy the upper bound on $f_a^{qcd}$  \\emph{without} fine-tuning the intial misalignement angle $\\theta_i^{qcd}$, one has to either invoke a very large cycle on which the QCD gauge group is localized, and/or have a sufficiently low string/11D scale, implying a large overall volume of the extra dimensions ${\\cal V}$) \\cite{Svrcek:2006yi}. Having a very large QCD cycle to make $f_a$ respect the above upper bound makes the strong gauge coupling $\\alpha_{QCD}$ far too weakly coupled at the compactification scale and hence does not seem very natural. So the only reasonable option without fine-tuning the misalignment angle, is to have a sufficiently small string/11D scale $M_s/M_{11} \\ll M_{GUT}$ with a very large overall volume ${\\cal V}$ but a small (still within the supergravity regime) QCD cycle such that $\\alpha_{QCD}$ is reasonable \\cite{Svrcek:2006yi}. However, this implies that the standard unification scale $\\sim 10^{16}$ GeV is above the string scale. Also, a low fundamental scale will likely lead to conflicts with current bounds on the proton lifetime. Of course, it is possible that the initial misalignment angle is small ($\\ll 1$) for some reason, which can make a large unification scale ($\\sim 10^{16}$ GeV) compatible with the observational bound on $f_a$, assuming standard thermal cosmology. In fact, it has been argued that a small misalignment angle consistent with observations could be anthropically selected \\cite{Linde:1991km}, \\cite{Hertzberg:2008wr}. The $M$ theory models, as we will see, (and as is assumed in \\cite{Arvanitaki:2009fg}) have $f_a \\sim M_{GUT} \\sim 10^{16}$ GeV.  One of the key points of \\cite{Acharya:2008bk} is that, if the Hubble scale after inflation $H_I \\geq m_{3/2} \\sim 50$ TeV then the energy density of the Universe is dominated by moduli oscillations until just before BBN. This is a different cosmological history than the one that yields the upper limit $f_a \\leq 10^{12}$GeV. Moduli with masses of order $50$ TeV will decay just before BBN and, as pointed out in \\cite{Fox:2004kb}, the entropy released by the moduli decay can allow a larger $f_a \\leq 5 \\times 10^{14}$ GeV without fine tuning. Thus, the anthropic fine tuning of the initial axion mis-alignment angle is much smaller since the upper bound is much closer to the GUT scale. \\cite{Arvanitaki:2009fg} do not consider such moduli dominated cosmological histories (which are presumably generic in string theory) and in fact take the Hubble scale to be rather small after inflation $H_I \\leq 0.1$GeV. It is somewhat satisfying that the fine-tuning of the misalignment angles is greatly reduced in the (presumably) more generic case of moduli dominated, non-thermal cosmology. We thus extend the results of \\cite{Fox:2004kb} to the Axiverse and extend the Axiverse to the presumably more generic case of $H_I \\geq m_{3/2}$. We find that the non-thermal moduli dominated pre-BBN Universe is much less fine tuned than a standard radiation dominated pre-BBN Universe arising in low scale inflation. Extending the results of \\cite{Fox:2004kb}, which only considers the QCD axion, to the Axiverse, we emphasise that the observations of tensor modes by Planck, when combined with the existing bounds on axion induced isocurvature perturbations, would rule out the string Axiverse completely, thereby requiring axion masses to be `large'.  The plan of the paper is as follows. In section \\ref{stable}, the mechanism to stabilize the axions is outlined in detail and the spectrum of eigenvalues and eigenstates is computed. A simple formula for the axion masses is obtained. In section \\ref{qcdaxion}, the effects of QCD instantons are taken into account and the precise method of identifying the QCD axion candidate among all axions is explained. It is shown that such a candidate typically occurs  within the $M$ theory framework considered. Section \\ref{cosmo} is an analysis of the axion relic abundances within the `thermal' and `non-thermal' cosmological histories, followed by a discussion of the allowed parameter space after imposing current cosmological constraints in section \\ref{Cons}. This is followed in  section \\ref{observables} by a discussion of cosmological observables, which could either falsify or provide strong evidence for the framework, and also distinguish among the two cosmological histories. Technical details follow in Appendices \\ref{masses}, \\ref{toy} and \\ref{stats}. ",
        "conclusions": ""
    },
    "1004/1004.0692_arXiv.txt": {
        "abstract": "We present new spectroscopic and photometric observations of the transiting exoplanetary system WASP-3. Spectra obtained during two separate transits exhibit the Rossiter-McLaughlin (RM) effect and allow us to estimate the sky-projected angle between the planetary orbital axis and the stellar rotation axis, $\\lambda=3.3^{+2.5} _{-4.4}$ degrees. This alignment between the axes suggests that WASP-3b has a low orbital inclination relative to the equatorial plane of its parent star. During our first night of spectroscopic measurements, we observed an unexpected redshift briefly exceeding the expected sum of the orbital and RM velocities by 140~m~s$^{-1}$. This anomaly could represent the occultation of material erupting from the stellar photosphere, although it is more likely to be an artifact caused by moonlight scattered into the spectrograph. ",
        "introduction": "One of the hottest exoplanets yet discovered is WASP-3b, a giant planet with a mass of $1.76$~$M_{\\rm J}$ that orbits an F7-8V star at a distance of only 0.032~AU, giving the planet an equilibrium temperature of 1960~K \\citep{pollacco}. This close-in, transiting system is an ideal target for study with the Rossiter-McLaughlin effect because the host star is both relatively bright ($V=10.5$) and rapidly rotating ($v\\sin i_\\star= 13.4 \\pm 1.5$~km~s$^{-1}$), and the transit occurs with an impact parameter near 0.5. Observations of the Rossiter-McLaughlin effect yield valuable information about the system, namely the angle ($\\lambda$) between the sky projections of the planetary orbital axis and the stellar rotation axis.  This parameter is a basic geometric property of the system and a possible clue about the processes of planet formation and orbital migration. Detailed descriptions of the Rossiter-McLaughlin effect and its applications can be found in \\citet{ohta}, \\citet{gim06}, \\citet{gaudi}, and \\citet{fabrycky}, while early examples of such measurements for exoplanetary systems are given by \\citet{queloz} and \\citet{winn05}.  In this paper we report the results of new photometric and spectroscopic observations of WASP-3, which were conducted with the primary goal of determining $\\lambda$ and the secondary goal of refining estimates of other system parameters. In \\S~\\ref{sec:obs} we describe the observations and data reduction, and in \\S~\\ref{sec:analysis} we present the model that was used to fit the data and our results for the system parameters. These results are discussed in \\S~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  We have observed the Rossiter-McLaughlin effect during transits of WASP-3b and found that the projected spin-orbit angle is $\\lambda=3.3 ^{+2.5} _{-4.4}$ degrees, suggesting that the stellar rotation axis and the orbital axis are closely aligned.\\footnote{While the first draft of this manuscript was under review, we learned of a similar study of WASP-3 by \\citet{2009arXiv0912.3643S}.  Those authors also measured the RM effect, and found $\\lambda=15_{-9}^{+10}$~deg, also suggesting a relatively close spin-orbit alignment.} This alignment makes WASP-3 similar to the majority of other exoplanets whose alignment has been studied \\citep{fabrycky} and unlike the misaligned systems XO-3 \\citep{hebrard, winn-rm-xo3}, HD~80606 \\citep{pont09, winn-rm-hd80606}, WASP-14 \\citep{johnson-wasp14}, HAT-P-7 \\citep{narita, winn-rm-hatp7}, Corot-1 \\citep{pont09}, and WASP-17 \\citep{anderson}. It would seem that for WASP-3, the inward migration process did not disrupt the initial coplanarity of the system, or tidal effects have coplanarized the system (although the theoretically calculated rate of coplanarization is too slow to be relevant, \\citet{barker}).  We also used new RV measurements and transit light curves to determine the system parameters, and found agreement with previously published values, with the exception of the RV semi-amplitude $K$ reported by \\citet{pollacco}. In addition, we observed an intriguing RV anomaly during one of the transits. Our investigation of the spectra led us to the conclusion that this ``RV spike'' was caused by a small amount of moonlight that was admitted to the slit along with the light from WASP-3.  However, we cannot be completely certain in this interpretation, and in any case it is interesting to ask whether there are any astrophysical phenomena that could produce such an effect.  One speculative idea is that the spike represents the planet's passage in front of a starspot.  When a spot is hidden by a planet, the apparent RV of the star will jump to a larger or smaller value, because the RM-like effect of the spot temporarily vanishes.  When the spot is no longer hidden, the star's apparent RV returns to its expected transiting value.  The resulting ``spike'' in RV would be analogous to the photometric ``bumps'' or ``rebrightenings'' that have been seen in transit light curves and attributed to starspot crossings, as in \\citet{rabus} and \\citet{dittmann}, amongst others.  The time between the RV spike's initial point and its maximum ($\\approx 20$~mins) is consistent with the ingress or egress duration, as would be required of the spot-occultation hypothesis.  However, to match the $+140$~m~s$^{-1}$ amplitude of the RV spike, the starspot would need to present a {\\it velocity}\\, contrast with the photosphere, in addition to a possible intensity contrast. This can be understood as follows: under the starspot hypothesis the amplitude of the RV spike would have an amplitude \\begin{equation} \\Delta v \\approx v_\\star \\left(\\frac{R_s}{R_\\star}\\right)^2\\left(1 - \\frac{I_sv_s}{I_\\star v_\\star}\\right), \\label{eq:rv-noise} \\end{equation} assuming there is a single spot with radius $R_s$, intensity $I_s$, and radial velocity $v_s$ that may differ from the surrounding photosphere. In the absence of the spot, the material at the spot's location would have intensity $I_\\star$ and radial velocity $v_\\star$. Because the spike occurred near the transit egress, the spot would have been near the receding limb of the star, where $v_\\star = +v\\sin i_\\star\\sqrt{1-b^2} \\approx +12$~km~s$^{-1}$. Assuming $R_s\\approx R_p$ and setting $\\Delta v$ of Eqn.~(\\ref{eq:rv-noise}) equal to $140$~m~s$^{-1}$, we find \\begin{equation} \\frac{I_s v_s}{I_\\star v_\\star} \\approx -0.06. \\end{equation} Since intensity cannot be negative, the above relation implies that the spot's effective RV must be oppositely directed from the local photospheric RV.  The planet must cover a blueshifted spot in order to produce a redshift as large as 140~m~s$^{-1}$. Specifically, the spot's peculiar RV, defined as $v_s - v_\\star$, is approximately $-12.7$~km~s$^{-1}$~$(I_\\star/I_s)$.  Unless the spot is very bright ($I_s \\gg I_\\star$) the peculiar velocity just estimated would exceed the sound speed of the photosphere, which we estimate to be 7~km~s$^{-1}$. The hidden-spot hypothesis for the RV spike thus requires not an ordinary spot, but rather an erupting and possibly even supersonic spot.  What we are calling a ``spot'' would actually be more akin to an active site undergoing a stellar flare.  Thus, the RV anomaly could be rationalized as a consequence of vigorous stellar activity, a hypothesis which would also produce transit depth variations (see \\S~3.1). A serious objection to this hypothesis, however, is that by the usual measures WASP-3 is a quiescent main-sequence F7-8 dwarf star for which stellar activity is not expected.  \\citet{pollacco} presented no evidence for photometric variations, chromospheric emission lines, or any other indication of stellar activity. Likewise, there is no detectable Ca II H~\\&~K emission in any of our own Keck/HIRES spectra, including those obtained during the RV spike.  The measured Mt.~Wilson $S$ values had an RMS of only 0.0013, with no noticeable offset even between the measurements spanning different seasons.  Considering $R'_{HK}$, the ratio of chromospheric emission in Ca II H \\& K cores to the total stellar emission, we find that WASP-3's mean $\\log{R'_{\\rm HK}}=-4.9$, giving it an approximate age of 4 Gyr \\citep{noyes}.  By comparison, a very active star such as HD189733 has $\\log{R'_{\\rm HK}}=-4.5$, while a very quiet star such as HD9407 has $\\log{R'_{\\rm HK}}=-4.98$, \\citep{wright04}. Finally, one might expect an unusually active star to be an unusually rapid rotator, but the value of $v\\sin i_\\star=14.1$~km~s$^{-1}$ is typical of an F7V star. The activity hypothesis would therefore imply a high level of activity in a star that one would otherwise expect to be quiescent.  Although the activity hypothesis is not tenable for WASP-3, we predict that similar observations of other systems with more active stars will eventually reveal ``RV spikes'' when the planet occults active regions.  Simultaneous photometry and spectroscopy of these events will allow investigations of the intensity and velocity contrast of the active regions, relative to the surrounding photosphere.  Active stars are often excluded from Doppler planet surveys, precisely because the excess RV noise caused by stellar activity is a hindrance to planet detection.  However, magnitude-limited photometric transit surveys have no such selection, and in this regard it is unsurprising that planets around very active stars are being routinely discovered by these surveys, such as CoRoT-2 \\citep{lanza} and CoRoT-7 \\citep{leger}.  The Doppler follow-up for these systems will undoubtedly be more difficult than for quiescent stars.  But, perhaps these difficulties will be compensated to some degree by the prospect that active stars with eclipsing planets will eventually enable a deeper understanding of starspots, stellar activity, and stellar flares."
    },
    "1004/1004.3287_arXiv.txt": {
        "abstract": "Digital co-addition of astronomical images is a common technique for increasing signal-to-noise and image depth. A modification of this simple technique has been applied to the detection of minor bodies in the Solar System: first stationary objects are removed through the subtraction of a high-SN template image, then the sky motion of the Solar System bodies of interest is predicted and compensated for by shifting pixels in software prior to the co-addition step. This ``shift-and-stack'' approach has been applied with great success in directed surveys for minor Solar System bodies. In these surveys, the shifts have been parameterized in a variety of ways. However, these parameterizations have not been optimized and in most cases cannot be effectively applied to data sets with long observation arcs due to objects' real trajectories diverging from linear tracks on the sky. This paper presents two novel probabilistic approaches for determining a near-optimum set of shift-vectors to apply to any image set given a desired region of orbital space to search. The first method is designed for short observational arcs, and the second for observational arcs long enough to require non-linear shift-vectors. Using these techniques and other optimizations, we derive optimized grids for previous surveys that have used ``shift-and-stack'' approaches to illustrate the improvements that can be made with our method, and at the same time derive new limits on the range of orbital parameters these surveys searched. We conclude with a simulation of a future applications for this approach with LSST, and show that combining multiple nights of data from such next-generation facilities is within the realm of computational feasibility. ",
        "introduction": "Detecting faint objects in the Solar System is a technically challenging problem. The approaches taken in the past can be broken into two basic categories: inter-frame source detection and linking, and the ``shift-and-stack'' approach. Inter-frame detection, where transient sources are detected in individual exposures and motion is linked between multiple exposures, is widely employed and favored in planned proposed surveys like Pan-STARRS (Denneau et al. 2007) and LSST (Axelrod et al. 2009) due to its simplicity and robustness. This method, however, does not exploit imaging data to the fullest extent possible, since detections are subject to the completeness limit of individual frames.  Since Solar System objects are moving across the sky, their images trail and the noise contribution from the sky background increases. Trailing losses create an upper limit to individual frame exposure times for efficient detection of Solar System objects. It is advantageous to combine multiple exposures, each short enough that trailing effects are negligible. Co-addition is routinely done to improve the SN of stationary sources and to remove cosmic ray events, but in order to apply it to the detection of moving objects several other steps must be performed. First, contamination from stationary sources can be removed by subtracting a high-SN template image from each image in the set. Next, the sky motion of the objects of interest must be predicted and this motion compensated for by shifting each image's pixels in software. When the images are then combined only the flux from sources that moved at the predicted rates will be constructively added.  This method first appeared in literature in Tyson et al. (1992), though details of the application were limited. An early detailed description of this technique was presented by Cochran et al. (1995). A set of images from the \\textit{Hubble Space Telescope} (HST) \\textit{Wide Field Planetary Camera 2} (WFPC2) were combined to allow statistical detection of very faint TNOs. The sky motion of the sources was parameterized as angular rates ($\\dot{\\theta}$) and angles on the sky ($\\phi$). Gladman et al. (1997) used the same technique in order to combine a series of images from the \\textit{Palomar} 5m telescope in the hope of constraining the size distribution of Kuiper Belt Objects (KBOs) by discovering objects with smaller radii than had been detected in most previous wide-area surveys that employed inter-frame detection. This survey was followed by further searches from \\textit{Keck} (Luu \\& Jewitt 1998; Chaing \\& Brown 1999), \\textit{CTIO} (Allen et al. 2001 and 2002; Fraser et al. 2008), \\textit{Palomar} (Gladman et al. 1998), \\textit{VLT} and \\textit{CFHT} (Gladman et al. 2001), which employed the same technique to varying degrees of success. Recently, Fraser \\& Kavelaars (2009) and Fuentes, George, \\& Holman (2009) both used similar techniques to search \\textit{Subaru} data for KBOs as faint as $R \\sim 27$.   All of these surveys had observational arcs less than two days in length. Over periods this short, the motion of outer Solar System sources can be well approximated (with respect to ground-based seeing) by the linear trajectories that the method these surveys implemented implicitly assumes. For longer arcs or higher-resolution data, however, the approximation of a linear trajectory is no longer adequate. Bernstein et al. (2004) employed a more advanced approach, which they called ``digital tracking,'' to extremely high-resolution HST \\textit{Advanced Camera for Surveys} (ACS) data. Even in arcs on the order of a day in length, nonlinear components of motion were non-negligible relative to the ACS PSF. In order to stack frames such that distant moving sources add constructively in this regime it becomes necessary to generate nonlinear shift-vectors via full ephemerides from orbits of interest. They stacked arcs up to 24 hours in length using shift-vectors parameterized as a grid over three parameters: two linear components of motion and heliocentric distance $d$. By sampling densely over this grid, they were able to recover orbits with $25$ AU $\\leq d \\leq \\infty$ and $i < 45^{\\circ}$, and detect objects as faint as $R\\sim28.5$. However, this dense sampling, somewhat unconstrained discovery volume, and extremely small PSF required $\\sim7\\times10^{5}$ shift-vectors, which translated into $\\sim10^{14}$ pixels to be searched. Bernstein et al. found that stacking more than 24 hours of their data became computationally prohibitive.  In the age of ground-based gigapixel CCD mosaic imagers, how is data obtained over long arcs going to be efficiently exploited to search for faint sources? In the interest of exploring massive datasets with long observation arcs to as deep a limit as possible, we have created a simple probabilistic method for generating an optimized set of shift-vectors for any imaging survey and any interesting volume of orbital element space. By generating shift-vectors in a probabilistic way, exploring an explicitly demarcated region of orbital parameter space becomes straightforward. The targeted region of orbital parameter space is searched with uniform sensitivity, simplifying the accurate characterization of any detected sample of objects. This method also avoids searching non-physical or uninteresting parameter space, thereby maximizing scientific return for minimum computational cost.    In Section \\ref{MTE}, we define the maximum tracking error, the basic parameter that defines the density of a grid. Section \\ref{MC} describes previous analytical methods for search-grid parameterization, our probabilistic approach, and a grid-tree structure for optimizing multiple-night arcs. Section \\ref{COMP} compares the efficiency of our method to those used in previous surveys, and in Section \\ref{LSST} we show an example application of this method to the kind of data that may be acquired by the Large Synoptic Survey Telescope (LSST). ",
        "conclusions": "We have shown that by generating search-grids for pencil-beam surveys in a probabilistic way and implementing an optimized grid geometry, significant reductions ($25-75\\%$) can be made in the total number of operations required, and using a tree-of-grids structure can reduce the total pixel additions required even further. Besides improvements in computational requirements, this method also ensures uniform sensitivity to the entire targeted volume of orbital parameter space.  This method also provides a simple way to include non-linear components of motion necessary to track objects over long observational baselines. This combination of reduced computational requirements and simplicity of extending the observational arc makes the prospect of stacking multiple nights of data appear completely feasible. The additional depth and improved precision of measured orbital properties derived from multi-night stacks would be hugely beneficial if used in upcoming large-area surveys like LSST and Pan-STARRS. Simulations of the LSST deep observations show that, for certain TNO populations, it would require less computational effort to obtain the same depth as Bernstein et al. (2004) over 500 times the area by stacking multiple nights of LSST data with the methods described here.  Additionally, this method allows for simple characterization of the orbital sensitivity of existing pencil-beam surveys in literature. From our characterization of the surveys by Fraser et al. (2008), Fuentes et al. (2009), and Fraser \\& Kavelaars (2009), we have shown that the orbital sensitivity of these surveys varied from what was advertised. Their extra sensitivity to the motions of distant sources makes these surveys sensitive to populations like distant members of the Scattered Disk and Sedna-like objects, and their non-detection provide useful upper limits on these populations. We will explore these upper limits in future work.  Since the discovery of extremely high-inclination, low-pericenter objects like 2008 KV$_{42}$ ($i \\simeq 110^\\circ$), it is important to understand these surveys' sensitivity to such populations. We note that Fuentes et al. (2009) is the only survey characterized here that was sensitive to objects with the same inclination and distance at discovery as 2008 KV$_{42}$ (see Figure \\ref{All_surveys}, bottom-right panel), but we contend that it is unlikely that they would have recognized any detection as belonging to this retrograde population. Since follow-up has been performed only rarely for these deep surveys, orbital inclination and heliocentric distance usually remain degenerate, with the prograde solution lying at lower distances (though this degeneracy is rarely acknowledged). It is possible that detections labeled as high-inclination prograde objects may in fact be objects with retrograde orbits like 2008 KV$_{42}$ at greater distance. Only multi-night arcs or follow-up at later epochs can break this degeneracy and clearly identify objects belonging to rare TNO sub-classes."
    },
    "1004/1004.4658_arXiv.txt": {
        "abstract": "Through a detailed numerical investigation in three spatial dimensions, we demonstrate that long-lived time-dependent field configurations emerge dynamically during symmetry breaking in an expanding de Sitter spacetime. We investigate two situations: a single scalar field with a double-well potential and an SU(2) non-Abelian Higgs model. For the single scalar, we show that large-amplitude oscillon configurations emerge spontaneously and persist to contribute about 1.2\\% of the energy density of the universe. We also show that for a range of parameters, oscillon lifetimes are {\\it enhanced} by the expansion and that this effect is a result of  parametric resonance. For the SU(2) case, we see about $4\\%$ of the final energy density in oscillons. ",
        "introduction": "Spontaneous symmetry breaking plays a key role in our current understanding of particle physics and is expected to have been a major factor in determining the physical properties of the early universe \\cite{Vilenkin}. In cosmology, two aspects of symmetry breaking are of great interest: it typically happens far from thermal equilibrium and it is inherently nonlinear. In the context of the electroweak phase transition, for example, an initially thermalized state is tossed out of equilibrium as the Higgs evolves to acquire a nonzero expectation value. In inflation, a nonthermal state thermalizes to reheat the universe with an explosive energy transfer from the inflaton to other field modes. It is thus of great interest to study the dynamics of symmetry breaking in an expanding background numerically in order to isolate key features that may escape analytical techniques.  Here, we report results on 3d simulations for two situations: a single, self-interacting scalar field with a double-well potential, and an SU(2) non-Abelian Higgs model. In Refs. \\cite{graham_cos,farhi_cos} results have been obtained for the case of a single scalar in 1d. It was shown that long-lived, time-dependent field configurations known as oscillons \\cite{bogol,gleiser,copeland} emerged spontaneously and contributed an amazing 50\\% of the total energy density. These initial results triggered the present study in the context of more realistic models. There are two broad classes of scalar field oscillons that have been studied in the literature, small and large-amplitude. Small-amplitude oscillons do not probe the highly nonlinear domain of the potential, and typically have large spatial widths \\cite{fodor,farhi_cos,amin}. Their small amplitude makes it possible to study them using linearization techniques. Large-amplitude oscillons are harder to investigate analytically \\cite{dosc,sicilia,hertzberg}. Simulations of scalar models in static 2d and 3d backgrounds \\cite{oscnuc} and expanding 1d backgrounds \\cite{graham_cos,farhi_cos} indicate that mostly large-amplitude oscillons are excited during symmetry breaking. As we show next, this is also the case for an expanding 3d spacetime. The situation is different for SU(2) models, as we explain below.  This paper is organized as follows: in the next section, we introduce the scalar field model in an expanding universe and discuss its lattice implementation. We report our results for a double well potential, showing that oscillons contribute about 1.2\\% of the energy density. In section III we show that, contrary to naive expectation, for certain values of the expansion rate oscillons may have their lifetimes enhanced. We explain this result analytically by making use of parametric resonance. In section IV we introduce the SU(2) non-Abelian Higgs model and discuss its lattice implementation in an expanding universe. In section V we discuss the results for this model. In particular, we show that, as in the case of a real scalar field, oscillons contribute a nontrivial percentage of the total energy density. Furthermore, our results indicate that the cosmological expansion seems to favor the formation of oscillons for a wider range of parameters as compared to the static case, where oscillons were found only in a 2:1 mass ratio for the Higgs and gauge boson. In section VI, we briefly discuss possible application of oscillons in cosmology, which we hope to explore in forthcoming work, and conclude with a summary of our results. ",
        "conclusions": ""
    },
    "1004/1004.1557_arXiv.txt": {
        "abstract": "{This is one of results from our program of massive simulations of close encounters for all nearby stars. Epsilon Eridani is an extremely interesting star having one confirmed planet and multiple asteroid and debris belts. It should have a quite massive Oort cloud as well. Deltorn et al. searched for past Nemesis encounters of $\\varepsilon$ Eri. In this paper we show that, according to current astrometric data, an other famous nearby star Luyten 726-8AB (=BL/UV Ceti) will pass at $\\lesssim$ 0.93 ly from Epsilon Eridani in $\\approx$ 31.5 kyr. So, it will probably pierce through the outer part of the hypothetical Oort cloud of $\\varepsilon$ Eri. BL/UV Ceti has only about 20 percent of the solar mass. Nevertheless, it could influence directly some long-period comets of Epsilon Eridani. The duration of mutual transit of two star systems within 1 ly from each other is $\\gtrsim$ 4.6 kyr. Our simulations show that stellar encounters within 1 ly might be more frequent than previously thought. It could explain Proxima's peculiar trajectory with respect to $\\alpha$ Cen AB or even Sedna's trajectory in the solar system. ",
        "introduction": "Epsilon Eridani is a relatively young K-star having one confirmed planet, two asteroid belts at $\\sim3$ and $\\sim20$ AU resp.~as well as an analog of Kuiper belt starting at 35 AU from the star. It should also have the outer cloud of long-period comets more massive than the Oort cloud in our solar system. The litterature on $\\varepsilon$ Eri is very extensive, \\emph{cf.}~1159 references currently available in {\\footnotesize SIMBAD} database for additional details.  \\smallskip BL/UV Ceti was discovered by Luyten in 1948-49 as a star with a very high proper motion. UV Ceti is actually an archetype of flare stars famous for its extreme changes of brightness. Two components are at few AU from each other on highly elliptical orbits. The total mass of the system is about $0.2M_{\\odot}$. Luyten 726-8 is also the closest currently known neighbour of Epsilon Eridani (\\emph{cf.}~{\\footnotesize SIMBAD} database for an extensive list of references).  \\smallskip Deltorn and Kalas \\cite{D} searched for past close encounters of $\\varepsilon$ Eri. They have discovered that Epsilon Eridani had recently a close encounter with Kapteyn's star at the mutual distance of $\\approx$ 1 pc. So, as far as I know, it is somewhat unexpected that it will have a much closer encounter just with its currently nearest neighbour! ",
        "conclusions": "This and other examples from our simulations show that stellar encounters within 1 ly might be more frequent than previously thought. It could explain peculiar trajectories of some stars, brown dwarves, planets and minor celestial bodies. For instance, Proxima's trajectory with respect to Alpha Centauri AB might be the result of one or several close encounters. Such encounters could even exlain Sedna's trajectory in the solar system!"
    },
    "1004/1004.1885_arXiv.txt": {
        "abstract": "We present a visual determination of the number of bright points (BPs) existing in the quiet Sun, which are structures though to trace intense kG magnetic concentrations. The measurement is based on a  0\\farcs 1 angular resolution G-band movie obtained  with the Swedish Solar Telescope at the solar disk center. We find 0.97 BPs~Mm$^{-2}$, which is a factor three larger than any previous estimate. It corresponds to  1.2 BPs per solar granule. Depending on the details of the segmentation, the BPs cover between 0.9\\,\\% and 2.2\\,\\% of the solar surface. Assuming their field strength to be 1.5 kG, the detected BPs contribute to the solar magnetic flux with an unsigned flux density between 13~G and 33~G. If network and inter-network regions are counted separately, they contain 2.2 BPs~Mm$^{-2}$ and 0.85  BPs~Mm$^{-2}$, respectively. ",
        "introduction": "\\label{intro}  Our understanding of the quiet Sun magnetic magnetic fields has turned over during the last decade. Traditional magnetograms showed magnetic signals only at the network boundaries, occupying only a small fraction of the solar surface, and mainly produced by spatially unresolved kG magnetic concentrations \\citep[e.g.,][]{bec77b,sol93}. Nowadays magnetic signals are detected almost everywhere \\citep[e.g.,][]{dom03a,har07,lit08}, and there is no doubt on the presence of a volume filling component of dG and hG fields as inferred from the Hanle depolarization signals \\citep[e.g.,][]{fau93,tru04}. The intra-network (IN) magnetic fields have passed from hardly delectable signals in long-integration low spatial resolution magnetograms \\citep[e.g.,][]{mar88,wan95}, to ubiquitous features \\citep[e.g.,][]{lin99,san00}.% The abundance of new magnetic structures makes them  potentially important to understand the global magnetic properties of the Sun \\citep[][]{san04,tru04}, and also makes it unlikely that the quiet Sun magnetism results from the decay of active regions \\citep[e.g.,][]{san03b}. Since the quiet Sun magnetic fields do not seem to be debris from active regions, they are not automatically generated by the dynamo responsible for the solar cycle. A different production mechanism seems to be at work, probably an efficient turbulent dynamo driven by granular convection \\citep{pet93,cat99b}, an hypothesis corroborated by the latest numerical simulations of granular magneto-convection \\citep[][]{vog07,pie10}. The sensitivity of the current polarimeters allow us to detect weak linear polarization signals produced by the transverse component of the magnetic fields \\citep[][]{har07,lit08}, and to find out that they often emerge as short loops pooping up from the sub-photosphere \\citep[][]{cen07,mar09}. In agreement with the turbulent dynamo scenario, quiet Sun magnetic fields come with strengths in the full range from basically zero to 2\\,kG \\citep[][]{san00,dom06}. Even if they only fill a small fraction of the quiet photosphere, the part having strong kG fields may be particularly important from a physical point of view. Magnetic flux and magnetic energy scale as powers of the field strength, therefore, depending on their (uncertain) area coverage, kGs may surpass the contribution of  the more common but weaker fields \\citep[][]{san04}. Moreover, buoyancy makes kG concentrations vertical \\citep[e.g.,][]{sch86} and so, they can provide a mechanical connection between the photosphere and the upper atmosphere \\citep[e.g.,][]{bal98,san08}. Theoretical studies support the traditionally-ignored but potentially-important influence of the quiet Sun photosphere on the chromosphere, transition region, and corona \\citep[e.g.][]{sch03b,goo04,jen06}. Despite all these (and other) significant advances, our understanding of the quiet Sun magnetism is still incomplete. Each observational improvement reveals more magnetic structures, in a process showing no sign of having converged yet.   As a part of this movement to characterize the properties of the quiet Sun magnetism, we discovered the ubiquitous presence of bright points (BPs) in the intergranular lanes of the quiet Sun \\citep{san04a}. The plasma in the lanes is cold and therefore dark, but the BPs look bright because they trace intense kG magnetic concentrations \\citep[]{spr77}. The magnetic field provides most of the pressure required to maintain the structure in mechanical balance, which reduces the gas pressure, the density and, thus, the opacity. The kG concentrations resemble holes carved on the solar surface which allow us to peep up into the sub-photosphere, where deeper usually means hotter and brighter. This picture is confirmed by realistic numerical simulations \\citep[e.g.,][]{car04,kel04}, and thus BPs are regarded as proxies for kG magnetic concentrations. The presence of BPs everywhere on the quiet Sun has been verified in several works, e.g., \\citet[][]{dew05,dew08,bov08,san08}. However, none of these works find as many BPs as \\citet[][0.3 BPs\\,Mm$^{-2}$]{san04a}. This is partly due to the angular resolution of the images, taken with small telescopes, and also due to the use of automatic algorithms for BP identification which can hardly detect the faintest BPs. In an effort to calibrate  with visual detection one of these methods \\citep[][]{bov01,bov07}, we repeated the tedious but otherwise accurate analysis carried out by \\citet{san04a}. This time we use a G-band image of much better quality. The result of such detailed analysis is surprising, and we report it here. We find three times more G-band BPs than in the original study and, therefore, far more than the numbers reported in all the previous studies cited above.  Our density implies that the quiet Sun has more BPs than granules, and at least 1\\,\\% of the quiet solar surface is covered by BPs. ",
        "conclusions": "According to the current paradigm, the BPs trace intense magnetic concentrations so, our observation constrains the fraction of quiet Sun plasma having kGs. Our measurement should be regarded as a lower limit since we only include secure BPs  (\\S~\\ref{description}), the area of large BPs is underestimated (\\S~\\ref{results}),  and diverse arguments suggest that some kGs are not bright \\citep[e.g.,][]{san01,vog05,bec07}. Moreover, we miss BPs since the detectability of the small ones critically depends on the resolution \\citep[e.g.][]{tit96}, and quiet Sun BPs are often at the resolution of the observation \\citep[e.g.,][]{san04a,bov08}. In fact, differences of angular resolution partly explain the lower number density of BPs previously reported in the literature (\\S~\\ref{intro}). Another part of the difference is due to the use of automatic algorithms based on single frames to detect the BPs,  which are still unable to identify the smallest and dimmest BPs in the images (\\S~\\ref{results}).  Even a 1\\,\\% filling factor in the form of kG fields represents a big challenge. They are not yet produced by the current turbulent solar dynamo simulations \\citep[][]{vog07}, and they are overlooked in recent studies of the quiet Sun magnetic field strength distribution based on Stokes spectro-polarimetry \\citep[e.g.,][]{oro07,ase09}. Assuming that typical BPs have 1.5~kG field strengths, they contribute to the solar magnetic fields with an unsigned flux density between 13~G and 33~G depending on whether we adopt 0.9\\,\\% or 2.2\\,\\% for their filling factors. However,  \\citet{oro07} report only 9.5~G,  with field strengths dominated by hG rather than kG fields. The solution to this double puzzle may also have to do with insufficient spatial resolution. On the one hand, the amount of kG fields created by magneto-convection simulations increases with increasing spatial resolution \\cite[e.g.,][]{bus08}. On the other hand, the standard procedure of interpreting Stokes profiles does not acknowledge the existence of unresolved structuring. When it is acknowledged, then kG fields with filling factors compatible with the limits posed here are found \\citep[][]{vit10}.  Figure~\\ref{fullfov} shows BPs throughout. Their abundance can be emphasized by comparison with the number of granules. There are some 0.82~granules~Mm$^{-2}$ \\citep[e.g.,][]{mul00}, which implies that the quiet Sun has more BPs than granules, i.e., our image has 1.2~ {\\rm BPs~per~granule.} Yet another revealing reference to compare with is the largest sunspot-group ever photographed on the Sun. It seems to correspond to a complex active region observed on April 7, 1947, which occupied  $4.3\\times 10^3$ millionths of solar hemisphere \\citep{hog47}. This record-braking large sunspot yielded only $f\\simeq 0.2$\\,\\%, i.e., one would need several ARs like this one to carry the same unsigned flux as the quiet Sun BPs have at any time if they uniformly cover the solar surface."
    },
    "1004/1004.3308_arXiv.txt": {
        "abstract": "The open cluster M67 has solar metallicity and an age of about 4~Gyr. The turn-off mass is close to the minimum mass for which solar metallicity stars develop a convective core during main sequence evolution as a result of the development of hydrogen burning through the CNO-cycle. The morphology of the color-magnitude-diagram (CMD) of M67 around the turn-off shows a clear hook-like  feature,  direct  sign  that  stars  close  to  the  turn-off  have convective cores. VandenBerg et al. investigated the  possibility of using the morphology of the M67 turn-off to put constraints on the solar metallicity, particularly CNO elements, for which solar abundances have been revised downwards by more than 30\\% over the  last few years. Here, we extend their  work filling in the gaps in their analysis. To this aim, we compute isochrones appropriate for M67 using new (low metallicity) and old (high metallicity) solar abundances and study whether the characteristic turn-off in the CMD of M67 can be reproduced or not. We also study the importance of other constitutive physics on determining the presence of such a hook, particularly element diffusion, overshooting  and nuclear reaction  rates. We  find that  using the  new solar abundance determinations, with low CNO abundances, makes it more difficult to reproduce the characteristic  CMD of M67.  This result is  in agreement with  results by VandenBerg et al. However, changes in the constitutive physics of the models, particularly overshooting, can influence and alter this result to the extent that isochrones constructed with models using low CNO solar abundances can also reproduce the turn-off morphology in M67. We conclude that only if all factors affecting the turn-off morphology are completely under control (and this is not the case), M67 could be used to put constraints on solar abundances. ",
        "introduction": "\\label{s:intro}  The latest revision of the solar abundances by \\citet[AGS05]{ags:2005} resulted in a drastic reduction of the abundances of carbon, nitrogen, and oxygen, and thus in the total metallicity of the Sun. While $Z/X$ in   the    previous   standard   references    by   \\citet[GN93]{gn:93}   and \\citet[GS98]{gs:98} was 0.0245 respectively 0.0230, \\citet{ags:2005} determined a value of 0.0165. The recently published complete re-analysis of all elements by \\citet[AGS09]{ags:2009} corrected this number slightly upwards to 0.0183. It is interesting to realize that $Z/X$ has constantly gone down with time since \\citet{ag:89}, who gave a value of 0.0275.  While the analysis by  \\cite{ags:2005} and in particular by \\cite{ags:2009} is undoubtedly the most complete, coherent and technically advanced one, the results are far from being accepted. The reason is that standard solar models (SSM), calculated for the new metal abundances, agree much less with well established and accurately determined results from helioseismology: the sound speed and density profile, the depth of the convective envelope, and its helium content. \\citet{bbps:2005} investigated the new ``solar model problem'', which has been confirmed in many subsequent, independent publications. \\citet{ab:2006} further demonstrated that helioseismic data require  the solar envelope to have a composition close to the old photospheric abundances, and \\citet{csbev:2007} used low-degree p-modes penetrating to the solar core to arrive at a similar conclusion.  To cure the ``trouble in paradise'' \\citep{ags:2005} attempts were undertaken to confirm the previous, higher abundances of intermediate elements, in particular of oxygen \\citep[e.g.\\ ][]{csn:2008,clsab:2008}. On the other hand, the lower abundances were recovered, too (\\citealt{mcbb:2009} for nitrogen; \\citealt{snn:2007} for oxygen).  On the side of the solar model community, possible or necessary changes to the constitutional physics were discussed to restore the previous excellent agreement with helioseismology (see \\citealt{guzik:2008} for a summary). In particular, an increase of opacities has been considered the most promising approach \\citep{ba:2004,bbs:2005,ab:2006,cddmhp:2009}. But updated Rosseland opacities for the  solar interior are not enough to  solve the solar abundance problem  \\citep{bsb:2005}. Also, a  postulated upward  correction to  the neon abundance has not found observational support \\citep{young:2005}.  It is therefore reasonable to shift attention also to other stellar objects, since any revision of the solar metallicity yardstick leads to a corresponding change in the metallicity of all stars. This is in particular true for stars with $\\mathrm{[Fe/H]}\\approx 0$, since their metallicity is often determined in a strictly differential and therefore quite accurate way. \\citet{algdc:2007} found that with the lower AGS05 solar composition they could reproduce the properties of the pre-main sequence binary system RS~Cha, while it was impossible to fulfill all observational constraints with the higher GN93 abundances.  Another obvious test was performed by \\citet[VG07]{vgeef:2007}, using the CMD morphology of the open cluster M67, which has a metallicity of $\\mathrm{[Fe/H]}= 0.00\\pm0.03$ \\citep{rsppp:2006} and a turn-off (TO) age around  4~Gyr \\citep[e.g.\\ ][]{pietr:04}. At this age the turn-off mass is slightly higher than $1\\msun$ and close to the critical mass at which a convective core on the main-sequence appears due to the dominance of the CNO-cycle over the pp-chain \\citep{kw:90}. The CMD of M67 \\citep{mmm67:93,esm67:2004} shows a clear hook-like structure at the turn-off, indicative of a convective core for the turn-off stars. Since the efficiency of the CNO-cycle is directly proportional to the abundance of these elements, the reduced solar abundances of AGS05 could possibly result in a radiative core for stars at the TO, although in the fitting process the change in luminosity  (at given mass,  luminosity is  lower for  the AGS05  mixture) and effective temperature complicates simple predictions about the TO-mass and core structure. Indeed, VG07 found that the TO-hook disappeared for the new abundances, thereby strengthening the case for the GS98 metallicity scale.  However, VG07 did not fail to mention some caveats: Their models were calculated without atomic diffusion taken into account, a physical process that at least for solar models -- irrespective of the metallicity scale used -- is essential for the best possible agreement with seismic inferences. They correctly pointed out that diffusion helps to support a convective core at a stellar mass lower than in models ignoring this effect \\citep{mrrv:2004}, mainly due to the increase of CNO abundances in the stellar core by gravitational settling. The different  ages of M67 and  the Sun, but  identical present-day abundances could also imply different initial abundances, an effect that should be taken into account in careful and precise studies. Apart from these points, one should keep in mind that with the occurrence of a convective core the question about the amount of overshooting arises. Its treatment, and in particular the value of any free parameter in its practical implementation, cannot simply be taken for granted, as usually such ``calibration'' was based on stellar models resting on the higher solar metallicity scale. Finally, the occurrence of a convective core may also depend on other aspects and therefore any conclusion concerning the metallicity is valid only under the assumption that all other influences are under control.  Given the result by VG07 and the open questions raised above, we revisited M67, trying to complete and extend the pioneering study of VG07. In \\S~\\ref{s:moddat} we will introduce briefly the stellar evolution code we have used for  most of the calculations in  the present work, and  those aspects of the input physics, besides the solar metallicity, we will also investigate. In \\S~\\ref{s:fits} we demonstrate that we can recover the results of VG07, a fact that is not irrelevant in the light of the following section. We then present our preferred model, showing that in this case the CMD morphology of M67 can be reproduced for both solar abundance scales. In \\S~\\ref{s:fitn} we investigate, how other approaches and variations of the input physics may also influence the quality of  isochrone fits, and  also possible systematic  differences arising from the use of different stellar evolution codes. Our conclusions will follow in \\S~\\ref{s:conc}. ",
        "conclusions": "\\label{s:conc}   The open  cluster M67  has solar metallicity  and its  CMD shows a  clear hook around  the  TO,  evidence  that  stars  populating  the  TO  have  convective cores. The  mass of  these stars  is around 1.2~M$_\\odot$,  very close  to the critical mass  value for which a convective core developes  as a result of the efficient operation of  the CNO-cycle. Based on this,  VG07 have suggested that M67 can  be used to test the new revisions  of solar abundances presented by  Asplund and collaborators  \\citep{ags:2005,ags:2009}, since  CNO elements, catalysts in the CNO-cycle, have suffered the largest reductions in their solar abundance values. Here, we have reconsidered the viability of using M67 as a benchmark for solar abundances, by studying  in detail under which conditions stars with masses around 1.2~M$_\\odot$ develop (or not) a convective core towards the end of the  main sequence and how  this affects the morphology  of isochrones best fitting M67 CMD.  We  have first  focused on  reproducing the  results obtained  by  VG07. Using initial compositions for both AGS05  and GS98 obtained from calibrating solar models without  element diffusion, we  have successfully recovered  their basic results: if stellar models include a small amount of overshooting but no element diffusion,  then the hook in M67 is  not reproduced by isochrones computed with  the AGS05 solar  composition, while it  is present if  the GS98 composition is used  since, in this case, stars populating the  TO have a well developed convective core. This has been  presented by VG07 as support for the GS98 solar  composition; although  VG07 pointed out  to some caveats  in their models that may influence this conclusion.  We  then included  element diffusion  in our  models. On  the one  hand, solar models   calibrated   with   diffusion   yield   a   higher   solar   initial metallicity,  as summarized in  Table~\\ref{t:ssm}. Additionally, because gravitational settling  is the dominant  effect, metals tend to  accumulate in the core. Inclusion of diffusion has then a double effect towards compensating lower metallicities. Our  models with diffusion,  but still  using NACRE nuclear rates, lead to the formation of  a convective core for TO stars in M67 even  with  the AGS05  composition  and,  consequently,  the distinctive  hook observed  in  the  CMD  is  also  present  in  the  isochrones,  as  shown  in Figure~\\ref{f:ssmdiff}.  Under  these assumptions,  we find no  clear evidence favoring the older, higher, GS98 solar metal abundances over AGS05.  In our study, we have also identified two nuclear reactions, $\\Iso{14}{N}(p,\\gamma)\\Iso{15}{O}$   and   $\\Iso{17}{O}(p,\\alpha)\\Iso{14}{N}$, that are critical in determining the precise value of $M_\\mathrm{ccc}$. The first reaction is the bottleneck of the whole CNO-cycle, while the second one is the bottleneck of the CNO-II part. When the temperature is not high enough, the latter acts as a {\\em  sink} for $\\Iso{14}{N}$ nuclei, thereby slowing down the whole CNO-cycle.  Let us first  discuss results  related to  the first  reaction. In addition to the NACRE rate, we also used for $\\Iso{14}{N}(p,\\gamma)\\Iso{15}{O}$  the latest measurement  by the  LUNA group \\citep{mfn14:2008}, which is about a factor of 2 lower than the NACRE rate. As can be expected, the LUNA rate shifts $M_\\mathrm{ccc}$ to  larger values, by about 0.08~M$_\\odot$.  With this rate, even models with GS98 composition (not including diffusion) lack the characteristic hook observed in  M67 (Figure~\\ref{f:vg07n14}). This is not surprising in the light of results with AGS05  composition, since the effective operation of the CNO-cycle depends on the product of the abundance of catalysts and the $\\Iso{14}{N}(p,\\gamma)\\Iso{15}{O}$ rate. The effect on the CNO-cycle of reducing the rate for this reaction by about a factor of 2, for stars with masses close to $M_\\mathrm{ccc}$ is very similar as to reducing the total number of catalysts by a similar amount. A somewhat more unexpected result relates  to  the   second  reaction,  $\\Iso{17}{O}(p,\\alpha)\\Iso{14}{N}$.  The branching  between  the CNO-I  and  CNO-II  cycles  is almost  independent  of temperature for the temperature range that interests us in this work.  However, if  the temperature  is not high enough for  the CNO-II cycle to be fully active,  $\\Iso{17}{O}$ is  created at  the  expense of $\\Iso{14}{N}$ but  the feedback  to the  CNO-I cycle (the  only  one relevant  in terms  of energetics) is inefficient. As a consequence, the  total number of catalysts is reduced, with $\\Iso{14}{N}$ being stored in the form of $\\Iso{17}{O}$. For this rate we  have used, in addition to the NACRE  rate, the older rate by \\citet{cf:1988} and the new measurement by \\citet{moaz:07}. When we use the \\citet{cf:1988} rate, that is about 10\\% of the NACRE rate in the relevant temperature range (see Figure~\\ref{f:o17rate} for a comparison of the rates), we find that, even for GS98 and the NACRE rate for  $\\Iso{14}{N}(p,\\gamma)\\Iso{15}{O}$ there is no convective  core in a 1.2~M$_\\odot$ models  towards the  end of the  main sequence.  This translates into the absence of a hook in  the isochrones that would best fit M67 CMD. One can regard the last exercise as merely of academic interest since the low rate from  \\citet{cf:1988}  seems  now to be ruled  out  by  experiments.  However, the importance  of this  rate is  still worth  being taken  into account  with the newest results by \\citet{moaz:07}, even  if this rate agrees with NACRE within 40\\%.   By  using  the \\citet{moaz:07}  rate  in   the  models,  the 1.2~M$_\\odot$ model  with the  LUNA rate for the $\\Iso{14}{N}(p,\\gamma)\\Iso{15}{O}$ reaction recovers a small convective core. These results are summarized in Figure~\\ref{f:o17track}.  The development of convective cores in stars close to $M_\\mathrm{ccc}$ is sensitive to the detailed physical input in the models, as seen from the above results. Moreover, an appreciable amount  of freedom in the modeling is still present because  of the  lack of  a proper convection  theory to  account for, particularly,  the amount  of overshooting  occurring in  stellar cores.  It is widely accepted, however, that an overshooting region of about 0.2~$H_P$ is  needed in stars  with masses  above $\\sim  2$~M$_\\odot$ to  explain, among others, the width of main sequences in  stellar clusters. On the other hand, for stars below that mass, and particularly close to $M_\\mathrm{ccc}$, it is also known that overshooting has  to be limited to a smaller  region. However, a detailed understanding of  how core overshooting depends of  stellar parameters, namely the stellar mass, is not known. We have used here two different approaches to limit  overshooting,   a  geometric  cutoff   and  an efficient factor that increases linearly with stellar mass, Equations~\\ref{eqn:ove2}~and~\\ref{eqn:ove3} respectively, with  the geometric cutoff  being the  standard  choice.  In relation  to  core overshooting,  its importance is that it helps to reproduce the hook  in M67 because it facilitates stars to sustain a convective core during the main sequence and the growth of the core towards  the end  of the  main sequence  for stars with  masses around $M_\\mathrm{ccc}$, critical for explaining the TO morphology in M67. On the other  hand, overshooting cannot, by itself,  {\\em create} a convective core when there  is none.  This is  the case, for example, for our models with AGS05 composition and the LUNA  rate for the $\\Iso{14}{N}+p$ reaction when the geometric cutoff is applied, as shown in the left panel of Figure~\\ref{f:calibn14}.  In this case,  the convective core developed towards the end of the pre-main sequence evolution cannot be  sustained by stars with masses close to the TO mass in M67, even with the help of overshooting. This may be the result of our geometric cutoff being {\\em too restrictive}. When, instead, the linear ramp function  of Equation~\\ref{eqn:ove3} is  used, isochrones based  on AGS05 show   again    the   characteristic   hook    of   M67   (right    panel   in Figure~\\ref{f:calibn14}).   With  this   prescription  for  overshooting,  for stellar masses around  M67 TO masses, the convective  core survives during all main sequence evolution and  naturally grows  towards the  end of it,  due to  increased CNO burning, leading to the occurrence of a hook in M67 CMD. Overshooting provides a certain degree  of freedom in stellar models, and here we show that this freedom renders isochrone fitting to M67 uncertain,  in that depending on  the way overshooting is  modeled, we are able to find, or not, a hook around  the TO. It is important to note that both prescriptions for the limitation of overshooting for stellar masses around $M_\\mathrm{ccc}$ can be made consistent with current observational evidence.  We have  also tried to assess the possible  dependence of our results on the numerics, that is the  stellar evolution code. Therefore, we have used,  in addition  to  GARSTEC, two  other  evolution codes:  LPCODE and  the Dartmouth  code.  None of  these  codes share  a  common  background in  their development. For  this work, when possible,  we have tried to  match the input physics, but some  differences still remain. The most  important, probably, is that in the Dartmouth code overshooting  is accounted for as a fraction of the pressure scale height. For models  appropriate for M67, this fraction has been taken as 0.07~$H_P$. Comparison of results between the three codes is   highly   satisfactory.   Solar   model   callibrations,   summarized   in Table~\\ref{t:ssm},  yield   very  small  differences.   A  comparison  between evolutionary   tracks  computed   with  LPCODE   and  GARSTEC   is   shown  in Figure~\\ref{f:laplata}. The  agreement between results  from the two  codes is excellent. A comparison of isochrone fitting to the CMD of M67 with GARSTEC and the Dartmouth code is shown in Figure~\\ref{f:dartm} for both compositions used in this work. Again, results are similar with both codes, which are able to reproduce the observed morphology around the TO. For each composition, both codes also yield the same age for the cluster, 4.5 and 4.2~Gyr for AGS05 and GS98 compositions respectively. We conclude from these tests that systematics originating in the use of different numerical schemes, and even some differences in the implementation of the input physics, do not affect our conclusions.  Finally, we comment on two additional points that reinforce our main conclusions given below. As stated in \\S~\\ref{s:data}, the uncertainty in the metallicity of M67, according to \\citet{gm67:2000} is about 0.06~dex, i.e. 15\\%. We have not played with the metallicity of our models and have always assumed a solar value (where solar means GS98 or AGS05). However, in view of the dependence of the precise value of $M_\\mathrm{ccc}$ to the details of the constitutive physics, variations in the metallicity of our models can in principle affect the occurrence or not of a convective core for stars with M67 TO-mass. The last point is that in the new determination of solar abundances by \\citet{ags:2009} CNO elements have been revised upwards. The total number of CNO catalysts in the new \\citet{ags:2009} is about 8\\% higher than in the previous AGS05 solar abundance compilation. This increase would also contribute to facilitate the occurrence of a convective core at the TO-mass in M67.  The main conclusions we draw from our study are as follow. The different solar compositions, namely AGS05 and GS98  certainly have an impact on the predicted morphology of  the CMD in M67. The occurrence of the TO-hook is more difficult to achieve for the low CNO abundances of AGS05. This  is a confirmation of previous results by VG07. However, we  have also  found that  other constitutive physics  in the models,  e.g.  element  diffusion,  nuclear reactions,  prescription  of  core overshooting, also influence the stellar  mass at which convective cores start to develop. This translates, in the case of M67, into isochrones that may have a hook around the TO even with the AGS05 composition or, on the contrary, that show  no  hook  even  for  the  GS98 composition.  There  is  a  certain  {\\em degeneracy} in the constitutive physics  underlying the presence of the hook in   M67  that  cannot   be,  at   the  present   status  of   our  knowledge, disentangled. The morphology  of M67 could in principle be  used to test solar abundances,  but only  under  the  strong assumption  that  all other  factors affecting the TO morphology, more precisely the occurrence of a convective core in stars with masses around the TO-mass, are completely under control. Since we understand this is not the  case, we conclude that M67 CMD morphology does not present a strong argument against low CNO abundances in the Sun."
    },
    "1004/1004.4002_arXiv.txt": {
        "abstract": "We report on new measurements of the luminosity function (LF) and mass function (MF) of field low--mass dwarfs derived from Sloan Digital Sky Survey (SDSS) Data Release 6 (DR6) photometry.  The analysis incorporates $\\sim$ 15 million low--mass stars (0.1 $\\msun <$  $\\cal{M}$  $<  0.8 \\msun$), spread over 8,400 square degrees.   Stellar distances are estimated using new photometric parallax relations, constructed from $ugriz$ photometry of nearby low--mass stars with trigonometric parallaxes.  We use a technique that simultaneously measures Galactic structure and the stellar LF from $7 < M_r < 16$.  We compare the LF to previous studies and convert to a MF using the mass--luminosity relations of \\cite{2000A&A...364..217D}.  The system MF, measured over  $-1.0 <$ log ${\\cal M} / \\msun$ $< -0.1$, is well--described by a log--normal distribution with $\\cal{M}_{\\circ}$ = 0.25 $\\msun$.  We stress that our results should not be extrapolated to other mass regimes.  Our work generally agrees with prior low--mass stellar MFs and places strong constraints in future star--formation studies of the Milky Way. ",
        "introduction": "Low--mass dwarfs (0.1 $\\msun <$ $\\cal{M}$  $<  0.8 \\msun$) are, by number, the dominant stellar population of the Milky Way.   These long-lived \\citep{1997ApJ...482..420L} and ubiquitous objects comprise $\\sim70\\%$ of all stars, yet their diminutive luminosities ($L \\lesssim 0.05 L_{\\odot}$) have traditionally prohibited their study in large numbers.  However, in recent years, the development of large--format CCDs has enabled accurate photometric surveys over wide solid angles on the sky, such as the Two--Micron All Sky Survey \\citep[2MASS; ][]{2006AJ....131.1163S} and the Sloan Digital Sky Survey \\citep[SDSS; ][]{2000AJ....120.1579Y}.  These projects obtained precise ($\\sigma \\lesssim 5\\%$) and deep ($r \\sim 22, J \\sim 16.5$) photometry of large solid angles ($\\gtrsim 10^{4}$ sq$.$ deg$.$).   The resulting photometric datasets contain millions of low--mass stars, enabling statistical investigations of their properties. In particular, 2MASS photometry led to the discovery of two new spectral classes: L and T \\citep{1999ApJ...519..802K, 2002ApJ...564..421B}, and was used to trace the structure of the Sagittarius dwarf galaxy \\citep{2003ApJ...599.1082M} with M giants.  SDSS data led to the discovery the first field methane brown dwarf, which was the coolest substellar object known at the time of its discovery \\citep{1999ApJ...522L..61S}.  Other notable SDSS results include the discovery of new stellar streams in the Halo \\citep[e.g.][]{2003ApJ...588..824Y,2006ApJ...642L.137B} and new Milky Way companions \\citep[e.g.,][]{2005AJ....129.2692W,2007ApJ...654..897B}, as well as unprecedented $in$ $situ$ mapping of the stellar density \\citep[][]{2008ApJ...673..864J} and metallicity \\citep{Ivezic08} distributions of the Milky Way and confirmation of the dual-halo structure of the Milky Way \\citep{2007Natur.450.1020C}.  SDSS has proven to be a valuable resource for statistical investigations of the properties of low--mass stars, including their magnetic activity and chromospheric properties \\citep{2004AJ....128..426W, 2008AJ....135..785W}, flare characteristics \\citep{2009arXiv0906.2030K} and their use as tracers of Galactic structure and kinematics \\citep{2007AJ....134.2418B, 2009AJ....137.4149F}.  Despite the advances made in other cool--star topics, two fundamental properties, the luminosity and mass functions, remain uncertain.  The luminosity function (LF) describes the number density of stars as a function of absolute magnitude ($\\Phi(M) = dN / dM$).  The mass function (MF), typically inferred from the LF, is defined as the number density per unit mass ($\\psi({\\cal M}) = dN / d{\\cal M}$).   For low--mass stars, with lifetimes much greater than the Hubble time, the observed present-day mass function (PDMF) in the field should trace the initial mass function (IMF).  Following \\cite{1955ApJ...121..161S}, the IMF has usually been characterized by a power--law $\\psi({\\cal M}) = dN / d{\\cal M} \\propto {\\cal M}^{- \\alpha}$, with the exponent $\\alpha$ varying over a wide range, from 0.5 to 2.5.  However, some studies have preferred a log--normal distribution.  Previous investigations are summarized in Table \\ref{chap6:table:field_mfs} (MF) and Table \\ref{chap5:table:existing_studies} (LF), which show the total number of stars included and solid angle surveyed in each study.  Present uncertainties in the LF and MF can be attributed to disparate measurement techniques that are a result of trade--offs in observing strategy.  Previous investigations of the LF and MF have fallen in one of two categories:  nearby, volume--limited studies of trigonometric parallax stars; or pencil--beam surveys of distant stars over a small solid angle.  Tables \\ref{chap6:table:field_mfs} \\& \\ref{chap5:table:existing_studies} detail the techniques used by modern investigations of the field LF and MF.  In both cases, sample sizes were limited to a few thousand stars, prohibiting detailed statistical measurements.  The solid angle shown in each table distinguishes the two types of surveys. There is considerable disagreement between the nearby, volume--limited investigations \\citep{1997AJ....113.2246R} and pencil beam studies of distant stars \\citep{1998AJ....116.2513M, 2001ApJ...555..393Z, 2006AA...447..185S}.  It has been suggested that this discrepancy is due to the presence of unresolved binary stars in the pencil beam surveys \\citep{1993MNRAS.262..545K, 2003ApJ...586L.133C}, but this has not been shown conclusively \\citep{1997AJ....113.2246R}.  We investigate the effect of unresolved binaries in \\S \\ref{sec:binaries}.  Using a sample drawn from SDSS, 2MASS and Guide Star Catalog photometry, and supplemented with SDSS spectroscopy, \\cite{covey08} performed the largest field low--mass LF and MF investigation to date.  Covering 30 sq.\\ deg.\\ and containing $\\sim30,000$ low--mass stars, their sample measured the LF using absolute magnitudes estimated from photometric colors, and quantified the contamination rate by obtaining spectra of every red point source in a 1 sq.\\ deg.\\ calibration region.  The \\cite{covey08} sample serves as a calibration study for the present work, as it quantified the completeness, bias and contamination rate of the SDSS and 2MASS photometric samples.  While their study focused on a limited sightline, the current investigation expands to the entire SDSS footprint, increasing the solid angle by a factor of $\\sim300$.  In \\S \\ref{sec:observations}, we describe the SDSS photometry used to measure the field LF and MF.  The color--absolute magnitude calibration is discussed in \\S \\ref{sec:phot_pi}.  In \\S \\ref{sec:analysis}, we introduce a new technique for measuring the LF of large, deep photometric datasets, and compare to previous analyses.  The resulting ``raw'' LF is corrected for systematic effects such as unresolved binarity, metallicity gradients and changes in Galactic structure in \\S \\ref{sec:correction}.  The final LF and our MF is presented in \\S \\ref{sec:LF} and \\S \\ref{sec:MF}.  Our conclusions follow in \\S \\ref{sec:conclusions}.   \\begin{turnpage} \\begin{deluxetable*}{llllllll} \\tablewidth{0pt}  \\tablecaption{Major Low--Mass Field IMF Studies} \\tabletypesize{\\scriptsize}  \\tablehead{ \\colhead{Authors} & \\colhead{$N_{\\rm Stars}$} & \\colhead{$\\Omega$ (sq. deg.)} & \\colhead{Filter(s)} & \\colhead{Depth} & \\colhead{Mass Range} & \\colhead{$\\alpha$, ${\\cal M}_{\\circ}$}& \\colhead{Notes} } \\startdata \\cite{1955ApJ...121..161S} & \\nodata    &  \\nodata     &    $V$     &   \\nodata      &  $0.3\\msun$ - 10 $\\msun$   & $ \\alpha =$ 2.35   & Compiled LFs from contemporaries\\tablenotemark{a} \\\\ \\cite{1979ApJS...41..513M} & \\nodata    &  \\nodata     &   $V$      &   \\nodata      &   $0.1\\msun$ - 60 $\\msun$  &  ${\\cal M}_{\\circ} \\simeq 0.1 \\msun$  & Log-normal fit, Compilation of 3 field LFs\\tablenotemark{b}  \\\\ \\cite{1990MNRAS.244...76K} & \\nodata    &  \\nodata     &   $V$      &   \\nodata      &  $0.1\\msun$ - 0.9 $\\msun$  &     ${\\cal M}_{\\circ} \\simeq 0.23 \\msun$, $\\alpha = 0.70$ & Adopted LFs of \\cite{1986FCPh...11....1S} and \\cite{1989MNRAS.238..709S}  \\\\ \\cite{1993MNRAS.262..545K} & \\nodata    &   \\nodata    &  $V,I$     &  \\nodata       &   $0.08\\msun$ - 0.5 $\\msun$     &   0.70 $< \\alpha < 1.85$       & LF from \\cite{1983nssl.conf..163W} and \\cite{1989MNRAS.238..709S}  \\\\ \\cite{1993ApJ...414..279T} &  3,500   &  280 deg$^2$   &   $I,K$  &  $I \\lesssim 17.5$   &  $0.1\\msun$ - 0.5 $\\msun$   &  \\nodata       & Turnover at 0.25 $\\msun$  \\\\ \\cite{1997AJ....113.2246R} &  151     &  $\\delta~>~-30^\\circ$  & $V,I$   & $d < 8 $pc &   $0.08\\msun $- 1.2 $\\msun$     & $ \\alpha =$ 1.2    &  Solar Neighborhood\\tablenotemark{c} \\\\ \\cite{1998AJ....116.2513M} & 1,500 & 0.83 deg$^2$  & $V,R$   & $V \\lesssim 23.5$   & $0.1 \\msun$ - 0.6 $\\msun$ & 1.3 &   \\\\ \\cite{2001ApJ...555..393Z}  &  $\\sim 1,400$   &  $\\sim 0.4$ deg$^2$     &   $V,I$      & $18  \\lesssim I \\lesssim 24$     &   $0.1 \\msun$  - 0.6 $\\msun$    & $ \\alpha =$ 0.45  & HST observations   \\\\ \\cite{2002Sci...295...82K} &  \\nodata   &  \\nodata     &  $V,I$       &  \\nodata       &   $0.08\\msun$ - 0.50 $\\msun$     &  $ \\alpha =$ 1.3      &  Compiled contemporary LFs\\tablenotemark{d}  \\\\ \\cite{2002AJ....124.2721R} &  558   &  3$\\pi$ ster.    &   $B,V$  &  $d \\lesssim 20$pc &   $0.1 \\msun$ - 3.0 $\\msun$     &  $\\alpha \\simeq 1.3$  &  Solar Neighborhood Survey  \\\\ \\cite{2003PASP..115..763C} &  \\nodata   &  \\nodata      &   $V,K$      &  \\nodata    &   $0.1 \\msun$ - 1.0 $\\msun$     &   ${\\cal M}_{\\circ} = 0.22 \\msun$      & Review of contemporary field LFs\\tablenotemark{e} \\\\ \\cite{2006AA...447..185S} & 3,600  &  $\\sim3$ deg$^2$     &   $r^{\\prime},i^{\\prime}$   &  $i^{\\prime} \\sim 21$       &  ${\\cal M} <$ 0.25 $\\msun$      &  $ \\alpha =$ 2.5    & CFHT MegaCAM observations  \\\\ \\cite{covey08} &  $\\sim29 \\times 10^3$   &   30 deg$^2$    &   $i,J$      &   $J=16.2$      &    $0.1\\msun$ - 0.8 $\\msun$    & ${\\cal M}_{\\circ} = 0.29 \\msun$   & Matched SDSS \\& 2MASS observations  \\\\ This Study & $\\sim15 \\times 10^6$    &  8,400 deg$^2$     &  $r,i,z$       &  $16 < r < 22$       &  $0.1\\msun$ - 0.8 $\\msun$      &  ${\\cal M}_{\\circ} = 0.25 \\msun$        &  SDSS observations\\\\ \\enddata  \\tablenotetext{a}{Salpeter averaged luminosity functions from \\cite{1925PGro...38D...1V,1936PGro...47....1V} and \\cite{1939POMin...2..121L,1941NYASA..42..201L}.} \\tablenotetext{b}{Their adopted LF was averaged from the LFs of \\cite{1966VA......7..141M}, \\cite{1968MNRAS.139..221L} and \\cite{1974HiA.....3..395W}.} \\tablenotetext{c}{The ``8 parsec'' sample was compiled by \\cite{1997AJ....113.2246R}, with later additions from \\cite{1999ApJ...521..613R}, \\cite{2003AJ....125..354R}, and \\cite{2007AJ....133..439C}.} \\tablenotetext{d}{\\cite{2002Sci...295...82K} presents a comprehensive summary of MFs derived from the field and clusters over a wide mass range.  For low-mass stars in the field, he refers to\\\\ \\cite{1999ApJ...521..613R}, \\cite{1999ApJ...526L..17H}, \\cite{2001ApJ...554.1274C} and \\cite{2001ApJ...555..393Z}.} \\tablenotetext{e}{\\cite{2003PASP..115..763C} compared the  LFs of \\cite{1986AJ.....91..621D}, \\cite{1990ApJ...350..334H} and \\cite{2001ApJ...555..393Z}.}  \\label{chap6:table:field_mfs} \\end{deluxetable*} \\end{turnpage}  \\begin{turnpage} \\begin{deluxetable*}{lllllll} \\tablewidth{0pt}  \\tablecaption{Major Low-Mass Stellar Field LF Studies} \\tabletypesize{\\scriptsize }  \\tablehead{ \\colhead{Authors} & \\colhead{$N_{\\rm Stars}$} & \\colhead{$\\Omega$ (sq. deg.)} & \\colhead{Filter(s)} & \\colhead{Depth} & \\colhead{Distance Method} & \\colhead{Notes} } \\startdata \\cite{1989MNRAS.238..709S} & 178    &  18.88                     & $V,I$ & $ I < 16$       & Phot. $\\pi$ & $VI$ photometry        \\\\ \\cite{1994AJ....108.1437H} & 92     &  $\\delta > -25^\\circ$      & $V$   & $d = 5$ pc      & Trig. $\\pi$ &  Spec., CSN3 photometry\\\\ \\cite{1995AJ....110.1838R} & 520    &  $\\delta > -30^\\circ$      & $V$   & $d \\sim 20$pc   & Trig. $\\pi$ &  Spec., CNS3 photometry\\tablenotemark{a}\\\\ \\cite{1997AJ....113.2246R} & 151    &  $\\delta > -30^\\circ$      & $V$   & $d = 8$ pc      & Trig. $\\pi$ &  Spec., CNS3 photometry\\tablenotemark{b}\\\\ \\cite{1998AJ....116.2513M} & 4,005  & 0.83                       & $V$   & $V \\sim 23.5$   & Phot. $\\pi$ & $UBVRI$ photometry \\\\ \\cite{2001ApJ...555..393Z} & 1,413  & $\\sim 0.4$                 & $V$   & $I \\sim 26.5$   & Phot. $\\pi$ & HST photometry \\tablenotemark{c}\\\\ \\cite{2002AJ....124.2721R} & 558    &  $\\delta > -30^\\circ$      & $V$   & $d \\sim 20$pc   & Trig. $\\pi$ & Spec., CNS3 photometry \\\\ \\cite{2007AJ....133..439C} & 99     & 14,823                     & $J$   & $J \\sim 17$, $d \\sim 20$pc    & Phot. $\\pi$, Trig. $\\pi$ & Spec., 2MASS photometry   \\\\ \\cite{covey08}             & $\\sim 29 \\times 10^3$   & 29        & $J$   & $J = 16.2 $     & Phot. $\\pi$ & SDSS \\& 2MASS photometry \\\\ This Study &               $\\sim 15 \\times 10^6$     & 8,417     & $r,J$ & $r = 22$        & Phot. $\\pi$ & SDSS photometry \\\\ \\enddata  \\label{chap5:table:existing_studies} \\tablenotetext{a}{See \\cite{1991adc..rept.....G} for details on sources of photometry.} \\tablenotetext{b}{The 8 pc sample was further extended by \\cite{1999ApJ...521..613R} and \\cite{2003AJ....125..354R} and presented in $J$ by \\cite{2007AJ....133..439C}.} \\tablenotetext{c}{Some of the HST observations in this study were presented by \\cite{1996ApJ...465..759G} and \\cite{1997ApJ...482..913G}.} \\end{deluxetable*} \\end{turnpage} ",
        "conclusions": "\\label{sec:conclusions} We have assembled the largest set of photometric observations of M dwarfs to date and used it to study the low--mass stellar luminosity and mass functions.  Previous studies were limited by sample size, mostly due to the intrinsic faintness of M dwarfs.  The precise photometry of the SDSS allowed us to produce a clean, complete sample of M dwarfs, nearly two orders of magnitude larger than other studies.  To accurately estimate the brightness and distances to these stars, we constructed new photometric parallax relations from data kindly provided to us prior to publication \\citep{Golimowski08}.  These relations were derived from $ugrizJHK_s$ photometry of nearby stars with known trigonometric parallax measurements.  We compared our new relations to those previously published for SDSS observations.  We also introduced a method for measuring the LF within large surveys.  Previous LF investigations either assumed a Galactic profile (for pencil beam surveys, such as \\citealp{2001ApJ...555..393Z}) or a constant density (for nearby stars, \\citealp[i.e.][]{1997AJ....113.2246R}).  However, none of these samples have approached the solid angle or number of the stars observed in this study.  We solved for the LF and Galactic structure simultaneously, using a technique similar to \\cite{2008ApJ...673..864J}.  Our LF is measured in the $r$--band.  Using multiple CMRs, we investigated systematic errors in the LF, and computed the effects of Malmquist bias, unresolved binarity and Galactic structure changes using Monte Carlo models.  This allowed us to compare our results both to distant LF studies (which sampled mostly the system LF) and nearby star samples (which can resolve single stars in binary systems).  Finally, we computed mass functions for single stars and systems.  Low luminosity stars are more common in the single star MF, since they can be companions to any higher mass star.  We fitted both MFs with a broken power law, a form preferred by \\cite{2002Sci...295...82K}, and a log--normal distribution, which is favored by \\cite{2003PASP..115..763C}.  The log--normal distribution at low masses seems to be ubiquitous: it is evident both in the field \\citep[this study; ][]{2003PASP..115..763C} and open clusters, such as Blanco 1 \\citep{2007A&A...471..499M}, the Pleiades \\citep{2004A&A...426...75M}, and NGC 6611 \\citep{2009MNRAS.392.1034O}. The best fits for this study are reported in Tables \\ref{table:sys_mf_fits} and \\ref{table:sing_mf_fits} .  We stress the point first made in \\cite{covey08}, that comparing MF $data$ is preferable to comparing analytic fits, since the latter are often heavily swayed by slight discrepancies amongst the data.  We also caution the reader against extrapolating our reported mass function beyond  0.1 $\\msun < {\\cal M} <$ 0.8 $\\msun$, the masses that bound our sample. In the future, we plan to investigate the luminosity function in other SDSS bandpasses, such as $i$ and $z$.  Our system and single-star MFs represent the best current values for this important quantity for low--mass stars.   JJB would like to acknowledge the support of the UW Astronomy department over the course of his thesis study.  In particular, he would like to acknowledge Ivan King, Andrew Becker, Nick Cowan, Lucianne Walkowicz, Nathan Kaib and Richard Plotkin for fruitful and helpful discussions.  He also acknowledges the support and guidance of Adam Burgasser.  We also thank the anonymous referee, whose insightful comments, especially on star--galaxy separation, greatly strengthened this paper. We also gratefully acknowledge the support of NSF grants AST 02-05875 and AST 06-07644 and NASA ADP grant NAG5-13111.   We appreciate the serene and productive atmosphere fostered by Friday Harbor Labs, where this project first started.  Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S. Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS Web Site is http://www.sdss.org/.  The SDSS is managed by the Astrophysical Research Consortium for the Participating Institutions. The Participating Institutions are the American Museum of Natural History, Astrophysical Institute Potsdam, University of Basel, University of Cambridge, Case Western Reserve University, University of Chicago, Drexel University, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Joint Institute for Nuclear Astrophysics, the Kavli Institute for Particle Astrophysics and Cosmology, the Korean Scientist Group, the Chinese Academy of Sciences (LAMOST), Los Alamos National Laboratory, the Max-Planck-Institute for Astronomy (MPIA), the Max-Planck-Institute for Astrophysics (MPA), New Mexico State University, Ohio State University, University of Pittsburgh, University of Portsmouth, Princeton University, the United States Naval Observatory, and the University of Washington."
    },
    "1004/1004.4833_arXiv.txt": {
        "abstract": "A precise electromagnetic measurement of the sky coordinates and redshift of a coalescing black hole binary holds the key for using its gravitational wave (GW) signal to constrain cosmological parameters and to test general relativity.  Here we show that the merger of $\\sim 10^{6-7}M_\\odot$ black holes is generically followed by electromagnetic flares from tidally disrupted stars. The sudden recoil imparted to the merged black hole by GW emission promptly fills its loss cone and results in a tidal disruption rate of stars as high as $\\sim 0.1~{\\rm yr}^{-1}$.  The prompt disruption of a star within a single galaxy over a short period provides a unique electromagnetic flag of a recent black hole coalescence event, and sequential disruptions could be used on their own to calibrate the expected rate of GW sources for pulsar timing arrays or the proposed {\\it Laser Interferometer Space Antenna}. ",
        "introduction": "Recently, general relativistic simulations demonstrated that the coalescence of a black hole binary is accompanied by the anisotropic emission of gravitational radiation, causing a typical recoil of hundreds of ${\\rm km~s^{-1}}$ for the black hole remnant \\citep{Pre, Bak, Cam}. Binaries of supermassive black holes (SMBH) are a generic consequence of galaxy mergers \\citep{CD, Esc, May, Cal}, and the resultant gravitational waves (GWs) are potentially detectable with the proposed {\\it Laser Interferometer Space Antenna} (LISA) \\footnote{http://lisa.nasa.gov/} or existing Pulsar Timing Arrays (PTAs) such as NANOGrav \\footnote{http://nanograv.org/}.  LISA would be most sensitive to binary mergers with a total mass $M_{\\rm BH}\\sim 10^{6-7}M_{\\odot}$ \\citep{McW}, whereas the PTA sensitivity peaks at $\\sim 10^8 M_{\\odot}$ \\citep{SVV}.  PTAs have a significantly poorer localization ability, with a typical uncertainty\\citep{SV} of $\\sim 40$ compared to $\\lesssim 1$ square degrees for LISA.  These positional errors limit the precise determination of the luminosity distance to merging binaries.  An electromagnetic (EM) counterpart would greatly reduce the positional error to sub-arcsecond scales, and also determine the redshift of the source, which would enable its use as a ``standard siren'' (independent of the cosmic distance ladder) for precision measurements of the dark energy equation of state \\citep{HH, Aru, Sch, Bod}.  For these reasons, prompt EM signals are of primary importance for realizing the full potential of GW cosmology. The proposed prompt EM signals have so far all assumed the uncertain presence of a circumbinary accretion disk prior to coalescence.  Dissipation of GW energy in the disk might result in a weak EM transient shortly after the merger \\citep{KL}, re-equilibration of the inner edge of the disk could create an X-ray brightening on a timescale of $10$--$10^3 {\\rm ~yr}$ \\citep{MP}, and shocks produced by the GW-induced recoil might generate EM reverberations after the recoil \\citep{LFH} which may take $\\sim 10^4$ years to dissipate as enhanced infrared luminosity \\citep{SK}. It is not obvious whether these EM signals could be distinguished from the much more abundant sources of temporal variability in single SMBH quasars.  Moreover, the luminosity of any circumbinary disk is expected to be significantly reduced by the cavity associated with the decoupling of the binary from the inner edge of the disk in the final stage of inspiral \\citep{MP, SK}.  The disk is not expected to refill the cavity and return to its full luminosity for a time $\\sim 7(1+z)(M_{\\rm BH}/10^6M_{\\odot})^{1.32} \\rm ~yr$ after coalescence \\citep{MP}.  On longer timescales, the portion of the accretion disk that remains bound to the recoiled SMBH is expected to be detectable as a kinematically and eventually a spatially offset quasar \\citep{Loe, BSS, SB, KZL, Com, Shi}, although its lifetime is limited by the supply of gas that can remain gravitationally bound to it \\citep{BL}.  Here we show that the tidal disruption of stars provides a prompt EM flag that does not depend on the prior existence of a gaseous disk in the vicinity of the merging binary, and can result from mergers of gas-poor galaxies. Recent observations have demonstrated that tidal disruption events (TDEs) have a generic lightcurve and emission spectrum \\citep{Gez1, Gez2, Gez3, Don, Esq} that are distinguishable from normal quasar variability.  Moreover, we find that SMBH recoil results in a sequence of TDEs over a timescale of decades and potentially years, with a rate that is $\\sim 4$ orders of magnitude higher than the typical TDE rate in normal galaxies \\citep{Don}.  The existence of TDEs accompanying SMBH mergers has been studied in the past only for long time delays ($\\sim 10^6$--$10^9$yr) before \\citep{Iva, Che} or after \\citep{KM} the binary coalescence event.  While previous studies have focused on mechanisms to slowly feed stars into an empty loss cone, here we show that {\\it GW recoil will instantaneously shift the loss cone to a non-empty region of phase space}. ",
        "conclusions": "We find that merging galaxies with black hole masses $M_{\\rm BH}\\lesssim 10^7M_\\odot$ are likely to produce a tidal disruption flare a few decades after coalescence in the case of a gas-poor merger.  Multiple flares could possibly be seen on a timescale of years if resonant migration is effective. For gas-poor mergers, the peak rate is therefore at least $\\sim 10^4$ times higher than the typical TDE rate in galaxies \\citep{Don}. The total number of TDEs is maximized, and delay until the first TDE minimized, if the kick velocity is in the range of $200$--$1000 \\rm km~s^{-1}$.  Our minimal predictions concerning dry mergers could be dramatically enhanced if  resonant migration increases the number and frequency of post-kick tidal disruptions.  The case of wet mergers is substantially more complicated, and because star formation and disk migration can also significantly increase the population of the post-kick loss cone, a much more detailed study of star-disk interactions around binary SMBHs is needed to make robust predictions in the gas-rich case.  Moderate to high kick velocities in dry mergers will provide a robust EM counterpart to the GW signature of black hole coalescence within the LISA band, enabling accurate identification of the host galaxy and a precise measurement of cosmological parameters \\citep{HH} within a few decades of the initial GW signal.  With the advent of massive transient surveys, such as PTF \\footnote{http://www.astro.caltech.edu/ptf/}, Pan-STARRS \\footnote{http://pan-starrs.ifa.hawaii.edu/public/}, and LSST \\footnote{http://www.lsst.org/lsst}, it is possible that sequential tidal disruption flares could flag black hole recoil events without a GW signal, providing an independent test of the strong field regime of general relativity and a calibration of the expected event rate for LISA and PTAs.  \\bigskip"
    },
    "1004/1004.1411_arXiv.txt": {
        "abstract": "We present very deep Wide Field Camera 3 (WFC3)  photometry of a massive, compact galaxy located in the Hubble Ultra Deep Field. This quiescent galaxy has a spectroscopic redshift $z=1.91$ and has been identified as an extremely compact galaxy by \\cite{dad05}. We use new $H_{\\mathrm{F160W}}$ imaging data obtained with Hubble Space Telescope/WFC3 to measure the deconvolved surface brightness profile to $H \\approx 28$ mag arcsec$^{-2}$. We find that the surface brightness profile is well approximated by an $n=3.7$ S\\'ersic profile. Our deconvolved profile is constructed by a new technique which corrects the best-fit S\\'ersic profile with the residual of the fit to the observed image. This allows for galaxy profiles which deviate from a S\\'ersic profile.  We determine the effective radius of this galaxy: $r_e=0.42 \\pm 0.14$ kpc in the observed $H_\\mathrm{F160W}$-band. We show that this result is robust to deviations from the S\\'ersic model used in the fit.  We test the sensitivity of our analysis to faint ``wings'' in the profile using simulated galaxy images consisting of a bright compact component and a faint extended component. We find that due to the combination of the WFC3 imaging depth and our method's sensitivity to extended faint emission we can accurately trace the intrinsic surface brightness profile, and that we can therefore confidently rule out the existence of a faint extended envelope around the observed galaxy down to our surface brightness limit. These results confirm that the galaxy lies a factor $\\sim 10$ off from the local mass-size relation. ",
        "introduction": "A significant fraction of massive galaxies at $z\\approx 2$ are early-type galaxies containing quiescent stellar populations (e.g., \\citealt{fra03}; \\citealt{dad05}; \\citealt{kri06}). These galaxies must have formed very early in the universe's history and can therefore provide important constraints on galaxy formation and evolution models. Many of these quiescent galaxies have been found to be extremely compact, with effective radii a factor $\\sim 6$ smaller than their low-z counterparts (e.g., \\citealt{dad05}; \\citealt{tru06}; \\citealt{dok08}). This is quite puzzling, since these compact galaxies are passively evolving and are therefore not expected to change strongly in size or mass if they do not merge. We note that \\cite{man10} find some large massive quiescent galaxies at $z\\sim 1.5$, showing that not all massive quiescent galaxies at high redshift are compact.  Within the context of current-day models, galaxy mergers play an important role in galaxy evolution (e.g., \\citealt{whi91}). These mergers may cause compact $z\\sim 2$ galaxies to grow ``inside-out'', i.e., the mergers would increase the size of the galaxies (e.g., \\citealt{dok10}; \\citealt{hop09b}). Whether the resulting size growth is large enough, however, is uncertain (e.g., \\citealt{bez09}).  Several authors have emphasised that there are several systematic uncertainties that affect both radius and mass determinations. First, effective radii may be underestimated due to complex morphologies. Specifically, an extended low surface brightness component could remain undetected due to low signal-to-noise ratio (S/N), thereby lowering the observed size (e.g., \\citealt{hop09a}; \\citealt{man10}, but see \\citealt{dok08}; \\citealt{wel08}). Second, mass-to-light gradients may result in a luminosity-weighted effective radius that is smaller than the mass-weighted effective radius (e.g., \\citealt{hop09a}; \\citealt{hop09b}). Such gradients arise in certain models for the formation of massive ellipticals (e.g., \\citealt{rob06}; \\citealt{naa07}). Finally, the inferred stellar masses may be affected by incorrect assumptions regarding the initial mass function (IMF) and stellar evolution models (e.g., \\citealt{muz09}, and references therein).  In this Letter we use new very deep near-infrared (NIR) imaging data from the Hubble Space Telescope's Wide Field Camera 3 (HST/WFC3) to investigate the possibility of size underestimation due to the lack of S/N. We examine the possibility of a ``hidden'' faint extended component being present in $z\\approx 2$ compact quiescent galaxies, focusing on the most massive quiescent galaxy in the Hubble Ultra Deep Field (HUDF) (Beckwith et al 2006), which has previously been studied by \\citet{dad05}. We adopt the following values for cosmological parameters: $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m = 0.3$, and $\\Omega_\\Lambda = 0.7$. All stellar masses are derived assuming a Kroupa IMF \\citep{kro01}. All effective radii are circularized, unless noted otherwise. ",
        "conclusions": "\\begin{figure*} \\plotone{f4.ps}  \\caption{Relations between size and stellar mass ({\\it left}) and size and rest-frame $r$-band luminosity ({\\it middle}) for a sample of low-redshift galaxies, taken from \\cite{guo09}. The large symbol with error bars indicates the position of our galaxy. Low-redshift galaxies are much larger at similar stellar masses and luminosities. The arrow in the middle plot indicates the change in luminosity due to passive evolution to $z=0$. The size of the galaxy is smaller than the local equivalents by a factor of 10.  {\\it Right:} comparison of best-fit residual-corrected rest-frame V-band surface brightness profile (red curve) to elliptical galaxies in the Virgo cluster, from \\cite{kor09} (black curves). Virgo galaxies are plotted with $M_*/M_\\odot > 10^{11}$. The observed high-$z$ surface brightness profile has been corrected for cosmological surface dimming and passive M/L evolution from $z=1.91$ to $z=0$ (see text). Assuming the galaxy has a mass $> 10^{11} M_\\odot$ at $z=0$ its profile at large radii will evolve very strongly over the next 10 Gyr. The central surface brightness profile, on the other hand, shows much less evolution between $z\\approx 2$ and $z=0$.}\\label{fig:rel}  \\end{figure*}  We have found that the galaxy under consideration is indeed remarkably small.  We have fitted a S\\'ersic model to the observed flux distribution, and corrected the profile for the observed deviations. We have measured the galaxy's half light radius: $r_{e, deconv} = 0.42 \\pm 0.14$ kpc. This result is robust to changes in the imposed S\\'ersic profile. As a check of our data's sensitivity to a low surface brightness component we have constructed simulated galaxy images which include a faint extended component. We can reproduce the effective radii to 10\\% using our technique.  A possible cause for concern is that the galaxy might deviate strongly from a S\\'ersic profile.  We have incorporated the residuals in our fit to compensate for such errors, and we note that the residuals from our best S\\'ersic model fit are quite low ($<10\\%$). This implies that our model profile is close to the real profile. This, and the fact that varying $n$ has little influence on the derived half-light radius, suggests that our results are not strongly affected by this source of error.  Thus, our findings indicate that the small effective radius that has been found is not due to oversimplified modeling or a lack of S/N, and gives additional evidence that a strong evolution in size occurs from $z\\approx 2$ to $z=0$. It should be noted that our derived effective radius is $1.6$ times smaller than the radii derived by \\cite{dad05} in the $i$ and $z$ bands.  When we repeat our analysis on the ACS $z$-band data we obtain a slightly different value, $r_{e, deconv} \\approx 0.65$ kpc (uncircularized), closer to the deep $H$-band imaging, and somewhat smaller than the value derived by \\cite{dad05} (but consistent within the errors).  Hence all bands indicate a very small size.  Figure~\\ref{fig:rel} illustrates the difference in size and mass between our galaxy and the $z=0$ elliptical population; plotted in the first two panels are the compact galaxy we have studied and a sample of low-redshift central galaxies from groups and clusters in the Sloan Digital Sky Survey, analyzed by \\cite{guo09}. The compact $z\\approx 2$ galaxy lies far off from the $z=0$ mass-size relation. The middle panel shows the galaxy on the mass-luminosity relation.  We estimated the luminosity evolution of the compact galaxy from $z=1.91$ to $z=0$ in two ways: we first used the rest-frame $B-I$ color difference between low and high redshift to estimate the difference in mass-to-light ratio. Second we used the Fundamental Plane to estimate the evolution from $z=0$ to $z=1$ from \\cite{wel05}, and used the average evolution of the mass-to-light ratios of early-types in the CDFS at $z=1$ and the $z=1.91$ galaxy, both from F\\\"orster Schreiber et al. (in preparation).  The resulting evolution is 1.8-2.2 magnitudes.  As a result, the galaxy still lies off from the size-magnitude relation after correcting for evolution.  In the third panel of Figure~\\ref{fig:rel} we compare the surface brightness profile of this galaxy to those of elliptical galaxies in the Virgo cluster. The profile shown has been corrected for cosmological surface brightness dimming and passive luminosity evolution from $z=1.91$ to $z=0$. The total correction is $-3.5 + 2 \\approx -1.5$ magnitudes. Even though the galaxy has an average density $>100$ times larger than the average $z=0$ elliptical of the same mass, its surface brightness profile in the central kpc is actually rather similar to those of the most massive galaxies at $z=0$ - the average density measured at fixed physical radius is not that different.  This is consistent with results obtained by other authors (e.g., \\citealt{bez09}; \\citealt{hop09a}; \\citealt{fel09}; \\citealt{dok10}).  Thus, the main difference between $z=0$ and this $z\\approx 2$ galaxy is at larger radii where the $z\\approx 2$ galaxy has much lower surface brightness. Such a result could be explained by inside-out growth.  We also  note that there may be significant errors in the mass determination of $z\\approx 2$ compact galaxies, due to e.g.\\ incorrect assumptions about the IMF. Changes in the low mass end of the IMF affect both the masses of the high redshift and low redshift galaxies, and are nearly irrelevant. However, changes in the slope of the IMF will affect the derived passive evolution between $z=2$ and $z=0$, and will increase or decrease the size evolution. Changes in the IMF could thus have important consequences for evolution.  Future deep NIR spectroscopic data should provide direct information on the kinematics of these objects and will allow us to confirm their high masses (see e.g., \\citealt{dok09}).  Finally, it will be interesting to obtain similar deep data on other compact massive galaxies, so that their profiles can be analyzed to the same surface brightness limit. We note that stacking can also lead to a great increase in imaging depth; e.g., \\cite{cas09}, \\cite{dok08}, and \\cite{wel08} stack samples of compact galaxies and obtain very good constraints on their average surface brightness profile. However, with the new WFC3 data available in the coming years many more compact massive galaxies can be studied on an individual basis.  We acknowledge  support from NASA grant HST-GO-11563 and ERC grant 227749."
    },
    "1004/1004.0790.txt": {
        "abstract": "{We present new observations of a transit of the 111.4-day-period exoplanet \\cibleb. Due to this long orbital period and to the orientation of the eccentric orbit ($e=0.9$), the \\cibleb's transits last for about 12 hours. This makes practically impossible the observation of a full transit from a given ground-based observatory. Using the \\spitzer\\ \\textit{Space Telescope} and its IRAC camera on the post-cryogenic mission, we performed a 19-hour-long photometric observation of \\cible\\ that covers the full transit of 13-14 January 2010 as well as off-transit references immediately before and after the event. We complement this photometric data by new spectroscopic observations that we simultaneously performed with \\sophie\\ at Haute-Provence Observatory. This provides radial velocity measurements of the first half of the transit that was previously uncovered with spectroscopy. This new data set allows the parameters of this singular planetary system to be significantly refined. We obtained a planet-to-star radius ratio $R_\\mathrm{p}/R_* = 0.1001 \\pm 0.0006$ that is more accurate but slightly lower than the one measured from previous ground observations in the optical. We found no astrophysical interpretations able to explain such a difference between optical and infrared radii; we rather favor underestimated systematic uncertainties, maybe in the ground-based composite light curve. We detected a feature in the \\spitzer\\  light curve that could be due to a stellar spot. We also found a transit timing about 20~minutes earlier than the ephemeris prediction; this could be caused by actual transit timing variations due to an additional body in the system, or again by underestimated systematic uncertainties. The actual angle between the spin-axis of \\cible\\ and the normal to the planetary orbital plane is found to be near $40^{\\circ}$ thanks to the fit of the Rossiter-McLaughlin anomaly, with a sky-projected value $\\lambda = 42^{\\circ} \\pm 8^{\\circ}$. This allows scenarios with aligned spin-orbit to be definitively rejected. Over the twenty planetary systems with measured spin-orbit angles, a few of them are misaligned; this is probably the signature of two different evolution scenarios for misaligned and aligned systems, depending if they experienced or not gravitational interaction with a third body. As in the case of \\cible, most of the planetary systems including a massive planet are tilted; this could be the signature of a separate evolution scenario for massive planets in comparison with Jupiter-mass~planets.} ",
        "introduction": "Among the  more than 400 extrasolar planets that have been found so far, the giant planet orbiting the G5 star \\object{\\cible} is certainly a unique case. Its eccentricity is particularly high: $e=0.93$. Only one known planet possibly has a higher eccentricity, namely HD\\,20782b; however its high $e$-value is still to be confirmed as it stands on one measurement only (O'Toole et al.~\\cite{otoole09}). The comet-like orbit of \\cibleb\\ was well established in its discovery paper by Naef et al.~(\\cite{naef01}) and has been largely confirmed by subsequent observations. Together with its 111.4-day period, the high eccentricity of \\cibleb\\ put it on an extreme orbit: during its revolution, the planet experiences the strongly irradiated regime of a \"hot Jupiter\" at periastron (0.03~AU), and milder conditions at apoastron (0.87~AU), around which it spends most of its time, and where the planet approaches the inner boundary of the habitable zone (see Fig.~\\ref{fig_david}).  \\begin{figure}[h] \\begin{center} \\includegraphics[scale=0.417]{fig_sketch.ps} \\caption{Schematic view of the \\cible\\ system. The red dots show the positions of the planet each 24 hours.} \\label{fig_david} \\end{center} \\end{figure}   As long as the inclination $i$ of the orbit with the line of sight was unknown, those parameters originally implied a probability of 1\\% for the planet to be transiting. In spite of this tenuous probability, an amazing chance makes the planet \\cibleb\\ actually transits its parent star, as seen from Earth, each 111.4~days. This is particularly advantageous as numerous crucial studies can be performed using photometry or spectroscopy when a planet passes in front its parent star (planetary transits) or behind it (planetary eclipses), especially in this case where the host star is bright ($V=9.1$) and nearby ($d=60$~pc). The fortunate transiting nature of \\cibleb\\ was established in February~2009 from the detection of a transit reported from ground observations, independently by Moutou et al.~(\\cite{moutou09}) from photometric and spectroscopic data, and by Garcia-Melendo \\& McCullough~(\\cite{garcia09}) and Fossey et al.~(\\cite{fossey09}) from photometric measurements. Additional observations of transits were later reported by Winn et al.~(\\cite{winn09a}) and Hidas et al.~(\\cite{hidas10}). The February~2009 observations followed the planetary eclipse discovery reported a few months earlier by Laughlin et al.~(\\cite{laughlin09}) from \\spitzer\\ photometric observations at 8\\,$\\mu$m during a 30-hour interval around the periastron. Among the  known transiting planets, \\cibleb\\ has the longest period and the most eccentric orbit. The second most extreme transiting planet is HD\\,17156b ($P=21.2$~days and $e=0.67$). Furthermore, \\cibleb\\ is also the most distant planet from its parent stars when it transits: 0.3~AU, against 0.05~AU or less for all other known cases\\footnote{See also the 95-day-period transiting exoplanet CoRoT-9b announced after the submission of the present paper (Deeg et al.~\\cite{deeg10}).}.  In addition to the photometric detection of the \\cibleb's transit, Moutou et al.~(\\cite{moutou09}) presented its spectroscopic detection using the Rossiter-McLaughlin %(RM) effect, from radial velocities measured with the \\sophie\\ spectrograph at Haute-Provence Observatory (OHP), France. The Rossiter-McLaughlin effect is an apparent distortion of the stellar lines profile due to the transit of the planet in front of the rotating star. From the \\sophie\\ measurements, Moutou et al.~(\\cite{moutou09}) have shown the first evidence for a spin-orbit misalignment, i.e. the orbital plane of the planet \\cibleb\\ is not perpendicular to the spin-axis of its host star. Using additional photometric data of the February 2009 event allowing a better constraint on the transit duration together with a combined analysis of the whole data set, Pont et al.~(\\cite{pont09}) refined the parameters of the system. They confirmed the spin-orbit misalignment from the Rossiter-McLaughlin distortion detected with \\sophie\\ and provided a measurement of the sky-projected angle between the planetary orbital axis and the stellar rotation axis: $\\lambda\\sim50^{\\circ}$, with the confidence interval [$14^{\\circ}-111^{\\circ}$] -- see also Gillon~(\\cite{gillon09a}). Thanks to new photometric and spectroscopic observations of the June 2009 transit, Winn et al.~(\\cite{winn09a}) thereafter reduced even more the confidence interval to [$32^{\\circ}-87^{\\circ}$]. Thus, the spin-orbit misalignment of the \\cible\\ system is now well established. \\cible\\ is the component of a wide binary with HD\\,80607; the projected separation of the system is about 1000~AU. The peculiar orbit of \\cibleb\\ could thus result from Kozai mechanism and tidal dissipation (see, e.g., Wu~\\& Murray~\\cite{wu03}), which can pump the eccentricity and the inclination.  Due to the long orbital period and to the orientation of the eccentric orbit, the duration of the transit of \\cibleb\\ is about 12-hour long. This should be compared to the transit duration of other known transiting exoplanets, which typically lasts less than five hours. The transit duration of \\cibleb\\ is even longer than transits of Mercury or Venus through the Solar disk as seen from the Earth. It is thus practically impossible that a full transit of \\cibleb\\ matches the duration of an observation night from ground. In addition, data secured before and after the transit are mandatory to obtain an accurate transit light curve, so the full sequence for \\cible\\ lasts longer than a night for ground observations. Observing an entire transit of this exoplanet is thereby challenging and only portions of a transit could be observed from a given ground-based telescope. This was the case of all the observation campaigns reported above which covered only fractions of transits. These fragmented observations induce significant uncertainties in the parameters derived from their~fit.  We present here the first full photometric observation of a transit of \\cibleb. We secured it on 13-14 January 2010 with the \\spitzer\\ space observatory using the IRAC camera at 4.5\\,$\\mu$m in post-cryogenic mission. Thanks to its Earth-trailing heliocentric orbit (Scoupe et al.~\\cite{spot06}), \\spitzer\\ allowed us to continuously observe during 19~hours, enabling the coverage of the whole 12-hour-long transit, as well as off-transit references immediately before and after the event. We complement this photometric data by new spectroscopic observations that we simultaneously performed with \\sophie\\ at OHP. This provides radial velocity measurements of the first half of the transit, a part that was up to now uncovered  with spectroscopy. Indeed, observing the full 12-hour-long transit is even more difficult in spectroscopy than in photometry, as the amplitude of the Rossiter-McLaughlin for \\cible\\ is about $10$\\,\\ms\\ whereas Northern instruments able to measure radial velocities with the required accuracy are sparse. We also performed a ground-based photometric monitoring of \\cible\\ during January 2010. All together, the data of this observational campaign allows the parameters of this planetary system  to be additionally~refined.  The  observations and data reduction are presented in Sect.~\\ref{sect_spitzer} and~\\ref{section_sophie} for \\spitzer\\ and \\sophie\\  respectively. The ground-based photometry in presented in Sect.~\\ref{sect_ground_photom}. The analyses and the results are presented in Sect.~\\ref{sect_analysis}, before discussion and conclusion in Sect.~\\ref{sect_discussion} and ~\\ref{sect_conclusion}. ",
        "conclusions": "\\label{sect_conclusion}  We presented an observation of the 12-hour-long transit of the highly eccentric, 111.4-day-period exoplanet \\cibleb\\ performed in January~2010. The~transit light curve we present is among the first ones carried with the post-cryogenic \\spitzer. Its shows systematic effects stronger than those seen in the Cold-\\spitzer\\ but the accuracy remains clearly better that ground-based observations. Together with the \\sophie\\ measurements acquired at the same epoch, this is one of the first  observational campaigns performed simultaneously in radial velocities and space-based high-accuracy photometry. With previously available data sets, this allows the parameters of this system to be significantly refined thanks to combined fits. There is a possible detection of a variation in the transit timing, which has to be confirmed by additional observations of \\cibleb\\ transits, from ground or space. A dark spot was also possibly detected on the surface of this inactive star. The spin-orbit misalignment is clearly confirmed and the $\\lambda$ angle is accurately measured. As most of the massive planets for which this angle is measured, the orbit of \\cibleb\\ is misaligned with the equatorial plan of its host star. This suggests a separate evolution scenario for massive planets in comparison with Jupiter-mass planets."
    },
    "1004/1004.5232_arXiv.txt": {
        "abstract": "There are a total of 1451 gamma-ray emitting objects in the \\emph{Fermi} Large Area Telescope First Source Catalogue. The point source location accuracy of typically a few arcminutes has allowed the counterparts for many of these sources to be found at other wavelengths, but even so there are 630 which are described as having no plausible counterpart at $80\\%$ confidence. In order to help identify the unknown objects, we have cross-correlated the positions of these sources with the \\emph{Rosat} All Sky Survey Bright Source Catalogue. In this way, for \\emph{Fermi} sources which have a possible counterpart in soft X-rays, we can use the, much smaller, \\emph{Rosat} error box to search for identifications. We find a strong correlation between the two samples and calculate that there are about 60 sources with a \\emph{Rosat} counterpart. Using the \\emph{Rosat} error boxes we provide tentative associations for half of them, demonstrate that the majority of these are either blazars or blazar candidates and give evidence that most belong to the BL Lac class. Given that they are X-ray selected and most are high synchrotron peaked objects, which indicates the presence of high energy electrons, these sources are also good candidates for TeV emission, and therefore good probes of the extragalactic background light. ",
        "introduction": "A key strategic objective of the \\emph{Fermi} mission is a survey of the sky at gamma-ray energies, making use of the large area and field of view of the LAT instrument \\citep{atwood}. The telescope allows the detection of sources with an angular resolution of about 0.6 degrees (68\\% at 1 GeV) and a point source location accuracy (PSLA) varying from around 1 to 6 arcminutess, depending on the detection significance. In the first \\emph{Fermi} catalogue (\\citealt{abdo1}), comprising data from the initial 11 months of the science phase of the mission, there are 1451 objects listed, of which 821 have been associated with known sources at other wavelengths. These identified sources comprise both extragalactic and galactic objects with the former including blazars (flat spectrum radio QSOs (FSRQ) or BL Lacs), a few radio galaxies and 4 normal galaxies, while the latter is made up of pulsars, pulsar wind nebulae, supernova remnants, globular clusters and a few binaries. Some peculiar objects are also found but their associations are less secure.  No plausible counterparts have been found for the remaining 630 objects and therefore these cannot yet be associated with any known class of gamma-ray emitting objects. The \\emph{Fermi}/LAT AGN catalogue (\\citealt{abdo2}, hereafter A1) lowers the confidence limit to 50\\% for high latitude ($|b| > 10^{\\circ}$) sources  and thereby provides possible counterparts for another 26 of these sources, and for 104 more there are potential associations ('affiliations') but with unquantified confidence.  Searching for counterparts of these new high energy sources is a primary objective of the survey work but it is made very difficult by the large, with respect to other wavelengths, \\emph{Fermi} error boxes. This uncertainty in their locations  means that a positional correlation with a known object is usually not enough to identify a \\emph{Fermi} source and instead, a multiwavelength approach, using X-ray, optical and radio data of likely counterparts must be used in order to understand their nature and to evaluate the likelihood of their association with the \\emph{Fermi} detections. Searches for X-ray counterparts are particularly useful in finding a positionally-correlated, highly-unusual object with the special parameters that might be expected to produce gamma rays. X-ray surveys are well suited for this type of search because they offer 3 great advantages: a) they allow a full coverage of the \\emph{Fermi} error box, b) they provide arcsecond location accuracy and c) they give information in an energy band quite close to that in which \\emph{Fermi} operates. Cross correlation analysis using X-ray catalogues can therefore be a useful tool with which to restrict the positional uncertainty of the objects detected by \\emph{Fermi} and so to facilitate the identification process.  Herein, we report on the strong level of positional correlation between the unassociated \\emph{Fermi} sources and the \\emph{Rosat} All Sky Survey Bright Source Catalogue  leading to evidence for the association of a number of GeV sources with a soft X-ray counterpart, better positions for all correlated objects and hence the possibility of optical follow-up work. We find that most of these associations are with BL Lac or BL Lac candidates. ",
        "conclusions": "We have shown that, as expected, there is a strong correlation between the \\emph{Fermi} survey source list and the \\emph{Rosat} All Sky Survey  Bright Source Catalogue, finding that there should be about 60 sources common to both lists. By placing a maximum correlation distance of 160 arcseconds in order to minimise chance associations we have a sample of 30 objects in which only $1 \\pm 1$ should be by chance alignment. We can use the \\emph{Rosat} error box to help find optical counterparts for these sources, which in all  cases is sufficiently small to allow the identification of one single object.  Most of the associations are of extragalactic nature, with only a few being galactic. We have also shown that this cross correlation analysis appears to preferentially select BL Lac objects. These X-ray selected objects  are often HSP BL Lacs and are thereffore good candidates for TeV emission. Clearly only optical spectroscopy of the \\emph{Rosat} counterparts can confirm this suggestion, but cross correlation with other catalogues may provide a different selection criterion and hence different source types."
    },
    "1004/1004.0284_arXiv.txt": {
        "abstract": "We study how well the mass of the graviton can be constrained from gravitational-wave (GW) observations of coalescing binary black holes. Whereas the previous investigations employed post-Newtonian (PN) templates describing only the inspiral part of the signal, the recent progress in analytical and numerical relativity has provided analytical waveform templates coherently describing the inspiral-merger-ringdown (IMR) signals. We show that a search for binary black holes employing IMR templates will be able to constrain the mass of the graviton much more accurately ($\\sim$ an order of magnitude) than a search employing PN templates. The best expected bound from GW observatories ($\\lambdag > 7.8\\times10^{13}$ km from Adv.~LIGO, $\\lambdag > 7.1\\times10^{14}$ km from Einstein Telescope, and $\\lambdag > 5.9\\times10^{17}$ km from LISA) are several orders-of-magnitude better than the best available \\emph{model-independent} bound ($\\lambda_g > 2.8 \\times 10^{12}$ km, from Solar system tests). ",
        "introduction": "\\label{sec:Intro}  General relativity (GR) has proven to be very successful in predicting and explaining a variety of gravitational phenomena. Over the last century, basic ingredients of GR have been tested in many different ways and in many different settings. In GR, gravitational interactions propagate with the speed of light, which means that the hypothetical quantum particle of gravity, the ``graviton'', has no rest mass. A non-zero graviton mass $m_g$ would produce several interesting effects: for e.g., it will cause the gravitational potential to take the Yukawa form, effectively cutting off gravitational interactions at distances greater than the Compton wavelength $\\lambdag \\equiv h/m_gc$ of the graviton. Indeed, the absence of such effects in the solar system has provided a lower bound on $\\lambdag$ (hence an upper bound on $m_g$). But a graviton mass with corresponding Compton wavelength much larger than the size of the solar system has not yet been ruled out in a model-independent way\\footnote{It has been argued that ``dark energy'' \\emph{might} be explained by a graviton-mass-like effect, with associated Compton wavelength comparable to the radius of the visible universe. See~\\cite{Goldhaber:2008xy} for a review.}. Also, if the graviton has non-zero mass, the gravitational waves (GWs) will have extra degrees of freedom (such as longitudinal modes), and will travel with a frequency-dependent speed, different from the speed of light.  Whereas the current experimental bounds on theories of gravity are from the weak-field limit, GW observations provide excellent test beds in the dynamical, strong-field regime  (see \\cite{lrr-Will:2006} and \\cite{lrr-2009-2} for recent reviews). Indeed, these are exciting times for the word-wide GW community. The Initial LIGO~\\cite{Abbott:2007kv} detectors have completed their fifth science run, at design sensitivity. The Virgo~\\cite{VirgoStatus-GWDAW2008} and GEO\\,600 detectors ran concurrently with LIGO for part of that run. Although a direct detection of GWs is yet to be made, a number of interesting astrophysical upper limits are constructed based on this data (see e.g.,~\\cite{Abbott:2008fx, Collaboration:2009ws}). After the ongoing sixth science run at a further improved sensitivity, the LIGO-Virgo detectors will undergo a commissioning break with the target of achieving much improved sensitivities with Advanced LIGO and Advanced Virgo.  Implementing a number of advanced technologies including the use of squeezed light, the GEO\\,600 is already undergoing commissioning work for the high-frequency detector configuration GEO-HF~\\cite{Willke:2006uw}. Also, the design study for a third generation detector, called the Einstein Telescope (ET) is ongoing in Europe. With the ever closer approach of the era in which GW observations become routine, it is interesting to see and update what information could be gathered about the Universe from these observations.  In this paper we investigate the possible bounds that can be put on theories with an effective mass in the propagation of GWs (massive-graviton theories\\footnote{ It has to be mentioned that so far there is no complete theory of gravity in which the graviton can have a mass. Constructing such a theory in a ghost-free manner has proven to be nontrivial; see, e.g.,~\\cite{Yunes:2009ke} and references therein.}, in short) from GW observations of coalescing binary black holes (BBHs). In massive-graviton theories, the speed of propagation of GWs depends on the wavelength. If the GW emission is accompanied by electromagnetic (EM) emission (such as the case of a core-collapse supernova or a white-dwarf binary), a bound on the mass of the graviton can be placed by comparing the time of arrival of the GW signal with the EM signal~\\cite{Will98}, or, by correlating the EM and GW signals~\\cite{Larson:1999kg,Cutler:2002ef,Kocsis:2008}.  But, since many of these sources are not very well understood in terms of their emission mechanisms and the delay between EM and GW emissions, this could introduce significant uncertainties in such measurements.  Coalescing BBHs provide us with the potential to constrain the mass of the graviton without relying on the presence of an EM counterpart.  In the case of BBHs, since the frequency of gravitational radiation sweeps from lower to higher frequencies, the speed of the emitted gravitons will vary, from lower speeds to higher speeds. This will cause a distortion in the observed phasing of the GWs. Since BBHs can be accurately modelled using analytical/numerical solutions of Einstein's equations, any deviation from GR can be parametrized and measured (see, e.g.,~\\cite{Yunes:2009ke} for a recent discussion on such a general framework). A framework for testing this possibility by measuring the distortion of GWs was originally developed by Will~\\cite{Will98}.  Will found a dispersive effect that appeared as an additional term in the post-Newtonian (PN) expansion of the GW phase, and showed that a bound on the mass of the graviton can be placed from the GW observations by applying appropriate matched filters.  Will's original work was performed using \\emph{restricted} PN waveforms describing the \\emph{inspiral} stage of non-spinning coalescing compact binaries, the phase of which was expanded to 1.5PN order. Recent work has elaborated on this by incorporating more accurate detector models, and by including more physical effects such as effects rising from the spin angular momentum of the compact objects, from the eccentricity of the orbit, and from the inclusion of higher harmonics rising from the contribution of the higher multipoles \\cite{Will:2004xi,BBW2004, Jones:2004uy,Berti:2005qd, Stavridis:2009mb, YagiTanaka09, ArunWill09}.  Since the PN formalism has enabled us to compute accurate waveforms from the \\emph{inspiral} stage of the coalescence, these analyses have focused on the information gained from the observation of the inspiral stage. The last few years have witnessed a revolution in the numerical simulations of compact binaries. In particular, numerical relativity was able to obtain exact solutions for  the ``binary-black-hole problem'' \\cite{Pretorius1, Baker05a, Campanelli:2005dd}.  Concomitant with this great leap has been significant progress in analytical relativity in the computation of high order PN terms and the inclusion of various effects arising from spins, higher harmonics etc. Combining the analytical and numerical results, different ways of constructing inspiral-merger-ring-down (IMR) waveforms have been proposed \\cite{Ajith:2007kx,Buonanno:2007pf,Damour:2007yf}.  It has been widely recognized that these IMR waveforms will significantly improve the sensitivity and distance reach of the searches for BBHs (see, e.g., \\cite{FlanHugh98, Buonanno:2006ui, Ajith:2007kx}) as well as the accuracy of the parameter estimation (see, e.g., \\cite{Ajith:2009fz, Aylott:2009ya, AylottVeitchVecchio:2009}).  In this paper, we estimate the bounds that can be placed on the mass of graviton from the GW observations of BBHs using IMR templates. This is motivated by the previous observations (see e.g.~\\cite{Ajith:2009fz}) that the IMR waveforms will significantly improve the accuracy of the parameter estimation by breaking the degeneracies between the different parameters describing the signal, including the parameter describing the mass of the graviton.  Due to the intrinsic randomness of the noise in the GW data, the estimated parameters of the binary (including the one parameter describing the mass of the graviton) will fluctuate around their mean values.   In the limit of high signal-to-noise ratios (SNRs), the spread of the distribution of the observed parameters --- the accuracy of the parameter estimation --- is quantified by the inverse of the \\emph{Fisher information matrix}~\\cite{Cramer46,Rao45}.  We employ the Fisher matrix formalism to estimate the expected bounds on the mass of the graviton using the \\emph{non-spinning} limit of the IMR waveform model proposed by Ref.~\\cite{Ajith:2009bn}. This is a frequency-domain waveform family describing the leading harmonic of the IMR waveforms from BBHs.  Indeed, in the absence of a complete massive-graviton theory, we consider only the propagation effects of a massive graviton, and assume that the wave generation is correctly given by GR, apart from negligible corrections.  This will introduce a systematic error on the possible constraints, which depends on the specific massive graviton theory\\footnote{If stable black holes exist in such theory, then we might expect that the major effect comes from the propagation, since the length scales involved in the propagation are typically much larger than the length scales involved in the wave generation. However, it is difficult to address this properly in the absence of a complete theory of massive graviton.}. Since we are neglecting spins and higher harmonics in the (general relativistic) signal model, this introduces another set of systematic errors; but these can be circumvented by including these effects in the signal model. We remind the reader that this paper only tries to quantify the noise-limited statistical errors in the observations.  The main findings of the paper are summarized below (Section~\\ref{sec:summary}).  The following sections present the details of the analysis. Section~\\ref{sec:dispcbc} briefly reviews the effect of massive graviton on the dispersion of GWs, and summarizes the existing bounds on the graviton mass. In Section~\\ref{sec:gwdispmgbounds}, we compute the expected upper bounds that can be placed on the mass of the graviton using the observations of IMR signals. In that section, we review the signal and detector models used, provide the details of the computation and present a discussion of the results and the limitations of this work.  \\subsection{Summary of results} \\label{sec:summary}  An executive summary of results is presented in Fig.~\\ref{fig:summaryplot} for the case of ground-based detectors Adv.~LIGO and ET as well as the space-borne detector LISA.  For ground-based detectors, the binary is assumed to be located optimally oriented at 1 Gpc, and for LISA, the binary is located at 3 Gpc. For the case of Adv.~LIGO (with low-frequency cutoff, $\\flow$ = 10 Hz), the best bound ($\\lambdag > 7.8\\times10^{13} \\mathrm{km} \\simeq 2.5 \\mathrm{pc}$; $m_g < 1.6\\times10^{-23}$~eV) using IMR templates is obtained from the observation of binaries with total mass $M \\simeq 360 M_\\odot$. This is significantly better than the \\emph{best} bound obtained using restricted 3.5PN templates ($\\lambdag > 7.9\\times10^{12}$ km for $M\\simeq 18 M_\\odot$).  For ET (with $\\flow$ = 1 Hz), the best bound using IMR templates ($\\lambdag > 7.1\\times10^{14} \\mathrm{km} \\simeq 23 \\mathrm{pc}$; $1.7\\times10^{-24}$~eV) is obtained from binaries with $M\\simeq 3000 M_\\odot$, while the best bound for PN templates ($\\lambdag > 1\\times10^{14}$ km) is obtained for binaries with $M\\simeq 65 M_\\odot$. For LISA observation of supermassive black-hole (BH) binaries, the best bound ($\\lambdag > 5.9\\times10^{17} \\mathrm{km} \\simeq 19 \\mathrm{kpc}$; $2.1\\times10^{-27}$~eV) using IMR templates is obtained from binaries with $M \\simeq 4.8\\times10^{7}M_\\odot$, while the best bound using PN templates ($\\lambdag > 6.3\\times10^{16}$ km) is obtained from binaries with $M\\simeq 1.9\\times10^{6} M_\\odot$. In summary, the \\emph{best} bounds using IMR templates are roughly an order of magnitude better than the \\emph{best} bounds using restricted PN templates. The improvement is partly due to the higher SNR and partly due to the extra information harnessed from the post-inspiral stages.  The best expected bound from ground-based observatories is over two orders-of-magnitude better than the best available model-independent bound ($\\lambdag > 2.8 \\times 10^{12}$ km) given by monitoring the orbit of Mars (see Sec.~\\ref{sec:otherbounds}). Additionally, LISA will be able to improve on those bounds by several orders of magnitude.  Most importantly, GW observations will provide the first constraints from the highly dynamical, strong-field regime of gravity. Table~\\ref{tab:boundsGWsummary} summarizes the bounds that can be placed with future GW observations employing different signal models / analysis methods.  \\smallskip  \\noindent \\emph{Effect of low frequency cutoff:---} The best bound that can be placed with Adv.~LIGO noise spectrum with low-frequency cutoff, $\\flow$ = 10 Hz will be $\\sim$ 25\\% better than the same obtained with $\\flow$ = 20 Hz. For ET, the best bounds using a configuration with $\\flow$ = 1 Hz is $\\sim$ 70\\% better than the same obtained with $\\flow$ = 10 Hz. We hope that this information can contribute to weighing the scientific case of different configurations of advanced detectors.  \\noindent \\emph{Detectability of the signals by GR-based templates:---} If we assume the constraints on the graviton mass given by solar system tests, the ``mismatch'' between the signal and GR-based templates is unacceptably high (7--50\\%). However, if the mass of the graviton is a factor of three smaller than the solar system limit, then the mismatch is negligible. This provides a rough estimate of the detectability of the deformed signals (due to the graviton mass) using GR-based templates.  \\begin{table}[tbh] \\begin{center} \\begin{tabular}{ccccc} \\hline \\hline Signal model / method &\\vline& Adv.~LIGO       & ET & LISA  \\\\ \\hline \\hline \\parbox[t][0.13in]{1.2in}{\\footnotesize 1.5PN inspiral~\\cite{Will98}} &\\vline& 0.6 (20) &  & 0.7 ($2\\times10^7$)  \\\\ \\parbox[t][0.26in]{1.3in}{\\footnotesize 1.5PN inspiral, more accurate detector model~\\cite{Will:2004xi}} &\\vline&  &  & \t0.5 ($10^7$)  \\\\ \\parbox[t][0.26in]{1.2in}{\\footnotesize 3.5PN inspiral, no spin [this paper]} &\\vline& 0.8 (18) & \t10 (65) & \t0.6 ($1.9 \\times 10^6$)  \\\\ \\parbox[t][0.26in]{1.2in}{\\footnotesize 3.5PN inspiral, higher harmonics, no spin~\\cite{ArunWill09}} &\\vline& 0.7 (60) & \t10 (400) & \t0.5 ($2\\times10^6$)  \\\\ \\parbox[t][0.26in]{1.2in}{\\footnotesize 3.5PN inspiral, 2PN spin, no precession~\\cite{BBW2004}} &\\vline&  & \t & \t0.5 ($2\\times10^6$)  \\\\ \\parbox[t][0.26in]{1.2in}{\\footnotesize 2PN inspiral, spin precession~\\cite{Stavridis:2009mb}} &\\vline&  & \t & \t0.7 ($2\\times10^7$)  \\\\ \\parbox[t][0.26in]{1.3in}{\\footnotesize 2PN inspiral, simple spin precession, eccentricity~\\cite{YagiTanaka09}} &\\vline&  & \t & \t0.4 ($10^7$)  \\\\ \\parbox[t][0.13in]{1.2in}{\\footnotesize IMR, no spin [this paper]} &\\vline& 8 (360) & \t70 (3000) & \t6 (4.8 $\\times 10^7$)  \\\\ \\parbox[t][0.39in]{1.2in}{\\footnotesize Measuring phase of arrival of diff.~harmonics from eccentric binaries~\\cite{Jones:2004uy}} &\\vline&  & \t & \t0.3 ($10^6$)  \\\\ \\parbox[t][0.39in]{1.2in}{\\footnotesize Correlating GW and EM observations from white dwarf binaries~\\cite{Cutler:2002ef}} &\\vline&  & \t & \t$10^{-3}$ (2.8)  \\\\ \\parbox[t][0.52in]{1.3in}{\\footnotesize Correlating GW and EM observations from supermassive BH binaries~\\cite{Kocsis:2008}} &\\vline&  & \t & \t$9 \\times 10^{-2}$ ($10^{7}$)  \\\\ \\hline \\hline \\end{tabular} \\caption{Expected bounds on $\\lambdag$ (in units of $10^{13}$ km for the case of Adv.~LIGO and ET, and $10^{17}$ km for LISA) to one significant digit using future GW observations employing different signal models / analysis methods. Total mass (in $M_\\odot$) giving the bound are shown in brackets. It should be noted that different investigations have used slightly different source- and detector models, and hence, are not strictly comparable.} \\label{tab:boundsGWsummary} \\end{center} \\end{table} ",
        "conclusions": "An executive summary of results is presented in Fig.~\\ref{fig:summaryplot} for the case of ground-based detectors Adv.~LIGO and ET as well as the space-borne detector LISA.  For ground-based detectors, the binary is assumed to be located optimally oriented at 1 Gpc, and for LISA, the binary is located at 3 Gpc. For the case of Adv.~LIGO (with low-frequency cutoff, $\\flow$ = 10 Hz), the best bound ($\\lambdag > 7.8\\times10^{13} \\mathrm{km} \\simeq 2.5 \\mathrm{pc}$; $m_g < 1.6\\times10^{-23}$~eV) using IMR templates is obtained from the observation of binaries with total mass $M \\simeq 360 M_\\odot$. This is significantly better than the \\emph{best} bound obtained using restricted 3.5PN templates ($\\lambdag > 7.9\\times10^{12}$ km for $M\\simeq 18 M_\\odot$).  For ET (with $\\flow$ = 1 Hz), the best bound using IMR templates ($\\lambdag > 7.1\\times10^{14} \\mathrm{km} \\simeq 23 \\mathrm{pc}$; $1.7\\times10^{-24}$~eV) is obtained from binaries with $M\\simeq 3000 M_\\odot$, while the best bound for PN templates ($\\lambdag > 1\\times10^{14}$ km) is obtained for binaries with $M\\simeq 65 M_\\odot$. For LISA observation of supermassive black-hole (BH) binaries, the best bound ($\\lambdag > 5.9\\times10^{17} \\mathrm{km} \\simeq 19 \\mathrm{kpc}$; $2.1\\times10^{-27}$~eV) using IMR templates is obtained from binaries with $M \\simeq 4.8\\times10^{7}M_\\odot$, while the best bound using PN templates ($\\lambdag > 6.3\\times10^{16}$ km) is obtained from binaries with $M\\simeq 1.9\\times10^{6} M_\\odot$. In summary, the \\emph{best} bounds using IMR templates are roughly an order of magnitude better than the \\emph{best} bounds using restricted PN templates. The improvement is partly due to the higher SNR and partly due to the extra information harnessed from the post-inspiral stages.  The best expected bound from ground-based observatories is over two orders-of-magnitude better than the best available model-independent bound ($\\lambdag > 2.8 \\times 10^{12}$ km) given by monitoring the orbit of Mars (see Sec.~\\ref{sec:otherbounds}). Additionally, LISA will be able to improve on those bounds by several orders of magnitude.  Most importantly, GW observations will provide the first constraints from the highly dynamical, strong-field regime of gravity. Table~\\ref{tab:boundsGWsummary} summarizes the bounds that can be placed with future GW observations employing different signal models / analysis methods.  \\smallskip  \\noindent \\emph{Effect of low frequency cutoff:---} The best bound that can be placed with Adv.~LIGO noise spectrum with low-frequency cutoff, $\\flow$ = 10 Hz will be $\\sim$ 25\\% better than the same obtained with $\\flow$ = 20 Hz. For ET, the best bounds using a configuration with $\\flow$ = 1 Hz is $\\sim$ 70\\% better than the same obtained with $\\flow$ = 10 Hz. We hope that this information can contribute to weighing the scientific case of different configurations of advanced detectors.  \\noindent \\emph{Detectability of the signals by GR-based templates:---} If we assume the constraints on the graviton mass given by solar system tests, the ``mismatch'' between the signal and GR-based templates is unacceptably high (7--50\\%). However, if the mass of the graviton is a factor of three smaller than the solar system limit, then the mismatch is negligible. This provides a rough estimate of the detectability of the deformed signals (due to the graviton mass) using GR-based templates.  \\begin{table}[tbh] \\begin{center} \\begin{tabular}{ccccc} \\hline \\hline Signal model / method &\\vline& Adv.~LIGO       & ET & LISA  \\\\ \\hline \\hline \\parbox[t][0.13in]{1.2in}{\\footnotesize 1.5PN inspiral~\\cite{Will98}} &\\vline& 0.6 (20) &  & 0.7 ($2\\times10^7$)  \\\\ \\parbox[t][0.26in]{1.3in}{\\footnotesize 1.5PN inspiral, more accurate detector model~\\cite{Will:2004xi}} &\\vline&  &  & \t0.5 ($10^7$)  \\\\ \\parbox[t][0.26in]{1.2in}{\\footnotesize 3.5PN inspiral, no spin [this paper]} &\\vline& 0.8 (18) & \t10 (65) & \t0.6 ($1.9 \\times 10^6$)  \\\\ \\parbox[t][0.26in]{1.2in}{\\footnotesize 3.5PN inspiral, higher harmonics, no spin~\\cite{ArunWill09}} &\\vline& 0.7 (60) & \t10 (400) & \t0.5 ($2\\times10^6$)  \\\\ \\parbox[t][0.26in]{1.2in}{\\footnotesize 3.5PN inspiral, 2PN spin, no precession~\\cite{BBW2004}} &\\vline&  & \t & \t0.5 ($2\\times10^6$)  \\\\ \\parbox[t][0.26in]{1.2in}{\\footnotesize 2PN inspiral, spin precession~\\cite{Stavridis:2009mb}} &\\vline&  & \t & \t0.7 ($2\\times10^7$)  \\\\ \\parbox[t][0.26in]{1.3in}{\\footnotesize 2PN inspiral, simple spin precession, eccentricity~\\cite{YagiTanaka09}} &\\vline&  & \t & \t0.4 ($10^7$)  \\\\ \\parbox[t][0.13in]{1.2in}{\\footnotesize IMR, no spin [this paper]} &\\vline& 8 (360) & \t70 (3000) & \t6 (4.8 $\\times 10^7$)  \\\\ \\parbox[t][0.39in]{1.2in}{\\footnotesize Measuring phase of arrival of diff.~harmonics from eccentric binaries~\\cite{Jones:2004uy}} &\\vline&  & \t & \t0.3 ($10^6$)  \\\\ \\parbox[t][0.39in]{1.2in}{\\footnotesize Correlating GW and EM observations from white dwarf binaries~\\cite{Cutler:2002ef}} &\\vline&  & \t & \t$10^{-3}$ (2.8)  \\\\ \\parbox[t][0.52in]{1.3in}{\\footnotesize Correlating GW and EM observations from supermassive BH binaries~\\cite{Kocsis:2008}} &\\vline&  & \t & \t$9 \\times 10^{-2}$ ($10^{7}$)  \\\\ \\hline \\hline \\end{tabular} \\caption{Expected bounds on $\\lambdag$ (in units of $10^{13}$ km for the case of Adv.~LIGO and ET, and $10^{17}$ km for LISA) to one significant digit using future GW observations employing different signal models / analysis methods. Total mass (in $M_\\odot$) giving the bound are shown in brackets. It should be noted that different investigations have used slightly different source- and detector models, and hence, are not strictly comparable.} \\label{tab:boundsGWsummary} \\end{center} \\end{table}   \\end{center} \\end{table}  We emphasize that currently there is no observational evidence for stellar-mass black holes having mass greater than $\\sim 30\\, M_\\odot$. Although a recently discovered ultra-luminous X-ray source \\cite{2009Natur.460...73F} represents a possible detection of a black hole with mass $\\sim 500\\, M_\\odot$, observational evidence for intermediate-mass ($\\sim 100 - 10^4\\, M_\\odot$) black holes is still under debate \\cite{MillerColbert:2004, miller-2008}. However, population synthesis studies have suggested a number of channels through which black-hole binaries with total mass $\\sim 30 - 10^4 M_\\odot$ could form~\\cite{2010ApJ...715L.138B,Sadowski:2008, Belczynski:2009xy,Brown:2007, Mandel:2007hi, imbhlisa-2006,Amaro:2006imbh}. Although the coalescence rates of these binaries are highly uncertain, advanced detectors \\emph{might} be able to detect several of such events per year~\\cite{2010arXiv1003.2480L}.  Another concern is the detectability of signals deformed by the propagation effect of the massive graviton by GR-based templates. In order to address this question (indeed partially), we have computed the \\emph{fitting factor} (fraction of optimal SNR recovered by a suboptimal template family)~\\cite{Apostolatos:1995pj} of the GR-based IMR templates with signals from equal-mass binaries (in the mass range $20-2000 M_\\odot$, located at 1Gpc) in Adv.~LIGO, assuming different values for the graviton mass. If we assume the constraints on the graviton mass given by solar system tests ($\\lambdag = 2.8 \\times 10^{12}$ km), the fitting factor is unacceptably low (0.5 -- 0.93). However, if the mass of the graviton is a factor of three smaller than the solar system limit, then the fitting factors are greater than 0.97. This dramatic change is due to the fact that the deformation of the observed signal is proportional to the square of the Compton wavelength (see Eq.~\\ref{eq:phaseDistortion}).    \\subsection{Limitations of this work} \\label{sec:limitations}  Due to the lack of a complete massive graviton theory, we have considered only the propagation effects of a massive graviton, and have assumed that the wave generation is correctly given by GR. If BHs in the true massive-graviton theory are significantly different from the BHs in GR, this could affect our estimates. Additionally, we have considered only the leading-harmonic gravitational waveforms (in GR) produced by non-spinning BH binaries.  For the case of binaries with high mass ratios, a considerable fraction of the emitted energy is deposited in the higher harmonics (see, e.g., \\cite{Berti:2007fi}). Also, the current understanding is that most of the BHs in nature may be spinning, with possibly high spin magnitudes (see, e.g., \\cite{Volonteri:2004cf, Gammie:2003qi,Shapiro05}).  This imperfect description of the model waveforms can introduce systematic errors in the estimated parameters (see, e.g., ~\\cite{Ajith:ParamEstimSyst}), including in the massive graviton bounds, which this paper does not try to address.  In this paper, the error bounds are estimated by means of the Fisher information matrix approach. This approach has a number of limitations (see e.g., \\cite{Vallisneri:2007ev} and \\cite{Zanolin:2009} for a discussion), including the fact that this is valid only in the limit of large SNRs, does not include various priors in the parameters (such as $\\eta \\in (0, 0.25]$), nor the known issues involved in real data analysis pipelines. These limitations are better addressed by techniques like Markov-Chain Monte-Carlo methods."
    },
    "1004/1004.4573_arXiv.txt": {
        "abstract": "We investigate the prospects for indirect detection of fermion WIMPless dark matter at the neutrino telescopes IceCube and DeepCore. The dark matter annihilating in the Sun is a hidden sector Majorana fermion that couples through Yukawa couplings to a connector particle and a visible sector particle, and it exhibits only spin-dependent scattering with nuclei via couplings to first generation quarks. We consider cases where the annihilation products are taus, staus, or sneutrinos of the three generations. To evaluate the muon fluxes incident at the detector, we propagate the neutrino spectra through the solar medium and to the Earth and account for the effects of neutrino oscillations, energy losses due to neutral- and charged-current interactions, and tau regeneration. We find that for the stau and sneutrino channels, a 5 yr 3$\\sigma$ detection of dark matter lighter than about 300~GeV is possible at IceCube for large Yukawa couplings or for dark matter and connector particles with similar masses. The tau channel offers far better detection prospects. However, due to its lower energy threshold and better muon background rejection capability, DeepCore is able to detect signals in all annihilation channels and for a wider range of dark matter masses. ",
        "introduction": "There has long been interest in the use of neutrino experiments for indirect dark matter (DM) detection, with a particular focus on neutrino signals from DM annihilation in the Sun. The basic idea is that DM particles can be captured by the Sun if they lose sufficient kinetic energy through elastic scattering from solar nuclei.  This leads to an enhanced DM number density in the solar core, where DM can annihilate to standard model (SM) products. These products in turn decay, producing a neutrino flux which can be detected by Earth-based detectors. Because the Sun is (for most models under consideration) in equilibrium, the DM capture rate is equal to twice the annihilation rate.  Thus, the neutrino flux on Earth is determined by the capture rate, which in turn is determined (up to a few ${\\cal O}(1)$ factors related to solar physics) by the DM mass $m_{\\chi}$ and the DM-nucleon scattering cross section $\\sigma$. This is one of the key advantages of neutrino-based indirect searches for DM: the event rate is directly related to the scattering cross section, without many of the astrophysical uncertainties which attend other indirect detection strategies.  Moreover, it is in principle possible to directly compare search results at neutrino experiments with those at direct-detection experiments, which also measure the DM-nucleon scattering cross section; for a recent discussion, see \\cite{Wikstrom:2009kw}.  This allows one to corroborate a signal at one type of experiment with a signal at a very different type of experiment, which is desirable.  The downside, however, is that for much of the $(m_\\chi,\\sigma)$-parameter space, sensitivities for current and future neutrino detectors have already been surpassed by direct-detection bounds, implying that neutrino detectors may have difficulty providing new input to DM studies.  However, there are two scenarios where neutrino experiments are expected to shine. One is at low-mass ($m_\\chi \\sim 4-10~\\gev$).  In this mass range, bounds from direct-detection experiments tend to become significantly worse because the nuclear recoil energies often fall below the experimental threshold. Experiments such as Super-Kamiokande (Super-K) and other proposed water Cherenkov or liquid scintillator-type detectors can provide the best bounds in this range~\\cite{lightDMdetect}, and could be sensitive to spin-independent (SI) scattering cross sections $\\sigma_{SI} \\sim 10^{-5}~\\pb$. DM in this range can potentially explain the annual modulation signal seen by the DAMA experiment (\\cite{lightDMDama}, though see also \\cite{Fairbairn:2008gz}), as well as unexplained events recently reported by the CDMS and CoGeNT Collaborations~\\cite{Ahmed:2009zw,Aalseth:2010vx,Cogentcdmstheory}. There is thus considerable interest to cross-check these results with neutrino detectors.  Another scenario for which neutrino experiments can excel is for models where the DM-nucleon scattering is largely spin-dependent (SD), rather than SI~\\cite{sigmaSDneutrino}. The reason is that DM can then be captured by SD scattering from hydrogen in the Sun, leading to significant neutrino fluxes.  Current direct-detection experiments provide much less sensitivity to SD scattering and the best bounds are set by indirect detection experiments, even at large $m_\\chi$.  This has been the focus of large neutrino detectors, such as AMANDA or IceCube.  For example, with 1800 live days of running, IceCube could be sensitive to ${\\cal O} (100~{\\rm GeV})$ DM with $\\sigma_{SD} \\sim 10^{-5}~{\\rm pb}$~\\cite{Braun:2009fr}. Although the relatively high threshold energy of these detectors limits their utility in studying the low-mass DM, their large size makes them ideal for studying higher-mass DM ($m_{\\chi} \\geq 100~{\\rm GeV}$) with $\\sigma_{SD} \\gg \\sigma_{SI}$.  Direct detection is much more sensitive to $\\sigma_{SI}$ due to the possibility of coherent scattering in the nucleus.  Direct-detection experiments, including CDMS and XENON100, currently have a sensitivity to $\\sigma_{SI}$ in the ${\\cal O} (100~{\\rm GeV})$ mass range of about 3 orders of magnitude~\\cite{Ahmed:2009zw,XENON100limits} greater than the IceCube projected 1800 d sensitivity to $\\sigma_{SD}$. But for models with $\\sigma_{SI} \\ll \\sigma_{SD}$, complementary coverage from neutrino detectors becomes important.  This complementary coverage is especially important  if one goes beyond models of weakly interacting massive particles (WIMPs). The recently proposed WIMPless model of dark matter~\\cite{Feng:2008ya,Feng:2008dz} is notable because it provides a very generic hidden sector DM candidate which naturally has about the right relic density to match astronomical observations, regardless of the mass of the DM particle or the details of the hidden sector.  This versatility suggests the possibility of specific WIMPless models with almost exclusively SD scattering, for which the best detection prospects would lie at neutrino detectors.  We present a model of WIMPless dark matter where the DM candidate is a Majorana fermion in the hidden sector.  At tree level, it will exhibit only SD scattering with SM nuclei.  We find that significant neutrino fluxes on Earth arise from models where DM annihilates either to tau, stau, or sneutrino pairs.  We consider the detection prospects for this type of model at DeepCore and IceCube.  The paper is organized as follows: In Sec.~2, we review WIMPless dark matter and describe the properties of the fermion candidate DM particle. In Sec.~3, we obtain the neutrino spectra arising from the decay of various annihilation products. In Sec.~4, we discuss the predicted event rates at DeepCore and IceCube, and we present our conclusions in Sec.~5. ",
        "conclusions": "We investigated the prospects for indirect detection of fermion WIMPless DM at IceCube and DeepCore. We considered a hidden sector Majorana DM particle of mass $m_{X}$ that couples through Yukawa couplings in the superpotential to a connector of mass $m_{Y}$ and visible sector particles.  These models are especially interesting in the context of IceCube/DeepCore because they exhibit only SD nuclear scattering, for which IceCube/DeepCore is expected to soon provide the greatest experimental sensitivity for $m_X \\sim {\\cal O} (\\textrm{few } 100~{\\rm GeV})$.   We focused our attention on DM annihilation to taus, staus and sneutrinos. In order to be captured by the Sun the DM particle needs to couple also to up and down quarks and first generation squarks through the superfield. Annihilation to light fermions does not produce an observable signal because it is chirality suppressed. As for annihilation into squarks, in the GMSB-inspired WIMPless scenario, squarks are generally very heavy and the resulting neutrino spectra are strongly dependent on the features of the sparticle spectrum. Therefore, we assumed that the dominant channels for annihilation are exclusively leptonic. We assumed that if the DM candidate annihilates to SM particles, these are predominantly taus. If it annihilates to supersymmetric particles, these belong to lepton superfields and are the NNLSP, with the lightest neutralino being the NLSP. In the cases of stau and sneutrino annihilation channels we focused our attention on the distinct situations where either the DM candidate couples predominantly to the third family (staus), or it couples with equal strength to degenerate sneutrinos of all three families. We believe this allows us to consider a range of interesting possibilities.  We propagated the neutrino spectra originated by DM annihilation at the center of the Sun through the solar medium. We took into account energy losses due to NC~interactions, CC~interactions, and tau~regeneration. We also implemented neutrino oscillations in the propagation through the Sun.  In order to calculate the event rate at the IceCube detector, we considered upward and contained events, taking into account muon energy losses due to ionization, bremsstrahlung, pair production, and photonuclear effects. The technical details of our analyses are provided in the appendices.  We found that it is not possible to obtain 3$\\sigma$~detection in upward or contained events at IceCube for the stau and sneutrino channel for DM masses $m_{X}\\lesssim 150$~GeV with 5~yrs of data. Moreover, a 300~GeV DM particle would not produce observable events for the sparticle channels, unless the Yukawa couplings are large or the connector mass is not more than about 30\\% heavier than the DM mass. Indirect detection of these relatively light DM particles is favored in the tau~channel, due to the broader neutrino spectra produced in this channel. We quantified the improvement DeepCore brings to the detection prospects in all channels, especially for low $m_{X}$. In particular, even in the cases mentioned above, where the steep decrease in flux below the detector energy threshold makes signal detection highly unlikely at IceCube, such limitations do not apply to DeepCore. While the performance of the IceCube detector varies significantly between different channels and different DM masses, thus requiring a case by case analysis, we showed that DeepCore can comfortably produce a 5~yr 3$\\sigma$~detection for any analyzed channel, without strong dependence on the Yukawa couplings or connector masses. We thus find good prospects for probing models of Majorana fermionic WIMPless DM at IceCube, including the DeepCore extension."
    },
    "1004/1004.4603_arXiv.txt": {
        "abstract": "The Rastall's theory is a modification of the General Relativity theory leading to a different expression for the conservation law in the matter sector compared with the usual one. It has been argued recently that such a theory may have applications to the dark energy problem, since a pressureless fluid may lead to an accelerated universe. In the present work we confront the Rastall's theory with the power spectrum data. The results indicate a configuration that essentially reduces the Rastall's theory to General Relativity, unless the non-usual conservation law refers to a scalar field, situation where other configurations are eventually possible. ",
        "introduction": "A large number of cosmological observational data requires the existence of two exotic components in the matter content of the universe, dark matter and dark energy \\cite{kamion}. Both constitute the so-called dark sector of the cosmic budget. Dark matter is necessary, for example, to explain the formation of structures in the expanding universe. Dark energy is necessary, on the other hand to explain the present stage of the accelerated expansion of the universe. Dark energy must have a negative pressure in order to induce the acceleration of the universe. This is a quite strange property, and one of the most important question today in theoretical physics concerns the nature of such an exotic fluid. All candidates to describe the dark energy component face serious problems and drawback. \\par Another possibility to explain the observational data is to consider that General Relativity is not the true gravitational theory. Some modifications of General Relativity may lead to an accelerated universe even if only usual types of matter are taked into account. A very fashion proposal is the $f(R)$ theories, a non-linear generalization of the Einstein-Hilbert action . For a review of the $f(R)$ theories and their present status, see reference \\cite{felice}. In fact, the unusual properties of dark energy motivate the search of many other alternatives to explain the observational data. \\par Recently, it has been investigated if the Rastall's theory may be a viable alternative to the introduction of dark energy in the General Relativity context \\cite{capone}. The Rastall's theory is a modification of the General Relativity, proposed in 1972 \\cite{rastall}. It implies a change of the Einstein's equation that mounts out to a modification of the conservation law for the energy-momentum tensor. One of the motivations to this proposal is the fact that the usual conservation law is only firmly tested in special relativity. The theory contains a free parameter $\\lambda$ such that $\\lambda = 1$ implies General Relativity. In this way, it can be considered as a deformation of General Relativity. The modification in the conservation law can lead to new effects compared with General Relativity. Depending on the value of the parameter $\\lambda$, a pressureless matter, for example, can induce an accelerated expansion. \\par An important drawback of Rastall's theory is the absence of the Langrangian formulation. But, it is possible that it can obtained from an action principle using Weyl's geometry \\cite{smalley}. It has been argued in reference \\cite{lee} that the Rastall's theory may appear even in the context of Riemannian geometries by a redefinition of the energy-momentum tensor. In any case, the Rastall's theory implies that a given fluid, with a specific equation of state, may have a very different effective equation of state when interpreted in the context of General Relativity. This properties has called the attention to this proposal as an alternative to dark energy. \\par In reference \\cite{capone}, the Rastall's theory has been confronted agains the supernova type Ia data. The result indicates that it is competitive with respect to the $\\Lambda$CDM model, but the mass density is very high, around $\\Omega_0 \\sim 0.5$. This suggests that the theory may face troubles when tested against the mass agglomeration phenomena. \\par In the present work, we intend to test the Rastall's theory using the power spectrum observational data. This implies a perturbative study. A perturbative study of the Rastall's theory has already been performed in reference \\cite{kerner}. The main conclusion was that the theory behaves also at perturbative level like General Relativity but with an effective equation of state. In particular, if the fluid has an effective equation of state necessary to induce the accelerated expansion of the universe, it will present a negative squared effective sound speed. Hence, the resulting scenario is plagued with instabilities. Moreover, in order to compare the theoretical predictions of the Rastall's theory with the power spectrum data, it seems unavoidable to introduce a two fluid model. This represents a problem in the original Rastall's theory. A possible way out is to consider a fluid which obeys the conservation law of Rastall's theory, with another one that conserves separately in the usual way. This may be justified if the former fluid is in fact a field, like a scalar field, that obeys a modified Klein-Gordon equation. In fact, if both fluids have a hydrodynamical representation, the power spectrum predicts $\\lambda = 1$ as it will be described later in this work - that is, the theory reduces to General Relativity. If the Rastall's fluid is a scalar field, the situation is more complex, and even if $\\lambda = 1$ remains favored, other possibility appears. \\par This paper is organized as follows. In next section, we present the Rastall's theory, deriving some cosmological relations. In section III, the matter power spectrum is determined in the hydrodynamical representation. In section IV, the matter power spectrum is determined for the case where one of the fluids is representing by a self-interacting scalar field. In section V we present our conclusions. \\par ",
        "conclusions": "In this work we have investigate the Rastall's proposal of modification of General Relativity with respect to the problem of structure formation. This proposal is equivalent to change the usual conservation law of General Relativity. In some sense, it can be seen as a modification of the equation of state of a given fluid from the dynamical point of view. In this sense, the Rastall's theory can be interesting in order to obtain an accelerated expansion of the universe without introducing exotic fluids. For example fluids with positive or null pressure may induce a dynamics typical of fluids of negative pressure. This behaviour has already been remarked, for example, in the Brans-Dicke theory \\cite{rose}. \\par Our results indicate that the Rastall's theory faces many problems at perturbative level. Considering a homogenous and isotropic universe for example, the effective equation of state of the background is entirely reproduced at perturbative level, leading to high instabilities when this effective equation of state implies negative pressure. One way out is to consider a two-fluid model, one of the fluids obeying the usual conservation law and the other one following the Rastall's prescription. The results point out to a configuration which still reduces the Rastall's theory to General Relativity. \\par When we keep the framework of two components which obeys different conservation laws, but with a self-interacting scalar field playing the role of the ``Rastall's fluid'', more interesting features appear. In particular, the Rastall's proposal can be seen as a modification of the Klein-Gordon equation, similarly to what happens in the K-Essence theories \\cite{mukhanov}. Again the configuration corresponding to General Relativity is favored, but other configurations with $\\gamma \\sim 0$ are also possible even if some statistical subtleties appear. All these considerations seem to indicate to the specificity of General Relativity and rule out the Rastall's theory. \\par A possible way out to save the Rastall's proposal is to consider the modification of the usual conservation law as a manifestation of quantum effects (like particle production) in the spirit of reference \\cite{ademir}. The particle production in an expanding universe may lead to new terms in the usual conservation law. Such possibility remains to explore. In any case, it seems clear that structure formation asks for a fluid we behaves in the background and at perturbative level with zero effective pressure (otherwise no mass agglomeration can effectively occur), and this poses a serious problem in the original framework of the Rastall's theory. In some sense, this has already been remarked in reference \\cite{capone}, forcing the authors of that work to use just one fluid to fit the supernova type Ia data. But, such procedure seems to be impossible concerning the structure formation problem. \\par Other observational methods, like CMB, may be used to constrain better the Rastall's theory. But, the results obtained in the present work seem to us to be strong enough to point out the difficulties that the Rastall's theory face when confronted with observational data. \\par \\noindent {\\bf Acknowledgements:} We thank Oliver Piattela, Winfried Zimdahl and Hermano Velten for their remarks and criticisms on the text. We thank also CNPq (Brazil) and FAPES (Brazil) for partial financial support."
    },
    "1004/1004.2739_arXiv.txt": {
        "abstract": "{ It has been established earlier that sharp features like the base of the convective zone or the second helium ionisation zone inside a star give rise to sinusoidal oscillations in the frequencies of pulsation. The acoustic depth of such features can be estimated from this oscillatory signal in the frequencies. We apply this technique for the CoRoT\\thanks{The CoRoT space mission was developed and is operated by the French space agency CNES, with the participation of ESA's RSSD and Science Programmes, Austria, Belgium, Brazil, Germany, and Spain.} frequencies of the solar-type star \\hdf. This is the first time that such analysis has been done of seismic data for any star other than the Sun.  We are able to determine the acoustic depth of both the base of the convective zone and the HeII ionisation zone of \\hdf\\ within $10\\%$ error from the second differences of the frequencies. The locations of these layers using this technique is in agreement with the current seismic models of \\hdf.  } ",
        "introduction": "\\label{sec:intro}  A discontinuity in the derivatives of the sound speed in the stellar interior introduces a characteristic oscillatory signature in the frequencies of low degree modes of the form \\citep{gough90} \\[ \\sin (4\\pi\\nu\\tau + \\phi) \\] where $\\tau$ is the acoustic depth of the layer, lying at a radius $r$, given by \\[ \\tau = \\int_r^R \\mathrm{d}r/c \\] and $\\nu$ is the frequency of oscillation. $\\phi$ is a phase factor.  In solar-type stars, the base of the convective envelope and the HeII ionisation zones are the most prominent layers of such discontinuity. Each of these layers contributes to a sinusoidal variation in the frequencies of pulsation with the periodicity corresponding to the respective acoustic depth of the layer.  The resulting combined signal in the frequencies can be fitted to a suitable functional form to extract the acoustic depths of these layers separately.  Here we apply this technique to estimate the acoustic depths of these regions of the solar-type star \\hdf\\ from its frequencies obtained from the CoRoT satellite \\citep{baglin06}. Although this method has been used with great success for the Sun \\citep[see e.g.,][]{ban94,mct94}, this is the first instance of an application to a distant star. This was one of the main objectives of the CoRoT seismology programme \\citep{michel06}.  The significance of this method is that it does not depend on theoretical stellar models -- the position of the sharp features are estimated directly from the frequencies themselves. In fact, such independent determination of these layers can help in the modelling of the star, as illustrated by \\citet{mazumdar05}. We describe the technique used in this work and the results in the next two sections. ",
        "conclusions": "\\label{sec:summary}  We have determined the acoustic positions of the base of the convective envelope and the second helium ionisation zone in the star \\hdf\\ from the oscillatory signal in its frequencies observed with CoRoT. The method involves fitting a suitable function to represent the two oscillatory signals arising from the two zones to the second differences of the frequencies. We find almost similar results for two data sets, one corresponding to only the initial 60-day run of CoRoT and another a 180-day combined time-series obtained with a long run of CoRoT. This is the first instance where such a technique has been successfully applied to a star other than the Sun.  Although these results are independent of any stellar modelling, they match well with a seismic model of the star. In fact the acoustic depths of such layers of sharp changes in the sound speed inside the star extracted in this method can help to model the star \\citep{mazumdar05}.  The same technique can be used to estimate the acoustic depths of such layers in other stars. We have established that a time series of moderate length, as obtained by CoRoT, is sufficient for the successful application of this technique. A longer time series, as expected from the Kepler mission, for example, would further reduce the errors on the extracted acoustic depths."
    },
    "1004/1004.5119_arXiv.txt": {
        "abstract": "\\begin{center} {\\large\\bf Abstract} \\end{center} ``Quasi-stationary'' states are approximately time-independent out of equilibrium states which have been observed in a variety of systems of particles interacting by long-range interactions. We investigate here the conditions of their occurrence for a generic pair interaction $V(r \\rightarrow \\infty) \\sim 1/r^\\gamma$  with $\\gamma > 0$, in $d>1$ dimensions. We generalize analytic calculations known for gravity in $d=3$ to determine the scaling parametric dependences of their relaxation rates due to two body collisions, and report extensive numerical simulations testing their validity. Our results lead to the conclusion that, for $\\gamma < d-1$, the existence of quasi-stationary states is ensured by the large distance behavior of the interaction alone, while for $\\gamma > d-1$ it is conditioned on the short distance properties of the interaction, requiring the presence of a sufficiently large soft core in the interaction potential. ",
        "introduction": " ",
        "conclusions": ""
    },
    "1004/1004.0399_arXiv.txt": {
        "abstract": "We have discovered a large number of circular and elliptical shells at 24$\\mu$m around luminous central sources with the MIPS instrument on-board the {\\sl Spitzer Space Telescope}. Our archival follow-up effort has revealed 90\\% of these circumstellar shells to be previously unknown. The majority of the shells is only visible at 24$\\mu$m, but many of the central stars are detected at multiple wavelengths from the mid- to the near-IR regime. The general lack of optical counterparts, however, indicates that these sources represent a population of highly obscured objects. We obtained optical and near-IR spectroscopic observations of the central stars and find most of these objects to be massive stars. In particular, we identify a large population of sources that we argue represents a narrow evolutionary phase, closely related or identical to the LBV stage of massive stellar evolution. ",
        "introduction": "}  The {\\sl Spitzer Space Telescope} has provided us with an unprecedented view of the mid-IR sky in terms of both resolution and sensitivity. Invariably, the opening of every new wavelength window carries with it the potential for the discovery of previously unrecognized phenomena and source populations, which have gone unnoticed at other wavelengths. Motivated by our discovery of an unusual well-defined elliptical shell -- lacking a central energizing source and only visible at 24$\\mu$m -- in some of the first data from the {\\sl Spitzer} mission \\citep{morris06}, we executed a search for similar objects that might shed light on its still puzzling origin and nature.  We utilized the publicly available 24$\\mu$m data products (delivery Version 3.0) of the MIPSGAL Legacy project \\citep{carey09} obtained with the Multiband Imaging Photometer for {\\sl Spitzer} (MIPS, \\citealt{rieke04}). We conducted our search by eye, inspecting all the images of the inner 248 square degrees of the Galactic plane covered by the MIPSGAL survey. Although there was no prediction of whether 24$\\mu$m shells would be rare or common occurrences, we were still surprised by our discovery of a large number of prominent, more or less ring-shaped 24$\\mu$m objects: we find over 200 such shells in the MIPSGAL survey area. Unlike the shell that originally motivated the search, however, the majority of the objects have some indication of a central source, at times quite luminous even at 24$\\mu$m.  Follow-up investigation utilizing the complementary data from the GLIMPSE {\\sl Spitzer} Legacy project \\citep{benjamin03} which covers the same survey area with the IRAC 3.5--8.0$\\mu$m bands, as well as 2MASS and the Digitized Sky Survey (DSS), shows that about 60\\% of these shells are {\\it only} detected at 24$\\mu$m (see Figure~\\ref{f-charts}). This is somewhat unusual, as we generally expect a strong emission component due to polycyclic aromatic hydrocarbons (PAHs) at 8$\\mu$m if we are simply observing warm dust continuum emission. In the case of the shell discussed by \\cite{morris06}, the lack of 8$\\mu$m detection is explained by a pure line emission spectrum with a dominant [OIV] 25.89$\\mu$m line. We speculate that some of our 24$\\mu$m--only shells might exhibit similar mid-IR spectral characteristics.  Ring-like shell morphology is observed frequently in planetary nebulae or in material associated with a variety of massive stars, such as Wolf-Rayet (WR) stars and red and blue supergiants \\citep{chu91}. So far, most known circumstellar rings and shells have largely been studied using narrow band optical filters (e.g.\\ H$\\alpha$, [OIII]) and, despite being labeled as ring-shaped, many of these nebulae appear quite amorphous. For our {\\sl Spitzer}--discovered shell sources, a SIMBAD search within 2\\arcmin\\ of the shell locations revealed that most (90\\% of the shells, 80\\% of the central sources) of these objects have not been previously studied. We present here the results of our efforts to determine the nature of the central stars that produce the 24$\\mu$m shells and to characterize the properties of this substantial population of new sources. ",
        "conclusions": "We have discovered a large number of prominent, highly symmetric 24$\\mu$m shells surrounding bright central sources. We were able to determine the nature of the apparent central sources for 45 of the 62 shells, 10 based on literature searches and 35 from dedicated spectroscopic follow-up. Our infrared and optical spectroscopy has revealed a mix of early and late type stars. However, the prevalence of rare, massive stars stands out. We find six bona fide WR stars in our sample (four of which are new identifications), as well as four stars which we tentatively classify as Oe/WN. Most strikingly, our sample of shells contains four known LBV type stars (one confirmed, three candidates), and we classify an additional six sources as LBV candidates based on the similarities between the spectra. We realize that the classification of a particular source as an LBV involves a host of criteria, including photometric and spectroscopic variability, that still need to be investigated for these new sources. However, given the small numbers of known LBVs, the presence of a shell combined with the spectroscopic characteristics make this relatively large sample of new candidates an exciting discovery. We also identify five new Be stars in our sample. The distinction between these and what we characterize as LBV candidates is somewhat arbitrary and simply conveys less peculiar spectral properties. It is entirely possible that these stars harbor additional LBV candidates, especially considering LBV spectral variability and the fact that the LBVs AG Car and HR Car have been classified as Be stars.  The detailed evolutionary sequence for the most massive stars is still a matter of intense debate. The standard model of massive star evolution (see, e.g., the review by \\citealt{crowther07}) expounds that O stars with initial masses of 25--40$M_\\odot$ will undergo an LBV and RSG phase on their way to becoming a WR star. More massive stars on the other hand, are thought to avoid the RSG phase. \\cite{smithowocki06} contend that substantial mass loss during an LBV stage is critical to the formation of WR stars. Analysis of the nebulae surrounding the known LBVs support the view that the nebulae formed during a BSG and not during a RSG phase \\citep{lamers01}. \\cite{smithconti08} strongly argue in favor of the notion that H-rich WN type stars precede the LBV stage, while H-poor WN stars represent a post-LBV evolutionary phase. In this context we point out that our Oe/WN classified sources all show strong H emission, and hence could be interpreted as pre-LBV stars. In fact, we speculate that a large number of the central stars associated with our most well-defined, circular shells represent a narrow evolutionary phase, closely related or identical to the LBV phase. In depth analysis of the shell emission is needed to further investigate this hypothesis. Finally, we predict that our upcoming observations of the remainder of our sample will reveal additional members of this pre-LBV/LBV group among the as yet unclassified shell central stars."
    },
    "1004/1004.0020.txt": {
        "abstract": "We present uniform CFHT Megacam $g$ and $r$ photometry for 34 X-ray selected galaxy clusters drawn from the X-ray Multi-Mirror (XMM) Large Scale Structure (LSS) survey and the Canadian Cluster Comparison Project (CCCP). The clusters possess well determined X-ray temperatures spanning the range $1<kT(\\rm keV)<12$.  In addition, the clusters occupy a relatively narrow redshift interval ($0.15<z<0.41$) in order to minimize any redshift dependent photometric effects. We investigate the colour bimodality of the cluster galaxy populations and compute blue fractions using criteria derived from Butcher and Oemler (1984). We identify a trend to observe increasing blue fraction versus redshift in common with numerous previous studies of cluster galaxy populations.  However, in addition we identify an environmental dependence of cluster blue fraction in that cool (low mass) clusters display higher blue fractions than hotter (higher mass) clusters.  Finally, we tentatively identify a small excess population of extremely blue galaxies in the coolest X-ray clusters (essentially massive groups) and note that these may be the signature of actively star bursting galaxies driven by galaxy-galaxy interactions in the group environment. ",
        "introduction": "The observed properties of galaxy populations reflect the environment in which they are located.  Comparisons of galaxy populations drawn from low (the field) and high density (rich galaxy clusters) environments indicate that the population distribution described using measures such as current star formation rate (e.g. \\citealt{balogh99}, \\citealt{pogg06}), integrated colour (e.g. \\citealt{blanton05}) and morphology (e.g. \\citealt{dressler97}, \\citealt{treu03}) varies as a function of changing environment. From studies such as these it is clear that galaxies located in the cores of rich clusters display lower star formation rates, redder colours and more bulge dominated morphologies compared to galaxies located in the field.  Studies of the fraction of blue galaxies contributing to a galaxy cluster provided some of the first direct evidence for the physical transformation of galaxies in cluster environments. \\cite{bo84} reported an increase in the fraction of blue galaxies in 33 rich galaxy clusters out to $z\\sim0.5$ compared to local clusters. However, subsequent studies designed to expand upon this initial discovery highlighted the many complexities associated with this relatively straightforward technique, including varying intrinsic cluster properties with redshift (e.g. X-ray luminosity as discussed by \\citealt{andreon99}), the use of $k$-corrections to determine rest frame colour distributions \\citep{andreon05} and the challenge of obtaining large samples with uniform photometry (successfully overcome by \\citealt{loh08}).  In addition to the Butcher-Oemler (BO) effect measured employing optical galaxy colours, analogous BO-type effects have been reported as either a morphological BO effect (increasing spiral fraction in clusters with increasing redshift; \\citealt{pogg99}) and an infra-red (IR) BO effect (increasing fraction of dust enshrouded star forming falaxies in clusters with increasing redshift; \\citealt{duc02}; \\citealt{saintonge08}) to name two examples.  An alternative approach is to consider an environmental Butcher-Oemler effect whereby one attempts to determine the variation of blue fraction as a function of varying intrinsic cluster properties selected over narrow redshift intervals.  This has been achieved by comparing blue fractions within clusters at increasing clustercentric radii (e.g. \\citealt{ellingson01}; \\citealt{loh08}) or by considering blue fractions measured between clusters of differing X-ray luminosities (e.g. \\citealt{wake05}).  A number of studies report the decrease of the fraction of blue galaxies with decreasing scaled clustercentric radius, e.g. the virial radius determined employing either cluster dynamics \\citep{ellingson01}, correlation properties \\citep{loh08}, or extrapolated from X-ray properties \\citep{wake05}.  The currently favoured explanation for these observed trends is that infalling field galaxies are processed physically as they travel from the field, through the cluster outskirts and virialise in the central cluster region \\citep{berrier09}.  However, the extent to which cluster galaxies were ``pre-processed'' by physical effects occuring during an earlier residence in a galaxy group\\footnote{Note that we generally refer both groups and clusters of galaxies as \"clusters\" in this paper and quantify our description using the X-ray temperature.} remains contested \\citep{li09}.  Numerous physical processes have been suggested as the agents of this apparent transformation of galaxy populations.  However, the dominant physical process(es) to which an infalling galaxy is subject remains unclear.  The two principal examples of such processes are ram pressure stripping and galaxy galaxy interactions.  Ram pressure stripping describes the effective force experienced by the diffuse gas component of the infalling galaxy as it travels through the hot, dense intra-cluster medium (ICM; \\citealt{gunn72}).  Both hot and cold gas may be stripped from the infalling galaxy leading to the exhaustion of the available gas supply that will otherwise cool and form stars (\\citealt{abadi99}; \\citealt{kawata08}; \\citealt{mccarthy08}). Hydrodynamical simulations have indicated that the ICM associated with galaxy group and cluster environments will strip of order 70\\% to 100\\% respectively of hot gas from a typical infalling spiral galaxy within one crossing time (\\citealt{kawata08}; \\citealt{mccarthy08}). The effect is manifest as a sharp decline in the galaxy star formation rate effective on a timescale comparable to the rate at which the unstripped cold gas supply is consumed.  The observations of a population of red (i.e. passive) spiral galaxies contributing to the red sequence would appear to support the view that some fraction of galaxies experience the stripping of disc gas via ram pressure effects \\citep{wolf09}.  The observation of galaxies displaying extended H{\\sevensize I} tails in the Virgo cluster \\citep{chung07} is nominally consistent with the expectations of ram pressure stripping. However, we note that the authors also report the presence of close companions to these galaxies in a number of cases and comment that a combination of ram pressure and tidal stripping provides a more compelling explanation.  When referring to galaxy-galaxy interactions we note that this may indicate one of a wide range of encounters.  Interactions may be predominantly tidal between close neighbours resulting in halo gas being moved outward where it is more readily stripped (e.g. \\citealt{chung07}).  High speed encounters (either referred to as harrassment or threshing in the literature) may also result in tidal stripping of halo gas \\citep{moore96}.  Finally, infalling galaxies may merge with existing cluster members. The products of such merger encounters may be predicted by considering the mass ratio of the merging galaxies: large mass ratios result in enhanced star formation in the satellite galaxy yet may not lead to a star burst in the more massive companion \\citep{cox08}.  Equal mass mergers on the other hand result in the complete disruption of the infalling spiral galaxy (for example) to form a bulge dominated system accompanied by a significant central star burst \\citep{matteo07}.  The internal disruption associated with such strong interactions can result in a short term enhancement in star formation followed by a rapid exhaustion of the available cold gas supply.  The potential effect of such strong galaxy-galaxy interactions is of interest to studies of the environmental dependence of galaxy evolution as the merging rate (the product of relative velocity and interaction cross section) is predicted to be a strong function of environment: the cross section for disruptive encounters increases as the cluster velocity dispersion approaches the internal velocity of the infalling galaxy \\citep{makino97}.  Compelling evidence for enhanced galaxy-galaxy interactions in rich cluster environments is provided by Hubble Space Telescope (HST) observations indicating a high merger fraction in such environments compared to field comparison samples (\\citealt{dressler94}; \\citealt{vandokkum99}).  However, such observations must be contrasted with mid-infrared selected moderate starburst galaxies located in rich cluster environments whose optical morphologies resemble undisturbed spiral galaxies (\\citealt{geach09}; \\citealt{oemler09}).  A picture is therefore emerging whereby multiple mechanisms (ICM stripping, merger induced star formation and tidally induced star formation) may participate in the physical processing of infalling galaxies.  Currently unanswered questions focus upon whether more than one physical process acts upon a typical galaxy falling into a dense environment (and which might be considered dominant) and whether the relative importance of each of the suggested physical processes changes as a function of the global properties of the group/cluster environment into which the galaxy falls.  In this paper we attempt to answer the related question of whether the effectiveness of galaxy processing can be determined as a function of global cluster environment. Our approach is to compute the blue fraction as defined by \\cite{bo84} for a sample of X-ray selected galaxy clusters and to determine whether the blue fraction displays a significant trend versus X-ray temperature (here employed as a proxy for the global cluster mass).  Throughout this paper, values of $\\Omega_{\\rm M,0}=0.3$, $\\Omega_{\\Lambda,0}=0.7$ and $H_0=70 \\rm \\, kms^{-1} \\, Mpc^{-1}$ are adopted for the present epoch cosmological parameters describing the evolution of a model Friedmann-Robertson-Walker universe. All magnitude information is quoted using AB zero point values. ",
        "conclusions": "\\label{sec_conc}  We have presented an analysis of the fraction of blue galaxies in a large sample of X-ray clusters spanning a wide range of X-ray temperature with uniform optical photometry.  We have noted that the computed blue fraction values display a trend to decrease with both decreasing redshift and increasing temperature.  The trend of decreasing cluster blue fraction versus redshift has been observed previously and has been interpreted as a global effect driven by the decreasing rate of field galaxy infall onto the cluster environment (e.g. \\citealt{ellingson01}).  Coupled with the increasing cluster mass with decreasing redshift, the specific infall rate per unit cluster mass decreases at an even faster rate.  Infall effects are further compounded with the sharply declining global star formation rate as a function of decreasing redshift \\citep{hopkins06} which may be viewed as a decrease in the available gas supply of infalling galaxies.  Looking beyond the redshift dependent component of the Butcher-Oemler effect, the decreasing blue fraction with increasing X-ray temperature identified in this paper is consistent with an environmental component to the Butcher-Oemler effect.  The trend of decreasing blue fraction versus temperature is most apparent whan comparing the across XMM-LSS and CCCP samples rather than within either sample. It remains possible that this difference is due in part to some subtle selection effect arising during the construction of the two samples rather than as a result of some physical process linked to the global environment.  However, it is not clear what form such a bias would take.  At this point we recall that the cool clusters may display a mild bias to higher gas densities or more centrally concentrated systems compared to the average population. One would naively expect such systems to be more effective at processing infalling field galaxies compared to the average (i.e. lower X-ray luminosity) cluster over the same temperature interval \\---\\ effectively introducing a mild bias toward lower blue fractions in this temperature sub-sample.  One could therefore argue that access to a more complete census of the low-temperature cluster population would only amplify the trend of blue fraction versus temperature.  It has also been noted that optically selected clusters display different optical properties compared to X-ray selected samples (see Haines et al. 2009 for a discussion).  However, X-ray samples of bright clusters are typically complete in terms of the optical properties they sample (mainly because bright X-ray clusters are rare and thus it is possible to select complete samples). Moreover, clusters selected by single band optical observations (i.e. purely based upon the projected overdensity of bright galaxies) are more likely to be biased toward high blue fraction values due to the high mass-to-light ratios of blue star forming galaxies.  We therefore note that any possible bias between the XMM-LSS and CCCP samples would be naively be expected to generate the opposite trend to that observed in this paper and we remain confident that the trend of blue fraction versus temperature we have identified is not a result of sample selection effects.  A currently favoured explanation for the observed trend of blue fraction versus temperature is that infalling field galaxies are processed physically as they interact with the cluster environment. This interaction may take the form of ram pressure stripping whose effectiveness is a relatively simple function of the mass scale represented by the group or cluster into which the galaxies are falling.  An alternative explanation is that infalling galaxies are processed via galaxy-galaxy interactions and therefore respond more readily to a combination of local rather than global velocity dispersion and galaxy density.  Whatever the cause, it is clear that the process by which blue galaxies are processed to appear red is more complete in hotter (more massive) environments compared to cooler (less massive) environments.  Disentangling the above two effects is difficult as, to first order, the effects upon star formation in the infalling galaxy depend on environment in the same manner: ICM stripping and subsequent SF suppression are expected to be weakest (though not absent) in group environments \\---\\ galaxy populations in these environments will display a larger fraction of blue galaxies relative to more massive clusters \\---\\ while SF enhancements associated with galaxy-galaxy interactions will be greater on group scales.  Each effect would naively generate the same observed trend of blue fraction versus temperature.  The blue fraction in each temperature sub-sample displays a very similar trend versus scaled radius and is consistent with an infalling field population albeit the cool sample is offset to higher blue fraction \\citep{ellingson01}.  However, the cool sample displays a marked increase in blue fraction with decreasing magnitude.  The increase in the cool sample blue fraction with increasing magnitude results from both a moderate increase in the number of blue galaxies but also in a relative deficit in the numbers of red galaxies compared to those in the mid and hot samples.  We will address the red sequence luminosity function of the cool, mid and hot samples in a subsequent paper. However, we note that the relative deficit of red galaxies in cool clusters compared to the mid and hot samples is consistent with the scenario where the processing of blue galaxies to red galaxies during infall is relatively incomplete.  Having subtracted the red sequence contribution from each colour histogram the blue galaxy colour distribution for each temperature sub-sample appears very similar.  We do note an excess of extremely blue galaxies in the cool sample compared to the mid and hot samples. The excess is small (an extremely blue fraction of 3\\% compared to $<1\\%$) yet significant. In addition the fraction of extremely blue galaxies increases in the cool sample from $\\sim3\\%$ to $\\sim 5\\%$ for galaxies displaying $r<21.5$ and $21.5<r<23$ respectively.  One may tentatively associate these galaxies as actively star-bursting galaxies potentially driven by galaxy-galaxy interactions (e.g. \\citealt{prop03}).  \\cite{prop03} comment that the absence of a strong Butcher-Oemler effect in near-infrared selected cluster galaxy populations compared to optically selected populations may be due to an increasing contribution from faint blue dwarfs whose optical brightness is boosted by recent star formation.  The extremely blue population identified in the cool sample is nominally consistent with this explanation yet clearly forms only a small component of the overall blue cluster galaxy population.  The small relative contribution of such potentially star bursting galaxies to the cluster population may be understood in terms of the short timescale over which a recent star burst will affect the integrated colour of an established stellar population.  This issue has been addressed by \\cite{barger96} who simulate the colour evolution with time of an underlying passive (elliptical-type) and continuous star formation (spiral-type) stellar population experiencing a short ($\\sim 100$~Myr), secondary ($\\sim 10\\%$ by mass) burst of star formation associated with an interaction resulting from the cluster environment. The shift towards bluer colours as a result of a moderate burst of star formation lasts typically only as long as the burst duration before the luminosity weighted colour reverts to that of the underlying pre-burst stellar population.  The effect is more pronounced when one applies the condition that all star formation ceases after the secondary burst, in which case the continuously star forming stellar population (spiral-type) reddens to approximately the same colour as the passive population within $\\sim 1$~Gyr of the burst.  Our analysis has identified a clear environmental dependence in the blue fraction of galaxies. However, the cause of this environmental Butcher-Oemler effect \\---\\ whether it be ram pressure stripping or galaxy-galaxy interactions \\---\\ cannot be discerned unambiguously on the basis of the colour distribution of the blue galaxy population in each sample.  Clearly further information is required, most notably the combination of colour information with morphological information. One would naively expect the above two environmental processes to affect galaxy morphology in markedly different ways, e.g. ram pressure effects hould not disrupt galaxy disks whereas a strong disruptive effect is expected from galaxy-galaxy interactions. We will present this investigation in a future paper."
    },
    "1004/1004.0957_arXiv.txt": {
        "abstract": "We study the coherent temperature and polarization patterns produced in homogeneous but anisotropic cosmological models. We show results for all Bianchi types with a Friedman-Robertson-Walker limit (i.e. Types I, V, VII$_{0}$, VII$_{h}$ and IX) to illustrate the range of possible behaviour. We discuss the role of spatial curvature, shear and rotation in the geodesic equations for each model and establish some basic results concerning the symmetries of the patterns produced. We also give examples of the time-evolution of these patterns in terms of the Stokes parameters $I$, $Q$ and $U$. ",
        "introduction": "Precise measurements of the temperature anisotropies of the Cosmic Microwave Background (CMB), particularly those from the Wilkinson Microwave Anisotropy Probe (WMAP) \\cite{WMAP1,WMAP2}, form the sturdiest foundations of the current (``concordance'') cosmological model describing a universe dominated by cold dark matter (CDM) and a cosmological constant \\cite{Coles}, and therefore known as $\\Lambda$CDM for short. It is an essential component of this model that the primordial metric perturbations that gave rise to the galaxies and large-scale structure we observe around us today should be Gaussian and statistically homogeneous (i.e. stationary) \\cite{Guth1982,Starobinskij1982,Bardeen1983} and this means that the temperature fluctuations observed in the CMB should be Gaussian and statistically isotropic.  Detailed analysis of the WMAP data has shown that any departures from the standard framework are small and of uncertain statistical significance. Although some anomalous behaviour has been reported {\\cite{Yadav2008}, there remains no clear evidence of primordial non-Gaussianity, but there are several indications of departures from statistical anistropy across the CMB sky. Among the interesting phenomena revealed by detailed analyses of the pattern of CMB temperature fluctuations are an extremely cold spot \\cite{Vielva,Cruz1,Cruz2,Cruz3,Cruz4,Cayon,Cspot}, unusual alignments between large-scale harmonic modes of the temperature pattern \\cite{Schwarz2004, Copi2004,Katz2004,Land2005a,Land2005b, Land2005c, Land2005d, Land2005e,Land2007, Copi2006, Copi2007} (sometimes dubbed ``The Axis of Evil'') and a global hemispherical power asymmetry \\cite{Eriksen2004a,Park2004,Eriksen2007,Hoftuft2009,Hansen2009}.  One must not get carried away with the these features because they are - almost without exception - based on {\\em a posteriori} evaluations of statistical significance. If the occurrence of an anomalous feature $A$ in a given model $M$ is  $P(A|M)\\simeq 0.001$ (say) then that does not necessarily mean that the probability of the model given that the anomaly is observed $P(M|A)$ is similarly low. One has to take into account all the other data that are {\\em not} anomalous before deciding on a true measure of the significance of a departure. Attempts to do this rigorously have generally confirmed that current observations are not sufficiently compelling to suggest that the standard model needs to be abandoned \\cite{HL09,GAWE,HL10,zb10}. Moreover, other analyses suggest the further possibility that the WMAP temperature fluctuations may be affected by foregrounds or other systematic problems \\cite{ccno07,cnc07,FP09,sc10}. Even slight effects of this type could seriously hamper our attempts to uncover evidence of physics beyond the standard model.  But, although the evidence for the examples of global asymmetry discussed above is by no means overwhelming, taken together they do at least suggest the possibility that we may live in universe which is described by a background cosmology that is globally anisotropic, i.e. one not described by a Friedmann-Robertson-Walker (FRW) model and it is therefore incumbent upon us to consider alternatives to the standard model in order to learn best how to use the data to confirm or rule out variations on the standard cosmology.  The approach we follow here is to study the Bianchi models, i.e. cosmological models based on exact solutions to Einstein's General Theory of Relativity involving homogeneous but not necessarily isotropic spatial sections. The Bianchi classification groups  all possible spatially homogeneous but anisotropic relativistic cosmological models into types depending on the symmetry properties of their spatial hypersurfaces \\cite{Gris,Ellis1}; we discuss these models in more detail later in the paper.  The Bianchi models are not particularly strongly motivated from the point of view of fundamental physics, but do nevertheless represent a promising and potentially profitable first step away from the standard cosmological framework. For example, it has been known for some time that interesting localized features in the CMB temperature pattern can occur in Bianchi models with negative spatial curvature \\cite{Collins1,Dautcourt1,Tolman1,Tolman2,Tolman3,Barrow1}.  The physical origin of such features lies in the focussing effect of space on the geodesics that squeezes the pattern of the anisotropic radiation field into a small region of the sky. The observed lack of large-scale asymmetries in the CMB temperature has in the past been used to place constraints on the global rotation and shear allowed in Bianchi models \\cite{Barrow1,BFS96,khb97}. More recently, however, attention has shifted to the possibility of using the additional parameters available in such models to reproduce a cold spot such as that claimed to exist in the WMAP data. Since we now know that our Universe is {\\em close} to isotropic, attention has focussed on the subset of the  Bianchi types that contain the FRW model as a limiting case.The model which appears to best able to reproduce the anomalous cold spot is the Bianchi VII$_h$ case \\cite{Jaffe1,Jaffe2,McEwan1,McEwan2,Bridge}, although Bianch V also has negatively curved spatial sections and can therefore, in principle, also produce localised features \\cite{sung1}.   However, as well as forming distinctive features in the temperature pattern, anisotropic cosmological models also generate characteristic signatures in the polarized component of the background radiation. Thomson scattering generates polarization as long as there is a quadrupole anisotropy in the temperature field of the radiation incident upon the scattering particle. In the concordance cosmology the temperature and polarization patterns are (correlated) stochastic fields arising from their common source in scalar and tensor perturbations arising from inflation. In a Bianchi cosmology, however, the patterns are coherent and have a deterministic relationship to one another owing to their common geometric origin. It has recently been shown \\cite{sung1,Pontzen1,pontzen} that only in special cases are the properties of the polarization field produced in Bianchi VII$_h$ consistent with the latest available WMAP polarization data \\cite{WMAPPol} because such models generally involve a large odd-parity (B-mode) \\cite{Kam,Hu1} contribution that exceeds the experimental upper limit. For a discussion of CMB polarization in the context of the so-called Axis of Evil, see ref. \\cite{FE10}.  But their ability to produce localized features is not the only reason to be interested in the temperature and polarization patterns produced in Bianchi models. For example, history provides a host of connections between these models and the understanding of the interaction between electromagnetic radiation and  gravitational waves \\cite{rees,nanos,NS80,BP80,Pol85,fpc94}.  Our aim in this paper is to give a relatively gentle introduction to the Bianchi models and show how to calculate the temperature and polarization properties of the radiation field expected to arise within them, using methods outlined in a previous short paper \\cite{sung1}. Our aim in doing this is neither to provide an exhaustive set of alternatives to the standard cosmology nor to perform a detailed statistical analysis of the patterns we calculate. Instead we plan to elucidate some of the general properties of the radiation field in anisotropic cosmological models belonging to the different Bianchi types.  The outline is as follows. In the next section we describe the Bianchi classification in general terms for the benefit of the non-expert. In Section 3 we introduce the specific formalism of Bianchi cosmological models, i.e. exact solutions of the Einstein equations involving space-times with symmetries described by the various Bianchi types. In Section 4 we explain how we compute the radiation patterns in these models and then, in Section 5,  show examples of the results obtained. We summarize and present our conclusions in Section 6. ",
        "conclusions": "The goal of this work was to compute the temperature and polarization patterns produced in anisotropic relativistic cosmologies described by various Bianchi types. We focussed on those types that contain the standard homogeneous and isotropic FRW background as a limiting case. We constructed an appropriate description of the radiation field in terms of spin-0 and spin-2 components representing the unpolarized and polarized parts, respectively. We integrated the equations for the radiation field numerically, and presented illustrative examples here.  The basic point that emerges from this study is related to the physical origin of CMB polarization: radiation affected by Thomson scattering from an electron in a radiation field possessing a quadrupole anisotropy will inevitably be partially linearly polarized. In the context of standard cosmological models, the environment of different electrons varies owing to the presence of density inhomogeneities and a background of gravitational waves. These sources of variation are stochastic so the variations in the polarized component of the radiation field, though correlated with the temperature variations, are essentially incoherent. In Bianchi cosmologies, however, global homogeneity requires that each electron sees {\\em the same} quadrupole. The polarized part of the radiation field is therefore coherent, and is in a fixed relationship to the temperature variation (once the model is specified).  We have exploited the formalism used to generate the examples shown here to produce an extensive test-bed of CMB temperature and polarization maps derived from Bianchi universe models of the types discussed in this paper. In future work we will analyse these maps using a variety of statistical measures of anisotropy. It is important to understand how standard statistical techniques, designed to be applied to stationary stochastic fluctuations, perform when applied to patterns which are neither stochastic nor stationary. Our Bianchi models, for example, yield patterns which have highly correlated spherical harmonic components in contrast to the case of a homogeneous Gaussian random field in which the harmonic modes are independent. It is important to understand how CMB fluctuations arising from Bianchi-type universes impact on orthodox analysis procedures and whether they produce characteristic signatures when analysed in this way, particularly in the presence of additional stochastic fluctuations. Higher-order statistics will be necessary to provide a full characterization of the coherent anisotropic fluctuations we have studied here. Moreover, present and future CMB experiments are largely aimed at improving the precision of polarization measurements in order to find evidence of a stochastic background of primordial gravitational waves. Techniques advocated to measure polarization pattern over the celestial sphere are therefore also generally tuned to detect signals of an incoherent nature. The presence of coherent signals in the CMB could be an indication of physics beyond the standard cosmological framework. This will hopefully lead us to better ways of searching for departures from the concordance cosmology using the next generation of datasets."
    },
    "1004/1004.3059_arXiv.txt": {
        "abstract": "We present new optical broadband colors, obtained with the Keck 1 and Vatican  Advanced Technology telescopes, for  six objects in the inner classical Kuiper Belt. Objects in the inner classical Kuiper Belt are of interest as they may represent the surviving members of the primordial Kuiper Belt that formed interior to the current position of the 3:2 resonance with Neptune, the current position of the plutinos, or, alternatively, they may be objects formed at a different heliocentric distance that were then moved to their present locations. The six new colors, combined with four previously published, show that the 10 inner belt objects with known colors form a  neutral clump and  a reddish clump in B$-$R color. Nonparametric statistical tests show no significant difference between the B$-$R color distribution of the inner disk objects compared to the color distributions of Centaurs, plutinos, or scattered disk objects. However, the B$-$R color distribution of the inner classical Kuiper belt objects {\\it does} differ significantly from the distribution of colors in the cold (low inclination) main classical Kuiper belt. The cold main classical objects are predominately red, while the inner classical belt objects are a mixture of neutral and red. The color difference may reveal the existence of a gradient in the composition and /or surface processing history in the primordial Kuiper Belt, or indicate that the inner disk objects are not dynamically analogous to the cold main classical belt objects. ",
        "introduction": "Ongoing discovery and orbital and physical characterization of Kuiper belt objects (KBOs) and related outer solar system bodies promise new insight into the formation and evolution of our solar system. However, the evolving dynamical picture shows that the orbits of many, if not most, objects have been substantially changed over time. Thus it may be difficult to study any  differences between objects that formed at different heliocentric distances, as the current heliocentric distances of individual objects may differ radically from the heliocentric distances at the time of their formation.  \\citet{Gomes03} hypothesized a dynamical mechanism to move the dynamically hot (high values of orbital inclination (\\iital) and eccentricity (\\eital)) KBOs from a massive protoplanetary disk truncated at 30$-$35 AU from the Sun outward to the present Kuiper Belt region. This mechanism would not work for the cold KBO population (low values of orbital inclination (\\iital) and eccentricity (\\eital)) \\citep{SSBNMorbiAJ}. In the Gomes model, the cold KBO population would either have formed {\\it in situ} or would have been moved by a different mechanism. In the context of the Nice model of the outer solar system, a mechanism has been proposed that could move even cold KBO objects from a disk truncated at $\\sim$ 35 AU to their present distances \\citep{LevisonNice08}. In this model, the dynamically cold KBO objects would have formed in the outer regions of the massive protoplanetary disk, perhaps in the region around 30 AU.  The vast majority of the known dynamically cold classical belt objects have semimajor axes (\\aital) between 42 and 48 AU. There are also low \\eital, low \\iital classical objects with \\aital $<$ 39.4 AU. This semimajor axis is less than the current position of the Neptune 3:2 resonance, the semimajor axis at which Pluto and the plutino KBOs are currently found. The population of these objects, known as inner classical belt KBOs (ICKBOs), with \\aital between about 36 and 39.4 AU, is much lower than that of the main classical belt objects ( \\aital between 42 and 48 AU). Because of the relative dearth of known classical belt objects interior to the Neptune 3:2 resonance, no published physical studies have specifically targeted these objects. This class of objects extends the \\aital range of objects  compared to the \\aital range spanned by the main belt alone. Thus it is potentially useful in the study of any gradients of chemical or surface processing history with formation distance for minor bodies in the outer solar system. Here we present new optical colors of a sample of inner classical KBOs. We compare the color distribution of these objects with that of other classes of outer solar system objects. ",
        "conclusions": "The B$-$R color distribution of a sample of 10 inner disk Kuiper Belt objects is shown to be not inconsistent with the color distribution of samples of plutinos, Centaurs and scattered disk objects. The current inner disk sample has both red and neutral colored objects, reminiscent of the Centaurs, which have a bimodal B$-$R color distribution. However, we cannot claim that the ICKBOs have a statistically significant bimodal signal with the present data.  The average B$-$R color of the ICKBOs are inconsistent, at the 98\\% level or higher, depending on specific sample and statistical test used, with the colors of a sample of cold (low inclination) classical KBOs with semimajor axes between 42 and 48 AU. Possible conclusions of this are:  1) The inner disk objects we observed, even though they are of low \\eital and \\iital, are members of a population analogous to the {\\it hot} classical KBOs. The hot classical KBOs do not show a predominately red optical color distribution, as do the cold classical objects.  2) If the inner disk objects are in fact members of a cold inner disk population which is a sunward continuation of the cold classical objects between 42 and 48 AU, then there is a radial color gradient in the colors of this cold disk population of KBOs."
    },
    "1004/1004.3329_arXiv.txt": {
        "abstract": "SN 2007if was the third over-luminous SN Ia detected after 2003fg and 2006gz. We present the photometric and spectroscopic observations of the supernova and its host by ROTSE-III, HET and Keck. From the H$_\\alpha$ line identified in the host spectra, we determine a redshift of 0.0736. At this distance, the supernova reached an absolute magnitude of -20.4, brighter than any other SNe Ia ever observed. If the source of luminosity is radioactive decay, a large amount of radioactive nickel ($\\sim$1.5~\\msun) is required to power the peak luminosity, more than can be produced realistically in a Chandransekar mass progenitor. Low expansion velocity, similar to that of 2003fg, is also measured around the maximum light. The observations may suggest that SN 2007if was from a massive white dwarf progenitor, plausibly exploding with mass well-beyond 1.4~\\msun. Alternatively, we investigate circumstellar interaction that may contribute to the excess luminosity. ",
        "introduction": "Type Ia supernovae (SNe Ia) are commonly understood to be thermonuclear explosions of carbon/oxygen white dwarfs (CO WDs). In the most likely scenario, a single WD approaches its Chandrasekhar mass limit by accreting from a non-degenerate binary companion (single-degenerate (SD) model, \\citealp{1973ApJ...186.1007W, 1982ApJ...253..798N}). The other probable channel is two WDs that merge due to gravitational radiation (double-degenerate (DD) model, \\citealp{1984ApJS...54..335I, 1984ApJ...277..355W}). In either case, the explosion results in the total obliteration of the entire star (although see \\citealp{1985A&A...150L..21S,1991ApJ...367L..19N} for an alternative outcome of accretion-induced collapse to a neutron star). A range of peak luminosities have been observed for SNe Ia, which is believed to be directly related to the amount of radioactive $^{56}$Ni synthesized \\citep{1982ApJ...253..785A}. A number of factors can potentially affect the scatter, e.g. the progenitor age \\citep[e.g. ][]{1996AJ....112.2391H,2001ApJ...554L.193H,2006ApJ...648..868S,2008ApJ...685..752G} and metallicity \\citep[e.g. ][]{2000AJ....120.1479H,2003ApJ...590L..83T,2009ApJ...691..661H}. Geometric effects, such as an asymmetric ignition \\citep{ssrh07, 2008ARA&A..46..433W, 2010ApJ...710..444H}, may at least account for part of the dispersion.  Although a large fraction of SN Ia seem to occupy a relatively small parameter space that allows parameterizing of their peak luminosities based on one or two light curve and spectral characteristics \\citep{npbbh95,plssh99,bcmta05,bdb09}, the number of exceptional events is rising with the increasing number of well-observed SNe. Among these, two recent cases, SN 2003fg (also known as SNLS-03D3bb, \\citealp{hsnec06}) and SN 2006gz \\citep{hgpbd07} brought attention to the possibility of an explosion that exceeded the Chandrasekhar mass. Although an asymmetric explosion model offers an alternative explanation \\citep{hsr07}, SN 2003fg was so luminous that it is difficult to explain by geometric effects alone. Such extreme cases might be intrinsically rare, but especially valuable for constraining the nature of the progenitors, and helping to understand how to reliably use SNe Ia as tools to measure the cosmic expansion history for a larger distance range.  We present in this paper a more recent addition to this category, SN 2007if. The supernova was discovered independently by the 0.45m ROTSE-IIIb telescope at McDonald Observatory, Texas, on Aug 19.28 UT and by the Nearby Supernova Factory (SNfactory) on Aug 25.4 UT \\citep{atel1212,cbet1059}. At redshift z$\\sim$0.074, SN 2007if reached a peak absolute magnitude exceeding -20 and is one of the most luminous Type Ia supernova ever observed. The exceptional brightness and the spectral features might be consistent with a super-Chandrasekhar mass explosion, but we consider other possibilities.  In the following sections, we first summarize the photometric and spectroscopic observations of SN 2007if and its host in $\\S$ 2. We then discuss the possible interpretation of the observed features in $\\S$ 3. This is followed by our conclusions in $\\S$ 4. Throughout the paper, we assume a standard cosmology model with the Hubble parameter H$_{\\rm 0}$=70~km~s$^{-1}$ Mpc$^{-1}$ and the density parameters $\\Omega_{\\rm m}$=0.3, $\\Omega_{\\Lambda}$=0.7. All quoted errors are 1-sigma (68\\% confidence), unless otherwise stated. ",
        "conclusions": "SN 2007if was one of the first discoveries of the ROTSE Supernova Verification Project (RSVP) \\citep{ya08}, which extends the efforts of the Texas Supernova Search (TSS, \\citealp{quimby_thesis}) and uses all four ROTSE-III telescopes.  The peak unfiltered magnitude of SN 2007if (calibrated to SDSS r-band) measured by ROTSE-IIIb, corresponds to an absolute magnitude of -20.4 at redshift 0.074. This is by far the brightest SN Ia observed. If powered by $^{56}$Ni decay, the exceptional brightness of SN 2007if requires a total $^{56}$Ni mass close to or even exceeding the Chandrasekhar mass of a non-rotating white dwarf. The photospheric expansion velocity derived around maximum is comparable to the estimate for SN 2003fg and suggests a relatively low kinetic energy. After two weeks post-maximum, the spectral evolution of SN 2007if become indistinguishable from a normal SN Ia, except for ``muted'' features and weak absorptions due to Si and Ca. Late time observations show that SN 2007if occured in a low luminosity host, similar to that of SN 2003fg.  The observed properties of SN 2007if may suggest that it had a massive progenitor, well above 1.4~\\msun, as proposed for a few other luminous SNe Ia. Super-Chandrasekhar explosions have been investigated for systems with rapid rotations \\citep{unh03,yl05}. Both a single WD or a merger of two degenerate stars can sustain a larger mass if rapidly rotating.  CSM interaction may provide another interpretation for the excess luminosity. In the presence of such emission above the photosphere, shallow spectral features are predicted, consistent with our observation of SN 2007if. However, it is noted that the spectra of other luminous SNe, such as 2003fg and 2006gz are not similarly ``muted.'' Further studies are needed to see whether the observed light curve can be reproduced.  So far, the few most luminous SNe Ia show diverse behaviors. In general, the events follow a positive correlation between peak brightness and stretch factor, except for SN 2003fg \\citep{hsnec06}, which was among the brightest but showed a relatively fast decay. SN 2009dc \\citep{ykkti09} and 2007if had similarly slow photospheric velocity derived from Si~II 6355 as for SN 2003fg. SN 2006gz \\citep{hgpbd07} and 2004gu \\citep{chpfs09} appear as close cousins, with relatively high photospheric velocity. Suppression of the silicon velocity was also noticed at early times for SN 2006gz \\citep{hgpbd07}. Although rare, over-luminous SNe Ia show significant deviations in brightness when the normal parametrization method is applied in cosmological studies \\citep{fpbch10}. More luminous events, polarimetric observations, X-ray monitoring and detailed simulations in the future may help to understand the category as a whole."
    },
    "1004/1004.4215_arXiv.txt": {
        "abstract": "A small fraction of core collapse supernovae (SNe) show evidence that the outgoing blast wave has encountered a substantial mass $\\sim 1-10M_{\\sun}$ of circumstellar matter (CSM) at radii $\\sim 10^{2}-10^{3}$ AU, much more than can nominally be explained by pre-explosion stellar winds.  In extreme cases this interaction may power the most luminous, optically-energetic SNe yet discovered.  Interpretations for the origin of the CSM have thus far centered on explosive eruptions from the star just $\\sim$ years$-$decades prior to the core collapse.  Here we consider an alternative possibility that the inferred CSM is a relic disk left over from stellar birth.  We investigate this hypothesis by calculating the evolution of proto-stellar disks around massive stars following their early embedded phase using a self-similar accretion model.  We identify a brief initial gravitationally-unstable (``gravito-turbulent'') phase, followed by a much longer period of irradiation-supported accretion during which less effective non-gravitational forms of angular momentum transport dominate.  Although external influences, such as the presence of a wide binary companion, may preclude disk survival in many systems, we find that massive ($\\sim 1-10M_{\\sun}$) disks can {\\it preferentially survive around the most massive stars}.  Reasons for this perhaps counter-intuitive result include (1) the shorter stellar lifetimes and (2) large photo-evaporation radii ($\\sim 10^{3}$ AU) of very massive stars; (3) suppression of the magneto-rotational instability due to the shielding from external sources of ionization; and (4) relative invulnerability of massive disks to lower mass stellar collisions and luminous blue variable eruptions.  Furthermore, disks with radii $\\sim 10^{2}-10^{3}$ AU are picked out by the physics of the embedded stage of accretion and the requisite conditions for subsequent disk survival.  The disk mass, radius, and scale-height at core collapse typically result in a $\\sim 10$ per cent efficiency for converting the kinetic energy of the exploding star into radiation, potentially producing a total electromagnetic output $\\sim 10^{50}-10^{51}$ ergs.  We identify two regimes of disk-supernova interaction, which are distinguished by whether the shocked disk radiates its thermal energy before being engulfed by the expanding SN ejecta.  This dichotomy may explain the difference between very luminous supernova which show narrow H line emission and those which show no direct evidence for hydrogen-rich CSM interaction.  Because very luminous SNe are rare, testing the relic disk model requires constraining the presence of long-lived disks around a small fraction of very massive stars. ",
        "introduction": "Sensitive, wide-field transient surveys are revolutionizing our understanding of the breadth of time-dependent astrophysical phenomena.  The discovery of extremely sub-luminous, rapidly-evolving Type I supernovae (SNe) (e.g. \\citealt{Foley+09}; \\citealt{Valenti+09}; \\citealt{Perets+09}; \\citealt{Poznanski+10}) may, for instance, indicate a wider diversity of white dwarf-related explosions than previously anticipated (\\citealt{Bildsten+07}; \\citealt{Metzger+09}; \\citealt{Shen+10}).  In the opposite extreme, a number of core-collapse SNe have recently been discovered that are unusually bright and/or optically-energetic.  These events, collectively known as very luminous SNe (VLSNe), include SN 2006gy (\\citealt{Ofek+07}; \\citealt{Smith+07}), 2006tf (\\citealt{Smith+08}), 2005ap (\\citealt{Quimby+07}), 2003ma (\\citealt{Rest+09}), 2007bi (\\citealt{Gal-Yam+09}), 2008fz (\\citealt{Drake+09}), 2008es (\\citealt{Gezari+09}; \\citealt{Miller+09}) and 2008iy (\\citealt{Miller+10}).  The mysterious event SCP 06F6 (\\citealt{Barbary+09}) and the luminous, high-redshift transients discovered by the Palomar Transient Factory \\citep{Quimby+09} may be related phenomena.  VLSNe radiate a total energy $\\gtrsim 10^{50}-10^{51}$ ergs and can reach peak absolute magnitudes approaching $M_{\\rm V} \\sim -23$ (e.g. 2005ap).  Though they represent only a small fraction of massive stellar deaths, VLSNe are thus particularly conspicuous because they can remain bright for months$-$years and are detectable even in the distant universe (\\citealt{Cooke+09}; \\citealt{Quimby+09}).  The optical emission from most Type I and some Type II SNe is powered by the radioactive decay sequence $^{56}$Ni $\\rightarrow$ $^{56}$Co $\\rightarrow$ $^{56}$Fe, suggesting that VLSNe may result from an anomalously high Ni yield (e.g. at least several solar masses are required to explain SN 2006gy; \\citealt{Kawabata+09}).  Though such a large $^{56}$Ni mass  is difficult to produce in a normal core collapse explosion, it is a natural byproduct of theoretically predicted ``pair-instability'' SNe (PI-SNe; \\citealt{Barkat+67}; \\citealt{Bond+84}).  Although PI-SNe are a plausible origin for some events including SN 2006gy \\citep{Smith+07} and SN 2007bi \\citep{Gal-Yam+09}, they are an unlikely explanation for all VLSNe since PI-SNe are expected to occur only from extremely massive stars with low metallicities \\citep{Heger&Woosley02}, as would form primarily in the early universe.  Another way to power an anomalously bright SN light curve is to increase the explosion's prompt radiative efficiency; indeed, the optical output of most core collapse SNe ($\\sim 10^{49}$ ergs; e.g. \\citealt{Bersten&Hamuy09}) represents only a small fraction of the total $\\gtrsim 10^{51}$ ergs of kinetic energy that is generally available.  Low radiative efficiency is typical because the stellar progenitors of core-collapse SNe are relatively compact: radii for Wolf-Rayet and red giants are $R \\sim 0.01$ and $\\lesssim 10$ AU, respectively.  Much of the thermal energy generated by the shocked stellar envelope is therefore lost to PdV work before the ejecta expands sufficiently to become transparent.  The radiative efficiency can be enhanced, however, if the kinetic energy of the ejecta is thermalized at much larger radii ($\\gtrsim 100$ AU), in a region above or closer to the photosphere (e.g. \\citealt{Falk&Arnett73,Falk&Arnett77}; \\citealt{Chevalier&Fransson94}; \\citealt{Smith&McCray07}).  This can occur via shock interaction with a dense circumstellar medium (CSM).  Indeed, the Type IIn class of SNe (e.g. \\citealt{Schlegel90}; \\citealt{Filippenko97}), which are characterized in part by narrow H emission lines suggestive of slower material, are generally believed to be powered by CSM interaction (e.g. \\citealt{Chugai+04}).  Although many VLSNe are indeed classified as Type IIn, others show little or no evidence for CSM interaction (such as SN 2005ap; \\citealt{Quimby+07}) or hydrogen in their spectra.  Such behavior may nevertheless be consistent with a CSM-powered luminosity if the CSM is especially massive and opaque, since this would delay the time required for the shocked thermal energy to escape \\citep{Smith&McCray07} and for the SN to begin displaying more typical IIn properties (e.g. \\citealt{Smith+08}).  We describe a specific example of this effect in $\\S\\ref{sec:tworegimes}$.  In order to explain luminous IIn SNe ranging from events like 2008iy \\citep{Miller+10} to extreme cases such as SN 2006gy via CSM interaction requires the presence of $\\sim 0.1-10M_{\\sun}$ on radial scales $\\sim 100-1000$ AU.  If one assumes that the narrow H line widths ($\\sim 10^{2}-10^{3}$ km s$^{-1}$) commonly observed in SN IIn represent the unshocked outflow speed of the {\\it same} CSM responsible for powering the SN luminosity, the inferred mass-loss rates range from $\\sim 10^{-3}M_{\\sun}$ yr$^{-1}$ up to $\\sim M_{\\sun}$ yr$^{-1}$, much higher than can be supplied by standard, line-driven winds (e.g. \\citealt{Gayley+95}; \\citealt{Owocki+04}; \\citealt{Smith&Owocki06}).  When combined with evidence for minima or ``gaps'' in the mass distribution between the star and the CSM in some cases (e.g. \\citealt{Chugai&Danziger94}; \\citealt{Miller+10}), these inferences have motivated the interpretation that the CSM is composed of shells of material ejected just years$-$decades before the SN explosion (e.g. \\citealt{Chugai+04}; \\citealt{Smith+07}; \\citealt{Dessart+09}).  Proposed explanations for these pre-SN outbursts thus far include luminous blue variable (LBV)-like eruptions analogous to the 19th century eruption of $\\eta$-Carinae (e.g.~\\citealt{Smith+07}) and pulsational pair instabilities (i.e. sub-energetic, non-terminal PI-SN analogs; \\citealt{Woosley+07}).  However, neither of these explanations is altogether satisfactory.  If LBV eruptions do occur just prior to stellar death, this suggests that very massive stars can die in their LBV phase, nominally thought to occur prior to core He burning (e.g. \\citealt{Langer+94}), and with their hydrogen envelopes intact (e.g. \\citealt{Smith08a}).  Pair-instabilities, on the other hand, occur only in stars with masses $\\gtrsim 95 M_{\\sun}$ that retain their He envelope following H burning; this again likely requires low metallicity, an unusual circumstance in the present-day universe.  It has been argued that both the hydrogen-rich deaths of very massive stars and modern-day PI-SNe become more likely if empirically-calibrated stellar mass-loss rates have been systematically overestimated (\\citealt{Gal-Yam+07}; \\citealt{Smith08b}) due to the insensitivity of $n^{2}$ emission diagnostics to wind clumping (e.g.~\\citealt{Fullerton+06}).  This does not, however, mitigate the fact that both LBV and pair instability models also appear to require fine-tuning: the delay between shell ejection and core collapse must be well-timed to produce a collision at the optimal radii ($\\sim 10^{2}-10^{3}$ AU) for producing bright emission.  Although this ``coincidence'' could in part result from a selection effect, the fact remains that the pre-SN eruption and core collapse {\\it must be preferentially synchronized to within a timescale $\\sim$ years$-$decades (which is orders of magnitude shorter than stellar transport or evolutionary timescales) or VLSNe would be even more rare than is observed}.  Although such a correlation could in principle arise from the time interval between pair-instabilities, this depends sensitively on uncertain details such as how long the core takes to contract following the first pulse (\\citealt{Heger&Woosley02}; \\citealt{Woosley+07}).  In this paper we consider an alternative origin for CSM interaction in luminous core collapse SNe: a gaseous proto-stellar disk left over from stellar birth.  The lifetimes of the most massive stars are only $\\sim 3$ Myr (e.g. \\citealt{Bond+84}; \\citealt{Maeder&Meynet87}), comparable to the observed lifetimes of disks around low-mass proto-stars (e.g. \\citealt{Strom95}).  Observations of moderately massive ($\\sim 10 M_{\\sun}$) proto-stars typically suggest much shorter disk lifetimes (e.g. \\citealt{Natta+00}; \\citealt{Fuente+06}; see, however, \\citealt{Manoj+07}), indicating that the ``disk dispersal'' processes at work following star formation may be more effective for higher mass stars (e.g. \\citealt{Hollenbach+00}).  VLSNe are, however, very rare and likely originate from a small subset of very massive stars.  The PTF events discovered by \\citet{Quimby+09} represent only a tiny fraction $\\sim 10^{-4}$ of core collapse SNe.  Although Type IIn SNe as a whole account for a few to $\\sim 10$ per cent of core collapse events (\\citealt{Cappellaro+99}), they are a highly inhomogeneous population and the majority may represent a different phenomena than the most luminous events.  If a fraction $f \\sim 10^{-4}-10^{-2}$ of massive stars must retain a massive disk in order to be consistent with the rate of VLSNe, then only a few to several hundred relic disk systems should be present in Milky Way-type galaxies at any time.  Testing this prediction is nontrivial because although a relic disk would likely be conspicuous, massive stars are necessarily very young and, hence, typically far away and often obscured; observations and their interpretation are thus particularly challenging due to limited sensitivity, angular, and spectral resolution (e.g. \\citealt{Cesaroni+07}).  Although there is no evidence for long-lived massive disks around unobscured ZAMS O stars, relic disk systems may not appear as normal or unobscured, and the local census of massive obscured stars and their environments is incomplete (e.g.~\\citealt{Wachter+10}).  Indeed, the rate of VLSNe is sufficiently low that if only a very small fraction of high mass stars retain a disk, it would have important observable implications.  The present theoretical study can broadly be divided into two parts: a detailed analysis of (1) the conditions under which a massive proto-stellar disk can survive until the core collapse of its host star; (2) the interaction of the relic disk (with properties determined self-consistently) with the outgoing supernova ejecta and the resulting emission.  In detail, the following sections are organized as follows.  In $\\S\\ref{sec:stages}$ we discuss the stages in the formation and evolution of massive proto-stellar accretion disks.  Relying heavily on the results from $\\S\\ref{sec:stages}$, in $\\S\\ref{sec:models}$ we construct time-dependent proto-stellar disk models which focus on the isolated stages of evolution following the embedded accretion phase, neglecting for the moment all sources of disk mass loss except accretion.  In $\\S\\ref{sec:destruct}$ we examine the susceptibility of the disk to destructive processes such as outflows and stellar collisions in order to determine whether and under what conditions disk survival is most probable.  In $\\S\\ref{sec:diskshock}$ we explore the interaction of the disk with the supernova ejecta following core collapse and its implications for the origin of luminous CSM-powered SNe.  In $\\S\\ref{sec:discussion}$ we discuss the implications of our results for long-lived disks around massive stars ($\\S\\ref{sec:disks}$), Type IIn (and {\\it non-}IIn) VLSNe ($\\S\\ref{sec:type2n}$) and ``hybrid'' Type I/IIn SNe ($\\S\\ref{sec:hybrid}$).  We conclude in $\\S\\ref{sec:conclusions}$. ",
        "conclusions": "\\label{sec:conclusions}  Circumstellar interaction is a promising mechanism for powering some of the most luminous and energetic SNe yet discovered.  However, if the requisite CSM is the result of a pre-SN stellar eruption, up to several solar masses of material must be ejected just years$-$decades prior to explosion.  Although in principle giant LBV-like eruptions or pair instabilities are capable of producing such prodigious mass loss, there is currently no explanation for why the eruption and the stellar explosion are so nearly synchronized.  In this paper we have developed an alternative hypothesis that the bulk of the CSM is not an outflow (thus avoiding the ``coincidence'' problem), but rather a pre-existing, long-lived disk left over from stellar birth.  One virtue of our model is that a CSM mass up to $\\sim 10 M_{\\sun}$ is explained as the maximum disk mass that can be supported against self-gravity by stellar irradiation (eq.~[\\ref{eq:diskmassratio_irr}]).  The CSM radii $\\sim 10^{2}-10^{3}$ AU required by observations are also explained naturally by (1) the maximum outer disk radius ($\\sim 100-200$ AU) that is stable to fragmentation during the embedded phase (\\citealt{Kratter&Matzner06}; $\\S\\ref{sec:embed}$); and (2) the limited subsequent viscous evolution experienced by disks around massive stars ($\\S\\ref{sec:irrad}$; $\\S\\ref{sec:models}$).  Indeed, we find that massive disks may preferentially survive until core collapse around the most massive stars ($\\S\\ref{sec:summary}$) for the following combination of reasons (see also Fig.~\\ref{fig:summary}): \\begin{itemize} \\item{Massive stars live short lives, thereby allowing less time for accretion and other disk dispersal processes to act.} \\item{Massive stars begin their irradiation-supported phase of evolution with more massive disks (eq.~[\\ref{eq:diskmassratio_irr}]), which generally take longer to disperse.} \\item{Disks around massive stars have larger surface densities at the onset of irradiation-supported accretion (eq.~[\\ref{eq:sigma0}]).  This shields the midplane from external sources of ionization, resulting in extensive and sustained dead-zones in which the magneto-rotational instability is quenched and angular momentum transport is inefficient.  The powerful stellar winds from massive stars appear likely to exclude interstellar cosmic rays, leaving X-rays (which penetrate to a shallower depth) as the chief source of surface ionization (eq.~[\\ref{eq:rloverr}]).} \\item{Due to their slower expected viscous spreading and deeper gravitational potential, disks around massive stars are less likely to be dispersed via photo-evaporation ($\\S\\ref{sec:photo}$). } \\item{Massive disks can be denser than the ejecta from giant LBV eruptions, suggesting they are more likely to remain bound to the star (\\S\\ref{sec:LBV}).} \\end{itemize} Our conclusion that the probability of disk survival rapidly increases with stellar mass is consistent with evidence that very luminous IIn SNe originate from the most massive stars (\\citealt{Gal-Yam+07}; \\citealt{Smith+07}).  We do not, however, suggest that the disk survival fraction approaches unity in any stellar mass range, as external influences, such as the presence of a massive binary companion or a particularly dense stellar environment, may preclude disk survival in the majority of cases.  Several of the observed properties of VLSNe can be explained by the collision between the supernova ejecta and a relic disk (see Figs.~\\ref{fig:predisk} and \\ref{fig:diskfig}).  The nominal radiative efficiency of the disk-ejecta interaction is $\\sim 10$ per cent (eq.~[\\ref{eq:erad}]), resulting in a total electromagnetic output $\\sim 10^{50}-10^{51}$ ergs for a SN with kinetic energy $\\sim 10^{51}-10^{52}$ ergs, substantially more optically-energetic than a normal Type II SN.  For lower mass or radially-extended disks, the light curve directly traces the shock-generated thermal energy ($\\S\\ref{sec:fastcool}$; Fig.~\\ref{fig:lc1}), plausibly resulting in copious ``intermediate'' width line emission, as observed in SN IIn.  For more massive, compact disks the shock-generated thermal energy must instead diffuse out of the entire stellar envelope ($\\S\\ref{sec:slowcool}$; Fig.~\\ref{fig:lc2}), which may explain the most luminous events (e.g. 2005ap), which show little or no direct evidence for CSM interaction.  Although there is no observational evidence for long-lived disks around massive stars, distinguishing between a truly young stellar object proto-star and an older system with a relic disk may be nontrivial.  Ultimately testing whether bona fide disks exist around evolved stars in sufficient numbers to explain VLSNe will require acquiring unbiased samples of embedded evolved stars and hyper-compact HII regions, distinguishing evolved from youthful systems, and, eventually, pinning down definitive disk or outflow signatures with e.g. atomic \\citep{Ercolano&Owen10} and molecular line diagnostics (e.g.~\\citealt{Krumholz+07}; \\citealt{Carr&Najita08}).  Though challenging, such a feat will be aided by the improved capabilities of present and upcoming observatories such as the EVLA, ALMA, {\\it Herschel}, and the {\\it James Webb Space Telescope}."
    },
    "1004/1004.0600_arXiv.txt": {
        "abstract": "{  We compute the one-point probability distribution function of small-angle cosmic microwave background temperature fluctuations due to curved cosmic (super-)strings with a simple model of string network by performing Monte Carlo simulations. Taking into account of the correlation between the curvature and the velocity of string segments, there appear non-Gaussian features, specifically non-Gaussian tails and a skewness, in the one-point pdf. The obtained sample skewness for the conventional field-theoretic cosmic strings is $g_1\\approx -0.14$, which is consistent with the result reported by Fraisse et al.~\\cite{Fraisse:2007nu}. We also discuss the dependence of the pdf on the intercommuting probability. We find that the standard deviation of the Gaussian part increases and non-Gaussian features are suppressed as the intercommuting probability decreases. For sufficiently small intercommuting probability, the skewness is given by $\\lesssim \\text{(a\\ few)}\\times 10^{-2}$.  } ",
        "introduction": "Cosmic strings are line-like topological defects formed in the early universe through a spontaneous symmetry breakdown. It was claimed that the formation of cosmic strings at the end of inflation is a generic feature of supersymmetric grand unified theories~\\cite{Jeannerot:2003qv}. The string tension $\\mu$ is directly related to the symmetry breaking energy scale. Also recent studies of stringy cosmology have revived interest in cosmic strings, because it was pointed out that brane inflation models may produce another class of string objects, called cosmic superstrings~\\cite{Sarangi:2002yt,Davis:2005dd,Copeland:2009ga,Majumdar:2005qc}. Cosmic superstrings may have different properties from those of conventional field-theoretic cosmic strings. One of the observationally interesting differences is the intercommuting probability $P$. It can be significantly smaller than unity for cosmic superstrings~\\cite{Jackson:2004zg}, while normally $P=1$ for field-theoretic strings~\\cite{Eto:2006db} (but see \\cite{Salmi:2007ah}).  The imprint of cosmic strings on the cosmic microwave background (CMB) has been widely studied. Although cosmic strings were excluded as a dominant source of the observed large angular scale anisotropy ~\\cite{Pogosian:2008am,Bevis:2007gh,Pogosian:2003mz} (see also \\cite{Perivolaropoulos:2005wa} for a review), a signal due to cosmic strings could still be observed at small angular scales~\\cite{Fraisse:2007nu,Pogosian:2008am} with future arcminutes experiments such as South Pole Telescope~\\cite{Ruhl:2004kv} or Atacama Cosmology Telescope~\\cite{Kosowsky:2004sw}. On small angular scales, the primary fluctuations are damped and only the integrated Sachs-Wolfe (ISW) effect is relevant. The Gott-Kaiser-Stebbins (GKS) effect~\\cite{Kaiser:1984iv,Gott:1984ef} is the most characteristic signal of the ISW effect due to cosmic strings. The GKS effect is caused by a discontinuity of the gravitational potential across a moving string segment when photons passes by the moving string. If photons are scattered by a number of moving string segments, the observed temperature fluctuations appear as a superposition of the discontinuities.  Since the network of strings is a highly nonlinear object, its non-Gaussian features may help us distinguish cosmic string signals from other secondary effects and hence may enhance the observability. A skewness, which is a measure of the asymmetry of a probability distribution function, is one of the simplest features of non-Gaussianity~\\cite{Komatsu:2001rj}. The skewness of temperature fluctuations is defined by \\begin{eqnarray} g_1=\\frac{\\overline{(\\Delta -\\bar{\\Delta})^3}}{\\sigma_{\\Delta}^{3}}\\,, \\label{eq:skew def} \\end{eqnarray} where $\\Delta \\equiv (T-\\bar{T})/T=\\Delta T/T$, the bar `` $\\bar{~}$ \" denotes the statistical average over a CMB map and $\\sigma_{\\Delta}$ the standard deviation. Recently, Fraisse et al.~\\cite{Fraisse:2007nu} found that the one-point probability distribution function (pdf) of the temperature fluctuations due to conventional cosmic strings has non-Gaussian tails and a negative skewness, $g_1\\approx -0.23$.  Non-Gaussian features would also appear in the bispectrum~\\cite{Gangui:2001fr}. Some authors~\\cite{Hindmarsh:2009qk,Hindmarsh:2009es,Regan:2009hv} calculated the bispectrum and the trispectrum by using the formula derived by Hindmarsh~\\cite{Hindmarsh:1993pu}. In particular, Hindmarsh et al.~\\cite{Hindmarsh:2009qk} pointed out that the numerically measured amplitude corresponds to a $|f_{\\rm NL}|\\approx 10^3$ for the local-type of primordial non-Gaussianity.  In \\cite{Takahashi:2008ui}, we computed analytically the one-point pdf of small-scale temperature fluctuations with a simple model of long straight segments and kinks. We derived basic equations for the evolution of string segments in cosmological backgrounds by extending the velocity-dependent one-scale model with the intercommuting probability $P\\neq1$~\\cite{Takahashi:2008ui,Martins:2000cs,Martins:1996jp,Avgoustidis:2005nv}. It was found that the obtained one-point pdf consists of a Gaussian component due to frequent scatterings by long straight segments and non-Gaussian tails due to close encounters with kinks. The dispersion of the Gaussian component is consistent with the result in \\cite{Fraisse:2007nu} and the non-Gaussian tails also can be fitted to that in \\cite{Fraisse:2007nu} by using two phenomenological parameters. The effect of the intercommuting probability $P$ was also studied there. It was shown that the non-Gaussian tails diminish as $P$ decreases. However the obtained one-point pdf is symmetric for positive and negative temperature fluctuations, hence did not reproduce the non-zero skewness reported in \\cite{Fraisse:2007nu}.  In this paper, we study in depth the non-Gaussian features, especially skewness, of the CMB temperature fluctuations due to cosmic (super-)strings, and interpret numerical results obtained in \\cite{Fraisse:2007nu}. For this purpose, we adopt a simple model of string network and perform Monte Carlo simulations. We consider the correlation between the curvature and the velocity of string segments as the origin of skewness and study the dependence of the skewness on the intercommuting probability $P$.  This paper is organized as follows. In section \\ref{sec:string dynamics}, we give a brief review of the derivation of basic equations for the evolution of string segments incorporating the intercommuting probability $P$~\\cite{Takahashi:2008ui}. In section \\ref{sec:HSV formula and correlations} we introduce general formulae for the temperature fluctuations due to cosmic strings~\\cite{Hindmarsh:1993pu,Stebbins:1994ng,Stebbins:1987va}. We also discuss the correlation between the curvature and velocity of string segments which gives rise to the skewness in the one point pdf. Then, in section \\ref{sec:numerical study}, we give the result of our Monte Carlo simulations. We show that the obtained one-point pdf has a negative skewness when the correlations are taken into account. The skewness is plotted as a function of $P$. Finally, we summarize our results in section \\ref{sec:summary}. ",
        "conclusions": "\\label{sec:summary}  In this paper, we have computed the one-point pdf of small-angle CMB temperature fluctuations due to the ``curved'' cosmic (super-)string segments by performing Monte Carlo simulations. Our purpose was to examine the numerical result obtained in \\cite{Fraisse:2007nu} with a simple model of the string network under a few reasonable assumptions, i.e., the velocity-dependent one-scale (VOS) model. We have found that the presence of the nonzero momentum parameter $k$ corresponds to the correlation between the curvature and the velocity of segments, and therefore the Hubble expansion is essential for a nontrivial correlation (see Eq.~(\\ref{eq:k def}) and Eq.~(\\ref{eq:curvature-velocity correlation})). We have also found another type of correlation between the curvature and the velocity induced by the effective curvature along the intersection of a string worldsheet and the observer's past light-cone (``light-cone effect''). The light-cone effect always exists independent from the properties of string configuration and the cosmic volume expansion.  Combining these two effects, we have calculated the one-point pdf numerically and showed that it reproduces features of the numerical simulations in \\cite{Fraisse:2007nu} very well. In particular, we have found that the one-point pdf has a negative skewness. It was also found that as $P$ decreases the standard deviation of the Gaussian part increases and the non-Gaussian tails and the skewness are suppressed. Our result suggests that, for the sufficiently small $P$, the skewness approaches $\\lesssim \\text{(a\\ few)}\\times 10^{-2}$.  Note that the observed temperature fluctuations are given by contributions not only from cosmic (super-)strings but also from primary and other secondary anisotropies. It is known that conventional cosmic strings can contribute at most $10\\%$ of the total power at $\\ell =10$ in the CMB fluctuations~\\cite{Bevis:2007gh}. Thus, in real world, contributions from primary anisotropies make the skewness significantly small on large angular scales. In order to suppress this effect of the large-scale anisotropies, it would be effective to choose an observation area of a sufficiently small angular size ($\\sim O(10){\\rm arcmin}^2$) where the primary CMB fluctuations have been completely damped out. This would remove the contribution of the primary fluctuations to the denominator of the skewness formula (see Eq.~\\eqref{eq:skew def}), hence enhance the observability of the skewness.  It is known that ${\\cal O}(0.1)$ skewness would indicate nonlinear parameter $|f_{\\rm NL}|\\approx {\\cal O}(10^3)$~\\cite{Cayon:2002jq,Hindmarsh:2009qk}. In this paper, we have shown that, for smaller intercommuting probability $P$, non-Gaussian features become more difficult to find. To discuss the relation between the skewness and the non-Gaussian parameter $f_{\\rm NL}$ in the bispectrum, it is important to compute the angular power spectrum in our model. The work along this direction is in progress~\\cite{Takahashi:future}. Since the non-Gaussian features significantly depend on the intercommuting probability, they could help us to distinguish field theoretic cosmic strings and cosmic superstrings."
    },
    "1004/1004.2928_arXiv.txt": {
        "abstract": "{% As it is long known, the presence of a convective region creates a discontinuity in the chemical profile of a star, which in turn translates into a sharp variation of the adiabatic sound speed. This variation produces an oscillatory behavior of the pulsation frequencies, to which low degree p-modes are sensitive. We investigate the possibility of detecting the signature related to the presence of a convective core in the frequency spectrum of low-mass stars by means of suitable frequency combinations (such as separations and ratios).} ",
        "introduction": "Asterosesimology has shown to be the most powerfull tool to study the inner properties of stars through the observation of stellar oscillations. Many efforts have been conducted to understand the interior physics of stars by understanding the influence that different physical phenomena have on the frequency spectra. In particular, it is known that sharp variations on the adiabatic sound speed produce an oscillatory signal in the p-mode spectra whose period is related to the position of such a discontinuity (e.g. \\cite{sv88}, \\cite{ap94}). Bearing this in mind, in this study we apply different asteroseismic tools to a computational model of a star containing a convective core and try to extract its position by means of the expected signatures on the frequency spectra. ",
        "conclusions": "The diagnostic potential of the long known periodic component of frequency perturbations due to a discontinuity or steep gradient in the adiabatic sound speed has been used to estimate the position within the acoustic cavity of such sharp feature. In particular, we have estimated the location of the convective core using a sine wave fit to the large oscillation pattern observed in two different frequency combinations, showing that the accuracy in the position determination is higher when the ratios are considered than using the small separations. We intend to continue this study and extend it to other mass ranges and different types of discontinuities, such as the ones produced by overshooting and semiconvection (\\cite{vs10}). We are aware that the potential of this diagnosis depends on the ranges of radial orders considered for the sine wave fitting of the data, for which a more elaborate analysis of the turning points of the dipole modes will be applied (e.g \\cite{mt06}). Finally, we plan to complement this study by also applying other asteroseismic tools developed to detect small convective cores to our models (e.g \\cite{cm07}, \\cite{am06}), and explore the possibilities of detecting these mixed zones with the current data being obtained by the CoRoT and Kepler missions."
    },
    "1004/1004.0093_arXiv.txt": {
        "abstract": "{We report on the discovery of the X-ray luminous cluster XMMU\\,J100750.5+125818 at redshift 1.082 based on 19 spectroscopic members, which displays several strong lensing features. SED modeling of the lensed arc features from multicolor imaging with the VLT and the LBT reveals likely redshifts $\\sim$2.7 for the most prominent of the lensed background galaxies. Mass estimates are derived for different radii from the velocity dispersion of the cluster members, $M_{\\rm 200} \\simeq 1.8\\times 10^{14}$\\,\\msun, from the X-ray spectral parameters, $M_{\\rm 500} \\simeq 1.0 \\times 10^{14}$\\,\\msun, and the largest lensing arc, $M_{\\rm SL} \\simeq 2.3 \\times 10^{13}$\\,\\msun. The projected spatial distribution of cluster galaxies appears to be elongated, and the brightest galaxy lies off center with respect to the X-ray emission indicating a not yet relaxed structure. XMMU\\,J100750.5+125818 offers excellent diagnostics of the inner mass distribution of a distant cluster with a combination of strong and weak lensing, optical and X-ray spectroscopy. } ",
        "introduction": "A crucial prerequisite for the use of clusters of galaxies for cosmological studies is the knowledge of the scaling relations  between cluster mass and an observable proxy, such as X-ray luminosity, or X-ray gas temperature. These relations must be calibrated at varying redshifts, using theoretical modeling and hydrosimulations, or simply empirical relations. Clusters at low and intermediate redshifts are bright enough in X-rays to measure the total gravitating mass by deprojecting their temperature and density profiles. This is currently difficult for clusters beyond $z\\sim\\! 1$ and subject to large statistical errors \\citep[e.g.][]{rosati+04}. In this high redshift regime, independent mass estimates based on gravitational lensing (both in the weak and the strong regimes) have long been invoked as the most effective method for calibrating cluster masses. In particular, mass measurements based on multiple strong lensing features are highly robust. Although such measurements are possible only in the central region of the cluster, they are independent of model asssumptions and can help to break the mass-sheet degeneracy of parameter-free weak-lensing mass-reconstructions.  To obtain a well-defined sample of clusters at high redshift, $z > 0.8$, we have started the XDCP  \\citep[XMM-Newton Distant Cluster Project, ][]{boehringer+05,fassbender+08,mullis+05} based on archival XMM-Newton X-ray data. The XDCP has been very successful, so far providing 18 clusters at $z>0.8$ and 10 clusters with $z>1$ including five redshift confirmations from other projects \\citep[for selected results see e.g.][]{fassbender+08,santos+09,rosati+09}.  Here we report the discovery of a distant strong lensing cluster of galaxies at redshift $z=1.082$ which shows a giant arc and other arc-like features interpreted as gravitationally lensed images of more distant, background objects. We study the properties of the lensing cluster and the lensed galaxies on the basis of available optical and X-ray imaging and spectroscopy.  Throughout this paper we use a standard $\\Lambda$CDM  cosmology with parameters $H_0=71$\\,km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\rm  M}=0.27$, and $\\Omega_{\\Lambda}=0.73$, which gives a scale of 8.188\\,kpc arcsec$^{-1}$ at redshift 1.082. All magnitudes in this paper are given in the AB system.  \\begin{figure*}[t] \\begin{minipage}{0.62\\hsize}{} \\resizebox{\\hsize}{!}{\\includegraphics{14430_f1a.ps}} \\end{minipage} \\hfill \\begin{minipage}[]{0.37\\hsize}{} \\resizebox{\\hsize}{!}{\\includegraphics{14430_f1b.ps}} \\caption[]{{\\it} (left) UBVR image of the field around \\xmmu1007\\ with X-ray contours overlaid and tentatively identified lensing features indicated. The cross indicates the X-ray position and the square the size of the color image shown right. {\\it (right)} Color composite of the central region of \\xmmu1007\\ with mean (Rz)V(UB) images in the RGB channels. North is top and East to the left in both frames. \\label{f:opt_composit} } \\end{minipage} \\end{figure*} ",
        "conclusions": "We have presented results of an initial study of the X-ray selected cluster \\xmmu1007. We determine a redshift of $z=1.082$ based on 19 spectroscopic confirmed members. As most prominent property, the cluster shows several strong lensing features. It is an optically rich cluster; the absence of a dominant BCG and the elongated distribution of member galaxies suggest a not yet relaxed structure. The bolometric X-ray luminosity is  \\(L_\\mathrm{X} \\approx 1.3 \\times 10^{44} \\mathrm{ergs/s}\\).  Estimates of the cluster mass were obtained via strong lensing, X-ray spectroscopy and the velocity dispersion of member galaxies; the values obtained at different radii are summarized in Fig.~\\ref{f:nfw}. One obtains overlapping error bars for the mass at $\\sim$$R_{200}$ ($\\sim$virial radius) and at $\\sim$$R_{500}$ ($\\sim$X-ray aperture) due to insufficient data and uncertainties in the scaling relations. If one takes $M_{\\rm vir} \\simeq M_{200} = 1.8\\times10^{14}$\\,\\msun\\ at face value, our strong lensing result seems to be discrepant with an NFW-profile. All data can be made compliant at a mass of about $4\\times10^{14}$\\,\\msun, higher than but not excluded by the velocity dispersion and the X-ray temperature, whereas a weak-lensing mass is key to fix the halo profile at large radii, a deep X-ray observation is necessary to determine the X-ray morphology and the gas temperature. The modelling of the strong lensing system is important to probe the inner density profile.  \\begin{figure}[t] \\resizebox{\\hsize}{!}{\\includegraphics[clip=]{14430_f3.ps}} \\caption[]{Comparison of different mass estimates for \\xmmu1007\\ in units of $10^{14}$\\,\\msun. Labels indicate results based an scaling relations by \\citet{vikhlinin+09,pratt+09,carlberg+97a,evrard+08}. The dashed and solid lines indicate universal density profiles \\citep{nfw97} with concentration parameters of 3.86 and 3.45 for virial masses 1.8 and $4.6\\times 10^{14}$\\,\\msun\\ at an assumed $R_{\\rm vir} = 780$\\,kpc, respectively \\citep{bullock+01}. } \\label{f:nfw} \\end{figure}  Interestingly, in a study of a complete sample of 12 MACS clusters \\citet{zitrin+10} find the observed Einstein radii to be larger and hence the central density of cluster higher than predicted by simulations and interpret their finding as a challenge to cluster formation in a $\\Lambda$CDM model. They invoke the possibility that the formation of clusters started at earlier epochs than currently assumed, leading to higher central Dark Matter concentrations. Also, \\citet{jee+09} suggest this scenario as a possible explanation for the discovery of unexpectedly massive clusters at $z>\\sim$1 in the moderate survey volume probed by the XMM-Newton serendipitous surveys \\citep[like e.g.~XMMU\\,J2235-2557 and  2XMM\\,J083026+524133,][]{mullis+05,lamer+08,rosati+09}. However, in a recent study of the strong lensing clusters in the {\\sc Mare Nostrum Universe}, Meneghetti et al.~(2010, subm.~to A\\&A) find the concentration and the X-ray luminosity to be biased high, and an excess of kinetic energy within the virial radius among the strong lensing clusters. \\xmmu1007\\ is a highly suited  target to confront theory and observation.  The elongated distribution of the cluster member galaxies may point to the principal direction of merging or accretion of \\xmmu1007\\ \\citep{dubinski98}. The merging hypothesis is consistent with the system being a strong lens \\citep[see e.g.~the chain distribution of the brightest member galaxies in the strong-lensing main component of the Bullet Cluster,][]{bradac+06}. On the other hand, the chain distribution of the bright member galaxies of \\xmmu1007\\ may suggest that the infall of these galaxies happens along filaments with small impact parameters, so that dynamical friction is particularly efficient in dragging them to the cluster center and originate a dominant BCG within a short time \\citep{donghia+05}. In analogy with the formation of fossil galaxy groups studied by the latter authors, the small impact parameters of the filaments may lead to an early assembly of the gas in the center, thus to an early start of its cooling. The resulting system may exhibit an enhanced X-ray luminosity with respect to the optical one. This efficient, early formation may offer an alternative explanation to merging for the reason why particularly X-ray luminous clusters with chain distribution of their bright members are detected at high redshifts."
    },
    "1004/1004.1338_arXiv.txt": {
        "abstract": "{} {We present an automatised fitting procedure for the IR range of AGB star spectra. Furthermore we explore the possibilities and boundaries of this method.} {We combine the radiative transfer code DUSTY with the genetic algorithm PIKAIA in order to improve an existing spectral fit significantly. } {In order to test the routine we carried out a performance test by feeding an artificially generated input spectrum into the program. Indeed the routine performed as expected, so, as a more realistic test set-up, we tried to create model fits for ISO spectra of selected AGB stars. Here we were not only able to improve existing fits, but also to show that a slightly altered dust composition may give a better fit for some objects.} {The use of a genetic algorithm in order to automatise the process of fitting stellar spectra seems to be very promising. We were able to improve existing fits and further offer a quantitative method to compare different models with each other. Nevertheless this method still needs to be studied and tested in more detail. } ",
        "introduction": "Asymptotic Giant Branch (AGB) stars \\citep[e.g.][]{agbbook} are one of the main sources of dust in the universe. According to their evolutionary state they produce different amounts and species of dust. At the onset of their mass loss stars with solar metallicity exhibit a mainly oxygen-rich dust mineralogy, which consists of oxides and silicates \\citep[e.g.][]{posch99, jaeger03}. At later stages of their evolution, the star's chemical composition changes depending on its mass, which may lead to a mineralogy dominated by carbon-rich dust such as amorphous carbon, carbides, sulphides and nitrides \\citep[e.g.][]{hony04}.  Dust formation is of immense importance for the overall behavior of these stars. It has an effect on the dynamics of their atmospheres and thus on their mass loss. As a consequence, the spectral appearance changes as well. The dust can produce prominent features in the infrared regime of a spectrum. Many of those dust species might be still unknown, others produce spectral features in the same wavelength region and thus are difficult to distinguish. Another aspect is the dependence of the feature shape on the dust grain shape itself.  Up to know, the fitting of those spectra has mostly been done manually. A routine to automatise this process, which also allows to quantify the results would be a big asset in this field of research. In this work we are taking first steps into this direction. By combining a genetic algorithm with a widely used radiative transfer code we want to show that it is possible to veer away from the traditional way of hand fitting to a more independent procedure. ",
        "conclusions": "AGB stars are known to exhibit strong dust features in the IR range of their spectra. The carriers of those features have to be identified in order to study the dust composition of these shells and hence the mass loss process of AGB stars.  The goal of this work was to study an automatised routine to fit this IR part of the spectrum of an AGB star and to explore the possibilities and more important the boundaries for a routine like this. Therefor we have focused in this work only on the improvement of  fits which have already been carried out on the respective objects. This goal was definitely reached. We could show that it is indeed possible to improve fits done by various researchers using a genetic algorithm for fine tuning of the parameters. We could also show that in cases of Z Cyg, SV Peg and TY Dra a more elaborate or even slightly different dust composition leads to better results than with the composition taken from  \\citet{poschPhD}, which was used as a starting point for our calculation. For CE And, o Cet, S Pav and SV Peg our results showed only minor difference from the dust compositions taken from \\citet{heras05}, which we used as starting parameters for our models for these objects.  Furthermore this routine also provides a quantitative method, namely the fitness value supplied by the genetic algorithm, for comparison of the model spectra. This is connected with the need of independence from an all hands on method done by a scientist, which is, even when carefully trying to avoid it, always biased by the individual experience on the respective field of research. This method can be seen as a little step away form this type of fitting process to a more automatised and independent fitting process.  Nevertheless there are a few things which need to be kept in mind while using PIKAIA and DUSTY in order to fit a spectrum. First, the starting parameters for the fits of all stars mentioned in this paper  were taken from the literature. However, one will not always work on objects that have already been the subject of investigations beforehand. This leaves us with the question  of how to determine the basic stellar parameters for each object needed as input for our automatic fitting method. For a star's temperature the most likely approach is carrying out JHK-photometry in order to retrieve these values. Another logical approach would be using only similar object of a certain group, which then should also offer some hints for the approximate values.  The dust temperature is also a value which might not be found  in the literature. In order to estimate the temperature a simple black body fit \\citep[e.g.][] {kerschbaumPhD} of the respective wavelength region should offer a sufficient starting value for our calculations.  Still, the by far most difficult part in choosing the starting values for a model is the selection of suitable dust species and their respective amount. So far the only reliable method is an experienced scientist. Although this step could be executed by a program doing a comparison of local maxima in the dust spectra with the star's spectrum, an accordance in the peaks unfortunately does not necessarily mean that a single species or a certain combination applies to the star. For now human intervention in this step seems to be unavoidable and a performance efficient  black box fitting procedure out of reach.  This leads also directly to the problem of the run time of our routine. \\citet{metcalfe00} as well as \\citet{puls05} used the parallelized version of PIKAIA to run their simulations. In this paper we used only a serial version of PIKAIA in order to study the overall behaviour of the routine in combination with DUSTY in a setting as simple as possible. Thus a full optimisation procedure took on an average work station up to a day at most.  The most time during the fitting procedure is consumed by  DUSTY  and not by PIKAIA itself. Thus, our attempts to speed up the routine focused on DUSTY. We reduced the wavelength grid as much as possible, skipped unnecessary command line outputs and used the basic black bodies, inbuilt in DUSTY, as a background spectrum instead of a more elaborated one. Another problem which increases the runtime, is that DUSTY models with a good fitness, that continue to the next generation are not stored anywhere, but are calculated again.  All those points together will likely lead to an extension of the runtime by up to a factor of two. In the current version of our code we did take these points into account as we were aiming at testing the general fitting possibilities of this approach first. Solutions to most of these points are straightforward and their implementation is foreseen for a forthcoming version of our code currently under development. The results shown in this paper are very promising, thus we expect that our method, keeping in mind its limitations, will finally provide a valuable tool for fitting dust spectra of AGB stars."
    },
    "1004/1004.3421_arXiv.txt": {
        "abstract": "We present HI observations performed at the GMRT of the nearby dwarf galaxy NGC 1560. This Sd galaxy is well-known for a distinct ``wiggle'' in its rotation curve. Our new observations have twice the resolution of the previously published HI data. We derived the rotation curve by taking projection effects into account, and we verified the derived kinematics by creating model datacubes. This new rotation curve is similar to the previously published one: we confirm the presence of a clear wiggle. The main differences are in the innermost $\\sim 100 \\arcsec$ of the rotation curve, where we find slightly ($\\lesssim 5$ km s$^{-1}$) higher velocities. Mass modelling of the rotation curve results in good fits using the core-dominated Burkert halo (which however does not reproduce the wiggle), bad fits using the a Navarro, Frenk \\& White halo, and good fits using MOND (Modified Newtonian Dynamics), which also reproduces the wiggle. ",
        "introduction": "Rotation curves of spiral galaxies are one of the most important tools to investigate the content and distribution of dark matter in galaxies. They have been used for a variety of purposes, in particular to investigate their systematic properties (Persic, Salucci \\& Stel 1996, Salucci et al. 2007, Gentile 2008), to test the validity of the predictions of the Cold Dark Matter (CDM) theory (e.g., de Blok et al. 2001, Marchesini et al. 2002, Gentile et al. 2004, 2005, 2007a, Kuzio de Naray et al. 2006, Corbelli et al. 2010), or to study the connection between the distributions of dark and luminous matter (Broeils 1992, McGaugh et al. 2000, McGaugh 2005a, Gentile et al. 2009, Donato et al. 2009).  Testing the validity of the CDM predictions is a very important issue, because (CDM-only) simulations result in dark matter halos with an almost universal density profile (the details of how universal the profile actually is have extensively been discussed in the literature), which is well described by the NFW (Navarro, Frenk \\& White 1996) halo, characterised by a central density cusp (the density $\\rho$ is proportional to $r^{-1}$ when the radius $r$ tends to zero, but see also Section \\ref{sect_results}), whereas observations tend to indicate the presence of a constant density core. The influence of baryons on the distribution of dark matter is a crucial point; however, there is no general consensus about what the dominant effect(s) is(are). A non-exhaustive list of ways by which baryons can change the distribution of dark matter includes: adiabatic contraction (Blumenthal et al. 1986, Gnedin et al. 2004, Sellwood \\& McGaugh 2005), the influence of bars (e.g., Weinberg \\& Katz 2002, McMillan \\& Dehnen 2005, Sellwood 2008), or the influence of gas (Mashchenko, Couchman \\& Wadsley 2006, Governato et al. 2010). Because of the additional complication brought by baryons, dwarf galaxies and low surface brightness (LSB) galaxies are better suited for deriving the properties of dark matter in galaxies.  An alternative explanation to dark matter in galaxies is MOND (Modified Newtonian Dynamics, introduced by Milgrom 1983), where the effective gravitational acceleration becomes stronger than expected in the Newtonian case, when the gravitational acceleration falls below a critical value, $a_0 \\sim 1.2 \\times 10^{-8}$ cm s$^{-2}$. MOND explains very well the observed kinematical properties of galaxies: LSB galaxies (McGaugh \\& de Blok 1998), tidal dwarf galaxies (Gentile et al. 2007b), the Milky Way (Famaey, Bruneton \\& Zhao 2007a, McGaugh 2008, Bienaym\\'e et al. 2009), early-type spiral galaxies (Sanders \\& Noordermeer 2007), elliptical galaxies (Milgrom \\& Sanders 2003, Tiret et al. 2007), and galaxy scaling relation in general, such as the baryonic Tully-Fisher relation (McGaugh 2005b). However, let us note that the MOND prescription is not sufficient to explain the observed discrepancy between luminous and dynamical mass in galaxy clusters (e.g. Angus et al. 2007).  The dwarf galaxy NGC 1560, whose rotation curve was derived by Broeils (1992, hereafter B92) based on WSRT (Westerbork Synthesis Radio Telescope) observations, is a nearby Sd galaxy with an absolute B-band magnitude of $M_{\\rm B}=-16.6$ (Krismer et al. 1995, assuming a distance of 3.5 Mpc). Estimates of its distance vary significantly from one study to another: values from 2.5 Mpc (Lee \\& Madore 1993) to 3.7 Mpc (Sandage 1988) can be found. B92 assumed a distance of 3 Mpc. Karachentsev et al. (2003) find $3.45 \\pm 0.36$ Mpc from the tip of the red giant branch method, using HST data. We assume this value unless stated otherwise, because it is one of the most accurate to date.  NGC 1560 is a very interesting galaxy to study because it is the stereotypical galaxy displaying what is known as ``Renzo's rule'' (from Sancisi 2004): for every feature in the distribution of visible matter there is a corresponding feature in the total distribution of matter. In the rotation curve of NGC 1560, as derived by B92, there is a clear ``wiggle'' in the total rotation velocity, which corresponds very closely to a similar wiggle in the gas contribution to the rotation curve. Mass models such as MOND naturally reproduce the feature, whereas models that include a dominant spherical (or triaxial) halo are too smooth to do so. This motivated us to reobserve NGC 1560 at higher resolution, to independently trace the rotation curve and probe the region around the velocity wiggle.  In the present paper, we present an analysis of HI observations performed with the GMRT (Giant Metrewave Radio Telescope) which have a spatial resolution almost two times better than the data presented in B92. We re-derive the rotation curve and make mass models, with various assumptions concerning the distribution of dark matter.    \\begin{figure*} \\begin{center} \\includegraphics[scale=1.15]{channels_hires_reduced.eps} \\end{center} \\caption{Observed channel maps of NGC 1560. The heliocentric radial velocity (in km s$^{-1}$) is indicated at the top right corner of each channel map. Contours are -2.5, 2.5 ($\\sim 3 \\sigma$), 5, 10, 20, 40 mJy beam$^{-1}$. The cross shows the location of the galaxy centre. The synthesized beam is 8.1$ \\arcsec \\times 6.4 \\arcsec $. } \\label{channels_hires} \\end{figure*} ",
        "conclusions": "NGC 1560 is a nearby dwarf Sd galaxy, whose rotation curve has a very distinct ``wiggle''. We observed NGC 1560 in HI with the GMRT, achieving a two times better resolution than the previous data of Broeils (1992), which were obtained with the WSRT.  We re-derived the rotation curve of NGC 1560 by taking projection effects into account (because of its high inclination angle, $\\sim 80^\\circ$), and checked the reliability of our findings by creating model datacubes, which were compared to the observations.  The new rotation curves is similar to the one derived by Broeils (1992), the main differences being in the innermost $\\sim 100 \\arcsec$: at those radii we find slightly ($\\lesssim 5$ km s$^{-1}$) higher velocities than Broeils (1992). Also, we confirm the presence of a ``wiggle'' in the rotation curve, at around 300$\\arcsec$.      The rotation curve was then corrected for asymmetric drift and used as input for mass modelling. The contribution of the stellar disk to the rotation curve was derived from NIR (I-band) data. The core-dominated Burkert halo gives a good fit to the observed rotation curve, but it does not manage to explain the wiggle. The NFW halo gives a bad fit, greatly overpredicting the velocity in the innermost regions and slightly underpredicting the outermost ones; using an Einasto halo would only slightly change the fits, making them marginally even worse. MOND gives a very good account of the data, particularly of the wiggle.  HI observations at about twice the spatial resolution of the previous ones confirmed thus that NGC 1560 is a nice example of the connection between baryons and total kinematics in galaxies (an expression of which is MOND).      \\begin{figure*} \\begin{center} \\includegraphics[scale=0.82]{combined_separately_simple.eps} \\end{center} \\caption{Rotation curve fits using MOND, fitting separately the two sides of the galaxy and using the simple interpolation function. On the left is the northern (receding) side of the galaxy, on the right is the southern (approaching) side. Lines and symbols are like those in Fig. \\ref{rcfit.bur}. } \\label{combined_separately_standard} \\end{figure*}"
    },
    "1004/1004.3798_arXiv.txt": {
        "abstract": "Magnetohydrodynamics (MHD) provides the simplest description of magnetic plasma turbulence in a variety of astrophysical and laboratory systems. MHD turbulence with nonzero cross helicity is often called \\emph{imbalanced}, as it implies that the energies of Alfv\\'en fluctuations propagating parallel and anti-parallel the background field are not equal. Recent analytical and numerical studies have revealed that at every scale, MHD turbulence consists of regions of positive and negative cross helicity, indicating that such turbulence is inherently locally imbalanced. In this paper, results from high resolution numerical simulations of steady-state incompressible MHD turbulence, with and without cross helicity are presented.  It is argued that the inertial range scaling of the energy spectra ($E^\\pm$) of fluctuations moving in opposite directions is independent of the amount of cross-helicity. When cross helicity is nonzero, $E^+$ and $E^-$ maintain the same scaling, but have differing amplitudes depending on the amount of cross-helicity. ",
        "introduction": "Magnetohydrodynamic (MHD) equations govern the dynamics of a plasma when the spatial scales of interest greatly exceed any intrinsic plasma scale. The MHD model is relevant in understanding waves, instabilities, and turbulence in a variety of astrophysical systems. For instance, magnetohydrodynamic (MHD) turbulence plays a key role in the theoretical modeling of the spectra of velocity and magnetic fluctuations in the Solar Wind \\citep[e.g.,][]{coleman66,coleman68,belcher71,marsch91,goldstein95} as well as electron density fluctuations in the interstellar medium \\citep[e.g.,][]{armstrong95,lithwick01}.   Theoretical modeling of the spectra in strong MHD turbulence has developed on different fronts: phenomenology, statistical closures and more recently, high-resolution numerical simulations. The first phenomenological model, a la Kolmogorov~\\cite{K41,obukhov41}, of the energy spectrum was developed independently by Iroshnikov~\\citep{iroshnikov63} and Kraichnan~\\citep{kraichnan65}. In these works they realized that the magnetic field of the large scale eddies acts in much the same way as a guide field that supports smaller scale Alfv\\'enic fluctuations, and the turbulent energy transfer takes place as the result of many cumulative collisions between counter-propagating Alfv\\'en wave packets traveling along the local magnetic field.  One deficiency of this phenomenology is that it is based on the assumption of an isotropic spectral transfer, in contradiction with more recent results that reveal the anisotropic character of MHD turbulence~\\cite[e.g.,][]{montgomery81,shebalin83}. To account for anisotropy in the strong turbulence regime, Goldreich and Sridhar~\\citep{GS95}, hereafter referred as GS, proposed a new phenomenology in which they introduced the so-called {\\em critical balance} condition, stating that the turbulence is strong when the time scale for wave propagation along the magnetic field matches the nonlinear interaction time. In the past decade, constantly increasing massively parallel computer clusters have become large enough to allow one to simulate the inertial range spectra of strong MHD turbulence, leading to an extensive volume of work addressing universal scaling laws in MHD turbulence \\citep[e.g.,][]{cho00,maron01,haugen03,biskamp,muller05, mason06,mason08,mininni07,beresnyak08,perez08,perez09}.  These simulations have motivated new phenomenological models and have resulted in a renewed interest in the fundamentals of MHD turbulence.  One aspect of MHD turbulence, which has recently drawn significant attention is the role that cross-helicity plays in the turbulent cascade~\\citep{lithwick07,chandran08,beresnyak09,perez09,podesta10,perez10}. Both total energy $E=\\frac{1}{2}\\int (v^2 +b^2)d^3 x$ and total cross-helicity $H^c=\\int ({\\bf v}\\cdot {\\bf b})d^3 x$ are conserved by nonlinear interactions; where ${\\bf v}$ and ${\\bf b}$ are velocity and magnetic fields. In MHD turbulence both energy and cross helicity cascade from large to small scales as a result of the nonlinear interaction of counter-propagating Alfv\\'en fluctuations. If we denote $E^\\pm=\\frac{1}{4}\\int ({\\bf v}\\pm {\\bf b})^2\\, d^3 x$, which are the energies carried by Alfv\\'en fluctuations propagating in opposite directions (the details are given below), the total energy and cross-helicity of the system are given by $E=E^++E^-$ and $H_c=E^+-E^-$, respectively. Therefore, when the latter is nonzero, the energies of waves propagating along and against the guide field are not equal, and in that sense the turbulence is called \\emph{imbalanced}.  Independent numerical simulations by different groups have demonstrated that strong MHD turbulence is always locally imbalanced, in the sense that in the steady state, the turbulence develops correlated regions of positive and negative cross-helicity, irrespective of the overall cross-helicity of the system.  Imbalance turbulence is also present in the solar wind, where velocity and magnetic fluctuations show high correlations of a preferred sign, that is, the normalized \\emph{cross-helicity} $\\sigma_c = {\\left\\langle\\vec v\\cdot\\vec b\\right\\rangle}/{E}={H_c}/{E}$ is close to unity. The preferred positive sign of $\\sigma_c$ indicates that there is more energy in Alfv\\'en waves propagating outwards from the Sun than propagating inwards~\\cite{belcher71}.  Recently, several phenomenological models have addressed steady-state strong imbalanced MHD turbulence~\\cite{lithwick07,beresnyak08,chandran08,perez09,podesta10}, some with support from numerical simulations and observations. However, these models have lead to conflicting predictions. For instance, the theories by Lithwick et al \\cite{lithwick07} and Beresnyak and Lazarian~\\cite{beresnyak08} conclude that in imbalanced regions the Els\\\"asser spectra have the same scalings $E^+(k_\\perp)\\propto E^-(k_\\perp)\\propto k_\\perp^{-5/3}$. The theory by Chandran~\\cite{chandran08} proposes that the spectra of $E^+(k_\\perp)$ and $E^-(k_\\perp)$ are different depending of the degree of imbalance. Finally, the analysis by Perez \\& Boldyrev~\\cite{perez09} and Podesta \\& Bhattacharjee~\\cite{podesta10} finds that the spectra of $E^+(k_\\perp)$ and $E^-(k_\\perp)$ have different amplitudes but the same scalings $E^+(k_\\perp)\\propto E^-(k_\\perp)\\propto k_\\perp^{-3/2}$.  In this work we present results from high resolution direct numerical simulations of strong and steadily driven MHD turbulence, aimed to elucidate the role of cross-helicity and resolve the controversies among different theories.  Hereafter, strong turbulence is defined in the sense of Golreich and Sridhar~\\cite{GS95}. This paper is organized as follows. Section \\ref{modeleqns} briefly describes the MHD equations and the relevance of using the Reduced MHD model for strong turbulence simulations. Section \\ref{numstrategy} describes in detail the proposed numerical strategy and discusses important numerical aspects that need to be considered when simulating strong imbalanced MHD turbulence. In section \\ref{numresults}, the most relevant results from an extensive number of numerical simulations are presented. In section \\ref{pheno} a phenomenological model for the simulation results is proposed and section \\ref{conclusion} presents the conclusions. ",
        "conclusions": "\\label{conclusion} We have presented a detailed numerical setup based on the Reduced MHD model (RMHD), and results from high resolution simulations of balanced and imbalanced MHD turbulence in steady state. We have used this numerical set up to address currently existing controversies regarding the spectra of imbalanced MHD turbulence. The simulations are consistent with the theories and observations predicting same scaling for both Els\\\"asser fields~ \\citep{beresnyak08,lithwick07,perez09,podesta10} and less consistent with the models predicting different scalings for $E^{\\pm}$, \\citep[e.g.,][]{chandran08}. The measured scaling exponent in the simulations is close to the $-3/2$ supported by phenomenological models based on dynamic alignment~\\cite{perez09,podesta10}. The analysis of our simulation results may also explain somewhat puzzling numerical findings by \\citet{beresnyak08,beresnyak09}, who report different spectra for the $E^{\\pm}$ Els\\\"asser fields, and the intersection of the spectra rather than pinning at the dissipation scale. According to our results, the explanation might lie in the fact that in these simulations the imbalance was typically high, up to $\\gamma^2=(z^+)^2/(z^-)^2 \\sim 1000$, and, therefore, the universal regime of imbalanced MHD turbulence was not reached, see our analysis in section~\\ref{re}.   The phenomenology of scale-dependent dynamic alignment can be applied to explain the observed spectra. In this phenomenology, the configuration space splits into eddies (domains) with highly aligned and anti-aligned magnetic and velocity fluctuations, where nonlinear interactions are reduced, as in left panel of Fig.~\\ref{domains_corr}. Even when the turbulence is balanced overall it still can be imbalanced locally, creating domains of positive and negative cross-helicity. In each of these regions the picture of imbalanced turbulence presented above applies. When averaged over all the regions, the spectra of balanced turbulence are reproduced.    This work was supported by the U.S. DOE Junior Faculty grant DE-FG02-07ER54932, by the DOE grant DE-SC0001794, and by the NSF Center for Magnetic Self-Organization in Laboratory and Astrophysical Plasmas at the University of Wisconsin-Madison. High Performance Computing resources were provided by the Texas Advanced Computing Center (TACC) at the University of Texas at Austin under the NSF-Teragrid Project TG-PHY070027T."
    },
    "1004/1004.2541.txt": {
        "abstract": "{} %AIM {The brown dwarf (BD) formation process has not yet been completely understood. To shed more light on the differences and similarities between star and BD formation processes, we study and compare the disk fraction among both kinds of objects over a large angular region in the Taurus cloud. In addition, we examine the spatial distribution of stars and BD relative to the underlying molecular gas} %METHOD %{ More than 40 young BDs are known in the Taurus cloud, spread across a wide surface area of more than 200 square degrees. %%(Guieu et al., 2006; Luhman, 2006) ; %%These surveys should be 75\\% complete for $A_V<4$. %% (Luhman 2006) %%   along the molecular gas filaments in a 32 square degrees area. %This substantial population has been the subject of many investigations. Guieu et al. (2007) have published a study of the circumstellar environment of 23 brown dwarfs in the Taurus cloud and found that $\\approx 48\\%$  have disks, a result consistent with previous work by Luhman et al. (2006). Guieu et al. (2007) also found that the objects with and without disks may possibly be distributed differently across the Taurus cloud. {In this paper, we present new and updated  photometry data from the Infrared Array Camera (IRAC) aboard the {{\\em Spitzer Space Telescope}} on 43 BDs in the Taurus cloud, and recalculate of the BD disk fraction in this region. %We look for  BDs in transition from the class~II to class~III stages in the Taurus cloud. We also useed recently available CO mm data  to study the spatial  distribution of stars and BDs relative to the cloud's molecular gas.} %RESULTS { We find that the disk fraction among BDs in the Taurus cloud is $41\\pm12\\%$, a value statistically consistent with the one among TTS  ($58\\pm9\\%$). We find that BDs in transition from a state where they have a disk to a diskless state are rare, and we study one isolated example of a transitional disk with an inner radius of $\\approx 0.1\\,AU$ (CFHT BD Tau 12, found via its relatively small mid-IR excess compared to most members of Taurus that have disks. We find that BDs are statistically found in regions of similar molecular gas surface density to those associated with stars. %, although there is a factor of two excess of BD in directions with very low H$_2$ surface density. Furthermore, we find that the gas column density distribution is almost identical for stellar and substellar objects with and without disks. % We conclude that the possible excess of BD disks in the northern part of the Taurus cloud might result from a faster disk evolution for BD than for stars and, if this is correct, we predict that the southern part of the Taurus cloud could be some $10^4-10^5\\,$yr older than the northern part. } {} %  study bd formation models and BD distribution in the Taurus cloud.\\\\ %    study the BD spatial distribution and their CS environment to the light of new {\\em Spitzer} data.\\\\ % mostly updated values for : - BD disk frequency, - BD+disk repartition vs CO cloud. ",
        "introduction": "%The Taurus molecular cloud is a nearby (140 pc, see e.g. Kenyon et al., 1994) prototypical low mass star formation region. %It is young (1-3 Myr) so the dynamical effects remain limited and yet it harbors objects ranging from ccks %s~0 sources to near zero-age main sequence. The Taurus stellar population is loosely spread ($1-10\\,$star\\,pc$^{-3}$), making it a very different environment e.g. from the dense region of the Trapezium, so that mutual interactions due to the in and out of star formation activity remain limited; there are no bright stars there to irradiate and disturb the stellar surroundings. %Taurus has a large spatial extension (over a hundred square degrees) so it can be studied for spatial distribution effects in the YSO population. Optical large scale surveys are a needed preliminary data to perform a wide-range search for YSOs in Taurus. However, until recently, such high sensitivity large scale surveys were lacking. On the same hand, apart from 2MASS, no real sensitive NIR survey data are available on the Taurus cloud. The recent availability of large cameras on the CFHT has allowed to reverse this trend. Such surveys should allow to study YSOs with comparable sensitivity in the Optical and in the near infrared as provided by 2MASS.    The brown dwarf (BD\\footnote{see Section~\\ref{sub:bdsample} for a practical definition of what is referred to as a brown dwarf in this paper.}) formation process has not yet been completely understood. Two main classes of models are currently envisioned to explain the origin of BDs. In the \"stellar\" formation scenario (eg. Padoan \\& Nordlund 2002; 2004), BDs form like stars via gravitational collapse and fragmentation of very low-mass cores, followed by significant disk accretion. In the \"ejection\" formation scenario (Reipurth \\& Clarke 2001), BDs are ejected from their parent core and end up being starved for molecular material. In this model, accretion disks around BDs are radially truncated by stripping during the dynamical ejection (see eg. Bate 2009). There is currently no definitive evidence available to reject one or the other scenario either from objects' global properties or from study of individual objects. Studying the disk proportion and structure around BDs may be a key to determining their formation model.  %Clues have been found to support both scenarios, from statistical studies as well as from studies of individual objects. %The fact that brown dwarfs bear disks and display T Tauri associated phenomena in similar proportion as TTS argues in favor of a common formation model  to that of stars. On the one hand, Luhman et al. (2005; 2006) have studied {\\em Spitzer Space Telescope} (Werner et al. 2004) photometry of young objects in Taurus, and from their {\\em Spitzer} $3.6$ to $8\\,\\mu$m colors, found a BD disk fraction of $\\approx 40\\%$, a value close to that of stars. Guieu et al. (2007) have studied and modeled  the SEDs of 23 substellar objects in the Taurus cloud %(with $A_V<10$) and find that $48\\%$ of BDs have disks, a fraction similar to that for T~Tauri stars (TTS) in the same region. In Chamaeleon~I, Damjanov et al. (2007) find no dependence of disk frequency on stellar mass in the K3-M8 spectral type range, well into the BD regime at the age of the cloud ($\\approx 2\\,$Myr, Luhman 2004). The disk phenomenon extends into the planetary mass regime ($M\\lesssim 15\\,M_{\\rm Jup}$). Circumstellar disks have been detected around BDs within the mass range of giant planets ($<15 M_{Jup}$), by Luhman et al. (2005, 2008a,b); Scholz \\& Jayawardhana (2008). % Luhman et al. (2005, 2008a, 2008b) have found circumstellar disks around planetary-mass objects;  Scholz \\& Jayawardhana (2008) find evidence of disks around isolated planetary mass objects in the $\\sigma\\,$Orionis cluster. Concerning individual objects, Bouy et al. (2008) have published an extensive data set on the substellar object 2MASS~J04442713+2512164, from the visible to the radio range, including the first photometric measurement of a BD disk at 3.7\\,mm, and allowing a detailed analysis of the disk properties. Their analysis shows that this BD has all the characteristics of a star with a disk. %Joergens (2006) found that the RV dispersion of BD in Cham~I is consistent within the error bars with that of T~Tauri stars, a result in line with the absence of mass dependence in models of ejection scenario for the formation of BD.  On the other hand, Luhman et al. (2008a), used {{\\em Spitzer}} photometry on a much larger number of objects than Damjanov et al. (2007) in Cha~I, to find a dependence on stellar mass in the disk fraction. In Upper Scorpius, Scholz et al. (2007) find a disk frequency  of 37\\%+/-9\\%, a value higher than what has previously been derived for K0-M5 stars in the same region (at a $1.8\\,\\sigma$ confidence level), suggesting a mass-dependent disk frequency. \\\\ Luhman et al. (2007) have presented a disk model for the BD 2MASS J04381486+2611399 showing that the disk has an inner hole and has an outer radius  ($R_{out} \\approx 20-40\\,$AU), a size similar to the predictions from hydrodynamical models ($R_{out} \\lesssim 10-40\\,$AU) by Bate et al. 2003  and Bate 2009. \\\\ Thies and Kroupa (2007), studying recent data on the multiplicity properties of stars and BDs show them to have different binary distribution functions. They uncovered a discontinuity of the multiplicity-corrected mass distribution in the very low mass star (VLMS) and BD mass regime, discarding a continous IMF with a high confidence level, and suggesting that VLMS and BDs on the one hand, and stars on the other, are disjoint populations with different dynamical histories.\\\\ %Morrow et al. (2008) have studied disk evolution around BDs in the TW~Hya association and find that disk evolution via grain growth around BDs could be faster  than around stars, or the radius of mid-IR emitting region in BD disks could be smaller, with grains growing faster at closer radii where the density is higher. A smaller mid-IR emitting region could be consistent with a smaller disk, as a consequence of disk stripping during ejection. Another explanation would be that BDs being colder objects, the warm emission region (300\\,K, $10\\,\\mu$m) of their circumstellar environment has a smaller extent. %in the Upper Scorpius OB association, Bouy \\& Martin (2009) found  that a significant fraction of confirmed young objects  show discrepant proper motions. None of these outlier show signs of accretion, a result consistent with the fact that ejection actually removes the disks. In the Taurus cloud, Guieu et al. (2007) find that BD disk fraction appears to vary with the object position within the Taurus filaments, unlike the known young star population.  %These two apparently contradictory sets of results could suggest that BDs share some properties in their formation process with stars (initial fragmentation and further disk accretion) but are more influenced by dynamical and ejection interactions. %The young brown dwarfs properties have been studied in a variety of star formation regions (see eg. Luhman et al., 2007; Guieu et al., 2006, 2007 ; Allen et al., 2007). %It is still  unclear whether the disk presence and fate depends more on the objects' central mass or on the surrounding environment (Damjanov et al. 2007).  Disk fraction is therefore one of the keys of the very low mass objects formation models. With this purpose in mind, we present in this paper new data about BD disks in the Taurus cloud. %, in order to shed some more light on the BD  formation model controversy. We obtained new or updated {\\em Spitzer} Infrared Array Camera (IRAC; Fazio et al. 2004) photometry data on 43 BDs in the Taurus cloud (Section~\\ref{sec:obs} and \\ref{sec:disk-frac}) and provide an update on the BD disk fraction in this region, complementing the study started by Guieu et al. (2007). We compare the BD disk fraction to the TTS disk fraction in the same region. We present a complete study of the only transition disk found among our BD sample (section~\\ref{sec:cfht12}). We also make use of recent CO data (Goldsmith et al. 2008) to study the distribution of stellar and substellar objects relative to the molecular gas (Section~\\ref{sec:*bd/gas}) and the distribution of objects with and without disks (section~\\ref{sec:obj-dnd}). %A discussion and the conclusions are presented in sections~\\ref{sec:disc} and \\ref{sec:concl}.  %The Taurus molecular cloud is a nearby (140 pc, see e.g. Kenyon et al., 1994) prototypical low mass star formation region. %It is young (1-3 Myr) so the dynamical effects remain limited and yet it harbors objects ranging from class~0 sources to near zero-age main sequence. The Taurus stellar population is loosely spread ($1-10\\,$star\\,pc$^{-3}$), making it a very different environment e.g. from the dense region of the Trapezium, so that mutual interactions due to the in and out of star formation activity remain limited; there are no bright stars there to irradiate and disturb the stellar surroundings. ",
        "conclusions": "\\label{sec:concl} We have presented new photometric measurements in the $3-8\\,\\mu$m domain from {\\em Spitzer} IRAC  instrument, complementing previous IR photometry by Guieu et al. (2007) for more than 40 BDs in the Taurus cloud. We associate these IR measurements with new H$_2$ surface density computations from Goldsmith et al. (2008) to study and compare the distribution of stellar and substellar objects relative to the underlying molecular gas. Based on this larger sample, we update the disk fraction for stars and BDs in the Taurus cloud and find that the disk fraction around BDs is similar to the disk fraction in  their TTS counterparts.  % This could be evidence that object formation via collapse and disk accretion could be mass dependent. We find that transition disks are very rare among BDs, a common result with TTS showing that the disk's disappearance timescale is also very short for substellar objects. These results call for further theoretical models of dust coagulation and settling as a function of the central object's mass. We model the only transition disk we find around CFHT~BD~Tau~12, an object that could have experienced more dust evolution than other members in Taurus. Note that Luhman et al (2006) find a slightly earlier spectral type than us for this object (M6 instead of M6.5) so CFHT~BD~Tau~12 could be an object at the border between VLMS and BD rather than a bona fide young BD.  We study the distribution of stars and BDs relative to the H$_2$ molecular gas surface density computed from CO observations and find that both distributions are likely drawn from the same distribution, even if there appears to be an excess of BDs found at low H$_2$ surface density. A similar study for objects with and without disks shows that they have the same distribution relative to the underlying gas. %Thus it is unlikely that the surrounding environment (molecular density) plays an important role in determining wether an object formed will bear  a disk or not. It could also mean that the ejection process is rather mild and does not systematically strip BD disks.  Finally, the finding that the fact that BD with disks are  less numerous in the southern than in the northern portion of the Taurus filaments could be a reflection of an older age for the southern part of the cloud combined with a shorter lifetime for BD disks. A further determination of the ages of the Taurus filaments via HR diagram for instance is difficult due to the large dispersion one obtains in the ages. Disk fraction measurements in Cha~I or IC 348 show no evidence that the lifetimes of substellar objects are shorter.  %, is either a statistical oddity, or has nothing to do with the surrounding environment. This result could have two consequences : i) the southern part of the Taurus cloud is slightly older and, if real, this provides a constraint on the global Taurus complex formation history, ii) disks around BD evolve more rapidly than around more massive objects, either through their disk lifetime and / or their disks' dissipation lifetime, a result consistent with previous results by Morrow et al. (2008) on silicate emission from BD."
    },
    "1004/1004.3084_arXiv.txt": {
        "abstract": "We perform direct numerical simulations of forced and freely decaying 3D magnetohydrodynamic turbulence in order to model magnetic field evolution during cosmological phase transitions in the early Universe. Our approach assumes the existence of a magnetic field generated either by a process during inflation or shortly thereafter, or by bubble collisions during a phase transition. We show that the final configuration of the magnetic field depends on the initial conditions, while the velocity field is nearly independent of initial conditions. ",
        "introduction": "Astronomical observations show that galaxies have magnetic fields with a component that is coherent over a large fraction of the galaxy with field strengths of order $10^{-6}$ Gauss (G); see Refs.\\ \\cite{Widrow,Vallee,Giovannini2008,beck} and references therein.\\footnote{ See Ref.\\ \\cite{neronov} and references therein for a technique that might soon prove useful for measuring a larger-scale, cosmological, magnetic field.} Modeling the origin of these fields is a challenging pro\\-blem. Here we consider models in which the seed field is generated in the early Universe, see, e.g. \\cite{Widrow,Giovannini2008,subramanian}. There are a number of such models and different models result in different magnetic field structures.  It is well known that quantum-mechanical fluctuations during inflation can generate energy-density inhomogeneities that seed observed large-scale structure, see, e.g., Ref.\\ \\cite{fischler}. Quantum-mechanical zero-point fluctuations can also generate a seed magnetic field, provided the relevant abelian gauge field is such that the lagrangian density is not conformally invariant during inflation \\cite{turner,ratra,bamba}. A seed magnetic field generated during inflation has a correlation length today that can be very much larger than the current Hubble radius. A convenient phenomenological way of breaking conformal invariance during inflation is to couple the inflaton scalar field $\\phi$ to the vector abelian gauge field by $e^{\\alpha\\phi} F_{\\mu\\nu} F^{\\mu\\nu}$, where $F_{\\mu\\nu}$ is the vector field strength tensor and $\\alpha$ is a parameter \\cite{ratra}. Depending on the value of $\\alpha$, this can result in a sufficiently large seed magnetic field to explain the observed galactic magnetic fields. This is a classically consistent model, with vector field back-reaction during inflation being negligibly small \\cite{ratra,ratralong}. While the abelian gauge field becomes strongly coupled at early times \\cite{ratra,ratralong,backreaction}, this is not a problem for the effective, phenomenological, classical model, \\cite{ratra,ratralong,backreaction} and also see \\cite{subramanian,backreaction2}. Of course, just as one does not reject the effective, phenomenological, classical ``standard'' $\\Lambda$CDM cosmological model, one does not reject inflation-based seed magnetic field generation models for being quantum-mechanically inconsistent. Rather, much as the case for $\\Lambda$CDM, it is of great interest to understand how such a successful inflation-based seed magnetic field generation model might arise from a more fundamental underlying model or theory.  There are a number of other seed field generation models; see Ref.\\ \\cite{horava} for a discussion of seed field generation in Ho{\\v r}ava-Lifshitz gravity. Yet another interesting possibility is primordial magnetic field generation during a cosmological phase transition \\cite{pmf}. If the phase transition takes place at late times the magnetic field correlation length is smaller than the Hubble radius.  If the magnetic field was generated during or shortly after inflation (during a pre- or re-heating phase transition), and survived to the epoch of recombination (the last-scattering surface), it should leave observable signatures on the cosmic microwave background (CMB) radiation fluctuations, see Ref.\\ \\cite{Giovannini2006} for a review and Refs.\\ \\cite{cmb,faraday} for more recent discussions.\\footnote{A potential advantage of a primordial seed field that is coherent over scales larger than the current Hubble radius, as can be generated by inflation, is that it might be able to explain some of the large-scale oddities of the observed CMB temperature anisotropy \\cite{bernui}, including potential non-Gaussianity \\cite{seshadri}.}  The shape and magnitude of a primordial magnetic field can be constrained observationally. Magnetic field energy density scales like radiation (i.e.\\ relativistic) energy density.\\footnote{The ratio of the magnetic field energy density $\\rho_B$ and the energy density of radiation $\\rho_{\\rm rad}$ is constant during cosmological evolution, if the primordial magnetic is not damped by MHD (or other) processes and so stays frozen into the plasma.} During nucleosynthesis, any new form of radiation-like energy density is constrained observationally to be less than about $10\\%$ of the usual radiation energy density. Hence, agreement between the big bang nucleosythesis (BBN) model light nuclei abundance predictions and the observations yields a constraint on the primordial magnetic field energy density: $\\Omega_B h_0^2 \\leq 2.4 \\times 10^{-6}$ \\cite{GrassoR}. Here $\\Omega_B$ is the current value of the magnetic field energy density parameter and $h_0$ is the current value of the Hubble parameter in units of 100 km s$^{-1}$ Mpc$^{-1}$. Naturally, this limit holds only if the primordial magnetic field was generated prior to or during nucleosynthesis.  Other constraints arise through the effects of a primordial magnetic field on the propagation of CMB photons (Faraday rotation of the CMB polarization \\cite{faraday}, and magnetic field induced scalar, vector, and tensor modes of CMB anisotropies \\cite{Giovannini2006,cmb}). Available data limit the current value of a cosmological magnetic field to be less than about $10^{-8}$ G for a scale-invariant or homogeneous primordial magnetic field.  Another interesting signature of a primordial magnetic field is relic gravitational waves generated by the anisotropic magnetic stress \\cite{magnet}. The amplitude of the gravitational wave signal is determined by the efficiency of gravitational wave generation and the strength of the primordial magnetic field. The efficiency of gravitational wave production is small \\cite{ktr09} due to the small value of the Newton constant $G$ in the coupling between the primordial magnetic field and the gravitational waves. However, a magnetic field that satisfies the BBN limit and that is generated during a strong-enough first-order phase transition can lead to a detectable signal for LISA. Therefore, the direct measurement of the gravitational wave background can lead to an independent limit on a magnetic field generated in the early Universe \\cite{wang}.  On large scales, because of the high conductivity of the plasma in the early Universe, the magnetic field is treated as a frozen-in field with its evolution determined by the simple dilution of magnetic field lines, ${\\bf B}({\\bf x}, t) \\propto {\\bf B}({\\bf x})/ a^2(t)$, where $t$ is the physical cosmic time and $a(t)$ is the scale factor. In general, however, the evolution of a primordial magnetic field is a complex process influenced by MHD as well as by cosmological dynamics \\cite{banerjee}. In particular, the presence of a magnetic field can dramatically affect the behavior of primordial turbulence (for example, turbulence associated with phase transition bubble motions) \\cite{axel-1,axel-2,caprini2009}. Also, the presence of a magnetic field might itself lead to the development of turbulent motions, see Refs.\\ \\cite{B, Verma, axel} and Sec.~III below.  In this paper we study the evolution of a primordial magnetic field that is coupled, through the MHD equations, with the fluid, during a cosmological phase transition. We consider two different possibilities: (i) when the magnetic field energy has been injected in the plasma with no initial vorticity perturbations (no turbulent motions) at a given typical scale; and, (ii) when vorticity perturbations are present during the phase transition and couple to the magnetic field with a given spectrum. Both these assumptions can be justified. The magnetic field might be generated prior to the phase transition, with the generation mechanism not requiring turbulent motion. On the other hand, the presence of initial turbulent motions can also be justified since bubble nucleation, expansion, and collisions can lead to primordial kinetic turbulence \\cite{kos1}. Of course, in reality both magnetic field and kinetic turbulence generation can take place together, while in our approach we assume the a priori presence of magnetic energy and/or kinetic energy during the phase transition.  A causal, phenomenological description of MHD on large scales during the cosmological phase transition motivates a white noise spectrum for the magnetic spectral energy density, $E_M(k) \\propto k^2$, where $k$ is the wavenumber \\cite{hogan}. A steeper Batchelor spectrum with $E_M(k) \\propto k^4$ has been claimed to follow from cosmological causality and the divergenceless condition of the magnetic field \\cite{Durrer}. To understand the evolution of a cosmological magnetic field, and to be able to make observable predictions in this model, it is important to resolve this impasse. One would like to know whether the final large-scale magnetic energy spectrum evolves to something steeper than white noise, or to a spectrum closer to the Batchelor one.  In fact, the MHD process itself might establish a different spectrum: The possibility of generating a random magnetic field from isotropic turbulence was first proposed by Batchelor \\cite{batchelor}. He invoked the imperfect analogy between vorticity and magnetic fields that should generally imply a large scale distribution of the magnetic energy similar to the kinetic energy of turbulent motions $E_K(k) \\propto k^4$ \\cite{batchelor2}. Later, Kazantsev \\cite{kazantzev} was able to rigorously establish the possibility of small-scale dynamo action. He assumed that the velocity field varied only on large scales and found that weak magnetic fields are amplified mainly on the resistive scale. Initially, this was thought to be applicable only to turbulence at large magnetic Prandtl number \\cite{Scheko}, where the viscous cutoff scale is much larger than the resistive cutoff scale. Kazantsev found that the magnetic energy spectrum increases with wavenumber like $k^{3/2}$, which is slightly shallower than the white noise spectrum $k^2$. The emergence of a Kazantsev spectrum turned out to be much more ubiquitous and not only applicable at large magnetic Prandtl number. Simulations at magnetic Prandtl numbers of order unity also clearly showed the Kazantsev $k^{3/2}$ spectrum \\cite{axel-2}. Such a spectrum could be of interest for primordial magnetic field evolution, because it implies somewhat larger power at large scales than white noise.  The main goal of our paper is to determine, through MHD modeling, the evolution of the magnetic energy spectrum. To keep our study as general as possible we do not make any assumption about the physical process leading to primordial magnetic field generation. Also, we do not address the phase transition physics itself, keeping the total magnetic field energy density $\\rho_B$ as a free parameter whose maximal value is fixed by the BBN bound. Obviously, if the magnetic field is generated during a phase transition, $\\rho_B$ will depend sensitively on the amount of latent heat that is transformed to magnetic energy, i.e.\\ on the efficiency of the magnetic field generation process. One important issue addressed here is to determine the spectral shape of the magnetic field at large scales, assuming that magnetic energy and vorticity perturbations are closely coupled during the phase transition. Another important question is related to the duration of MHD turbulence and how long it takes to reach equipartition between kinetic and magnetic energy densities in the primordial plasma.  The structure of this paper is as follows. In Sec.\\ II we describe the model, defining the main model parameters (Sec.\\ IIA), characterizing the magnetic field spectrum (Sec.\\ IIB), and formulating initial conditions (Sec.\\ IIC). In Sec.\\ III we present results from direct 3D MHD simulations performed using the {\\sc Pencil Code} \\cite{pencil}. We conclude in Sec.\\ IV. We employ natural units with $\\hbar = 1 = c$ and gaussian units for electromagnetic quantities. To properly account for the expansion of the Universe we use comoving quantities with conformal time $t$ that is related to physical time $t_{\\rm phys}$ as $t \\propto t_{\\rm phys}^{1/2}$ during the radiation dominated epoch. ",
        "conclusions": "In this paper we have presented results from direct numerical simulations of primordial magnetic field evolution during cosmological phase transitions. These simulations account for the expansion of the Universe. Simulations were performed to model two different stages of primordial magnetic field evolution, the first stage when phase transition processes drive turbulence, and the second stage when free decay occurs.  We show that different types of initial conditions (drivers) lead to rapid development of a turbulent state close to equipartition. During the following stage we model the free decay of these turbulence configurations. We study the development during this free decay stage and analyze characteristic parameters of the slowly varying power spectrum of the turbulence. We assume that the properties of the turbulence configurations modeled in our numerical experiments are similar to those of cosmological primordial turbulence.  Our simulations allow us to estimate the spectral indices of the large-scale distribution of the kinetic and magnetic field energies in developed turbulence. It seems that the final configuration of primordial magnetic field depends on the initial conditions that must be determined by the phase transition physics. On the other hand, the kinetic energy density of the developed state does not retain any information about the initial conditions and hence is not sensitive to the details of the phase transition. The spectral index of the kinetic energy at large scales can be well approximated as 2 (white noise spectrum), while the magnetic energy density spectral index ranges between 2 and 4 depending on the initial conditions. Similar spectral indices for the large-scale magnetic field are well established for laboratory turbulence, see Ref.\\ \\cite{textbooks}. Note that simulations in a finite periodic domain may suffer from the fact that originally disconnected and causally independent regions come into causal contact within one box turnover time, which, based on the scale of the domain, is $(v_0 k_1)^{-1}$. For the case shown in Fig.~\\ref{Evolution}, this time is about $15\\,t_1$, but the effective time can be even shorter owing to the effects of acoustic and Alfv\\'en waves.  Our numerical results allow us to estimate typical time-scales for decaying free turbulence. It seems that cosmological turbulence decays slightly slower then classical grid turbulence in laboratory experiments. Although we took a rather small damping wavenumber, we cannot expect the establishment of a Kolmogorov-like spectrum at small scales. We see a fast decorrelation of turbulence at small scales. We note that the phenomenological approaches developed in Refs.~\\cite{ktr09,caprini2009} imply a fast Kraichnan-like decorrelation of turbulence. Consequently, only the large-scale or peak-scale magnetic field results may have cosmological significance and contribute observable signatures, such as gravitational wave generation and/or CMB anisotropy production.  We have considered the case of non-helical turbulence. The presence of weak initial helicity can significantly change the development of turbulence. This is because magnetic helicity is a conserved quantity in the limit of large magnetic Reynolds numbers and can inverse cascade to larger scales \\cite{Frisch,Biskamp}, which could be cosmologically significant. This process has been confirmed through numerous simulations \\cite{enqvist,axel-1,axel,banerjee}.   Summarizing, we find that the generation of a magnetic field at phase transition scales will lead to the development of turbulent motions, and, in the case of the electroweak phase transition, this turbulence has an observable signature in the form of a gravitational wave signal."
    },
    "1004/1004.2494_arXiv.txt": {
        "abstract": "{ It is textbook knowledge that open clusters are conspicuous members of the thin disk of our Galaxy, but their role as contributors to the stellar population of the disk was regarded as minor. Starting from a homogenous stellar sky survey, the \\ascc, we revisited the population of open clusters in the solar neighbourhood from scratch. In the course of this enterprise we detected 130 formerly unknown open clusters, constructed volume- and magnitude-limited samples of clusters, re-determined distances, motions, sizes, ages, luminosities and masses of 650 open clusters. We derived the present-day luminosity and mass functions of open clusters (not the stellar mass function in open clusters), the cluster initial mass function CIMF and the formation rate of open clusters. We find that open clusters contributed around 40 percent to the stellar content of the disk during the history of our Galaxy. Hence, open clusters are important building blocks of the Galactic disk.} ",
        "introduction": "Open clusters constitute an important part of a process transforming gas and dust into stars. They are observed as the most prominent parts in the regions of active star formation, or as tracers of the ceased star formation process in the general Galactic field. However, the role they are playing in this process has still not been fully understood. In spite of their prominence, there had been indications that classical open clusters contribute only 10\\% or even less input \\citep{1971A&A....13..309W,1978PASP...90..506M,2006A&A...445..545P} to the total stellar population of the Galactic disk. This discordance can be explained either by an early decay of a considerable fraction of newly formed star clusters \\citep[see e.g.][]{2005A&A...441..117L,2001MNRAS.321..699K,1978A&A....70...57T} or by an insufficient knowledge of cluster formation statistics. In this context, one should note that the most important items of cluster formation like the distribution of cluster masses at birth (i.e., the initial mass function of star clusters) and the cluster formation rate were still poorly known a decade ago.  In principle, basic parameters like distance, motion, age, and metallicity can be determined for an open cluster more accurately than for a single field star. Indeed, they are better tracers of large scale structures of the Galactic disk population than field stars. Nevertheless, the most comprehensive  studies of the Galactic cluster population are about 20 years old \\citep{1982A&A...109..213L,1988AJ.....95..771J}. They were based on the best data available at that time, the Lund Catalogue of Open Cluster Data \\citep[][hereafter, the Lund Catalogue]{lyn87} and its subset of  clusters with 3-colour photometry \\citep{1982ApJS...49..425J}. Although these studies represent an important step in our understanding of the general properties of the cluster population, they suffer from  incompleteness of the cluster samples and from inhomogeneity of the cluster parameters.  About 1200 clusters were known in the Lund catalogue by 1988. Only 400 of them had accurate, but heterogeneous UBV photometry, and photometric distances, reddening and age values. Although for almost all clusters apparent diameters were given in the Lund catalogue (estimated by eye from sky charts or defined by the size of detector field of view), only about 100 clusters were studied in a systematic way on the basis of star counts \\citep{1994A&AT....6...85D}.  Currently, the on-line list of open cluster data by \\citet[][DLAM hereafter]{2002A&A...389..871D}, which can be considered as a continuation of the Lund Catalogue, contains by a factor of 1.5 more clusters than its predecessor, but the degree of completeness of this list is still unknown. Since the cluster data in the DLAM list are taken from literature, the set of the derived parameters differ from cluster to cluster. Also, the parameters themselves are based on heterogeneous observations and different methods of cluster definition and of parameter determination. Whenever using these data for cluster population studies, one is confronted with problems caused by uncertain cluster statistics and data heterogeneity.  In the following sections we decribe our approach to construct a representative sample of open clusters in the solar neigbourhood, derive a homogeneous set of cluster parameters, especially age, mass and luminosity which are basic for estimating the role of open clusters as building blocks in the evolution of the Galactic disk. A previous review on this topic was given by \\citet{2009IAUS..254..221Z} at IAU Symposium 254. ",
        "conclusions": "Summing all this up, it is fair to say that this work on the population of open clusters in the solar neighbourhood that started from an all-sky survey has found that open clusters are larger, more massive, live longer and contribute more to the thin stellar disk of the Galaxy than was believed a decade before.  Although much progress in our knowledge of the statistical properties of the open cluster population has been made in the past decade, there is still a long way to go. The sample discussed here allows us to draw general conclusions, but in some cases the derived parameters are erroneous because of a severe undersampling of the cluster membership due to the bright magnitude limit of the \\linebreak \\ascc survey. New questions have appeared, such as: is the local cluster population representative for the whole disk? how to discern open clusters from compact associations? are these separate populations or can the latter be seen as the high-mass end of the cluster population? What exactly do we mean by ''the mass of a cluster''? Tidal, virial and star-counted masses may not necessarily measure the same mass. To get more insight into this problem, one  must carefully define what a ''cluster member'' is. This becomes especially important near the boundary of a cluster and consequently influences our understanding of ''mass''.  Progress on some of these topics may be expected from an exploitation of the new deep survey catalogue PPMXL \\citep{2010arXiv1003.5852R}, and, on a somewhat longer timescale from Gaia."
    },
    "1004/1004.0172_arXiv.txt": {
        "abstract": "New high energy emission features have been recently discovered by the Cherenkov telescopes from active galaxies e.g., a few minutes variability time scale of TeV emission from Mrk 501 and PKS 2155-304, sub-TeV $\\gamma$-ray emission from GeV peaked blazar 3C 279, and TeV emission from two nearby active galaxies, M87 and Cen~A, which jets are inclined at a relatively large angle to the line of sight. These results have put a new light on the high energy processes occurring in central parts of active galaxies stimulating more detailed studies of $\\gamma$-ray emission models. Here we report the results of a detailed analysis concerning the most general version of the model for the $\\gamma$-ray production by leptons injected in the jet which interact with the thermal radiation from an accretion disk (the so called {\\it external inverse Compton model}). We investigate the $\\gamma$-ray spectra produced in an anisotropic Inverse Compton (IC) $e^\\pm$ pair cascade in the whole volume above the accretion disk. The cascade $\\gamma$-ray spectra are obtained for different locations of the observer in respect to the direction of the jet. We also study the time evolution of this $\\gamma$-ray emission caused by the propagation of the relativistic leptons along the jet and the delays resulting from different places of the origin of $\\gamma$-rays above the accretion disk. We discuss the main features of such a cascade model assuming constant injection rate of electrons along the jet. We are investigating two models for maximum energies of injected electrons: with a constant value independent on the distance along the jet or limited by the synchrotron energy losses considered locally in the jet. The model is discussed in the context of blazars observed at small and large inclination angles taking as an example the parameters of the two famous sources Cen~A and 3C 279. ",
        "introduction": "Recent observations of active galaxies in the TeV $\\gamma$-rays confirmed earlier reports on an extremely variable nature of this emission. It can change on various time scales starting from a few minutes  (e.g. \\cite{gai96,al07,ah07}) up to months/years \\citep{hs09}. This variable emission can correspond to the time scales of, either the periods of enhanced energy generation in the disk-jet system, or acceleration time of particles, or the passing time of relativistic particles through the medium which creates the target for efficient $\\gamma$-ray production. Moreover, the time scales can be modified due to relativistic effects caused by the movement of the emission region towards the observer and, for very distant sources, cosmological propagation. The precise location of the emission site within the active nuclei remains unclear. Theoretical arguments suggest that the emission place is rather distant from the black hole. $\\gamma$-ray radiation produced closer cannot escape to the observer due to a very strong, soft radiation field originating from the inner part of the accretion disk. On the other hand, this strong radiation field can create a target on which $\\gamma$-rays could be produced. Recent observations of the nearby active galaxy, M87, whose jet is observed at a relatively large angle, suggest that the TeV $\\gamma$-ray emission region can likely appear within a few tens of Schwarzschild radius from the black hole \\citep{ah06,acc08,al08a}. It is likely that the emission region is close to the accretion disk inside the inner jet as suggested by the recently discovered connection between the TeV $\\gamma$-ray flare and the subsequent appearance of the radio blobs \\citep{acc09}. Also Cen~A, another nearby active galaxy inclined at large angle, has been recently detected as a TeV $\\gamma$-ray source \\citep{ah09}. This emission is as well interpreted as originated very close to the central black hole \\citep{ra09}.  Recently the LAT Telescope on board of the {\\it Fermi} observatory is operating with an order of magnitude better sensitivity than the EGRET telescope. So far, after 11 months of operations it observed 709 AGNs in the GeV range \\citep{ab10}. Among them 300 are BL Lac objects and 296 are flat-spectrum radio quasars. On top of this, 8 objects have been associated as misaligned blazars. Therefore, investigation of the anisotropic high energy radiation processes in the  active galaxies is very timely.  It is widely believed that $\\gamma$-rays are produced in active galaxies by particles (leptons or hadrons) accelerated in the jet launched perpendicularly to the plane of the accretion disk. The acceleration mechanism is not well known. Usually two idealized cases of angular distribution of particles in the jet are discussed, either isotropic in the jet frame or in the recti-linear acceleration regions generally directed along the jet. The first case corresponds to the acceleration of particles in a shock wave. Due to a mild relativistic motion of the emission region this scenario is sometimes referred to as a {\\it slow jet model}. The latter case may correspond to the acceleration in magnetic field reconnection regions. This scenario is sometimes called as {\\it a fast jet model} because of  an ordered, super-relativistic motion of particles.  $\\gamma$-rays can be produced in different radiation mechanisms. A few basic scenarios have been proposed soon after discovery of the GeV $\\gamma$-ray emission from blazars by the EGRET telescope on the CGRO. Among the famous leptonic models are: synchrotron self-Compton (SSC) model \\citep{mgc92}, external Compton (EC) model with scattering of the disk radiation \\citep{dsm92}, or radiation around the disk \\citep{sbr94}. Also hadronic models have been discussed, e.g. proton initiated cascade (PIC) model \\citep{mb92}. Starting from that discovery period, several modifications of these first more general scenarios have been discussed in detail with the applications to specific objects. A more realistic three dimensional cascade model occurring in the strong radiation of the accretion disk are starting to be considered only recently \\citep{sb10, rb10}. In this paper we are specially interested in the formation of the $\\gamma$-ray spectrum in terms of the external Compton accretion disk model (see for details of this model \\citet{ds93}).  It has been shown that for typical parameters of blazars detected in GeV-TeV $\\gamma$-ray energies, the optical depths for $\\gamma$-rays injected relatively close to the disk at an arbitrary place can be clearly above unity, see e.g. recent calculations for 3C 279 and 3C 273 \\citep{sb08} or for Cen~A \\citep{sb10}. Therefore, the cascade processes can be important in these sources provided that injection of primary particles occurs relatively close to the central engine as suggested by the recent observations of TeV $\\gamma$-ray emission from some objects (e.g. Cen~A, M87).  We have already investigated the $\\gamma$-ray production in blazars in a scenario of IC $e^\\pm$ pair cascade initiated by primary $\\gamma$-rays produced e.g., in the Synchrotron self-Compton model. These $\\gamma$-rays develop IC $e^\\pm$ pair cascade in the thermal radiation from the accretion disk in the whole volume above the disk~\\citep{sb10}. We have calculated the $\\gamma$-ray spectra produced in such a model for arbitrary location of the observer in respect to the direction of the jet. Furthermore the influence of a magnetic field above the accretion disk on the development of the IC $e^\\pm$ pair cascade has been also investigated.  In the present paper we assume that the cascade process in the accretion disk is initiated by leptons (electrons and positrons) accelerated in the jet as postulated by the external IC model \\citep{ds93}. These leptons produce the first generation of $\\gamma$-rays by scattering an anisotropic disk radiation. Therefore the angular distribution and spectrum of those primary $\\gamma$-rays is dependent on the disk/jet system properties. They significantly differ from the spectra obtained in terms of the previous calculations~\\citep{sb10} in which primary $\\gamma$-rays were injected isotropically in the blob frame. These first generation $\\gamma$-rays initiate anisotropic IC $e^\\pm$ pair cascade in the radiation field of the accretion disk. In a framework of this model we investigate spectral, angular and also time properties of the $\\gamma$-ray emission. Therefore, our present calculations follow significantly more complicated and complete model than offered by the previously considered scenario with the isotropic injection of primary $\\gamma$-rays produced in the SSC model. ",
        "conclusions": "We discuss the most complete version of the external IC model for the $\\gamma$-ray production in blazars. It is assumed that electrons are injected in the inner part of the jet and comptonize thermal radiation coming from the accretion disk surface. As a result a first generation of $\\gamma$-rays is produced. These $\\gamma$-rays initiate IC $e^\\pm$ pair cascade in the anisotropic disk radiation. We calculate the $\\gamma$-ray spectra escaping to the observer located at an arbitrary angle in respect to the direction of the jet. In the framework of this model we calculate also the $\\gamma$-ray light curves expected from blazars in which blobs containing relativistic electrons propagate with arbitrary velocities. The angular dependent $\\gamma$-ray light curves and spectra are calculated for blazars observed at large and small inclination angles. As an example we have applied the parameters of the two well known AGNs: Cen~A and 3C~279, which have been recently observed in the TeV $\\gamma$-rays. The first source is located relatively nearby, and has a jet which propagates at a relatively large angle in respect to the observer with a mild velocity. In contrast, the second source is more distant. It's inclined at a small angle and the jet propagates with a large Doppler factor.  In order to perform the cascade calculations with a limited number of free parameters, we made some simplified assumptions. At first, we assume that the disk radiation is the only source of soft photons with which relativistic electrons can interact. The calculations with other radiation fields (e.g. infrared radiation from the molecular torus or the disk and jet radiation scattered by the matter around the accretion disk) are straightforward. But the interpretation of the final results will be more difficult due to the large number of free parameters describing such more realistic model. Therefore, the analysis of such an anisotropic cascade model with a more complicated soft radiation field is left for a future paper. We also assume the simplest possible (i.e. constant along the jet) dependence of the injection rate of relativistic electrons along the jet. This allows us to study the basic features of this complicated cascade model. These features may change significantly for the more complicated forms of the injection rates along the jet. In fact such more complicated dependence for the injection rate is suggested by the observed $\\gamma$-ray light curves of flares observed from blazars at different energy ranges. The simultaneous future observations of $\\gamma$-ray light curves at different energy ranges, e.g. GeV and TeV, will put constraints on the detailed models of $\\gamma$-ray emission from specific blazars.  We considered two different models for the maximum energies of injected electrons. In the model I, the maximum energies of electrons are fixed on 10 TeV independent on the distance from the base of the jet. In the second model, electron maximum energies are determined by their synchrotron energy losses in the local magnetic field of the jet. The electrons have clearly lower energies in the second model.  Based on our calculations, we predict a few characteristic features of the IC $e^\\pm$ pair cascade model which can be tested by the future multiwavelength GeV and  TeV $\\gamma$-ray observations. In the case of blazars inclined at relatively large angles such as Cen A, we predict the $\\gamma$-ray spectra extending through the TeV energy range even at large inclination angles provided that the blob velocities are mildly relativistic (as observed in the case of Cen A). For slowly moving blobs, the GeV $\\gamma$-ray emission at large inclination angles can be even larger than for small angles. However, the TeV $\\gamma$-ray emission behaves in the opposite way. Therefore, the model predicts even higher chances of detection of nearby blazars with blobs moving with mild velocities and inclined at large angles in respect to those once inclined at small angles.  Assuming that the injection rate of primary electrons within the blob is constant during the propagation within the jet, the $\\gamma$-ray flares last for a few hours up to a few days depending on the range of injection distances from the accretion disk. However, the $\\gamma$-ray spectrum does not change significantly with the development of the flare. It is expected that at large inclination angles the $\\gamma$-ray flares observed at the GeV energies should precede the $\\gamma$-ray flares at TeV energies. This is caused by the strong absorption of TeV $\\gamma$-rays propagating at large inclination angles at the beginning of the flare (when the relativistic electrons in the blob are still relatively  close to the accretion disk).  In the case of blazars observed at relatively small inclination angles and with fast moving blobs (e.g 3C 279), the $\\gamma$-ray emission features differ significantly. Such blazars are typically located  at cosmological distances. So then, additional effects are caused by the propagation of $\\gamma$-rays through the Extragalactic Background Light. Let us at first concentrate on the $\\gamma$-ray emission features in the source frame (before propagating in the EBL). The $\\gamma$-ray emission during the initial stage has similar intensity (the plateau phase) which is followed by an exponential decay whose time scale depends on the range of injection distances along the jet and the velocity of the blob. The $\\gamma$-ray spectrum has similar shape during the plateau phase extending to larger energies with the development of the phase. However, during the exponential decay of the flare, the $\\gamma$-ray spectrum becomes steeper and cut-offs at lower energies. Therefore we predict significantly longer $\\gamma$-ray flares at the GeV energies than at the TeV energies. The $\\gamma$-ray spectra in the observer's reference frame (after propagation in the EBL) are drastically absorbed at energies above a few hundred GeV. Therefore, the shapes of the $\\gamma$-ray spectra at different stages of the flare become more similar to each other.  In order to envisage whether the considered model parameters are reasonable, keeping in mind the limitations of the simplicity of the scenario, we compare our results with the spectral features of two blazars visible at a small and large inclination angles, Cen A and 3C~279. The GeV-TeV $\\gamma$-ray spectrum observed from Cen A  can be well described by our model in the case of the observer located at the angle $<40^{\\rm o}$ for the electron injection spectrum with the spectral index $-3$ at the range of distances between $1-300$ r$_{\\rm in}$ from the base of the jet (see Fig.~3). Note that the estimates of the inclination angle of the jet in this source lay in the range $10^{\\rm o}-80^{\\rm o}$. So then, our calculations are consistent with the lower values of the inclination angle. The spectra calculated for larger angles are clearly too steep. Unfortunately, Cen A is a relatively weak source. Therefore, the $\\gamma$-ray spectral information at a flare state and the light curves are at present unavailable. We have to wait for the next generation Cherenkov instrument (e.g. CTA) and/or a strong flare from this object in order to test some predictions of our model mentioned above (e.g. hard to soft evolution of the high energy $\\gamma$-ray spectra with developing of the flare).  We also show satisfactory fit to the GeV-TeV $\\gamma$-ray spectrum reported by the EGRET and the MAGIC telescopes from 3C~279 during the flare state and predict how these spectra should evolve in time (see Fig.~8). The spectral information from distant blazars (such as 3C~279) is partially hidden by the propagation effects in the EBL. Therefore, it will be more difficult to confront predictions for the spectra (e.g. its evolution in time) and light curves (e.g. the plateau and tail emission) in the $\\gamma$-ray range expected in terms of our model with the observations.  In the case of Cen A, we modeled the GeV-TeV $\\gamma$-ray spectrum in terms of the cascade model initiated by primary electrons (pure external Compton $e^\\pm$ pair cascade model, this work) and the cascade model initiated by primary $\\gamma$-rays (produced in the jet, e.g. in the SSC model, see \\citet{sb10}). There aren't any definitive differences between these two scenarios. Both can describe the spectral features equally well (compare Fig.~3 in this paper with Fig.~4 in \\citet{sb10}). However, the injection spectrum of primary particles have to differ significantly. Moreover, it is expected that the application of the SSC model as a source of primary $\\gamma$-rays for considered here cascade scenario may have additional consequences. For example, SSC scenario may produce additional soft radiation field which may compete with the accretion disk radiation. As a result, other component in the broad band high energy spectrum might appear and some effects on the IC $e^\\pm$ cascade spectrum produced only in the disk radiation is expected due to an additional cooling mechanism of secondary leptons. They have not been included in the cascade calculations by \\citet{sb10}.  We have discussed the $\\gamma$-ray emission features in terms of the anisotropic, time dependent, IC $e^\\pm$ pair cascade model in the case of a simple, constant lepton injection for a range of parameters. The $\\gamma$-ray light curves can change significantly in the case of a more complicated model which includes, e.g. the change of the injection rate of electrons along the jet, the extended injection of electrons at the base of the jet, or the escape of electrons from the blob. Moreover, blobs can propagate along the jet in a more complicated manner than considered in this simple scenario. For example, the blob may decelerate during its propagation along the jet and its viewing angle may also change as postulated in some works. Therefore, detailed modeling of the development of $\\gamma$-ray flares from specific blazars in such a more complicated version of the external Compton model have to wait for a more complete knowledge on the processes within the jets of AGNs.  We would like to stress an importance of the simultaneous multiwavelength observations of blazars observable at different inclination angles. They would allow a more detailed modeling of the processes responsible for the production of $\\gamma$-rays in blazars. Simultaneous GeV and TeV observations of flares from blazars can efficiently limit the acceleration region and the dynamics of the injection rate."
    },
    "1004/1004.4617_arXiv.txt": {
        "abstract": "{X-rays from very young stars are powerful probes to investigate the mechanisms at work in the very first stages of the star formation and the origin of X-ray emission in very young stars.} {  We present results from a 500 ks long observation of the Rho Ophiuchi cloud with a { XMM-Newton} large program named DROXO, aiming at studying the X-ray emission of deeply embedded young stellar objects (YSOs).} {The data acquired during the DROXO program were reduced with SAS software, and filtered in time and energy to improve the signal to noise of detected sources; light curves and spectra were obtained.} { We detected $111$ sources, 61 of them associated with $\\rho$\\,Ophiuchi YSOs as identified from infrared observations with ISOCAM.  Specifically, we detected 9 out of 11 Class\\,I objects, 31 out of 48 Class\\,II and 15 out 16 Class\\,III objects. Six objects out of 21 classified Class\\,III candidates are also detected. At the same time we suggest that $15$ Class\\,III candidates that remain undetected at $\\log{L_{\\rm x}}\\,{\\rm [erg/s]} < 28.3$ are not related to the cloud population. The global detection rate is $\\sim64\\%$. We have achieved a flux sensitivity of $\\sim 5\\cdot 10^{-15}$ erg s$^{-1}$ cm$^{-2}$. The $L_\\mathrm X$ to $L_\\mathrm{bol}$ ratio shows saturation at a value of $\\sim-3.5$ for stars with T$_\\mathrm{eff} \\le 5000$ K  or 0.7 M$\\odot$ as observed in the Orion Nebula. The plasma temperatures and the spectrum absorption show a decline with YSO class, with Class I YSOs being hotter and more absorbed than Class II and III YSOs. In one star (GY~266) with infrared counterpart in 2MASS and Spitzer catalogs we have detected a soft excess in the X-ray spectrum, which is best fitted by a cold thermal component less absorbed than the main thermal component of the plasma. This soft component hints at plasma heated by shocks due to jets outside the dense circumstellar material.} {} ",
        "introduction": "{ Low-mass stars in the pre-Main Sequence (PMS) phase are characterized by intense X-ray emission. X-ray activity has not been yet firmly confirmed in the very initial protostellar cores that accrete material from a surrounding envelope (Class 0 objects), {  because X-rays are either} absent or completely absorbed by the dense circumstellar material. In Class I, II, and III {  objects} and until the Zero Age Main Sequence stage the X-ray emission is very strong when compared to solar-age stars and also characterized by impulsive variability \\citep{Feigelson99}. In the case of Class III objects all X-rays are thought to originate from a magnetized stellar corona,  which likely resembles a scaled-up version of the corona of late-type Main Sequence stars and the Sun. {  Additional X-ray production mechanisms are possibly at work in} Class I and II objects. {  At these young evolutionary stages} the magnetic field drives the material in-falling from the disk to the stellar surface {  where the matter becomes shocked}. Furthermore, outflows along the disk axis may interact with the circumstellar material and form shocks. {  The plasma in the shocks in both in- and outflows may become heated} up to a few million degrees, thus emitting X-rays \\citep{Pravdo01,Favata02,Kastner05, Giardino07, Sacco08}.  Significant effort has been devoted to the X-ray observation of star-forming regions, which provide natural laboratories where to study the X-ray emission from young stars and its implications for the mechanisms of star formation. The $\\rho$ Ophiuchi cloud is among the nearest star-forming regions (120 pc, \\citealp{Loinard08}) and has been extensively studied from the infrared (IR) to the X-rays bands. The $\\rho$ Ophiuchi cloud is shaped as a dense multicore structure, hosting more than 200 members comprising young stellar objects (YSOs) in all evolutionary stages from Class 0 to Class III \\citep{Wilking08}. While a dispersed population of optically visible stars associated with the Upper Scorpius OB association with an age of $\\sim 5$ Myr  is present in the region, the studies in the IR band have revealed a younger population  of Pre Main Sequence (PMS) stars and protostellar objects with ages of only 0.3--1 Myr \\citep{LR99}, which are thus younger than YSOs in other star-forming regions like the Taurus Molecular Cloud ($\\sim 1-5$ Myr) and the Orion Nebula ($\\sim 1$ Myr, \\citealp{Hillenbrand97}). {  The mid-IR survey with ISOCAM on board the ISO satellite} reported by \\citet{Bontemps01} (hereafter Bo01) allowed the detection of a population of 16 Class I, 123 Class II and 77 Class III YSOs, adding 71 previously unknown members of Classes I and II. X-ray studies carried out with the {\\em Einstein} and {\\em ROSAT} satellites had revealed several tens of embedded Class I, II, and III stars that are highly active in X-rays and are characterized by a very strong time variability \\citep{Montmerle83,Casanova95}. \\citet{Casanova95} tentatively identified several X-ray sources with Class I protostars with ROSAT/PSPC observations. However, X-ray observations with the ROSAT/High-Resolution Imager were needed to confirm X-ray emission from Class I protostars \\citep{Grosso97,Grosso01}. Recent observations with the XMM-Newton \\citep{Ozawa05} and Chandra satellites \\citep{Imanishi01,Flaccomio03} have increased the number of candidate PMS members of the region. These studies suggested that accreting YSOs have increasing X-ray activity going from Class~I to Class~III and decreasing plasma temperatures. The aim of the {\\it Deep Rho Ophiuchi XMM-Newton Observation} (DROXO) was to obtain a high-sensitivity survey in the Core F of $\\rho$~Ophiuchi Cloud by means of a long, almost continuous observation. % We report the data analysis and the X-ray properties of the PMS stars observed during the DROXO program. The structure of the paper is as follows: Sects. 2 and 3  describe the observation and the data analysis, Sects. 4 and 5 describe the sensitivity of the survey and the nature of the sources detected in DROXO. In Sect. 6 we discuss the X-ray emission of the sample of classified YSOs, in particular we explore the behavior of plasma temperatures, absorption, X-ray luminosity and ratio L$_\\mathrm X/$L$_\\mathrm{bol}$ with respect to effective temperatures, masses and evolutionary stages. In Sect. 7 we give a summary. Appendix A and B contain tables with the list of detected sources, the best-fit parameters of spectra, the upper limits to rate for undetected YSOs in the ISOCAM sample, the Spitzer counterparts to DROXO sources, and an example of one page of the atlas (published electronically only) with the EPIC spectrum and light curve of each source. } ",
        "conclusions": "{   We described the {\\em Deep Rho Ophiuchi XMM-Newton Observation} (DROXO), aimed at exploring at high sensitivity the X-ray emission of YSOs in the $\\rho$ Ophiuchi core F region. We detected 111 X-rays sources, and for 91 of them we obtained a model fit to their X-ray spectra. By using optical and IR data we estimated the photospheric parameters of most of the sample of X-ray detected YSOs.  We find 10 unidentified sources and a further 6 sources already detected in previous X-ray surveys, but without optical/IR counterpart. Three of them (src. num. 45, 48 and 95) show light curves with some impulsive variability. The spectra of src. num. 45 and 48 have a good fit with an absorbed power law, while src. num. 95 has a spectrum compatible with a thermal model. The lack of 2MASS counterparts could be a hint that they are of extragalactic nature, but, given the sensitivity of 2MASS catalog, it is not possible to rule out they are very low mass PMS stars or even brown dwarfs, especially src. num. 95.  The sensitivity of the survey is $f_\\mathrm X \\sim 5\\cdot10^{-15}$ erg s$^{-1}$ cm$^{-2}$ in flux and log $L_\\mathrm{X}$ (erg s$^{-1}) \\sim 27.9$ in luminosity, but it strongly depends on the local absorption, which is largely variable in the field-of-view. The sample of 96 classified YSOs given by Bo01 in the DROXO field of view has allowed us to explore the X-ray emission from Class I, II, and III YSOs of $\\rho$~Ophiuchi. The X-ray properties of $\\rho$~Ophiuchi PMS stars obtained from DROXO were compared with those of the Orion Nebula Cloud PMS stars studied in the COUP survey. When fitting the relation between mass and X-ray luminosity with a power law, the index that we find for DROXO sample is quite similar to that found for ONC and TMC \\citep{Preibisch05,Telleschi07}. With respect to COUP we find a lower normalization, suggesting that on average the PMS stars in the range of masses between 0.1--2 $M_\\odot$ in our sample are less luminous than their analogs in COUP. Young stellar objects in $\\rho$~Ophiuchi exhibit a saturation of the ratio L$_\\mathrm X / $L$_\\mathrm{bol}$ near the ``canonical'' value of  $\\sim 10^{-3}$ for masses between 0.5 and 1 M$_\\odot$, while stars with masses below 0.5 M$\\odot$ show a lower limit of saturation L$_\\mathrm X / $L$_\\mathrm{bol}\\sim10^{-3.5}$.  This is consistent with what was found by \\citet{Preibisch05} for the COUP sample.  We detected a large fraction of YSOs in the field of view, 94\\% of which are bona fide Class III stars, 77\\% of Class~II and 82\\% of Class~I YSOs, respectively. We confirm the high detection rate among Class~I YSOs found by \\citet{Imanishi01}. The detection rate in our Class I sample is higher than the analog rate found in the COUP survey by \\citet{Prisinzano08} (82\\% vs. 62\\%) despite the higher sensitivity of COUP with respect to DROXO. \\citet{Prisinzano08} make a distinction between Class Ia (characterized by a rising SED from $K$ to 8$\\mu$m) and Class Ib (rising SED up to 4.5$\\mu$m). The detection rates in these two subclasses are 44\\% and 82\\% respectively. X-ray luminosities of Class Ib objects are higher than those of Class Ia objects. The rate of X-ray detection of Class~Ib objects in Orion is identical to that of our sample of Class~I objects in $\\rho$~Ophiuchi. With support of Spitzer photometry we suggest that only a small fraction of our Class I stars are deeply embedded objects. From X-ray data the high detection rate suggests that our Class I sample  could be mainly formed by Class Ib YSOs as defined by Prisinzano et al. (2008). These objects in Orion have X-ray luminosities more similar to Class II objects. In $\\rho$~Oph we find that our nine detected Class I stars are on average more luminous than Class II stars.  \\citet{Bontemps01} list 39 stars defined as Class III candidates, 21 of which are surveyed in DROXO. Six out of 21 are detected in DROXO, providing  support for their PMS nature. The other 15 are undetected below $\\log L_\\mathrm X = 28.3$.  Likely they are not PMS stars and thus not members of the $\\rho$~Oph cloud. By means of XLFs we evaluated the level of emission of Class~I, II, and III YSOs. After excluding the suspect Class~III YSOs, we find that  Class I YSOs emit at the same level of Class III YSOs (median $\\log L_\\mathrm X \\sim {29.6}$ % and {29.7}, respectively) and both samples are more luminous than Class II YSOs  (median $\\log L_\\mathrm X$ $ \\leq {28.8}$). % In COUP Class III objects are more luminous than Class I and II for masses between 0.5 and 1.2 M$\\odot$, while for masses in 0.1-0.5 M$\\odot$ Class Ib sources are slightly more luminous than Class II and III (see Fig. 8 in \\citealp{Prisinzano08}).  The analysis of X-ray spectra leads us to conclude that the mean plasma temperatures of Class~I protostars are higher than in Class~II and Class~III YSOs. We find a clear trend of decreasing plasma temperatures passing from Class I to Class III objects, differently from what is observed in ONC where \\citet{Prisinzano08} do not find a significant evolution of temperatures from Class~I to Class~III objects. Four Class III stars and one Class II star  (GY~112, SR~12A, GY~296, GY~380 and GY~3) have spectra noticeably softer than the rest of YSOs with mean temperatures around 0.7-0.8 keV.  The absorption derived from X-ray spectra is higher in Class~I than in Class~II and III. Given the high absorption among Class~I YSOs, it is impossible to determine if a soft thermal component in their spectra is completely absorbed or is not present at all. This soft component would reduce the intrinsic $kT$ for Class\\,I objects, making them more similar to the ones in the ONC. However, higher absorption should lead to an underestimate of the luminosities of heavily absorbed spectra like those of Class~I objects, while we observe a luminosity for them higher than that of in Class~II and Class~III YSOs. Furthermore, Prisinzano et al. (2008) have found in extensive simulations that varying $N_\\mathrm H$ does not significantly influence the XLF.   The star GY~266 shows a peculiar X-ray spectrum composed by two thermal components that are differently absorbed. The hot and heavily absorbed component is similar to that found in other PMS stars, while the soft one, less absorbed, could arise from unresolved jet shocks."
    },
    "1004/1004.3347_arXiv.txt": {
        "abstract": "We study the small population of high-redshift ($z_\\mathrm{em}>2.7$) quasars detected by GALEX, whose far-UV emission is not extinguished by intervening \\ion{H}{1} Lyman limit systems. These quasars are of particular importance to detect intergalactic \\ion{He}{2} absorption along their sightlines. We correlate almost all verified $z_\\mathrm{em}>2.7$ quasars to the GALEX GR4 source catalog covering $\\sim 25000$\\,deg$^2$, yielding 304 sources detected at S/N$>3$. However, $\\sim 50\\%$ of these are only detected in the GALEX NUV band, signaling the truncation of the FUV flux by low-redshift optically thick Lyman limit systems. We exploit the GALEX UV color $m_\\mathrm{FUV}-m_\\mathrm{NUV}$ to cull the most promising targets for follow-up studies, with blue (red) GALEX colors indicating transparent (opaque) sightlines. Extensive Monte Carlo simulations indicate a \\ion{He}{2} detection rate of $\\sim 60$\\% for quasars with $m_\\mathrm{FUV}-m_\\mathrm{NUV}\\la 1$ at $z_\\mathrm{em}\\la 3.5$, a $\\sim$50\\% increase over GALEX searches that do not include color information. We regard 52 quasars detected at S/N$>3$ to be most promising for HST follow-up, with an additional 114 quasars if we consider S/N$>2$ detections in the FUV. Combining the statistical properties of \\ion{H}{1} absorbers with the SDSS quasar luminosity function, we predict a large all-sky population of $\\sim 200$ quasars with $z_\\mathrm{em}>2.7$ and $i\\la 19$ that should be detectable at the \\ion{He}{2} edge at $m_\\mathrm{304}<21$. However, SDSS provides just half of the NUV-bright quasars that should have been detected by SDSS \\& GALEX. With mock quasar photometry we revise the SDSS quasar selection function, finding that SDSS systematically misses quasars with blue $u-g\\la 2$ colors at $3\\la z_\\mathrm{em}\\la 3.5$ due to overlap with the stellar locus in color space. Our color-dependent SDSS selection function naturally explains the inhomogeneous $u-g$ color distribution of SDSS DR7 quasars as a function of redshift and the color difference between color-selected and radio-selected SDSS quasars. Moreover, it yields excellent agreement between the observed and the predicted number of GALEX UV-bright SDSS quasars. We confirm our previous claims that SDSS preferentially selects $3\\la z_\\mathrm{em}\\la 3.5$ quasars with intervening \\ion{H}{1} Lyman limit systems. Our results imply that broadband optical color surveys for $3\\la z_\\mathrm{em}\\la 3.5$ quasars have likely underestimated their space density by selecting IGM sightlines with an excess of strong \\ion{H}{1} absorbers. ",
        "introduction": "The intergalactic space is pervaded by a filamentary cosmic web of gas of almost primordial composition, the so-called intergalactic medium (IGM), seen in absorption against background sources \\citep{rauch98,meiksin09}. The absence of \\ion{H}{1} Ly$\\alpha$ absorption troughs in spectra of $z_\\mathrm{em}<6$ quasars signals that the hydrogen in the IGM is highly ionized \\citep{gunn65}. Instead, the plethora of narrow \\ion{H}{1} Ly$\\alpha$ absorption lines, known as the Ly$\\alpha$ forest, traces the tiny residual neutral hydrogen fraction of the IGM as the largest reservoir of baryons in the universe. The ionizing radiation of quasars and star-forming galaxies is filtered by the IGM, leading to the buildup of the UV background radiation field that determines the ionization state of the gas \\citep{haardt96,fardal98,faucher09}. The UV background changes in amplitude and spectral shape due to evolution in the source number density, cosmological expansion and structure formation \\citep[e.g.][]{dave99}. This is particularly important for the ionization state of helium, the second most abundant element in the IGM. Due to its 5.4 times higher recombination rate and 4 times higher ionization threshold, the reionization epoch of helium (\\ion{He}{2}$\\longrightarrow$\\ion{He}{3}) is expected to be delayed with respect to hydrogen.  The Ly$\\alpha$ transition of intergalactic \\ion{He}{2} at $\\lambda_\\mathrm{rest}=303.78$\\AA\\ is observable in the far UV (FUV) from space only at $z>2$ due to the Galactic Lyman limit. The determination of the \\ion{He}{2} reionization epoch via the \\ion{He}{2} Gunn-Peterson test towards high-redshift quasars has been a major goal in extragalactic UV astronomy since the launch of the Hubble Space Telescope \\citep[HST, e.g.][]{miralda-escude90,miralda-escude93}. However, the accumulated Lyman continuum (LyC) absorption of the \\ion{H}{1} absorber population severely attenuates the quasar flux in the FUV, rendering just a few percent of $z_\\mathrm{em}>3$ sightlines to be relatively transparent \\citep{moller90}. The combination of the rising LyC absorption and the declining quasar luminosity function results in a sharply dropping number of observable UV-bright quasars at $z_\\mathrm{em}>3$ \\citep{picard93,jakobsen98}.  Until very recently \\ion{He}{2} Ly$\\alpha$ absorption had been found only in a handful of sightlines despite considerable effort, since the UV fluxes of most targeted quasars had been unknown. HST observations of Q~0302$-$003 at $z_\\mathrm{em}=3.285$ \\citep{jakobsen94,hogan97,heap00} and PKS~1935$-$692 at $z_\\mathrm{em}=3.18$ \\citep{anderson99} revealed a high \\ion{He}{2} effective optical depth at $z\\ga 3$ that is consistent with a Gunn-Peterson trough ($\\tau_\\mathrm{eff,He\\,II}>3$). In contrast, the lines of sight towards HS~1700$+$6416 at $z_\\mathrm{em}=2.736$ \\citep{davidsen96,fechner06}, HE~2347$-$4342 at $z_\\mathrm{em}=2.885$ \\citep{reimers97,kriss01,smette02,zheng04,shull04} and HS~1157$+$3143 at $z_\\mathrm{em}=2.989$ \\citep{reimers05} show patchy \\ion{He}{2} absorption with voids ($\\tau_\\mathrm{eff,He\\,II}<1$) and troughs ($\\tau_\\mathrm{eff,He\\,II}>3$). At $z\\la 2.7$ this patchy absorption evolves into a \\ion{He}{2} Ly$\\alpha$ forest that has been resolved in high-resolution spectra obtained with the Far Ultraviolet Spectroscopic Explorer \\citep[FUSE, ][]{kriss01,zheng04,shull04,fechner06}.  The strong evolution of the \\ion{He}{2} absorption suggests a late reionization epoch of helium at $z\\sim 3$, when quasars have been sufficiently abundant to supply the required hard photons. The patch-work of absorption and transmission evokes a picture of overlapping \\ion{He}{3} zones around quasars that lie close to the sightline \\citep{reimers97,heap00,smette02}. Indeed, the \\ion{He}{3} proximity zones of quasars have been detected both along the line of sight \\citep{hogan97,anderson99} and in transverse direction \\citep{jakobsen03}. In the past few years, great progress has been made in developing the theoretical framework to interpret these observations. Both semi-analytic \\citep[e.g.][]{haardt96,fardal98,gleser05,furlanetto10} and numerical radiative transfer simulations \\citep{maselli05,tittley07,paschos07,mcquinn09a} indicate that the \\ion{He}{2} reionization process should be very inhomogeneous and extended over $3\\la z\\la 4$, since rare luminous quasars dominate the photoionizing budget of the overall quasar population. The few quasars contributing to the UV radiation field at the \\ion{He}{2} ionization edge at a given point likely give rise to fluctuations in the FUV background that can be tracked by the co-spatial absorption of \\ion{He}{2} and \\ion{H}{1} \\citep{bolton06,worseck06,worseck07,furlanetto09}. The UV background hardens as \\ion{He}{2} reionization proceeds \\citep{heap00,zheng04}, but $\\ga 10$~Mpc fluctuations are expected to persist even after its end \\citep{fechner07}.  Other, more indirect observations might suggest that \\ion{He}{2} reionization is ending at $z\\sim 3$. The IGM is reheated as the individual \\ion{He}{3} bubbles around quasars overlap, however the amplitude of this temperature jump is highly uncertain \\citep{bolton09,bolton09b,mcquinn09a}. Observationally, several studies indicated a jump in the IGM temperature at $z\\sim 3$ \\citep{ricotti00,schaye00,theuns02b}, whereas others are consistent with an almost constant IGM temperature at $2\\la z\\la 4$ \\citep{mcdonald01b,lidz10}. Moreover, photoionization models of metal line systems indicate a significant hardening of the UV background at $z\\la 3$ \\citep{agafonova05,agafonova07}. However, these observations are restricted to rare metal line systems showing various ions with a simple velocity structure.  At present, the five \\ion{He}{2} absorption sightlines studied at scientifically useful spectral resolution provide the best observational constraints on \\ion{He}{2} reionization. However, just one or two sightlines probe the same redshift range, and given the large predicted variance in the \\ion{He}{2} absorption, this small sample clearly limits our current understanding of \\ion{He}{2} reionization\\footnote{Ironically, the $z\\sim 6$ epoch has substantially better statistics.}. The Sloan Digital Sky Survey (SDSS) has dramatically increased the number of high-redshift quasars to search for the presence of flux at \\ion{He}{2} Ly$\\alpha$, yielding three $z_\\mathrm{em}>3.5$ quasars with detected \\ion{He}{2} Gunn-Peterson troughs \\citep{zheng04b,zheng05,zheng08}. More importantly, the almost completed first UV all-sky survey with the Galaxy Evolution Explorer (GALEX) enables the pre-selection of UV-bright quasars for follow-up UV spectroscopy, leading to the recent discovery of 22 new clear sightlines towards SDSS quasars at $3.1<z_\\mathrm{em}<3.9$ \\citep{syphers09b,syphers09a}. The available GALEX photometry dramatically increases the survey efficiency by almost an order of magnitude to $\\simeq 42$\\% in the \\citeauthor{syphers09b} survey.  The recently installed Cosmic Origins Spectrograph (COS) on HST offers unprecedented sensitivity to study \\ion{He}{2} reionization via \\ion{He}{2} Ly$\\alpha$ absorption spectra. With its confirmed throughput at $\\lambda>1105$\\AA\\ \\citep{mccandliss10} HST/COS is now able to probe \\ion{He}{2} Ly$\\alpha$ at $z>2.64$, thereby covering the full redshift range of interest for \\ion{He}{2} reionization. Very recently, \\citet{shull10} presented a high-quality COS spectrum of HE~2347$-$4342, dramatically improving on earlier FUSE data. In the near future, COS will be employed to both obtain follow-up spectroscopy of the recently confirmed \\ion{He}{2} sightlines, and to discover new ones. In this paper we introduce the quasar UV color measured by GALEX as a powerful discriminator to select the most promising sightlines for follow-up spectroscopy. Moreover, we significantly improve on earlier predictions on the number of UV-bright quasars \\citep{picard93,jakobsen98}, based on observational advances to characterize both the quasar luminosity function and the optically thick IGM absorber distribution. The structure of the paper is as follows: In \\S\\ref{galexsample} we will present our sample of verified high-redshift quasars detected by GALEX. Section \\ref{mcspectra} describes our Monte Carlo routine to compute \\ion{H}{1} absorption spectra and to perform mock GALEX and SDSS photometry. In \\S\\ref{results} we determine the expected number of UV-bright $z_\\mathrm{em}>2.7$ quasars and establish GALEX UV color selection criteria to select quasars with probable \\ion{He}{2}-transparent sightlines. We compare the observed and predicted number counts of UV-bright SDSS quasars in \\S\\ref{sdssgalex} before concluding in \\S\\ref{conclusions}. ",
        "conclusions": "\\label{conclusions}  We have correlated verified $z_\\mathrm{em}>2.7$ quasars to GALEX photometry to reveal the rare high-redshift quasars whose far-UV fluxes are not extinguished by intervening Lyman limit systems, with the goal to establish a sample of UV-bright quasars that likely show intergalactic \\ion{He}{2} absorption. We have used the GALEX UV color $m_\\mathrm{FUV}-m_\\mathrm{NUV}$ to cull the most promising targets for follow-up. Red UV colors indicate that the quasar flux is prematurely truncated redward of the \\ion{He}{2} edge, whereas the rare quasars with blue UV colors and significant FUV flux will likely show flux at \\ion{He}{2} Ly$\\alpha$.  We have performed extensive Monte Carlo simulations to estimate the UV color distribution of UV-bright quasars and their surface density on the sky. We predict that $\\sim 600$ ($\\sim 200$) quasars with $z_\\mathrm{em}>2.7$ and $m_{1450}<19$ should be detectable in the NUV (FUV) at $m_\\mathrm{UV}<21$ (Fig.~\\ref{allskycounts}; without considering sources near the Galactic plane). The number of UV-bright quasars strongly declines with redshift due to the declining quasar space density and the increasing \\ion{H}{1} Lyman continuum absorption experienced at the \\ion{He}{2} edge (Fig.~\\ref{taulycstat}). Nevertheless, there are enough targets within reach of HST/COS to significantly constrain \\ion{He}{2} reionization by \\ion{He}{2} absorption spectra, provided that the quasars are known and have been imaged with GALEX for efficient pre-selection.  Most confirmed $z_\\mathrm{em}<3.5$ \\ion{He}{2} quasars have blue UV colors and our simulations indicate a $\\sim 60$\\% \\ion{He}{2} detection rate of quasars at similar UV color (Fig.~\\ref{gr4zcolor}), a $\\sim$50\\% increase over approaches that do not include color information \\citep{syphers09b,syphers09a}. We have identified 166 additional quasars as prime targets for UV follow-up spectroscopy with HST/COS to significantly extend the sample of \\ion{He}{2} sightlines before the end of HST's mission (Tables~\\ref{candlist1} and \\ref{candlist2}). We have started a survey with HST/COS in Cycle~17 to obtain FUV follow-up spectra of 8 UV-bright quasars selected from the much smaller GALEX GR3 footprint (star symbols in Fig.~\\ref{gr4zcolor}).  We have reassessed the SDSS color selection efficiency by applying the SDSS quasar selection criteria to mock photometry of our Monte Carlo spectra. We find that SDSS preferentially misses UV-bright quasars due to their blue colors that make them indistinguishable from main-sequence stars (Figs.~\\ref{sdssselfunc} and \\ref{sdss_u_minus_g}). The observed $u-g$ colors of color-selected SDSS quasars are significantly redder than those of radio-selected ones at $3\\la z_\\mathrm{em}\\la 3.5$, and agree well with our color-dependent SDSS selection function (Fig.~\\ref{sdss_u_minus_g}). These missing quasars lack strong Lyman continuum absorption due to Lyman limit systems along their lines of sight that would redden the $u-g$ color (Fig.~\\ref{sdssigmtrans}).  The SDSS color bias has not been well studied previously. Figure~18 of \\citet{bernardi03} reveals that SDSS rarely selected blue quasars at $3.2<z_\\mathrm{em}<3.6$, but the authors did not investigate this further. \\citet{richards06} explored whether primarily radio-selected SDSS quasars have different color selection efficiencies, but due to low number statistics, they regarded the differences to be insignificant (their Fig.~10). For the first time we have been able to demonstrate the full effect and its consequences.  Since the UV-brightest quasars are among the bluest in SDSS $u-g$ at all epochs, we conclude that SDSS is inefficient in finding further promising targets for detecting intergalactic helium. Although about two dozen quasars in the SDSS database already have been confirmed to show \\ion{He}{2} \\citep{syphers09b,syphers09a}, we predict that the FUV-brightest quasars without strong Lyman continuum absorption are insensitive to standard color selection techniques.  Due to the restrictive SDSS selection criteria the statistics of high-column density IGM absorbers measured towards color-selected SDSS quasars will be biased high \\citep{prochaska09,prochaska10}. Our results also indicate that the incidence of DLAs based on SDSS samples \\citep[e.g][]{prochaska05} has been overestimated. At $3\\la z_\\mathrm{em}\\la 3.5$ the few radio-selected quasars are probably the only ones within SDSS that are truly unbiased in the statistics of high-column density absorbers.  The Lyman limit system bias will also affect the frequency of metal absorption lines that primarily occur at moderate to high \\ion{H}{1} column densities. Because these absorbers trace the large-scale structure of the IGM, the Lyman limit system bias might also impact analyses of the clustering properties and the power spectrum of the Ly$\\alpha$ forest. Presumably all $z_\\mathrm{em}\\ga 3$ quasars that have been first detected in broadband color surveys are affected by such a bias to some sort, depending on the exact color selection criteria and the number and effective wavelengths of the employed filters. Because the abundance of optically thick absorbers is far less constrained in reality than in 'completeness' simulations like ours, determinations of the optical quasar luminosity function at $3<z_\\mathrm{em}<3.5$ should invoke a variety of IGM models which will result in large systematic uncertainties in the luminosity function. Our study indicates that results based on color-selected high-redshift quasar samples are not as easy to interpret as previously thought due to the intertwined demographics of strong IGM absorbers and their background candles."
    },
    "1004/1004.0626.txt": {
        "abstract": "{}{}{}{}{} % 5 {} token are mandatory  \\abstract % context heading (optional) % {} leave it empty if necessary {} % aims heading (mandatory) {We have performed a comprehensive multiwavelength analysis of a sample of 20 starburst galaxies that show  a substantial population of very young massive stars, most of them classified as Wolf-Rayet (WR) galaxies. In this paper, the forth of the series, we present the global analysis of the derived photometric and chemical properties.} % methods heading (mandatory) {We compare optical/\\NIR\\ colours and the physical properties (reddening coefficient, equivalent widths of the emission and underlying absorption lines, ionization degree, electron density, and electron temperature) and chemical properties (oxygen abundances and N/O, S/O, Ne/O, Ar/O, and Fe/O ratios) with previous observations and galaxy evolution models. %In total, We compile 41 independent star-forming regions --with oxygen abundances between \\mbox{\\abox= 7.58} and 8.75--, of which 31 have a direct estimate of the electron temperature of the ionized gas.} %, and hence their element abundances were derived using the direct method.} % results heading (mandatory) %{We checked that all objects can be classified as pure star-forming galaxies. {According to their absolute $B$-magnitude, many of them are not dwarf galaxies, but they should be during their quiescent phase. We found that both \\CHb\\ and \\Wabs\\ increase with increasing metallicity. %The abundance ratios agree with previous observations. The differences in the N/O ratio is explained assuming differences in the star formation histories. We detected a high N/O ratio in objects showing strong WR features (HCG~31~AC, UM~420, IRAS~0828+2816, III~Zw~107, ESO~566-8 and NGC~5253). The ejecta of the WR stars may be the origin of the N enrichment in these galaxies. We compared the abundances provided by the direct method with those obtained through empirical calibrations, finding that (i) the Pilyugin method %(Pilyugin 2001a,b; Pilyugin \\& Thuan 2005) %\\citep{P01a,P01b,PT05} %, which consider the $R_{23}$ and the $P$ parameters, is the best suited empirical calibration for these star-forming galaxies, (ii) the relations %the relations between the oxygen abundance and the $N_2$ or the $O_3N_2$ parameters provided by Pettini \\& Pagel (2004) give acceptable results for objects with \\abox$>$8.0, and (iii) the results provided by empirical calibrations based on photoionization models %(McGaugh, 1991; Kewley \\& Dopita, 2002; Kobulnicky \\& Kewley, 2004) are systematically 0.2 -- 0.3 dex higher than the values derived from the direct method. %These differences are of the same order that the abundance discrepancy found between %recombination and collisionally excited lines. This may suggest the existence of temperature %fluctuations in the ionized gas, as exists in Galactic and other extragalactic \\HII regions. The O and N abundances and the N/O ratios are clearly related to the optical/\\NIR\\ luminosity; the dispersion of the data is a consequence of the differences in the star-formation histories. The $L$--$Z$ relations tend to be tighter when using \\NIR\\ luminosities, which facilitates distinguishing tidal dwarf galaxies candidates and pre-existing dwarf objects. Galaxies with redder colours tend to have higher oxygen and nitrogen abundances.} % conclusions heading (optional), leave it empty if necessary {Our detailed analysis is fundamental to understand the nature of galaxies that show strong starbursts, as well as to know their star formation history and the relationships with the environment. This study is complementary --but usually more powerful-- to the less detailed analysis of large galaxy samples that are very common nowadays.}  %  {We consider that it is fundamental to perform a detailed analysis %of both the photometric and the chemical properties of the star-forming galaxies %to understand the evolutionary stage of each system. Such study is absolutely %needed to globally analyse the properties of larger galaxy samples.}      \\titlerunning{Massive star formation in Wolf-Rayet galaxies IV: Colours and metallicities}  \\authorrunning{L\\'opez-S\\'anchez \\& Esteban} ",
        "introduction": "%*** copy-paste paper II  The knowledge of the chemical composition of galaxies, in particular of dwarf galaxies, is vital for understanding their evolution, star formation history, stellar nucleosynthesis, the importance of gas inflow and outflow, and the enrichment of the intergalactic medium. Indeed, metallicity is a key ingredient for modelling galaxy properties, because it determines \\UV, optical and \\NIR\\ \\mbox{colours} at a given age (i.e., Leitherer et al. 1999), nucleosynthetic yields (e.g., Woosley \\& Weaver 1995), the dust-to-gas ratio (e.g., Hirashita et al 2001), the shape of the interstellar extinction curve (e.g., Piovan et al. 2006), or even the properties of the Wolf-Rayet stars \\citep{Crowther07}.  The most robust method to derive the metallicity in star-forming and starburst galaxies is via the estimate of metal abundances and abundance ratios, in particular through the determination of the gas-phase oxygen abundance and the nitrogen-to-oxygen ratio. The relationships between current metallicity and other galaxy parameters, such as colours, luminosity, neutral gas content, star-formation rate, extinction or total mass, constrain galaxy-evolution models and give clues about the current stage of a galaxy. For example, is still debated whether massive star formation results in the instantaneous enrichment of the interstellar medium of a dwarf galaxy, or if the bulk of the newly synthesized heavy elements must cool before becoming part of the interstellar medium (ISM) that eventually will form the next generation of stars. Accurate oxygen abundance measurements of several \\HII regions within a dwarf galaxy will increase the understanding of its chemical enrichment and mixing of enriched material. The analysis of the kinematics of the ionized gas will also help to understand the dynamic stage of galaxies and reveal recent interaction features. Furthermore, detailed analyses of starburst galaxies in the nearby Universe are fundamental to interpret the observations of high-z star forming galaxies, such as Lyman Break Galaxies \\citep{EP03}, as well as quantify the importance of interactions in the triggering of the star-formation bursts, which seem to be very common at higher redshifts (i.e., Kauffmann \\& White 1993; Springer et al. 2005).  %In addition to providing information about previous episodes of star %formation and mass-loss, metallicity may provide clues into the dynamical %effects of interactions on gas flows in galaxy pairs. %Isolated disk galaxies often display metallicity gradients %%(see Shields 1990; Dinerstein 1996, for a review) %which may be altered during a merger as a result of tidally-driven gas flows. %Metallicity gradients in interacting galaxies could potentially provide a unique %method for finding interacting galaxies that have experienced strong tidal gas flows.  The comparison of the metallicity (which reflects the gas reprocessed by stars and any exchange of gas between the galaxy and its environment) with the stellar mass (which reflects the amount of gas locked up into stars) provides key clues about galaxy formation and evolution.  These analyses have shown a clear correlation between mass and metallicity. In practice, luminosity has been used as substitute of mass because of the difficulty of deriving reliable galaxy masses, yielding to the so-called metallicity-luminosity relation (i.e., Rubin et al. 1984; Richer \\& McCall 1995; Salzer et al. 2005), although in recent years mass-metallicity relations are also explored (i.e., Tremonti et al. 2004; Kewley \\& Elisson, 2008), and are studied even at high redshifts (i.e., Kobulnicky et al. 1999; Pettini et al. 2001; Kobulnicky \\& Kewley 2004; Erb et al. 2006; Liang et al. 2006). The evolution of such relationships are now predicted by semi-analytic models of galaxy formation within the $\\Lambda$-cold dark matter framework that include chemical hydrodynamic simulations (De Lucia et al. 2004; Tissera et al. 2005; De Rossi et al. 2006; Dav\\'e \\& Oppenheimer 2007). Ironically, today the main problem is not to estimate the mass of a galaxy but its real metallicity, so that different methods involving direct estimates of the oxygen abundance, empirical calibrations using bright emission-line ratios or theoretical methods based on photoionization models yield very different values (i.e., Yin et al. 2007; Kewley \\& Elisson, 2008).  Hence precise photometric and spectroscopic data, including a detailed analysis of each particular galaxy that allows conclusions about its nature, are crucial to address these issues. We performed such a detailed photometric and spectroscopic study in a sample of strong star-forming galaxies, many of them previously classified as dwarf galaxies. The majority of these objects are Wolf-Rayet (WR) galaxies, a very inhomogeneous class of star-forming objects which share at least %have as minimum common property an ongoing or recent star formation event that has produced stars sufficiently massive to evolve into the WR stage \\citep{SCP99}. %The definition of WR galaxy hides a loose concept %because of the wide variety of morphological types in which they are found. %Indeed, %%The term WR galaxy must be used with caution. %depending on the distance of the object and size of the %area analysed, the region of concern may be a single extragalactic %\\HII  region with a few WR stars in a galaxy, a massive star %cluster or the nucleus of a powerful starburst galaxy harbouring %numerous massive stars \\citep*{SCP99}. However, WR features in the spectra of a galaxy provides useful information about the star-formation processes in the system. As the first WR stars typically appear around 2 -- 3 Myr after the starburst is initiated and disappear within some 5 Myr \\citep{MeynetMaeder05}, their detection gives indications about both the youth and strength of the burst, offering the opportunity to study an approximately coeval sample of very young starbursts \\citep{SV98}. %WR galaxies also are ideal targets to study %the properties and numbers of O and WR stars at sub-solar metallicities, %giving strong constrains on the characteristics of the massive stars and the massive %end of the initial mass function (\\IMF) and their dependence on metallicity.  %the number of O and WR stars, the WN/WC ratio %and place strong constraints on the massive end of the initial mass function (\\IMF). %These are fundamental ingredients to build satisfactory population synthesis models %that reproduce the observed properties of unresolved stellar populations.  %Many WR galaxies were serendipitously discovered. The first compilation %was done by \\citet{C91} who listed 37 objects. \\citet*{SCP99} increased the %known WR galaxy list to 139. Subsequent analysis of WR galaxy samples were performed %by \\citet*{GIT00,FCCG04} and \\citet{Buckalew05}. %In a study of emission-line galaxies extracted from the Sloan %Digital Sky Survey (SDSS, York et al. 2000) WR features were identified %in many star-forming galaxies \\citep{Kniazev04,Zhang07} %The most recent survey for WR features using the SDSS Data Release 6 \\citep{Adelman-McCarthy08} %%(Adelman-McCarthy et al. 2008) %was performed by Brichmann, Kunth \\& Durret (2008, hereafter BKD08), %\\citet{BKD08}, %who detected 570 galaxies with significant %WR features and a further 1115 potential candidates.    The main aim of our study of the formation of massive stars in starburst galaxies and the role that the interactions with or between dwarf galaxies and/or low surface brightness objects have in its triggering mechanism. In Paper~I \\citep{LSE08} we described the motivation of this work, compiled the list of the 20 analysed WR galaxies (Table~1 of Paper~I), the majority of them showing several sub-regions or objects within or surrounding them, and presented the results of the optical/\\NIR\\ broad-band and \\Ha\\ photometry. In Paper~II \\citep{LSE09} we presented the results of the analysis of the intermediate resolution long-slit spectroscopy of 16 WR galaxies of our sample -- the results for the other four galaxies were published separately. In many cases, two or more slit positions were used to analyse the most interesting zones, knots or morphological structures belonging to each galaxy or even surrounding objects. Paper~III \\citep{LSE10} presented the analysis of the O and WR stellar populations within these galaxies. In this paper, the forth of the series, we globally compile and analyse the optical/\\NIR\\ photometric data (Sect.~2) and study the physical (Sect.~3) and chemical (Sect.~4) properties of the ionized gas within our galaxy sample. Thirty-one up to 41 regions have a direct estimate of the electron temperature of the ionized gas, and hence the element abundances were derived with the direct method. Section~4 includes the analysis of the N/O ratio with the oxygen abundance, a discussion of the nitrogen enrichment in WR galaxies, a study of the $\\alpha$-elements to oxygen ratio with the oxygen abundance, and the comparison of the results provided by the most common empirical calibrations with those derived following the direct method (the Appendix compiles all metallicity calibrations used in this work). Section~5 analyses the metallicity-luminosity relations obtained with our data. Section~6 discusses the relations between the metallicity and the optical/$\\NIR$ colours. Finally, we list our main conclusions in Sect.~7.  The final paper of the series (Paper~V) will compile the properties derived with data from other wavelengths (UV, FIR, radio, and X-ray) and complete a global analysis of all available multiwavelength data of our WR galaxy sample. We have produced the most comprehensive data set of these galaxies so far, involving multiwavelength results and analysed according to the same procedures. ",
        "conclusions": "We compiled and analysed globally the optical/\\NIR\\ colours and the physical and chemical properties of the ionized gas for a sample of 20 Wolf-Rayet galaxies. The individual analysis of the photometry of each galaxy was presented in Paper~I, while the individual spectroscopic analysis was discussed in Paper~II.  The metallicity of these galaxies lies between 7.58 and 8.75 --in units of \\abox--. The most important conclusions found in this study are \\begin{enumerate} \\item The colours estimated for our galaxy sample, which were corrected for both extinction and nebular emission using our spectroscopic data, agree quite well with the predictions given by evolutionary synthesis models, %(Bruzual \\& Charlot 2003; PEGASE.2: Fioc \\& Rocca-Volmerange, 1997; %STARBURST99: Leitherer et al. 1999; Vazquez \\& Leitherer, 2005), especially in compact and dwarf objects. Small discrepancies are explained because of the existence of several stellar populations within each galaxy and differences in their star-formation history. %The ages of both young and old stellar populations have been derived consistently. All galaxies show evidence of an old stellar population underlying the starburst, with ages older than 500 Myr. \\item %We analysed the physical properties of their ionized gas, %including reddening coefficient, \\CHb, equivalent width of the \\Hb\\ emission line, %\\WHb, equivalent width of the underlying stellar absorption in the Balmer lines, %\\Wabs, ionization degree, electron density, and electron temperature. We checked that all objects can be classified as pure star-forming galaxies. In total, we compiled 41 independent star-forming regions, of which 31 have a direct estimate of the electron temperature of the ionized gas, and hence their element abundances were derived using the direct method. We found that younger bursts have larger ionization budgets and are therefore capable to heat the ionized gas to higher electron temperatures. Both \\CHb\\ and \\Wabs\\ increase with increasing metallicity, as predicted by galaxy evolution models. %We showed that more metallic objects tend to have a higher underlying %stellar absorption, and hence affecting to the observed values of \\WHb. \\item We compiled the oxygen abundance and N/O, S/O, Ne/O, Ar/O, and Fe/O ratios in our WR galaxy sample. They generally agree well with previous observations. The N/O ratio is found to be rather constant for objects with \\abox$\\leq$7.6, has an important dispersion in galaxies with 7.6$\\leq$\\abox$\\leq$8.3, and increases with the metallicity in objects with \\abox$\\geq$8.3. This behaviour is explained assuming the very different star-formation histories that each individual system has experienced. \\item We detected a high N/O ratio in objects showing clear WR features (HCG~31~AC, UM~420, IRAS~0828+2816, III~Zw~107, ESO~566-8 and NGC~5253). The ejecta of the WR stars may be the origin of the N enrichment in these galaxies, but further detailed data comparing the chemical properties of a larger sample of both WR and non-WR galaxies, as well as careful analyses of galaxies showing a localized high N/O ratio, are needed. \\item The relative abundance ratios of the $\\alpha$-elements to oxygen are approximately constant, which is expected because all four elements are mainly produced by massive stars. We found indications of a moderate depletion of oxygen and iron onto grains in the most metal-rich galaxies. \\item We compared the abundances provided by the direct method with those obtained using the most common empirical calibrations: \\begin{itemize} \\item The Pilyugin-method of \\citet{P01a,P01b}, which considers the $R_{23}$ and the $P$ parameters and is updated by \\citet{PT05}, is nowadays the best suitable empirical calibration to derive the oxygen abundance of star-forming galaxies. The cubic fit to $R_{23}$ provided by \\citet*{Nagao06} is not valid for analysing these star-forming galaxies. \\item The relations between the oxygen abundance and the $N_2$ or the $O_3N_2$ parameters provided by \\citet{PP04} give acceptable results for objects with \\abox$>$8.0. \\item The results provided by empirical calibrations based on photoionization models \\citep{McGaugh91,KD02,KK04} are systematically 0.2 -- 0.3 dex higher than the values derived from the direct method. These differences are of the same order as the abundance discrepancy found between abundances determined from recombination and collisionally excited lines of heavy-element ions. This may suggest temperature fluctuations in the ionized gas, as they exist in Galactic and other extragalactic \\HII\\ regions. \\end{itemize} \\item We studied the optical/\\NIR\\ luminosity-metallicity relations for our sample galaxies. We found that our data generally disagree with previous relations, perhaps because the objects analysed here host strong starbursts, but maybe also because we used the direct method and not the empirical calibrations to derive the oxygen abundances. The $L$--$Z$ relations tend to be tighter when using \\NIR\\ luminosities, \\TDG\\ candidates are indeed easier detected using the $M_{Ks}$--$Z$ relation. \\item The nitrogen abundance also has a strong correlation with the luminosity, but %but we found two trends that are consequence of the different %star-formation histories of each galaxy. normal and dwarf non-starbursting galaxies show a higher N/O ratio than strong starbursting galaxies. %Surprisingly, the derived N/H-$M_B$ relation agrees quite well with the %relation provided by van Zee \\& Haynes (2006), that used non-starbursting systems. \\item We found that galaxies with redder colours tend to have a higher oxygen abundance. The N/O ratio also increases with redder colours. Both results agree with galaxy evolution models. The colours of the underlying component seem to be metallicity-independent, but more data are still needed to confirm this trend. %The dispersion of the data is explained by considering %different star-formation histories in every object. \\end{enumerate} We finally conclude that it is fundamental to perform a detailed analysis of both the photometric (optical and \\NIR\\ magnitudes and colours of both burst and underlying component) and the chemical (oxygen abundances, N/O ratios) properties of these star-forming galaxies to understand the evolutionary stage of every system. A larger galaxy sample following the ideas compiled in this work %(that will provide a better understanding of the star formation history of each individual object) will complement the results derived from huge databases (which do not distinguish the evolutionary state of the galaxies or the relative contribution of the burst/underlying populations, and for which the majority of the properties have been derived automatically and, in many cases, have considered empirical calibrations to determine the chemical abundances) when analysing the global properties and evolution of star-forming galaxies.   %We then conclude that it is fundamental to perform such detailed analysis of %the photometric and spectroscopic characteristics of this kind of star-forming galaxies, %including a proper separation of the properties of both the star-forming bursts and the %underlying stellar component, to understand the evolutionary stage of every system."
    },
    "1004/1004.1342_arXiv.txt": {
        "abstract": "We present optical emission-line spectra for outlying HII regions in the extended neutral gas disk surrounding the blue compact dwarf galaxy NGC 2915. Using a combination of strong-line R23 and direct oxygen abundance measurements, we report a flat, possibly increasing,  metallicity gradient out to 1.2 times the Holmberg radius. We find the outer-disk of NGC 2915 to be enriched to a metallicity of 0.4 Z$_{\\odot}$.  An analysis of the metal yields shows that the outer disk of NGC 2915 is overabundant for its gas fraction, while the central star-foming core is similarly under-abundant for its gas fraction. Star formation rates derived from very deep $\\sim14$ ks GALEX FUV exposures indicate that the low-level of star formation observed at large radii is not sufficient to have produced the measured oxygen abundances at these galactocentric distances.  We consider 3 plausible mechanisms that may explain the metal-enriched outer gaseous disk of NGC 2915: radial redistribution of centrally generated metals, strong galactic winds with subsequent fallback, and galaxy accretion. Our results have implications for the physical origin of the mass-metallicity relation for gas-rich dwarf galaxies. ",
        "introduction": "\\label{intro} The mass fraction of heavy elements in a galaxy is a fundamental physical property: it contains a record of the star formation history and therefore serves as an indicator of its evolutionary state. Numerous observational studies of gas metal abundances in a wide variety of galaxies, spanning a large range in mass and morphological types, have established a now well-known, yet poorly understood global relation in which the mean metallicity and stellar mass are directly correlated (e.g. Lequeux et al. 1979, Tremonti et al. 2004, Geha et al. 2009\\nocite{lequeux79,tremonti04, geha09}).  The physical processes that govern this mass-metallicity relation are under debate, yet there is some consensus that low-mass dwarf galaxies hold the key to understanding this relation. Using a sample of 53,400 star-forming Sloan Digital Sky Survey galaxies, \\cite{tremonti04} argue that metal loss via galactic winds is most likely to be responsible for the steep decline in the ``effective yield\" (a parameterization of the metallicity divided by the galaxy total gas fraction) with galaxy baryonic mass. The physical basis for these outflows would be the cumulative effect of supernovae in disk OB associations which leads to low-mass galaxies with shallow gravitational potential wells selectively losing their metals \\citep{maclow99,strickland04,brooks07}. Nonetheless, there are several other possible mechanisms for explaining the correlation between low metal yields and galaxy mass, including metal mixing in extended gaseous disks \\citep{tassis08}, and a lower effective stellar upper mass limit to the initial mass function in dwarf galaxies \\citep{koppen07, meurer09}.   The as-yet open question of whether the extended neutral gas disks of dwarf galaxies are metal-enriched is quite relevant to this discussion, since each proposed model makes a prediction for the metal distributions in galaxies. For instance, in the SN-blowout scenario, metals tend to get ejected into the IGM for galaxies less massive than 10$^{9}$ M$_{\\odot}$ \\citep{maclow99} or 10$^{11}$  M$_{\\odot}$ \\citep{strickland04}, while in the metal-mixing scenario, the central regions and outermost regions of the dwarf galaxy should have the same metal abundance. Therefore, abundance gradients (or a lack thereof) in the extended gaseous disks of dwarf galaxies may help to untangle the physical origin of the mass-metallicity relation.    \\begin{figure*}[t] \\epsscale{0.8} \\plotone{figure1.eps} \\caption{Top: A CTIO 1.5-m broadband optical V-Band image of NGC 2915 overlaid with ATCA HI contours. The HI contour levels are 5, 15, 25, 35, 45, 55, 65, 75, 85, and 95 percent of the peak HI column density, 1.71 $\\times$10$^{21}$ cm$^{-2}$.  The outer gaseous disk of NGC 2915 mapped with the ATCA appears to extend 5 times beyond the optical component of the galaxy which has a 2.3 kpc extent. On the image, we have boxed the 3 outlying HII regions included in this study, and contained in the lower zoomed-in panels. Bottom: H$\\alpha$ continuum-subtracted images of the core of NGC 2915 shown in two stretches: (A) to highlight the outer HII regions included in this study, and (B) to highlight the inner HII regions included in this study. } \\end{figure*}  NGC 2915 is one of the most extreme examples of a blue compact dwarf galaxy with an extended gaseous disk (Meurer et al. 1996; hereafter M96\\nocite{meurer96}).  This nearby (4.1 Mpc, Meurer et al. 2003 \\nocite{meurer03}) dwarf galaxy has an HI disk that extends 5 times beyond the optical stellar component (12 kpc for the gas; 2.3 kpc for the stars; see Figure 1, and Table 1 for a list of its full properties). \\begin{deluxetable}{llr} \\tablecaption{NGC 2915 Optical and HI Properties \\label{props}} \\tabletypesize{\\scriptsize} \\tablewidth{0pt} \\tablehead{\\colhead{Property} & \\colhead{} & \\colhead{Value}} \\startdata RA (J2000)& &09:26:11.5\\\\ dec (J2000)&&-76:37:35\\\\ distance&&4.1$\\pm$0.3 Mpc\\\\ Holmberg Radius&&2.3 kpc\\\\ M$_{\\star}$&&3.2$\\times$10$^{8}$ M$_{\\odot}$\\\\ M$_{gas}$&& 7.4$\\times$10$^{8}$ M$_{\\odot}$\\\\ E(B-V)&&0.275$\\pm$0.04\\\\ M$_{V}$& &-16.42\\\\ M$_{dyn}$/L & & 62* (solar)\\\\ R$_{HI}$&&11.9 kpc\\\\  \\enddata \\tablecomments{Properties from M03 and M96, in some cases updated for a revised distance measurement. (*) New HI data from \\cite{elson10} indicates that M$_{dyn}$/L may in fact be as high as 140 (solar units). Reddening comes from \\cite{schlegel98}.} \\end{deluxetable} \\begin{deluxetable*}{lccccc}[b] \\tablecaption{NGC 2915 HII Regions \\label{list}} \\tabletypesize{\\scriptsize} \\tablewidth{0pt} \\tablehead{\\colhead{HII Region} & \\colhead{RA} & \\colhead{dec} & \\colhead{R/R$_{Ho}$} & \\colhead{Log L$_{H\\alpha}$} & \\colhead{notes} \\\\ \\colhead{} & \\colhead{J2000} & \\colhead{J2000} & \\colhead{} & \\colhead{(ergs s$^{-1}$)} & \\colhead{}} \\startdata 1 & 09:26:00.3 & -76:35:25.0 & 1.2 & 36.2 & low ionization parameter; [OIII]$\\lambda$4959 not detected \\\\ 2 & 09:26:45.9 & -76:37:01.1 & 1.1 & 36.5 &  diffuse, outer HII region; [OII]$\\lambda\\lambda$3727 detected at 3$\\sigma$\\\\ 3 & 09:26:41.7 & -76:38:12.9 & 1.0 & 35.8  & only H$\\alpha$ and [OIII]$\\lambda$5007\\\\ 4 & 09:26:19.8 &  -76:38:02.0 & 0.34 & 38.4 & slit length = 30\\arcsec, on edge of central star formation in NGC 2915\\\\ 5 & 09:26:11.1 & -76:37:39.0 & 0.04 & 39.2 & central, luminous HII region; [OIII]$\\lambda$4363 detected\\\\ 6 & 09:26:07.3 & -76:37:12.9 & 0.23 & 37.8 &  bright HII region within galaxy\\\\ \\enddata \\tablecomments{Values of the total H$\\alpha$ luminosity are derived from the TTF image, and are corrected for internal extinction derived from \\cite{schlegel98}. Errors in the H$\\alpha$ luminosity are on the order of 10\\%\\, dominated by the flux calibration errors. } \\end{deluxetable*}  Its total baryonic mass (gas plus stars; $\\sim10^{9}$ M$_{\\odot}$) puts it on the high-mass end of the spectrum of dwarf galaxies, and its total dynamical mass gives it one of the highest-known mass-to-light ratios for a gas-rich galaxy (M96).  Recent H$\\alpha$ images of NGC 2915 have revealed several small pockets of star formation embedded in its extended gaseous disk at projected radii of $\\sim3$kpc that otherwise contains few stars (see Figure 1). In addition, new, very deep (t$_{exp}$ = $\\sim14$ ks) images from the {\\it{Galaxy Evolution Explorer}} Satellite ({\\it{GALEX}}) show very faint extended-UV (XUV) emission near the location of these HII regions.  Because of their strong emission-line spectra, HII regions are relatively straightforward probes of the oxygen abundance of surrounding gas. In the most luminous cases, a direct measurement of the oxygen abundance can be made using the temperature sensitive, faint [OIII]$\\lambda4363$ line, and in other cases, calibrated relations between strong lines (e.g. R23; Pagel et al. 1979; McGaugh 1991, hereafter M91 \\nocite{pagel79,mcgaugh91}) can provide reasonable estimates of  the oxygen abundance. Dwarf galaxies have been observed to have spatially homogeneous oxygen abundances, although the most recent studies have so far  been limited to well-within the Holmberg radius \\citep{lee06, croxall09}. By comparison, we report the oxygen abundances for HII regions out to 1.2 times the Holmberg radius (the isophotal radius at which the surface brightness is below 26.5 in the B-band) and discuss the implications of our results. ",
        "conclusions": "We have derived oxygen abundances from optical emission-line spectra of 5 HII regions in the extended gaseous disk of NGC 2915, in locations ranging from the central starburst to 1.2 times the Holmberg radius. Outer-disk HII regions were originally found using deep AAT H$\\alpha$ images, and appear to lie in a very faint extended-UV disk visible in recent deep GALEX images. Out to these large galactocentric distances, we find no evidence of a decreasing metallicity gradient. The central HII regions have a metallicity of 0.2 Z$_{\\odot}$  ($\\pm$ 0.15 dex), while the outer HI regions appear to be enriched at a level of 0.4 Z$_{\\odot}$  ($\\pm$ 0.25 dex). Based on calculations of metal yields and star formation rates in NGC 2915, we conclude that the outer disk is considerably more metal-rich than its ``expected\" value of 0.04 $-$ 0.08 Z$_{\\odot}$, and that the central region is similarly under-abundant compared to its ``expected\" value of 1.7 $-$ 1.9 Z$_{\\odot}$. These observations indicate that some process, other than ongoing star formation, has enriched the gas at large radii. We present 3 plausible (and non-exclusive) scenarios for the metal-enriched outer gas disk of NGC 2915: metal mixing, supernovae-driven winds entraining metals and falling back down at large radii, or a past interaction leading to the accretion of enriched gas. Measurements of metal abundances for the outer gaseous disks of additional galaxies will cast light on the processes that  redistribute metals and their effect on the mass-metallicity relation for galaxies."
    }
}